{
    "0809/0809.4857_arXiv.txt": {
        "abstract": "{$\\iota$~Dra (\\object{HIP~75458}) is a well-known example for a K giant hosting a substellar companion since its discovery by \\citeauthor{Frink02} \\citeyearpar{Frink02}. We present radial velocity measurements of this star from observations taken with three different instruments spanning nearly 8 years. They show more clearly that the RV period is long-lived and coherent thus supporting the companion hypothesis. The longer time baseline now allows for a more accurate determination of the orbit with a revised period of $P=511\\,$d and an additional small linear trend, indicative of another companion in a wide orbit. Moreover we show that the star exhibits low amplitude, solar like oscillations with frequencies around 3-4\\cyd  (34.7-46.3\\,$\\mu$Hz).} ",
        "introduction": "Up to now long period radial velocity (RV) variations have been detected in several K giants, e.g. $\\beta$~Gem \\citep{Hatzes06,Reffert06}, HD~47536 \\citep{Setiawan03}, HD~13189 \\citep{Hatzes05} and $\\iota$~Dra \\citep{Frink02}. Rotational modulation can be excluded as a cause of these variations for most of these giants due to lack of variations in photometry or line bisectors with the RV period \\citep{Hatzes98Bisector}. The most likely interpretation is orbiting, substellar companions.  Since the progenitor stars to planet hosting giant stars can have masses significantly larger than $1M_{\\odot}$, these giant stars offer a way to investigate planet formation around massive stars. In their evolved phase the Doppler induced effects of a planet are easier to detect than in their main-sequence phase, due to cooler effective temperatures and smaller rotational velocities.  But for evolved stars the determination of the mass is more difficult. The most practical way for determining the stellar mass is the use of evolutionary tracks, but those tracks converge in the region of the giant branch for stars over a wide range of masses and thus makes the derived masses uncertain.  Another possibility to measure the stellar mass is via asteroseismology. This requires the investigation of stellar oscillations. Photometric and RV variations with short periods from hours to days have already been detected in some K and late G giants (e.g. $\\alpha$~Boo: \\citealp{Retter03,Tarrant07}; $\\alpha$~Ari: \\citealp{Kim06}; $\\epsilon$~Oph: \\citealp{deRidder06}; $\\xi$~Hya: \\citealp{Frandsen02}, \\citealp{Stello06}; $\\beta$~UMi: \\citealp{Tarrant08}) and are consistent with solar-like p-mode oscillations. Such oscillations have also been measured for ensembles of red giants with photometry of star rich regions or clusters (e.g. \\citealp{Gilliland08}, \\citealp{Stello08}) which is a very efficient way to perform such asteroseismologic studies.  Solar like oscillations have also been discovered in planet hosting main sequence stars (e.g. $\\mu$~Arae: \\citealp{Bazot05}; $\\iota$~Hor: \\citealp{Vauclair08}) and in the planet hosting K giant $\\beta$~Gem for which \\cite{Hatzes07} gave, in combination with interferometric measurements of the angular diameter, an estimation for the stellar mass completely independent from evolutionary tracks. In a similar way we will do this here for $\\iota$~Dra. ",
        "conclusions": "We revised the orbit solution for the companion of $\\iota$~Dra. The orbital period is $P=511\\,$d, somewhat lower than the value in the discovery paper by \\cite{Frink02}. Furthermore there is a linear trend of $-13.8\\,$m/s/yr present, possibly caused by another companion.  An excess of power around 3.8\\cyd was found independently in two data sets, taken at different times and with very different sampling. The amplitude and the location of the power excess  are consistent with solar-like oscillations.  Our analysis of the short period oscillations indicates a somewhat high stellar mass of $M=2.2M_\\odot$, considerably higher than the $1.05M_\\odot$ of \\cite{Prieto99} or the stellar masses derived  from the Girardi track ($1.2-1.7M_\\odot$). However, this is a preliminary result based on limited data. Assuming $1.4M_\\odot$ for the mass $\\iota$~Dra yields a minimum mass of $10M_{\\mathrm{Jup}}$ for the companion.  Our RV measurements for $\\iota$~Dra indicate that it shows multi-periodic stellar oscillations. This means that an asteroseismic analysis can yield an accurate mass. This is best done by measuring the frequency splitting in the p-mode oscillation spectrum which requires more measurements and better sampling than the data presented here."
    },
    "0809/0809.1369_arXiv.txt": {
        "abstract": "{The International Celestial Reference Frame (ICRF), currently based on the position of 717 extragalactic radio sources observed by VLBI (Very Long Baseline Interferometry), is the fundamental celestial reference frame adopted by the International Astronomical Union (IAU) in 1997. Within the next 10 years, the European space astrometry mission Gaia, to be launched by 2011, will permit determination of the extragalactic reference frame directly in the visible for the first time. Aligning these two frames with the highest accuracy will therefore be very important in the future for ensuring consistency between the measured radio and optical positions.} {This paper is aimed at evaluating the current astrometric suitability of the individual ICRF radio sources which are considered appropriate for the alignment with the future Gaia frame.} {To this purpose, we cross-identified the ICRF and the optical catalog V\\'{e}ron-Cetty and V\\'{e}ron (2006), in order to identify the optically-bright ICRF sources that will be positioned with the highest accuracy with Gaia. Then we investigated the astrometric suitability of these sources by examining their VLBI brightness distribution.} {We identified 243 candidate ICRF sources for the alignment with the Gaia frame (i.e. with an optical counterpart brighter than the apparent magnitude 18), but only 70 of these (i.e. only 10\\% of the ICRF sources) are found to have the necessary high astrometric quality (i.e. a brightness distribution that is compact enough) for this link. Additionally, it was found that the QSOs (quasi stellar objects) that will have the most accurate positions in the Gaia frame tend to have less-accurate VLBI positions, most probably because of their physical structures.} {Altogether, this indicates that identifying other high-quality VLBI radio sources suitable for the alignment with the future Gaia frame is mandatory.} ",
        "introduction": "The International Celestial Reference System (ICRS) is a kinematical system that assumes that the Universe does not rotate, hence breaking the history of the inertial system materialization from observations attached to the apparent motion of the Sun \\citep{Arias1995}. The International Celestial Reference Frame (ICRF) is the realization at radio wavelengths of the ICRS, through very long baseline interferometry (VLBI) measurements of extragalactic radio source positions. It was adopted by the International Astronomical Union (IAU) as the fundamental celestial reference frame during the IAU XXIIIrd General Assembly in Kyoto, Japan, in August 1997. The initial realization was based on the positions measured with VLBI of 212 \\textit{defining} extragalactic radio sources (setting the direction of the ICRF axes). In addition, positions for 396 other sources (divided into 294 \\textit{candidate} and 102 \\textit{other} sources) were included to make the frame denser \\citep{Ma1998}. The ICRF was extended at later stages with the positions of another 109 sources commonly referred to as \\textit{new} sources \\citep{Fey2004}. Overall, the ICRF currently consists of a set of VLBI coordinates for 717 extragalactic radio sources, most of which show sub-milliarcsecond accuracy. The accuracy of the individual source positions depends on the number of observations but also on the compactness and positional stability of the sources that show brightness distributions that are often not point-like and time-variable (\\citealt{Ma1998}; \\citealt{Fey2000}).  The European space astrometry mission Gaia, to be launched by 2011, will survey about one billion stars in our Galaxy and throughout the Local Group along with 500\\,000 quasi stellar objects (QSOs), down to an apparent optical magnitude of 20 \\citep{Perryman2001}. Optical positions with Gaia will be determined with an unprecedented accuracy, ranging from a few tens of microarcseconds ($\\mu$as) at magnitude 15--18 (targeted accuracies are 16 $\\mu$as at 15 mag and 70 $\\mu$as at 18 mag) to about 200~$\\mu$as at magnitude 20 \\citep{Lindegren2008}. Unlike Hipparcos, Gaia will permit the construction of an extragalactic frame directly at optical wavelengths, based on the QSOs with the most accurate positions (i.e. the QSOs with magnitude brighter than 18; \\citealt{Mignard2003}). \\citet{Mignard2002} demonstrates that the residual spin of the Gaia frame can be determined to $0.5~\\mu$as/yr with a ``clean sample'' of about 10\\,000 such QSOs. A preliminary Gaia catalog is expected to be available by 2015 with the final version released by 2019.  In the future, aligning the ICRF with the Gaia celestial reference frame will be crucial for ensuring consistency between the measured radio and optical positions. This alignment will be important not only for guaranteeing the proper transition in the case that the fundamental reference frame is moved from the radio domain to the optical domain (as currently anticipated) but also for registering the radio and optical images of any celestial target with the highest accuracy. To align the two frames with the highest accuracy, it is desirable to have several hundred common objects (the more common objects, the more accurate the alignment) with a uniform sky coverage and very accurate radio and optical positions. Obtaining such accurate positions implies that the link sources must have (i) an apparent optical magnitude $V$ brighter than 18, for the highest Gaia astrometric accuracy \\citep{Mignard2003}, and (ii) no extended VLBI structures, for the highest VLBI astrometric accuracy \\citep{Fey2000}.  This paper is aimed at evaluating the suitability of the current individual ICRF extragalactic radio sources for the alignment with the future Gaia frame. First, the ICRF is cross-identified with the \\citet{Veron2006} optical catalog of QSOs in order to identify the ICRF sources with a proper optical counterpart for the Gaia link. The astrometric suitability of these sources is then investigated by examining their VLBI brightness distribution and their position accuracy in the ICRF. In this investigation, the structure index defined by \\citet{Fey2000} is used to identify the sources that have the most compact VLBI brightness distributions. This study results in a sample of 70 ICRF sources, which have appropriate compact structures on VLBI scales and are brighter than magnitude 18. These sources are at present the best candidates for the ICRF--Gaia alignment. ",
        "conclusions": "This study focused on the astrometric suitability of the current ICRF extragalactic radio sources for the alignment with the future Gaia frame. We identified 243 candidate sources for this alignment, but only 70 of these (10\\% of the ICRF) possess the high astrometric quality required to ensure the highest accuracy in the alignment. Accordingly, the current number of VLBI sources for accurately aligning the ICRF and the future Gaia frame is not sufficient. We also showed that the QSOs that will have the most accurate positions measured with Gaia are not those that have the best astrometric positions in the ICRF statistically.  To compensate for the lack of high astrometric quality extragalactic radio sources for this alignment, the VCS catalog, along with the future ICRF-2, will be of high interest. An additional direction would be to identify new appropriate VLBI sources from current deep radio surveys. These solutions mostly concern the northern hemisphere, since the VLBI arrays and observations are concentrated in this part of the world. However, to ensure a homogeneous sky coverage, a major effort is also necessary in the southern hemisphere, most specifically for the declinations below $-40^{\\circ}$. To this end, the Asia Pacific Telescope \\citep{Gulyaev2007} is likely to play a significant role."
    },
    "0809/0809.2556.txt": {
        "abstract": "We present the results of parsec-scale circular polarization measurements based on Very Long Baseline Array data for a number of radio-bright, core-dominated Active Galactic Nuclei obtained simultaneously at 15, 22 and 43 GHz. The degrees of circular polarization $m_c$ for the VLBI core region at 15 GHz are similar to values reported earlier at this wavelength, with typical values of a few tenths of a percent. %The origin of this %polarization is most likely the conversion of linear to circular %polarization during the propagation of the radiation through a magnetised %plasma, although some contribution from ``intrinsic'' circular polarisation %generated by the synchrotron mechanism cannot be ruled out. We find that $m_c$ as often rises as falls with increasing frequency between 15 and 22~GHz, while the degree of circular polarisation at 43~GHz is in all cases higher than at 22 and 15~GHz. This behaviour seems contrary to expectations, since the degree of circular polarization from both synchrotron radiation and Faraday conversion of linear to circular polarization -- the two main mechanisms considered thus far in the literature -- should {\\em decrease} towards higher frequencies if the source is homogeneous. The increase in $m_c$ at 43~GHz may be due to the presence of regions of both positive and negative circular polarisation with different frequency dependences (but decreasing with increasing frequency) on small scales within the core region;  alternatively, it may be associated with the intrinsic inhomogeneity of a Blandford--K\\\"onigl-like jet. In several objects, the detected circular polarization appears to be near, but not coincident with the core, although further observations are needed to confirm this.  We find several cases of changes in sign with frequency, most often between 22 and 43~GHz.  We find tentative evidence for transverse structure in the circular polarization of 1055+018 and 1334$-$127 that is consistent with their being generated by either the synchrotron mechanism or Faraday conversion in a helical magnetic field. Our results confirm the earlier finding that the sign of the circular polarization at a given observing frequency is generally consistent across epochs separated by several years or more, suggesting stability of the magnetic field orientation in the innermost jets. ",
        "introduction": "The radio emission of core-dominated, radio-loud Active Galactic Nuclei (AGN) is synchrotron radiation generated in the relativistic jets that emerge from the nucleus of the galaxy, presumably along the rotational axis of a central supermassive black hole. Synchrotron radiation can be highly linearly polarized, to $\\simeq 75\\%$ in the case of a uniform magnetic ({\\bf B}) field (Pacholczyk 1970), and linear polarization observations can yield unique information about the orientation and degree of order of the {\\bf B} field in the synchrotron source, as well as the distribution of thermal electrons and the {\\bf B}-field geometry in the immediate vicinity of the AGN (e.g., via Faraday rotation of the plane of polarization).  Techniques for deriving circular-polarization (CP) information on parsec scales were pioneered by Homan and his collaborators in the late 1990s (Homan \\& Wardle 1999; Homan, Attridge \\& Wardle 2001) using data taken on the NRAO\\footnote{The National Radio Astronomy Observatory of the USA is operated by Associated Universities, Inc., under co-operative agreement with the US NSF.} Very Long Baseline Array (VLBA). Recently, CP measurements for the first epoch of the MOJAVE project (monitoring of 133 AGN at 15 GHz with the VLBA) were published by Homan and Lister (2006). Circular polarization was detected in 34 of these objects at the 2$\\sigma$ level or higher. These results confirmed previously noted trends: the circular polarization is nearly always coincident with the VLBI core, with typical degrees of polarization $m_c$ being a few tenths of a percent.  Homan \\& Lister (2006) found no evidence for any correlation between $m_c$ and any of 20 different optical, radio and intrinsic parameters of the AGN. Interestingly, five of the 34 AGN displayed CP in their {\\em jets}, well outside the VLBI-core region, suggesting that the mechanism generating the circular polarization is capable of operating effectively in optically thin regions (although, strictly speaking, direct spectral-index measurements were not available at the corresponding epoch, since the MOJAVE measurements were made only at 15~GHz).  The two main mechanisms that are usually considered to be the most likely generators of the observed CP are the synchrotron mechanism and the Faraday conversion of linear to circular polarization [Legg \\& Westfold 1968, Jones \\& O'Dell 1977; see also the reviews by Beckert \\& Falcke (2002) and Wardle \\& Homan (2003)].  Although the intrinsic CP generated by synchrotron radiation may be able to reach a few tenths of a percent at 15~GHz for the magnetic-field strengths characteristic of the observed VLBI cores of AGN (typically $\\simeq 0.4$~G; Lobanov 1998, O'Sullivan \\& Gabuzda, in preparation), the highest observed $m_c$ values seem too high to plausibly be attributed to this mechanism. This suggests that Faraday conversion plays a role, possibly the dominant one, since it is expected to be more efficient at generating CP than the synchrotron mechanism for the conditions in radio cores (Jones \\& O'Dell 1977).  Faraday conversion (Jones \\& O'Dell 1977, Jones 1988) occurs because the component of the linear-polarization electric vector parallel to the conversion magnetic field, E$_{\\|}$, gives rise to oscillations of free charges in the conversion region, while the component orthogonal to this magnetic field, E$_{\\perp}$, cannot (the charges are not free to move orthogonal to the magnetic field). This leads to a delay of E$_{\\|}$ relative to E$_{\\perp}$, manifest as the introduction of a small amount of circular polarization. We recently obtained intriguing evidence that the CP in AGN may be generated by Faraday conversion in helical jet magnetic fields (Gabuzda et al. 2008). In this scenario, linearly polarized radiation emitted at the far side of the jet is partially converted to circular polarization as it passes through the magnetic field at the near side of the jet on its way to the observer. In the simple model considered, the sign of the generated CP is determined by the pitch angle of the helical field, which can be approximately deduced from the observed linear polarization structure, and the helicity of the helical field, which can be deduced from the direction of the Faraday-rotation gradient across the VLBI jet, due to the systematically changing line-of-sight component of the helical field. The inferred helicity depends on whether the longitudinal component of the field corresponds to a ``North'' or ``South'' pole of the central black hole.  After identifying 8 AGN with both detected VLBI-scale CP and transverse Faraday rotation gradients, Gabuzda et al. (2008) determined the CP signs that should be observed for the cores of these AGN based on the inferred pitch angles and the helicities for their helical fields, if the jets were emerging from the North or South magnetic pole of the black hole. The result was clearly not random: in all 8 cases, the observed CP signs agreed with the signs predicted by the simple helical-field model for the case of the longitudinal field corresponding to South polarity. Although the physical origin of this result is not clear, it strongly suggests a close connection between CP and the presence of helical jet magnetic fields; the probability of finding that 8 of 8 of these AGN had CP signs corresponding to one particular polarity is less than 1\\%. The apparent predominance of jets associated with South magnetic poles may have cardinal implications for the intrinsic magnetic-field structures of the central black holes; for example, one way to understand this result is if the central fields are quadrupolar, with a predomint tendency for the two South poles to correspond to the jets and the two North poles to lie in the plane of the accretion disc, for reasons that are not yet understood.  We present here new CP measurements for 41 AGN at 15, 22 and 43~GHz, as well as for an additional 18 AGN at 15~GHz only. Parsec-scale CP was detected at two of the three frequencies in 5 of the 41 AGN, and at all three frequencies in another 5 AGN. Our 15-GHz results are generally in excellent agreement with the first-epoch MOJAVE CP values (Homan \\& Lister 2006).  The results do not show any universal frequency dependence for the degree of CP $m_c$, and both rising and falling spectra with frequency are observed. Unexpectedly, we find $m_c$ to be higher at 43~GHz than at the lower two frequencies. We also find tentative evidence for transverse CP structure consistent with the CP being produced in a helical magnetic field geometry in two objects. ",
        "conclusions": "We have presented the results of parsec-scale circular polarization measurements for 41 AGN at 15+22+43~GHz and an additional 18 at 15~GHz alone. We have detected parsec-scale CP in 8 sources for the first time, and confirm previous detections in an additional 9 objects. In all 7 AGN in which 15-GHz CP was detected in both our measurements and in the MOJAVE first-epoch measurements (Homan \\& Lister 2006), the signs of the CP agree, demonstrating that a consistent picture is beginning to emerge from the collected results.  The first-epoch MOJAVE measurements revealed the presence of jet CP in five AGN (3C84, 3C273, 2128$-$123, 2134+004 and 2251+158). Vitrishchak \\& Gabuzda (2007) reported the presence of jet CP in 1334$-$127 and 3C279, and we have confirmed jet CP in these objects and added 2223$-$052 to this list.  These are among the very first multi-frequency VLBI CP measurements. We have measured CP at more than one frequency for 10 of the 41 AGN. No simple picture of the frequency dependence of the CP emerges. For example, comparable numbers of the AGN displaying measureable CP at both 15 and 22~GHz display higher $|m_c|$ values at the higher or lower of these frequencies. At the same time, there is a clear tendency for the degrees of CP to be higher at 43~GHz than at the two lower frequencies. While virtually all the signs for the CP measured at 15 and 22~GHz agree, the sign of the observed core-region CP often changes between 15/22 and 43~GHz (in 4 of 6 sources). This suggests the action of optical-depth effects, or some other mechanism giving rise to a systematic change in the CP sign with distance from the jet base.  None of the observed CP spectra correspond in a straightforward way to the frequency dependence expected for either intrinsic CP or CP generated by Faraday conversion operating in a homogeneous source, which both predict that the CP should decrease with increasing frequency. The fact that we have found $m_c$ to increase with frequency in several sources, and to universally be higher at 43~GHz compared to our lower frequencies, is intriguing in light of the scaling arguments presented by Wardle \\& Homan (2003), which point out that CP from either synchrotron radiation or Faraday conversion in a Blandford--K\\\"onigl jet could have an inverted CP spectrum, $m_c\\sim\\nu$. This suggests that our measurements are probing scales on which the intrinsic inhomogeneity of the jet plays an appreciable role in determining the observed CP and its spectrum.  In addition, the picture may be complicated by the fact that the core CP measurements may represent the superposition of CP contributions from several different regions in the core and innermost jet, as has been directly observed in the nearby radio galaxy 3C84 (Homan \\& Wardle 2004).  Although we cannot distinguish between various CP-generation mechanisms based on the observed spectra, we argue, as have previous authors, that the synchrotron mechanism has trouble explaining the growing number of measurements indicating degrees of CP of the order of 1\\% or more. Therefore, it remains most likely that the detected CP is generated primarily by Faraday conversion (possibly with a smaller contribution by the synchrotron mechanism in some objects). Gabuzda et al. (2008) have argued that the CP is generated predominantly by Faraday conversion in helical {\\bf B} fields inherently associated with the jets.  We have found tentative evidence for transverse structure in the CP distributions of 1055+018 and 1334$-$127. %In 1334$-$127, we observe one region %of CP of a single sign, well out in the jet and displaced toward the edge %of the jet. We have tested the reality of this feature using test data %and found it to be robust. %%This is %%consistent with the expectations for CP generated by Faraday conversion %%in a helical jet {\\bf B} field with a pitch angle of $\\simeq 40-45^{\\circ}$ %%viewed at an angle of $\\simeq 60-80^{\\circ}$ (or $\\simeq 100-120^{\\circ}$) %%in the jet rest frame. Based on the sign of the observed extended jet CP, %%we deduce that, in this model, the helical jet {\\bf B} field in 1334$-$127 %%is right-handed. %In 1055+018, we observe two regions of CP of opposite signs in the core %region, on either side of the innermost VLBI jet. Unfortunately, because %the peaks of the two inferred CP components are nearly the same and they %are located nearly symmetrically about the VLBI core region, we cannot %rule out the possibility that this structure is due to a small transverse %shift between the $RR$ and $LL$ images introduced by the phase calibration. %However, the fact that it is much more natural for such $RR/LL$ misalignments %to occur along the jet than orthgonal to the jet direction suggests that this %structure may well be real. In fact, transverse CP structures of the sort observed in 1334$-$127 and tentatively detected in 1055+018 would come about naturally if the CP is generated in helical jet {\\bf B} fields. Further, this is true for both intrinsic CP and CP generated by helical-field-driven Faraday conversion. In contrast, CP generated in jets in which the unidirectional {\\bf B} field giving rise to the CP is longitudinal would display transverse structure only by chance, due, for example, to bending of the jets. We thus consider the transverse CP structures tentatively revealed by these first 43-GHz CP measurements to provide further supporting evidence for the hypothesis that the CP of compact radio-loud AGN is closely tied to helical {\\bf B} fields that are organically related to the jets.  The interesting and unexpected new results we have presented here indicate that further CP analyses at 22 and 43~GHz are very worthwhile, despite the techical difficulties they present (weaker source fluxes, higher noise levels, particularly at 43~GHz). Further high-resolution observations for additional AGN will indicate how common transverse CP structures really are, and, in the case of symmetrical transverse CP structures with opposite signs on opposite sides of the core region, more conclusively demonstrate their reality. One possibility in this regard will be if similar structures can be detected further from the core, as the reality of these can be tested directly using the PHC procedure described above and by Homan \\& Lister (2006)."
    },
    "0809/0809.5256_arXiv.txt": {
        "abstract": "In this paper we study the determinants of starless core temperatures in the Perseus molecular cloud.  We use \\amm\\ (1,1) and (2,2) observations to derive core temperatures (\\tc) and data from the COMPLETE Survey of Star Forming Regions and the c2d Spitzer Legacy Survey for observations of the other core and molecular cloud properties.  The kinetic temperature distribution probed by \\amm\\ is in the fairly narrow range of $\\sim$ 9 - 15 K.  We find that cores within the clusters IC348 and NGC1333 are significantly warmer than ``field'' starless cores, and \\tc\\ is higher within regions of larger extinction-derived column density.  Starless cores in the field are warmer when they are closer to class O/I protostars, but this effect is not seen for those cores in clusters.  For field starless cores, \\tc\\ is higher in regions in which the $^{13}$CO linewidth and the 1.1mm flux from the core are larger, and \\tc\\ is lower when the the peak column density within the core and average volume density of the core are larger.  There is no correlation between \\tc\\ and $^{13}$CO linewidth, 1.1mm flux, density or peak column density for those cores in clusters.  The temperature of the cloud material along the line of sight to the core, as measured by CO or far-infrared emission from dust, is positively correlated with core temperature when considering the collection of cores in the field and in clusters, but this effect is not apparent when the two subsamples of cores are considered separately. ",
        "introduction": "Starless cores are the link between molecular clouds and Class 0 protostars. They are smaller and denser (r $\\sim$ 0.1 pc, n $\\gtrsim 10^4$ cm$^{-3}$) than, and kinematically distinct from, the molecular clouds in which they are found \\citep{Goodman98}.  Because starless cores have no source of internal heating, the temperature of a starless core is determined by external illumination and self-shielding.  The temperatures of starless cores are set by the interstellar radiation field, and in some cases luminous nearby stars, attenuated by the molecular cloud.  The temperature profile within a starless core is determined by the heating of its outer layers and the internal extinction, making the centers of starless cores colder than their outer layers \\citep{Evans01, Stamatellos07, Zucconi01}.  The temperatures of starless cores can be inferred from thermal emission from dust or from molecular lines.  Deriving core temperatures from continuum emission is complicated by uncertainties in the dust emission spectrum.  For instance, in cold and dense regions dust grains are expected to coagulate and become covered in icy mantles, creating variation in the emissivity spectral index \\citep{Ossenkopf94}.  In addition, below a density of $\\sim3 \\times 10^4$ the dust temperature is not coupled to the gas temperature \\citep{Goldsmith01, Galli02}, which is the more important quantity because the gas makes up $\\sim$99\\% of the mass.  One of the best tracers of gas temperature is the rotation-inversion lines of \\amm. Ammonia is particularly well suited for this type of study because 1) \\amm\\ does not seem to deplete onto dust grains in starless cores, and in fact may have an enhanced abundance where CO is depleted \\citep{Tafalla02} 2) \\amm\\ has hyperfine splitting of its inversion lines, making fits to the temperature, velocity and optical depth more certain 3) the \\amm\\ (1,1) and (2,2) transitions are separated by a small difference in frequency, so it is possible to simultaneously observe both sets of lines.  In this paper we use ammonia (1,1) and (2,2) observations \\citep{Rosolowsky08} to derive the temperature of starless cores.  We examine whether these data support the several models of the relationship between core properties and the molecular cloud where they are found \\citep[e.g.,][]{Evans01, Bethell04,Stamatellos04}  Perseus is one of the Gould's belt molecular clouds included in the Spitzer legacy project ``From Molecular Cores to Planet-Forming Disks'' (c2d) \\citep{Evans03} and the COMPLETE survey of star-forming regions \\citep{Ridge06}, at a distance of 250 pc \\citep{Cernis03, Belikov02}.  Because it is so well observed, there is a wealth of data on the cores and the cloud that enables us to study the influences of various internal and environmental factors on the temperature of 50 starless cores in Perseus. ",
        "conclusions": "\\label{CONCLUSIONS}  In this paper we investigate the temperature distribution of starless cores in Perseus.  Although, the range of temperatures derived from \\amm\\ (1,1) and (2,2) pointing is quite narrow, from $\\sim 9-15$ K, there are several interesting correlations and lack of correlations.  The location of a starless core does affect its temperature.  We find that those cores inside the two main star-forming regions in Perseus (IC348 and NGC1333) are warmer than those not in clusters.  Although there is no correlation between \\tc\\ and the distance to the nearest class 0/I protostar for those starless cores in clusters, starless cores in the field are colder the further away they are from protostars.  Certain properties of the molecular cloud are related to the temperatures of the starless cores embedded in it.  \\tc\\ is higher when the extinction-derived column density of the molecular cloud is larger, despite the additional shielding from the ISRF that the higher column density implies, which may be due to the increased pressure on the cores from the molecular cloud in higher column density regions. This correlation is not seen when column density along the line of sight in Perseus is measured from far-infrared dust emission.  After controlling for the cluster/field temperature difference, we find that \\tc\\ does not depend on the temperature of the molecular cloud, as measured by CO or dust emission.  The clusters have higher values of \\tco\\ and \\td\\ on average, so presumably the same processes that are heating the clusters are also heating the cores within them.  Those starless cores outside NGC1333 and IC348 are warmer in regions with higher $^{13}$CO linewidth, suggesting the the turbulence causing the high linewidths makes the cloud more porous to the ISRF, which in turn heats the embedded cores.  There is no significant correlation between $\\sigma_{CO}$ and \\tc\\ for those cores in the clusters.  Some internal properties of starless cores are correlated with their temperature. For those starless cores in the field, total 1.1mm flux is higher in cores that are warmer, and peak column density and average volume density are higher in cores that are colder.  A higher temperature naturally leads to more 1.1mm flux for a given core mass, and the correlations between \\tc\\ and peak core column density and volume density can be explained by the amount of shielding from the ISRF that the additional material provides.  However, it is surprising that none of these correlations exist for starless cores inside the clusters.  Support was provided to MLE by NASA through the Spitzer Space Telescope Fellowship Program."
    },
    "0809/0809.0190_arXiv.txt": {
        "abstract": "The concept of ``age'' as a parameter for the description of the    state of  development of  high energy showers in the atmosphere has been in use  in cosmic  ray  studies for  several decades. In  this  work we briefly  discuss     how this concept, originally  introduced to  describe the average behavior of  electromagnetic  cascades,  can be fruitfully applied to  describe  individual  showers   generated by  primary particles  of different  nature, including protons, nuclei and neutrinos. Showers with  the same age share three different important properties: (i)  their electron size  has the same fractional rate of   change  with increasing depth, (ii) the  bulk of  the electrons and  photons in  the shower (excluding  high  energy particles)  have energy spectra with shapes  and  relative normalization uniquely determined by  the age parameter, (iii) the electrons and photons in the shower  have also  the same angular  and lateral distributions  sufficiently far from the shower axis. In this  work we  discuss how  the properties associated with the shower  age can be  understood with simple arguments, and how the shapes of the  electron and photon spectra  and the  relative normalization   that correspond to a   certain age can be calculated  analytically. ",
        "introduction": "The concept of the ``age'' of a shower  has  been in use in the cosmic  ray community for more  than half a  century. The concept  first emerged \\cite{Rossi-Greisen}  in the study   of the average longitudinal   development of purely  electromagnetic  showers  generated   by photons or electrons. It was then also  applied \\cite{Nishimura-Kamata,Greisen1} to the lateral distribution of electrons around the shower axis. Soon,  it was also understood that  it is possible and useful to assign an ``age'' also to individual  showers, and that the concept is  applicable also to showers generated  by hadronic   primary particles  such as protons or nuclei.  Some  recents  works \\cite{Giller,Nerling,Gora} have   rediscussed the  concept of of shower  age for  the showers  generated  by ultra high energy cosmic rays in the Earth's atmosphere. Giller et al   \\cite{Giller}  and Nerling et al. \\cite{Nerling} have studied with  montecarlo methods the showers generated  by high energy  protons  and  nuclei  in air, and  have  observed  that the energy and angle distributions  of the electrons   (in this  work with ``electrons''  we   will refer to  the sum of electrons and positrons) in the showers have shapes that to a   good  approximation  are only determined  by an  age  parameter  $\\overline{s}$    defined  as: \\begin{equation} \\overline{s} (t, t_{\\rm max} ) = \\frac{3 \\,t}{t + 2 \\; t_{\\rm max}} ~. \\label{eq:s0} \\end{equation} where  $t$ is  the depth  in unit of radiation lengths and  $t_{\\rm max}$ is the depth where the shower  reaches its maximum  size. The shower size  is defined  as the  total number of charged particle integrated over all energy, and effectively coincides with the electron size. These results have been extended to the lateral distribution of  electrons by Gora et al. \\cite{Gora}. This property of ``universality''  is obviously very  important for the analysis and interpretation of  high energy cosmic  ray  observations, and it is therefore  desirable to have a deeper understanding of its origin and   of its limitations.  In this   work we want to review  critically the  concepts of shower ``age''  and ``universality''.  One of the main points we want to   make is to   argue that the  definition of age of equation (\\ref{eq:s0}), while reasonably accurate in most cases, is in general not correct, and should be  replaced by a    better  motivated and more accurate definition.  The  essence of the   concept of shower age  can be understood observing that    all  showers  at the maximum of their development are  in  an appropriate sense ``similar''  to each other  (that is have the same ``age''). This ``similarity'' is  represented by the  fact that in all showers at maximum   the energy spectra of ``most''  electrons and  photons have   the  same  shape and  the same relative normalization. These particles  also have  the same  angular  distributions (that are  obviously strongly correlated with energy) and, for  an equal  density profile of the medium where the  shower is propagating, also the same lateral distribution around the shower  axis.  The   idea  that all  showers at  maximum, that is at the  depth where the   derivative  of the shower size $N(t)$ vanishes, are ``similar'' independently from the   energy and nature of the primary particle, can be naturally be generalized, stating that all  showers   that have the same fractional    rate of  change with depth, that is the same ``size slope'' $\\lambda$: \\begin{equation} \\lambda =  \\frac{1}{N(t)} \\; \\frac{dN(t)}{dt} \\end{equation} are also ``similar''. This  means that to each  value of the  $\\lambda$ correspond well determined shapes of the electron and  photon spectra (again only valid   for ``most''  particles) and  a   well determined relative normalization for the two populations. It is  intuitive (and will be later  verified by  detailed calculations) that    the  energy  spectra of electrons and photons  become  progressively  softer as   $\\lambda$ decreases     going from   positive (when the shower  size  grows)  to  negative (when the shower size decreases)   values.  There is a  one to one mapping between the   values of the size slope  $\\lambda$  and    the values of the shower ``age'' $s$.  This  mapping is   encoded  by  a  function that, in the notation introduced by Rossi and Greisen  in their ``classic'' paper \\cite{Rossi-Greisen}, is  called  $\\lambda_1(s)$. The     general  definition of the shower age is therefore: \\begin{equation} s = \\lambda_1^{-1} ( \\lambda)  = \\lambda_1^{-1} \\left ( \\frac{1}{N(t)} \\; \\frac{dN(t)}{dt} \\right ) \\label{eq:age-def} \\end{equation} where  $\\lambda_1^{-1}$ is the inverse function  of $\\lambda_1(s)$. The   function $\\lambda_1(s)$ (that  will be discussed in more detail in the following) is  monotonically decreasing and   has a single  zero   at $s = 1$, therefore   according to equation (\\ref{eq:age-def}) showers  have   age   $s=1$ at maximum  and  age $s < 1$ ($s >1$)   before  (after) maximum.  The definition of  the age parameter (\\ref{eq:age-def})   may seem at first sight (and in some  sense actually is) arbitrary,   since the   size slope  $\\lambda$ itself or any  monotonic  function  of $\\lambda$ are also   perfectly  adequate to  identify ``similar''  showers. The choice of the particular mapping of    equation (\\ref{eq:age-def}) is   motivated by the fact that  one can attach a  direct physical  meaning   to the  quantity $s$. The shapes  of the electron and photon spectra   for  $E$ above the critical energy $\\varepsilon$ (the electron  critical energy $\\varepsilon$ is the average energy lost by an electron in a radiation length, and corresponds also to the electron energy for which radiative and collision losses are equal) and  $E \\ll E_0$    (with $E_0$ the primary particle energy)  are  well  represented by a power law: \\begin{equation} n_e (E) \\sim n_\\gamma (E) \\sim E^{-(s+1)} \\end{equation} These power  law  behaviors  stops when $E$ approaches (from above) the electron critical energy. For  energies  below $E \\sim \\varepsilon$ the electron spectrum has a sharp cutoff, while the photon spectrum has a ``knee''  and  takes the form $E^{-1}$. The precise  shapes of the  cutoff of the electron spectrum, of  the knee of the  photon spectrum (that is the transition    from the  form $E^{-(s+1)}$ to the   $E^{-1}$) and  the relative normalizations of the photon and electron spectra  are  all entirely determined  by the age  $s$ (or equivalently  by the size slope $\\lambda$) and can  be computed in   detail.  The  commonly used   definitions  of  age in equation (\\ref{eq:s0}) always coincides  with the general  definition (\\ref{eq:age-def}) at shower maximum and therefore,   by construction, it is a   reasonably good approximation  for  showers   sufficiently close to maximum. The motivation for the more  general   definition may appear as only a   formal question of  ``principle''. In fact in some  circumstances the  two definitions are significantly different, and the general  definition gives the correct  shower age. The definition (\\ref{eq:s0})  coincides  with the  correct one only for a particular  shape of  the shower longitudinal  that  is known as   the ``Greisen  profile'' \\cite{Greisen1}. In fact  Greisen profile and the age definition (\\ref{eq:s0}) are intimately connected, and can be seen  (with the mediation  of the function $\\lambda_1$) as  the integral  and the  derivative of each other. The Greisen profile (discussed   below  in section~\\ref{sec:greisen}) describes accurately the average development of purely electromagnetic  showers, but  is only a    rough approximation  for  the  description of individual  hadronic showers, and naturally fails  completely in the description of neutrino--induced showers. The deviations of the  definition (\\ref{eq:s0}) from the true age  (\\ref{eq:age-def}) are of  the same order of the deviations of the profile of a shower from  the Greisen  profile that has the same $t_{\\rm max}$.  The authors  in \\cite{Giller,Nerling} have  calculated with montecarlo methods the shape of the electron spectra in hadronic showers  of different age. The parametrizations of their results  are   essentially identical  to the  shapes  of  the electron spectra   in  showers of the same  age calculated    several  decades  ago  by Rossi and Greisen \\cite{Rossi-Greisen}. These  modern  works  have therefore effectively only ``rediscovered'' with montecarlo methods    what should  be called the ``Rossi--Greisen''  spectra. We want to  attract attention to this   fact  for  three reasons. The first one is  that it  is  obviously appropriate   to give credit  to  the  remarkable work of  the  pioneers. The second is that the works of   \\cite{Giller,Nerling} do  not include a discussion of   photon spectra. The shapes of these spectra and their relative normalization  with respect to the  electron ones are also  determined  unambiguously by the shower  age, and have been also  computed explicitely  by Rossi and Greisen. Finally the   derivation of   the spectral shapes  obtained by Rossi and Greisen with analytic  methods allows   physical  insights on  the  origin  and  limitations  of the ``universality'' of the spectra,  that are not    easily deducible from a  montecarlo calculation.   It should be stressed   that  the ``universality'' of properties for  cascades  of the same age has  clearly limitations  since it only applies to  ``most''  but not all  particles in  the  shower. For example, the ``similarity''  among showers at   the maximum of their  development, (that is at age $s=1$) does not imply that the showers simply differ   by the  absolute normalization  of their  electromagnetic  component. Considering at first the case of purely electromagnetic cascades, at maximum the  showers  generated by a photon of initial  energy $E_0$ contain (in essentially all cases) more high energy particles   than the showers generated by photons of lower energy. These   high energy particles are negligible in number   and  do not contribute   significantly  to the total size but in general  carry an important fraction of the shower energy, they  ``feed''  the shower development and  influence its  development. The showers generated by other  types of  primary particles have ``cores''  of  different  structure and particle content,  and    follow different   development profiles.  \\vspace{0.3 cm} This  work is  organized as follows: in the next   two sections  we review  a very well known  subject, discussing   the average longitudinal  evolution of  purely electromagnetic  showers first in ``approximation~A'',  that is  neglecting the  electron ionization losses, and then in ``approximation~B''. The concept of age  emerged   naturally   in these studies. In approximation~B   the shower equations have ``elementary solutions'' labeled  by the   parameter $s$, these  solutions correspond to the ``universal spectra'' of showers with age $s$. The following  section  discusses the well known ``Greisen profile'' that  describes the average longitudinal development of  purely  electromagnetic showers. Finally we   discuss the evolution of individual hadronic  showers, and  give some  conclusions. ",
        "conclusions": "The concept of   the shower age can be very useful   for the analysis  of high energy cosmic ray data.  The essence of the idea is  very  simple and can  be  summarized  in a nutshell saying that the $t$--slope  $\\lambda$ and the $E$--slope  $s$    of a shower are connected to each other  by a one to one mapping. The $t$--slope (or size slope)  is the fractional  rate of change of the shower size with increasing depth ($\\lambda = N^{-1} \\, dN/dt$).  The $E$--slope (or energy slope) is the  integral  slope of  the (power law)  energy spectra of photons or electrons   above the critical energy. The mapping   between   $\\lambda$ and $s$   is given by $\\lambda = \\lambda_1(s) \\simeq (s -1 - 3 \\, \\ln s)/2$. The spectra of photons and electrons have a more complex  shape  around and below  $\\varepsilon$ that is also determined  by $s$ (or $\\lambda$), and have a  relative normalization also determined by $s$ (or  $\\lambda$). It is remarkable that the  electron and photon spectra that correspond to different $s$ (or $\\lambda$)   have been calculated accurately  with analytic  methods   by  the pioneers Rossi and Greisen many decades  ago.  These properties of ``universality''   extend to the  angular and lateral  distributions of electrons and photons. This  crucially important  subject is not discussed  here (see appendix~\\ref{sec:lateral}  for  some remarks).  The  definition of age  discussed   here is independent  from  the  shape of the longitudinal development  of a  shower, and is therefore more general and accurate that the commonly used    definition $s \\simeq  3 \\, t/(t + 2 \\, t_{\\rm max})$ that is  correct  only when  the   shower development is described by the  ``Greisen profile''. The average shape (and the fluctuations around this  average) of  the longitudinal development of high energy cosmic  ray showers is determined by the nature of the primary particles and by the properties of   hadronic  interactions. The  observation of  these  shape  is  a very important subject for  future   experimental studies.  A  definition of age that depends only on the derivative of the shower  size  can be applied  also to neutrino--induced showers, and  more  generally  is expected to remain valid for the description of  all shower  where the size is dominated by  electrons. This includes showers  generated  by exotic primaries, or the presence of unexpected  physics (or unexpected fluctuations)  in the development of  the showers by  primary particles of known nature. The search for   events  that have  unusual longitudinal   developments, such  as   multiple maxima is  an  interesting  direction of   research. It is  likely (and at least the best   possible {\\em a priori} assumption) that the spectra  of the electromagnetic  component around and below the critical energy will, also  in these cases,  be controled   by  the shower  age.   These ideas   can be  useful in the analysis  of high energy  cosmic  ray  observations in several ways. As examples: (i) the knowledge of the   variations of the electron  energy spectrum   during the evolution of a  shower can be used to  obtain a better reconstruction of the longitudinal profile  of the shower in observations   that  use fluorescence and/or Cherenkov light   detectors (see \\cite{Unger:2008uq} for more discussion); (ii) the reconstruction of the  shower age  from the lateral  distribution  of  its electromagnetic component can in principle help in the reconstruction of the energy in surface  array  measurements; (iii) in case  of hybrid measurements   of the showers, the   redundant measurement of the age (from the  size longitudinal   development and the lateral distribution of the electromagnetic  component at the  ground)  can allow to   disentangle a muon component, allowing   composition measurements, or test of hadronic interaction  models (see   \\cite{Engel:2007cm} for more discussion).   It should  finally  be stressed  that the  ``universality'' in the electromagnetic  component of high energy showers, is clearly   an important  analysis  tool, but gives  only a partial information about the shower. Other information is  contained in  the shower muon component,   moreover the shower core, that  is  essentially  undetected in large area shower  arrays,   in most cases also contains  a significant amount of energy. The   energy contained in the core must be  ``infered'' from the  information obtained at large distances from the  shower axis,  introducing unavoidably some model  dependence in the energy reconstruction.  \\vspace{0.4 cm} \\noindent {\\bf Acknowledgments.}\\\\ It is a pleasure to  acknowledge fruitful discussions with Ralph Engel, Maurizio Lusignoli  and Silvia Vernetto.   \\newpage \\appendix"
    },
    "0809/0809.3310.txt": {
        "abstract": "A global lopsidedness in the distribution of the stars and gas is common in galaxies. It is believed to trace a non-equilibrium dynamical state caused by mergers, tidal interactions, asymmetric accretion of gas, or asymmetries related to the dark matter halo. We have used the Sloan Digital Sky Survey (SDSS) to undertake an investigation of lopsidedness in a sample of $\\sim$25,000 nearby galaxies ($z <$ 0.06). We use the $m=1$ azimuthal Fourier mode between the 50\\% and 90\\% light radii as our measure of lopsidedness. The SDSS spectra are used to measure the properties of the stars, gas, and black hole in the central-most few-kpc-scale region. We show that there is a strong link between lopsidedness in the outer parts of the galactic disk and the youth of the stellar population in the central region. This link is independent of the other structural properties of the galaxy. These results provide a robust statistical characterization of the connections between accretion/interactions/mergers and the resulting star formation. We also show that residuals in the galaxy mass-metallicity relation correlate with lopsidedness (at fixed mass, the more metal-poor galaxies are more lopsided). This suggests that the events causing lopsidedness and enhanced star formation deliver lower metallicity gas into the galaxy's central region. Finally, we find that there is a trend for the more powerful active galactic nuclei (the more rapidly growing black holes) to be hosted by more lopsided galaxies (at fixed galaxy mass, density, or concentration). However if we compare samples matched to have both the same structures {\\it and central stellar populations}, we then find no difference in lopsidedness between active and non-active galaxies. Indeed the correlation between the youth of the stellar population and the rate of black hole growth is stronger than the correlation between lopsidedness and either of these other two properties. This leads to the following picture. The presence of cold gas in the central region of a galaxy (irrespective of its origin) is essential for both star-formation and black hole growth. The delivery of cold gas is aided by the processes that produce lopsidedness. Other processes on scales smaller than we can probe with our data are required to transport the gas to the black hole. ",
        "introduction": "In the standard $\\Lambda CDM$ universe, galaxies continue to grow significantly at the current epoch, accreting an average of roughly 5\\% of their mass per Giga-year (e.g., \\citealt{k+05,k+06}). Minor mergers, tidal interactions with close companions, and asymmetric accretion of cold gas can all perturb the underlying structure of the dark matter halo, the stars, and the gas in galaxies for time-scales of-order a Giga-year (e.g., \\citealt{zr97,lev+98,jog99,kr+02,bou+05,dim07,map+08,cox08}).  A characteristic observational signature of these types of non-equilibrium situations is a global asymmetry (lopsidedness) in the distribution of the stars and/or gas. Indeed, galaxies commonly exhibit asymmetric structures. Asymmetry in the global \\ion{H}{1} 21cm emission-line profile is present in half or more of disk galaxies \\citep{ric+94,mat+98}. Lopsided distributions of the \\ion{H}{1} gas are revealed by spatial maps of galaxies in the field (e.g., \\citealt{swa+99}) and in groups \\citep{ang+06,ang+07}.  The stellar mass distributions of disk galaxies are often lopsided as well. \\citet{zr97} showed that $\\sim$30\\% of a sample of 60 field spiral galaxies had a significant $m=1$ azimuthal Fourier component in the stellar mass. \\citet{rud+98} showed that lopsidedness in the stellar mass distribution was common in early-type disk galaxies. Simulations of spiral galaxies have shown that the typical measured amplitudes of lopsidedness cannot be explained by internal dynamical mechanisms alone (e.g. \\citealt{bou+05}). The external processes described above are required.  It has long been recognized that tidal interactions and mergers can act as triggers for star-formation \\citep{too+72,lar+78}.  The resulting non-axisymmetric time-dependent gravitational potential can transport angular momentum out of the gas and also lead to shocks in which the gas is compressed and kinetic energy is ultimately converted into radiation and lost to the system (e.g., \\citealt{mih+96,dim07,cox08}). These processes lead to the inflow of gas, an increase in gas density (either globally or locally), and hence an increase in the star formation rate \\citep{ken98}.  \\citet{li+08a} have investigated the link between interactions and star formation in the largest sample of galaxies studied to date. They found a strong correlation between the proximity of a near neighbor (closer than 100 kpc) and the average specific star formation rate. This result pertains to the relatively early stages of an eventual merger. Using lopsidedness as a probe would allow us in principal to measure the history of the enhancement in star formation (in an time-averaged sense) all the way through to the post-merger phase. \\citet{zr97} exploited this idea, and found a correlation between lopsidedness and excess $B$-band luminosity emitted by a young stellar population.  \\citet{rud+00} showed that both recent (the past 1 Gyr) and ongoing star formation are correlated with lopsidedness, and estimated that as much as 10\\% of stellar mass in typical galaxies can be formed in such events. However, these papers investigated very small samples compared to the \\citet{li+08a} study.  An additional link between lopsidedness and star formation is indicated by the results in \\citet{k+05} and \\citet{bou+05}. The former authors use SPH simulations to conclude that the star formation history of galaxies is primarily regulated by the accretion rate of cold gas, rather than by mergers. The accretion of this cold gas will not occur in a spherically symmetric fashion, and \\citet{bou+05} argue that this will lead to long-lived lopsidedness in disk galaxies.  Large samples are crucial to the investigation of the connection between lopsidedness and the recent star formation history. In our recent paper (\\citealt{paper1}, hereafter Paper I) we reported on our analysis of multi-color Sloan Digital Sky Survey (SDSS) images of $\\sim$25000 low-redshift ($z < 0.06$) galaxies. We confirmed that Lopsidedness in the stellar mass distribution is indeed very common, and quantified its distribution as a function of the principal structural properties of the galaxies. We found that galaxies of lower mass, lower density, and lower concentration are systematically more lopsided. Similarly, \\citet{kau+03a,kau+03b} showed that galaxies with lower mass, density, and concentration have younger stellar populations on-average. Taking these results together, the causal connections between star formation, lopsidedness, and other galaxy structural properties are ambiguous. Unraveling this web of mutual correlations requires the careful analysis of a large and homogeneous sample.  The strong correlation between the stellar mass of a galaxy and the metallicity of its interstellar medium \\citep{tre+04} provides a powerful constraint on the chemical evolution of galaxies and the intergalactic medium (e.g., \\citealt{dal07}). If lopsidedness in a galaxy has been recently induced by either a minor merger with a gas-rich low mass companion galaxy or the accretion of cold intergalactic gas, then one signature would be a decrease in the metallicity of the interstellar medium in lopsided galaxies compared to other galaxies of the same mass. This can be tested with a large sample of galaxies.  Interactions and mergers are also believed to be an important mechanism for fueling the formation and growth of a central supermassive black hole. This idea dates back at least to Toomre \\& Toomre (1972), and is one of the cornerstones of contemporary models that attempt to explain the co-evolution of galaxies and black holes within the context of the hierarchical build-up of structure (e.g., \\citealt{hop+07,dim+07}). However, the observational verification of this idea remains insecure. At high-redshift the data are still too sparse to be conclusive, while at low-redshift there have been many past studies that have led to contradictory results. \\citet{li+08b} have undertaken an analysis of the largest low-redshift sample to date, based on $10^5$ galaxies in the SDSS. They find a strong link between close companions and star formation, but none between close companions and AGN. One possibility is that the black hole growth occurs only late in a merger, after the companion galaxy has been captured. This idea can be tested using lopsidedness and a similarly large sample since lopsidedness can be used to trace later stages of interaction than that seen in pair studies.  In this study, we use the large data-set described in Paper I to try to understand how lopsidedness may be linked to star formation, metallicity, and the AGN phenomenon. We describe our specific galaxy sample in \\S\\ref{sec:data} and summarize our measurements of the structural parameters, star formation indicators, metallicity, and AGN indicators. Next, in \\S\\ref{sec:sfh}, we compare the correlations between star formation and structure with an emphasis on separating out the dependence of star formation on lopsidedness versus other galaxy structural properties. Next, in \\S\\ref{sec:met} we explore whether the residuals in the mass-metallicity relation correlate with lopsidedness, and test whether this is a fundamental correlation, or only a secondary one induced through a mutual dependence of metallicity and lopsidedness on galaxy structure. Finally, in \\S\\ref{sec:agn}, we compare lopsidedness and the structural and stellar population properties for the host galaxies of AGN, to test for a possible link between interactions/mergers/accretion and the growth of black holes.  We discuss the implications of our conclusions in \\S\\ref{sec:conclusions}. ",
        "conclusions": "}  We have studied a sample of approximately 25000 low-redshift ($z < 0.06$) galaxies from the Sloan Digital Sky Survey (SDSS) Data Release 4 to investigate the links between large-scale asymmetries in the stellar mass distribution in galaxies (lopsidedness), star formation, the metallicity of the interstellar medium, and the presence of active galactic nuclei (AGN). Lopsidedness has been defined as the radially averaged $m=1$ azimuthal Fourier amplitude ($A_1$) measured between the radii enclosing 50\\% and 90\\% of the galaxy light in the SDSS images (i.e. in the outer part of the galaxy). We have previously shown that lopsidedness traces the distribution of the underlying stellar mass (Paper I). Lopsidedness is a signpost of a non-equilibrium global dynamical state, and can be induced by mergers, asymmetric accretion of cold gas, tidal interactions, or underlying asymmetries related to the dark matter halo. We have used spectra obtained through the SDSS 3\\arcsec diameter fibers (typical projected diameter of $\\sim$ 3 kpc) to characterize the stars, gas, and AGN in these central regions.  We have used the amplitude of the 4000 \\AAA break, the strength of the high-order Balmer absorption-lines, and the specific star formation rate derived from the nebular emission-line to characterize the stellar population in the central few-kpc-scale region. We found strong links between lopsidedness and recent/on-going central star-formation: galaxies with younger stellar populations are more lopsided and more lopsided galaxies have younger stellar populations. Starburst and post-starburst galaxies are the most lopsided on average.  We have previously shown that galaxies with lower surface mass density and mass are more lopsided (Paper I), and that galaxies with lower mass and lower density have younger stellar populations \\citep{kau+03b,bri+04}. Here, we have shown that there is still a strong correlation between lopsidedness and the age of the stellar population even after their mutual dependences on these other structural parameters have been removed. These results are consistent with other evidence that mergers and tidal interactions trigger central star-formation, but place these results on a firm statistical base. They are also consistent with the idea that star formation in galaxies today is regulated by the accretion of cold gas \\citep{k+05} which can excite lopsidedness \\citep{bou+05}. They imply that the timescale for the transport of gas into the central region can not be significantly longer than the timescale over which lopsidedness persists in the outer disk. This is consistent with recent numerical simulations of star formation in minor mergers and galaxy interactions \\citep{dim07,cox08}.  We have also shown that at fixed stellar mass, more lopsided galaxies have systematically lower gas-phase metallicities by about 0.1 dex on average. This suggests that the processes causing lopsidedness deliver lower metallicity gas into the central region, but the small amplitude of the effect rules out extreme events as being typical.  Finally, we found that there is a trend for the more rapidly growing black holes to be hosted by galaxies with higher average lopsidedness. This is true when comparing galaxies at fixed galaxy mass, surface mass density, and concentration. However, when the AGN hosts were compared to non-active galaxies that were also matched in the age of the stellar population in their central region, we found no excess lopsidedness in the AGN hosts. The strongest link is that between the youth of the stellar population and the growth rate of the black hole. The correlations of these two properties with lopsidedness are weaker. This suggests that the presence of cold gas in the central few kpc-scale region (which is facilitated through the processes that produce lopsidedness, and which leads to significant star formation) is a necessary but not sufficient condition for subsequent fueling of the growth of the black hole. Other processes are subsequently required to deliver the gas from scales of a few kpc all the way to the black hole accretion disk and these would not be directly related to the process(es) that produced lopsidedness.  Combining our results with the analysis by \\citet{li+08b} and Ellison et al. (2008a) of the role of close companion galaxies in driving star formation and fueling black holes, suggests that the period of black hole growth may be preferentially associated with the end stages of a minor merger. This idea will be tested in a future paper.  We would like to thank Christy Tremonti for reading a draft of this manuscript.  Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.     \\begin{deluxetable}{cclr} \\tabletypesize{\\small} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Partial Correlation Coefficients: Lopsidedness and Star Formation} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{Dependence} & \\colhead{Partial} \\\\ \\colhead{Par. 1} & \\colhead{Par. 2} & \\colhead{Removed} & \\colhead{Corr. Coeff.} } \\startdata $\\log A_1^i$ & $D_{4000}$ & \\nodata & $-0.58$ \\\\ $\\log A_1^i$ & $H\\delta_A$ & \\nodata & $ 0.52$ \\\\ $\\log A_1^i$ & $PC_1$ & \\nodata & $-0.60$ \\\\ $\\log A_1^i$ & $PC_2 - PC_1$ & \\nodata & $ 0.46$ \\\\ $\\log A_1^i$ & $\\log SFR/M_*$ & \\nodata & $ 0.33$ \\\\ $\\log A_1^i$ & $g-i$ & \\nodata & $-0.57$ \\\\ \\hline $\\log A_1^i$  & $D_{4000}$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $-0.30$ \\\\ $\\log A_1^i$  & $H\\delta_A$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $ 0.27$ \\\\ $\\log A_1^i$  & $PC_1$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $-0.33$ \\\\ $\\log A_1^i$  & $PC_2 - PC_1$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $ 0.22$ \\\\ $\\log A_1^i$  & $\\log SFR/M_*$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $ 0.15$ \\\\ $\\log A_1^i$  & $g-i$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $-0.21$ \\\\ \\hline $\\log A_1^i$ & $\\log M_*$ & $D_{4000}$ & $-0.14$ \\\\ $\\log A_1^i$ & $\\log \\mu_*$ & $D_{4000}$ & $-0.26$ \\\\ $\\log A_1^i$ & $C_i$ & $D_{4000}$ & $-0.17$ \\\\ \\hline $\\log A_1^i$ & $H\\delta_A$ & $D_{4000}$ & $ 0.11$ \\\\ $\\log A_1^i$ & $PC_2$ & \\nodata & $-0.20$ \\\\  \\enddata \\label{tab:parcor-sfr} \\end{deluxetable}    \\begin{deluxetable}{cclr} \\tabletypesize{\\small} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Partial Correlation Coefficients: Lopsidedness and Metallicity} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{Dependence} & \\colhead{Partial} \\\\ \\colhead{Par. 1} & \\colhead{Par. 2} & \\colhead{Removed} & \\colhead{Corr. Coeff.} } \\startdata $\\log A_1^i$ & $12 + \\log (O/H)$ & $\\log M_*$ & $-0.15$ \\\\ $\\log A_1^i$ &  $12 + \\log (O/H)$  & $\\log M_*$, $\\log \\mu_*$ & $-0.10$ \\\\ $\\log \\mu_*$ & $12 + \\log (O/H)$ & $\\log M_*$ & $ 0.20$ \\\\ $\\log \\mu_*$ &  $12 + \\log (O/H)$  & $\\log A_1^i$, $\\log M_*$ & $ 0.19$ \\\\ $\\log SFR/M_*$ & $12 + \\log (O/H)$ & $\\log M_*$ & $ 0.02$ \\\\ \\enddata \\label{tab:parcor-oh} \\end{deluxetable}  \\begin{deluxetable}{cclr} \\tabletypesize{\\small} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Partial Correlation Coefficients: Lopsidedness and AGN Activity} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{Dependence} & \\colhead{Partial} \\\\ \\colhead{Par. 1} & \\colhead{Par. 2} & \\colhead{Removed} & \\colhead{Corr. Coeff.} } \\startdata $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & \\nodata & $ 0.32$ \\\\ \\hline $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $\\log M_*$ & $ 0.30$ \\\\ $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $\\log \\mu_*$ & $ 0.23$ \\\\ $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $C_i$ & $ 0.19$ \\\\ $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $D_{4000}$ & $ 0.09$ \\\\ \\hline $\\log A_1^i$  & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $\\log M_*$, $\\log \\mu_*$, $C_i$ & $ 0.22$ \\\\ $\\log A_1^i$ &  $\\log (L$[\\ion{O}{3}]$/M_{BH})$  & $\\log M_*$, $D_{4000}$ & $ 0.09$ \\\\ $\\log A_1^i$ &  $\\log (L$[\\ion{O}{3}]$/M_{BH})$  & $\\log \\mu_*$, $D_{4000}$ & $ 0.07$ \\\\ $\\log A_1^i$ &  $\\log (L$[\\ion{O}{3}]$/M_{BH})$  & $C_i$, $D_{4000}$ & $ 0.05$ \\\\ $\\log A_1^i$ & $\\log (L$[\\ion{O}{3}]$/M_{BH})$ & $\\log M_*$, $\\log \\mu_*$, $C_i$, $D_{4000}$ & $ 0.06$ \\\\  \\enddata \\label{tab:parcor-agn} \\end{deluxetable} \\clearpage"
    },
    "0809/0809.4193_arXiv.txt": {
        "abstract": "We present the first {\\em Swift} Ultra-Violet/Optical Telescope (UVOT) gamma-ray burst (GRB) afterglow catalog. The catalog contains data from over $64,000$ independent UVOT image observations of 229 GRBs first detected by {\\em Swift}, the {\\em High Energy Transient Explorer~2} (HETE2), the {\\em INTErnational Gamma-Ray Astrophysics Laboratory} (INTEGRAL), and the Interplanetary Network (IPN). The catalog covers GRBs occurring during the period from 2005 Jan 17 to 2007 Jun 16 and includes $\\sim86\\%$ of the bursts detected by the {\\em Swift} Burst Alert Telescope (BAT). The catalog provides detailed burst positional, temporal, and photometric information extracted from each of the UVOT images. Positions for bursts detected at the $3\\sigma$-level are provided with a nominal accuracy, relative to the USNO-B1 catalog, of $\\sim0\\farcs25$. Photometry for each burst is given in three UV bands, three optical bands, and a `$white$' or open filter. Upper limits for magnitudes are reported for sources detected below $3\\sigma$. General properties of the burst sample and light curves, including the filter-dependent temporal slopes, are also provided. The majority of the UVOT light curves, for bursts detected at the $3\\sigma$-level, can be fit by a single power-law, with a median temporal slope ($\\alpha$) of 0.96, beginning several hundred seconds after the burst trigger and ending at $\\sim1\\times10^5 {\\rm ~s}$. The median UVOT $v$-band ($\\sim5500 {\\rm ~\\AA}$) magnitude at $2000 {\\rm ~s}$ for a sample of ``well\" detected bursts is 18.02. The UVOT flux interpolated to $2000 {\\rm ~s}$ after the burst, shows relatively strong correlations with both the prompt {\\em Swift} BAT fluence, and the {\\em Swift} X-ray flux at $11 {\\rm ~hours}$ after the trigger. ",
        "introduction": "The UVOT utilizes seven broadband filters during the observation of GRBs. The characteristics of the filters --- central wavelength ($\\lambda_c$), FWHM, zero points (the magnitudes at which the detector registers $1 {\\rm ~count\\,s^{-1}}$; $m_z$), and the flux conversion factors ($f_{\\lambda}$) --- can be found in Table~\\ref{tab5} \\citep{PTS2007,RPWA2005}. The flux density conversion factors are calculated based on model GRB power law spectra with a redshift ranging from $0.3 < z < 1.0$ \\citep{PTS2007}\\footnote{The most recent calibration data are available from the {\\em Swift} calibration database at http://swift.gsfc.nasa.gov/docs/heasarc/caldb/swift/}. The nominal image scale for UVOT images is $0\\farcs502 {\\rm ~pixel^{-1}}$ (unbinned). UVOT data is collected in one of two modes: event (or photon counting) and image. Event mode captures the time of the arriving photon as well as the celestial coordinates. The temporal resolution in this mode is $\\sim11\\,{\\rm ~ms}$. In image mode, photons are counted during the exposure and the position is recorded, but no timing information is stored except for the start and stop times of the exposure. Because the spacecraft has limited data storage capabilities, most UVOT observations are performed in image mode since the telemetry rate is significantly lower than event mode observations.  Since the launch of {\\em Swift}, the automated observing sequence of the UVOT has been changed a few times in order to optimize observations of GRBs. The automated sequence is a set of variables which includes, but is not limited to, the filters, modes, and exposure times. The basic automated sequence design consists of finding charts and a series of short, medium, and long exposures in various filters. The finding charts are typically taken in both event and image mode simultaneously, in both the $white$ and $v$ filters, and have exposure times ranging from $100-400\\,{\\rm ~s}$. A subset of these finding charts are immediately telemetered to ground-based telescopes to aid in localizing the GRB\\footnote{A discussion of the finding chart and simultaneous observations in event and image mode is beyond the scope of this paper. The reader is referred to \\citet{RPWA2005}. For observations between 2006 Jan 10 to 2006 Feb 13, and 2006 Mar 15 to the present, a second set of finding charts was included in the sequence. The second set of finding chart exposures are taken in the same way as the first set, except that exposures in the $white$ filter are taken in image mode only.}. After completion of the $white$ and $v$ finding charts, a series of short exposures is typically taken in event mode, in all seven broadband filters, and has exposure times ranging from $10-50\\,{\\rm ~s}$. A series of medium exposures is then taken in image mode, in all seven broadband filters, and has exposure times ranging from $100-200\\,{\\rm ~s}$. Finally, a series of long exposures is taken in image mode, in all seven broadband filters, and has typical exposure times of $900\\,{\\rm ~s}$.\\footnote{Between 2005 Jan 17 to 2006 Jan 9 and between 2006 Feb 24 to 2006 Mar 14, exposures taken in uvw2, uvm2, and uvw1 were taken in event mode.} In all cases, exposures can be cut short due to observing constraints.  This catalog covers UVOT observations of 229 GRB afterglows from 2005 Jan 17 to 2007 Jun 16. It includes bursts detected by {\\em Swift} BAT, HETE2, INTEGRAL, IPN and observed by UVOT. A total of 211 BAT-detected bursts were observed by the UVOT (after instrument turn on) representing 93\\% of the BAT sample. Those that were not observed by the UVOT were either too close in angular distance to a bright ($\\sim 6 {\\rm ~mag}$) source, or occurred during UVOT engineering observations. Not included in the catalog are nine bursts first detected by BAT and INTEGRAL and observed by UVOT but with no afterglow position reported by the XRT or ground based observers (see Table~\\ref{tab6}). Inspection of the UVOT images reveals no obvious afterglows for these bursts.  Hereafter, we adopt the notation $F(\\nu ,t) \\propto t^{-\\alpha} \\nu^{-\\beta}$ for the afterglow flux density as a function of time, where $\\nu$ is the frequency of the observed flux density, $t$ is the time post trigger, $\\beta$ is the spectral index which is related to the photon index $\\Gamma$ ($\\beta = \\Gamma - 1$) , and $\\alpha$ is the temporal decay slope. ",
        "conclusions": ""
    },
    "0809/0809.4470_arXiv.txt": {
        "abstract": "\\noindent  We present the results of an observational study of the efficiency of deep mixing in globular cluster red giants as a function of stellar metallicity. We determine [C/Fe] abundances based on low-resolution spectra taken with the Kast spectrograph on the 3m Shane telescope at Lick Observatory. Spectra centered on the $4300\\mbox{\\AA}$ CH absorption band were taken for 42 bright red giants in 11 Galactic globular clusters ranging in metallicity from M92 ([Fe/H]$=-2.29$) to NGC 6712 ([Fe/H]$=-1.01$). Carbon abundances were derived by comparing values of the CH bandstrength index $S_{2}(CH)$ measured from the data with values measured from a large grid of SSG synthetic spectra. Present-day abundances are combined with theoretical calculations of the time since the onset of mixing, which is also a function of stellar metallicity, to calculate the carbon depletion rate across our metallicity range. We find that the carbon depletion rate is twice as high at a metallicity of [Fe/H]$=-2.3$ than at [Fe/H]$=-1.3$, which is a result qualitatively predicted by some theoretical explanations of the deep mixing process. ",
        "introduction": "The observation that carbon abundance in globular cluster red giants declines continuously as the stars evolve has inspired a great deal of observational and theoretical study (e.g., Suntzeff 1981\\nocite{S81}, Carbon et al. 1982\\nocite{C82}, Trefzger et al. 1983\\nocite{T83}, Suntzeff \\& Smith 1991\\nocite{SS91}, Weiss \\& Charbonnel 2004\\nocite{WC04}, Denissenkov \\& Tout 2000\\nocite{DT00}, Smith \\& Briley 2006\\nocite{SB06}, and similar work). Canonically, abundances should be static  on the red giant branch after the first dredge-up because of the broad radiative zone between the hydrogen-burning shell and the surface. Progressive carbon depletion with rising luminosity on the giant branch is commonly interpreted as a sign of a non-convective ``deep mixing'' process that mixes carbon-depleted material from the hydrogen-burning shell, where the CN(O) cycle is acting, to the surface (e.g., Sweigart \\& Mengel 1979\\nocite{SM79}, Charbonnel 1995\\nocite{C95}, Charbonnel et al. 1998\\nocite{CBW98}, Bellman et al. 2001\\nocite{BBSC01}, Denissenkov \\& VandenBerg 2003\\nocite{DV03}, and others). This same depletion of surface carbon abundance is also observed in red giants in the halo field (e.g., Gratton et al. 2000\\nocite{GSCB00}), and is observed to occur at the same rate in the halo field as in globular clusters with halo-like metallicities (e.g., Smith \\& Martell 2003\\nocite{SM03}).  The process of deep mixing, inferred from observations of low [C/Fe], $\\log \\epsilon$(Li), and $^{12}$C/$^{13}$C, is only observed to occur in stars brighter than the red giant branch (RGB) luminosity function ``bump'' \\citep{CBW98}. There are indications (e.g., Langer et al. 1986\\nocite{L86}, Bellman et al. 2001\\nocite{BBSC01}) that there may be carbon depletion in stars fainter than the RGB bump in the metal-poor globular cluster M92, though there is not an obvious physical explanation for such a phenomenon. During the first dredge-up, in the subgiant phase, the base of the convective envelope drops inward to smaller radius as the stellar core contracts, mixing the partially-processed material of the stellar interior with the unprocessed material at the surface. As hydrogen shell burning progresses, low on the giant branch, the temperature gradient in the star steepens and the base of the convective envelope begins to move outward, leaving behind a sharp jump in mean molecular weight (the ``$\\mu$-barrier'') at the point of its furthest inward progress \\citep{I65}. As the convective envelope retreats outward, this steep $\\mu$ gradient finds itself within a radiative region between the hydrogen-burning shell and the base of the convective zone, where it can potentially hinder mass motions within the radiative zone.  The red giant branch bump is an evolutionary stutter that occurs when the hydrogen-burning shell, which is advancing outward in mass, encounters the $\\mu$ barrier. The sudden influx of hydrogen-rich material to the hydrogen-burning shell causes the star to become briefly bluer and fainter before it re-equilibrates and continues to evolve along the red giant branch (e.g., Iben 1968\\nocite{I68}, Cassisi et al. 2002\\nocite{C02}). In a collection of stars with equal age and composition, this evolutionary loop will result in an unexpectedly large number of stars at a particular magnitude, and a bump in the differential luminosity function. At a fixed mass, the base of the convective envelope sinks lower in higher-metallicity stars during the first dredge-up, meaning that the RGB bump occurs at a fainter luminosity on the RGB in high-metallicity globular clusters than in low-metallicity clusters (e.g., Zoccali 1999\\nocite{Z99}). However, because evolutionary timescales shorten as stellar mass rises, there is a maximum mass of $\\simeq 2$M$_{\\odot}$ for stars to experience this evolutionary loop: above that mass, the hydrogen-burning shell does not move outward far enough to cross the $\\mu$ barrier in the short time the star is on the RGB \\citep{G89}. The fact that deep mixing does not begin until after the hydrogen-burning shell crosses the $\\mu$ barrier is interpreted by, e.g., \\citet{C94} to mean that the gradient of mean molecular weight is the dominant factor in permitting or prohibiting deep mixing. Indeed, \\citet{DV03} point out that $\\nabla \\mu$ is ``the only physical quantity that changes significantly while approaching the hydrogen-burning shell'' in post-bump red giants. \\citet{CPT05} provide a somewhat different perspective: in their maximal-mixing models, the $\\mu$ gradient inhibits mixing on the upper giant branch, but rotational mixing processes are not strong enough on the lower giant branch to cause observable changes in surface abundances, regardless of whether there is a steep $\\mu$ gradient present.  The underlying physical reason for deep mixing is not clear, though rotation has commonly been implicated since \\citet{SM79} proposed meridional circulation as an explanation for CNO anomalies in red giants. Recent theoretical studies tend to focus on specific parametrizations and representations of the process; for example, Rayleigh-Taylor instability \\citep{EDL08}, diffusion \\citep{DV03}, and thermohaline mixing \\citep{CZ07}. \\citet{P06} demonstrated that meridional circulation, differential rotation, and shear turbulence do not create enough mixing to account for the observed variations in surface abundances, implying that additional hydrodynamical processes must be acting. The large study of surface abundances in field giants published by \\citet{GSCB00} is a key to differentiating between models of deep mixing, since it demonstrates clearly the progressive depletion of carbon on the giant branch, as well as the sharp drop in $^{12}$C/$^{13}$C and $\\log \\epsilon$(Li) that happens at the RGB bump.  The fundamental result from \\citet{GSCB00} that all current deep mixing models must reproduce is that deep mixing is universal among post-bump red giants. \\citet{CBW98} consider mixing in terms of the ``critical $\\mu$ gradient,'' the largest gradient in mean molecular weight that still permits deep mixing. In an observational study of seven mildly metal-poor red giants in the region of the RGB bump, they find that the critical $\\mu$ gradient is independent of composition or mass.  \\citet{DV03} use the formalism of diffusion, with mixing depth and a diffusion constant as the important parameters, to model deep mixing. They find that the mixing depth does not depend strongly on metallicity, which implies that all red giants with a mass less than $\\simeq 2$M$_{\\odot}$ will experience deep mixing. Their figures also show that the evolution of surface abundances of carbon and nitrogen are not particularly affected by metallicity, though reductions in $\\log \\epsilon$(Li) and $^{12}$C/$^{13}$C are more sensitive. \\citet{EDL08} show that the reaction $^{3}$He($^{3}$He, 2p)$^{4}$He causes a $\\mu$-inversion in the outer edge of the hydrogen-burning shell, and claim that the resulting Rayleigh-Taylor instability is important in driving deep mixing. \\citet{CZ07} argue that the more complex process of thermohaline convection will act in that $\\mu$-inversion region. They use the \\citet{U72} prescription to parametrize the thermohaline mixing as a diffusion process. In contrast to \\citet{DV03}, they find that the evolution of the surface abundances of carbon, nitrogen, and lithium are all affected by overall stellar metallicity, while the $^{12}$C/$^{13}$C ratio approaches its equilibrium value very quickly at all metallicities.  Although questions of deep mixing rate (e.g., Smith \\& Martell 2003\\nocite{SM03}) and depth (e.g., Charbonnel et al. 1998\\nocite{CBW98}) have been studied observationally by many authors, the results available in the literature can be difficult to synthesize into a single conclusion. Many studies focus on one or two particular clusters (e.g., Da Costa \\& Cottrell 1980\\nocite{DC80}, Suntzeff 1981\\nocite{S81}, Trefzger et al. 1983\\nocite{T83}, Lee 1999\\nocite{L99}), or attempt to correlate deep mixing with other cluster properties such as horizontal branch morphology \\citep{CN00}, stellar rotational velocity \\citep{CPT05} or cluster ellipticity \\citep{N87}. Individual authors and collaborations develop their own analysis tools, and the differences between spectral index definitions, model atmospheres, spectral synthesis engines, and abundance determination methods produce significant systematic differences in different authors' abundance scales, as is clear from literature-compilation studies such as \\citet{S02}.  One can construct a phenomenological picture of deep mixing from this heterogeneous information, and it goes roughly as follows: all stars with mass less than $\\simeq 2$M$_{\\odot}$ will at some point have their hydrogen-burning shell cross the $\\mu$-barrier. The $\\mu$-barrier is larger than the critical $\\mu$-gradient for deep mixing, so its destruction permits deep mixing to begin. Deep mixing occurs continuously, and involves all material outside the radius where the $\\mu$-gradient within the outer H-burning shell is critical. The onset of deep mixing happens lower on the giant branch for higher-metallicity clusters, because their $\\mu$-barrier is at smaller radius. However, in higher-metallicity stars the hydrogen-burning shell is more compact \\citep{SM79}, so that the radius where the $\\mu$-gradient is critical is relatively further out in the hydrogen-burning shell. This means that the material mixed to the surface in higher-metallicity stars is less processed than in low-metallicity stars. Various authors (e.g., Charbonnel et al. 1998\\nocite{CBW98}, Cassisi et al. 2002\\nocite{C02}) use this relation between metallicity and mixing efficiency to study the structure of the hydrogen-burning shell.  Our goal in this project is to determine the relative efficiency of deep mixing across a broad range of metallicity by measuring present-day carbon abundances and depletion rates from a homogeneous set of globular cluster red giants in similar evolutionary phases. An earlier example of this approach is the study of \\citet{BD80}, who found that [C/Fe] on the upper RGB of M3, M13 and NGC 6752 correlated with [Fe/H] metallicity. ",
        "conclusions": "In summary, the results of this work show that within the [Fe/H] range $-2.2$ to $-1.0$ dex the rate of deep mixing varies with metallicity. We measure CH bandstrength using the index $S_{2}(CH)$, which was designed to be valid across a broad range in [Fe/H] \\citep{MSB08b}. Carbon abundances are determined by matching CH bandstrengths measured from the data to bandstrengths measured from specifically-designed grids of SSG synthetic spectra. Under the assumption \\citep{CBW98} that deep mixing begins at the RGB luminosity function bump, we establish the carbon depletion rate for a given star as the change in its [C/Fe] from an initial solar value divided by the time elapsed since the onset of mixing. Since the RGB luminosity function bump is fainter in high-metallicity globular clusters than in lower-metallicity globular clusters, and higher-metallicity red giants evolve more slowly than lower metallicity red giants, higher-metallicity stars at a given absolute magnitude qualitatively ought to have spent more time mixing than their lower-metallicity counterparts. In order to quantify the time spent mixing as a function of metallicity and absolute V magnitude, we interpolate the \\citet{FP90} observational data on $M_{V}^{bump}$ as a function of metallicity to the metallicities of the clusters in our study, then use metallicity-appropriate Yale-Yonsei \\citep{D04} isochrones to convert $\\Delta M_{V}^{bump}$ for each individual star into $\\Delta t^{bump}$. As can be seen in Figure \\ref{fig:LBf10}, the lower-metallicity red giants do evolve more quickly, and therefore spend less time mixing than the higher-metallicity red giants in this study. However, dividing [C/Fe] by $\\Delta t^{bump}$ for each individual star produces the result that the deep mixing rate is roughly twice as large at [Fe/H]$=-2.0$ than at [Fe/H]$=-1.0$, as can be seen in Figure \\ref{fig:LBf11}.  This present analysis includes some assumptions worth noting, since they may need to be more carefully examined if this result is to provide constraints for theoretical models of deep mixing. To express $X_{C, env}$ as a straightforward exponential in time, we hold the mass of the stellar envelope constant, though it shrinks continuously as the hydrogen-burning shell proceeds outward \\citep{DV03}. We also hold the mixing rate $\\dot{M}$ constant, though it is controlled by the structure of the hydrogen-burning shell, and may well evolve. In addition, we implicitly assume in constructing Figure \\ref{fig:LBf11} that all stars have equal (and solar) initial abundances of carbon. A primordial depletion of $\\sim 0.3$ dex in a subset of our sample would thus be misinterpreted as an artificially high mixing rate in those stars. A study of carbon abundances in fainter post-bump red giants, or pre-bump giants, in the globular clusters included in this study, would allow for direct calculation of $\\Delta$[C/Fe]$/\\Delta t$ without assumptions about the initial carbon abundance. In addition, a study of carbon depletion rates in high-metallicity red giants, either in old open clusters or in high-metallicity globular clusters, would allow an estimation of the maximum metallicity at which deep mixing still operates. That maximum metallicity would be helpful for constraining the various models of deep mixing. Nevertheless, our fundamental result, that carbon depletion proceeds more quickly in low-metallicity globular cluster red giants than in their high-metallicity counterparts, is robust, and provides a clear affirmation of present theoretical models.  \\clearpage \\begin{deluxetable}{ l r r c c r c } \\tablewidth{0pt} \\tablehead{\\colhead{Cluster ID} & \\colhead{RA (J2000)} & \\colhead{$\\delta$ (J2000)} & \\colhead{$(m-M)_{V}$} & \\colhead{$E(B-V)$} & \\colhead{[Fe/H]} & \\colhead{$\\rm{N_{obs}}$}} \\startdata NGC 4147&12 10 06.2&+18 32 31&16.48&0.02&-1.83&4\\\\ NGC 5727 (M3)&13 42 11.2&-28 22 32&15.08&0.01&-1.39&3\\\\ NGC 5904 (M5)&15 18 33.8&+02 04 58&14.46&0.03&-1.29&3\\\\ NGC 6205 (M13)&16 41 41.5&+36 27 37&14.45&0.02&-1.54&5\\\\ NGC 6254 (M10)&16 57 08.9&-04 05 58&14.08&0.28&-1.52&8\\\\ NGC 6341 (M92)&17 17 07.3&+43 08 11&14.64&0.02&-2.29&2\\\\ NGC 6535&18 03 50.7&-00 17 49&15.22&0.34&-1.80&2\\\\ NGC 6712&18 53 04.3&-08 42 22&15.60&0.45&-1.01&3\\\\ NGC 6779 (M56)&19 16 35.5&+30 11 05&15.65&0.20&-1.94&5\\\\ NGC 7078 (M15)&21 29 58.3&+12 10 01&15.23&0.10&-2.25&1\\\\ NGC 7089 (M2)&21 33 29.3&-00 49 23&15.49&0.06&-1.62&6\\\\ \\enddata \\end{deluxetable}  \\begin{deluxetable}{ l l l r r r r r r } \\tablewidth{0pt} \\tablehead{\\colhead{Cluster ID} & \\colhead{Star ID} & \\colhead{Date obs.} & \\colhead{$M_{V}$} & \\colhead{[Fe/H]} & \\colhead{$S_{2}(CH)$} & \\colhead{$\\sigma_{S}$} & \\colhead{[C/Fe]} & \\colhead{$\\sigma_{C}$}} \\startdata M10&II-105&2006-06-03&-1.23&-1.52&1.753&0.0114&-0.625&0.0549\\\\ M10&III-73&2006-06-03&-1.22&-1.52&1.784&0.0011&-0.507&0.0327\\\\ M10&III-85&2006-06-03&-1.38&-1.52&1.784&0.5153&-0.544&0.0325\\\\ M10&IV-30&2006-06-03&-1.28&-1.52&1.803&0.0052&-0.456&0.0381\\\\ M10&I-15&2004-07-12&-1.22&-1.52&1.799&0.0047&-0.461&0.0365\\\\ M10&III-93&2004-07-12&-1.18&-1.52&1.802&0.0047&-0.448&0.0370\\\\ M10&III-97&2004-07-12&-1.41&-1.52&1.724&0.0047&-0.764&0.0375\\\\ M10&I-12&2005-07-12&-1.14&-1.52&1.819&0.0102&0.375&0.0513\\\\ M13&IV-53&2005-04-17&-1.77&-1.54&1.747&0.0019&-0.729&0.0336\\\\ M13&II-33&2006-06-01&-1.78&-1.54&1.764&0.0021&-0.647&0.0337\\\\ M13&III-52&2006-06-01&-1.78&-1.54&1.735&0.2635&-0.763&0.0325\\\\ M13&II-57&2006-06-02&-1.75&-1.54&1.703&0.0023&-0.910&0.0342\\\\ M13&III-18&2006-06-02&-1.68&-1.54&1.698&0.0020&-0.914&0.0338\\\\ M15&K77&2005-09-06&-1.53&-2.25&1.584&0.0016&-0.718&0.0386\\\\ M2&I-103&2005-07-13&-1.92&-1.62&1.775&0.0071&-0.582&0.0447\\\\ M2&I-104&2005-07-13&-1.52&-1.62&1.792&0.0042&-0.486&0.0359\\\\ M2&I-298&2005-09-07&-1.46&-1.62&1.755&0.0037&-0.601&0.0351\\\\ M2&I-190&2005-09-08&-1.75&-1.62&1.698&0.0024&-0.888&0.0340\\\\ M2&II-60&2005-09-09&-1.69&-1.62&1.772&0.0034&-0.570&0.0348\\\\ M2&II-71&2005-09-09&-1.64&-1.62&1.771&0.0042&-0.570&0.0361\\\\ M3&BC&2006-05-31&-1.24&-1.39&1.825&0.0010&-0.423&0.0333\\\\ M3&I-46&2006-05-31&-1.26&-1.39&1.809&0.0050&-0.488&0.0389\\\\ M3&V-80&2006-06-02&-1.61&-1.39&1.763&0.0031&-0.702&0.0361\\\\ M5&I-39&2006-06-01&-1.39&-1.29&1.838&0.0041&-0.342&0.0410\\\\ M5&IV-34&2006-06-02&-1.41&-1.29&1.791&0.0045&-0.584&0.0418\\\\ M5&IV-49&2006-06-03&-1.32&-1.29&1.713&0.0031&-0.917&0.0384\\\\ M56&I-10&2006-08-30&-1.53&-1.94&1.682&0.0022&-0.610&0.0304\\\\ M56&E-48&2006-08-31&-1.95&-1.94&1.694&0.0031&-0.678&0.0315\\\\ M56&I-141&2005-09-06&-1.22&-1.94&1.733&0.0025&-0.343&0.0301\\\\ M56&E-22&2005-09-07&-1.78&-1.94&1.701&0.0035&-0.597&0.0318\\\\ M56&I-66&2005-09-07&-1.78&-1.94&1.653&0.0015&-0.798&0.0297\\\\ M92&II-70&2006-05-31&-1.40&-2.29&1.621&0.0036&-0.398&0.0420\\\\ M92&IV-94&2006-05-31&-1.44&-2.29&1.549&0.0047&-0.823&0.0563\\\\ NGC 4147&II-14&2005-02-01&-1.150&-1.83&1.697&0.0123&-0.580&0.0539\\\\ NGC 4147&II-30&2005-02-01&-1.760&-1.83&1.754&0.0027&-0.519&0.0313\\\\ NGC 4147&II-45&2005-02-02&-1.960&-1.83&1.668&0.0093&-0.913&0.0516\\\\ NGC 4147&IV-13&2005-02-02&-1.390&-1.83&1.713&0.0066&-0.585&0.0381\\\\ NGC 6535&13&2004-07-13&-1.71&-1.80&1.738&0.2743&-0.583&0.0308\\\\ NGC 6535&19&2004-07-13&-0.990&-1.80&1.781&0.0068&-0.288&0.0379\\\\ NGC 6712&KC564&2005-07-13&-1.430&-1.01&1.850&0.0147&-0.335&0.1187\\\\ NGC 6712&LM11&2005-07-13&-1.200&-1.01&1.838&0.0139&-0.397&0.0974\\\\ NGC 6712&B66&2006-08-31&-1.57&-1.01&1.808&0.0110&-0.676&0.1069\\\\ \\enddata \\end{deluxetable}  \\clearpage \\begin{deluxetable}{l c c c l} \\tablewidth{0pt} \\tablehead{\\colhead{Index} & \\colhead{Blue comparison band (\\mbox{\\AA})} & \\colhead{Science band (\\mbox{\\AA})} & \\colhead{Red comparison band (\\mbox{\\AA})} & \\colhead{Reference}} \\startdata $s_{CH}$&4220-4280&4280-4320&-&\\citet{BS93}\\\\ $m_{CH}$&4080-4130&4270-4320&4420-4470&\\citet{S96}\\\\ $CH(G)$&4230-4260&4270-4320&4390-4420&\\citet{L99}\\\\ $S(CH)$&4050-4100&4280-4320&4330-4350&\\citet{MSB08}\\\\ $S(4243)$&-&4290-4318&4314-4322&\\citet{B90}\\\\ $S_{2}(CH)$&4212-4242&4297-4317&4330-4375&\\citet{MSB08b}\\\\ \\enddata \\end{deluxetable}"
    },
    "0809/0809.0631.txt": {
        "abstract": "We present an up-to-date review of Big Bang Nucleosynthesis (BBN). We discuss the main improvements which have been achieved in the past two decades on the overall theoretical framework, summarize the impact of new experimental results on nuclear reaction rates, and critically re-examine the astrophysical determinations of light nuclei abundances. We report then on how BBN can be used as a powerful test of new physics, constraining a wide range of ideas and theoretical models of fundamental interactions beyond the standard model of strong and electroweak forces and Einstein's general relativity. ",
        "introduction": "\\label{sec:introduction} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% A remarkable scientific achievement in the second half of the 20$^{\\rm th}$ century has been the establishment of ``Standard Models'' of Particle Physics (SMPP) and Cosmology (SMC). In particular,  the latter has been possible thanks to an incredibly fast growth of  the amount and quality of observations over the last couple of decades. The picture revealed is at the same time beautifully simple and intriguingly mysterious: on one hand, known gauge interactions and Einstein's general relativity seem able to explain a huge wealth of information in terms of a few free parameters specifying the composition/initial conditions of the Universe; on the other hand, these numbers are not explained in terms of dynamical processes involving the known fields and interactions. This is the case of the ``dark energy'' density (consistent with a cosmological constant), of the non-baryonic dark matter, of the baryon-antibaryon asymmetry, the flatness, homogeneity and isotropy of the universe on large scales, etc.  The very success of the cosmological laboratory is thus providing much indirect evidence for physics beyond the SMPP. On the other hand, advances in particle physics (a very recent example being the phenomenology of massive neutrinos) have an impact at cosmological level. This interplay has proven extremely fertile ground for the development of `astroparticle physics', especially since many theories beyond the  SMPP predict new phenomena far beyond the reach of terrestrial laboratories, but potentially testable in astrophysical and cosmological environments. In this respect, the nucleosynthesis taking place in the primordial plasma plays a twofold role: it is undoubtedly one of the observational pillars of the hot Big Bang model, being indeed known simply as ``Big Bang Nucleosynthesis'' (BBN); at the same time, it provides one of the earliest direct cosmological probes nowadays available, constraining the properties of the universe when it was a few seconds old, or equivalently at the MeV temperature scale. Additionally, it is special in that all known interactions play an important role: gravity sets the dynamics of the ``expanding cauldron'', weak interactions determine the neutrino decoupling and the neutron-proton equilibrium freeze-out, electromagnetic and nuclear processes regulate the nuclear reaction network.  The basic framework of the BBN emerged in the decade between  the seminal Alpher-Bethe-Gamow (known as $\\alpha\\beta\\gamma$) paper in 1948 \\citep{Alp48} and the essential settlement of the paradigm of the stellar nucleosynthesis of elements heavier than $^7$Li with the B$^2$FH paper \\citep{Bur57}. This pioneering period---an account of which can be found in \\citep{Kra96}---established the basic picture that sees the four light-elements $^2$H, $^3$He, $^4$He and $^7$Li as products of the early fireball, and virtually all the rest produced in stars or as a consequence of stellar explosions.  In the following decades, the emphasis on the role played by the BBN has evolved significantly.  In the simplest scenario, the only free parameters in primordial nucleosynthesis are the baryon to photon ratio $\\eta$  (equivalently, the baryon density of the universe) and the neutrino asymmetry parameters, $\\eta_{\\nu_\\alpha}$ (see Section~\\ref{nuasymm}). However, only neutrino asymmetries larger than $\\eta$ by many orders of magnitude have appreciable effects. This is why the simple case where all $\\eta_{\\nu_\\alpha}$'s  are assumed to be negligibly small (e.g., of the same order of $\\eta$) is typically denoted as Standard BBN (SBBN). Since several species of `nuclear ashes' form during BBN, SBBN is an over-constrained theory whose self-consistency can be checked comparing predictions with two or more light nuclide determinations. The agreement of predicted abundances of the light elements  with their measured abundances (spanning more than nine orders of magnitude!) confirmed the credibility of BBN as cosmological probe. At the same time,  the relatively narrow range of $\\eta$ where a consistent picture emerged was the first compelling argument in favor of the non-baryonic nature of the ``dark matter'' invoked for astrophysical dynamics.  The past decade, when for the first time a redundancy of determinations of $\\eta$ has been possible, has stressed BBN as a consistency tool for the SMC. Beside BBN, one can infer the density of baryons from the Lyman-$\\alpha$ opacity in quasar spectra due to  intervening high redshift hydrogen clouds \\citep{Mei93,Rau96,Wei97}; from the baryon fraction in clusters of galaxies, deduced from the hot x-ray emission \\citep{Evr97}; most importantly, from the height of the Doppler peak in the angular power spectrum of the cosmic microwave background anisotropy (see \\citep{Dun08} for the latest WMAP results). These determinations are not only mutually consistent with each other, but the two most accurate ones (from the CMB and BBN) agree within $5\\div$10\\%. While losing to CMB the role of ``barometer of excellence'', BBN made possible a remarkable test of consistency of the whole SMC.  It comes without surprise that this peculiar `natural laboratory' has inspired many investigations, as testified by the numerous reviews existing on the subject, see e.g. \\citep{Mal93,Cop95,Oli00,Sar96,Sch98,Ste07}. Why then a new review? In the opinion of the authors, a  new BBN review seems worthy because at present, given the robustness of the cosmological scenario, the attention of the community is moving towards a new approach to the BBN. On one hand, one uses it as a precision tool in combination with other cosmological information to reduce the number of free parameters to extract  from multi-parameter fits. On the other hand, BBN is an excellent probe to explore the very early universe, constraining scenarios beyond the SMPP. The latter motivation is particularly intriguing given the perspectives of the forthcoming LHC age to shed light on the TeV scale. A new synergy with the Lab is expected to emerge in the coming years, continuing a long tradition in this sense. Finally, a wealth of data from nuclear astrophysics and neutrino physics have had a significant impact on BBN, and it is meaningful to review and assess it. In particular, the  recent advances in the neutrino sector  have made obsolete many exotic scenarios popular in the literature still a decade ago and improved numerous constraints, providing a clear example of the synergy we look forward to in the near future.\\\\ This review is structured as follows: in Section \\ref{sec:cosm_overview} we summarize the main cosmological notions as well as most of the symbols used in the rest of the article. Section \\ref{sec:bbn_overview} is devoted to the description of the Standard BBN scenario. Section \\ref{sec:obsabund} treats the status of observations of light nuclei abundances, which in Section \\ref{sec:BBNanalysis} are compared with theoretical predictions. The following Sections deal with exotic scenarios: Section \\ref{sec:BBN_nuphys} with neutrino properties, Section \\ref{sec:IBBN} with inhomogeneous models, Section \\ref{sec:constr_fundam} with constraints to fundamental interactions and Section \\ref{s:massive-particle} with massive particles. In Section \\ref{sec:conclusions} we report our conclusions. Although this article is a review, many analyses have been implemented ex-novo and some original results are presented here for the first time. Due to the large existing literature and to space limitation, we adopt the criterion to be as complete as possible in the post-2000 literature, while referring to previous literature only when pertinent to the discussion or when still providing the most updated result. Also, we adopt a more pedagogical attitude in introducing arguments that have rarely or never entered previous BBN reviews, as for example extra dimensions or variation of fundamental constants in Section \\ref{sec:constr_fundam}, while focusing mainly on new results (as opposed to a `theory review') in subjects that have been extensively treated in the past BBN literature (as in SUSY models leading to cascade nucleosynthesis, the gravitino `problem', etc.). Other topics, which for observational or theoretical reasons have attracted far less interest in the past decade in relation to BBN bounds, are only briefly mentioned or omitted completely (this is the case of technicolor or cosmic strings). Older literature containing a more extensive treatment of these topics can be typically retraced from the quoted reviews. In the following, unless otherwise stated, we  use natural units $\\hbar=c=k_B=1$, although conventional units in the astronomical literature (as parsec and multiples of it) are occasionally used where convenient for the context.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In this review we have reported the current status of BBN, focusing in the first part on precision calculations possible in the standard scenario, which provide a tool for current cosmological framework, in the second part on the constraints to new physics, which become particularly important in the forthcoming LHC era. The ``classical parameter\" constrained by BBN is the baryon to photon ratio, $\\eta$, or equivalently the baryon abundance, $\\Omega_B h^2$. At present, the constraint is dominated by the deuterium determination, and we find $\\Omega_B h^2=0.021\\pm 0.001$(1 $\\sigma$). This determination is consistent with the upper limit on primordial $^3$He/H (which provides a lower limit to $\\eta$), as well as with the range selected by $^4$He determinations, which however provides a constraint almost one order of magnitude weaker. The agreement within 2 $\\sigma$ with the WMAP determination, $\\Omega_B h^2=0.02273\\pm 0.00062$, represents a remarkable success of the Standard Cosmological Model. On the other hand, using this value as an input, a factor $\\gsim 3$ discrepancy remains with $^7$Li determinations, which can hardly be reconciled even accounting for a conservative error budget in both observations and nuclear inputs. Even more puzzling are some detections of traces of $^6$Li at a level far above the one expected from the Standard BBN. If the observational determinations are solid, both nuclides indicate that either their present observations do not reflect their primordial values, and should thus be discarded for cosmological purposes, or that the early cosmology is more complicated and exciting than the Standard BBN lore. Neither a non-standard number of massless degrees of freedom in the plasma (parameterized via $\\neff$) or a lepton asymmetry $\\xi_e$ (all asymmetries assumed equal) can reconcile the discrepancy. Current bounds on both quantities come basically from the $^4$He measurement, $\\neff=3.2\\pm 0.4\\,(1\\,\\sigma)$ and $\\xi_e=-0.008\\pm 0.013\\,(1\\,\\sigma)$.  On the other hand, other exotic proposals have been invoked to reconcile this discrepancy. Typically they involve massive meta-stable particles with weak scale interactions, which should be soon produced at the LHC. In Supersymmetric scenarios, long-lived particles are possible whenever the Next to Lightest Supersymmetric Particle (NLSP) decays into the Lightest Supersymmetric Particle (LSP) are gravity-mediated, or ``disfavored'' by phase space arguments, with a modest mass splitting between NLSP and LSP. Cases frequently considered in the recent literature are neutralino $\\to$ gravitino decays, for example, or stau $\\to$ gravitino. The phenomenology associated with the catalysis of reactions due to bound states of charged particles (as the stau) with ordinary nuclei is a particularly new topic in recent investigations. Also, the importance of a possible primordial origin of the $\\lisix$ measured in a few systems of the $\\li7$ plateau has been recognized: first,  the bounds in parameter space tighten significantly if lithium constraints are used, especially $\\lisix$ \\citep{Hol96,Jed00,Hol99,Jed04a}; second, because these exotic BBN scenarios may accommodate for a cosmological origin for $\\lisix$ while solving the $\\li7$ excess problem as well, their phenomenology is very appealing. Although these links among primordial nucleosynthesis, dark matter, and perhaps SUSY phenomenology are quite fascinating, it is worth stressing that BBN bounds on cascade decays or annihilations of massive particles apply well beyond the restricted  class of SUSY-inspired models. For example, the electromagnetic cascades following heavy sterile neutrino decays are constrained by these kinds of  arguments, as well as decays of massive pseudo Nambu-Goldstone bosons, as considered in \\citep{Mas97,Mas04}.  There are two directions along which we can expect the BBN field to develop in the future. On one hand, BBN is an important tool for precision cosmology, especially if its priors are used in combination with other cosmological observables. Already BBN provides the best bounds on parameters as $\\neff$ and $\\xi_e$ (and bounds on $\\eta$ comparable to the CMB); yet, since theoretical uncertainties are at the moment well below observational ones, there is surely room to refine its power, provided that significantly greater efforts are devoted to determine light element abundances, and in particular $Y_p$. It is instructive in this sense to look back to what S. Sarkar wrote in his review \\citep{Sar96} thirteen years ago:  {\\it Thousands of person years of effort have been invested in obtaining the precise parameters of the $Z^0$ resonance in $e^{+}-e^{-}$ collisions, which measures the number of light neutrino species (and other particles) which couple to the $Z^0$. In comparison, a modest amount of work has been done, by a few small teams, on measuring the primordial light element abundances, which provide a complementary check of this number as well as a probe of new superweakly interacting particles which do not couple to the $Z^0$.}  Despite the improvements reported in this article, we feel that unfortunately insufficient attention has been devoted to this problem, if compared to other areas of observational cosmology. In particular, the $\\He4$ determination is still plagued by systematic uncertainties. Although their importance has been recently recognized and assessed more carefully, the fact that this reanalysis was triggered after the independent determination of $\\eta$ from CMB (and its agreement with the ``low deuterium determinations\" in QSO spectra) shows that there is still a long way to go towards a precision era for primordial elements. On the other hand, a significant improvement has taken place in assessing and reducing theoretical uncertainties, mostly related to nuclear reaction data. BBN has benefit from a wealth of new nuclear astrophysics measurements at low energies and covering large dynamical ranges. Given the much larger observational uncertainties, in this sector an effort in reassessing the systematic errors in older datasets might be more useful in reducing remaining discrepancies in the nuclear rates error budget. This is in particular the case for reactions involving $^7$Be.  The other direction of development follows from the interplay with Lab experiments. Neutrinos have reserved many surprises, and it is not excluded that exotic properties may show up in future experiments with important implications for BBN, as we illustrated in Section \\ref{sec:BBN_nuphys}. However, it is in particular from LHC that one expects a better understanding of high energy scales, and thus of the cosmology at earlier times and higher temperatures. Most theories that go beyond the Standard Model of Particle Physics require new states to appear at or above the electroweak scale and, as already reported, they might have implications for the phenomenology at the BBN epoch. If the LHC should provide indication for the existence of the SMPP Higgs and nothing else, there will be no natural scale to explore. In this case, albeit sad, BBN and other cosmological tools might be the only practical means to explore very high energy phenomena leaving their imprint on the cosmos. One example treated here is the effect of variations of fundamental ``constants\"  on cosmological time-scales that emerge in extra dimensional scenarios, possibly embedded in grand unified theories or string theories. If, as hopefully more likely, the LHC will reveal  new dynamics above the electroweak scale, we might be able to infer from the empirical evidence the presence of cosmological effects before the BBN epoch. A new Standard Cosmological Model would emerge as well, perhaps making the BBN one more step in the ladder back to the Big Bang, rather than the first one.  \\ack We would like to thank C. Abia, A.D. Dolgov, J. Lesgourgues, S. Pastor and G.G. Raffelt for valuable comments and suggestions, and G.L. Fogli for having particularly encouraged this work. We also thank  M. Kamimura, and especially K. Jedamzik, for suggestions and clarifying remarks which much improved the manuscript, and K. Jedamzik for providing also the updated version of some Figures. F. Iocco is supported by MIUR through grant PRIN-2006, and acknowledges hospitality at Fermilab during some stage of this work. G. Miele acknowledges supports by Generalitat Valenciana (Grant No.\\ AINV/2007/080) and by the Spanish MICINN (grants SAB2006-0171 and FPA2005-01269). G. Mangano, G. Miele, and O. Pisanti acknowledge supports by INFN - I.S. FA51 and by PRIN 2006 ``Fisica Astroparticellare: Neutrini ed Universo Primordiale\" of Italian MIUR. P.D. Serpico is supported by the US Department of Energy and by NASA grant NAG5-10842. Fermilab is operated by Fermi Research Alliance, LLC under Contract No.~DE-AC02-07CH11359 with the United States Department of Energy."
    },
    "0809/0809.0315_arXiv.txt": {
        "abstract": "Star forming regions are expected to show linear proper motions due to the relative motion of the Sun with respect to the region. These proper motions appear superposed to the proper motions expected in features associated with mass ejection from the young stellar objects embedded in them. Therefore, it is necessary to have a good knowledge of the proper motions of the region as a whole in order to correctly interpret the motions associated with mass ejection. In this paper we present the first direct measurement of proper motions of the NGC~1333 star forming region. This region harbors one of the most studied Herbig-Haro systems, HH 7-11, whose exciting source remains unclear. Using VLA A configuration data at 3.6 cm taken over 10 years, we have been able to measure the absolute proper motions of four thermal sources embedded in NGC~1333. From our results we have derived the mean proper motions of the NGC~1333 star forming region to be $\\mu_{\\alpha}\\cos\\delta$ = 9 $\\pm$ 1 mas yr$^{-1}$ and $\\mu_{\\delta}$ = $-$10 $\\pm$ 2 mas~yr$^{-1}$. In this paper, we also discuss the possible implications of our results in the identification of the outflow exciting sources. ",
        "introduction": "The stars in nearby regions of star formation are expected to show linear proper motions in the order of milliarcseconds (mas) per year. These proper motions are the result of the relative motion between the Sun and the stars studied, and can be obtained, to first order, by comparing the observed position of the star with respect to the reference frame of the remote quasars at different epochs.  In the last few years, Very Large Array (VLA) observations taken with time baselines of up to two decades have been used to measure these proper motions for regions like Taurus (Loinard et al. 2003), L1527 (Loinard et al. 2002), L1551 (Rodr\\'{\\i}guez et al. 2003), Ophiuchus (Curiel et al. 2003), and Orion (G\\'omez et al. 2005).  Even when the VLA observations lack sufficient angular resolution to detect more subtle motions (i.e. geometric parallax), they can provide not only the average proper motion of the region but also reveal the presence of stars moving with peculiar velocities with respect to the region, such as those found in the Orion BN/KL region (Rodr\\'{i}guez et al. 2005, G\\'omez et al. 2005), that suggest a runaway nature for some of the sources. Much more accurate astrometry can be achieved using Very Long Baseline Interferometry (VLBI) techniques (e.g. Loinard et al. 2007), but these observations are restricted to very compact and bright non thermal radio stars and cannot be applied to thermal sources, where the emission is relatively extended and its brightness temperature is not expected to exceed $10^4$~K, far below the sensitivity of present day VLBI.  \\begin{deluxetable*}{cccccrc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Observations Parameters \\label{tabla1}} \\startdata \\hline \\hline &             &         &   Bootstrapped\t\t    &\t\t\t &\t\t\t\t    &\t                         \\\\ &\t      &  VLA    &   Flux Density of\t\t    & \\multicolumn{2}{c}{Synthesized Beam\\tablenotemark{b}} &\t                         \\\\ \\cline{4-5} & Observation & Project & Phase Calibrator\\tablenotemark{a} &\t      HPBW\t &   \\multicolumn{1}{c}{P.A.}\t    & rms Noise\\tablenotemark{b} \\\\ Epoch  & Date        &  Code   & (Jy)  \t\t\t    &  (arcsec) \t & \\multicolumn{1}{c}{(deg)}\t    & ($\\mu$Jy beam$^{-1}$)\t \\\\ \\hline 1989.1  & 89-Jan-14   & AR202   &      1.343 $\\pm$ 0.003\t    & 0.29 $\\times$ 0.24 &         62                       &    28                      \\\\ 1994.3\t& 94-Apr-23   & AR277   &      1.686 $\\pm$ 0.006\t    & 0.28 $\\times$ 0.20 &      $-$68                       &    27                      \\\\ 1996.9\t& 96-Dec-22   & AR277   &      1.305 $\\pm$ 0.009\t    & 0.35 $\\times$ 0.30 &      $-$88                       &    24                      \\\\ 1998.2\t& 98-Mar-27   & AA218   &      1.476 $\\pm$ 0.009\t    & 0.26 $\\times$ 0.24 &       $-$5                       &    18                      \\\\ 1998.4\t& 98-May-26   & AA218   &      1.610 $\\pm$ 0.010\t    & 0.26 $\\times$ 0.24 &       $-$7                       &    18                      \\\\ 1999.5  & 99-Jul-03   & AA239   &      1.681 $\\pm$ 0.008\t    & 0.29 $\\times$ 0.25 &         39                       &    17                      \\\\ \\hline \\enddata  \\tablenotetext{a}{The phase calibrator used in all the epochs was 0333+321.}  \\tablenotetext{b}{From naturally weighted maps.}  \\end{deluxetable*}  \\begin{deluxetable*}{crrcrr} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Positions of the Sources\\tablenotemark{a} \\label{tabla2}} \\startdata \\hline \\hline &                       \\multicolumn{2}{c}{VLA~2}                           & &                   \\multicolumn{2}{c}{VLA~3}                               \\\\ \\cline{2-3} \\cline{5-6} Epoch     & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} & & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} \\\\ \\hline 1989.1    & 03 29 01.9535 $\\pm$ 0.0003      & 31 15 38.307 $\\pm$ 0.005  \t   & & 03 29 03.369 $\\pm$ 0.002 \t & 31 16 01.89 $\\pm$ 0.02\t    \\\\ 1994.3    &       01.9586 $\\pm$ 0.0004      &       38.221 $\\pm$ 0.004  \t   & &       03.375 $\\pm$ 0.002 \t &\t 01.82 $\\pm$ 0.02\t    \\\\ 1996.9    &       01.9605 $\\pm$ 0.0004      &       38.285 $\\pm$ 0.008  \t   & &       03.376 $\\pm$ 0.001 \t &\t 01.75 $\\pm$ 0.02\t    \\\\ 1998.2    &       01.9606 $\\pm$ 0.0002      &       38.188 $\\pm$ 0.005  \t   & &       03.376 $\\pm$ 0.001 \t &\t 01.72 $\\pm$ 0.02\t    \\\\ 1998.4    &       01.9604 $\\pm$ 0.0002      &       38.211 $\\pm$ 0.006  \t   & &       03.373 $\\pm$ 0.002 \t &\t 01.78 $\\pm$ 0.02\t    \\\\ 1999.5    &       01.9622 $\\pm$ 0.0002      &       38.183 $\\pm$ 0.003  \t   & &       03.373 $\\pm$ 0.001 \t &\t 01.78 $\\pm$ 0.01\t    \\\\ \\hline \\\\ &                       \\multicolumn{2}{c}{VLA~4A}                          & &                       \\multicolumn{2}{c}{VLA~4B}                          \\\\ \\cline{2-3} \\cline{5-6} Epoch     & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} & & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} \\\\ \\hline 1989.1    & 03 29 03.730 $\\pm$ 0.002\t    & 31 16 04.02 $\\pm$ 0.02\t\t  & & 03 29 03.751 $\\pm$ 0.002         & 31 16 04.08 $\\pm$ 0.02     \\\\ 1994.3    &       03.730 $\\pm$ 0.002\t    &\t    03.95 $\\pm$ 0.02\t\t  & &\t    03.758 $\\pm$ 0.004         &       04.01 $\\pm$ 0.02     \\\\ 1996.9    &       03.731 $\\pm$ 0.002\t    &\t    03.94 $\\pm$ 0.02\t\t  & &\t    03.757 $\\pm$ 0.002         &       04.01 $\\pm$ 0.02     \\\\ 1998.2    & \\multicolumn{1}{c}{$\\cdots$\\tablenotemark{b}} & \\multicolumn{1}{c}{$\\cdots$\\tablenotemark{b}} & & \\multicolumn{1}{c}{$\\cdots$\\tablenotemark{b}} & \\multicolumn{1}{c}{$\\cdots$\\tablenotemark{b}} \\\\ 1998.4    &       03.735 $\\pm$ 0.002\t    &\t    04.00 $\\pm$ 0.03\t\t  & &\t    03.760 $\\pm$ 0.002         &       03.99 $\\pm$ 0.02     \\\\ 1999.5    &       03.734 $\\pm$ 0.002\t    &\t    03.96 $\\pm$ 0.02\t\t  & &\t    03.757 $\\pm$ 0.001         &       03.91 $\\pm$ 0.02     \\\\ \\hline \\enddata  \\tablenotetext{a}{Positions derived from Gaussian ellipsoid fits in the naturally weighted maps. Units of right ascension are hours, minutes and seconds, and units of declination are degrees, arcminutes, and arcseconds.}  \\tablenotetext{b}{The components of the binary system VLA~4 could not be resolved in this epoch.}  \\end{deluxetable*}  \\begin{figure*} \\epsscale{0.98} \\plotone{f1.eps}  \\caption{\\footnotesize{Naturally weighted VLA~3.6 cm continuum map of the thermal sources VLA~2, VLA~3, VLA~4A and VLA~4B at the 1998.4 epoch. Contours are $-3$, 3, 4, 5, 6, 8, 10, 12, 14, 16, 18, 20, 25 and 30 times the rms of the map, 18 $\\mu$Jy beam$^{-1}$. The synthesized beam, shown in the right hand panel, is 0$\\farcs$26 $\\times$ 0$\\farcs$24; P.A.=$-$7$^\\circ$}. The three panels are plotted at the same scale.}  \\label{fig1} \\end{figure*}    The NGC~1333 star forming region, located at a distance of 235 $\\pm$ 18 pc (Cernis 1990; Hirota et al. 2008) in the Perseus complex, harbors the  classical bright Herbig-Haro (HH) system HH~7-11, first reported by Herbig (1974) and by Strom et al. (1974). The optically visible star SVS~13, discovered as a near-IR source (Strom et al. 1976), is roughly aligned  with the chain of HH objects and was proposed as the powering source of  this HH system. This association was questioned by Rodr\\'{\\i}guez et al. (1997), who discovered a cm radio source (VLA 3) located $6''$ to the SW  of SVS~13, and argued that this new object is a better candidate to drive the HH outflow. SVS~13 is also associated with cm emission (source VLA~4  of Rodr\\'{\\i}guez et al. 1997) and mm emission (Looney et al. 2000). Bachiller et al. (2000), through interferometric observations, found that SVS~13 is associated with an ``extremely'' high velocity molecular outflow, although ``standard'' velocity gas was found in the vicinity of both SVS~13  and VLA~3. Subarcsecond VLA observations at 3.6 cm and 7 mm by Anglada et al. (2000, 2004) revealed that the radio source associated with SVS~13 is  actually a close binary with two components, VLA~4A and 4B, separated by  0$\\farcs$3 ($\\sim$65 AU). Interestingly, a detailed analysis of the positions and spectral energy distribution in the cm-mm range suggests that the two sources have very different properties: VLA~4B appears to be associated with the observed strong mm emission,  probably arising from a circumstellar dust disk, while VLA~4A is the  counterpart of the visible star SVS~13 with dust emission absent or much  less significant in this component.  Since high velocity water masers are indicators of outflow activity and can be observed with high angular resolution, Rodr\\'{\\i}guez et al. (2002) analyzed archive VLA data of the water masers associated with SVS~13 in order to investigate which of the two stars of the binary system was the most likely candidate to drive the outflow. It was found that the water masers appear segregated in two groups, according to their position and LSR velocity. A group with the LSR velocity similar to that of the ambient cloud is associated with VLA~4A, and a blueshifted velocity group is associated with VLA~4B. This result was interpreted as favoring VLA~4B as the outflow driving source. However, the study of Rodr\\'{\\i}guez et al. (2002) was restricted to the line-of-sight component of the velocity. Recently, Hirota et al. (2008) presented VLBI astrometric observations with VERA and derived absolute positions and proper motions of the order of 15-20 mas~yr$^{-1}$ (corresponding to velocities of 16-22 km~s$^{-1}$) of the water masers associated with SVS~13. These observations confirm the positional association with VLA~4A of the masers with an LSR velocity close to that of the ambient cloud. However, the proper motions (velocity and direction) relative to the sources could not be firmly established because of a poor knowledge of the absolute proper motions of the region. In fact, the optical data of Herbig \\& Jones (1983) suggest that the stars projected upon heavy obscuration (presumably associated with the NGC 1333 cloud) have an average proper motion of $\\sim$10 mas yr$^{-1}$ relative to the stars near the edge of the obscuration (presumably background stars). This suggests that the absolute proper motion of the region could have a significant contribution to the proper motions of the masers observed by Hirota et al. (2008).  In this paper we present an absolute astrometry analysis and proper motion calculations over 10 years of four thermal radio sources in the NGC~1333 region. There are no reports of absolute proper motions for embedded young stars in this region and our study provides the first determination of these motions. ",
        "conclusions": "We have presented the results of multi-epoch high angular resolution VLA observations at 3.6 cm of the NGC~1333 star forming region. From these observations we have measured the absolute positions at different epochs of four thermal sources (VLA~2, VLA~3, VLA~4A and VLA~4B) tracing YSOs embedded in this region. All these sources show an absolute proper motion to the SE, which we interpret as a proper motion of the NGC~1333 region as a whole. From our data we obtained the average values $\\mu_{\\alpha}\\cos\\delta$ = 9 $\\pm$ 1 mas~yr$^{-1}$ and $\\mu_{\\delta}$ = $-$10 $\\pm$ 2 mas~yr$^{-1}$, that we interpret as representative of the proper motion of the molecular cloud associated with the NGC~1333 region as a whole.  Our results allow us to correct the proper motion measurements of water maser spots associated with VLA~4A (SVS 13) obtained by Hirota et al. (2008) in order to obtain the proper motions relative to the NGC~1333 cloud. We conclude that feature 2 of Hirota et al. (2008) is practically stationary with respect to the cloud, and that the small residual proper motions of feature 1 indicate motions to the east, suggesting that none of these features is associated with the HH 7-11 system or the molecular outflow.  Our data also suggest that VLA~3 is a member of the NGC~1333 molecular cloud, since this source shows a proper motion similar to that of the others sources embedded in the region. Since the position of VLA~3 appears close to the base of the HH 7-11 system, we conclude that this proximity is physically real, and not a projection effect.  \\emph{Acknowledgements.} C.C.-G. acknowledges support from MEC (Spain) FPU fellowship. G.A., C.C.-G., M.O., and J.M.T. acknowledge support from MEC (Spain) grants AYA 2005-08523-C03 and AYA 2008-06189-C03 (including FEDER funds), and from Junta de Andaluc\\'{\\i}a (Spain). L.F.R. acknowledges the support of DGAPA, UNAM, and of CONACyT (M\\'exico). We thank Tomoya Hirota for useful comments on this paper."
    },
    "0809/0809.0912_arXiv.txt": {
        "abstract": "We report the discovery of five gravitationally lensed quasars from the Sloan Digital Sky Survey (SDSS). All five systems are selected as two-image lensed quasar candidates from a sample of high-redshift ($z>2.2$) SDSS quasars. We confirmed their lensing nature with additional imaging and spectroscopic observations. The new systems are SDSS~J0819+5356 (source redshift $z_s=2.237$, lens redshift $z_l=0.294$, and image separation $\\theta=4\\farcs04$), SDSS~J1254+2235 ($z_s=3.626$, $\\theta=1\\farcs56$), SDSS~J1258+1657 ($z_s=2.702$, $\\theta=1\\farcs28$), SDSS~J1339+1310 ($z_s=2.243$, $\\theta=1\\farcs69$), and SDSS~J1400+3134 ($z_s=3.317$, $\\theta=1\\farcs74$). We estimate the lens redshifts of the latter four systems to be $z_l=0.2-0.8$ from the colors and magnitudes of the lensing galaxies. We find that the image configurations of all systems are well reproduced by standard mass models. Although these lenses will not be included in our statistical sample of $z_s<2.2$ lenses, they expand the number of lensed quasars which can be used for high-redshift galaxy and quasar studies. ",
        "introduction": "Gravitationally lensed quasars are unique astronomical and cosmological tools, as described in the review of \\cite{kochanek06}. We can study the mass distributions of lensing objects from individual mass modeling, as well as the substructures in lensing objects \\citep[e.g.,][]{kochanek91,mao98}. We can also investigate their interstellar media from dust extinctions \\citep[e.g.,][]{falco99,munoz04} or absorption lines appearing in spectra of multiple quasar images \\citep[e.g.,][]{curran07}. The statistics of lensed quasars and the measurement of time delays between lensed images are useful tools to constrain cosmological parameters \\citep[e.g.,][]{refsdal64,turner90,fukugita90}. In addition, lensed quasars sometimes provide opportunities to study the central structures of quasar host galaxies in detail through microlensing events \\citep[e.g.,][]{richards04,poindexter08}.  Motivated by these ideas, astronomers have searched for lensed quasars using various methods and wavebands. Roughly 100 lensed quasars have been identified to date \\citep{kochanek06}. A number of homogeneously selected samples have been constructed \\citep[e.g.,][]{maoz93}, allowing statistical studies to be done. For example, the Cosmic Lens All Sky Survey \\citep[CLASS;][]{myers03,browne03} has created a sample of 22 lensed objects selected from $\\sim$16,000 radio sources. This sample has been used to obtain a variety of cosmological and astrophysical results \\citep[e.g.,][]{rusin01,mitchell05,chae06}.  The Sloan Digital Sky Survey \\citep[SDSS;][]{york00} has discovered $\\sim80,000$ spectroscopically identified quasars \\citep{schneider07}. We are conducting a survey of lensed quasars selected from the large dataset of the SDSS. The survey, the SDSS Quasar Lens Search \\citep[SQLS;][]{oguri06,oguri08a,inada08} has discovered more than 30 lensed quasars \\citep[e.g.,][and references therein]{kayo07,oguri08b}, making it the current largest lensed quasar survey. The SQLS also recovered nine previously known lensed quasars included in the SDSS footprint \\citep{walsh79,weymann80,surdej87,bade97,oscoz97,schechter98, myers99,morgan01,magain88}. The first statistical sample of 11 SQLS lenses \\citep{inada08} was constructed from the SDSS Data Release 3 quasar catalog \\citep[4188 deg${}^{2}$;][]{schneider05}, and used to constrain dark energy \\citep{oguri08a}.  The SQLS restricts the statistical lens sample to ${z_s}<2.2$ because we cannot make a well-defined quasar sample for homogeneous lens surveys at higher redshifts. The SDSS quasars at ${z_s}>2.2$ are required to be point sources \\citep[see][]{richards02}, and therefore they have a strong bias against the homogeneous lens candidate selection \\citep{oguri06,inada08}. However, the SQLS candidate finding algorithm can easily be extended to locate higher redshift lensed quasars \\citep{inada08}. Such high-redshift lensed quasars can be used as astronomical and cosmological tools to study (high-redshift) lensing galaxies \\citep[e.g.,][]{kochanek00} and constrain the Hubble constant \\citep[e.g.,][]{oguri07}. They are also useful for detailed studies of (lensed) high-redshift quasars. In this paper, we report the discoveries of five lensed quasars with high source redshifts (${z_s}=2.237$--$3.626$). They were selected as lensed quasar candidates from the SDSS data, and were confirmed as lenses with the observations at the University of Hawaii 2.2-meter telescope (UH88), the Astrophysical Research Consortium 3.5-meter telescope (ARC 3.5m), and the 3.58-meter Telescopio Nazionale Galileo (TNG 3.6m). All five candidates are confirmed to be double-image lensed quasars, with image separations of 1\\farcs28--4\\farcs04.  The structure of this paper is as follows. Brief descriptions of the SDSS data and our lens candidate selection algorithm are presented in \\S~\\ref{sec:sdss}. We present the results of imaging and spectroscopic observations to confirm the lensing hypotheses for the five objects and estimate the redshifts of the lensing galaxies in \\S~\\ref{sec:observation}. We model the five lensed quasars in \\S~\\ref{sec:model} and summarize our results in \\S \\ref{sec:summary}. We use a standard cosmological model with matter density $\\Omega_M=0.27$, cosmological constant $\\Omega_\\Lambda=0.73$, and Hubble constant $h=H_0/100{\\rm km\\,sec^{-1}Mpc^{-1}}=0.71$ \\citep[e.g.,][]{spergel03} throughout this paper. ",
        "conclusions": "\\label{sec:summary}  We discovered five high-redshift ($z_s>2.2$) lensed quasars, SDSS~J0819+5356, SDSS~J1254+2235, SDSS~J1258+1657, SDSS~J1339+1310, and SDSS~J1400+3134 from the SDSS. They were confirmed to be lenses by the imaging and spectroscopic observations at the UH88, ARC 3.5m, and TNG 3.6m telescopes. All five objects are two-image lensed quasars, with image separations of 1\\farcs28--4\\farcs04. The source redshifts range from 2.24 to 3.63. The lens redshift of SDSS~J0819+5356 is measured to be $z_l=0.294$ from the \\ion{Ca}{2} H\\&K absorption lines, whereas the lens redshifts of the other four objects are estimated to be 0.2--0.8 from the colors and magnitudes of the lensing galaxies. The image configurations and fluxes of all the lenses are well reproduced by standard lens models. We find signatures of strong external shears for SDSS~J1258+1657 and SDSS~J1339+1310, presumably coming from nearby galaxies whose redshifts are estimated to be similar to that of the lensing galaxy.  The statistical lensed quasar sample of the SQLS is restricted to $z_s<2.2$, and therefore all the lensed quasars discovered here will not be included in the SQLS statistical sample. The reason is that the SDSS quasars at ${z_s}>2.2$ are selected only from point sources and therefore the SDSS-selected quasars have a strong bias against our ``morphological selection''. Thus the five lenses will be included in a statistical sample when a homogeneous catalog with quasars at ${z_s}>2.2$ is constructed. However, the five lenses will definitely be useful for detailed future studies, such as deep spectroscopy for the lensing galaxies to measure their redshifts and velocity dispersions \\footnote{Currently, measurement of velocity dispersions is probably possible only for the lensing galaxy of SDSS~J0819+5356.} and for the quasar images to study the transverse structure in the Ly$\\alpha$ forest, and high-resolution imaging to see the structure of the systems. In addition, monitoring observations to measure time delays and microlensing events will provide useful opportunities to study the central structures of the quasars and to constrain the Hubble constant. These high-redshift quasar lenses will also be important to extend the redshift range of the lens applications; only about 30 objects out of the $\\sim$100 lensed quasars\\footnote{CASTLES webpage (C.~S.~Kochanek et al., http://cfa-www.harvard.edu/castles/.)} are identified to be lenses at ${z_s}>2.2$."
    },
    "0809/0809.0253_arXiv.txt": {
        "abstract": "Following Prendergast we study the relativistically expanding electromagnetic fields generated by an axisymmetric explosion of magnetic energy in a small volume. The magnetic field expands uniformly either within a cone or in all directions and it is therefore accompanied by an electric field. In the highly conducting plasma the charges move to annul the electric field in the frame of the moving plasma. The solutions presented are analytical and semi-analytical. We find that the time-scale for the winding up of the initial magnetic field is crucial, as short time-scales lead to strong radiant fields. Assuming a magnetic field of $10^{13}Gauss$ emerging from a magnetosphere of $10^{9}cm$ we end with a jet when confined by a pressure environment that falls more slowly than $r^{-4}$. The jet carries energy of $10^{51}erg$, which is mostly due to differential rotation at the base. ",
        "introduction": "Observations of a wide variety of astronomical objects suggests the existence of magnetic fields in relativistic environments. There have been many studies of magnetic fields emanating from differentially rotating systems. Some of them confine themselves to the non-relativistic regime in which the displacement currents can be neglected. Even then few of them are analytic e.g. \\cite{1976Natur.262..649L}, \\cite{1994MNRAS...267..146L}, \\cite{ 1994A&A...287..893S}, \\cite{2006MNRAS.369.1167L}, while most of them are computational e.g. \\cite{1970ApJ...161..541L}, \\cite{1995MNRAS.277.1327B}, \\cite{1997Natur.385..409O}. In the force-free case it was shown that the time evolution of the magnetic field arises solely from the time dependence of the boundary conditions so that the exact dynamical evolution can be calculated from the time dependent sequence of static models \\citep{2006MNRAS.369.1167L}. Those in turn can be derived from the energy principle. However, that simplification depends on both the force-free condition and the neglect of the displacement currents which is only valid when the velocities are much less than $c$. Relativistic problems are harder as that approximation is invalid, so most studies are purely computational e.g. \\cite{1992ApJ...394..459L}, \\cite{2002MNRAS.336..759K}, \\cite{2003ApJ...589..444G}, \\cite{2005ApJ...620..878D} and \\cite{2008MNRAS.388..551T} or semi-analytical e.g. \\cite{1995ApJ...446...67C}.  Despite this, \\cite{2005MNRAS.359..725P} was able to find an exact solution to the relativistic MHD problem of a point magnetic explosion. He derived a time dependent relativistic analogue of the Grad-Shafranov equation that governs axially symmetric force-free MHD by assuming that the radial coordinate and the time since the explosion only appear in the dimensionless combination $v=r/ct$. Strictly speaking such equations are only valid when the length scale and $c \\times$ the time scale of the region of the explosion are much smaller than $r$ and $ct$. The resulting Prendergast equation is non-linear but becomes linear in a special case. It was this special case that Prendergast studied in detail. However, he found that there are spherical nodes where the radial magnetic field is zero and these nodes occur before the highly relativistic regime $r \\approx ct$ is reached. This has the unfortunate consequence that magnetic field lines emanating from small $r$ turn back before they reach the extremely relativistic region where the displacement currents are very important. The field lines in that region are an appendage unattached to their origin. Very similar effects are well known when one takes the analogous case, $\\mathbf{j}= \\alpha \\mathbf{B}$ with $\\alpha$ constant, in non-relativistic MHD in spherical coordinates. When $\\alpha r$ becomes large, the system gives way to oscillating solutions with a series of nodes. This difficulty occurs because $\\alpha^{-1}$ has dimension $L$ and the field has to vary on this fixed length scale even at large r. At large distances there is too much current for unit field and the smaller scale is then reflected in the scale of oscillation.  It is known that non-linear ansatzes for the non-relativistic Grad-Shafranov equation can avoid this, which we now recognise as a bad consequence of a mathematically simple linear approximation which is not generally justified in the physics of the problem. We shall therefore study the non-linear Prendergast equation in all its glory!   We generalise the idea of force-free fields by considering a configuration where for each point there exists a Lorentz frame in which the electric field is zero and the current is along the magnetic field. The force density $\\mathbf{f}$ in the frame fixed at the origin is then:  \\begin{eqnarray} \\mathbf{f}=\\rho(\\mathbf{E}+\\mathbf{v} \\times \\mathbf{B})=0, \\end{eqnarray}  where $c\\mathbf{v}(\\mathbf{r},t)$ is the velocity of the moving frame. In this paper, we study such fields by solving the equation proposed by \\cite{2005MNRAS.359..725P}. However, before stepping to a solution, we extract as much information as possible about the fields that is independent of the detailed form of the solution. Then we solve the equations semi-analytically as it impossible to achieve general analytical solutions. The solutions simplify in the non-relativistic limit and converge to those of \\cite{1994MNRAS...267..146L}. There are some other cases that are  interesting and exactly soluble, namely the current-free magnetic dipole and the linear force-free field of Prendergast. The introduction of currents in the system allows the existence of a toroidal component of the field at the cost of making analytical solutions much harder, however it is still possible to design analytical solutions for this structure. ",
        "conclusions": "In this paper we have solved analytically the equation proposed by Prendergast for relativistically expanding axisymmetric self-similar force-free fields. This equation is a relativistically moving extension of the Grad-Shafranov equation for static force-free fields with axial symmetry.  To excellent accuracy our solution is given everywhere by equations (62) and (63) with the $f(\\mu)$ function given by equation (57) or more accurately (60). The resulting electromagnetic fields are given by equations (8) and (9). All solutions of the Prendergast equation need to be averaged over a small time $\\Delta \\tau$ that corresponds to the light crossing time across the magnetic explosion. His equation assumes a point explosion. Exact solutions of Prendergast's equation give singular fields where the expansion speed is that of light but we have shown how averaging over $\\Delta \\tau$ yields genuine solutions of Maxwell's equations without singularities. Nevertheless this $\\Delta \\tau$ may be quite short and the resulting fields close to $r=ct$ are strong, carrying around $1\\%$ of the system's energy. This may be related to $\\gamma$-ray burst precursors. However, the fields at the top require further modification as the ram pressure of their relativistic motion is retarded by the interstellar medium. Considerations as those of \\citep{2005ApJ...628..847R} are clearly important there, whatever drives the shock waves and it is there that the eponymous $\\gamma$-rays themselves are generated.  While this paper has given a basic mechanism that produces highly relativistic motion there is a big gap between that and the phenomenology of $\\gamma$-ray bursts.  On the way to understanding the solutions of the fully relativistic Prendergast equation we were led to give purely poloidal solutions of Maxwell's equations for time-dependent multipoles, interesting in their own right. We have also seen how in the non-relativistic limit our problem is related to the solutions found in \\cite{2006MNRAS.369.1167L} and \\cite{1994MNRAS...267..146L}.  Finally we have seen how small opening angles of the cone correspond to greater coiling of the magnetic field and thus to mush greater concentration  of energy into the jet.  Setting some values for the magnetic field quoted from the stronger magnetars we find that we can power jets up to the energies of $\\gamma$-ray bursts, for which the isotropic energy is between $10^{51}$ to $10^{54}$ erg \\citep{2005RvMP...76.1143P}. In the case studied we found that the energy inside the jet is of order $10^{51} erg$, where the isotropic energy is two orders of magnitude greater. The very strong magnetic field found at the top of the magnetic configuration can oppose the strong ram pressure because of expansion and swipe material out of the way for the jet to expand. This feature may be related to the precursor observed in some $\\gamma$-ray bursts e.g. \\cite{1991Natur.350..592M}, as it is the first to break out and carries a small fraction of the energy of the $\\gamma$-ray burst.  Jets of narrow opening angles can be more twisted than those of wide opening angles, giving evidence that collimation increases with twist.  Finally a strong stagnation pressure is found at the head of the jet. It prevents the jet from being relativistic near the base but as it expands outwards and encounters underdense material it becomes extremely relativistic as is indeed observed in $\\gamma$-ray bursts."
    },
    "0809/0809.2583_arXiv.txt": {
        "abstract": "I present an overview of important results obtained using high-resolution very long baseline interferometry (VLBI) observations of X-ray binary systems.  These results derive from both astrometric observations and resolved imaging of sources, from black holes to neutron star and even white dwarf systems.  I outline a number of upcoming developments in instrumentation, both new facilities and ongoing upgrades to existing VLBI instruments, and I conclude by identifying a number of important areas of investigation where VLBI will be crucial in advancing our understanding of X-ray binaries. ",
        "introduction": "Very Long Baseline Interferometry (VLBI) has provided us with some of the highest resolution observations to date of radio-emitting systems throughout the visible Universe.  X-ray binary systems, accreting compact objects in orbit with less-evolved donor stars, emit synchrotron radiation from their relativistic jets.  During outbursts, the radio emission from these sources becomes bright, reaching levels of several Jy in the brightest sources, and is often resolved with high-resolution observations, making such sources ideal targets for study with VLBI arrays.  Being Galactic systems, X-ray binaries are much closer than their scaled-up extragalactic analogues, the Active Galactic Nuclei (AGN). However, since they are typically $10^6$--$10^8$ times less massive than AGN, but only $10^4$ times closer (a few kpc compared to several tens to hundreds of Mpc), it is, counterintuitively, possible to probe closer to the compact object (in units of gravitational radii) in AGN than in X-ray binary systems, so the Galactic systems are in fact less useful for studying the close-in regions where the jets are accelerated and collimated.  Furthermore, since the number of X-ray binaries with resolved radio emission is still small, studies of these sources are to some extent limited by the peculiarities of individual objects.  However, studies of X-ray binary systems do have some advantages over AGN.  Their proximity means that fundamental parameters such as distance and space velocity can be determined by astrometric measurements.  And since lengths and timescales close to the compact object scale with black hole mass, it is possible to observe the evolution of the jets in X-ray binaries on timescales of hours to days, rather than years to decades, following the evolution of the source through its entire duty cycle in a matter of months, and watching the evolution of the radio jets in real time as they move out and interact with their environments.  Furthermore, since the class of X-ray binaries comprises accreting neutron stars as well as black hole systems, they allow us to probe the importance of a deep potential well and the existence of a solid surface and a stellar magnetic field in the acceleration and collimation of relativistic jets.  In this article I will attempt to provide a brief overview of recent high-resolution studies of X-ray binary systems.  In this definition, I choose to include all accreting Galactic compact objects, from black holes through neutron stars, and even white dwarf systems.  I will outline new and upcoming developments in instrumentation, as well as their potential applications in solving some of the important open questions in the field of X-ray binary research. ",
        "conclusions": ""
    },
    "0809/0809.5148_arXiv.txt": {
        "abstract": "In this contribution to the conference ``Beyond Einstein: Historical Perspectives on Geometry, Gravitation and Cosmology in the Twentieth Century'', we give a critical status report of attempts to explain the late accelerated expansion of the universe by modifications of general relativity. Our brief review of such alternatives to the standard cosmological model addresses mainly readers who have not pursued the vast recent literature on this subject. ",
        "introduction": "The phenomenologically very successful cosmological `concordance model', within the framework of general relativity (GR), leaves us with the mystery of dark energy (DE). Since no satisfactory explanation of DE has emerged so far\\footnote{See, e.g., \\cite{CST}, \\cite{NS}, and references therein.}, it is certainly reasonable to investigate whether possible modifications of GR might change the late expansion rate of the universe. After all, GR has not yet been tested on cosmological scales.  Modified gravity models have to be devised such that they pass the stringent Solar System tests, and are compatible with the rich body of cosmological data that support the concordance model ($\\Lambda$CDM model). At the same time, the theories should be consistent on a fundamental level. Since we are dealing with higher spin equations, possible acausalities are, for instance, a serious issue.  Apart from all that, one should not forget that the old profound vacuum energy problem \\cite{NS} and the cosmic coincidence problem remain, and thus extreme fine tuning is unavoidable. This holds, of course, also for all dynamical models of DE \\cite{CST}.  In my brief review I shall mainly concentrate on so called $f(R)$ gravity. This is the simplest modification. Moreover, there have been some recent developments that I find interesting. After some generalities, many of you know very well, I shall discuss the weak field limit and Solar System tests. For some time there was a lot of confusion on this issue, with conflicting statements, but was eventually clarified. We shall, however, see that the weak field approximation may break down, and a so-called Chameleon mechanism can be at work that hides a scalar degree of freedom of the theory on solar system scales.  There are $f(R)$ models that pass the solar system tests and are cosmologically almost indistinguishable from the successful $\\Lambda$CDM model. Recently it was, however, discovered by Kobayashi and Maeda that these models are in serious trouble in the strong-field regime.  Some of the other modified gravity theories are even in greater difficulties. This will be briefly discussed in a final part. ",
        "conclusions": "A positive aspect of the largely negative outcome of the previous discussion seems to me that the distinguished role of GR among large classes of gravity theories has once more become apparent. We know, of course, that GR is an effective theory, and that quantum theory will produce all sorts of induced terms (a phenomenon that is well-known from QED), but stopping any expansion after a few terms will hardly lead to a consistent theory that agrees with observations on all scales.  Some of the modified gravity theories, such as $f(R)$ or braneworld models, may perhaps be of limited use for testing GR on cosmological scales. Guided by such models\\footnote{Especially from the evolution of linear cosmological perturbations for such models \\cite{SHS}, \\cite{BBP}, \\cite{PS}.}, there have recently been some interesting attempts to develop a parameterized post-Friedmann description of gravity that parallels the parameterized post-Newtonian description of Solar System tests (discussed by C. Will at this meeting). In contrast to the latter, there appear unavoidably some free functions, instead of just a bunch of parameters, in the description of the evolution of inhomogeneities \\cite{Hu2}, \\cite{Ber}. It will, therefore, be difficult to discriminate between dark energy and modified gravity, but this remains a major goal for years to come. One can hope that this will eventually become possible with better data on the CMB background, weak gravitational lensing, and the growth of large scale structures."
    },
    "0809/0809.2597_arXiv.txt": {
        "abstract": "{The Copernican principle, stating that we do not occupy any special place in our universe, is usually taken for granted in modern cosmology. However recent observational data of supernova indicate that we may live in the under-dense center of our universe, which makes the Copernican principle challenged. It thus becomes urgent and important to test the Copernican principle via cosmological observations. Taking into account that unlike the cosmic photons, the cosmic neutrinos of different energies come from the different places to us along the different worldlines, we here propose cosmic neutrino background as a test of the Copernican principle. It is shown that from the theoretical perspective cosmic neutrino background can allow one to determine whether the Copernican principle is valid or not, but to implement such an observation the larger neutrino detectors are called for.} \\begin{document} ",
        "introduction": "Based on the cosmological principle, the recent observations of type Ia Supernova indicate that either our present universe is dominated by some new exotic matter called dark energy or a modification of general relativity is needed on cosmic scales at least in an effective sense, which raises profound questions on fundamental physics. A conservative way out is to engage in the speculation on the validity of the cosmological principle. As an assumption of the isotropy and homogeneity throughout our universe, the cosmological principle is known to be partly satisfied. The observation of near-isotropy of cosmic microwave background(CMB) spectrum implies that our universe is very nearly isotropic. In addition, our universe is observed to be approximately homogeneous on large scales\\cite{Hogg,YBPS}. However, the radial homogeneity on scales of Gpc remains to be confirmed. In fact, it has been theoretically shown that the spherically symmetric but radially inhomogeneous Lemaitre-Tolman-Bondi(LTB) cosmological models can explain the supernova data very well without introducing the dark energy or resorting to a modification of general relativity\\cite{Celerier,Moffat1,Moffat2,Garfinkle,Enqvist,ABNV}. Different from the case of the Friedmann-Robertson-Walker(FRW) cosmological models, in the context of the LTB cosmological models, we are constrained in a special position, i.e., in or near the center of a void where the local matter density is relatively low, which violates the Copernican principle. So whether or not the Copernican principle is valid plays a pivotal role in our understanding of the genuine mechanism underlying the evolution of our universe. Although the Copernican principle may be widely accepted by fiat, it should be observationally tested without an a priori bias.  To date, there have been some ideas proposed to serve as observational tests of the Copernican principle\\cite{CS,UCE,CBL,CFL,YKN,BW}. Note that in all of these proposals the observational data come from the past light cone. On the other hand, neutrinos are believed to be massive, however small the mass is. Thus as illustrated in Fig.\\ref{shooting}, they will freely travel to us from inside of the past light cone after decoupling, which definitely brings more information about the structure of our universe. Keeping this in mind, we here put forward cosmic neutrino background(CNB) as a new test of the Copernican principle if observed.  In next section, we shall provide a brief review of the LTB cosmological models, where a specific parametrization is also chosen. In subsequent section, we will work out the cosmic neutrino background spectrum in the chosen LTB and FRW cosmological models, and examine the feasibility of our proposal by demonstrating the spectrum difference between the LTB and FRW cases. We then finish the paper with some discussions.  Unless otherwise is explicitly stated, Planck units are used here, i.e., $c=G=\\hbar=k=1$.  \\begin{figure}[htb!] \\includegraphics[scale=0.5]{shooting.eps} \\caption {\\small\\sl Different from the cosmic photons, the cosmic neutrinos of different energies come from the different places on the surface of constant $t_L$ and travel to us along the different worldlines.} \\label{shooting} \\end{figure} ",
        "conclusions": "Motivated by the intuitive observation that the cosmic neutrinos of different energies travel to us from the different places, we have investigated the feasibility of cosmic neutrino background as a test of the Copernican principle. As a result, in the region of small momenta, cosmic neutrino background shows us the definite signal to determine whether our universe is homogeneous or not. However, because this signal shows up at low energies, it seems invisible in our current neutrino detectors. A more detailed evaluation of the possibility to carry out our proposal in the future neutrino telescopes is worthy of further investigation, which goes beyond the scope of this paper. But we expect that our theoretical analysis here manifests the very advantage of neutrinos over photons in bringing extra information about our universe, thus provides another stimulus to such endeavors of construction of large neutrino telescopes.  We conclude with some caveats. Note that our result of $n\\sim p$ relation depends on the isothermal distribution assumption of cosmic neutrinos at the time $t_L$. So a deviation from the assumed distribution will enhance or suppress the signal. In addition, so far our discussions have been restricted to the specific LTB cosmological model. More general LTB cosmological models may give us somewhat different results. Although a more realistic neutrino distribution is believed to make the signal enhanced, and the results are expected to display the same qualitative behavior for other viable LTB cosmological models, a careful investigation is needed. Last but not least, if neutrinos are heavy enough, then the local gravitational clustering of cosmic neutrinos will become significant such that the $n\\sim p$ relation will be distorted\\cite{RW}. It is thus important to see how our result is influenced by this clustering effect. We expect to report these subtle issues elsewhere."
    },
    "0809/0809.0247_arXiv.txt": {
        "abstract": "Warm dark matter (WDM) may resolve the possible conflict between observed galaxy halos and the halos produced in cold dark matter (CDM) simulations. Here we present an extension of MSSM to include WDM by adding a gauge singlet fermion, $\\overline{\\chi}$, with a portal-like coupling to the MSSM Higgs doublets. This model has the property that the dark matter is {\\it necessarily warm}. In the case where $M_{\\overline{\\chi}}$ is mainly due to electroweak symmetry breaking, the $\\overline{\\chi}$ mass is completely determined by its relic density and the reheating temperature, $T_R$. For $10^2 \\GeV \\; ^{<}_{\\sim} \\; T_{R} \\; ^{<}_{\\sim} \\; 10^{5} \\GeV$, the range allowed by $\\ochi$ production via thermal Higgs annihilation, the $\\overline{\\chi}$ mass is in the range 0.3-4 keV, precisely the range required for WDM. The primordial phase-space density, $Q$, can directly account for that observed in dwarf spheroidal galaxies, $Q \\approx 5 \\times 10^{6} {\\rm (eV/cm^3)/(km/s)^3}$, when the reheating temperature is in the range $T_R \\approx 10-100$ TeV, in which case $M_{\\overline{\\chi}} \\approx 0.45$ keV. The free-streaming length is in the range 0.3-4 Mpc, which can be small enough to alleviate the problems of overproduction of galaxy substructure and low angular momentum of CDM simulations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2132_arXiv.txt": {
        "abstract": "We report on observations of correlated behavior between the prompt $\\gamma$-ray and optical emission from GRB 080319B, which (i) strongly suggest that they occurred within the same astrophysical source region and (ii) indicate that their respective radiation mechanisms were most likely dynamically coupled. Our preliminary results, based upon a new cross-correlation function (CCF) methodology for determining the \\emph{time-resolved} spectral lag, are summarized as follows. First, the evolution in % the arrival offset of prompt $\\gamma$-ray photon counts between Swift-BAT 15-25 keV and 50-100 keV energy bands \\emph{(intrinsic $\\gamma$-ray spectral lag)} appears to be anti-correlated with the arrival offset between prompt 15-350 keV $\\gamma$-rays and the optical emission observed by TORTORA \\emph{(extrinsic optical/$\\gamma$-ray spectral lag)}, thus effectively partitioning the burst into two main episodes at $\\sim T+28\\pm2$ sec. Second, prompt optical emission is nested within intervals of (a) trivial intrinsic $\\gamma$-ray spectral lag ($\\sim T+12\\pm2$ and $\\sim T+50\\pm2$ sec) with (b) discontinuities in the hard to soft evolution of the photon index for a power law fit to 15-150 keV Swift-BAT data ($\\sim T+8\\pm2$ and $\\sim T+48\\pm1$ sec), both of which coincide with the rise ($\\sim T+10\\pm1$ sec) and decline ($\\sim T+50\\pm1$ sec) of prompt optical emission. This potential discovery, robust across heuristic permutations of BAT energy channels and varying temporal bin resolution, provides the first observational evidence for an implicit connection between spectral lag and the dynamics of shocks in the context of canonical fireball phenomenology. ",
        "introduction": "Swift's unique dynamic response and spatial localization precision, in conjunction with correlative ground-based follow-up efforts, has resulted in the collaborative broad-band observations of GRB 080319B \\cite{Racusin:2008c}. In the context of an analysis focused on confronting the lag-luminosity relation \\cite{Norris:2000b} in the Swift era \\cite{Stamatikos:2008b}, a correlation was observed between the evolution of time-resolved spectral lag and the behavior of the extraordinarily well-sampled prompt optical emission light curve associated with GRB 080319B. In general, the spectral lag is determined via either a peak pulse fit \\cite{Norris:2005b} or cross-correlation function (CCF) analysis \\cite{Band:1997b}. Previous studies have reported on the variability of spectral lag throughout burst emission \\cite{Chen:2005}, as well as its correlation to pulse evolution \\cite{Hakkila:2008}. In this work, we develop a new method to calculate the time-resolved spectral lag via a modification to the traditional CCF approach. In this manner, we are able to explore for the first time the evolution and correlation of the intrinsic $\\gamma$-ray spectral lag with prompt optical emission. We interpret these correlated behaviors as strong observational evidence that the prompt optical and $\\gamma$-ray emission of GRB 080319B took place within the same astrophysical source region, with indications that their respective radiation mechanisms where dynamically coupled throughout the prompt phase of the burst. ",
        "conclusions": "The generic agreement of the overall temporal coincidence and morphology between the prompt $\\gamma$-ray and optical light curves (Figure~\\ref{plot}, Panels A \\& B) suggests that they arose from a common source region. However, separate radiation mechanisms were most likely responsible since the extrapolated $\\gamma$-ray flux density to the optical band was deficient by $\\sim$4 orders of magnitude when compared to observation \\cite{Racusin:2008c,Kumar:2008,Yu:2008}. The steep rise/decline, short duration and lack of increasing pulse width of the prompt optical emission disfavor external forward/reverse shocks. Hence, internal shocks have been suggested as the source region with synchrotron emission responsible for the optical and inverse Compton scattering/synchrotron self Compton for the $\\gamma$-rays, with associated GeV photon emission \\cite{Racusin:2008c,Kumar:2008}. Alternatively, it has been suggested that non-relativistic forward internal shocks generated the prompt optical emission, while relativistic reverse internal shocks generated the prompt $\\gamma$-rays, with sub-GeV/MeV photon emission \\cite{Yu:2008}. Such high energy emission may be tested by Fermi (formally known as GLAST) via joint analyses with Swift-BAT \\citep{Stamatikos:2008c}.  Our preliminary results are illustrated in Figure~\\ref{plot}, Panels A-D. There are several observations that lead to effectively separating the burst's duration into two main episodes partitioned roughly at the midpoint of $\\sim T+28\\pm2$ sec. The first is that the bimodal evolution of the intrinsic $\\gamma$-ray spectral lag increases at t $>\\sim T+28\\pm2$ sec, which appears to be anti-correlated with the extrinsic (optical/$\\gamma$-ray) spectral lag observed via the smoothed Gaussian fits, where the spectral lag is of the order of a few seconds when t < $\\sim T+28\\pm2$ sec (Panels A \\& B). Beyond this common midpoint, the optical and $\\gamma$-rays do not correlate as well or at least are ambiguously correlated, i.e. extrinsic (optical/$\\gamma$-ray) spectral lag is either zero or negative at later times (ambiguity due to peak misalignment). Hence, the intrinsic time-resolved $\\gamma$-ray spectral lag is maximum at t $>\\sim T+28\\pm2$ sec, while the extrinsic time-resolved (optical/$\\gamma$-ray) spectral lag is maximum at t $<\\sim T+28\\pm2$ sec. In addition, an independent analysis \\cite{Margutti:2008} of BAT 15-150 keV light curves has revealed that the characteristic variability timescale of GRB 080319B was $\\sim$100 ms for t < $\\sim T$+28 sec and $\\sim$1 sec for t > $\\sim T$+28 sec. Furthermore, time-resolved spectral analysis by Konus-Wind illustrated that E$_{peak}$ decreased from $751\\pm26$ keV to $537\\pm28$ keV at $\\sim T+24\\pm2$ sec \\cite{Racusin:2008c}.  The hard to soft evolution of the photon index for time-resolved power law fits to 15-150 keV Swift-BAT data occurs in steps at $\\sim T+8\\pm2$ and $\\sim T+48\\pm1$ sec, which coincide with the respective rise ($\\sim T+10\\pm1$ sec) and decline ($\\sim T+50\\pm1$ sec) of prompt optical emission (see Panels A \\& C). This is consistent with Konus-Wind time-resolved spectral analysis and hardness ratios \\cite{Racusin:2008c}. We also note that intervals of trivial intrinsic spectral lag coincide with the rise and decline of prompt optical emission (see Panels A \\& D), at $\\sim T+12\\pm2$ and $\\sim T+50\\pm2$ sec, respectively. We interpret these correlated behaviors as strong observational evidence that the prompt optical and $\\gamma$-ray emission took place within the same astrophysical source region, which until now has only been conjecture. We also find indications for a dynamical coupling between the radiation mechanisms, perhaps via the processes mentioned above \\cite{Racusin:2008c,Kumar:2008,Yu:2008}.  This potential discovery, provides the first observational evidence for an implicit connection between spectral lag and the dynamics of internal shocks in the context of canonical fireball phenomenology. A full theoretical analysis of this result is currently in preparation. Future work includes an application of our methodology to observations of a subset of bursts with prompt optical emission to probe either the unique or ubiquitous nature of these observations. Ultimately, understanding the mechanism(s) responsible for spectral lag may reveal a fundamental and unprecedented view from within the GRB fireball.    \\begin{theacknowledgments} The authors are grateful to Craig Markwardt and David Band for very fruitful discussions in regards to this analysis. M. Stamatikos is supported by an NPP Fellowship at NASA-GSFC administered by ORAU. \\end{theacknowledgments}"
    },
    "0809/0809.2418_arXiv.txt": {
        "abstract": "{} {We present a dynamical analysis of the galaxy cluster Abell~1942 based on a set of 128 velocities obtained at the European Southern Observatory.} {Data on individual galaxies are presented and the accuracy of the determined velocities is discussed as well as some properties of the cluster. We have also made use of publicly available Chandra X-ray data.} {We obtained an improved mean redshift value  $  z = 0.22513 \\pm 0.0008$ and velocity dispersion $\\sigma= 908^{+147}_{-139}\\kms$. Our analysis indicates that inside a radius of $\\sim 1.5 h_{70}^{-1}\\,$Mpc ($\\sim 7\\,$arcmin) the cluster is well relaxed, without any remarkable feature and the X-ray emission traces fairly well the galaxy distribution. Two possible optical substructures are seen at $\\sim 5\\,$arcmin from the centre towards the Northwest and the Southwest direction, but are not confirmed by the velocity field. These clumps are however, kinematically bound to the main structure of Abell~1942. X-ray spectroscopic analysis of Chandra data resulted in a temperature $kT = 5.5 \\pm 0.5\\,$keV and metal abundance $Z = 0.33 \\pm 0.15 Z_{\\odot}$. The velocity dispersion corresponding to this temperature using the $T_{X}$--$\\sigma$ scaling relation is in good agreement with the measured galaxies velocities. Our photometric redshift analysis suggests that the weak lensing signal observed at the south of the cluster and previously attributed to a ``dark clump'', is produced by background sources, possibly distributed as a filamentary structure.} {} ",
        "introduction": "\\label{Introduction}  In the hierarchical $\\Lambda$CDM scenario for structure formation, clusters of galaxies are the largest coherent and gravitationally bound structures in the Universe, growing by accretion of nearby galaxy groups or even other clusters. These newcomers are often observed as substructures in the galaxy distribution and, indeed, substructures have been detected in a significant fraction of galaxy clusters \\citep[e.g.,][]{flin06}. Clusters can then be used to trace the cosmological evolution of structure with time and to constrain cosmological parameters \\citep[e.g.,][]{richstone92,kauffmann93}.  However, clusters comprise a diverse family, presenting a large range of structural behaviour and, in order to be useful as cosmological probes, the structural and dynamical properties of individual systems should be determined. Clusters are also complex entities, containing both baryonic and non-baryonic matter. In the case of the former, most of the baryons occupy the cluster volume in the form of a hot gas emitting in X-rays. Consequently, studies of galaxy clusters aiming to unveil their actual properties are greatly benefited by multiwavelength observations, in particular in X-rays for the gas and in the optical for the galaxies and even the dark matter.  Here we present a study of the cluster of galaxies Abell~1942. It has richness class 3 and Bautz-Morgan type III. It is of particular interest since it has at first glance a quite symmetrical morphology, similar to Abell~586 which can be used as a laboratory to test different mass estimators and to analyze its dynamics \\citep[as in][]{Cypriano05}. This cluster was observed in X-ray by several satellites, ASCA, ROSAT, and Chandra (see Section \\ref{x-ray}, below). Finally, A1942 has receive some attention because a putative mass concentration would have been detected by its shear effect at 7~arcmin southward from the cluster centre, with no obvious concentration of bright galaxies at this location \\citep{Erben00, Linden06}: the dark clump. A ROSAT-HRI image was also analyzed by the same authors, showing that the brightest peak of the X-ray emission corresponds with the cluster centre and its central galaxy. A weak secondary source was also detected at 1~arcmin from the mass concentration. Using ASCA data, \\citet{White00} gives 2 temperatures for this cluster: 5.6~keV for a broad band single temperature fit, and 15.6~keV for a cooling-flow fit, which would corresponds to a velocity dispersion $\\simeq 1800~\\kms$ and also a huge cooling flow of 817 $M_{\\odot}$/year.  From the optical data, only 2 velocities of galaxy members were available, one being the radio-source PKS 1435+038 \\citep{Kristian78}. Moreover, a deep image of the cluster centre \\citep{Smail91} shows the existence of a few lensed arcs, with one being close to the central galaxy. Up to now, no detailed lens model is available to study the central mass distribution.  In this paper, we analyze Abell~1942 from its photometric, spectroscopic and X-ray properties. In Section~\\ref{sec:photo}, we give evidence of the structure and substructures of the cluster from its photometric data. Section~3 presents the spectroscopic survey of the cluster galaxies in order to study the velocity dispersion in the cluster centre, as well as its variations with the radius until the measured limit of the shear up to 8~arcmin (equivalent to a radius of $1.7 h^{-1}_{70}$~Mpc at the cluster redshift). In Section~4 we analyze the X-ray data. The velocity analysis is detailed in Section~5. With such a set of velocities we build in Section~6 a detailed image of the cluster dynamics and mass distribution. Moreover we analyze the velocity distribution of the galaxies located close to the mass concentration area in order to know its nature as distant cluster or concentration of matter associated with the main cluster. We adopt here, whenever necessary, $H_{0}= 70\\, h_{70}\\kms$Mpc$^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We have investigated the cluster of galaxies Abell~1942. More than a hundred new spectroscopic redshifts were measured in a $14 \\times 14$ region around its centre. Together with X-ray archive data from Chandra, and photometric data from SDSS, we were able to achieve the dynamical and kinematical analysis of this cluster for the first time.  -- We found that about half of the observed galaxies are kinematical members of the cluster. We have also found some kinematical evidence for the presence of nearby groups of galaxies whose spatial counterparts however have not been confirmed. The cluster is situated at redshift 0.22513.  -- Our analysis indicates that inside a radius of $\\sim 1.7 h_{70}^{-1}\\,$Mpc ($\\sim 7\\,$arcmin) the cluster galaxy distribution is well relaxed without any remarkable feature and with a mean velocity dispersion of $\\sigma= 908^{+147}_{-139}\\kms$.  -- We have analyzed archival Chandra data and derived a mean temperature $kT = 5.5 \\pm 0.5\\,$keV and metal abundance $Z = 0.33 \\pm 0.15 Z_{\\odot}$. The velocity dispersion corresponding to this temperature using the $T_{X}$--$\\sigma$ scaling relation is in good agreement with the measured galaxies velocities. The X-ray emission traces fairly well the galaxy distribution.  -- We derive dynamical mass estimates of the cluster, assuming hydrostatic equilibrium of the (isothermal) intracluster X-ray emitting gas. At a radius equivalent to $r_{500}$ we obtained $M_{\\rm dyn}(< r_{500}) \\sim 5 \\times 10^{14} M_{\\odot}$. An estimate of the dynamical mass using the gravitational arc found at about $ \\sim 8.2$ arcsec from the cluster centre showed it to be consistent with that derived from the hydrostatic equilibrium hypothesis.  -- We do not confirm the mass concentration 7~arcmin south from the cluster centre from our dynamical and X-ray analysis. However, we do see a concentration of background galaxies towards these regions which may be at the origin of the weak lensing signal detected before."
    },
    "0809/0809.3134_arXiv.txt": {
        "abstract": "Substructures, expected in cold dark matter haloes, have been proposed to explain the anomalous flux ratios in gravitational lenses. About 25\\% of lenses in the Cosmic Lens All-Sky Survey (CLASS) appear to have luminous satellites within \\mbox{$\\sim$ 5 $ h^{-1}$ kpc} of the main lensing galaxies, which are usually at redshift $z\\sim 0.2-1$.  In this work we use the Millennium Simulation combined with galaxy catalogues from semi-analytical techniques to study the predicted frequency of such satellites in simulated haloes.  The fraction of haloes that host bright satellites within the (projected) central regions is similar for red and blue hosts and is found to increase as a function of host halo mass and redshift. Specifically, at $z = 1$, about $11$\\% of galaxy-sized haloes (with masses between $10^{12} h^{-1}$ M$_\\odot $ and $ 10^{13} h^{-1}$ M$_\\odot$) host bright satellite galaxies within a projected radius of 5 $ h^{-1}$ kpc.  This fraction increases to about 17\\% (25\\%) if we consider bright (all) satellites of only group-sized haloes (with masses between $10^{13} h^{-1}$ M$_\\odot $ and $10^{14} h^{-1}$ M$_\\odot$).  These results are roughly consistent with the fraction ($\\sim 25\\%$) of CLASS lensing galaxies observed to host luminous satellites.  At \\mbox{$z = 0$}, only $\\sim 3$\\% of galaxy-sized haloes host bright satellite galaxies.  The fraction rises to $\\sim 6\\%\\, (10\\%)$ if we consider bright (all) satellites of only group-sized haloes at $z = 0$.  However, most of the satellites found in the inner regions are `orphan' galaxies where the dark matter haloes have been completely stripped.  Thus the agreement crucially depends on the true survival rate of these `orphan' galaxies.  We also discuss the effects of numerical resolution and cosmologies on our results. ",
        "introduction": "In the hierarchical scenario of structure formation, structures in the Universe are assumed to have grown from tiny quantum fluctuations (generated during an inflationary period) through gravitational instability.  Due to the shape of the power spectrum, structures form hierarchically - larger structures form via accretion and merging of smaller structures.  Dense cores of the smaller structures often survive the merging process and manifest as subhaloes in the primary haloes.  If substantial star formation occurs in these subhaloes (or their progenitors), then they will appear as satellite galaxies.  In the Milky Way, hundreds of subhaloes are predicted, starting from earlier semi-analytical studies (\\citealt{bib:Kauffmann93}), to more recent high-resolution simulations (\\citealt{bib:Klypin99}; \\citealt{bib:Moore99}; \\citealt{bib:Gao04_gal,bib:Gao04_dm}; \\citealt{bib:Diemand07}).  A few years ago, there were only a dozen or so satellites known, far fewer than the predicted number of subhaloes.  However, very recently, a new population of satellites has been discovered in the Sloan Digital Sky Survey data (e.g. \\citealt{bib:Belokurov07}).  It should be noted though that these satellites are compact and, in general, much fainter than the previously known ones, thus it is likely that even this new population of satellite galaxies cannot completely remove the discrepancy between simulations and observations (\\citealt{bib:Madau08}).  It is possible that many subhaloes are dark due to inefficient star formation, for example, due to its suppression by the UV-background radiation (e.g. \\citealt{bib:Doroshkevich67}; \\citealt{bib:Couchman86}; \\citealt{bib:Efstathiou92}).  Such dark substructure can, potentially, be detected through several means, for example through gamma-ray radiation due to annihilations of dark matter particles (\\citealt{bib:Stoehr03}; \\citealt{bib:Die07}).  Gravitational lensing is, in principle, another way to detect dark (and luminous) substructure.  Flux anomalies (\\citealt{bib:Mao98}), astrometric perturbations (\\citealt{bib:Chen07}) and time-delays (\\citealt{bib:Keeton08}) can be used to infer the presence of substructure in strong gravitational lenses.  The results of these studies are so far inconclusive (e.g. \\citealt{bib:Kochanek04, 2004ApJ...604L...5M}). If all substructure is equally efficient in affecting the flux ratios, then it is clear that there is more than sufficient mass in subhaloes to explain the flux anomalies.  Unfortunately, most subhaloes are in the outer part of the galaxy halo, which means they will have relatively little impact on the flux anomalies occurring in the central parts of lensing galaxies.  Curiously, as emphasised by Schneider (2007, private communication), 3 of the 6 radio lenses studied by \\cite{bib:Kochanek04} exhibit luminous satellite galaxies close to the primary lensing galaxy, namely MG0414+0534, B1608+656 and B2045+265.  A question naturally arises: are such luminous satellite galaxies expected this frequently in the current structure formation theory?  The lensing cross-section is dominated by elliptical galaxies, thus for lensing applications it is important to divide galaxies into different types, and see whether the subhalo populations are different. Furthermore, we explore in more detail the evolution of satellite galaxies as a function of redshift.  If the evolution is slow, we can more conveniently use studies of nearby galaxies to infer the properties of the luminous satellite population of galaxies at intermediate redshift (between 0.5 and 1, where most lensing galaxies lie).  These are the two specific aspects of the subhalo population that we will address in this paper.  For this purpose, we will use the largest cosmological simulation combined with semi-analytical catalogues to select haloes and study their satellite populations.  We compare our results to the CLASS survey (\\citealt{bib:Browne03}; \\citealt{bib:Myers03}).  The plan of the paper is as follows.  In Section \\ref{sims}, we describe the Millennium Simulation and the semi-analytical galaxy catalogue we use.  Our main results are presented in Section \\ref{results}, and we finish with a summary in Section \\ref{summary}. ",
        "conclusions": "\\label{summary} In summary, for the CLASS survey, approximately 5 of the 22 primary lensing galaxies appear to have a faint companion within the projected central 5 $ h^{-1}$ kpc. The companions have luminosities of about $2 - 40$\\% of the primary galaxy. We have studied host galaxies covering a comparable range of luminosities and host-to-satellite separations to the CLASS lenses, and found that the predicted fraction of galaxy-(group-) sized haloes hosting central luminous satellites ($\\sim$ 3\\% (6\\%)  at $z = 0$; $\\sim$ 11\\% (17\\%) at $z = 1$) is slightly lower than (but possibly consistent with) the observed value.  While this fraction is largely independent of galaxy type, it is shown to increase with redshift. The Poisson probability of detecting luminous substructure in 5 out of 22 lenses, given that $3\\%$ of haloes host luminous substructure, is $\\sim 5 \\times 10^{-4}$.  If $17\\%$ of haloes host luminous substructure, the probability of such a detection is $\\sim 0.14$. Our prediction of the redshift and mass dependence appears to be roughly consistent with the data: three lenses with luminous satellites are in groups (see Section \\ref{sec:observations}), and all appear have redshifts close to 1, higher than the median redshift ($\\sim$ 0.6) of all CLASS lenses (see the right panel of Fig. \\ref{obs}). One possibility that we have not considered, is whether lensing galaxies are biased tracers of substructure; such bias may arise if substructure enhances the lensing cross-sections significantly. Previous studies, on cluster scales, for giant arcs indicates that the bias is small (\\citealt{bib:Hennawi07}); it remains to be seen whether this holds true for galaxy-scale lenses.  Observationally, the Sloan Lens ACS Survey (SLACS) seems to indicate that the lensing galaxies at z $\\sim$ 0.2 are typical early-type galaxies (\\citealt{bib:Treu08}). Another possibility is that some of the luminous `satellites' are not associated with lensing galaxies at all, but just happen to be along the line of sight (\\citealt{2005ApJ...629..673M}).  \\cite{bib:Shin08} recently studied the effect of satellite galaxies on gravitational lensing flux ratios using analytic expressions for the host potential and the satellite galaxies.  They use a spherically symmetric galaxy distribution, and assume that the three-dimensional number density falls off like $r^{-3.5}$, comparable to the Milky Way.  They show that the probability of a finding a large dwarf is about 10\\%  within two Einstein radii and about 3\\% within one Einstein radius.  We find that our $z = 0$ results are consistent with this.  The Millennium Simulation assumes a power-spectrum normalisation of $\\sigma_8 = 0.9$, slightly higher than the latest \\emph{WMAP} five-year result (\\citealt{bib:wmap5}), where $\\sigma_8 = 0.8$. A lower value of $\\sigma_8$ will mean that haloes are expected to form later and be less concentrated.  However, the impact of this parameter on our results is complicated.  The semi-analytic models allow some fine-tuning of parameters to match observations.  For example, \\cite{bib:Wang08} found no significant difference in the galaxy populations (at the redshift range relevant here) created from semi-analytic models based on the \\emph{WMAP} one-year (\\citealt{bib:wmap1_03}) and \\emph{WMAP} three-year (\\citealt{bib:Spergel07}) $\\sigma_8$ values of 0.9 and 0.722, provided suitable galaxy formation parameters were chosen (the difference becomes significant at high redshift).  We find that our results do not change significantly when based on the \\emph{WMAP}3 galaxy catalogue produced by \\cite{bib:Wang08} when the same merger timescale is adopted (as in their model C).  However, in their model B (which has the same star formation efficiency but a shorter merger timescale than the \\citealt{bib:DeLucia07} catalogue) we find a factor of $\\sim$ 2 fewer haloes with central substructure.  To summarise, while we find that the fraction of luminous satellites in group-sized haloes at $ z \\sim 1$ is roughly consistent with the observational data we caution that a firm conclusion can only be reached with higher-resolution simulations involving a realistic treatment of the gas processes.  At the same time, a larger sample of gravitational lenses will also be beneficial to constrain these models and allow more definitive conclusions on the properties of the substructure to be made."
    },
    "0809/0809.4008_arXiv.txt": {
        "abstract": "We explore the consequences of the existence of a very large number of light scalar degrees of freedom in the early universe.    We distinguish between {\\em participator\\/} and {\\em spectator\\/} fields. The former have a small mass, and can contribute to the inflationary dynamics; the latter are either strictly massless or have a negligible VEV.    In N-flation and generic assisted inflation scenarios, inflation is a co-operative phenomenon driven by $N$ participator fields, none of which could drive inflation on its own.  We   review upper bounds on $N$, as a function of the inflationary Hubble scale $H$.  We then consider stochastic and eternal inflation in models with $N$ participator fields showing that individual fields may evolve stochastically while the whole ensemble behaves deterministically, and that a wide range of eternal inflationary scenarios are possible in this regime. We then compute one-loop quantum corrections to the inflationary power spectrum. These are largest with $N$ spectator fields and a single participator field, and the resulting bound on $N$ is always {\\em weaker\\/} than those obtained in other ways. We find that loop corrections to the N-flation power spectrum do not scale with $N$, and thus place no upper bound on the number of participator fields. This result also implies that, at least to leading order, the theory behaves like a composite single scalar field.  In order to perform this calculation, we address a number of issues associated with loop calculations in the Schwinger-Keldysh ``in-in'' formalism. ",
        "introduction": "Many candidate theories of fundamental physics predict the existence of large numbers of scalar degrees of freedom.  Classically, if these modes are not excited they play no role in the cosmological dynamics.  Quantum mechanically, however, light scalar modes fluctuate with an amplitude set by the Hubble scale $H$ and, since everything couples to the graviton, they contribute to loop corrections. Thus, while adding $N$ light scalar modes need not change the classical dynamics of the early universe, we expect quantum  contributions that scale with $N$.  In addition, if these fields are in thermal equilibrium with the rest of the universe their contribution to the effective number of degrees of freedom modifies the relationship between density and temperature.  For a given $H$ we can therefore find an upper limit on  $N$, above which these modes dominate the cosmological evolution.  Consider a scenario with $N$ scalar fields, \\begin{equation} S =  \\int d^4 x \\sqrt{g}\\left[\\frac{M_p^2}{2}R+ \\sum_I \\left(-\\frac{1}{2}(\\partial{\\phi_I})^2 + V_I(\\phi_I)\\right) \\right] \\, ,\\label{eqn:Nflationaction} \\end{equation} where the total potential $V = \\sum_I V_I(\\phi_I)$ is a sum of $N$ uncoupled potential terms.\\footnote{We use the reduced Planck mass $M_p^2 = 1/8\\pi G$ throughout.}  A field  $\\phi_I$  is ``light''  if   $d^2V_I/d\\phi_I^2  /2\\equiv m_I^2 \\ll H^2$.  We make a distinction between {\\em participator\\/} and {\\em spectator\\/} fields.  The latter are either massless because $V_I$ is strictly zero, or their vacuum expectation value (VEV) is very small. Conversely, participator fields have small but non-zero mass and  sufficient VEV for them to help drive inflation via their contribution to the overall density.  Perhaps the simplest route to a meaningful bound on $N$ is to  note that all these fields undergo quantum mechanical fluctuations of order $\\delta \\phi_i \\sim H/2\\pi$. During inflation these fluctuations freeze out as classical perturbations at scales larger than the Hubble length,  $1/H$.  Each field has gradient energy  $(\\nabla \\phi)^2/2$, which counts towards the overall energy density.  The  gradient energy thus scales like $N(\\delta \\phi /\\delta x)^2 /2 \\sim N H^4/8 \\pi^2$.  Given $H$, the energy density is provided by the $0-0$ Einstein equation, $H^2 = \\rho/3{M_p^2}$.  For self-consistency,  the gradient  contribution must be much smaller than other contributions to $\\rho$, so \\begin{equation} N \\ll \\frac{M_p^2}{H^2}. \\label{eqn:gradientbound} \\end{equation} If  inflation occurs at the GUT scale, then $M_p/H\\sim 10^5$ and $N\\ll 10^{10}$. This bound can be derived by many routes  (e.g. \\cite{Huang:2007zt,Dvali:2008sy}), and appears to be robust.\\footnote{In \\cite{Ahmad:2008vy,Ahmad:2008eu} this limit on $N$ is derived in the context of N-flation but in reality it applies to {\\em any\\/} scenario with $N$ light fields.}  One can also consider bounds on $N$ from loop corrections to the gravitational constant. Since gravity (and thus the graviton) couples to all fields, matter loop corrections can renormalize its value \\cite{Zee:1981mk,Adler:1980bx}. Veneziano \\cite{Veneziano:2001ah} argued that in order for the effective value of Newton's constant to be greater than zero we need \\begin{equation} N \\ll \\frac{M_p^2}{\\Lambda^2}, \\label{eqn:speciesbound} \\end{equation} where $\\Lambda$ is the scale of the invariant UV cut-off, for example the mass scale of the $N$ stable fields.  Likewise, Veneziano \\cite{Veneziano:2001ah} points  out that this bound prevents potential violations of the holographic bound on the entropy density that can be encountered when the total number of species grows without limit   \\cite{Unruh:1982ic, Sorkin:1981wd}, the so-called ``Species Problem''. More recently, Dvali \\cite{Dvali:2007hz} noted that in the presence of a large number of species, the Veneziano mechanism weakens the gravitational coupling by a factor of $1/N$.  Setting  $N\\sim 10^{32}$  and working at the TeV scale to satisfy the bound of equation~(\\ref{eqn:speciesbound})  solves the  hierarchy problem, provided some ultraviolet completion of the standard mode can produce the requisite  value of $N$.  Dvali also provides an alternate non-perturbative derivation of equation (\\ref{eqn:speciesbound}) based on the consistency of black hole physics.  In Section~ \\ref{sect:eternal} we begin by  considering stochastic inflation \\cite{Vilenkin:1983xq,Linde:1986fd,Guth:2000ka} with multiple degrees of freedom. Many simple inflationary potentials possess a range of field values for which the potential is safely sub-Planckian, while the stochastic motion of the field  dominates the semi-classical rolling.  With a single field, a stochastic phase is necessarily eternal, since the inflationary domains perpetually reproduce themselves. However, once $N>1$, more complicated scenarios become possible.  Specifically, with  $N$ participator  fields we can implement {\\em assisted\\/} eternal inflation without any single field acquiring a super-Planckian VEV. This scenario is a variety of ``assisted inflation'' where the inflaton is a composite of many individual fields  \\cite{Liddle:1998jc}.  Secondly, we find solutions where {\\em individual\\/} fields move stochastically, but  a well-defined composite field rolls smoothly towards its minimum. Finally, we find models where a field (or fields) evolves semi-classically, along with other fields that move stochastically. If these fields have symmetry breaking potentials they yield an apparent ``multiverse'' of disconnected bubbles. However, the stochastic phase ends globally, as the semi-classically evolving field gives a natural cut-off to what would otherwise be an eternally inflating universe, providing a potential toy model for studying the well-known measure problem in eternal inflation \\cite{Linde:1993xx,Guth:2000ka,Garriga:2005av,Easther:2005wi, Bousso:2006ev, Aguirre:2006ak,Easther:2007sz}.  All fields couple to the inflaton gravitationally, and thus contribute loop corrections to the inflation potential, via a coupling of order $(H/M_p)^2$. This is necessarily a small number during the last 60 e-folds of inflation, given that $H$ fixes the scale of tensor fluctuations and is bounded  by the absence of an observed B-mode in the microwave background.  Consequently, any single loop makes a tiny  correction to the  inflaton propagator. However, this contribution is amplified by $N$, the number of species that can flow round the loops, and in Section \\ref{sect:loop} we compute the relevant one-loop corrections to the inflaton propagator.   We consider two limits -- $N$ spectator fields with a single inflaton, and the N-flation case, with $N$ participator fields  \\cite{Dimopoulos:2005ac,Easther:2005zr}.  With $N$ spectator fields we find an upper bound on $N$  similar to those of  \\cite{Veneziano:2001ah, Dvali:2007wp}.    We describe these results in  Section \\ref{sect:loop}, relegating many details to Appendix \\ref{app:1loop}. We work in the ``in-in\" formalism \\cite{Schwinger:1961,Keldysh:1964ud,Jordan:1986ug,Calzetta:1986ey}, which has been turned into an extremely powerful tool for studying higher order corrections to cosmological correlations by Maldacena \\cite{Maldacena:2002vr}  and  Weinberg \\cite{Weinberg:2005vy}.  This calculation requires us to develop the fourth order interaction Hamiltonian for a theory of inflation with $N$ uncoupled scalar fields,  which we present in Appendix \\ref{App:4ptderivation}, and which will have applications beyond the current calculation. Moreover, it turns out that there are some subtleties with the computation  of loop corrections in the in-in formalism which we also clarify in the Appendices.   Finally we conclude in Section \\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}    In this paper we explore modifications to the dynamics of inflationary models in the presence of $N$ light scalar degrees of freedom. We make a distinction between spectator fields, which do not contribute to the background energy density, and participator fields, whose potential terms contribute to the inflationary background.  As summarized in the introduction, a number of very general arguments place finite upper bounds on $N$, which typically take the form $N\\ll M_p^2/H^2$. We first consider the dynamics of stochastic inflation in the presence of large number of light fields, and show that when $N$ is large there is a distinction between stochastic and eternal inflation which does not apply in the single field case. In particular, with a large number of participator fields we show that there is a regime where the individual field motion is dominated by quantum fluctuations and thus stochastic, while the overall evolution of the universe is deterministic.  Moreover, if the stochastic fields have symmetry breaking potentials, then one can create a large number of apparent ``pocket'' universes, while retaining the ability to end inflation globally, and controlling the divergences characteristic of scenarios in which inflation is genuinely eternal.    Secondly, we explore loop corrections to the 2-point correlation function that provides the inflationary perturbation spectrum. Since any light field can run round a loop, these will typically scale with $N$. We analyze two subcases -- a single inflaton with $N$ spectator fields, and $N$ participator fields. In the former case, we find an explicit bound, but one which is weaker (by one over a factor that must be small during slow roll) than the simple form described above, namely \\begin{equation} N\\lsim \\frac{M_p^2}{H^2}\\frac{1}{\\epsilon}. \\end{equation} On the other hand, with $N$ participator fields ($N$-flation type scenarios), the loop correction is small and independent of $N$.  We can understand this result by recasting the action in terms of a composite single scalar field.  Finally, in the course of this work, we have had need to look closely at the computation of loop corrections in the in-in formalism.  These calculations raise a number of subtle issues, and we give details of our approach in the Appendices.  One might  whether bounds on $N$ actually rule out otherwise realistic fundamental theories. Within string theory, Vafa \\cite{Vafa:2005ui} argued that $T$ and $S$ dualities ensure that the volume of scalar moduli space is finite, and that  there is a strict upper bound on the number of matter fields.   We are not aware of explicit stringy constructions which would saturate the large $N$ bounds described above  in any reasonable compactification of string theory and with a finite non-zero Newton constant.  As we have noted, many independent arguments put limits on the allowed value of $N$.      For example, reference \\cite{Watanabe:2007tf} notes that if $N$ is too large, the inflaton can decay into a large number of  species during reheating, which may cause phenomenological problems in the later universe, while  \\cite{Leblond:2008gg} suggests that validity of the perturbative expansion itself provides a constraint on $N$.    We have chosen to work in the simplest models of multi-field inflation, where the fields have uncoupled potentials. One can  multifield hybrid inflation \\cite{Linde:1993cn} with large number of coupled fields, and it would be interesting to ask if such models which are not ruled out by radiative corrections to their potentials but by the loop corrections to their power spectrum.   Also, while higher correlation functions  themselves do not seem to impose any bound on $N$ \\cite{Battefeld:2006sz}, one could check whether this was also true of their quantum corrections."
    },
    "0809/0809.3244_arXiv.txt": {
        "abstract": "Since their discovery, cosmic crystalline silicates have presented several challenges to understanding dust formation and evolution. The mid-infrared spectrum of IRAS\\,17495$-$2534, a highly obscured oxygen-rich asymptotic giant branch (AGB) star, is the only source observed to date which exhibits a clear crystalline silicate absorption feature. This provides an unprecedented opportunity to test competing hypotheses for dust formation. Observed spectral features suggest that both amorphous and crystalline dust is dominated by forsterite (Mg$_2$SiO$_4$) rather than enstatite (MgSiO$_3$) or other silicate compositions. We confirm that high mass-loss rates should produce more crystalline material, and show why this should be dominated by forsterite. The presence of Mg$_2$SiO$_4$ glass suggests that another factor (possibly C/O) is critical in determining astromineralogy. Correlation between crystallinity, mass-loss rate and initial stellar mass suggests that only the most massive AGB stars contribute significant quantities of crystalline material to the interstellar medium, resolving the conundrum of its low crystallinity. ",
        "introduction": "\\label{intro}   One of the most exciting recent developments in astronomy was the discovery of crystalline silicate stardust by the Infrared Space Observatory \\citep[ISO;][]{isoref}. Crystalline silicates were initially discovered around evolved intermediate mass stars at far-infrared (IR) wavelengths \\citep{waters96}, but have since been detected in % young stellar objects \\citep{herbig}, comets \\citep{comets}, and Ultra Luminous Infrared Galaxies \\citep{ulirg}. In the interstellar medium (ISM) silicate grains in crystalline form are rare \\citep[$\\sim$1\\% crystalline by mass;][]{min07,kemper04}. However, Asymptotic Giant Branch (AGB) stars and their successors (pre-planetary and planetary nebulae) are the major contributors of material to the ISM \\citep{kwok04}. It has been suggested that AGB stars could contain \\gapprox 10\\% crystalline silicates \\citep{kemper04}, so it is vital to understand the formation of dust around AGB stars to resolve this apparent discrepancy.     AGB stars are luminous cool giants, with high mass-loss rates ($10^{-7} < \\dot{M} < 10^{-3}$\\,M$_{\\odot}$/yr), which increase over time \\citep{mloss}. They are highly evolved descendants of low- to intermediate-mass (0.8--8\\,M$_{\\odot}$) stars. Mass loss produces a circumstellar envelope, where dust forms. The order in which different species condense from the outflowing gas depends on the physical conditions within the envelope. Once $\\dot{M}$ is very high the circumstellar shell becomes optically thick, even at IR wavelengths, and the silicate emission features absorb themselves. The dust mineralogy is dictated by the dust condensation sequence and depends strongly on $\\dot{M}$ \\citep{dijk05}. \\citet{blommaert07} showed that the mineralogy depends only on the age of the star for a given initial stellar mass.   To date, crystalline silicate features have been seen clearly only in emission. It has been proposed that crystalline silicate absorption at 11$\\mu$m contributes to the broadening of the classic 10$\\mu$m amorphous silicate absorption feature associated with OH/IR stars \\citep{sylvester98,vanhollebeke}. However, IRAS\\,17495$-$2534 (hereafter I17495) is the first object to show a distinct crystalline 11$\\mu$m absorption feature.      I17495 is located in the Galactic Plane (gal.\\ coord.\\ = 003.6844$+$00.3880) at a distance of $\\sim$4kpc \\citep{loup93}, about half way to the Galactic Center. The IRAS LRS spectrum clearly exhibits an 11.1$\\mu$m absorption feature superposed on the classic 10$\\mu$m amorphous silicate feature (Fig.~\\ref{SED}). This is the first clear crystalline silicate absorption feature seen to date in any object.  $\\dot{M}$ for this object is estimated to be $\\sim 2 \\times 10^{-4} < \\dot{M} < 5\\times 10^{-4}$\\,M$_{\\odot}$/yr based on  mid-IR ([25]-[12]) color \\citep{loup93} and the range of optical depths ($15 < \\tau_{10 \\rm \\mu m} < 45$) consistent with our model results (see \\S~\\ref{RTmod}). Both $\\dot{M}$ and expansion velocity \\citep[$v_{\\rm exp}$, 16km/s;][]{loup93} are at the high end of normal for AGB stars. Most extreme O-rich AGB stars are OH/IR stars, exhibiting OH-maser emission. I17495 is one of a handful of visibly obscured O-rich AGB stars not exhibiting OH-maser emission \\citep[dubbed color-mimics;][]{bmlewis1}. However, the spectra of other color mimics more closely resemble those of OH/IR stars, i.e., when observed, the putative 11$\\mu$m crystalline silicate feature is merely a hump superposed on the regular amorphous silicate absorption feature. This is also true for OH/IR stars with similar mass-loss rates in the Galactic Bulge (GB). Even the GB source with the highest $\\dot{M}$ \\citep[$3\\times 10^{-4}$M$_\\odot$/yr; IRAS\\,17276$-$2846;][]{vanhollebeke} shows the 11$\\mu$m feature as a hump, rather than the distinct 11$\\mu$m feature seen in Fig.~\\ref{SED}.  While forsterite (crystalline Mg$_2$SiO$_4$) is expected to have a peak near 11.3$\\mu$m \\citep[e.g.][]{koike}, the peak can shift to shorter wavelengths. In the laboratory, \\citet{jaeger2} found that synthetic forsterite peaks closer to 11.2$\\mu$m, and matches the position of the hump observed in GB OH/IR stars \\citep{vanhollebeke}. \\citet{fabian} investigated grain shape effects. Their mass absorption coefficients for ellipsoidal forsterite grains are included in Fig.~\\ref{SED}, and demonstrate that the crystalline forsterite feature can peak at 11.1$\\mu$m. \\citet{boersma} attributed an 11.1$\\mu$m emission feature in the spectrum of a Herbig Ae/Be star to forsterite. \\citet{tamanai} found that, for non-embedded free-flying particles of forsterite, the feature actually peaks at 11.06$\\mu$m. Composition, grain-shape and grain-size can significantly shift the position of this feature. ",
        "conclusions": "We have presented the first crystalline silicate absorption feature observed to date. The spectrum of I17495 is enigmatic and suggests that its dust is dominated by Mg$_2$SiO$_4$ in both the crystalline and amorphous phases. We have confirmed that high mass-loss rates should produce more crystalline material than low mass-loss rates; and that the crystalline mineralogy should be increasingly dominated by forsterite as mass-loss rates increase. The most likely factor controlling both crystallinity and mineralogy is the C/O ratio. We suggest that the correlation between crystallinity, mass-loss rate and initial stellar mass mitigates the problem of the very different crystal fractions observed in some AGB stars and in the ISM. Only the rarer, higher mass AGB stars contribute a significant amount of crystalline material to the ISM."
    },
    "0809/0809.4286.txt": {
        "abstract": "%%% Abstract to run on from here. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0698_arXiv.txt": {
        "abstract": "We present an astrometry and photometry catalogue of globular cluster (GC) candidates detected with HST WFPC2 in a sample of 19 early-type galaxies, appropriate for comparison to the low-mass X-ray binary (LMXB) populations observed with \\chandra. In a companion paper, we present the \\chandra\\ data and investigate the relation between these populations. We demonstrate that, although there is little evidence of a colour-magnitude correlation for the GCs, after estimating mass and metallicity from the photometry, under the  assumption of a single age simple stellar population, there is a significant positive correlation between mass and metallicity. We constrained ${\\rm [Z/H] = (-2.1\\pm0.2)+(0.25\\pm0.04)log_{10}M}$, with a 1-$\\sigma$ intrinsic scatter of 0.62~dex in metallicity. If GCs are bimodal in metallicity this relation is consistent with recent suggestions of a mass-metallicity relation only for metal-poor clusters. Adopting a new technique to fit the GC luminosity function (GCLF) accounting for incompleteness and the Eddington bias, we compute the V-band local GC specific frequency (\\Sn) and specific luminosity (\\Sl) of each galaxy. We show that \\Sl\\ is the more robust measure of the richness of a GC population where a significant fraction is undetected due to source detection incompleteness. We find that the absolute magnitude of the GCLF turnover exhibits intrinsic scatter from galaxy to galaxy of $\\sim$0.3~mag (1-$\\sigma$), limiting its accuracy as a standard distance measure. ",
        "introduction": "Globular clusters (GCs) are found in galaxies of all morphological types and sizes. As primarily old stellar systems, their distribution and properties provide crucial insights into the way in which structure forms within the Universe. For a recent review we refer the reader to \\citet{brodie06a}. An intriguing characteristic of these objects, which provides valuable clues as to their structure and internal dynamics, is their association with low-mass X-ray binaries (LMXBs). Although it has long been recognized that LMXBs are overabundant per unit optical light by a factor $\\sim$100 in Milky Way GCs as compared to the field \\citep{fabian75a,clark75a}, the small number of sources has limited what this phenomenon can tell us about LMXB formation. With the advent of \\chandra, however, it has become feasible to resolve individual LMXBs in galaxies outside the Local Group, and thus to begin to assemble larger samples of LMXB-hosting GCs \\citep[\\eg][]{sarazin03}. Given their typically rich GC populations and their clean, old stellar populations which prevent contamination of the X-ray sources with high-mass X-ray binaries, massive early-type galaxies, in particular, provide an ideal environment in which to conduct such studies \\citep[\\eg][]{kundu02a,sarazin03,humphrey04a}.  Even with the small samples of galaxies in which the GC-LMXB connection has been investigated to date, some intriguing trends have already been observed. By comparing the numbers of GCs and the total luminosity of LMXBs in early-type galaxies \\citet{irwin05a} argued that a significant fraction of the LMXBs observed in the field form in the field, despite the stellar population being old. \\citet{juett05a} reported a similar result although, since they did not strictly compare the numbers of LMXBs and GCs within the same apertures, the strength of their correlation has been called into question \\citep[][]{kim05a}. Typically $\\sim$4\\% of GCs are observed to harbour LMXBs, with brighter and redder (more metal rich) GCs preferentially likely to contain them \\citep{kundu02a,sarazin03,kim05a}. The GC luminosity dependence is broadly consistent with the probability of harbouring an LMXB being proportional to the mass or to the stellar capture cross-section in the GC \\citep[\\eg][]{jordan04a,smits06a,sivakoff06a}. The origin of the colour dependence is unclear and several different possible explanations have been proposed \\citep[for a review of some of these, see][]{jordan04a}. \\citet{maccarone04a} suggested it may relate to different patterns of mass-transfer in metal-rich and metal-poor LMXBs. Alternatively, it may arise due to variations in the IMF between metal-poor and metal-rich clusters \\citep{grindlay87a}, or the efficiency of magnetic braking \\citep{ivanova06c}. So far, however, most studies of the LMXB-GC connection have been relatively small and so strong general conclusions are difficult to draw.  The properties of the GC populations in early-type galaxies themselves are also of considerable interest since they provide a unique insight into how such galaxies form. Although the numbers of GCs do not simply scale with the total stellar mass of a galaxy, the GC luminosity function (GCLF) is observed to be remarkably uniform \\citep[\\eg][]{harris91a}. Typically it can be well-fitted by a log-normal distribution, the absolute shape of which varies only weakly with the galaxy properties \\citep[\\eg][]{kundu01b,jordan06a}. The origin of this shape may be related to the dynamical destruction of low-mass GCs \\citep[\\eg][]{vesperini03a}. The stability of the GCLF peak (``turnover'') has led to its adoption as a standard distance measure for nearby galaxies \\citep{harris91a,jacoby92}, although there has been some debate as to its reliability \\citep[\\eg][]{ferrarese00a,kundu01b}. A recent study of early-type galaxies in Virgo (which span $\\sim$7~magnitudes in B) by \\citet{jordan07a} found that the width and, possibly, the turnover magnitude vary with the galaxy absolute magnitude, in the sense that the least-massive galaxies have narrower, fainter GCLFs. The authors did, however, find significant scatter about these relations for galaxies at a given magnitude.  The GC colour distributions of early-type galaxies are observed to be almost universally bimodal, \\citep[\\eg][]{zepf93a,gebhardt99a}, probably indicating ubiquitous metal-rich and metal-poor sub-populations \\citep[although see][]{richtler06a,yoon06a}. The GC distribution, in particular that of the metal-poor (\\ie\\ blue) population, is observed to be substantially more extended than the optical light \\citep{geisler96a}. The average colour of the GCs in a galaxy strongly correlates with the total galaxy magnitude \\citep{brodie91a}, reflecting an increasing fraction of red GCs in more massive galaxies coupled with a correlation between mean GC metallicity and galaxy mass for both blue and red clusters \\citep{peng06a}. There is, however, no correspondingly strong correlation between the colour and luminosity of the GCs themselves \\citep[\\eg][]{larsen01a,kundu01b}. A weak colour-magnitude relation has been observed, but only in  metal-poor GCs, in recent observations of nearby galaxies \\citep[][]{harris06a,strader06a,mieske06a}. These clues point to different origins for the two sub-populations, with the metal-poor clusters possibly forming in low-mass halos in the very early Universe and being accreted during hierarchical structure formation. In contrast, metal rich clusters may form during the gas-rich mergers which assemble the massive galaxy \\citep[][and references therein]{brodie06a}.  In this paper, we present a study of the GC populations in 19 early-type galaxies observed with \\hst\\ WFPC2. In particular, we present a photometric catalogue of GC candidates which can be directly compared to the properties of observed LMXBs. In a companion paper, \\citet[][hereafter \\lmxbpaper]{humphrey06b}, we present the \\chandra\\ data for these galaxies and investigate the relation between these populations, significantly expanding the sample of galaxies in which the LMXB-GC connection has been investigated. In the present work we compute self-consistently derived properties of the GC populations, such as the GC specific frequency, \\Sn, specific luminosity, \\Sl, and the GCLF turnover, \\mturn. The self-consistent computation of such parameters is essential for our analysis in Paper~I. The galaxies used in this study were chosen from the \\chandra\\ and \\hst\\ archives as having sufficiently deep ACIS and  WFPC2 data to enable a significant fraction of the LMXBs and GCs to be individually resolved. To enable a consistent comparison between the galaxies and to mitigate against the rapidly-degrading \\chandra\\ PSF off-axis, we focus here only on WFPC2 pointings of the centre of each galaxy. The galaxies, and the details of the archival observations used are shown in Table~\\ref{table_obs}. All errors quoted are 90\\% confidence regions, unless otherwise stated in the text. ",
        "conclusions": "In order to facilitate a direct comparison between LMXBs in early-type galaxies and possible GC candidates, we have presented a catalogue of GC candidates and key derived properties based on the \\hst\\ WFPC2 centrally-pointed observations of 19 nearby early-type galaxies. The complete source catalogue, including coordinates and photometry, is given in Appendix~\\ref{sect_sourcelist}. The total fraction of the GC light detected and GC specific luminosities are given for each galaxy in Table~\\ref{table_derived}. In \\lmxbpaper\\ we combine these data with archival \\chandra\\ observations of early-type galaxies to show that the majority of LMXBs must have formed in GCs.  We self-consistently derived various properties of the GC populations, including fits to the GCLFs. We report specific frequency, \\Sn, and specific luminosity, \\Sl, of the GC populations. Since the luminosity is dominated by clusters brighter than the turnover, we found that  \\Sl\\ is much less sensitive than \\Sn\\ to uncertainties in the GCLF turnover or width, making it a more robust indicator of the richness of the GC population where a significant fraction of the GCs are undetected due to source detection incompleteness. We find that the absolute magnitude of the GCLF turnover (\\Mturn) exhibits intrinsic scatter from galaxy to galaxy of $\\sim$0.3--0.4~mag, when compared to SBF distance moduli, consistent with the estimate of \\citet{ferrarese00a}. \\citet{kundu01b} argued the intrinsic uncertainty in the GCLF distance measure is $\\sim0.14$~mag, but only when using a small subset of the galaxies in their sample with the smallest error-bars on \\mturn. However, there is a significant correlation between (\\mturn-\\mturn$_{\\rm model}$) and the error on \\mturn\\ (as expected, since, the fainter the turnover, the less complete the data), and so this systematically selects against galaxies with faint \\Mturn, possibly reducing the scatter. Provided their error-estimates are reasonably representative, one can adopt the procedure outlined in \\S~\\ref{sect_mass_metallicity} to estimate the intrinsic scatter for the  25 galaxies in their full sample which overlap the  SBF sample of \\citet{tonry01}. Using this approach, we estimate an intrinsic scatter of $\\sim$0.38~mag (with a 90\\% lower limit of 0.23~mag), consistent with our results. \\citet{jordan06a} \\citep[see also][]{jordan07a} fitted Gaussian models to the GCLFs of early-type galaxies in Virgo, and found a trend of increasingly faint turnover magnitude with fainter $M_B$. For their galaxies with $M_B$\\ltsim -20, we estimate by eye an intrinsic scatter of $\\sim$0.3~mag, also consistent with the estimate from our data. To some extent, this scatter is driven by the presence of significant populations of diffuse star clusters in two systems; it may be that the presence of similar populations in the galaxies we consider here have contributed to the significant scatter we observe.  In keeping with past studies of GCs using V and I-band photometry we did not find any evidence of a significant correlation between colour and luminosity.  Converting the colour and magnitude of each GC to mass and metallicity, however, we found evidence of a weak mass-metallicity correlation for the GC populations as a whole. This trend cannot simply be attributed to observational biases, such as the systematic failure to detect faint (low-mass), red (metal-rich) GCs, since that would actually result in a stronger correlation in the colour-magnitude diagram. Similarly, our assumptions of a constant age for all GCs or a blue horizontal branch for all GCs do not appear to bias our results.  The apparent contradiction of no colour-magnitude correlation implying a mass-metallicity relation can be understood in terms of the weak dependence of the SSP mass-to-light ratio on the metallicity. To put this another way, if there was no correlation between mass and metallicity, due to the dependence of the mass-to-light ratio on metallicity, we would expect there to be a significant colour-magnitude correlation. To investigate this further, we  simulated an artificial dataset comprising a similar number of GCs to that observed. For each artificial GC, we randomly assigned a mass and metallicity value from, respectively, uncorrelated log-normal and normal distributions which approximately matched the observed data. Using the appropriate V and I-band mass-to-light ratios from \\citet{maraston05a}, these data were converted to V and I-band photometry, and estimated experimental noise was added. There was strong evidence of a weak correlation (the probability of no correlation was $<10^{-40}$) between colour and magnitude, with the fainter GCs being, on average redder. Conversely, we have also used similar simulations to generate colour-magnitude diagrams from artificial GC datasets which obey the observed mass-metallicity dependence. Exactly as expected, we find that there is only weak evidence of a correlation in the resulting data (prob of no correlation $\\sim$0.1\\%, as compared to $\\sim 10^{-37}$ in colour-metallicity space), demonstrating that the metallicity-dependent mass-to-light ratio can, indeed, destroy a colour-magnitude relation.   The galaxies in our sample span a range of magnitudes and so it is important to consider whether the known correlation between the galaxy mass and the peak of the GC colour distribution may affect our results. In part the relation reflects the increasing importance of red clusters in higher-mass galaxies \\citep[\\eg][]{peng06a} and it is difficult to imagine how that could introduce a spurious GC mass-metallicity correlation if one is not already present. More problematically though, \\citeauthor{peng06a} report a correlation between the mean metallicity of the {\\em blue} GCs and the galaxy mass and a similar trend for the red clusters. However, lower-mass galaxies are observed to have systematically narrower and fainter GCLFs \\citep{jordan07a}, which would be consistent with a trend for brighter GCs to be redder. In any case, for the magnitude range of the galaxies we consider here, the scatter about each correlation is at least as large as the trend.  After this work was submitted for publication, we became aware of a paper by \\citet{kundu08a} which suggests that an apparent mass-metallicity relation may arise due to observational biases. He proposes that a weak mass-radius relationship, coupled with differing errors in the photometric aperture correction (since larger clusters may be marginally resolved) in different wavebands may lead to a spurious correlation. To investigate whether this could plausibly explain our observations, we first estimated the error in the aperture correction for both V and I band by convolving King models with varying half light radii (\\rh) with PSF images (as described in \\S~\\ref{sect_reduction}). Since the vast majority of the GCs in our study were detected on the WF CCDs and in galaxies more distant than 15~Mpc, we computed the correction for our WF photometry only and adopted a 15~Mpc distance while allowing \\rh\\ to vary from 1 to 8~pc \\citep[consistent with observations, \\eg][]{spitler06a}. Although the size of the cluster can affect the photometry in either band by as much as $\\sim$0.3~mag, the effect was very similar in each filter so that V-I only changed by $\\ll 0.1$~mag. Such a small effect is not enough to wipe out the expected V-I {\\em versus} V correlation in the case that mass and metallicity are uncorrelated. We have verified this by adding a colour-size relation into our Monte-Carlo simulations detailed above. We adopted a pathological monotonic linear relation between GC radius and V-band magnitude, which is much stronger than the weak trend observed. For our simulated colour-magnitude diagram under the assumption of no mass-metallicity relation, we found that, although the mass-radius effect slightly reduces the resulting colour-magnitude correlation, the probability the data are consistent with no correlation is still $\\sim 10^{-32}$.   While our results depend to some degree on the adopted SSP models, it is clear that two independent sets of models--- those of \\citet{maraston05a} and those of \\citet{fioc97a} both imply a significant correlation between the mass and metallicity in these objects. The different normalizations and slopes for these models simply arise from differences in the predicted relations between the V and I-band mass-to-light ratios and the metallicity between the sets of models. However, {\\em any} prediction of a metallicity-dependent mass-to-light ratio will clearly give rise to a mass-metallicity correlation of this kind.  It is also interesting to compare our best-fitting mass-metallicity relation with that observed in early-type galaxies. We show in Fig~\\ref{fig_composite_cmd} an extrapolation of the best-fitting straight-line relation for early-type galaxies found by \\citet{thomas05a}. Intriguingly, we find a reasonable qualitative agreement, although the \\citeauthor{thomas05a} relation is slightly less steep and has a slightly higher normalization. The observed trend for GCs may arise from a number of effects, such as self-enrichment \\citep[\\eg][]{mieske06a}, in which the deeper potential wells of more massive GCs are more able to hold onto metals ejected from the very first generations of stars within them than in less-massive systems. Alternatively, the mass of a forming GC may be related to the mass (and metallicity) of its parent gas cloud \\citep[\\eg][]{harris06a}. Whatever processes actually underlie the mass-metallicity relation, the remarkable amount of intrinsic scatter (0.6~dex in metallicity) implies that stochastic processes actually dominate in determining the metallicity of an individual GC. Alternatively if there are, in fact, metal-rich and metal-poor sub-populations that cannot be distinguished in our data due to the statistical noise, and which exhibit different relations, the measured scatter will actually substantially overestimate the magnitude of such effects.    Recent (B,I)-band photometry, which allows better separation of metal-rich and metal-poor clusters, has revealed colour-magnitude correlations for the metal-poor (blue) sub-population of GCs in some early-type galaxies, but no clear trend for the red population. Transforming a composite colour-magnitude diagram of three giant Virgo ellipticals onto the mass-metallicity plane, \\citet{mieske06a} found little evidence of bimodality, but a trend of increasing average metallicity with mass, broadly consistent with our observed mass-metallicity dependence. Although bimodal colour distributions need not necessarily imply bimodal metallicity distributions \\citep[\\eg][]{richtler06a}, it is unclear whether the absence of bimodality in the mass-metallicity plane is simply an artefact of noise introduced during the transformation from colour-magnitude space. In particular, for our data, the error-bars on [Z/H] were typically so large that two putative distinct populations may have been blurred together. This is particularly relevant when we consider whether the ``blue tilt'' and, in particular, the lack of a colour-magnitude relation for the metal-rich clusters found by previous authors, are consistent with our data. To assess this, we simulated a set of artificial GC photometry data-points in the (g,z) photometric bands, with random z-band absolute magnitudes ranging from -7 to -11 and colours obeying the ``blue tilt'' relation found for blue GCs by \\citet{strader06a}. Transforming the data to the mass-metallicity plane using the models of \\citet{maraston05a}, the data were approximately consistent with a simple linear relation of the form [Z/H]=-4.43+0.50$\\log_{10}(M/M_\\odot)$. Performing the same analysis for the red GCs found by \\citeauthor{strader06a}, which we approximate as having g-z colours of 1.4, we found [Z/H]$\\simeq$-0.22. For a population of both blue and red GCs, we would expect the mean mass-metallicity relation to be simply a weighted average of these two relations. The existing best-fit slope and intercept are actually reproduced fairly well (0.25 and -2.3, respectively) for a (reasonable) blue GC fraction of $\\sim$50\\%. This excellent agreement with our results indicates that we may be observing the same trend as was observed in the (g,z) band, albeit the different trends for the red and blue populations are averaged together.   \\appendix"
    },
    "0809/0809.0521_arXiv.txt": {
        "abstract": "We have imaged 45 quasars from the Sloan Digital Sky Survey (SDSS) with redshifts $1.85 < z < 4.26$ in $JHK_s$ with the KPNO SQIID imager. By combining these data with optical magnitudes from the SDSS we have computed the restframe optical spectral indices of this sample and investigate their relation to quasar redshift.  We find a mean spectral index of $\\langle\\alpha_o\\rangle = -0.55\\pm0.42$ with a large spread in values. We also find possible evolution of the form $\\alpha_o = (0.148\\pm0.068)z - (0.964\\pm0.200)$ in the luminosity range $-28.0 < M_i < -26.5$. Such evolution suggests changes in the accretion process in quasars with time and is shown to have an effect on computed quasar luminosity functions. ",
        "introduction": "The recent publication of large samples of quasars from the Sloan Digital Sky Survey \\citep[SDSS;][]{sch07} and the Two Degree Field QSO Redshift Survey \\citep[2QZ;][]{cro04} allows the study of the evolution of the quasar population spanning long lookback times with large, homogenous catalogs of objects and the characterization of the quasar luminosity function (QLF). The \\citet{ric06a} QLF formed from the SDSS Third Data Release Quasar Catalog (DR3QLF) is defined from 15,343 quasars over 1622 deg$^2$, spanning redshifts from $z=0$ to $5$ and reaching luminosities down to $M_i < -23$ in their lowest redshift bins.  They find a peak in type 1 quasar activity between $z=2.2$ and $2.8$ and a general flattening of the slope of the bright-end QLF with increasing redshift. \\citet{hop07} compute a bolmetric QLF, combining many quasars surveys, including the DR3QLF and the 2QZ.  They find a peak in the quasar luminosity density at $z=2.15$ and an evolving QLF bright-end slope that becomes flatter at redshifts $z>3$.  One major difficulty in characterizing the evolution in quasar space densities is finding an appropriate way to calculate the absolute magnitudes of a survey sample and compare these magnitudes with other surveys, or even to objects within the same survey if it spans a large redshift range, i.e., how to calculate the appropriate {\\it K} correction.  Traditionally, absolute magnitudes were corrected back to the restframe $B$-band \\citep{sch83,boy87}, as most early surveys for quasars included the measurement of flux around $4400\\AA$, by assuming a power-law form for quasar spectra redward of Lyman-$\\alpha$ of the form $f_\\nu \\propto \\nu^{\\alpha}$ and assuming an average spectral index, $\\alpha$, for the survey sample.  As surveys move to higher redshifts, beyond $z=3$, the observed $B$-band ceases to sample the quasar continuum effectively, as the Lyman-$\\alpha$ line moves into and redward of that filter.  In the recent past, research groups have adopted several methods for computing absolute magnitudes and have referenced their magnitudes to different restframe wavelengths. For instance, \\citet{sch95} and \\citet{ken95} adopt an average spectral index of $\\alpha=-0.5$ in their computations of $M_B$, while \\citet{war94} avoid the adoption of a spectral index by computing the QLF as a function of $M_{C(1216)}$, the flux in the continuum under the Lyman-$\\alpha$ line, from their direct measurements of the flux at that point in their survey spectra.  More recently, \\citet{cro04} present their QLF from the 2QZ at $0.4 < z < 2.1$ in terms of $M_{b_J}$, and \\citet{ric06a} assume an average spectral index, but correct $i$-band magnitudes to $z=2$.   The correction to $z=2$ is prompted by the desire to minimize the effects of extrapolating the assumed quasar powerlaw over large wavelengths, as suggested by \\citet{wis00}.  Such differences in the methods for computing absolute magnitudes has led to some discrepancies in computed quasar space densities in past surveys \\citep{ken97}. The initial goals of this near-infrared (NIR) imaging program were to measure the restframe $B$-band flux for a set of quasars in several redshift ranges, to compare the differences in the absolute magnitudes, $M_B$, as computed from extrapolations of flux values at lower wavelengths to flux values measured directly, and to see how this changed with redshift.  This is equivalent to measuring the restframe optical spectral index of the quasars from optical and NIR photometry.  In this paper we report the initial results of a program to measure the restframe optical spectral index, $\\alpha_o$, of a subset of SDSS quasars at $1.85 < z < 4.26$ using their reported optical magnitudes and NIR photometry obtained at the KPNO 2.1m telescope using the Simultaneous Quad Infrared Imaging Device (SQIID). In \\S 2 we present the program design.  The data used in the project, both archived and new observations, are described in \\S 3. In \\S 4 we give the program results, including computations of spectral indices and the implications for the resulting absolute magnitudes.  We discuss our results further in \\S 5. ",
        "conclusions": "It has long been recognized that adopting an average spectral index for the power-law form of quasar spectral energy distributions can effect the evolution in the QLF \\citep[e.g.][]{gia92,fra96,ken97,wis98,ric06a}. However, the adoption of an average spectral index persists in most studies of the QLF, with $\\alpha=-0.5$ being the most common value used.  Recent work in the X-ray region has led to a general consensus that there is a dependence of quasar SED on luminosity \\citep[see][for an exception]{tan07}. As for dependence on $z$, some groups find no evidence for the evolution of X-ray spectral indices, $\\alpha_{ox}$, with redshift \\citep[e.g.][]{ste06}, while others do see a linear dependence of optical to X-ray spectral indices with $z$ \\citep{kel07}.  X-ray emission in AGN is generally taken to arise from a hot corona of optically thin gas heated by Compton scattering of thermal photons from a thick accretion disk, the likely source of the optical/UV emission.  While emission from these two regions is likely related, there is little evidence for a direct link between the X-ray and optical/UV emission, and \\citet{kel07} find no evidence for a correlation between $\\alpha_{UV}$ and $\\alpha_{ox}$.  Previous attempts to study the optical spectral energy distributions of quasars include \\citet{fra96} who find $\\alpha=-0.46\\pm0.30$ from a sample of LBQS quasars using optical and NIR photometry, \\citet{cri90} who compute {\\it K} corrections as a function of redshift in $UBV$ and find $\\alpha\\approx-0.7$ at $1000-5500\\AA$ from their composite quasar spectra, \\cite{van01} who construct a composite quasar spectrum from an SDSS quasar sample and find $\\alpha=-0.46$ for the region Ly-$\\alpha$ to H$\\beta$, and \\citet{pen03} who find $\\langle\\alpha\\rangle = -0.57 \\pm 0.33$ from optical and NIR photometry of 45 $z > 3.60$ SDSS quasars. There is considerable spread in the values of $\\alpha$ within each sample, and the distribution of $\\alpha$ found here and by the groups using optical and NIR photometry \\citep{fra96,pen03} are remarkably similar, each having a peak around -0.5 to -0.3 with a tail to the red that shifts the mean of the samples to steeper values.  This could be consistent with the findings of \\citet{web95} in their comparison of optical to NIR colors of radio quiet and radio loud quasars, who predict quasars should have an SED with $\\alpha=-0.3$ and that the distribution is caused by dust reddening in the host galaxies. Perhaps more important than the evolution of $\\alpha_o$ with $z$ is the distribution of the spectral index values.  \\citet{laf97} have pointed out that, if there is a spread in the spectral slope, there will be a corresponding slower luminosity evolution and steeper QLF's.  In general, the most desirable way to present the QLF is in terms of the bolometric luminosity. \\citet{ric06b} demonstrate that computing bolometric luminosities from optical luminosities assuming a single mean quasar SED can lead to errors as large as 50\\%. \\citet{hop07} have determined the bolometric QLF by combining the results of over two dozen quasar surveys from the hard X-ray to the mid-IR.  They construct a model SED but allow for the distribution in the power-law components of the model, stressing that there is no ``effective mean'' SED. They also adopt the luminosity dependent value of $\\alpha_{ox}$ of \\citet{ste06}.  However, they do not adopt a value of $\\alpha_{ox}$ that depends on redshift as has been suggested by \\citet{kel07} and is supported by our findings in the restframe optical.  Given the brightness of our sample, we have not attempted to correct for the contribution of the host galaxy to the SED.  The optical and NIR data were taken several years apart, and we have not considered variability in this sample. We have obtained NIR observations for $\\sim100$ more SDSS quasars in these redshift ranges, along with nearly simultaneous ($\\sim 1$ month) optical imaging with the aim of adding to our survey sample and addressing the issue of variability. \\citet{ric03} have stressed the need to correct for redshift dependent color effects when computing photometric spectral indexes, as failing to do so can lead to effects systematic with redshift. We have neglected the presence of emission lines in our photometric passbands.  However, due to the nature of the sample - each quasar is sampled at essentially the same points in their SED's as is clearly seen in Figure 5 - the apparent evolution would persist.  The generally accepted view of quasar activity is ascribed to the release of gravitational energy by accretion of matter on to a supermassive black hole.  The UV/optical flux is seen as arising from a thin, optically thick accretion disk \\citep[see][for a review]{kor99}, so a correlation between $\\alpha_{o}$ and $z$ implies evolution in the accretion process. More studies of the form and possible evolution of quasar SED's are obviously still needed to constrain theoretical models of AGN structure and energy production mechanisms. While ever larger samples of quasars are discovered \\citep[e.g.][]{sch07}, the form of the QLF cannot be fully characterized until we have a better understanding of quasar energy distributions and how they are affected by luminosity, redshift, and environment.  Since we are now coming to understand how quasar activity might be related to galaxy formation and evolution, quantifying the shape of the QLF and its evolution remains a pressing problem.  Here we have presented some evidence that quasar SED's evolve with cosmic time and have shown that this has a direct effect on the evolution of their luminosity function."
    },
    "0809/0809.0717_arXiv.txt": {
        "abstract": "Galaxy Zoo is the first study of nearby galaxies that contains reliable information about the spiral sense of rotation of galaxy arms for a sizeable number of galaxies.  We measure the correlation function of spin chirality (the sense in which galaxies appear to be spinning) of face-on spiral galaxies in angular, real and projected spaces.  Our results indicate a hint of positive correlation at separations less than $\\sim$ 0.5 Mpc at a statistical significance of 2-3 $\\sigma$. This is the first experimental evidence for chiral correlation of spins.  Within tidal torque theory it indicates that the inertia tensors of nearby galaxies are correlated. This is complementary to the studies of nearby spin axis correlations that probe the correlations of the tidal field.  Theoretical interpretation is made difficult by the small distances at which the correlations are detected, implying that substructure might play a significant role, and our necessary selection of face-on spiral galaxies, rather than a general volume-limited sample. ",
        "introduction": "Understanding the creation and evolution of the angular momentum of dark matter halos and galaxies is a crucial building block of a comprehensive theory of galaxy evolution. \\cite{hoyle49} was first to propose that the galaxy spin can be ascribed to the gravitational coupling with the surrounding galaxies. This idea has been formalised and extended in subsequent work \\citep{1969ApJ...155..393P,Dorosspin,1984ApJ...286...38W,1988MNRAS.232..339H,1996MNRAS.282..436C} into the modern theory of the evolution of galaxy spin, known as the tidal torque theory (TTT; see \\cite{2008arXiv0808.0203S} for a review). This theory asserts that protohalos acquire most of their angular momentum in the early stages of their formation, from the lowest non-vanishing contribution from the linear Lagrangian theory, that is, a coupling of the quadrupole of the local mass distribution to the external gravitational shear. Compared to $N$-body simulations, theory produces qualitatively correct results, although there are still significant discrepancies at a more quantitative level. Moreover, it seems that at present there are no clear theoretical directions for improving analytical models \\citep{1987ApJ...319..575B,Porciani:2001db,2005ApJ...627..647B}.   On the observational side, most of the work has been done using spiral galaxies. These are characterised by a rotating disk of baryonic matter. The line perpendicular to the plane of the disk determines the axis of rotation, while the spiral arms in most galaxies encode the sense of rotation, i.e the difference between left-hand screw and right-hand screw sense of rotation.  For spiral galaxies seen in projection, one can measure the observed galaxy ellipticities, which constrain the \\emph{axis} of the galaxy spin \\citep{2000ApJ...543L.107P,Lee:2001vz,Lee:2007jx,2006ApJ...640L.111T}.  This axis is known to within two-fold degeneracy associated with the tilt of the galactic plane with respect to the plane of the sky.  Since the vector can point in two directions on the same axis, the ellipticities constrain the spin vector within a total of four-fold degeneracy. Note that chiral information, \\emph{viz.} information about the actual directions of the spin vectors as opposed to spin axis, is completely absent in the study of galactic ellipticities.  However, this information contains important clues about the details of the emergence of the spin. As we will explain later in the text, the detection of chiral correlation function implies that the local inertia tensors must be correlated. This lends experimental support to the theoretical expectations that the inertia and gravitational shear tensor are correlated \\citep{Porciani:2001er}.    We now have a unique tool to study the chiral properties of galaxy spins.  Through an online project called Galaxy Zoo\\footnote{www.galaxyzoo.org} ~\\citep{2008arXiv0804.4483L}, members of the public have visually classified the morphologies and spin orientations for the entire spectroscopic sample of the Sloan Digital Sky Survey (SDSS from now on) \\citep{2000AJ....120.1579Y} Data Release 6 (DR6) \\citep{2008ApJS..175..297A}. The data and its reduction is extensively discussed in \\cite{2008arXiv0804.4483L}.  Spiral galaxies in the Galaxy Zoo sample are classified as clockwise, anti-clockwise or edge on. The spin direction convention used here is such that clockwise and anti-clockwise rotations correspond to galaxies whose arms are rotating in the sense of the letters Z and S respectively \\citep{1995MNRAS.276..327S}. For each face-on galaxy we thus receive one bit of information corresponding to the sign of the galaxy spin vector projected along the line of sight. It is important to note that this information is independent of the tilt of the plane of the galaxy.  We will refer to this one-bit information simply as galaxy spin. By the galaxy spin vector we mean the unit vector that defines the apparent spin of the galaxy: it is perpendicular to the disk plane and points in the direction the right turn screw would move if turned following the spiral arms inwards.  This quantity is strongly correlated with the real angular momentum of the gas. The correlation, however, is not perfect and observations show that the angular momentum vector of the gas points in the opposite direction in about four percent of systems \\citep{1982Ap&SS..86..215P}. In turn, there are theoretical expectations that there is a strong, but not perfect correlation between the angular momentum vector of gas and that of the dark matter halos hosting the galaxy \\citep{2002ApJ...576...21V}.  A detection of correlation in the galaxy spins would therefore imply a correlation in the dark matter spin vectors.  Conversely, a non-detection of the spin correlation can be used to put upper limits on the correlation between angular momentum vectors of dark matter halos.  This paper is structured as follows.  In Section \\ref{sec:TTT} we shortly review the tidal torque theory and its main results. Section \\ref{theory} will connect the correlation function $\\eta$ to an observable correlation function of spins $c$, while the Section \\ref{datamethod} will introduce our data and measurement technique. We present our results and discuss systematics in Section \\ref{sec:results}. Finally, we discuss our results and conclude in the last Section \\ref{sec:end}. ",
        "conclusions": "\\label{sec:end}   We are now in position to make a synthesis of our results.  The redshift-space results show that there is a significant correlation of $c(r)\\sim 0.15$ in the projected galaxy spins up to the $\\sim$ 0.5 $h$/Mpc. The angular correlation, shows larger correlations of $c(\\theta) \\sim 0.4$ that are significant correlations up to 30 arc-seconds, which roughly corresponds to projected distances of about 0.03 Mpc/$h$, since mean redshift of $0.08$ correspond to distance of $\\sim 230$ Mpc/$h$. This is consistent with the redshift-sample results - both exponential and Gaussian fits do predict $c(r)$ to raise to $\\sim$ 0.3--0.4 as $r$ goes to zero. A consistent picture is therefore the following: The angular sample detect correlations at the shortest distances, where majority of pairs are physical associations, but these get diluted at larger distances due to interlopers. The redshift-space correlations track these correlations to larger physical distances. Projected space pairs do not have enough signal-to-noise to detect these correlations at high significance.    How do these results compare to theoretical predictions? Simple models as those suggested in \\cite{2000ApJ...543L.107P} (equations (\\ref{eq:pry}) and (\\ref{eq:prx})) predict a vanishing $\\eta$ and hence we have directly detected a non-random distribution of inertia tensors. Within the standard model, the reason for correlations of moments of intertia are the correlations of these with the (slowly varying) tidal field. On the other hand, if moments of intertia are \\emph{perfectly} aligned with the tidal field, the tidal torque cannot produce any angular momentum and therefore the resulting angular momentum is due to the residual 10\\% of misaligment \\citep{Porciani:2001er}. The stunning outcome of our result, if confirmed, is that even these 10\\% misaligments are correlated from (sub-)halo to (sub-)halo.  What is also interesting is, however, that in \\cite{Porciani:2001db} $\\eta$ correlations have not been detected in simulations at $z=0$ at all separations. In particular, $\\eta<0.02$ at $r=0.5$ Mpc/$h$. A virtually identical result has been found by \\cite{2005ApJ...627..647B}, who also find $\\eta<0.02$ at $r=0.5$Mpc/$h$ (our $\\eta$ is their $\\xi_{LL}$). This is in tension with our results even after conversion factors in Equation (\\ref{eq:conv}) are taken into account. There are many reasons that explain why our results are not directly comparable to the above work. First, they are comparing individual dark matter halos. In our case, we see the signal at pair separations of less than 1 Mpc. At such distances, one-halo pairs (pairs of galaxies that reside in the same dark matter halo), dominate over two halo pairs (pairs in which two galaxies occupy two different halos).  By selecting spiral galaxies, we are essentially selecting pairs that are composed of satellites residing in the same halo, rather than pairs compromised of central halo galaxies. The latter are bright ellipticals and hence inaccessible using our method. Unfortunately, not very much theoretical work has been done for spin correlations of substructure. The most relevant paper in the literature is \\cite{2005ApJ...629L...5L}, which, however, still uses the chirality agnostic model of \\cite{2000ApJ...543L.107P} and does not calculate the chiral correlation function. More work on the theoretical side and $N$-body side is required to understand the implications of our results. Hopefully, the results could be turned around and help us understand what kind of substructure spiral galaxies occupy in a typical dark matter halo.  It is therefore imperative that our observational protocol is simulated on a large enough $N$-body simulation, for example the \\emph{Millennium Simulation} \\citep{Springel:2005nw} or \\emph{MareNostrum Universe} \\citep{2007ApJ...664..117G} simulations. There, halos and sub-halos hosting spiral galaxies can be identified and those, whose inclination with respect to a given observer is small enough to be considered face-on, should be correlated. This would  result in a quantity $c(r)$ that is directly comparable to the observables that we constrain with the Galaxy Zoo data.  Another interesting aspect of our results is that, for spiral galaxies, we essentially exclude large and random misaligments between gas and dark matter angular momenta. Since the dark matter is dynamically dominant, gas angular momenta can only be correlated if they are so due to correlations between dark matter.   Finally, it is tempting to combine our measurements with the ellipticity measurements to improve signal-to-noise and remove some systematic. Note, however, that our 1-bit signal divides a four-fold degeneracy into a two-fold one and thus this is a non-trivial task which will be left for the future.   To conclude: We have tentatively detected a chiral correlation function in the spins of spiral galaxies. This correlation function vanishes in the simplest models based on tidal torque theory. Our results indicate, that moments of intertia of protohalos that end up hosting spiral galaxies are correlated at distances less than $\\sim$ 0.5 Mpc/$h$.These short distances imply that these protohalos are often likely to be substructures of massive halos.  More work is required to understand these results at a quantitative level."
    },
    "0809/0809.4628_arXiv.txt": {
        "abstract": "\\noindent Scalar field dynamics may give rise to a nonzero cosmological variation of fundamental constants. Within different scenarios based on the unification of gauge couplings, the various claimed observations and bounds may be combined in order to trace or restrict the time history of the couplings and masses. If the scalar field is responsible for a dynamical dark energy or quintessence, cosmological information becomes available for its time evolution. Combining this information with the time variation of couplings, one can determine the interaction strength between the scalar and atoms, which may be observed by tests of the Weak Equivalence Principle. We compute bounds on the present rate of coupling variation from experiments testing the differential accelerations for bodies with equal mass and different composition and compare the sensitivity of various methods. In particular, we discuss two specific models of scalar evolution: crossover quintessence and growing neutrino models. ",
        "introduction": "Several positive and null results on the time variation of fundamental constants have been reported. In a previous paper \\cite{Part1} we reviewed the relevant observational evidence and bounds, and introduced several distinct scenarios with unification of gauge couplings (GUT) \\cite{Marciano}. Each scenario is characterized by specific relations between the variations of different Standard Model (SM) couplings: specifically, all fractional variations in SM couplings were taken to be proportional to the variation of one underlying unified coupling or a fixed linear combination of several unified couplings. The different scenarios have different consequences for the interpretation of any signal of nonzero variation. We also divided the observations into several cosmological ``epochs'' and discussed the consistency between them within each epoch, as well as whether a monotonic time evolution between epochs could fit the observations. Up to this point no particular dynamical mechanism for the coupling variations was considered.  In quantum field theory, a possible time variation of couplings must be associated to the time variation of a field. This field may describe a new ``fundamental particle'' or stand for the expectation value of some composite operator. In this paper we discuss the consequences of the simplest hypothesis, namely that the field is a scalar, such that its time varying expectation value preserves rotation and translation symmetry locally, as well as all gauge symmetries of the Standard Model.  Any observation of coupling variations in the cosmological history from nucleosynthesis up to now would imply that the expectation value of this scalar field changes appreciably during this cosmological epoch. Such a ``late time evolution'' (in contrast to the electroweak or QCD phase transition, GUT-phase transition or even inflation) is characteristic for models with a dynamical dark energy or quintessence. Indeed, the potential and kinetic energy of the scalar ``cosmon'' field would contribute a homogeneous and isotropic energy density in the Universe, and therefore lead to dynamical dark energy \\cite{Wetterich88}. The time variation of ``fundamental couplings'' is a generic prediction of such models \\cite{Wetterich:1987fk}. In fact, it has to be explained why the relative variation is tiny during a Hubble time, while the simplest models for a scalar cosmon field arising from superstrings or other unified frameworks typically lead to relative variations of the order one. Efficient mechanisms have been proposed for explaining the smallness of variations, for example based on the approach to a fixed point \\cite{Wetterich:2002wm,Wetterich:2008sx}. Unfortunately, the strength of such mechanisms is not known, such that the overall strength of the scalar coupling to matter remains theoretically undetermined. Several investigations have explored the consequence of scalar quintessence models for the time variation of couplings, as well as models specifically designed to account for coupling variations \\cite{Gasperini:2001pc, Dvali:2001dd, Olive:2001vz, Sandvik:2001rv, Chiba:2001er, Barrow:2002ed, Wetterich:2002ic, WetterichCrossoverQ, Farrar:2003uw, Parkinson:2003kf, Doran:2004ek, Lee:2004vm, Wetterich:1994bg}.  The cosmon coupling to atoms determines both the outcome of tests of the Weak Equivalence Principle and the time variation of couplings in the most recent cosmological epoch, including the present. For the latter, one also needs information for the rate of change of the expectation value of the scalar field, which can be expressed in terms of cosmological observables, namely the fraction in dark energy, $\\Omega_h$, and its equation of state, $w_h$. In consequence, the differential acceleration $\\eta$ between two bodies of different composition, and therefore different specific ``cosmon charge'', can be related to the present time variation of couplings and cosmological parameters as \\cite{Wetterich:2002ic} \\beq \\label{eq:DifferentialAccelerationParam} \\eta \\simeq 3.8 \\times 10^{-12} \\left( \\frac{\\dot{\\alpha}/\\alpha}{10^{-15}{\\rm y}^{-1}} \\right)^2 \\frac{ F }{\\Omega_h (1 + w_h)}\\, . \\eeq Here the present dark energy fraction is $\\Omega_h \\approx 0.73$ and $\\dot{\\alpha}/\\alpha$ is the present relative variation of the fine structure constant. % The ``unification factor'' $F$ encodes the dependence on the precise unification scenario (defined in Section \\ref{sec:unif}) and on the composition of the test bodies. In the notation of \\cite{Wetterich:2002ic} we have $F = \\Delta R_Z (1+\\tilde{Q})$, where $\\Delta R_Z = \\Delta (Z/(Z+N))$ is the difference in the proton fraction $Z/A$ between the two bodies, which takes a value $\\Delta R_Z \\lesssim 0.05$ for typical experimental tests of the equivalence principle. The present dark energy equation of state $w_h$ is close to $-1$; we take $1 + w_h \\lesssim 0.1$. The factor $F$ will be calculated for our different unified scenarios in Section~\\ref{sec:UnificationFactor}: for typical test mass compositions we find $1 \\leq F\\leq \\rm{few}\\times 10^2$.  Once $F$ is fixed, the relation \\eqref{eq:DifferentialAccelerationParam} allows for a direct comparison between the sensitivity of measurements of $\\eta$ versus the measurements of $\\dot{\\alpha}/\\alpha$ from laboratory experiments, or bounds from recent cosmological history, such as from the Oklo natural reactor or the isotopic composition of meteorites.  Our paper is organized as follows. After a brief summary of our previous results concerning evolution factors for different unified scenarios, we introduce in Section~\\ref{sec:Scalar} a cosmologically evolving scalar field as the source of the dynamics of a possible coupling variation. The scalar has both an effect on large-scale cosmological evolution as a consequence of its homogeneous potential and kinetic energy, and may also produce local gravitational effects, violating the Weak Equivalence Principle (WEP) due to its interactions with matter. In Sec.~\\ref{sec:WEP} we use both aspects to set further constraints on a possible time variation of couplings. We determine, for each unified scenario, which observational methods are most sensitive to detect late-time or present-day variations. This may enable us to distinguish between scenarios if there is a non-zero signal.  In Section~\\ref{sec:QuintessenceModels} we determine for each scenario the evolution factors that a smooth time evolution of the scalar must satisfy to be consistent with observations, in particular with a detectable nonzero variation at $z\\sim 1\\,$--$\\,3$. We discuss whether a simple model of monotonically varying ``crossover quintessence'' can produce the required behaviour. Section~\\ref{sec:Growing} is devoted to a recently proposed type of model \\cite{Amendola_grownu,Wett_grownu} where the scalar field evolution is halted by its coupling to neutrinos whose mass increases. Such models solve the ``coincidence problem'' of dark energy and give rise to an interesting scalar field evolution with damped oscillations at recent epochs. These models need not obey bounds derived for a monotonic time variation, and we use instead global fits over all data for some specific choices of cosmological model. In Section~\\ref{sec:Conclude} we summarize and discuss our findings. ",
        "conclusions": "\\label{sec:Conclude} This paper demonstrates how a clear observation of time variation of fundamental couplings would not only rule out a constant dark energy, but also put important constraints on the time evolution of a dynamical dark energy or quintessence.  In Sec.~\\ref{sec:WEP} we have seen how the comparison of a varying coupling with the bounds from tests of the equivalence principle can put a lower nonzero bound on the combination $\\Omega_h (1+w_h)$, according to \\beq \\label{eq:BoundOnOmegaw} \\Omega_h (1+w_h) \\gtrsim 3.8 \\times 10^{18} F (\\partial_t \\ln \\alpha)^2 \\eta_{\\rm max}^{-1} \\, , \\eeq with $\\partial_t \\ln \\alpha$ in units of y$^{-1}$, and where $\\eta_{\\rm max} \\simeq 1.8 \\times 10^{-13}$ is the current experimental limit on the differential acceleration of two test bodies of different composition. Thus, if $|\\partial_t \\ln \\alpha |$ is nonzero and not too small, $w_h$ {\\em cannot}\\/ be arbitrarily close to $-1$ (a cosmological constant); nor can the contribution of the scalar to the dark energy density be insignificant.  The precise bound depends on the ``unification factor'' F, which differs between different scenarios of unification and also depends on the composition of the experimental test bodies. We find $1\\leq F \\leq \\rm{few}\\times 10^{2}$ for the Be-Ti masses used for the best current bound on $\\eta$ and for representative cases of unified scenario.  Conversely, assuming that the scalar field responsible for the varying couplings is the cosmon, whose potential and kinetic energy account for a dynamical dark energy in the Universe, the bounds on $\\eta$ can be used to set bounds on the time variation of couplings in the present epoch. For the different unified scenarios we compare the relative sensitivity of these bounds as compared to laboratory measurements or bounds from the Oklo natural reactor and the composition of meteorites in Table~\\ref{tab:PresentVariationErrors}.  For a given unified scenario, the bounds on the time variation of various couplings in different cosmological epochs strongly restrict the possible time evolution of the cosmon field, once at least one irrefutable observation of some coupling variation at some redshift becomes available. We have demonstrated this by an analysis that implicitly assumes a nonzero variation, considering both general features and specific quintessence models. We are aware that the actual values for the evolution factors $l(z)$ from this analysis may be premature, since the observational situation is unclear and on moving grounds. For example, taking the recent reanalysis of the variation of the proton to electron mass ratio $\\mu$ in Ref.~\\cite{King:2008ud} instead of the results in Ref.~\\cite{Reinhold:2006zn} used in this paper, would strongly influence the values of the evolution factors. We have demonstrated this in a somewhat different way by investigating the change in the evolution factors if some claimed observations of varying couplings are omitted. Needless to say that without a clear signal any analysis may become obsolete. For the time being our analysis remains a useful tool for the comparison of the relative sensitivity of different experiments and observations, and for a judgement about mutual consistency of different claimed variations and bounds from other observations.   \\subsection*{Acknowledgements} We acknowledge useful discussions with M.~Doran, G.~Robbers, M.~Pospelov, J.~Donoghue and P.~Avelino, and correspondence with V.~Flambaum. T.\\,D. thanks the Perimeter Institute for hospitality during the workshop ``Search for Variations in Fundamental Couplings and Mass Scales''. T.\\,D. is supported by the {\\em Impuls- and Vernetzungsfond der Helmholtz-Gesellschaft}."
    },
    "0809/0809.4302_arXiv.txt": {
        "abstract": "Text{% The zonal flow in Jupiter's upper troposphere is organized into alternating retrograde and prograde jets, with a prograde (superrotating) jet at the equator. Existing models posit as the driver of the flow either differential radiative heating of the atmosphere or intrinsic heat fluxes emanating from the deep interior; however, they do not reproduce all large-scale features of Jupiter's jets and thermal structure. Here it is shown that the difficulties in accounting for Jupiter's jets and thermal structure resolve if the effects of differential radiative heating and intrinsic heat fluxes are considered together, and if upper-tropospheric dynamics are linked to a magnetohydrodynamic (MHD) drag that acts deep in the atmosphere and affects the zonal flow away from but not near the equator. Baroclinic eddies generated by differential radiative heating can account for the off-equatorial jets; meridionally propagating equatorial Rossby waves generated by intrinsic convective heat fluxes can account for the equatorial superrotation. The zonal flow extends deeply into the atmosphere, with its speed changing with depth, away from the equator up to depths at which the MHD drag acts. The theory is supported by simulations with an energetically consistent general circulation model of Jupiter's outer atmosphere. A simulation that incorporates differential radiative heating and intrinsic heat fluxes reproduces Jupiter's observed jets and thermal structure and makes testable predictions about as-yet unobserved aspects thereof. A control simulation that incorporates only differential radiative heating but not intrinsic heat fluxes produces off-equatorial jets but no equatorial superrotation; another control simulation that incorporates only intrinsic heat fluxes but not differential radiative heating produces equatorial superrotation but no off-equatorial jets. The proposed mechanisms for the formation of jets and equatorial superrotation likely act in the atmospheres of all giant planets.} ",
        "introduction": "The zonal flow in Jupiter's upper troposphere has been inferred by tracking cloud features, which move with the horizontal flow in the layer between about 0.5 and 1~bar atmospheric pressure \\citep{Ingersoll04,West04,Vasavada05}. In this layer, the zonal flow is organized into a strong prograde (superrotating) equatorial jet and an alternating sequence of retrograde and prograde off-equatorial jets (Fig.~\\ref{f:winds}a). This flow pattern has been stable at least between the observations by the Voyager and Cassini spacecrafts in 1979 and 2000, with some variations in jet speeds, for example, a slowing of the prograde jet at $21^\\circ$N planetocentric latitude by $\\about 40\\,\\mathrm{m\\,s^{-1}}$ \\citep{Porco03,Ingersoll04}. The zonal flow in layers above the clouds has been inferred from the thermal structure of the atmosphere, using the thermal wind relation between meridional temperature gradients and vertical shears of the zonal flow. Meridional temperature gradients and thus vertical shears between 0.1 and 0.5~bar are generally small: meridional temperature contrasts along isobars do not exceed $\\about 10\\,\\mathrm{K}$ \\citep{Conrath98,Simon-Miller06}. The thermal stratification in the same layer is statically stable \\citep{Simon-Miller06,Read06}.  About the zonal flow in lower layers, it is only known that at one site at $6.4^\\circ$N planetocentric latitude, where the Galileo probe descended into Jupiter's atmosphere, it is prograde and increases with depth from $\\about 90\\,\\mathrm{m\\,s^{-1}}$ at 0.7~bar to $\\about 170\\,\\mathrm{m\\,s^{-1}}$ at 4~bar; beneath, it is relatively constant up to at least $\\about 20$~bar \\citep{Atkinson98}. At the same site, the thermal stratification is statically stable but approaches neutrality with increasing depth between 0.5 and 1.7~bar; beneath, it is statically nearly neutral or neutral up to at least $\\about 20$~bar \\citep{Magalhaes02}. These are the large-scale features (if the Galileo probe data are representative of large scales) that a minimal model of Jupiter's general circulation should be able to reproduce. \\begin{figure}[!htb] \\centering\\includegraphics{figure1} \\caption{Zonal flow in Jupiter's upper troposphere and in simulations. (a) Zonal velocity on Jupiter, inferred by tracking cloud features from the Cassini spacecraft \\citep{Porco03} (orange) and in Jupiter simulation at 0.65~bar (blue). (b) Zonal velocity at 0.65~bar in control simulations: with intrinsic heat fluxes but with uniform insolation at the top of the atmosphere (magenta), and with differential insolation but without intrinsic heat fluxes (light blue).  Zonal velocities in simulations are zonal and temporal means in statistically steady states (over 1500 days for Jupiter simulation and over 900 days for control simulations); differences between the (statistically identical) hemispheres here and in subsequent figures are indicative of sampling variability. For Jupiter, latitude here and throughout this paper is planetocentric; the simulated planets are spherical, so planetocentric and planetographic latitudes are identical.} \\label{f:winds} \\end{figure}  There are two plausible energy sources for Jupiter's general circulation. First, $\\about 8\\,\\mathrm{W\\, m^{-2}}$ of solar radiation are absorbed in Jupiter's atmosphere \\citep{Hanel81}, where the time-mean insolation at the top of the atmosphere, given Jupiter's small obliquity of $3^\\circ$, varies approximately with the cosine of latitude. Second, $\\about 6\\,\\mathrm{W\\, m^{-2}}$ of intrinsic heat fluxes emanate from Jupiter's deep interior \\citep{Ingersoll04,Guillot04,Guillot05}; observations of convective storms \\citep{Gierasch00,Porco03,Sanchez-Lavega08} and the neutral or nearly neutral thermal stratification along much of the Galileo probe descent path show that the intrinsic heat fluxes are at least partially convective. Existing models of Jupiter's zonal flow posit as the driver of the general circulation either differential radiative heating of the atmosphere or intrinsic convective heat fluxes \\citep[e.g.,][]{Busse76,Busse94,Williams79,Ingersoll04,Vasavada05}; however, they do not reproduce all large-scale features of the jets and thermal structure. For example, in parameter regimes relevant for Jupiter, they generally do not produce equatorial superrotation, unless artifices are employed such as imposing an additional heat source near the equator \\citep{Williams03a} or assuming excessively viscous flow and permitting intrinsic heat fluxes several orders of magnitude stronger than Jupiter's \\citep{Heimpel05,Heimpel07}.  Here we show that the difficulties in accounting for Jupiter's jets and thermal structure resolve if the effects of differential radiative heating and intrinsic heat fluxes are considered together, and if upper-tropospheric dynamics are linked to drag that acts deep in the atmosphere and affects the zonal flow away from but not near the equator. The key is to distinguish the different ways in which eddies can be generated near the equator and away from it, and to consider their role and the role of drag at depth in the balance of angular momentum (the angular momentum component about the planet's spin axis).  First we describe how eddies can be generated near the equator and away from it, how convectively generated equatorial waves can lead to equatorial superrotation, and how the angular momentum balance of the upper troposphere is linked to drag and constrains the flow at depth (sections~\\ref{s:am_fluxes}--\\ref{s:deep_flow}). Then we use simulations with a three-dimensional general circulation model (GCM) of a thin shell in Jupiter's outer atmosphere, with an idealized representation of effects of drag deep in the atmosphere, to demonstrate that the mechanisms proposed can account for Jupiter's observed jets and thermal structure (sections~\\ref{s:gcm}--\\ref{s:controls}). ",
        "conclusions": "Based on the theory and simulations and consistent with the available observational data, we propose that baroclinic eddies generated by differential radiative heating are responsible for Jupiter's off-equatorial jets, and that Rossby waves generated by intrinsic convective heat fluxes are responsible for the equatorial superrotation. Mean meridional circulations adjust entropy gradients and the zonal flow in lower layers of Jupiter's atmosphere, such that the zonal flow is in thermal wind balance with the entropy gradients, and convergence or divergence of angular momentum fluxes in the upper troposphere and the MHD drag on the zonal flow at depth balance upon averaging over cylinders concentric with the planet's spin axis. As demonstrated by the Jupiter simulation, the resulting view of how the zonal flow and general circulation are generated and maintained is consistent with observed large-scale features of Jupiter's jets and thermal structure, such as the zonal flow and meridional temperature variations in the upper troposphere and the thermal stratification of the upper troposphere and layers beneath. It is also consistent with the observed eddy angular momentum fluxes and with energetic constraints indicating that these fluxes are confined to a relatively thin atmospheric layer. As demonstrated by the control simulations, differential radiative heating alone can account for the off-equatorial jets, and intrinsic convective heat fluxes can account for the prograde equatorial jet. However, intrinsic convective heat fluxes alone do not necessarily lead to formation of off-equatorial jets.  The theory and simulations predict aspects of the general circulation that have not been observed but that are or soon will be observable. For example, the transition in energy-containing eddy scale between the equatorial region and regions away from the equator (Fig.~\\ref{f:eddy_scales}), pointing to different mechanisms of eddy generation, should be observable by tracking cloud features. And we predict that the measurements of NASA's upcoming Juno mission to Jupiter will be consistent with zonal jets that extend deeply into the atmosphere at all latitudes, away from the equator up to depths at which the MHD drag acts. As already observed in the upper troposphere, we expect that also at lower layers the speeds of the prograde off-equatorial jets decrease with depth, and that there are associated equatorward temperature gradients along isobars: qualitatively as in Fig.~\\ref{f:vertical_structure}a and b, but likely not quantitatively so because the jets are expected to extend to much greater depth than in our simulations and thus likely have weaker shear. Depending on the strength of the MHD drag and on the depth at which it acts, the speeds of the retrograde off-equatorial jets may decrease or increase with depth, with weaker shear than the prograde jets, and with associated poleward (reversed) or weaker equatorward temperature gradients. If the shear of the zonal flow can be inferred from measurements, it can be used together with observations of the eddy angular momentum transport and with the implied transfer of kinetic energy from eddies to the mean flow to constrain the strength and depth of the MHD drag and thus to elucidate dynamics of the deep atmosphere that are not amenable to direct measurement.  The proposed mechanisms are generic and likely act in the atmospheres of all giant planets. They suggest, for example, that the reason that Saturn's prograde equatorial jet is wider and stronger than Jupiter's may be that Saturn's tropospheric gravity wave speed and equatorial Rossby radius are greater. The greater depth at which MHD drag on Saturn is estimated to act implies a wider region over which there is effectively no drag on the zonal flow \\citep{Liu08}, thus making a wider equatorial jet possible. The proposed mechanisms also suggest that the reason Uranus and Neptune do not exhibit equatorial superrotation may be that their intrinsic heat fluxes are not sufficiently strong to lead to convection penetrating into the upper troposphere.  \\paragraph"
    },
    "0809/0809.4134_arXiv.txt": {
        "abstract": "In this letter we present a new N-flation model constructed by making use of multiple scalar fields which are being described by their own DBI action. We show that the dependence of the e-folding number and of the curvature perturbation on the number of fields changes compared with the normal N-flation model. Our model is also quite different from the usual DBI N-flation which is still based on one DBI action but involves many moduli components. Some specific examples of our model have been analyzed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.3839_arXiv.txt": {
        "abstract": "\\fontsize{10}{10.6}\\selectfont Circularly polarized 3.5 cm continuum emission was detected toward three radio sources in the R CrA region using the Very Large Array. The Class I protostar IRS 5b persistently showed polarized radio emission with a constant helicity over 8 yr, which suggests that its magnetosphere has a stable configuration. There is a good correlation between the Stokes $I$ and Stokes $V$ fluxes, and the fractional polarization is about 0.17. During active phases the fractional polarization is a weakly decreasing function of Stokes $I$ flux, which suggests that IRS 5b is phenomenologically similar to other types of flare stars such as RS CVn binaries. The variability timescale of the polarized flux is about a month, and the magnetosphere of IRS 5b must be very large in size. The Class I protostar IRS 7A was detected once in circularly polarized radio emission, even though IRS 7A drives a thermal radio jet. This detection implies that the radio emission from the magnetosphere of a young protostar can escape the absorption by the partially ionized wind at least once in a while. The properties of IRS 7A and IRS 5b suggests that Class I protostars have organized peristellar magnetic fields of a few kilogauss and that the detectability of magnetospheric emission may depend on the evolutionary status of protostar. Also reported is the detection of circularly polarized radio emission toward the variable radio source B5. ",
        "introduction": "Magnetic fields play an important role in star formation (Andr{\\'e} 1996; Feigelson et al. 2007; Bouvier et al. 2007). This is especially true in the protostellar accretion/evolution processes, and peristellar magnetic fields are expected or even required in many theoretical models: Protostellar outflows are generated through magnetocentrifugal ejection mechanism (e.g., Pudritz et al. 2007). The magnetorotational instability may provide the viscosity required to explain the dynamics of accretion disks (Balbus \\& Hawley 1991). Magnetospheres may force the protostellar spin to be coupled with the rotation of circumstellar disk, and the inner edge of accretion disks may be truncated by the magnetosphere near the corotation radius (Camenzind 1990; K{\\\"o}nigl 1991). The transfer of material from the accretion disk to the protostar may be done through magnetospheric accretion (Bouvier et al. 2007). Despite all these expectations, observational evidence of peristellar magnetic fields of protostars is very difficult to come by.  Much of our knowledge about the peristellar magnetic fields of young stellar objects comes from the observations of T Tauri stars, especially radio continuum observations at centimeter wavelengths (Andr{\\'e} et al. 1992; Andr{\\'e} 1996). The key characteristics are the variability and circular polarization of emission from nonthermal electrons. The variability timescale is usually hours to days, much longer than solar-type flares, which implies that large-scale magnetic fields are involved. Moderate degrees of circular polarization were observed, suggesting that the emission mechanism is gyrosynchrotron radiation from mildly relativistic electrons. The detection of circular polarization at $\\sim$5 GHz implies a dipole-like large-scale magnetosphere with a surface field strength of $\\sim$1 kG (G{\\\"u}del 2002).  X-ray surveys of pre-main-sequence stars suggest that the level of activity decays with stellar age (Preibisch \\& Feigelson 2005), and we expect that Class I protostars may be magnetically more active than T Tauri stars. However, the free-free absorption by thermal wind is supposed to obscure the magnetosphere of protostars in most cases (Andr{\\'e} et al. 1992; Andr{\\'e} 1996). The youngest object detected in circularly polarized radio emission is IRS 5 in the R CrA region (Feigelson et al. 1998). In fact, IRS 5 is the only known protostellar object persistently displaying circularly polarized emission (Forbrich et al. 2006; Miettinen et al. 2008). Therefore, IRS 5 provides a wonderful opportunity to study the magnetic activity of protostars.  To obtain high-quality images of the R CrA region, we observed in the centimeter continuum with the Very Large Array (VLA) of the National Radio Astronomy Observatory. The main results were presented in Paper I (Choi et al. 2008). In this paper, we present the polarimetry of the R CrA region. We describe the data in \\S~2. In \\S~3 we briefly report the results of the polarimetry. In \\S~4 we discuss the star forming activity of the radio sources showing circularly polarized centimeter continuum emission. A summary is given in \\S~5. ",
        "conclusions": ""
    },
    "0809/0809.3652_arXiv.txt": {
        "abstract": "The  Serpens cloud  has received  considerable attention  in  the last years, in particular the small  region known as the Serpens cloud core where a plethora  of star formation related phenomena  are found. This review summarizes  our current  observational knowledge of  the cloud, with emphasis on the core. Recent results are converging to a distance for  the cloud  of  $\\sim 230  \\pm 20$  pc,  an issue  which has  been controversial over the  years. We present the gas  and dust properties of  the  cloud core  and  describe  its  structure and  appearance  at different wavelengths. The core contains a dense, very young, low mass stellar cluster with more than 300 objects in all evolutionary phases, from collapsing  gaseous condensations to pre-main  sequence stars. We describe  the  behaviour and  spatial  distribution  of the  different stellar  populations (mm  cores, Classes  0,  I and  II sources).  The spatial concentration and the fraction number of Class 0/Class I/Class II  sources is considerably  larger in  the Serpens  core than  in any other  low  mass star  formation  region,  e.g.  Taurus, Ophiuchus  or Chamaeleon, as also stated  in different works. Appropriate references for coordinates and fluxes of  all Serpens objects are given. However, we provide  for the first time  a unified list of  all near-IR sources which have  up to now been  identified as members of  the Serpens core cluster; this list includes some  members identified in this review. A cross-reference  table of  the near-IR  objects with  optical, mid-IR, submillimeter, radio  continuum and X-ray  surces is also  provided. A simple analysis has allowed us to identify a sample of $\\sim 60$ brown dwarf  candidates among  the 252  near-IR objects;  some of  them show near-IR excesses and, therefore,  they constitute an attractive sample to study very young substellar  objects. The review also refers to the outflows associated with the young  sources. A section is dedicated to the  relatively small  amount  of works  carried  out towards  Serpens regions outside the core. In particular, we refer to ISO and to recent Spitzer  data.  These  results  reveal  new  centers  of  active  star formation in the Serpens cloud and the presence of new young clusters, which deserve  follow-up observations  and studies to  determine their characteristics  and  nature in  detail.  Finally,  we  give a  short, non-exhaustive list of individually interesting Serpens objects. ",
        "introduction": "\\label{introduction} A considerable amount  of work has been dedicated  to the region known as the  Serpens dark cloud  (Galactic coordinates $l^{II} =  32\\deg, ~ b^{II} = 5\\deg$)  since it was recognized by Strom et  al. (1974) as a site  of active  star  formation. The  cloud  extends several  degrees around the  young variable  star VV  Ser and forms  part of  the large local  dark cloud  complex  called  the Aquila  Rift,  which has  been extensively mapped  in several molecular  line surveys (e.g.   Dame \\& Thaddeus 1985,  Dame et al.  1987,  2001, see the chapter  by Prato et al.). A  large scale extinction  map has been presented  by Cambr\\'esy (1999) and Dobashi et  al.  (2005). Serpens is seen  on optical images as  a  large area  irregularly  obscured  by  large amounts  of  dust; Fig. \\ref{Serpens} reproduces the DSS  $R$ band photograph of a region $2\\deg  \\times 2\\deg$  in  size. Several  reflection nebulosities  are distinguished, the most  prominent of which are S  68 (Sharpless 1959, see  also Dorschner  \\& G\\\"urtler  1963,  van den  Bergh 1966,  Bernes 1977), illuminated  by HD 170634,  and a very red  bipolar, reflection nebulosity,  usually referred to  as the  Serpens object,  the Serpens nebula, or the Serpens reflection nebulosity (SRN), illuminated by the pre-main  sequence star  SVS 2  (Strom et  al. 1974,  1976,  Worden \\& Grasdalen 1974, King  et al.  1983, Warren-Smith et  al. 1987, G\\'omez de Castro et al. 1988). In this work, we will refer to that nebulosity as SRN.    \\begin{figure}[!ht] \\begin{center} \\scalebox{0.60}{ \\includegraphics{dss_serpens1.ps} } \\caption{The  Serpens molecular  cloud  as seen  on  the DSS  $R$-band plates. Large  scale irregular dark  structures are clearly  seen. The reflection nebulosities  S 68  and SRN are  indicated, as well  as the position  of VV  Ser  and of  the  H$\\alpha$ emission  line stars  Ser G3/G6.   Field   size  is   $2\\deg   \\times   2\\deg$.  Field   centre: $\\alpha_{2000}    =   18^h$    30$^m$,    $\\delta_{2000}   =    0\\deg$ $50\\arcmin$. North  to the top, East  to the left. The  image has been downloaded    from     the    Canadian    Astronomy     Data    Center (http://cadcwww.dao.nrc.ca/dss).} \\label{Serpens} \\end{center} \\end{figure}  Relatively little  work has been  devoted to studying the  large scale properties  of  the Serpens  cloud,  having  been mainly  concentrated around SRN  and VV  Ser, which  are located towards  the North  of the large  extinction map  analysed by  Cambr\\'esy (1999).  Cohen  \\& Kuhi (1979) and Chavarr\\'\\i  a et al.  (1988) conducted optical  studies of the region  and identified several  H$\\alpha$ emission line  stars and some B,  A and  later spectral type  stars associated  with reflection nebulae. Loren et al. (1979) carried out  a 2.6 mm line CO map over an area of  $\\sim$ 300  square arc minutes  and detected a  dense H$_2$CO core of  $\\sim ~3\\farcm3$  in radius  centred in SRN.  IRAS maps  of a $\\sim ~3\\deg  \\times 3\\deg$  region around the  cloud core  show large scale  extended   far-IR  emission  and   several  point-like  sources associated   with  some   visible  and   invisible  stars   (Zhang  et al. 1988b). The cloud core with several point-like sources embedded in extended emission shows  up very prominently in all  IRAS bands (Zhang et  al. 1988a).  More  recently  Schnee et  al.  (2005) have  analysed recalibrated  IRAS data  to estimate  dust column  densities  and dust temperatures.   Assuming a  distance of  700 pc  to Serpens,  Zhang et al. (1988b) estimate a total  far-IR luminosity from the whole Serpens cloud of  $\\sim$ 3200  L$_\\odot$ and the  dust mass inferred  from the extended  emission is  $\\sim$ 290  M$_\\odot$. The  actual  distance to Serpens is most likely significantly smaller than the value assumed by Zhang et  al. and,  therefore, the luminosity  and dust  emitting mass should  be corrected.  In fact,  the distance  to Serpens  has  been a controversial  point over  the years,  with suggested  figures  in the range  from $\\sim$  200 pc  to $\\sim$  700 pc.  At present,  a smaller distance value is commonly accepted. We analyse this issue in the next section.  After  the discovery  of the  cloud core,  observational  efforts have concentrated in that area. The  reason is the extreme richness in star formation  activity found  in the  core, with  all  relevant phenomena taking place  in a small region  of just few arc  minutes in diameter: pre-protostellar     gaseous    condensations;     pre-stellar    dust condensations;  Class  0, Class  I  and  Class  II objects;  different evidences of  infalling gas; disks; molecular and  atomic outflows and jets;  clustering  of  young   stellar  objects  (YSOs)  in  different evolutionary stages, etc. The  Serpens core indeed represents a unique laboratory  for studies  of  star formation  processes and  observable phenomena, and the inter-relation between them.  A first review  article of the most relevant  observational results up to  $\\sim  1991$  was  written   by  us  (Eiroa  1991).  In  this  new contribution we concentrate mainly on results achieved during the last 15  years. We  review  observations  related to  the  Serpens core  in particular, since this part of the cloud has continued to be the focus of  the  large observational  effort  carried  out  by many  different research groups.  Nonetheless, a section  is also dedicated  to recent work carried  out in Serpens areas  outside the core.  The new results confirm and emphasize the role of Serpens as a key region for detailed studies  of  star  formation  and  the understanding  of  all  related phenomena. ",
        "conclusions": ""
    },
    "0809/0809.2780_arXiv.txt": {
        "abstract": "We present observations of redshifted CO(1-0) and CO(2-1) in a field containing an overdensity of Lyman break galaxies (LBGs) at $z=5.12$. Our Australia Telescope Compact Array observations were centered between two spectroscopically-confirmed $z=5.12$ galaxies. We place upper limits on the molecular gas masses in these two galaxies of M(H$_2$)$<1.7 \\times 10^{10}$ M$_\\odot$ and $<2.9 \\times 10^{9}$ M$_\\odot$ (2\\,$\\sigma$), comparable to their stellar masses.  We detect an optically-faint line emitter situated between the two LBGs which we identify as warm molecular gas at $z=5.1245\\pm0.0001$. This source, detected in the CO(2-1) transition but undetected in CO(1-0), has an integrated line flux of $0.106 \\pm 0.012$ Jy km\\,s$^{-1}$, yielding an inferred gas mass M(H$_2$)=$(1.9\\pm0.2)\\times 10^{10}$\\,M$_\\odot$.  Molecular line emitters without detectable counterparts at optical and infrared wavelengths may be crucial tracers of structure and mass at high redshift. ",
        "introduction": "\\label{sec:intro}  Lyman-break galaxies have given us our first detailed look at star formation at high redshifts ($z>4$).  However, our picture of the distant universe is limited because these objects are selected for strong rest-frame ultraviolet emission arising from substantial unobscured starbursts in the galaxies. Individually, they tell us little about darker baryons at the same redshift, whether the systems containing the baryons are intrinsically less UV-luminous or are obscured by dust.  The Lyman-break galaxy (LBG) population has now been probed from $z=2$ \\citep{2004ApJ...604..534S} up to $z=6$ \\citep{2007MNRAS.376..727S}, with tentative detections reported at redshifts as high as $z=10$ \\citep{2005ApJ...624L...5B}. The photometric and spectroscopic properties of LBGs evolve with increasing redshift. Typical LBGs at $z\\approx5$ have stellar masses of a few times 10$^9$\\,M$_\\odot$, are predominantly young ($<50$\\,Myr), and are sub-solar but not primordial in their metallicity (Z$\\approx$0.2\\,Z$_\\odot$), where these properties are deduced from their rest-frame optical and near-infrared spectral energy distribution \\citep{2007astro.ph..1725V}.  However something is missing in our picture of the universe at $z>5$. The youth of Lyman-break objects when compared to the range of lookback times probed in a given survey suggests that only a small fraction of the overall population is undergoing a phase of strong unobscured star formation, and hence is detectable in the rest UV, at a given time. Fossil stellar populations in the most massive galaxies at the present time suggest that the bulk of their stars were formed at $z>3$ \\citep{2004Natur.428..625H}, but the mass assembly history of the universe as traced by LBGs implies that only 1\\% of the baryonic material that will eventually form today's galaxies was in collapsed systems at $z=5$. Where is the other 99\\%? The bulk of the baryonic material at these high redshifts must be in either low mass systems or the gas phase.  Damped Lyman-alpha systems in the spectra of distant quasars have shown that dense neutral gas is ubiquitous at high redshift, holding a third of the HI atoms at $z=5$ and appearing in 10\\% of sight-lines per unit cosmological distance \\citep{2005ApJ...635..123P}. These sources represent reservoirs of neutral gas for star formation, yet are known to be deficient in H$_2$ relative to local galaxies. Molecular gas is instead more tightly concentrated in high-density, relatively dusty regions \\citep{2006ApJ...643..675Z}.  Securely identifying this material is challenging.  At $z\\approx5$, the far-infrared lines emitted by cool and obscured material are redshifted to millimeter wavelengths, but not necessarily into an atmospheric window. At these extreme luminosity distances lines are faint, requiring long integration times to secure even a tentative detection. The fields of view of existing millimeter interferometers are small, and their correlators have, until recently, been limited to bandwidths of a few hundred MHz (equivalent to $\\Delta z/z < 0.01$). The resulting small volumes made survey work at high redshift unfeasible. As a result, the few detections of line emission from molecular and atomic gas at high redshift have been towards active galactic nuclei - quasars to $z=6.4$ \\citep{2003A&A...409L..47B}, radio galaxies to $z=5.2$ \\citep{2005ApJ...621L...1K} - and hence probed highly atypical environments.  In this paper we present the results of a pilot programme to explore the cool gas associated with a large scale structure at $z>5$ marked out by Lyman-break galaxies rather than by an active galaxy. Our 40 arcmin$^2$ target field contains seven UV-luminous sources with spectroscopically confirmed redshifts in a narrow range ($\\Delta z < 0.1$), a 6$\\sigma$ excess over the typical density of such sources in our survey, but shows no clear spatial clustering\\footnote{The detailed properties of this field, and of others probed in our ESO Remote Galaxy Survey (ERGS) will be discussed in a forthcoming paper by Douglas et al.}. Given the redshift dispersion and spatial distribution of these galaxies, it is clear that their star formation cannot be a response to a common triggering event, and yet the probability of such a configuration arising by chance is small. The simplest explanation is that there is a large mass of hidden baryons in this field at this redshift. At any one time this structure is traced by the small fraction of that material undergoing comparatively short-term UV-luminous star formation events and appearing as LBGs. While the stars within the LBGs are likely to end up in the most massive current-day galaxies after subsequent mergers, the total stellar mass in even this overdensity of LBGs is an order of magnitude too small to account for the stars in the most massive spheroids. Consequently, there should be far more baryons in the vicinity of the LBGs than revealed in the UV. Our pilot programme was designed to probe the cool gas mass in a small section of the structure, centered on a pair of LBGs with Lyman-$\\alpha$ emission redshifts separated by $\\Delta z=0.004$. The observations target not only the two LBGs, but any darker/fainter systems in the underlying large scale structure.  The results for lower redshift LBGs \\citep[e.g.][]{2004ApJ...604..125B}, suggested that we were unlikely to detect the two known galaxies directly.  All magnitudes in this paper are quoted in the AB system \\citep{1983ApJ...266..713O}.  We adopt a $\\Lambda$CDM cosmology with ($\\Omega_{\\Lambda}$, $\\Omega_{M}$, $h$)=(0.7, 0.3, 0.7). ",
        "conclusions": "\\label{sec:conclusions}  Our main conclusions can be summarized as follows:  (i) We have investigated the molecular gas content, as traced by low-order transitions of CO, in a large scale structure identified in Lyman break galaxies at $z=5.12$.  (ii) Neither of two known Lyman break galaxies in our survey area show associated molecular line emission, suggesting that their molecular gas content does not significantly exceed their stellar content.  (iii) We have identified a source with line emission at 37.642\\,GHz, which we identify as emission in the CO(2-1) transition at $z=5.1245\\pm0.0001$. This source is currently undetected at any other wavelength.  (iv) The inferred gas mass in this single source is M(H$_2$)$=(1.9\\pm0.2) \\times 10^{10}$ M$_\\odot$, comparable to the total stellar mass of UV-luminous sources in the same region.  (v) Further studies of UV-faint sources at high redshift are essential to characterize such systems, and could make substantial progress towards balancing the baryon budget at early times."
    },
    "0809/0809.2263_arXiv.txt": {
        "abstract": "The existence of cusps on non-periodic strings ending on D-branes is demonstrated and the conditions, for which such cusps are generic, are derived. The dynamics of F-, D-string and FD-string junctions are investigated. It is shown that pairs of FD-string junctions, such as would form after intercommutations of F- and D-strings, generically contain cusps.  This new feature of cosmic superstrings opens up the possibility of extra channels of energy loss from a string network. The phenomenology of cusps on such cosmic superstring networks is compared to that of cusps formed on networks of their field theory analogues, the standard cosmic strings.   \\vspace{.2cm} \\noindent ",
        "introduction": "Fundamental (F) strings and Dirichlet branes with one non-compact spatial dimension (D-strings) are generically formed~\\cite{sarangi-tye,ma,jonesetal,dvalietal} at the end of brane inflation~\\cite{dvali-tye99} within the context of string inspired cosmological models.  Such strings, known as cosmic superstrings, are of cosmological size and could play the r\\^ole of cosmic strings~\\cite{Vilenkin_shellard,ms-cs07}, false vacuum remnants formed generically at the end of hybrid inflation within Grand Unified Theories~\\cite{jrs03,ms-cs08}. Cosmic superstrings have gained a lot of interest, particularly since it is believed that they may be observed in the sky, providing both a means of testing string theory and a hint for a physically motivated inflationary model (for a recent reviews, see e.g. Ref.~\\cite{sakellariadou2008,Davis2005,Davis2008}).  The most significant difference between cosmic superstrings and the field theory cosmic strings we are more familiar with, is the existence of three string junctions, the presence of which could strongly effect the dynamics of the string network.  Understanding these new dynamical effects is critical if there is to be any hope of differentiating cosmic superstrings from their solitonic analogues. A number of analytical~\\cite{Copeland:2006if,Copeland:2006eh,Copeland:2007nv} and numerical~\\cite{Sakellariadou:2004wq,Avgoustidis:2004zt,Copeland:2005cy,Hindmarsh:2006qn,Rajantie:2007hp,Urrestilla:2007yw,Sakellariadou:2008ay} studies have addressed cosmic superstring dynamics. We note that, in principle, cosmic superstring dynamics ought to be studied using the Dirac-Born-Infeld action, the low-energy effective action for many varieties of strings arising in the context of string theory.  In what follows, we investigate (generic) string solutions ending on parallel Dirichlet branes as a pedagogical example of the effects possible at a three string junction. We then look at the dynamics of a junction made up of an F-string, a D-string, and an FD-string, which are expected to form through intercommuting of initial configurations composed from F- and D-string networks.  We show that cusps are generic features for such strings, opening up a new energy loss mechanism for the network, in addition to the formation and subsequent decay of closed loops and the formation of bound states~\\cite{Sakellariadou:2008ay}. Studies of the phenomenological implications of cosmic superstrings, particularly the gravitational~\\cite{Damour:2004kw,Siemens:2006vk,Siemens:2006yp,Urrestilla:2007yw,hogan} and Ramond-Ramond~\\cite{Sakellariadou:2004wq,Firouzjahi:2007dp} radiation emitted from cosmic superstrings --- predominantly from cusps and to some extent from kinks --- are then justified. Some of the phenomenological consequences of cosmic superstring dynamics are discussed here. ",
        "conclusions": "We have shown how to characterise classical solutions to low energy effective string actions in an intuitive geometric manner. We have shown how the boundary conditions for strings ending on D-branes can lead to cusps, points with luminous velocity, in such solutions, in an analogous manner to cusps of standard cosmic string loops. In particular, we have demonstrated that cusps would be a generic feature of an F-string ending on two (parallel and stationary) D-strings. We have then considered general three string junctions that are possible in DBI-actions and have shown that the boundary conditions of such a junction can similarly be characterised and understood in a geometric setting, for the case of an F- and D-string meeting an FD-string. We find that this system exactly reproduces the situation of an F-string ending on two D-strings and hence a pair of such junctions would generically include cusps. We have shown that this remains true even when the D-strings are moving.  The relevance of such a scenario is that networks of F- and D-strings are expected to form at the end of brane inflation. Collisions of such networks would lead to pairs of three string junctions, each of which would then be expected to have cusps, opening up a new energy loss mechanism for such networks. Our new feature is in addition to cusps on string loops. Importantly the formation and existence of cusps is expected to be most significant early in the evolution of the network. The fact that the radiation from such cusps should include all available fields present in the low energy string theory, makes it possible that signatures of their presence would appear in baryogenesis or big bang nucleosynthesis.  The extreme nature of cusps means that they are possible targets for observations of string networks and here we have shown that it is, in principle, possible to distinguish between standard cosmic strings formed during cosmological phase transitions and cosmic superstring, relics of brane inflation. The observation of three string junctions would provide strong evidence  for string theory. In future it might be possible to observe the emission of radiation from cusps and the distribution of such cusps.  \\ack It is a pleasure  to thank M. Green and J. Polchinski for enlightening discussions.  The work of M.S. is partially supported by the  European Union through  the Marie Curie Research  and Training Network {\\sl UniverseNet} (MRTN-CT-2006-035863). This work is supported in part by STFC.  \\vskip1.truecm"
    },
    "0809/0809.4029_arXiv.txt": {
        "abstract": "AXTAR is an X-ray observatory mission concept, currently under study in the U.S., that combines very large collecting area, broadband spectral coverage, high time resolution, highly flexible scheduling, and an ability to respond promptly to time-critical targets of opportunity.  It is optimized for submillisecond timing of bright Galactic X-ray sources in order to study phenomena at the natural time scales of neutron star surfaces and black hole event horizons, thus probing the physics of ultradense matter, strongly curved spacetimes, and intense magnetic fields.  AXTAR's main instrument is a collimated, thick Si pixel detector with 2--50~keV coverage and 8~m$^2$ collecting area.  For timing observations of accreting neutron stars and black holes, AXTAR provides at least an order of magnitude improvement in sensitivity over both RXTE and Constellation-X.  AXTAR also carries a sensitive sky monitor that acts as a trigger for pointed observations of X-ray transients and also provides continuous monitoring of the X-ray sky with 20$\\times$ the sensitivity of the RXTE ASM. AXTAR builds on detector and electronics technology previously developed for other applications and thus combines high technical readiness and well understood cost. ",
        "introduction": "The natural time scales near neutron stars (NSs) and stellar mass black holes (BHs) are in the millisecond range.  These time scales characterize the fundamental physical properties of compact objects: mass, radius, and angular momentum. For example, the maximum spin rate of a NS is set by the equation of state of the ultradense matter in its interior, a fundamental property of matter that still eludes us. Similarly, orbital periods at a given radius near a BH are set by the BH's mass, angular momentum, and the laws of relativistic gravity. Although it was recognized for decades that the measurement of such time scales would provide unique insight into these compact stars and their extreme physics, it was not until the 1995 launch of the Rossi X-ray Timing Explorer (RXTE; \\citep{brs93}) that oscillations on these time scales were actually detected from accreting NSs and stellar-mass BHs.  This conference celebrates RXTE's discovery of millisecond oscillations that trace the spin rate of accreting NSs.  RXTE has also discovered millisecond oscillations from accreting BHs with frequencies that scale inversely with BH mass and are consistent with the orbital time scale of matter moving in the strongly curved spacetime near the BH event horizon.  While RXTE revealed the existence of these phenomena, it lacks the sensitivity to fully exploit them in determining the fundamental properties of NSs and BHs.  Here, we describe the Advanced X-ray Timing Array (AXTAR), a new mission concept boasting an order of magnitude improvement in sensitivity over RXTE\\footnote{A similar mission concept was previously discussed by Phil Kaaret and collaborators \\citep{kgl+01,kaa04}.}.  AXTAR was initially proposed as a medium-class probe concept for NASA's 2007 Astrophysics Strategic Mission Concept Studies.  A modified version of AXTAR is also being studied as a NASA MIDEX-class mission concept.   \\begin{table} \\begin{tabular}{p{1.5in}p{3in}} \\hline \\tablehead{1}{c}{b}{Science Objective} & \\tablehead{1}{c}{b}{AXTAR Observations} \\\\ \\hline NS mass, radius, EOS  & X-ray burst oscillations. kHz QPOs. Accreting ms pulsars. Asteroseismology with magnetar oscillations. \\\\ \\hline Strongly curved spacetimes & BH oscillations. Broad Fe lines. Phase-resolved spectroscopy of low-freq QPOs. AGN monitoring. \\\\ \\hline Physics of nuclear burning & Thermonuclear X-ray bursts and superbursts. \\\\ \\hline Physics of accretion  & Accreting msec pulsars. kHz QPOs in NSs. \\\\ \\hline Physics of jets       & NS and BH transients.  \\\\ \\hline Physics of mass transfer & X-ray pulsar orbital evolution.  \\\\ \\hline Multipolar magnetic field components of pulsars & Hard X-ray cyclotron lines in high-mass binaries. Magnetar pulse profiles. \\\\ \\hline \\end{tabular} \\caption{Selected AXTAR Science Topics} \\label{tab:a} \\end{table} ",
        "conclusions": ""
    },
    "0809/0809.3115_arXiv.txt": {
        "abstract": "In this paper we extend the idea suggested previously by \\citet{Petri2005a,Petri2005b} (paper~I and II) that the high frequency quasi-periodic oscillations (HF-QPOs) observed in low-mass X-ray binaries (LMXBs) may be explained as a resonant oscillation of the accretion disk with a rotating asymmetric background (gravitational or magnetic) field imposed by the compact object. Here, we apply this general idea to black hole binaries.  It is assumed that a test particle experiences a similar parametric resonance mechanism such as the one described in paper~I and II but now the resonance is induced by the interaction between a spiral density wave in the accretion disk, excited close to the innermost stable circular orbit, and vertical epicyclic oscillations. We use the Kerr spacetime geometry to deduce the characteristic frequencies of this test particle. The response of the test particle is maximal when the frequency ratio of the two strongest resonances is equal to 3:2 as observed in black hole candidates. Finally, applying our model to the microquasar GRS~1915+105, we reproduce the correct value of several HF-QPOs.  Indeed the presence of the 168/113/56/42/28~Hz features in the power spectrum time analysis is predicted.  Moreover, based only on the two HF-QPO frequencies, our model is able to constrain the mass~$M_{\\rm BH}$ and angular momentum~$a_{\\rm BH}$ of the accreting black hole. We show the relation between $M_{\\rm BH}$ and $a_{\\rm BH}$ for several black hole binaries. For instance, assuming a black hole weakly or mildly rotating, i.e.  $a_{\\rm BH} \\le 0.5 \\, G\\,M_{\\rm BH}/c^2$, we find that for GRS~1915+105 its mass satisfies $13 \\, M_\\odot \\le M_{\\rm BH} \\le 20 \\,M_\\odot$.  The same model applied to two other well-known BHCs gives for GRO~J1655-40 a mass $5 \\, M_\\odot \\le M_{\\rm BH} \\le 7 \\, M_\\odot$ and for XTE~J1550-564 a mass $8 \\, M_\\odot \\le M_{\\rm BH} \\le 11 \\, M_\\odot$. This is consistent with other independent estimations of the black hole mass.  Finally for H1743-322, we found the following bounds, $9 \\, M_\\odot \\le M_{\\rm BH} \\le 13 \\, M_\\odot$. ",
        "introduction": "High frequency quasi-periodic oscillations (HF-QPOs) are common features to all accreting compact objects, be they neutron stars, white dwarfs or black holes. A number of recent observations have revealed the existence of these HF-QPOs in several black hole binaries \\citep{Strohmayer2001a,MacClintock2003,Remillard2006}. Usually, a pair of HF-QPOs appears in a 3:2 ratio. If these oscillations are connected to the orbital motion of the accretion disk at its inner edge as predicted by several models, these QPOs become a useful test of gravity in the strong field regime.  The 3:2 ratio was first noticed by \\cite{Abramowicz2001}. In order to explain this ratio, they introduced a resonance mechanism between orbital and epicyclic motion around Kerr black holes therefore leading to an estimate of their mass and spin. \\cite{Kluzniak2004a} showed that the twin kHz-QPOs is explained by a non linear resonance in the epicyclic motion of the accretion disk.  \\cite{Rebusco2004} developed the analytical treatment of these oscillations.  This parametric epicyclic resonance model of hydrodynamical modes in the accretion disk was applied by \\cite{Kluzniak2002} to some microquasars. They pointed out the 3:2 ratio in the HF-QPOs observed in GRO J1655-40, XTE J1550-564 and GRS1915+105. Moreover, in this latter black hole binary, \\cite{Strohmayer2001b} reported another pair of frequencies, namely 69.2~Hz and 41.5~Hz, which are in a 5:3 ratio as noticed by \\cite{Kluzniak2002} and supports the parametric resonance model. Furthermore, \\cite{Rezzolla2003} suggested that the HF-QPOs in black hole binaries are related to p-mode oscillation in a non Keplerian torus.  Nevertheless, the propagation of the emitted photons in curved spacetime can also produce some intrinsic peaks in the Fourier spectrum of the light curves \\citep{Schnittman2004}. In combination with a vertically oscillating torus, the gravitational lensing effect can also reproduce the 3:2 ratio \\citep{Bursa2004,Schnittman2005}.  Resonances in the geodesic motion of a single particle have been investigated by \\cite{Abramowicz2003}. The specific coupling force between radial and vertical oscillation was left unspecified. Their results for non-linear resonance were applied to accreting neutron stars.  In this paper, we describe a coupling between spiral density waves in the accretion disk and epicyclic motions of test particles. It is divided in two parts. In Sec.~\\ref{sec:Model} we specify the perturbation pattern used in our model. Then the equation of motions due to the perturbation are derived leading to some resonance conditions.  In Sec.~\\ref{sec:Discussion}, the results are applied to GRS~1915+105 for which several components in the Fourier time analysis are predicted in agreement with the HF-QPOs detected for this object. We also put some constrain on the mass-spin relation for several BHCs. ",
        "conclusions": "\\label{sec:Conclusion}  In this paper, the consequences of an outgoing spiral density wave on the evolution of a test particle orbiting around a Kerr black hole have been explored.  The twin peaks ratio around 3:2 for the kHz-QPOs is naturally explained by the parametric resonance model. The connection to lower frequency QPOs is also clearly demonstrated. It is an extension to accreting black hole for the model already suggested in neutron star and white dwarf binaries.  From the analysis of HF-QPOs, we are also able to constrain the mass and the spin of the hole as already suggested by \\cite{Torok2005}.  In the case of GRS~1915+105, it appears that the 56~Hz feature in the Fourier time analysis is a fundamental frequency of the black hole from which the other QPOs can be derived. This feature should therefore be related to the intrinsic properties of the black, namely, its mass and its angular momentum. The way to associate $\\Omega_{\\rm w}$ with $a_{\\rm BH}$ and $M_{\\rm BH}$ remains unclear and needs further investigation but spiral density wave excitation close to the ISCO at nearly the maximum of the radial epicyclic frequency as suggested by~\\citet{Mao2008} is a interesting idea that needs further investigations."
    },
    "0809/0809.3323_arXiv.txt": {
        "abstract": "{ For thermonuclear flashes to occur on neutron-star surfaces, fuel must have been accreted from a donor star. However, sometimes flashes are seen from transient binary systems when they are thought to be in their quiescent phase, during which no accretion, or relatively little, is expected to occur. We investigate the accretion luminosity during several such flashes, including the first-ever and brightest detected flash from \\object{Cen\\,X-4} in 1969. We infer from observations and theory that immediately prior to these flashes the accretion rate must have been between about 0.001 and 0.01 times the equivalent of the Eddington limit, which is roughly 2 orders of magnitude less than the peak accretion rates seen in these transients during an X-ray outburst and 3--4 orders of magnitude more than the lowest measured values in quiescence. Furthermore, three such flashes, including the one from \\object{Cen\\,X-4}, occurred within 2 to 7 days followed by an X-ray outburst. A long-term episode of enhanced, but low-level, accretion is predicted near the end of the quiescent phase by the disk-instability model, and may thus have provided the right conditions for these flashes to occur. We discuss the possibility of whether these flashes acted as triggers of the outbursts, signifying a dramatic increase in the accretion rate. Although it is difficult to rule out, we find it unlikely that the irradiance by these flashes is sufficient to change the state of the accretion disk in such a dramatic way. } ",
        "introduction": "In low-mass X-ray binaries (LMXBs) a neutron or black hole accretes matter via an accretion disk from a less massive Roche-lobe filling companion star. In many LMXBs the accretion from the disk on the compact object is transient. After a transient outburst (hereafter referred to as outburst), lasting typically a few months, the system settles in quiescence for a few months to several decades. The disk-instability model (DIM; Osaki 1974, see Lasota 2001 for a review) predicts that, if in quiescence the disk extends down to the stellar surface (or the last stable Keplerian orbit), the accretion rate is negligibly low ($\\sim$10$^{5}$\\,g\\,s$^{-1}$; Lasota et al.\\ 2008, see Eq.~(\\ref{eq:mdotcrit}) in Sect.~4). However, in general the X-ray emission of quiescent transient sources corresponds to accretion rates that cannot be qualified as negligible (e.g., van Paradijs et al.\\ 1987, Campana et al.\\ 2004).  In the case of neutron-star transients there has been a controversy on the origin of the quiescent X-ray luminosity. On the one hand, quiescent-disk truncation easily explains the luminosity (Lasota et al.\\ 1996, Dubus et al.\\ 2001, Narayan \\&\\ McClintock 2008), is also applicable to black-hole systems, and has been confirmed by observations (e.g., Done 2002). On the other hand, Brown et al.\\ (1998) propose that quiescent X-rays have their origin in heating of the neutron-star crust by nuclear reactions and not in accretion. Both models have their difficulties. The observed rapid variability of the quiescent X-ray flux and the presence of a substantial power-law component (as seen in, e.g., the LMXB transient \\object{Cen\\,X-4}, see Campana et al.\\ 1997, 2004, Rutledge et al.\\ 2001), as well as a very low quiescent X-ray luminosity (less than about several times 10$^{30}$\\,erg\\,s$^{-1}$ for \\object{1H\\,1905+000}; Jonker et al.\\ 2007), are difficult to reconcile with the deep crustal heating model (see, e.g., Jonker 2008 for a recent discussion). The lack of a well understood disk-truncation mechanism (see, however, Liu et al.\\ 2002) and the overpredicted ratio of neutron-star to black-hole quiescent X-ray luminosities (Menou et al.\\ 1999) are the weaknesses of its competitor. Therefore, any independent estimate of the quiescent accretion rate in transient systems would be a great help in resolving the controversy.  Type I X-ray bursts (Grindlay et al.\\ 1975, Belian et al.\\ 1976, Hoffman et al.\\ 1978) are thermonuclear flashes at the surface of a neutron star (Joss 1977, Maraschi \\&\\ Cavaliere 1977, Lamb \\&\\ Lamb 1978; for reviews see, e.g., Lewin et al.\\ 1993, Strohmayer \\&\\ Bildsten 2006). We hereafter refer to these events as flashes. If the energy release during a flash is fast and large enough, the local luminosity on the neutron star surface can surpass the Eddington limit, resulting in a lift-up of the photosphere. Such flashes are referred to as photospheric radius-expansion X-ray bursts. During the expansion phase the inferred temperature decreases, whereas the inferred emitted area increases (see, e.g., Lewin et al.\\ 1993, for a review).  The first observed flash, in retrospect, was the event detected on July 7, 1969, with {\\it Vela 5B} from \\object{Cen\\,X-4} (Belian et al.\\ 1972; see Sect.~2). The event lasted for about 10\\,min, and is still the brightest ever observed with a peak flux of about 60\\,Crab\\footnote{The commonly used Crab unit is equivalent to about 2$\\times$10$^{-8}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ in the classical 2--10\\,keV photon-energy band.} (3--12\\,keV). Two days after the flash \\object{Cen\\,X-4} went into an X-ray outburst, which peaked at about 25\\,Crab (3--12\\,keV) and lasted for about 80 days (Conner et al.\\ 1969; Evans et al.\\ 1970; see also Sect.~2). Except for the flash, no other X-ray emission was detected from \\object{Cen\\,X-4} before the outburst (Belian et al.\\ 1972). Interestingly, a similar situation recently occurred in another source: a $\\simeq$40\\,s long flash was detected from \\object{IGR\\,J17473$-$2721} about 2 days prior to the X-ray outburst (Del Monte et al.\\ 2008, Markwardt et al.\\ 2008). In an other case three flashes separated by 2--3~days were seen during the beginning of an X-ray outburst of \\object{2S\\,1803$-$245} (Cornelisse et al.\\ 2007). Our inspection of the {\\it RXTE} All Sky Monitor (ASM) light curve shows that the first flash, which lasted for about 40--50\\,s, occurred when there was no detectable X-ray emission. The following two flashes occurred when the source showed X-ray emission at a slightly elevated level. About a week later the source developed into a bright X-ray outburst.  A few other LMXB transients have shown flashes after and/or in between X-ray outbursts when their X-ray emission was below the detection thresholds and the systems were inferred to be in their quiescent phase: \\object{2S\\,1711$-$339}, \\object{SAX\\,J1808.4$-$3658} and \\object{GRS\\,1747$-$312} (Cornelisse et al.\\ 2002a, in 't Zand et al. 2001, 2003b, respectively). Also \\object{Cen\\,X-4} may have shown a flash $\\sim$2 years after its 1969 X-ray outburst (Gorenstein et al.\\ 1974; see Appendix \\ref{apollo}). Related examples are various flashes seen from sources without detectable pre-flash emission, the so-called burst-only sources (Cocchi et al.\\ 2001, Cornelisse et al. 2002a,b, and references therein).  For all the afore-mentioned flashes, accretion must have been ongoing prior to these events at a level that is orders of magnitude lower than that achieved during an outburst, but, interestingly, higher than quiescent levels. This brings about prospects for new constraints on models for the emission during the quiescent phase. We explore in Sect.~3 the observed and expected mass-accretion rates around the times of the flashes seen for the sources presented above. Furthermore, the flashes which are within 2--7~days followed by X-ray outbursts suggest the presence of a physical process which has thus far not been discussed in the literature\\footnote{Belian et al.\\ (1972) called \\object{Cen\\,X-4}'s 1969 flash a probable precursor to its subsequent X-ray outburst. They suggested the two events to be associated, but no physical scenario was discussed.}: that the flashes serve as triggers for accretion-disk instabilities resulting in X-ray outbursts. We study in Sect.~4 the viability of this idea in the context of \\object{Cen\\,X-4}. ",
        "conclusions": "\\label{discussions}  The detection of a flash at the end of the quiescent phases in \\object{Cen\\,X-4}, \\object{IGR\\,J17473$-$2721} and \\object{2S\\,1803$-$245}, in between recurrent X-ray outbursts of \\object{SAX\\,J1808.4$-$3658} and \\object{GRS\\,1747$-$312}, and after an X-ray outburst of \\object{2S\\,1711$-$339}, shows that accretion was (temporarily) ongoing onto the neutron star. Such accretion must be at a level several orders of magnitudes higher than that inferred for quiescence, and the rates required are consistent with the truncated disk model which predicts an accretion-rate enhancement before the onset of the outburst. Therefore, it seems unlikely that crustal heating is the main source of luminosity in all of quiescence. Thus, care must be taken when considering the properties of a cooling neutron star when an outburst is over.  The DIM-predicted enhancement in accretion rate near the end of quiescence of about 10$^{13}$ to 10$^{14}$\\,g\\,s$^{-1}$ might have provided the right conditions for the flashes to occur in just before the outbursts of \\object{Cen\\,X-4}, \\object{IGR\\,J17473$-$2721} and \\object{2S\\,1803$-$245}. The uncertainties inherent in the DIM do not allow us to test in detail the hypothesis that the flashes triggered or accelerated the start of, the outbursts (althought we regard it as unlikely), nor the 2--7-day delay between the flash and the X-ray outburst."
    },
    "0809/0809.1326_arXiv.txt": {
        "abstract": "We present an efficient method to generate large simulations of the Epoch of Reionization (EoR) without the need for a full 3-dimensional radiative transfer code. Large dark-matter-only simulations are post-processed to produce maps of the redshifted 21cm emission from neutral hydrogen. Dark matter haloes are embedded with sources of radiation whose properties are either based on semi-analytical prescriptions or derived from hydrodynamical simulations. These sources could either be stars or power-law sources with varying spectral indices. Assuming spherical symmetry, ionized bubbles are created around these sources, whose radial ionized fraction and temperature profiles are derived from a catalogue of 1-D radiative transfer experiments. In case of overlap of these spheres, photons are conserved by redistributing them around the connected ionized regions corresponding to the spheres. The efficiency with which these maps are created allows us to span the large parameter space typically encountered in reionization simulations. We compare our results with other, more accurate, 3-D radiative transfer simulations and find excellent agreement for the redshifts and the spatial scales of interest to upcoming 21cm experiments. We generate a contiguous observational cube spanning redshift 6 to 12 and use these simulations to study the differences in the reionization histories between stars and quasars. Finally, the signal is convolved with the LOFAR beam response and its effects are analyzed and quantified. Statistics performed on this mock data set shed light on possible observational strategies for LOFAR. ",
        "introduction": "The history of our Universe is largely unknown between the surface of last scattering ($z\\approx1100$) down to a redshift of about 6. Because of the dearth of radiating sources and the fact that we know very little about this epoch, it is often referred to as the ``dark ages''. Theoretical models suggest that around redshifts 10 -- 20, the first sources of radiation appeared that subsequently reionized the Universe. Two different experiments provide the bounds for this epoch of reionization (EoR); the high polarization component at large spatial scales of the temperature-electric field (TE) cross-polarization mode of the cosmic microwave background (CMB) providing the upper limit for the redshift at $z \\approx 11$ \\citep{page07} and the rapid increase in the Lyman-$\\alpha$ optical depth towards redshift 6, observed in the spectrum of high redshift quasars \\citep{fan06}, the lower limit. Although the redshifted 21cm hyperfine transition of hydrogen was proposed as a probe to study this epoch decades ago \\citep{sz75}, the technological challenges to make these observations possible are only now being realised. In the meantime, theoretical understanding of the EoR has improved greatly \\citep{hogan79,scott90,madau97}.  Over the past few years there have been considerable efforts in simulating the 21cm signal from the Epoch of Reionization. Almost all of the methods employed in simulating the 21cm involve computer intensive full 3-D radiative transfer calculations \\citep{otvet,crash,simpX,rsph,ftte,art,zeus,flash,c2ray,zahn07,mesinger07,pawlik08}.  Theories predict that the process of reionization is complex and sensitively dependent on many not-so-well-known parameters. Although stars may be the most favoured of reionization sources, the role of mini-quasars (miniqsos), with the central black hole mass less than a few million solar masses, are debated \\citep{nusser05,zaroubi05,kuhlen05,rajat08}.  Even if the nature of the sources of radiation would be relatively well constrained, there are a number of ``tunable'' parameters like the photon escape fraction, masses of these first sources, and so on, that are not well constrained.  In a couple of years, next generation radio telescopes like LOFAR\\footnote{www.lofar.org}\\footnote{www.astro.rug.nl/\\~~LofarEoR} and MWA\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} will be tuned to detect the 21cm radiation from the EoR. Although the designs of these telescopes are unprecedented, the prospects for successfully detecting and mapping neutral hydrogen at the EoR critically depends on our understanding of the behaviour and response of the instrument, the effect of diffuse polarized Galactic \\& extra-galactic emission, point source contamination, ionospheric scintillations, radio frequency interference (RFI) and, not least, the characteristics of the desired signal. A good knowhow of the above phenomena would enable us to develop advanced signal processing/extraction algorithms, that can be efficiently and reliably implemented to extract the signal.  In order to test and confirm the stability and reliability of these algorithms, it is imperative that we simulate, along with all the effects mentioned above, a large range of reionization scenarios.  Fig.~\\ref{fig:bigpic} shows the basic building blocks of the simulation pipeline being built for the LOFAR-EoR experiment. This paper basically constitutes the first block, i.e., simulation of the cosmological 21cm EoR signal. This then passes through a sequence of blocks like the foreground simulation \\citep{jelic08}, instrument response and extraction (Lambropoulous et al., \\emph{in prep}). The extracted signal is then compared with the original signal to quantify the performance of the extraction scheme. This process needs to be repeated for various reionization scenarios to avoid any bias the extraction scheme would have if only a subsample of all possible signal characteristics were used.  Simulating observing windows as large as the Field of View (FoV) of LOFAR ($\\sim 5^{\\circ} \\times 5^{\\circ} $) and for frequencies corresponding to redshift 6 to 12 is a daunting task for conventional \\hbox{3-D} radiative transfer codes because of multiple reasons such as a requirement for high dynamic range in mass for the sources of reionization, their large number towards the end of reionization and the size of the box which strains the memory of even the largest computer cluster. In order to facilitate the simulation of such large mock data sets for diverse reionization scenarios, we need to implement an approximation to these radiative transfer methods that mimic the ``standard'' \\hbox{3-D} simulations to good accuracy.  It was clear from the onset that the details of the ionization fronts like its complex non-spherical nature will not be reproduced by the semi-analytical approach that we propose here. But the argument towards overlooking this discrepancy is that when the outputs of our semi-analytical approach and that of a \\hbox{3-D} radiative transfer code are passed through the machinery of the LOFAR-EoR pipeline, they are experimentally indistinguishable. The reason being the filtering nature of the telescope's point spread function (PSF) across the sky and the substantial bandwidth averaging along the frequency/redshift direction that is needed to recover the signal, smoothes out the structural details captured by \\hbox{3-D} codes.  Recently, several authors \\citep{zahn07,mesinger07} have proposed schemes to reduce the computational burden of generating relatively accurate 21cm maps.  These methods do fairly well, although there are some caveats, like for example the intergalactic medium (IGM) ionization being treated as binary, i.e., the IGM is either ionized or neutral \\citep{mesinger07}. Although this might be the case for stellar-like sources, others with a power-law component could exhibit an effect on the IGM wherein the ionizing front is extended and hence this assumption need not hold \\citep{zaroubi05,rajat08}. Added to this, the schemes presented in order to compute the 21cm maps, make the assumption that the spin temperature of hydrogen, $T_s$ is much larger than the CMB temperature. Towards the end of reionization ($z < 8 $) this might very well be valid. But the dawn of reionization would see a complex spatial correlation of IGM temperatures with the sources of radiation, its clustering and spectral energy distributions \\citep{venky01,zaroubi07,rajat08,pritchard07}. In the current paper we have assumed that the spin temperature is coupled to the kinetic temperature and that they are much higher than the CMB temperature. At higher redshifts this need not be a valid assumption. The effects of heating by different types of radiative sources on the IGM and the coupling (both Ly-$\\alpha$ and collisional) between the spin and kinetic temperature will be simulated accurately in a follow-up paper \\emph{Thomas et al., in prep}, using the same scheme, but now applied to heating.   In this paper we propose a method of post processing numerical simulations in order to rapidly generate realistic 21cm maps. Briefly, the algorithm consists of simulating the ionization fronts created by the ``first'' radiative sources for a range of parameters which include the power spectrum, source mass function and clustering. We then identify haloes in the outputs of N-body simulations and convert them to a photon count using semi-analytical prescriptions, or using the photon count derived from a Smoothed Particle Hydrodynamics (SPH) simulation.  Depending on the photon count and the spectrum, we embed a sphere around the centre-of-mass (CoM) of the halo whose radial profile matches that of a profile from the table created by the 1-D radiative transfer code of \\citet{rajat08}. Appropriate operations are carried out to conserve photon number. Since, the basic idea is to expand bubbles around locations of the sources of radiation, we call this method BEARS (Bubble Expansion Around Radiative Sources). For the sake of brevity and comparison with full \\hbox{3-D} radiative transfer codes, we restrict ourselves to monochromatic radiative transfer with a fixed temperature.  A following paper will include a full spectrum along with the temperature evolution. The results of our semi-analytic scheme will be compared to those obtained with the full \\hbox{3-D} Monte Carlo radiative transfer code CRASH \\citep{crash,maselli03}.  In \\S\\ref{sec:simulations} we describe the various steps involved in implementing the BEARS algorithm on the outputs of N-body simulations. We describe the specifications of the N-body simulations, the 1-D radiative transfer code used to produce the catalogue of ionization profiles, the algorithm employed to embed the sources and finally an illustrative example of the procedure to correct for the overlap of ionized bubbles. In \\S\\ref{sec:comparison} the fully \\hbox{3-D} cosmological radiative transfer code CRASH is summarized and the results of its qualitative and statistical comparison with BEARS are discussed.  \\S\\ref{sec:datacube} describes the method of generating the cube with maps of the brightness temperature ($\\delta T_b$) at all the frequencies that will be observed by an EoR experiment.  In \\S\\ref{sec:starsVquasars} we use the simulation of the cube, with maps of the sky at different frequencies, to study the difference between two popular sources of reionization, i.e., stars and quasars.  These maps in the cube are then filtered through the LOFAR antenna response to output the final data cube in \\S\\ref{sec:instrumental}. Finally in \\S\\ref{sec:conclusions} we summarize our results and outline further improvements that need to be made in our approach in order to start exploring the large parameter space involved in reionization studies.  \\begin{figure} \\centering \\hspace{0cm} \\includegraphics[width=.5\\textwidth]{fig1.ps} \\vspace{-3cm} \\caption{The big picture: This simple flow diagram encapsulates the essence of the ``LOFAR EoR simulation pipeline''. Starting from the generation of the the cosmological EoR signal, the pipeline includes the addition of foreground contaminants like Galactic synchrotron and free-free radiation and other point sources, and the LOFAR antenna response. This ``mock'' data set is then used to extract the signal using various inversion algorithms and the result is then compared to the original ``uncorrupted'' signal in order to study the accuracy and stability of the inversion schemes employed. } \\label{fig:bigpic} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions} We emphasized the need to expand the scenarios that need to be explored for reionization and their role in building a reliable and robust pipeline, thus necessitating fast realizations of different scenarios.  We built a scheme in which a range of 1-D ionization profiles were catalogued for a number of luminosities, redshifts, densities and source spectra and which were subsequently coupled to an $N$-body simulation to obtain an approximation of a ``standard'' \\hbox{3-D} radiative transfer code. The results obtained were validated using CRASH, a full \\hbox{3-D} radiative transfer code with ray tracing. The agreement between the two methods was excellent for early redshifts ($z>8$), but as expected, the discrepancy grew towards lower redshifts. Several visual comparisons of the slices of different boxes were made along with three different statistical measures of the similarity between the simulations.  Many snapshots ($\\approx 75$) were used between a redshift of 15 and 6 and the radiative transfer done on them as described in the preceding sections. These snapshots were then used to make a contiguous cube running from redshift 6 to 12. In terms of the observing frequency, the cube spans 115 $\\mathrm{MHz}$ to 200 $\\mathrm{MHz}$ with a frequency resolution matching that of LOFAR, about one $\\mathrm{MHz}$. Cubes were generated for scenarios involving only stars or quasars and some diagnostics are provided to quantitatively differentiate between them. For both models the variance in $\\delta T_b$ peaked at around 160 $\\mathrm{MHz}$. Although neither may reflect reality, these scenarios demonstrate the use of the techniques we have developed to span large parameter spaces of variables. The PDF of $\\delta T_b$ (see Fig.~\\ref{fig:dtb_pdf}) provides a statistical discriminant between the different source scenarios and could be used in the future to look for the statistical detection of the signal.  The cubes generated provide $\\delta T_b$ as a function of frequency. The cube was then averaged over a bandwidth of 1~$\\mathrm{MHz}$ and convolved with the beam pattern of LOFAR to understand the distortions caused by incomplete sampling by an interferometer. Even if the images were blurred by the operation, the overall characteristics of the signal remain detectable. Although the behaviour of the variance of the signal before and after the convolution remained the same (except that the latter had lower values on average), the image/slice entropy showed a very different behaviour. In the former case, i.e., before the convolution, the image entropy remained almost flat throughout the frequency range whereas in the latter the entropy steeply rises at around 160 $\\mathrm{MHz}$.  As a note, it is important to mention that although in principle the ionized bubbles do move with a peculiar velocity $v_r(z) = v_r(0)(1+z)^{-1/2}$, where $v_r(0)$ is the typical peculiar velocity of galaxies at redshift zero, assuming $v_r(0) \\approx 600$km/s, a redshift 10 object would have a peculiar velocity of 200 km/s. For a typical lifetime of the source considered, i.e., 10 Myr, this corresponds to motion of about a couple of kpc. This is an order of magnitude less than the resolution of the simulation box at that redshift. Therefore we ignore this effect. On the other hand we have taken into account the effect of redshift distortions whose effects are relatively more important.  The simulations and comparisons in this paper have focused on purely stellar or quasars sources, but it is plausible that the early sources of reionization were a mixture of stars and quasars or other yet unknown sources. It is therefore important to simulate reionization by a mixture of these sources, taking into account their clustering properties. The simulations presented in this paper did not take into account the contraints on the population of ionizing sources imposed by various measurements like the infrared excess in the case of stars and the soft X-ray excess for quasars.  Apart from the ionization patterns induced by these sources, the $\\delta T_b$ maps will also depend on the kinetic temperature which is coupled to the spin temperature via collisions or Ly-$\\alpha$ pumping. Hence, it is imperative that we include these temperature effects on the IGM. We will incorporate the mixture of sources and the effect of the temperature in an upcoming paper (Thomas et al., \\emph{in prep}).  One of the main astrophysical hurdles for the detection of the EoR signal is the existence of prominent Galactic and extragalactic foregrounds. Typically, the difference between the mean amplitude of the EoR signal and the foregrounds is expected to be 4 to 5 orders of magnitude, but an interferometer like LOFAR measures only the fluctuations which in this case are expected to be different by `only' three orders of magnitude \\citep{jelic08}. In subsequent papers we will explore the effects of foregrounds and their removal strategies."
    },
    "0809/0809.3862_arXiv.txt": {
        "abstract": "{} {In the frame of the search for extrasolar planets and brown dwarfs around early-type stars, we present the results obtained for the F-type main-sequence star HD\\,60532 (F6V) with \\harps.} {Using 147 spectra obtained with \\harps~at {\\small La Silla} on a time baseline of two years, we study the radial velocities of this star.} {HD\\,60532 radial velocities are periodically variable, and the variations have a Keplerian origin. This star is surrounded by a planetary system of two planets with minimum masses of 1 and 2.5\\,\\Mjup~and orbital separations of 0.76 and 1.58\\,AU respectively. We also detect high-frequency, low-amplitude (10\\,\\ms~peak-to-peak) pulsations. Dynamical studies of the system point toward a possible 3:1 mean-motion resonance which should be confirmed within the next decade.} {} ",
        "introduction": "Radial-velocity (RV) surveys have lead to the detection of nearly 300 planets\\footnote{Jean Schneider, http://exoplanet.eu} during the past decade. They mainly focus on solar and late-type main-sequence (MS) stars ($\\ga$\\,F7) which exhibit numerous lines with low rotational broadening. It is often thought that planets around more massive MS stars are not accessible to radial-velocity techniques, as they present a small number of stellar lines, usually broadened and blended by stellar rotation. However, we recently showed (\\cite{galland05a} 2005a, Paper\\,I) that with a new radial-velocity measurement method that we developed, it is possible to detect planets even around early A-type MS stars with large rotational velocities (typically 100\\,\\kms). Finding planets around such massive MS stars is of importance, as this allows to test planetary formation and evolution processes around a larger variety of stars, in terms of stellar mass and time scales of evolution processes. This approach is complementary with the one which intends to detect planets around evolved intermediate-mass stars (\\eg \\cite{sato05} 2005, \\cite{lovis07} 2007). In this case, close-in planets have been wiped out but the stellar variability is in principle less intense.  We performed a radial-velocity survey dedicated to the search for extrasolar planets and brown dwarfs around a volume-limited sample of A--F main-sequence stars with the \\harps~spectrograph (\\cite{mayor03} 2003) installed on the 3.6-m ESO telescope at La Silla Observatory (Chile). We monitored a sample of 185 MS stars with \\bv~ranging between $-0.1$ and 0.6. From the measured RV jitter, we computed the minimum detectable masses with \\harps, and showed that in 100 cases, planets with periods smaller than 100 days can be detected, even around stars with early spectral types (down to $\\sim$0.1\\,\\Mjup~at 100 days around slow-rotating late-F stars). Given the data at hand, we also provided the achieved detection limits on the individual targets (\\cite{lagrange08} 2008).  In the course of this survey, we identified a few stars which RV variations could be attributed to planets. Most of these stars are still under follow-up. We present here the detection of a planetary system around one of these stars, HD\\,60532. Section~2 provides the stellar properties of this star, the measurement of the radial velocities and their relevance; we also present a Keplerian solution associated to the presence of two planets. ",
        "conclusions": "We have shown that HD\\,60532, an F6IV--V, 1.44\\,\\Msun~star hosts two planets with minimum masses of 1 and 2.5\\,\\Mjup~and orbital separations of 0.76 and 1.58\\,AU respectively, in a possible 3:1 resonance which needs to be confirmed within the next 10 years. Noticeably sofar, only one other multiple system had been reported around a MS star more massive than 1.3\\,\\Msun~(\\object{HD\\,169830}; 1.4\\,\\Msun). Note also that the low metallicity of HD\\,60532 is not common for stars harbouring Jupiter-mass planets; the relation between the star's metallicity and the presence of massive planets has still to be investigated further."
    },
    "0809/0809.1867_arXiv.txt": {
        "abstract": "Predicting the colors of Luminous Red Galaxies (LRGs) in the Sloan Digital Sky Survey (SDSS) has been a long-standing problem. The $g,r,i$ colors of LRGs are inconsistent with stellar population models over the redshift range $0.1<z<0.7$. The $g-r$ colors in the models are on average redder than the data (of the order 0.1 mag) while the $r-i$ colors in the models are bluer (by 0.05 mag) towards low redshift. Beyond redshift 0.4, the predicted $r-i$ color becomes instead too red, while the predicted $g-r$ agrees with the data. We provide a solution to this problem, through a combination of new astrophysics and a fundamental change to the stellar population modeling. We find that the use of the empirical library of Pickles (1998), in place of theoretical libraries based on model atmosphere calculations, modifies the evolutionary population synthesis predicted colors exactly in the way suggested by the data, i.e., gives a redder $r-i$ color, and a bluer $g-r$ color, in the observed frame at $z=0.1$. The reason is a lower flux in the empirical libraries, with respect to the theoretical ones, in the wavelength range $5500-6500$~\\AA. The discrepancy increases with decreasing effective temperature independently of gravity. This result has general implications for a variety of studies from globular clusters to high-redshift galaxies. The astrophysical part of our solution regards the composition of the stellar populations of these massive Luminous Red Galaxies. We find that on top of the previous effect one needs to consider a model in which $\\sim~3\\%$ of the stellar mass is in old metal-poor stars. Other solutions such as substantial blue Horizontal Branch at high metallicity or young stellar populations can be ruled out by the data. Our new model provides a better fit to the $g-r$ and $r-i$ colors of LRGs and gives new insight into the formation histories of these most massive galaxies. Our model will also improve the {\\it k}- and evolutionary corrections for LRGs which are critical for fully exploiting present and future galaxy surveys. ",
        "introduction": "Age-dating the stellar populations of galaxies provides astronomers with a cosmic timescale which, being ruled by stellar evolution, is independent of cosmological models and allows the use of galaxies as cosmological probes \\citep[e.g.][]{jimloe02}. The interpretation of galaxy spectra in terms of stellar populations is also the only effective way of reconstructing their star-formation histories, which allows one to get clues on the still poorly-known process of galaxy formation. There is a long history of research in this area. We provide here a selection of recent works \\citep{kauetal03,thoetal05,neletal05,beretal06,panetal07,jimetal07} and refer the reader to Renzini (2006) for a comprehensive review of studies at both low and high redshift.  Studies of galaxy evolution based on the Sloan Digital Sky Survey (SDSS) data of Luminous (massive) Red Galaxies (LRGs) have highlighted a potentially fundamental problem with the rest-frame optical of stellar population models. As first noticed by \\citet{eisetal01}, the models are systematically too red in the $g-r$~observed-frame in the redshift range 0.1 to 0.4 (see Figure~1, left-hand panel). This discrepancy could not be cured by adopting more complex stellar population models with various star formation histories, different stellar population model codes or empirical galaxy spectra. To be able to analyse the data, \\cite{eisetal01} applied a shift of 0.08 mag to the models.  \\cite[][W06]{waketal06} extended this study by including the $r-i$ color as a further constraint and by adding the 2dF SDSS LRG and Quasar (2SLAQ) survey sample \\citep{canetal06} to increase the redshift range to $z\\sim 0.8$, demonstrating further discrepancies between the models and data. While the models are too red in the observed frame $g-r$ at low redshift, the $r-i$ color turns out to be too blue (see Figure~1, left-hand panel). This pattern changes with increasing redshift such that the predicted $r-i$ color becomes too red, while the predicted $g-r$ matches the data (all observed frame). Figure~1 in W06 demonstrates that the addition of star formation does not solve the problem. Indeed, the light of young, blue stars while curing the $g-r$ syndrome at the lowest redshifts, worsens the discrepancy in the $r-i$ color at low redshifts and in the $g-r$ color at high redshifts, such that colors in the model become too blue. The origin of this mismatch has remained a puzzle so far, and has serious consequences for the interpretation of current observations and the planning of future galaxy surveys. In this paper we provide a solution to this problem and present a working stellar population model.  The paper is arranged as follows. In Section 2 we recall the features of the luminous red galaxy sample. In Section 3 we present the working solution and we comment on other discarded options. A conclusive discussion is placed in Section 4. \\begin{figure*} \\centering\\includegraphics[width=0.75\\textwidth]{maraston_fig1.ps} \\caption{The $g-r$ and $r-i$ colors of LRGs as functions of redshift (points; data from W06). The median is given by the green line. Typical errors as function of redshift indicated by the error bars. {\\it Left panels}.  A solar-metallicity passively evolving, single-burst model with an age of 12 Gyr at redshift zero (red line). {\\it Middle panels}. Same data as in the left-hand panel. The stellar population model uses the Pickles (1998) empirical spectral library instead of the theoretical one (see text). {\\it Right panels}. Same data as in the left-hand panel with a composite model with $3\\%$ by mass of metal-poor stars. Both the metal-rich and the metal-poor component are 12 Gyr old at redshift zero. The metal-rich component uses the Pickles (1998) empirical spectral library. \\label{colors}} \\end{figure*} ",
        "conclusions": "In this paper we provide a working stellar population model for luminous red galaxies (LRG), which provides a substantial improvement to a long-standing mismatch with the data from the SDSS survey \\citep{eisetal01,waketal06}.  The new model adopts empirical stellar spectra from the library of Pickles~(1998) in place of the theoretical ones at solar metallicity in the evolutionary population synthesis. We found that an excess flux around $6000$~\\AA\\ rest-frame in the theoretical spectra was responsible for making the synthetic $r-i$~color too blue at redshift 0.1, and the $g-r$~ color too red. This finding impacts on a variety of studies which involve the $B-V$ and $V-R$ rest-frame and might be the cause for a mismatch between stellar population models and Milky Way Globular Cluster data in $B-V$ as highlighted in M05.  The new model further includes the addition of a small (3\\% by mass), metal-poor subcomponent, being coeval to the old and metal-rich dominant population. The metal-poor old subcomponent can better match the observed color and color evolution because of its slow evolution with redshift in comparison to a young sub-population, whose color evolution is much stronger. This model was constrained using data in the redshift range 0.1 to $\\sim~0.8$. An early version of this model was used in Cool et al.~(2008) for studying the luminosity function up to a redshift 0.8. Though the model allowed for a better match towards the highest redshifts, the synthetic $g-r$ was still too red at the lowest redshifts (cfr. Figure~13 in Cool et al. 2008), and the $r-i$ was not well recovered. This early model did not include the empirical spectra. For this reason a larger metallicity for the dominant population needed to be assumed, and as a consequence a higher fraction by mass of metal-poor stars. The use of the empirical spectra has solved the remaining problems.  For the discussion of the possible origins of a metal-poor subcomponent in massive galaxies we refer to \\cite{martho00}. It needs to be assessed whether the requirement that 3$\\%$~by mass of stars in LRGs are very metal-poor is just the reflection of the metallicity gradient known to be present in massive galaxies or is indeed caused by the presence of a very metal-poor subcomponent. The latter could originate from accretion of metal-poor dwarf satellites during the evolution of the galaxy. Future investigations on this question will be very valuable.  The models are available at www.maraston.eu."
    },
    "0809/0809.1783_arXiv.txt": {
        "abstract": "{ Low-mass satellites, like asteroids and comets, are expected to be present around the black hole at the Galactic center. We consider small bodies orbiting a black hole, and we study the evolution of their orbits due to tidal interaction with the black hole. } {In this paper we investigate the consequences of the existence of plunging orbits when a black hole is present. We are interested in finding the conditions that exist when capture occurs.} {Earlier  analysis of the evolution of classical Keplerian orbits was extended to relativistic orbits around a Schwarzschild black hole.} {The main difference between the Keplerian and black hole cases is in the existence of plunging orbits. Orbital evolution, leading from bound to plunging orbits, goes through a ``final\" unstable circular orbit. On this orbit, tidal energy is released on a characteristic black hole timescale.} {This process may be relevant for explaining how small, compact clumps of material can be brought onto plunging orbits, where they may produce individual short duration accretion events. The available energy and the characteristic timescale are consistent with energy released  and the timescale typical of Galactic flares.} ",
        "introduction": "Most galaxies, if not all, harbor massive black holes at their center. Some of them manifest their presence in a violent way and show up as active galactic nuclei; others, like the one in our Galaxy, are quiescent. The mass of the black hole at the center of the Milky Way has been estimated as $M_{\\rm bh}=3.61\\pm 0.32 \\times 10^6 M_\\odot$ from observations with SINFONI \\citep{Eisenh05} confined within a radius of 45 AU. The Galactic center is only $\\approx 8$ kpc from the Sun and therefore it is the nearest laboratory where we can study the environment of super-massive black holes in detail, see e.g.~\\citet{2005PhR...419...65A}, \\citet{2007Goldwurm}.  The standard assumption for AGNs is that gas accretes onto the black hole forming an accretion disc and releasing energy. On the other hand, some non-active galaxies exhibit short-timescale, low-intensity flaring activity \\citep{2003Natur.425..934G,2006JPhCS..54..420B,2006A&A...455....1E,2006A&A...460...15M}, which could be the result of individual accretion events of small bodies onto the central black hole.  It seems reasonable to assume that stars at the Galactic center are surrounded by planets and by other small orbiting bodies, like asteroids and comets. Such low-mass satellites (LMS) may be stripped off their parent stars by tidal interaction, while they approach the black hole. In this way they contribute to a distribution of low-mass objects that cluster the central black hole. According to \\citet{1995ApJ...455..342C}, the Edgeworth-Kuiper belt of our Solar System may still contain as many as $2\\times 10^8$ objects with radii $\\lesssim$ 10~km. One might therefore expect that LMS will be copiously present all the way down to the black hole.  In this paper we study the evolution of orbits of LMS caused by tidal interaction with the black hole. We extend the analysis presented by \\citet{1980A&A....92..167H,1981A&A....99..126H,1982A&A...110...37H} concerning the tidal evolution of close binary systems in the weak friction model in the Newtonian approximation. We show that orbital evolution of these objects leads to capturing orbits by the black hole and argue that tidal interaction can rapidly inject enough energy to explain the energy source of Galactic flares. ",
        "conclusions": "In this paper we have investigated the fate of small bodies, like asteroids and comets, that may find themselves on highly eccentric orbits in the vicinity of the Galactic black hole. Extending Hut's analysis of the tidal evolution of close binary systems, we find that they experience very high tides at their periastra, which heat and eventually liquefy them. Those objects that are electrically conductive are likely to have enough magnetic rigidity to efficiently transfer angular momentum from orbit to spin, so that the loss of orbital angular momentum leads to smooth transition from bound to plunging orbits. We propose this as a mechanism that could bring relatively small and condensed clumps of material to the vicinity of the black hole. Such clumps can produce individual accretion events characterized by the time scale of the last circular orbit. The tidal energy released during this process can reach up to $0.1\\ mc^2$. This means that both the available energy and the characteristic timescale are consistent with the energy and timescale characteristic of Galactic flares."
    },
    "0809/0809.3265_arXiv.txt": {
        "abstract": "{Circumstellar shells around AGB stars are built over long periods of time that may reach several million years. They may therefore be extended over large sizes ($\\sim$ 1 pc, possibly more), and different complementary tracers are needed to describe their global properties.} {We set up a program to explore the properties of matter in the external parts of circumstellar shells around AGB stars and to relate them to those of the central sources (inner shells and stellar atmospheres).} {In the present work, we combined 21-cm \\HI and CO rotational line data obtained on an oxygen-rich semi-regular variable, RX Lep, to describe the global properties of its circumstellar environment.} {With the SEST, we detected the CO(2-1) rotational line from RX Lep. The line profile is parabolic and implies an expansion velocity of $\\sim$ 4.2 \\kms ~and a mass-loss rate $\\sim$ 1.7\\,$\\times$\\,10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$ (d $=$ 137 pc). The \\HI line at 21 cm was detected with the Nan\\c{c}ay Radiotelescope on the star position and at several offset positions. The linear shell size is relatively small, $\\sim$ 0.1 pc, but we detect a trail extending southward to $\\sim$ 0.5 pc. The line profiles are approximately Gaussian with an FWHM $\\sim$ 3.8 \\kms ~and interpreted with a model developed for the detached shell around the carbon-rich AGB star Y CVn. Our \\HI spectra are well-reproduced by assuming a constant outflow ($\\rm \\dot{M}$ $=$ 1.65 $\\times$ 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$) of $\\sim$ 4 $\\times$ 10$^{4}$ years duration, which has been slowed down by the external medium. The spatial offset of the \\HI source is consistent with the northward direction of the proper motion measured by Hipparcos, lending support to the presence of a trail resulting from the motion of the source through the ISM, as already suggested for Mira, RS Cnc, and other sources detected in \\HI. The source was also observed in SiO (3 mm) and OH (18 cm), but not detected. } {A detached shell, similar to the one around Y CVn, was discovered in \\HI around RX Lep. We also found evidence of an extension in the direction opposite to the star proper motion. The properties of the external parts of circumstellar shells around AGB stars should be dominated by the interaction between stellar outflows and external matter for oxygen-rich, as well as for carbon-rich, sources, and the 21-cm \\HI line provides a very useful tracer of these regions.} ",
        "introduction": "\\label{intro} Evolved stars on the asymptotic giant branch (AGB) are often surrounded by circumstellar shells. The material in these shells is flowing outwards with velocities from a few~\\kms~up to 40 km s$^{-1}$ (Nyman et al. 1992). The observed mass-loss rates range from $\\sim$ 10$^{-8}$ to a few 10$^{-4}$ M$_{\\odot}$ yr$^{-1}$ (e.g. Knapp \\& Morris 1985; Olofsson et al. 2002), the lower limit being probably set by detectability. In this phase of the stellar life, the evolution is dominated by mass loss rather than nuclear processes (Olofsson 1999). The history of mass loss over the full AGB is complex and the details of this process are currently not well known (e.g., Lafon \\& Berruyer 1991; Habing 1996; Le{\\~a}o et al. 2006). A general picture, however, has arisen from both theoretical and observational findings that - on average - the mass-loss rate increases towards the end of the AGB phase, leading in some cases to the formation of a planetary nebula (e.g., Renzini 1981; Hrivnak \\& Bieging 2005).  The validity of this simple picture may depend on the parameters of the star, e.g., on its initial mass. Schr{\\\"o}der et al. (1999) combine mass-loss rates derived from consistent wind models with stellar evolution calculations and find that the mass-loss rate should increase along the AGB for stars with initial masses greater than 1.3 M$_{\\odot}$. Stars with lower initial mass would experience a single short-lived ($\\sim$ 1000 yr) episode of high mass loss only, which would leave behind a very narrow detached shell as observed in the case of, e.g., TT Cyg (Olofsson et al. 2000). On the other hand, the mass-loss phenomenon appears to be highly variable on even shorter time scales as indicated e.g. by concentric arcs observed in scattered light around the prototype carbon Mira IRC~+10216 (Mauron \\& Huggins 2000) and around some proto-planetary nebulae (e.g., Hrivnak et al. 2001). The time scale of these mass-loss variations would be around a few 10$^2$\\,yr. The physical mechanism responsible for these variations still needs to be identified, although different possibilities have already been proposed: e.g., interaction between gas and dust within stellar outflows (Simis et al. 2001), or solar-like magnetic cycle (Soker 2002). In contrast to these later phases of AGB mass loss at rather high rates ($\\sim$ 10$^{-5}\\rm M_{\\odot}$ yr$^{-1}$), information about the mass-loss process on the early AGB, is even scarcer.  To unravel the processes involved in the mass-loss phenomenon, we have to find suitable tracers. One of the most studied among these tracers is the CO molecule, because so far it has been considered to provide the best estimate of the mass-loss rate for AGB stars (Ramstedt et al. 2008). Not only can it be used to estimate this mass-loss rate, but it also yields important parameters of the AGB wind (e.g.: expansion velocity, central star velocity, etc.). Since CO is photodissociated by UV radiation from the interstellar radiation field (ISRF), it can only probe the inner parts (r $\\leq$ 10$^{-3}$ - 10$^{-1}$ pc, Mamon et al. 1988) of circumstellar shells (CSs). Therefore, the CO emission is only related to ``recent'' (i.e. a few 10$^{3}$ - 10$^4$ years) mass-loss episodes.  On the other hand, \\HI is in general protected from photoionization by the surrounding interstellar medium (ISM). As a result, \\HI can be used to probe the external parts of circumstellar shells and can give indications on the mass-loss on longer timescales (a few 10$^5$ years: Libert et al. 2007). Hence, CO and \\HI complement each other nicely to describe the history of the mass-loss rate of an AGB star.  The drawbacks of \\HI circumstellar observations are that hydrogen is ubiquitous in the Galaxy and that the genuine stellar \\HI must be separated from the ambient H\\,{\\sc {i}}. Ideal cases would be bright \\HI sources, with relatively high velocity with respect to the local standard of rest (LSR) and reasonably far above the Galactic plane. For the other sources, the interstellar \\HI should be studied with care. In the present paper, we analyze this confusion with a new approach that consists in a 3D-mapping using \\HI spectra.  The mass-loss phenomenon is different from one AGB star to another and may vary highly with time. Nevertheless, observing in \\HI provides a global view of the CS behavior and, on timescales of about 10$^5$ years, small variations in the mass-loss rate may be flattened out. Thus, we have developed a model of the circumstellar gas, based on a scenario already proposed by Young et al. (1993), in which CSs are the result of a constant outflow eventually slowed down by the surrounding medium. This deceleration produces a snowplough effect around the source, resulting in a detached shell of compressed matter originating from the star and the external medium (Lamers \\& Cassinelli 1999, Chap. 12).  A schematic view of this model can be pictured as follows: a wind is flowing outward from the star, in free expansion with a constant velocity (V$\\rm _{exp}$) and a constant rate. It encounters a shock at a radius r$_1$ (termination shock), due to the slowing down by the surrounding matter. Between r$_1$ and r$\\rm _f$ (contact discontinuity), the stellar matter is compressed. Between r$\\rm _f$ and until a second shock at r$_2$ (bow shock), the interstellar matter has been swept up by the wind of the AGB star. Finally, beyond r$_2$, the external matter is considered to be at rest.  Recently, we successfully applied this model to a carbon-rich star: Y CVn (Libert et al. 2007). In \\HI at 21 cm, this star exhibits a composite profile, made of a broad, rectangular component and a narrow, Gaussian-shaped one. In our description, the broad component is the signature of the freely expanding wind, whereas the narrow component is produced by the \\HI compressed in the snowplough between r$_1$ and r$_2$. Our model provides a simple explanation for some of the so-called ``detached dust shells'' observed in the far infrared (Izumiura et al. 1996). If this approach is correct, then it should also apply to detached shells around oxygen-rich AGB stars. In this paper we present \\HI and CO data that we obtained on an oxygen-rich AGB star, RX Lep, and interpret them with the model that we developed for Y CVn. The \\HI interstellar confusion in the direction of RX Lep is moderate and we illustrate, in that case, our new approach to extract a genuine \\HI spectrum.  In this simplified description we assume spherical symmetry. However, recent H\\,{\\sc {i}}, far-infrared and UV data (G\\'erard \\& Le\\,Bertre 2006; Matthews \\& Reid 2007; Ueta et al. 2006; Martin et al. 2007) have shown that the AGB star motion with respect to the ISM may lead to a distortion, and eventually a disruption, of the circumstellar environment. Previous, and more recent, numerical modelings (Villaver et al. 2003; Wareing et al. 2007) are in line with this interpretation. The circumstellar environment of RX Lep might provide a new illustration of this phenomenon, and we will discuss this possibility. ",
        "conclusions": "We detected CO(2-1) and \\HI line emissions from the semi-regular oxygen-rich E-AGB star, RX Lep. These emissions indicate a stellar outflow at a velocity $\\sim$ 4.2 \\kms~and a rate $\\sim$\\,1.7\\,$\\times$\\,10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$, with a duration of 4 $\\times$ 10$^{4}$ years. The \\HI source has a size of $\\sim$ 2$'$ ($\\approx$ 0.08 pc) in the east-west direction and possibly 15$'$ ($\\approx$ 0.6 pc) in the north-south direction.  The modeling of the \\HI line profiles obtained at different positions suggests that the outflow is slowed down by the interaction with the ambient ISM, and that the external part of RX Lep circumstellar shell is made of compressed material, at $\\sim$ 200 K, as in the well-known detached shell around Y CVn.  The elongated shape of the RX Lep \\HI source is compatible with the direction of its proper motion, as in the cases of Mira and RS Cnc, which have already been studied at high angular resolution with the VLA."
    },
    "0809/0809.5055_arXiv.txt": {
        "abstract": "We discuss how lensing by magnetic fields in galaxy clusters affects ultrahigh energy cosmic ray (UHECR) observations. As specific example, we use Virgo together with the cluster magnetic fields obtained earlier in a constrained simulation of structure formation including MHD processes. We find that, if M87 is the single source of UHECRs from Virgo, the emitted flux is strongly anisotropic in the most interesting energy range, (50--100)\\,EeV, and differs from the average value by a factor five or more for a significant fraction of observers. Since magnetic lensing is energy dependent, the external energy spectrum as seen by different observers varies strongly too. These anisotropies are averaged out in the case that all active galactic nuclei in Virgo emit UHECRs. In both cases, the anisotropies of the emitted UHECR flux may introduce an important bias in the interpretation of UHECR data like, e.g., the determination of the source density $n_s$ and the source energy spectrum of UHECRs. ",
        "introduction": "The identification of the sources of ultrahigh energy cosmic rays (UHECRs) is one of the most longstanding problems of astrophysics. Only recently, evidence has been accumulating that we are at the dawn of ``charged particle astronomy.'' A first piece of evidence have been the observation of the Greisen-Zatsepin-Kuzmin cutoff~\\cite{GZK} in the UHECR spectrum  with $5 \\sigma$ significance by the HiRes experiment~\\cite{spectrum_HiRes} which was recently confirmed by the AUGER experiment with $6\\sigma$ significance~\\cite{auger_spectrum2008}. Secondly,  anisotropies on medium scales have been found previously combining all available data of ``old'' CR experiments~\\cite{msc}. The data from the Pierre Auger Observatory (AUGER) presented in Ref.~\\cite{ICRC} confirmed these findings, showing also a surplus of clustering compared to an isotropic distribution in the broad range from 7 to 30~degrees. Such anisotropies were predicted as a consequence of the observed large-scale structure of matter~\\cite{old} and favor therefore together with the presence of the Greisen-Zatsepin-Kuzmin cutoff the extragalactic origin of UHECRs. Comparing the expected angular auto-correlation function of different sources with the data, the authors of Ref.~\\cite{ATF} suggested that the sources of UHECRs are active galactic nuclei (AGN) or another strongly clustered sub-sample of galaxies. Finally, the Auger collaboration reported  evidence for an anisotropy of the UHECR arrival directions, observing a correlation of the arrival directions of UHECRs with the positions of AGN or similarly clustered matter in the nearby Universe~\\cite{auger_corr}. At present the significance of this correlation is only a the $3\\sigma$ confidence level and is not confirmed by the data from the HiRes experiment~\\cite{HiRes}.  Most (auto-) correlation analyses assume identical sources, emitting isotropically UHECRs. Here we show that these assumptions are not fulfilled in the case that UHECR sources are located inside regions with relatively strong magnetic fields as in the core of galaxy clusters. As a concrete example we use the Virgo cluster and investigate how lensing effects of magnetic fields extending on cluster scales influence the emitted flux and energy spectrum both  for isotropic and jet-like emission of single sources. After a brief discussion of cluster magnetic fields in Sec.~2, we examine in Sec.~3 first the anisotropies in the case of a single isotropic source, then of jet-like emission by a single source, of several isotropic sources and finally the resulting modifications of the observed CR energy spectrum. We conclude in Sec.~4. ",
        "conclusions": "The amplification of the UHECR flux by magnetic lensing was discussed previously in the context of the Galactic regular and turbulent field~\\cite{galaxy,Kachelriess:2005qm}. For instance, Ref.~\\cite{Kachelriess:2005qm} calculated exposure maps for three different models of the Galactic regular magnetic field and found also significant anisotropies at $E=40\\,$EeV. The strength of the magnetic field close to the core of galaxy clusters is smaller by a factor 30--100 than in the Milky Way. The weakness of cluster fields is, however, over-compensated by their much larger extension compared to the Galactic disc. While the deflection of UHECRs in cluster fields leads even for the nearest cluster only to maximal angular differences between the source and the arrival direction of order $1.5^\\circ$ and is thus negligible, the resulting anisotropies of the emitted CR flux may introduce an important bias in the interpretation of CR data.   We have found that assuming magnetic fields as in Ref.~\\cite{weak} the flux from M87 differs from the average value by a factor five for a significant fraction of observers in the most interesting energy region, (50--100)\\,EeV. This may be the explanation for the non-observation of UHECR from the Virgo cluster, if M87 is the dominant UHECR source in the Virgo cluster. More generally, anisotropies induced by magnetic fields result in a reduction of the effective number of sources, leading to a bias in auto-correlation studies. Vice versa, if the source type of UHECRs and thus their number density is known, a comparison with the estimated value from auto-correlation studies informs us about the importance of anisotropies and/or beaming effects. Note that magnetic lensing will affect also sources behind the Virgo cluster.   The second important consequence of the anisotropies induced by the cluster magnetic fields is the resulting modulation of the energy spectrum of UHECRs leaving the galaxy cluster compared to their generation spectrum. A similar energy modulation is be introduced by the Galactic magnetic field and both effects will complicate the reconstruction of generation spectrum, if UHECRs sources are discovered.   Finally we want to stress that we obtained all our results for two specific realizations of magnetic fields obtained within the simulations of Ref.~\\cite{weak}. Although we expect that the main features of our results are robust, quantitative differences will appear using different simulations of cluster magnetic fields: For instance, if the volume fraction filled by relatively large magnetic fields around clusters is as large as in Ref.~\\cite{strong}, we expect that the same phenomena happen as found in this work, but extending over a larger energy range."
    },
    "0809/0809.2505_arXiv.txt": {
        "abstract": "{} {We investigate the mineralogy and dust processing in the circumbinary discs of binary post-AGB stars using high-resolution TIMMI2 and SPITZER infrared spectra. } {We perform a full spectral fitting to the infrared spectra using the most recent opacities of amorphous and crystalline dust species. This allows for the identification of the carriers of the different emission bands. Our fits also constrain the physical properties of different dust species and grain sizes responsible for the observed emission features.  } {In all stars the dust is oxygen-rich: amorphous and crystalline silicate dust species prevail and no features of a carbon-rich component can be found, the exception being EP\\,Lyr, where a mixed chemistry of both oxygen- and carbon-rich species is found. Our full spectral fitting indicates a high degree of dust grain processing. The mineralogy of our sample stars shows that the dust is constituted of irregularly shaped and relatively large grains, with typical grain sizes larger than 2\\,$\\mu$m. The spectra of nearly all stars show a high degree of crystallinity, where magnesium-rich end members of olivine and pyroxene silicates dominate. Other dust features of e.g. silica or alumina are not present at detectable levels. Temperature estimates from our fitting routine show that a significant fraction of grains must be cool, significantly cooler than the glass temperature. This shows that radial mixing is very efficient is these discs and/or indicates different thermal conditions at grain formation. Our results show that strong grain processing is not limited to young stellar objects and that the physical processes occurring in the discs are very similar to those in protoplanetary discs.} {} ",
        "introduction": "Post-AGB stars are low- and intermediate-mass stars that evolve rapidly from the Asymptotic Giant Branch (AGB) towards the Planetery Nebula (PNe) phase, before cooling down as white dwarves. During the previous AGB phase, the star has undergone severe mass loss, leaving behind a slowly expanding circumstellar dust shell. Depending on the dilution in the line-of-sight, the expected spectral energy distribution (SED) is double-peaked, with the first UV-visible peak indicative of the central star and the second infrared peak coming from thermal emission of the cool dust in the circumstellar environment (CE). The release of the IRAS All-Sky Survey mission made a large-scale identification of post-AGB stars on the basis of their infrared colours possible \\citep{lloydevans85,hrivnak89,oudmaijer92}. More recent surveys on post-AGB stars include \\citet{szczerba07}, who present the latest catalogue of Galactic post-AGB stars with almost 400 objects.  IRAS colour-colour diagrams revealed a large sample of stars that did not show the expected double-peaked SED \\citep{lloydevans85,deruyter06}. Instead, they show a strong near-IR excess, pointing to the presence of hot dust in the system, while the central stars are currently too hot to have an ongoing dusty mass loss. Correlations with optical databases of variable stars point to the presence of pulsating stars that also show this specific IR-excess. These RV\\,Tauri stars are luminous evolved stars that cross the Population II Cepheid instability strip and populate a well-defined part of the IRAS colour-colour diagram \\citep{lloydevans85,raveendran89,lloydevans99}. Recent interferometric studies \\citep{deroo07b,deroo07a} prove that the circumstellar emission originates from a very compact region and that the SEDs of these stars are well modelled with a passive 2D disc model \\citep{dullemond04,deroo07b,gielen07}. The spectral slope of the submillimetre SED points to the presence of large, up to $\\mu$m and cm sized, grains in these discs \\citep{deruyter06,gielen07}, which have settled to form a cool disc midplane. The inner radius of these discs is determined by the dust sublimation radius. The hot inner rim is puffed up by gas presssure from the central star and radiates mainly in the near-IR, while the outer parts of the disc can be strongly flared and are responsible for the strong absorption and re-radiation of the stellar light. In a few cases the direct confirmation of the Keplerian kinematics of the circumstellar discs comes from resolved interferometric CO measurements \\citep{bujarrabal88,bujarrabal05,bujarrabal07}. Recent studies have shown that these post-AGB stars are very likely all binaries \\citep{maas02,maas03,vanwinckel99,vanwinckel07,gielen07}.  Famous examples of these objects include HD\\,44179, the central star of the Red Rectangle. It is a carbon-rich post-AGB object for which the binarity was proven by \\citet{vanwinckel95}. The star is surrounded by an oxygen-rich disc \\citep{waters98} and is resolved in ground-based high spatial resolution imaging \\citep{roddier95,menshchikov02} as well as in HST optical images \\citep{cohen04}. The disc is also resolved in interferometric CO measurements, where the Keplerian velocity of the disc was detected \\citep{bujarrabal05}.  Infrared spectroscopy is the ideal tool to study the physico-chemical characteristics of the circumstellar material, since it samples resonances of the most dominant dust species and important ro-vibrational bands of dominant molecules. Spectroscopic data obtained with the Infrared Space Observatory (ISO) allowed for the first time to study the mineralogy of circumstellar environments of objects in different evolutionary stages \\citep[e.g.][]{waters96,waelkens96}. More advanced studies of the CE of young stellar objects \\citep[e.g.][]{vanboekel05,lisse07} and mass-losing objects \\citep{molster02a,molster02b,molster02c} further identified the dominant dust species and grain sizes of the circumstellar dust.  Infrared spectral studies of post-AGB objects \\citep[][and references therein]{vanwinckel03} allowed for the detection of different CE chemistries. The more typical C-rich post-AGB stars are characterised by a strong amorphous SiC feature at 11.3\\,$\\mu$m. In these stars, a strong unidentified 21\\,$\\mu$m feature can sometimes be found \\citep{kwok89}. In addition to the detection of O-rich post-AGB stars, dominated by silicate species, surprisingly also mixed chemistries were found, showing features of both O-rich and C-rich dust species.  Young stellar objects are the natural environment to study circumstellar disc physics, but previous studies of the mineralogy of discs around the few brightest infrared evolved objects  with ISO \\citep{molster99,molster02a,molster02b,molster02c} show strong dust processing, with oxygen-rich and highly crystalline dust. \\citet{molster02a,molster02b,molster02c} found that a high degree of crystallisation is indicative for the presence of a stable disc, and not a dusty outflow.  In our pilot study \\citep{gielen07}, based on SPITZER-IRS and ISO data, we investigate the mineralogy and spectral energy distribution of two post-AGB stars, RU\\,Cen and AC\\,Her. The spectra of these stars are extremely similar and show a strong resemblance to the infrared spectrum of the solar-system comet Hale-Bopp and young stellar objects, such as HD\\,100546. The observed crystalline emission features are well modelled using magnesium-rich crystalline silicates. The grains are irregularly shaped, with grain sizes of 0.1\\,$\\mu$m and 1.5\\,$\\mu$m. Both hot and cool grains are necessary to reproduce the observed spectrum. The spectral energy distributions of both stars were fitted using a 2D passive disc model \\citep{dullemond04}. The discs surrounding these objects start at 35\\,AU from the central star, which is well beyond the dust sublimation radius. The outer radius is less well constrained, but extends till about 300\\,AU. The discs are strongly flared.  In this work we study the mineralogy of the discs in a large sample of post-AGB binaries and perform a detailed fitting of the observed emission features. We have observed 21 binary post-AGB objects with the SPITZER-IRS spectrograph and look for relations between the dust parameters and stellar characteristics.  The outline of this paper is as follows: in Sects.~\\ref{programstars} and ~\\ref{observations} we introduce our programme stars and the different observation and reduction strategies used. Section~\\ref{sed} contains the construction of the spectral energy distributions and the total extinction determination. In Sect.~\\ref{general} we give a general overview of the spectra and the observed emission features. The profile and full spectral fitting is presented in Sects.~\\ref{profilefit} and ~\\ref{fullfit}. The discussion and conclusion of our analysis are given in Sects.~\\ref{discussion} and \\ref{conclusions}. ",
        "conclusions": "  \\begin{itemize}  \\item Almost all discs display only O-rich spectral signatures. The noticeable exception is EP\\,Lyr, which shows a very similar spectrum as the central star of the Red Rectangle. PAH emission features dominate the spectrum till 20\\,$\\mu$m, at larger wavelengths crystalline silicates start to dominate.  \\item Our mineralogy study indicates the dominance of Mg-rich amorphous and crystalline silicate dust in the disc. The high crystallinity and the large fraction of large grains, as deduced from our full spectral fitting, show strong dust grain processing in the discs.  \\item The temperature estimates from our fitting routine show that a significant fraction of crystalline grains must be cool. This shows that radial mixing is efficient is these discs or indicate a different thermal history at grain formation.  \\item Trend analysis of our fitting parameters show no clear correlation with stellar characteristics. For the moment it is not clear if and how the observed diversity in observed spectra relates to specific structural elements of the disc, the star and/or the orbits or whether we witness directly an evolutionary change between different sources.  \\end{itemize}  To further improve our understanding of these circumbinary discs, as a next step, we will combine our photometric and spectroscopic data with interferometric measurements. Comparing spatial information from the MIDI and AMBER interferometric instruments with a realistic disc model \\citep{dullemond04}, constrained by photometric and spectroscopic data, will allow us, not only to study the mineralogy, but also the structure of the discs. So far, interferometric measurements for five of our sample stars have been obtained \\citep{deroo06,deroo07c}.  Probing the dust processing in the discs around evolved objects proves to be an excellent complement to study physics in planet-forming young discs."
    },
    "0809/0809.3735_arXiv.txt": {
        "abstract": "% M16 (the Eagle Nebula) is a striking star forming region, with a complex morphology of gas and dust sculpted by the massive stars in NGC\\,6611. Detailed studies of the famous ``elephant trunks'' dramatically increased our understanding of the massive star feedback into the parent molecular cloud. A rich young stellar population (2$-$3\\,Myr) has been identified, from massive O-stars down to substellar masses. Deep into the remnant molecular material, embedded protostars, Herbig-Haro objects and maser sources bear evidence of ongoing star formation in the nebula, possibly triggered by the massive cluster members. M\\,16 is a excellent template for the study of star formation under the hostile environment created by massive O-stars. This review aims at providing an observational overview not only of the young stellar population but also of the gas remnant of the star formation process. ",
        "introduction": "\\begin{figure}[ht] \\begin{center} \\includegraphics[scale=0.65,angle=0]{eagle_nebula_f1_s.eps} \\end{center} \\caption{Wide-field image (approximately 40\\,arcmin$\\times$41\\,arcmin) of the Eagle Nebula, showing the blister H\\,{\\sc ii} region created by the massive stars in NGC\\,6611. Towards the center of the image the dusty pillars are clearly seen (labelled from I to V). This image was taken with the 0.9\\,m telescope at the Kitt Peak Observatory with the NOAO Mosaic CCD camera and it combines H$\\alpha$\\, (green), [O\\,{\\sc iii}] (blue) and [S\\,{\\sc ii}] (red) images. Credit: T.A.\\,Rector, B.A.\\,Wolpa and NOAO.} \\label{m16} \\end{figure}  The cluster Messier\\,16 (M\\,16), in the constellation Serpens Cauda, was first discovered in 1745 by Jean-Philippe Loys de Cheseaux, a Swiss astronomer from Lausanne, who in 1746 presented to the French Academy of Science a list of clusters and nebulae, including M16. On June 3, 1764, Charles Messier independently discovered the cluster, noting its nebulous nature, and gave it the number M16 by which it is now known. It was later included in John Herschel's catalogue of clusters and nebulae as h2006, and eventually received the designation NGC\\,6611 when listed in Dreyer's New General Catalogue. In the popular literature, it is often known as the Eagle Nebula.  The massive stars in the young cluster NGC\\,6611 \\citep{walker61} are responsible for the ionization of the H\\,{\\sc ii} region identified as Sh2$-$49 \\citep{sharpless59}, Gum\\,83 \\citep{gum55} or RCW\\,165 \\citep*{rodgers60} in the optical or W\\,37 \\citep{westerhout58} as observed at radio wavelengths. These stars have been photoevaporating the surrounding parent molecular cloud and sculpting the overdense molecular cores into the famous ``elephant trunks'' (Fig.\\,\\ref{m16}). First identified by \\citet{duncan20}, these structures were revealed in their full glory in the iconic HST images of \\citet{hester96}; they appear in the optical images as opaque fingers or columns of dense obscuring material projected against the diffuse background nebular emission. Even if not as dense as originally thought \\citep*{thompson02}, these columns harbour small protrusions on or near their surface, the EGGs, a fraction of which seem to be associated with embedded Young Stellar Objects (YSOs), signs of a recent episode of star formation \\citep[e.g.][]{mccaughrean02}.  NGC\\,6611 cluster members are distributed over a region of $\\sim$14\\,arcmin radius, with a higher concentration in the largely unobscured 4\\,arcmin radius central area \\citep{belikov00,kharchenko05}. The cluster contains a large number of massive stars as well as a large population of pre-main sequence (PMS) stars \\citep{hillenbrand93}. These optically visible members of NGC\\,6611, with masses in the range 2 to 85\\,M$_{\\odot}$, have an average age of 2$-$3\\,Myr \\citep{hillenbrand93,belikov00}. More recently a rich low-mass PMS population (down to substellar masses) was identified by \\citet{oliveira05} and \\citet{oliveira08}.  The Eagle Nebula is a superb example of the interaction of massive stars with their environment, showcasing both its destructive power and its potential for triggering star formation at the periphery of the remnant molecular cloud. ",
        "conclusions": ""
    },
    "0809/0809.4265_arXiv.txt": {
        "abstract": "The {\\it HST} proper motion (PM) measurements of the Clouds have severe implications for their interaction history with the Milky Way (MW) and with each other. The Clouds are likely on their first passage about the MW and the SMC's orbit about the LMC is better described as quasi-periodic rather than circular. Binary L/SMC orbits that satisfy observational constraints on their mutual interaction history (e.g. the formation of the Magellanic Bridge during a collision between the Clouds $\\sim$300 Myr ago) can be located within 1$\\sigma$ of the mean PMs. However, these binary orbits are not co-located with the Magellanic Stream (MS) when projected on the plane of the sky and the line-of-sight velocity gradient along the LMC's orbit is significantly steeper than that along the MS.  These combined results ultimately rule out a purely tidal origin for the MS: tides are ineffective without multiple pericentric passages and can neither decrease the velocity gradient nor explain the offset stream in a polar orbit configuration. Alternatively, ram pressure stripping of an extended gaseous disk may naturally explain the deviation. The offset also suggests that observations of the little-explored region between RA 21$^\\mathrm{h}$ and 23$^\\mathrm{h}$ are crucial for characterizing the full extent of the MS. ",
        "introduction": "The recent high-precision proper motion (PM) measurements of the L/SMC determined by \\cite[Kallivayalil et al. (2006a, 2006b - hereafter K06a and K06b; see also these proceedings)]{K06a} imply that the Magellanic Clouds are moving $\\sim$100 km/s faster than previously estimated and now approach the escape velocity of the Milky Way (MW). \\cite[Besla et al. (2007)]{B07} (hereafter B07) re-examined the orbital history of the Clouds using the new PMs and a $\\Lambda$CDM-motivated MW model and found that the L/SMC are either on their first passage about the MW or, if the mass of the MW is $> 2\\times10^{12}M_\\odot$, that their orbital period and apogalacticon distance are a factor of three larger than previously estimated. This means that models of the Magellanic Stream (MS) need to reconcile the fact that although the efficient removal of material via tides and/or ram pressure requires multiple pericentric passages through regions of high gas density, the PMs imply that the Clouds did not pass through perigalacticon during the past $\\ge$5 Gyr (this is true even if a high mass MW model is adopted). While the most dramatic consequence of the new PMs is the limit they place on the interaction timescale of the Clouds with the MW, there are a number of other equally disconcerting implications: the relative velocity between the Clouds has increased such that only a small fraction of the orbits within the PM error space allow for stable binary L/SMC orbits (K06b, B07); the velocity gradient along the orbit is much steeper than that observed along the MS; and the past orbits are not co-located with the MS on the plane of the sky (B07). In these proceedings the listed factors are further explored and used to argue that the MS is not a tidal tail. ",
        "conclusions": "The new PMs have dramatic implications for phenomenological studies of the Clouds that assume they have undergone multiple pericentric passages about the MW and/or that the SMC is in a circular orbit about the LMC. The orbits deviate spatially from the current location of the MS on the plane of the sky and the velocity gradient along the orbit is much steeper than that observed. These results effectively rule out a purely tidal model for the MS and lend support for hydrodynamical models, such as ram pressure stripping.  The offset further suggests that the Clouds have travelled across the little explored region between RA 21$^\\mathrm{h}$ and 23$^\\mathrm{h}$ (i.e. the region spanned by the blue lines in Fig.\\ \\ref{fig3}). \\cite[Putman et al. (2003)]{P03} detected diffuse HI in that region that follow similar velocity gradients as the main stream (their Fig. 7), but otherwise material in that region has been largely ignored by observers and theorists alike. The offset orbits suggest that the MS may be significantly more extended than previously believed and further observations along the region of sky they trace are warranted."
    },
    "0809/0809.1986_arXiv.txt": {
        "abstract": "{For many years modeling of double-mode pulsation of classical pulsators was a challenging problem. Inclusion of turbulent convection into pulsation hydrocodes finally led to stable double-mode models. However, it was never analysed, which factor of turbulent convection is crucial. We show that the double-mode behaviour displayed in the computed models results from incorrect assumptions adopted in some of the pulsation hydrocodes, namely from the neglect of buoyant forces in convectively stable layers. This leads to significant turbulent energies and consequently to strong eddy-viscous damping in deep, convectively stable layers of the model. Resulting differential reduction of fundamental and first overtone amplitudes favours the occurrence of double-mode pulsation. Once buoyant forces in convectively stable regions are taken into account (as they should), no stable double-mode behaviour is found. The problem of modeling double-mode behaviour of classical pulsators remains open.}{hydrodynamics -- convection -- overshooting -- stars: oscillations  -- stars: variables: Cepheids -- methods: numerical} ",
        "introduction": "%  Among classical pulsators (Cepheids and RR~Lyrae stars) double-mode pulsators (or beat pulsators) represent a small, but very interesting group. These stars pulsate simultaneously in two, low-order radial modes of pulsation, either in fundamental and first overtone (F/O1) or in two lowest overtones (O1/O2). Their astrophysical importance results from the fact, that two radial modes of pulsation simultaneously observed strongly constrain the physical parameters of the star. More than 30 years ago Petersen (1973) showed, that simple linear theory may be used to calculate the masses of such pulsators. Disagreement between these so-called beat masses and masses derived from evolutionary calculations was one of the factors motivating revision of the opacity tables (Simon 1982). Indeed, with improved physics, new opacity tables solved the beat-mass discrepancy problem (Moskalik, Buchler \\& Marom 1992). Beat pulsators proved to be a powerful diagnostic tool for our knowledge of physical processes acting in stars. Recently Kov\\'acs (2000a,b) used double-mode Cepheids and RR~Lyrae stars to derive the distance modulus to the Magellanic Clouds. More detailed study of the Petersen diagram for double-mode RR~Lyrae stars in Magellanic Clouds was performed by Popielski, Dziembowski \\& Cassisi (2000). Beaulieu \\etal (2006) used beat Cepheids to study the metallicity distribution in M33 galaxy. Their methods were further developed by Buchler \\& Szab\\'o (2007), who show how the metallicity of beat Cepheid may be constrained, using the Petersen diagram.  Results mentioned in the previous paragraph, are based on computations done with linear pulsation hydrocodes. In these codes stability of the model against small perturbations is studied. As a result the pulsation eigenmodes, their periods and growth rates are calculated. Necessary condition for beat pulsation to occur is that the two pulsation modes are simultaneously linearly unstable, that is their linear growth rates, $\\gamma$, are positive. For F/O1 beat pulsations one requires that $\\gamma_0>0$ and $\\gamma_1>0$. These conditions are satisfied in a wide range of astrophysical parameters, covering a significant part of the instability strip. However, the relative scarcity of the beat pulsators in comparison to single-mode pulsators indicate, that the double-mode pulsation domain is much smaller, and is determined by nonlinear effects. Nonlinear computations may further constrain the astrophysical parameters of the beat pulsators. Therefore, from the very beginning of nonlinear pulsation computations, modeling of the beat phenomenon was one of the major objectives.  Early attempts to model the beat phenomenon with direct time integration, radiative hydrocodes were inconclusive. Limited computer resources didn't allow to perform long-lasting integrations. Mixed-mode states persisting for tens of periods were observed, however, their permanent double-mode nature could not be confirmed. Introduction of relaxation technique (Baker \\& von Sengbush 1969, Stellingwerf 1974) greatly facilitated the search for double-mode phenomenon. Relaxation scheme allows for fast convergence to a limit cycle (monoperiodic, finite-amplitude oscillation) and additionally provides information about limit cycle stability through the Floquet stability coefficients. Two stability coefficients are of interest, when one searches for F/O1 double-mode pulsation. $\\eta_{0,1}$ measures the stability of the first overtone limit cycle against the fundamental mode perturbation (switching rate toward fundamental mode), while $\\eta_{1,0}$ measures the stability of the fundamental mode limit cycle against the first overtone perturbation (switching rate toward first overtone). Positive value of a Floquet coefficient means, that the respective mode is unstable and will switch into the other mode. If both $\\eta_{0,1}$ and $\\eta_{1,0}$ are positive, double-mode state is unavoidable.  Using relaxation technique several surveys of radiative models were performed to search for the beat pulsation. Examples of resonant, F/O1 doubly-periodic, however, triple-mode models were found by Kov\\'acs \\& Buchler (1988) (RR~Lyrae models), Buchler, Moskalik \\& Kov\\'acs (1990) and Smolec (2008) ($\\delta$~Cephei models). In all these cases, origin of the doubly-periodic pulsation is connected with the 2:1 resonance between the linearly unstable fundamental mode, and linearly damped higher order overtone (third overtone in case of RR~Lyrae models and second overtone in case of $\\delta$~Cephei models). Physical reasons for resonant doubly-periodic pulsation were analysed by Dziembowski \\& Kov\\'acs (1984) with the help of the amplitude equations formalism. Linearly damped, resonant overtone acts as an energy sink for the linearly unstable fundamental mode, leading to the decrease of its amplitude. As a result, otherwise dominant fundamental mode, is no longer able to saturate the pulsation instability on its own. This allows the growth of the first overtone mode. Due to nonlinear phase-lock, resulting triple-mode pulsation is doubly-periodic. Described resonant F/O1 doubly-periodic models are mainly of theoretical interest, as they do not fulfill the observational constraints.  The search for non-resonant F/O1 (or O1/O2) double-mode pulsation with radiative codes failed, although partial success was achieved by Kov\\'acs \\& Buchler (1993). They obtained several non-resonant double-mode RR~Lyrae models, by playing with artificial viscosity dissipation. Their models however, are sensitive to numerical details, and do not satisfy all observational constraints. Experiments with artificial viscosity motivated the development of less dissipative codes with physical dissipation instead of artificial one.  Inclusion of turbulent convection into pulsational hydrocodes finally led to robust double-mode behaviour in Cepheid (Koll\\'ath \\etal 1998) and in RR~Lyrae (Feuchtinger 1998) models. More detailed surveys of the double-mode pulsation models were performed by the Florida-Budapest group (Koll\\'ath \\& Buchler 2001, Koll\\'ath \\etal 2002, Szab\\'o, Koll\\'ath \\& Buchler 2004). These analysis showed, that the occurrence of the double-mode behaviour does not depend strongly on the parameters of the convection model used, and is mainly controlled by physical parameters of the stellar model. Detailed comparison with observed double-mode pulsators was not performed, however. Also, particular effect of turbulent convection, responsible for the occurrence of double-mode phenomenon, was not identified.  In our recent paper (Smolec \\& Moskalik 2008, hereafter Paper~I), we described convective hydrocodes we have recently developed. In our hydrocodes, we adopt convection model based on Kuhfu\\ss{} (1986) work. Essentially the same convection model was used by Florida-Budapest group, however with modified treatment of the convectively stable layers. In these layers buoyant forces were neglected (Koll\\'ath \\etal 2002), which is unphysical. Also, convective flux was neglected in convectively stable regions. In Paper~I we showed the consequences of these neglects for the single-mode Cepheid models. In the present paper we show dramatic consequences for the double-mode models.  The structure of this paper is as follows. In Section~2 we briefly discuss the turbulent convection models we use in this paper. In Section~3 we describe the methods of modal selection analysis we adopt, and discuss typical scenarios for two convection models considered, with buoyant forces present and buoyant forces neglected in convectively stable zones. Reasons for the occurrence of double-mode behaviour in the latter case are explained in Section~4. In Section~5 we describe the extensive survey of models we have computed in search for the double-mode behaviour with convection model including buoyant forces in convectively stable regions. We discuss our results and conclusions in Section~6. ",
        "conclusions": "%  We have clearly identified the reasons for stable double-mode behaviour, observed in models ignoring negative buoyancy effects (PP models). Double-mode pulsation results from the neglect of the buoyant forces in convectively stable zones. As a consequence, turbulent velocities are not braked effectively below the envelope convective zone, which leads to artificial overshooting. The range of such artificial overshooting is very large, significant turbulent energies extend to more than 6 local pressure scale heights below the envelope convective zones (Paper~I). These turbulent energies are high enough to produce a significant eddy-viscous damping in the internal convectively stable regions of the model. This damping acts differentially on the pulsation modes, having stronger effect on the fundamental mode than on the first overtone, since the latter has significantly lower amplitude in the deep envelope. Without this damping (in NN models which take negative buoyancy into account) amplitude of the fundamental mode pulsation is significantly higher than the amplitude of the first overtone pulsation. Therefore, fundamental mode is able to saturate the pulsational instability of both modes. Its limit cycle is firmly stable. In PP models, amplitude of the fundamental mode is strongly reduced and it is no longer able to saturate the pulsation instability alone. This allows the first overtone to grow, and consequently double-mode pulsations arise. When convectively stable regions are treated properly and buoyant forces are included, no stable double-mode behaviour can be found (at least in the parameter range considered in this analysis).  Most of the convective double-mode models published up to date, were computed with the Florida-Budapest hydrocode. In this code PP convection model was adopted: buoyant forces as well as convective flux were neglected in convectively stable zones. This assumption however, was never discussed by the Florida-Budapest group. One case of double-mode pulsation in RR~Lyrae models was published by Feuchtinger (1998). The model was computed with the Vienna pulsation hydrocodes (Wuchterl \\& Feuchtinger 1998, Feuchtinger 1999). Although it was not explicitly stated, that turbulent source function was neglected in convectively stable regions in the Vienna hydrocode, we suspect that it was so. Our believe is based on the results concerning first overtone Cepheid models computed with both Florida-Budapest and Vienna hydrocodes (Feuchtinger, Buchler \\& Koll\\'ath 2000). These authors state, that both codes give essentially the same results. If different treatments of source function in convectively stable zones are adopted in these codes, we would expect significant differences in the computed amplitudes of the single-mode models, as we have shown in Paper~I. No such differences were reported in this paper. This indicates that also the double-mode RR~Lyrae model computed by Feuchtinger (1998) may be physically not correct, although computations of RR~Lyrae models, similar to presented in this analysis are necessary to confirm this suspicion. Many convective models of both RR~Lyrae stars and Cepheids were published by Italian group (\\eg Bono \\& Stellingwerf 1994, Bono, Marconi \\& Stellingwerf 1999). We do not know whether any systematic search for double-mode behaviour was performed with this code. However, to our best knowledge, no nonlinear double-mode models were published. The source function in the Italian code has different form, as the convection model is based on the Stellingwerf (1982) model. In the original Stellingwerf model it was assumed that $S\\sim \\sqrt{Y}$ instead of $S\\sim Y$ as in the Kuhfu\\ss{} model. Such form of the source term, particularly the fact, that in convectively stable regions source term cannot damp the convective motions, was criticized by Gehmeyr \\& Winkler (1992). This drawback however, was removed in the Italian code, through setting $S\\sim\\mathrm{sgn}(Y)\\sqrt{|Y|}$ (Bono \\& Stellingwerf 1992, 1994). Therefore, Italian code is void of artificial eddy-viscous damping responsible for double-mode behaviour in PP models (see Paper~I, and work integrals \\eg in Bono, Marconi \\& Stellingwerf 1999). In the light of our analysis, the lack of double-mode models in the computations of Italian group is not surprising.  Our extensive search for F/O1 double-mode Cepheid models with NN convection model (which include negative buoyancy effects) yielded null result. For 24 studied sets of convective and physical parameters, double-mode behaviour was not found. Instead, an either-or domain, where both F and O1 single-mode pulsations are stable, separates the first overtone pulsation domain at the hot side of the instability strip and fundamental mode pulsation domain at the cool side. What is more important, soon after fundamental mode becomes linearly unstable, its pulsation amplitude grows and strongly exceeds the pulsation amplitude of the first overtone. Consequently, fundamental mode limit cycle is firmly stable, across the significant part of the instability strip. It is the main factor preventing the occurrence of the double-mode behaviour.  Nevertheless, the F/O1 double-mode Cepheids do exist in nature. Therefore, a question arises, what is missing or what physical effect is treated incorrectly in our pulsation hydrocodes. Double-mode Cepheid models computed with PP convection, although unphysical, may provide a hint. We need a mechanism that differentially reduces the modal amplitudes. It is also likely that this mechanism should act in the deep interior of the model, where properties of the fundamental and first overtone modes do differ. In our opinion, two issues should be considered in more detail. First, the treatment of turbulent convection in our models and second, the opacities.  Turbulent convection models adopted in pulsation hydrocodes are very simplified. At the present moment, due to limited computer resources, only simple one equation models are suitable for extensive nonlinear computations. These models contain many free parameters and very simplified description of non-local phenomena, such as overshooting. Of these theories, the Kuhfu\\ss{} model seems most consistent and correct. Unfortunately, as we have shown, it is not able to reproduce the observed modal selection.  As envelope convection zones are clearly present in classical pulsators, turbulent convection is a necessary component of pulsation hydrocodes. However, it is not sure at the the moment, whether convection is crucial in bringing up the double-mode behaviour at all. Non-resonant beat phenomenon is found in radiative, and numerically robust $\\beta$~Cephei models (Smolec \\& Moskalik 2007). Also radiative non-resonant double-mode RR~Lyrae models (Kov\\'acs \\& Buchler 1993) and $\\delta$~Cephei models (Smolec, unpublished) were computed. Although, these models were obtained through playing with artificial viscosity, they are just indication of the physics missing in our hydrocodes. This turns our attention into opacities.  Recently there is a growing evidence that the opacity computations as well as solar chemical composition are still uncertain. This concerns specially the heavier elements and opacity computations in the hot temperature regimes, corresponding to iron opacity bump ($T\\sim 2\\cdot 10^5$K). Solar chemical composition was recently revised by Asplund \\etal (2005). With new solar mixture however, the standard helioseismic solar model is in serious trouble (see \\eg Montalban \\etal 2006). This indicates that solar chemical composition or opacity computations need further revisions. The strength of the opacity bump is still under debate and recent asteroseismic models indicate that enhancement of opacities is desired (Pamyatnykh, priv. comm.). Also, our computations of $\\beta$~Cephei models (Smolec \\& Moskalik 2007) indicate, that modal selection strongly depends on the opacities being used (OP \\vs OPAL opacities). Enhancement of iron-group opacities may have an effect on modal selection for Cepheid models. It affects the convective properties of our models, producing additional convection zone in the deep interior. For models considered in this analysis, that is of masses around 4.5\\MS and with solar metallicity, iron opacity bump is too weak to produce a convection zone, and hence eddy-viscous damping in the deep interior. With enhanced iron opacity bump, we expect that convective zone will develop in the deep interior, producing eddy-viscous damping. To check these effects  we have computed additional sequence of models, AZ (Table~1) with all parameters of set A, but with increased metallicity to $Z=0.03$. In these models we observe a convection zone connected with the iron bump. It produces an eddy-viscous damping, that acts on fundamental mode, but has negligible effect on the first overtone. However, this effect is weak, as the convection zone is very narrow, and confined to typically $2-4$ zones. Double-mode solution is not found.  One may draw conclusions about double-mode pulsators (\\eg their parameters, such as mass or metallicity, as well as properties of the stellar systems they are in) through linear as well as nonlinear computations. Our results invalidate the nonlinear studies done with PP convection model. However, conclusions and results obtained with linear analysis of models computed with PP convection are correct. Significant turbulent energies and consequently strong eddy-viscous damping in convectively stable zones of full amplitude models are clearly nonlinear effect. Crucial eddy-viscous terms do not affect the computed static structure of the models. If turbulent flux is turned off, turbulent energies are negligible in convectively stable regions of the static model (see Paper~I). As we describe in Paper~I, static structure computed with both PP and NN convection models is very similar. Hence, computed periods and period ratios agree very well. Also linear growth rates are similar, slightly differing for models with turbulent flux turned on. Therefore, the main outcome of the linear analysis, domains of simultaneous instability of fundamental and first overtone modes, as well as periods and period ratios, obtained with PP convection model are correct. Thus, we do not expect significant changes in the analysis of Petersen diagram in the context of beat Cepheid metallicities (Buchler \\& Szab\\'o 2007, Buchler 2008), if NN convection model is used instead of PP model.  We have analysed the modal selection in classical Cepheid models, considering fundamental and first overtone modes only. Hence, our conclusions strictly apply to F/O1 double-mode Cepheids. We haven't searched for possible O1/O2 double-mode Cepheid models. We just note, that no systematic survey of such models was performed up to date. According to Buchler \\& Koll\\'ath (2000), some O1/O2 models were found by Florida-Budapest group, among models computed with P\\'eclet correction, that accounts for radiative losses, however, no details had been given. Since we haven't performed any analysis of possible overtone double-mode models, we cannot say, whether the same unphysical mechanism is responsible for these double-mode overtone models or not. We note however, that O1 and O2 are more confined to surface layers of the model. In the deep interior both these modes have significantly smaller amplitudes than the fundamental mode. The difference between O1 and O2 modes is less pronounced. Therefore, unphysical eddy-viscous dissipation, causing the differential reduction of modal amplitudes, crucial in F/O1 double-mode PP models, should play a lesser role (if any) in case of O1/O2 double-mode models. The search of O1/O2 double-mode behaviour with our pulsation hydrocodes (NN convection) is planned.     \\Acknow{We are grateful to prof. Wojciech Dziembowski for fruitful discussions and commenting the manuscript. Alosza Pamyatnykh is acknowledged for the permission to use the opacity interpolating subroutines and computation of opacity tables with new solar mixture. This work has been supported by the Polish MNiSW Grant No. 1 P03D 011 30.}"
    },
    "0809/0809.3721_arXiv.txt": {
        "abstract": "We have constructed a thermally compensated field-widened monolithic Michelson interferometer that can be used with a medium-resolution spectrograph to measure precise Doppler radial velocities of stars. Our prototype monolithic fixed-delay interferometer is constructed with off-the-shelf components and assembled using a hydrolysis bonding technique. We installed and tested this interferometer in the Exoplanet Tracker (ET) instrument at the Kitt Peak 2.1m telescope, an instrument built to demonstrate the principles of dispersed fixed delay interferometry. An iodine cell allows the interferometer drift to be accurately calibrated, relaxing the stability requirements on the interferometer itself. When using our monolithic interferometer, the ET instrument has no moving parts (except the iodine cell), greatly simplifying its operation. We demonstrate differential radial velocity precision of a few m s$^{-1}$ on well known radial velocity standards and planet bearing stars when using this interferometer. Such monolithic interferometers will make it possible to build relatively inexpensive instruments that are easy to operate and capable of precision radial velocity measurements. A larger multi-object version of the Exoplanet Tracker will be used to conduct a large scale survey for planetary systems as part of the Sloan Digital Sky Survey III (SDSS III). Variants of the techniques and principles discussed in this paper can be directly applied to build large monolithic interferometers for such applications, enabling the construction of instruments capable of efficiently observing many stars simultaneously at high velocity-precision. ",
        "introduction": "Since the discoveries of the first extrasolar planets \\citep{Wolszczan92, Mayor95} more than 290 planets have been discovered using radial velocity techniques, transit searches, and microlensing. The majority of the known extrasolar planets have been discovered using precision radial velocity measurements that detect the gravitational influence of the planet on the parent star. The presence of the planet induces a periodic change in the line of sight velocity of the star and this velocity is measured using the small Doppler shift in the spectral lines of the star. Ongoing radial velocity surveys with high resolution echelle spectrographs have achieved velocity precisions of $1-3$ m s$^{-1}$ \\citep{Butler96, Rupprecht04}, allowing the detection of $m\\sin{i}= 7-20$ Earth mass planets close to their host star \\citep{Santos04, Rivera05}. These echelle spectrographs are generally large and expensive instruments which are not easily duplicated, leading to the availability of such instruments becoming a limiting factor in intensifying high-precision radial velocity surveys. High precision radial velocity observations of more stars are necessary in order to discover lower mass planets, explore different regimes of stellar mass and evolution (Johnson et al. 2007), planet formation around young stars (Setiawan et al. 2007), and to discover more planets around low-mass M stars (e.g. Endl et al. 2007).  Over the past few years we have developed the dispersed fixed-delay interferometer (DFDI) technique \\citep*{Erskine00, Ge02a, Ge02b} into viable instruments capable of achieving high precision radial velocity measurements on stars and detecting extrasolar planets. Such instruments are substantially less expensive than echelle spectrographs since they do not require high spectral resolutions or large-format gratings. At the heart of these instruments is a stable, field-widened, fixed-delay Michelson interferometer. One such instrument, the Exoplanet Tracker (ET) instrument at the Kitt Peak 2.1m has been used to confirm known exoplanets (van Eyken et al. 2004) and detect a hot Jupiter around HD102195 (Ge et al. 2006). The ability to reliably construct monolithic Michelson interferometers has the potential to make such instruments stable, easy to use, and relatively easy to duplicate. This will enable multiple versions of these instruments to easily be built to complement existing echelle spectrographs in order to meet the need for high precision radial velocity measurements to support upcoming space missions like {\\it KEPLER} (Borucki et al. 2003) and {\\it SIM} (Shao et al. 2007). An additional advantage of the DFDI technique is the ability to observe a large number of stars simultaneously using a wide field-of-view telescope. This technique will be utilized to conduct MARVELS (Multi-object Apache point observatory Radial Velocity Exoplanet Large-area Survey), a multi-object survey for exoplanets, as part of SDSS III (Ge et al. 2007). MARVELS will be able to simultaneously monitor 120 F,G, and K stars in the 3 degree field of view of the Sloan 2.5m telescope, efficiently searching for Jupiter-mass planets \\citep{Ge07}.  The Michelson interferometer is a crucial component of the ET and MARVELS instruments and needs to be stable in order to reliably recover relative velocity drifts of a few m s$^{-1}$. Phase drifts of the interferometer during an exposure will tend to wash out the fringes, reducing the measured fringe visibility and the velocity precision. Large drifts between observing runs make it difficult to achieve coherent velocity links over large timespans. The current Michelson interferometer setup for the ET instrument at Kitt Peak uses a piezoelectric transducer (PZT) to control the path difference. This stabilizes the interferometer fringe and reduces the interferometer fringe phase drifts. The need to run a closed-loop system on the most critical component also makes observations difficult, since the observer has to continually monitor the phase, and restart the PZT after a software or computer failure. In terms of a long duration survey for exoplanets, failure of the PZT can lead to the inability to coherently link radial velocities. The PZT control also becomes increasingly more difficult for multi-object instruments based on ET that use much larger beamsplitters due to the need to accommodate more fibers.  To address some of these problems and to make ET-like instruments more stable and substantially easier to use, we have been pursuing a program to design and build an inexpensive monolithic interferometer that satisfies our criteria of field widening and stability (see Mahadevan et al. 2004). Monolithic interferometers have also been discussed by Mosser et al. (2007) for use in asteroseismology from DOME C in Antarctica. Polarizing Michelson interferometers similar to our prototype have been in use for solar oscillation measurements for a number of years in the Global Oscillation Network Group (GONG, Harvey et al. 1995) and Solar and Heliospheric Observatory Michelson Doppler Imager (SOHO MDI) instruments and are the design chosen for the Helioseismic and Magnetic Imager (HMI) instrument (Graham et al. 2003) on the planned Solar Dynamics Observatory. Monolithic interferometer assemblies have also been used to measure the solar resonance fluorescence of OH (Englert et al. 2007). Polarizing interferometers are usually designed to observe a single absorption line, while interferometers built for spatial heterodyne spectroscopy are only effective over a small wavelength region. The monolithic interferometer for the ET instrument needs to be able to function in the wavelength range 5000-6000 \\AA, be non-polarizing, and have a slight tilt in one mirror to created the fringes observed with ET.  In this paper we describe the construction of an inexpensive prototype monolithic interferometer and present results demonstrating the ability to acquire precise stellar radial velocities using the monolithic interferometer with the Exoplanet Tracker (ET) instrument. ",
        "conclusions": "Using off-the-shelf components and a novel bonding technique we have built a monolithic fixed-delay Michelson interferometer for the ET instrument. With this inexpensive prototype we have effectively demonstrated the ability of such an interferometers to recover precise radial velocities. The results for the stable and planet bearing stars demonstrate that the prototype monolithic interferometer, developed using only off-the-shelf components, is capable of allowing us to achieve a velocity precision of $\\sim 10$ m s$^{-1}$ or better over observing runs as long as two weeks. Monolithic interferometers are significantly more insensitive to vibrations, and do not need active locking using a PZT device. The temperature-compensated, field widened interferometer also makes the ET instrument significantly easier to use since the observer no longer has to constantly monitor the phase locking software. The performance of one of our prototype interferometers in the lab is substantially better than the one tested at KPNO. We believe that the large drifts seen at Kitt-Peak were caused by the warping on the copper ring. Nevertheless, we have demonstrated the ability of such a monolithic interferometer to aid in the precise measurement of radial velocities using a suitable reference to track the instrument and interferometer drifts.  The hydrolysis catalyzed bonding technique we have employed is transparent over a wide wavelength region, and stable over a large range in temperatures. The next goal in our development efforts is a larger and very stable field-widened and temperature-compensated fixed-delay interferometer that can be used in a multi-object instrument to observe many stars simultaneously. We expect that the issues related to the warping of the ring can be avoided if a second glass material were used instead of an airgap. Better designs of such interferometers, coupled with better environmental control, may make it possible to build interferometers that are intrinsically stable enough to not require simultaneous iodine calibration in order to measure to radial velocities to a high precision. With high enough stability and good thermal control it may become possible to achieve acceptable velocity precision by bracketing the stellar observations with a velocity reference, or using another reference fiber to track the star fiber. With minor variations, such interferometers can potentially be employed in a variety of dispersed fixed delay interferometer instruments including large multi-object instruments planned for upcoming surveys like MARVELS."
    },
    "0809/0809.5107_arXiv.txt": {
        "abstract": "Using {\\it Spitzer} IRAC and MIPS observations of the Large Magellanic Cloud, we have identified 13 objects that have extremely red mid-IR colors.  Follow-up {\\it Spitzer} IRS observations of seven of these sources reveal varying amounts of SiC and C$_2$H$_2$ absorption as well as the presence of a broad MgS feature in at least two cases, indicating that these are extreme carbon stars. Preliminary estimates find these objects have luminosities of 4--11$\\times$10$^3$~$L_\\odot$ and preliminary model fitting gives mass-loss rates between 4$\\times$10$^{-5}$ and 2$\\times$10$^{-4}$~$M_\\odot$~yr$^{-1}$, higher than any known carbon-rich AGB star in the LMC.  These spectral and physical properties require careful reconsideration of dust condensation and mass-loss processes for carbon stars in low metallicity environments. ",
        "introduction": "We have undertaken a study of star formation in the Large Magellanic Cloud \\citep[LMC;][]{GC08}, using archival {\\it Spitzer Space Telescope} observations, such as those of Surveying the Agents of a Galaxy's Evolution \\citep[SAGE;][]{Metal06}. In the process of identifying young stellar objects (YSOs) we noticed 13 bright mid-IR sources (Table 1) that all had similar photometric properties with spectral energy distributions (SEDs) markedly different from those of evolved stars or YSOs: (1) they have extremely red mid-IR colors, [4.5]-[8.0]$>$4.0; (2) their SEDs, peaking between 8 and 24~$\\mu$m, can be moderately well fit by blackbodies with effective temperatures of $\\sim$230--320~K; (3) they all fall in a narrow range of brightness, 7.0$>$[8.0]$>$8.5; and (4) none of the sources have counterparts in the Two Micron All Sky Survey Point Source Catalog \\citep[2MASS PSC;][]{Setal06} or in the Digitized Sky Survey.  We dubbed these sources Extremely Red Objects (EROs).  The EROs are not likely asteroids because of their high brightnesses and lack of proper motions in SAGE observations from two epochs separated by $\\sim$3 months.  They cannot be background galaxies or Galactic sources, as the {\\it Spitzer} Wide-area IR Extragalactic Survey \\citep[SWIRE;][]{Letal03} and the Galactic Legacy IR Mid-Plane Survey Extraordinaire \\citep[GLIMPSE;][]{Betal03} do not have counterparts with similar mid-IR colors and brightnesses.  These EROs are most likely associated with the LMC, and their high luminosities imply that their Galactic counterparts would have saturated in the SWIRE and GLIMPSE Surveys.  Queries of the SIMBAD database found that some EROs had been detected previously in either {\\it MSX} or {\\it IRAS} observations \\citep[e.g.,][]{S89}. Five have been suggested as possible obscured AGB stars by \\citet{Letal97}; however, many of our EROs have also been suggested to be YSO candidates \\citep{Wetal08}. As these EROs have never been confirmed or rejected as YSOs spectroscopically, we included seven among our follow-up observations of massive YSOs in the LMC using the {\\it Spitzer} InfraRed Spectrograph \\citep[IRS;][]{Hetal04}. These IRS observations show unambiguously that the EROs are carbon stars; furthermore, the spectra reveal silicon carbide (SiC) absorption features. While SiC emission features are fairly common in both Galactic and LMC carbon stars, this is the first clear detection of SiC {\\it absorption} for LMC carbon stars. Preliminary analysis of the IRS spectra suggests that these are extraordinary carbon stars with very high mass-loss rates.  This paper reports their discovery. In \\S{2} we describe their basic photometric properties and their possible optical and near-IR counterparts.  In \\S{3} we introduce our IRS observations and in \\S{4} we discuss the results. ",
        "conclusions": ""
    },
    "0809/0809.0208_arXiv.txt": {
        "abstract": "Radio astronomical observations have very poor signal to noise ratios, unlike in other disciplines. On the other hand, it is possible to observe the object of interest for long time intervals as well as using a wider bandwidth. Traditionally, by averaging  in time and in frequency, it has been possible to improve the signal to noise ratio of astronomical observations to improve the dynamic range. This is possible due to the inherent assumption that the object of interest in the sky is invariant over time and the frequency range of observation. However, in reality this assumption does not hold, due to intrinsic variation of the sky as well as due to errors generated by the instrument. In this paper, we shall discuss  an alternative to averaging of images, without ignoring subtle changes in the observed data over time and frequency, using subspace decomposition. By separation of data to signal and noise subspaces, not only would this improve the quality of the data, but also enable us to detect faint artifacts due to calibration errors, interference etc. ",
        "introduction": "Astronomical observations have very weak signal to noise ratios (SNR). In radio astronomy (interferometry), this weak SNR is improved mainly by  prolonging the observation time. This not only improves the SNR, but also sampling of the $uv$ plane. A typical observation has a bandwidth of a few MHz, divided into smaller narrowband channels. Normally the bandwidth of a single channel is about a few kHz. After interference mitigation, calibration and imaging (also deconvolution) each of these narrow band channels, it is normal procedure to average them, on the image plane. The resultant averaged image has an improved SNR compared to images made by each of the narrowband channels, assuming all systematic errors have been removed. This is the typical process by which current radio telescopes like the Westerbork Synthesis Radio Telescope (WSRT), is able to achieve dynamic  ranges in the order of 1,000,000 to 1.   However, there is a significant drawback in this averaging procedure, because it is based on a erroneous assumption: that the source as well as the effects of the instrument remains invariant in all averaged data. Only the noise varies and averaging significantly decreases the noise power. The invariance assumption is accurate to some extent, but closer scrutiny reveals that there is significant amount of variation: \\begin{itemize} \\item There is intrinsic variation of the  sky (or the sources) with frequency (non zero spectral indices). \\item The $uv$ points scale with frequency, thus the point spread function (PSF) change in the images. \\item Transient sources as well as radio frequency interference (RFI) create temporal variation. \\item Calibration errors, small as they might be, create subtle variations in time as well as in frequency. \\end{itemize}  Indeed, by averaging, both noise and the aforementioned effects get suppressed. However, it is to our advantage that we detect such variations in our images. For instance, we could detect trace amounts of RFI and calibration errors etc., that appear only in certain time or frequency ranges. This could also enable us to discard certain parts of the data, where such effects are dominant, and thus improve the quality of our averaged, final image.  Subspace techniques have previously being used in image compression \\cite{Andrews76, Yang95}, as well as image comparison. However, to the best of our knowledge, this has not been used in (radio) astronomical image processing.  The rest of the paper is organized as follows: In the next section, we give a formal representation of images and subspace decomposition. Next, we investigate various uses of the image subspace decomposition. Later, we consider possible extensions of the proposed method, by taking into account the data weights as well as the affect of PSF. Finally, we give our conclusions.   Notation: All bold lowercase letters represent a column vector, all bold uppercase letters represent a matrix. ",
        "conclusions": "We have proposed a technique to improve the quality of the data in astronomical observations using subspace decomposition. Application of this technique to real data has shown us that it is superior to the normally used averaging techniques."
    },
    "0809/0809.5041_arXiv.txt": {
        "abstract": "{}{Emission lines in polars show complex profiles with multiple components that are typically ascribed to the accretion stream, threading region, accretion spot, and the irradiated secondary-star. In low-state polars the fractional contribution by the accretion stream, and the accretion spot is greatly reduced offering an opportunity to study the effect of the secondary-star irradiation or stellar activity. We observed VV Pup during an exceptional low-state to study and constrain the properties of the line-forming regions and to search for evidence of chromospheric activity and/or irradiation.} {We obtained phase-resolved optical spectra at the ESO VLT+FORS1 with the aim of analyzing the emission line profile and radial velocity as a function of the orbital period.  We also tailored irradiated secondary-star models to compare the predicted and the observed emission lines and to establish the nature of the line-forming regions.} {Our observations and data analysis, when combined with models of the irradiated secondary-star, show that, while the weak low ionization metal lines (FeI and MgI) may be  consistent with irradiation processes, the dominant Balmer H emission lines, as well as NaI and HeI, cannot be reproduced by the irradiated secondary-star models. We favor the secondary-star chromospheric activity as the main forming region and cause of the observed H, NaI and He emission lines, though a threading region very close to the L1 point cannot be excluded. } {} ",
        "introduction": "Different  types  of cataclysmic  variables  (CVs,  i.e. dwarf  novae, polars, etc)  in different states (outburst or  quiescence) have shown Balmer emission lines forming  close to the secondary-star.  Depending on  the  system,  plausible  origins  of  these  emission lines  have  been identified as  the irradiated  secondary hemisphere (e.g.   Steeghs et al. 2001; Araujo-Betancor  et al. 2003; Thoroughgood et  al. 2005), or the accretion stream (e.g. Cowley  et al. 1982; Mukai 1988; Schwope et al.   1997).  More  recently, it  has  been proposed  that the  Balmer emission  lines  in  low-state  polars could  be  the  signature  of chromosphere activity on the secondary-stars in AM Her, ST LMi, and EF Eri (Kafka et  al.  2006, 2007, Howell et  al.  2006b).  This scenario is  particularly intriguing  if  we consider  that V471~Tau  H$\\alpha$ emission line,     previously      thought     to     be      caused     by irradiation\\footnote{Young   et  al.    (1988)  use   the   word  {\\it fluorescence}  in the  broader context  of conversion  of  high energy photons to low energy ones.}  (Young et al. 1988), is best explained by stellar activity (Rottler et al.  2002).  In  order to  understand the  source  and formation  mechanism of  the observed Balmer emission lines in  low-state polars, we  analyzed time-resolved optical spectra of VV~Pup  obtained during a low-state.  The spectra  were secured  at  the  ESO VLT+FORS1  and  have been  already presented in Mason  et al. (2007), where we focused on the white dwarf magnetic field signatures and, in particular, the first detection of the Zeeman absorptions. Here we present  in Sec.~\\ref{uno} the data sample, in Sec.~\\ref{due} our radial  velocity study  and the  emission line  profile  analysis.  In   Sec.~\\ref{tre} the  irradiated atmosphere model tailored for the VV~Pup system is compared to the observations. Sec.\\ref{quattro} summarizes our results and conclusions. ",
        "conclusions": "\\label{quattro}  Phase-resolved optical spectroscopy of the polar VV~Pup during a low-state has shown the presence of Balmer H, FeI, MgI, NaI and HeI emission lines, which form on the side of the secondary-star facing the white dwarf. Their radial velocity curves are roughly consistent with each other but different from the secondary-star radial velocity curve measured by Howell et al. (2006a). Also, the application of the inverse K-correction is insufficient to produce agreement between our and the Howell et al. radial velocity.  The emission line fluxes show a bell shaped modulation across the orbit during epoch 1 observations, which is often taken as irradiation of the secondary-star. We produced irradiated secondary-star models tailored for VV~Pup and, though irradiation predicts low ionization metal lines, it does not explain the Balmer, Na and HeI emission lines strengths we observed. In particular, to produce these emission lines % would require a $>$20000K white dwarf in VV Pup, contrary to the UV observations. The possible presence of a hot spot (20000 to 24000K) on the magnetic white dwarf, similar to that observed in low-state polars such as EF Eri and AM Her, could alternatively provide some Balmer, Na and possibly He emission lines. However, in epoch 2 observations taken while VV~Pup was experiencing an increase in mass transfer, we observed the same lines to strengthen and show a more scattered flux distribution with the orbital period. This also disfavors irradiation of the secondary-star.  Our observations are better explained by a scenario in which the H, Na and HeI emission lines form mainly on  the white dwarf facing side of the secondary-star either in the threading region close to the L1 point, or in chromospheric prominence-loops possibly shocked near the location of the white dwarf magnetospheric radius. We note that the NaI doublet and  the H$\\alpha$ emission  lines are known to be good indicators of chromosphere activity in M and L dwarfs (Andretta et  al.  1997, Schmidt et al.  2007). Flux ratios of the HeI$\\lambda$5875/HeI$\\lambda$6678$\\sim$2.5-3.5, as measured by us in the ``bursting'' spectra of epoch 2, are  consistent with  low density gas regions  having temperatures  higher than 8000K  as in the stellar chromosphere or chromosphere-like regions (Giampapa  et al.   1978; Howell  et  al. 2006b). In addition, the strengthening of the chromospheric activity line indicators is consistent with stellar activity phenomena which, depending on the secondary-star dynamo magnetic field, might trigger episodes of enhanced mass transfer rate as observed in our epoch 2 ``bursting'' spectra."
    },
    "0809/0809.2382_arXiv.txt": {
        "abstract": "A kinetic theory for spin plasmas is put forward, generalizing those of previous authors. In the model, the ordinary phase space is extended to include the spin degrees of freedom. Together with Maxwell's equations, the system is shown to be energy conserving. Analysing the linear properties, it is found that new types of wave-particle resonances are possible, that depend directly on the anomalous magnetic moment of the electron. As a result new wave modes, not present in the absence of spin, appear. The implications of our results are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.4667_arXiv.txt": {
        "abstract": "\\noindent Gravitino dark matter, together with thermal leptogenesis, implies an upper bound on the masses of superparticles. In the case of broken R-parity the constraints from primordial nucleosynthesis are naturally satisfied and decaying gravitinos lead to characteristic signatures in high energy cosmic rays. We analyse the implications for supergravity models with universal boundary conditions at the grand unification scale. Together with low-energy observables one obtains a window of superparticle masses, which will soon be probed at the LHC, and a range of allowed reheating temperatures. ",
        "introduction": "Standard thermal leptogenesis \\cite{Fukugita:1986hr} provides a simple and elegant explanation of the origin of matter. It is a natural consequence of the seesaw mechanism, and it is perfectly consistent with the small neutrino masses inferred from neutrino oscillation data \\cite{Buchmuller:2005eh}.  Thermal leptogenesis works without and with supersymmetry. In the latter case, however, there is a clash with the `gravitino problem' \\cite{Weinberg:1982zq, Ellis:1984er,Kawasaki:2004yh}: the large temperature required by leptogenesis exceeds the upper bound on the reheating temperature from primordial nucleosynthesis (BBN) in typical supergravity models with a neutralino as lightest superparticle (LSP) and an unstable gravitino. If the gravitino is the LSP, the condition that relic gravitinos do not overclose the universe yields an upper bound on the reheating temperature \\cite{Moroi:1993mb}. Furthermore, the next-to-lightest superparticle (NLSP) is long lived, and one has to worry about the effect of NLSP decays on nucleosynthesis.  It is remarkable that, despite these potential problems, a large leptogenesis temperature of order $10^{10}$~GeV can account for the observed cold dark matter in terms of thermally produced relic gravitinos \\cite{Bolz:1998ek}. Requiring consistency with nucleosynthesis  yields constraints on the superparticle mass spectrum. Due to improved analyses of BBN, the original proposal of a higgsino NLSP is no longer viable, and also other possible NLSPs are strongly constrained. The case of a stau NLSP is cornered by bounds following from catalyzed production of $^6{\\rm Li}$ \\cite{Pospelov:2006sc}, with the possible exception of a large left-right mixing in the stau sector \\cite{Ratz:2008qh}. In some models a sneutrino \\cite{Kanzaki:2006hm} or a stop \\cite{DiazCruz:2007fc} can still be a viable NLSP.  Recently, it has been shown that in the case of small R-parity and lepton number breaking, such that the baryon asymmetry is not erased by sphaleron processes \\cite{Campbell:1990fa}, thermal leptogenesis, gravitino dark matter and primordial nucleosynthesis are naturally consistent \\cite{Buchmuller:2007ui}. Although the gravitino is no longer stable, its decay into standard model (SM) particles is doubly suppressed by the Planck mass and the small R-parity breaking parameter. Hence, its lifetime exceeds the age of the universe by many orders of magnitude, and it remains a viable dark matter candidate \\cite{Takayama:2000uz}. Gravitino decays lead to characteristic signatures in high energy cosmic rays. The produced flux of gamma-rays \\cite{Takayama:2000uz,Buchmuller:2007ui,Bertone:2007aw, Ibarra:2007wg, Ishiwata:2008cu} and positrons \\cite{Ibarra:2008qg,Ishiwata:2008cu} may explain the observed excess in the EGRET \\cite{Strong:2004ry} and HEAT \\cite{Barwick:1997ig} data. This hypothesis will soon be tested by the satellite experiments FGST and PAMELA.  In this paper we study the implications of leptogenesis and gravitino dark matter with broken R-parity on the mass spectrum of superparticles. Since the unification of gauge couplings in the minimal supersymmetric extension of the standard model (MSSM) is one of the main motivations for low-energy supersymmetry, we shall focus on versions of the MSSM with universal boundary conditions for scalar and gaugino masses at the grand unification (GUT) scale. As we shall see, the corresponding spectrum of superparticle masses will be fully covered at the LHC. This is the main result of our analysis.  After some comments on R-parity violation in Section~2, we discuss the lower bound on the reheating temperature from leptogenesis and the upper bound on the NLSP mass from gravitino dark matter in Section~3. Section~4 deals with constraints on MSSM parameters from low-energy observables, and the results of our numerical analysis are presented in Section~5, followed by some conlusions in Section~6. ",
        "conclusions": "\\label{Sec:Conclusion}  We have studied the implications of thermal leptogenesis and gravitino dark matter for the mass spectrum of superparticles. In the case of broken R-parity the constraints from nucleosynthesis are naturally fulfilled, and universal gaugino masses at the GUT scale are possible, contrary to the case of stable gravitinos.  As an illustration, we have considered two boundary conditions which lead to a bino-like NLSP and a stau NLSP, respectively. Low-energy observables and gravitino dark matter together with thermal leptogenesis yield upper and lower bounds on NLSP and gluino masses, which in both cases lie within the discovery range of the LHC. It is encouraging that the supersymmetric explanation of the muon $g-2$ anomaly favours smaller masses within these mass windows.  A cosmology with leptogenesis and gravitino dark matter also leads to the prediction of a maximal temperature in the early universe. In the case of universal gaugino masses at the unification scale we find the upper bound $T_R^{\\rm max}\\simeq 6\\times 10^9$~GeV, which is somewhat relaxed for large scalar masses. This bound has been obtained under the assumption of thermal equilibrium, which appears unlikely for a maximal temperature. Nevertheless, it is intriguing that the temperature $T_R^{\\rm max}$ is of the same order of magnitude as the critical for the destabilization of compact dimensions in higher-dimensional supersymmetric theories \\cite{Buchmuller:2004xr}. The effect of the reheating process on the stabilization of extra dimensions and the relation to baryogenesis and dark matter require futher investigations.  Gravitino decays produce a flux of photons and positrons, which can significantly contribute to the EGRET and HEAT anomalies for a lifetime $\\tau_{3/2} \\sim 10^{26}$~s. If these anomalies are indeed related to gravitino decays, the satellite experiments FGST and PAMELA should soon detect characteristic features in the photon and positron spectrum, respectively. Observation of a line in the gamma-ray spectrum by FGST and a rise with sharp cutoff in the positron spectrum by PAMELA would lead to a determination of the gravitino mass. This would considerably tighten the predictions for superparticle mass windows which will be probed at the LHC."
    },
    "0809/0809.3337_arXiv.txt": {
        "abstract": "We study the bulk viscosity of neutron star matter including $\\Lambda$ hyperons in the presence of quantizing magnetic fields. Relaxation time and bulk viscosity due to both the non-leptonic weak process involving $\\Lambda$ hyperons and direct Urca processes are calculated here. In the presence of a strong magnetic field of $10^{17}$ G, the hyperon bulk viscosity coefficient is reduced whereas bulk viscosity coefficients due to direct Urca processes are enhanced compared with their field free cases when many Landau levels are populated by protons, electrons and muons. \\pacs{97.60.Jd, 26.60.-c, 04.40.Dg} ",
        "introduction": "R-mode instability plays an important role in regulating spins of newly born neutron stars as well as old and accreting neutron stars in low mass x-ray binaries \\cite{Nar}. Gravitational radiation drives the r-mode unstable due to Chandrasekhar-Friedman-Schutz mechanism \\cite{Chan,Frie,Kok,And01,And98,Fri98,Lin98,And99,Ster}. R-mode instability could be a promising source of gravitational radiation. It would be possible to probe neutron star interior if it is detected by gravity wave detectors.  Like gravitational radiation, electromagnetic radiation also drives the r-mode unstable through Chandrasekhar-Friedman-Schutz mechanism. There exists a class of neutron stars called magnetars \\cite{Tho} with strong surface magnetic fields $10^{14}-10^{15}$ G as predicted by observations on soft gamma-ray repeaters and anomalous x-ray pulsars \\cite{Kou,Vas}. The effects of magnetic fields on the spin evolution and r-modes in protomagnetars were investigated by different groups \\cite{Rez1,Lai,Rez2}. On the one hand, it was shown that the growth of the r-mode due to electromagnetic and Alfv\\'en wave emission for strong magnetic field and slow rotation could compete with that of gravitational radiation \\cite{Lai}. On the other hand, it was argued that the distortion of magnetic fields in neutron stars due to r-modes might damp the mode when the field is $\\sim 10^{16}$ G or more \\cite{Rez1,Rez2}.  The evolution of r-modes proceeds through three steps \\cite{Owen}. In the first phase, the mode amplitude grows exponentially with time. In the next stage, the mode saturates due to nonlinear effects. In this case viscosity becomes important. Finally, viscous forces dominate over gravitational radiation driven instability and damp the r-mode. This shows that viscosity plays an important role on the evolution of r-mode. Bulk and shear viscosities were extensively investigated in connection with the damping of the r-mode instability \\cite{Nar,Jon1,Jon2,Lin02,Dal,Dra,Chat1,Chat2,Chat3,Han,And06,Mad92,Mad00,Don1,Don2,Pan,Bas,Schm,Alf,Bas2}. In particular, it was shown that the hyperon bulk viscosity might effectively damp the r-mode instability \\cite{Chat3}. However all these calculations of viscosity were performed in the absence of magnetic fields. The only calculation of bulk viscosity due to Urca process in magnetised neutron star matter was presented in Ref.\\cite{anand}. This motivates us to investigate bulk viscosity due to non-leptonic process involving hyperons in the presence of strong magnetic fields. It is to be noted that the magnetic field in neutron star interior might be higher by several orders of magnitude than the surface magnetic field \\cite{Lai2}. Further it was shown that neutron stars could sustain strong interior magnetic field $\\sim 10^{18}$ G \\cite{Car,Bro}.  The paper is organised in the following way. In Section \\ref{eos} we describe hyperon matter in strong magnetic fields. We calculate bulk viscosity due to the non-leptonic process involving $\\Lambda$ hyperons and due to leptonic processes in Section \\ref{bv}. We discuss results in Section \\ref{rd} and a summary is given in Section \\ref{sum}. ",
        "conclusions": ""
    },
    "0809/0809.4517_arXiv.txt": {
        "abstract": "% We report preliminary $VRI$ differential photometric and spectroscopic results for KBS~13, a recently discovered non-eclipsing sdB+dM system. Radial velocity measurements indicate an orbital period of $0.2923 \\pm 0.0004$ days with a semi-amplitude velocity of 22.82 $\\pm$ 0.23 ${\\rm km\\,s^{-1}}$. This suggests the smallest secondary minimum mass yet found.  We discuss the distribution of orbital periods and secondary minimum masses for other similar systems. ",
        "introduction": "%  SdB and dM binaries are fairly rare (Green et al.\\ 2005), even though it is photometrically straightforward to detect an M dwarf secondary by its `reflection effect' for orbital periods less than a day or two. Still, even though the peak in the orbital period histogram is close to one day, there are about 20 to 30 times as many known post-common envelope sdB+white dwarf binaries as there are sdB+dM binaries.  The latter are particularly interesting because they will eventually evolve into cataclysmic variables (CV). Understanding the pre-CV evolution may shed additional light on, for example, the CV period gap.  KBS~13 was selected from a list of blue objects in the fields of the Kepler mission. D.~Sing's spectroscopic survey (priv.~communication) of Kepler Blue Stars (KBS) identified it as an sdB. Exploratory lightcurves for several KBS sdB candidates in November 2005 showed a reflection effect for KBS~13.  Its 2MASS colors constrain the main sequence companion to be no brighter than a mid-M dwarf.  Follow-up photometry in May and September 2006 confirmed the reflection effect, and showed that KBS~13 does not eclipse. ",
        "conclusions": ""
    },
    "0809/0809.2936_arXiv.txt": {
        "abstract": "{We present updated values for the mass-mixing parameters relevant to neutrino oscillations, with particular attention to emerging hints in favor of $\\theta_{13}>0$. We also discuss the status of absolute neutrino mass observables, and a possible approach to constrain theoretical uncertainties in neutrinoless double beta decay. Desiderata for all these issues are also briefly mentioned.}   \\normalsize\\baselineskip=15pt ",
        "introduction": " ",
        "conclusions": "Since the atmospheric $\\nu$ oscillation discovery 10 years ago, important pieces of information are being slowly added to the puzzle of absolute $\\nu$ masses. We have discussed the most recent oscillation and non-oscillation updates in the field, as presented at this NO-VE Workshop---and updated after the recent  Neutrino~2008 Conference. Oscillation parameters are robustly constrained, and an intriguing indication for  $\\theta_{13}>0$ appears to emerge. Concerning non-oscillation observables, despite some recent experimental and theoretical progress, a coherent picture remains elusive. In particular, the $0\\nu2\\beta$ claim is still under independent experimental scrutiny, and it may be compatible or incompatible with the cosmological bounds, depending on data selection (especially Ly$\\alpha$). Reduction of nuclear matrix element uncertainties is also crucial to improve the comparison of different $0\\nu2\\beta$ results. A confident assessment of the $\\nu$ mass scale will require converging evidence from at least two of the three observables $(m_\\beta,\\,m_{\\beta\\beta},\\,\\Sigma)$ within the narrow limits allowed by oscillation data."
    },
    "0809/0809.0354.txt": {
        "abstract": "In clusters of galaxies, the specific entropy of intracluster plasma increases outwards.  Nevertheless, a number of recent studies have shown that the intracluster medium is subject to buoyancy instabilities due to the effects of cosmic rays and anisotropic thermal conduction. In this paper, we present a new numerical algorithm for simulating such instabilities. This numerical method treats the cosmic rays as a fluid, accounts for the diffusion of heat and cosmic rays along magnetic field lines, and enforces the condition that the temperature and cosmic-ray pressure remain positive.  We carry out several tests to ensure the accuracy of the code, including the detailed matching of analytic results for the eigenfunctions and growth rates of linear buoyancy instabilities.  This numerical scheme will be useful for simulating convection driven by cosmic-ray buoyancy in galaxy cluster plasmas and may also be useful for other applications, including fusion plasmas, the interstellar medium, and supernovae remnants.   %%%%% %This article is the first of a series of two papers. The ultimate goal %of these papers is to carry out 3D high-resolution cooling-flow %cluster simulations including cosmic rays, magnetic fields, and %anisotropic transport.  In this paper, we present the numerical %schemes that we have developed to simulate a composite of plasma plus %cosmic rays undergoing anisotropic thermal conduction and cosmic-ray %diffusion, as well as several numerical tests.  %We base our code on an existing total variation diminishing (TVD) %scheme, which we extend to include a second fluid component of cosmic %rays. We then implement a new, second-order method for purely %anisotropic transport. This method involves only a small amount of %numerical cross-field diffusion, and also ensures that temperature %minima are heated while temperature maxima are cooled, thereby %preventing unphysical results such as negative temperatures or %cosmic-ray pressures, even in the presence of sharp temperature and %cosmic-ray-pressure gradients.  %We carry out several tests to ensure the accuracy of the code, %including a shock-tube test for a composite of plasma and cosmic %rays. We test the anisotropic transport scheme for static, circular %field lines, and find a ratio of perpendicular to parallel conduction %as low as 0.0002 for 100 by 100 grid. We also compare numerical %simulations of small-amplitude convective instabilities to the %eigenfunctions and growth rates from the analytic theory of linear %convective instabilities in a composite of plasma and cosmic rays %undergoing anisotropic transport. In all cases, we find that the code %agrees very well with theory. We note that in the presence of magnetic %fields and anisotropic transport, the convective stability criterion %is strongly modified from the Schwarzchild one, and negative %temperature gradients and/or negative cosmic-ray-pressure gradients %can trigger Parker/convective modes even in the presence of a positive %entropy gradient. Negative temperature gradients are observed in the %outer regions of galaxy clusters, and negative cosmic-ray pressure %gradients are expected in galaxy-cluster cores. The Parker/convective %instability may thus play an important role in clusters, a point that %we will address in more detail in the second paper of this series. We %also note that our numerical scheme may be useful for other %applications, including fusion plasmas, the interstellar medium, and %supernovae remnants. %%%%%  %This article is the first of a series of 2 papers. The ultimate goal of these papers is to carry-out 3D high-resolution cooling-flow cluster simulations including cosmic rays, magnetic fields and anisotropic transport. In this paper, we methodically present the numerical schemes implemented to evolve a composite of plasma plus cosmic rays undergoing anisotropic diffusion and thermal conduction. In each part, a linear and a non-linear test are presented, the last part showing a linear test for the overall system as well as a physical insight into the non-linear regime of the Parker/convective instability (PCI). Our methodology is the following. We first use the TVD MHD code to solve the magnethohydrodynamics (MHD) equations. Cosmic rays are then evolved using the same total variation diminishing (TVD) method. The implementation is successfully tested with the shock-tube test for a composite of plasma plus cosmic rays. Finally, a great effort has been done to develop a new, second order, positive and accurate method for purely anisotropic transport. This is indeed essential to avoid unphysical results such as negative temperature. The idea is to consider small magnetic flux-tubes and to compute temperature gradients and heat flux directly along the field lines. This flux-tube scheme has been tested on circular field lines and shows a ratio of perpendicular to parallel conduction as low as 2/10000 for a 100 by 100 grid points simulation. In presence of such physical ingredients the convective stability criterion is strongly modified from the Schwarzchild one. Negative temperature gradients as well as negative non-thermal pressure-gradients (expected in clusters) could trigger Parker/convective modes even in presence of a positive entropy gradient. We therefore compare carefully the eigenfunctions of such convective mode with the numerical results to test our overall implementation. Afterwards, in the non-linear regime, the magnetic and kinetic energy of the fluctuations saturate. The heat transport by conduction, diffusion and convection tends to homogenize the temperature and cosmic-ray pressure. The behavior is therefore very similar to the magnetothermal instability. Since the Parker/convective instability is a key ingredient of the AGN-driven convection model of galaxy-cluster plasmas, this is an obvious future application. However, the code could find various other applications ranging from fusion plasma, to interstellar medium or supernov{\\ae} remnants. ",
        "introduction": "The hierarchical model of galaxy formation succesfully predicts the evolution of baryons in the universe over a wide range of scales assuming that supernov{\\ae} feedback is taken into account \\citep{kauffmann99,somerville99,cole00,hatton03,springel03,rasera06}. However, the baryon budget remains inaccurate in large scale structures (large galaxies, groups or clusters) where the total amount of cold gas and stars is overestimated. This overcooling problem is particularly critical for galaxy clusters, in which the cooling time near the center of a cluster is often much shorter than a cluster's age. In the absence of heating, one would expect cooling flows to form in these clusters, with large amounts of plasma cooling and flowing in towards the center. However, the star formation rate in cluster cores is typically 10-100 times lower than the predictions of the cooling-flow model \\citep{mcnamara04}, and line emission from plasma at temperatures lower than one third of the virial temperature of the cluster is weak \\citep{peterson03,mcnamara04}. The inconsistency between the cooling-flow model and these observations is known as the ``cooling-flow problem''.  A promising hypothesis to solve this puzzle is heating by active galactic nuclei (AGN) in cluster cores. Two main arguments support this idea. First, AGN power is expected to be a decreasing function of the specific entropy at a cluster's center and therefore tends naturally towards a self-regulated state in which heating balances cooling \\citep{nulsen04,boehringer04}. Second, almost all cooling core clusters possess active central radio sources \\citep{eilek04}. However, one important problem remains: how is AGN power transferred to the ambiant plasma? Over the last decade, a number of numerical simulations have been carried out to answer this question. In the first simulations \\citep{churazov01,quilis01,bruggen02a,bruggen02b}, thermal energy was injected near the center of a 2D or 3D cluster-like hydrostatic profile. This resulted in hot and underdense bubbles, which then rose buoyantly. By agitating the surrounding medium, these bubbles were able to reduce the cooling while achieving some correspondence with the observations of X-ray cavities seen in roughly one-fourth of the clusters of the Chandra archive \\citep{birzan04}.  Subsequent simulations extended these earlier works to include new physical ingredients, such as viscosity \\citep{ruszkowski04a,ruszkowski04b,reynolds05,bruggen05,sijacki06}. It was found that viscous dissipation contributed to the energy transfer, and that viscosity helped to prevent bubbles from breaking up. Other studies \\citep{reynolds01,reynolds02, omma04a,omma04b, cattaneo06, heinz06} injected not only thermal energy but also kinetic energy in subrelativistic bipolar jets. This approach also leads to cavities, but the dynamics are different than in the previous works because of the initial momentum of the bubbles and because the energy is deposited over a more narrow range of angles. In this context, the importance of turbulence, magnetohydrodynamics effects, and plasma transport processes has been underlined by \\citet{vernaleo06}, who suggested that these ingredients could prevent the heating from being highly concentrated along the jet axis, as is the case for one-fluid pure-hydrodynamics simulations of jets in clusters that are initially at rest. %However, one important problem remains: how to transfer the AGN power to the ambiant plasma? In this context, many hypotheses have been invoked: MHD-wave-mediated heating from cosmic rays \\citep{boehringer88,rosner89,loewenstein91}, heating by shocks from central jets \\citep{omma04,soker05,brighenti06,cattaneo07}, cosmic-ray bubbles and the associated turbulence \\citep{loewenstein90,chandran04,churazov04}, sound waves \\citep{fabian03, ruszkowski04,ruszkowski04b, fabian05}, and work that they generate \\citep{churazov00,churazov01,ruszkowski02}, etc.   The above simulations treated the intracluster medium (ICM) as a single fluid.  In single-fluid simulations, when AGN-heated plasma at temperature $T_{\\rm hot}$ mixes with ambient intracluster plasma at temperature~$T_0$, the result is a Maxwellian plasma with a temperature intermediate between $T_0$ and $T_{\\rm hot}$.  Although this approach is valid in clusters if~$T_{\\rm hot}$ is not too large, it breaks down if the hot particles are relativistic or transrelativistic, because then Coulomb collisions do not have sufficient time to bring the hot particles into thermal equilibrium with the ambient intracluster plasma.  If we focus on hot protons, the type of collision that brings such protons most rapidly into thermal equilibrium with the background plasma is collisions with background electrons.  The time scale for thermal electrons to remove energy from a hot proton is \\citet{gould72}, \\begin{eqnarray} \\tau_\\epsilon=\\frac{(\\gamma-1) m_p m_e v_p c^2}{4 \\pi e^4 n_e} \\left[\\textrm{ln}\\left( \\frac{2 m_e c v_p p}{\\hbar (4 \\pi e^2 n_e/m_e)^{\\frac{1}{2}}} \\right)-\\frac{v_p^2}{2 c^2}\\right], \\label{eq:taue} \\end{eqnarray} with $n_e$ the electron density, $\\gamma$ the Lorentz factor, $v_p$ the proton velocity, $e$ and $m_e$ the electon charge and mass, $p$ and $m_p$ the proton momentum and mass, and $\\hbar$ the reduced Plank constant.  %\\begin{equation} %\\tau_\\epsilon \\simeq 8\\times 10^9 E_{\\rm GeV}^{3/2} %[(10^{-2} \\mbox{ cm}^{-3})/n_e] \\mbox{ yr} %\\label{eq:taue} %\\end{equation} %(assuming a Coulomb logarithm of 35), where $E_{\\rm GeV}$ is the %proton energy in~GeV. (YANN: would you CHECK THIS NUMBER, and look for %a value that extends to the relativistic regime? BOOK 1990) For a typical proton energy of $E \\simeq 1$~GeV (transrelativistic regime) and a typical cluster-core electrons density of $n_e=0.01$~cm$^{-3}$, the thermalization time scale is $\\tau_\\epsilon \\simeq 7$~Gyr, which is much larger than the time for protons to escape the cluster core via diffusion or convection. In this case, the ICM is essentially a two-fluid system similar to the interstellar medium of the galaxy, with a thermal background plasma plus a population of high-energy particles (cosmic rays).  There are a few problems with treating a mix of cosmic rays and thermal plasma as a single Maxwellian fluid. One is that the single-fluid approximation to the temperature contains the cosmic-ray contribution to the energy density, and thus overestimates the actual temperature of the thermal plasma. If the cosmic-ray energy density is a significant fraction of the total energy density, the single-fluid model is unable to accurately predict the temperature profile of a cluster. In addition, since the thermal conductivity depends sensitively on the temperature ($\\kappa_T \\propto T^{5/2}$), and since conduction can make an important contribution to the heating of a cluster core \\citep{zakamska03} errors in the temperature profile can also lead to significant secondary errors in the thermal balance of the ICM.  A more subtle difficulty in applying a one-fluid model to a cosmic-ray/thermal-plasma mixture concerns the convective stability of intracluster plasma. It turns out that a radial gradient in the cosmic-ray energy density is much more destabilizing than a radial gradient in the thermal plasma energy density when the plasma mass density decreases outwards (see Eq.\\ref{criterion} below). A correct accounting of the fraction of the total pressure contribution by cosmic rays is thus essential for understanding the convective stability of clusters. A more extensive discussion of this point is given by \\citet{chandran06}.   %It turns out that the convective stability of %low-density intracluster plasmas depends critically on the way in %which heat and particles diffuse along magnetic field lines \\citep{balbus00, balbus01, parrish05, parrish07, chandran05, chandran06}. In stellar interiors where heat is conducted %isotropically by photons, convective stability is governed by the %Schwarzchild criterion, and requires that the specific entropy %increase outwards.  The situation is very different in a low-density %plasma in which heat and cosmic rays diffuse primarily along magnetic %field lines.  It turns out that the Schwarzchild criterion does not %apply, and the actual convective stability criterion makes convection %in clusters likely \\citep{chandran05, chandran06}. %Moreover, the way in which cosmic rays and thermal plasma enter the %convective stability criterion is very different (see, e.g. equation \\ref{criterion} %below).  The essential reason for this is that thermal conduction %equalizes the temperature along magnetic field lines, whereas %cosmic-ray diffusion equalizes the cosmic-ray pressure along field %lines.  [See \\citet{chandran06} for a more detailed %  discussion.]  As a result, a modest cosmic-ray pressure fraction can destabilize the intracluster medium to convection \\citep{chandran05,chandran06,chandran07}. A correct accounting of the fraction of the total pressure %contributed by cosmic rays is thus essential for understanding the %convective stability of clusters.   A more accurate treatment of the ICM, in which the cosmic rays are treated as either a second fluid or as collisionless particles, is thus needed. In this paper, we present a new numerical algorithm for simulating the ICM that treats the ICM as a two-fluid (cosmic-ray plus thermal-plasma) system. We also present the results of a suite of tests for our code. Our numerical approach is similar to that of Mathews \\& Brighenti (2008), who carried out two-fluid simulations of cosmic-ray bubbles in the ICM. However, in contrast to this latter study, we take thermal conduction and cosmic-ray diffusion to occur almost entirely along magnetic field lines (cross-field transport arising only from numerical diffusion). Such anisotropic transport arises in clusters because the Coulomb mean free paths of thermal particles in clusters are much larger than their gyroradii, and the scattering mean free paths of cosmic rays are much larger than their gyroradii. The effects of magnetic fields on conduction are some times taken into account in when considering thermal conduction over length scales much larger than the correlation length of the (tangled) intracluster magnetic field, $l_B \\simeq 1-10$~kpc \\citep{kronberg94,taylor01,taylor02,vogt03}.  In this case, the conductivity~$\\kappa_T$ is effectively isotropic \\citep{rechester78,chandran98} with a value that is $\\simeq 0.1-0.2$ times the Spitzer thermal conductivity for a non-magnetized plasma \\citep{narayan01, chandran04b, maron04}.  However, on scales~$\\lesssim l_B$, the anisotropy of the thermal conductivity has a powerful effect on the convective stability of the intracluster medium \\citep{balbus00, balbus01, parrish05,chandran06, parrish07, parrish08, quataert08}, in such a way as to make convection much more likely than when the conductivity is treated as isotropic.  This is true even if the magnetic field is so weak that the Lorentz force is negligible. In order to simulate buoyancy instabilities and convection in clusters, it is thus essential to incorporate anisotropic transport.  %In any case, realistic simulations have to be carried out in order to get rigourous conclusion. It means that, in addition to the plasma, non-thermal components have to be taken into account. Indeed, \\citet{ensslin04, dunn06,vernaleo06} emphasize the important role of magnetic field and cosmic rays in cluster cores. They contribute to a non-negligible part of the pressure. Also, two-fluid approach is very important to not overestimate the plasma temperature and the associated conduction. Finally, it modifies the convective stability criterion \\citep{chandran01,chandran05,chandran06,dennis07}. Another fundamental ingredient which is often neglected in cooling flow studies with AGN heating is the anisotropic conduction. It has been shown that isotropic conduction with a conductivity of 10-100\\% of the Spitzer value could balance cooling in large clusters \\citep{fabian02, voigt02,zakamska03}. Even though cpmduction is unable to balance radiative cooling in small clusters, it indicates that conduction should not be neglected. However, in all this studies the conductivity is quite uncertain since it depends on assumption about the geometry of the underlying magnetic field. This is why a proper and self-consistent modelisation needs to take into account the highly anisotropic character of conduction in clusters. Here again, this will modify the convective instability criterion \\citep{balbus00,parrish05}.  %Our objective in this series of two articles is to carry-out high-resolution cooling-flow cluster simulations in order to undersand the role of the non-thermal components and anisotropic transport.    The remainder of this paper is organized as follows. In section~\\ref{sec:equations} we present the basic equations of our two-fluid model. In section~\\ref{sec:notransport} we present the total-variation-diminishing (TVD) code that we use to solve these equations as well as several numerical tests, focusing on the case in which there is no conduction or diffusion. In section~\\ref{TVD anisotropic transport} we present the standard numerical discretization method for anisotropic conduction. We show how it can lead to negative temperature as emphasized before by \\citet{sharma07}. We then describe our new method that does not suffer from negative temperature problems. Tests such as the circular conduction test and Sovinec-test are also presented. Finally, in section~\\ref{sec:buoyancy} we present results for the linear buoyancy instabilities involving cosmic rays and anisotropic transport and compare our numerical solutions to analytic results.  % and compare these results to the magnetothermal instability \\citep{parrish05}. ",
        "conclusions": "This article could be viewed as a test guide for those who want to implement cosmic-ray and anistropic transport routines, which are essential ingredients to simulate cooling-flow clusters. We indeed used many linear and non-linear tests for each physical ingredient as well as a new linear test for the full system. Our contribution could be divided into three parts.  First, we showed that the TVD method of \\citet{pen03} can be used to evolve the cosmic rays and the plasma simultaneously. The shock tube problem for a composite of plasma and cosmic rays is an example of the successful implementation. Second, we insisted on the importance of having a positive implementation of the anisotropic conduction in order to ensure physical results even in the presence of sharp gradients. We therefore presented a new flux-tube method which has two important properties: positivity and accuracy. This is important for clusters of galaxies since the conduction is highly anisotropic as opposed to a very diffusive scheme. Moreover, the random magnetic fields and potential large temperature and cosmic-ray-pressure gradients in cluster cores may cause negative temperatures in a non-positive scheme. Third, the linear regime of the cosmic ray magnetothermal instability (CRMTI) provides a new, sensitive test, of the overall implementation. The main interest is that this instability has a broader range of applications since the criterion is $n k_{B}dT/dz+dp_{cr}/dz+de_{mag}/dz>0$. One interesting future application of this code concerns the cores of clusters of galaxies. Indeed, in these regions, negative radial gradients of cosmic-ray pressure may trigger convection. Moreover, possible large gradients of cosmic-ray pressure near the edges of X-ray cavities (cosmic-ray bubbles) require positive implementation of the anisotropic transport.  It is worth noting that even though our discussion has focused mainly on clusters of galaxies, the flux-tube method that we have developed could be used to simulate a variety of physical systems, including the interstellar medium, fusion plasmas, and supernovae-remnants."
    },
    "0809/0809.2107_arXiv.txt": {
        "abstract": "It has been suggested that the accelerated expansion of the Universe is due to backreaction of small scale density perturbations on the large scale spacetime geometry. While evidence against this suggestion has accumulated, it has not yet been definitively ruled out. Many investigations of this issue have focused on the Buchert formalism, which computes spatial averages of quantities in synchronous comoving gauge. We argue that, for the deceleration parameter of this formalism to agree with observations, the spatial average of the three dimensional Ricci scalar (spatial curvature) must be large today, with an $\\Omega_k$  in the range of $1 \\le  \\Omega_k  \\le 1.3$. We argue that this constraint is difficult to reconcile with observations of the location of the first Doppler peak of the CMBR. We illustrate the argument with a simple toy model for the effect of backreaction, which we show is generically incompatible with observations. ",
        "introduction": "Measurements of luminosity distance  as function of redshift  for type Ia supernovae, as well as measurements of inhomogeneities in the cosmic microwave background radiation, indicate that the expansion of the Universe is accelerating today \\cite{Riess,Perlmutter,Bennett}. Explanations of this phenomenon usually involve either an introduction of \"dark energy\" -- a form of matter with negative pressure, or  a modification of general relativity. Recently, a different explanation has been put forward \\cite{Rasanen,Notari,Kolb1,Kolb2,lisch1,lisch2,Wiltshire,lnw,buchert2}, where the  acceleration is a consequence of subhorizon density perturbations. According  to this idea, small scale cosmological density perturbations evolve in a nonlinear manner to produce backreaction that affects the large scale spacetime geometry and modifies  the expansion of the Universe. This explanation is controversial and many authors  have argued that backreaction can not explain the current acceleration of the Universe  \\cite{sf,iw,bbr}. Our viewpoint is that backreaction is likely to be too small to produce a significant modification to the large scale expansion of the universe. However, it deserves  to be investigated in detail since the backreaction explanation has not yet been definitively refuted.  To quantify the rate of expansion of an inhomogeneous  Universe, Buchert \\cite{buchert1,buchert2} introduced a particular method of taking a spatial average of the Einstein equations. He specialized to comoving synchronous gauge, and on each surface of constant time, denoted by $t$, he considers a spatial domain  $D(t)$ such that the boundary of $D(t)$ is comoving [i.e. $D(t)$ is independent of time in the synchronous comoving coordinates]. Defining  $V_D(t)$ to be the proper volume of this domain, the effective scale factor $a_D(t)$ is given by \\[ \\frac{4\\pi}{3} a^3_D(t)=V_D(t)\\,. \\] By averaging the Einstein equations for an irrotational dust Universe, Buchert derived the following evolution equations for $a_D(t)$ \\begin{eqnarray}\\label{ad}\\label{aprime} &&\\left(\\frac{a'_D}{a_D}\\right)^2 =\\frac{8\\pi}{3}\\rho_{eff} \\,,\\\\ \\label{adbprime} &&- \\frac{a''_D}{a_D} = \\frac{4\\pi}{3}(\\rho_{eff}+ 3p_{eff})\\,, \\end{eqnarray} where prime denotes differentiation with respect to cosmic time $t$, and throughout this paper we use geometrized units where  $G=c=1$. Equations (\\ref{aprime},\\ref{adbprime}) have the same form as the Friedmann equations, except that their sources are an effective density and an effective pressure, $\\rho_{eff}$ and  $p_{eff}$, respectively, which are defined by \\begin{eqnarray}\\label{rhoeff} \\rho_{eff}\\equiv\\langle \\rho \\rangle_D -\\frac{1}{16\\pi}(\\langle R_3 \\rangle_D + \\langle Q \\rangle_D) \\,,\\\\ \\label{peff} p_{eff} \\equiv -\\frac{1}{16\\pi}(\\langle Q \\rangle_D -\\frac{1}{3} \\langle R_3 \\rangle_D)\\,. \\end{eqnarray} Here  $\\rho $ denotes the matter density, and $R_3$ denotes the spatial three-dimensional Ricci scalar. The brackets $\\langle...\\rangle_D$ denote an average over the domain $D(t)$, for example \\[ \\langle R_3 \\rangle_D\\equiv \\int_D R_3 \\sqrt{ \\det(g_{ij})} dV\\ /\\ \\int_D \\sqrt{\\det(g_{ij})}dV= V_D^{-1}\\int_D R_3 \\sqrt{ \\det(g_{ij})} dV\\,, \\] where ${\\det(g_{ij})}$ denotes the determinant of  the spatial 3-dimensional induced metric, and $dV$ denotes the  three-dimensional coordinate volume-element. The quantity denoted $\\langle Q \\rangle_D$ is defined by \\begin{equation} \\langle Q \\rangle_D \\equiv \\frac{2}{3}\\langle(\\theta-\\langle \\theta \\rangle_D)^2 \\rangle_D-\\langle \\sigma_{\\alpha\\beta} \\sigma^{\\alpha\\beta} \\rangle_D\\,. \\end{equation} Here $\\theta$ denotes the dust expansion parameter and $\\sigma_{\\alpha\\beta}$  denotes the shear tensor. Notice that the quantity $\\langle Q \\rangle_D$  vanishes for a homogeneous and isotropic Universe, but becomes nonzero if one includes density perturbations. Furthermore, Eq. (\\ref{adbprime}) implies that  a sufficiently large value of $\\langle Q \\rangle_D$ could produce a negative value for $\\rho_{eff}+ 3p_{eff}=\\langle \\rho \\rangle_D -({1}/{4\\pi})\\langle Q \\rangle_D$, and by virtue of Eq. (\\ref{adbprime}) give rise to an accelerated expansion  $a''_D>0$.  While the effective scale factor $a_D$ is a mathematically well defined quantity, its relation  to  cosmological observations is not completely clear. The physical interpretation of $a_D$ faces three main difficulties. First, $a_D$  is a quantity defined on a  spacelike hypersurface, and so, in general, it  can not be directly related  to cosmological observations which are determined by quantities on the  past lightcone of the  observer. Second, the time evolution of $a_D$ does not provide sufficient information to allow calculation of cosmological observables such as luminosity distance as function of redshift, which requires a metric for its calculation. Third, defining a quantity related to a constant time hypersurface is  somewhat  arbitrary, since one is free to choose a different  time coordinate that defines a different  spacetime foliation. Despite these difficulties,  it has been argued that if the Buchert formalism predicts an effective declaration parameter $q_D$ which is approximately $-1/2$, then it is likely that the predicted value of the actual declaration will also be large and negative. In this paper we will adopt this point of view, and ignore the above mentioned difficulties with the interpretation of $a_D$.  In Ref. \\cite{lisch1}, Buchert's formalism is used to calculate the evolution of $a_D$ in a perturbed Friedmann Robertson Walker (FRW) irotational dust Universe. In particular the term $\\langle Q \\rangle_D$  which presumably drives the accelerated expansion of the Universe is calculated to second order in perturbation theory, and is found to be a boundary term, depending only on the metric perturbations on the boundary of the domain $D(t)$. This result generalizes other previous computation using Newtonian cosmological perturbation theory, which calculates a quantity related to $\\langle Q \\rangle_D$ which is also found to be a boundary term \\cite{aei}. However,  there is a difficulty in reconciling this property  of $\\langle Q \\rangle_D$ at second-order with the interpretation of  $\\langle Q \\rangle_D$ as the source of backreaction. To see this, suppose that the  Universe is spatially compact, and that the domain $D$ is chosen to be the complete  space. In this case any boundary term must vanish identically, and can have no affect on the time evolution of the Universe. While there is no observational evidence for a compact Universe, the fact that  an FRW Universe has a particle horizon implies that a noncompact Universe is observationally indistinguishable from a spatially compact one as long as the scale of compactness  is larger then the observer's particle horizon at decoupling. This bound  translates into a lower bound on  the scale of compactness today of about  twice the size of the horizon. We are  therefore free to choose the scale of compactness to be roughly twice the horizon size today, without changing anything we can measure. This implies that   $\\langle Q \\rangle_D=0$ for $D\\approx 2\\times{\\rm horizon}$. Now Buchert's formalism is valid for all choices of $D$ and gives no guidelines as to what choice of $D$ to make. This ambiguity is part of the overall problem of relating the Buchert formalism to observations. Yet, it seem plausible  that the correct answer (if it exist) should be roughly $D\\approx {\\rm horizon\\  size\\  today}$. It seems reasonable that the value  of $\\langle Q \\rangle_D$ should  not change much if we reduce $D$ from $2\\times{\\rm horizon}$ to roughly the horizon size, and if so, for this choice of $D$, the Buchert formalism predicts that expansion of the Universe is unaffected by the vanishing $\\langle Q \\rangle_D$ at second-order. Nevertheless, there remains the possibility that third order and higher order perturbations could produce a large backreaction effects \\cite{Notari} which can not be represented by a boundary term so the backreaction issue is not settled.  In this paper we argue that general considerations suggest that it is hard to reconcile a large cosmological  backreaction described by the Buchert formalism with observational constraints coming from measurements of luminosity distance and angular-diameter distance as functions of redshift. Observations of the the first Doppler peak of the CMBR together with baryon acoustic oscillation and supernovae data has been used to severely constrain the spatial curvature of a  ${\\rm\\Lambda CDM}$   Universe. By combining these observations with the assumptions that the dynamics of Universe is governed by the FRW metric and that the effect of density perturbations is negligible it has been found that spatial curvature satisfies $\\Omega_K=-0.0052\\pm0.0064$ ($68\\%$ CL) \\cite{WMAP}, where  $\\Omega_K=-R_3(t_0)/(6H_0^2)$, $t_0$ denotes the current time and $H_0$ denotes the current Hubble rate. In this paper we shall consider a more general theoretical framework that include perturbations and possibly large backreaction. One might expect that these observations should also place constraints on the average curvature $\\langle R_3 (t_0)\\rangle_D$. In order for the Buchert formalism to reproduce the desired accelerated expansion from backreaction alone, it must  have a significant averaged curvature with $0.975 \\le  \\Omega_k  \\le 1.294 $ (see Sec. \\ref{k0}), where we have defined  $\\Omega_k\\equiv-\\langle R_3 (t_0)\\rangle_D/[6H_D^2(t_0)]$,  and denoted the effective Hubble rate by  $H_D=a'_D/a_D$. This large averaged curvature seems to be  hard to reconcile with the flat Universe implied by observations.  One possible avenue for evading this observational constraint in spatial curvature, suggested in Ref. \\cite{buchert2}, is the fact that in the Buchert formalism the effective energy density in the spatial curvature need not have the standard scaling $\\propto a^{-2}$. It is not clear whether the strong constraints coming from CMBR are more sensitive to the low redshift curvature or high redshift curvature. If the constraint principally applies to high redshift curvature, then an evolving curvature that is negligible at high redshift could evade the CMBR constraints. In this paper  we shall argue that this avenue for evading the constraint is unlikely. Any  nonstandard time evolution of  the spatial curvature is quite constrained, since at  high redshifts  the density perturbations evolve linearly and the Universe is accurately described by  a weakly perturbed CDM Universe. Non-standard time evolution must therefore be confined to the low redshifts, where nonlinear effects are presumably  important. However, in this regime the spatial  curvature is constrained by the requirement that the backreaction formalism reproduces the correct  luminosity distance as function of redshift that agrees with supernovae data. These observations constrain the time evolution of the metric. Therefore, a nonstandard time evolution of the curvature in this regime would require a nonstandard evolution of the metric such  that  supernovae data observations are reproduced despite  the large spatial curvature. In this scenario  the full time evolution of the metric has two regimes. In the first low redshift regime, the metric evolves in a highly non-standard manner, and in the second high redshift regime, it evolves according to a standard weakly perturbed CDM Universe. The difficulty  that it is not guaranteed that this time evolution reproduces  the correct angular  power spectrum as measured by WMAP.  In this paper we construct a simple toy model that illustrates the observational difficulties that arise in models with a large value of averaged spatial curvature today, even allowing for nonstandard evolution of that curvature. For this purpose, we adopt the point of view of the Buchert backreaction formalism, and assume that we can replace the actual spacetime geometry by a set of averaged quantities. To be able to make predictions, we assemble these quantities and construct an averaged metric that allows us to calculate observables. Here we should make the following remarks.   First, the spatial averaged curvature in the Buchert formalism need not be dominated by low spatial frequency components, it may be mostly high spatial frequency components. Nevertheless, CMBR photons traveling along our past lightcone experience some sort of average curvature along  their way. While this average is different from that of the Buchert formalism, it is plausible that  they are not too much different. We will not address this issue in this paper. Second, in this paper we shall calculate an average spatial curvature using an expression for an averaged metric. However,  this calculation is in general different from an average of the curvature of the true metric. We shall ignore this discrepancy in this paper.  We construct  the following averaged metric toy model \\begin{equation}\\label{metric} d{s}^2=a^2(\\eta)\\left[-d\\eta^2+\\frac{dr^2}{1-k(\\eta)r^2}+r^2d\\Omega^2\\right]\\,. \\end{equation} The  time coordinate $t$ is related to $\\eta$ by $dt^2=a^2 d\\eta^2$. It should be emphasized that this metric should be thought of as an averaged metric and so it does not have to satisfy the Einstein field equations. Here, $a(\\eta)$ and $k(\\eta)$ are certain functions of the conformal time $\\eta$, where we set the present value of the scale factor to unity $a(\\eta_0)=1$. This averaged metric is designed to allow for  a time evolution of the averaged spatial curvature to mimic what is presumably produced by backreaction. Notice that for every constant time hypersurface the induced three dimensional metric obtained from (\\ref{metric}) coincides with a corresponding induced three-metric of an FRW constant time slice, and so this three-metric is isotropic and homogeneous about every point. However, for a generic function $k(\\eta)$, the overall four-dimensional spacetime is not maximally symmetric. Finally, we should mention here that after completing this work we  learned that the form (\\ref{metric}) of an averaged metric has been suggested before in Refs. \\cite{buchert2}, see also Refs. \\cite{mwmk,Rasanen2}.  Recently the toy model (\\ref{metric}) was studied in detail by Larena et. al. \\cite{buchert3}.  This study claims that there is a good agreement between this toy model and data from WMAP and supernovae observations, while we reach the opposite conclusion. We believe that the reason for this discrepancy originates from the fact that Larena et. al.  use a different expression for  the redshift in terms of the scale factor and the function $k(\\eta)$. In their study  it is argued that under some approximation the relation between redshift and scale factor is the standard $1+{z}\\propto a^{-1}$  relation [see  their Eq. (31)]. Using the standard definition of redshift (\\ref{rshift1}) we show that the nonstandard time evolution of the spatial curvature significantly changes this relation, and the correct relation is  given by Eq. (\\ref{rshift3}). As we show, this nonstandard expression for the redshift has a significant effect on the calculation of observables in this model.   Our goal in this paper is to confront the model (\\ref{metric}) with observations. For this purpose we calculate the luminosity distance $D_L(z)$ and angular diameter distance $D_A(z)$ as functions of redshift, using the metric (\\ref{metric}) and compare the results with observational constraints\\footnote{In practice, it is sufficient to calculate $D_A(z)$, since $D_L(z)$ can be obtained from  the relation $D_L(z)=(1+z)^2 D_A(z)$ which is valid in any spacetime, see Refs. \\cite{Ellis,Etherington}. } . We start by choosing a set of functions $k(\\eta)$ parametrized by two parameters [see Eq. (\\ref{ketadef}) below]. For each set of values of the parameters we then choose $a(\\eta)$ to enforce the equation $D_L(z)=D^{{\\rm\\Lambda}CDM}_L(z)$ at low redshifts, where $D^{{\\rm\\Lambda}CDM}_L(z)$ is the luminosity distance derived from a $\\Lambda CDM$ FRW model with parameter values agreeing with supernovae observations. This equation is enforced up to a maximum redshift. Using this requirement we calculate $a(\\eta)$  for this low redshift part of the spacetime. We focus attention only on those metrics in which the the spatial curvature vanishes at a large redshift. Once the averaged spatial curvature vanishes  the backreaction effect should vanish  as well, and so in this high redshift regime we assume that $a(\\eta)$ follows  the standard evolution of a $CDM$ cosmology without a cosmological constant. Using this law of evolution we calculate the function $a(\\eta)$ for the remaining part of spacetime. Once we have  calculated $a(\\eta)$, we use the metric (\\ref{metric}) to  compare  the characteristic angular scale of the CMBR power spectrum as derived from our model with  observation of  WMAP. We find that generically the characteristic angular scale of our model is at odds with WMAP  observations.  This paper is organized as follows. In Sec. \\ref{findfunctions} we explain in detail how we determine the the functions $a(\\eta)$ and $k(\\eta)$ of our model . In Sec. \\ref{firstpeak} we calculate the sound horizon that determines the location of first CMBR peak in our model. In Sec. \\ref{results}  we  explore various values of the parameters which determines the function $k(\\eta)$ and describe the results.  \\section {Construction of the backreaction model} \\label{findfunctions}  To be able to interpolate between an initially vanishing averaged spatial curvature that  corresponds to a weakly perturbed FRW Universe, and  a current large value of averaged spatial curvature needed for the backreaction picture, we  assume that the function  $k(\\eta)$ in the metric (\\ref{metric}) takes the following form \\begin{equation}\\label{ketadef} k(\\eta) = \\begin{cases} H_0^2 \\bar{k}\\frac{f^2} {f^2+1} & , \\eta\\ge\\bar{\\eta} \\\\ 0 & , \\eta\\le\\bar{\\eta}  \\end{cases} \\end{equation} where \\[ f\\equiv\\frac{H_0(\\eta - \\bar{\\eta})} {w}    \\,. \\] Here $H_0=(a^{-1} \\frac{da}{dt})_{t_0}$ is the value of the Hubble constant today, and $\\bar{k}$, $w$ and $H_0 \\bar{\\eta}$ are dimensionless parameters. The parameter $\\bar{\\eta}$ marks the conformal time of the transition between a conformally flat spacetime and a conformally curved spacetime, and $w$ governs the rapidity of this transition. Using the definition $\\Omega_k\\equiv-\\langle R_3 (t_0)\\rangle_D/[6H_D^2(t_0)]$ together with  Eq. (\\ref{ketadef}) and identifying $H_0$ with $H_D(t_0)$ we find that \\begin{equation}\\label{omegak} \\Omega_k=-\\bar{k}\\left(1+\\frac{w^2}{(\\eta_0 - \\bar{\\eta})^2}\\right)^{-1}\\,, \\end{equation} where $\\eta_0$ is the value of the conformal time today. Below we explore various values for the parameters $w$ and $\\bar{\\eta}$, while $\\bar{k}$ is determined from the requirement that the Buchert formalism gives an equation of state parameter of dark energy near $-1$ [see  Sec. \\ref{k0}].  To calculate  the scale factor $a(\\eta)$, we demand that the luminosity distance as function of redshift,  $D_L^{model}(z)$, in our model matches observational data. Since the luminosity distance of a ${\\rm\\Lambda}CDM$  cosmology  matches observational data we impose \\begin{equation}\\label{fitdl} D^{model}_L(z)=D^{{\\rm\\Lambda}CDM}_L(z)\\,, \\end{equation} where  throughout the superscripts 'model' and $'{\\rm\\Lambda}CDM'$ refer to our model and to a flat ${\\rm\\Lambda}CDM$ cosmology, respectively. We impose the condition (\\ref{fitdl}) only at low redshifts, in the range of values of $\\eta$ given by $\\eta\\ge \\bar{\\eta}$. At high redshifts, we switch to imposing the Friedmann equation, since backreaction should be negligible at high redshifts. The evolution of $a(\\eta)$ for $\\eta\\le \\bar{\\eta}$ is determined from an Einstein- de Sitter model for which the scale factor satisfies \\[ \\dot{a}=h \\sqrt{ a}\\,, \\] where an overdot denotes differentiation with respect to $\\eta$. We determine the constant $h$ by demanding continuity of $\\dot{a}$ at the transition time $\\bar{\\eta}$. The solution of this equation is given by \\begin{equation} a(\\eta)=\\left[\\frac{h}{2}(\\eta-\\bar{\\eta})-\\sqrt{a(\\bar{\\eta})}\\right]^{2}\\ ,\\ \\eta\\le \\bar{\\eta} \\,. \\end{equation} where  we used the continuity of   $a({\\eta})$  at $\\eta=\\bar{\\eta}$ .   \\subsection{Matching luminosity distances as function of redshift}  In this section we describe how we  calculate $a(\\eta)$  in practice from the matching  requirement (\\ref{fitdl}) . In a general spacetime the luminosity distance $D_L(z)$ is related to the angular diameter distance $D_A(z)$ by \\cite{Etherington,Ellis} \\begin{equation}\\label{dlda} D_L(z)=D_A(z)(1+z)^2\\,. \\end{equation} Using  Eq. (\\ref{dlda}), the matching  requirement (\\ref{fitdl}) takes the form of \\begin{equation}\\label{fitda} D^{model}_A(z)=D^{{\\rm\\Lambda}CDM}_A(z)\\,. \\end{equation} In  a flat ${\\rm\\Lambda}CDM$ cosmology the right hand side of Eq. (\\ref{fitda}) is given by \\begin{equation}\\label{dafrw} D^{{\\rm\\Lambda}CDM}_A(z)=\\frac{1}{(1+z)H_0}\\int_0^z [\\Omega_m(1+z')^3+\\Omega_\\Lambda]^{-1/2}dz'\\,. \\end{equation} Here the parameters $\\Omega_m$ and $\\Omega_\\Lambda$ satisfy $\\Omega_m+\\Omega_\\Lambda=1$, and the contribution from radiation energy-density has been neglected since we  confine the discussion to the epoch after recombination.  We now consider the left hand side of Eq. (\\ref{fitda}) and derive an expression for  $D^{model}_A$. Suppose that an observer views a sizeable distant object (e.g. a distant galaxy or a structure of the CMB anisotropy) that has  a  transverse proper cross sectional area $\\delta A$, and subtends a small solid angle $\\delta{\\Omega}$. From these quantities the observer can determine the angular diameter distance \\[ D_A=\\sqrt{\\frac{\\delta A}{\\delta\\Omega}}\\,. \\] Since the wavelength of the electromagnetic radiation  is typically much  smaller than  the spacetime curvature, we can safely use the geometric optics approximation and describe the electromagnetic radiation as a bundle of light rays that trace a congruence of null geodesics. We assume that  the light rays converge at  an event $p$ at the location of the observer, which is chosen to be at the origin, so  that $r(p)=0$ and  $\\eta(p)=\\eta_0$. Since the metric  (\\ref{metric}) is isotropic about the origin, the light rays trace radial null geodesics from the source at $r(\\eta)$, where $\\eta<\\eta_0$, to the observer; and  the angular diameter distance is given by \\begin{equation}\\label{daar} {D}^{model}_A=a(\\eta)r(\\eta) \\,. \\end{equation} Substituting Eqs. (\\ref{dafrw}) and (\\ref{daar}) into Eq. (\\ref{fitda}) and differentiating with respect to $\\eta$ gives \\begin{equation}\\label{fit2} H_0 \\frac{d}{d\\eta}[(1+z)a r]=\\frac{dz}{d\\eta} [\\Omega_m(1+z)^3+\\Omega_\\Lambda]^{-1/2} \\,. \\end{equation} Our goal to solve Eq. (\\ref{fit2}) for $a(\\eta)$. As a preliminary step, we first calculate $r(\\eta)$ and $z(\\eta)$ and then substitute these functions into Eq. (\\ref{fit2}).  The  calculation of  $r(\\eta)$  for radial null geodesics follows directly from the metric (\\ref{metric}), which gives $\\dot{r}^2=1-k(\\eta)r^2$, where dot denotes differentiation with respect to $\\eta$. Later we shall assume that $k(\\eta)<0$ and so $\\dot{r}^2>0$ implies that the null geodesics have  no turning points. Since $\\eta$ is a monotonically decreasing function of $r$  we have \\begin{equation}\\label{rdot} \\dot{r}=-\\sqrt{1-k(\\eta)r^2} \\,. \\end{equation} Below we use Eq. (\\ref{ketadef}) to specify $k(\\eta)$ and solve Eq. (\\ref{rdot}) numerically together with the initial condition $r(\\eta_0)=0$.  We now consider the calculation of the redshift $z(\\eta)$. By definition the redshift is given by \\begin{equation}\\label{rshift1} 1+{z}=\\frac{{({k}^\\alpha {u}^\\beta {g}_{\\alpha\\beta})}_{source}}{{({k}^\\alpha {u}^\\beta {g}_{\\alpha\\beta})}_{observer}}\\,, \\end{equation} where $k^\\alpha$ is the 4-momentum of the photon, and $u^\\alpha$ is the 4-velocity of the cosmological fluid. For simplicity we assume that both the observer and the source have four-velocities of the form ${u}^\\alpha=a^{-1}\\delta^\\alpha_\\eta $ meaning that their peculiar velocities  vanish. We normalize $k^\\alpha$ by demanding that ${k}^\\alpha {u}^\\beta g_{\\alpha\\beta}=-1$ at the observer. Some simplification is gained by  considering a conformal transformation  of the form \\begin{equation}\\label{conformal} g_{\\alpha\\beta}=a^2 \\hat{g}_{\\alpha\\beta} \\ \\ , \\ \\ {k}^\\alpha=a^{-2}\\hat{k}^\\alpha\\,. \\end{equation} Combining our choices for normalization and velocities with Eqs. (\\ref{rshift1}) and (\\ref{conformal}) gives \\begin{equation}\\label{rshift2} {z}=a^{-1}\\hat{k}^\\eta-1\\,. \\end{equation} The conformal null  vector field $\\hat{k}^\\eta$  satisfies a geodesic equation in the conformal spacetime where the metric is $\\hat{g}_{\\mu\\nu}$, and so it is independent of $a(\\eta)$. Using this geodesic equation together with Eq.(\\ref{rdot})  we obtain \\begin{equation}\\label{ketadot} (\\hat{k}^\\eta)^{-1}  \\frac{d}{d\\eta}\\hat{k}^\\eta =-\\frac{r^2\\dot{k}}{2(1-kr^2)}\\,. \\end{equation} Integrating Eq. (\\ref{ketadot}) and using  Eq. (\\ref{rshift2}) we find that the red shift is given by \\begin{equation}\\label{rshift3} {z}=a^{-1}e^{1/2\\int_{\\eta}^{\\eta_0} r^2\\dot{k} {(1-kr^2)}^{-1} d\\eta' }-1 \\,. \\end{equation} Notice that for $k={\\rm const}$ Eq. (\\ref{rshift3}) reduces to the FRW relation $z+1\\propto 1/a$. Eq. (\\ref{rshift3}) is at odds with Eq. (30) of Ref. \\cite{buchert3} which seems to be inconsistent with the standard definition of redshift (\\ref{rshift1}). Below we solve for  $r(\\eta)$ by specifying $k(\\eta)$ and solving Eq. (\\ref{rdot}). Using $r(\\eta)$ we evaluate the integral in Eq. (\\ref{rshift3}) and obtain  $z(\\eta)$. Both $r(\\eta)$ and $z(\\eta)$  are then substituted into  Eq. (\\ref{fit2}) which is solved to give $a(\\eta)$. Finally $a(\\eta)$ and $r(\\eta)$ are inserted into Eq. (\\ref{daar}) to obtain ${D}^{model}_A(\\eta)$, and this is combined with $z(\\eta)$ to obtain the angular diameter distance $D_A^{model}(z)$ as function of redshift $z$. All these calculations are done numerically.  \\subsection{Constraining the toy model parameters}\\label{k0}  In this section we use observational constraints  on the cosmological parameters to  place constraints on the parameters of our toy model. The calculation is based on Buchert's formalism  \\cite{buchert1}  which was summarized in Sec. \\ref{intro}. In this formalism  the equation of state parameter of dark energy is given by \\begin{equation}\\label{eqofstate} w_{de}(\\eta)=\\frac{p_{de}(\\eta)}{\\rho_{de}(\\eta)}\\,, \\end{equation} where   $\\rho_{de}$ denotes the dark energy density, and  $p_{de}$ denotes the dark energy pressure. These quantities are related to the effective density  $\\rho_{eff}$ and effective pressure $p_{eff}$, which are given by  Eqs. (\\ref{rhoeff}) and (\\ref{peff}), trough the relations  $\\rho_{eff}=\\langle \\rho \\rangle_D+\\rho_{de}$ and $p_{eff}=p_{de}$. Eq. (\\ref{eqofstate}) together with Eqs. (\\ref{rhoeff},\\ref{peff}) give \\begin{equation}\\label{wde} w_{de}=\\frac{-(1/3)\\langle R_3\\rangle_D+\\langle Q\\rangle_D}{\\langle Q\\rangle_D+\\langle R_3\\rangle_D}\\,. \\end{equation} The current value of this parameter is in the range  $-1.1 \\le w_{de}(\\eta_0) \\le -0.9$ \\cite{woodetal}.  Using  this constraint together with Eq. (\\ref{wde}) gives \\begin{equation}\\label{qr} -\\frac{57}{17}\\langle Q  (\\eta_0)\\rangle_{D}\\le \\langle R_3 (\\eta_0) \\rangle_{D} \\le -\\frac{63}{23}\\langle Q  (\\eta_0)\\rangle_{D} \\,. \\end{equation} We define the effective deceleration parameter  by \\begin{equation}\\label{qdef} q_D=-\\frac{a''_D}{a_D H_D^2}\\,, \\end{equation} and substitute Eq. (\\ref{adbprime}) into Eq. (\\ref{qdef}) and use  Eqs. (\\ref{rhoeff}) and (\\ref{peff}). This gives \\begin{equation}\\label{qd} q_D(\\eta_0)=\\frac{1}{2}\\Omega_{m(D)}-\\frac{\\langle Q(\\eta_0)\\rangle_D}{3H_D^2}\\,, \\end{equation} where $\\Omega_{m(D)}=8\\pi\\langle \\rho \\rangle_D/3H_D^2$. We demand that this expression be equal to the current deceleration  of a standard flat ${\\rm \\Lambda}CDM$ cosmology, where the deceleration parameter is given by \\begin{equation}\\label{q0frw} q_{\\rm \\Lambda CDM}(\\eta_0)=\\frac{1}{2}\\Omega_m-\\Omega_\\Lambda\\,. \\end{equation} where  $\\Omega_m$ and $\\Omega_\\Lambda$ are the ${\\rm \\Lambda CDM}$ densities of matter and dark energy, respectively. Assuming that we can substitute   $\\Omega_m$ in place of $\\Omega_{m(D)}$ we find from Eq. (\\ref{qd}) and Eq.(\\ref{q0frw}) that \\begin{equation}\\label{qom} \\frac{\\langle Q(\\eta_0)\\rangle_D}{3H_D^2}=\\Omega_{\\Lambda}\\,, \\end{equation} The 5-year WMAP data \\cite{WMAP} reveals that  the dark energy density  is in the range  $\\Omega_{\\Lambda}=0.742\\pm 0.030$. We use the WMAP constraint on  $\\Omega_{\\Lambda}$ and  Eqs. (\\ref{qom}), (\\ref{qr}) together with the definition  $\\Omega_k\\equiv-\\langle R_3 (t_0)\\rangle_D/[6H_D^2(t_0)]$, to  obtain \\begin{equation}\\label{findk0} 0.975 \\le  \\Omega_k  \\le 1.294 \\,. \\end{equation} By combining $ \\Omega_k\\approx1.1$ together with  Eq.  (\\ref{omegak}) and Eq. (\\ref{findk0}) we determine  the parameter $\\bar{k}$ once the parameters $w$ and $\\eta_0$ have been specified. ",
        "conclusions": "In this paper we studied a toy model of a backreaction mechanism. In this model the averaged spatial curvature grows at low redshifts so that the expansion of the Universe presumably induced by backreaction could be consistent with supernovae data. In the high redshift regime,  we assumed that Universe evolves according to a standard weakly perturbed CDM Universe. We showed that this model  alters the predictions for the sound horizon at decoupling and that it is generically inconsistent with the power spectrum as measured by WMAP."
    },
    "0809/0809.0334_arXiv.txt": {
        "abstract": "Inhomogeneous cosmological models have recently become a very interesting alternative to standard cosmology. This is because these models are able to fit cosmological observations without the need for dark energy. However, due to inhomogeneity and pressure-less matter content, these models can suffer from shell crossing singularities. These singularities occur when two shell of dust collide with each other leading to infinite values of the density. In this {\\em Letter} we show that if inhomogeneous pressure is included then these singularities can be prevented from occurring over the period of structure formation. Thus, a simple incorporation of a gradient of pressure allows for more comprehensive studies of inhomogeneous cosmological models and their application to cosmology. ",
        "introduction": "There has been a significant amount of recent interest in inhomogeneous cosmological models as an alternative to the $\\Lambda$CDM concordance model [for a review see \\citet{C07}].  These inhomogeneous models are able to fit many cosmological observations without the need for dark energy  \\citep*{DH98,C00,GSS04,AAG06,AA06,CR06,AA07,EM07,ABNV07,BTT07a,BTT07b,BMN07,MKMR07,KKMA08,BN08,B08,GH08a,E08,YKN08}.  One principle method underlying these schemes is to use inhomogeneous solutions of the Einstein field equations rather than the standard Friedmann-Robertson-Walker solutions used in $\\Lambda$CDM models. The implementation of these models suggests that we need to live close to the centre of the Gpc-scale void. There have already been several methods suggested as a test for these types of models. They are based on the time drift of cosmological redshift \\citep*{UCE08}, spectral distortions of the CMB power spectrum \\citep{CS08}, the kinematic Sunyaev--Zel'dovich effect \\citep{GH08b}, future measurements of supernova in the redshift range of 0.1-0.4 \\citep{CFL08} and future Baryon Acoustic Oscillations measurements \\citep{BW08}. Based on current observations, the alternative of a Gpc-scale underdensity is indistinguishable from the dark energy scenario.   However, there remain a number of concerns regarding these inhomogeneous cosmologies.  In particular, it is known that pressure-free Lema\\^itre--Tolman \\citep{L33,T34} models can evolve to form shell crossing singularities.  This is an additional singularity to the Big Bang that occurs when two shells of matter collide with each other, leading to infinite values of the density.  These singularities were first discussed in the context of astrophysical applications of gravitational collapse, where it has been shown that they can form without being hidden inside an horizon, and are therefore globally naked \\citep{YSM73}.  However, they are not considered as real singularities for a number of reasons.  First of all because they are weak in the sense that the spacetime can be extended through the singularity \\citep{N86,FK95,N03,PK}.  Secondly, as was shown by \\cite{J93}, an object sent through the singularity would not be crushed (i.e. they would not be focused onto a surface or a line).  Finally, in spherically symmetric, asymptotically flat Einstein-Vlasov systems, these types of singularities have been proved not to occur \\citep{RRS95}.  However, this has not been proved in cosmological models which are not asymptotically flat.  In gravitational collapse scenario's, shell crossings generally do not require consideration because the initial conditions are generally unrealistic.  However, this is not the case in cosmological models.  \\cite{BKH05} showed that a large class of realistic models of voids exhibit shell crossing singularities when the galaxy wall surrounding the void began to form.  In many applications, initial conditions are chosen such that the subsequent evolution does not exhibit shell crossings \\citep{HL85}, however this is extremely inconvenient and prevents the study of a wide class of models.  In this {\\em Letter} we show that the simple incorporation of pressure gradients leads to an elimination of shell crossings over the time scales involved with structure formation.  Anomalies in the evolution associated with the formation of acoustic oscillations (see Sec. \\ref{acousticosc}) prevent us from determining conclusively whether this simple incorporation of inhomogeneous pressure can permanently prevent the shell crossing singularities.  The structure of the {\\em Letter} is as follows; In Sec. \\ref{ssist} the Lema\\^itre model is presented. Sec. \\ref{setup} gives an explicit specification of the models being considered.  Sec. \\ref{evolution} presents the evolution of both models and shows that sufficiently large pressure gradients can prevent shell crossing singularities. Sec. \\ref{acousticosc} then discuss the occurrence of acoustic oscillations due to non-zero pressure gradient. ",
        "conclusions": "In this {\\em Letter} we studied the shell crossing singularities which occur in pressure-free spherically symmetric cosmological models. These models have recently become very popular because they can fit cosmological observations without the need for dark energy. However, because of shell crossing singularities, which occur when two shells of matter collide, the full parameter space cannot be considered. We have investigated the incorporation of pressure into these models. Adding a gradient of pressure implies the equations can no longer be solved analytically, although conceptually they are no more difficult -- see eq. (\\ref{mr}) -- and they are still simple to compute numerically. We showed that the incorporation of a pressure gradient can prevent shell crossing singularities occurring over the period of structure formation. This will enable the unlimited investigation of these inhomogeneous cosmological models without the occurrence of these anomalous singularities.  We also showed in this {\\em Letter} that, in some cases, when the density contrast is sufficiently high, the existence of pressure gradients leads to acoustic oscillations. At this stage the numerical investigation becomes troublesome. However, by choosing a harder form of the equation of state we can postpone the occurrence of shell crossing singularities so they do not occur  over the time of structure evolution. It should be noted that such harder equations of state require that the constant $K$ in (\\ref{peos}) be unnaturally large -- by several orders of magnitude larger that the realistic value. This means that to fully describe the process of acoustic oscillations we need to perform a more comprehensive analysis. In particular, an application of more realistic fluids which allows for an anisotropic pressure and heat flow is essential."
    },
    "0809/0809.1738.txt": {
        "abstract": "We present the results of a pair of 100 ksec {\\it Chandra} observations in the Small Magellanic Cloud to survey HMXBs, stars and LMXBs/CVs down to $L_x$ = 10$^{32}$ erg/s. The two SMC deep-fields are located in the most active star forming region of the bar. Deep Field-1 is  positioned at the most pulsar-rich location (identified from previous surveys). Two new pulsars were discovered in outburst: CXO J004929.7-731058 (P=894s),  CXO J005252.2-721715  (P=326s), and 14 candidate quiescent pulsars were identified from their timing and spectral properties. Out of 12 previously known pulsars in the fields, 9 were detected, with pulsations seen in five of them. This demonstrates for the first time that a significant fraction (at least 60\\%) of these systems have appreciable accretion driven X-ray emission during quiescence. Two known pulsars in the field were not detected, with an upper limit of $L_x$ $<$ 5 $\\times$10$^{32}$. The full catalog of 394 point-sources is presented along with detailed analyses of timing and spectral properties. Future papers will report associated observations obtained with HST and Magellan to identify optical counterparts.  \\end {abstract} ",
        "introduction": "The Small Magellanic Cloud has proven to be an incomparable laboratory for astrophysics. The SMC is dwarf irregular satellite of the Milk Way, and unlike all other members of that group is experiencing an era of intense star  formation. Situated in a part of the sky unobstructed by the galactic plane, the SMC affords a ringside seat to the unfolding drama. The combination of low extinction ($nH\\sim$10$^{21}$) and a small physical dimensions in relation to its distance from earth (53 kpc ) effectively puts the entire population at a common distance. A substantial fraction of the SMC can be observed simultaneously thanks to its compact size, which facilitates population studies.  Monitoring surveys with e.g.  RXTE \\cite{laycock2005} only see pulsars in outburst at $>$10$^{36}$ yet most X-ray binary pulsars spend the majority of the time in quiescence. In order to obtain a measure of the true population, it is necessary to observe and count X-ray pulsars in their low luminosity state. The very concept of a quiescent low-luminosity state is an open hypothesis, so we target a densely populated environment at known distance and low nH with sufficient sensitivity (10$^{33}$erg/s) to detect the unseen majority. Based on observational evidence and theoretical calculations we predicted that Chandra could detect these systems, or place strong restrictions on the quiescent emission scenarios if none were found.  The focus of this paper is the search for X-ray pulsars in a quiescent state. That is, exhibiting emission (whether pulsed or not), at times other than during a type I or \"normal\" periastron outburst (and excluding type II, \"giant\" outbursts). In this work ~\\ref{sect:known} we present new evidence for quiescent emission in previously known pulsars, whose expected outburst dates are predicted by ephemerides. However if the SMC harbors a large population of systems that spend most of the time in quiescence it will not be possible to find their orbital parameters by X-ray monitoring. Instead these systems can only be found by deep pointed X-ray observations with sufficient angular resolution to guide searches for optical counterparts. The orbital parameters can then be measured by photometric monitoring or radial velocity studies. There are two compelling scenarios that  lead us to expect that such a population exists.  (1) The physics of magnetically channelled accretion (see for example \\cite{king}) predicts the existence of a centrifugal barrier. Magnetic field lines emanating from the neutron star's polar regions co-rotate with the NS itself, causing plasma attached to the field lines beyond the co-rotation radius (where the keplerian orbital period equals the NS rotation period) to be flung away by centrifugal force. For accretion to occur, the ram-pressure of the infalling matter must drive the magnetosphere boundary inside the co-rotation radius ($r_c$). Matter orbiting inside $r_c$ has angular velocity equal to or greater than co-rotation, and is channelled onto the NS polar caps, spinning up the pulsar in the process. This condition is a strong function of pulse period because the $r_c$ gets progressively closer to the NS surface for increased spin rate. The most rapidly rotating neutron stars are therefor only expected to switch on as pulsars at very high mass transfer rates.  Observational evidence for this picture is mounting, as many more outbursts are seen in SMC pulsars with periods longer than a few tens of seconds, compared to shorter period systems.  (2)The majority of X-ray pulsars, and all but one of those in the SMC with identified optical counterparts, have been found in binary systems with Be-star companions (See  \\cite{mcbride2008} for a recent review of the evidence). The accreted material powering X-ray emission from the pulsar comes from a circumstellar disk around the Be star. The exact mechanism of disk formation is unclear at present, but progress has been made in understanding the disk-NS interaction \\cite{okazaki2001}. Several lines of observational and theoretical evidence point to the role of orbital eccentricity in producing regular ``normal'' (type-I) X-ray outbursts in Be-HMXB. Systems with circular orbits (or nearly so) are thought to exhibit only occasional giant (type-II) outbursts. It is further speculated that these systems can experience constant low-level accretion, yielding X-ray luminosites not detectable by the pre-Chandra generation of instruments. In the Milky-Way, the Be/X system X Persei is the prototype of this class with eccentricity=0.11, and there are at least two other galactic examples \\cite{delgado2001}. X Per fortuitously lies somewhat out of the galactic plane otherwise it would be hidden by dust/gas, as we presume are many of its brethren. The X-ray luminosity of X Per varies, but is seldom above $L_x$=10$^{35}$ erg/s.  Our expectation is that X-ray binaries containing fast-rotating neutron stars, and those with circular orbits are under-represented in current catalogs. In both cases due to low quiescent luminosity and the infrequent occurrence of giant outbursts. One of our aims is to provide a complete sample of X-ray binaries, with greatly reduced luminosity bias compared to the existing sample. Stellar population synthesis models will greatly benefit from hard numbers upon which to base the relative probabilities of various evolutionary paths. For example in modelling the supernovae that occur in formation of XRBs, one needs to reproduce the actual distribution of orbital periods and eccentricities. If current catalogs under represent systems with large orbital separations and low $\\eta$ the input to these models will be flawed.  Studies of archival Chandra data for a large number of galaxies by \\cite{grimm2003} have found an apparently universal correlation between star formation rate (SFR) and the integrated luminosity of HMXB. The absolute number normalization of this relationship is however unknown, because sufficiently faint fluxes have not been probed to discover the underlying numbers of HMXBs. It is for example physically relevant to know the number of neutron stars formed in binaries per unit SFR, which one might not be able to tell purely from the luminosity function, if it is missing a significant sub-class of HMXB.  We present the complete X-ray source catalog as an electronic table accompanying this paper. We describe below in detail the properties of the brightest sources, and give an analysis of the timing and spectral properties of the sample, focussing on X-ray pulsars. An optical survey is being conducted in parallel \\cite{antinou2008} to enable more robust and complete classifications of X-ray binaries, Cataclysmic variables, stellar coronae and background active galactic nuclei (AGN). Initial results are reported by \\cite{zezas2008}. ",
        "conclusions": "The complete source catalog is available on-line in electronic form, it contains positions, count-rates, fluxes and quantiles for 394 X-ray sources (211 in DF1 \\& 183 in DF2).  We have detected 16  X-ray pulsar candidates in the SMC Chandra Deep Fields (7 in DF-1, 9 in DF-2) based primarily on periodogram analysis of the lightcurves. Two of the newly discovered pulsars were in outburst, and their periods were measurable to a high degree of accuracy. We have therefore assigned the names SXP892 \\& SXP326 following the convention used by the SMC RXTE monitoring project \\cite{laycock2005} , \\cite{galache2008}.  An additional three candidates were detected with sufficient S/N to place them unambiguously on the quantile diagram in the region inhabited by pulsars, in tight agreement with the powerlaw spectral model ($\\Gamma$=1, nH$\\sim$10$^{21}$) characteristic of SMC pulsars. The remaining 11 pulsar candidates were identified from periodogram detections with blind search significance levels between 90\\%--95\\%, or in sources with fewer than 250 net counts.  We report these candidate pulsars without an SXP name due to the uncertainty attached to their period determinations, for example we cannot be confident that the fundamental is the most significant peak in the periodogram, or that a small fraction ($\\sim$10\\% for the 90\\% significance threshold) of these periods could be spurious. In the following discussion we assume that all the candidate pulsars are real detections.  There were 12 X-ray pulsars known in the survey region prior to these observations, 7 in DF1 and 5 in DF2. Thus we have increased the sample of pulsars in the region to a total of  28 ([7+7] + [5+9]). Most interesting is the fact that we were able to detect all the known pulsars in DF2 and most (4/7) in DF1. One of the  primary objectives of this project was to find out if HMXB pulsars could be detected between outbursts, and to determine their quiescent luminosity.  Using the latest X-ray ephemerides produced by the RXTE monitoring program \\cite{galache2008} we can calculate the (presumed orbital) phase of each pulsar at the time of the Chandra Deep Field observations. The ephemerides were constructed from pulsed-flux lighturves spanning about 10 years of weekly RXTE observations. As such they predict the most likely time of X-ray outburst, which is presumed to reflect orbital modulation of the mass-transfer rate from the mass-donor star to the neutron star (the actual orbital ephemeris is not known for any SMC pulsars except SMC X-1).   Of the 10 pulsars with an X-ray ephemeris fromn RXTE, 6 were detected far from their normal range of outburst-phase and with luminosity well below outburst levels. Two pulsars (SXP59, SXP7.78) were detected during their normal range of outburst phase. Three known pulsars were not detected, which for a requirement of 5 counts in 100ksec gives an upper limit of L$_x <2.6\\times$10$^{32}$erg/s. Of the undetected pulsars, two have X-ray ephemerides and were both observed outside their normal range of outburst-phase.  By detecting 6/10 pulsars for which there are ephemerides,  outside of their predicted outburst phase, we have demonstrated that a significant fraction (at least 60\\%) of these systems have appreciable accretion driven X-ray emission during quiescence. To the six confirmed quiescent pulsars we can add the 14 candidate quiescent pulsars, whose 2-10 keV luminosities are substantially lower than normal Be-pulsar outbursts. This new sample is consistent with a large population -about twice the size of the known pulsar population- of quiescent \"normal\" pulsars. It is also possible that some are in a perpetually low luminosity state, such systems are known in the Galaxy, for example X Per. The theoretical work of \\cite{okazaki2001} informs us that Be/NS systems with low orbital eccentricity accrete at a very low constant rate. The Be star's circumstellar disk in such systems is permanently tidally truncated (at the 3:1 resonance) well within the neutron star's orbital radius, except during large mass-ejections (the assumed cause of type II \"giant\" outbursts). This picture neatly explains how more than 50\\% of the X-ray pulsar population has remained hidden from RXTE and is only now coming to light.  Pulse profiles created for every pulsar detected in the Chandra data show a range of forms. The typical profile shows significant flux above the un-pulsed component for most of the pulsation cycle. A qualitative distinction exists between profiles that are broad and asymmetric (e.g. SXP326), and those that are strongly peaked which tend to be narrower and more symmetric (e.g. SXP172).  These differences are likely due to differences in geometry of the emission regions (polar caps). The literature documents several efforts to model X-ray pulse profiles in terms of combinations of fan and pencil beams (See for example \\cite{parmar1989}). Broad and cuspy profiles are attributed to fan beam geometry, where the X-ray beam has the form of a (somewhat irregular) hollow cone,  due to obscuration of the central part of the polar cap by the incoming accretion stream, or reflection/re-processing effects. Narrow profiles more closely resemble radio pulsars and are accordingly interpreted in the \"light-house\" model, as collimated beams (pencil-beams) emanating directly from the polars caps.  Including only sources with S/N$>$5, and counting the sources populating regions defined on the quantile diagram (Figures~\\ref{fig:quantiles1}b \\& ~\\ref{fig:quantiles2}b), we estimate of the proportion of  candidate HMXB, Stars, AGN \\& obscured AGN in each field. These numbers are then used to estimate the absolute number of XRB, stars \\& AGN in the Deep Fields. We multiply by the fraction of high S/N sources in each category by the total size of the catalog down to the completeness limit of our dataset. Uncertainties are computed by taking  $\\sqrt{N}$ of the number of 5$\\sigma$ sources, propagated to the full sample. We defined a flux limit corresponding 10 count source at the Chandra/ACIS-I aimpoint in our 100 ksec stacked dataset, to approximate the completeness limit of our sample. We do this calculation using a flux limited sample because our flux calculations include correction for Chandra's changes in sensitivity and spectral response with distance from the aim-point. If a count limit were used instead, it would not be uniform across the field. Each source has a 0-5-8 keV flux computed based on its quantile values and the exposure map, a 10 count limit at the aimpoint translates to a flux limit $F_{x} > $10$^{-15}$ erg/cm$^{2}$/s.  In Deep-Field-1 we find 10.7\\% Stars, 53.3\\% AGN (of which 9\\% are obscured), and 35.7\\% XRB. Multiplying these percentages by the number (131) of sources down to our completeness limit  yields 14$\\pm$6 stars, 71$\\pm$17 AGN, and 46$\\pm$13 HMXB.  The same analysis of Deep-Field-2 yields 10.5\\% Stars, 50\\% AGN (of which 7.9\\% are obscured), and 39.5\\% XRB.  Multiplying these percentages by the number (107) of sources down to our completeness limit  yields 11$\\pm$6 Stars, 54$\\pm$16 AGN, and 42$\\pm$14 XRB.  For comparison with our observations, we first estimated the expected number of background AGN in our survey based on the work of \\cite{kim2007}. For the same flux limit of $F_{x} > $10$^{-15}$ erg/cm$^{2}$/s, \\cite{kim2007} Figure 4 shows of order 10$^3$ AGN per sq. degree. Thus for an ACIS-I field of view (16'$\\times$16') one expects to see $\\sim$71 AGN per field.  The purpose of this comparison is to determine whether our proposed large expansion in the X-ray binary population of the SMC is reasonable in the light of what is known about the background AGN  population. The AGN number for DF1 is remarkably consistent with our prediction based on the AGN luminosity function of  \\cite{kim2007} . We note that DF2 contains a large X-ray bright supernova remnant which likely explains the slightly (23\\%) lower AGN numbers in this field. Systematic errors were estimated by varying the boundary of the region in the quantile plane used to classify the AGN. Moving the boundary  between $\\Gamma$=1.2 --1.5  changed AGN vs XRB fractions by an insignificant amount compared to the statistical error.  Based on the above results, we argue for a large quiescent population of High Mass X-ray binaries (HMXB) in the SMC. Only a small fraction of these systems are in outburst at a given instant in time. The work of \\cite{grimm2003} demonstrates a universal relation between the integrated X-ray luminosity of HMXBs in a galaxy and the instantaneous star formation rate. This result now must be seen in the light of a similarly universal duty cycle distribution for HMXB  populations. At any instant, the proportion of systems in outburst is roughly the same. Although more numerous, the quiescent systems evidently do not contribute significantly to the total luminosity of the galaxy.  We note that Cataclysmic variables and LMXB are expected to occupy the same quantile parameter-space as the AGN, due to their soft power-law and/or hot thermal spectra. Given that our AGN-count is consistent with predictions, there is not a great deal of room for additional source populations. Thus we can already constrain the total number of LMXB and CVs to be within the error bars of the AGN estimate.  The further  goals of this project are as follows, and are being pursued in a series of papers in collaboration with our co-proposers. (1) Classify all the sources by finding and studying their optical counterparts. Discriminating the different source classes is a requisite for aims 3-5 below. Our optical survey is being carried out in parallel with the X-ray survey using the Hubble Space Telescope and Magellan 6.5m telescope. With high resolution imaging/photometry to R,I$<$26 we will be able to independently classify the optical counterparts to most of the sample.  The optical counterpart identifications will enable items 2-4 below.  (2) Construct the luminosity distribution down to faint X-ray fluxes, this will test whether multiple point source populations may be present, vs a monolithic HMXB ensemble.  (3) Search for coronal X-ray emission from the most active stars in the SMC.  (4) Constrain the incidence of LMXBs and CVs in the SMC. CVs should be present throughout the SMC which is dominated by an old population of red giants. Outbursts (Novae and dwarf novae) are extremely infrequent and fainter in X-rays than HMXB. Thus they will be harder to find.  No LMXBs or CVs are known in the SMC, and extrapolation from the total mass in stars predicts only a handful of LMXBs , thus an observational constraint (or positive discoveries of LMXB)  will provide important input to population synthesis models."
    },
    "0809/0809.2088_arXiv.txt": {
        "abstract": "NGC~1407 is the central elliptical in a nearby evolved group of galaxies apparently destined to become a galaxy cluster core. We use the kinematics of globular clusters (GCs) to probe the dynamics and mass profile of the group's center, out to a radius of 60 kpc ($\\sim$~10 galaxy effective radii)---the most extended data set to date around an early-type galaxy. This sample consists of 172 GC line-of-sight velocities, most of them newly obtained using Keck/DEIMOS, with a few additional objects identified as dwarf-globular transition objects or as intra-group GCs. We find weak rotation for the outer parts of the GC system ($v/\\sigma \\sim 0.2$), with a rotational misalignment between the metal-poor and metal-rich GCs. The velocity dispersion profile declines rapidly to a radius of $\\sim$~20~kpc, and then becomes flat or rising to $\\sim$~60~kpc. There is evidence that the GC orbits have a tangential bias that is strongest for the metal-poor GCs---in possible contradiction to theoretical expectations. We construct cosmologically-motivated galaxy+dark halo dynamical models and infer a total mass within 60~kpc of $\\sim 3\\times10^{12} M_{\\odot}$, which extrapolates to a virial mass of $\\sim 6\\times 10^{13} M_{\\odot}$ for a typical $\\Lambda$CDM halo---in agreement with results from kinematics of the group galaxies. We present an independent {\\it Chandra}-based analysis, whose relatively high mass at $\\sim$~20~kpc disagrees strongly with the GC-based result unless the GCs are assumed to have a peculiar orbit distribution, and we therefore discuss more generally some comparisons between X-ray and optical results. The group's $B$-band mass-to-light ratio of $\\sim$~800 (uncertain by a factor of $\\sim$~2) in Solar units is extreme even for a rich galaxy cluster, much less a poor group---placing it among the most dark matter dominated systems in the universe, and also suggesting a massive reservoir of baryons lurking in an unseen phase, in addition to the nonbaryonic dark matter. We compare the kinematical and mass properties of the NGC~1407 group to other nearby groups and clusters, and discuss some implications of this system for structure formation. ",
        "introduction": "\\label{sec:intro}  In our current understanding of the universe, stellar systems larger than a globular cluster (GC; baryon mass $\\sim 10^6 M_{\\odot}$) are generally embedded in massive halos of dark matter (DM). Their systemic properties versus increasing DM halo mass show that there must be shifts between the various regimes of physical processes shaping the baryonic properties (e.g., \\citealt{1977MNRAS.179..541R}; \\citealt{2003ApJ...599...38B}; \\citealt{2004MNRAS.355..694M}; \\citealt{2006MNRAS.368....2D,2007MNRAS.381..389K,2007MNRAS.382.1481N,2008MNRAS.385L.116G}). In particular, somewhere between the scales of individual galaxies and clusters of galaxies, the mass of a halo's primary galaxy stops scaling with the halo mass, and reaches a universal upper limit (see e.g., \\citealt{2004A&A...415..931R}; \\citealt{2005ApJ...627L..85C}; \\citealt{2006AJ....131.2018B}; \\citealt{2007arXiv0709.2192K}; \\citealt{2007arXiv0710.3780H}), while the bulk of the baryons become spread out as a hot intergalactic medium (IGM). The high-mass systems behave less like galaxies with their own enveloping halos and unique histories, and more like extended halos shaping the destinies of their captive galaxian members (e.g., \\citealt{2006PASA...23...38F}).  In collapsed galaxy clusters, the dominant processes affecting the larger galaxies involve cooling, heating, and stripping by the IGM, while some level of harassment and stripping also occurs by galaxy fly-bys within the cluster. However, it is thought that processes such as gas strangulation and major galaxy mergers occurred earlier, within the groups that went on to conglomerate into the cluster (e.g., \\citealt{1985MNRAS.215..517B,1998ApJ...496...39Z,2001MNRAS.326..637K,2001ApJ...563..736C,2004PASJ...56...29F,2006MNRAS.370.1223B}, hereafter B+06a; \\citealt{2007ApJ...671.1503M,2008ApJ...672L.103K,2008ApJ...680.1009W}; \\citealt{2008ApJ...688L...5K,2008arXiv0812.2021P}). In this picture, the characteristic properties of cluster galaxies (dominated by early-types) would primarily arise by ``pre-processing'' in groups. Thus by studying galaxy groups, one can learn not only about the ecosystems of most of the galaxies in the universe (\\citealt{2004MNRAS.348..866E}), but also about the pre-history of the rarer but interesting galaxy clusters.  The nearby group surrounding the elliptical galaxy NGC~1407, sometimes called the ``Eridanus A'' group (\\citealt{1989AJ.....98.1531W}; \\citealt{1990AJ....100....1F}; \\citealt{2006MNRAS.369.1351B}, hereafter B+06b), is a potentially informative transitional system. It has the X-ray and optical luminosities of a group (with 1--3 galaxies of at least $L^*$ luminosity, and highly dominated by early-types), but possibly the mass of a cluster. An early study of its constituent galaxy redshifts suggested a group mass-to-light ratio of $\\Upsilon \\sim$~2500~$\\Upsilon_{\\odot}$  ($B$-band; \\citealt{1993ApJ...403...37G}), and while subsequent studies have brought this estimate down to $\\sim$~300--1100 (partially through longer distance estimates; \\citealt{1994A&A...283..722Q,2005ApJ...618..214T}; B+06b; \\citealt{2006MNRAS.369.1375T,2006MNRAS.372.1856F}; \\citealt{2007ApJ...656..805Z}, hereafter Z+07), these are still high values more reminiscent of clusters than of groups.  The NGC~1407 group may provide insight into the origins of the properties of galaxy cluster cores and of brightest cluster galaxies. The group is part of a larger structure or ``supergroup'' that should eventually collapse to form a cluster with NGC~1407---and possibly its associated group---at its core (B+06b; \\citealt{2006MNRAS.369.1375T}). Major mergers are prone to scramble the memories of their progenitors' prior properties, but the future assembly of the NGC~1407 cluster appears to involve low mass ratios of $\\sim$1:5 to 1:10. Thus it is likely that the NGC~1407 group will remain relatively unscathed at the core of the resulting cluster, and we can make a fair comparison of its current dynamical properties with those of existing cluster cores (such as M87 in Virgo and NGC~1399 in Fornax) to see if these systems are part of the same evolutionary sequence.  NGC~1407, like many bright elliptical galaxies, is endowed with a swarm of GCs visible as bright sources extending far out into its halo, beyond the region of readily observable galaxy light. GCs are invaluable as tracers of major episodes of star formation in galaxies \\citep{1995AJ....109..960W,2000A&A...354..836L,2006ARA&A..44..193B}. In NGC~1407, \\citet[hereafter C+07]{2007AJ....134..391C} have used them to study the galaxy's early metal enrichment history, and to search for signs of recent gas-rich mergers. Here our focus is on using GCs as bountiful kinematical tracers, useful for probing both the dynamical properties of their host galaxy's outer parts, and the mass distribution of the surrounding group.  The use of GCs as mass tracers is well established: see \\citet{2006ARA&A..44..193B} and \\citet{2009gcgg.book..433R} for reviews, and \\S\\ref{sec:massprof} of this paper for specific examples. GCs are so far less widely exploited as {\\it orbit tracers}---using the internal dynamics of a globular cluster system (GCS) to infer its formation history and its connections with other galaxy subcomponents. Studies of GCS rotation have implicitly treated the GCs as proxies for the halo field stars or for the DM particles (e.g., \\citealt{1998AJ....116.2237K}), but the connections between these different entities have not yet been established. \\citet[hereafter B+05]{2005MNRAS.363.1211B} presented pioneering work in this area, deriving predictions for GCS, stellar, and DM kinematics from galaxy merger simulations.  We will make some comparisons of these predictions to the data.  The 10-meter Keck-I telescope with the LRIS multi-slit spectrograph \\citep{1995PASP..107..375O} has produced a legacy of the first large samples of extragalactic GC kinematics outside of the Local Group (\\citealt{1997ApJ...486..230C,2000AJ....119..162C,2003ApJ...591..850C}, hereafter C+03). Now we present the first use of DEIMOS on Keck-II \\citep{2003SPIE.4841.1657F} to study extragalactic GCS kinematics, representing a significant advance in this field. Its 80~arcmin$^2$ field-of-view is one of the most expansive available on a large telescope, and its stability, red sensitivity, and medium resolution make it highly efficient for acquiring accurate velocities for GCs (which are intrinsically red objects) even in gray skies, using the Ca~{\\small II} absorption line triplet (CaT). Our results here inaugurate the SAGES Legacy Unifying Globulars and Galaxies Survey of the detailed photometric and spectroscopic properties of $\\sim$~25 representative galaxies and their GCSs.  NGC~1407 is a moderately-rotating E0 galaxy, neither boxy nor disky, with a ``core-like'' central luminosity profile, a weak central AGN, and absolute magnitudes of $M_B = -21.0$ and $M_K=-24.8$ (B+06b). Its GCS displays a typical color bimodality of ``blue'' and ``red'' GCs (\\citealt{1997AJ....113..895P}; \\citealt{2006MNRAS.366.1230F}; \\citealt{2006ApJ...636...90H}, hereafter H+06a; \\citealt{Harris09}). Its overall specific frequency of GCs per unit luminosity was previously found to be a modest $S_N \\sim 4$, but a new deep wide-field study has produced a richer estimate of $S_N \\sim 6$ (L.~Spitler et al., in preparation, hereafter S+09). NGC~1407 hosts bright, extended X-ray emission that extends as far out as 80~kpc into the surrounding group (B+06b; Z+07), although there is nothing about its X-ray properties that marks it as a particularly unusual galaxy (\\citealt{2004MNRAS.350.1511O}, hereafter OP04). We adopt a galaxy stellar effective radius of \\Reff$=57''=$~5.8~kpc (see \\S\\ref{sec:assumpt}), with a distance to the galaxy and group of 20.9~Mpc \\citep{2006MNRAS.366.1230F}---discussing later the effects of the distance uncertainty.  We present the spectroscopic GC observations in \\S\\ref{sec:obs}, and convert these into kinematical parameters in \\S\\ref{sec:kin}. In \\S\\ref{sec:mass}, we analyze the group mass profile using GC dynamics and X-ray constraints, and review independent mass constraints more generally. We place our findings for NGC~1407 into the wider context of galaxy groups in \\S\\ref{sec:context}, and \\S\\ref{sec:conc} summarizes the results. ",
        "conclusions": ""
    },
    "0809/0809.2794_arXiv.txt": {
        "abstract": "We measure the logarithmic scatter in mass at fixed richness for clusters in the maxBCG cluster catalog, an optically selected cluster sample drawn from SDSS imaging data.  Our measurement is achieved by demanding consistency between available weak lensing and X-ray measurements of the maxBCG clusters, and the X-ray luminosity--mass relation inferred from the 400d X-ray cluster survey, a flux limited X-ray cluster survey. We find $\\sigma_{\\ln M|N_{200}}=0.45^{+0.20}_{-0.18}$ ($95\\%$ CL) at $N_{200}\\approx 40$, where $N_{200}$ is the number of red sequence galaxies in a cluster.  As a byproduct of our analysis, we also obtain a constraint on the correlation coefficient between $\\ln \\Lx$ and $\\ln M$ at fixed richness, which is best expressed as a lower limit, $r_{L,M|N} \\geq 0.85\\ (95\\%\\ \\mbox{CL})$.  This is the first observational constraint placed on a correlation coefficient involving two different cluster mass tracers. We use our results to produce a state of the art estimate of the halo mass function at $z=0.23$ --- the median redshift of the maxBCG cluster sample --- and find that it is consistent with the WMAP5 cosmology.   Both the mass function data and its covariance matrix are presented. ",
        "introduction": "The space density of galaxy clusters as a function of cluster mass is a well-known cosmological probe \\citep[see e.g.][]{holderetal01,haimanetal01,rozoetal04,limahu04}, and ranks among the best observational tools for constraining $\\sigma_8$, the normalization of the matter power spectrum in the low redshift universe \\citep[see e.g.][]{frenketal90,henry91, schueckeretal03,gladdersetal07,rozoetal07a}.\\footnote{$\\sigma_8$ is formally defined as the variance of the linear matter density averaged over spheres with radius $R=8\\ h^{-1}\\ \\Mpc$.}  The basic idea is this: in the high mass limit, the cluster mass function falls off exponentially with mass, with the fall-off depending sensitively on the amplitude of the matter density fluctuations. Observing this exponential cutoff can thus place tight constraints on $\\sigma_8$. In practice, however, the same exponential dependence that makes cluster abundances a powerful cosmological probe also renders it susceptible to an important systematic effect, namely uncertainties in the estimated masses of clusters.  Because mass is not a direct observable, cluster masses must be determined using observable mass tracers such as X-ray emission, SZ decrements, weak lensing shear, or cluster richness (a measure of the galaxy content of the cluster).   Of course, such mass estimators are noisy, meaning there can be significant scatter between the observable mass tracer and cluster mass.  Since the mass function declines steeply with mass, up-scattering of low mass systems into high mass bins can result in a significant boost to the number of systems with apparently high mass \\citep{limahu05}.  If this effect is not properly modeled, the value of $\\sigma_8$ derived from such a cluster sample will be overestimated.  One approach for dealing with this difficulty is to employ mass tracers that have minimal scatter, thereby reducing the impact of said scatter on the recovered halo mass function. For instance, \\citet{kravtsovetal06} introduced a new X-ray mass estimator, $Y_X=M_{gas}T_X$, which in their simulations exhibits an intrinsic scatter of only $\\approx 8\\%$, independent of the dynamical state of the cluster. Use of a mass estimator with such low scatter should lead to improved estimates of $\\sigma_8$ from X-ray cluster surveys  \\citep{pierpaolietal01,rb02,schueckeretal03, henry04,staneketal06}.  Such tightly-correlated mass tracers are not always available. In such cases, determination of the scatter in the mass-observable relation is critical to accurately inferring the mass function and thereby determining cosmological parameters. Of course, in practice, it is impossible to determine this scatter to arbitrary accuracy, but since the systematic boost to the mass function is proportional to the square of the scatter \\citep{limahu05} (i.e. the variance), even moderate constraints on the scatter can result in tight $\\sigma_8$ constraints.  In this paper, we use optical and X-ray observations to constrain the scatter in the mass--richness relation for the maxBCG cluster catalog presented in \\citet{koesteretal07a}. Specifically, we use observational constraints on the mean mass--richness relation, and on the mean and scatter of the $\\Lx-$richness relation, to convert independent estimates of the scatter in the $\\Lx-M$ relation into estimates of the scatter in the mass--richness relation.  An interesting byproduct of our analysis is a constraint on the correlation coefficient between mass and X-ray luminosity at fixed richness.  To our knowledge, this is the first time that a correlation coefficient involving multiple cluster mass tracers has been empirically determined.  The layout of the paper is as follows. In section \\ref{sec:notation} we lay out the notation and definitions used throughout the paper.  Section \\ref{sec:data} presents the data sets used in our analysis.  In section \\ref{sec:rough} we present a pedagogical description of our method for constraining the scatter in the richness-mass relation, while section \\ref{sec:formalism} formalizes the argument.  Our results are found in section \\ref{sec:results}, and we compare them to previous work in section \\ref{sec:other_work}. In section \\ref{sec:mf}, we use our result to estimate the halo mass function in the local universe at $z=0.23$, the median redshift of the maxBCG cluster sample, and we demonstrate that our recovered mass function is consistent with the latest cosmological constraints from WMAP \\citep{wmap08}.  A detailed cosmological analysis of our results will be presented in a forthcoming paper (Rozo et al., in preparation).  Our summary and conclusions are presented in section \\ref{sec:conclusions}.  \\subsection{Notation and Conventions} \\label{sec:notation}  We summarize here the notation and conventions employed in this work. Given any three cluster mass tracers (possibly including mass itself) $X,Y,$ and $Z$, we make the standard assumption that the probability distribution $P(X,Y|Z)$ is a bivariate lognormal.  The parameters $A_{X|Z}$, $B_{X|Z}$, and $\\alpha_{X|Z}$ are defined such that \\begin{eqnarray} \\avg{\\ln X|Z} & = & A_{X|Z}+\\alpha_{X|Z}\\ln Z \\\\ \\ln \\avg{X|Z} & = & B_{X|Z} + \\alpha_{X|Z}\\ln Z. \\end{eqnarray} Note the slopes of the mean and logarithmic mean are the same, as appropriate for a log-normal distribution.  The scatter in $\\ln X$ at fixed $Z$ is denoted $\\sigma_{X|Z}$, and the correlation coefficient between $\\ln X$ and $\\ln Y$ at fixed $Z$ is denoted $r_{X,Y|Z}$.  {\\it We emphasize that all quoted scatters are the scatter in the natural logarithm, not in dex.} Note these parameters are simply the elements of the covariance matrix specifying the Gaussian distribution $P(\\ln X,\\ln Y|\\ln Z)$. Under our lognormal assumption for $P(X,Y|Z)$, the parameters $A_{X|Z}$ and $B_{X|Z}$ are related via \\begin{equation} B_{X|Z} = A_{X|Z}+\\frac{1}{2}\\sigma_{X|Z}^2. \\end{equation}  In this work, the quantities of interest are cluster mass $M$, X-ray luminosity $\\Lx$, and cluster richness $N$.  Unless otherwise specified, cluster mass is defined as $M_{500c}$, the mass contained within an overdensity of 500 relative to critical.  $\\Lx$ is the total luminosity in the rest-frame $0.5-2.0\\ \\keV$ band, and $N$ is the maxBCG richness measure $N_{200}$, the number of red sequence galaxies with luminosity above $0.4L_*$ within an aperture such that the mean density within said radius is, on average, $200\\Omega_m^{-1}$ times the mean galaxy density assuming $\\Omega_m=0.3$. Likewise, unless otherwise stated all parameters governing the relations between $M$, $\\Lx$, and $N$ assume that $M$ is measured in units of $10^{14}\\ \\msun$, $\\Lx$ is measured in units of $10^{43}\\ \\mathrm{ergs}/\\mathrm{s}$, and $N$ is measured in ``units'' of $40$ galaxies. For instance, including units explicitly, the mean relation between cluster mass and richness reads \\begin{equation} \\frac{\\avg{M|N}}{10^{14}\\ \\msun} = \\exp(B_{M|N}) \\left (\\frac{N}{40}\\right)^{\\alpha_{M|N}}. \\end{equation} A Hubble constant parameter $h=0.71$ is assumed through out.\\footnote{For other values of $h$, our weak lensing masses scale as $M \\propto h^{-1}$ and the X-ray luminosities as $\\Lx \\propto h^{-2}$.}  In addition, the weak lensing data presented in this analysis assumed a flat $\\Lambda$CDM cosmology with $\\Omega_m=0.27$. The recovered mass function has the standard hubble parameter degeneracy. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that by combining the information in the maxBCG richness function, the mean richness-mass relation, the mean and scatter of the $\\Lx-$richness relation, and the mean and scatter of the $\\Lx-M$ relation, we can constrain both the scatter in mass at fixed richness for maxBCG clusters, as well as the correlation coefficient between mass and $\\Lx$ at fixed richness.  We find \\begin{eqnarray} \\sigma_{M|N} & = & 0.45^{+0.20}_{-0.18}\\ (95\\%\\ \\mbox{CL}) \\\\ r_{L,M|N} & \\geq & 0.85\\ (95\\%\\ \\mbox{CL}). \\end{eqnarray} These constraints are dominated by uncertainties in the $\\Lx-M$ relation, and can be significantly tightened if our understanding of the $\\Lx-M$ relation improves.  We also found our results are consistent with those presented in \\citet{beckeretal07} once miscentering of maxBCG clusters is taken into account.  Our lower limit on the correlation between $M$ and $\\Lx$ at fixed richness constitutes the first observational constraint on a correlation coefficient involving two different halo mass tracers.  Note that the large correlation between $\\Lx$ and $M$ implies that $\\Lx$ - even without core exclusion - is a significantly better mass tracer than the maxBCG richness estimator (i.e. at fixed richness, over-luminous cluster are nearly always more massive).  This is an important result, which we use in a concurrent paper to help us define new richness estimators that are better correlated with cluster mass~\\citep{rozoetal08b}.  Using our results, and assuming $\\Omega_m=0.27$ and $h=0.71$, we have estimated the halo mass function at $z=0.23$, corresponding to the median redshift of the cluster sample.  We find that our recovered mass function is in good agreement with the mass function predicted by \\citet{tinkeretal08} for the WMAP5 cosmology \\citep{wmap08}. A detailed cosmological analysis will be presented in a forthcoming paper (Rozo et al, in preparation).  Our work sheds new light on the interrelationship of bulk properties of massive halos. We have used weak lensing, X-ray luminosities, and optical richness estimates to constrain the scatter in the richness-mass relation, which can lead to improved cosmological constraints.  In principle, one could also turn this question around, and, assuming cosmology, we could constrain the scatter in the richness-mass relation, which would then allow us to place constraints on the amplitude, slope, and scatter of the $\\Lx-M$ relation.   Such an analysis would be interesting in that, by doing so, one could compare the predicted amplitude of the $\\Lx-M$ relation to that derived from hydrostatic mass estimates, thereby directly probing the amount of non-thermal pressure support in galaxy clusters. Note that even though this question can also be directly addressed by comparing weak lensing and X-ray mass estimates of individual clusters, the analysis suggested here would benefit from having small uncertainties, whereas projection effects result in rather noisy weak lensing mass estimates for individual systems."
    },
    "0809/0809.0791_arXiv.txt": {
        "abstract": "Long-term precise timing of Galactic millisecond pulsars holds great promise for measuring the long-period (months-to-years) astrophysical gravitational waves. Several gravitational-wave observational programs, called Pulsar Timing Arrays (PTA), are being pursued around the world.  Here we develop a Bayesian algorithm for measuring the stochastic gravitational-wave background (GWB) from the PTA data. Our algorithm has several strengths: (1) It analyses the data without any loss of information, (2) It trivially removes systematic errors of known functional form, including quadratic pulsar spin-down, annual modulations and jumps due to a change of equipment, (3) It measures simultaneously both the amplitude and the slope of the GWB spectrum, (4) It can deal with unevenly sampled data and coloured pulsar noise spectra. We sample the likelihood function using Markov Chain Monte Carlo (MCMC) simulations. We extensively test our approach on mock PTA datasets, and find that the algorithm has significant benefits over currently proposed counterparts. We show the importance of characterising all red noise components in pulsar timing noise by demonstrating that the presence of a red component would significantly hinder a detection of the  GWB  Lastly, we explore the dependence of the signal-to-noise ratio on the duration of the experiment, number of monitored pulsars, and the magnitude of the pulsar timing noise. These parameter studies will help formulate observing strategies for the PTA experiments. ",
        "introduction": "At the time of this writing several large projects are being pursued in order to directly detect astrophysical gravitational waves. This paper concerns a program to detect gravitational waves using pulsars as nearly-perfect Einstein clocks.  The practical idea is to time a set of millisecond pulsars (called the ``Pulsar Timing Array'', or PTA) over a number of years \\citep{Foster}. Some of the millisecond pulsars create pulse trains of exceptional regularity. By perturbing the space-time between a pulsar and the Earth, the gravitational waves (GWs) will cause extra deviations from the periodicity in the pulse arrival times \\citep{Estabrook, Sazhin, Detweiler}. Thus from the measurements of these deviations (called ``timing-residuals'', or TR), one may measure the gravitational waves. Currently, several PTA project are operating around the globe. Firstly, at the Arecibo Radio Telescope in North-America several millisecond pulsars have been timed for a number of years. These observations have already been used to place interesting upper limits on the intensity of gravitational waves which are passing through the Galaxy \\citep{Kaspi, Lommen}. Together with the Green Bank Telescope, the Arecibo Radio Telescope will be used as an instrument of NANOGrav, the North American PTA. Secondly, the European PTA is being set up as an international collaboration between Great Britain, France, Netherlands, Germany, and Italy, and will use 5 European radio telescopes to monitor about 20 millisecond pulsars \\citep{Stappers}. Finally, the Parkes PTA in Australia has been using the Parkes multi-beam radio-telescope to monitor 20 millisecond pulsars \\citep{Manchester}. Some of the Parkes and Arecibo data have also been used to place the most stringent limits on the GWB to date \\citep{Jenet-2006}.  One of the main astrophysical targets of the PTAs is the stochastic background of the gravitational waves (GWB). This GWB is thought to be generated by a large number of black-hole binaries which are thought to be located at the centres of galaxies \\citep{Begelman, Phinney, Jaffe, Wyithe, Sesana}, by relic gravitational waves \\citep{Grishchuk}, or, more speculatively, by cusps in the cosmic-string loops \\citep{Damour}. This paper develops an algorithm for the optimal PTA measurement of such a GWB.  The main difficulty of such a measurement is that not only Gravitational Waves  create the pulsar timing-residuals. Irregularities of the pulsar-beam rotation (called the ``timing noise''), the receiver noise, the imprecision of local clocks, the polarisation calibration of the telescope \\citep{Britton}, and the variation in the refractive index of the interstellar medium all contribute significantly to the timing-residuals, making it a challenge to separate these noise sources from the gravitational-wave signal. However, the GWB is expected to induce correlations between the timing-residuals of different pulsars. These correlations are of a specific functional form [given by Eq.~(\\ref{eq:alphaab}) below], which is different from those introduced by other noise sources \\citep{Hellings}.  \\citet[hereafter J05]{Jenet-2005} have invented a clever algorithm which uses the uniqueness of the GWB-induced correlations to separate the GWB from other noise sources, and thus to measure the magnitude of the GWB.  Their idea was to measure the timing-residual correlations for all pairs of the PTA pulsars, and check how these correlations depend on the sky-angles between the pulsar pairs. J05 have derived a statistic which is sensitive to the functional form of the GWB-induced correlation; by measuring the value of this statistic one can infer the strength of the GWB.  While J05 algorithm appears robust, we believe that in its current form it does have some drawbacks, in particular: \\newline (1) The statistic used by J05 is non-linear and non-quadratic in the pulsar-timing-residuals, which makes its statistical properties non-transparent.\\newline (2) The pulsar pairs with the high and low intrinsic timing noise make equal contributions to the J05 statistic, which is clearly not optimal.\\newline (3) The J05 statistic assumes that the timing-residuals of all the PTA pulsars are measured during each observing run, which is generally not the case.\\newline (4) The J05 signal-to-noise analysis relies on the prior knowledge of the intrinsic timing noise, and there is no clean way to separate this timing noise from the GWB.\\newline (5) The prior spectral information on GWB is used for whitening the signal; however, there is no proof that this is an optimal procedure. The spectral slope of the GWB is not measured.  In this paper we develop an algorithm which addresses all of the problems outlined above.  Our method is based on essentially the same idea as that of J05: we use the unique character of the GWB-induced correlations to measure the intensity of the GWB. The algorithm we develop below is Bayesian, and by construction uses optimally all of the available information. Moreover, it deals correctly and efficiently with all systematic contributions to the timing-residuals which have a known functional form, i.e.~the quadratic pulsar spin-downs, annual variations, one-time discontinuities (jumps) due to equipment change, etc. Many parameters of the timing model (the model popular pulsar timing packages use to generate TRs from pulsar arrival times) fall in this category.   The plan of the paper is as follows. In the next section we review the theory of the GWB-generated timing residuals and introduce our  model for other contributions to the timing residuals.  In Sec.~\\ref{sec:bayes} we explain the principle of Bayesian analysis for GWB-measurement with a  PTA, and we evaluate the Bayesian likelihood function. There we also show how to analytically marginalise over the contributions of known functional form but unknown amplitude (i.e., annual variations, quadratic residuals due to pulsar spin-down, etc.). The details of this calculation are laid out in Appendix A. Section~\\ref{sec:montecarlo} discusses the numerical integration technique which we use in our likelihood analysis: the Markov Chain Monte Carlo (MCMC). In Sec.~\\ref{sec:results} we show the analyses of mock PTA datasets.  For each mock dataset, we construct the probability distribution for the intensity of the GWB, and demonstrate its consistency with the input mock data parameters. We study the sensitivity of our algorithm for different PTA configurations, and investigate the dependence of the signal-to-noise ratio on the duration of the experiment, on redness and magnitude of the pulsar timing noise, and on the number of clocked pulsars. In Sec.~\\ref{sec:conclusions} we summarise our results. ",
        "conclusions": "\\label{sec:conclusions} In this paper we have introduced a practical Bayesian algorithm for measuring the GWB using Pulsar Timing Arrays. Several attractive features of the algorithm should make it useful to the PTA community:\\newline (1) the ability to simultaneously measure the  amplitude and slope of GWB,\\newline (2) its ability to deal with unevenly sampled datasets, and \\newline (3) its ability to treat systematic contributions of known functional form. From the theoretical point of view, the algorithm is guaranteed to extract information optimally, provided that our parametrization of the timing noise is correct.  Test runs of our algorithm have shown that the experiments signal-to-noise (S/N) ratio strongly decreases with the redness of the pulsar timing noise, and strongly increases with the duration of the PTA experiment. We have also charted the S/N dependence on the number of well-clocked pulsars and the level of their timing noise. These charts should be helpful in the design of the optimal strategy for future PTA observations."
    },
    "0809/0809.2198_arXiv.txt": {
        "abstract": "At low metallicity, B-type stars show lower loss of mass and, therefore, angular momentum so that it is expected that there are more Be stars in the Magellanic Clouds than in the Milky Way. However, till now, searches for Be stars were only performed in a very small number of open clusters in the Magellanic Clouds. Using the ESO/WFI in its slitless spectroscopic mode, we performed a Halpha survey of the Large and Small Magellanic Cloud. Eight million low-resolution spectra centered on Halpha were obtained. For their automatic analysis, we developed the ALBUM code. Here, we present the observations, the method to exploit the data and first results for 84 open clusters in the SMC. In particular, cross-correlating our catalogs with OGLE positional and photometric data, we classified more than 4000 stars and were able to find the B and Be stars in them.  We show the evolution of the rates of Be stars as functions of area density, metallicity, spectral type, and age. ",
        "introduction": "Emission line stars (ELS) range from young to evolved stars (TTauri, Herbig Ae/Be, WR, Planetary Nebulae, etc), from hot to cool stars (Classical Be star, Oe, Supergiant star, Mira Ceti, Flare stars, etc). Among the ELS, here, we focus on the classical Be stars. They are non-supergiant B type stars, which have displayed at least once emission lines in their spectra, mainly in the Balmer series of the hydrogen. The emission lines come from a circumstellar decretion disk formed by episodic matter ejection of the central star. It appears that the Be phenomenon is related to fast rotation and probably additional properties such as non-radial pulsation or magnetic fields.  For a comprehensive review of Be stars in the Milky Way we refer the reader to \\cite[Porter \\& Rivinius (2003)]{}. It seems also that the low metallicity plays a role (\\cite[Kudritzki et al. 1987]{}): at low metallicity, typical of the Small Magellanic Cloud (SMC), the stellar radiatively driven winds are less efficient than at high metallicity (typical of the Milky Way [MW]), thus the mass loss is lower and the stars keep more angular momentum. As a consequence, B-type stars rotate faster in the SMC/LMC than in the MW (\\cite[Martayan et al. 2007]{}).  It is then expected that the metallicity has also an effect on the Be-phenomenon itself as reported by \\cite[Maeder et al. (1999)]{} or \\cite[Wisniewski \\& Bjorkman (2006)]{}, while the evolutionary phase could also play a role (\\cite[Fabregat \\& Torrej\\'on 2000]{}). In most cases, the study of open clusters was done by using photometric observations, in the MW. The work by \\cite[McSwain \\& Gies (2005)]{} is a typical example. To test these issues and improve the comparisons between SMC and MW, spectroscopic observations of stars in open clusters have to be done. In the SMC, a survey of emission line objects was performed by \\cite[Meyssonnier \\& Azzopardi (1993)]{} using photographic plates. However, they were not able to study the stars in open clusters but mainly in the field. This paper deals with our slitless spectroscopic survey of ELS and Oe/Be/Ae stars in SMC open clusters, while in the whole SMC 3 million spectra were obtained.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=3.1in, angle=-90]{marta1fig1.ps} \\includegraphics[width=3.1in, angle=-90]{marta1fig2.ps} \\caption{Absolute V magnitude vs. dereddened colour (B-V)-top or (V-I)-bottom for SMC stars of our sample. Blue crosses correspond to definite ELS, blue diamonds correspond to candidate ELS, and red + to absorption stars.} \\label{figs1} \\end{center} \\end{figure} ",
        "conclusions": "Preliminary studies have shown a trend of the increase of the fraction of Be stars with decreasing metallicity. However, up to now, the studies were only performed on a very limited sample of open clusters (less than 10 in the SMC) and  with photometric data. Using the ESO/WFI in its slitless spectroscopic mode, we observed 84 open clusters in the SMC. Thanks to different codes and OGLE data, we were able to find and classify the emission line stars and absorption stars ($\\sim$4300 stars). The ratios of Be stars to B stars in the SMC and MW were studied. The comparison allows to quantify the increase of the number of Be stars with decreasing metallicity. Be stars in the SMC are 2 to 4 times more abundant than in the MW depending on the spectral types. It seems also that early Be stars follow the same distribution in the SMC and MW with a maximum at the spectral type B2. About the stellar phases at which Be stars appear, from our preliminary results, it seems that certain Be stars could be born as Be stars, while others could appear during the main sequence depending also probably on the metallicity and spectral types of stars."
    },
    "0809/0809.0362_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{s0}  Lately, relativistic astrophysics has seen an increasing interest in papers where solutions with traversable wormholes (WHs) are discussed. This interest is caused, among other things, by projecting and constructing high-precision radiointerferometers which will make it possible to discriminate WHs from black holes.  In this paper we consider a static and spherically symmetric WH solution which only slightly differs from that for the extremal charged Reissner-Nordstr\\\"{o}m black hole (see~\\cite{3} or \\cite{7}). As is shown in \\cite{2}, if the equation-of-state parameters of matter change their values the WH solution can smoothly turn into the extremal Reissner-Nordstr\\\"{o}m solution with a horizon and the WH stops being traversable.  WH sustains its exotic properties thanks to phantom matter violating the null energy condition\\footnote{This condition means satisfying the following inequality (NEC): ${T_{ik}\\zeta^i\\zeta^k>0}$, $T_{ik}$ being the energy-momentum tensor of the matter and $\\zeta^i$ -- the null 4-vector of photon. Physically, the exoticness of the phantom matter violating the NEC leads to the possibility of choosing a reference frame where observed energy density is negative.} and surrounding a WH's throat -- see, for example,~\\cite{1} or \\cite{2}.  What can be practically interesting is the WH that only slightly differs from the extremally (whether electrically or magnetically) charged Reissner-Nordstr\\\"{o}m black hole with the charge $q$. In this particular case the amount of the phantom matter can be taken arbitrarily small. ",
        "conclusions": "\\label{s-zakl}  In spite the result obtained stating that the light distribution in the WH throat is homogeneous for each WH model, it is worth noting that in the real universe the number of visible stars is finite, though big. This implies that if angular resolution of the observer's instrument in our Universe is high enough they will be able to discover the changing star density in the throat ${\\bf J(h)}$. The left panel of Fig.~\\ref{R1} shows this plot for ${\\delta =0.001}$. Sharp minima on the plot correspond to zeros of the sine in expression (\\ref{2-1}). This is because at sufficiently large impact parameters the light rays are deflected by large angles (${\\theta>\\pi}$) so that in the vicinities of the points ${\\theta =\\pi n}$ abrupt declines in distribution arise. But near these declines the observed stellar brightness tends to infinity (lensing), which ultimately provides the (on average) uniform light flow over the WH throat (see the right panel of Fig.~\\ref{R1}).  Positions of the declines depend on the value of $\\delta$. Hence, registering them makes it possible to determine the equation-of-state parameters of the WH matter and features of the WH model (which is highly analogous to processing the light spectra).  Thus, in this paper we have proposed a technique of calculating the deflection of light passing through wormholes as well as methods of observing distinctive features of specific WH models."
    },
    "0809/0809.0012_arXiv.txt": {
        "abstract": "The transfer of polarized radiation in magnetized and non-magnetized relativistic plasmas is an area of research with numerous flaws and gaps. The present paper is aimed at filling some gaps and eliminating the flaws. Starting from a Trubnikov's linear response tensor for a vacuum wave with ${\\bf k}=\\omega/c$ in thermal plasma, the analytic expression for the dielectric tensor is found in the limit of high frequencies. The Faraday rotation and Faraday conversion measures are computed in their first orders in the ratio of the cyclotron frequency $\\Omega_0$ to the observed frequency $\\omega$. The computed temperature dependencies of propagation effects bridge the known non-relativistic and ultra-relativistic limiting formulas. The fitting expressions are found for high temperatures, where the higher orders in $\\Omega_0/\\omega$ cannot be neglected. The plasma eigenmodes are found to become linearly polarized at much larger temperatures than thought before. The results are applied to the diagnostics of the hot ISM, hot accretion flows, and jets. ",
        "introduction": "We learn much of our information about astrophysical objects by observing the light they emit. Observations of the polarization properties of light can tell us the geometry of the emitter, strength of the magnetic field, density of plasma, and temperature. The proper and correct theory of optical activity is essential for making accurate predictions. While the low-temperature propagation characteristics of plasma are well-established \\citep{landau10}, the theory of relativistic effects has not been fully studied. In this paper I discuss the propagation effects through a homogeneous magnetized relativistic plasma. A non-magnetized case emerges as a limit of the magnetized case. The discussion is divided into three separate topics.  Two linear plasma propagation effects are Faraday rotation and Faraday conversion \\citep{mueller}. Traditionally, these effects are considered in their lowest orders in the ratio $\\beta$ of the cyclotron frequency $\\Omega_0$ to the circular frequency of light $\\omega,$ id est in a high-frequency approximation. The distribution of particles is taken to be thermal \\begin{equation}\\label{distrib} dN=\\frac{n \\exp (-\\gamma/T)}{4\\pi m^2 T^2 K_2(T^{-1})}d^3p \\end{equation} with the dimensionless temperature $T$ in the units of particle rest mass temperature $m c^2/k_B.$ The Faraday rotation measure $RM$ and conversion measure are known in a non-relativistic $T\\ll 1$ and an ultra-relativistic $T\\gg 1$ limits \\citep{melrosec}. I derive a surprisingly simple analytic expression for arbitrary temperature $T.$  The smallness of $\\beta=\\Omega_0/\\omega,$  $\\beta\\ll1$ in the real systems led some authors \\citep{melrosea} to conclude that the high-frequency approximation will always work. However, there is a clear indication that it breaks down at high temperatures $T\\gg 1.$ It was claimed that the eigenmodes of plasma are linearly polarized for high temperatures $T\\gg 1$ \\citep{melrosec}, because the second order term $\\sim\\beta^2$ becomes larger than the first order term $\\sim \\beta$ due to the $T$ dependence. The arbitrarily large $T$-factor may stand in front of higher order expansion terms in $\\beta$ of the relevant expressions. I find the generalized rotation measure as a function of $\\beta$ and $T$ without expanding in $\\beta$ and compare the results with the known high-frequency expressions. The high-$T$ behavior of the plasma response is indeed significantly different.  Plasma physics involves complicated calculations. This led to a number of errors in the literature \\citep{melrosec}, some of which have still not been fixed. In the article I check all the limiting cases numerically and analytically and expound all the steps of derivations. Thus I correct the relevant errors and misinterpretations made by previous authors, hopefully not making new mistakes. The analytical and numerical results are obtained in Mathematica 6 system. It has an enormous potential in these problems \\citep{mathem_blog}.  The paper is organized as follows. The formalism of plasma response and calculations are described in \\S \\ref{method}. Several applications to observations can be found in \\S \\ref{applications}. I conclude in \\S \\ref{discussion} with a short summary and future prospects. ",
        "conclusions": "\\label{discussion} This paper presents several new calculations and amends the previous calculations of propagation effects in uniform magnetized plasma with thermal particle distribution equation (\\ref{distrib}). The expression (\\ref{main}) for the correct response tensor is given in a high-frequency approximation. The exact temperature dependence (\\ref{diag}) and (\\ref{nondiag}) is found in first orders in $\\Omega_0/\\omega$ in addition to the known highly-relativistic and non-relativistic results. The higher order terms may be important for relativistic plasmas in jets and hot accretion flows. The fitting expressions (\\ref{gaunt_diag}) and (\\ref{gaunt_nondiag}) are found for the dielectric tensor components (\\ref{diagdifffit}) and (\\ref{nondiagfit}) at relatively high temperatures.  The results of numerical computations are given only when the corresponding analytical formulas are found. One can always compute the needed coefficients numerically for every particular frequency $\\omega$, plasma frequency $\\omega_p,$ cyclotron frequency $\\Omega_0,$ and distribution of electrons. However, the analytic formulas offer a simpler and faster way of dealing with the radiative transfer for a non-specialist. The eigenmodes were not considered in much detail, since radiative transfer problems do not require a knowledge of eigenmodes. However the knowledge of eigenmodes is needed to compute the self-consistent response tensor (see \\S~\\ref{exact}).  The response tensor in the form (\\ref{main}) can be expanded in $\\Omega_0/\\omega$ and $\\omega_p/\\omega.$ This expansion is of mathematical interest and will be presented in a subsequent paper as well as the expressions for a power-law electron distribution. Propagation through non-magnetized plasmas will also be considered separately."
    },
    "0809/0809.2826_arXiv.txt": {
        "abstract": "{A brief overview of the $r$-process is given with an emphasis on the observational implications for this process. The conditions required for the major production of the heavy $r$-process elements ($r$-elements) with mass numbers $A>130$ are discussed based on a generic astrophysical model where matter adiabatically expands from a hot and dense initial state. Nucleosynthesis in the neutrino-driven winds from nascent neutron stars is discussed as a specific example. Such winds readily produce the elements from Sr to Ag with $A\\sim 88$--110 through charged-particle reactions in the $\\alpha$-process but appear incapable of making the heavy $r$-elements. Observations of elemental abundances in metal-poor stars have provided many valuable insights into the $r$-process. They have demonstrated that the production of the heavy $r$-elements must be associated with massive stars evolving on short timescales, provided evidence strongly favoring core-collapse supernovae over neutron star mergers as the major source for these elements, and shown that this source cannot produce any significant amount of the low-$A$ elements from Na to Ge including Fe. A self-consistent astrophysical model of the $r$-process remains to be developed, and it appears well worthwhile to carry out more comprehensive studies on the evolution and explosion of the massive stars of $\\sim 8$--$11\\,M_\\odot$ that undergo O-Ne-Mg core collapse and produce a very insignificant amount of the low-$A$ elements.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1, 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "This paper discusses some of the progress that has been made in the understanding of the $r$-process since the publication of \\cite{cowa91}, a major review on this process, in 1991. More detailed discussion can be found in \\cite{qian03,qw07}. Readers are referred to \\cite{arno07} for a review of the $r$-process that covers a wider range of topics and has different emphases.  For convenience of discussion, an $r$-process event can be separated into two phases. During the first phase, relatively heavy nuclei are produced mostly via charged-particle reactions (CPRs). These nuclei become the seed nuclei to capture the neutrons during the second phase. The crucial result of the first phase is a large (number) abundance ratio of the neutrons to the seed nuclei. The second phase involves neutron-rich nuclei far from stability participating in a number of types of nuclear reactions, which include neutron capture, photo-disintegration, $\\beta$-decay, and possibly fission. The properties of such nuclei play an essential role in determining the exact abundance pattern produced by the $r$-process event. This important aspect of the $r$-process will not be discussed here. For the presentation below, the abundance pattern resulting from the second phase is characterized by the average mass number $\\langle A_r\\rangle$ of the corresponding $r$-process nuclei. Assuming that all the neutrons are captured, the results from the first and second phases are related by mass conservation as \\begin{equation} \\langle A_s\\rangle Y_s+Y_n=\\langle A_r\\rangle Y_r, \\label{eq-ns} \\end{equation} where $\\langle A_s\\rangle$ is the average mass number of the seed nuclei, $Y_s$ is their initial total (number) abundance, $Y_n$ is the initial neutron abundance, and $Y_r$ is the final total abundance of the $r$-process nuclei produced. In the absence of fission, the total number of heavy nuclei is conserved. For this case $Y_r=Y_s$ and Equation~(\\ref{eq-ns}) can be rewritten as \\begin{equation} \\langle A_s\\rangle + \\frac{Y_n}{Y_s}=\\langle A_r\\rangle, \\label{eq-ns2} \\end{equation} which relates the average $r$-process nuclei to the average seed nuclei by the neutron-to-seed ratio $Y_n/Y_s$. Extensive fission cycling occurs for $Y_n/Y_s\\gg\\langle A_s\\rangle$, in which case Equation~(\\ref{eq-ns}) approximately becomes \\begin{equation} Y_n\\approx\\langle A_r\\rangle Y_r. \\end{equation} For this extreme case, a regular abundance pattern is expected from the steady-state nature of the nuclear flow between the fission products and the fissioning nuclei, and the final outcome is governed by the initial neutron abundance but not the exact abundance or form of the seed nuclei.  It can be seen from the above discussion that the neutron-to-seed ratio $Y_n/Y_s$ is crucial to the success of an $r$-process event. Before discussing how this ratio is determined in the astrophysical models of these events, some general comments regarding the $r$-process based on observations are in order. Figure~\\ref{fig1} shows the $\\log\\epsilon$ values\\footnote{For an element E with a (number) abundance ratio (E/H) relative to hydrogen, $\\log\\epsilon({\\rm E})\\equiv\\log({\\rm E/H})+12$.} for the elements from Sr to Ag with mass numbers $A\\sim 88$--110 and those from Ba to Au with $A>130$ observed in four metal-poor stars CS~22892--052 \\cite{sned03}, HD~221170 \\cite{ivan06}, CS~31082--001 \\cite{hill02}, and BD~$+17^\\circ3248$ \\cite{cowa02} with ${\\rm [Fe/H]}\\equiv\\log{\\rm (Fe/H)}-{\\rm (Fe/H)}_\\odot=-3.1$, $-2.2$, $-2.9$, and $-2.1$, respectively. The solid curves in Fgiures~\\ref{fig1}a and \\ref{fig1}b represent the so-called solar ``$r$-process'' abundance pattern ($r$-pattern) that is translated to pass through the Eu data for CS~22892--052 and CS~31082--001, respectively. This pattern was derived by subtracting the $s$-process contributions calculated in \\cite{arla99} from the total solar abundances. It can be seen from Figure~\\ref{fig1} that the data for the elements with $A>130$ (Ba to Au with atomic numbers $Z=56$--79) in all four stars follow the translated solar $r$-pattern rather well, especially for the two stars in Figure~\\ref{fig1}a. These elements will be referred to as the heavy $r$-process elements ($r$-elements) hereafter. However, the data for the elements with $A\\sim 88$--110, especially Sr and Ag ($Z=38$ and 47, respectively), exhibit significant to large deviations. As will be discussed in Section~\\ref{sec2}, these elements are routinely produced by the CPRs in the neutrino-driven wind from a newly-formed neutron star, and therefore, are not the true $r$-elements produced by rapid neutron capture. They will be referred to as the CPR elements hereafter. The remarkable agreement between the solar $r$-pattern and the data for the heavy $r$-elements in four metal-poor stars demonstrates that the $r$-process already occurred in the early Galaxy, and therefore, must be associated with massive stars evolving on short timescales. The rather regular abundance pattern exhibited by the heavy $r$-elements also suggests that fission cycling may be involved in their production.  \\begin{figure} \\vskip -0.75cm \\begin{center} \\includegraphics[angle=270,width=0.52\\textwidth]{f1a.eps}\\hspace{-0.61cm}% \\includegraphics[angle=270,width=0.52\\textwidth]{f1b.eps} \\caption{Data on the $\\log\\epsilon$ values for the CPR elements from Sr to Ag ($Z=38$--47, $A\\sim 88$--110) and the heavy $r$-elements from Ba to Au ($Z=56$--79, $A>130$) for (a) CS~22892--052 (squares with error bars; \\cite{sned03}) and HD~221170 (circles with error bars; \\cite{ivan06}) and (b) CS~31082--001 (squares with error bars; \\cite{hill02}) and BD~$+17^\\circ3248$ (circles with error bars; \\cite{cowa02}). The solid curves in (a) and (b) represent the solar ``$r$-pattern'' \\cite{arla99} that is translated to pass through the Eu data for CS~22892--052 and CS~31082--001, respectively. Note that the data on the heavy $r$-elements follow the translated solar $r$-pattern rather well, especially for the two stars in (a). However, the CPR elements, especially Sr and Ag, exhibit significant to large deviations. As discussed in Section~\\protect\\ref{sec2}, these elements are not the true $r$-elements produced by rapid neutron capture.\\label{fig1}} \\end{center} \\end{figure} ",
        "conclusions": "There has been significant progress in the understanding of the $r$-process through both theoretical and observational studies over the last few decades. The conditions required for the major production of the heavy $r$-elements have been derived for a generic astrophysical model where matter adiabatically expands from a hot and dense initial state. Nucleosynthesis in the neutrino-driven winds from nascent neutron stars has been studied extensively. While such winds readily produce the elements from Sr to Ag through charged-particle reactions in the $\\alpha$-process, they appear incapable of making the heavy $r$-elements. The insights provided by observations of elemental abundances in metal-poor stars are especially important. These observations have demonstrated that the production of the heavy $r$-elements must be associated with massive stars evolving on short timescales, provided evidence strongly favoring CCSNe over NSMs as the major source for these elements, and shown that this source cannot produce any significant amount of the low-$A$ elements from Na to Ge including Fe. A self-consistent astrophysical model of the $r$-process remains to be developed, and it appears well worthwhile to carry out more comprehensive studies on the evolution and explosion of the massive stars of $\\sim 8$--$11\\,M_\\odot$ that undergo O-Ne-Mg core collapse and produce a very insignificant amount of the low-$A$ elements."
    },
    "0809/0809.1708_arXiv.txt": {
        "abstract": "When the accretion rate is more than a small fraction of Eddington, the inner regions of accretion disks around black holes are expected to be radiation-dominated.  However, in the $\\alpha$-model, these regions are also expected to be thermally unstable.  In this paper, we report two 3-d radiation MHD simulations of a vertically-stratified shearing box in which the ratio of radiation to gas pressure is $\\sim 10$, and yet no thermal runaway occurs over a timespan $\\simeq 40$ cooling times. Where the time-averaged dissipation rate is greater than the critical dissipation rate that creates hydrostatic equilibrium by diffusive radiation flux, the time-averaged radiation flux is held to the critical value, with the excess dissipated energy transported by radiative advection. Although the stress and total pressure are well-correlated as predicted by the $\\alpha$-model, we show that stress fluctuations precede pressure fluctuations, contrary to the usual supposition that the pressure controls the saturation level of the magnetic energy.  This fact explains the thermal stability.  Using a simple toy-model, we show that independently-generated magnetic fluctuations can drive radiation pressure fluctuations, creating a correlation between the two while maintaining thermal stability. ",
        "introduction": "As first demonstrated by \\citet{sha73}, radiation pressure must dominate gas pressure in the innermost parts of black hole accretion disks radiating at any non-negligible fraction of the Eddington limit.  Although the original argument made use of the assumption that the stress and the pressure are related by a fixed ratio $\\alpha$, the critical radius within which radiation pressure exceeds gas pressure depends so weakly on the value of $\\alpha$ that this conclusion would be virtually unaltered even if the ratio of stress to pressure varied substantially from place to place.  In this sense, the expectation of radiation dominance in the inner rings of brightly-radiating accretion disks rests only on the supposition that stress and pressure are comparable, the dimensional analysis foundation from which the $\\alpha$ model began.  It is therefore crucial that we understand the role that radiation pressure plays in the physics of accretion disks if we are to understand luminous black hole sources.  Yet, for the past thirty years the canonical internal equilibrium of disks in the radiation-dominated limit has been believed to be unstable \\citep{lig74,sha76,pir78}.  In contrast to other physical regimes, the vertical structure of a stationary disk is in some ways very tightly constrained if radiation pressure truly dominates the hydrostatic support of the disk against the tidal gravitational field of the hole.  Neglecting relativistic corrections for simplicity, one finds that the outward radiation flux in a geometrically thin disk as a function of height $z$ above the midplane is \\begin{equation} F(z)={c\\over\\kappa}g(z)={c\\over\\kappa}\\Omega^2z, \\label{eqfz} \\end{equation} where $c$ is the speed of light, $\\kappa$ is the opacity (usually dominated by electron scattering), and the orbital angular velocity $\\Omega$ determines the vertical tidal gravitational acceleration $g$.  Because energy conservation in a time-steady disk demands (if we neglect inner boundary condition effects) that the flux emerging from the disk surface is $(3/8\\pi)\\dot M \\Omega^2$, the half-thickness $H$ of the disk depends only on the accretion rate $\\dot{M}$ and is independent of radius and the nature of the turbulent stress: $H=3\\kappa\\dot{M}/(8\\pi c)$.  Moreover, equation (\\ref{eqfz}) combined with the condition of radiative equilibrium immediately implies that the dissipation rate per unit volume $Q$ is constant as a function of height above the disk midplane: $Q = dF/dz = c\\Omega^2/\\kappa$.  Given that $Q$ must, on average, be given by the turbulent stress $\\tau_{r\\phi}$ times the rate of strain $r|d\\Omega/dr|$, we also have a tight constraint on the vertically-averaged stress: \\begin{equation} \\tau_{r\\phi}={2c\\Omega\\over3\\kappa}, \\end{equation} a conclusion first reached in slightly different form by \\citet{sha76}.  On the other hand, other aspects of radiation-dominated disks are left totally unconstrained.  Because $\\kappa$ (when it is dominated by electron scattering) is independent of density, the vertical density profile $\\rho(z)$ is completely irrelevant to both the hydrostatic and the thermal equilibrium of the disk.  In the absence of additional information about how the dissipation rate depends on density, it is therefore completely undetermined, and even within Shakura-Sunyaev models it must be specified in some {\\it ad hoc} manner.  Often it is assumed that the dissipation rate per unit mass is constant, which leads to a vertically constant density out to the photospheres.  This element of arbitrariness is in addition, of course, to the {\\it ad hoc} ``$\\alpha$-viscosity'' prescription for the stress used in such models.  The freedom to choose arbitrary values of $\\alpha$ in that stress prescription has always been one of its unfortunate aspects, but in the radiation pressure dominated regime we also have freedom in the choice of which pressure with which to scale the stress.  The standard \\citet{sha73} assumption is that the stress $\\tau_{r\\phi}$ should be proportional to the total (gas plus radiation) pressure, and this choice leads to both thermal (e.g. \\citealt{sha76}) and ``viscous'' (more properly, ``inflow\"; \\citealt{lig74}) instabilities where the radiation to gas pressure ratio exceeds 3/2.  The thermal instabilities generally have faster growth rates than the inflow modes, and are the ones of primary interest here.  The origin of the thermal instability is most simply expressed in terms of the different dependences of the heating and cooling rates on the midplane temperature $T$, assuming constant surface mass density $\\Sigma$ \\citep{pir78}.  In the limit of radial wavelength long compared to the disk thickness, radiative diffusion implies that the cooling rate per unit area $F^-\\sim4caT^4/(3\\kappa\\Sigma)\\propto T^4/\\Sigma$, where $a$ is the radiation density constant.  This must balance the heating rate per unit area $F^+\\sim2H\\tau_{r\\phi}r|d\\Omega/dr|$.  Combining these two expressions with the $\\alpha$-prescription $\\tau_{r\\phi}=\\alpha aT^4/3$ and the condition of hydrostatic equilibrium, we find $H=2aT^4/(3\\Sigma\\Omega^2)$ and $F^+\\sim2\\alpha a^2T^8/(3\\Omega\\Sigma)\\propto T^8/\\Sigma$.  Thermal instability follows because positive temperature perturbations lead to greater increases in heating than in cooling.  It is also possible to express these relationships in terms of the total radiation energy content per unit area $U\\sim2aT^4 H$, in which case $F^-\\propto U^{1/2}/\\Sigma^{1/2}$ and $F^+\\propto U$; once again, instability is indicated.  However, it has never been clear whether the assumption that the stress is proportional to the radiation pressure is correct, and the existence of the thermal instability has therefore always been suspect.  Alternative prescriptions that make the stress proportional to either the gas pressure alone or some combination of the gas and radiation pressure have been advocated over the years, and give rise to more stable configurations (e.g. \\citealt{sak81,ste84,bur85,szu90,mer02,mer03}). Moreover, all of these prescriptions refer to vertically-integrated or one-zone vertical structure models of the accretion disk.  The actual vertical distribution of dissipation may also be important.  If much of the local accretion power is actually dissipated in optically thin regions above and below the disk, then the disk itself can become supported by gas pressure and there would then be no thermal instability \\citep{sve94}.  It is now widely believed that the turbulence in black hole accretion disks resembles that seen in simulations of the nonlinear development of the magnetorotational instability (MRI, e.g. \\citealt{bal98}).  Such simulations, if they include radiation physics, could in principle help resolve these uncertainties by calculating the spatial and temporal structure of the turbulent dissipation within the disk, and checking to see whether a long-lived, and therefore presumably stable, equilibrium exists in the radiation pressure dominated regime.  It may also be possible to investigate the scaling of average turbulent stresses with local averages of gas and/or radiation pressure using the data from these simulations.  Well-resolved, global simulations of optically thick accretion disks that include radiation transport have not yet been published, but it is computationally feasible to investigate the thermal physics locally within the disk using a stratified shearing box geometry \\citep{bra95,sto96,mil00}. The necessity of (sheared)-periodic radial boundary conditions zeroes any net accretion rate through the box, precluding any investigation of inflow instabilities, but thermal instabilities can still be studied.  These same boundary conditions also preclude studying any modes with wavelengths longer than twice the box width, except for infinite wavelength modes---the sheared periodic radial boundary conditions place no restriction on modes that are completely independent of radius.  The first attempt at this was made by \\citet{tur04}, who performed a radiation magnetohydrodynamic (MHD) simulation of a stratified shearing box resulting in an average midplane radiation to gas pressure ratio of 14. Despite this large ratio, the simulation produced no obvious thermal instability over eight thermal times.  The robustness of this result is somewhat suspect, however, given that the computational methods employed precluded the photospheres (top and bottom) from being included within the simulation domain.  Moreover, $27\\%$ of the work done on the box disappeared due to numerical losses, and half the mass was lost from the simulation domain during a heating/expansion phase.  More recently, we \\citep{hir06} developed numerical techniques that conserve energy with high accuracy, place both photospheres within the problem volume, and retain nearly all the initial mass on the grid. Solving the total energy equation enables us to capture all grid scale losses of magnetic and kinetic energy and convert them to internal energy in the gas. Furthermore, by simultaneously solving the internal energy equation, we can compute the instantaneous local losses of magnetic and kinetic energy, i.e. the energy dissipation rate.  The only violations of energy conservation come from the small amounts of energy (generally $\\sim 0.1$\\% of the total dissipated energy\\footnote{0.09\\% in ``1112a'' and 0.15\\% in ``1126b''; see section \\ref{sec:initial_condition} for the labels.}) artificially injected by the action of the density floor and similar numerical limiters. Using this code, we have previously studied a case in which gas pressure dominated over radiation pressure \\citep{hir06} and, in \\citet{kro07,bla07}, one in which the gas and radiation pressures were comparable.  Both of the simulations studied in our previous papers achieved a thermal balance between heating and cooling on time scales comparable to the thermal time.  However, no real steady state was achieved in the simulation with comparable gas and radiation pressure.  Instead, there were long term ($\\sim7$ thermal times) up and down variations in total energy content by factors of three to four.  The two hottest epochs in this simulation marginally violated the thermal instability criterion of \\citet{sha76}, but no thermal runaway occurred.  In this paper we present results of two simulations that are in the fully radiation pressure dominated regime, with box-integrated radiation to gas pressure ratios $\\sim10$.  The two simulations shared the same parameters, boundary conditions, and (almost) the same initial conditions, differing only in the small noise fluctuations imposed on the initial state to seed the MRI.  Despite being in gross violation of the standard thermal instability criterion, there is no evidence for such an instability in either of their time evolutions, which in both cases extended over $\\sim 40$ thermal times.  This paper is organized as follows.  In section 2 we briefly summarize how the code works and discuss the parameters of the new simulations. Quantitative results about the thermodynamics, internal vertical structures, and variability of the simulations are presented in section 3.  In section 4 we discuss a simple model that might explain the thermal stability we observe, and we summarize our findings in section 5. ",
        "conclusions": "In this paper we have presented the results of two simulations that each followed the evolution of a vertically-stratified shearing box with the same surface density and orbital frequency for $\\sim 600$~orbits, or $\\sim 40$~thermal times.   The initial conditions of the two simulations differed only in being given different realizations of the small amplitude noise that was imposed on their otherwise smooth initial state.  Because the magneto-rotational instability generates MHD turbulence, these systems are chaotic; differing small amplitude noise therefore leads to order unity contrasts in their subsequent evolution.  The two simulations should thus be viewed as two different realizations of the many evolutions possible for these parameters.  Their qualitative aspects are, nonetheless, similar.  For example, just as we have found in previous simulations studying shearing boxes with $p_r \\simeq 0.2p_g$ \\citep{hir06} and $p_r \\sim p_g$ \\citep{kro07}, most of the disk mass is found near the midplane, where the magnetic field energy is $\\sim 10^{-2}$ times the combined gas and radiation pressure (see also \\citealt{tur04}). Likewise, in all cases the upper layers of the disk are magnetically-dominated, and the photosphere lies within this region.  Unlike the earlier simulations, in these two the box-averaged radiation pressure is $\\sim 10$ times greater than the gas pressure at all times.  Most importantly, in both cases, the energy content of the box undergoes order unity fluctuations over timescales of many tens of orbits, but these fluctuations have no long-term trend.  This result directly challenges the prediction made in \\citet{sha76} that radiation pressure-dominated disk segments should be thermally unstable. Note, however, that the prediction by \\citet{lig74} of inflow instability in radiation-dominated disks remains to be investigated.  In our simulations, the time-averaged dissipation rate in the disk body is roughly equal to the characteristic value $c\\Omega^{2}/\\kappa_{\\rm es}$, the rate at which diffusive radiation flux maintains hydrostatic balance \\citep{sha76}.  However, in some places ($|z|\\sim H$), the local dissipation rate  exceeds the characteristic value by more than $30\\%$.  If this heat went into radiation flux, hydrostatic balance would be disrupted. We find that the excess is exactly compensated by radiative advection associated with an acoustic breathing mode.  At the same time, mechanical work done on the fluid by the shearing boundaries in excess of the local dissipation rate is transported outward by Poynting flux and radiation pressure work associated with the breathing mode.  Both the dissipated energy carried by radiation advection and (non-dissipated) energy carried by Poynting flux and radiation pressure work are finally deposited and dissipated at $|z|\\sim 2H$, and from there all the way through the rest of the structure radiation diffusion overwhelms all other energy fluxes.  To explain the thermal stability of radiation-dominated disks, we argue that the comparability of stress and pressure inferred from dimensional analysis (which underlies the $\\alpha$-model, and the prediction of instability) is due to dissipation of magnetic turbulence (which produces the stress) providing the heat that is then transformed into radiation pressure.  Consequently, fluctuations in magnetic energy drive fluctuations in pressure, and not (as has been commonly assumed) the other way around.  Our claim is supported by two lines of evidence: First, crosscorrelation analysis of simulation data demonstrates that magnetic fluctuations {\\it lead} radiation energy fluctuations by 5--15 orbits, a little less than a thermal time.  Pressure fluctuations cannot, then, drive magnetic fluctuations. Second, we constructed a simple model realizing this picture, and this model reproduces two other important features observed in the simulation data: The system undergoes thermal fluctuations closely resembling in amplitude and timescale those seen in the simulations.  In addition, although {\\it no} correlation between magnetic energy and radiation pressure is built into the model, thermal balance automatically creates one after the fact.  Thus, correlations between stress and pressure are due to the dissipation of magnetic energy supplying thermal energy, not to the pressure defining a characteristic scale for the stress.  A logical consequence of this point of view, in which stress determines pressure, rather than the other way around, is that the fundamental independent variables are surface density and orbital frequency.  That the orbital frequency is independent of disk parameters is obvious, so long as the disk mass is small compared to the central mass.  So long as the inflow timescale is long compared to the thermal timescale, the surface density must likewise be regarded as an independent parameter with respect to thermal and dynamical fluctuations.   These two independent parameters, through the intertwined and nonlinear processes of MHD instability, tapping the energy reservoir of orbital shear, magnetic dissipation, thermal radiation, and radiative diffusion, with all of these taking place under conditions of vertical (as well as radial) gravitational confinement, combine to determine the magnetic field strength, both its mean saturation level and its fluctuations.  Orbital shear fixes the stress exerted by this field, while the turbulent cascade sets the dissipation rate.  These two are not entirely separate, of course, as the time-averaged accretion rate is equivalent to either one. The pressure follows from the heating rate, as regulated by photon diffusion, and, in turn, closes the loop by determining the disk thickness.  Despite all these complications, at bottom, everything is still determined by only two variables, surface density and orbital frequency.  We are grateful to Shane Davis, Jim Stone, and Neal Turner for very useful discussions and comments. This work was partially supported by NSF Grant AST-0507455 and NASA ATP Grant NNG06GI68G (JHK) and NSF Grants AST-0307657 and AST-0707624 (OMB).  The computations were performed on the SX8 at the Yukawa Institute for Theoretical Physics of Kyoto University and the VPP5000 at the Center for Computational Astrophysics of the National Astronomical Observatory of Japan."
    },
    "0809/0809.4407_arXiv.txt": {
        "abstract": "{} { We present new extraction and identification techniques for supernova (SN) spectra developed within the Supernova Legacy Survey (SNLS) collaboration.} {The new spectral extraction method takes full advantage of photometric information from the Canada-France-Hawa\\\"{\\i} telescope (CFHT) discovery and reference images by tracing the exact position of the supernova and the host signals on the spectrogram. When present, the host spatial profile is measured on deep multi-band reference images and is used to model the host contribution to the full (supernova + host) signal. The supernova is modelled as a Gaussian function of width equal to the seeing. A $\\chi^2$ minimisation provides the flux of each component in each pixel of the 2D spectrogram. For a host-supernova separation greater than $\\gesim$ 1 pixel, the two components are recovered separately and we do not use a spectral template in contrast to more standard analyses. This new procedure permits a clean extraction of the supernova separately from the host in about 70\\% of the 3rd year ESO/VLT spectra of the SNLS. A new supernova identification method is also proposed. It uses the SALT2 spectrophotometric template to combine the photometric and spectral data. A galaxy template is allowed for spectra for which a separate extraction of the supernova and the host was not possible.} {These new techniques have been tested against more standard extraction and identification procedures. They permit a secure type and redshift determination in about 80\\% of cases. The present paper illustrates their performances on a few sample spectra.} {} ",
        "introduction": "Observing Type Ia supernovae (SNe~Ia) for the purpose of constraining cosmological parameters is now a mature activity. Large-scale projects of detection and spectrophotometric follow up of hundreds of \\Iae from ground-based telescopes reach completion today \\citep{Astier06,Wood-Vasey07}, and future \\Iae surveys will bring this number up to several thousand. These current and future surveys yield and will yield new samples of \\Iae observations of unprecedented number and quality. Extracting the most possible out of these exceptional sets of \\Iae data is a challenge for the teams involved in these projects.  At redshifts greater than $z\\gesim 0.2$, the problem of spectroscopically identifying supernovae is more challenging than it is at low redshift. At low redshift, galaxies have a larger angular size, so at the position of the supernova (SN) the spectrograph slit usually contains little contaminating host galaxy light. As redshift increases, apparent galaxy sizes decrease and the galaxy light can be comparable to, or dominate, the supernova light at a given position. For observations of high-redshift supernovae, new techniques are therefore required to efficiently recover the supernova signal (see e.g., Blondin at al. 2005).  Currently, the \\Iae spectroscopic sample of the Supernova Legacy Survey (SNLS) represents more than 500 spectra taken from 8-10m diameter telescopes, both in the Southern and Northern hemispheres (VLT, Gemini N and S, Keck-I and -II) during large observing programmes of hundreds of hours run from 2003 up to the present day.  The spectra obtained with such programmes sample a large fraction of the wavelength space (depending on the instrument used) in the visible. They have restframe phases\\footnote{Throughout this paper, the phase $\\phi$ is the restframe age of the supernova in days with respect to the B-band maximum light.} comprised between -15 and +30 days and sufficient signal-to-noise ratios (S/N) for an unambiguous identification in $\\sim 80$\\% of cases.  Besides their use for cosmological purposes -- that is identification of the supernova type and determination of their redshift for their inclusion in a Hubble diagram, these spectra contain valuable information to understand the nature and the diversity of SNe~Ia and test the adequacy of using them as ``standardisable\", if not standard, candles. As an example, \\citet{Guy07} have implemented a number of SNLS spectra in the training set of the spectrophotometric template SALT2, one of the fitters used to fit the light curves of \\Iae in the 3rd year SNLS Hubble diagram. Moreover, comparison of the spectral properties of these high-redshift ($<z>\\sim 0.5-0.6$) supernovae with their low-$z$ counterparts (from, e.g.,  the spectral time series currently collected by the SN Factory experiment, see \\citealt{Aldering02} and \\citealt{Matheson08}) should help in understanding possible evolutionary effects in \\Iae populations \\citep{Bronder08,Balland06,Balland07,Blondin06,Garavini07,Foley08}. Another interesting particularity of the SNLS spectral sample is that, due to the high average redshift of the SNLS SNe~Ia, the ultraviolet (UV) part of the supernovae spectra down to $\\approx 3200$ $\\AA$ is accessible for a large number of supernovae at various phases (see, e.g., \\citealt{Ellis08} for a study of the UV features of 36 SNLS-Keck spectra). This is of particular interest as it has been suggested that possible evolution with redshift, due to different progenitor metallicities at low and high redshift, might be imprinted in the spectral features of this region of the spectrum \\citep{Hoeflich98, Lentz00}. Optimal extraction of SNLS spectra and their identification are thus not only a crucial step in building the Hubble diagram but are also crucial for their subsequent use in studies of the physics of SNe~Ia.  Standard analyses of spectra are often based on the Horne's extraction method \\citep{Horne86} that uses the spatial profile of the source light to perform an optimal (minimum-variance unbiased) spectral extraction. A drawback of this approach is that only mild geometrical distortions due, for example, to flexure of the instrument are corrected for (however, see \\citet{Marsh89} for an extension of Horne's method suited for highly distorted spectra). Another major problem of this approach is that it does not permit the separate extraction of the supernova from the host signal, except if the two components (SN and host galaxy) are well separated on the spectrogram (and in this case, correlations between the host and the SN spectra are usually unavoidable, even if the extraction window is adjusted to maximise the supernova signal and minimise the host contamination). In most cases, the SN/host separation is performed {\\it a posteriori} by fitting a two-component model (built from a supernova and a galaxy template), the contribution of each component to the total model being evaluated by a $\\chi^2$ minimisation \\citep{Howell02,Howell05,Balland06,Balland07}. As the model is built from a (limited) set of templates, the SN and host spectra obtained by this procedure are strongly model-dependent and provide at best a hint of the host contribution to the full spectral {\\it model}. \\citet{Ellis08} have developed an improved separation technique based on fitting synthetic galaxy spectra to the host photometry measured on reference images. By doing so, they still rely on a spectral model to estimate the host contribution to the full signal. \\citet{Blondin05} use a convolution technique to simultaneously recover a host-free PSF-like SN component and its background galaxy spectrum. This method does not use any spectral template for the host modeling and thus improves over more classical extraction techniques. However, this technique is sensitive to the width of the  spatial resolution {\\em Gaussian} kernel used to recover the background (host) component. The width has to be tuned by the user and highly depends on the spatial extension of the host galaxy. Moreover, the method does not use photometric information that potentially goes along with the spectrum.  To circumvent these drawbacks and improve over standard extraction methods, we have developed a dedicated reduction and extraction pipeline, called PHASE (PHotometry Assisted Spectral Extraction). This new technique uses photometric priors from Canada-France-Hawa\\\"{\\i} Telescope (CFHT) detection images obtained with {\\sc megacam} \\citep{Boulade03} to derive the exact trace of the supernova on the spectrogram. This trace is then used as a guide for extraction. Here, \"true\" photometric information of the host contribution in each pixel, obtained from CFHT Legacy Survey (CFHTLS) deep reference images in various photometric bands, is used, when possible, at the stage of the extraction. In the favourable cases of large enough host/SN separation, we do not use a spectral template to model the host galaxy contribution. The spatial profile of the host, measured on the reference image along columns parallel to the slit, is used instead.  Technical choices sometimes different from the standard ones have been investigated to select the most efficient and robust ones necessary for a clean and final extraction. We have automated the reduction pipeline as much as possible to reduce the operator work load as well as unavoidable human errors. Besides the technical aspects of this new approach, in all the reduction and extraction processing, one of our main concerns is to control and propagate the noise level estimation from raw data up to the final reduced spectrum in the cleanest way possible. In particular, re-sampling of data that introduce correlations among pixels is kept to a minimum and is done as late as possible in order to preserve as long as possible the independence of pixels. This allows us to extract a clean, optimal\\footnote{In fact, our extraction is not strictly optimal because of the PSF modelling as a Gaussian (see below). However, we are close to optimality (within a few percent of minimum variance). In the following, the word 'optimal' stands for 'close to optimal'.} signal-to noise spectrum, and in about 70\\% of the currently treated spectra (3rd year VLT spectra of the SNLS, see \\citealt{Balland08}), this is a  host-free supernova spectrum. We discuss in this paper the improvement in terms of S/N over more standard extractions.  Identification of the supernova type is then performed by combining spectral and light curve informations using the spectrophotometric template SALT2 developed by \\citet{Guy07}. In this approach, all available information on a given supernova (both photometric and spectral) is used to ease the identification. This contrasts with the more straightforward method of a two component spectral model. As underlined in, e.g., \\citet{Hook05,Lidman05,Balland06,Balland07}, a difficulty of this approach is the possible confusion of \\Iae spectra past maximum with Type Ic supernovae (SNe~Ic) at an earlier phase. Using the photometric phase as a constraint is crucial to alleviate this degeneracy.  In this paper, we present the techniques developed for the spectral extraction and the identification of our spectra. {We illustrate our results with a few sample spectra obtained at VLT, as part of two large spectroscopic programmes running from June 2003 to September 2007, that are being processed using this technique  \\citep{Balland08}}. In Sections 2, 3 and 4, we describe the PHASE technique and test it on simulated data. In Section 5, we illustrate PHASE results and discuss the improvements over standard extractions. In Section 6, the identification method of SNLS spectra using the SALT2 spectral template is described and results are discussed on a few examples. Discussion and conclusion are in Section 7 and Section 8, respectively. ",
        "conclusions": "We have developed new techniques for both supernova spectral extraction and identification. These new tools have been developed for the purpose of making use of the deep imaging obtained at CFHT to obtain an homogeneous set of SNLS supernova spectra.  Considerable effort has been put into extracting, in the most efficient way, VLT spectra in order to get a clean set of supernova spectra at redshifts ranging from $z=0.1$ to $z\\approx 1$. One key feature of our so-called PHASE extraction is that the extraction is guided by computing the trace of the supernova on the spectrogram. This is done by using deep CHTLS reference images, in various photometric bands, to recover the profile of the non-transient objects present in the slit, from which a multi-component model is built by adding the (point-like) supernova flux as a Gaussian of width equal to the seeing. Fitting the model to the spectrogram light yields the flux of each component in each pixel. As we have shown, this improves the host-supernova separation over more standard techniques. A second key feature of the PHASE extraction is that it avoids re-sampling the data and correlating pixels along the procedure. This ensures a clean estimation of the error level associated with the signal used later for the identification.  We use the spectrophotometric model of SALT2 \\citep{Guy07} as an help to determine the type of supernova candidates. Taking advantage of an ever increasing training set as new supernovae are discovered and followed up by the SNLS and other collaborations, the model is able to adequately reproduce \\Ia features in a wide range of cases (\\citealt{Balland08} show that it is true from phases as early as -10,-15 days up to a few weeks after maximum light and for a wide spectral range, including the UV region down to 2100 $\\AA$ for the most distant supernovae). The combined fit of light curves and spectrum tightens the constraints on the parameters that describe the supernova. Non \\Iae candidates are identified on a case-by-case basis as their spectrophotometric best-fitting parameters deviate from the average properties of the SNe~Ia training sample. Uncertainty however remains for the most host-contaminated (host fraction $>$ 70-80\\%) and most distant ($z\\gesim 1$) cases.  PHASE extractions have been tested against simulations, and SALT2 based identifications have been cross-checked with template fitting identifications performed on the VLT data using the technique described in \\citet{Howell05}.  We find that using SALT2 for the purpose of identification improves the reliability of the type determination over standard techniques. For VLT spectra, PHASE extraction combined to SALT2 identification takes the most possible out of the data to secure a clean type and redshift determination."
    },
    "0809/0809.1364_arXiv.txt": {
        "abstract": "We examine the biases induced on cosmological parameters when the presence of secondary anisotropies is not taken into account in Cosmic Microwave Background analyses. We first develop an exact analytical expression for computing the biases on parameters when any additive signal is neglected in the analysis. We then apply it in the context of the forthcoming \\emph{Planck} experiment. For illustration, we investigate the effect of the sole residual thermal Sunyaev--Zel'dovich signal that remains after cluster extraction. We find in particular that analyses neglecting the presence of this contribution introduce on the cosmological parameters $n_{\\rm s}$ and $\\tau$ biases, at least $\\sim 6.5$ and $2.9$ times their one $\\sigma$ confidence intervals. The $\\Omega_{\\rm b}$ parameter is also biased to a lesser extent. ",
        "introduction": "Future Cosmic Microwave Background (CMB) experiments, which are designed to be cosmic variance limited, will allow us to determine the cosmological parameters with a relative precision of the order of, or better than, one percent.  It will be made possible in particular through the measurement of CMB temperature and polarisation anisotropies with unequalled sensitivities, exquisite angular resolution and optimal frequency coverage.  In this context, additional contributions to the signal (galaxies, point sources, secondaries arising from the interaction of CMB photons with matter after decoupling, etc.)  can no more be neglected. More specifically, a precise quantification of the biases on the parameters and of their sources is now needed.   The study of biases in cosmology is receiving growing attention. In the context of weak lensing, \\citet{amara07} have derived a method based on a Fisher matrix type analysis for quantifying systematic biases.  In CMB analyses, \\citet{Miller08} examined the biases introduced by beam systematics for five upcoming experiments that will measure the B-mode polarisation (\\emph{Planck}, PolarBeaR, Spider, Q/U Imaging Experiment (QUIET)+Clover and CMBPol).  Similarly, using \\emph{Planck} characteristics, previous studies have estimated the biases induced by the contribution from patchy reionisation \\citep{zahn05}. They concluded that the biases, depending on the model of reionisation, can be as high as a few in units of the one sigma error.  More recently, \\citet{serra08} have focused on the contribution from clustered IR point sources and its effect on the cosmological parameters. Those two studies used quite different approaches. While the \\citet{serra08} analysis was based on Monte Carlo Markov Chains (MCMC), \\citeauthor{zahn05} used an approximate analytic computation of the biases \\citep[see also][]{huterer06}.  In this study, we present an analytical derivation of the biases on the cosmological parameters that goes beyond the aforementioned approximations.  It is a method valid when the primary signal and secondary contribution (astrophysical or systematic) are additive and it is exact as it can be applied even when the secondary signal is dominant over the primary. The method presented here is applied to the estimate of the cosmological parameters with the CMB power spectrum. We furthermore focus on one single source of additional anisotropies: those associated with undetected clusters. The Sunyaev--Zel'dovich (SZ) effect of galaxy clusters \\citep{SunyaevZeldovich72} is indeed the major source of secondary temperature anisotropies \\citep[][ and references therein]{Aghanimrevue08}.  The SZ effect is two-fold: the thermal SZ due to the inverse Compton scattering of photons off the hot electrons in the intra-cluster gas \\citep[e.g.][]{Rephaeli95, Itoh98}, and the kinetic effect due to the Doppler shift caused by the proper motion of the clusters in the CMB reference frame.  The upcoming large multi-frequency surveys will be able to detect and extract galaxy clusters using their specific thermal SZ spectral signature. We nevertheless expect some level of residual SZ contribution from undetected clusters in the temperature anisotropy maps. Such a residual signal might be the cause of the excess of power measured by small scale CMB experiments like BIMA (Berkeley Illinois Maryland Association Array), CBI, ACBAR \\citep{bima, cbi, acbar, cbi2}. The excess could be also due to unremoved point sources \\citep[][]{toffolatti05,Marian06}.    Our article is organised as follows, we present in Section 2 the method to calculate the biases on cosmological parameters. We then apply our method to estimate the biases induced by the SZ residuals. We present our results in Section 3 and discuss them in the following section. Finally, we conclude in Section 5.  Throughout the article, we assume a flat $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmological model and use for the reference model the \\emph{Wilkinson Microwave Anisotropy Probe5} (WMAP5) cosmological parameters \\citep{wmap5}, unless otherwise stated. ",
        "conclusions": "In this study, we develop an analytical method to calculate the biases on the cosmological parameters. Our method applies to any contamination provided the primary signal and the contamination are additive. Additionally, it is an exact derivation to the second order with the advantage of being applicable even when the contaminant dominates over the primary signal. The next generation of CMB experiments will measure, with a high sensitivity, the signal at small angular scales where secondary contributions intervene. We apply our method to the case of a contribution from undetected thermal SZ clusters to the primary CMB signal (assuming no contribution from foregrounds or point sources).  For illustration, we take the characteristics (noise, beam size) of \\emph{Planck} and compute the residual SZ signal from the undetected clusters assuming all clusters above 3 or 5$\\sigma_Y$ are detected simultaneously in the channels 100, 143 and 353~GHz. The residual SZ signal contributes more than 10 per cent of the total signal at multipoles higher than 1000. The higher the SZ cluster detection threshold, the higher the contamination.  We perform a bias estimation simultaneously on six cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm b}$, $H_0$, $n_{\\rm s}$, $\\sigma_8$ and $\\tau$) using temperature and polarisation anisotropy TT, TE and EE power spectra. This quantifies the effect of fitting the data, that include a residual contribution, with a model that ignores it.  We then compare the biases to the expected 1$\\sigma$ errors on each parameter. We find that the biases induced by the thermal SZ residual signal on $\\Omega_\\Lambda$ and $H_0$ are negligible. At most, they are of the order of 0.08 and 0.6 in units of the 1$\\sigma$ error on the parameters. On the contrary, the determination of $\\Omega_{\\rm b}$, $n_{\\rm s}$ and $\\tau$ is significantly altered by the residual SZ signal. The biases are 2.4, 6. and 10.4$\\sigma$ respectively. This is easily understood as they are the most sensitive parameters to an excess of power at small scales and, moreover, they are degenerate.  We point out the importance of taking into account the SZ residuals in the analysis of the small scale high sensitivity CMB data. The SZ residual power spectrum depends on the cosmological parameters and on the cluster selection function. A joint analysis of primary and secondary CMB signal will provide additional constraints on the cosmological parameters and thus reduce the biases arising from the SZ residual power excess at high multipoles. A coherent analysis, including full cosmological parameters dependency of the primary and secondary signal, allow one to use the whole range of multipoles, including the highest ones."
    },
    "0809/0809.3157_arXiv.txt": {
        "abstract": "The recent discovery of a new class of recurrent and fast X--ray transient sources, the Supergiant Fast X--ray Transients, poses interesting questions on the possible mechanisms responsible for their transient X--ray emission. The association with blue supergiants, the spectral properties similar to those of accreting pulsars and the detection, in a few cases, of X-ray pulsations, confirm that these transients are High Mass X-ray Binaries. I review the different mechanisms proposed to explain their transient outbursts and the link to persistent wind accretors. I discuss the different models in light of the new observational results coming from an on-going monitoring campaign of four Supergiant Fast X--ray Transients with \\emph{Swift}. ",
        "introduction": "An unusual X--ray transient in the direction of the Galactic centre region, XTE~J1739--302, was discovered in 1997 (Smith et al. 1998). It  displayed peculiar properties, particularly related with the duration of the outburst: this source remained active only one day. Since the X--ray spectrum was well fitted with a thermal bremsstrahlung with a temperature of $\\sim$20~keV, resembling the spectral properties of accreting pulsars, it was at first classified as a peculiar Be/X-ray transient with an unusually short outburst (Smith et al. 1998).  The INTEGRAL satellite has been performing a monitoring of the Galactic plane since its launch in October 2002 (Bird et al. 2007). Thanks to these observations, several other sources have been discovered, displaying similar properties to XTE~J1739--302: they show sporadic, recurrent, bright and short flares (with a typical duration of a few hours; Sguera et al. 2005, 2006;  Negueruela et al. 2006a). The X-ray spectra resemble the typical shape of High Mass X--ray Binaries (HMXBs) hosting X--ray pulsars, with a flat hard power law below 10 keV, and a high energy cut-off at about 15-30~keV, sometimes strongly absorbed at soft energies (Walter et al., 2006; Sidoli et al., 2006).  Follow-up X--ray observations (e.g. Kennea et al. 2005, Tomsick et al. 2006) allowed to refine the source positions at the arcsec level, and to perform IR/optical observations, which permitted to associate these transients with OB supergiant companions (e.g. Halpern et al. 2004, Pellizza et al. 2006, Masetti et al. 2006, Negueruela et al. 2006b, Nespoli et al. 2008). These two main characterizing properties (the unusually short transient X--ray emission together with the association with blue supergiant companions) suggested that these sources define a new class of HMXBs, later named Supergiant Fast X--ray Transients (SFXTs). Indeed, HMXBs with supergiant companions were previously known to show only persistent X--ray emission (e.g. Nagase, 1989). Thanks to the observations of several new SFXTs performed with INTEGRAL (e.g., Sguera et al. 2005, 2006, 2007a; Negueruela et al. 2006a; Walter \\& Zurita Heras 2007; Blay et al. 2008) and, at softer energies, with \\emph{Chandra} (in't Zand 2005), XMM-Newton (Gonzalez-Riestra et al. 2004; Walter et al. 2006) and archival observations with ASCA (Sakano et al. 2002), other important properties have been observed: SFXTs display a high dynamic range, spanning 3 to 5 orders of magnitudes, from a quiescent emission at 10$^{32}$~erg~s$^{-1}$ (characterized by a very soft spectrum, likely thermal; IGR~J17544--2619, in't Zand 2005; IGR J08408--4503, Leyder et al. 2007) up to a peak emission in outburst of 10$^{36}$--10$^{37}$~erg~s$^{-1}$. At least two SFXTs are X--ray pulsars: IGR~J11215--5952 ($P_{\\rm spin}=186.78\\pm0.3$\\,s, Swank et al. 2007) and AX~J1841.0--0536 ($P_{\\rm spin}$$\\sim$4.7\\,s, Bamba et al. 2001). To date, eight are the confirmed members of the class of SFXTs (IGR~J08408--4503, IGR~J11215--5952, IGR~J16479--4514, XTE~J1739--302, IGR~J17544--2619, SAX~J1818.6--1703, IGR~J18410-0535/AX~J1841.0-0536, IGR~J18483--0311), with other $\\sim$15 candidates (see e.g., the new INTEGRAL sources web page at http://isdc.unige.ch/$\\sim$rodrigue/html/igrsources.html) which showed short transient flaring activity, but with no confirmed association with an OB supergiant companion. The field is rapidly evolving, with an increasing number of  new transients discovered by instruments with a wide field of view (INTEGRAL/IBIS or Swift/BAT), so we expect that the whole class will grow in the near future.  The first SFXT displaying periodic outbursts is IGR~J11215--5952 (Sidoli et al. 2006, Smith et al. 2006) which undergoes an outburst about every  165~days (Romano et al. 2007b). This periodically recurrent X--ray activity is probably related to the orbital period of the system. Thanks to the predictability of the outbursts, an observing campaign was planned in February 2007 with \\emph{Swift} (Romano et al. 2007a) of the fifth known outburst from IGR~J11215--5952. Thanks to the combination of sensitivity and time coverage, \\emph{Swift} observations are a unique data-set and allowed a study of  this SFXT from outburst onset to almost quiescence. These observations demonstrated (see the \\emph{Swift}/XRT lightcurve in Fig.~\\ref{fig:igr11215}) that short duration and bright flares are actually part of a longer accretion phase at a lower level, lasting days, not only hours. ",
        "conclusions": ""
    },
    "0809/0809.5151_arXiv.txt": {
        "abstract": "{ RAVE, the RAdial  Velocity Experiment, is an ambitious  program to conduct a survey to measure the  radial velocities, metallicities and abundance ratios for  up to  a million  stars using  the 1.2-m  UK Schmidt  Telescope  of the Anglo-Australian Observatory (AAO), over the  period 2003 - 2010. The survey represents a  giant leap forward in  our understanding of our  own Milky Way galaxy, providing  a vast stellar  kinematic database larger than  any other survey proposed  for this coming  decade.  The main  data product will  be a southern  hemisphere survey  of about  a million  stars.  This  survey would comprise  0.7 million  thin disk  main  sequence stars,  250\\,000 thick  disk stars,  100\\,000 bulge  and  halo stars,  and  a further  50\\,000 giant  stars including some out to  10 kpc from the Sun. RAVE will  offer the first truly representative  inventory  of  stellar   radial  velocities  for  all  major components of the  Galaxy.  Here we present the  first scientific results of this survey as  well as its second data release which  doubles the number of previously released radial velocities.  For the first time, the release also provides atmospheric  parameters for a  large fraction of the  2$^{nd}$ year data making  it an unprecedented  tool to study  the formation of  the Milky Way.} ",
        "introduction": "It is now widely accepted that the Milky Way galaxy is a suitable laboratory to study the formation and evolution of galaxies.  Despite the fact that the Galaxy  is one unique  system, understanding  its formation  holds important keys to understand  the broader context of disc  galaxy formation. Thanks to the past and ongoing large surveys  such as Hipparcos, SDSS, 2MASS or DENIS, we have access  to data which contains a wealth  of information which enable us to refine our knowledge of Galaxy formation.  However, with the exception of  the  SDSS survey,  the  full  description of  the  6D  phase space,  the combination of the position and velocity spaces, is not available due to the missing radial velocity and/or distance.   With the advent of multi-fiber  spectroscopy, combined to the large field of view of  Schmidt telescopes,  it is  now possible to  acquire spectra  for a large sample of  stars representative of the populations of  the Galaxy in a reasonable amount of time. Spectroscopy  enables us to measure the generally missing  radial  velocity, which  allows  us to  study  the  details of  the dynamics of a system.  Spectroscopy also permits to measure the abundance of chemical elements in  a stars atmosphere which holds  important clues on the chemical  evolution of  its parent  system.  Measuring  this  missing radial velocity and the chemical  abundances to complement the existing astrometric and photometric catalogues is the main purpose of the RAVE project, with the driving goal to understand how the Milky Way formed.   RAVE uses the 6dF, the multi--fiber spectroscopic facility at the UK Schmidt telescope   of   the  Anglo--Australian   Observatory   in  Siding   Spring, Australia. The  6dF enables us  to collect up  to 150 spectra in  one single pointing, with a  resolution of 7\\,500 in the  Calcium triplet region around 8\\,600\\AA. This medium resolution  allows the measurement of accurate radial velocities  ($\\sim$2~km/s)  as well  as  atmospheric parameters  ($T_{eff}$, $\\log g$, [M/H]) and chemical abundances.  In Section 2 we present the first scientific outcomes obtained  using the RAVE data, while  Section 3 presents the second data release of the  RAVE project (DR2) and discusses the current status and prospects of the project. ",
        "conclusions": "RAVE released its second catalogue in July 2008, reporting radial velocities for 51\\,829  spectra and 49\\,327  different stars, randomly selected  in the magnitude range of $9<I<12$ and located more than 25$^{\\circ}$ away from the Galactic plane.  This  release doubles the size of  the previously published catalogue. The typical error on the published radial velocity is between 1.3 and 1.7~km/s.  In addition to radial velocities,  the catalogue contains for the first time atmospheric    parameters     measurements    for    more     than    20\\,000 spectra. Uncertainties for a typical RAVE  star are of 400 K in temperature, 0.5 dex in gravity, and 0.2 dex in metallicity, but the error depends on the S/N and can vary by a factor of $\\geq 2$ for stars at the extreme of the S/N range.\\\\  The survey continues  to collect spectra and is  scheduled to continue until 2010.   Currently more  than 300\\,000  spectra have  been collected  and are currently  being processed.   Acquisition of  new calibration  data  is also underway, enabling  us to refine our atmospheric  parameters measurements as well  as our radial  velocities, and  represents an  important task  for the ongoing  effort. Other  ongoing  efforts in  the  collaboration include  the measurement of chemical  abundances and distances which will  be released as companion catalogues in the future.\\\\  The RAVE catalogue can be used to  study the formation of the Milky Way and, for  example,  the collaboration  succeeded  in  refining significantly  the measurement of the local escape velocity using the high velocity stars found in  the RAVE  catalogue. Further  studies carried  out by  the collaboration include the search for vertical  stream of matter in the Solar neighborhood, a  new determination of  the vertical  structure of  the Galactic  disc, the study  of  diffuse  interstellar  bands  in the  Galactic  plane  and  their correlation  to  the  interstellar  extinction  or the  inclination  of  the velocity ellipsoid towards the Galactic plane at 1~kpc below the plane.\\\\  As the survey  progresses and new data become  available, the RAVE catalogue will enable  us to  refine our view  on the  structure and formation  of the Milky Way, paving the way for new insights on the formation of galaxies.\\\\"
    },
    "0809/0809.0404_arXiv.txt": {
        "abstract": "We calculate the spin dependent Fermi liquid parameters (FLPs), single particle energies and  energy densities of various spin states of polarized quark matter. The expressions for the incompressibility($K$) and sound velocity ($c_1$) in terms of the spin dependent FLPs and polarization parameter $(\\xi)$ are derived. Estimated values of $K$ and $c_1$ reveal that the equation of state (EOS) of the polarized matter is stiffer than the unpolarized one. Finally we investigate the possibility of the spin polarization (ferromagnetism) phase transition. ",
        "introduction": "One of the important research areas of the contemporary high energy physics has been the study of matter under extreme conditions. Such a matter, in the laboratory can be produced by colliding heavy ions at ultra-relativistic energies. Due to asymptotic freedom of quantum chromodynamics (QCD), it is predicted that the hadronic matter at high temperature and/or density can undergo a series of phase transitions like confinement-deconfinement and/or chiral phase transition {\\cite{hwa_book,niegawa05}}. In the high density regime QCD predicts the existence of color superconducting state {\\cite{rajagopal,inui07,nakano03}}. These apart, the possibility of spin polarized quark liquid {\\em i.e.} the existence of ferromagnetic phase in dense quark system has also been suggested recently with which we are presently concerned {\\cite {niegawa05,tatsumi00}}. The properties of dense quark system are particularly relevant for the study of various astrophysical phenomenon.  The part of the motivation to study the ferromagnetic phase transition in dense quark matter (DQM), as mentioned in \\cite{tatsumi00} is provided by the discovery of `magnetars' {\\cite{tatsumi01}} where an extraordinarily high magnetic field $\\sim 10^{15} G$ exists \\cite{tatsumi00,son08}. In \\cite{tatsumi00}, it is argued, that the origin of such a high magnetic field can be attributed to the existence of spin polarized quark matter {\\cite{tatsumi06}}. To examine the possibility of ferromagnetism in DQM in ref.\\cite{tatsumi00} a variational calculation is performed where it is observed that there exists a critical density below which spin polarized quark matter is energetically favorable than unpolarized state. Subsequently various other calculations were also performed to investigate this issue {\\cite{son08,niegawa05,inui07,nakano03,tatsumi01,tatsumi06,ohnishi07}}. For example, in \\cite{nakano03} it is shown that there is no contradiction between color superconductivity and ferromagnetism and both of these phase can co-exist. In \\cite{ohnishi07}, the same problem was studied in the large $N_c$ and $N_f$ limit while keeping $N_c/N_f$ fixed where it was shown that spin polarized state can exist, however, in presence of magnetic screening, color superconductivity or dense chiral waves disappear. It might be mentioned that such screening is now supported by the lattice calculation {\\cite{hand06, ohnishi07}}. In{\\cite{nakano03}} it is analytically shown that, if quarks are massless, ferromagnetism does not appear which is consistent with the conclusion drawn in \\cite{ohnishi07}. Ref.\\cite{son08} shows that ferromagnetism might appear in quark matter with Goldstone boson current where the magnetization is shown to be related to triangle anomalies.  In the present work, we apply relativistic Fermi liquid theory (RFLT) to study the possibility of para-ferro phase transition in DQM. The relativistic Fermi liquid theory was developed by Baym and Chin \\cite{baym76} where it has been shown how the  various physical quantities like chemical potential ($\\mu$), incompressibility ($K$), sound velocity ($c_1$) etc. can be expressed in terms of the Landau parameters (LPs) calculated relativistically. However, the formalism developed in \\cite{baym76} is valid for unpolarized matter and LPs calculated there are spin averaged.  In this paper we extend the formalism of RFLT and the required LPs are calculated by retaining their explicit spin dependencies.  As a result, here various combination of parameters like $f_{0,1}^{++} $, $f_{0,1}^{+-}$, $f_{0,1}^{-+} $ and $f_{0,1}^{--}$  corresponding to scattering involving up-up, up-down, down-up or down-down spins are appear \\cite{baym76}.  Once determined, these parameters are  used to calculate quantities like chemical potentials for the spin up and spin down quarks or the total energy density of the system as a function of ${\\xi}= (n^+_q-n^-_q)/n_q$ and $n_q$ together with various other quantities as we shall see. Here $n_q^+$ and $n_q^-$ correspond to densities of spin up, down quarks respectively and $n_{q}=n_{q}^{+}+n_{q}^{-}$, denotes total quark density \\cite{tatsumi00}. We, also compare some of our results with those presented in \\cite{tatsumi00} where more direct approach was adopted to calculate the total energy density from the loop. In addition, the present work is extended further to estimate incompressibility and sound velocity in dense quark system for a given fraction of spin-up or down quarks.   Furthermore, in dealing with the massless gluons, we find that naive series expansion fails and one has to use hard density loop (HDL) corrected gluon propagator to get the finite result for the LPs involving scattering of like spins {\\cite{tatsumi08}}. This however does not cause any problem for the calculation of various physical quantities like chemical potential, exchange energy, incompressibility etc. We shall see, even though  $f_0$ and $f_1$ (suppressing spin indices) individually remain divergent, what appears in our case is the particular combination of these parameters where such divergences cancel.  The plan of the paper is as follows. In Sec.II, as mentioned before, we extend the formalism of RFLT to include explicit spin dependence. In Sec.III, we derive spin dependent LPs due to one gluon exchange (OGE) for polarized quark matter. Subsequently, we calculate chemical potential and energy density. We find the density dependence of incompressibility $(K)$ and first sound velocity $(c_1)$ with arbitrary spin polarization $(\\xi)$. To compare with ref.{\\cite{tatsumi00}}, we present ultra-relativistic and non-relativistic results and studied para-ferro phase transition of quark matter. Sec.IV is devoted to summary and conclusion. In Appendix, we calculate various LPs for unlike spin states of scatterer.  \\vskip 0.4in \\section {Formalism}  In FLT total energy density $E$ of an interacting system is the functional of occupation number $n_{p}$ of the quasi-particle states of momentum $p$. The excitation of the system is equivalent to the change of occupation number by an amount $\\delta n_{p}$. The corresponding energy density of the system is given by {\\cite{baym_book,baym76}}, \\beq{\\label {total_energy}} E&=&E^{0}+\\sum_{s}\\int\\frac{d^3{p}}{(2\\pi)^3} \\varepsilon_{ps}^{0}\\delta n_{ps} +\\frac{1}{2}\\sum_{ss'}\\int\\frac{d^3{p}}{(2\\pi)^3}\\frac{d^3{p'}}{(2\\pi)^3} f_{ps,p's'} \\delta n_{ps}\\delta n_{p's'}, \\eeq where $E^0$ is the ground state energy density and $s$ is the spin index, and the quasi-particle energy can be written as, \\beq\\label{quasi_energy} \\veps_{ps}=\\veps_{ps}^{0}+\\sum_{s'}\\int\\frac{d^3{p'}}{(2\\pi)^3}f_{ps,p's'} \\delta n_{p's'}, \\eeq where $\\veps_{ps}^{0}$ is the non-interacting single particle energy. The interaction between quasi-particles is given by $f_{ps,p's'}$, which is defined to be the second derivative of the energy functional with respect to occupation functions, \\beq\\label{quasi_interac} f_{ps,p's'}=\\frac{\\delta^{2}E}{\\delta{n}_{ps}~\\delta{n}_{p's'}}. \\eeq  Since, the quasiparticles are well defined only near the Fermi surface, one assumes \\beq \\veps_{ps}&=&\\mu^{s}+v_{f}^{s}(p-p_{f}^{s}). \\eeq   In FLT, the interaction parameter, $f_{ps,p's'}$, is expanded on the basis of Legendre polynomials, $P_l$ \\cite{baym_book,baym76}. The coefficients of this expansion are known as FLPs, which are given by  \\beq\\label{landau_para} f_l^{ss'}=(2l+1)\\int\\frac{d\\Omega}{4\\pi}P_{l}(\\cos\\theta)f_{ps,p's'}, \\eeq  where $\\theta$ is the angle between $p$ and $p'$, both taken to be on the Fermi surface, and the integration is over all directions of $p$ {\\cite{baym76}}. Note that unlike \\cite{baym_book,baym76}, here we retain explicit spin indices without performing spin summation. We restrict ourselves for $l\\le 1$ i.e. $f_{0}^{s}$ and $f_{1}^{s}$, since higher $l$ contribution decreases rapidly as the scattering is dominated by the small angles and the series converges, here, $f_{l}^{s}=\\frac{1}{2}\\sum_{s'}f_{l}^{ss'}$ {\\cite{mig_book}}.  The Landau Fermi liquid interaction $f_{ps,p's'}$ is related to the two particle forward scattering amplitude via {\\cite{baym_book,baym76}},  \\beq\\label{int_para} f_{ps,p's'}&=&\\frac{m_{q}}{\\veps_{p}^0}\\frac{m_{q}}{\\veps_{p'}^0} {\\cal M}_{ps,p's'}, \\eeq  where $m_{q}$ is the mass of the quark and the Lorentz invariant matrix ${\\cal M}_{ps,p's'}$ consists of the usual direct and exchange amplitude, which may, therefore be evaluated by conventional Feynman rules. The dimensionless LPs are defined as $F_{l}^{s}=N^{s}(0)f_{l}^{s}$ {\\cite{baym76}}, where $N^s(0)$ is the density of states at the Fermi surface is given by,  \\beq\\label{dens_of_state} N^{s}(0)&=&\\int\\frac{\\rm d^3{p}}{(2\\pi)^3} \\delta(\\veps_{ps}-\\mu^{s})\\nn\\\\ &=&\\frac{g_{deg}p_{f}^{s^2}}{2\\pi^2}\\left(\\frac{\\del p}{\\del\\veps_{ps}} \\right)_{p=p_{f}^{s}}\\nn\\\\ &\\simeq&\\frac{g_{deg}p_{f}^{s}\\veps_{f}^{s}}{2\\pi^2}. \\eeq Here $g_{deg}$ is the degeneracy factor. In our case $g_{deg}=N_{c}N_{f}$ where $N_{c}$ and $N_{f}$ are the color and flavor index for quark matter. For spin up $(+)$ and spin down $(-)$ quark, density of states will be change accordingly. In the above expression $(\\del p/\\del\\veps_{ps})_{p=p_{f}^{s}}$ is the inverse Fermi velocity $(1/v_{f}^{s})$ related to the FL parameter $F_{1}^{s}$,  \\beq\\label{inv_vel} \\frac{1}{v_{f}^{s}}=(\\del p/\\del\\veps_{ps})_{p=p_{f}^{s}}=(\\mu^{s}/p_{f}^{s}) (1+F_{1}^{s}/3). \\eeq  With Eq.(\\ref{dens_of_state}) and Eq.(\\ref{inv_vel}) one reads the general relation as {\\cite{matsui81}}  \\beq\\label{relative} \\veps_{f}^{s}=\\mu^{s}(1+\\frac{1}{3}F_{1}^{s}). \\eeq  The compression modulus or incompressibility $(K)$ of the system is defined by the second derivative of total energy density $E$ with respect to the number density $n_{q}$, is given by \\cite{matsui81,chin77,brown02,negele_book,aguirre07}  \\beq\\label{compres1} K&=&9n_{q}\\frac{\\del^2 E}{\\del n_{q}^2}. \\eeq  Now we introduce a polarization parameter ${\\xi}$ by the equations, $n_{q}^{+}=n_{q}(1+{\\xi})/2$ and $n_{q}^{-}=n_{q}(1-{\\xi})/2$  under the condition $0\\le{\\xi}\\le 1$ {\\cite{tatsumi00}}. The Fermi momenta in the spin-polarized quark matter then are $p_{f}^{+}=p_{f}(1+{\\xi})^{1/3}$ and $p_{f}^{-}=p_{f}(1-{\\xi})^{1/3}$, where $p_{f}=(\\pi^2n_{q})^{1/3}$, is the Fermi momentum of the unpolarized matter $({\\xi}=0)$. So, there are two Fermi surfaces corresponding to spin-up $(+)$ and spin-down $(-)$ states, such that $E\\equiv E(n_{q}^{+},n_{q}^{-})$. We have  \\beq\\label{delE_deln} \\frac{\\del E}{\\del n_{q}}&=&\\frac{\\del E}{\\del n_{q}^{+}} \\frac{\\del n_{q}^{+}}{\\del n_{q}}+\\frac{\\del E}{\\del n_{q}^{-}} \\frac{\\del n_{q}^{-}}{\\del n_{q}}\\nn\\\\ &=&\\frac{1}{2}\\left[(1+{\\xi})\\mu^{+}+(1-{\\xi})\\mu^{-}\\right] \\eeq Using Eq.(\\ref{delE_deln}), the incompressibility becomes{\\cite{aguirre07}} \\beq\\label{compres} K&=&\\frac{9n_{q}}{4}\\left[(1+{\\xi})^2\\frac{\\del \\mu^{+}}{\\del n_{q}^{+}} +(1-{\\xi})^2\\frac{\\del \\mu^{-}}{\\del n_{q}^{-}}\\right]\\nn\\\\ &=&\\frac{9n_{q}}{4}\\left[(1+{\\xi})^2\\left(\\frac{1+F_{0}^{+}}{N^{+}(0)}\\right) +(1-{\\xi})^2\\left(\\frac{1+F_{0}^{-}}{N^{-}(0)}\\right)\\right], \\eeq where {\\cite{baym76}}  \\beq \\frac{\\del \\mu^{s}}{\\del n_{q}^{s}}&=&\\frac{1+F_{0}^{s}}{N^{s}(0)}. \\eeq   Similarly, the relativistic first sound velocity is given by the first derivative of pressure $P$ with respect to energy density $E$. Since $P={\\sum}_{s}\\mu^{s}n_{q}^{s}-E$ {\\cite{aguirre07, matsui81}}, we have,  \\beq\\label{soundvel} c_{1}^{2}=\\frac{\\del P}{\\del E} &=&\\frac{\\del P}{\\del n_{q}}\\frac{\\del n_{q}}{\\del E}\\nn\\\\ &=&\\left[\\frac{(1+{\\xi})n_{q}^{+}\\frac{\\del \\mu^{+}}{\\del n_{q}^{+}} +(1-{\\xi})n_{q}^{-}\\frac{\\del \\mu^{-}}{\\del n_{q}^{-}}} {(1+{\\xi})\\mu^{+}+(1-{\\xi})\\mu^{-}}\\right]\\nn\\\\ &=&\\frac{n_{q}}{2[(1+{\\xi})\\mu^{+}+(1-{\\xi})\\mu^{-}]} \\left[(1+{\\xi})^2\\left(\\frac{1+F_{0}^{+}}{N^{+}(0)}\\right) +(1-{\\xi})^2\\left(\\frac{1+F_{0}^{-}}{N^{-}(0)}\\right)\\right]. \\eeq In the above Eq.(\\ref{compres}) and Eq.(\\ref{soundvel}), $N^{\\pm}(0)$ and $F_{0}^{\\pm}$ corresponds to density of states at Fermi surface and dimensionless LP for spin up $(+)$ and spin down $(-)$ quark respectively. For unpolarized matter, $\\xi=0$ implying $\\mu^+=\\mu^-$, $F_0^+=F_0^-$ and $N^+(0)=N^-(0)$. From Eq.(\\ref{compres}) and (\\ref{soundvel}) we have the well known result as $K=9n_{q}\\frac{\\del \\mu}{\\del n_{q}}$ {\\cite{matsui81}} and $c_1^2=\\frac{n_{q}}{\\mu}\\frac{\\del\\mu}{\\del n_{q}}$ {\\cite{baym76}}.    \\vskip 0.4in ",
        "conclusions": "To summarize and conclude, in this work we have applied RFLT to study the properties of dense quark matter. Accordingly, we calculate the FLPs by retaining their explicit spin dependencies. We also show how the physical quantities like chemical potential of spin up and spin down states, their energy densities and the quantities like incompressibility, sound velocity for polarized quark matter can be expressed in terms of these spin dependent RFLPs.  For the scattering involving like spin states, the LPs $f_{0,1}^{++}$ and $f_{0,1}^{--}$ are found to diverge. However, we show that the combination in which they appear in the calculation of the physical quantities such divergences cancel. For the scattering involving unlike spin states no such divergence appear. The appearance of such divergences is related to the unscreened gluonic interaction between the quarks which might be cured by invoking hard dense loop corrected gluon propagator. We do not perform such calculation here and postpone this for future investigation. As far as the equation of state (EOS) is concerned, we in the present model find that the EOS for the polarized quark matter is stiffer than the unpolarized one. In addition, we also show that there exists a metastable state which disappear at higher density, although it seems that the effect is tiny.   We reconfirm that DQM can exhibit ferromagnetism at low density as was originally suggested in {\\cite{tatsumi00}}. However, the density at which the spin polarized ferromagnetic state in the present model might appear depends strongly on the quark mass. The critical density increases with increasing mass. In Fig.~(\\ref{engtot}), we observe that states with $\\xi$ appear only below or around normal nuclear density where deconfined quark matter is not likely to exist. We cannot, however, ascertain the critical density from the present analysis where we restrict ourselves only to OGE diagrams and one flavor system. In this regime, multi-gluon exchange processes {\\cite{ohnishi07}} might play an important role. Furthermore, the correlations as given by the ring diagrams can also change the conclusion. Further work therefore is necessary to understand the existence of ferromagnetic quark matter in real multi flavor system which might appear in astrophysics."
    },
    "0809/0809.2637_arXiv.txt": {
        "abstract": "We used the red clump stars from the Optical Gravitational Lensing Experiment  (OGLE II) survey and the the Magellanic Cloud Photometric Survey (MCPS), to estimate the line of sight depth. The observed dispersion in the magnitude and colour distribution of red clump stars is used to estimate the line of sight depth, after correcting for the contribution due to other effects. This dispersion due to depth, has a range from minimum dispersion that can be estimated, to 0.46 mag (a depth of 500 pc to 10.44 Kpc), in the LMC. In the case of SMC, the dispersion ranges from minimum dispersion to 0.35 magnitude (a depth of 665 pc to 9.53 Kpc). The thickness profile of LMC bar indicates that it is flared. The average depth in the bar region is 4.0$\\pm$1.4 kpc.  The halo of the LMC (using RR Lyrea stars) is found to have larger depth compared to the disk/bar, which supports the presence of inner halo for the LMC. The large depth estimated for the LMC bar and the disk suggests that the LMC might have had minor mergers. In the case of SMC, the bar depth (4.90$\\pm$1.23 Kpc) and the disk depth (4.23$\\pm$1.48 Kpc) are found to be within the standard deviations. We find evidence for increase in depth near the optical center (up to 9 kpc). On the other hand, the estimated depth for the halo (RR Lyrea stars) and disk (RC stars) for the bar region of the SMC is found to be similar. Thus, increased depth and enhanced stellar as well as HI density near the optical center suggests that the SMC may have a bulge. ",
        "introduction": "The Magellanic Clouds were believed to have interactions with our Galaxy as well as between each other \\cite[(Westerlund 1997)]{west97}. The N-body simulations by \\cite[Weinberg (2000)]{wein00} predicted that the LMC's evolution is significantly affected by its interactions with Milky Way and the tidal forces will thicken and warp the LMC disk. \\cite[Alves and Nelson (2000)]{an00} studied the carbon star kinematics and found that the scale height, h, increases from 0.3 to 1.6 kpc over the range of radial distance, R, from 0.5 to 5.6 kpc and hence concluded that the LMC disk is flared. Using an expanded sample of carbon stars \\cite[van der Marel et al. (2002)]{vdm02} also found that the thickness of LMC disk increases with the radius. There has not been any direct estimate of the thickness or the line of sight depth of the bar and disk of the L\\&SMC so far.  \\cite[Mathewson, Ford \\& Visvanathan (1986)]{mfv86} found that SMC Cepheids extend from 43 to 75 Kpc with most Cepheids found in the neighbourhood of 59 Kpc. Later, the line of sight depth of SMC was estimated \\cite[(Welch et al. 1987)]{welch87} by investigating the line of sight distribution and period - luminosity relation of Cepheids and found the line of sight depth of SMC to be around 3.3 Kpc. \\cite[Hawkins et al. (1989)]{haw89}, estimated the line of sight depth in the outer regions of SMC to be around 10-20 kpc. Measurements of thickness in the central regions of Magellanic Clouds, especially LMC, is of strong interest to understand the contribution of LMC's self lensing to the observed microlensing events from this Galaxy.\\\\  Red Clump (RC) stars are core helium burning stars which are metal rich and slightly more massive counter parts of horizontal branch stars.  In the this paper, we use the dispersions in the colour and magnitude distribution of RC stars for the depth estimation. The dispersion in colour is due to a combination of observational error, internal reddening (reddening within the disk of the LMC/SMC) and population effects. The dispersion in magnitude is due to internal disk extinction, depth of the distribution, population effects and photometric errors associated with the observations. By deconvolving other effects from the dispersion of magnitude, we can estimate the depth of the disk. The advantage of choosing RC stars as proxy is that there are large number of these stars available to determine the dispersions in their distributions with good statistics. ",
        "conclusions": ""
    },
    "0809/0809.4011_arXiv.txt": {
        "abstract": "In the CDM scenario, dark matter halos are assembled hierarchically from smaller subunits. A long-standing problem with this picture is that the number of sub-halos predicted by CDM simulations is orders of magnitudes higher than the known number of satellite galaxies in the vicinity of the Milky Way. A plausible way out of this problem could be that the majority of these sub-halos somehow have so far evaded detection. If such ``dark galaxies'' do indeed exist, gravitational lensing may offer one of the most promising ways to detect them. Dark matter sub-halos in the 10$^6$ -- 10$^{10} M_{\\odot}$ mass range should cause strong gravitational lensing on (sub-)milliarcsecond scales. We study the feasibility of a strong lensing detection of dark sub-halos by deriving the image separations expected for density profiles favoured by recent simulations and comparing these to the angular resolution of both existing and upcoming observational facilities. We find that there is a reasonable probability to detect sub-halo lensing effects in high resolution observations at radio wavelengths, such as produced by the upcoming VSOP-2 satellite, and thereby test the existence of dark galaxies. ",
        "introduction": "A multitude of recent cosmological studies suggest that most of the matter in the universe consists of non-baryonic particles which are believed to constitute the cold dark matter (CDM) \\cite{Komatsu}. While the CDM scenario has been very successful in explaining the formation of large-scale structures in the Universe (see e.g., \\cite{Primack}, for a review), its predictions on the scales of individual galaxies have not yet been confirmed in any convincing way. One particularly interesting feature of the CDM model is the high level of halo substructure generated, seemingly inconsistent with the known number of satellite galaxies in the vicinity of the Milky Way (e.g., \\cite{klypin99,moore99a,diemand07b}). This discrepancy persists even when considering the newly discovered ultra-faint dwarf galaxies in the Sloan Digital Sky Survey and correcting for its sky coverage \\cite{simon07}. One possible way out of this problem is to assume that most of these low-mass halos correspond to so-called dark galaxies \\cite{Verde}, i.e., objects of dwarf-galaxy mass which either do not contain baryons or in which the baryons have not formed many stars.  Gravitational lensing may in principle offer a route to detecting even completely dark galaxies. Dark matter sub-halos in the 10$^6$--10$^{10} M_{\\odot}$ range are expected to cause gravitational milli-lensing, i.e., image splitting at a characteristic separation of milliarcseconds \\cite{wambsganss92,yonehara03}. To put the CDM sub-halo predictions to the test, it has been suggested that one should target quasars which are already known to be gravitationally lensed on arc second scales, as one can then be sure that there is a massive halo well-aligned with the line of sight, which substantially increases the probability for sub-halo milli-lensing \\cite{yonehara03}. Indeed, the magnification associated with milli-lensing has long been suspected to be the cause of the flux ratio anomalies seen in such systems \\cite{mao98,kochanek04}. Sub-halo milli-lensing has also been advocated as an explanation for strange bending angles of radio jets \\cite{metcalf02} and image positions which smooth halo models seem unable to account for \\cite{biggs04}.  Here, we take a critical look at the prospects for strong-lensing detections of dark sub-halos in the dwarf-galaxy mass range. In $\\S$\\ref{imsep}, we derive the image separations expected for sub-halo density profiles favoured by recent simulations. In $\\S$\\ref{lensprob}, we estimate the overall sub-halo lensing probabilities for point sources as well as extended sources. Finally, we discuss several effects that could affect our predictions and present our conclusions in $\\S$\\ref{concl}.  Throughout the paper, we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\Lambda}$ = 0.762, $\\Omega_{\\mathrm{M}}$ = 0.238 and $h = 0.73$ ($H_0=100h$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$) in concordance with the \\textit{WMAP} 3-year data release \\cite{spergel}. ",
        "conclusions": "\\label{concl}  Our results indicate that the detection of sub-halos through gravitational image-splitting is likely to be considerably more challenging than suggested in previous studies, due to the smaller image separations predicted for sub-halo density profiles more realistic than the SIS models often adopted. In fact, no currently planned telescope will be able to resolve the image separations produced by sub-halos with density profiles of the type suggested by the most realistic simulations recently available (H03 \\& K04). If the sub-halos would have steeper central density slopes (e.g., of M99 type), these would give rise to image separations that could be resolved even with existing telescopes. A recent simulation with the highest resolution currently available \\cite{diemand08}, obtained central density profiles of the form $\\rho(r) \\propto r^{-\\alpha}$ with central density slopes $\\alpha \\approx 1.24$ (i.e., between the unfavourable NFW with $\\alpha = 1$ and the optimistic M99 profile with $\\alpha = 1.5$). If strong lensing effects produced by sub-halos which such profiles could be observed with any present or future telescope, is something we will be looking into in the near future.  Despite the somewhat bleak detection prospects presented here, there are at least two effects that can potentially improve the detectability of image splitting by sub-halos: baryon cooling and the presence of intermediate mass black holes.  Baryon cooling would cause the sub-halo (or its progenitor halo) to contract (e.g., \\cite{Gnedin,Maccio,Gustafsson,Kampakoglou}), thereby increasing the central density and boosting the image separation. It is, however, not clear how strong this effect is likely to be, since this depends on the details of the mechanism that prevents the dark sub-halo from forming stars. Baryon cooling is usually assumed to be associated with star formation, and most attempts to explain the lack of bright sub-halos therefore propose that the baryons have either been lost early in the history of the Universe, or are prevented from cooling by the ultraviolet background provided by re-ionization \\cite{Barkana,Read}. Neither mechanism is likely to result in any significant halo contraction. Nonetheless, claims of dark baryons in the form of cold gas in galactic disks have been made \\cite{Bournaud}, implying that there may be some route for gas to cool without forming stars or being detected by the usual H$_2$ tracers \\cite{Pfennigera,Pfennigerb}. One may therefore speculate that there could be alternative scenarios, in which sub-halos are kept dark even though baryon cooling has taken place.  Intermediate mass black holes (IMBHs, see e.g., \\cite{vanderMarel,Zhao,Noyola}) with masses of $\\sim 10^2$--$10^4\\ M_\\odot$ would also boost the image separations, if present in the centres of dark sub-halos. This could for instance be the case if the empirical relations between the mass of a super-massive black hole and its dark matter halo \\cite{Ferrarese} would extend into the dwarf-galaxy mass range. Even in the case of $\\kappa=\\gamma=0$, an IMBH of mass $10^4 \\ M_\\odot$, would give an image separation of $\\approx 4\\times 10^{-4}$$^{\\prime\\prime}$ for $z_\\mathrm{l}=0.5$ and $z_\\mathrm{s}=2.0$. This is sufficient to allow VSOP-2 to resolve the image splitting, regardless of the density profile of the sub-halo hosting the IMBH.  Of course, even if the Ferrarese relations would hold for luminous dwarf galaxies, they do not necessarily do so for dark ones, since this depends on the formation details of IMBHs. Moreover, a population of halo IMBHs {\\it not} associated with sub-halos could form an undesired background of milli-lensing events that would obfuscate attempts to study sub-halos through image-splitting effects.  We have shown that the optical depth $\\tau$ for sub-halo lensing of point sources is lower than previously predicted. Even for the most favorable conditions it barely exceeds 0.01. We conclude that it is currently not feasible to use this technique to search for strong lensing signatures in point sources as e.g. quasars in the optical.  If one instead targets extended sources, such as quasars in the radio wavelength regime, there is a high probability for sub-halo lensing of sources of sufficient size. For source sizes $r_{\\mathrm{s}} \\gtrsim 1$\\,kpc, this is valid even at rather large projected distance of the source to the host halo center. This allows for a different search strategy than those previously proposed. Instead of only targeting multiply-imaged quasar systems, even quasar-galaxy pairs with a separation of several tens of arc seconds should show effects of strong lensing by substructures in the lens galaxy halo.  However, these effects will strongly depend on the sub-halo mass and density profile. Even for the most favorable density profiles, there will be minimum masses for which the image separation drops below the resolution of any present or planned observational facility. E.g., for the currently available EVN, the minimum sub-halo mass which could be resolved under optimal conditions is approximately $4 \\times 10^6 M_{\\odot}$ in case of an SIS profile but increases to $3 \\times 10^7 M_{\\odot}$ for an M99 profile. Figure \\ref{fig5} shows how the expected number of intervening substructures $N$ depends on the minimum sub-halo mass $M_{\\mathrm{sub\\_min}}$. For a host halo of mass $M = 1.8 \\times 10^{12} M_{\\odot}$ at $z_{\\mathrm{l}}$ = 0.5 and a source of size $r_{\\mathrm{s}} = 1$\\,kpc at $z_{\\mathrm{s}}$ = 1, we compute $N$ as a function of projected radius from the host halo center for minimum sub-halo mass $M_{\\mathrm{sub\\_min}} = 4 \\times 10^6 M_{\\odot}$, $10^8 M_{\\odot}$ and $10^9 M_{\\odot}$, respectively. Since the sub-halo mass function predicts about equal mass within each logarithmic mass interval, most sub-halos will be of low mass ($n(M_{\\mathrm{sub}}) \\propto M_{\\mathrm{sub}}^{-0.9}$). Thus, the expected number of intervening substructures is very sensitive to the minimum sub-halo mass that can be resolved.  \\begin{figure}[t] \\includegraphics[clip,width=0.60\\textwidth]{fig5.eps}\\hfill% \\begin{minipage}[b]{0.38\\textwidth}\\caption{\\label{fig5} Average number of substructures covering a source with $r_{\\mathrm{s}} = 1$\\,kpc at $z_{\\mathrm{s}} = 1$ as a function of projected radius for different minimum sub-halo masses $M_{\\mathrm{sub\\_min}}$. For a host halo at $z_{\\mathrm{l}} = 0.5$ with $M_{\\mathrm{host}} = 1.8 \\times 10^{12} M_{\\odot}$, we vary $M_{\\mathrm{sub\\_min}}$ with $4 \\times 10^6 M_{\\odot}$ (dotted line), $10^8 M_{\\odot}$ (dashed line) and $10^9 M_{\\odot}$ (solid line), respectively. The peak has been cut with respect to a maximum magnification factor of 30 at the Einstein radius of the host halo.} \\end{minipage} \\end{figure}  Therefore, angular resolution will be crucial when attempting to detect CDM substructure via its lensing effects on background quasars and sub-milliarcsecond-resolution facilities will be required. We assess that the prospects for such a detection with the upcoming VSOP-2 satellite are good, with the condition that the ground array should be extensive enough. This is not a problem especially at 8.4 GHz where a number of large antennas are available from the geodetic VLBI network. However, one must be careful not to confuse internal structures found in quasars observed at radio wavelengths with sub-halo lensing signals. Quasars already macro-lensed on arc second scales can be used to test that it is possible to distinguish between the two, since internal structures should be mapped in all of the macro-images while sub-halo lensing will only affect one of the images. Taking the above into account, one should be able to use this technique to put constraints on dark matter sub-halos predicted by simulations in the near future.  \\ack TR acknowledges support from the HEAC Centre funded by the Swedish Research Council. EZ acknowledges research grants from the Swedish Research Council, the Royal Swedish Academy of Sciences and the Academy of Finland."
    },
    "0809/0809.1772_arXiv.txt": {
        "abstract": "We present Giant Metrewave Radio Telescope (GMRT) observations for three (viz., DDO~68, SDSS~J2104--0035 and UGC~772) of the six most metal-deficient actively star-forming galaxies known. Although there is a debate as to whether these galaxies are undergoing their first episode of star formation or not, they are `young' in the sense that their ISM is chemically unevolved. In this regard, they are the nearest equivalents of young galaxies in the early Universe.  All three galaxies, that we have observed, have irregular \\HI\\ morphologies and kinematics, which we interpret as either due to tidal interaction with neighbouring galaxies, or the consequences of a recent merger. The remaining three of the six most metal-deficient galaxies are also known to have highly disturbed \\HI\\ distributions and are interacting. It is interesting because these galaxies were chosen solely on the basis of their metallicity and not for any particular signs of interaction. In this sense (i.e., their gas has not yet had time to settle into a regular disc), one could regard these extremely metal deficient (XMD) galaxies as `young'. The current star formation episode is likely to have been triggered by interaction/merger. It is also possible that the tidal interaction has lead to enhanced mixing with metal-poor gas in outer disc, and hence to a low gas-phase metallicity in the central star-forming regions.  We also try to determine the threshold gas-density for star-formation in our sample of galaxies, and find that in general these galaxies do not show a one-to-one correspondence between regions of high \\HI\\ column density and regions with current star formation. However, to the extent that one can define a threshold density, its value ($\\sim$10$^{21}$~atoms~cm$^{-2}$) is similar to that in galaxies with much higher metallicity. The highest column densities that we detect in regions far outside star-forming regions (i.e., a lower limit to the star-formation threshold) are $\\sim2\\times$10$^{21}$~atoms~cm$^{-2}$. ",
        "introduction": "\\label{sec:intro}  Blue Compact Galaxies (BCGs) were first identified by \\cite{zwicky1965} as unusual members of the extragalactic zoo. A subset of them were later recognized to be nearby low-mass galaxies. \\cite{sargent1970} found that the giant \\HII\\ regions in BCGs have low metal-content and blue colours, typical of young ($<$10 Myr) stellar populations. As luck would have it, the first well studied BCGs, viz., I~Zw~18 and II~Zw~40 turned out to be fairly unrepresentative of the BCG class. Even after 30 years of study of BCGs, I~Zw~18 continued to hold the record for the lowest metallicity, while II~Zw~40 is a galaxy pair in an advanced state of merger, which in general is not characteristic of BCGs. Partly because of the extreme properties of these two BCGs, and also because of the other unusual properties of the BCG class as a whole, it was hypothesized that some BCGs, including I~Zw~18, may be undergoing their first burst of star formation. However, this was quickly recognized as being unlikely, since it was found that the outer parts of many BCGs have red colours, typical of a normal old stellar population.  \\begin{figure*} \\includegraphics[width=6.5cm]{lowresddo68HI.eps} \\includegraphics[width=6.5cm]{ddo68lowresvel.eps} \\caption{(a.) The integrated \\HI\\ emission map of DDO~68 (in contours), at an angular resolution of $\\sim$39~$\\times$~35~arcsec$^{2}$, is overlaid on $B$-band $DSS$-II optical image (grey-scale). The contours are at \\HI\\ column densities of 7.8, 18.3, 42.8, 99.0 and 233.9 ~$\\times$~10$^{19}$~\\atoms. The grey-scale is in arbitrary units. (b.) The \\HI\\ velocity field at same resolution (contours) overlaid on same grey-scale image as in (a). The contours are from 462 to 541 \\kms, (from south to north) placed at regular intervals of 6.6 \\kms.} \\label{fig:lowresddo68HI} \\end{figure*}  None the less the idea that a small fraction of BCGs, particularly those with very low metallicity (i.e., the extremely Metal-Deficient (XMD) galaxies, with 12~+~log(O/H)~$<$~7.65, or Z$<$Z$_{\\odot}$/10; see \\cite{kunth2000} for a review), could be undergoing their first burst of star formation remained popular. In particular, deep photometry of the unresolved stellar population in outer parts of several XMD BCGs shows no evidence for an old stellar population, e.g., I~Zw~18 (\\cite{papaderos2002}), SBS~0335--052~E, W (\\cite{papaderos1998}; \\cite{pustilnik2004}, but see also \\cite{ostlink2001}), Tol~65 (\\cite{papaderos1999}) and DDO~68 (Pustilnik, Kniazev \\& Pramskij 2005;  Pustilnik, Tepliakova \\& Kniazev 2007). However, resolved CMD are essential to convincingly establish that the current episode of star formation is the first one. The case of I~Zw~18 has been particularly tortuous. From $HST$ WFPC2 (\\cite{aloisi99}) and NICMOS (\\cite{ostlin2000}) observations of this galaxy, it was argued that I~Zw~18 contains AGB stars older than 1~Gyr and is hence not undergoing its first episode of star formation. However using deeper HST ACS observations, \\cite{izotov2004} concluded that I~Zw~18 lacked RGB stars and was a bona fide young galaxy. Reanalysis of the same data by other groups indicated that RGB stars were, in fact, present (\\cite{momany2005}; \\cite{tosi2006}). More recently, \\cite{aloisi2007} analysed deeper {\\it HST} images of I~Zw~18 and convincingly established the presence of an RGB population, and hence that I~Zw~18 has been forming stars for $\\gtrsim$~2~Gyr. One more XMD BCG, SBS~1415+437, which was initially claimed to be a",
        "conclusions": "\\label{sec:dis}  In Table~\\ref{tab:mainpar}, we summarise the main observed and derived parameters of the studied galaxies. Their $B$-band luminosities all fall within a factor of 2.5, with the average M$_{\\rm B} \\sim$--14.3, while their M(\\HI) and M(\\HI)/L$_{\\rm B}$ ratios differ by less a factor of 1.5 and 1.7, respectively. In this respect, they form a much more homogeneous sample than the sample of 22 XMD galaxies studied by Pustilnik \\& Martin (2007). Of course a sample of 3 is too small to draw any meaningful statistical conclusion, this homogeneity is suggestive of similar evolutionary path-ways for all three galaxies. All of the three galaxies that we have observed also have peculiar \\HI\\ morphology and kinematics. UGC~772 is located in a loose group, and has easily identifiable neighbours which could be the cause of the observed disturbances. DDO~68, on the other hand, is located in a low-density region, in the periphery of the Lynx-Cancer void (Pustilnik et al. 2005). The Hubble flow around DDO~68 is very quiet. The radial velocities of all six dwarf galaxies within 500~kpc of it differ by less than 40~\\kms. The nearest neighbour of DDO~68, the dwarf galaxy UGC~5427, is located at a projected distance of $\\sim$200~kpc and a radial velocity separation of 5~\\kms. UGC~5427 also has a somewhat disturbed appearance with a chain of bright knots in a spiral-arm like feature, which makes it tempting to suggest that these two galaxies are tidally interacting. However, given its relatively large projected separation, it seems unlikely that tidal interaction with UGC~5427 is the cause of DDO~68's disturbed morphology. For example, if we assume that their true spatial separation and relative velocity are 200~kpc and $\\sim$50~\\kms, respectively, then their close encounter could occur $\\sim$4~Gyr ago. This is large compared to the ages ($\\sim$0.01--1~Gyr (Pustilnik et al. 2005; Pustilnik et al. 2007)) of the star-formation episodes in DDO~68.  If interaction with UGC~5427 is not the cause of DDO~68's disturbed morphology, then a remaining possibility is that DDO~68 represents the late-stage merger of two gas-rich progenitors, in which the gas is now settling down to form a disc. In this context, it is interesting to note that numerical simulations (e.g., \\cite{springel}) show that a disc galaxy can form in merger of extremely gas-rich progenitors. Mergers are not very rare among BCGs. E.g., \\cite{ostlin2001} also interpret, from the disturbed optical morphology and H$\\alpha$ kinematics of their sample of BCGs, that they are most likely the results of mergers. Similarly, \\cite{pustilnik2001BCG} assign $\\sim$15~per~cent of their sample BCGs as having a merger-like morphology.  We use a very approximate calculation below to show that the ages of the oldest stellar populations in DDO~68 ($\\sim$1~Gyr), as derived by Pustilnik et al. (2007), match with the time of first encounter between its progenitors. Figure~11 of \\cite{springeletal} indicates that `short' tails, like those of DDO68, are seen during later phases of merger (in comparison to the extended ones produced during the first encounter). Their ages are comparable to the rotation period of merging components and one-third of the time since the first encounter. Assuming that the velocity of tail relative to the systemic velocity of DDO~68 is $\\sim$50~\\kms and its length is $\\sim$5~kpc,  age of tail is $\\sim$200~Myr, which corresponds to a time since the first encounter of $\\sim$0.6~Gyr.  We note, however, that although the morphology and large-scale kinematics of DDO~68 suggest that it could be an ongoing merger, its large gas-mass fraction, and the lack of evidence for an enhanced velocity dispersion are in conflict with the predictions of numerical models of mergers (Springel et al. (2005) and Elmegreen et al. (1993), respectively). In the Springel et al. (2005) models, the gas-fraction of the merger product is (f$_{\\rm gas} \\sim$0.2--0.3), while DDO~68 is gas-rich ($M_{HI}/L_{B} \\sim$2.9 (Table~\\ref{tab:mainpar}); f$_{gas} \\sim$0.95 (Pustilnik et al. 2007)). These differences may be, in part, due to the fact that the masses of the putative progenitors of DDO~68 are roughly two orders of magnitude smaller than those of the model galaxies in the Elmegreen et al. (1993) and Springel et al. (2005) simulations, and in part, due to limitations in the star formation and feed back recipes used in the simulations. Given the importance of gas-rich low-mass mergers in the early universe, a detailed numerical model of DDO~68 would be particularly interesting.  J2104--0035 is also relatively isolated, and the only other galaxy within 500~kpc and 500~\\kms\\ from it is the compact dwarf emission-line galaxy SDSS~J210347.23--004949.7 (M$_{\\rm B}$~=~--14.49), which is at a projected distance of 22.3~arcmin ($\\sim$130~kpc for our assumed distance of 20~Mpc) and has, within cited error bars of $\\sim$5~\\kms, the same redshift. This is the only likely candidate for having caused the observed disturbance in J2104--0035. Alternatively we conjecture that J2104--0035 may be a case of an advanced merger.  It is interesting to note that all three of our sample galaxies have disturbed morphology and kinematics. In this regard, they closely resemble the three lowest metallicity BCGs, viz., I~Zw~18 and SBS~0035--052~E,~W. \\cite{pustilnik2001} show that SBS~0035--052 consists of an interacting pair of dwarf galaxies with a common \\HI\\ envelope, while \\cite{zee98IZw} show that I~Zw~18 has a complex \\HI\\ distribution, which they describe as a `fragmenting \\HI\\ cloud in the early stages of galaxy formation'. Thus, although tidal interaction was not specifically selected for when searching for XMD BCGs, all six of the most metal-deficient galaxies  have highly disturbed morphologies and kinematics. Indeed this seems to be a common feature for XMD BCGs (e.g., \\cite{chengalur95}; \\cite{chengalur2006}; Ekta et al. 2006).  Our, and past, results suggest that interactions/mergers could play an important role in triggering star formation in XMD galaxies. There are two factors that we can speculatively suggest as being responsible for this. The first is that intense SF episodes may occur only due to strong tidal disturbances (particularly, if the gas discs are less susceptible to internal instabilities, see the discussion below). The second is that tidal interaction could lead to a more efficient large-scale mixing of the gas. This brings lower-metallicity gas from the  outer parts of the galaxy to its centre, resulting in a lower metallicity in the central star-forming regions of the gas distribution. In principle, one could check the validity of this hypothesis by looking for statistical differences in the metallicity of BCGs with and without tidally induced star formation, or by detailed modelling of gas-rich mergers (e.g., Bekki 2008).  Regardless of how the SF is triggered, as far as the relationship between the gas distribution and regions of current SF is concerned, XMD BCGs do not appear to be qualitatively different from other late-type dwarf galaxies with low metallicities. For a sample of dwarf irregular galaxies (with metallicities mostly in XMD regime), \\cite{begum2006} find, much as we discuss above, that although SF generally occurs in regions of relatively high \\HI\\ column density, there is no one-to-one relation between gas at high column density and H$\\alpha$ emission. This lack of correspondence could be due to several effects, for example, ionization of the gas by $UV$ radiation from \\HII\\ regions could reduce \\HI\\ column density in regions with strong H$\\alpha$ flux. Further, since the massive star formation that the H$\\alpha$ emission traces are short-lived, stochasticity in SF would also reduce the correlation between the instantaneous observed H$\\alpha$ flux and the \\HI\\ column density. It is also worth noting that, to the extent that one can determine a threshold density for SF, there does not seem to be any significant difference between XMD galaxies and the more metal-rich dwarfs in the Taylor et al. (1994) sample.  On the other hand, in the model of \\cite{schaye2004} a lower limit to the threshold column density can be derived from the largest column densities outside the sites of current SF. In our galaxies this gas-density (corrected for inclination and with account of He) is $\\sim$2$\\times$10$^{21}$~atoms~cm$^{-2}$ for DDO~68 and UGC~772, and $\\sim$1.2--2.0$\\times$10$^{21}$~atoms~cm$^{-2}$ for J2104--0035 (as inclination of corresponding region is uncertain, gas-density may lie in the given range). For a proper comparison with more metal-rich galaxies, however, one would need to compare the column densities at a similar linear resolution.  Finally, we note that the baryonic masses of the XMD galaxies in our sample are all sufficiently large (greater than several 10$^8$~M$_{\\odot}$), such that current star burst is not strong enough to cause a significant loss of metals through escape of metal-enriched gas (\\cite{fragile2004}). The low ISM metallicity of these galaxies could be due to a generally subdued star-formation rate in the past (in case of LSB galaxy progenitor). For those XMDs for which no tracers of old stellar population have so far been found in outer parts, the option of this interaction inducing the first SF episode remains a possibility."
    },
    "0809/0809.1882_arXiv.txt": {
        "abstract": "We consider the possibility that masses and gravitational potentials of galaxy cluster, estimated at X-ray wavelengths, could be explained without assuming huge amounts of dark matter, but in the context of $f(R)$-gravity. Specifically, we take into account the weak field limit  of such theories and show that the corrected gravitational potential allows to  estimate the total mass of a sample of $12$  clusters of galaxies. Results show that such a gravitational potential provides a fair fit to the mass of visible matter (i.e. gas + stars) estimated by X-ray observations, without the need of additional dark matter while the size of the clusters, as already observed at different scale for galaxies, strictly depends on the interaction lengths of the corrections to the Newtonian potential. ",
        "introduction": "Since the pioneering work by \\cite{Zwi33}, the problems of high mass-to-light ratios of galaxy clusters and of the rotation curves of spiral galaxies  have been faced by asking for huge amounts of unseen matter in the framework of Newtonian theory of gravity. It is interesting to stress the fact that Zwicky addressed such an issue dealing with {\\it missing matter} and not with {\\it dark matter}.  Later on, several versions of the the cold dark matter model (CDM) have been built starting from the assumption that a large amount of non-baryonic matter (i.e. matter non-interacting with the electromagnetic radiation) could account for the observations in the framework of the standard Newtonian dynamics. Besides, by adding a further ingredient, the cosmological constant $\\Lambda$ \\cite{CarLam,Sahni}, such a model (now $\\Lambda$CDM) has become the new cosmological  paradigm usually called the {\\it concordance model}.  In fact,  high quality data coming from the measurements of cluster properties as the mass, the correlation function and the evolution with redshift of their abundance \\cite{eke98,vnl02,bach03,bb03}, the Hubble diagram of Type Ia Supernovae \\cite{Riess04,ast05,clo05}, the optical surveys of large scale structure \\cite{pope04,cole05,eis05}, the anisotropies of the cosmic microwave background \\cite{Boom,WMAP}, the cosmic shear measured from weak lensing surveys \\cite{vW01,refr03} and the Lyman\\,-\\,$\\alpha$ forest absorption \\cite{chd99,mcd04} are evidences toward a spatially flat universe with a subcritical matter content and undergoing a phase of accelerated expansion. Interpreting all this information in a self-consistent model is the main task of modern cosmology and $\\Lambda$CDM model provides a good fit to the most of the data \\cite{Teg03,Sel04,sanch05} giving a reliable snapshot of the today observed universe.  Nevertheless, it is affected by serious theoretical shortcomings that have motivated the search for alternative candidates generically referred to as {\\it dark energy} or {\\it quintessence}. Such models  range from  scalar fields rolling down self interaction potentials to phantom fields, from phenomenological unified models of dark energy and dark matter to alternative gravity theories  \\cite{PB03,Pad03,tsu1}.  Essentially, dark energy (or any alternative component) has to act as a negative pressure fluid which gives rise to an overall acceleration of the Hubble fluid. Despite of the clear mechanisms generating the observed cosmological dynamics, the nature and the fundamental properties of  dark energy remain essentially unknown notwithstanding the great theoretical efforts made up to now.  The situation for dark matter is similar: its clustering and distribution properties are fairly well known at every scale but its nature is unknown, up to now, at a fundamental level.  On the other hand, the need of  unknown components as dark energy and dark matter could be considered nothing else but as a signal of the breakdown of Einstein General Relativity at astrophysical (galactic and extragalactic) and cosmological scales.  In this context, Extended Theories of Gravity could be, in principle, an interesting alternative to explain cosmic acceleration without any dark energy and large scale structure without any dark matter. In their simplest version, the Ricci curvature scalar $R$, linear in the Hilbert-Einstein action, could be replaced by a generic function $f(R)$ whose true form could be \"reconstructed\" by the data. In fact, there is no a priori reason to consider the gravitational Lagrangian linear in the Ricci scalar while observations and experiments could contribute to define and constrain the \"true\" theory of gravity. For a discussion on this topic, see \\cite{capfra,odi,faraoni}.  In a cosmological framework, such an approach lead to modified Friedmann equations that can be formally written in the usual form by defining an effective  {\\it curvature fluid} giving rise to a negative pressure which drives the cosmic acceleration \\cite{capozzcurv,noirev,cdtt,capfra}. Also referred to as $f(R)$ theories, this approach is recently  extensively studied both from the theoretical  and the observational point of view (see e.g.  \\cite{noiijmpd,ccf,borowiec}). Moreover, this same approach has been also proposed, in early cosmology, as a mechanism giving rise to an inflationary era without the need of any inflaton field \\cite{Star80}.  All this amount of work has been, essentially,  concentrated on the cosmological applications  and have convincingly demonstrated that extended theories of gravity (in particular $f(R)$ gravity) are indeed able to explain the cosmic speed up and fairly fit the available data-sets and, hence, represents a viable alternative to the dark energy models \\cite{cct}.  Changing the gravity sector could have consequences not only at cosmological scales, but also at the galactic and cluster scales so that it is mandatory to investigate the low energy limit of such theories. A strong debate is  open with different  results arguing in favor \\cite{dick,sotiriou,cembranos,navarro,allrugg,ppnantro} or against \\cite{dolgov,chiba,olmo} such models. It is worth noting that, as a general result, higher order theories of gravity cause the gravitational potential to deviate from its Newtonian $1/r$ scaling \\cite{b10,hj,hjrev,cb05,sobouti,Mendoza} even if such deviations may  be  vanishing.  In  \\cite{CapCardTro07},  the Newtonian limit of power law $f(R) = f_0 R^n$ theories has been investigated, assuming that the metric in the low energy limit ($\\Phi/c^2 << 1$) may be taken as Schwarzschild\\,-\\,like. It turns out that a power law term $(r/r_c)^{\\beta}$ has to be added to the Newtonian $1/r$ term in order to get the correct gravitational potential. While the parameter $\\beta$ may be expressed analytically as a function of the slope $n$ of the $f(R)$ theory, $r_c$ sets the scale where the correction term starts being significant. A particular range of values of $n$  has been investigated so that the corrective term is an increasing function of the radius $r$ thus causing an increase of the rotation curve with respect to the Newtonian one and offering the possibility to fit the galaxy rotation curves without the need of further dark matter components.  A set of low surface brightness (LSB) galaxies with extended and well measured rotation curves has been considered \\cite{dbb02,db05}. These systems are supposed to be dark matter dominated, and successfully fitting data without dark matter is a strong evidence in favor of the approach (see also \\cite{Frigerio} for an independent analysis using another sample of galaxies). Combined with the hints coming from the cosmological applications, one should have, in principle, the possibility to address both the dark energy and dark matter problems resorting to the same well motivated fundamental theory \\cite{prl,Koivisto,lobo}. Nevertheless, the simple power law $f(R)$ gravity is nothing else but a toy-model which fail if one tries to achieve a comprehensive model for all the cosmological dynamics, ranging from the early universe, to the large scale structure up to the late accelerated era \\cite{prl,Koivisto}.  A fundamental issue is related to  clusters and  superclusters of galaxies. Such structures, essentially, rule the large scale structure, and are the intermediate step between galaxies and cosmology. As the galaxies, they appear dark-matter dominated but the distribution  of dark matter component seems clustered and organized in a very different way with respect to galaxies. It seems that dark matter is ruled by the scale and also its fundamental nature could depend on the scale. For a comprehensive review see \\cite{Bah96}.  In the philosophy of Extended Theories of Gravity, the issue is now to reconstruct the mass profile of clusters {\\it without} dark matter, i.e. to find out corrections to the Newton potential producing the same dynamics as dark matter  but starting from a well motivated theory. This is the goal of this paper.  As we will see, the problem is very different with respect to that of galaxies and we  need different  corrections to the gravitational potential in order to consistently fit the cluster mass profiles.  The paper is organized as follows. In Section~\\ref{sec:ext_theories}, also referring to other results, we discuss the weak field limit of $f(R)$-gravity showing that an $f(R)$ theory is, in principle, quite fair to address the cluster problem \\cite{lobo}. In Section~\\ref{sec:ext_systems}, starting from the corrected potential previously derived for a point-like masses,  we generalize the results  to  extended spherically symmetric systems  as required for well shaped cluster models of the Bautz and Morgan classification \\cite{Bau70}. Section~\\ref{sec:mass_profiles} is devoted to the description of the general properties of galaxy clusters. In Section~\\ref{sec:sample}, the sample of galaxy clusters which we are going to fit is discussed in details. Results are presented in Section~\\ref{sec:results}, while Section~\\ref{sec:discussion} is devoted to the discussion and the conclusions. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have investigated the possibility that the high observational   mass-to-light ratio of galaxy clusters  could be addressed by $f(R)$- gravity  without assuming huge amounts of dark matter. We point out  that this proposal comes out from the fact that, up to now, no definitive candidate for dark-matter has been observed at fundamental level and then  alternative solutions to the problem should be viable. Furthermore, several results in $f(R)$-gravity seem to confirm that valid alternatives to $\\Lambda$CDM can be achieved in cosmology. Besides, as discussed in the  Introduction, the rotation curves of spiral galaxies can be explained in the weak field limit of $f(R)$-gravity.  Results of our analysis go in this direction.  We have chosen a sample of  relaxed galaxy clusters  for which accurate spectroscopic temperature measurements and gas mass profiles are available. For the sake of simplicity, and considered the sample at our disposal, every cluster has been modelled as a self bound gravitational system with spherical symmetry and in hydrostatic equilibrium.  The mass distribution has been described by a corrected  gravitational  potential obtained from a generic analytic $f(R)$-theory. In fact, as soon as $f(R)\\neq R$, Yukawa-like exponential corrections emerge in the weak field limit while the standard Newtonian potential is recovered only for $f(R)=R$, the Hilbert-Einstein theory.  Our goal has been to analyze if the dark-matter content of clusters can be addressed by these  correction potential terms. As discussed in detail in the previous section and how it is possible to see by a rapid inspection of Figs.~\\ref{fig:plot_fin_1}-~\\ref{fig:plot_fin_12}, all the clusters of the sample are consistent with the proposed model at $1\\sigma$ confidence level.  This shows, at least \\textit{qualitatively},  that the high mass-to-light ratio of clusters can be explained by using a modified gravitational potential. The good agreement is achieved on distance scales starting from $150$ kpc up to $1000$ kpc. The differences observed at smaller scales can be ascribed  to non-gravitational phenomena, such as cooling flows, or to the fact that the gas mass is not a good tracer at this scales. The remarkable result is that we have obtained a consistent agreement with data only using the corrected gravitational potential in a large range of radii. In order to put in evidence this trend, we have plotted the baryonic mass vs radii considering, for each cluster,  the scale where the trend is clearly evident.  In our knowledge, the fact that $f(R)$-gravity could work at these scales  has been only supposed but never achieved by a direct fitting with data (see \\cite{lobo} for a review). Starting from the series coefficients $a_{1}$ and $a_{2}$, it is possible to state that, at cluster scales, two characteristic sizes emerge from the weak field limit of the theory.  However, at smaller scales, e.g. Solar System scales, standard Newtonian gravity has to be dominant in agreement with observations.  In conclusion, if our considerations are right, gravitational interaction depends on the scale and the {\\it infrared limit} is led by the series coefficient of the considered effective gravitational Lagrangian. Roughly speaking, we expect that starting from cluster scale to galaxy scale, and then down to smaller  scales as Solar System or Earth, the terms of the series lead the clustering of self-gravitating systems beside other non-gravitational phenomena. In our case, the Newtonian limit is recovered for $a_{1} \\rightarrow 3/4$ and $L(a_{1},a_{2})\\gg r$ at small scales and for $L(a_{1},a_{2})\\ll r$ at large scales. In the first case, the gravitational coupling has to be redifined, in the second $G_{\\infty}\\simeq G$. In these limits, the linear Ricci term is dominant in the gravitational Lagrangian and the Newtonian gravity is restored \\cite{Qua91}. Reversing the argument, this could be the starting point to achieve a theory  capable of explaining the strong segregation in masses and sizes of gravitationally-bound systems."
    },
    "0809/0809.3139_arXiv.txt": {
        "abstract": "{The connection of cluster mergers with the presence of extended, diffuse radio sources in galaxy clusters is still debated. An interesting case is the rich, merging cluster Abell 520, containing a radio halo. A recent gravitational analysis has shown in this cluster the presence of a massive dark core suggested to be a possible problem for the current cold dark matter paradigm.}  {We aim to obtain new insights into the internal dynamics of Abell 520 analyzing velocities and positions of member galaxies.}{Our analysis is based on redshift data for 293 galaxies in the cluster field obtained combining new redshift data for 86 galaxies acquired at the TNG with data obtained by CNOC team and other few data from the literature.  We also use new photometric data obtained at the INT telescope. We combine galaxy velocities and positions to select 167 cluster members around $z\\sim0.201$.  We analyze the cluster structure using the weighted gap analysis, the KMM method, the Dressler-Shectman statistics and the analysis of the velocity dispersion profiles. We compare our results with those from X-ray, radio and gravitational lensing analyses.}{We compute a global line--of--sight (LOS) velocity dispersion of galaxies, $\\sigma_{\\rm v}=1066_{-61}^{+67}$ \\kss.  We detect the presence of a high velocity group (HVG) with a rest--frame relative LOS velocity of ${\\rm v_{\\rm rf}}\\sim 2000$ \\ks with respect to the main system (MS). Using two alternative cluster models we estimate a mass range $M(<1 \\hhh)=(4.0-9.6)$\\mquaa.  We also find that the MS shows evidence of subclumps along two preferred directions.  The main, complex structure ${\\cal NE}1$+${\\cal NE}2$ (with a velocity comparable to that of the MS) and the ${\\cal SW}$ structure (at ${\\rm v_{\\rm rf}}\\sim +1100$ \\kss) define the NE--SW direction, the same of the merger suggested by X--ray and radio data.  The ${\\cal E}$ and ${\\cal W}$ structures (at ${\\rm v_{\\rm rf}}\\sim -1150$ and ${\\rm v_{\\rm rf}}\\sim -300$ \\kss) define the E--W direction. Moreover, we find no dynamical trace of an important structure around the lensing dark core. Rather, the HVG and a minor MS group, having different velocities, are roughly centered in the same position of the lensing dark core, i.e. are somewhat aligned with the LOS.}{ We find that Abell 520 is definitely a very complex system.  Our results suggest that we are looking at a cluster forming at the crossing of three filaments of the large scale structure. The filament aligned with the LOS and projected onto the center of the forming cluster might explain the apparent massive dark core shown by gravitational lensing analysis.} ",
        "introduction": "\\label{intr}  Clusters of galaxies are by now recognized to be not simple relaxed structures, but rather they are evolving via merging processes in a hierarchical fashion from poor groups to rich clusters. Much progress has been made in recent years in the observations of the signatures of merging processes (see Feretti et al. \\cite{fer02b} for a general review). A recent aspect of these investigations is the possible connection of cluster mergers with the presence of extended, diffuse radio sources: halos and relics. Cluster mergers have been suggested to provide the large amount of energy necessary for electron reacceleration and magnetic field amplification (Feretti \\cite{fer99}; Feretti \\cite{fer02a}; Sarazin \\cite{sar02}). However, the question is still debated since the diffuse radio sources are quite uncommon and only recently we can study these phenomena on the basis of a sufficient statistics (few dozen clusters up to $z\\sim 0.3$, e.g., Giovannini et al. \\cite{gio99}; see also Giovannini \\& Feretti \\cite{gio02}; Feretti \\cite{fer05}).  Growing evidence of the connection between diffuse radio emission and cluster merging is based on X--ray data (e.g., B\\\"ohringer \\& Schuecker \\cite{boh02}; Buote \\cite{buo02}). Studies based on a large number of clusters have found a significant relation between the radio and the X--ray surface brightness (Govoni et al. \\cite{gov01a}, \\cite{gov01b}) and connections between the presence of radio--halos/relics and irregular and bimodal X--ray surface brightness distribution (Schuecker et al. \\cite{sch01}).  Optical data are a powerful way to investigate the presence and the dynamics of cluster mergers (e.g., Girardi \\& Biviano \\cite{gir02}), too.  The spatial and kinematical analysis of member galaxies allow us to detect and measure the amount of substructure, to identify and analyze possible pre--merging clumps or merger remnants.  This optical information is really complementary to X--ray information since galaxies and intra--cluster medium react on different time scales during a merger (see, e.g., numerical simulations by Roettiger et al. \\cite{roe97}). In this context we are conducting an intensive observational and data analysis program to study the internal dynamics of radio clusters by using member galaxies. Our program concerns both massive clusters, where diffuse radio emissions are more frequently found (e.g., Barrena et al. \\cite{bar07b} and refs. therein), and low--mass galaxy systems (Boschin et al. \\cite{bos08}\\footnote{please visit the web site of the DARC (Dynamical Analysis of Radio Clusters) project: http://adlibitum.oat.ts.astro.it/girardi/darc.}).  During our observational program we have conducted an intensive study of the massive cluster Abell 520 (hereafter A520). This cluster shows a radio halo, discovered by Giovannini et al. (\\cite{gio99}), having a low surface brightness with a clumpy structure  slightly elongated in the NE--SW direction (Govoni et al. \\cite{gov01b}; \\cite{gov04}, see Fig.~\\ref{figradio}).  A520, also known as MS 0451+02 in the EMSS catalog (Gioia et al. \\cite{gio90}), is a fairly rich, X--ray luminous, and hot cluster, with a galaxy population characterized by a high velocity dispersion: Abell richness class $=1$ (Abell et al. \\cite{abe89}), $L_\\mathrm{X}$(0.1--2.4 keV)=14.20$\\times 10^{44} \\ h_{50}^{-2}$ erg\\ s$^{-1}$ (Ebeling et al.~\\cite{ebe96}); $T_\\mathrm{X}= 7.1\\pm0.7$ keV (Chandra data, Govoni et al. \\cite{gov04}); $\\sigma_{\\rm v}=(988\\pm 76)$ \\ks (Carlberg et al. \\cite{car96}).  First hints about the young dynamical status of this cluster came from both X--ray and optical data (Le Fevre et al. \\cite{lef94}; Gioia \\& Luppino \\cite{gio94} and refs. therein). The complexity of its structure was confirmed by analyses of ROSAT and Chandra X--ray data (Govoni et al. \\cite{gov01b}; \\cite{gov04}). In particular, new unprecedent insights were recovered from deep Chandra observations by Markevitch et al. (\\cite{mar05}). They revealed a prominent bow shock indicating a cluster merger where a SW irregular structure consists of dense, cool pieces of a cluster core that has been broken up by ram pressure as it flew in from the NE direction (see Fig.~\\ref{figradio}).  The overall structure of the radio halo seems connected with the cluster merger and may even suggest two distinct components, a mushroom with a stem and a cap, where the main stem component goes across the cluster along the NE--SW direction and the cap ends at the bow shock (Govoni et al. \\cite{gov01b}; Markevitch et al. \\cite{mar05}).  The complex structure of A520 was also confirmed by gravitational lensing analysis of Dahle et al. (\\cite{dah02}), Mahdavi et al. (\\cite{mah07}, hereafter M07) and Okabe \\& Umetsu (\\cite{oka08}). Okabe \\& Umetsu (\\cite{oka08}, based on Subaru data) found a general good agreement between mass and galaxy luminosity distribution. However, the detailed study of M07 based on the same Subaru data and additional CFHT data pointed out a less clear situation. M07 found four very significant peaks in the lensing mass distribution. Among these, peaks No. 1, 2 and 4 correspond to peaks in the galaxy distribution and give usual values for the mass--to--light ratio.  Peak No. 3 corresponds to the central X-ray emission peak, but is largely devoid of galaxies. This peak is characterized by a very large mass--to--light value; thus to be referred as a ``massive dark core''. A region characterized by a somewhat low mass--to--light ratio exists, too (less significant peak No. 5). This displacement between galaxy and mass (i.e. dark matter, for the most part) remains very puzzling. In fact, galaxies and cold dark matter (CDM), being both treated as collisionless components, are expected to have similar behavior during a cluster merger.  If confirmed by better observations, this situation would be difficult to explain within the widely accepted CDM paradigm of cosmological structure formation (see M07 for further discussions).   \\begin{figure*} \\centering \\includegraphics[width=16cm]{f1.eps} \\caption{ Multiwavelength picture of A520 (North is at the top and East to the left). A smoothed Chandra 0.5--2 keV image (orange and yellow colors) of the central region of A520 (courtesy of M. Markevitch - Markevitch et al. \\cite{mar05}, X--ray point sources are removed) is superimposed to a $r^\\prime$--band image taken with the WFC camera of the INT. The contour levels of a VLA radio image at 1.4 GHz (courtesy of F. Govoni - Govoni et al. \\cite{gov01b}) are shown, too.  Main structures recovered by our analysis are highlighted (see Sect.~\\ref{dyn} for more details).  Label HVG indicates the center of the high velocity group having a relative LOS velocity of ${\\rm v_{\\rm rf}} \\sim 2000$ \\ks with respect to the main system (MS).  Blue circles and numbers highlight the positions of the five peaks in the lensing mass distribution found by M07. The size of the circles indicate the regions where we find evidence for an individual, dynamically important structure of the MS. The name of each structure is indicated by the label close the corresponding M07 peak number and the blue small square indicates the central, luminous galaxy (i.e. galaxies IDs 204, 170, 106 and 205 for ${\\cal NE}1$, ${\\cal NE}2$, ${\\cal SW}$ and ${\\cal E}$, respectively). Finally, the green square indicates a head--tail radiogalaxy.  } \\label{figradio} \\end{figure*}  As for the analysis of the internal dynamics based on member galaxies, Proust et al. (\\cite{pro00}) found some evidence of substructure using a sample of 21 galaxies, while the large data sample constructed by the Canadian Network for Observational Cosmology (hereafter CNOC) team (Carlberg et al.  \\cite{car96}; Yee et al. \\cite{yee96}) is still not exploited a part from few individual galaxies in M07.  Recently, we have carried out spectroscopic observations at the TNG telescope giving new redshift data for 86 galaxies in the field of A520, as well as photometric observations at the INT telescope. Our present analysis is based on these optical data as well as on the large data sample obtained by CNOC.  This paper is organized as follows.  We present our new optical data in Sect.~2 and the complete redshift catalog with the addition of CNOC and a few other data in Sect.~3. We present our results about global properties and substructure in Sect.~4.  We furtherly analyze and discuss the dynamical status of A520 in Sect.~5. We draw our conclusions in Sect.~6.  Unless otherwise stated, we give errors at the 68\\% confidence level (hereafter c.l.). Throughout this paper, we use $H_0=70\\ h_{70}$  km s$^{-1}$ Mpc$^{-1}$ in a flat cosmology with $\\Omega_0=0.3$ and $\\Omega_{\\Lambda}=0.7$. In the adopted cosmology, 1\\arcm corresponds to $\\sim 199$ \\kpc at the cluster redshift.  \\begin{figure*} \\centering \\includegraphics[width=12cm]{f2.eps} \\caption{INT $r^\\prime$--band image of A520 (West at the top and North to the left). Circles and boxes indicate cluster members and non--member galaxies, respectively (see Table~\\ref{catalogA520}).} \\label{figottico} \\end{figure*} ",
        "conclusions": ""
    },
    "0809/0809.4005_arXiv.txt": {
        "abstract": "We study energy dissipation and heating by supersonic MHD turbulence in molecular clouds using Athena, a new higher-order Godunov code.  We analyze the dependence of the saturation amplitude, energy dissipation characteristics, power spectra, sonic scaling, and indicators of intermittency in the turbulence on factors such as the magnetic field strength, driving scale, energy injection rate, and numerical resolution.  While convergence in the energies is reached at moderate resolutions, we find that the power spectra require much higher resolutions that are difficult to obtain.  In a $1024^3$ hydro run, we find a power law relationship between the velocity dispersion and the spatial scale on which it is measured, while for an MHD run at the same resolution we find no such power law.  The time-variability and temperature intermittency in the turbulence both show a dependence on the driving scale, indicating that numerically driving turbulence by an arbitrary mechanism may not allow a realistic representation of these properties.  We also note similar features in the power spectrum of the compressive component of velocity for supersonic MHD turbulence as in the velocity spectrum of an initially-spherical MHD blast wave, implying that the power law form does not rule out shocks, rather than a turbulent cascade, playing a significant role in the regulation of energy transfer between spatial scales. ",
        "introduction": "Observed non-thermal line widths in molecular clouds (MCs), where all star formation in the Galaxy takes place, point to the presence of supersonic turbulence in such regions (Falgarone \\& Philips 1990). The properties of the turbulent medium, such as Mach number and magnetic field strength, may determine the products of the star formation process.  As an important source of heating within molecular clouds (Stone et al.~1998, hereafter S98), turbulent energy dissipation may also play a role.  Much effort has been directed toward numerically simulating turbulent media in order to better understand the link between turbulence and star formation (see Elmegreen \\& Scalo 2004, MacLow \\& Klessen 2004, and McKee \\& Ostriker 2007 and references therein).  As we have not yet identified the turbulent driving mechanism, there remain many unanswered questions about the evolution of molecular clouds.  Is the turbulence periodically re-energized, or does it simply decay away?  How much impact do magnetic fields have on the properties of the turbulence?  Crutcher (1999), using observations of Zeeman splitting, found magnetic fields in some clouds strong enough that one cannot safely neglect their effects.  Although it has been shown that magnetic fields do not appreciably lengthen the turbulent decay time (S98; Mac Low 1999), they do create anisotropy within the clouds (e.g.~Vestuto et al.~2003, hereafter V03; Esquivel et al.~2003), which may have important observational and evolutionary consequences.  For example, molecular clouds are often observed to be filamentary (e.g.~Mizuno et al.~1995; Churchwell et al.~2004).  In this paper, we will investigate the energy dissipation properties of supersonic hydrodynamic and MHD turbulence with a variety of magnetic field strengths using Athena, a new higher-order Godunov code.  An important goal of this study is to investigate the effect of the assumed driving mechanism on the properties of the resulting turbulence, such as power spectra and intermittency indicators.  Our analysis utilizes data from high-resolution numerical simulations with twenty-five different parameter sets.  In recent years, a variety of results have been reported on the properties of supersonic MHD turbulence, including energetics (e.g.~S98; Mac Low 1999; Ostriker et al.~2001), power spectra (e.g.~Cho \\& Lazarian 2003, 2005; V03; Padoan et al.~2007, hereafter P07), and probability distribution functions (e.g.~Padoan et al.~1997; Passot \\& Vazquez-Semadeni 1998; Ostriker et al.~2001; Kowal et al.~2007; Kritsuk et al.~2007, hereafter K07). Where possible, we identify differences in our methods and results as compared to those of other groups.  In a separate letter (Lemaster \\& Stone 2008, hereafter Paper I), we have reported the results of an investigation of the variation of the probability distribution function (PDF) of density with Mach number.  Our primary result in that paper was that the intermittent behavior of turbulence could be responsible for the large cloud-to-cloud variation in the observed star formation rate per solar mass.  The primary tool available to investigate the properties of highly nonlinear, supersonic MHD turbulence is direct numerical simulation. To date, most results have been computed using a few numerical algorithms, such as ZEUS (e.g.~S98; V03; Mac Low 1999; Ostriker et al.~1999; Ostriker et al.~2001), the PENCIL code (e.g.~Haugen \\& Brandenburg 2004), the Stagger code (e.g.~P07), and ENO methods (e.g.~Cho \\& Lazarian 2003).  There has been concern expressed in the literature that previous results may be strongly affected by numerical dissipation.  Thus, it is worth re-examining the problem with more accurate methods.  This study represents one of the first applications of higher-order Godunov methods to the study of supersonic MHD turbulence.  Without knowing the driving mechanism, we are left with many possible methods of generating turbulence in simulations.  The hope is that the choice of method for the simulated driving will have an insignificant effect on the results.  Federrath, Klessen, \\& Schmidt (in prep.), however, have shown that compressive and solenoidal forcing produce dramatically different turbulent states. For some diagnostics, such as intermittency, the time-variability of the turbulent state is critical.  Even if an array of driving methods produce the same mean state, do the instantaneous states have the same distribution about the mean?  We will investigate the dependence on driving scale of various diagnostics, given our particular driving method, which is very similar to that employed by, e.g., S98 and V03.  For power spectra of simulated turbulence to be valuable, the resolution needs to be high enough that the driving and numerical dissipation scales are well separated, allowing the inertial range to be studied.  When magnetic fields are taken into account, simulating turbulence at these resolutions can be prohibitively expensive. Another goal of this study is to investigate whether properties of the power spectra and other diagnostics are converged at the numerical resolutions feasible at the moment, up to $1024^3$.  We begin in \\S\\ref{sec:methods} by describing our numerical methods in detail.  We proceed, in \\S\\ref{sec:results}, to present our results. \\S\\ref{sec:satresults} includes a convergence study and Mach number scaling study of saturation amplitudes, \\S\\ref{sec:specresults} includes a power spectrum analysis, \\S\\ref{sec:sonicresults} analyzes the sonic scale in $1024^3$ runs, and \\S\\ref{sec:mittresults} considers time-variability and temperature intermittency in the turbulence.  Finally, we summarize our results in \\S\\ref{sec:concl} and discuss the implications. ",
        "conclusions": "\\label{sec:concl}   The saturation energies and dissipation timescales we find support the conclusion of S98 that supersonic turbulence dissipates rapidly even in the presence of magnetic fields.  The results of our Godunov code agree with ZEUS, an operator-split method that relies on artificial viscosity to capture shocks, but reach convergence at slightly lower resolutions.  At $512^3$ the difference in the total energy in fluctuations between Athena and ZEUS for strong-field MHD is very small, indicating that these two codes converge to similar turbulent states at sufficiently high resolution.  The convergence resolution of our simulations depends both on the driving scale and the presence of a magnetic field.  While hydro turbulence driven at \\pk{4} converges by $64^3$, strong-field MHD turbulence driven at \\pk{8} does not converge until $512^3$. Although very high resolutions are needed to capture the inertial range of the turbulence, lower resolutions are often adequate for studying the energy dissipation characteristics of the turbulence as well as other volume-integrated quantities.  At high Mach number, the ratio of the dissipation timescale to the flow crossing time at the driving scale increases with increasing magnetic field strength; however, it does not exceed unity even for strong-field MHD turbulence.  This ratio is independent of driving scale.  The fractions of the kinetic energy in solenoidal and compressive modes are also independent of driving scale, with the compressive fraction generally being more than twice as high for hydrodynamic than for strong-field MHD turbulence.   Our spherically-integrated velocity power spectra for decaying, subsonic hydrodynamic turbulence show evidence of the bottleneck effect in the velocity power spectrum, consistent with the findings of Sytine et al.~(2000).  We find more power at high wavenumber in our driven, supersonic hydro and MHD turbulence simulations than was found by V03, but it is unclear, particularly in the MHD case, whether or not this is due to a bottleneck.  Resolutions exceeding $1024^3$ will be necessary to draw firm conclusions about the slope of the inertial range.  The cylindrically-averaged velocity power spectrum for driven MHD turbulence is very anisotropic; it has a slope that approximates that of the hydrodynamic case parallel to the magnetic field, while perpendicular to the field it is much more shallow.  The compressive component of velocity has an isotropic power spectrum, contrary to what was found in V03.  We find the compressive component of velocity in driven, supersonic MHD turbulence to have a power spectrum that is difficult to distinguish from the velocity spectrum of an initially-spherical MHD blast wave.  This calls into question the long-held assumption that supersonic turbulence power spectra result from an energy cascade facilitated by interactions local in Fourier space.  The analysis of structure functions may be useful in determining the source of the power law spectrum in supersonic turbulence, either a Fourier-space cascade as in incompressible turbulence, or an ensemble of shocks as in Burgers turbulence.   For hydrodynamic turbulence, we find a power law scaling of the velocity dispersion with spatial scale, $\\sigma(l) \\propto l^{0.58}$, for scales where the velocity dispersion is supersonic.  However, we find no such power law for strong-field MHD turbulence, where the sonic scale is not the only scale of interest.  In this case, we find that the velocity dispersion drops off more rapidly than a power law as one approaches smaller scales.   We see time-variability in the saturation energies comparable to that of the equivalent runs in S98 despite the impulsive driving employed therein.  Our time-variability increases when we apply our turbulent driving at larger scales, but remains lower than that shown in K07 even with \\pk{2}.  It is possible that the acceleration-based driving method of K07 is responsible for the difference.  At this large driving scale, the time-variability increases with Mach number for both hydro and strong-field MHD.  The method used to drive the turbulence appears to have a substantial impact on the resulting turbulent state.  Further investigation should be conducted to determine how much of the increase in time-variability with driving scale is due to the limited range of scales over which an inverse cascade can occur.  If the level of time-variability is determined to be a result of simulation setup and not representative of the physical system we are trying to simulate, better diagnostics or simulation methods need to be developed in order to quantify intermittency in turbulent media. Although we have very poor statistics in our temperature intermittency analysis, it appears that this, too, increases with driving scale. Considering the difference that the driving method makes on the results, it may be more realistic to study decaying turbulence instead of driving it arbitrarily.   For identical turbulent data cubes, we find that the post-processing methods chosen significantly influence the results.  The turbulent Mach number changes by $\\sim 4\\%$ depending on the method used in its computation.  The means of computing the power spectrum also has a strong influence on the power at wavenumbers overlapping with the inertial range.  When comparing results published by different groups, one should keep in mind that this, as well as the code, driving method, and initial conditions, can affect the quantities being compared.   Although the saturation energies and energy dissipation characteristics of the turbulence converge at resolutions within our current computational capabilities, the power spectra appear to require much higher resolutions to provide valuable information.  Also considering that turbulent power spectra can be approximated by non-turbulent phenomena, it would seem that, for the time being, our focus should be on other diagnostics.   In conducting these numerical simulations, we have made many simplifications in order to make the problem more tractable.  These assumptions, however, may prove to significantly impact the results. Future studies should consider non-ideal MHD in order to model low-ionization molecular clouds, a non-isothermal equation of state in order to study heating and cooling, and self-gravity in order to follow the collapse of the bound clumps that form in the turbulent medium."
    },
    "0809/0809.4252.txt": {
        "abstract": "IC 2144 is a small reflection nebula located in the zone of avoidance near the Galactic anticenter.  It has been investigated here largely on the basis of Keck/HIRES optical spectroscopy (R $\\approx$ 48,000) and a SpeX spectrogram of the near-IR (R = 2000) obtained at the NASA IRTF.  The only star in the nebula that is obvious in the optical or near-IR is the peculiar emission-line object MWC 778 ($V$ = 12.8), which resembles a T Tauri star in some respects.  What appear to be  F- or G-type absorption features are detectable in its optical region under the very complex emission line spectrum; their radial velocity agrees with the CO velocity of the larger cloud in which IC 2144 is embedded.  There are significant differences between the spectrum of the brightest area of the nebula and of MWC 778, the presumed illuminator, an issue discussed in some detail. The distance of IC 2144 is inferred to be about 1.0 kpc by reference to other star-forming regions in the vicinity.  The extinction is large, as demonstrated by $[$\\ion{Fe}{2}$]$ emission line ratios in the near-IR and by the strength of the diffuse interstellar band spectrum; a provisional value of $A_{V}$ of 3.0 mag was assumed.  The SED of MWC 778 rises steeply beyond about 1 $\\mu$m, with a slope characteristic of a Class I source.  Integration of the flux distribution leads to an IR luminosity of about 510 $L_{\\sun}$.  If MWC 778 is indeed a F- or G-type pre--main-sequence star several magnitudes above the ZAMS, a population of faint emission H$\\alpha$ stars would be expected in the vicinity.  Such a search, like other investigations that are recommended in this paper, has yet to be carried out. ",
        "introduction": "On conventional photographs such as the Palomar Sky Atlas, IC 2144 appears as a roughly rectangular object about 25\\arcsec $\\times$ 16\\arcsec \\ that might be mistaken for a small irregular galaxy like M82, or an \\ion{H}{2} region.  It is in fact a complex reflection nebulosity, apparently illuminated by an off-center star of $V = 12.8$, and embedded in a small, rather inconspicuous dark cloud.  That star\\footnote{$\\alpha$ = 5$^{h}$\\, 50$^{m}$\\, 13$^{s}$.5, $\\delta$ = +23$^{\\circ}$\\, 52$'$\\, 17$''$ (J2000); $l$ = 184$^{\\circ}$.9, $b$ = $-$1$^{\\circ}$.7 } was found by \\citet{mer49} to have H$\\alpha$ as well as lines of \\ion{Fe}{2} and [\\ion{Fe}{2}] in emission.  They named it MWC 778 and assigned type Bpe.  In what follows, we continue to distinguish between the star MWC 778 and the nebulosity IC 2144.  IC 2144 is reproduced in Figure 1 from an exposure obtained with the SuprimeCam instrument behind an H$\\alpha$ filter at the Subaru Telescope\\footnote{The Subaru Telescope is operated by the National Astronomical Observatory of Japan.} and kindly made available by Bo Reipurth. MWC 778 is located near the white spot in the overexposed nebulosity. The nebula appears to be much larger (about 120\\arcsec $\\times$ 100\\arcsec) and more complex than apparent on earlier images. Figure 2 is the same image plotted on a logarithmic intensity scale.  The internal structure is shown in Figures 3a, b, and c, which are $R$-band images of IC 2144, extracted from 2 s snapshots (obtained in poor seeing) of the slit plane of the Keck I HIRES spectrograph during one of the spectroscopic exposures to be described later.  Very bright nebulosity immediately surrounds the star and extends about 5\\arcsec\\ to the northwest (the image of MWC 778 itself is slightly elongated in the same direction), where it splits into two curving forks, best seen in Figure 3b, which is the same image as Figure 3a, on a logarithmic intensity scale.  At longer wavelengths the surface brightness of IC 2144 becomes so high and so structured as to confuse the 2MASS $J$, $H$, and $K_{s}$ images, which have a resolution of about 4\\arcsec.  Optically, the appearance of the interior of IC 2144 is defined by what appears to be a band of dust in the shape of a reversed question mark that begins north of MWC 778, curves as it crosses the bright background of the nebula to the east, and appears to unwind as it disappears against the faintly luminous background south of MWC 778. There is a faint (R magnitude about 17) star (arrowed in Fig. 3b) about 4\\arcsec\\ northwest of MWC 778, at the beginning of this feature, and another about 3  \\arcsec southwest. It is unknown whether they are related to the nebula,  or if the brighter of the two northeast forks is really the illuminated base of the dust band.  \\citet{allen74} noted the emission lines of H, \\ion{Fe}{2}, [\\ion{Fe}{2}] and [\\ion{S}{2}] in MWC 778, and performed the first near-IR photometry out to 18 $\\mu$m.  He commented that ``its almost linearly rising energy distribution [in the $\\lambda$F$_{\\lambda}$, $\\lambda$ plane] is unique.\"  The IRAS fluxes (it is the point source 05471+2351 and the Small Scale Structure Source X0547+238) reinforce that impression: they have not begun to decline as far as the 100 $\\mu$m point, although the IRAS catalog notes that both the 60 and 100 $\\mu$m fluxes are uncertain because of the bright background.  Published photometry of the star is collected in Table 1, but it is not always clear what aperture size was used, or if allowance was made for a contribution by the bright nebulosity.  IC 2144 is near the Galactic anticenter, so the LSR velocity of its associated CO (about +2 km s$^{-1}$: $\\S$ 2) is too small to establish a kinematic distance.  On the sky, there is no well-defined obscuring cloud nearby. Within a radius of 3$\\arcdeg$--4$\\arcdeg$ of IC 2144 there are two distant \\ion{H}{2} regions (S 242, at about 2 kpc, and S 243), two distant star clusters (NGC 2129 at 2 kpc and Be 21 at 5--6 kpc), and the bright B star 139 Tau (distant 0.9--1.1 kpc).  There is no reason to believe that IC 2144 is associated with any of these objects.  There are two star-forming regions in that general direction, but both are about 3$\\arcdeg$ distant.  To the northwest is the small cloud containing the HAeBe star RR Tau and a few T Tauri Stars (TTS); its distance has been estimated as 380 pc by \\citet{rost99}. The contrast between the star density projected upon that cloud and that of its background is large,  unlike the situation at IC 2144 where the contrast is small. This is to be expected if IC 2144 is considerably more distant.  To the south is the small cloud CB34, studied most recently by \\citet{khan02}.  Its distance is usually taken to be 1.5 kpc on the assumption that it is associated with the Gem OB1 association, which lies still farther east. On the basis of all these samples, we assume for the purposes of this paper that the distance of IC 2144 is 1.0 kpc, but emphasize the uncertainty of this estimate.  \\section {Optical Spectroscopy}  Spectrograms of the object were obtained on 2003 December 13 and 2004 November 21 with the HIRES spectrograph at the Keck I telescope.\\footnote{The W. M. Keck Observatory is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA.  The Observatory was made possible by the generous financial support of the W. M. Keck Foundation.  The spectrograph is described in some detail in http://www2.keck/hawaii.edu/inst/hires/.} The first series covers the region 4350--6740 \\AA\\ at a resolution of about 48,000 with slit dimensions of 0\\farcs86 $\\times$ 7$''$.  The seeing then was poor: star images have a FWHM of about 1\\farcs5.  The first exposure of that series was centered on MWC 778. The second crossed the brightest portion of the nebulous arc about 5$''$ northwest of the star; its precise location is shown in Figure 3c.  The 2004 exposure was also centered on MWC 778, but covered 4400--8800 \\AA\\ at slightly higher resolution.  Reductions were carried out with standard IRAF procedures.  The average S/N per pixel in the continuum of the two (reduced) exposures of the star averaged about 50 in the 5000 \\AA\\ region, 100 near 6200 \\AA, and 130 near 7700 \\AA.  The spectrum of MWC 778 is dominated by a strong ($W \\simeq 170$ \\AA\\ ) H$\\alpha$ emission, followed by many emission lines of \\ion{Fe}{2}, [\\ion{Fe}{2}], [\\ion{O}{1}], [\\ion{S}{2}],  and \\ion{Si}{2}, and a few of [\\ion{Ni}{2}] and [\\ion{Cr}{2}].  The profiles of H$\\alpha$ and H$\\beta$ are shown in the upper section of Figure 4, with velocities\\footnote{The radial velocities in this paper are heliocentric, unless noted otherwise.  The correction of heliocentric to LSR at the position of IC 2144 is $-$11.8 km s$^{-1}$. } of several features marked.  \\citet{viei03}, at a resolution of 9,000, also noted the off-center reversal in H$\\alpha$ and the presence of [\\ion{S}{2}] and [O\\, I]; they gave the type as ``B1?.\"  The spectrum between 3900 and 6700 \\AA\\ at moderate resolution was reproduced by \\citet{su06}.  Only emission lines are apparent. Their classification was ``BQ[\\, ].\"  The lines of $[$\\ion{S}{2}$]$ and $[$\\ion{N}{2}$]$ $\\lambda$6583 scatter in velocity between +14 and +17 km s$^{-1}$, in agreement with $[$\\ion{Fe}{2}$]$. However, both [\\ion{O}{1}] $\\lambda$6300 and $\\lambda$6363 are double, with peak velocities of +13 and +29 km s$^{-1}$.  That region (on the 2004 spectrum) is shown in Figure 5, where the splitting of the $[$\\ion{O}{1}] lines is apparent (see the expanded profile of $\\lambda$6300 in Fig. 5), as is the striking difference between the widths of permitted and forbidden lines.  Despite their duplicity, the FWHM of the combined $[$\\ion{O}{1}] lines is only about 60 km s$^{-1}$, as compared to 220 km s$^{-1}$ for \\ion{Si}{2} $\\lambda$6347.  The \\ion{Fe}{2} lines are double, with a splitting of about 72 km s$^{-1}$. This is most obvious in the weaker \\ion{Fe}{2} lines such as $\\lambda\\lambda$ 5234 and 5362, shown in Figure 6 (plotted on a velocity scale), while in the strongest, such as $\\lambda\\lambda$ 4923 and 5018, it is apparent only as unresolved structure in the shortward wing.  The [\\ion{Fe}{2}] lines, on the other hand, are narrow and single (see two examples in Fig. 6), at a mean velocity of +16.5 $\\pm$ 1 km s$^{-1}$, in clear contrast to \\ion{Fe}{2} at $-28 \\pm 2$ and +44 $\\pm$ 2 km s$^{-1}$.  The spectrum of MWC 778 is cluttered with line emission, but an absorption-line spectrum is dimly discernible in the clearer regions.  No real classification is possible, but most of the detectable lines are those expected of an F- or G-type star, although a late A type cannot be ruled out. Those lines are shallow, probably because of veiling, and are not sharp; their widths correspond to a $v$ sin $i$ of 30--40 km s$^{-1}$.  This is apparent in Figure 7,  a plot of the 6080--6160 \\AA\\ region in MWC 778 (top), compared to that region in the G0\\, V star HR 8314 (bottom), and (middle) the same HR 8314 spectrum broadened to $v$ sin $i$ = 35 km s$^{-1}$.  Note that the vertical scales in the several panels of Figure 7 are very different in order to display each spectrum at approximately the same line amplitude.  The Li\\, I 6707 \\AA\\ line is present in MWC 778, as expected for a pre--main-sequence star, at an equivalent width of 47 m\\AA.  The velocity of the absorption  spectrum, from 23 measurable lines, is +13 $\\pm$ 2 km s$^{-1}$  The B-type absorption spectrum suspected by \\citet{viei03} is not in evidence.  The CO cloud velocity at IC 2144 is +14.1 $\\pm$ 0.3 km s$^{-1}$ according to \\citet{blitz82} and +13.8 $\\pm$ 0.1 according to \\citet{wout89}. In the optical region of MWC 778,  strong interstellar \\ion{Na}{1} D$_{1,2}$ lines are present on the HIRES spectrograms (equivalent widths 0.85 and 0.74 \\AA ) at +13 km s$^{-1}$.  The one unblended interstellar \\ion{K}{1} line $\\lambda$7698 is at +12 km s$^{-1}$.  Two sets of much weaker IS lines fall in the shortward wings of the D lines, at  $-$4 and  $-$15 km s$^{-1}$. The diffuse interstellar band spectrum is strong; several diffuse interstellar bands (DIBs) are indicated in Figure 5. Equivalent widths and velocities of some of the more prominent DIBs are given in Table 2; their mean velocity is +12 $\\pm$ 1 km s$^{-1}$.  The agreement of cloud CO and MWC 778 absorption line velocities is as expected if the star formed within that cloud.  The agreement with the optical interstellar velocities (\\ion{Na}{1}, \\ion{K}{1}) is compatible with the idea that that material also is associated with the CO cloud, although at that longitude only a weak dependence of velocity upon distance is expected.  The emission spectrum of MWC 778 is very similar to that of another high-luminosity pre--main-sequence object, LkH$\\alpha$-101, in which the \\ion{Fe}{2} lines also are double, the $[$\\ion{Fe}{2}$]$ lines are single at an intermediate velocity, the $[$\\ion{O}{1}] lines are double, and the \\ion{Si}{2} lines are very broad \\citep{herb04}.  In that case, the \\ion{Fe}{2} splitting was interpreted as produced in a tipped annulus in rotation around a 15 $M_{\\sun}$ star. ",
        "conclusions": ""
    },
    "0809/0809.3625_arXiv.txt": {
        "abstract": " ",
        "introduction": "%   There has been a dramatic progress in very-high-energy \\cite{TeV} gamma-ray astronomy in the last decade. With the advent of imaging atmospheric Cherenkov telescopes, the number of gamma-ray sources detected in the TeV energy range has increased significantly in the last several years, and now it exceeds seventy, according to the rapporteur talk by Jim Hinton at the International Cosmic Ray Conference in 2007 \\cite{Hin07}. These gamma-ray sources form a new class of high-energy objects in the Universe, including active galactic nuclei, radio galaxies, galactic binaries, pulsar wind nebulae, in addition to supernova remnants which are assumed to be the origin of cosmic-rays for a long time. Exploring the emission mechanism from these objects is a big challenge in astrophysics. Non-thermal nature of emission inherently needs multiwavelength observations to study the phenomena, involving astronomers working in other wavelength. Furthermore, many unidentified gamma-ray sources have been reported and are posing new mysteries.  In this review, highlights of recent findings from ground-based very-high-energy gamma-ray observations are summarized, and up-to-date problems in this research field are presented. We concentrate on recent results from Cherenkov telescopes here, but new results from particle shower arrays are also exciting (G.~Sinnis and M.~Takita, in these proceedings.) Recent review on this field is also found elsewhere \\cite{DeA07}. ",
        "conclusions": "Now the very-high-energy window of the Universe is open, and Cherenkov telescopes are providing additional 2--3 decades to the photon spectra of high-energy objects in the sky. We can see wider variety of source classes among seventy TeV sources than anticipated, indicating cosmic accelerators are {\\it ubiquitous} in the Universe. However, understanding the underlying physical processes is still at a preliminary stage. Together with GLAST (P.~Michelson, in these proceedings.), to be launched in 2008 to observe GeV gamma-rays from space, we can expect many findings from the exploration of high-energy Universe with next-generation ground-based gamma-ray detectors."
    },
    "0809/0809.0760_arXiv.txt": {
        "abstract": "We have observed cosmic-ray electrons from 10~GeV to 800~GeV by a long duration balloon flight using Polar Patrol Balloon (PPB) in Antarctica. The observation was carried out for 13 days at an average altitude of 35~km in January 2004. The detector is an imaging calorimeter composed of scintillating-fiber belts and plastic scintillators inserted between lead plates with 9 radiation lengths. The performance of the detector has been confirmed by the CERN-SPS beam test and also investigated by Monte-Carlo simulations. New telemetry system using a commercial satellite of Iridium, power supply by solar batteries, and automatic level control using CPU have successfully been developed and operated during the flight. From the long duration balloon observations, we derived the energy spectrum of cosmic-ray electrons in the energy range from 100~GeV to 800~GeV. In addition, for the first time we derived the electron arrival directions above 100~GeV, which is consistent with the isotropic distribution. ",
        "introduction": "\\label{sec:introduction}  High-energy cosmic-ray electrons cannot propagate far from the sources, because the electrons lose rapidly energy with an energy loss rate of the square of energy through the synchrotron and inverse Compton process. Kobayashi et al. (2004) \\cite{kobayashi04} suggests that the energy spectrum of cosmic-ray electrons might have a structure in the energy region over several 100GeV due to the discrete effect of local sources. This means that we can identify cosmic-ray electron sources from the electron spectrum over several 100GeV together with anisotropies in arrival direction. In addition, there is a possibility that Weakly Interacting Massive Particles (WIMPs) annihilate directly into electron-positron pairs and produce mono-energetic electrons and positrons \\cite{cheng02}. Although the propagation through the Galaxy would broaden the line spectrum, the observed electron and positron spectrum could still have a distinctive feature. Since there are no other known production mechanisms that would produce an electron and positron peak at energies of 100~GeV - 10~TeV, such a distinctive feature clearly indicates the existence of WIMP dark matter in the Galactic halo. Thus, the observations of high-energy electrons bring us unique information about sources and propagation of cosmic rays, and enable us to search for WIMP dark matter.    Although the cosmic-ray electrons have been observed with many kinds of detectors borne on balloons, most observations are limited below $\\sim$100 GeV. As an exceptional instance, the observation with the emulsion chambers (ECC) has achieved to detect electrons beyond 1TeV by the large amount of exposures of $\\sim14$ days in total for 30 years \\cite{nishimura80, kobayashi02}. The difficulty of observation originates from that the electron flux itself is very small and decreases with energy much more rapidly than that of protons because of the electro-magnetic energy loss by radiation. The electron flux is estimated to be $\\sim$1\\% of protons around 10GeV and less than $\\sim$0.2\\% of protons above 100~GeV from the observed energy spectra with a power-law index of $-3.0$ $\\sim$ $-3.3$ for electrons \\cite{torii01,kobayashi02} and $-2.7$ for protons \\cite{haino04}. Therefore, for the electron observations above 100~GeV, we need a long duration observation by a detector with an excellent capability of large geometrical factor and powerful background rejection.   In order to meet the requirements for the electron observation, we have developed a highly granulated imaging calorimeter, the Balloon-borne Electron Telescope with Scintillating fibers (BETS), that preserves the superior qualities of both electronic detectors and emulsion chambers. We have successfully observed cosmic-ray electrons in the energy range from 10 to 100~GeV at Sanriku in Japan \\cite{torii01}. The BETS was improved to observe atmospheric gamma rays at mountain and balloon altitudes for calibrating the atmospheric neutrino flux calculations \\cite{kasahara02}. Furthermore, for observing the electrons over 100GeV by a long duration balloon experiment, we have achieved the development of an advanced detector, PPB-BETS, flown by PPB in Antarctica. The PPB has a capability to realize a flight for $2-4$ weeks in one round of the Antarctica at an altitude of $\\sim$35~km. The PPB project is carried out by the inter-university collaboration with Institute of Space and Astronautical Science (ISAS), JAXA and National Institute of Polar Research (NIPR) \\cite{nishimura94,kadokura02}.   In this paper, we describe the instrumentation of PPB-BETS, the calibration of the detector with accelerator beams, the observations in Antarctica, and the results of our experiments. ",
        "conclusions": "\\label{sec:result}   \\subsection{Energy spectrum of electrons} Figure~\\ref{fig:espec} presents the electron energy spectrum multiplied by the cube of energy, in comparison with other electron observations \\cite{tang84, golden84, boezio00, duvernois01, torii01, aguilar02, kobayashi02, chang05}. Table~\\ref{tab:elec_flux} presents the flux of electrons at the top of atmosphere in each energy interval. Our spectrum agrees well with the extrapolated spectrum of the BETS\\cite{torii01} at 100~GeV. The combined spectrum of PPB-BETS and BETS can be represented by a single power-law function of \\begin{equation} J_{e} = (1.82{\\pm}0.13){\\times}10^{-4} (\\frac{E}{\\rm 100GeV})^{-3.05{\\pm}0.05} ({\\rm m^{-2} s^{-1} sr^{-1} GeV^{-1}}) \\label{eq:elec_spec} \\end{equation} in the energy range of 10~GeV to 800~GeV. The reduced chi-square value is 1.605 for 11 degrees of freedom. A power-law spectrum is acceptable at the 95~\\% confidence level of statistical errors.    The energy spectrum exceeding 100~GeV is crucial to detect the nearby SNRs as discussed by Kobayashi et al. (2004) \\cite{kobayashi04}, and electron-positron pairs from Kaluza-Klein dark matter annihilations \\cite{cheng02}. Although the data statistics of our results are insufficient to discuss the details of the contribution of nearby SNRs and/or WIMP dark matter, our energy spectrum may indicate a sign of a structure in the several 100~GeV region, as shown in Fig.~\\ref{fig:espec}. Similar structure in the electron energy spectrum is reported by the ATIC-2 observations \\cite{chang05}. Therefore, although the energy spectrum with PPB-BETS cannot reject a single power-law function, both the observations with PPB-BETS and ATIC-2 may indicate a significant spectral structure. For the future observations, we are planning to increase greatly our statistics by experiments such as CALET with the geometrical factor of nearly 1~m$^{2}$sr and the observation time of 3 years on the International Space Station \\cite{torii06}. It is expected that CALET can detect these distinctive signatures from the nearby SNRs and WIMP annihilations \\cite{torii08}.     \\subsection{Arrival directions of electrons}  Incident directions of electrons on the PPB-BETS detector are determined by using the shower axis with an accuracy of around 0.5~degree. Since the attitudes of PPB-BETS instrument are determined with sun aspect sensors and geomagnetic aspect sensors, we can determine arrival directions of cosmic-ray electrons with an accuracy of several degree. For the determination of attitudes with geomagnetic aspect sensors, we referred to the International Geomagnetic Reference Field (IGRF) model \\cite{igrf05}. From these capabilities of PPB-BETS, for the first time we derived the arrival directions of high-energy electrons. Figure~\\ref{fig:obs2isotropy} presents a ratio of the observed distribution above $\\sim$100~GeV to the isotropic distribution along the Galactic longitude. As shown in Fig.~\\ref{fig:obs2isotropy}, the arrival directions of electrons are consistent with the isotropic distribution.    Because the rate of energy loss due to the synchrotron and inverse Compton processes is much higher for electrons than for nuclei, the degree of anisotropy of high-energy cosmic ray electrons is expected to be higher than that of the nuclear component \\cite{shen71, ptuskin95}. The magnitude of anisotropy is expected to be only around 1~\\% in the 100~GeV -- 1~TeV in the direction of Vela at $(\\ell, b) = (-96^{\\circ}.1, -3^{\\circ}.3)$ by using the calculated results by Kobayoashi et al. (2004) \\cite{kobayashi04, yoshida08b}. The first result of the electron arrival directions by PPB-BETS shows no significant anisotropy within statistical errors. For the future observations, the detection of anisotropy toward the sources will enable us to identify the cosmic-ray electron sources, together with the unique signatures of the energy spectrum described above."
    },
    "0809/0809.4249_arXiv.txt": {
        "abstract": "{We have determined carbon abundances for 51 dwarf stars and manganese abundances for 95 dwarf stars in two distinct and well defined stellar populations - the Galactic thin and thick disks. As these two populations have different chemical histories we have been able to, through a differential abundance analysis using high-resolution spectra, constrain the formation sites for carbon and manganese in the Galactic disk(s).  The analysis of carbon is based on the forbidden [C\\,{\\sc i}] line at 872.7\\,nm which is an abundance indicator that is insensitive to errors in the stellar atmosphere parameters. Combining these data with our previously published oxygen abundances, based on the forbidden [O\\,{\\sc i}] line at 630.0\\,nm, we can form very robust [C/O] ratios that we then used to investigate the origin of carbon and the chemical evolution of the Galactic thin and thick disks. We find that the [C/Fe] versus [Fe/H] abundance trends for the thin and thick disks are totally merged and being flat for sub-solar metallicities. The thin disk that extends to higher metallicities then shows a shallow decline in [C/Fe] from $\\rm [Fe/H]\\approx0$ and up to $\\rm[Fe/H]\\approx+0.4$. The [C/O] versus [O/H] trends are well separated between the two disks (due to differences in the oxygen abundances) and bear great resemblance with the [Fe/O] versus [O/H] trends. Our interpretation of our abundance trends is that the sources that are responsible for the carbon enrichment in the Galactic thin and thick disks have operated on a time-scale very similar to those that are responsible for the Fe and Y enrichment (i.e., SN\\,Ia and AGB stars, respectively).  For manganese, when comparing our Mn abundances with O abundances for the same stars we find that the abundance trends in the stars with kinematics typical of the thick disk can be explained by metallicity dependent yields from SN II. Furthermore, the [Mn/O] versus [O/H] trend in the halo is flat. We conclude that the simplest interpretation of our data is that manganese most likely is produced in SN II and that the Mn yields for such SNae must be metallicity dependent. }  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "We have used a two well-defined samples of thin and thick disk stars, both with their own chemical histories, to put further constraints on the cosmic production sites of carbon and manganese. Our findings indicate that carbon in the Milky Way disk mainly comes from low and intermediate mass stars (SN\\,Ia and AGB stars) while manganese is likely to depend on metallicity dependent yields from massive stars. For further details regarding the analysis methods and a full discussion of the results we refer the reader to Bensby \\& Feltzing~(2006) and Feltzing et al.~(2007)."
    },
    "0809/0809.3005_arXiv.txt": {
        "abstract": "We present simultaneous \\chandra-{High Energy Transmission Gratings} (\\hetg) and {Rossi X-ray Timing Explorer} (\\rxte) observations of a moderate flux `soft state' of the black hole candidate \\fu.  These spectra, having a minimally discernible hard X-ray excess, are an excellent test of modern disk atmosphere models that include the effects of black hole spin.  The \\hetg\\ data show, by modeling the broadband continuum and direct fitting of absorption edges, that the soft disk spectrum is only very mildly absorbed with ${\\rm N_H} = 1$--$2 \\times 10^{21}~{\\rm cm}^{-2}$.  These data additionally reveal $\\lambda\\lambda13.449$ \\ion{Ne}{9} absorption consistent with the warm/hot phase of the interstellar medium.  The fitted disk model implies a highly inclined disk around a low mass black hole rapidly rotating with normalized spin $a^* \\approx 1$.  We show, however, that pure Schwarzschild black hole models describe the data extremely well, albeit with large disk atmosphere ``color-correction'' factors. Standard color-correction factors can be attained if one additionally incorporates mild Comptonization. We find that the \\chandra\\ observations do not uniquely determine spin, even with this otherwise extremely well-measured, nearly pure disk spectrum. Similarly, {XMM-Newton}/\\rxte\\ observations, taken only six weeks later, are equally unconstraining.  This lack of constraint is partly driven by the unknown mass and unknown distance of \\fu; however, it is also driven by the limited bandpass of \\chandra\\ and {XMM-Newton}.  We therefore present a series of 48 \\rxte\\ observations taken over the span of several years and at different brightness/hardness levels.  These data prefer a spin of $a^* \\approx 1$, even when including a mild Comptonization component; however, they also show evolution of the disk atmosphere color-correction factors. If the rapid spin models with standard atmosphere color-correction factors of ${\\rm h_d} = 1.7$ are to be believed, then the \\rxte\\ observations predict that \\fu\\ can range from a 3\\,$\\msun$ black hole at 10\\,kpc with $a^*\\approx 0.83$ to a 16\\,$\\msun$ black hole at 22\\,kpc with $a^* \\approx 1$, with the latter being statistically preferred at high formal significance. ",
        "introduction": "\\label{sec:intro}  \\setcounter{footnote}{0}  Despite having an X-ray brightness comparable to and occasionally exceeding that of \\lmcxo\\ or \\lmcxt, and despite also being one of the few persistent black hole candidates (BHC), \\fu\\ has received relatively little attention. Similar to \\lmcxt\\ \\citep[see, e.g.,][]{wilms:01a}, \\fu\\ is usually in a soft state that shows long term (hundreds of days) variations \\citep{nowak:99d,wijnands:02c}. Long term variations also have been seen in its optical lightcurve: modulation over its 9.33\\,hr orbital period has ranged from $\\pm 10\\%$ and sinusoidal \\citep{thorstensen:87a} to $\\pm 30\\%$ and complex \\citep{hakala:99a}.  \\citet{hakala:99a} interpret these changes in the optical lightcurve as evidence for an accretion disk with a large outer rim, possibly due to a warp, being nearly edge-on and partly occulted by the secondary.  The lack of any X-ray evidence for binary orbital modulation \\citep{nowak:99d,wijnands:02c} indicates that the orbital inclination cannot exceed $\\approx 75^\\circ$.  Observations with {Ginga} showed a soft spectrum that could be fit with a multi-temperature disk blackbody \\citep[{\\tt diskbb};][]{mitsuda:84a} plus a power law tail with photon index $\\Gamma \\approx 2$--3 \\citep{yaqoob:93a}.  The power-law component in those observations comprised up to 25\\% of the 1-18\\,keV flux.  Later \\rxte\\ observations also were consistent with a disk blackbody model, but showed no evidence of either a hard component or any X-ray variability above background fluctuations \\citep{nowak:99d}. \\citet{wijnands:02c} conducted a series of Target of Opportunity observations designed to catch \\fu\\ at the high end of its count rate as determined by the {All Sky Monitor} (\\asm) on-board of \\rxte.  They found that at its highest flux, \\fu\\ exhibited both a hard tail and mild X-ray variability.  For all of the above cited X-ray observations, the disk parameters had two attributes in common: the best fit inner disk temperatures were high (up to 1.7\\,keV), and the fitted normalizations were low ($\\sim 10$).  Nominally, the disk blackbody normalization corresponds to $(R_{\\rm km}/D_{\\rm 10})^2 \\cos\\theta$, where $R_{\\rm km}$ is the disk inner radius in {\\rm km}, $D_{\\rm 10}$ is the source distance in units of 10\\,kpc, and $\\theta$ is the inclination of the disk.  In physical models, such large temperatures with such low normalizations normally can be achieved only by a combination of large distance, high inclination, low black hole mass, high accretion rate, and possibly rapid black hole spin.  Disk temperatures increase as the 1/4 power of fractional Eddington luminosity, but decrease as the 1/4 power of black hole mass.  Thus, lower mass serves to increase the temperature and decrease the normalization via decreasing $R_{\\rm km}$.  Large distances and high inclination also serve to reduce the normalization.  (Again, the lack of X-ray eclipses limits $i\\aproxlt 75^\\circ$.)  Finally, rapid spin yields higher efficiency in the accretion disk (and thus higher temperatures), and a decreased inner disk radius.  As pointed out by a number of authors, however, the {\\tt diskbb} model is not a self-consistent description of a physical accretion disk \\citep[see, for example,][]{li:05a,davis:05a,davis:06a}.  Several authors have developed more sophisticated models that incorporate a torque (or lack thereof) on the inner edge of the disk, atmospheric effects on the emerging disk radiation field, the effects of limb darkening and returning radiation (due to gravitational light bending), and the effects of black hole spin.  Two of these models are the {\\tt kerrbb} model of \\citet{li:05a} and the {\\tt bhspec} model of \\citet{davis:05a}.  Using these models on soft-state BHC spectra with minimal hard tails, various claims have been made as to their ability to observationally constrain black hole spin \\citep[e.g.,][]{shafee:06a,davis:06a,mcclintock:06a,middleton:06a}. Given the apparent dominance of a soft spectrum in \\fu, and the typical weakness of any hard tail, \\fu\\ becomes an excellent testbed for assessing the ability of these models to constrain physical parameters.  The outline of this paper is as follows.  First, we describe the analysis procedure for the \\chandra, {XMM-Newton}, and \\rxte\\ data (\\S\\ref{sec:data}).  We then discuss fits to the \\chandra\\ data with both phenomenological (i.e., {\\tt diskbb}) and more physically motivated (i.e., {\\tt kerrbb}) models (\\S\\ref{sec:broad}).  We then discuss the {XMM-Newton} data, and specifically look at the absorption edge regions in both the \\chandra\\ gratings spectra and the {XMM-Reflection Gratings Spectrometer} ({RGS}) spectra (\\S\\ref{sec:edge}).  We then discuss spectral fits to 48 observations from the \\rxte\\ archives (\\S\\ref{sec:rxte_spectra}), and we consider the variability properties of these observations (\\S\\ref{sec:rxte_vary}).  Finally, we present our conclusions and summary (\\S\\ref{sec:discuss}). ",
        "conclusions": "\\label{sec:discuss}  We have presented observations of the black hole candidate \\fu\\ performed with the \\chandra, \\xmm, and \\rxte\\ X-ray observatories.  All of these observations point toward a relatively simple and soft spectrum, indicative of a classic BHC disk-dominated soft state.  The \\chandra\\ and \\xmm\\ spectra, especially at high resolution, indicate a remarkably unadulterated disk spectrum.  That is, the absorption of the spectrum is very low at only ${\\rm N_H} = 1$--$2\\times 10^{21}\\,{\\rm cm}^{-2}$, and with the possible exception of an unidentified line at 21.8\\,\\AA, there is very little evidence of spectral complexity intrinsic to the source.  From that vantage point, \\fu\\ may be the cleanest disk spectrum with which to study modern disk atmosphere models.  \\begin{figure} \\epsscale{1} \\plotone{f14_col.ps} \\caption{An approximately 25\\,Hz quasi-periodic oscillation seen in the 4.5-27.5\\,keV \\fu\\ lightcurve from ObsID 50128-01-09-00, the observation with the smallest disk contribution to the observed 3-18\\,keV flux.  The top panel shows the power spectrum with the Poisson noise level unsubtracted and normalized to $\\approx 2$.  The bottom panel shows Fourier frequency times the power spectrum, now with the Poisson noise level subtracted and the PSD normalized to (root mean square)$^2$/Hz. \\label{fig:qpo}} \\end{figure}  \\input table5  Perhaps the one indication of additional, unmodeled spectral complexity is the fact that the broadband fits and the more localized, high resolution edge and line fits yield (low) neutral columns that differ from each other by a factor of two.  This, along with the 21.8\\,\\AA\\, absorption feature might indicate the presence of other weak, unmodeled spectral components at very soft X-ray energies.  Even if this were the case, this is to be compared to other BHC to which such atmosphere models have been applied, e.g., GRO J1655$-$40 \\citep{shafee:06a}, where there is both a larger neutral column (${\\rm N_H} = 7 \\times 10^{21}\\,{\\rm cm^{-2}}$), a complex Fe K$\\alpha$ region (modeled with both emission and smeared edges by \\citealt{shafee:06a}) and high resolution observations of intermittently present, complex X-ray spectral features that are local to the system \\citep[i.e.,][]{miller:06a,miller:08a}. In comparison, the \\fu\\ spectrum, even with its uncertainties, is much simpler.  The observations presented here strengthen the case that \\fu\\ is a black hole system.  The spectrum closely matches a classic, soft state disk spectrum (well-modeled by {\\tt diskbb}), with a contribution from a harder, non-disk component that generally increases as the flux increases.  Occasionally (a half dozen observations), the contribution from this hard component increases, the X-ray variability increases, a possible QPO is observed, and we even observe lightcurve `dipping' and `flip-flop' behavior \\citep{miyamoto:91a,miyamoto:93a}.  All of this is behavior familiar from studies of other BHC \\citep[][and references therein]{homan:01a}.  Unfortunately, observation of these behaviors provides little strong input to estimates of the system parameters, i.e., mass, distance, and inclination.  `Flip-flop' behavior has often been associated with luminosities of a few tens of percent \\citep{miyamoto:91a,miyamoto:93a,homan:01a}, which is higher than we have presumed for the fits described here. The threshold luminosity for such behavior, however, is not firm, and we further note that high accretion rate/large distance fits to the \\pca\\ spectra failed (\\S\\ref{sec:rxte_spectra}).  The fact that we do not see a transition to a BHC hard state--- although we may see some indications from increases in the fitted disk radius at low luminosity--- strongly suggest that the mass is $>3\\,\\msun$ and the distance is $>10$\\,kpc.  Despite the fact that we do not have a firm estimate of mass and distance for \\fu, the relativistic disk models still strongly statistically prefer rapid spin solutions.  Essentially, this is because the {\\tt diskbb} models fit very high temperatures -- as high as 1.7\\,keV -- with fairly low normalizations.  We saw, however, that for the \\chandra\\ observation, or for the \\rxte\\ observations where the power-law/Comptonization component strengthens, the need for spectral hardening of the disk, either via a color-correction factor (${\\rm h_d}$) or rapid spin, was greatly reduced.  Phenomenologically, a coronal component can mimic the effects of rapid spin.  This naturally leads to the question of whether or not one is sure that any residual coronal component completely vanishes at low flux.  Does one ever observe a ``bare'' disk spectrum?  It is a rather remarkable fact that the \\chandra\\ observation, and the low flux \\rxte\\ observations, are mostly described via three parameters: absorption, {\\tt diskbb} temperature, and {\\tt diskbb} normalization.  Even more remarkably, changes in the spectrum, both in terms of amplitude and overall shape, are mainly driven by variations of only \\emph{one} parameter: {\\tt diskbb} temperature.  The relativistic disk models, on the other hand, must ascribe essentially this one parameter to a combination of four parameters: $a^*$, disk inclination (which affects the appearance of relativistic features), $({\\rm \\dot M/M}^2)$, and ${\\rm h_d}$.  In principle, any of the latter three can vary among observations.  (Disk inclination can vary due to warping; see \\citealt{pringle:96a}.)  Unfortunately, there is no truly unique, sharp feature in the high resolution spectrum that completely breaks degeneracies among these parameters. The overall magnitude of the {\\tt diskbb} temperature and shape of the \\fu\\ spectra, however, seem best reproduced with the most rapid spin models; simply increasing accretion rates or spectral hardening factors is insufficient to model the spectra fully.  We do not see the essential question as being whether or not \\fu\\ is rapidly spinning.  Instead, we view the more observationally motivated and perhaps more fundamental question as being: why is the characteristic disk temperature of \\fu\\ so high?  Rapid spin is one hypothesis.  Another would be that there is indeed a residual, low temperature corona, even at low flux.  The possibly high inclination of this source ($\\approx 75^\\circ$ to be consistent with optical lightcurve variations, while being consistent with the lack of X-ray eclipses) would mean that we are viewing the disk through a larger scattering depth than would be usual for most BHC in the soft state. We also note that there exists very little, self-consistent theoretical modeling of low temperature Comptonizing coronae with high temperature (1--2\\,keV) seed photon temperatures.  Most energetically balanced, self-consistent models \\citep[e.g.,][]{dove:97a} have focused on high coronal temperatures (50--200\\,keV) with low seed photon temperatures (100--300\\,eV).  If one is to take the rapid spin hypothesis for the high {\\tt diskbb} temperature at face value, then these observations offer something of a prediction.  Statistically, they do prefer rapid spin, and based upon observations of other BHC we would have expected to observe a hard state if the faintest observations were $\\aproxlt 3\\%\\,{\\rm L_{Edd}}$.  If the theoretically preferred value of the spectral hardening factor is indeed ${\\rm h_d}=1.7$, future measurements should find \\fu\\ to be an $\\approx 16\\,\\msun$ black hole at $\\approx 22$\\,kpc.  Whether or not this turns out to be the case, given the nature of \\fu\\ as perhaps the simplest, cleanest example of a BHC soft state, observations to independently determine this system's parameters (mass and distance) are urgently needed."
    },
    "0809/0809.0006_arXiv.txt": {
        "abstract": "We report about the detection of 10 clusters of galaxies in the ongoing {\\it Swift}/BAT all-sky survey. This sample, which comprises mostly merging clusters, was serendipitously detected in the 15--55\\,keV band. We use the BAT sample to investigate the presence of excess hard X-rays above the thermal emission. The BAT clusters do not show significant (e.g. $\\geq$2\\,$\\sigma$) non-thermal hard X-ray emission. The only exception is represented by Perseus whose high-energy emission is likely due to NGC 1275. Using XMM-Newton, Swift/XRT, Chandra and BAT data, we are able to produce upper limits on the Inverse Compton (IC) emission mechanism which are in disagreement with most of the previously claimed hard X-ray excesses. The coupling of the X-ray upper limits on the IC mechanism to radio data shows that in some clusters the  magnetic field might be larger than 0.5\\,$\\mu$G. We also derive the first log $N$ - log $S$ and luminosity function distribution of  galaxy clusters above 15\\,keV. ",
        "introduction": "\\label{intro}  Galaxy clusters are potentially powerful observational probes of dark matter and dark energy. However, the use of clusters to measure cosmological parameters becomes accessible only when astrophysical uncertainties are well understood and controlled. Indeed the non-thermal pressure due to Cosmic Rays (CRs), magnetic fields and turbulence, is a source of systematic bias when cluster masses are estimated using the assumption of hydrostatical equilibrium \\citep[e.g.][]{ensslin97}. The detection of clusters' X-ray emission above $\\sim$20\\,keV is a fundamental step towards the  firm grasp of  these processes.   It is well understood that clusters of galaxies contain a large amount of hot gas, called intracluster medium (ICM), that comprises 10--15\\,\\% of their total mass. Already the first X-ray observations indicated the presence of this optically thin plasma, characterized by an atomic density of about $10^{-4}$--$10^{-2}$\\,cm$^{-3}$ and temperatures of the order of $10^7$--$10^8$\\,K \\citep[e.g.][]{fel66,cat72}. Also well established is the fact  that the observed X-ray radiation from clusters of galaxies is primarily due to the thermal bremsstrahlung emission of  such  diffuse hot plasma \\citep{sar88,petrosian01}.  However, evidences gathered at different wavelengths point to the existence of a non-thermal component. In particular the detections of an extended synchrotron radio emission \\citep[e.g.][]{wil70,har78,gio93,gio00,kem01,thierbach03} and, more recently, of a possibly non-thermal extreme ultraviolet (EUV) excess \\citep{lieu96,bow99,bon01,dur02} and soft excess \\citep[e.g.][]{wer07} suggest the existence of a non-thermal X-ray component originating from a population  of relativistic electrons. This scenario is confirmed by the detection of non-thermal emission in the hard X-ray spectra of a few galaxy clusters \\citep[see e.g.][for a complete review]{kaastra08,rephaeli08}. Still its actual presence and origin remain controversial \\citep{renaud06,fus07,wer07,lutovinov08}.   A non-thermal component could arise from a population of point sources \\citep[e.g. AGN as in][]{katz76,fabian76,fujita07} or from inverse Compton (IC) scattering of cosmic microwave background (CMB) photons by relativistic electrons \\citep[e.g.][]{rep79,sar99}. Other possible mechanisms are non-thermal bremsstrahlung \\citep[e.g.][]{sar99,sar00} and synchrotron emission from ultra-relativistic electrons \\citep[][]{tim04,ino05,eck07}. If the origin of the high-energy emission is IC scattering, then the presence of a large population of relativistic electrons (Lorentz factor $>>1000$) is required. This population could have been accelerated in shocks of different origin. Indeed it could be associated to merger shocks \\citep[e.g.][]{fuj03,bru04}, dark matter bow shocks \\citep[e.g.][]{byk00}, ram-pressure stripping of infalling galaxies \\citep[e.g.][]{dep06}, jets, Active Galactic Nuclei outbursts \\citep[][in the case of radio mini-halos such as in Perseus cluster]{fujita07}, accretion shocks \\citep[e.g.][]{ino05}. Non-thermal electrons loose energy on short timescales (below $1$ Gyr). Therefore some models consider a continued supply of primary accelerated electrons (i.e. via first order Fermi mechanism), while others assume a constant in-situ re-acceleration via CR collisions or second order Fermi mechanism.   If clusters are a large reservoir of non-thermal particles, then they should emit at higher energies, up to the $\\gamma$-rays. Indeed, if CRs acceleration is taking place at the shock fronts then $\\gamma$-rays can be produced via IC, non-thermal bremsstrahlung and $\\pi^0$  decay \\citep[e.g.][]{rep79,dar95,rei03,reimer04,blasi07}. A statistical upper limit on the flux above $>100$ MeV was obtained by \\cite{rei03}, analyzing the emission from $58$ clusters observed with EGRET.  The role of CRs in the formation and evolution of clusters of galaxies has been much debated. \\cite{chu07} suggest that in massive galaxy clusters hydrostatic equilibrium is satisfied reasonably well, as long as the source has not experienced a recent major merger. However, in non-relaxed clusters the non-thermal pressure due to CRs, magnetic fields and micro-turbulence can affect  the mass estimates based on hydrostatic equilibrium \\citep[e.g.][]{mir95,nag07}. This would lead to a higher baryonic to total mass ratio. Knowing the importance of CRs, the mechanisms that heat the ICM and the frequency at which it is shocked, is crucial for the upcoming X-ray and Sunyaev-Zeldovich surveys \\citep[see][]{and07}.   In this paper we report the {\\it Swift}/BAT all-sky detection of 10 galaxy  clusters in the 15--55\\,keV band. This constitutes the first complete sample  so far detected at these energies. We use this sample to investigate the role of non-thermal processes in clusters. The structure of the paper is the following. In $\\S$~\\ref{subsec:batsurvey} we describe the {\\it Swift}/BAT observations and discuss the properties of each individual cluster ($\\S$~\\ref{subsec:individual}). In  $\\S$~\\ref{subsec:non-thermal}, we provide, for all the clusters, constraints on the non-thermal emission as well as an estimate of the clusters' magnetic fields ($\\S$~\\ref{subsec:magnetic}). The cluster source count distribution and the luminosity function are  derived in  $\\S$~\\ref{sec:pop}. We discuss  the results of our analysis  in $\\S$~\\ref{sec:disc}, while $\\S$~\\ref{sec:concl} summarizes our findings.    We adopt a Hubble constant of $H_0 = 70$ h$_{70}$\\,km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.3 and $\\Omega_\\Lambda$ = 0.7. Unless otherwise stated errors are quoted at the 90\\,\\% confidence level (CL) for one interesting parameter and solar abundances are determined using the meteoritic values provided in \\cite{anders89}. ",
        "conclusions": "\\label{sec:concl} BAT is the first instrument to detect above 15\\,keV an all-sky sample of galaxy clusters\\footnotemark{}. \\footnotetext{We are aware of an independent work \\citep{okajima08} based on an alternative analysis of BAT survey data which reaches conclusions consistent with this analysis.} The BAT energy range (15--200\\,keV) is the best one to investigate the presence of non-thermal emission whose detection   remained so far controversial. The results of our investigation can be summarized as follows: \\begin{itemize} \\item Perseus is the only cluster among the 10 BAT objects which displays an high-energy non-thermal component which extends up to 200\\,keV. It is very likely that the central AGN NGC 1275 is responsible for such emission. This claim is supported by several evidences: 1) the  variability seen with BeppoSAX \\citep{nevalainen04}, 2) the XMM-Newton spectral analysis \\citep{churazov2003}, and 3) our combined BAT--XRT--XMM-Newton analysis which shows that the nucleus has a typical AGN spectrum.   \\item The BAT spectra of the remaining 9 galaxy clusters is well fitted by a simple thermal model that constrains non-thermal flux to be below 1\\,mCrab in the 50--100\\,keV band. \\item Assuming that IC scattering is the main mechanism at work for producing non-thermal high-energy flux, it is possible to estimate the magnetic field using Radio data and the upper limits derived above. We obtain that  in all the BAT clusters the (average) magnetic field is $>0.1$\\,$\\mu$G. These (rather uncertain) values are in disagreement (if the magnatic field intensities are close to the lower limits) with the, also uncertain, Faraday rotation measurements which show that the magnetic field is in the $\\sim$\\,$\\mu$G range. Our low magnetic field values would imply that the magnetic field is far from equipartition. \\item The stacked spectrum of the BAT clusters (except Perseus and Coma) confirms once again the absence of any non-thermal high-energy component. The $\\sim$56\\,Ms stacked spectrum constrains any non-thermal flux to be below 0.3\\,mCrab (or 1.9$\\times 10^{-12}$\\,erg cm$^{-2}$s$^{-1}$) in the 50--100\\,keV band. \\item  Using Swift/XRT, XMM-Newton and Chandra, in addition to BAT data, we were able to produce X--ray cluster spectra which extend more than 3 decades in energy (0.5--50\\,keV). In all cases, but Perseus and Abell 0754, the broad-band X--ray spectrum is well approximated by a single-temperature thermal model. These spectra allowed us to put constrains on the IC emission mechanism which are a factor $>$5 lower than those derived using BAT data alone. This would in turn imply a larger intensity of the magnetic field. For both Perseus and Abell 0754 an additional power-law component is statistically required, but several evidences confirm that two X-ray point sources (NGC 1275 and  2MASS\t09091372-0943047) account for the total non-thermal emission. \\item The cluster centroid shift in different wavebands, the morphology and the complex temperature maps (available in literature), show that 8 out of 10 clusters are in the middle of a major merging phase. Shocks, which are revealed by XMM-Newton and Chandra images, are actively heating the ICM as the BAT high temperatures testify. The BAT observations and limits on the non-thermal emissions can help to calibrate the large scale structure formation simulations focusing in particular on the treatment of non-thermal particle emission and cooling. \\item We have produced the first cluster source count (also known as log {\\it N} - log {\\it S}) distribution above 15\\,keV. This shows that, at the limiting fluxes sampled by BAT, the surface density of clusters is $\\sim$5\\,\\% of that one of AGNs. Moreover, we find that the contribution of clusters to the Cosmic X-ray background is of the $\\sim$0.1\\,\\% order in the 15--55\\,keV band. The BAT log {\\it N} - log {\\it S} can be used to predict the cluster surface density for future hard X-ray instruments. \\item The X-ray luminosity function of the BAT clusters, the first derived above 15\\,keV, is in excellent agreement with the ROSAT luminosity function derived in the 0.1--2.4\\,keV band. \\end{itemize}"
    },
    "0809/0809.0695_arXiv.txt": {
        "abstract": "We review the ensemble of anticipated gravitational-wave (GW) emission processes in stellar core collapse and postbounce core-collapse supernova evolution.  We discuss recent progress in the modeling of these processes and summarize most recent GW signal estimates. In addition, we present new results on the GW emission from postbounce convective overturn and protoneutron star $g$-mode pulsations based on axisymmetric radiation-hydrodynamic calculations. Galactic core-collapse supernovae are very rare events, but within $3-5\\;\\mathrm{Mpc}$ from Earth, the rate jumps to 1 in $\\sim 2$ years. Using the set of currently available theoretical gravitational waveforms, we compute upper-limit optimal signal-to-noise ratios based on current and advanced LIGO/GEO600/VIRGO noise curves for the recent SN 2008bk which exploded at $\\sim$3.9 Mpc.  While initial LIGOs cannot detect GWs emitted by core-collapse events at such a distance, we find that advanced LIGO-class detectors could put significant upper limits on the GW emission strength for such events.  We study the potential occurrence of the various GW emission processes in particular supernova explosion scenarios and argue that the GW signatures of neutrino-driven, magneto-rotational, and acoustically-driven core-collapse SNe may be mutually exclusive. We suggest that even initial LIGOs could distinguish these explosion mechanisms based on the detection (or non-detection) of GWs from a galactic core-collapse supernova. ",
        "introduction": "\\label{section:intro}  Ever since the very first experimental efforts to detect gravitational waves (GWs), core-collapse supernovae (SNe) have been considered as potential astrophysical emission sites.  There are very strong indications from theory and observation that multi-D dynamics play a prominent and probably decisive role in core-collapse SNe (see, e.g., \\cite{burrows:00nature,janka:07}).  GWs are emitted at lowest order by an accelerated mass-energy quadrupole moment. Hence, by their intrinsic multi-D nature, GWs, if detected from a core-collapse event, will very likely prove powerful messengers that can provide detailed and live dynamical information on the intricate multi-D dynamics occurring deep inside collapsing massive stars.  Massive stars ($8$--$10\\;\\mmsun \\lesssim M \\lesssim 100\\;\\mmsun$ at zero-age main sequence [ZAMS]) form electron-degenerate cores composed primarily of iron-group nuclei in the final stages of their exoergic nuclear burning. Once such an iron core exceeds its effective Chandrasekhar mass (see, e.g., \\cite{baron:90,bethe:90}) it grows gravitationally unstable.  Collapse ensues, leading to dynamical compression of the inner core material to nuclear densities. There, the nuclear equation of state (EOS) stiffens, resulting in the rebound of the inner core (``core bounce''). A hydrodynamic shock wave is launched at the outer edge of the inner core and propagates outward in mass and radius, slamming into the still infalling outer core. Owing to the dissociation of heavy nuclei and to energy losses to neutrinos that stream away from the postshock region, the shock quickly loses energy, stalls and must be \\emph{revived} to plow through the stellar envelope, blow up the star, and produce a SN explosion, leaving behind a neutron star.  Without shock revival, black-hole (BH) formation is inevitable and even with a successful explosion, a BH may still form via fall-back accretion.  Iron core collapse is the most energetic process in the modern universe, liberating some $10^{53}\\;\\mathrm{erg} = 100\\;\\mathrm{B}$ (Bethe) of gravitational energy.  Most of this energy, $\\sim 99\\,\\%$, is emitted in neutrinos as the protoneutron star (PNS) contracts and cools over a timescale of $\\sim 100\\,\\mathrm{s}$. Only $\\sim 1\\,\\%$ goes into the asymptotic energy of the SN explosion and becomes visible in the electromagnetic spectrum. The fundamental question that core-collapse SN theory has been facing for the past $\\sim 45$ years is how exactly the necessary fraction of gravitational energy is transferred to revive the shock and ultimately unbind the stellar envelope. Shock revival must occur sooner than $1$--$1.5\\;\\mathrm{s}$ after bounce (depending on progenitor star structure setting the rate of mass accretion) in order to produce a compact remnant that obeys observational and theoretical neutron star upper baryonic mass limits around $\\sim 1.5$--$2.5\\;\\mmsun$ (see \\cite{lattimer:07} and references therein).   The SN \\emph{explosion mechanism} may involve (a combination of) heating of the postshock region by neutrinos, multi-dimensional hydrodynamic instabilities of the accretion shock, in the postshock region, and in the PNS, rotation, PNS pulsations, magnetic fields, and nuclear burning (for a recent review, see \\cite{janka:07}, but also \\cite{burrows:07bethe}). Three SN mechanisms are presently discussed in the literature. The \\emph{neutrino mechanism} has the longest pedigree \\cite{bethewilson:85,bethe:90,janka:07}, is based on postbounce neutrino energy deposition behind the stalled shock, and appears to require \\cite{buras:06b,marek:07,murphy:08} convection and the standing-accretion-shock instability (SASI, see, e.g., \\cite{scheck:08} and references therein) to function in all but the very lowest-mass massive stars which may explode even in spherical symmetry~\\cite{burrows:07c,kitaura:06}. Recent detailed 2D neutrino-radiation hydrodynamics simulations by the Garching group produced explosions in particular $11.2$-$\\mmsun$ and $15$-$\\mmsun$ progenitor models \\cite{buras:06b,marek:07}. However, it is not yet clear how the neutrino mechanism's efficacy varies with progenitor ZAMS mass and structure and what its detailed dependence on the high-density nuclear EOS may be.  The \\emph{magneto-rotational (or MHD) mechanism}, probably operating only in the context of rapid progenitor rotation, depends on magnetic-field amplification during collapse and at postbounce times. It leads to explosions that develop in jet-like fashion along the axis of rotation \\cite{leblanc:70,symbalisty:84,yamada:04,burrows:07b,sawai:08} and may reach hypernova energies of $\\sim 10\\,\\mathrm{B}$~\\cite{burrows:07b}. The MHD mechanism may also be relevant in the context of long-soft gamma-ray bursts~(GRBs, see, e.g., \\cite{wb:06}) and could be a precursor, setting the stage for a later GRB~(e.g., \\cite{burrows:07b}, but see \\cite{dessart:08a}).  The \\emph{acoustic mechanism} for core-collapse SNe, as recently proposed by Burrows~et~al.~\\cite{burrows:06,burrows:07a,ott:06prl, burrows:07bethe}, requires the excitation of large-amplitude PNS pulsations (primarily $g$-modes) by turbulence and SASI-modulated accretion downstreams. These pulsations damp by the emission of strong sound waves that steepen to shocks and deposit energy in the postshock region, eventually leading to late explosions at $\\gtrsim 1\\;\\mathrm{s}$ after bounce. This mechanism appears to be sufficiently robust to blow up even the most massive and extended progenitors, but has so far not been confirmed by other groups (see also \\cite{yoshida:07,weinberg:08}).  Constraining the core-collapse SN mechanism via astronomical observations is difficult. The intricate pre-explosion dynamics of the SN core deep inside the supergiant presupernova star are inaccessible by the traditional means of astronomy. Theoretical models of the SN mechanism can currently be tested via secondary observables only, including the asymptotic explosion energy, ejecta morphology, nucleosynthesis products, compact remnant mass and proper motion, and pulsar spin/magnetic fields.  GWs and neutrinos are the only messengers with the potential of delivering first-hand information on the physical processes leading to explosion: Both are emitted deep inside the SN core and travel to observers on Earth practically without interaction with intervening material.  A small number of neutrinos were detected from SN 1987A in the Large Magellanic (distance $D \\approx 50\\;\\mathrm{kpc}$; see, e.g., \\cite{bethe:90} and references therein). GWs have not yet been observed directly, but the advent of GW astronomy has begun. An international network of broad-band light-interferometric GW observatories is active, encompassing the US LIGOs \\cite{ligo}, the British-German GEO600 \\cite{geo600}, the French-Italian VIRGO~\\cite{virgo} and the Japanese TAMA 300 \\cite{tama300}. The three LIGO interferometers have recently reached their design sensitivities, and, in their S5 science run, have taken a year worth of data, partly in coincidence with GEO600 and VIRGO. In addition, there are a number of active resonant bar/sphere GW detectors in operation, including the four bar detectors of the International Gravitational Event Collaboration (IGEC-2), ALLEGRO, AURIGA, EXPLORER, and NAUTILUS~\\cite{astone:07}, and the resonant spheres MiniGrail~\\cite{minigrail} and Schenberg~\\cite{schenberg}. The current status of ground-based GW detection was recently summarized by Whitcomb~\\cite{whitcomb:08}.  GWs from astrophysical sources are weak and notoriously difficult to detect (e.g., \\cite{thorne:87}).  Hence, in order to disentangle an astrophysical GW signal from the mostly overwhelming detector noise, GW astronomy does not only require sensitive detectors, but also extensive processing and analysis of the detector output on the basis of reliable theoretical estimates for the GW signals presently expected from astrophysical sources. The latter must, in most cases, be obtained via detailed numerical modeling of the dynamics responsible for the GW emission in a given source.  In iron core collapse and postbounce SN evolution, the emission of GWs is expected primarily from rotating collapse and bounce, nonaxi\\-symmetric rotational instabilities, postbounce convective overturn/SASI, and PNS pulsations. In addition, anisotropic neutrino emission, global precollapse asymmetries in the iron core and surrounding burning shells, aspherical mass ejection, magnetic stresses, and the late-time formation of a black hole may contribute to the overall GW signature.  The aim of this topical review is to summarize the recent significant progress in the modeling of the various GW emission processes in core-collapse SNe with a focus on the early postbounce, pre-explosion SN evolution up to $\\sim1-2\\;\\mathrm{seconds}$ after bounce. We do not cover the GW emission from the collapse of supermassive primordial or very massive population III stars, dynamical fission processes, late-time fall-back accretion on black holes, nuclear phase-transitions in PNSs, or from late-time postbounce secular instabilities such as $r$-modes. Other reviews of the GW signature of core-collapse SNe are those of Kotake~et~al.~\\cite{kotake:06a} and Fryer~and~New~\\cite{new:03}.  In \\sref{section:history}, we provide a concise historical overview on early work and go on to discuss computational core-collapse SN modeling and the various ways in which GR and GW extraction are treated in \\sref{section:snmodel}. Section~\\ref{section:rotcoll} covers the most extensively modeled GW emission mechanism, rotating core collapse and bounce. The potential for and the GW emission from nonaxisymmetric rotational instabilities is the topic of \\sref{section:rotinst}. Section~\\ref{section:convsasi} is devoted to the emission of GWs from convective overturn and SASI, while we discuss the GW signal from PNS core pulsations in \\sref{section:puls}. To both of these sections we add new, previously unpublished results obtained via 2D Newtonian radiation-hydrodynamics calculations.  In \\sref{section:neutrinos}, we discuss the emission of GWs from anisotropic neutrino radiation fields and in \\sref{section:others} we summarize the GW signals associated with rapid aspherical outflows, precollapse global asymmetries, strong magnetic stresses, and PNS collapse to a black hole.  While the SN rate in the Milky Way and the local group of galaxies is rather low and probably less than 1 SN per two decades (e.g., \\cite{vdb:91}), there may be 1 SN occurring about every other year between $3-5\\;\\mathrm{Mpc}$ from Earth~\\cite{ando:05}. The recent SN 2008bk, which exploded roughly $3.9\\;\\mathrm{Mpc}$ away, is an example SN from this region of space. Thus, in \\sref{section:2008bk}, we present optimal single-detector matched-filtering signal-to-noise ratios for LIGO/GEO600/VIRGO and advanced LIGOs for a subset of the gravitational waveforms reviewed in this article and with an assumed source distance corresponding to that of SN 2008bk. We find that initial LIGO-class detectors had no chance of detecting GWs from SN 2008bk. Advanced LIGOs, however, could put some constraints on the GW emission strength, but still would probably not allow detailed GW observations. We wrap up our review in section~\\ref{section:conclusions} with a critical summary of the subject matter and discuss in which way the various GW emission processes can be linked to particular SN explosion mechanisms. We argue that the GW signatures of the neutrino, MHD, and acoustic SN mechanisms may be mutually exclusive and that the mere detection, or, in fact non-detection, of GWs from a nearby core-collapse SN can constrain significantly the core-collapse SN explosion mechanism.   In this review article, all values of the dimensionless GW signal amplitudes $h_+$ and $h_\\times$ are given for optimal source-observer orientation and a source distance of $10\\;\\mathrm{kpc}$ is typically assumed. In most figures showing GW signals, the waveforms are plotted as $h_{+,\\times}\\, D$, rescaled by distance $D$ and in units of centimeters. Summaries of GW extraction methods can be found in \\cite{thorne:87,moenchmeyer:91,zwerger:97,ott:04,ott:06phd} and a comparison of Newtonian and general relativistic methods were presented in \\cite{shibatasekiguchi:03,baiotti:08b}. ",
        "conclusions": ""
    },
    "0809/0809.5079_arXiv.txt": {
        "abstract": "We present optical photometry and spectroscopy of SN~2005ip for the first 3~yr after discovery, showing an underlying Type II-L supernova (SN) interacting with a steady wind to yield an unusual Type IIn spectrum.  For the first $\\sim$160~d, it had a fast linear decline from a modest peak absolute magnitude of about $-$17.4 (unfiltered), followed by a plateau at roughly $-$14.8 for more than 2~yr.  Initially having a normal broad-lined spectrum superposed with sparse narrow lines from the photoionized circumstellar medium (CSM), it quickly developed signs of strong CSM interaction with a spectrum similar to that of SN~1988Z.  As the underlying SN~II-L faded, SN~2005ip exhibited a rich high-ionization spectrum with a dense forest of narrow coronal lines, unprecedented among SNe but reminiscent of some active galactic nuclei.  The line-profile evolution of SN 2005ip confirms that dust formation caused its recently reported infrared excess, but these lines reveal that it is the first SN to show clear evidence for dust in {\\it both} the fast SN ejecta and the slower post-shock gas.  SN~2005ip's complex spectrum confirms the origin of the strange blue continuum in SN~2006jc, which also had post-shock dust formation.  We suggest that SN~2005ip's late-time plateau and coronal spectrum result from rejuvenated CSM interaction between a sustained fast shock and a clumpy stellar wind, where X-rays escape through the optically thin interclump regions to heat the pre-shock CSM to coronal temperatures. ",
        "introduction": "Core-collapse supernovae (SNe) show a variety of spectral properties (see Filippenko 1997 for a review) based primarily on the amount of mass shed by the progenitor, causing the outer layers of the star to be stripped to different chemical layers at the time of its explosion. In the Type~IIn subclass, substantial mass stripping has occurred recently, leaving dense hydrogen gas in the circumstellar medium (CSM) into which the SN blast wave propagates.  In this interaction, the blast wave is decelerated and the CSM is illuminated, giving rise to relatively narrow emission lines from the post-shock gas and the unshocked CSM.  The Type~IIn subclass shows particularly wide diversity in both luminosity and spectral features, depending on density, speed, and how long before explosion the H-rich material was shed by the star. SNe~IIn can be among the most luminous SNe observed if the CSM is very dense, as in the case of SN~2006tf (Smith et al.\\ 2008b) --- but they can also be among the faintest SNe observed as in the case of the so-called ``supernova impostors'' and related objects, which are not yet clearly understood (e.g., Van Dyk et al.\\ 2005; Thompson et al.\\ 2008).  In order for the CSM interaction luminosity to compete with the main peak of the SN (arising from the diffusion of radioactive decay luminosity and shock-deposited energy in the SN ejecta), the CSM must be dense.  At lower densities where the conversion of shock energy into visual light is less efficient, signs of weaker CSM interaction might become more prominent in the spectrum once the underlying SN fades; indeed, here we suggest that SN~2005ip was an example of the latter.  SN~2005ip was discovered (Boles 2005) on 2005 Nov.\\ 5.163 (UT dates are used throughout this paper), located 2$\\farcs$8~E and 14$\\farcs$2~N of the center its host Scd galaxy NGC~2906 ($\\sim$2.1 kpc in projection on the sky).  Fox et al.\\ (2008) show images of SN~2005ip in its host galaxy.  With an apparent redshift of $z = 0.00714$ (de Vaucouleurs et al.\\ 1991), NGC~2906 is located at a distance of roughly 29.7 Mpc (adopting $H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$).  Modjaz et al.\\ (2005) reported that in a spectrum obtained $\\sim$1~d after discovery, SN~2005ip appeared to be a normal Type~II event a few weeks after explosion, with a blue continuum and broad absorption features indicating a SN-ejecta expansion speed of roughly 15,400 km s$^{-1}$.  Recently, Fox et al.\\ (2008) have shown that SN~2005ip had near-infrared (IR) excess emission from hot dust, and suggested that grains formed in the post-shock region, analogous to SN~2006jc (Smith et al.\\ 2008a).  We present spectra showing that SN~2005ip also had narrow H$\\alpha$ emission indicative of a Type~IIn classification, even on day 1, plus an unusually rich forest of narrow coronal emission lines that dominate the spectrum at later times.  Immler \\& Pooley (2007) reported a high X-ray luminosity at late times, implying a level of CSM interaction comparable to that of other SNe~IIn like SN 1988Z.  We present a brief discussion of the remarkable light curve and focus on the spectral evolution of SN~2005ip, interpreting it as the result of rejuvenated late-time CSM interaction in a steady wind with relatively poor efficiency in converting shock energy to visual light.  We confirm the suggestion by Fox et al.\\ (2008) that dust formed, but our analysis of line profiles reveals that dust formed in both the post-shock shell and the fast SN ejecta at different times.    \\begin{figure} \\epsscale{0.99} \\plotone{fig1.eps} \\caption{The unfiltered (approximately $R$-band) absolute magnitude light curve of SN~2005ip measured by KAIT, compared to that of SN~1988Z (from Turatto et al.\\ 1993).  The absolute magnitude of SN~2005ip is shown for a distance of 29.72 Mpc and an extinction of $A_R = 0.126$ mag (see text).  The small triangles at the bottom mark dates for which we secured spectra of SN~2005ip, and the ``X'' marks the date (day 466) for which Immler \\& Pooley (2007) observed its X-ray emission.  The gray curve shows a normal SN~II-P (SN~1999em) from our database for comparison, and the dot-dash line is $^{56}$Co decay luminosity for 0.08~M$_{\\odot}$ of $^{56}$Ni.} \\label{fig:one} \\end{figure}  \\begin{deluxetable}{lcc}\\tabletypesize{\\scriptsize} \\tablecaption{Photometry of SN 2005\\lowercase{ip}} \\tablewidth{0pc} \\tablehead{ \\colhead{JD} &\\colhead{mag.}  &\\colhead{$\\sigma$} } \\startdata 2453694.05  &15.15  &0.13  \\\\ 2453711.03  &15.62  &0.13  \\\\ 2453716.96  &15.73  &0.13  \\\\ 2453747.91  &16.20  &0.13  \\\\ 2453758.93  &16.42  &0.13  \\\\ 2453770.91  &16.70  &0.13  \\\\ 2453776.82  &16.79  &0.13  \\\\ 2453788.84  &17.06  &0.13  \\\\ 2453818.76  &17.52  &0.13  \\\\ 2453852.70  &17.77  &0.13  \\\\ 2454046.05  &17.62  &0.13  \\\\ 2454070.04  &17.69  &0.13  \\\\ 2454071.06  &17.62  &0.13  \\\\ 2454074.97  &17.74  &0.13  \\\\ 2454087.99  &17.63  &0.13  \\\\ 2454110.93  &17.75  &0.13  \\\\ 2454116.93  &17.67  &0.13  \\\\ 2454123.88  &17.72  &0.13  \\\\ 2454150.90  &17.55  &0.13  \\\\ 2454168.80  &17.65  &0.13  \\\\ 2454177.78  &17.70  &0.13  \\\\ 2454185.71  &17.70  &0.18  \\\\ 2454192.75  &17.67  &0.13  \\\\ 2454200.74  &17.76  &0.13  \\\\ 2454232.69  &17.70  &0.13  \\\\ 2454402.02  &17.65  &0.13  \\\\ 2454410.07  &17.66  &0.13  \\\\ 2454422.07  &17.73  &0.13  \\\\ 2454431.00  &17.75  &0.13  \\\\ 2454443.04  &17.67  &0.13  \\\\ 2454447.98  &17.68  &0.13  \\\\ 2454454.06  &17.71  &0.13  \\\\ 2454466.02  &17.77  &0.13  \\\\ 2454478.00  &17.58  &0.15  \\\\ 2454482.93  &17.55  &0.13  \\\\ 2454504.89  &17.74  &0.13  \\\\ 2454510.87  &17.57  &0.13  \\\\ 2454522.91  &17.62  &0.13  \\\\ 2454528.78  &17.77  &0.13  \\\\ 2454546.79  &17.84  &0.13  \\\\ 2454551.73  &17.70  &0.13  \\\\ 2454557.77  &17.87  &0.13  \\\\ 2454564.75  &18.01  &0.13  \\\\ 2454573.72  &17.66  &0.13  \\\\ 2454590.69  &17.68  &0.13  \\\\ 2454596.68  &18.05  &0.13  \\\\ \\enddata \\tablecomments{KAIT unfiltered photometry; roughly $R$ band.} \\end{deluxetable}   \\begin{deluxetable}{lclccc}\\tabletypesize{\\scriptsize} \\tablecaption{Spectroscopy of SN~2005\\lowercase{ip}} \\tablewidth{0pt} \\tablehead{ \\colhead{UT Date} &\\colhead{Day\\tablenotemark{a}} &\\colhead{Inst.\\tablenotemark{b}} &\\colhead{$\\lambda$/$\\Delta\\lambda$}  &\\colhead{Airmass} &\\colhead{Exp.\\ (s)} } \\startdata 2005 Nov. 06.65 &1    &DEIMOS &2500  &1.07 &100  \\\\ 2006 Jan. 05.40 &61   &Kast   &600   &1.18 &1200 \\\\ 2006 Feb. 06.40 &93   &Kast   &700   &1.19 &1800 \\\\ 2006 Feb. 22.36 &109  &Kast   &600   &1.18 &1500 \\\\ 2006 Apr. 27.26 &173  &LRIS   &1100  &1.03 &504  \\\\ 2006 Dec. 23.56 &413  &DEIMOS &2500  &1.03 &500  \\\\ 2007 Nov. 11.63 &736  &LRIS   &1500  &1.10 &600  \\\\ 2008 Apr. 28.26 &905  &LRIS   &1500  &1.03 &600  \\\\ \\enddata \\tablenotetext{a}{Days after discovery.} \\tablenotetext{b}{Kast, Lick 3-m; LRIS, Keck I; DEIMOS, Keck II.} \\end{deluxetable} ",
        "conclusions": "We present and analyze visual-wavelength photometry and spectroscopy of SN~2005ip obtained at the Lick and Keck Observatories during a period of $\\sim$3 yr after explosion.  Our main conclusions are as follows.  (1) The data show that SN~2005ip was composed of an underlying broad-lined Type II-L event that dominated the luminosity decline during the first $\\sim$160~d, superposed with a Type IIn spectrum arising from constant-luminosity CSM interaction that caused a remarkably steady late-time plateau.  SN~2005ip was not unusually luminous compared with other SNe~IIn, and its initial mass of $^{56}$Ni was less than about 0.1~M$_{\\odot}$.  (2) As the underlying SN~II-L faded, the spectrum of SN~2005ip came to exhibit a forest of narrow coronal emission lines, the number and strength of which dominated the spectrum to an unprecedented degree. Ionization levels as high as [Fe~{\\sc xiv}] are seen.  These forbidden coronal lines arise exclusively in the pre-shock CSM, showing no broader components from the post-shock gas or SN ejecta.  Only a subset of the narrow coronal lines was present in the day 1 spectrum.  (3)  Following the discovery by Fox et al.\\ (2008) that SN~2005ip showed near-IR excess emission attributable to freshly synthesized dust, we confirm evidence for this dust formation from its influence on the evolution of visual-wavelength emission lines.  Components of both He~{\\sc i} $\\lambda$7065 and H$\\alpha$ have red wings that fade faster than the constant blue wings, as expected if new dust grains preferentially block the far side of the object.  However, these two tracers reveal dust at different places at different times as the SN evolves.  At very late times (2--3 yr after explosion), intermediate-width components of He~{\\sc i} lines reveal dust forming in a post-shock shell analogous to the post-shock dust formation in SN~2006jc (Smith et al.\\ 2008a), as proposed by Fox et al.\\ (2008). At early times, however, the situation is qualitatively different: we see no evidence for dust in the intermediate-width components of these lines, but instead, we see the same fading of the red wing in the {\\it broad} component of H$\\alpha$.  This means that at early times, either (a) the dust formed directly in the fast SN ejecta, or (b) the dust formed in the post-shock gas of individual clumps, but was quickly incorporated into the rapidly expanding SN ejecta as the parent clumps were ablated and destroyed.  At both epochs the dust in both locations is heated primarily by the constant luminosity from CSM interaction.  (4)  The photometric and spectroscopic evolution of SN~2005ip was most similar to that of the strongly interacting and X-ray/radio-bright SN~IIn 1988Z, especially with regard to the evolution of its H$\\alpha$ profile.  The close comparison holds with the exceptions that SN~2005ip was less luminous, had a higher $L_X$/$L_{Bol}$ ratio than SN~1988Z, and exhibited stronger narrow coronal lines; moreover, the fastest speeds seen in SN~2005ip persisted longer than in SN~1988Z.  (5)  We propose a model for SN~2005ip that is similar to that of Chugai \\& Danziger (1994) for SN~1988Z, wherein a SN~II plows into a steady but clumpy progenitor wind.  Dense clumps in the CSM persist into the post-shock zone as they are overtaken by the forward shock passing through the low-density region between them, while much slower shocks are driven into individual dense clumps giving rise to the intermediate-width components of the H$\\alpha$ and He~{\\sc i} lines. In this model, the differences between SNe~2005ip and 1988Z noted in point (4) above can be explained if the progenitor of SN~2005ip had a mass-loss rate a factor of $\\sim$5 lower than that of SN~1988Z, making its CSM interaction region less massive and less optically thick.  The lower optical depth allows X-rays generated in the CSM interaction to thoroughly ionize unshocked CSM gas, giving rise to the coronal spectrum.  (6) Given the rarity of coronal spectra like that of SN~2005ip, and in light of the similarity of its inferred progenitor's wind to the extreme conditions observed in the winds of the most luminous known RSGs such as VY~CMa, the most straightforward interpretation seems to be that moderate-luminosity SNe~IIn like SN~2005ip arise from extreme RSGs having relatively high initial masses of 20--40 M$_{\\odot}$. Since the fairly modest luminosity of SN~2005ip is apparently near the limit of what can be achieved by interaction with the densest steady stellar winds known, this finding underscores the requirement that exceptionally luminous SNe~IIn such as SNe 2006tf and 2006gy require much more extreme pre-SN mass ejections analogous to massive LBV eruptions (Smith et al.\\ 2007, 2008b).  At the same time, SN~2005ip demonstrates that not necessarily {\\it all} core-collapse SNe~IIn require this type of LBV-like mass ejection."
    },
    "0809/0809.1092_arXiv.txt": {
        "abstract": "Estimation of the angular power spectrum of the Cosmic Microwave Background (CMB) on a small patch of sky is usually plagued by serious spectral leakage, specially when the map has a hard edge. Even on a full sky map, point source masks can alias power from large scales to small scales producing excess variance at high multipoles. We describe a new fast, simple and local method for estimation of power spectra on small patches of the sky that minimizes spectral leakage and reduces the variance of the spectral estimate. For example, when compared with the standard uniform sampling approach on a $8$ degree $\\times$ $8$ degree patch of the sky with $2\\%$ area masked due to point sources, our estimator halves the errorbars at $\\ell=2000$ and achieves a more than fourfold reduction in  errorbars at $\\ell=3500$. Thus, a properly analyzed experiment will have errorbars at $\\ell=3500$ equivalent to those of an experiment analyzed with the now standard technique with $\\sim 16-25$ times the integration time. ",
        "introduction": "Cosmic Microwave Background (CMB) is a statistically isotropic \\citep{HS06}  and Gaussian \\citep{Komatsu} random field. If we ignore secondary effects, all of the information in high resolution CMB maps is encoded in the angular correlation function or equivalently, in the angular power spectrum, $\\cl$. The angular power spectrum is widely used to estimate the cosmological parameters.  Accurate measurements of  angular power spectrum are needed for precise estimation of cosmological parameters.\\par Over the past decade, CMB power spectrum has been measured over a large range of multipoles, $\\ell$,  by various groups \\citep{COBE, Nolta, Hinshaw, ACBAR, QUAD}. And more experiments are under way to measure the $C_\\ell$ on smaller scales with high accuracy \\citep{ACT, SPT, Planck}. Most power spectrum analyses use uniform or noise weighted maps. This performs reasonably well for power spectra that have nearly equal power in equal logarithmic intervals of multipoles,  {\\it i.e.} $\\ell \\leq 1000$ for the CMB. For smaller scales, (larger $\\ell$) this method is non-optimal, as we show in section \\ref{master}. CMB power spectrum estimated from an incomplete sky map is the underlying full-sky  power spectrum convolved with the power spectrum of the mask. This leads to coupling of modes in the estimated power spectrum. For high resolution experiments such as ACT and SPT which will map the small scale  anisotropies of the CMB on small patches of the sky, this mode-mode coupling will be a serious problem. The reason is that CMB power spectrum is very red on those scales (it falls off as $\\ell^{-4}$ at large $\\ell$) and hence is highly vulnerable to the leakage of power due to mode-mode coupling. There are two methods to remedy this:  to taper the map near the sharp edges, and to pre-whiten the CMB power spectrum. In order to minimize the loss of information due to applying a taper to the map, we use the multitaper method \\citep{Percival+93}. This method involves weighting the map with a set of orthonormal functions which are space limited but maximally concentrated in the frequency domain.  Power spectrum of each of these tapered maps is a measurement of the power spectrum of that map with a different amount of mode coupling. Final power spectrum is obtained from a particular linear combination of these tapered power spectra that minimizes the bias in the estimated power spectrum. The use of multiple tapers also reduces the error-bars in the measured power spectrum. \\par Mode coupling is less harmful if the map has a nearly white power spectrum.  Traditionally, an inverse covariance matrix weighting is used in analysis to prewhiten the maps \\citep{Tegmark}. This method works well, but is a computationally expensive, non-local operation and may be complicated to implement, specially for  high resolution experiments \\citep{Dore, Smith}.  We propose a simple and local prewhitening operator in real space (\\S~\\ref{prewhiten}) that is fast to implement and reduces the bias due to the leakage of power. This method prevents the unnecessarily large error bars at $\\ell \\gtrsim 1500$ due to the point source masks. Usually masks  have sharp edges and holes at the positions of point sources. This leads to a mode-coupled power spectrum that is highly biased at large $\\ell$.  Deconvolution of the mode-coupled power spectrum is a well-studied problem in the CMB data analysis literature \\citep{2002:Hivon} and has been applied to many experiments. But deconvolution of a highly biased power spectrum leads to large error bars in the final power spectrum at large $\\ell$. The mode coupling problem will be worse for the upcoming set of CMB experiments as bright point sources will be much more of a limiting foreground at high resolution.\\par As we show in \\S~\\ref{master}, prewhitening followed by the multitaper method for power spectrum estimation reduces the error-bars (specially at large $\\ell$) in the decoupled power spectrum (cf. Figs. \\ref{errorCompare} \\&~\\ref{PWMasteredErrors}).\\par We begin with a review of the multitaper method in one-dimension in \\S~\\ref{review} and discuss the salient features of the method, generalizing it to the two-dimensional case. Next, we discuss the statistical properties of multitaper spectrum estimators. As a simple application,  we demonstrate the method in context of CMB power spectrum estimation in \\S~\\ref{application}. Next, we formulate the prewhitening method (\\S~\\ref{prewhiten}) and apply it to the case of CMB power spectrum estimation in presence of masks. In \\S~\\ref{master}, we describe the algorithm for deconvolving the power spectrum and the implications of the multitaper method and prewhitening in its context. We summarize and conclude in \\S~\\ref{conclusion}. ",
        "conclusions": "} Power spectrum estimation on small sections of the CMB sky is a non-trivial problem due to spectral leakage from the finite nature of the patch, which is further compounded by the application of point source masks. The direct application of standard decorrelation techniques,  like the MASTER algorithm \\citep{2002:Hivon}, to obtain an unbiased estimate of the power spectrum leads to unnecessarily large uncertainties at high multipoles due to the highly biased nature of the pseudo power spectrum at those multipoles. We have put forward two techniques to reduce the uncertainties in the deconvolved power spectrum. First, we have formulated a two-dimensional adaptive multitaper method (AMTM) which produces nearly unbiased pseudo power spectra for maps without point source masks, by minimizing the leakage of power due to the finite size of the patch. This is achieved at the cost of lowered spectral resolution. The deconvolution of the pseudo power spectrum so produced, leads to an unbiased estimate of the true power spectrum that has several times smaller error bars at high multipoles than the deconvolved periodogram. In presence of point source masks, however, this method becomes non-optimal because the pseudo power spectrum estimated even by AMTM is no longer unbiased. To deal with the point source mask, we have put forward a novel way of prewhitening a CMB map, with manifestly local operations which has simple representations in the Fourier space. This operation produces a map, the power spectrum of which has several orders of magnitude lower dynamic range than the original map. This renders the leakage of power due to holes and edges a relatively benign issue for the prewhitened map. If the prewhitening operation can be tuned to make the power spectrum of the map nearly white, a pseudo power spectrum obtained via a simple periodogram may be nearly unbiased and therefore, can be deconvolved to give a precisely unbiased estimate of the power spectrum, thereby avoiding unnecessarily large error bars at large multipoles. If the map cannot be made sufficiently white for a periodogram, an AMTM method can be applied to the prewhitened map to guard against leakage and achieve the same result. We have shown that by applying these methods, one can reduce the error bar in the small scale power spectrum by a factor of $\\sim 4$ at $\\ell\\sim 3500$. This can be translated into a many-fold reduction in the required integration time of a CMB experiment to achieve some target uncertainty on the small scale power spectrum than that dictated by standard techniques. \\appendix \\begin{widetext}"
    },
    "0809/0809.4024_arXiv.txt": {
        "abstract": "We investigate the occurrence of a CME-driven coronal dimming using unique high resolution spectral images of the corona from the Hinode spacecraft. Over the course of the dimming event we observe the dynamic increase of non-thermal line broadening in the 195.12\\AA{} emission line of \\ion{Fe}{12} as the corona opens. As the corona begins to close, refill and brighten, we see a reduction of the non-thermal broadening towards the pre-eruption level. We propose that the dynamic evolution of non-thermal broadening is the result of the growth of \\alfven{} wave amplitudes in the magnetically open rarefied dimming region, compared to the dense closed corona prior to the CME. We suggest, based on this proposition, that, as open magnetic regions, coronal dimmings must act just as coronal holes and be sources of the fast solar wind, but only temporarily. Further, we propose that such a rapid transition in the thermodynamics of the corona to a solar wind state may have an impulsive effect on the CME that initiates the observed dimming. This last point, if correct, poses a significant physical challenge to the sophistication of CME modeling and capturing the essence of the source region thermodynamics necessary to correctly ascertain CME propagation speeds, etc. ",
        "introduction": "Coronal dimmings, or ``transient coronal holes'' as they are sometimes known, have provoked great curiosity in the solar physics community since their initial observation with Skylab \\citep[][]{Rust1976, Rust1983}. They were first noticed as a rapid intensity reduction of the soft X-Ray corona around active regions, and have subsequently been connected to coronal mass ejections \\citep[CMEs; see, e.g.,][]{Forbes2000, Kahler2001}. Indeed, coronal dimmings are now viewed as the residual footprint of the CME in the corona \\citep[e.g.,][]{Thompson2000}, the radio and plasma signatures of which are observed in interplanetary space \\citep[e.g.,][]{Cane1984, Neugebauer1997, Attrill2008}. The connection between dimming and CME has been quantitatively fortified by the recent statistical surveys of \\citet{Reinard2008} and \\citet{Bewsher2008} using EUV instrumentation on the {\\em SOHO} spacecraft \\citep[][]{Fleck1995}. The former of these surveys indicates that at least 50\\% of front-sided CMEs have associated dimming regions, while the latter stressed that the relationship between the two phenomena is one that grows considerably when only narrowband spectroscopic observations are considered. Therefore, rigorously establishing the poorly understood physical connection between CMEs and coronal dimmings using detailed spectroscopic measurement is a must.  We focus our analysis on observations of NOAA AR~10930 from the Extreme-ultraviolet Imaging Spectrometer \\citep[EIS;][]{Culhane2007} on {\\em Hinode} \\citep[][]{Kosugi2007} between 19:00UT December 14 2006 and 06:00UT December 15 2006. This time period saw an X-Class flare and a $\\sim$1000km/s halo CME\\footnote{The CME properties were automatically derived from {\\em SOHO}/LASCO data by the NASA/GSFC CDAW ({\\url http://cdaw.gsfc.nasa.gov/}) and the Royal Observatory of Belgium/SIDC CACTUS \\citep[][\\-- {\\url http://www.sidc.be/cactus/}]{Robbrecht2004} catalogues.} emanating from this complex active region at around 20:12~UT \\citep[relative to SOHO/EIT imaging; ][]{Boudine1995}. In this Letter, we expand on the analysis of \\citet{Harra2007}, exploiting rare detailed spectroscopic measurements of a dimming region. EIS provides a tantalizing look at the dynamic behavior of EUV emission lines, and their non-thermal line widths in particular, over the course of the eruption. The interpretation of the dynamic evolution of the non-thermal line widths presented in this Letter forms a challenge to the rapidly increasing sophistication of numerical CME models, in that they need to cope with the complex thermodynamics of the CME source region.  \\begin{figure} \\epsscale{0.65} \\plotone{f1.eps} \\caption{Contextual pre-CME images of NOAA AR 10930 from {\\em SOHO} and {\\em Hinode}. We show the closest EIT 195\\AA{} image to the start of the first EIS raster - shown inset as the peak intensity of the \\ion{Fe}{12} 195.12\\AA{} emission line. See the online edition of the journal to see a movie of the dimming evolution. \\label{f1}} \\end{figure} ",
        "conclusions": "\\label{discuss} \\citet{Harra2007} reported on the considerable (relative) blue-shifted Doppler velocities ($\\sim$40km/s) that emanated from the dimming region and are indicative of significant coronal outflows while the region is open. We believe that these outflows are real, are consistent with the measurements discussed above, and reinforce our belief that the transient dimmings are really short lived coronal holes.  Unfortunately, the relative nature of the EUV Doppler measurements leaves some ambiguity in the absolute magnitude of these flows. However, we have demonstrated that the non-thermal line widths of these hot coronal emission lines also evolve dynamically, and are very responsive to the large-scale CME-related intensity fluctuations that historically constitute a coronal dimming event.  The rapid growth of the coronal non-thermal line widths appears to be tied to the post-CME evacuation of the dimming region. Further, this rapid growth of non-thermal line widths is followed by a slow decrease, with values beginning to approach their pre-eruption levels - a result of the closing and gradual filling of the corona as the CME cuts its magnetic ties with the Sun \\citep[][]{Reinard2008, Attrill2008}.  In an effort to explain this evolution, we invoke the ubiquitous presence of \\alfvenic{} plasma motions in the chromosphere \\citep[e.g.,][]{DePontieu2007} and corona \\citep[e.g.,][]{Tomczyk2007} and their likely connection to the non-thermal line widths measured in the upper atmosphere \\citep[e.g.,][]{Tomczyk2007,McIntosh2008}. The observed increase in non-thermal line width in the dimming region is consistent with sub-resolution Doppler broadening resulting from the increased amplitude of \\alfven{} waves in an increasingly rarefied plasma. For a constant volumetric energy flux in a uniform, unchanging, magnetic field, the amplitude of undamped \\alfven{} waves should grow, as the density drops, by a factor of $\\delta\\rho^{-1/4}$. Further, as the corona begins to refill, the wave amplitudes should shrink back to their nominal (pre-event closed magnetic topology) value, precisely as we have observed in this instance. We stress that the dynamic behavior of the 195.12\\AA{} non-thermal line widths is mimicked in the other spectrally isolated EIS lines studied for this event, but not shown in this Letter \\citep[][]{McIntosh2009}.  The likely role of \\alfven{} waves in the acceleration of the fast solar wind \\citep[][]{Suzuki2005, DePontieu2007, Cranmer2007, Verdini2007} that originates in coronal holes and their observed connection following the passage of CMEs in interplanetary space \\citep[e.g.,][]{Neugebauer1997} is evocative. If a coronal dimming is indeed a transiently evolving coronal hole, then it {\\em must} carry some, if not all, of the same spectral, thermodynamic and compositional characteristics of its longer lived brethren. This implies that dimming regions are sources of fast wind streams blowing behind--in the magnetic envelopes of--CMEs. Such wind streams should ``switch-on'' on a timescale commensurate with the \\alfven{} crossing time of the region and ``blow'' for a length of time commensurate with that required for the open magnetic flux behind the CME to close again. If this conjecture is indeed true, we would expect that the amount of unsigned magnetic flux in the dimming region should correlate strongly with the speed of the eventual Interplanetary CMEs seen in situ - scaling with the amount of open wave-carrying magnetic flux \\citep[][]{Schwadron2008, McComas2008} - and observed by \\citet{Chen2006}.  The potential link between CMEs, coronal dimmings and induced fast solar wind streams will require an increase in the sophistication of current numerical CME models to capture the rapidly evolving complex thermodynamic state of the lower boundary. Only then will we be able to assess the impact of the following wind stream on the CME. Such measures are essential to correctly estimate the propagation speed of CMEs and necessary for accurate space weather predictions.  We surmise that the combination of the observed effect dictates that the evolution of the CME and the nearly instantaneous release of a fast wind stream are intrinsically coupled. This goes some way to explaining the frequent in situ observations of chromospheric compositional characteristics and the significant \\alfvenic{} modulation that closely follows the CME in interplanetary space \\citep[][]{Neugebauer1997, Skoug2004}. The influence of the dimming region-initiated wind stream on the flight characteristics of the CME remains to be seen. The hypothesis presented is one that we hope to directly test when the Coronal Multi-channel Polarimeter \\citep[CoMP;][]{Tomczyk2007,Tomczyk2008} instrument begins regular observation from the Mees Solar Observatory on Haleakala early in 2009."
    },
    "0809/0809.2602_arXiv.txt": {
        "abstract": "We study a sample of 43 early-type galaxies, selected from the Sloan Digital Sky Survey (SDSS) because they appeared to have velocity dispersions of $\\sigma\\ge$350~kms$^{-1}$. High-resolution photometry in the SDSS $i$ passband using the High-Resolution Channel of the Advanced Camera for Surveys on board the Hubble Space Telescope shows that just less than half of the sample is made up of superpositions of two or three galaxies, so the reported velocity dispersion is incorrect. The other half of the sample is made up of single objects with genuinely large velocity dispersions.  None of these objects has $\\sigma$ larger than $426\\pm 30$~km~s$^{-1}$. These objects define rather different size-, mass- and density-luminosity relations than the bulk of the early-type galaxy population:  for their luminosities, they are the smallest, most massive and densest galaxies in the Universe. Although the slopes of the scaling relations they define are rather different from those of the bulk of the population, they lie approximately parallel to those of the bulk {\\em at fixed $\\sigma$}. This suggests that these objects are simply the large-$\\sigma$ extremes of the early-type population -- they are not otherwise unusual. These objects appear to be of two distinct types: the less luminous ($M_r>-23$) objects are rather flattened, and their properties suggest some amount of rotational support. While this may complicate interpretation of the SDSS velocity dispersion estimate, and hence estimates of their dynamical mass and density, we argue that these objects are extremely dense for their luminosities, suggesting merger histories with abnormally large amounts of gaseous dissipation. The more luminous objects ($M_r<-23$) tend to be round and to lie in or at the centers of clusters.  Their circular isophotes, large velocity dispersions, and environments are consistent with the hypothesis that they are BCGs.  Models in which BCGs form from predominantly radial mergers having little angular momentum predict that they should be prolate.  If viewed along the major axis, such objects would appear to have abnormally large velocity dispersions for their sizes, and to be abnormally round for their luminosities. This is true of the objects in our sample once we account for the fact that the most luminous galaxies ($M_r<-23.5$), and BCGs, become slightly less round with increasing luminosity.  Thus, the shapes of the most luminous galaxies suggest that they formed from radial mergers, and the shapes of the most luminous objects in our big-$\\sigma$ sample suggest that they are the densest of these objects, viewed along the major axis. ",
        "introduction": "The most massive galaxies may place interesting constraints on models of galaxy formation (e.g. De Lucia et al. 2006; Almeida et al. 2007). But which observable one should use as a proxy for mass is debatable. If luminosity is a good proxy, then the Brightest Cluster Galaxies should be the most massive galaxies; this has led to considerable interest in their properties (Scott 1957; Sandage 1976; Thuan \\& Romanishin 1981; Malumuth \\& Kirschner 1981; Hoessel et al. 1987; Schombert 1987, 1988; Oegerle \\& Hoessel 1991; Lauer \\& Postman 1994; Postman \\& Lauer 1995; Crawford et al. 1999; Laine et al. 2003; Lauer et al. 2007; Bernardi et al. 2007a).  \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig1.ps} \\caption{Surface brightness isophotes of the objects in our sample, which provide the basis for determining which objects are singles.  The two panels for each object show a larger $5\\times 5$~arcsec region with logarithmically spaced isophotes, and a $1\\times 1$~arcsec region with linearly spaced isophotes.  Thick solid circle shows the size of the SDSS fiber; structure within this circle is likely to have contributed to the estimated velocity dispersion, whereas structures outside it have not.} \\label{A1} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig2.ps} \\caption{Continued from previous figure.  The assymetry in in the image which is second from top on the right is due to dust.} \\label{A2} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig3.ps} \\caption{Continued from previous figure.} \\label{A3} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig4.ps} \\caption{Continued from previous figure.} \\label{A4} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig5.ps} \\caption{Continued from previous figure.} \\label{A5} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=1.33]{bernardi_A_fig6.ps} \\vspace{-8cm} \\caption{Continued from previous figure.} \\label{A6} \\end{figure*}  However, velocity dispersion is sometimes used as a surrogate for mass (the virial theorem has mass $\\propto R\\sigma^2$), so it is interesting to ask if a sample selected on the basis of velocity dispersion also contains the most massive galaxies. Such a sample is also interesting in view of the fact that black hole mass correlates strongly and tightly with velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000), so galaxies with large $\\sigma$ are expected to host the most massive black holes.  With this in mind, Bernardi et al. (2006) culled a sample of $\\sim 100$ objects with $\\sigma > 350$~km~s$^{-1}$ from the First Data Release of the Sload Digital Sky Survey (SDSS DR1; Abazajian et al. 2003).  The total area from which these objects were selected is about 2000~deg$^2$; for a spatially flat Universe with $\\Omega_m=0.3$ and Hubble constant $H_0 = 70$~km~s$^{-1}$, which we assume in what follows, this corresponds to a comoving volume of about $3.34\\times 10^8$~Mpc$^3$ out to $z=0.3$.  Some of these objects turned out to be objects in superposition, evidence for which came primarily from the spectra (see Bernardi et al. 2006 for details). A random subset of the others, 43 objects in all, was observed with the High Resolution Camera of the Advanced Camera System on board the Hubble Space Telescope (HST-ACS HRC). On the basis of this high-resolution imaging, we have been able to separate out the true singles from those which are objects in superposition.  As a result, we are now able to study the properties of objects with large $\\sigma$.  We describe the HST observations, our identification of the truly single objects, our classification of their isophotal shapes, and how we combine the HST imaging with SDSS data to determine the environments of these objects in Section~\\ref{hst}. The isophotal shapes of these objects are compared with those of BCGs in  Section~\\ref{shapes}. Various scaling relations, size-luminosity, the fundamental plane, etc. are presented in Section~\\ref{scaling}.  While these relations are different from those defined by the bulk of the early-type galaxy population, we show that they can actually be derived from the early-type scaling relations simply by studying what these relations look like at fixed velocity dispersion. These analyses show that our sample appears to contain two distinct types of objects:  the more luminous objects tend to be round, have large sizes as well as velocity dispersions, and tend to be in crowded fields.  Indeed, they are rounder than other objects of similar luminosities, so it is possible that they are prolate, and viewed along the long axis. The less luminous objects are not round, have abnormally small sizes and are not necessarily in crowded fields; even if rotational motions have compromised the SDSS velocity dispersion estimates, these objects may be amongst the densest galaxies for their luminosities. A final section summarizes our findings and discusses some implications of the bimodal distribution we have found, as well as of the fact that no galaxy appears to have a velocity dispersion larger than $426\\pm30$~km~s$^{-1}$.  In a companion paper (Hyde et al. 2008), we study the surface brightness profiles of the objects classified as singles in much more detail, placing them in the context of other HST-based studies of early-type galaxies (e.g. Laine et al. 2003; Ferrarese et al. 2006; Lauer et al. 2007). ",
        "conclusions": "\\label{discuss} We used HST imaging (Figures~\\ref{A1}--\\ref{A6}) to separate genuinely single objects from those which are superpositions in a sample drawn from the SDSS DR1 and chosen to have $\\sigma\\ge 350$~km~s$^{-1}$.   The abundance of these large $\\sigma$ objects that are singles is consistent with that given by extrapolating fits to the SDSS velocity function to higher $\\sigma$ -- there is no `toe' at $\\sigma > 400$~km~s$^{-1}$ (Figure~\\ref{phiv}).  The scaling relations (size-$L$, mass-$L$ and density-$L$) defined by these objects are different from those defined by the bulk of the early-type galaxy population:  for a given luminosity, these objects are amongst the most massive and densest early-type galaxies (Figures~\\ref{sizeL} and~\\ref{OvsL}).  However, these differences can be understood by thinking of this sample as being the large-$\\sigma$ tail of the early-type galaxy population, but not being otherwise unusual.  Table~2 lists the redshifts, luminosities, sizes, velocity dispersions, colors, Mg$_2$ abundances, shapes, and environments of these objects.  (The luminosities, sizes and shapes have been corrected for known problems with SDSS sky-subtraction; Section~\\ref{photocorr}).  These single galaxies with $\\sigma\\ge 350$~km~s$^{-1}$ appear to be of two types:  the more luminous objects ($M_r<-23.5$ or so) are round ($b/a> 0.7$), whereas the less luminous objects are flatter (Figure~\\ref{ba}).  In addition, Hyde et al. (2008) show that the HST-based surface brightness profiles of the low and high luminousity objects are `power-laws' and `cores' (in the language of Faber et al. 1997); the trend with luminosity is consistent with previous HST work on the centres of early-type galaxies, which shows that power-law galaxies may have significant rotation. The cores, in the galaxies which have them, are about a factor of ten smaller than the half-light radii.  However, the cores are not unusually small for the total $L$ or $\\sigma$ (Hyde et al. 2008); this is in contrast to the half-light radii, which are (Figure~\\ref{sizeL}). On the other hand, if the objects we classify as power-laws actually have cores that are below our resolution limit, then the upper limits we can set to the core-size are already at the low-end of the expected sizes for their $\\sigma$, and perhaps also for their $L$ (Hyde et al. 2008).  What do our observations imply for the formation histories of these two populations? At large $L$, these objects tend to be found in crowded fields so they may be BCGs.  If so, then it is likely that they formed from merging or accreting smaller galaxies, and this would explain why their inner profiles are shallow.  The puzzle is to explain why these objects are so much denser than the average galaxy or BCG of the same $L$ (bottom panels of Figures~\\ref{OvsL} and~\\ref{FP}), especially since the sizes of the inner core radii appear to be normal -- they do not appear to be small for their $L$ or $\\sigma$ (Hyde et al. 2008). One possibility is that these objects formed from predominantly radial mergers with little angular momentum.  If the prolate objects which result are viewed along the long axis, this would produce slightly smaller half-light radii and slightly larger velocity dispersions.  In this context, it is worth noting that the trend for the mean shape to become increasingly round as luminosity increases (e.g., Vincent \\& Ryden 2005) appears to reverse at $M_r<-23.5$ (Figures~\\ref{bazoom} and~\\ref{baLS}).  This appears to be true for galaxies in the main sample, and for BCGs -- and is in qualitative agreement with models which postulate that radial mergers are common at large luminosities.  Compared to this decrease in $b/a$ at large $L$, the luminous objects in our big-$\\sigma$ sample appear to be rounder than expected (Figure~\\ref{bazoom}), consistent with the hypothesis that projection effects have resulted in smaller sizes and larger velocity dispersions. Nevertheless, even if one accounts for this effect, these objects are amongst the densest BCGs for their luminosities.  Whereas projection effects may be important at large $L$, they almost certainly cannot account for the flattened shapes we see at low $L$. We can think of two plausible models for these flattened objects.  \\\\ i)  These are objects which contain a substantial component that is rotationally supported.  \\\\ ii) These are objects in which gaseous dissipation has been most efficient.\\\\ Option (i) must matter for the objects with $b/a\\le 0.6$, since anisotropic dispersions cannot produce extremely flat shapes. Indeed, the fact that we see small $b/a$ only at small $L$ (Figure~\\ref{ba}) suggests that these are examples of the low luminosity, fast-rotator population which has received considerable recent attention from e.g., the SAURON group (Cappellari et al. 2007). If so, then the SDSS velocity dispersions are artificially broadened by rotational motions, thus affecting the mass and density estimates. Reducing the mass estimates by the appropriate $b/a$-dependent factor would bring the flattened objects closer to the relation defined by the bulk of the population; it would also bring the objects closer to the $\\kappa_1+\\kappa_2=8$ boundary in $\\kappa-$space (Figure~\\ref{FP}).  Further indirect evidence for rotation comes from Hyde et al. (2008) who show that these objects have `power-law' inner profiles, and significant amounts of dust.  Previous work has shown that such objects often have a significant rotational component (e.g. Laine et al. 2003). Determining if this is the case for our sample requires spatially resolved kinematics.  Nevertheless, we argue that even if one accounts for rotational motions, these objects are likely to remain the densest for their luminosities.  Thus, it appears that both (i) and (ii) are true for these objects (e.g., Figure~\\ref{fig:mg2se} and related discussion).  If the lower luminosity objects are indeed contaminated by rotation whereas the higher luminosity objects are not, then one might ask if the fact that both low and high luminosity objects define the same power-law scaling relations (Figures~\\ref{sizeL}--\\ref{kappa}) is simply fortuitous.  However, both rotational and random motions contribute to the kinetic energy in the virial theorem.  E.g., Appendix~B of Bender, Burstein \\& Faber (1992) suggests that using only the true $\\sigma$ of an isotropic oblate rotator underestimates the true mass by 35\\% if $b/a=0.6$.  In addition, recent work on the velocity dispersion estimates of more distant objects suggests that including the effect of rotation on mass estimates may indeed be important (e.g. van der Vel \\& van der Marel 2008).  So it may be that the contribution of ordered motions to SDSS velocity dispersion estimates helps to keep the scaling relations power-laws. In this regard, when comparing our sample of high-velocity dispersion galaxies with higher redshift z$\\sim$1.5 samples of passive galaxies (Figure~19 in Cimatti et al. 2008), the lower-luminosity galaxies in our sample populate a similar locus in the size, mass, surface density plane as the superdense z$\\sim$1.5 passive galaxies. It is possible that our low-redshift high-density galaxies are the rare examples of the high-redshift superdense galaxies which have not undergone any dry merging. This scenario is supported by the fact that the low luminosity galaxies in our sample are in low-density environments and have intact power-law centers. So it would be interesting to check if the superdense z$\\sim$1.5 galaxies are ``fast-rotators''.  It is common to predict black hole abundances by transforming an observed luminosity or velocity dispersion function using an assumed scaling relation between black hole mass $M_\\bullet$ and galaxy $L$ or $\\sigma$ (e.g. Lauer et al. 2007; Tundo et al. 2007; Shankar et al. 2008).  If one ignores the fact that there is scatter around the mean $M_\\bullet-L$ or $M_\\bullet-\\sigma$ relations, then one might conclude that the big-$\\sigma$ sample studied here would predict higher black hole masses from their $\\sigma$ than from their $L$s.  However, the scatter is significant: the analysis in Bernardi et al. (2007b) shows how to include the possibility that the scatter in the $\\sigma-L$ relation is correlated with scatter in the $M_\\bullet-L$ and $M_\\bullet-\\sigma$ relations. See their Section 2.3 for a discussion of the effect of selecting objects with large $\\sigma$ for their $L$.  Finally, we note that although we have focussed on the single objects in this sample, the superpositions are interesting in their own right. Because they are close superpositions in both angle and redshift, in which the spectra show little or no sign of recent star formation, and because they provide information about smaller scales than is possible with ground based data, they can be combined with other HST-based samples of early-type galaxies (e.g. Laine et al. 2003; Lauer et al. 2007) to constrain dry-merger rates more precisely than previously possible (e.g. Bell et al. 2006; Masjedi et al. 2007; Wake et al. 2008). Such combined samples can also be used to build more realistic models of the expected configurations of multiple-lens systems. These studies are in progress.   \\small \\begin{table*} \\caption[]{Properties of the 23 objects identified as singles. Superscript $c$ means the inner profile is shallow core (from Hyde et al. 2008).  Magnitudes, sizes, color and $b/a$ were computed by Hyde et al. (2008) on SDSS r-band images, while velocity dispersions and Mg$_2$ index-strenghts are from Bernardi et al. (2006).  } \\begin{tabular}{ccccccccccccccc} \\hline &&&\\\\ ID$_{\\rm S}$ & $z$ & $M_r$ & $e_M$ & $g-r$ & $e_{g-r}$ & log$_{10} R$ & $e_R$ & $\\sigma$ & $e_\\sigma$ & Mg$_2$ & $e_{{\\rm Mg}_2}$ & $b/a$ & $e_{b/a}$  & Env \\\\ & & [mag] & [mag] & [mag] & [mag] & [kpc] & [kpc] & kms$^{-1}$ & kms$^{-1}$ & [mag] & [mag] & & & \\\\ \\hline &&&\\\\ 1 & 0.23949 & -23.40 & 0.06 & 0.74 & 0.04 &  1.00 &  0.03 & 360 &  37 &         -- &       -- &  0.74 &  0.04 & 0\\\\ $2^c$ & 0.19827 & -23.31 & 0.04 & 0.84 & 0.03 &  0.93 &  0.02 & 367 &  28 &         -- &       -- &  0.78 &  0.03 & 1\\\\ $3^c$ & 0.16773 & -23.59 & 0.03 & 0.84 & 0.02 &  1.09 &  0.02 & 367 &  29 &      0.307 &    0.010 &  0.80 &  0.02 & 0\\\\ $4^c$ & 0.32787 & -24.16 & 0.06 & 0.78 & 0.05 &  1.32 &  0.03 & 366 &  52 &         -- &       -- &  0.78 &  0.04 & 1\\\\ $5^c$ & 0.27775 & -24.09 & 0.05 & 0.83 & 0.04 &  1.24 &  0.03 & 371 &  30 &      0.324 &    0.009 &  0.84 &  0.04 & 1\\\\ $6^c$ & 0.21356 & -23.67 & 0.04 & 0.83 & 0.03 &  0.89 &  0.02 & 374 &  22 &      0.314 &    0.008 &  0.78 &  0.02 & 0\\\\ $7^c$ & 0.26271 & -24.38 & 0.04 & 0.77 & 0.04 &  1.36 &  0.02 & 372 &  29 &      0.310 &    0.009 &  0.73 &  0.03 & 1\\\\ $8^c$ & 0.24639 & -23.63 & 0.05 & 0.78 & 0.04 &  1.04 &  0.03 & 377 &  33 &         -- &       -- &  0.79 &  0.03 & 0\\\\ 9 & 0.20533 & -23.09 & 0.06 & 0.88 & 0.04 &  0.86 &  0.03 & 380 &  34 &         -- &       -- &  0.63 &  0.03 & 0\\\\ 10 & 0.22792 & -23.08 & 0.05 & 0.86 & 0.04 &  0.71 &  0.03 & 382 &  27 &         -- &       -- &  0.69 &  0.04 & 0\\\\ 11 & 0.15939 & -22.39 & 0.04 & 0.96 & 0.03 &  0.72 &  0.02 & 383 &  41 &         -- &       -- &  0.48 &  0.02 & 0\\\\ 12 & 0.15335 & -22.10 & 0.05 & 0.80 & 0.03 &  0.39 &  0.03 & 383 &  28 &      0.321 &    0.008 &  0.47 &  0.03 & 0\\\\ $13^c$ & 0.26275 & -24.20 & 0.04 & 0.81 & 0.03 &  1.24 &  0.02 & 383 &  34 &      0.316 &    0.009 &  0.63 &  0.02 & 1\\\\ $14^c$ & 0.23073 & -23.99 & 0.03 & 0.81 & 0.03 &  1.06 &  0.02 & 384 &  32 &      0.314 &    0.009 &  0.85 &  0.03 & 1\\\\ 15 & 0.21930 & -23.00 & 0.06 & 0.81 & 0.04 &  0.71 &  0.03 & 387 &  41 &         -- &       -- &  0.67 &  0.04 & 1\\\\ 16 & 0.28489 & -23.62 & 0.07 & 0.80 & 0.05 &  0.96 &  0.04 & 390 &  44 &         -- &       -- &  0.73 &  0.04 & 0\\\\ 17 & 0.12705 & -21.63 & 0.05 & 0.82 & 0.03 &  0.34 &  0.02 & 400 &  28 &      0.338 &     0.010 &  0.44 &  0.02 & 0\\\\ $18^c$ & 0.26965 & -24.20 & 0.05 & 0.77 & 0.03 &  1.20 &  0.03 & 399 &  35 &      0.315 &    0.009 &  0.87 &  0.03 & 1\\\\ 19 & 0.11610 & -21.85 & 0.03 & 0.83 & 0.02 &  0.20 &  0.02 & 412 &  27 &      0.351 &    0.009 &  0.61 &  0.02 & 1\\\\ 20 & 0.16037 & -22.44 & 0.04 & 0.88 & 0.03 &  0.53 &  0.02 & 405 &  26 &      0.371 &    0.009 &  0.60 &  0.03 & 0\\\\ $21^c$ & 0.29718 & -24.35 & 0.04 & 0.82 & 0.04 &  1.10 &  0.02 & 412 &  27 &      0.327 &    0.007 &  0.93 &  0.03 & 1\\\\ $22^c$ & 0.13343 & -22.74 & 0.02 & 0.81 & 0.01 &  0.63 &  0.01 & 423 &  31 &      0.367 &     0.010 &  0.70 &  0.02 & 1\\\\ $23^c$ & 0.25026 & -24.41 & 0.04 & 0.74 & 0.03 &  1.35 &  0.02 & 424 &  30 &      0.326 &    0.006 &  0.79 &  0.02 & 1\\\\ \\hline &&&\\\\ \\label{tab:singles} \\end{tabular} \\end{table*} \\normalsize"
    },
    "0809/0809.2091_arXiv.txt": {
        "abstract": "Recent galaxy evolution models suggest that feedback from Active Galactic Nuclei (AGN) may be responsible for suppressing star formation in their host galaxies and the subsequent migration of these systems onto the red sequence. To investigate the role of AGN in driving the evolution of their hosts, we have carried out a study of the environments and optical properties of galaxies harboring X-ray luminous AGN in the Cl1604 supercluster at $z\\sim0.9$.  Making use of \\emph{Chandra}, \\emph{HST}/ACS and \\emph{Keck}/DEIMOS observations, we examine the integrated colors, morphologies and spectral properties of nine moderate-luminosity ($L_{\\rm X} \\sim 10^{43}$ erg s$^{-1}$) type 2 Seyferts detected in the Cl1604 complex. We find that the AGN are predominantly hosted by luminous spheroids and/or bulge dominated galaxies which have colors that place them in the valley between the blue cloud and red sequence in color-magnitude space, consistent with predictions that AGN hosts should constitute a transition population.  Half of the hosts have bluer overall colors as a result of blue resolved cores in otherwise red spheroids and a majority show signs of recent or pending interactions.  We also find a substantial number exhibit strong Balmer absorption features indicative of post-starburst galaxies, despite the fact that we detect narrow [OII] emission lines in all of the host spectra. If the [OII] lines are due in part to AGN emission, as we suspect, then this result implies that a significant fraction of these galaxies (44\\%) have experienced an enhanced level of star formation within the last $\\sim1$ Gyr which was rapidly suppressed.  Furthermore we observe that the hosts galaxies tend to avoid the densest regions of the supercluster and are instead located in intermediate density environments, such as the infall region of a massive cluster or in poorer systems undergoing assembly. Overall we find that the properties of the nine host galaxies are generally consistent with a scenario in which recent interactions have triggered both increased levels of nuclear activity and an enhancement of centrally concentrated star formation, followed by a rapid truncation of the latter, possibly as a result of feedback from the AGN itself.  Our finding that the hosts of moderate-luminosity AGN within the Cl1604 supercluster are predominantly a transition population suggests AGN feedback may play an important role in accelerating galaxy evolution in large-scale structures. ",
        "introduction": "Several studies have demonstrated that galaxies segregate by rest-frame color, with passively evolving early-type galaxies forming a well defined ``red sequence'' and gas-rich, star forming systems located in a more diffuse ``blue cloud'' (e.g.~Baldry et al.~2004).   This observed color bimodality is commonly thought to be the result of an evolutionary sequence, such that galaxies in the blue cloud quickly migrate onto the red sequence following the termination of star formation, leading to a sparsely populated transition zone between the two regions (Faber et al.~2007; although see Driver et al.~2006 for an alternative view). While simple passive evolution likely contributes to the migration of galaxies from blue to red, recent findings that galaxies exhibit reduced star formation rates as they assemble into regions of even moderate density (Lewis et al.~2002; Gomez et al.~2003) suggest that processes prevalent in richer environments must also play a role in terminating star formation and driving galaxy evolution.  Recently merger-driven feedback from Active Galactic Nuclei (AGN) has been suggested as one such process (Hopkins et al.~2005, 2007; Somerville et al.~2008).  In this scenario strong gravitational torques produced as a result of galaxy mergers funnel material to the nuclear region of a system, leading to elevated accretion onto the central black hole and enhanced star formation in the form of a nuclear starburst (Barnes \\& Hernquist 1991; Mihos \\& Hernquist 1996).  If sufficiently energetic, the AGN eventually drives outflows that disrupt the host, effectively quenching ongoing star formation by removing the galaxy's gas supply (Springel et al.~2005; Hopkins et al.~2005).   As the last generation of newly formed stars fade, the galaxy quickly migrates from the blue cloud to the red sequence, with the activity of the central black hole gradually declining over the same period.  This feedback mechanism provides a method by which the overall properties of a galaxy can be regulated by its central black hole. Indeed there is a growing body of evidence that suggests AGN are directly linked to the evolution of their host galaxies.  This includes the observed correlation between central black hole mass and the stellar velocity dispersion of a host galaxy's bulge (Gebhardt et al.~2000; Tremaine et al.~2002) and the coeval decline in both the cosmic star formation rate and nuclear activity since $z\\sim1$ (Boyle \\& Terlevich 1998).  If AGN are directly responsible for driving the evolution of their hosts, then signs of recent or ongoing transformation should be present in these galaxies. This includes disturbed morphologies due to recent interactions, post-starburst signatures resulting from rapidly quenched star formation following a merger-induced enhancement, and colors consistent with the sparsely populated transition zone in the color-magnitude diagram. While many of these features are observed in the hosts of luminous quasars (QSO), studies of moderate-luminosity Seyferts ($L_{\\rm X}\\sim10^{43}$ erg s$^{-1}$) in this regard have produced mixed results.   For example, while such AGN are predominantly found in massive, early-type galaxies which have younger stellar populations than their non-active counterparts (e.g. Kauffmann et al.~2003), only a small fraction show signs of recent merger activity (De Robertis et al.~1998; Grogin et al.~2005; Pierce et al.~2007).   Likewise, although AGN are found in post-starburst galaxies more often than in normal galaxies (Heckman 2004, Yan et al.~2006, Goto et al.~2006), the fraction of Seyferts associated with post-starburst hosts at low redshift remains fairly small; on the order of $4\\%$ according to a recent study by Goto et al.~(2006).  A clearer picture of the prevalence of AGN feedback as an evolutionary driver may be obtained by studying host galaxies in and around large-scale structures at high redshift.  In addition to a considerably higher AGN density and galaxy merger rate at $z\\sim1$ relative to the present epoch (Barger et al.~2005; Kartaltepe et al.~2007), there are indications that AGN are triggered more often in structures than in the field (Gilli et al~2003; Cappelluti et al~2005; Eastman et al.~2007; Kocevski et al.~2008), presumably due to the presence of richer, dynamic environments where galaxy interactions are more common (Cavaliere et al.~1992). Several recent wide-field surveys have found that the colors of higher-redshift host galaxies are consistent with a population in transition (Sanchez et al.~2004; Nandra et al.~2007; Georgakakis et al.~2008; Silverman et al.~2008; although see also Westoby et al.~2007) and there are indications the fraction of AGN hosts exhibiting post-starburst signatures is higher at $z\\sim0.8$ (Georgakakis et al.~2008).  Furthermore, Silverman et al.~(2008) found that a majority of host galaxies associated with two large-scale structures in the Extended \\emph{Chandra} Deep Field South (E-CDF-S) at $z\\sim0.7$ have colors consistent with the transition zone of the color-magnitude diagram.  Coupled with the increased incidence of AGN in such regions, this finding implies that AGN related feedback may play an important role in accelerating galaxy evolution in large-scale structures.  In this study we examine the environments and optical properties of galaxies hosting X-ray selected, moderate-luminosity AGN in the Cl1604 supercluster at $z=0.9$.  The system is the largest structure mapped at redshifts approaching unity, consisting of eight spectroscopically confirmed galaxy clusters and groups and a rich network of filamentary structures.  The system spans roughly 10 $h_{70}^{-1}$ Mpc on the sky and 100 $h_{70}^{-1}$ Mpc in depth. The complex structure of the supercluster, as mapped by extensive spectroscopic observations, is described in Gal et al.~(2008; hereafter G08) and our X-ray observations of the system are discussed in Kocevski et al.~(2008; hereafter K08).  The Cl1604 supercluster is well suited for this study as it provides a diverse set of  structures within which AGN may be preferentially found and a wide range of environments and local conditions that can help constrain the mechanisms responsible for triggering their activity.  In the following sections we begin by describing our X-ray and optical observations of the Cl1604 system (\\S2) and the process by which we identify AGN in the supercluster (\\S3).  This is followed by an examination of the morphology (\\S4.1), optical colors (\\S4.2), spectral properties (\\S4.3), and environments (\\S4.4) of the galaxies found to host AGN.  We also determine the fraction supercluster members which harbor AGN and compare this to low-redshift results in \\S5.  We discuss our results and their implications for the AGN feedback model in driving galaxy evolution in the Cl1604 system in \\S6.  Finally we summarize our results in \\S7.  Throughout this paper all quoted line equivalent widths are in the rest frame, magnitudes are in the AB system and we assume a $\\Lambda$CDM cosmology with $\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_{0} = 70$ $h_{70}$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In summary, we find that the X-ray selected AGN detected in the Cl1604 supercluster are largely hosted by luminous ($M_{V} < -21.1$), spheroidal and/or bulge dominated galaxies which are bluer than similar galaxies on the system's red sequence. The integrated colors of a majority of the hosts (5/8) place them in the valley between the supercluster red sequence and the more diffuse blue cloud. In half of the hosts, the bluer overall colors are the result of blue resolved cores in otherwise red spheroids.  Given that the supercluster AGN luminosities do not reach the QSO level and the lack of obvious AGN contamination in their optical spectra, we interpret the blue cores as the result of recent star formation.  Many of these galaxies do show signs of starburst activity within the last $\\sim1$ Gyr as $\\sim44\\%$ (4/9) of the host spectra exhibit either strong H$-\\delta$ absorption or a pronounced Balmer series.  We also find a majority of hosts (6/9) show signs of recent or pending interactions, a possible indication of their triggering mechanism.  With regard to environment, we observe that the hosts tend to avoid the densest regions of the supercluster and are instead located in intermediate density environments, such as the infall region of a massive cluster or in poorer systems undergoing assembly.  Finally, we measure an increased fraction of supercluster members that harbor AGN relative to low-redshift ($z<0.3$) clusters, but we do not observe a substantial change from the fraction reported at $z\\sim0.6$.  In general our observations are fairly consistent with theories that link AGN activity to galaxies undergoing transformation from star forming systems in the blue cloud to passively evolving galaxies on the red sequence.   This is evidenced by the fact that (i) a significant fraction of the AGN hosts are located in the transition zone of the color-magnitude diagram, which is normally sparsely populated, (ii) several hosts show evidence of a recent enhancement of star formation that was abruptly terminated and (iii) there are indications of recent merger activity in many of the hosts, which are located in moderate-density environments thought to be conducive to galaxy interactions. An alternative theory that is also consistent with our observations is one of episodic star formation and AGN activity in red sequence galaxies triggered by galaxy interactions.  In this scenario, AGN hosts in the supercluster are not in fact transitioning from blue to red, but instead from low to high luminosity (mass) on the red sequence.  Many of these conclusions, especially those regarding the post-starburst-AGN connection, depend on our assumption that the [OII] emission lines seen in the host galaxies are due in part to AGN emission.  We plan to test this assumption in the near future with near-infrared spectroscopic observations of each AGN host in order to obtain H$\\alpha$ and [NII] EWs and line ratios.  The observations of one host which have already been completed have confirmed that AGN emission contributes to the galaxies [OII] line flux, leading to an overestimated SFR.  If we find similar results for the remaining hosts, this will help confirm that a significant fraction of AGN in the Cl1604 supercluster have experienced a recent and rapid truncation of their star formation activity."
    },
    "0809/0809.4742_arXiv.txt": {
        "abstract": "The first set of supermassive black hole mass estimates, published from 1977 to 1984 by \\'{E}. A. Dibai, are shown to be in excellent agreement with recent reverberation-mapping estimates.  Comparison of the masses of 17 AGNs covering a mass range from about  $10^6$ to $10^9 M_{\\sun}$ shows that the Dibai mass estimates agree with reverberation-mapping mass estimates to significantly better than $\\pm 0.3$ dex and were, on average, only 0.14 dex ($\\sim 40$\\%) systematically lower than masses obtained  from reverberation mapping.  This surprising agreement with the results of over a quarter of a century ago has important implication for the structure and kinematics of AGNs and implies that type-1 AGNs are very similar. Our results give strong support to the use of the single-epoch-spectrum (Dibai) method for investigating the co-evolution of supermassive black holes and their host galaxies. ",
        "introduction": "It is exactly 100 years this year since the publication of the first evidence of nuclear activity in galaxies \\citep{fath1908}. Over the last half century or so, active galactic nuclei (AGNs) have been the subject of increasingly intensive study which has resulted in tens of thousands of papers being published. \\citet{zeldovich64} and \\citet{salpeter64} proposed that the huge energy release lasting for millions of years from a typical AGN could be explained by the accretion of matter onto a supermassive black hole. In such accretion, the energy output efficiency may reach 43\\% of the accreting matter's rest mass energy.  This gave the first estimates of lower limits to the masses of AGNs because the luminosity, $L$, cannot go much above the Eddington limit, $L_{Edd} \\sim 1.3 \\times 10^{38} (M/M_{\\sun})$ erg/s, so the mass of the black hole in an AGN has to be of the order of several million to several billion solar masses, depending on the luminosity of the AGN \\citep{zeldovich+novikov64}.   While the Eddington limit gives a lower limit to the mass, $M_{BH}$, of the black hole, the actual mass could be orders of magnitude greater than this. To understand the working of AGNs, $M_{BH}$ needs to be determined observationally, so estimating $M_{BH}$ has always been considered to be a matter of utmost importance in AGN studies. If the motions of gas clouds are dominated by gravity, masses can be estimated in principle from the virial theorem if we know a velocity and an appropriate distance from the center (e.g., \\citealt{woltjer59}). The velocity along the line of sight can easily be determined from the Doppler broadening of emission lines, but determining the distance of the emitting material from the central object is difficult.  It was not until the work of \\citet{dibai77,dibai78,dibai80,dibai81,dibai84a,dibai84b} that an attempt was made at a consistent spectroscopic determination of the masses of the central objects for a large sample of AGNs.  This enabled Dibai to determine black hole masses and Eddington ratios, $L/L_{Edd}$, for dozens of AGNs for the first time \\citep{dibai80,dibai84b}, and to begin investigations in what promised to be (and indeed has turned out to be) a very fruitful area: the relationships between these quantities and other AGN parameters (see, for example, \\citealt{dibai84a} and \\citealt{dibai+zasov85}). Unfortunately, this work was cut short by Dibai's premature death, and in the two decades after his 1977 paper there were only a few papers by others using the Dibai method to estimate black hole masses (e.g., \\citealt{joly+85,wandel+yahil85,padovani+rafanelli88}).  In the last decade, however, there has been an enormous growth of interest in determining masses by Dibai's method because of the close relationship between the masses of black holes and the masses of the bulges of their host galaxies (see \\citealt{kormendy+gebhardt01} for a review). It is only through the Dibai method that large numbers of black hole masses in AGNs can currently be determined, and the method has already been used in making tens of thousands of mass estimates for AGNs of all redshifts in the Sloan Digital Sky Survey (e.g., \\citealt{mcclure+dunlop04,salviander+07,greene+ho07,shen+08}). It is therefore of interest to revisit the original Dibai mass estimates and see how they compare with more recent independent estimates.  \\citet{dibai+pronik67} showed that the emitting regions of the broad and narrow components of AGN optical emission lines (what we now call the BLR and NLR respectively) had to originate in different locations in space, and that the emitting gas, in both cases, had to have a cloudy structure, that is, to fill only a small fraction, $\\epsilon$, of the volume of the corresponding region. It has long been recognized (e.g., \\citealt{bahcall+72}; see also \\citealt{bochkarev+pudenko75}) that BLR variability timescales could be used to estimate the effective distance of the emitting gas from the black hole, and hence to get masses via the virial theorem.  \\citet{pronik+chuvaev72} presented the first long-term H$\\beta$ light curve for an AGN.  \\citet{lyutyi+cherepashchuk72} and \\citet{cherepashchuk+lyutyi73} made narrow-band observations of three AGNs over several months and published the first actual estimates of BLR sizes from the time lag between continuum variability and line variability. \\citet{bochkarev+antokhin82}, \\citet{blandford+mckee82}, \\citet{capriotti+82}, and \\citet{antokhin+bochkarev83} independently developed methods for recovering information about the BLR from the response of the lines to continuum variations, a subject which has now become known as ``reverberation mapping''. Determining BLR sizes became practical and widespread with the introduction of cross-correlation techniques \\citep{gaskell+sparke86,gaskell+peterson87} a few years later.  Once BLR radii were being reliably obtained from cross-correlation studies, the main remaining problem in estimating masses was in establishing that the gas motions were dominated by gravity.  The dominant belief from the early days of AGN studies (e.g., \\citealt{burbidge58}) was that emission-line gas was outflowing from AGNs, and the virial theorem obviously cannot be applied to an outflowing wind. The idea of determining BLR kinematics diagnostics by line-profile variations was first brought forward by S.~N. \\citet{fabrika80}, then Dibai's graduate student. The first velocity-resolved reverberation mapping (\\citealt{gaskell88,koratkar+gaskell89}; see also \\citealt{shapovalova+01a, shapovalova+01b}) showed that the BLR was not outflowing, but instead showing some net inflow in combination with Keplerian and/or chaotic motion (see \\citealt{gaskell+goosmann08}). This discovery immediately permitted the first reverberation-mapping determinations of black hole masses \\citep{gaskell88}.  Despite the promising results of the pioneering observations of \\citet{lyutyi+cherepashchuk72} and \\citet{cherepashchuk+lyutyi73}, it was emphasized by \\citet{bochkarev84}, confirming earlier calculations of \\citet{bochkarev+antokhin82} and \\citet{antokhin+bochkarev83}, that reliable results could only be obtained from of long time-sequences of well-sampled, high-accuracy, spectral and photometric observational data obtained through a large-scale international project.  The urgent need for starting such collaborations was discussed in detail by \\citet{bochkarev87a,bochkarev87b}.  The following decades saw the progress of the {\\it International AGN Watch} ({\\it IAW}) program (see \\citealt{clavel+91,peterson+91}, et seq.) aimed at determining the BLR size and structure from reverberation mapping. That project proved to be one of the largest global astronomical monitoring programs to date. More than 200 astronomers from 35 countries cooperated for 15 years in accumulating long, densely-sampled, UV, optical, and other wavelength time series for many AGN. As a result of this, and of additional optical monitoring campaigns, reverberation-mapping estimates of the AGN central black hole mass have now been obtained for over 40 AGNs (see \\citealt{peterson+04} and \\citealt{vestergaard+peterson06}).  In this paper we make a comparison of AGN central object mass values from reverberation mapping campaigns with the masses obtained by Dibai over a quarter of a century ago.  In Section 2 we briefly describe the assumptions made by Dibai.  In Section 3 we compare the results yielded by the two methods. Section 4 discusses the implications of the comparison. ",
        "conclusions": ""
    },
    "0809/0809.4268_arXiv.txt": {
        "abstract": "We review our understanding of the kinematics of the LMC and the SMC, and their orbit around the Milky Way. The line-of-sight velocity fields of both the LMC and SMC have been mapped with high accuracy using thousands of discrete traces, as well as HI gas. The LMC is a rotating disk for which the viewing angles have been well-established using various methods. The disk is elliptical in its disk plane. The disk thickness varies depending on the tracer population, with $V/\\sigma$ ranging from $\\sim 2$--10 from the oldest to the youngest population. For the SMC, the old stellar population resides in a spheroidal distribution with considerable line-of-sight depth and low $V/\\sigma$. Young stars and HI gas reside in a more irregular rotating disk. Mass estimates based on the kinematics indicate that each Cloud is embedded in a dark halo. Proper motion measurements with HST show that both galaxies move significantly more rapidly around the Milky Way than previously believed. This indicates that for a canonical $10^{12} \\Msun$ Milky Way the Clouds are only passing by us for the first time. Although a higher Milky Way mass yields a bound orbit, this orbit is still very different from what has been previously assumed in models of the Magellanic Stream. Hence, much of our understanding of the history of the Magellanic System and the formation of the Magellanic Stream may need to be revised. The accuracy of the proper motion data is insufficient to say whether or not the LMC and SMC are bound to each other, but bound orbits do exist within the proper motion error ellipse.\\looseness=-2 ",
        "introduction": "\\label{s:intro}  The Magellanic Clouds are two of the closest galaxies to the Milky Way, with the Large Magellanic Cloud (LMC) at a distance of $\\sim 50 \\kpc$ and the Small Magellanic Cloud (SMC) at $\\sim 62 \\kpc$. Because of their proximity, they are two of the best-studied galaxies in the Universe. As such, they are a benchmark for studies on various topics, including stellar populations and the interstellar medium, microlensing by dark objects, and the cosmological distance scale. As nearby companions of the Milky Way with significant signs of mutual interaction, they have also been taken as examples of hierarchical structure formation in the Universe. For all these applications it is important to have an understanding of the kinematics of the LMC and the SMC, as well the kinematics (i.e., orbit) of their center of mass with respect to the Milky Way and with respect to each other. These topics form the subject of the present review. Other related topics, such as the more general aspects of the structure of the LMC and SMC, the nature of the LMC bar, the possible presence of fore- or background populations, and the large radii extent of the Clouds are not discussed here. The nature, origin, and models of the Magellanic Stream are touched upon only briefly. All these topics are reviewed in others papers in this volume by, e.g., Harris, Majewski, Besla, Bekki, and others. ",
        "conclusions": "\\label{s:conc}  The kinematics of the LMC are now fairly well understood, with velocities of thousands of individual tracers of various types having been fitted in considerable detail with (thick) disk models. Questions that remain open for further study include the reality and origin of kinematical differences between different stellar tracer populations, the differences between the gaseous and stellar kinematics, and the amount and origin of non-equilibrium features in the kinematics. The kinematics of the SMC are understood more poorly, but appear generally consistent with being a spheroidal system of old stars with an embedded irregular disk of gas and young stars.  The HST PM work has provided the most surprising results in recent years, with important implications for both the history of the Magellanic System and the origin of the Magellanic Stream. Of course, it is natural in discussions about this to wonder about the robustness of the observational results. It should be noted in this context that many experimental features and consistency checks are built in that support the general validity of the HST PM results. These include: (1) the use of random telescope orientations causes systematic errors tied to the detector frame to cancel out when averaging over all fields; (2) the final PM errors are based on the observed scatter between fields, with no assumptions about the source and nature of the underlying errors; (3) two groups used different methods to analyze the same data and obtained consistent results; (4) P08 managed to measure a PM rotation curve for the LMC that is broadly consistent with expectation, which would have been impossible if the PM errors were in reality larger than claimed; (5) the difference between the LMC and SMC PMs is more or less consistent with expectation for a binary orbit, which would not generally have been the case if the measurements suffer from unknown systematics; (6) the LMC PM is consistent with expectation based on the line-of-sight velocity field of carbon stars (see Section~\\ref{ss:vtrans}); and (7) the LMC PM leads to an HI velocity field with a straight zero-velocity curve, by contrast to previously assumed values (see Section~\\ref{ss:vtrans}).  One interesting feature in the observational PM results is that with the P08 PM values, there are no bound LMC-SMC orbits, given their different $\\mu_W$ and smaller error bars for the SMC compared to the K06b results. However, the SMC PM is significantly less certain that that for the LMC, due to the smaller number of fields observed with HST, and the fact that most of them were observed at a similar telescope orientation (which implies that potential systematic errors that are fixed in the detector frame do not average out when the results from different fields are combined). This underscores the need for additional PM observations. A third epoch of observations for most fields has already been obtained with HST/WFPC2, and preliminary analysis supports the validity of the results based on the first two epochs (Kallivayalil et al., these proceedings). A fourth epoch is planned with HST/ACS and HST/WFC3 in 2009. With the increased time baselines and use of multiple different instruments it will be possible to further reduce random errors and constrain possible systematic errors. In turn, this will allow new scientific problems to be addressed, such as the internal proper motion kinematics of the Clouds, and their rotational parallax distances (the distances obtained by equating the line-of-sight and proper motion rotation curves)."
    },
    "0809/0809.4574_arXiv.txt": {
        "abstract": "Observations have evidenced that passively evolving massive galaxies at high redshift are much more compact than local galaxies with the same stellar mass. We argue that the observed strong evolution in size is directly related to the quasar feedback, which removes huge amounts of cold gas from the central regions in a Salpeter time, inducing an expansion of the stellar distribution. The new equilibrium configuration, with a size increased by a factor $\\gtrsim 3$, is attained after $\\sim$40 dynamical times, corresponding to $\\sim 2$ Gyr. This means that massive galaxies observed at $z\\geq 1$ will settle on the Fundamental Plane by $z\\sim 0.8$--1. In less massive galaxies ($M_{\\star}\\lesssim 2 \\times 10^{10}\\,M_{\\odot}$), the nuclear feedback is subdominant, and the mass loss is mainly due to stellar winds. In this case, the mass loss timescale is longer than the dynamical time and results in adiabatic expansion that may increase the effective radius by a factor of up to $\\sim 2$ in 10 Gyr, although a growth by a factor of $\\simeq 1.6$ occurs within the first 0.5 Gyr. Since observations are focused on relatively old galaxies, with ages $\\gtrsim 1\\,$Gyr, the evolution for smaller galaxies is more difficult to perceive. Significant evolution of velocity dispersion is predicted for both small and large galaxies. ",
        "introduction": "Several recent observational studies have found that massive, passively evolving, galaxies at $z>1$ are much more compact than local galaxies of analogous stellar mass (Ferguson et al. 2004; Trujillo et al. 2004, 2007; Zirm et al. 2007; Cimatti et al. 2008; Damjanov et al. 2008). Since similarly superdense massive galaxies are extremely rare or absent at $z\\simeq 0$ (Shen et al. 2003) a strong size evolution, by a factor $\\sim 3$ or more, is indicated. No convincing mechanism able to account for such size evolution has been proposed so far. In the following we show that an expansion consistent with the observed one is naturally expected as a consequence of feedback from active nuclei (Silk \\& Rees 1998; Granato et al. 2001, 2004), which is now widely recognized as a crucial ingredient of semi-analytic models (see, e.g., Di Matteo et al. 2008).  In the local Universe spheroidal galaxies occupy a quite narrow region, the Fundamental Plane, in the 3-dimensional space identified by the effective or half-light radius $r_e$, by the central velocity dispersion $\\sigma_0$, and by the mean surface brightness within $r_e$ (Djorgovski \\& Davis 1987). The tight color-magnitude relation, the color-velocity dispersion relation, and spectral line indices imply that the bulk of stars of elliptical galaxies formed at $z\\geq 1.5$. The enhancement of $\\alpha$-elements abundances with respect to iron in massive elliptical galaxies entails that most of their stars formed within the first Gyr of their life (see Renzini 2006 for a comprehensive discussion).  An additional key result is the generic presence of a Super Massive Black Hole (SMBHs) in the center of local elliptical galaxies. Its mass is directly proportional to the mass of the old stellar population, $M_{\\rm BH}\\sim 2\\times 10^{-3}\\,M_{\\star}$ (Magorrian et al 1998; see Ferrarese \\& Ford 2005 for a review), implying that quasars and spheroidal galaxies form and evolve in strict relation and with mutual feedback. Specifically, it has been suggested that the central SMBH  grows until until its feedback unbinds the residual gas in the host galaxy and sweeps it out through a high velocity wind, thus halting both the star formation and its own fueling and establishing a relationship between the SMBH mass and the stellar velocity dispersion (Silk \\& Rees 1998). Strong AGN feedback also appears to be the only viable mechanism to explain the exponential cut-off at the bright end of the galaxy luminosity function (Croton et al. 2006) and the observed bimodality in the color-magnitude diagram of galaxies at $z\\gtrsim 1.5$ (Menci et al. 2006). Direct observational indications of massive outflows close to high redshift quasars, consistent with this scenario, have been reported (e.g. Simcoe et al. 2006; Prochaska \\& Hennawy 2008).  As shown below, the amount of gas rapidly stripped from the central regions of the galactic halo can be large enough to drive a large increase of the galaxy size. The puffing up of a system by rapid mass loss is a well known phenomenon, extensively studied both analytically and through numerical simulations, with reference to galaxies (Biermann \\& Shapiro 1979), and, especially, to globular clusters (Hills 1980; Goodwin \\& Bastian 2006). Slower, adiabatic expansion is caused by mass loss due to stellar winds or supernova explosions (Hills 1980; Richstone \\& Potter 1982).  In this {\\it Letter} we first summarize the effect of mass loss on the galaxy size evolution (\\S\\,2), then we present quantitative estimates on the evolution of the effective radius and of the central stellar velocity dispersion as a function of galactic mass, using the Granato et al. (2004) model as a reference (\\S\\,3), and finally, in \\S\\,4, we summarize and discuss our results. ",
        "conclusions": "The variation of the effective radius  implies that elliptical galaxies spend the initial part of their life outside the 'local' fundamental plane and only subsequently settle on it.  For massive galaxies ($M_\\star\\geq 2\\times 10^{10}\\,M_{\\odot}$), for which the expansion is mostly due to quasar-driven super-winds, Fig.~\\ref{fig:re} displays three steps in the evolution: {\\it i)} the end of the quasar phase at cosmic time $t_{\\rm QSO}\\simeq t_{\\rm vir}+\\Delta t_{\\rm burst}$, which corresponds to the end of the rapid mass loss; {\\it ii)} the intermediate time $t_{\\rm int}$ when the stellar component reaches a new virial equilibrium with a larger size; {\\it iii)} the present time $t_0$, at which a further, minor increase of the size is achieved due to stellar winds. The initial effective radius has been computed from eq.~(\\ref{eq:re}) with $f_{\\sigma}\\approx 1.3$.  Equation~(\\ref{eq:fast}) shows that the size expansion can be extremely large as the fraction of expelled gas approaches 50\\% of the initial mass. However, at variance with the case of star clusters, confined by their own gravitation field, the presence of the DM halo, dynamically dominant at large radii, prevents the disruption and restricts the range of possible final structures of large elliptical galaxies.  The time lapse required to reach the new virial equilibrium can be inferred from numerical simulations of stellar clusters, which show that the system reaches the new equilibrium after about 40 dynamical times (computed for the initial configuration), independently of the initial mass (Geyer \\& Burkert 2001; Goodwin \\& Bastian 2006). By scaling up this result, massive galaxies are expected to reach a new equilibrium  after $\\sim 2\\,$Gyr, at $t_{\\rm int}\\sim t_{\\rm QSO}+2\\,$Gyr (asterisks in Fig.~\\ref{fig:re}). Since then only minor adiabatic mass losses occur, producing slight changes of radius and velocity dispersion (from asterisks to squares). We recall that the quasar activity statistically reaches its maximum at redshift $z\\sim 2$, corresponding to a cosmic time $t_{\\rm cosm}\\sim 3\\,$Gyr. The rapid mass loss is then expected to statistically peak at the same redshift, triggering size variations that stabilize after $\\sim 2\\,$Gyr. Thus the massive galaxies should on the average reach the local size-mass relation at $t_{\\rm cosm}\\sim 5\\,$Gyr, corresponding to $z\\sim 0.8$. This result is in keeping with sizes observed at $z \\lesssim 1$ (Cimatti et al. 2008) and with results of studies of the Fundamental Plane at $z\\leq 0.8$ (Renzini 2006). In the redshift interval $0.8 \\leq z \\leq 3$ a large scatter in the size--stellar mass relation is expected and indeed observed.  This scenario may appear to be contradicted by observations of Ferguson et al. (2004), who found an average size $r_e\\simeq 1.7\\,$kpc for luminous Lyman Break Galaxies (LBGs) at $z\\simeq 4$, while we  would expect  $r_e\\sim 0.5\\,$kpc, inserting the appropriate masses in eq.~(\\ref{eq:re}). However as pointed out by Joung, Chen \\& Bryan (2008), the apparent effective radius could be a factor of about 3 larger than the intrinsic one, because of the presence of dust in the central star forming regions.  In the case of rapid mass loss, the velocity dispersion of the stars at virial equilibrium scales as $\\sigma_{\\star}\\propto r_e^{-1/2}$. As shown by Fig.~\\ref{fig:sigma}, the size evolution illustrated by Fig.~\\ref{fig:re} brings the quite high velocities inferred from the observed values of $M_\\star$ and $r_e$ in the compact phase, to within the locally observed range.  The structural evolution of low mass E galaxies follows a different path, driven by the slow mass loss due to galactic winds and supernova explosions, because their nuclear activity is of low power. Adopting a Chabrier (2005) Initial Mass Function (IMF), we find a mass loss of a factor of $\\sim 2$ on a time scale of several Gyrs (the factor is only $\\sim 1.4$ for a Salpeter (1955) IMF); after eq.~(\\ref{eq:slow}) the size increases by the same factor. However most of the expansion occurs when these galaxies are young. For a Chabrier IMF, already half a Gyr after the baryon collapse the size has increased by a factor of $\\sim 1.6$ (triangles in Fig.~\\ref{fig:re}) and the subsequent expansion is limited to a factor of $\\sim 1.3$ (squares). Thus such galaxies observed at ages $\\gtrsim 0.5\\,$Gyr should exhibit a size quite close to that of local spheroidal galaxies with the same stellar mass. As a consequence, high-$z$ galaxies with $M_\\star < 2\\times 10^{10}\\,M_{\\odot}$ should also exhibit a smaller scatter in the effective radius--stellar mass relation. They are predicted to have significantly more compact sizes only at ages $\\lesssim 0.5\\,$Gyr, when their star formation rates are of tens $M_\\odot$/yr (see Fig.~1 of Lapi et al. 2006), typical of relatively low luminosity, high-$z$ LBGs. In the adiabatic expansion case the velocity dispersion scales as $\\sigma_{\\star}\\propto r^{-1}$. Again, this brings the high initial velocities within the locally observed range.  As apparent from Fig.~\\ref{fig:re}, we predict significant evolution for galaxies with $M_{\\star}\\geq 2\\times 10^{10}$ M$_{\\odot}$, well below the threshold of $M_{\\star}\\geq 5\\times 10^{11}\\,M_{\\odot}$ implied by the semi-analytic model of Khochfar \\& Silk (2006).  In conclusion we suggest that the rapid mass loss driven by the quasar feedback is the main agent of the size and velocity dispersion evolution of massive spheroidal galaxies. Lower-mass galaxies experience a weaker, but non negligible evolution (of amplitude depending on the adopted IMF), due to mass loss mainly powered by supernova explosions; most of it occurs during their active star-formation phase. Although our calculations have been carried out in the framework of the Granato et al. (2004) model, this evolutionary behaviour is a generic property of all models featuring large mass loss due to quasar and/or supernova feedback. Observational evidences, some of which are briefly summarized in \\S\\,\\ref{sect:size}, suggest that quasar driven high velocity outflows may have removed from the central regions of massive galaxies a gas mass comparable to the mass in stars on a timescale shorter than the dynamical time. If so, simple, model independent, physical arguments, presented in \\S\\,\\ref{sect:loss}, imply a swelling of the stellar distribution and a decrease of the stellar velocity dispersion. The model contributes detailed, testable, predictions on the evolution of the effective radius and of the velocity dispersion as a function of the galactic age and of the halo mass. It specifically predicts different evolutionary histories for galaxies with present day stellar masses above and below $M_{\\star}= 2\\times 10^{10}\\, M_{\\odot}$.  The energy injected by dry mergers into the stellar systems can also enlarge their size. However mergers increase the mass in stars too, and, to first order, move galaxies roughly parallel to the size-mass relation (van Dokkum et al. 2008; Damjanov et al. 2008), while the data show size evolution at fixed mass in stars.  The simple analysis presented in this {\\it Letter} is intended to be a first exploration, sketching a promising scenario for interpreting challenging observational results. This scenario can be tested, on one side, by numerical simulations with appropriate time resolution, properly taking into account the interactions between baryons and DM, and, on the other side, by observations of size and velocity dispersion especially of lower mass high-$z$ galaxies, that are expected to show an evolutionary behaviour different from that of massive galaxies because of the different mass loss history."
    },
    "0809/0809.3448_arXiv.txt": {
        "abstract": "We model the nonlinear saturation of the r-mode instability via three-mode couplings and the effects of the instability on the spin evolution of young neutron stars.  We include one mode triplet consisting of the r-mode and two near resonant inertial modes that couple to it. We find that the spectrum of evolutions is more diverse than previously thought. We start our evolutions with a star of temperature $\\sim 10^{10}$ K and a spin frequency close to the Kepler break-up frequency. We assume that hyperon bulk viscosity dominates at high temperatures (T $\\sim$ $10^9-10^{10}$ K) and boundary layer viscosity dominates at lower temperatures ($\\sim$ a few $\\times$ $10^8$ K). To explore possible nonlinear behavior, we vary properties of the star such as the hyperon superfluid transition temperature, the strength of the boundary layer viscosity, and the fraction of the star that cools via direct URCA reactions. The evolution of the star is dynamic and initially dominated by fast neutrino cooling.  Nonlinear effects become important when the r-mode amplitude grows above its first parametric instability threshold. The balance between neutrino cooling and viscous heating plays an important role in the evolution. Depending on the initial r-mode amplitude, and on the strength of the viscosity and of the cooling this balance can occur at different temperatures. If thermal equilibrium occurs on the r-mode stability curve, where gravitational driving equals viscous damping, the evolution may be adequately described by a one-mode model. Otherwise, nonlinear effects are important and lead to various more complicated scenarios. Once thermal balance occurs, the star spins-down oscillating between thermal equilibrium states until the instability is no longer active. The average evolution of the mode amplitudes can be approximated by quasi-stationary states that are determined algebraically. For lower viscosity we observe runaway behavior in which the r-mode amplitude passes several parametric instability thresholds. In this case more modes need to be included to model the evolution accurately. In the most optimistic case, we find that gravitational radiation from the r-mode instability in a very young, fast spinning neutron star within about 1 Mpc of Earth may be detectable by advanced LIGO for years, and perhaps decades, after formation. Details regarding the amplitude and duration of the emission depend on the internal dissipation of the modes of the star, which would be probed by such detections. ",
        "introduction": "\\footnotetext[1]{Current affiliation: Center for Gravitational Wave Physics, Department of Physics, Pennsylvania State University, University Park, PA 16802, USA} Neutron stars are believed to be born in the aftermath of core-collapse supernova explosions as the stellar remnant becomes gravitationally decoupled from the stellar ejecta. Two interesting and timely question are: (1) Do neutron stars spin at millisecond periods at birth or do they spin closer to the observed periods of young pulsars?; (2) Do they emit gravitational radiation that is detectable by interferometers on Earth? %  Theoretically, conservation of angular momentum in the core collapse of a $8 - 30 M_\\odot$ progenitor can lead to a newborn neutron star with a period of $\\sim 1$ ms or shorter. Observationally, the fastest known young pulsar is in the Large Magellanic Cloud supernova remnant N157B, and has a rotation period of 16 ms \\cite{16msDiscovery}. The Crab pulsar is the next fastest neutron star in a supernova remnant with an age $\\sim 10^3$ yr.  Its current  rotation period is 33 ms. Its initial period is estimated to be $\\sim 19$ ms \\cite{WangLai05} by assuming the rotational spin-down is well described by a power law $\\dot{\\Omega} \\propto - \\Omega^n$ with braking index $n = 2.51 \\pm 0.01$ \\cite{lyne}. One way to predict the distribution of initial pulsar periods is through population synthesis studies. These studies generally use present day observations with some assumption of their time evolution to reconstruct the birth distribution periods and magnetic fields of the pulsar population. Current studies favor initial periods in the range of several tens to several hundreds of milliseconds \\cite{kaspi2006, perna2007}. The apparent discrepancy between the theoretically possible fast rotation rates and the observed slow rotation rates of young neutron stars could be reconciled if the r-mode instability or some other mechanism could spin neutron stars down efficiently, preventing them from maintaining millisecond periods.   R-modes are quasi-toroidal oscillations in rotating fluids that occur because of the Coriolis effect. These modes are driven unstable by gravitational radiation reaction via the Chandrasekhar-Friedman-Schutz (CFS) mechanism \\cite{C,FS}. In the absence of fluid dissipation, the CFS mechanism causes any mode that is retrograde in the co-rotating frame, but prograde in the inertial frame, to grow as gravitational radiation is emitted \\cite{nils, Sharon}. The most unstable r-mode is the $n=3, m=2$ mode, called the r-mode throughout the rest of the paper (Here $n$ and $m$ label the degree and order of the Legendre functions associated with the mode.) The gravitational driving equals viscous dissipation for this mode along a critical curve in the angular velocity - temperature ($\\Omega-T$) phase space.  Above the critical curve the r-mode is linearly unstable and the r-mode amplitude grows exponentially. Once the r-mode amplitude passes its first parametric instability threshold, other near-resonant inertial modes are excited via energy transfer from the r-mode. At this point nonlinear effects become important. The parametric instability threshold amplitude depends on internal neutron star physics and changes with angular velocity and temperature.   The first investigation that modeled the spin-down of a young neutron star due to the r-mode instability was performed by Lindblom {\\it et al.\\ }\\cite{LOM} (See also Owen {\\it et al.}~\\cite{OwenEtAl} for a more detailed analysis.) They used a simple one-mode evolution model that assumes  the r-mode amplitude saturates because of nonlinear effects at some arbitrarily fixed value. The saturation amplitude was chosen to be of order 1. In this model, once the instability is saturated, the star spins down at fixed r-mode amplitude. Lindblom {\\it et al.\\ }estimated that a newborn neutron star would cool to approximately $10^9$ K and spin down from a frequency close to the Kepler frequency to about 100 Hz in $\\sim$ 1 yr. In their calculation they included the effects of shear viscosity and bulk viscosity for ordinary neutron star matter composed of neutrons, protons and electrons and assumed modified URCA cooling.  Jones  \\cite{Jones} and Lindblom \\& Owen \\cite{LO} pointed out that if the star contains exotic particles such as hyperons, internal processes could lead to a very high bulk viscosity in the cores of neutron stars. They predicted that for young neutron stars this viscosity would either eliminate the instability altogether or leave a short window of instability of up to a day or so for modified URCA cooling \\cite{LO,mohit} that would render  the gravitational radiation undetectable. Andersson, Jones, and Kokkotas \\cite{AJK} found that, in the case of strange stars, young neutron stars can evolve along the r-mode stability curve reaching a quasi-adiabatic equilibrium at low r-mode amplitudes. Similar conclusions were reached by Reisenegger \\& Bonacic \\cite{nonlinearbulk2} for large hyperon bulk viscosity. We include nonlinear effects and find that the spectrum of possible evolutions is more diverse. Quasi-equilibrium evolutions along the r-mode stability curve are just one scenario.  Schenk {\\it et al.\\ }\\cite{Schenk} developed a formalism to study nonlinear interactions of the r-mode with other inertial modes. They assumed a small r-mode amplitude and treated the oscillations of the modes with weakly nonlinear perturbation theory via three-mode couplings. This assumption was tested by Arras {\\it et al.\\ } \\cite{arras} and Brink {\\it et al.} \\cite{Jeandrew1, Jeandrew2, Jeandrew3,JeandrewThesis}. Arras {\\it et al.} proposed that a turbulent cascade will develop in the strong driving regime. They estimated that the r-mode amplitude was small and could have values between $10^{-4} - 10^{-1}$. Brink {\\it et al.\\ }computed the interactions of about 5000 modes via approximately $1.3 \\times 10^6$ couplings among modes with $n \\le 30$. They modeled the star as incompressible and calculated the coupling coefficients analytically. The couplings were restricted to near resonant modes with a fractional detuning of $\\delta \\omega/(2 \\Omega) < 0.002$.  Brink {\\it et al.\\ }showed that the nonlinear evolution saturates at amplitudes comparable with the lowest parametric instability threshold. They did not include spin or temperature evolution in their calculation.  In a previous paper \\cite{us} we investigated the nonlinear saturation of the r-mode instability for neutron stars in LMXBs and included temperature and spin evolution. In this paper we begin a study of the nonlinear development of the r-mode instability in newborn neutron stars.  We use an effective three-mode treatment in that we include one triplet of modes with a statistically relevant coupling and detuning coefficient and treat it as the lowest parametric instability threshold. The triplet consists of the $n=3$, $m=2$ r-mode and two near-resonant inertial modes with $n=14$ and $n=15$ that couple to it.  The exact inertial modes that are excited when the r-mode grows above its parametric instability threshold will change as the star spins down. However, since the nearest resonance is statistically expected to be at a detuning of  $\\delta \\omega/2 \\Omega \\sim 10^{-4}$ and the mode coupling coefficient for the lowest parametric instability threshold is typically of order one~\\cite{JeandrewThesis}, a mode triplet using these values should provide a qualitatively correct picture.  The model also includes neutrino cooling via a combination of fast and slow processes, viscous heating due to hyperon bulk viscosity and boundary layer viscosity, and spin-down due to gravitational radiation and magnetic dipole radiation. In oder to explore possible nonlinear behaviors we vary: (1) the hyperon superfluid transition temperature $T_{\\rm h}$, which is believed to be $\\sim 10^9-10^{10}$ K; (2) the strength of the hyperon bulk viscosity;  (3) the boundary layer viscosity via the slippage factor $S_{\\rm ns}$ that parametrizes the interaction between oscillating fluid core and the elastic crust \\cite{LU, GA}; and (4) the fraction of the star that cools via direct URCA reactions $f_{\\rm dU}$.  We find a variety of scenarios that depend on these parameters and on the initial r-mode amplitude. The star cools until the cooling is balanced by viscous heating. It then follows a quasi-adiabatic evolution either on the r-mode stability curve or on other Cooling = Heating curves on which the neutrino cooling is balanced by viscous heating from the three modes. For low hyperon bulk viscosity or when we include only slow cooling, we find runaway behavior in which the energy dissipated by the two inertial modes is not sufficient to stop the growth of the r-mode amplitude and  several parametric instability thresholds are passed. Modeling such behavior accurately requires the inclusion of multiple mode triplets, which is beyond the scope of this paper.  The three mode problem is a natural first step in studying nonlinear effects for the r-mode instability. This model is valid as long as the r-mode amplitude does not grow above parametric instability thresholds that are higher than the lowest threshold. The thresholds are functions of the angular velocity and temperature of the star and depend on the viscous damping rates of the modes and model details. The star cools at constant angular velocity until the cooling is stopped by viscous heating and then spins down oscillating between thermal equilibrium states until the r-mode is no longer unstable.  The r-mode amplitudes we find are still fairly low $\\sim 10^{-2} - 10^{-3}$ and the spin-down torque due to gravitational radiation reaction is typically lower than that due to magnetic dipole radiation. So, the spin down timescale, which is approximately the timescale on which the r-mode is unstable, is dominated by the magnetic dipole radiation timescale for $B \\sim 10^{13}$ G.  Our evolutions are determined by competitions between cooling and heating, gravitational driving and viscous damping, and magnetic dipole and gravitational spin-down with the competition between cooling and heating playing the most prominent role. We expect that some of this behavior will hold for more sophisticated models as well.  Our evolutions are determined by competitions between neutrino cooling and viscous heating, gravitational driving and viscous dissipation, neutrino cooling and gravitational driving, neutrino cooling and magnetic spin-down of the star with the competition between cooling and heating placing the most prominent role. The three-mode model is adequate as long as the mode amplitudes do not grow high enough to excite more modes in the star. Stable three-mode evolutions are more likely to happen for high viscosity such that due hyperons. %  The results are summarized in Sec.\\ \\ref{summary}. Model details are presented in Sec.\\ \\ref{MB}. Stable evolutions are examined in detail Sec.~\\ref{stable}. The runaway evolutions are discussed in Sec.~\\ref{unstable}. The prospects of detecting gravitational radiation are considered in Sec.~\\ref{detection}. Limitations of the model are discussed in Sec.\\ \\ref{limitations}. Concluding remarks are presented in Sec.\\ \\ref{conclusion}. % ",
        "conclusions": "\\label{conclusion} This paper is the first treatment of the r-mode instability that uses a physical model for nonlinear saturation in newborn neutron stars. We use one triplet of modes: the $n=3, m=2$ r-mode, and two near-resonant inertial modes that couple to it. Nonlinear effects become important when the r-mode amplitude grows above its parametric instability threshold. This threshold provides a physical cutoff to the r-mode instability by energy transfer to other inertial modes in the system.  The behavior we find is richer than expected from previous work. We find a variety of evolutions that depend on internal neutron star physics, which is still very uncertain. The competition between viscous heating and neutrino cooling plays the most prominent role. Once the cooling is stopped by dissipative heating,  the star spins down oscillating around quasi-steady thermal equilibrium.  Our model has some schematic aspects. We use typical detuning and coupling coefficients that will change as the model becomes more sophisticated. The detuning is roughly inversely proportional to the number of direct couplings to the r-mode, which is a function of the number of modes included. Some evolutions are independent of this frequency difference between modes. When the inertial mode viscosity is larger than the detuning, the parametric instability threshold becomes independent of the detuning.  The threshold in the strong viscosity limit coincides with its value for zero detuning among the frequencies of the mode triplet and still leads to significant nonlinear effects.  Some of our evolutions lead to a gravitational radiation amplitude that is detectable by Advanced LIGO. Assuming a broad band Advanced LIGO curve with two detectors and an integration time of about two weeks, we find that gravitational radiation would be detectable within about 1 Mpc of Earth. The sources would have to be young neutron stars within years or perhaps decades after formation. The gravitational wave amplitude and the duration of the emission depend on the internal dissipation of the modes of the star. Significant spin-down of the star over short periods of time ($\\sim$ weeks) would make detection challenging. However, if fast spinning young neutron stars exist and are detected, the gravitational wave amplitude could give unique information on internal neutron star physics."
    },
    "0809/0809.1217_arXiv.txt": {
        "abstract": "{Extremely metal-deficient [\\oh$\\la$7.6] emission-line galaxies in the nearby universe are invaluable laboratories of extragalactic astronomy and observational cosmology since they allow us to study collective star formation and the evolution of galaxies under chemical conditions approaching those in distant protogalactic systems. However, despite intensive searches over the last three decades, nearby star-forming (SF) galaxies with strongly subsolar metallicity remain extremely scarce. } {We searched the Sloan Digital Sky Survey (SDSS) and the Six-Degree Field Galaxy Redshift Survey (6dFGRS) for promising low-metallicity candidates using a variety of spectroscopic criteria. } {We present long-slit spectroscopy with the 3.6m ESO telescope of eight \\h2\\ regions in seven emission-line dwarf galaxies, selected from the Data Release 4 of SDSS (six galaxies) and from 6dFGRS (one galaxy). In addition, we use SDSS imaging data to investigate the photometric structure of the sample galaxies. } {From the 3.6m telescope spectra, we determine the oxygen abundance of these systems to be \\oh$\\la$7.6, placing them among the most metal-poor star-forming galaxies ever discovered. Our photometric analysis reveals a moderately blue, stellar host galaxy in all sample galaxies.} {The detection of a stellar host in all galaxies studied here and all previously studied extremely metal-deficient SF galaxies implies that they are unlikely to be forming their first generation of stars. With regard to the structural properties of their host galaxy, we demonstrate that these systems are indistinguishable from blue compact dwarf (BCD) galaxies. However, in contrast to the majority ($>$90\\%) of BCDs that are characterised by red elliptical host galaxies, extremely metal-poor SF dwarfs (hereafter XBCDs) reveal moderately blue and irregular hosts. This is consistent with a young evolutionary status and in the framework of standard star formation histories implies that several XBCDs formed most of their stellar mass in the past $\\sim$2 Gyr. A large fraction of XBCDs reveal a \\emph{cometary} morphology due to the presence of intense SF activity at one edge of an elongated host galaxy with a gradually decreasing surface brightness towards its antipodal end. } ",
        "introduction": "} The identification and detailed studies of chemically unevolved, star-forming (SF) galaxies in the nearby universe, of almost pristine chemical composition, is a major task for contemporary observational cosmology. Some important aspects of these studies are the following.  First, systematic studies of SF galaxies with strongly subsolar metal abundances are indispensable for placing tight observational constraints on the primordial $^4$He abundance $Y_{\\rm p}$ \\citep[][and references therein]{ITS2007,Peimbert2007}.  Secondly, the metallicity plays a key role in virtually all aspects of star- and galaxy evolution since it influences e.g. a) the production rate of Lyman continuum photons and efficiency of radiative winds in massive stars \\citep[see e.g.][]{L92,SchaererdeKoter1997,Kudritzki2002}, including the properties of Wolf-Rayet stellar populations \\citep{Izotov1997-IZw18,ShaererVacca1998,Guseva2000-WR,Crowther2006,Brinchman2008}, as well as b) the cooling efficiency of the hot ($\\sim 10^7$ K) X-ray emitting gas \\citep{BoehringerHensler1989}, recognised to be an ubiquitous component of starburst galaxies \\citep[see e.g.][]{Heckman1995-NGC1569,PapaderosFricke1998a,Martin2002-NGC1569,Ott2005b}. Spatially resolved studies of extremely metal-poor nearby SF galaxies may therefore yield crucial insights into the early evolution of faint protogalactic systems forming out of primordial or almost metal-free gas in the young universe.  The most suitable nearby objects for exploring these issues are blue compact dwarf (BCD) galaxies. These galaxies form a morphologically heterogeneous class of intrinsically faint ($M_B$ $\\ga$ --18 mag) extragalactic systems undergoing intense starburst activity on a spatial scale of typically $\\sim$1 kpc \\citep[see][hereafter P02 and references therein]{Papaderos2002-IZw18}. In $\\sim$90\\% of the local BCD population, SF activity proceeds in one or several luminous \\h2\\ regions within the central part of a more extended, low-surface brightness (LSB) elliptical host galaxy \\citep[][hereafter LT86 and P96a, respectively]{LooseThuan1986,Papaderos1996a}. This red (0.8$\\la B-R\\,\\,{\\rm (mag)}\\la$1.2; P96a) galaxy host contains, on average, one half of the optical emission of a BCD \\cite[][hereafter P96b]{Papaderos1996b} and typically dominates the line-of-sight intensity of this system for surface brightness levels fainter than 24.5 $B$ \\sbb\\ (P96a, P02). Deep surface photometry in the optical and near infrared \\citep{Noeske2003-NIR,Cairos2003-NIR,GildePazMadore2005} and colour-magnitude-diagram analyses \\citep[e.g.,][]{SchulteLadbeck1999-7Zw403,Tosi2001-NGC1705} confirmed the earlier conclusion (LT86, P96b) that BCDs are, overwhelmingly, evolved gas-rich dwarfs undergoing recurrent starburst activity.  BCDs are the most metal-deficient emission-line galaxies known in the nearby universe. However, while all BCDs show subsolar chemical abundances, it is notoriously difficult to find extremely metal-deficient systems in the range \\oh$\\la$7.6. The BCD abundance distribution peaks at \\oh$\\approx$8.1 with a sharp drop-off at lower values \\citep{Terlevich1991,Thuan1995,IT98a,KunthOstlin2000}. One of the first BCDs discovered, I\\,Zw\\,18 \\citep{SargentSearle1970} with \\oh=7.17$\\pm$0.01 \\citep{Izotov1997-IZw18}, held the record as the most metal-deficient SF galaxy known for more than three decades. Only very recently was this system replaced in the metallicity ranking by the BCD SBS 0335--052\\,W with an oxygen abundance \\oh=7.12$\\pm$0.03 \\citep{Izotov2005-SBS0335W}.  Despite large observational efforts over the past three decades and systematic studies of several $10^5$ catalogued emission-line galaxies, less than 20 BCDs with \\oh$\\la$7.6 (hereafter XBCDs) were identified in the nearby ($z\\la 0.04$) universe until a few years ago \\citep[][and references therein]{KunthOstlin2000}. Since then substantial progress has been achieved, and more than one dozen further XBCDs discovered \\citep{K03,K04b,Guseva2003-SBS1129,Guseva2003-HS1442, Pustilnik2005-DDO68,Pisano2005,Pustilnik2006-HS2134, Izotov2006-2XBCD,Papaderos2006-2dF,IzotovThuan2007-MMT,Kewley2007-SDSS0909}. At higher redshift ($0.2 \\la z \\la 0.8$), dedicated surveys utilising 10m-class telescopes and optimised search strategies led to the discovery of about ten more SF dwarf galaxies with \\oh$\\la$7.6 and an absolute rest-frame magnitude in the range of BCDs \\citep{Kakazu2007}. In spite of the extreme scarcity and intrinsic faintness of XBCDs, there is therefore a tangible prospect of unveiling a significant number of these systems in the years to come.  \\begin{table*} \\caption{Coordinates of the sample galaxies (J2000.0) \\label{tab1}} \\begin{tabular}{lccccccc} \\hline Name       & R.A.   & DEC  & redshift & Distance$^{\\rm a}$ (Mpc) & airmass & P.A. ($^{\\circ}$) & exp.time (sec)\\\\ \\hline J 0133$+$1342   & 01 33 52.6 & $+13$ 42 09  & 0.00879 & 37.8 & 1.37 & 170.4  & 3600 (3 $\\times$ 1200) \\\\ g0405204-364859 & 04 05 20.4 & $-36$ 48 59  & 0.00268 & 11.4 & 1.04 & 122.0  & 2400 (3 $\\times$ 800) \\\\ J 1044+0353a    & 10 44 57.8 & $+03$ 53 13  & 0.01274 & 51.2 & 1.19 & 101.5  & 3600 (3 $\\times$ 1200) \\\\ J 1044+0353b    & 10 44 58.0 & $+03$ 53 13  & 0.01284 & 51.2 & 1.19 & 101.5  & 3600 (3 $\\times$ 1200) \\\\ J 1201+0211     & 12 01 22.3 & $+02$ 11 08  & 0.00327 & 14.0 & 1.21 & 119.0  & 2400 (2 $\\times$ 1200) \\\\ J 1414-0208     & 14 14 54.1 & $-02$ 08 23  & 0.00528 & 23.0 & 1.16 & 155.3  & 3600 (3 $\\times$ 1200) \\\\ J 2230-0006     & 22 30 36.8 & $-00$ 06 37  & 0.00559 & 24.9 & 1.33 & 156.0 & 3600 (3 $\\times$ 1200) \\\\ J 2302+0049     & 23 02 10.0 & $+00$ 49 39  & 0.03302 & 134.6 & 1.17 & 156.0 & 2400 (2 $\\times$ 1200) \\\\ \\hline \\end{tabular}  \\vspace*{0.5ex}\\parbox{16cm}{ $^{\\rm a}$: Distance, derived after correction of the measured redshifts for the motion relative to the Local Group centroid and the Virgocentric flow, assuming a Hubble constant of 75 km s$^{-1}$ Mpc$^{-1}$.} \\end{table*}   \\smallskip Third, recent work provides strong observational support to the idea that some XBCDs in the nearby universe formed most of their stellar mass in the past $\\sim$1 Gyr, which implies that they are cosmologically young systems \\citep{Papaderos1998-SBS0335,Vanzi2000-SBS0335E,Guseva2001-SBS0940,Papaderos2002-IZw18, Guseva2003-SBS1129,Guseva2003-HS1442,Guseva2003-SBS1415,Hunt2003-IZw18,Pustilnik2004-SBS0335}. The youth hypothesis for XBCDs is further supported by a number of evolutionary synthesis and colour-magnitude diagram (CMD) studies that indicate an upper age of 0.1-2 Gyr for the stellar component of some of those systems \\citep{Izotov1997-SBS0335E,Thuan1997-SBS0335E,Izotov2001-IZw18, Fricke2001-T1214,IzotovThuan2004-IZw18,IzotovThuan2004-UGC4483}.\\\\ In the case of the XBCD I Zw 18, the conclusion that this system started forming stars not earlier than $\\la$0.5 Gyr ago \\citep{IzotovThuan2004-IZw18}, was recently disputed by the subsequent CMD studies of \\citet{Aloisi2007-IZw18}.  However, apart from the question of when the first stars in a XBCD were formed, there appears to be broad consensus that many of these systems underwent the dominant phase of their formation only recently. If so, XBCDs may be regarded as convenient laboratories to explore the main processes driving dwarf galaxy formation, as long as their morphological and dynamical relics have not had time to be erased in the course of the secular galactic evolution, through e.g. two-body relaxation, galaxy merging and subsequent star formation episodes. Low-mass protogalaxies in the distant universe, once detected, will remain due to the cosmological dimming and their intrinsic faintness and compactness challenging to study with sufficient resolution and accuracy, even with the next generation of extremely large telescopes. Young XBCD candidates provide in this respect a bridge between near-field and high-redshift observational cosmology and invaluable laboratories of extragalactic research.  This paper investigates the spectroscopic, photometric, and morphological properties of a new sample of XBCDs. It is organised as follows: the sample selection, data acquisition, and reduction are described in Sect. \\ref{obs}. Our spectroscopic and photometric analysis is presented in Sect. \\ref{results}, and in Sect. \\ref{objects}, we discuss the properties of individual sample galaxies. The photometric and morphological properties of XBCDs are reviewed in Sect. \\ref{discussion}, and in Sect. \\ref{summary}, we summarise our conclusions. ",
        "conclusions": "} Including the seven galaxies studied here, the total number of XBCDs has increased to $\\sim$35. It is therefore timely to review their general photometric and morphological properties, in order to explore possible trends and better coordinate future searches for these systems in the nearby universe and at higher redshift.  \\begin{figure*} \\hspace*{0.0cm}\\psfig{figure=aa10028f33.ps,angle=-90.0,width=11.5cm,clip=} \\hspace*{-2.2cm}\\psfig{figure=aa10028f34.ps,angle=-90.0,width=11.5cm,clip=} \\caption{Comparison of the structural properties of the host galaxy of XBCDs, BCDs, dIs, dwarf ellipticals (dEs) and Low-Surface Brightness (LSB) galaxies. Data for iE/nE BCDs are compiled from \\citet{Cairos2001a}, \\citet{drinkwater91}, \\citet{Marlowe1997}, \\citet{PapaderosFricke1998a}, \\citet{Noeske2000-cometary}, P96a, \\citet{Papaderos1998-PhD} and P02. Data for other types of dwarf galaxies are taken from \\citet{BinggeliCameron91}, \\citet{BinggeliCameron93}, \\citet{Bothun91}, \\citet{Caldwell87}, \\citet{Carignan89}, \\citet{Hopp91}, \\citet{PT1996}, \\citet{Vigroux86} and \\citet{vanZee2000}. Photometric quantities in the SDSS $g$ and $V$ band were transformed into $B$ band assuming $B - g = 0.3$ mag and $B - V = 0.5$ mag (see discussion in Sect. \\ref{photo2} and in P06). The lines show the shift of the data points caused by a change of the Hubble constant from 75 to 50 and 100 km s$^{-1}$ Mpc$^{-1}$. {\\bf (a)} Central surface brightness $\\mu _{\\rm E,0}$ vs. absolute $B$ magnitude $M_{\\rm host}$ of the LSB component. {\\bf (b)} Logarithm of the exponential scale length $\\alpha $ in pc vs. $M_{\\rm host}$. Data for XBCDs are compiled from \\citet{Papaderos1998-SBS0335}, \\citet{Papaderos1999-Tol65}, \\citet{Kniazev2000-HS0822}, \\citet{Guseva2001-SBS0940}, \\citet{Fricke2001-T1214}, P02, \\citet{Guseva2003-SBS1129}, \\citet{Guseva2003-HS1442}, \\citet{Guseva2003-SBS1415}, P06 and \\citet{Pustilnik2005-DDO68}. Five further systems with an oxygen abundance slightly above \\oh=7.6 are also included: \\object{2dF 169299, UM 570, UM 559} and \\object{2dF 84585} [\\oh=7.68, 7.71, 7.72, 7.66, respectively; P06] and \\object{Pox 186} \\citep[\\oh=7.74, ][]{Guseva2004-Pox186}. } \\label{structure} \\end{figure*}  A first important conclusion from the present and previous studies is that, similar to other dwarf galaxies, the host galaxy of XBCDs can be approximated by an exponential fitting law in its outer parts, i.e. for galactocentric radii $3\\,\\alpha \\la $\\rr$\\la 6\\,\\alpha$. Additionally, several XBCDs display extended central flat cores in their host galaxies, reaching in some cases out to \\rr$\\sim$3$\\alpha$. Centrally flattened exponential profiles have also been observed in a sizeable fraction of early and late-type dwarfs spanning a wide range in absolute magnitude \\citep[][P96a]{BinggeliCameron91,Noeske2003-NIR}. The XBCD -- BCD connection is further supported by the fact that both object classes populate roughly the same locus in the $\\mu_{\\rm E,0}$ versus $M_{\\rm host}$ and log($\\alpha$) versus $M_{\\rm host}$ parameter space (Fig. \\ref{structure}) and are comparable in their effective radius (P06).  These lines of evidence suggest that XBCDs do not represent peculiar cases of dwarf galaxy evolution, reflected in strongly distinct structural properties (e.g. an abnormally diffuse or an ultra-compact LSB host), but that they share similar structural properties and are therefore likely also to lie on a common evolutionary track with the main population of more metal-rich BCDs with \\oh$\\ga$8. Of course, this conclusion may not be free of selection biases, given that XBCDs are mostly detected by their high EW(\\hb), blue colours, and compactness. It is therefore conceivable that a substantial population of relatively quiescent and more diffuse metal-poor SF dwarfs remains strongly under-represented in current XBCD surveys. The irregular SF dwarf \\object{J0113+0052}, discovered by \\citet{Izotov2006-2XBCD}, may be regarded as an example of this kind. Additionally, it might be questioned that optical spectroscopic searches for XBCDs, based on forbidden line measurements in \\h2\\ regions, can uncover systems with a gas-phase metallicity significantly lower than \\oh$\\sim$7.0 \\cite[$\\sim$\\zsun/60, adopting a solar abundance of 8.76,][]{Caffau2008}, thus better constrain the minimum level of chemical enrichment in the warm ISM of SF galaxies. This is because, below some abundance level and ionization parameter, oxygen forbidden lines become too weak to be measurable. For example, calculations with the photoionization code CLOUDY \\citep{Ferland1998_CLOUDY90} show that, at a metallicity \\oh=7.0, electron density $N_{\\rm e}$=100 cm$^{-3}$ and ionization parameter $U$=10$^{-3}$, the oxygen [O {\\sc iii}]$\\lambda$4363 and [O {\\sc iii}]$\\lambda$5007 line intensities decrease to $<$1\\% and 38\\% of that of the H$\\beta$ line, respectively, a determination of the electron temperature based on the [O {\\sc iii}]$\\lambda$4363 line becomes therefore practically impossible.  \\smallskip It is important to note that the detection of a \\emph{stellar} LSB host in all known XBCDs, including \\object{SBS 0335-052E} \\citep[][P98]{Izotov1997-SBS0335E,Thuan1997-SBS0335E} and \\object{I Zw 18} \\citep[][P02]{Izotov2001-IZw18,Papaderos2001-IZw18} suggests that none of these systems are currently forming \\emph{in situ} its first generation of stars. This is also indicated by spatially resolved evolutionary synthesis studies \\citep{Izotov1997-SBS0335E,Vanzi2000-SBS0335E,Guseva2001-SBS0940,Guseva2003-SBS1415,Hunt2003-IZw18} and colour-magnitude diagram analyses of selected systems \\citep[e.g.,][]{IzotovThuan2004-IZw18,IzotovThuan2004-UGC4483,OstlinMouhcine2005-IZw18,Aloisi2007-IZw18}.  On the other hand, the existing work consistently suggests that the host galaxies of XBCDs are less evolved than those of BCDs. This is indicated by deep surface photometry studies that reveal uniformly blue colours in the XBCD hosts, with a $V-I$ and $g-i$ index in the range between $\\sim$0.1 and $\\la$0.5 mag \\citep[e.g., P02;][and references therein]{Guseva2003-SBS1415} and $<$0.4 mag (this paper), respectively. By contrast, the LSB hosts of iE/nE BCDs, representing $\\sim$90\\% of the BCD population (LT86) show typically red ($\\sim$1 mag) $B-R$ and $V-I$ colours \\citep[P96b;][]{Cairos2001b,GildePazMadore2005}.  Another important, largely overlooked aspect concerns the morphology of XBCDs. The hosts of these systems reveal in their majority conspicuous deviations from axis-symmetry suggesting a little degree of dynamical relaxation. By this, they again significantly differ from the \\emph{bona fide} old, more metal-rich iE/nE BCDs whose defining property is a smooth \\emph{elliptical} LSB host galaxy.  \\begin{figure*} \\hspace*{0.7cm}\\psfig{figure=aa10028f35.ps,angle=0.0,width=17.0cm,clip=} \\caption{ Comparison of systems in the range of XBCD metallicity with the main class of the more metal-rich, \\emph{bona fide} old BCDs \\object{Mkn 996} and \\object{Mkn 86}. These two systems may be regarded as prototypical examples of the dominant nE/iE morphological BCD type (cf. LT86). All galaxy images were scaled to a common distance to better facilitate a comparison. Those images that are displayed magnified or downsized for the sake of better visibility are labeled with the respective scaling factor. Off-center low-metallicity \\h2\\ regions are marked with crosses for clarity. Within brackets we indicate the oxygen abundance of the lowest-metallicity \\h2\\ region in each XBCD. Metallicities are taken from the literature sources given in pages \\pageref{irr_xbcds1} and \\pageref{irr_xbcds2}. The images are from P06 (\\object{2dF 169299, 2dF 115901, 2dF 171718, 2dF 84585}), P98 (\\object{SBS 0335-052E}, see also Thuan et al. 1997, HST program 5408), \\citet[][\\object{SBS 0335-052W}]{Papaderos2006-IAU}, \\citet[][\\object{SBS 0940+544}]{Guseva2001-SBS0940}, \\citet[][\\object{SBS 1129+576}]{Guseva2003-SBS1129}, \\citet[][\\object{SBS 1415+437}, HST program 5408]{Guseva2003-SBS1415}, \\citet[][\\object{I Zw 18}, $R$ band HST WFPC2 exposure with the \\ha\\ emission subtracted out, HST program 5309]{Papaderos2001-IZw18}, the HST archive (\\object{Tol 65, Tol 1214-277, Mkn 996}; HST programs 5408 and 6678), Papaderos \\& Noeske (2008, in preparation, \\object{HS 0822+3542, UGC 4483, Mkn 86}), and from the SDSS (\\object{J1414-0208, J1044+353, J2302-0049, J1201+211, J2104-0035, J0301-0052, J0911+3135, DDO68, J2238+1400}). } \\label{XBCD-morph} \\end{figure*}   Faint, non-axisymmetric distortions are present in roughly one half of the XBCDs studied here. The irregular outer morphology of those systems is certainly not due to insufficient detection of their underlying galaxy host. This is because, as is evident from Fig. \\ref{colour_maps}, SDSS images allow us to interpolate contours down to a $g$ band surface brightness $\\mu_g \\geq 25.5$ \\sbb, corresponding to $\\mu_B \\approx 25.8$ \\sbb. At such intensity levels the elliptical host of evolved BCDs dominates the light \\citep[P96b, ][P02]{Cairos2001b} and should have been detected if it were present. Furthermore, with the possible exception of \\object{J1201+0211}, stacked $gri$ images do not reveal tidal features in any of our sample SDSS galaxies, ruling out \\emph{strong} gravitational interactions or galaxy merging as the origin of the observed morphological distortions. Widespread SF activity in the LSB host can also be excluded from $g-i$ colour maps and long-slit spectra.  Likewise, most of the XBCDs investigated previously on the basis of deep surface photometry are as well characterised by irregular host galaxies. \\label{irr_xbcds1} Such examples are \\object{SBS 0335-052W} \\citep[\\oh=7.12, ][see P98 and Papaderos et al. 2006c for photometry]{Izotov2005-SBS0335W}, \\object{Tol 65} \\citep[\\oh=7.54, ][see Papaderos et al. 1999 for photometry]{Izotov2004-T1214-T65}, \\object{SBS 1415+437} \\citep[\\oh=7.6,][]{Thuan1999-SBS1415,Guseva2003-SBS1415}, \\object{Tol 1214-277} \\citep[\\oh=7.55, ][see Fricke et al. 2001 for photometry]{Izotov2004-T1214-T65}, \\object{I Zw 18} \\citep[\\oh=7.17, ][see Papaderos et al. 2001 and P02 for photometry]{Izotov1997-IZw18}, \\object{SBS 0940+544} \\citep[\\oh=7.46, ][]{Guseva2001-SBS0940}, \\object{SBS 1129+576} \\citep[\\oh=7.36, ][]{Guseva2003-SBS1129}, \\object{HS 0837+4717} \\citep[\\oh=7.64,][]{Pustilnik2004-HS0837}, \\object{DDO 68} \\citep[\\oh=7.21\\dots 7.13, ][respectively]{Pustilnik2005-DDO68,IzotovThuan2007-MMT}, \\object{HS 2134+0400} \\citep[\\oh=7.44,][]{Pustilnik2006-HS2134,Guseva2007-BJ}, \\object{J2104-0035} and \\object{J0113+0052} \\citep[\\oh$\\approx$7.2 and 7.26, respectively,][]{Izotov2006-2XBCD}, \\object{2dF 171716, 2dF 115901} (\\oh=7.5, 7.57, respectively; P06) and \\object{J0301-0052}, \\object{J0911+3135}, \\object{J2238+1400} \\citep[\\oh=7.52, 7.51 and 7.56, respectively, ][]{IzotovThuan2007-MMT}. This is also the case for several systems with an oxygen abundance slightly above \\oh=7.6, such as \\object{2dF 84585} and \\object{2dF 169299} (\\oh=7.66 and 7.68, respectively; P06).  \\smallskip With regard to the morphological properties of the star-forming component of XBCDs (e.g. multiplicity, luminosity distribution, and the degree of the confinement of SF regions to the center of the XBCD host), these have never been studied in a systematic manner before, and a comparison with iE/nE BCDs would therefore be premature at this point. However, two trends are apparent. First, in a significant fraction of XBCDs, star-forming activities are not strongly confined to the geometrical center of the host galaxy and, second, the global SF process in several of these systems appears to be largely driven by propagation.  Figure \\ref{XBCD-morph} contains several examples of XBCDs with off-center SF activity, including \\object{2dF 115901}, \\object{SBS 0335-052W}, \\object{J1414-0208}, and \\object{J0911+3135}. The most impressive instances of off-center SF activity among XBCDs are \\emph{cometary} systems (iI,C type in the BCD classification scheme devised by LT86). These objects contain a luminous SF region at the one tip of an elongated blue and irregular host galaxy with a gradually decreasing surface brightness towards its antipodal end. The hypothesis that these systems are edge-on disks with a dominant SF complex in their outermost periphery can be dismissed on statistical grounds. Some iI,C XBCDs (for example, \\object{SBS 1415+437} and \\object{SBS 0940+544}) display signatures of low-level ongoing or recent star formation along their major axis or colour gradients suggesting propagation of SF activities from the far-end side of their elongated host galaxy towards the young, dominant SF region \\citep[see e.g. P98,][]{Guseva2001-SBS0940,Guseva2003-SBS1415}.  That the formation of the stellar component in a XBCD may largely be driven by SF propagation has been suggested from an analysis of \\label{irr_xbcds2} \\object{SBS 0335-052E} \\cite[\\oh=7.3\\dots 7.2,][respectively]{Izotov1997-SBS0335E,Papaderos2006-SBS0335}, for which P98 have estimated a SF propagation velocity of $\\sim$20 km/sec, of the order of the sound speed in the warm ISM. It is unclear for how long this process can continue; however, as long as gas supply of the appropriate temperature and density is available ahead of the SF front, no self-limiting mechanism should exist. Propagating star formation with a constant speed $u$ over a period $\\tau$ may naturally lead to a cometary morphology on a linear scale $l \\sim$ 10 kpc $\\times (u/10\\,\\,{\\rm km/s}) \\times (\\tau/1\\,\\,{\\rm Gyr})$. As evident from Fig. \\ref{XBCD-morph}, $l$ is of the order of the projected major-axis of several cometary XBCDs.  We note that the iI,C morphology is not uniquely observed in XBCDs and that examples of this kind also exist among BCDs of higher metallicity \\citep{Noeske2000-cometary,Cairos2001a,Noeske2003-NIR,GildePaz2003}. The essential trend, however, is that whereas systems with cometary morphology or strongly off-center SF activity comprise less than 10\\% of the BCD population (LT86), they apparently dominate the XBCD population. In this context, it is worth pointing out that studies of iI,C BCDs in the metallicity range between $\\approx$7.8 and 8.0 by \\cite{Noeske2000-cometary} suggest that these systems are younger ($\\la$4 Gyr) than the main class of iE/nE BCDs, they are therefore likely to represent intermediate stages of BCD evolution. This conjecture is consistent with the high incidence of cometary systems among young XBCD candidates discussed here.  It is unclear whether or not cometary morphology is linked to galaxy interactions, as neither observations nor numerical simulations presently provide tight constraints in this respect. However, it is known that a significant fraction ($>$30\\%) of BCDs are not truly isolated but have optically faint nearby companions \\citep{Noeske2001,Pustilnik2001-env}. Dwarf galaxy encounters with a wide range of impact parameters are therefore likely to play an important role in BCD evolution. The frequency of strong collisions and, eventually, subsequent merging of BCD progenitors is still not constrained well. However, several notable examples (iI,M BCDs in the classification scheme of LT86) exist in the samples of \\cite{Cairos2001a} and \\cite{GildePaz2003}, and there is growing evidence that the major fraction of intrinsically luminous ($M_B<$--18 mag) Blue Compact Galaxies are of merger origin \\citep{Ostlin2001-BCGII}. In addition, the morphology of some XBCDs (for example, \\object{2dF 169299}, \\object{2dF 115901}, and \\object{2dF 171716}; cf. Fig. \\ref{XBCD-morph}) is consistent with the merger interpretation. The importance of gravitational interactions to BCD evolution is indicated further by radio interferometry that reveals HI clouds in the close vicinity ($\\la$100 kpc) of numerous systems \\citep{Taylor1993,THL2004}.  The hypothesis that \\emph{strong} gravitational interactions or galaxy merging are the origin of cometary morphology does not however appear to be tenable. If cometary BCDs were indeed forming by galaxy merging, one would expect tidal features to protrude far beyond their Holmberg radius and be readily detectable at surface brightness levels of $\\mu\\sim$24 $B$ \\sbb, in a similar way to merging disk galaxies. Since the visibility timescale of tidally ejected stellar and gaseous matter is long \\cite[$\\sim$1 Gyr;][and references therein]{HibbardMihos1995}, of the order of the luminosity-weighted age of the XBCD galaxy host, these features should be almost ubiquitous in iI,C systems. The probability that both merging counterparts are metal poor and retain their gas-phase metallicity of a level \\oh$\\la$7.6, even after a strong, merging-induced starburst, is also low.  A more viable interpretation involves \\emph{weak} interactions with low-mass stellar or gaseous companions. These may have a twofold effect, leading to a bar-like gas distribution and triggering SF activities that, by propagation along the direction of maximum gas density, could subsequently produce a cometary BCD morphology. Alternatively, a propagating shock wave induced by gas-cloud infall onto a quiescent late-type dwarf might generate a similar star formation pattern.  In summary, the hypothesis that SF propagation is the main process driving the formation of cometary XBCDs is not in conflict with the hypothesis that the evolution of these systems is largely influenced by interactions. However, lacking theoretical guidence and robust statistics on the gas distribution and kinematics of these systems, the possible role of interactions cannot be reliably assessed. However, numerical simulations of increasing sophistication continue to reproduce important properties of SF dwarfs \\cite[e.g.][]{Noguchi2001,Recchi2002,Pelupessy2004,Pelupessy2006, Hensler2004,Bekki2008}, and promise to provide key insights into the star formation history and morphological evolution of cometary dwarfs in the near future. Dedicated observational studies of these extremely metal-poor cometary galaxies will also hopefully provide important constraints and incentives for theoretical work.  \\smallskip As mentioned above, a compelling argument that XBCDs have undergone previous evolution derives from the detection of an extended \\emph{stellar} host galaxy. The blue, luminosity-weighted colours of this component, interpreted in the framework of simple SFH parametrisations (e.g., {\\lvss SFH2}) suggest that XBCDs form a heterogeneous class of predominantly cosmologically young objects, with examples among them which have formed most of their stellar mass in the past few $10^8$ yrs and are possibly still experiencing the major phase of their dynamical assembly (e.g. \\object{SBS 0335-052E}, \\object{I Zw 18}) to moderately evolved cometary systems that likely formed most of their stellar mass within the last $\\sim$2 Gyr (\\object{Tol 65}, \\object{Tol 1214-277}, \\object{SBS 1415+437}, \\object{SBS 0940+544}, \\object{J1044+0353}). This conclusion is supported further by detailed evolutionary synthesis models \\cite[e.g.][]{Izotov2001-IZw18,Guseva2001-SBS0940,Guseva2003-SBS1129,Guseva2003-SBS1415} involving more complex SFHs and which aim to account self-consistently for a variety of observables, such as the EWs of Balmer emission and absorption lines, the slope of the spectral energy distribution (SED), and spatially resolved colours.  In the case of the XBCD \\object{I Zw 18}, CMD studies based on HST ACS images do not converge into a broadly accepted interpretation about the age of its stellar component, yielding values between $\\leq$0.5 Gyr \\citep{IzotovThuan2004-IZw18} and $>$1 Gyr \\citep{Aloisi2007-IZw18}. This might be partly due to the specific properties of this system. \\cite{Papaderos2002-IZw18} showed that the exponential LSB host of \\object{I Zw 18}, which was previously thought to be dominated by stellar emission, is due entirely to extended, patchy ionized gas emission, therefore extreme caution is required when using CMDs and surface photometry to place constraints on the formation history of its stellar component. The \\emph{stellar} component of \\object{I Zw 18} is by a factor of approximately 2 more compact that the ionized gas halo and, in contrast to the main class of evolved BCDs, has uniformly blue colours down to a surface brightness level of $\\mu \\sim$ 26 $B$ \\sbb. Additionally, as found by \\cite{Cannon2002-IZw18}, \\object{I Zw 18} shows a highly inhomogeneous extinction pattern, a fact that further complicates CMD analyses. The nature of point sources detected in the presense of dominant ionized gas emission in \\object{I Zw 18}, at a distance of 18.2 Mpc \\citep{Aloisi2007-IZw18}, which corresponds to a projected area of $\\sim$20 pc$^2$ per HST ACS pixel, is undoubtedly an important question in XBCD research.  The formation history of the XBCD host is clearly a crucial and outstanding issue. XBCDs may all contain a faint substrate of stars of cosmological age (see discussion in P98), quite similar to the ancient metal-poor stellar population in Local Group dwarf spheroidals \\citep{GrebelGallagher2004}. As pointed out in P98, this putative ancient stellar background cannot be entirely ruled out on spectrophotometric grounds, since its effect on the observed SED might be barely detectable. It is only possible to infer an upper bound to the mass fraction of this hypothetical old stellar population that, by necessity, is tied to simplifying assumptions about the SFH and intrinsic extinction. Current upper limits for a few thoroughly analysed XBCDs range between $\\sim$15\\% and $<$50\\% \\citep{Vanzi2000-SBS0335E,Guseva2001-SBS0940, Guseva2003-SBS1415,Hunt2003-IZw18,Pustilnik2004-SBS0335}.  Arguably, a low mass fraction of ancient stars possibly present in XBCDs do not rule out the young evolutionary status of these systems. This would only be the case if the formation epoch of those first stars were coeval with the dominant phase of XBCD formation, and XBCDs/BCDs were invariably forming in a single, short (1--3 Gyr) SF episode that converted most of their gaseous reservoir into stars. This interpretation is untenable, however, inter alia, because of the gas-richness and recurrent starburst activity of these systems \\citep[see e.g.][]{Thuan1991,MasHesseKunth1999}. As a matter of fact, dwarf galaxies of a given baryonic mass may have followed quite diverse evolutionary pathways. This is manifestated in e.g. the Local Group galaxies for which we observe an impressive manyfold of SFHs and morphologies, comprising both ancient metal-poor dwarf spheroidals and metal-rich dEs \\citep[e.g. \\object{NGC 185, NGC 205}][see also Mateo 1998]{Dolphin2005} as well as late-type irregulars that underwent the dominant phase of their build-up only $\\sim$1 Gyr ago \\citep[e.g. \\object{Sextans A}][]{Dolphin2003-SexA}. The inherent variety of SFHs in dwarf galaxies is further enhanced by the role of the environment, which is recognised to be an important factor in galaxy evolution \\citep{BO1984,Dressler1980,Dressler1997,Poggianti1997,Pustilnik2001-env}. For example, it is well established that the EW(\\hb) of SF dwarfs increases towards the periphery of galaxy clusters \\citep{Vilchez1995} and that XBCDs preferentially populate the extreme field \\citep{Pustilnik2001-env}. These findings are consistent with the idea that the formation timescale of relatively isolated late-type dwarfs is the longest, making low-density regions promising sites to search for young XBCDs candidates. In view of such considerations, the existence of a small number of unevolved XBCDs in the nearby universe is unsurprising, and a manifestation of the diversity in the SFHs of dwarf galaxies."
    },
    "0809/0809.1021_arXiv.txt": {
        "abstract": "{Gravitational lensing is predicted by general relativity and is found in observations. When a gravitating body is surrounded by a plasma, the lensing angle depends on a frequency of the electromagnetic wave due to refraction properties, and the dispersion properties of the light propagation in plasma. The last effect leads to dependence, even in the uniform plasma, of the lensing angle on the frequency, what resembles the properties of the refractive prism spectrometer. The strongest action of this spectrometer is for the frequencies slightly exceeding the plasma frequency, what corresponds to very long radiowaves. } ",
        "introduction": "An ordinary theory of the gravitational lensing is developed for the light propagation in the vacuum. Gravi\\-ta\\-tio\\-nal lensing in the vacuum is achromatic because the deflection angle for the photon does not depend on the frequency of the photon \\cite{GL}. In the limit of a weak lensing in vacuum by a body with a mass $M$ the deflection angle for the photon (Einstein angle) is $\\hat{\\alpha} = 4GM/c^2 b = 2 r_g/b$, under condition $b \\gg r_g$, where $b$ is impact parameter, $r_g$ is the Schwarzschild radius \\cite{GL}.  Propagation of the light in the medium at presence of the gravity field was considered by many authors \\cite{Muhl1966}, \\cite{Muhl1970}, \\cite{Noonan}, \\cite{B-Minakov} and references there. If we consider inhomogeneous medium (without gravity) the light rays move along the curved trajectories in this medium. In the papers concerning gravitational deflection  the inhomogeneous medium was considered. The deflection due to the gravitation, and the deflection due to the inhomogenity of the medium had been considered separately, without an account of the influence of the dispersion in plasma on the light propagation in the gravitational field. In this work we show that even in the homogeneous medium the dispersion in the plasma leads to dependence of the light deflection  angle on the wavelength, what is different from the constant deflection angle in the vacuum.  It have been shown  \\cite{GL}, \\cite{B-Minakov}, \\cite{Fok}, \\cite{LL2}, \\cite{Myoller}, that light propagation in the gravitational field in the vacuum may be formulated as  its propagation in a inhomogeneous medium with an effective refraction index $n_g$, depending on metric. It have been shown also, that in presence of the medium in the gravitation field, such analogy is valid too. In this case we should use the refraction index, which is a multiplication of the effective gravitational refraction index $n_g$ and usual refraction index $n$, determined by the physical properties of the medium: $n_{eff} = n_g n$ \\cite{B-Minakov}, \\cite{Noonan}. In the case when both $n_g$ and $n$ are close to unity, the combined deviation of the refraction index from unity is reduced to the sum of both separate effects \\cite{Muhl1966}, \\cite{Muhl1970}, \\cite{B-Minakov}.  In this work we consider the gravitational lensing in a homogeneous plasma. Plasma is a dispersive medium, where the refraction index depends on the frequency of the photon. Therefore in the plasma the photons with different frequencies move with different velocities, namely the photons with smaller frequency (or bigger wavelength) move with smaller group velocity of the light signal. We obtain here, that in a homogeneous plasma, in the presence of gravity, the deflection angle of the photon depends on the frequency of the photon, and discuss observational effects of this phenomenon.  In the works, where the effect of the light dispersion in the plasma was not taken into account, the dependence of the gravitational deflection on frequency was connected only with the plasma inhomogeneity, and disappeared in the uniform plasma. In the book of   Synge \\cite{Synge}, the geometrical optics in the medium with gravity was considered in great details, and he had derived equations for the photon propagation in an arbitrary medium with gravity. Here we calculate, on the basis of equations of Synge \\cite{Synge}, the deflection angle for the photon wave packet, moving in the gravitational field in the presence of a uniformly distributed plasma. We obtain the dependence of the deflection angle on the frequency for the case of a week field of the Schwarzschild black hole metric. The deflection angle increases with decreasing of the frequency, and when the frequency is approaching the electron plasma frequency $\\omega_e^2=\\frac{4\\pi e^2 n_e}{m_e}$, all photons are falling into the black hole, if their impact parameter is less than the critical one, depending on frequency, $b < b_c(\\omega)$. When $\\omega$ approaching $\\omega_e$, the critical $b_c$ formally goes to $\\infty$. ",
        "conclusions": "Let us consider the case of a weekly inhomogeneous medium without gravity, with the refraction index $n = n_0 + n_1$, $n_0 =$const, $n_1 \\ll n_0$. The system of equation for the trajectory of the photon in this case follows from equations (\\ref{syst-gen}) with the flat metric $g_{ik}=\\eta_{ik}$, see \\cite{LL8},  \\begin{equation} \\label{inhom-n} \\frac{d x^\\alpha}{dz} = p^\\alpha, \\; \\; \\frac{d p^\\alpha}{dz} = - \\frac{1}{2} \\, \\eta^{\\beta \\gamma}_{, \\alpha} p_\\beta p_\\gamma + \\frac{1}{2} \\left(n^2 \\chi^2\\right)_{, \\alpha} = \\frac{1}{2 n_0^2} \\frac{\\partial n^2}{\\partial x^{\\alpha}} \\simeq \\frac{1}{n_0} \\frac{\\partial n}{\\partial x^{\\alpha}} \\, . \\end{equation} The light propagation in a week gravitational field may be considered, as a propagation in a flat space with the \"gravitational\" refraction index, which for the Schwarzschild metric is written as \\cite{GL, Muhl1966, Muhl1970, Noonan, B-Minakov, Fok}  \\begin{equation} \\label{ng} n_g=1+\\frac{r_g}{r}, \\quad \\frac{r_g}{r} \\ll 1. \\end{equation} The total effective refraction index $n_{eff}$, in presence of plasma with the proper refraction index $n$, is determined \\cite{Noonan},\\cite{B-Minakov} as $n_{eff}=nn_g$. When plasma in a week gravitational field has a refraction index close to unity, it follows from the definition of $n_{eff}$, that in the linear approximation the effective refraction index is obtained as a sum of two different additions to the unity \\cite{Muhl1966}, \\cite{Muhl1970},\\cite{B-Minakov}:  \\begin{equation}\\label{ntot} n_{eff} = 1 + \\frac{r_g}{r} - \\frac{\\omega_e^2(r)}{2\\omega^2}, \\; \\; \\frac{r_g}{r} \\ll 1, \\; \\; \\frac{\\omega_e^2(r)}{\\omega^2} \\ll 1. \\end{equation} From (\\ref{inhom-n}) we obtain the deflection angle (\\ref{angle-fin-BM}). Note that when $|n^2-1|$ is not small, $n_{eff}$ cannot be represented in the form (\\ref{ntot}), and such case was considered in the subsection 3.2.  The main new result of this work is obtaining of the dependence, of the lensing angle on the frequency in a homogeneous plasma in the gravitational field. This effect has a relativistic nature, and is connected with the dispersive properties of plasma. It is interesting, that, as shown in the subsection 3.1, in the medium without dispersion, the trajectories of photons with different frequencies (energies) are exactly the same as in the vacuum, while their velocities are less that the vacuum light velocity $c$.    Observational effect of such frequency dependence is easy to explain on the example of lensing by the Schwarzschild point-mass lens. This lens gives two images of source, on the opposite sides from lens. Angular positions of images depend on the Schwarzschild radius of lens and positions of source, lens and observer. The dependence of the deflection angle on the frequency in plasma lead to the phenomenon, that instead of two concentrated images with complicated spectra, we will have two line images, formed by the photons with different frequencies, which are deflected by different angles (Fig.1,2).  The description of the mass distribution as a point mass (Schwarzschild lens) is rarely sufficient for gravitational lensing considerations \\cite{GL}. In reality the gravitational lenses have more a complicated structure, and position of images different from that of the point-mass lens. But all standard models of gravitational lenses are based on the same Einstein deflection angle $2r_g/b$, which should be modified for sufficiently long waves, according to our formula (\\ref{main-res}), in the presence of plasma. Note also, that taking into account of plasma effects on the gravitational lensing may influence the spectrum of the microwave background radiation, leading to the dependence of power spectrum of the fluctuations on the photon wavelength.   The light signals are propagated with the group velocity. It follows from (\\ref{main-res}), that the smaller group velocity (smaller frequency and bigger wavelength) corresponds to a larger deflection angle. Hence the effect of difference in the gravitational deflection angles is significant for larger wavelengths, when $\\omega$ is approaching $\\omega_e$, what is possible only for the radio waves. Therefore, the gravitating center in plasma is acting as a radiospectrometer. The longest radiowaves are registered \\cite{Braude} in the band $\\sim 3 \\cdot 10^3$ cm, corresponding to $\\nu \\simeq 10^7$ Hz $= 10$ MHz and $\\omega = 2 \\pi \\nu \\simeq 6 \\cdot 10^7$ sec$^{-1}$. The spectroscopic effects of lensing will be important when $N_e \\geq 3 \\cdot 10^5$ cm$^{-3}$, corresponding to 10 $\\%$ difference in the lensing angles. Such electron densities may be expected around the supermassive black holes, or during lensing at earlier stages of the universe expansion at $z \\geq 10^3$.   \\begin{figure} \\centerline{\\hbox{\\includegraphics[width=0.6\\textwidth]{fig1.eps}}} \\caption{Lensing of the point source by the Schwarzschild point-mass lens. Instead of two point images due to lensing in the vacuum we have two line images. The pairs of images, corresponding to the same photon frequency, are indicated by the same numbers. Two images with number 1 correspond to the vacuum lensing. } \\end{figure}   \\begin{figure} \\centerline{\\hbox{\\includegraphics[width=0.6\\textwidth]{fig2.eps}}} \\caption{Axis on lensing by the Schwarzschild point-mass lens. The case of the Einstein ring. Instead of a thin ring corresponding to the vacuum lensing (the inner circle of the ring) we have a thick ring, formed by the photons of different frequencies.} \\end{figure}"
    },
    "0809/0809.0161_arXiv.txt": {
        "abstract": "{The radial-orbit instability is a collective phenomenon that has heretofore only been observed in spherical systems. We find that this instability occurs also in triaxial systems, as we checked by performing extensive $N$-body simulations whose initial conditions were obtained by sampling a self-consistent triaxial model of a cuspy galaxy composed of luminous and dark matter. $N$-body simulations show a time evolution of the galaxy that is not due to the development of chaotic motions but, rather, to the collective instability induced by an excess of box-like orbits. The instability quickly transforms such models into a more prolate  configuration, with $0.64<b/a<0.77$ and $0.6<c/a<0.7$ for the dark halo and $0.64<b/a<0.77$  and $0.59<c/a<0.67$ for the luminous matter. Stable triaxial, cuspy galaxies with dark matter halos are obtained when the contribution of radially-biased  orbits to the solution is reduced. These results constitute the first evidence of the radial-orbit instability in triaxial galaxy models.} ",
        "introduction": "The aim of the current work is to test, by extensive $N$-body simulations, the stationarity of the triaxial galaxy models realized by Capuzzo Dolcetta et al. (2007, hereafter CLMV07). These  models were constructed by  means of the standard orbital superposition method introduced by Schwarzschild (1979), and represent triaxial, cuspy elliptical galaxies embedded in triaxial dark halos in which the dark  and the luminous components have the same axial ratios ($a=1, b=0.86, c=0.7$, i.e., maximal triaxiality). In CLMV07 these models were referred to as MOD1 and MOD1-bis; the only difference between the two solutions was the maximum integration time permitted during the orbital integration, which was longer in MOD1-bis (5 Hubble times) than in MOD1 (2 Hubble times), i.e., MOD1-bis represents a more stationary solution. ",
        "conclusions": "Our study has important implications for future construction of self-consistent models of realistic galaxies. In particular, it has been shown that the dynamical properties of solutions cannot be directly deduced by the Schwarzschild method itself since the temporal development of models can be strongly affected by the distributions in velocity space. In this context, $N$-body tests of the stability of the models seems to be indispensable tool for drawing conclusions about general properties of models constructed via orbital superposition.  Regarding the ROI, because a complete understanding of its mechanism is still lacking, it would be interesting to  investigate more in depth its role  in the case of triaxial potentials. Our simulations  show  dynamical features which cannot be detected in the case of spherical symmetry. The method exploited here to change the semi-radial orbit contribution to the models may be useful for constructing other radially-unstable triaxial systems."
    },
    "0809/0809.5058_arXiv.txt": {
        "abstract": "We present results from a new ultra-deep $\\approx$400~ks \\chandra\\ observation of the SSA22 protocluster at $z=3.09$.  We have studied the \\xray\\ properties of 234 $z \\sim 3$ Lyman break galaxies (LBGs; protocluster and field) and 158 $z = 3.09$ Ly$\\alpha$ emitters (LAEs) in SSA22 to measure the influence of the high-density protocluster environment on the accretion activity of supermassive black holes (SMBHs) in these UV-selected star forming populations.  We detect individually \\xray--emission from active galactic nuclei (AGNs) in six LBGs and five LAEs; due to small overlap between the LBG and LAE source population, ten of these sources are unique.  At least six and potentially eight of these sources are members of the protocluster.  These sources have rest-frame \\hbox{8--32~keV} luminosities in the range of \\hbox{$L_{\\rm 8-32~keV} =$~(3--50)~$\\times 10^{43}$~\\xlum} and an average observed-frame \\hbox{2--8~keV} to \\hbox{0.5--2~keV} band-ratio of $\\approx$0.8 (mean effective photon index of $\\Gamma_{\\rm eff} \\approx$~1.1), suggesting significant absorption columns of \\hbox{$N_{\\rm H} \\simgt 10^{22}$--$10^{24}$~cm$^{-2}$}.  We find that the fraction of LBGs and LAEs in the $z = 3.09$ protocluster harboring an AGN with \\hbox{$L_{\\rm 8-32~keV} \\simgt 3 \\times 10^{43}$~\\xlum} is $9.5^{+12.7}_{-6.1}$\\% and $5.1^{+6.8}_{-3.3}$\\%, respectively.  These AGN fractions are somewhat larger (by a mean factor of $6.1^{+10.3}_{-3.6}$; significant at the $\\approx$95\\% confidence level) than $z \\sim 3$ sources found in lower-density ``field'' environments.  Theoretical models imply that these results may be due to the presence of more actively growing and/or massive SMBHs in LBGs and LAEs within the protocluster compared to the field.  Such a result is expected in a scenario where enhanced merger activity in the protocluster drives accelerated galaxy and SMBH growth at \\hbox{$z \\simgt$~2--3}.  Using \\spitzer\\ IRAC imaging we found that the fraction of IRAC detected LBGs is significantly larger in the protocluster than in the field (by a factor of 3.0$^{+2.0}_{-1.3}$).  From these data, we constrained the median rest-frame $H$-band luminosity in the protocluster to be \\hbox{$\\simgt$1.2--1.8} times larger than that for the field.  When combined with our \\xray\\ data, this suggests that both galaxies and SMBHs grew more rapidly in protocluster environments. ",
        "introduction": "Over the last decade, investigations have revealed that galaxies with spheroid components (i.e., elliptical galaxies, lenticulars, and spiral galaxy bulges) ubiquitously contain supermassive black holes (SMBHs) in their cores (e.g., Kormendy \\& Richstone 1995; Magorrian \\etal\\ 1998).  These studies have also confirmed the existence of a tight relationship between the mass of the SMBHs and the stellar mass of the spheroid, suggesting a causal connection between the growth of these two galactic components (e.g., Gebhardt \\etal\\ 2000).  In addition, theories of large-scale structure formation in a CDM Universe predict that galaxy formation is accelerated in high-density regions (Kauffmann~1996; de~Lucia \\etal\\ 2006).  Observational studies have provided convincing support for this hypothesis, showing that there is a strong relationship between galaxy stellar age and local environment in the nearby universe (e.g., Smith \\etal\\ 2008 and references therein); the most evolved and massive galaxies reside in the highest density regions of local clusters, while more typical galaxies that are undergoing significant star formation are generally found in lower density environments (e.g., Lewis \\etal\\ 2002).  Studies of distant galaxy populations at $z \\approx$~1 are finding that the star-formation activity occurs in higher density environments than seen locally (e.g., Geach \\etal\\ 2006; Elbaz \\etal\\ 2007; Heinis \\etal\\ 2007; Cooper \\etal\\ 2008; Poggianti \\etal\\ 2008).  These studies indicate that a reversal in the star-formation rate (SFR)/galaxy density relation occurs at $z \\simgt 1$, where the most intense galaxy growth is expected to occur in the highest density clusters.  The progenitors of the highest density clusters in the local universe are also expected to be the highest density structures at $z \\simgt$~2--3 and should be undergoing vigorous star formation during their assemblage (e.g., Governato \\etal\\ 1998).  These protoclusters are identified through overdense redshift ``spikes'' in high-redshift galaxy surveys of blank fields (e.g., Adelberger \\etal\\ 1998; Steidel \\etal\\ 2003) and in the vicinity of certain powerful radio galaxies (e.g., Venemans \\etal\\ 2007).  It is therefore plausible to expect that if the growth of galaxies and their central SMBHs are causally linked, then the highest density structures will also be the sites of significant SMBH accretion, identifiable as active galactic nuclei (AGNs).  To detect and study typical AGN ($L_{\\rm X} \\simgt 10^{43-44}$~\\xlum) in $z \\simgt$~2--3 protocluster galaxies requires significant optical spectroscopic and \\xray\\ observational investments.  Therefore, few programs have yet explored how AGN activity varies as a function of environment at these redshifts.  Nonetheless, initial studies of high-density regions and clusters in the \\hbox{$z \\simgt$~0.5--3} universe have provided suggestive evidence for an elevation in the AGN activity in such high-density environments compared to the field (see, e.g., Pentericci \\etal\\ 2002; Gilli \\etal\\ 2003; Johnson \\etal\\ 2003; Smail \\etal\\ 2003; Croft \\etal\\ 2005; Eastman \\etal\\ 2007; Silverman \\etal\\ 2008).  However, a rigorous quantification of such an enhancement in AGN activity has yet to be obtained for actively forming protoclusters that are precursors to rich clusters at $z = 0$.  The $z = 3.09$ SSA22 protocluster was originally identified by Steidel \\etal\\ (1998) as a significantly overdense region (by a factor of $\\approx$4--6) through spectroscopic follow-up observations of $z \\sim 3$ candidate Lyman break galaxies (LBGs).  Theoretical modelling indicates that the protocluster should collapse into a $z = 0$ structure resembling a rich local cluster with a total mass $\\simgt$$10^{15}$~\\msol\\ (e.g., Coma; see Steidel \\etal\\ 1998 for details).  Since its discovery, the protocluster has been found to contain a factor of $\\approx$6 overdensity of Ly$\\alpha$ emitters (LAEs; Steidel \\etal\\ 2000; Hayashino \\etal\\ 2004; Matsuda \\etal\\ 2005) and several remarkable bright extended Ly$\\alpha$-emitting blobs (LABs; Steidel \\etal\\ 2000; Matsuda \\etal\\ 2004), which are hypothesized to be sites of either cooling flows or starburst/AGN outflows (e.g., Bower \\etal\\ 2004; Geach \\etal\\ 2005, Wilman \\etal\\ 2005; Geach \\etal\\ in preparation).  Therefore, SSA22 is an ideal field for studying how SMBH growth depends on environment in the $z \\simgt$~2--3 universe.  \\begin{figure} \\centerline{ \\includegraphics[width=9cm]{./fig01.ps} } \\caption{ Adaptively-smoothed \\hbox{0.5--8~keV} \\chandra\\ image (with boundaries shown as a {\\it dashed polygon\\/}) of the SSA22 field.  X-ray detected LBGs and LAEs have been highlighted with open circles and squares, respectively.  The dotted regions show the extent of the Steidel \\etal\\ (2003) ``SSA22a'' and ``SSA22b'' LBG surveys.  We note that the Hayashino \\etal\\ (2004) sample of LAEs covers a region larger than that shown here.  LAE source-density contours are shown in gray (computed using our LAE catalog and a spatially-varying density extraction circle with radius of 3\\farcm0) and have levels of 1.0, 1.3, 1.7, and 2.0 LAEs arcmin$^{-2}$.  For comparison the LAE source-density in the \\ecdfs\\ field comparison sample is $\\approx$0.3 LAEs arcmin$^{-2}$.  We note that most of the \\xray\\ sources lie in regions of high LAE density.  For reference, we have highlighted the location of \\hbox{J221739.04+001330.1}, which is both an LBG and LAE. } \\end{figure}  In this paper, we utilize new ultra-deep $\\approx$400~ks \\chandra\\ observations of the $z = 3.09$ SSA22 protocluster region to identify luminous AGNs indicative of accreting SMBHs.  We use this data to place constraints on how the AGN properties (e.g., luminosity and \\xray\\ spectra) and frequency in the protocluster compares with AGNs identified in lower density field regions of the \\chandra\\ Deep Field-North (\\cdfn) and Extended \\chandra\\ Deep Field-South (\\ecdfs).  The Galactic column density for SSA22 is \\hbox{$4.6 \\times 10^{20}$~cm$^{-2}$} (Stark \\etal\\ 1992).  All \\xray\\ fluxes and luminosities quoted in this paper have been corrected for Galactic absorption.  The coordinates throughout this paper are J2000.0.  \\hbox{$H_0$ = 70~km s$^{-1}$ Mpc$^{-1}$}, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7 are adopted (e.g., Spergel \\etal\\ 2003), which give the age of the Universe as 13.5~Gyr and imply a $z=3.09$ look-back time and spatial scale of 11.4~Gyr and 7.6~kpc~arcsec$^{-1}$, respectively. ",
        "conclusions": ""
    },
    "0809/0809.2487_arXiv.txt": {
        "abstract": "{This paper presents morphological type, membership, and $U-V$ color for a sample of galaxies in the Coma cluster direction, complete down to $M_B=-15.00$ mag and extending down to $M_B=-14.25$ mag. We have examined 1155 objects from the GMP 1983 catalog on B and V images of the CFH12K camera, and obtained the Hubble type in most cases. Coma cluster membership for 473 galaxies was derived using morphology, apparent size, and surface brightness, and, afterward, redshift. The comparison among morphology- and redshift- memberships and among luminosity functions derived from this morphologically-selected sample, or by using statistical members or spectroscopic members, all show that the morphological membership provided here can be trusted. For the first time, the morphological classification of Coma galaxies reaches magnitudes that are faint enough to observe the whole magnitude range of the giant types, E, S0, and spiral stages. The data presented in this paper makes our sample the richest environment where membership and morphology for complete samples down to faint magnitudes ($M_B\\sim-15$ mag) are available, thereby enlarging the baseline of environmental studies.} ",
        "introduction": "Morphological classification is probably the oldest tool used in galaxy studies, as recently reviewed in Sandage (2005). Galaxies populate only a limited number of forms. With these forms are associated, or correlated, a number of such important properties as luminosities, internal dynamics, and material composition, i.e. stars and ISM. In the study of such systems of galaxies as clusters, one would like to obtain statistically significant data, including of course the morphological types of all objects. But as emphasized by Buta (2000) the art of morphology is strongly dependent on the resolution of the images of galaxies that span such a large domain of intrinsic sizes and distances. The {\\em nec plus ultra} in galaxy morphology still is the catalog of Binggeli et al. (1985 or BST2) for the Virgo cluster. An important novelty of this work was the use of morphology, apparent size, and surface brightness to ascertain the membership of faint galaxies in the cluster. It is indeed not feasible to extend redshift measurements to all detected objects, and the addition of appearance as a tool of distance estimates is welcome. The observation of new radial velocities (Binggeli et al. 1993) then confirmed nearly all the 1985 results.  The Coma cluster is one of the best-studied such objects because of its relative proximity and because it may be a prototype of dynamically relaxed clusters, although significant substructures are present (Colless \\& Dunn, 1996, Biviano et al. 1996). This is in marked contrast to other clusters such as Virgo with its complex geometry and important substructure (Binggeli \\& Huchra, 2000). Compared to the fair amount of photometric, colorimetric, and spectrographic information available for Coma cluster galaxies, morphological data are far from complete for these. Dressler (1980a, 1980b) described the overpopulation of early-type, E, and S0, galaxies in the Coma and other clusters in relation to their density, and introduced the morphology-density relation that is also discussed by Whitmore et al. (1993) and many other later works. These later studies used detailed photometry to refine the classification of relatively bright objects (Jorgensen \\& Franx, 1994; Michard, 1995; Andreon et al. 1996, 1997; Andreon 1996; Guttierez et al. 2004; Aguerri et al. 2004).  The aim of the present paper is to classify as many galaxies as it is feasible in the Coma cluster, increasing the number of members detected from velocity data by considering object morphology, apparent size, and surface brightness. In subsequent papers of this series, we will study the resulting member galaxies populations. These can be considered from two aspects: in-cluster variations and cluster-to-cluster comparisons. For this second approach we chose the Virgo cluster because of the superior quality and completeness of the BST2 data.  In the present paper, Sect. 2 summarizes the technical aspects of the classification, describes its output catalogs and presents an independent test of the validity of the morphological membership in the Coma cluster, i.e. the determination of the membership based on galaxy morphology, apparent size, and surface brightness. Section 3 presents our essential data: the catalog, the histograms of the various Hubble-types for Coma and the Virgo comparison sample, and a table showing the sample size of each morphological type in both clusters. We also introduce an auxiliary catalog of V magnitudes and U$-$V colors derived from the data by Terlevich et al. (2001) (TCB01) and others. Section 4 briefly summarizes the results and gives some highlights of a later paper(s) in this series.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{studiedfield.eps}} \\caption[ ]{ B-band image of the studied portion of the Coma clusters. The image is highly binned ($16 \\times 16$) for display purposes. The 3 zones used in the discussion of radial population changes are marked. North is at the top, east to the left. } \\label{studiedfield} \\end{figure} ",
        "conclusions": "A catalog of morphological types were obtained for 1155 objects in the Coma clusterfield, and memberships estimated from their appearance, i.e. morphology, surface brightness, rarely color, and, aftewards, redshifts. Four hundred seventy-three cluster members were identified, and it has been shown that our memberships can be trusted both when individual redshifts are available (Sect. 2.3.2), and when membership is only known in a statistical sense (Sect. 2.3.3). We obtained from the literature V$_t$/U$-$V data for 2/3 of our Coma members and corrected them for aperture effects.  For the first time, the morphological classification of Coma galaxies reaches magnitudes faint enough to observe the full magnitude range of the giant types, E, S0, and spiral stages. Tables with morphological types, membership, and U-V colors are given. Our sample is the richest environment where membership and morphology for complete samples down to faint magnitudes ($M_B\\sim-15$ mag) are available, thereby enlarging the baseline of environmental studies.    Some new results  will be presented in detail in future papers. a) Dwarf galaxies are more abundant in Virgo than in Coma. b) It appears that late-type dwarfs in Coma are less numerous, less bright, and less active in star formation than in Virgo; c) We derived LF parameters of the giant galaxies in Coma, distinguishing between various Hubble-types and covering the full luminosity range for each type. This has not been possible up to now, except for the Virgo cluster in the classical work of Sandage et al. (1985), which we follow closely in our study. The LF of dwarfs can also be studied paying the necessary attention to sample incompleteness."
    },
    "0809/0809.2170_arXiv.txt": {
        "abstract": "{The discovery of photospheric absorption lines in \\emph{XMM-Newton} spectra of the X-ray bursting neutron star in \\exo\\ by Cottam and collaborators allows us to constrain the neutron star mass-radius ratio from the measured gravitational redshift. A radius of $R=9-12$\\,km for a plausible mass range of $M=1.4-1.8$\\,M$_\\odot$ was derived by these authors.} {It has been claimed that the absorption features stem from gravitationally redshifted ($z=0.35$) $n=2-3$ lines of H- and He-like iron. We investigate this identification and search for alternatives.} {We compute LTE and non-LTE neutron-star model atmospheres and detailed synthetic spectra for a wide range of effective temperatures (\\Teffw{1-20}) and different chemical compositions.} {We are unable to confirm the identification of the absorption features in the X-ray spectrum of \\exo\\ as $n=2-3$ lines of H- and He-like iron (\\ion{Fe}{xxvi} and \\ion{Fe}{xxv}). These are subordinate lines that are predicted by our models to be too weak at any \\Teff. It is more likely that the strongest feature is from the $n=2-3$ resonance transition in \\ion{Fe}{xxiv} with a redshift of $z=0.24$. Adopting this value yields a larger neutron star radius, namely $R=12-15$\\,km for the mass range $M=1.4-1.8$\\,M$_\\odot$, favoring a stiff equation-of-state and excluding mass-radius relations based on exotic matter. Combined with an estimate of the stellar radius $R > 12.5$\\,km from the work of \\\"Ozel and collaborators, the $z=0.24$ value provides a minimum neutron-star mass of $M > 1.48$\\,M$_\\odot$, instead of $M > 1.9$\\,M$_\\odot$, when assuming $z=0.35$.} {The current state of line identifications in the neutron star of \\exo\\ must be regarded as highly uncertain. Our model atmospheres show that lines other than those previously thought must be associated with the observed absorption features.} ",
        "introduction": "\\label{sect:intro}  \\citet[][hereafter CPM02]{cpm2002} identified discrete absorption features corresponding to electronic transitions in highly ionized iron in the burst spectra of the neutron star in \\exo\\ observed with \\emph{XMM-Newton}. They identified $n=2-3$ absorption features of H-like (\\ion{Fe}{xxvi}) and He-like iron (\\ion{Fe}{xxv}) in early and late phases of the bursts, respectively. These features provided a redshift measurement of $z=0.35$, corresponding to a mass-radius ratio of $M/R=0.152$\\,M$_\\odot$\\,km$^{-1}$. Using this redshift and an estimate for the stellar radius of $R>13.8$\\,km, \\citet{ozel06} inferred a neutron star mass of $M>2.10$\\,M$_\\odot$. However, using \\\"Ozel's numbers and formulae, we obtain slightly different values ($R>12.5$\\,km, $M>1.9$\\,M$_\\odot$).  The identification of only a few observed lines in burst spectra of NS with unknown redshift is potentially ambiguous. It is the aim of our paper to confirm the proposed line identifications in \\exo. To this end, we performed LTE and non-LTE model-atmosphere calculations in a wide parameter range using two independently developed stellar atmosphere modeling codes. In this systematic study, we elaborate on our earlier suspicion that an alternative line identification is more likely \\citep{wea2007}.  In the past two decades, spectral analysis of hot, compact stars by means of fully line-blanketed NLTE model atmospheres \\citep[e.g.][]{r2003} has achieved a high level of sophistication. For our analyses, the T\\\"ubingen NLTE Model Atmosphere Package \\citep[\\emph{TMAP}\\footnote{http://astro.uni-tuebingen.de/$^\\sim$rauch/TMAP/TMAP.html},][]{wea2003, rd2003} was used to calculate plane-parallel NLTE model atmospheres that are in radiative and hydrostatic equilibrium.  Such model atmospheres were used successfully in the analysis of hot white dwarfs, e.g\\@.  \\lsv\\ \\citep[\\TeffwkK{95},][]{rea2007} and  \\kpd\\ \\citep[\\TeffwkK{200},][]{wrk2008}.  \\emph{TMAP} models were also calculated for the extremely hot super-soft X-ray source \\vsgr\\ \\citep[\\TeffwkK{610},][]{rea2005}.  The \\emph{TMAP} NLTE model atmospheres can also be employed in the analysis of neutron stars with low magnetic fields,  i.e\\@. in the range where the magnetic field strength has no significant impact on atomic data  ($B \\sla 10^{12}\\,\\mathrm{G}$). Since magnetic fields in low-mass X-ray binaries (LMXBs) are believed to be small, X-ray spectra of the neutron star in \\exo\\ can be compared with our synthetic spectra.  We calculated \\emph{TMAP} models for the relevant \\Teff\\ range and investigated their \\Teff-dependence \\sK{sect:models}. We describe results of a comparison of LTE and NLTE  model-atmosphere fluxes in \\se{sect:LTEvsNLTE}. A comparison of \\exo\\ X-ray observations with our models follows in \\se{sect:exo0748-676} and we conclude in \\se{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have performed model-atmosphere calculations to describe the X-ray spectra of thermal radiation from neutron stars. We have compared our computed spectra with X-ray burst spectra of the LMXB \\exo. We have been unable to confirm the line identification by CPM02 as being due to subordinate transitions of H- and He-like iron. These line features were too weak at any \\Teff\\ and iron content. Our models suggested that a more likely identification was the resonance line of Li-like iron. As a consequence, the measured line redshift was $z=0.24$ rather than $z=0.35$. This implied a larger neutron star radius of $R=12-15$\\,km for the mass range $M=1.4-1.8$\\,M$_\\odot$.  We compared results from two entirely different model codes, a LTE and a NLTE code, and concluded that NLTE effects were less important than uncertainties in the chemical composition of the bursting neutron-star atmospheres. Given the current state of observations, NLTE effects on continuum shape and spectral line profiles are negligible. On the other hand, we investigated the influence of the various chemical compositions and found that the derived conclusion concerning line identification does not depend on the chemical composition.  We investigated the relevance of Compton electron scattering and found that it was unimportant for solar composition models with $T_{\\rm eff} < 15\\,\\mathrm{MK}$.  \\begin{figure}[ht!] \\resizebox{\\hsize}{!}{\\includegraphics{xm_15.eps}} \\caption[]{Comparison of redshifted ($z=0.24$) NLTE model-atmosphere fluxes calculated from \\Teffw{8 - 14} models of solar iron abundance. Note that for clarity these lines have not been convolved with either the instrument's resolution or any rotational profile (cf\\@. \\ab{fig:nature}). } \\label{fig:8-14MK} \\end{figure}  The comparison of model-atmosphere spectra with blackbody flux distributions \\sA{fig:10MK} has shown that model-atmosphere spectra peak at higher energies and have a higher peak intensity. A determination of \\Teff\\ by assuming blackbody spectra, as performed by CPM02, therefore overestimates \\Teff\\ \\citep[cf\\@.][]{rea2005}.   \\begin{figure}[ht!] \\resizebox{\\hsize}{!}{\\includegraphics{xm_16.eps}} \\caption[]{Comparison of unredshifted NLTE model-atmosphere fluxes calculated from \\Teffw{10} models with different iron content (thin: solar, thick: 30\\,\\% Fe content, short dashes: 9\\,\\% Fe) The long-dashed line represents a blackbody spectrum. } \\label{fig:10MK} \\end{figure}   \\begin{figure}[ht!] \\resizebox{\\hsize}{!}{\\includegraphics{xm_17.eps}} \\caption[]{Allowed values for $M$ and $R$ of \\exo\\ for redshifts $z=0.24$ and $z=0.35$ (straight lines; thick portions of the graphs denote the mass range $1.4-1.8$~M$_\\odot$) compared to various theoretical $M-R$ relations \\citep{Hae06}. The thick dot on the $z=0.35$ line denotes the minimum $M$ and $R$ derived by \\citet{ozel06}. The arrow indicates the shift of this result when we assume $z=0.24$. A description of the theoretical mass-radius relations is given in the text.} \\label{fig:eos} \\end{figure}"
    },
    "0809/0809.2036_arXiv.txt": {
        "abstract": "We study six groups and clusters of galaxies suggested in the literature to be `fossil' systems (i.e. to have luminous diffuse X-ray emission and a magnitude gap of at least 2 mag-R between the first and the second ranked member within half of the virial radius), each having good quality X-ray data and SDSS spectroscopic or photometric coverage out to the virial radius. The poor cluster AWM\\,4 is clearly established as a fossil system, and we confirm the fossil nature of four other systems (RX\\,J1331.5+1108, RX\\,J1340.6+4018, RX\\,J1256.0+2556 and RX\\,J1416.4+2315), while the cluster RX\\,J1552.2$+$2013 is disqualified as fossil system.  For all systems we present the luminosity functions within 0.5 and 1 virial radius that are consistent, within the uncertainties, with the universal luminosity function of clusters. For the five {\\em bona fide} fossil systems, having a mass range $2\\times10^{13}-3\\times10^{14}~\\mathrm{M_\\odot}$, we compute accurate cumulative substructure distribution functions (CSDFs) and compare them with the CSDFs of observed and simulated groups/clusters available in the literature. We demonstrate that the CSDFs of fossil systems are consistent with those of normal observed clusters and do not lack any substructure with respect to simulated galaxy systems in the cosmological $\\mathrm{\\Lambda}$CDM framework. In particular, this holds for the archetype fossil group RX\\,J1340.6$+$4018 as well, contrary to earlier claims. ",
        "introduction": "Early numerical simulations suggested that the most compact galaxy groups could merge to form a single elliptical galaxy (hence a `fossil group') in a few billion years \\citep{barnes_89}.  An elliptical galaxy formed by the merger of such a group retains its X-ray emitting halo of hot gas, which is unaffected by merging \\citep{ponman_bertram_93}.  Following this indication, \\cite{ponman_etal_94} discovered the archetype fossil group RX\\,J1340.6$+$4018.  \\cite{vikhlinin+99} and \\cite{jones+03}, on the basis of ROSAT observations, have suggested that fossil groups constitute a considerable population of objects. Their X-ray extent, bolometric X-ray luminosity ($L_{\\mathrm{X, bol}} > 10^{42} h_{50}^{-2}~\\mathrm{erg}~\\mathrm{s}^{-1}$), dark matter dominated total mass, and mass in the diffuse hot gas component are comparable to those of bright groups and poor clusters of galaxies ($\\sim 10^{13}$--$10^{14}~h_{70}^{-1}~\\mathrm{M}_{\\sun}$). The brightest member of a fossil group has an optical luminosity comparable to that of a cluster cD galaxy (i.e., $M_\\mathrm{R} < -22.5 + 5 \\mathrm{log}~h_{50}$) and dominates the galaxy luminosity function of the system. The observational definition of a fossil system lies in the detection of extended, very luminous X-ray emission from the hot gas of the intracluster medium (ICM), and in the existence of an R-band magnitude gap $\\Delta m_{12} > 2$ between the brightest and second brightest members within $0.5 R_\\mathrm{vir}$ \\citep{jones_ponman_forbes_00}.  As noted by \\cite{vikhlinin+99}, fossil groups may represent the ultimate examples of hydrostatic equilibrium in virialised systems, since they must have been undisturbed for a very long time if they are the result of galaxy merging within a group.  High-resolution hydrodynamical cosmological simulations in the $\\mathrm{\\Lambda}$CDM framework have shown that fossil groups have already assembled half of their final DM mass at redshifts $z \\ge 1$, and subsequently they typically grow by minor mergers only, whereas non-fossil systems of similar masses on average form later \\citep{donghia+05}.  The early assembly of fossil groups leaves sufficient time for objects with luminosities close to the characteristic luminosity (i.e., $L \\sim L^{\\star}$) to merge into the central galaxy by dynamical friction, producing the magnitude gap which defines a fossil system.  In addition, the simulated fossil groups were found to be over-luminous in X-rays relative to non-fossil groups of the same optical luminosity, in qualitative agreement with observations \\citep[cf.][]{vikhlinin+99,jones_ponman_forbes_00}.  In a recent paper, \\cite{vonbendabeckmann+08} showed that many galaxy groups may undergo a fossil phase in their lives but may not necessarily stay fossil down to $z=0$, owing to renewed infall of $L \\sim L^{\\star}$ galaxies from the large-scale environment. Such infall episodes are statistically more likely for more massive systems, so that the fraction of quasi-fossil systems (i.e. those with a large luminosity gap between the central galaxy and the most luminous satellite) is lower among clusters than among groups \\citep{milosavljevic+06,yang+08}.  Fossil groups have become a puzzling problem to cosmology since \\citet[DL04]{donghia_lake_04} showed that, with respect to state-of-the-art predictions on the frequency of substructures in cold dark matter (CDM) halos \\citep{delucia+04a}, a virialised system like RX\\,J1340.6$+$4018 lacks galaxies nearly as luminous as the Milky Way. Conversely, the same numerical simulations are able to accurately describe the frequency of substructures in galaxy clusters as massive as Virgo or Coma to well below the circular velocity of a Milky Way-size dark halo (i.e., with $\\vcirc \\le 220~\\mathrm{km}~\\mathrm{s}^{-1}$).  In this respect, fossil groups appeared to exacerbate the so-called `small-scale crisis' of CDM universes \\citep{klypin+99b}. In fact, DL04 concluded that the missing substructure problem affects systems up to the scale of groups (typically with $\\mathrm{k} T_\\mathrm{X} \\le 1~\\mathrm{keV}$). However, this result has been challenged by \\cite{sales+07}, who find that the abundance and luminosity function of simulated fossil systems are in reasonable agreement with the few available observational constraints.  Given the great interest of this cosmological issue, some optical studies have aimed at better characterizing the mass and luminosity function (LF) of already known fossil systems \\citep[e.g.][]{mendesdeoliveira+06,cypriano+06} or identifying new ones \\citep[e.g.][]{santos+07}, although identifying low-mass fossil systems is hampered by the fact that groups are under-represented in existing X-ray catalogs. In spite of the increasing quality of the data and number of candidates, no new cumulative substructure distribution function (CSDF)\\footnote{The cumulative substructure distribution function gives the number of sub-halos with circular velocity $\\vcirc$ larger than a given fraction of the circular velocity of the parent halo $V_\\mathrm{parent}$.}  has so far been produced for a fossil system to compare with that obtained by DL04 for the archetype fossil group RX\\,J1340.6$+$4018.  The determination of the CSDF for fossil systems with a range of masses is the objective of the present study.  In the following, we adopt a $\\mathrm{\\Lambda}$CDM cosmological model ($\\Omega_{\\mathrm{m}} = 0.3$, $\\Omega_{\\mathrm{\\Lambda}} = 0.7$) with $H_0 = 70~h_{70}^{-1}~\\mathrm{km}~\\mathrm{s}^{-1}~\\mathrm{Mpc}^{-1}$. This model is consistent with the main {\\em WMAP5} results \\citep[cf.][]{hinshaw+08}.  This work is mainly based on Sloan Digital Sky Survey \\citep[SDSS]{SDSS} data from the sixth data release \\citep[DR6, and references therein]{DR6}. ",
        "conclusions": "We have studied six galaxy groups and clusters suggested in the literature to be fossil systems, and having high quality X-ray and SDSS spectroscopic or photometric information. Among them, we have established AWM\\,4 as fossil system. We confirm three other systems (RX\\,J1331.5$+$1108, RX\\,J1340.6+4018 and RX\\,J1416.4+2315) to be fossil. RX\\,J1256.0+2556 has characteristics very close to a genuine fossil system, although observational errors are still too large for a robust classification as fossil. Finally we demonstrate that RX\\,J1552.2$+$2013 does not match the magnitude gap requirement to qualify as a fossil.  For each system we have computed the luminosity functions within 0.5 and 1 virial radius. Although with uncertainties, that are remarkably large for the least massive systems, they are consistent with the universal luminosity function of clusters derived by \\cite{popesso+04_Xscale}.  We have derived detailed cumulative substructure distribution functions (CSDFs) for the six galaxy groups and clusters. Our motivation was to produce CSDFs for systems with a range of masses in the group and poor cluster regime, for comparison with the original CDSF derived for the archetype fossil system RX\\,J1340.6$+$4018 by \\citet[DL04]{donghia_lake_04}. In fact, these authors claimed a lack of substructure for RX\\,J1340.6$+$4018, suggesting that the so-called `small-scale crisis' of Milky-Way-size haloes \\citep{klypin+99b} exists up to the scale of X-ray bright groups. In addition we wanted to compare the CSDFs of our fossil systems with results from normal groups and clusters identified in the SDSS \\citep{desai+04c}, and in cosmological simulations \\citep{desai+04c,sales+07}.\\\\ Our conclusions are as follows: \\vspace{-0.2cm} \\begin{itemize} \\item The CSDF of AWM\\,4, based on spectroscopic data, is completely consistent both with \\cite{desai+04c}'s data and with simulations. AWM\\,4 also matches the simulations by \\cite{sales+07} of isolated bright galaxies in terms of central galaxy luminosity vs. magnitude gap between the first and the tenth ranked member. \\item The photometrically derived CSDFs of the other four {\\em bona fide} fossil systems are all completely consistent with the CSDF envelope derived from observed normal groups and clusters \\citep{desai+04c}. With respect to numerically simulated systems in \\cite{desai+04c} none of our fossil systems appears to lack any substructure. \\end{itemize} \\vspace{-0.2cm}  We therefore conclude that no evidence can be provided that the `small-scale crisis' is occurring on the scale of fossil systems.  In other words, the presence of a large magnitude gap between the first and second ranked members of a group/cluster does not imply anything special about its substructure, which can be fully accounted for by existing LCDM simulations.   Interestingly, we also observe a systematic shift of the CSDFs toward lower $\\vcirc/V_\\mathrm{parent}$ at increasing $V_\\mathrm{parent}$. This can be reproduced using the scaling relations of first and second brightest members vs. halo mass derived for a representative sample of groups and clusters by \\cite{yang+08}. Once more, this reinforces the idea that fossil systems are not anomalous as far as scaling relations are concerned.  Further deep, wide-field spectroscopic observations of fossil systems are required to confirm the results we have obtained in the present work.  This is particularly necessary at the low end of the mass range.  On the theoretical side, analyses which mimic observational approaches such as that described here would allow us to compare real data and simulations on an equal footing. This paper is the first in a series which will characterise many of these aspects of fossil groups."
    },
    "0809/0809.3801_arXiv.txt": {
        "abstract": "{We report the 2006-2008 light curves obtained with the REM telescope in VRIJHK bands for the two BL Lac objects PKS 0537-441 and PKS 2155-304}  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Since 2004 we are using the Rapid Eye Mount robotic telescope (REM, la Silla Chile, Zerbi et al 2004), for an intensive Blazar monitoring program in near-infrared (JHK) and optical (IRV) bands (see Tosti, these proceedings). The remarkable activities of PKS 0537- 441 (z=0.896), and PKS 2155-304 (z=0.116) observed with REM in 2004-2005 are described in Dolcini et al. (2005, 2007) and in Pian et al (2007).\\\\ Here we report preliminary results obtained for these two sources in 2006-2008. ",
        "conclusions": ""
    },
    "0809/0809.3495_arXiv.txt": {
        "abstract": "The eastern tip region of the Carina Nebula was observed with the Suzaku XIS for 77 ks to conduct a high-precision spectral study of extended X-ray emission. XMM-Newton EPIC data of this region were also utilized to detect point sources. The XIS detected strong extended X-ray emission from the entire field-of-view with a 0.2--5 keV flux of $0.7\\sim4\\times10^{-14}$ erg s$^{-1}$ arcmin$^{-2}$. The emission has a blob-like structure that coincides with an ionized gas filament observed in mid-infrared images. Contributions of astrophysical backgrounds and the detected point sources were insignificant. Thus the emission is diffuse in nature. The X-ray spectrum of the diffuse emission was represented by a two-temperature plasma model with temperatures of 0.3 and 0.6 keV and an absorption column density of 2$\\times10^{21}$ cm$^{-1}$. The X-ray emission showed normal nitrogen-to-oxygen abundance ratios and a high iron-to-oxygen abundance ratio. The spectrally deduced parameters, such as temperatures and column densities, are common to the diffuse X-ray emission near $\\eta$ Car. Thus, the diffuse X-ray emission in these two fields may have the same origin. The spectral fitting results are discussed to constrain the origin in the context of stellar winds and supernovae. ",
        "introduction": "\\label{sec:intro}  Diffuse X-ray emission extending over several to tens pc has been reported in many massive star-forming regions, such as NGC 2024 \\citep{Ezoe2006a}, the Orion Nebula \\citep{Guedel2008}, the Rosette Nebula, \\citep{Townsley2003}, M17 \\citep{Townsley2003,Hyodo2008}, RCW 38 \\citep{Wolk2002}, NGC 6334 \\citep{Ezoe2006b}, the Carina Nebula \\citep{Hamaguchi2007}, W49 \\citep{Tsujimoto2006}, NGC 3603 \\citep{Moffat2002}, the Arches Cluster \\citep{Yusef-Zadeh2002,Tsujimoto2007}, and the Quintuplet Cluster \\citep{Wang2002}. This diffuse component contributes a considerable fraction of the total X-ray emission and shows different spectral characteristics among these regions. Diffuse X-ray emission can be roughly classified into three types: thin-thermal plasma emission with a temperature $kT\\sim$ 0.1--1 keV, higher-temperature plasma emission with $kT\\sim$ 2--10 keV, and possibly non-thermal emission with a photon index of 1-1.5. These phenomena can be explained by plasma heating and particle acceleration in strong shocks by fast stellar winds from young OB stars \\citep{Townsley2003,Ezoe2006a,Ezoe2006b,Guedel2008} and/or past supernova remnants (SNRs) \\citep{Wolk2002,Hamaguchi2007}. The precise origin of diffuse X-ray emission, however, is often unclear.  Recently \\citet{Hamaguchi2007} suggest that the origin of diffuse X-ray emission can be constrained by plasma diagnostics or measurements of elemental abundances. While main-sequence late-O to early B stars have nearly solar abundances (e.g., \\cite{Cunha1994,Daflon2004}), evolved stars show non-solar elemental compositions due to the CNO cycle. For instance, the plasma will be overabundant in nitrogen if its origin is the wind from a nitrogen-rich Wolf-Rayet star. On the other hand, the plasma will be overabundant in oxygen, neon, and silicon if it is produced by a Type II SNR (e.g., \\cite{Tsujimoto1995}). The low-temperature ($kT\\sim$0.1--1 keV) type of diffuse X-ray emission is ideal for such diagnostic studies, because a variety of K-shell lines exist in the 0.2--2 keV range.  The Carina Nebula is an excellent site to investigate plasma diagnostics of diffuse X-ray emission. At a distance of 2.3 kpc, it is one of the most active massive star forming regions in the Galaxy. It contains eight massive stellar clusters: Trumpler 14, 15, 16, Collinder 228, Bochum 10, 11, NGC 3293, and NGC 3324. In total, there exist more than 64 O stars \\citep{Feinstein1995, Smith2006}, including the extreme-type luminous blue variable $\\eta$ Car and four Wolf-Rayet stars. The age of the nebula is estimated to be $\\sim3$ Myr based on the most evolved stars and the size of the HII region \\citep{Smith2000}. The young massive stars that are still enshrouded in gas and dust have been observed in optical, infrared, and radio wavelengths (e.g., \\cite{Harvey1979, Smith2000, Yonekura2005}). The number counts of the most massive O stars, e.g., O3 stars, suggest that the star-formation activity of the Carina Nebula rivals those of the most active regions, such as NGC 3603 at $D=6.9$ kpc and W49 at $D=11.4$ kpc. The proximity of the Carina Nebula, compared with NGC 3603 and W49, makes it the best target to study diffuse X-ray emission resulting from star-forming activities.  Seward et al.\\ (1979) first suggested the existence of possible diffuse soft X-ray emission in the Carina Nebula with a luminosity of $\\sim$10$^{35}$ erg s$^{-1}$ and a spatial extent of several pc, using {Einstein} observations. Although the limited energy and spatial resolution of Einstein hindered the determination of precise plasma parameters and the contribution from point sources, \\citet{Seward1982} postulated that the extended X-rays were from hot gas with $T\\sim10^{7}$ K. With {Chandra}, Evans et al.\\ (2003) confirmed the existence of diffuse X-ray emission near $\\eta$ Car in addition to point sources; however, the limited photon statistics and high background prevented detailed spectral analysis. Hamaguchi et al.\\ (2007) obtained spectra of the diffuse X-ray emission around $\\eta$ Car with {Suzaku}. Owing to the good low-energy response and the low background of the X-ray CCD onboard {Suzaku}, % the spectra extracted from the regions north and south of $\\eta$ Car can be modeled with high precision; both spectra are best represented by three-temperature plasma models with $kT\\sim0.2$, $\\sim0.6$, and $\\sim5$ keV. Analyzing the {Suzaku} data in conjunction with XMM-Newton and Chandra data, Hamaguchi et al.\\ (2007) concluded that the 0.2 keV and 0.6 keV components most likely originated from diffuse plasma, but the 5 keV component could not be easily distinguished from the unresolved point sources. They found that the iron and silicon abundances were significantly different in the north and south regions, and that the nitrogen-to-oxygen abundance ratios in both regions were far lower than those of stellar winds from evolved massive stars such as $\\eta$ Car and WR25 in this field. From these results, they concluded that the diffuse X-ray emission near $\\eta$ Car originated from one or multiple SNRs.  We have studied extended X-ray emission from an eastern tip region of the Carina Nebula. Located $\\sim$ 30$'$ ($\\sim$ 20 pc) from $\\eta$ Car, this region is less contaminated by X-ray emission from OB stars than the regions near $\\eta$ Car. Previous {Einstein} observations have revealed strong extended X-ray emission in this region, although no massive stars earlier than B3 are known here \\citep{Seward1982}. It is not known whether this emission is truly diffuse and, if so, whether its origin is similar to that of the regions near $\\eta$ Car. In this paper we report the first detailed spectral analysis of the extended X-ray emission in this eastern tip region of the Carina Nebula using Suzaku observation.  To augment the limited angular resolution of Suzaku, the analysis also made use of {XMM-Newton} data. ",
        "conclusions": "In the present paper, we have investigated the properties of diffuse X-ray emission associated with the eastern tip region of the Carina Nebula using the Suzaku and XMM-Newton data. Our conclusion is as follows.  (1) Strong extended X-ray emission was detected from the entire field-of-view of Suzaku with a 0.2--5 keV surface brightness of 0.7$\\sim$4$\\times10^{-14}$ erg s$^{-1}$ arcmin$^{-2}$. Comparisons with the estimated contamination from astrophysical backgrounds and point sources suggest that most of the emission is diffuse in nature.  (2) The observed absorption column density and temperature are consistent with those in the $\\eta$Car region, suggesting the same origin as the diffuse X-ray emission in the vicinity of the Carina Nebula.  (3) Estimated physical properties of the plasma such as pressure, total energy, and mass can be explained by stellar winds from OB stars in the Carina Nebula or young SNR(s) with the age less than $\\sim$10$^4$ yr. The SNR interpretation can provide the necessary energy and mass more easily.  (4) Absolute abundance values are strongly affected by abundances of metals fixed in spectral fits, allowing both the stellar-wind and SNR interpretations. The low nitrogen-to-oxygen and high iron-to-oxygen ratios derived from the spectral fits may support that the diffuse plasma heated up by stellar winds from main-sequence OB stars. The abundance ratios can be produced by a mixture of stellar winds and SNRs, as well.  The authors acknowledge discussion with Y. Hydo. K.H. is supported by the NASA Astrobiology Program under RTOP 344-53-51.  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.7\\textwidth,){fig1av1p.eps} \\FigureFile(0.7\\textwidth,){fig1bv1.eps} \\end{center} \\caption{(a) % An MSX band-A 8.28 $\\mu$m image of the Carina nebula taken from the NASA/IPAC Infrared Science Archieve. Yellow contours show the 12CO (J$=$2-1) map \\citep{Yonekura2005}. Two white boxes represent field-of-views of this (upper left) and previous Suzaku observations (lower right, \\cite{Hamaguchi2007}). (b) An EPIC MOS mosaic of the same nebula with XMM-Newton taken from the XMM-Newton Image Gallery. The red, green and blue colors show soft (0.4--0.7 keV), medium (0.7--1.3 keV) and hard (2--7 keV) X-ray emission, respectively. Circles and numbers indicate point sources detected in our analysis (see \\S \\ref{sec:ps}). }\\label{fig:overview} \\end{figure}  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.8\\textwidth,){fig2v2.eps} \\end{center} \\caption{Suzaku BI images of the eastern tip region of the Carina nebula in the (a) 0.2--2 keV and (b) 2--10 keV bands, displayed on the J2000.0 coordinates. For clarity, images are binned by a factor of 8 and smoothed by a Gaussian of $\\sigma$ = 3 pixels. % Vignettings are corrected (see \\S \\ref{sec:obs}). The unit of intensity in the greyscale wedge is arbitrary. Solid black lines mark regions utilized in the spectral analysis. The strong emission at the bottom-left and -right parts of the panel (b) corresponds to calibration 55Fe sources. }\\label{fig:xisimage} \\end{figure}  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.55\\textwidth,){fig3v2.eps} \\end{center} \\caption{Background-subtracted BI (black) and FI (red) spectra of (a) the blob, (b) the east, and (c) the nw regions. Center energies of emission lines are shown in the panel (a). For comparison, the vertical axis is normalized by the angular size of each region. }\\label{fig:xisdiffspec} \\end{figure}  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.8\\textwidth,){fig4v0.eps} \\end{center} \\caption{XMM-Newton EPIC-pn spectra of the bright point sources (No.1, 2, 4, and 7). The solid lines show the best-fit absorbed plasma models. The dotted lines in the panel (a) show two plasma components. The bottom panels exhibit residuals from the best-fit models. }\\label{fig:xmmsrcspec} \\end{figure}  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.55\\textwidth,){fig5v2.eps} \\end{center} \\caption{Background-subtracted BI spectra of (a) the blob, (b) the east, and (c) the nw regions. Solid lines show estimated contamination from X-ray sources. }\\label{fig:bgdest} \\end{figure}   \\begin{figure}[p] \\begin{center} \\FigureFile(0.55\\textwidth,){fig6v2.eps} \\end{center} \\caption{Best-fit spectral results of the diffuse X-ray emission in (a) the blob, (b) the east, and (c) the nw regions. The best-fit model is shown in a solid black line. The model components for the XIS1 spectrum are shown in solid colored lines (cyan: the LHB component, magenta and red: the two-temperature plasma component, blue: the power-law component). See tables \\ref{tbl:fit1} and \\ref{tbl:fit2} for the obtained parameters. }\\label{fig:fit12} \\end{figure}  \\clearpage   \\begin{figure}[p] \\begin{center} \\FigureFile(0.55\\textwidth,){fig7v0.eps} \\end{center} \\caption{XMM-Newton MOS fitting result of the diffuse X-ray emission in the blob region. Line colors are the same as figure \\ref{fig:fit12} but for the MOS1 and MOS2 data. See table \\ref{tbl:fit3} for the obtained parameters. }\\label{fig:fit3} \\end{figure}  \\bigskip   \\begin{figure}[p] \\hspace*{-10mm} \\FigureFile(1.1\\textwidth,){fig8v3.eps} \\bigskip \\caption{Abundance distributions in (a) the blob region in the eastern tip region (this paper), (b) the north and south regions in the $\\eta$ Car region \\citep{Hamaguchi2007}, and (c) sub-regions in M17 \\citep{Hyodo2008}. The filled marks represent the fixed values. }\\label{fig:abd} \\end{figure}  \\clearpage   \\begin{table}[p] \\caption{Properties of XMM-Newton point sources.}\\label{tbl:srclist} \\begin{center} \\begin{tabular}{lcccccccc}  \\hline No. & R.A.$^{\\rm a}$  & Decl.$^{\\rm a}$ & $C_{\\rm MOS}$$^{\\rm b}$ & $C_{\\rm pn}$$^{\\rm b}$ & Flux$^{\\rm c}$  & Var$^{\\rm d}$ & Counterpart$^{\\rm e}$ \\\\ \\hline  1   & 10:47:31.68  & -59:32:33.7 & 199$\\pm13$       & 411$\\pm30$ &   3.8     &  No & 2MASS J 10473148-5932345       \\\\  2\t  & 10:49:28.08  & -59:34:55.2 &  68$\\pm8$        & 124$\\pm$17 &   1.1     &  No & 2MASS J 10492780-5935000       \\\\  3$^{\\rm f}$\t  & 10:48:52.80  & -59:40:39.4 &  62$\\pm9$        &  22$\\pm13$ &   0.07    &  Yes & 2MASS J 10485248-5940392       \\\\  4\t  & 10:47:22.80  & -59:23:37.3 &   66$\\pm8$       &  174$\\pm$19 &    0.38   &  No & 2MASS J 10472233-5923392       \\\\  5\t  & 10:46:45.84  & -59:30:13.0  % &   30$\\pm7$       &  83$\\pm15$ &    0.27   &  No & \\begin{tabular}{c} 2MASS J 10464595-5930130       \\\\ GSC 0862602325\\\\ \\end{tabular}\\\\  6\t  & 10:47:05.28  & -59:30:27.0  % &    15$\\pm7$      &  61$\\pm17$ &   0.20    &  No & 2MASS J 10470505-5930247       \\\\  7$^{\\rm g}  $& 10:46:09.12& -59:43:08.0  % &    ---           &  336$\\pm28$ &   0.79   &  No & 2MASS J 10460901-5943025       \\\\  8  & 10:48:10.80  & -59:41:44.9  % &   47$\\pm10$      &  73$\\pm22$ &   0.23   &  No & 2MASS J 10480990-5941489       \\\\  9  & 10:48:30.96  & -59:42:14.8  % &   26$\\pm8$       &   78$\\pm19$ &   0.25   &  No & 2MASS J 10483064-5942144       \\\\  10  & 10:48:14.16  & -59:43:32.2 % &   15$\\pm9$      &  81$\\pm19$ &   0.26   &  No & ---      \\\\        \\hline \\end{tabular} \\end{center} {\\noindent $^{\\rm a}$ Source positions in J2000 coordinates.\\\\ $^{\\rm b}$ Background-subtracted photon counts detected with EPIC MOS and pn in 0.4--10 keV. $C_{\\rm MOS}$ is the average counts of MOS1 and MOS2. Errors are 1$\\sigma$.\\\\ $^{\\rm c}$ The 0.4--10 keV X-ray flux in 10$^{-13}$ erg s$^{-1}$ cm$^{-2}$ which corresponds to 6$\\times10^{31}$ erg s$^{-1}$ at 2.3 kpc. Fluxes of No.1, 2, 4, and 7 are derived from the spectral fittings, while the other are estimated from $C_{\\rm pn}$ assuming the thin-thermal plasma model (see text).\\\\ $^{\\rm d}$ Time variability based on the $\\chi^2$ statistics.\\\\ $^{\\rm e}$ Counterpart candidates within 10 arcsec of the X-ray position, based on searches of the 2MASS and AGAST catalogs.\\\\ $^{\\rm f}$ This source falls in the CCD gap of pn.\\\\ $^{\\rm g}$ This source is outside the FOV of MOS.\\\\ } \\end{table}   \\begin{table}[p] \\caption{Results of the spectral fits to the point sources$^{\\rm a}$.} \\label{tbl:srcfit}  \\begin{center} \\begin{tabular}{lccccc}  \\hline No. & $N_{\\rm H}$$^{\\rm b}$ & $kT$$^{\\rm c}$ & Normalization$^{\\rm d}$ & $L_{\\rm X}$$^{\\rm e}$ & $\\chi^2$/d.o.f. \\\\ \\hline  1 & 0.17$_{-0.12}^{+0.57}$ & 0.42$_{-0.18}^{+0.45}$, $>32$ & 4.8$_{-2.9}^{+4.1}$$\\times10^{-5}$, 2.5$_{-0.6}^{+0.5}$$\\times10^{-4}$ & 3 & 19.5/25 \\\\  2 & 9.6$_{-3.8}^{+6.3}$ & 1.8$_{-0.8}^{+2.3}$ & 1.0$\\pm{0.1}$$\\times10^{-3}$ & 6 & 3.5/7 \\\\  4 & $<0.95$ & 0.70$_{-0.44}^{+0.10}$ & 4.4$\\pm0.7$$\\times10^{-5}$ & 0.3 & 7.1/9 \\\\   7 & $0.72_{-0.16}^{+0.17}$ & $0.12_{-0.03}^{+0.04}$ & 8.4$_{-4.5}^{+4.6}$$\\times10^{-2}$ & 80 & 37.6/23 \\\\  \\hline \\end{tabular} \\end{center} {\\noindent  $^{\\rm a}$ A single plasma model is assumed for No. 2, 4, and 7, while a two temperature model is used for No.1. A metal abundance is fixed at 0.3 solar value. Errors refer to the 90\\% confidence range. \\\\ $^{\\rm b}$ Hydrogen column density in 10$^{22}$ cm$^{-2}$.\\\\ $^{\\rm c}$ Plasma temperature in keV.\\\\ $^{\\rm d}$ Normalization factor of the APEC model, representing 10$^{-14}$/4$\\pi D^2$ $EM$, where $D$ is a distance to the Carina nebula and $EM$ is an emission measure in cm$^{-3}$.\\\\ $^{\\rm e}$ Absorption-corrected 0.4--10 keV luminosity in 10$^{32}$ erg s$^{-1}$ assuming a distance of 2.3 kpc.\\\\  } \\end{table}   \\begin{table}[p] \\caption{Results of the two-temperature plasma model fit to the diffuse X-ray emission in the blob region.} \\label{tbl:fit1}  \\begin{center} \\begin{tabular}{lccccc}  \\hline\\hline Model$^{\\rm a}$                               &  1           &   2          & 3            & 4            & Typical error$^{\\rm b}$ \\\\ \\hline  Two-temperature plasma component$^{\\rm c}$\\\\  $N_{\\rm H}$ (10$^{22}$ cm$^{-2}$)              & 0.23         & 0.25         & 0.22         & 0.26         & 0.01\\\\ % $kT_1$  (keV)                                & 0.25         & 0.24         & 0.25         & 0.24         & 0.01  \\\\ % $kT_2$  (keV)                                & 0.55         & 0.56         & 0.55         & 0.54         & 0.01  \\\\ % C       (solar)                              & 0.3 (fixed)  & 1.0 (fixed)  & 0.0 (fixed)  & 0.0 (fixed)  & $-$   \\\\ N       (solar)                              & 0.3 (fixed)  & 1.0 (fixed)  & 0.0 (fixed)  & 0.0 (fixed)  & $-$   \\\\ O       (solar)                              & 0.24         & 0.55         & 0.16         & 0.10         & 0.01  \\\\ % Ne      (solar)                              & 0.46         & 0.93         & 0.36         & 0.21         & 0.01  \\\\ % Mg      (solar)                              & 0.44         & 0.94         & 0.32         & 0.24         & 0.01  \\\\ % Al      (solar)                              & 0.3 (fixed)  & 1.0   (fixed)& 0.075 (fixed)& 0.029 (fixed)& $-$   \\\\ Si      (solar)                              & 0.54         & 1.1          & 0.40         & 0.37         & 0.02  \\\\ % S       (solar)                              & 0.74         & 1.5          & 0.56         & 0.57         & 0.1   \\\\ % Ar      (solar)                              & 0.3 (fixed)  & 1.0 (fixed)  & 0.0 (fixed)  & 0.13 (fixed) & $-$   \\\\ Ca      (solar)                              & 0.3 (fixed)  & 1.0 (fixed)  & 0.0 (fixed)  & 0.0 (fixed)  & $-$   \\\\ Fe      (solar)                              & 0.32         & 0.71         & 0.23         & 0.19         & 0.01  \\\\ % Ni      (solar)                              & 0.3 (fixed)  & 1.0 (fixed)  & 0.089 (fixed)& 0.78 (fixed) & $-$   \\\\ log$EM_1$ (cm$^{-3}$ arcmin$^{-2}$)           & 54.9         & 54.6         & 55.0         & 55.4         & 0.02 \\\\ % log$EM_2$ (cm$^{-3}$ arcmin$^{-2}$)           & 54.5         & 54.2         & 54.6         & 54.7         & 0.02 \\\\ % Flux1 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 1.8          & 1.8          & 1.8          & 2.3          & 0.1 \\\\ % Flux2 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 2.2          & 2.2          & 2.2          & 1.8          & 0.1 \\\\ % \\hline  Power-law component$^{\\rm d}$        \\\\ Flux (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$) & 0.56         & 0.57         & 0.56         & 0.56         & 0.02  \\\\ % \\hline  LHB component$^{\\rm e}$              \\\\ Flux (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$) & 0.10         &  0.11        &  0.10        &  0.10        & 0.02  \\\\ % \\hline  $\\chi^2$/d.o.f.                              &  1.24        & 1.31         & 1.21         & 1.22         &       \\\\ d.o.f.                              &  1231        & 1231         & 1231         & 1231         &       \\\\ \\hline   \\end{tabular}  \\end{center} {\\noindent  $^{\\rm a}$ Fitting models with different fixed abundances.  \\\\ $^{\\rm b}$ Typical fitting errors at the 90\\% confidence level.\\\\ $^{\\rm c}$ A commonly-absorbed plasma model. Arabic numbers 1 and 2 denote the two temperature components. Parameter definitions are the same as those in table \\ref{tbl:srcfit}. Fluxes are calculated in 0.2--5 keV.\\\\ $^{\\rm d}$ A power-law model representing CXB, GRXE, and point sources. A photon index is fixed at 1.5. The same absorption for the two-temperature plasma is assumed. Normalization is photon keV$^{-1}$ cm$^{-2}$ at 1 keV.\\\\ $^{\\rm e}$ A single-temperature plasma model representing LHB. A plasma temperature is fixed at 0.1 keV.\\\\  } \\end{table}   \\begin{table}[p] \\caption{Results of the two-temperature plasma model fit to the diffuse X-ray emission in the east and nw regions$^{\\rm a}$.} \\label{tbl:fit2}  \\begin{center} \\begin{tabular}{lccccc}  \\hline\\hline Region                                         &  east                  &   nw                      \\\\ \\hline  Two-temperature plasma component\\\\  $N_{\\rm H}$ (10$^{22}$ cm$^{-2}$)              & 0.21$_{-0.07}^{+0.09}$ & 0.32$_{-0.07}^{+0.11}$     \\\\ % $kT_1$  (keV)                                  & 0.20$_{-0.02}^{+0.04}$ & 0.19$_{-0.03}^{+0.02}$     \\\\ % $kT_2$  (keV)                                  & 0.54$_{-0.07}^{+0.04}$ & 0.41$_{-0.08}^{+0.10}$     \\\\ % C       (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ N       (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ O       (solar)                                & 0.15$_{-0.06}^{+0.12}$ & 0.07$_{-0.02}^{+0.04}$     \\\\ % Ne      (solar)                                & 0.33$_{-0.14}^{+0.31}$ & 0.27$_{-0.05}^{+0.14}$     \\\\ % Mg      (solar)                                & 0.30$_{-0.14}^{+0.29}$ & 0.25$_{-0.10}^{+0.16}$     \\\\ % Al      (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ Si      (solar)                                & 0.43$_{-0.20}^{+0.37}$ & 0.96$_{-0.43}^{+0.57}$     \\\\ % S       (solar)                                & 0.74 (fixed)           & 0.74 (fixed)               \\\\ % Ar      (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ Ca      (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ Fe      (solar)                                & 0.16$_{-0.05}^{+0.10}$ & 0.17$_{-0.06}^{+0.03}$     \\\\ % Ni      (solar)                                & 0.3 (fixed)            & 0.3 (fixed)                \\\\ log$EM_1$ (cm$^{-3}$ arcmin$^{-2}$)            & 54.3$_{-0.5}^{+0.4}$   & 55.0$\\pm{0.7}$       \\\\ % log$EM_2$ (cm$^{-3}$ arcmin$^{-2}$)            & 54.0$\\pm{0.2}$         & 54.3$_{-0.3}^{+0.8}$       \\\\ % Flux1 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 0.23$_{-0.15}^{+0.41}$ & 0.41$\\pm{0.08}$     \\\\ % Flux2 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 0.48$_{-0.20}^{+0.31}$ & 0.45$_{-0.20}^{+0.11}$     \\\\ % \\hline  Power-law component        \\\\ Flux  (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 0.56$\\pm{0.06}$          & 0.55$\\pm{0.09}$          \\\\ % \\hline  LHB component              \\\\ Flux  (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 0.087$_{-0.032}^{+0.031}$& 0.15$_{-0.06}^{+0.05}$  \\\\ % \\hline  $\\chi^2$/d.o.f.                                & 0.77                     & 0.60                    \\\\ d.o.f.                                & 245                      & 111                     \\\\ \\hline \\end{tabular}  \\end{center} {\\noindent  $^{\\rm a}$ Notations and symbols are the same as table \\ref{tbl:fit1}.  } \\end{table}  \\clearpage  \\begin{table}[p] \\caption{Result of the two-temperature plasma model fit to the XMM MOS spectra of the blob region$^{\\rm a}$.} \\label{tbl:fit3}  \\begin{center} \\begin{tabular}{lccccc}  \\hline\\hline Model                                         &  1                   \\\\ \\hline  Two-temperature plasma component\\\\  $N_{\\rm H}$ (10$^{22}$ cm$^{-2}$)             & 0.19$_{-0.02}^{+0.03}$  \\\\ % $kT_1$  (keV)                                                    & 0.24$\\pm{0.01}$              \\\\ % $kT_2$  (keV)                                                    & 0.58$\\pm{0.01}$              \\\\ %  C       (solar)                                                        & 1.2$_{-0.9}^{+0.4}$            \\\\ N       (solar)                                                        & 0.3 (fixed)                          \\\\ O       (solar)                                  & 0.23$_{-0.03}^{+0.02}$      \\\\ % Ne      (solar)                                 & 0.44$\\pm{0.07}$                  \\\\ % Mg      (solar)                                & 0.46$\\pm{0.07}$                   \\\\ % Al      (solar)                                  & 0.3 (fixed)                               \\\\ Si      (solar)                                  & 0.48$_{-0.08}^{+0.07}$      \\\\ % S       (solar)                                  & 0.48$_{-0.18}^{+0.20}$      \\\\ % Ar      (solar)                                  & 0.3 (fixed)                              \\\\ Ca      (solar)                                & 0.3 (fixed)                               \\\\ Fe      (solar)                                 & 0.32$_{-0.04}^{+0.05}$        \\\\ % Ni      (solar)                                 & 0.3 (fixed)                                 \\\\ log$EM_1$ (cm$^{-3}$ arcmin$^{-2}$)               & 54.8$\\pm{0.1}$                    \\\\ % log$EM_2$ (cm$^{-3}$ arcmin$^{-2}$)               & 54.5$\\pm{0.1}$                    \\\\ % Flux1 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 1.7$\\pm{0.2}$               \\\\ % Flux2 (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & 2.7$_{-0.4}^{+0.2}$     \\\\ % \\hline  Power-law component        \\\\ Flux  (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)   & 0.84$\\pm{0.05}$                    \\\\ % \\hline  LHB component              \\\\ Flux  (10$^{-14}$ erg s$^{-1}$ arcmin$^{-2}$)  & $<0.06$  \\\\ % \\hline  $\\chi^2$/d.o.f.                                & 1.20                       \\\\ d.o.f.                                       & 337                        \\\\ \\hline \\end{tabular}  \\end{center} {\\noindent  $^{\\rm a}$ Notations and symbols are the same as table \\ref{tbl:fit1}.  } \\end{table}  \\clearpage   \\begin{table}[p] \\caption{Physical properties of the diffuse plasma in the blob region$^a$.} \\label{tbl:plasma}  \\begin{center} \\begin{tabular}{lccccc}  \\hline\\hline Parameter                                    &  Scale Factor                &       $T_1$                     &      $T_2$\\\\ \\hline  \\multicolumn{4}{c}{Observed X-ray Properties}\\\\ \\hline  $kT$ (keV)                                    &   $-$                               &        0.3                          &      0.6 \\\\ $L_{\\rm X}$ (ergs s$^{-1}$)       &   $-$                               &   2$\\times10^{34}$     &     1$\\times10^{34}$\\\\ $V$ (cm$^{3}$)                            &   $\\eta$                          &   1$\\times10^{57}$     &     1$\\times10^{57}$\\\\ \\hline  \\multicolumn{4}{c}{Estimated X-ray Plasma Properties}\\\\ \\hline   $n_{\\rm e}$ (cm$^{-3}$)             &   $\\eta^{-1/2}$            &        0.3                          &      0.4 \\\\ $P/k$ (K cm$^{-3}$)                    &   $\\eta^{-1/2}$            &   2$\\times10^{6}$       &     5$\\times10^{6}$\\\\ $U$ (ergs)                                     &   $\\eta^{1/2}$             &   4$\\times10^{47}$     &     1$\\times10^{48}$\\\\ $t_{\\rm cool}$ (yr)                         &  $\\eta^{1/2}$             &   1$\\times10^{6}$       &     4$\\times10^{6}$\\\\ $M_{\\rm plasma}$ ($M_\\odot$) &  $\\eta^{1/2}$             &   0.2                               &     0.2 \\\\  \\hline  \\end{tabular}  \\end{center} {\\noindent  $^{\\rm a}$ $\\eta$ is a filling factor for the volume of the plasma. $T_1$ and $T_2$ indicate the two temperature plasma component in table \\ref{tbl:fit1} model 1.  } \\end{table}  \\clearpage"
    },
    "0809/0809.3947_arXiv.txt": {
        "abstract": "{} {We present the first direct comparison of the distribution of the gas, as traced by the [\\ion{O}{i}] 6300\\AA\\ emission, and the dust, as traced by the 10 $\\mu$m emission, in the planet-forming region of proto-planetary disks around three intermediate-mass stars: HD 101412, HD 135344 B and HD 179218.} {$N$-band visibilities were obtained with VLTI/MIDI. Simple geometrical models are used to compare the dust emission to high-resolution optical spectra in the 6300\\,\\AA\\ [\\ion{O}{i}] line of the same targets.} {HD 101412 and HD 135344 B show compact ($<$ 2 AU) 10 $\\mu$m emission while the [\\ion{O}{i}] brightness profile shows a double peaked structure. The inner peak is strongest and is consistent with the location of the dust, the outer peak is fainter and is located at 5-10 AU. In both systems, spatially extended PAH emission is found. HD 179218 shows a double ring-like 10 $\\mu$m emission with the first ring peaking at $\\sim$ 1 AU and the second at $\\sim$ 20 AU. The [\\ion{O}{i}] emitting region is more compact, peaking between 3 -- 6 AU.} {The disks around HD 101412 and HD 135344 B appear strongly flared in the gas, but self-shadowed in the dust beyond $\\sim$ 2 AU. The difference in the gas and dust vertical structure beyond 2 AU might be the first observational evidence of gas-dust decoupling in protoplanetary disks. The disk around HD 179218 is flared in the dust. The 10 $\\mu$m emission emerges from the inner rim and from the flared surface of the disk at larger radii. No dust emission is detected between $\\sim$ 3 -- 15 AU. The oxygen emission seems also to come from a flared structure, however, the bulk of this emission is produced between $\\sim$ 1 -- 10 AU. This could indicate a lack of gas in the outer disk or could be due to chemical effects which reduce the abundance of OH -- the parent molecule of the observed [\\ion{O}{i}] emission -- further away from the star. It may also be a contrast effect if the [\\ion{O}{i}] emission is much stronger in the inner disk. We suggest that the three systems, HD 179218, HD 135344 B and HD 101412, may form an evolutionary sequence: the disk initially flared becomes flat under the combined action of gas-dust decoupling, grain growth and dust settling.} ",
        "introduction": "Many pre-main-sequence stars are characterized by excess infrared emission above the stellar photospheric level which, depending on the evolutionary state of the system, may start in the near infrared (1 -- 5~$\\mu$m) or at longer wavelengths. The dust distributed around the young star is responsible for this emission. The dust particles absorb a large fraction of the short wavelength stellar photons and re-emit them at infrared wavelengths. This dust is believed to be confined to a disk-like structure which forms in the early phase of star formation as a result of the conservation of the angular momentum of the parental cloud. Disk formation is followed by a longer phase of disk accretion during which disk material accretes onto the young star at a typical accretion rate of $10^{-7} - 10^{-10}$ M$_{\\sun}$ yr$^{-1}$. The interstellar dust and gas that forms the disk undergoes changes in its composition and size. Infrared surveys show that after a mean age of 3 Myr the inner part of the disk is cleared of dust. Viscous accretion, photo-evaporation and planet formation are the likely mechanisms responsible for this phenomenon (see review in Henning \\cite{henning}).  \\noindent  \\begin{table*} \\caption{Properties of the programme stars. The column \"Meeus group\" gives the classification in {\\it flared} (group I) and {\\it self-shadowed} (group II) disks by Meeus et al. (\\cite{meeus}). Column ``[\\ion{O}{i}] extent'' reports the extent (minimum and maximum radius) of the [\\ion{O}{i}] emitting region as derived in paper I. The disk position angle (PA) and inclination are taken from paper I.} \\label{tab:prop} \\centering \\begin{tabular}{lllllllllll} \\hline\\hline Star        & RA      & DEC     & Sp.T & M$_{star}$   & F$_{12\\mu m}$ & Distance & Meeus Group & PA & Inclination & [\\ion{O}{i}] extent\\\\ & (J2000) & (J2000) &      & [M$_{\\sun}$] & [Jy]        & [pc]     &             & [\\degr] & [\\degr]          & [AU]      \\\\ \\hline HD 101412   & 11:39:44.46 & -60:10:27.7 & A0IIIe & 2.3 & 3.22       & 160      & II        &  --     & 30               & 0.15 - 10 \\\\ HD 135344 B & 15:15:48.44 & -37:09:16.0 & F4Ve   & 1.7 & 1.59       & 140      & I         & 100     & 45               & 0.1 - 100 \\\\ HD 179218   & 19:11:11.25 & +15:47:15.6 & B9e    & 2.7 & 23.4       & 240      & I         & 10      & 40               & 0.4 - 50  \\\\ \\hline\\hline \\end{tabular} \\end{table*}  A circumstellar disk is believed to be the locus where planet formation takes place. The large number of recently discovered exo-planetary systems and the variety of such systems raised many new questions about the structure and evolution of {\\it protoplanetary} disks. Of particular interest for the understanding of disk evolution and planet formation is the coupling of gas and dust in disks. A main assumption underlying essentially all proto-planetary disk models is that dust and gas are thermally coupled. It is an open question whether this assumption holds on the disk surface and efforts have been made to improve disk models by taking into account the dust-gas decoupling (e.g. Kamp \\& Dullemond \\cite{kd}). While gas and dust are thermally coupled in the disk interior, in the low density environments of the disk surface layer, the two components may decouple.  \\smallskip \\noindent  The detailed structure of the disk surface temperature in the presence of gas-dust decoupling was studied by Jonkheid et al. (\\cite{jonkheid}), Kamp \\& Dullemond (\\cite{kd}) and Nomura \\& Millar (\\cite{nomura}). Different heating/cooling processes act at different heights in the disk. In particular, very high in the atmosphere, with particle densities as low as {\\it n} $< 10^5 {\\rm cm}^{-3}$ (A$_V$ $\\lesssim$ 10$^{-3}$ mag), the gas temperature is set by the balance of photoelectric heating and fine structure line cooling of neutral oxygen (Jonkheid et al. \\cite{jonkheid}, Kamp \\& Dullemond \\cite{kd}). In the upper layers of protoplanetary disks the gas temperature may exceed the dust temperature. At small radii ($<$ 50 AU) the temperature of the gas above the disk photosphere may reach $\\sim 10^4$ K. At larger radii ($>$ 50 AU) the gas can become as hot as a few hundred Kelvin (Kamp \\& Dullemond \\cite{kd}). However, it is not obvious that such a hot and tenuous disk atmosphere can remain as a stable structure of a disk. The difference in gas and dust temperature as well as other processes (e.g. disk wind, disk evaporation, dust coagulation and settling) may lead to a physical decoupling of gas and dust in disks.  \\smallskip \\noindent  In this paper we present the first direct comparison of the dust and gas emission of three pre-main-sequence stars: \\object{HD 101412}, \\object{HD 135344 B} and \\object{HD 179218}. These are intermediate-mass (1.7 - 2.7 M$_{\\sun}$) stars belonging to both group I (flaring disk: HD 135344 B and HD 179218) and group II (flattened disk: HD 101412) according to the classification of Meeus et al. (\\cite{meeus}). In a previous paper (van der Plas et al. \\cite{plas}, hereafter paper I) we presented high resolution spectroscopy ($\\frac{\\lambda}{\\Delta\\lambda} = 77000$) of the optical [\\ion{O}{i}] 6300.304\\AA\\ line with VLT/UVES of these three stars. Here we present $N$-band interferometric observations obtained with VLTI/MIDI aimed at spatially resolving the mid-infrared emitting region of the disk. We will also compare the thermal coupling between dust and gas in the disks surroundings our three targets by comparing the MIDI observations to a simple phenomenological model derived from the [\\ion{O}{i}] data. The properties of the programme stars are reported in Table \\ref{tab:prop}. The paper is organized as follow: section 2 briefly summarizes the results of paper I; section 3 contains information about observations and data reduction; in section 4 and 5 we report respectively an analysis of the interferometric measurements and the comparison with the optical data of paper I. The discussion and conclusion are given in section 6. ",
        "conclusions": "In this paper we presented the first direct comparison of gas and dust emission from the surface of three protoplanetary disks. The comparison of high optical spectroscopy with infrared interferometric observations gives some insight into the relative size and shape of the gas and the dust emitting region.  \\smallskip \\noindent  For HD 179218 the new MIDI results presented here and the UVES results of paper I are in good agreement with a flared disk structure. The 10 $\\mu$m emission comes from two separate regions: an inner part located between $\\sim$ 0.3 -- 3 AU (likely the disk inner rim) and an outer part between $\\sim$ 15 -- 23 AU (flared disk surface). No dust emission is detected between 3 -- 15 AU. The [\\ion{O}{i}] brightness profile shows instead a single strong peak centered at $\\sim$ 3 -- 6 AU. The difference between gas and dust emitting regions for HD 179218 may reflects a real different structure for the two components or might be caused by other effects such as: 1) the action of chemical processes that reduce the abundance of OH in the outer part of the disk (in this case plenty of gas may still exist but we are not able to see it); 2) a contrast effect of the [\\ion{O}{i}] emission. In the latter case the [\\ion{O}{i}] emission is much stronger in the inner disk and it may outshine the [\\ion{O}{i}] from the outer disk. The presence of a dust gap needs further investigation as it may indicate possible ongoing planet formation.  \\smallskip \\noindent  HD 135344 B and HD101412 show a compact 10 $\\mu$m emitting region. The oxygen instead has a double peak brightness radial profile with the first peak coincident with the dust emission and a second peak further away from the star at $\\sim$ 5 -- 10 AU. In the case of HD 135344 B extended dust emission is found at 20.5 $\\mu$m (Doucet et al. \\cite{doucet}). Modelling the SED of HD 135344 B, Brown et al. (\\cite{brown}) propose a ``cold-disk'' geometry where the disk is mostly void of dust between 0.4 -- 45 AU. Dust is present within the inner 0.4 AU and outside 45 AU. This is in agreement with our result in the sense that we detect a compact mid-infrared dust emission; the dust at larger radii is cold and radiates at wavelengths longer than those traced by our observations. We suggest that the second [\\ion{O}{i}] peak arises either from a dust gap or from a dust-free layer above the shadow of the inner rim.  Following the phenomenological classification of Meeus et al. (\\cite{meeus}), HD 101412 is a group II source, i.e. the disk is flat and self-shadowed in the dust. At the dust sublimation radius the disk is puffed-up and shadows the outer disk, which is not reached by the stellar UV radiation. The second peak of the [\\ion{O}{i}] emission (Fig. 8) is not expected for such a self-shadowed geometry. We suggest that this is the result of a flat disk geometry for the dust and a flared geometry for the gas. This might be the signature of a different scale height, and hence vertical structure, between gas and dust beyond the inner rim. This means that -- in this object -- gas and dust are physically decoupled in the surface layers of the disk allowing the gas to emerge from the shadow of the inner rim. If this is confirmed, it may have important consequences for the disk evolution: once gas and dust decouple the dust grains settle faster towards the mid-plane. The dust settling favors the physical separation between gas and dust until all the dust settles towards the mid-plane (Dullemond \\& Dominik \\cite{dullemond04}) eventually leaving a region of gas in the upper layers of the disk. We thus suggest that HD 179218, HD 135344 B and HD 101412 form an evolutionary sequence where the disk, initially flared, becomes flat under the combined action of gas-dust decoupling, grain growth and dust settling. Interestingly, Acke et al. (\\cite{acke04}) already found that cold grains in the mid-plane of the disk have grown to considerably larger sizes in Meeus et al., group II sources than in group I sources, suggesting that disks may evolve from a flared to a self-shadowed geometry. Similarly, Bouwman et al. (\\cite{bouwman}) find a correlation between the strength of the amorphous silicate feature and the shape of the SED, consistent with the settling of dust as a consequence of grain growth. What we add to this picture is a completely independent piece of evidence supporting this and perhaps the identification of the driving mechanism of this process.  \\begin{appendix}"
    },
    "0809/0809.5200_arXiv.txt": {
        "abstract": "We report on our attempt for the first non-LTE modeling of gaseous metal disks around single DAZ white dwarfs recently discovered by G\\\"ansicke \\etal and thought to originate from a disrupted asteroid. We assume a Keplerian rotating viscous disk ring composed of calcium and hydrogen and compute the detailed vertical structure and emergent spectrum. We find that the observed infrared \\ion{Ca}{ii} emission triplet can be modeled with a hydrogen-deficient gas ring located at $R$\\,=\\,1.2\\,R$_\\odot$, inside of the tidal disruption radius, with \\Teff$\\approx$\\,6000\\,K and a low surface mass density of $\\approx$\\,0.3\\,g/cm$^{2}$. A disk having this density and reaching from the central white dwarf out to $R=1.2$\\,R$_\\odot$ would have a total mass of $7\\cdot 10^{21}$\\,g, corresponding to an asteroid with $\\approx$160\\,km diameter. ",
        "introduction": "More than two decades ago Zuckerman \\& Becklin (1987) announced the discovery of an IR excess around the DAZ white dwarf G29$-$38. The white dwarf itself is enriched in metals. Considering the short sedimentation timescales in the photosphere this implies that the star is accreting matter at a relatively high rate (Koester \\etal 1997). Since no cool companion has been found at G29$-$38, the hypothesis was put forward that a dust cloud around the white dwarf causes the IR excess. In fact, the presence of dust has been confirmed by \\emph{Spitzer} observations (Reach \\etal 2005). Graham \\etal (1990) concluded that the dust is located in the equatorial plane. Subsequently, further DAZ white dwarfs with potential dust disks were found (Becklin \\etal 2005, Kilic \\etal 2005, 2006). As a possible origin of these disks tidally disrupted comets were discussed (Debes \\& Sigurdsson 2002) and, more likely because of the absence of H and He, disrupted asteroids (Jura 2003). ",
        "conclusions": "The infrared \\ion{Ca}{ii} emission triplet in the spectrum of the DAZ white dwarf \\sdss\\ can be modeled with a geometrically and optically thin, Keplerian  viscous gas disk ring at a distance of 1.2~$\\rsun$ from the WD, with \\Teff\\,$\\approx$\\,6000\\,K and a low surface mass density $\\Sigma$\\,$\\approx$\\,0.3\\,g/cm$^{2}$. One serious open problem is the unknown disk-heating mechanism. The disk is hydrogen-deficient (H\\,$\\leq$\\,1\\% by mass) and it is located within the tidal disruption radius ($R_{\\rm tidal}$\\,=1.5\\,\\rsun). If one assumes that the disk reaches down to the WD and that it has uniformly this surface density, then its total mass would be 7$\\cdot\\,10^{21}$\\,g. A rocky asteroid ($\\bar{\\rho}$\\,=\\,3\\,g/cm$^{3}$) with this mass would have a diameter of about 160\\,km. An asteroid of this size in our solar system is, e.g., 22~Kalliope.  Future work will include other abundant metal species, with a composition appropriate for asteroids. However, at the moment there are only two \\ion{Fe}{ii} lines detected. Real progress for an abundance analysis of the gas disk around \\sdss\\ is only possible with future UV spectroscopy.  While the determination of the composition of accreted material from abundance analyses in DAZ atmospheres is an indirect method that depends on our theoretical knowledge about metal settling timescales, the analysis of the gas disks has the obvious advantage that is a direct composition measurement.  \\ack We thank Boris G\\\"ansicke for sending us his \\sdss\\ spectrum in electronic form and for useful discussions. T.R. is supported by the \\emph{German Astrophysical Virtual Observatory} project of the Federal Ministry of Education and Research (grant 05\\,AC6VTB)."
    },
    "0809/0809.1753_arXiv.txt": {
        "abstract": "We present VLT and Magellan spectroscopy and NTT photometry of nine faint cataclysmic variables (CVs) which were spectroscopically identified by the Sloan Digital Sky Survey. We measure orbital periods for five of these from the velocity variations of the cores and wings of their H$\\alpha$ emission lines. Four of the five have orbital periods shorter than the 2--3\\,hour period gap observed in the known population of CVs. SDSS J004335.14$-$003729.8 has an orbital period of $\\Porb = 82.325 \\pm 0.088$\\,min; Doppler maps show emission from the accretion disc, bright spot and the irradiated inner face of the secondary star. In its light curve we find a periodicity which may be attributable to pulsations of the white dwarf. SDSS J163722.21$-$001957.1 has $\\Porb = 99.75 \\pm 0.86$\\,min. By combining this new measurement with a published superhump period we estimate a mass ratio of $q \\approx 0.16$ and infer the physical properties and orbital inclination of the system. For SDSS J164248.52$+$134751.4 we find $\\Porb = 113.60 \\pm 1.5$\\,min. The Doppler map of this CV shows an unusual brightness distribution in the accretion disc which would benefit from further observations. SDSS J165837.70$+$184727.4 had spectroscopic characteristics which were very different between the SDSS spectrum and our own VLT observations, despite only a small change in brightness. We measure $\\Porb = 98.012 \\pm 0.065$\\,min from its narrow H$\\alpha$ emission line. Finally, SDSS J223843.84$+$010820.7 has a comparatively longer period of $\\Porb = 194.30 \\pm 0.16$\\,min. It contains a magnetic white dwarf and, with $g = 18.15$, is brighter than the other objects studied here. These results continue the trend for the fainter CVs identified by the SDSS to be almost exclusively shorter-period objects with low mass transfer rates. ",
        "introduction": "\\label{sec:intro}  Cataclysmic variables (CVs) are interacting binary stars containing a white dwarf primary component in a close orbit with a low-mass secondary star which fills its Roche lobe. In most of these systems the secondary component is hydrogen-rich and transfers material to the white dwarf via an accretion disc. Comprehensive reviews of the properties of CVs have been given by \\citet{Warner95book} and \\citet{Hellier01book}.  The evolution of CVs is thought to be governed primarily by the loss of orbital angular momentum due to gravitational radiation \\citep{Paczynski67aca} and magnetic braking \\citep{VerbuntZwaan81aa,Rappaport++82apj}. These effects \\reff{are predicted to} cause CVs to evolve towards shorter orbital periods until a minimum value of about 80\\,min, at which point the secondary stars become degenerate and the CVs evolve back to longer periods \\reff{(e.g.\\ \\citealt{Patterson98pasp})}.  The distribution of orbital periods of the observed population of CVs has a characteristic `period gap' \\citep{WhyteEggleton80mn,Knigge06mn} in the interval between 2.2 and 3.2 hours. The deficiency in the number of systems in this period range is thought to result from the sudden cessation of magnetic braking due to structural changes in CV secondary stars \\citep{SpruitRitter83aa}. The number of CVs shortwards of this period gap are observed to be roughly equal to the number which are longward of the gap \\citep[e.g.][]{Downes+01pasp,RitterKolb03aa}.  Unfortunately, theoretical studies of the population of CVs have consistently predicted that the vast majority of these objects should have short periods ($\\Porb \\la 2$\\,hours) due to their relatively longer evolutionary timescale \\citep{Dekool92aa, DekoolRitter93aa, Kolb93aa, KolbDekool93aa, Politano96apj, Politano04apj, KolbBaraffe99mn, Howell++01apj, Willems+05apj}, culminating in a strong `spike' in the population at a minimum period of about 65\\,min. The remarkable differences between the predicted and observed populations of CVs have not yet been satisfactorily explained, although it is clear that observational selection biases have a lot to answer for.  To understand these selection biases, and to discover what the properties of the intrinsic population of CVs are, we are conducting a research program to characterise the sample of CVs identified by the Sloan Digital Sky Survey (SDSS\\footnote{\\tt http://www.sdss.org/}; \\citealt{York+00aj}). A total of 212 of these objects have been found spectroscopically from an initial selection based on single-epoch photometric colour indices \\citep{Szkody+02aj, Szkody+03aj, Szkody+04aj, Szkody+05aj, Szkody+06aj, Szkody+07aj}. This sample is therefore not biased towards CVs which are variable or strong X-ray emitters, and also has a much wider coverage of colour space than previous large-scale surveys \\citep{Green++86apjs,Chen+01mn,Aungwerojwit+06aa}. Results and further discussion of our project to measure the orbital periods of SDSS CVs can be found in \\citet{Gansicke+06mn}, \\citet[][hereafter Paper\\,I]{Me+06mn}, \\citet{Me+07mn,Me+07mn2,Me++08mn}, \\citet{Dillon+08mn,Dillon+08}, and \\citet{Littlefair+06mn, Littlefair+06sci, Littlefair+07mn, Littlefair+08}.  In this work we present time-resolved spectroscopy and photometry of nine CVs, and measure orbital period for five of these. We shall abbreviate the names of the targets to SDSS\\,J0043, SDSS\\,J0337, SDSS\\,J1601, SDSS\\,J1637, SDSS\\,J1642, SDSS\\,J1658, SDSS\\,J1659, SDSS\\,J2232 and SDSS\\,J2238. Their full names and $ugriz$ apparent magnitudes are given in Table\\,\\ref{tab:iddata}. In Fig.\\,\\ref{fig:sdssspec} we have plotted their SDSS spectra for reference.  \\begin{table*} \\begin{center} \\caption{\\label{tab:iddata} Apparent magnitudes of our targets in the SDSS $ugriz$ passbands. $g_{\\rm spec}$ are apparent magnitudes we have calculated by convolving the SDSS flux-calibrated spectra with the $g$ passband function. They are obtained at a different epoch to the $ugriz$ magnitudes measured from the imaging observations, but are less reliable as they are affected by `slit losses', and any errors in astrometry or positioning of the spectroscopic fibre entrance.} \\begin{tabular}{lllccccccc} \\hline SDSS name                  & Short name  & Reference           & $u$ & $g$ & $r$ & $i$ & $z$ & $g_{\\rm spec}$  \\\\ \\hline SDSS J004335.14$-$003729.8 & SDSS\\,J0043 & \\citet{Szkody+04aj} & 20.18 & 19.86 & 19.81 & 19.97 & 19.84 & 19.95 \\\\ SDSS J033710.91$-$065059.4 & SDSS\\,J0337 & \\citet{Szkody+07aj} & 19.64 & 19.54 & 19.72 & 19.97 & 20.14 & 23.27 \\\\ SDSS J160111.53$+$091712.6 & SDSS\\,J1601 & \\citet{Szkody+06aj} & 19.96 & 20.11 & 20.12 & 20.22 & 19.74 & 20.39 \\\\ SDSS J163722.21$-$001957.1 & SDSS\\,J1637 & \\citet{Szkody+02aj} & 16.83 & 16.60 & 16.59 & 16.75 & 16.84 & 20.57 \\\\ SDSS J164248.52$+$134751.4 & SDSS\\,J1642 & \\citet{Szkody+07aj} & 18.48 & 18.64 & 18.50 & 18.42 & 18.21 & 18.03 \\\\ SDSS J165837.70$+$184727.4 & SDSS\\,J1658 & \\citet{Szkody+06aj} & 20.51 & 20.07 & 20.12 & 20.09 & 19.68 & 19.71 \\\\ SDSS J165951.68$+$192745.6 & SDSS\\,J1659 & \\citet{Szkody+06aj} & 16.84 & 16.73 & 16.78 & 16.86 & 16.96 & 17.12 \\\\ SDSS J223252.35$+$140353.0 & SDSS\\,J2232 & \\citet{Szkody+04aj} & 17.74 & 17.66 & 17.80 & 17.88 & 17.97 & 23.16 \\\\ SDSS J223843.84$+$010820.7 & SDSS\\,J2238 & \\citet{Szkody+03aj} & 18.29 & 18.15 & 18.08 & 18.18 & 18.17 & 18.27 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}    \\begin{table*} \\begin{center} \\caption{\\label{tab:obslog} Log of the observations presented in this work. The acquisition magnitudes were measured from the VLT/FORS2 acquisition images and are discussed in Section\\,\\ref{sec:obs:vltphot}. The passbands for the acquisition magnitudes are indicated, where `$V$' denotes the Johnson $V$ band and `$Wh$' indicates that no filter was used.} \\begin{tabular}{lcccccccc} \\hline Target    & Date & Start time & End time & Telescope and      & Optical         &  Number of   & Exposure & Mean          \\\\ & (UT) &  (UT)      &  (UT)    &  instrument        & element         & observations & time (s) & magnitude     \\\\ \\hline SDSS\\,J0043 & 2007 08 16 & 07:47 & 10:25 & VLT\\,/\\,FORS2      & 1200R grism                  &  21 &   400  & $V = 19.9$  \\\\ SDSS\\,J0043 & 2007 08 17 & 07:13 & 09:03 & VLT\\,/\\,FORS2      & 1200R grism                  &  14 &   440  & $V = 19.8$  \\\\ SDSS\\,J0043 & 2005 10 08 & 05:05 & 06:49 & APO\\,3.5m\\,/\\,DIS  & unfiltered                   & 175 & 15--20 & $V = 19.8$  \\\\ [2pt] SDSS\\,J0337 & 2007 08 17 & 09:20 & 10:18 & VLT\\,/\\,FORS2      & 1200R grism                  &   6 &   600  & $V = 21.4$  \\\\ [2pt] SDSS\\,J1601 & 2007 08 10 & 23:25 & 00:10 & NTT\\,/\\,SUSI2      & unfiltered                   &  36 &    60  & $Wh = 20.1$ \\\\ [2pt] SDSS\\,J1637 & 2007 08 16 & 01:18 & 04:06 & VLT\\,/\\,FORS2      & 1200R grism                  &  20 &600--400& $V = 20.3$  \\\\ SDSS\\,J1637 & 2007 08 16 & 23:28 & 00:40 & VLT\\,/\\,FORS2      & 1200R grism                  &   8 &   480  & $V = 20.6$  \\\\ [2pt] SDSS\\,J1642 & 2007 08 06 & 23:38 & 02:55 & NTT\\,/\\,SUSI2      & $V$ filter                   & 262 &    30  & $V = 19.5$  \\\\ SDSS\\,J1642 & 2007 08 07 & 23:21 & 04:21 & NTT\\,/\\,SUSI2      & $V$ filter                   & 414 &    28  & $V = 19.6$  \\\\ SDSS\\,J1642 & 2007 08 17 & 00:48 & 03:41 & VLT\\,/\\,FORS2      & 1200R grism                  &  29 &   300  & $V = 18.5$  \\\\ [2pt] SDSS\\,J1658 & 2007 08 15 & 00:17 & 03:23 & VLT\\,/\\,FORS2      & 1200R grism                  &  19 &   480  & $Wh = 19.5$ \\\\ SDSS\\,J1658 & 2007 08 15 & 23:33 & 01:01 & VLT\\,/\\,FORS2      & 1200R grism                  &  12 &   400  & $V = 20.1$  \\\\ [2pt] SDSS\\,J1659 & 2007 08 11 & 00:19 & 02:45 & NTT\\,/\\,SUSI2      & $V$ filter                   & 166 & 30--60 & $V = 17.0$  \\\\ [2pt] SDSS\\,J2232 & 2005 08 11 & 00:10 & 05:34 & NOT\\,/\\,ALFOSC     & unfiltered                   & 177 &    60  & $Wh = 22.1$ \\\\ SDSS\\,J2232 & 2006 07 21 & 02:03 & 05:12 & NOT\\,/\\,ALFOSC     & unfiltered                   & 109 &    60  & $Wh = 21.8$ \\\\ SDSS\\,J2232 & 2007 08 15 & 03:37 & 06:10 & VLT\\,/\\,FORS2      & 1200R grism                  &  14 &   600  & $V = 21.4$  \\\\ [2pt] SDSS\\,J2238 & 2007 08 14 & 07:27 & 09:44 & Magellan\\,/\\,IMACS & 600\\,$\\ell$\\,mm$^{-1}$ grism &  22 &   300  &             \\\\ SDSS\\,J2238 & 2007 08 15 & 07:24 & 09:51 & Magellan\\,/\\,IMACS & 600\\,$\\ell$\\,mm$^{-1}$ grism &  17 &   300  &             \\\\ SDSS\\,J2238 & 2007 08 16 & 07:39 & 09:42 & Magellan\\,/\\,IMACS & 600\\,$\\ell$\\,mm$^{-1}$ grism &  19 &   300  &             \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{figure} \\includegraphics[width=0.48\\textwidth,angle=0]{plotsdss.eps} \\\\ \\caption{\\label{fig:sdssspec} SDSS spectra of the CVs studied in this work. For this plot the flux levels have been smoothed with 10-pixel Savitsky-Golay filters. The units of the abscissae are $10^{-21}$\\,W\\,m$^{-2}$\\,nm$^{-1}$, which corresponds to $10^{-17}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$\\,\\AA$^{-1}$.}\\end{figure} ",
        "conclusions": "\\label{sec:conclusion} \\label{sec:discussion}  \\begin{table} \\begin{center} \\caption{\\label{tab:result} Summary of the orbital periods obtained for the objects studied in this work.} \\setlength{\\tabcolsep}{4pt} \\begin{tabular}{l r@{.}l@{\\,$\\pm$\\,}r@{.}l l} \\hline Object       & \\multicolumn{4}{c}{Period (min)}       & Notes                                   \\\\ \\hline SDSS\\,J0043  &  82&325  &  0&088                      & VLT spectroscopy                        \\\\ SDSS\\,J0337  & \\multicolumn{4}{c}{}                   & Faint, no RV motion noticed             \\\\ SDSS\\,J1601  & \\multicolumn{4}{c}{short period}       & More photometry needed                  \\\\ SDSS\\,J1637  &  97&04   &  0&19                       & VLT spectroscopy                        \\\\ SDSS\\,J1642  & 113&6    &  1&5                        & VLT spectroscopy                        \\\\ SDSS\\,J1658  &  98&012  &  0&065                      & VLT spectroscopy, low state             \\\\ SDSS\\,J1659  & \\multicolumn{4}{c}{}                   & NTT photometry                          \\\\ SDSS\\,J2232  & \\multicolumn{4}{c}{}                   & VLT spec., period may be $\\sim$4.5\\,hr  \\\\ SDSS\\,J2238  & 194&30   &  0&16                       & Magellan spectroscopy                   \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  We have presented time-resolved photometry and spectroscopy of nine faint CVs which were identified by the SDSS. For five of these systems we have determined orbital periods (Table\\,\\ref{tab:result}), and four of these are shorter than the 2--3\\,hour period gap apparent in the known population of CVs. This work brings the total number of SDSS CVs with measured orbital periods to approximately 110 of the total population of 212 objects.  From VLT spectroscopy of SDSS\\,J0043 we found an orbital period of $\\Porb = 82.325 \\pm 0.088$\\,min, placing this object close to the observed minimum period for hydrogen-rich CVs. Its spectrum shows a strong contribution from the white dwarf in the system, indicating that the accretion disc is faint and the mass transfer rate is low. We have used Doppler tomography to decompose the spectra into a Doppler map in velocity space. The map shows a circular accretion disc and an inner emission peak. The latter feature can be attributed to the irradiated inner face of the secondary star, if the velocity amplitude measured from the H$\\alpha$ emission line overestimates the motion of the white dwarf by a factor of two. The Doppler map shows enhanced emission from two bright regions on the accretion disc. If the inner emission peak does indeed come from the secondary star, one of these bright regions is in the correct position to be a bright spot caused by the mass transfer stream impacting the disc. This interpretation is supported by its presence in a Doppler map of the \\ion{He}{I} 6678\\,\\AA\\ line. The second region of enhanced emission is in an unusual position in velocity space and its origin is not straightforwardly explicable. A short light curve of SDSS\\,J0043 shows evidence of a variation with a period of $207 \\pm 1$\\,s and an amplitude of $17 \\pm 5$\\,mmag. The white dwarf component may be a ZZ\\,Ceti-type pulsating star.  SDSS\\,J1637 was observed with the VLT during quiescence, resulting in an orbital period measurement of $\\Porb = 97.01 \\pm 0.19$\\,min. This object is a dwarf nova which has previously been observed in superoutburst, when superhumps with a period of $99.75 \\pm 0.86$\\,min were detected. Using the calibrations presented by \\citet{Patterson98pasp} and \\citet{Knigge06mn}, we find $q \\approx 0.16$ from these two period measurements. Assuming a white dwarf mass of 0.8\\Msun\\ gives an estimated secondary mass of 0.13\\Msun. This is in good agreement with the expected properties of a CV with $\\Porb = 97$\\,min, and together with the velocity amplitude from the H$\\alpha$ emission line point to the system having an orbital inclination of roughly 40$^\\circ$.  SDSS\\,J1642 was studied both photometrically with the NTT and spectroscopically with the VLT. The VLT data yield an unambiguous period of $\\Porb = 113.6 \\pm 1.5$\\,min. The NTT photometry shows a number of features and results in a periodogram with a small forest of peaks at frequencies below 15\\cd. The spectroscopic period is in best agreement with the peak corresponding to a period of $110.60 \\pm 0.16$\\,min, but we prefer the spectroscopic value for our final orbital period measurement. Doppler maps of the H$\\alpha$ and \\ion{He}{I} emission lines reveal a clear accretion disc and bright spot as well as weak emission from the secondary star. More surprisingly, the \\ion{He}{I} maps show an accretion disc with much higher velocities than those for H$\\alpha$, indicating that the emitting region for \\ion{He}{I} is physically smaller and closer to the white dwarf than for H$\\alpha$. The H$\\alpha$ map also shows strong emission arising from the part of the disc opposite the bright spot, but the short duration of our spectroscopic observations means this unusual feature could be caused by brightness variation which differ from orbit to orbit.  SDSS\\,J1658 is perhaps the most unusual system studied in this work. Its SDSS spectrum shows the strong emission lines characteristic of a short-period CV whose light is dominated by a hydrogen-rich accretion disc. However, our VLT spectra show a much weaker and narrower central emission line, and broad absorption features arising from both the white dwarf and secondary star, but no flux which could be unambiguously assigned to an accretion disc. The object was only 0.4\\,mag fainter during our observations than when the SDSS spectrum was taken, so the strong emission lines in that spectrum were accompanied by only weak continuum flux from the accretion disc. The velocity variation of the narrow H$\\alpha$ emission in our VLT spectra yields a period measurement of $\\Porb = 98.012 \\pm 0.065$\\,min, confirming that this object is a short-period binary star system. The velocity amplitude (125\\kms) is too large for the white dwarf, \\reff{so we have attempted to assign it to the secondary star. The observational constraints can be satsified in this scenario if the white dwarf is massive ($\\la$1\\Msun) and the orbital inclination is low (10--30$^\\circ$).}  A modest number of photometric and spectroscopic observations of SDSS\\,J2232 suggest that this is a dwarf nova with an orbital period close to 4.5\\,hr. \\reff{From the flux contribution of the secondary star we infer a distance of roughly 3\\,kpc, corresponding to distance of 2\\,pkc from the Galactic plane.} Further investigation is needed to prove its membership of this old stellar population.  We obtained 59 Magellan spectra of SDSS\\,J2238, which is the brightest of the five objects for which we determine orbital periods in this work. Velocity measurements of its H$\\alpha$ emission line give a period of $\\Porb = 194.30 \\pm 0.16$. This value is in good agreement with a published photometric period, which also confirmed that the system contains a magnetic white dwarf with a rotational period of 6.7284\\,min.  These observations provide further confirmation that the faintest of the CVs identified by the SDSS have predominantly short orbital periods. Theoretical population studies of CVs predict a huge population of faint short-period CVs which has not previously been detected. Our observations are now uncovering this `quiet majority' of the CV population. The unusual characteristics of several of the objects studied in this work show that even the short-period CVs demonstrate impressively varied behaviour, many aspects of which cannot easily be explained in the standard picture of CV structure. Our project to study the SDSS CV population is invaluable for extending our knowledge of these fascinating objects."
    },
    "0809/0809.0506_arXiv.txt": {
        "abstract": "We use a Press-Schechter-like calculation to study how the abundance of voids changes in models with non-Gaussian initial conditions.  While a positive skewness increases the cluster abundance, a negative skewness does the same for the void abundance.  We determine the dependence of the void abundance on the non-Gaussianity parameter $\\fnl$ for the local-model bispectrum---which approximates the bispectrum in some multi-field inflation  models---and for the equilateral bispectrum, which approximates the bispectrum in e.g. string-inspired DBI models of inflation.  We show that the void abundance in large-scale-structure surveys currently being considered should probe values as small as $\\fnl \\lesssim 10$ and $\\fnl^{\\mathrm{eq}}\\lesssim 30$, over distance scales $\\sim10$~Mpc. ",
        "introduction": "The paradigm of cosmological structure formation from a spectrum of primordial perturbations like those predicted by inflation has now been fairly well established by cosmic microwave background (CMB) experiments \\cite{cmbexperiments}. We are thus now motivated to test more precisely the predictions of inflation and to look for possible deviations.  One of several such possibilities is measurement of departures from Gaussianity of the initial perturbations (see, e.g., Ref.~\\cite{NGReview} and references therein). The simplest slow-roll single-field models of inflation predict that primordial perturbations should be very closely Gaussian \\cite{inflation}, but with predictably small departures from Gaussianity \\cite{localmodel,Equil}.  Multi-field \\cite{larger} models, such as the curvaton model \\cite{curvaton}, and string-inspired DBI \\cite{Dvali:1998pa} inflationary models can produce larger deviations from non-Gaussianity.  Departures from primordial Gaussianity can be sought in the CMB \\cite{Luo:1993xx,Verde:1999ij,Komatsu:2001rj}, large-scale structure (LSS) \\cite{LSS}, and the abundances and properties of the most massive gravitationally-bound objects in the Universe today or at high redshift \\cite{abundances,MVJ00,Verde:2000vr}. The CMB provides a more powerful  and clean probe of primordial non-Gaussianity than direct measurement of the bispectrum in low-redshift LSS surveys in models with scale-invariant non-Gaussianity \\cite{Verde:1999ij}, although biasing may amplify the effects of non-Gaussianity on LSS to the level where they may be comparable in detectability to the CMB % \\cite{Dalal:2007cu, MV08, Slosar:2008hx,CVM08}.  Measurements of the cluster abundance may do better than the CMB and LSS if the non-Gaussianity is not scale-invariant \\cite{LoVerde:2007ri}, as may occur in DBI models.  What is clear, however, is that the thorny systematic effects that enter in all of these approaches will require that a variety of complementary avenues be taken to establish a robust detection of non-Gaussianity.  Voids have been considered as probes of cosmology, but no systematic study has been carried out for voids as probes of primordial non-gaussianity \\cite{grossivoids}. In this paper, we consider the abundance of voids as a test of the distribution of the primordial perturbations.  Galaxy clusters form at the highest overdensities of the primordial density field and thus probe the high-density tail of the primordial density distribution function.  Similarly, voids form in low-density regions and should thus probe the low-density tail of the distribution function.  If there is a large negative skewness, the void-size distribution function will be increased at the largest void sizes and decreased at smaller void sizes, opposite to the effect on the cluster mass function.  In Section \\ref{sec:PSabundance}, we develop a Press-Schechter (PS) estimate of the void abundance for Gaussian initial conditions.  This PS-like calculation is far from state of the art \\cite{Sheth:2003py,Furlanetto:2005cc} for Gaussian initial conditions.  However, it is easily generalized to non-Gaussian initial conditions and should be sufficiently reliable to estimate the {\\it fractional} effects of non-Gaussianity on the void abundance (as it describes  well the halo abundance \\cite{paper_in_prep}.  In Section \\ref{sec:skewness}, we discuss the relation between the skewness and the non-Gaussian parameter $\\fnl$ for the local model \\cite{localmodel}, which approximates the non-Gaussianity in multi-field models, and the equilateral model \\cite{Equil}, which approximates that in e.g. string-inspired DBI models. In Section \\ref{sec:results}, we provide results of the void-abundance calculation, and we estimate the smallest $\\fnl$, for the local model and for the equilateral model (with and without scale dependence) that should be detectable in several surveys currently under study. In Section \\ref{sec:conclusions}, we make some concluding remarks and outline further steps that must be taken before the void abundance can be used to probe non-Gaussianity. ",
        "conclusions": "\\label{sec:conclusions}  The bottom line of our analysis is that the abundance of voids in LSS surveys that are currently being considered can probe values of the non-Gaussianity parameter in the local model as small as $\\fnl\\sim10$ and in the equilateral model (with no scale dependence) as small $\\fnl^{\\mathrm{eq}}\\sim 30$.  This probe of non-Gaussianity may be competitive with those coming from the CMB, LSS, and cluster abundances \\cite{Verde:1999ij,Verde:2000vr,LoVerde:2007ri}. They will complement CMB constraints by (a) providing a different avenue with different systematic effects, and (b) probing non-Gaussianity primarily over distance scales 2--60 Mpc, scales generally smaller than those that will be probed by the CMB. The fact that voids constraints on  the different types of non gaussianity are comparable  is particularly interesting. In fact the large-scale bias effect due to  equilateral non-gaussianity  is  orders of magnitude smaller than the effect for the local case. Thus  the combination of the two measurements will  help discriminate among different type of non-gaussian initial conditions.  Finally, we note that we have here done no more than estimate the smallest $\\fnl$ detectable with the void abundance. To do so, we have taken a simple Press-Schechter approach to estimate the fractional change in the void abundance. While this approach should provide reasonable rough estimates to the fractional change in the abundance, it will not reliably provide the abundances themselves.  Before the void-abundance probe of $\\fnl$ can be implemented, we will therefore require significantly more sophisticated calculations, which will probably require numerical simulations, to accurately model not only the void growth, but also the systematic effects that will arise in any realistic void-identification algorithm.  We hope that our results motivates this type of future work."
    },
    "0809/0809.0409.txt": {
        "abstract": "We present the results of optical panoramic and long-slit spectroscopy of the nebula MF16 associated with the Ultraluminous X-ray Source NGC6946~ULX-1. More than 20 new emission lines are identified in the spectra. Using characteristic line ratios we find the electron density $n_e \\sim 600\\,\\cmc$, electron temperature in the range from $ \\sim 9\\,000~\\rm K$ to $ \\sim 20\\,000\\,\\rm K$ (for different diagnostic lines) and the total emitting gas mass $M \\sim 900 \\Msun$. We also estimate the interstellar extinction towards the nebula as $A_V \\simeq 1\\magdot{.}54$ somewhat higher than the Galactic absorption. Observed line luminosities and ratios appear to be inconsistent with excitation and ionization by shock waves so we propose the central object responsible for powering the nebula. We estimate the parameters of the ionizing source using photon number estimates and {\\it Cloudy } modelling. Required EUV luminosity ($\\sim 10^{40}$\\ergl) is high even if compared with the X-ray luminosity. We argue that independently of their physical nature ULXs are likely to be bright UV and EUV sources. It is shown that the UV flux expected in the {\\it GALEX} spectral range (1000$\\div$3000$\\AAA$) is quite reachable for UV photometry. Measuring the luminosities and spectral slopes in the UV range may help to distinguish between the two most popular ULX models. ",
        "introduction": "\\subsection{ULXs and their Optical Nebulae}  A point-like X-ray source in an external galaxy is considered an Ultraluminous X-ray source (ULX) if its luminosity exceeds $10^{39} \\ergl$ and it is not an active galactic nucleus. It also makes sense to exclude X-ray bright SNe and young (such as several years) Supernova Remnants (SNR) that can shine as bright as $10^{41}\\ergl$ in X-rays like the remnant of SN1988Z \\citep{FaTer}. There are more than 150 ULXs known at the present time \\citep{swartz} but very little is clear yet about the physical nature of these objects.  ULXs are widely accepted as a class of accreting compact objects violating the Eddington luminosity limit for a conventional stellar mass black hole (about $1.3 \\times 10^{39}$\\ergl\\ for $10\\Msun$). High luminosity may be a consequence of a higher accretor mass, supercritical accretion, mild geometrical collimation, relativistic beaming \\citep{koer} or some combination of these effects. Two models are the most popular: {\\it (i)} Intermediate-Mass Black Holes (IMBHs) with masses in the range $10^2\\div 10^4 \\Msun$ \\citep{IMBH,MR2001} accreting at sub-Eddington rates; {\\it (ii)} supercritical accretion discs like that of SS433 around stellar mass black holes observed at low inclinations $i \\lesssim 20^\\circ$ \\citep{katz86,king,FaMes}. Some authors \\citep{soria_kuncic} propose ``hybrid'' models involving $20\\div 100~\\Msun$ black holes accreting in a moderately supercritical regime. Supercritical accretion discs are considered by different authors either in the wind-dominated regime first introduced by \\citet{SS73} or in the advection-dominated regime known as slim disc \\citep{ACLS1988,Okajima2007}.  Large number of ULXs are surrounded by large-scale bubble nebulae \\citep{pakull03}, probably shock-powered. ULX nebulae (ULXNe), however, do not form a homogeneous class of objects. In \\citet{list} we review the observational properties of 8 ULXNe including the object under consideration. We conclude that only about 50\\% of ULXNe may be considered shock-powered shells. In some of the observed ULXNe high-excitation lines such as \\oiii$\\lambda$4959,5007 are enhanced. In some cases HeII$\\lambda$4686 emission is detected as bright as 0.2$\\div$0.3 in H$\\beta$ units \\citep{lehmann}. The line may have stellar origin \\citep{kuntz} as well as nebular. Measuring line widths may help to distiguish between these two cases. %The question about the sources of ionization in ULX nebulae remains %open. % and may be solved differently for different objects. %In \\citet{list} we review the observational %properties of 8 well-studied ULX nebulae, including the object under %consideration. %They appear to be diverse in their observational properties. %In different cases either shocks or photoionization seem %to play the main role. The nebulae may clarify the nature of corresponding X-ray sources in two ways: {\\it (i)} probing the photoionizing radiation of the central object via gas excitation and ionization conditions; {\\it (ii)} detecting and measuring jet/wind activity by kinematical effects and shock emission.  Dynamically disturbed gas may be considered an evidence against the IMBH hypothesis because a standard accretion disc is unlikely to produce strong jets or wind. Supercritical accretion, on the other hand, is supposed to provide a massive outflow in the form of a strong wind carrying practically all the accreted mass and having kinematical luminosity $\\sim 10^{38}$\\ergl\\ in the case of SS433 \\citep{ss2004}. Power comparable with the Eddington luminosity ($\\sim 10^{39}$\\ergl) may be provided in  the form of relativistic jets. Their mechanical luminosity acting for $\\sim 10^5$ years is sufficient to produce a wind-blown bubble with properties close to those of some of the existing ULX nebulae \\citep{pakull03,list}.  ULXNe exhibit very diverse observational properties such as size, morphology and line ratios \\citep{list,pakull03}. Integral luminosities in Balmer lines are in some cases of the order $10^{38}$\\ergl\\ or higher that requires an energy source with luminosity $\\gtrsim 10^{39}$\\ergl\\ %Existence of ULX nebulae sometimes as bright as $\\gtrsim 10^{38}$\\ergl %\\citep{pakull03,list} in optical emission lines, indicates that ULXs are indicating that ULXs are really powerful objects. Therefore strong beaming effects are excluded as the reason of apparently high luminosities of ULXs. For some objects like HoIX~X-1 and IC342~X-1 the nebula is clearly shock-powered. In other cases (HoII~X-1, M101-P098) low intensity of low-excitation lines, bright \\oiii\\ and \\heii\\ emissions and quiet kinematics\\footnote{For HoII~X-1 \\citet{lehmann} report velocity dispersion $\\sim 13 \\kms$ in the vicinity of the X-ray source.} suggest photoionization to be the main energy source. %Supercritical accretors are likely to provide both outflows with the %power close to the Eddington luminosity and EUV luminosity of the same %order or higher.  The only proven example of a persistent supercritical accretor is the peculiar binary SS433 \\citep{ss2004}. Its accretion rate is of the order $10^{-4}\\Msun\\,\\rm yr^{-1}$, indicating mass transfer on the thermal timescale of a massive star. %Significant part of the outflow %energy is depleted in the form of a pair of relativistic jets, powering a %bright elongated radio nebula W50. Most of the accreting material is ejected in the form of optically-thick wind with velocities $V \\simeq 1000 \\div 2000 $\\kms. The total kinetic power of the jets of SS433 is of the order of $10^{39}$\\ergl. SS433 is surrounded by a large (roughly $40\\times 120$\\pc) radionebula W50 \\citep{dubner}. The nebula harbors several optical emission-line filaments situated in the propagation direction of the jets. % and originating most probably in the secondary shocks produced % by jet activity. Large difference in angular size and heavy absorption in the case of W50 make comparison with ULXNe difficult. %The total optical %emission-line luminosity of the nebula is highly uncertain due to the large angular size of the nebula %($\\sim 1^\\circ$).  \\subsection{NGC6946~ULX-1}  The X-ray source under consideration is known as NGC6946~ULX-1 \\citep{swartz}, NGC6946~X-8 \\citep{MF16_lira}, N58 \\citep{MF16_Holt} or NGC6946~X-11 \\citep{RoCo}. It was first detected by {\\it ROSAT} \\citep{schlegel94} %% first detection! and identified with an optical nebula by % \\citet{BF_94} and a marginally resolved radiosource by \\citet{vandyk}.  %PA  The nebula is known as MF16 -- Matonick and Fesen~16 \\citep{MF_cat}. It was considered a SNR luminous in X-rays, in fact the most luminous in X-rays among the optically bright SNRs \\citep{dunne}, till it was proved by \\citet{RoCo} that the X-ray emission originates from a much more compact source in the center of the nebula.  {\\it Chandra} source coordinates for J2000 epoch are: $ \\alpha = 20^h 35^m 00^s.74$, $\\delta = +60^\\circ\\, 11\\arcmin{\\,} 30\\farcs6$. %{\\it Chandra} coordinates have absolute accuracy about $0\\farcs6$ %\\citep{swartz} mainly originating from systematic errors.  The host galaxy NGC6946 is a late-type spiral with active star formation \\citep{degioia_sfr} and the greatest number of detected supernovae (no less than 8, according to \\citet{SN_n6946}). The distance to the galaxy is estimated by different authors as 5.1 \\citep{dist_5.1}, 5.5 \\citep{dist_5.5}, 5.7 \\citep{dist_5.7} and 5.5$\\,\\Mpc$ \\citep{dist_5.9} therefore we assume $D=5.5\\,\\Mpc$. The divergence in distance estimates leads to a 15\\% uncertainty in all the luminosities inferred. Spatial scale for $D=5.5\\,\\Mpc$ is about $27\\,\\pc$ per arcsecond.  %The nebula is known as MF16 -- Matonick\\&Fesen~16 \\citep{MF_cat}. It was %considered a SNR luminous in X-rays, in fact the %most luminous in X-rays among the optically %bright SNRs \\citep{dunne}, till it was proved by \\citet{RoCo} %that the X-ray emission originates from a compact %source in the center of the nebula.  A point source in the center of MF16 was detected by \\citet{BFS} in {\\it HST} data. The source (star {\\bf d} as denoted by \\citet{BFS}) has a V magnitude of $22\\magdot{.}64$ and a colour index of $B-V = 0\\magdot{.}46$. The object is significantly redder than %in comparison with %relative to the nearby stars \\textbf{a}-\\textbf{c} offset to the North-West by several seconds. %(stars \\textbf{a}-\\textbf{c} as denoted by \\citet{BFS}). %If the intrinsic absorption is negligible If one accounts only for Galactic interstellar extinction (see discussion in section \\ref{sec:charlines}) equal to $A^{(GAL)}_V = 1\\magdot{.}14$ according to \\citet{schlegel_abs} in the direction of MF16, the point-like counterpart resembles an early A~Ia supergiant with $M_V \\simeq -7\\magdot{\\,}$ and $B-V \\simeq 0\\magdot{\\,}$. % \\citep{leng}. Additional intrinsic absorption $A_V \\sim 0\\magdot{.}5$ (reported by \\citet{BFS}) places the object slightly above the Ia sequence in the upper left part of the HR diagram.  Optical emission-line spectrum of the nebula is generally consistent with the suggestion of shock heating but also contains high-excitation lines such as \\heii$\\lambda4686$ and \\oiii$\\lambda5007,4959$ doublet \\citep{BF_94, BFS} too intensive for optically-bright SNRs \\citep{MFB_opt}. Partially radiative shocks were proposed to explain the emission-line spectrum of MF16 \\citep{BFS} but they require a very powerful energy source (see discussion in section \\ref{sec:balance}). %, as it will be shown %below in section~\\ref{sec:balance}.  The object is much brighter in the optical than a usual SNR. Its H$\\alpha$ luminosity is about $2\\times 10^{38}$\\ergl, an order of magnitude higher than the upper limiting H$\\alpha$ luminosity for optically-bright SNRs in nearby galaxies \\citep{braunm31}. \\citet{dunne} detected two-component structure in the emission lines of the object in high-resolution echelle spectra. Broader components suggest expansion velocities about $250$\\kms, narrower components have velocity dispersion $\\sim 20\\div 40$\\kms.  MF16 is also known as a radio source as bright as $1.3\\rm mJy$ at $20\\rm cm$ \\citep{vandyk} marginally resolved by VLA observations as a $\\sim 1\\arcsec{\\,}$ size object. The radio source appears to be displaced relative to the X-ray object and the optical source {\\bf d} by $\\sim 0\\farcs5$ (see figure ~\\ref{fig:FOV}). Luminosity of the radio counterpart is about 20 times higher than that of W50 at 20$\\rm cm$ \\citep{dubner}. The size of the radiosource is poorly known but the optical nebula is about 5 times more compact than W50.  MF16 is practically isolated from other HII regions. The nearest one with comparable brightness is situated about $200\\pc$ away \\citep{BFS}. The inner brighter part of the nebula has a shape of a shell $20 \\pc \\times 34 \\pc$ elongated in East to West direction with a brighter Western loop. In the deep {\\it HST} images a faint asymmetric halo may be seen \\citep{BFS}. %(figure \\ref{fig:FOV}).  In the following section we describe the observational data. In section~\\ref{sec:spectra} we present the results of spectral analysis. In section~\\ref{sec:balance} we analyze the excitation and ionization sources of the nebula and model the observed spectrum using {\\it Cloudy}. In section~\\ref{sec:trep} we discuss our results and the perspectives of UV observations of ULXs. ",
        "conclusions": "In our optical spectra of MF16 we detect more than 30 %% PA: emission lines including several high-excitation lines of \\heii\\ and \\ariv\\ and a rich spectrum of moderately high excitation \\feiii\\ lines. \\heii\\ lines are narrow (broadened by $\\lesssim 300\\kms$). We also do not detect any Wolf-Rayet features. Therefore we suggest that the \\heii$\\lambda$4686 line is of nebular origin.  Large intensities of iron lines indicate moderately low depletion of the element into dust ($30\\div 50\\%$ in the gas phase), probably due to dust destruction in shock waves. Destruction is probably incomplete because {\\it Cloudy} models overestimate the intensities of \\feiii\\ lines.  We find the electron density in the nebula $n_e = 570\\, \\pm 60\\,\\cmc$ (\\sii$\\lambda$6717,6731), electron temperatures for different ions are $T(\\oiii) = 17\\,700\\pm1\\, 200\\,\\rm K$, % (\\oiii\\ diagnostic lines), $T(\\nii) = 15\\,600\\pm 2\\,000\\,\\rm K$ % (\\nii\\ lines) and $T(\\sii) = 9\\,000\\pm 1000\\,\\rm K$ % (\\sii\\ lines) indicating the presence of regions with different electron temperatures. Total hydrogen emitting gas mass is $M \\sim 900\\,\\Msun$. Interstellar absorption is $A_V \\simeq 1\\magdot{.}34$ in traditional H$\\alpha$/H$\\beta$=3 assumption. More realistic value  H$\\alpha$/H$\\beta$=2.8 results in a higher extinction value $A_V = 1\\magdot{.}54$.  The observed line luminosities and diagnostic line ratios appear to be inconsistent with excitation and ionization by shock waves, therefore we suggest an EUV source responsible for powering the nebula. Photoionization modelling with {\\it Cloudy} as well as Zanstra estimates suggest the central source must be ultraluminous not only in X-rays but also in the UV/EUV range emitting about $10^{40}$\\ergl\\ in the spectral range $100-1000\\AAA$. Using the observed X-ray spectrum, ionizing flux estimates from \\heii, \\hei\\ and H lines and the optical point-like counterpart (star {\\bf d}) we reconstruct the SED of the central object from X-rays to the optical. The derived spectrum is roughly flat, $\\nu L_{\\nu} \\sim const$.  Both most popular models of ULXs (IMBHs and supercritical accretion discs) predict high UV/EUV luminosities. However these two models predict different spectral slopes in the UV region. We conclude that measuring the spectral slope will help to distinguish between the two models as well as to determine the parameters of the successfull model such as the black hole mass in the case of IMBH or accretion rate in the case of a supercritical accretion.  Actual excitation and ionization conditions in MF16 may be much more complicated. Though introducing a bright EUV source is a possible way to explain the spectrum, we suggest that other explanations such as high pre-shock density and multiple shocks with different velocities are not completely excluded. Observations with higher angular and spectral resolution (but also higher S/N ratio than that used by \\citet{dunne}) may provide additional information about the enigmatic nebula MF16.  %% Acknowledgements: \\bigskip  We thank V. Afanasiev and N. Borisov for assistance with the observations and the anonymous referee for valuable comments and suggestions. This work was supported by the Russian RFBR grants 06-02-16865 and 07-02-00909 and the RFBR/JSPS grant 05-02-19710. TK is supported by a 21st Century COE Program at Tokyo Tech ``Nanometer-Scale Quantum Physics'' by the Ministry of Education, Culture, Sports, Science and Technology.  This work is supported by the Japan-Russia Research Cooperative Program of Japan Society for the Promotion of Science.  \\newpage"
    },
    "0809/0809.2997_arXiv.txt": {
        "abstract": "We report the detection of the cool, Jovian-mass planet MOA-2007-BLG-400Lb. The planet was detected in a high-magnification microlensing event (with peak magnification $A_{\\rm max} = 628$) in which the primary lens transited the source, resulting in a dramatic smoothing of the peak of the event.  The angular extent of the region of perturbation due to the planet is significantly smaller than the angular size of the source, and as a result the planetary signature is also smoothed out by the finite source size.  Thus the deviation from a single-lens fit is broad and relatively weak ($\\sim$ few percent).  Nevertheless, we demonstrate that the planetary nature of the deviation can be unambiguously ascertained from the gross features of the residuals, and detailed analysis yields a fairly precise planet/star mass ratio of $q=2.6\\pm 0.4 \\times 10^{-3}$, in accord with the large significance ($\\Delta\\chi^2=1070$) of the detection.  The planet/star projected separation is subject to a strong close/wide degeneracy, leading to two indistinguishable solutions that differ in separation by a factor of $\\sim 8.5$.  Upper limits on flux from the lens constrain its mass to be $M<0.75\\,M_\\odot$ (assuming it is a main-sequence star). A Bayesian analysis that includes all available observational constraints indicates a primary in the Galactic bulge with a mass of $\\sim 0.2-0.5 M_\\odot$ and thus a planet mass of $\\sim 0.5-1.3 M_{\\rm Jup}$.  The separation and equilibrium temperature are $\\sim 0.6-1.1~{\\rm AU}$ ($\\sim 5.3-9.7~{\\rm AU}$) and $\\sim 103~{\\rm K}$ ($\\sim 34~{\\rm K}$) for the close (wide) solution. If the primary is a main-sequence star, follow-up observations would enable the detection of its light and so a measurement of its mass and distance. ",
        "introduction": "}  In the currently favored paradigm of planet formation, the location of the snow line in the protoplanetary disk plays a pivotal role.  Beyond the snow line, ices can condense, and the surface density of solids is expected to be higher by a factor of several relative to its value just inside this line.  As a result of this increased surface density, planet formation is expected to be most efficient just beyond the snow line, whereas for increasing distances from the central star the planet formation efficiency drops, as the surface density decreases and the dynamical time increases \\citep{lissauer87}.  In this scenario, gas-giant planets must form in the region of the protoplanetary disk immediately beyond the snow-line, as the higher surface density is required to build icy protoplanetary cores that are sufficiently massive to accrete a substantial gaseous envelope while there is remaining nebular gas \\citep{pollack96}.  Low-mass primaries are expected to be much less efficient at forming gas giants because of the longer dynamical times and lower surface densities at the snow lines of these stars \\citep{laughlin04,ida05,kennedy08}. Migration due to nebular tides and other dynamical processes can then bring the icy cores or gas giants from their formation sites to orbits substantially interior to the snow line \\citep{lin96,ward97,rasio96}.  The precise location of the snow line in protoplanetary disks is a matter of some debate (e.g., \\citealt{lecar06}), and is even likely to evolve during the epoch of planet formation, particularly for low-mass stars \\citep{kennedy06,kennedy08}.  The condensation temperature of water is $\\sim 170~{\\rm K}$, and a fiducial value for the location of the snow line in solar-mass stars motivated by observations in our solar system is $\\sim 2.7~{\\rm AU}$.  This may scale linearly with the stellar mass $M$, since the stellar luminosity during the epoch of planet formation scales as $\\sim M^2$ for stars with $M \\la M_\\odot$ \\citep{burrows93,burrows97}.  Whereas the radial velocity and especially transit methods are most sensitive to planets that are close to their parent star at distances well inside the snow line, the sensitivity of the microlensing method peaks at planetary separations near the Einstein ring radius of the primary lens \\citep{mp91,gould92}, which is $\\sim 3.5~{\\rm AU}(M/M_\\odot)^{1/2}$ for typical lens and source distances of $6~{\\rm kpc}$ and $8~{\\rm kpc}$, respectively.  This corresponds to a peak sensitivity at equilibrium temperatures of $T_{\\rm eq} \\sim 150~{\\rm K} (M/M_\\odot)$ for a mass-luminosity relation of the form $L \\propto M^{5}$, and distances relative to the snow line of $\\sim 1.3 (M/M_\\odot)^{-1/2}$ if the location of the snow line at the epoch of planet formation scales as $M$. Thus microlensing is currently the best method of probing planetary systems in the critical region just beyond the snow line \\citep{gould92}.  Planetary perturbations in microlensing events come in two general classes.  The majority of planetary perturbations are expected to occur when a planet directly perturbs one of the two images created by the primary lens, as the image sweeps by the planet during the microlensing event \\citep{gould92}.  Although these perturbations are more common, they are also unpredictable and can occur at any time during the event.  Early microlensing planet searches focused on this class of perturbations, as it was the first to be identified and explored theoretically \\citep{gould92,bennett96,gaudi97}.  The second class of planetary perturbations occurs in high-magnification events, in which the source becomes very closely aligned with the primary lens \\citep{griestsafi}.  In such events, the two primary-lens images become highly distorted and sweep along nearly the entirety of the Einstein ring \\citep{liebes}.  These sweeping images probe subtle distortions of the Einstein ring caused by nearby planets, which will give rise to perturbations within the full-width half-maximum of the event \\citep{bond02,rattenbury02}.  Although high-magnification events are rare and so contribute a minority of the planetary perturbations, they are individually more sensitive to planets because the images probe nearly the entire Einstein ring. Furthermore, since the perturbations are localized to the peak of the event which can be predicted beforehand, they can be monitored more efficiently with limited resources than the more common low-magnification events.  For these reasons, current microlensing planet searches tend to deliberately focus on high-magnification events.  Thus, of the seven prior microlensing planets discovered to date \\citep{ob03235,ob05071,ob05390,ob05169,ob06109,mb07192}, five have been found in high-magnification events, with peak magnifications ranging from $A=40$ to $A=800$.  However, despite the fact that this planet-search strategy has proven to be so successful, the properties of the planetary perturbations generated in high-magnification events are less well-understood than those in low-magnification events.  Most of the studies of the properties of planetary perturbations in high-magnification events have focused on the properties of the caustics, the locus of points defining one or more closed curves, upon which the magnification of a point source is formally infinite.  The morphology and extent of the region of significant perturbation by the planetary companion can be largely understood by the shape and size of these caustic curves.  Planetary perturbations in high-magnification events are caused by a central caustic located near the position of the primary, and thus several authors have considered the size and shape of these central caustics as a function of the parameters of the planet \\citep{griestsafi,dominik99,chung05}. These and other authors have identified several potential degeneracies that complicate the unique interpretation of central caustic perturbations.  The first to be identified is a degeneracy such that the caustic structure (and so light curve morphology) of a planet with mass ratio $q \\ll 1$ and projected separation in units of the Einstein ring $d$ not too close to unity is essentially identical under the transformation $d \\leftrightarrow d^{-1}$ \\citep{griestsafi}.  A second degeneracy arises from the fact that very close or very wide roughly equal-mass binaries also produce perturbations near the peak of the light curves.  These give rise to perturbations that have the same gross observables as planetary perturbations.  The severity of these degeneracies depends on both the specific parameters of the planetary/binary companion, as well as on the data quality and coverage. \\citet{griestsafi} and \\cite{chung05} demonstrated that the $d \\leftrightarrow d^{-1}$ degeneracy is less severe for more massive planets with separations closer to the Einstein ring ($d \\sim 1$).  Empirically, this degeneracy was broken at the $\\Delta \\chi^2 \\sim 4$ level for the relatively large mass-ratio planetary companion OGLE-2005-BLG-071Lb \\citep{dong08}, for which the light curve was well-sampled, but was essentially unresolved for the low mass-ratio planetary companion MOA-2007-BLG-192Lb \\citep{mb07192}, for which the planetary perturbation was poorly sampled.  For the planetary/equal-mass binary degeneracy, \\citet{han08} argued that, although the gross features of central caustic planetary perturbations can be reproduced by very close or very wide binary lenses, the morphologies differ in detail, and thus this degeneracy can be resolved with reasonable light curve coverage and moderate photometric precision. Indeed for every well-sampled high-magnification event containing a perturbation near the peak (and that is not in the \\citet{cr1} limits), this degeneracy has been resolved \\citep{albrow02,ob05169,dong08}.  Even for the relatively poorly-sampled light curve of MOA-2007-BLG-192Lb, an equal-mass binary model is ruled out at $\\Delta\\chi^2\\sim 120$ \\citep{mb07192}, although in this case this is partially attributable to the exquisite photometric precision ($<1\\%$).  One complication with searching for planets in high-magnification events is that, the higher the magnification, the more likely it is that the primary lens will transit the source.  When this happens, the peak of the event is suppressed and smoothed out, as the lens strongly magnifies only a small portion of the source. If the source is also larger than the region of significant perturbations due to a planetary companion (roughly the size of the central caustic), then the planetary deviations will also be smoothed out and suppressed \\citep{griestsafi,han07}. These finite source effects have potential implications for both the detectability of central caustic perturbations, as well as the ability to uniquely determine the planetary parameters, and in particular resolve the two degeneracies discussed above.  In practice, the caustic structures of all four high-magnification planetary events (containing five planets) were larger than the source.  Hence, while there were detectable finite-source effects in all cases (which helped constrain the angular Einstein radius and so the physical lens parameters), the planetary perturbations were in all four cases quite noticeable.  Thus the effect of large sources on the the detectability and interpretation of central-caustic perturbations has not been explored in practice.  Theoretical studies of detectability of central caustic perturbations when considering finite source effects have been performed by \\citet{griestsafi}, \\citet{chung05}, and \\citet{han07}.  These authors demonstrated that the qualitative nature of planetary perturbations from central caustics is dramatically different for sources that are larger than the caustic.  In particular, the detailed structure of the point-source magnification pattern, which generally follows the shape of the caustic, is essentially erased or washed out.  Rather, the perturbation structure is characterized by a roughly circular region of very low-level, almost imperceptible deviation from the single-lens form that is roughly the size of the source and centered on the primary lens.  This region is surrounded by an annular rim of larger deviations that has a width roughly equal to the width of the caustic \\citep{griestsafi,chung05}. Finally, there are less pronounced deviations that extend to a few source radii. Planets are detectable even if their central caustics are quite a bit smaller than the source, provided that the $\\chi^2$ deviation is sufficiently high.  (Often $\\Delta\\chi^2>60$ is adopted, although $\\Delta\\chi^2>150$ may be more realistic.)  The magnitude of these perturbations decrease as the ratio between the source size and caustic size increases, making it difficult to detect very small planets for large sources \\citep{chung05,han07}.  Although central caustics may formally be detectable when the source is substantially larger than the caustic, it remains a significant question whether these very washed out caustics can be recognized in practice, and even if they can, whether they can be uniquely interpreted in terms of planetary parameters.  Indeed, it is unknown whether a washed out central caustic due to a planet can actually be distinguished from one due to a binary companion.  This question is especially important with regard to low-mass planets.  The size of the central caustic scales as the product of the planet/star mass ratio and a definite function of planet-star separation.  Hence, taken as a whole, smaller planets produce smaller caustics, meaning that events of higher magnification are required to detect them. These are just the events that are most likely to have their peaks washed out by finite-source effects.  Here we analyze the first high-magnification event with a buried signature of a planet, in which the source size is larger than the central caustic of the planet.  The caustic is indeed so washed out that the event appears unperturbed upon casual inspection.  However, the residuals to a point-lens fit are clear and highly significant. We show that one can infer the planetary (as opposed to binary) nature of the perturbation from the general pattern of these residuals, and that a detailed analysis constrains the mass ratio of the planet quite well, but leaves the close/wide ($d \\leftrightarrow d^{-1}$) degeneracy intact. Hence, at least in this case, the fact that the caustic is buried in the source does not significantly hinder one's ability to uncover the planet and measure its mass ratio. ",
        "conclusions": "}  MOA-2007-BLG-400 is the first high-magnification microlensing event for which the central caustic generated by a planetary companion to the lens is completely enveloped by the source. As a comparison, the planetary caustic of OGLE-2005-BLG-390 \\citep{ob05390} is smaller than its clump-giant source star in angular size. When the planetary caustics is covered by the source, the finite-source effects broaden the ``classic'' \\citet{gould92} planetary perturbation features \\citep{gaudi97}. By contrast, planet-induced deviations in MOA-2007-BLG-400 are mostly obliterated, rather than being broadened, because the source crosses the central caustic rather than the planetary caustic. We showed, nevertheless, that the planetary character of the event can be inferred directly from the light-curve features and that the standard microlensing planetary parameters $(d,q)=(2.9,2.6\\times 10^{-3})$ can be measured with good precision, up to the standard close/wide $d\\leftrightarrow d^{-1}$ degeneracy.  We demonstrated that, in this case, the close/wide degeneracy is quite severe, and the wide solution is only preferred by $\\Delta\\chi^2 = 0.2$. This is unfortunate, since the separations of the two solutions differ by a factor of $\\sim 8.5$. We argued that the severity of this degeneracy was primarily related to the intrinsic parameters of the planet, rather than being primarily a result of the large source size.  Although the mass ratio alone is of considerable interest for planet formation theories, one would also like to be able to translate the standard microlensing parameters to physical parameters, i.e., the planet mass $m_p = qM$, and planet-star projected separation $r_\\perp = d\\theta_\\e D_L$.  Clearly this requires measuring the lens mass $M$ and distance $D_L$. In this case, the pronounced finite source effects have already permitted a measurement of the Einstein radius $\\theta_\\e=0.32\\,\\mas$, which gives a relation between the mass and lens-source relative parallax (eq.~[\\ref{eqn:mpirel}]). This essentially yields a relation between the lens mass and distance, since the source distance is close enough to the Galactic center that knowing $D_L$ is equivalent to knowing $\\pi_{\\rm rel}$. Therefore, a complete solution could be determined by measuring either $M$ or $D_L$, or some combination of the two.  One way to obtain an independent relation between the lens mass and distance is to measure the microlens parallax, $\\pi_\\e$. There are two potential ways of measuring $\\pi_\\e$.  First, one can measure distortions in the light curve arising from the acceleration of the Earth as it moves along its orbit.  Unfortunately, this is out of the question in this case because the event is so short that these distortions are immeasurably small.  Second, one can measure the effects of terrestrial parallax, which gives rise to differences between the light curves simultaneously observed from two or more observatories separated by a significant fraction of the diameter of the Earth.  Practically, measuring these differences requires a high-magnification event, which would appear to make this event quite promising.  Unfortunately, although we obtained simultaneous observations from two observatories separated by several hundred kilometers during the peak of the event, one of these datasets suffers from large systematic errors and an unknown time zero point, rendering it unusable for this purpose.  The only available alternative for breaking the degeneracy between the lens mass and distance would be to measure the lens flux, either under the glare of the source or, at a later date, to separately resolve it after it has moved away from the line of sight to the source \\citep{alcock01,koz07}. Panels (e) and (f) of Figure~\\ref{fig:prop} show the Bayesian estimates of the lens brightness in $I$-band and $H$-band, respectively. If the lens flux is at least 2\\% of the source flux, then the former kind of measurement could be obtained from a single epoch {\\it Hubble Space Telescope} observation, provided it were carried out in the reasonably near future.  At roughly 99\\% probability, the blended light would be either perfectly aligned with the source (and so associated with the event) or well separated from it. {\\it HST} images can be photometrically aligned to the ground-based images using comparison stars with an accuracy of better than 1\\%.  Hence, photometry of the source+blend would detect the blend, unless it were at least 4 mag fainter than the source.  In principle, the blend could be a companion to either the source or lens.  Various arguments can be used to constrain either of those scenarios.  We do not explore those here, but see \\citet{dong08}. If the lens is not detectable by current epoch {\\it HST} observations (or no {\\it HST} observations are taken), then it will be detectable by ground-based AO $H$-band observations in about 5 years.  This is because the lens-source relative proper motion is measured to be $\\mu_\\rel = 8\\,\\mas\\,\\rm yr^{-1}$, and the diffraction limit at $H$ band on a 10m telescope is roughly 35 mas.  If the lens proves to be extremely faint, then a wider separation (and hence a few years more time baseline) would be required.  In the absence of additional observational constraints, we must rely on a Bayesian analysis to estimate the properties of the host star and planet, which incorporates priors on the distribution of lens masses, distances, and velocities \\citep{dominik06,ob04343}. This is a standard procedure, which we only briefly summarize here. We adopt a \\citet{han95} model for the Galactic bar, a double-exponential disk with a scale height of 325 pc, and a scale length of 3.5 kpc, as well as other Galactic model parameters as described in \\citet{bennett02}. We incorporate constraints from our measurement of the lens angular Einstein radius $\\theta_\\e$ and the event timescale, as well as limits on the microlens parallax and $I$-band magnitude of the lens.  In practice, only the measurements of $\\theta_\\e$ and $t_{\\e}$ provide interesting constraints on these distributions. In addition, we include the small penalty on the close solution, $\\exp(-\\Delta \\chi^2/2)$, where the wide solution is favored by $\\Delta\\chi^2 = 0.2$. For the estimates of the planet semimajor axis, we assume circular orbits and that the orbital phases and cos(inclinations) are randomly distributed.  The resulting probability densities for the physical properties of the host star, as well as selected properties of the planet, are shown in Figure~\\ref{fig:prop}.  The Bayesian analysis suggests a host star of mass $M=0.30_{-0.12}^{+0.19} M_\\odot$ at distance of $D_L=5.8_{-0.8}^{+0.6}~{\\rm kpc}$.  In other words, given the available constraints, the host is most likely an M-dwarf, probably in the foreground Galactic bulge.  Given that the planet/star mass ratio is measured quite precisely, the probability distribution for the planet mass is essentially just a rescaled version of the probability distribution for the host star mass. We find $m_p=0.82_{-0.33}^{+0.52}~M_{\\rm Jup}$.  The close/wide degeneracy is apparent in the probability distribution for the semimajor axis $a$. We estimate $a_{\\rm close}= 0.72_{-0.16}^{+0.38}~{\\rm AU}$ for the close solution, and $a_{\\rm wide}=6.5_{-1.2}^{+3.2}~{\\rm AU}$ for the wide solution.  The equilibrium temperatures for these orbits are $T_{\\rm eq., close}=103_{-26}^{+28}~{\\rm K}$ and $T_{\\rm eq., wide}=34\\pm{9}~{\\rm K}$ for the close and wide solutions, respectively.  Thus our Bayesian analysis suggests that this system is mostly likely a bulge mid-M-dwarf, with a Jovian-mass planetary companion.  The semimajor axis of the planetary companion is poorly constrained primarily because of the close/wide degeneracy, but the implied equilibrium temperatures are cooler than the condensation temperature of water. Specifically we find that $T_{\\rm eq} \\la 173~{\\rm K}$ at $2\\,\\sigma$ level.  Alternatively, if we assume the snow line is given by $a_{\\rm snow}=2.7~{\\rm AU}(M/M_\\odot)$, we find for this system a snow line distance of $\\sim 0.81~{\\rm AU}$, very close to the inferred semimajor axis of the close solution.  Thus this planet is quite likely to be located close to or beyond the snow line of the system.  Although we cannot distinguish between the close and wide solutions for the planet separation, theoretical prejudice in the context of the core-accretion scenario would suggest that a gas-giant planet would be more likely to form just outside the snow line, thus preferring the close solution.  However, we have essentially no observational constraints on the frequency and distribution of Jupiter-mass planets at the separations implied by the wide solution ($\\sim 5.3-9.7~{\\rm AU}$), for such low-mass primaries.  Unfortunately, the prospects for empirically resolving the close/wide degeneracy in the future are poor. The only possible method of doing this would be to measure the radial velocity signature of the planet. Given the faintness of the host star (see \\S\\ref{sec:blend} and Fig.\\ \\ref{fig:prop}), this will likely be impossible with current or near-future technology.  The mere existence of a gas-giant planet orbiting a mid-M-dwarf is largely unexpected in the core-accretion scenario, as formation of such planets is thought to be inhibited in such low-mass primaries \\citep{laughlin04}.  Observationally, however, although the frequency of Jovian companions to M-dwarfs with $a\\la 3~{\\rm AU}$ does appear to be smaller than the corresponding frequency of such companions to FGK dwarfs \\citep{endl06,johnson07,cumming08}, several Jovian-mass companions to M dwarfs are known (see \\citealt{dong08} for a discussion), so this system would not be unprecedented.  Furthermore, it must be kept in mind that the estimates of stellar (and so planet) mass depend on the validity of the priors, and even in this context have considerable uncertainties.  Most of the ambiguities in the interpretation of this event would be removed with a measurement of the host star mass and distance, which could be obtained by combining our measurement of $\\theta_\\e$ with a measurement of the lens light as outlined above.  The Bayesian analysis informs the likelihood of success of such an endeavor. This analysis suggests that, if the host is a main-sequence star, its magnitude will be $I_L=23.9_{-1.0}^{+0.8}$ and $H_L=21.4_{-1.0}^{+0.7}$, which corresponds to $0.6\\%$ and $1.7\\%$ of the source flux, respectively. If initial efforts to detect the lens fail, more aggressive observations would certainly be warranted: microlensing is the most sensitive method for detecting planets around very low-mass stars simply because it is the only method that does not rely on light from the host (or the planet itself) to detect the planet. And given equation~(\\ref{eqn:mpirel}), even an M dwarf at the very bottom of the main sequence $M=0.08\\,M_\\odot$, would lie at $D_L=3.5\\,\\kpc$ and so would be $H\\sim 24$."
    },
    "0809/0809.0992_arXiv.txt": {
        "abstract": "We showed that the part of strings could be detected by optical method is only $20\\%$ from the total available amount of such objects, therefore the gravitational lensing method has to be \"`completed\"' by CMB one. We found the general structure of the CMB anisotropy generated by a cosmic string for simple model of straight string moving with constant velocity. For strings with deficit angle 1-2 arcsec the amplitude of generated anisotropy has to be $15-30 \\mu K$ (the corresponding string linear density is     $G \\mu \\propto 10^{-7}$ and energy is GUT one, $10^{15}$GeV). To use both radio and optical methods the deficit angle has to be from 0.1 arcsec to 5-6 arcsec. If cosmic string can be detected by optical method, the length of corresponding brightness spot of anisotropy has to be no less than 100 degrees. ",
        "introduction": "Cosmic strings are linear topological defects could form in the early Universe during face transitions. These objects have been introduced in theoretical cosmology by \\cite{k}, \\cite{z}, \\cite{v}.  Among all possible types of topological defects, cosmic strings are particularly interesting and their existence finds support in superstring theories, both in compactification models and in theories with extended additional dimensions \\cite{h}, \\cite{vs}, \\cite{dk}. The energy scale for cosmic strings is the scale of GUT or less. The energy scale for fundamental strings is Planck scale. Such heavy strings do not exist in our Universe today, and cannot  have played any role in cosmological evolution except in the first few Planck times. Now we know models with large compact dimensions, in which  the string scale may be much lower, down to the GUT scale or even less. Branes, which now play key role in superstring theory, can collide and generate cosmic strings.  At the current moment there are no observational evidences to prove the existence of cosmic strings, neither to disprove it. The status of cosmic strings as theory related with observational data is based on the possibility to restrict  number, linear density and spatial distribution of such objects. These topological defects can influence on the CMB (Cosmic Microwave Background) anisotropy: according to WMAP 5-years data they are not primary source of primordial density perturbations and not responsible for large scale structure formation, therefore their density has to be hard restricted and be small enough.  From theoretical point of view all characteristics of cosmic strings are resulted mostly by extensive multiparameter computer simulations of string networks. But it is very important to note that these simulations are oriented to search string networks only, not individual strings as we now proposed in our work, because no theory tells us how many strings could exist in the Universe.  The modern methods of cosmic string detection can be divided in three main parts: \\begin{itemize} \\item by optical surveys looking for gravitational lensing events, \\item by radio surveys investigating the structure of CMB anisotropy, \\item looking for model depended and rare features, as gravitational radiation from string loops and from straight string, interactions of strings and black holes, decay of heavy particles emitted by string, interaction of two strings etc. \\end{itemize} The superstring theory permits existence of cosmic strings in wide range of their parameters: with linear density varying from the scale of electroweak unification to GUT scales and with velocity from zero to speed of light. Also there are permissible strings possessed of curvature and loop like structures. All strings share two properties which are model independent: the extremely long cosmological length and a negligibly small cross-section. The investigation of the structure of CMB anisotropy looking for string footprints and search for gravitational lensing events induced by strings are so attractive because these effects should exist for all theoretically possible cosmic strings.  Gravitational lensing events should have particularly and easy distinguishing with high resolution instruments (as HST) features. Cosmic string always generates gravitational lensed pairs of images of background sources, therefore direct detection of such pairs should be direct proof of cosmic string existence. On 11th January 2006 the russian-italian group (\\cite{s0},\\cite{s1}) received and completely analyzed the data from HST on the object CSL-1 which was unique candidate to be produced by a cosmic string, and it was proved that this is not a case of cosmic string lensing.  A base of this experience it was elaborated the technique of cosmic string search through gravitational lensing using ultra-deep optical surveys with high resolution instruments (see, for example, \\cite{s1}). Analizing the results we found that the gravitational lensing method possesses very serious lack of sky covering both in terms of area and depth. The method of searching of gravitational lensing events is optical one. Modern optical surveys cover only $1/6$ part of the whole sky and their red shift depth is not more than z=7. Therefore, approximately only $3\\%$ of possible strings could be detected by this method (if we suppose simple homogeneous distribution of straight strings in the Universe). Taking into account the predicted number of cosmic strings it was estimated that gravitational lensing events are very rare ones: the number of so-called \"chains\" of gravitational lensed pairs is from 0.3 up to 3 in the most optimistical models. This estimation can explain unsuccessful previous experience and shows that additional methods of cosmic string search are urgently required because even utilization of ultra deep galaxy surveys is not enough.  The investigation of CMB anisotropy seems to be the most fruitful and natural in cosmic string search, because of two following reasons. Firstly, this method operates with one of the model independent  properties of cosmic strings, their possibility to form conical space-time. Secondly, the red shift depth of radio surveys is z=1000 and covers the whole sky. Therefore using the WMAP 5-years data on CMB anisotropy we can make search of cosmic strings in all volume inside the surface of last scattering and could detect $100\\%$ strings, not $3\\%$ as in optical surveys, or well-groundedly disprove existence of such objects for wide range of their parameters.  In the present work it was elaborated the method of searching for cosmic strings based on analysis of the CMB anisotropy. Moving straight cosmic string was shown to generate distinctive structures of enhanced and reduced temperature fluctuations on the surface of last scattering. It was analized the conditions under which a cosmic string could be detected by both CMB anisotropy and gravitational lensing. ",
        "conclusions": "Cosmic string could produce distinctive features in distribution of CMB: the temperature would have step-like discontinuities. It was analyzed in details the geometry of such step-like structures, depending on the distance to a cosmic string and its velocity. We analyzed the simplest model of string dynamic with isotropy CMB.  An appeared anisotropy induced by single cosmic string represents a sequence of zones of decreased and increased temperature: the cold spot in front of moving string, then the step-like jump and appearance of hot spot and than a cold spot follows again (see Fig.1 --- Temperature distribution over the sky sphere (Molweide projection). The equator coincides with the horizontal line. The string lies along the axis connecting the poles. The angle $\\phi$ is measured from the string front from right to left. There is a small cold spot in front of the string; a sharp temperature jump occurs at the front followed by an extended hot spot, which again gives way to an indistinct cold spot. The temperature distribution is typical of a moving string and is virtually independent of its parameters, only the spot width and the temperatures at a maximum and local minima change.). The amplitude of these discontinuities and their extension depend on position of the string with respect of an observer, on the string velocity and its direction, and on the string linear density. But the structure always remains the same.  For a cosmic string with deficit angle from 1 to 2 arcsec the amplitude of generated anisotropy would be from 15 to 30 $\\mu$K. It was obtained the structure and amplitude of expected anisotropy induced by cosmic string as function of its parameters. It was shown that in order to detect cosmic strings by both methods of gravitational lensing and CMB anisotropy the range of deficit angle has to be from 0.1 arcsec to 6 arcsec, which covers the huge amount of theoretically predicted strings. If cosmic string can be detected by optical method, the length of corresponding brightness spot of anisotropy has to be no less than 100 degrees.  Therefore our current investigations achieved the level which can certainly indicate the direct way to find cosmic strings in the whole Universe for wide range of string parameters or for the first time prove that such topological defects do not exist. According to superstring approach, cosmic strings are the most preferable objects to be formed in early Universe among other topological defects. There are also deep connections between cosmic strings and brane-world scenario based on superstrings. Long superstrings - macroscopic fundamental strings - may be stable and may appear at the same energy as GUT scale cosmic strings. They could have been produced in the early Universe and then grown to macroscopic size with the expansion of the Universe.  But more work is required. To really detect the string on the sky we have to distinguish the CMB anisotropy induced by scalar density perturbation from the cosmic string one; the problem is a low signal to noise ratio. The amplitude of noise (adiabatic perturbations) is larger then expected signal from cosmic strings. To detect a cosmic string with parameters set by GUT we have to detect anisotropy with signal to noise ratio much less then unity. For detection of the signal which is less then noise we intent to apply two method: wavelet and curvelet analysis and detection by optimal filtering."
    },
    "0809/0809.4263_arXiv.txt": {
        "abstract": "In HST Cycles 11 and 13 we obtained two epochs of ACS/HRC data for fields in the Magellanic Clouds centered on background quasars. We used these data to determine the proper motions of the LMC and SMC to better than 5\\% and 15\\% respectively.  The results had a number of unexpected implications for the Milky Way-LMC-SMC system. The implied three-dimensional velocities were larger than previously believed and close to the escape velocity in a standard $10^{12}$ solar mass Milky Way dark halo, implying that the Clouds may be on their first passage. Also, the relative velocity between the LMC and SMC was larger than expected, leaving open the possibility that the Clouds may not be bound to each other. To further verify and refine our results we requested an additional epoch of data in Cycle 16 which is being executed with WFPC2/PC due to the failure of ACS. We present the results of an ongoing analysis of these WFPC2 data which indicate good consistency with the two-epoch results. ",
        "introduction": "The Large and Small Magellanic Clouds (LMC \\& SMC) at distances of $\\sim 50$ kpc from the Sun, and $\\sim 25$ kpc from the Galactic Plane, provide one of our best probes of the composition and properties of the Galactic dark halo, and have long been upheld as the poster-child for a strongly interacting system, both with each other and with the Milky Way (MW). It has commonly been assumed that the Clouds have made multiple pericentric passages about the MW, and indeed current formation theories for the Magellanic Stream, which may involve tidal or ram-pressure forces, require multiple pericentric passages in order to be viable stripping mechanisms (\\cite[Gardiner \\& Noguchi 1996]{GN96}; hereafter GN96, \\cite[Connors \\etal \\ 2006]{Connors06}, \\cite[Mastropietro \\etal \\ 2005]{Mastropietro05}, \\cite[Yoshizawa \\& Noguchi 2003]{YN03}, \\cite[Lin \\etal \\ 1995]{Lin95}, \\cite[Moore \\& Davis 1994]{MD94}, \\cite[Heller \\& Rohlfs 1994]{HR94}, \\cite[Lin \\& Lynden-Bell 1982]{LL82}, \\cite[Murai \\& Fujimoto 1980]{MF80}).  However, recent high-precision proper motion measurements for the Clouds made by our group with two epochs of ACS High Resolution Camera (HRC) data in Cycles 11 and 13, where we measured the proper motion of LMC stars relative to background quasars, imply that the LMC tangential velocity is $\\sim 370 \\kms$, approximately $100 \\kms$ higher than previously thought (\\cite[Kallivayalil \\etal \\ 2006a]{K1}, \\cite[Kallivayalil \\etal \\ 2006b]{K2}; hereafter K1 \\& K2). The proper motion values of GN96, which have been adopted in all theoretical models of the formation of the Stream thus far, are not consistent with the new HST result. In particular, for the LMC there is a 7-$\\sigma$ difference. The values for the SMC are in more acceptable agreement (3-$\\sigma$ difference). The new measurements also indicate a significant relative velocity between the LMC \\& SMC of $105 \\pm 42$ ${\\rm km \\ s^{-1}}$. This has been assumed to be closer to $\\sim 60 \\kms$ in theoretical models, i.e., approximately the value for the SMC to be on a circular orbit around the LMC.  These results have surprising physical implications which require a reconsideration of the formation mechanism for the Stream. These include the possibility that the LMC may only be on its first passage about the MW (\\cite[Besla \\etal \\ 2007]{Besla07}; hereafter B07). B07 demonstrated this by studying the past orbital paths of the LMC using our observed proper motions and errors in a $\\Lambda$CDM-motivated dark halo with a NFW profile (\\cite[Navarro \\etal \\ 1996]{Navarro96}). This gave rise to starkly different trajectories for the LMC than those produced in a simple isothermal halo potential: even in the `best case' scenario (proper motion in the west direction, $\\mu_W, \\ +4 \\sigma$), the LMC only completes 1 orbit within 10 Gyr and reaches an apogalacticon distance of 550 kpc (see Figure~4 in B07). Subsequently, this has led to a series of papers exploring whether the LMC is indeed \\textit{bound} to the MW (e.g., \\cite[Shattow \\& Loeb 2008]{SL08}, \\cite[Wu \\etal \\ 2008]{Wu08}). Perhaps even more provocative is the possibility that that the LMC \\& SMC may have only recently become a binary system (K2, although see Besla \\etal \\ these proceedings).   Such large motions were unexpected in light of our previous understanding of the MW-LMC-SMC system, and it is therefore crucial that they be verified and further improved through the acquisition of additional data.  Because of the large distances involved, even small differences in the proper motions can give vastly different orbits for the Clouds ($1 \\masyr$ $\\approx 238$ ${\\rm km \\ s^{-1}}$ at the distance of the LMC). We thus applied for and were successful in getting a third epoch of snapshot imaging for our quasar-fields in Cycle 16. These executed with WFPC2 due to the failure of ACS. In this paper we present the preliminary results of an on-going analysis of these WFPC2 data. In \\S~2 we present the WFPC2 data and analysis strategy with special attention to the relative size of the position errors vis-a-vis ACS. We describe the main sources of systematic error in both the ACS and WFPC2 data. In \\S~3 we present results based on simple cuts aimed at minimizing these systematic errors, and discuss the expected overall improvement that this additional epoch affords. Since our analysis is still in the process of being refined, we present the results only in comparative fashion, both to our ACS-only two-epoch results (K1 \\& K2), and to those of \\cite[Piatek \\etal \\ 2008]{Piatek08} (hereafter P08), who recently re-analyzed our ACS data using their own methods to obtain results that are consistent with ours. A brief summary and future prospects are presented in \\S~4. ",
        "conclusions": "We have presented an ongoing analysis of the proper motions of the Magellanic Clouds using a third epoch of WFPC2 data centered on background quasars. At present the RMS error in the position of the quasar is roughly 3 times as large for WFPC2 as for ACS ($\\sim 0.021$ pixels versus 0.008 pixels). However, with an improved method to deal with CTE and magnitude-related effects, and with the increase in time-baseline from 2 to 5 years, we expect final error bars for the proper motions that are smaller by a factor of $\\sim 2$ from the two-epoch analysis: the two-epoch error bars in the north and east directions are $(0.05, 0.08 \\masyr)$, while we expect three-epoch error bars of $(0.04,0.04 \\masyr)$. This will have several benefits, not the least of which is an important consistency check on our earlier results.  This will allow us to better understand the orbit of the Clouds around each other. Combined with our understanding of the properties of the Magellanic Stream, this will allow us to better constrain the MW dark halo potential, as well as weigh into whether the LMC is indeed bound to the MW. Finally, we have a fourth epoch of observations scheduled in Cycle 17 with ACS and WFC3. The expected improvement in accuracy will continue to provide fundamental new insights into the unique and enigmatic Milky Way-LMC-SMC system."
    },
    "0809/0809.1500_arXiv.txt": {
        "abstract": "{Molecular hydrogen (H$_2$) is the most abundant molecule in the circumstellar (CS) environments of young stars, and is a key element in giant planet formation. The measurement of the H$_2$ content provides the most direct probe of the total amount of CS gas, especially in the inner warm planet-forming regions of the disks. } {Most Herbig Be stars (HBes) are distant from the Sun and their nature and evolution are still debated. We therefore conducted mid-infrared observations of H$_2$ as a tracer of warm gas around HBes known to have gas-rich CS environments.} {We report a search for the H$_2$ S(1) emission line at 17.0348 $\\mu$m in the CS environments of 5 HBes with the high resolution spectroscopic mode of \\visir\\ ({\\it ESO VLT Imager and Spectrometer for the mid-InfraRed}). } {No source shows evidence for H$_2$ emission at 17.0348 $\\mu$m. Stringent 3$\\sigma$ upper limits on the integrated line fluxes are derived. Depending on the adopted temperature, limits on column densities and masses of warm gas are also estimated. These non-detections constrain the amount of warm ($>150~K$) gas in the immediate CS environments of our target stars to be less than $\\sim 1-10 M_{\\rm Jup}$.} {} ",
        "introduction": "Molecular hydrogen (H$_2$) is the most abundant molecule in the environment of young stars. It remains optically thin up to high column densities ($\\sim$10$^{23}$ cm$^{\u22122}$), and does not freeze onto dust grain surfaces. It furthermore self-shields efficiently against photodissociation by far-ultraviolet (FUV) photons. Therefore, H$_2$ is the only molecule that can directly constrain the mass reservoir of molecular gas in the circumstellar (CS) environment of pre-main sequence stars. On the other hand, H$_2$ is one of the most challenging molecules to detect. Electronic transitions occur in the UV to which the Earth's atmosphere is opaque, and rotational and rovibrational transitions at infrared (IR) wavelengths are faint because of their quadrupolar origin. FUV spectroscopic observations of H$_2$ absorption lines have provided evidence for the presence of cold and warm excited H$_2$ around numerous Herbig Ae/Be stars (HAeBes) but have not allowed us to determine the spatial distribution of the observed gas \\citep[e.g.][]{klr08a}. Detections of H$_2$ line emission from disks around young stars with {\\it ISO} \\citep{Thi01} were contradicted by ground-based observations \\citep{Richter02, Sako05}, which showed that the observed emission could be due to the surrounding cloud material. \\cite{Carmona08} modeled the mid-IR H$_2$ lines originating in a gas-rich disk, seen face-on, surrounding a Herbig Ae (HAe) star at 140 pc from the Sun. By assuming that the gas and dust were well-mixed in the disk, a gas-to-dust ratio of about 100, and that $T_{gas}=T_{dust}$, those authors demonstrated that mid-IR H$_2$ lines could not be detected with the existing instruments. Indeed, they were unable to detect any H$_2$ mid-IR emission line in their sample of 6 HAe stars. At the present time, observations of H$_2$ mid-IR emission lines have been reported in 8 of 76 T\\,Tauri stars observed with {\\it Spitzer} \\citep{Lahuis07}. Although mid-IR windows are strongly affected by sky and instrument background emission, the advent of high spectral and spatial resolution spectrographs allows us to study the H$_2$ emission from the ground, as demonstrated in the cases of two HAeBes, \\object{HD~97048} \\citep{klr07a} and AB Aur \\citep{Bitner07}, for which particular conditions are required in their CS disks. The analyses of these data usually assumes that the H$_2$ excitation is in local thermodynamic equilibrium (LTE) and can thus be characterized by a single excitation temperature, which should be close to the gas temperature because of the low critical densities. The AB Aur observations infer an H$_2$ gas temperature that is significantly higher than the dust temperature. \\begin{table*} \\begin{center} \\caption{Astrophysical parameters of the sample stars (columns 2 to 6).  \\vrad\\ is the radial velocity of the star in the heliocentric rest frame. Columns 7 to 12 summarize the observations. The airmass and seeing intervals are given from the beginning to the end of the observations. \\vspace{-0.4cm}} \\begin{tabular}{lcccccccccccccccccc} \\hline \\hline Star      & Sp.    & $T_{eff}$     & \\Av         &  \\vrad        & $d$            &  $t_{exp}$ & Airmass & Optical      & Standard & Airmass & Optical\\\\ & Type     & (K)          & (mag)        & (\\kms)       & (pc)            &  (s)      &         & Seeing   &  Star     &        & Seeing \\\\ &     &           &         &       &             &       &         & ('')  &       &        &  ('') \\\\ \\hline HD~98922  & B9       & 10470$^{(1)}$ & 0.34$^{(1)}$ & -15$^{(2)}$   & $>$540$^{(1)}$      & 2700     & 1.14-1.15  & 0.76-1.41  & HD~25025  & 1.024-1.025 & 0.69-0.71 \\\\ HD~250550 & B7       & 12800$^{(3)}$ & 0.57$^{(3)}$ & +31$^{(4)}$   & 606$\\pm$367$^{(5)}$ & 3600     & 1.32-1.50  & 0.83-1.50  & HD~93813  & 1.018-1.023 & 0.91-1.05  \\\\ HD~259431 & B5       & 15900$^{(3)}$ & 0.88$^{(3)}$ & +43$^{(4)}$   & 290$\\pm$84$^{(5)}$  & 2250     & 1.23-1.32  & 0.75-1.28  & HD~39425  & 1.12-1.14   & 1.51-1.62 \\\\ HD~76534  & B2       & 20000$^{(6)}$ & 0.80$^{(7)}$ & +17$^{(4)}$   & $>$160$^{(1)}$      &  3600     & 1.05-1.14  & 0.85-1.25  & HD~93813  & 1.014-1.017 & 0.92-0.96  \\\\ HD~45677  & B2/B1    & 21400$^{(1)}$ & 0.87$^{(1)}$ & +21.6$^{(8)}$ & 500$^{(9)}$         &  3450     & 1.02-1.08  & 0.69-1.40  & HD~25025  & 1.024-1.025 & 0.69-0.71   \\\\ \\hline \\end{tabular} \\begin{list}{}{} References: $^{(1)}$ \\cite{VdA98b}; $^{(2)}$ \\cite{Acke05}; $^{(3)}$ \\cite{JCB03b}; $^{(4)}$ \\cite{FINK_JAN84}; $^{(5)}$\\cite{Brittain07}; $^{(6)}$ \\cite{klr04b}; $^{(7)}$ \\cite{valenti00}; $^{(8)}$ \\cite{Evans67}; $^{(9)}$ \\cite{Zorec98} as quoted in \\cite{Cidale01}.  \\vspace{-0.8cm} \\\\ \\end{list} \\label{param} \\end{center} \\end{table*}  In contrast to HAes, the more distant Herbig Be stars (HBes) earlier than B9 type have been poorly studied. As a consequence, the nature and evolution of the CS material surrounding HBes is still a subject of controversy. The Spectral Energy Distributions (SEDs) are very different from one HBe to another \\citep[e.g.][]{HILL92}. Some have SEDs that are comparable to those observed for HAes, and have been classified as group I or group II stars haboring disks according to the classification of \\citet[][]{Meeus01}. However, numerous HBes SEDs with little IR excess have been modeled successfully by free-free emission originating in a gaseous CS envelope, and compared with rapidly rotating classical Be stars. In addition, the mid-IR emission observed in HBes is generally not confined to optically thick disks but originates in more complex environments such as remnant envelopes or halos \\citep{LEINERT01, POLOMSKI02, Vinkovic06}. This implies that there are structural differences between HAes and the more luminous HBes \\citep{NATTA00, LEINERT01}. These results are fully consistent with the more rapid evolution of the more massive HAeBes, which should still be partially embedded in their natal cloud. On the other hand, near-IR interferometric data of HBes have been interpreted using flared disks models \\citep[e.g.][]{Kraus07}.  We selected five HBe stars whose CS environments are rich in gas as observed by the \\fuse\\ satellite \\citep{klr08a}.  The selected stars have IR excesses that are assumed to represent the presence of CS dusty disks \\citep[e.g.][]{HILL92, POLOMSKI02}, but the disk scenario has not yet been clearly established. The aim of our study is to detect the H$_2$ pure rotational emission line at 17.0348 $\\mu$m in the vicinity of these stars with the high resolution mode of the VLT/\\visir\\ spectrograph \\citep{Lagage04} and constrain the warm CS gaseous component. ",
        "conclusions": "\\label{concl}  We observed five HBe stars with the high resolution spectroscopic mode of \\visir\\ to search for H$_2$ pure rotational S(1) emission at 17.0348 $\\mu$m. As probed by previous \\fuse\\ observations, the CS environments of these stars provide evidence of large reservoirs ($N({\\rm H}_2) \\sim 10^{21}$~cm$^2$) of cold H$_2$ ($T \\sim 100$~K), and also evidence for the presence of warm/hot excited H$_2$ ($T \\geq 500$~K). The H$_2$, observed by \\fuse, probably corresponds to the remnant of the molecular cloud in which the stars were formed and the observed H$_2$ FUV lines do not originate in a CS disk. \\cite{klr08a} modeled the excitation conditions of H$_2$ and the \\fuse\\ spectra of HBes stars using the {\\it Meudon PDR Code} and showed that the observed gas was located in relatively diffuse regions ($10-3000$ cm$^{-3}$) at significant distances from the central stars ($0.03-1.5$ pc).  None of the four targets with spectra of sufficiently high S/N showed any evidence for H$_2$ emission. From the 3$\\sigma$ upper limits to the emission line flux, we calculated upper limits on the column density and mass of H$_2$ for each star. We found that the column densities should be lower than $\\sim$10$^{23}$ cm$^{-2}$ at 150~K, and lower than $\\sim$10$^{21}$ cm$^{-2}$ at 1000~K, which does not contradict the column densities derived from the \\fuse\\ observations. Upper limits to the masses of warm gas have been estimated to be in the range from $\\sim$10$^{-2}$ to $\\sim$6 $M_{\\rm Jup}$ (1 $M_{\\rm Jup}$ $\\sim$ 10$^{-3}$ $M_{\\odot}$), assuming LTE excitation, and depending on the adopted temperature.  It should be pointed out that mid-IR H$_2$ lines only probe warm gas located in the surface layers of the observed media because of the opacity of the interior layers when dust is present. The surface layers of any potential disks surrounding our target stars do not therefore contain sufficient warm gas to enable it to be detected in emission at mid-IR wavelenghs. Following the calculations of \\cite{Carmona08}, the H$_2$ mid-IR lines produced by a gas-rich disk with $T_{gas} = T_{dust}$ should not be observable with existing instruments. As shown by \\cite{Bitner07} and \\cite{klr07a}, the warm H$_2$ can be detected in the CS disk of a Herbig star, but to explain the detection, particular conditions have to be assumed for the gas and dust, such as $T_{gas} > T_{dust}$, which may be created by gas heating by X-rays or UV photons. Our non-detections therefore imply either that the gas and dust in any potential disks are well mixed and have almost equal temperatures, or that these disks are simply not present at all."
    },
    "0809/0809.4922_arXiv.txt": {
        "abstract": "{To investigate the joint evolution of active galactic nuclei and star formation in the Universe.} {In the 1.4\\,GHz survey with the Australia Telescope Compact Array of the Chandra Deep Field South and the European Large Area ISO Survey\\,-\\,S1 we have identified a class of objects which are strong in the radio but have no detectable infrared and optical counterparts. This class has been called Infrared-Faint Radio Sources, or IFRS. 53 sources out of 2002 have  been classified as IFRS. It is not known what these objects are.} {To address the many possible explanations as to what the nature of these objects is we have observed four sources with the Australian Long Baseline Array.} {We have detected and imaged one of the four sources observed. Assuming that the source is at a high redshift, we find its properties in agreement with properties of Compact Steep Spectrum sources. However, due to the lack of optical and infrared data the constraints are not particularly strong.} {} ",
        "introduction": "Infrared-Faint Radio Sources (IFRS) were recently discovered as a class by \\cite{Norris2006a}, and may be related to the Optically Invisible Radio Sources (OIRS) identified by \\cite{Higdon2005}. IFRS are radio sources which have no counterparts in infrared images from the {\\it Spitzer} Wide-Area Extragalactic Survey (SWIRE) between 3.6\\,\\um\\ and 24\\,\\um, and are discovered in arcsec-scale radio observations. They are unexpected because it was thought that any galaxy which is detected in radio observations should be detected in the infrared with relatively short integrations. Assuming the SED of known classes of galaxy, a 5\\,mJy radio source in the local Universe should produce a detectable Spitzer source, regardless of whether it is generated by star formation or AGN (active galactic nuclei). Similarly a normal L$_*$ galaxy at $z<1$, whether spiral or elliptical, should be visible in our Spitzer and I-band data.  \\cite{Norris2006a} and \\cite{Middelberg2008a} together identified 53 such sources out of 2002 (2.7\\,\\%) detected in the ATLAS survey, co-located with the SWIRE survey. Most of these sources have flux densities of only a few hundred microjansky, but some are strong and have flux densities of more than 20\\,mJy. Stacking 3.6\\,\\um\\ {\\it Spitzer} images at the positions of 22 IFRS, \\cite{Norris2006a} were unable to make a detection in the averaged image and so demonstrated that IFRS are well below the detection threshold of the SWIRE survey.  The nature of IFRS and the reason for their faintness at infrared wavelengths is unclear. Possible explanations are that (i) these sources are extremely redshifted Active Galactic Nuclei (AGN); (ii) they are dust-rich, extremely obscured galaxies which makes them invisible in the infrared; (iii) they are lobes of nearby, unidentified radio galaxies; or (iv) they are an unknown type of galactic or extragalactic object. Because IFRS have so far only been detected at radio wavelengths it is not possible to measure their redshifts, as spectroscopy requires sub-arcsecond positional accuracy, which the radio observations cannot provide. Also, the comparatively low resolution of the radio images makes it difficult to select the correct optical counterpart, because the corresponding optical observations are deep, and hence confusion-limited.  A promising route to find out more about IFRS is radio observations with Very Long Baseline Interferometry (VLBI). VLBI observations are sensitive only to very compact structures with brightness temperatures of the order of $10^6$\\,K or more, which are unambiguous signposts of AGN activity. They also yield, when astrometric calibrators are used, positions accurate to milliarcseconds, and milliarcsecond-scale morphologies which can be interpreted in terms of the emission mechanism at work.  \\cite{Norris2007b} have observed two IFRS with VLBI and discovered one. Unfortunately, their $(u,v)$ coverage was too poor to make a reliable image of the detected source. They concluded that the VLBI observations were consistent with a radio-loud AGN at high redshift, or with a lower-power AGN at lower redshift in an abnormally obscured galaxy.  Here we present VLBI observations of four IFRS discovered in the ATLAS/ELAIS field (\\citealt{Middelberg2008a}). We selected S427 and S509 because they were the strongest IFRS in the ATLAS/ELAIS field, and S775 because it is very extended on arcsecond scales, showing structures reminiscent of lobes and jets frequently seen in AGN. After the observations it was discovered that the weak IFRS S433 was located only 24.6\\,arcsec north-east of S427 (Fig.~\\ref{fig:S427+ir}), and was well within the field of view of the VLBI array. The details of the sources observed are listed in Table~\\ref{tab:sources}.  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{0454fig1.eps} \\caption{Contour plot of the ATCA 20\\,cm image of S427 superimposed on the 3.6\\um\\ {\\it Spitzer} image made as part of the SWIRE survey (\\citealt{Lonsdale2003}). Contours are drawn at 0.1\\,mJy$\\times$(1, 2, 4, ...) and the restoring beam was 10.3$\\times$7.2\\,arcsec. The nearest infrared source, located towards the south-east, is more than 6\\,arcsec away, making it very unlikely to be the infrared counterpart. Source S433, visible as a two-contour object 24.6\\,arcsec north-east of S427, also was classified as an IFRS. It has a flux density of 245\\,\\uJy.} \\label{fig:S427+ir} \\end{figure} ",
        "conclusions": "We present the first VLBI image of an Infrared-Faint Radio Source and report the non-detection of three more IFRS discovered in the ATCA 1.4\\,GHz observations of the ATLAS/ELAIS field. The main result of our observations is that S427 harbours an AGN, and that it is not simply a radio lobe of an unidentified radio galaxy. The size, spectrum, and radio and IR luminosity of the detected IFRS S427 are consistent with those of a high-redshift Compact Steep-Spectrum source, and are inconsistent with a standard L$_*$ galaxy at z$<$1, or a FR\\,I/II galaxy at any redshift.  Together with the two IFRS observed by Norris et al. (2007) the number of IFRS observed with VLBI is now 6, and the number of detections is 2, showing that at least some fraction of IFRS are associated with AGN. We plan further VLBI studies to determine the nature of these enigmatic objects."
    },
    "0809/0809.3733_arXiv.txt": {
        "abstract": "This paper presents detailed analysis of large-scale peculiar motions derived from a sample of $\\sim 700$ X-ray clusters and cosmic microwave background (CMB) data obtained with WMAP. We use the kinematic Sunyaev-Zeldovich (KSZ) effect combining it into a cumulative statistic which preserves the bulk motion component with the noise integrated down. Such statistic is the dipole of CMB temperature fluctuations evaluated over the pixels of the cluster catalog (Kashlinsky \\& Atrio-Barandela 2000). To remove the cosmological CMB fluctuations the maps are filtered with a Wiener-type filter in each of the eight WMAP channels (Q, V, W) which have negligible foreground component. Our findings are as follows: The thermal SZ (TSZ) component of the clusters is described well by the Navarro-Frenk-White profile expected if the hot gas traces the dark matter in the cluster potential wells. Such gas has X-ray temperature decreasing rapidly towards the cluster outskirts, which we demonstrate results in the decrease of the TSZ component as the aperture is increased to encompass the cluster outskirts. We then detect a statistically significant dipole in the CMB pixels at cluster positions. Arising exclusively at the cluster pixels this dipole cannot originate from the foreground or instrument noise emissions and must be produced by the CMB photons which interacted with the hot intracluster gas via the SZ effect. The dipole remains as the monopole component, due to the TSZ effect, vanishes within the small statistical noise out to the maximal aperture where we still detect the TSZ component. We demonstrate with simulations that the mask and cross-talk effects are small for our catalog and contribute negligibly to the measurements. The measured dipole thus arises from the KSZ effect produced by the coherent large scale bulk flow motion. The cosmological implications of the measurements are discussed by us in Kashlinsky et al (2008). ",
        "introduction": "In the popular gravitational instability picture for growth of the large scale structure in the Universe, peculiar velocities on large cosmological scales probe directly the peculiar gravitational potential and provide important information on the underlying mass distribution in the Universe [e.g. see review by Kashlinsky \\& Jones 1991]. Previous attempts to measure the peculiar flows in the local Universe mostly used empirically established (but not well understood theoretically) galaxy distance indicators. While very important, such methods are subject to many systematic uncertainties [e.g. see reviews by \\cite{strauss-willick,willick}] and lead to widely different results.  Early measurements by \\cite{rubin-ford} indicated large peculiar flows of $\\sim$700 km/sec. A major advance was made using the ``Fundamental Plane\" (FP) relation for elliptical galaxies \\cite{7s-di,djorgovski} with the implication that elliptical galaxies within $\\sim 60h^{-1}$Mpc were streaming at $\\sim 600$ km/sec with respect to the rest frame defined by the cosmic microwave background (CMB) \\cite{7s-motion}. Mathewson et al (1992) used the Tully-Fisher (TF) relation for a large sample of spiral galaxies suggesting that the flow of amplitude 600 km/sec does not converge until scales much larger than $\\sim 60 h^{-1}$ Mpc. This finding was in agreement with a later analysis by \\cite{willick99}. Employing brightest cluster galaxies as distance indicators \\cite{lauer-postman} measured a bulk flow of $\\sim$700 km/sec for a sample 119 rich clusters of galaxies on scale of $\\sim$150$h^{-1}$Mpc suggesting significantly larger amount of power than expected in the concordance $\\Lambda$CDM model. However, a re-analysis of these data \\cite{hudson-ebeling} taking into account the correlation between the luminosities of brightest-cluster galaxies and that of their host cluster found a bulk flow in a greatly different direction and at a smaller amplitude. Using the FP relation for early type galaxies in 56 clusters \\cite{hudson} find a bulk flow of a similarly large amplitude of $\\sim 630$ km/sec to \\cite{lauer-postman} on a comparable scale, but in a different direction. On the other hand, a sample of 24 SNIa shows no evidence of significant bulk flows out to $\\sim 100 h^{-1}$ Mpc \\cite{riess} and similar conclusion is reached with the TF based survey of spiral galaxies by \\cite{courteau}. The directions associated with each bulk-flow measurement are equally discrepant.  The current situation with measurements based on the various distance indicators is confusing and it is important to find alternative ways to measure the large scale peculiar flows. One way to achieve this is via the kinematic component of the Sunyaev Zeldovich (SZ) effect produced on the CMB photons from the hot X-ray emitting gas in clusters of galaxies ([see review by \\cite{birkinshaw}]. The kinematic SZ (KSZ) effect is independent of redshift and measures the line-of-sight peculiar velocity of a cluster in its own frame of reference. For each individual cluster the KSZ temperature distortion will be small and difficult to measure. Attempts at measuring the peculiar velocities of individual clusters from the KSZ effect using the current generation of instruments lead to uncertainties of $\\gsim 1000$ km/sec per cluster [see review by \\cite{carlstrom}]. On the other hand, as proposed by \\cite{kab} (hereafter KA-B) for many clusters moving at a coherent bulk flow one can construct a measurable quantity using data on CMB temperature anisotropies which will be dominated by the bulk flow KSZ component, whereas the various other contributions will integrate down. This quantity, {\\it the dipole of the cumulative CMB temperature field evaluated at cluster positions}, is used in this investigation on the 3-year WMAP data in conjunction with a large sample of X-ray clusters of galaxies to set the strongest to-date limits on bulk flows out to scales $\\sim 300 h^{-1}$Mpc.  In the accompanying Letter (Kashlinsky et al 2008) we summed the results and their cosmological implications. These are obtained using the KA-B method applied to 3-year WMAP CMB data and the largest all-sky X-ray cluster catalog to date. This paper provides the details relevant for the measurement and is structured as follows: Sec \\ref{steps} summarizes the KA-B method and the steps leading to the measurement. Sec. \\ref{catalog} describes the cluster X-ray catalog used in this study and Sec. \\ref{cmb} outlines the CMB data processing. Sec. \\ref{errors} discusses the methods to estimate the errors followed by Sec \\ref{results} with the results on the dipole measurement. Sec. \\ref{tsz} shows why the measured dipole arises from the KSZ component due to the cluster motion and Sec. \\ref{calibration} dicusses the translation of the measured dipole in $\\mu$K into velocity in km/sec and its uncertainty. Future prospects foreseeable at this time to improve this measurement are discussed in Sec. \\ref{future}. We summarize our results in Sec. \\ref{summary}. ",
        "conclusions": "\\label{summary}  We now summarize the main conclusions from this study:  $\\bullet$ Our measurements indicate the existence of the residual CMB dipole evaluated over the CMB pixels associated with the hot SZ producing gas in clusters of galaxies. The dipole is measured at high-signifance level ($\\sim 8\\sigma$ in the outer bins) and persists out the limit of our cluster catalog $z_{\\rm median}\\simeq 0.1$. Its direction is not far off the direction of the \"global CMB dipole\" measured from the entire unprocessed maps.  $\\bullet$ We show with detailed simulation that the CMB mask and/or cluster sample discreteness induced cross-talk effects are negligible and cannot mimic the measured dipole.  $\\bullet$ The dipole originates exclusively at the cluster pixels and, hence, cannot be produced by foregrounds or instrument noise. It must originate from the CMB photons that have passed through the hot gas in the catalog clusters.  $\\bullet$ We prove that the signal arises from the hot SZ producing cluster gas because we demonstrate that in the unfiltered CMB maps there remains statistically significant temperature {\\it decrement} as expected from the TSZ effect. Its profile is consistent with the NFW profile out the largest aperture where we still detect hot gas ($\\sim 30^\\prime$). At larger radii the dipole begins to decrease as expected.  $\\bullet$ In the filtered maps, designed to reduce the cosmological CMB fluctuations, the dipole is isolated simultaneously as the monopole component vanishes. This proves that its origin lies in the KSZ component. The monopole vanishes (within the noise) because for the NFW profile the gas in hydrostatic equilibrium must have a strong decrease in the X-ray temperature in the outer parts. This decrease is consistent with the available direct X-ray measurements, but more importantly is demonstrated empirically in AKKE.  $\\bullet$ With the current cluster catalog we determine that the amplitude of the dipole corresponds to bulk flow of 600-1000 km/sec. This conversion factor, $C_{\\rm 1,100}$, may however have some systematic offset related to our current cluster modelling. However, this possible uncertainty only affect the amplitude of the motion, not its coherence scale or existence.  $\\bullet$ The cosmological implications are discussed in Kashlinsky et al (2008). We show there that the concordance $\\Lambda$CDM model cannot account for this motion at many standard deviations. Instead, it is possible that this motion extends all the way to the current cosmological horizon and may originate from the tilt across the observable Universe from far away pre-inflationary inhomogeneities \\cite{ktf,turner}.  This work is supported by the NASA ADP grant NNG04G089G in the USA (PI - A. Kashlinsky) and by the Ministerio de Educaci\\'on y Ciencia and the ''Junta de Castilla y Le\\'on'' in Spain (FIS2006-05319, PR2005-0359 and SA010C05, PI - F. Atrio-Barandela). We thank Gary Hinshaw for useful information regarding the WMAP data specifics. FAB thanks the University of Pennsylvania for its hospitality when part of this work was carried out. We thank Carlos Hernandez-Monteagudo for spotting a technical correction in the SZ energy distribution, eq. 4 and Fig. 7.  {\\it NOTE ADDED IN PROOF}: Our results have received recently additional support from an independent study by Watkins, Feldman and Hudson (arXiv:0809.4041; 2009, MNRAS, 392, 743) . The Watkins et al. study compiled all major peculiar velocity surveys to date to determine bulk flows within a 100 $h^{-1}$Mpc sphere. Although the scales involved are much smaller than, and the method completely different from ours, Watson et al find that the galaxies within a $\\sim 50-100 h^{-1}$Mpc sphere are moving at a significant velocity in the same direction as found in our work. The amplitude of their motion, at $\\sim400-500$ km/sec, appears somewhat smaller, but still overlaps within $<$ 2 standard deviations with our velocity assuming the calibration above. We anticipate that recalibrating the cluster sample as described in Sec. 8 will further decrease the difference between the measured velocity amplitudes."
    },
    "0809/0809.2609_arXiv.txt": {
        "abstract": "We used the {\\em Advanced Camera for Surveys} on board the {\\em Hubble Space Telescope} to obtain high resolution $i$-band images of the centers of 23 single galaxies, which were selected because they have SDSS velocity dispersions larger than $350$~km~s$^{-1}$.  The surface brightness profiles of the most luminous of these objects ($M_i<-24$) have well-resolved `cores' on scales of 150-1000~pc, and share similar properties to BCGs.  The total luminosity of the galaxy is a better predictor of the core size than is the velocity dispersion.  The correlations of luminosity and velocity dispersion with core size agree with those seen in previous studies of galaxy cores. Because of high velocity dispersions, our sample of galaxies can be expected to harbor the most massive black holes, and thus have large cores with large amounts of mass ejection. The mass-deficits inferred from core-Sersic fits to the surface-brightness profiles are approximately double the black-hole masses inferred from the $M_\\bullet-\\sigma$ relation and the same as those inferred from the $M_\\bullet-L$ relation.  The less luminous galaxies ($M_i>-23$) tend to have steeper `power-law' inner profiles, higher-ellipticity, diskier isophotes, and bulge-to-total ratios of order 0.5 -- all of which suggest that they are `fast-rotators' and rotational motions could have contaminated the velocity dispersion estimate.  There are obvious dust features within about 300~pc of the center in about 35\\% of the sample, predominantly in power-law rather than core galaxies.  \\\\ ",
        "introduction": "This is the third in a series of papers about the galaxies with the highest velocity dispersions ($\\sigma>350$~km s$^{-1}$) in the low redshift Universe ($z<0.3$).  Using imaging and spectroscopy from the Sloan Digital Sky Survey (SDSS) (Abazajian et al. 2003), Bernardi et al. (2006, hereafter Paper I) constructed a sample of 70 objects which were likely to be single galaxies with high ($>350$ km/s) velocity dispersion. However, in this sample there still remained the strong possibility that some of these measured velocity dispersions were contaminated by the presence of another galaxy along the line of sight. In a companion to this work, Bernardi et al. (2008, hereafter Paper II) resolve this problem with observations which take advantage of the superior angular resolution of the Hubble Space Telescope (HST) to identify superpositions, and exclude them from the sample of true high velocity dispersion galaxies.  Of the original sample of 70 objects, HST observed 43 (as this was a SnapShot program, not all 70 targets were observed), and 23 of these appear to be truly single galaxies.  Paper II shows that these appear to be of two types: luminous, round galaxies which share many properties with BCGs (e.g., Laine et al. 2003, Bernardi et al. 2007), and fainter, higher ellipticity, disky galaxies which point to having over-estimated measured velocity dispersions because of rotational velocities.  One of the goals of this paper is to see if a detailed photometric analysis of the HST data supports those conclusions.  However, our analysis of the HST photometry also allows us to place our sample in the context of other HST-based studies of early-type galaxies.  In particular, the centers of early-type galaxies have been studied extensively with HST by the Nuker Team (e.g., Lauer et al. 1995, Faber et al. 1997, Laine et al. 2003, Lauer et al. 2005, Lauer et al. 2007) using WFPC1 and WFPC2, and by the Virgo Cluster Survey (VCS) (e.g., Ferrarese et al. 2006a, Ferrarese et al. 2006b, C{\\^o}t{\\'e} et al. 2006) using ACS WFC. Both groups have found that the centers of early-type galaxies can contain nuclei (corresponding to a point-like increase in the light profile), follow a single power-law or Sersic (continuously changing power-law) light distribution, or have a break radius inside of which the surface brightness follows a shallow or flat power law. It has been established that bright early-types have round isophotes and shallow central slopes while fainter early-types have elongated, disky isophotes and steep central slopes (Ferrarese et al. 1994, Lauer et al. 1995, Faber et al. 1997, Ravindranath et al. 2002, Ferrarese et al. 2006a). There is active debate regarding whether the inner light profile slopes of early-type galaxies form a bimodal distribution of ``cores'' and ``power-law galaxies'' (Lauer et al. 2007, Rest et al. 2001), or a continuous distribution (Graham et al. 2003, Ferrarese et al. 2006a). But there is general agreement that the connection between the properties of the centers of core galaxies and their overall structure can be explained by the presence of a central supermassive black hole binary system. Since there is a  correlation between central black hole mass and velocity dispersion (Ferrarese \\& Merritt 2000, Gebhardt et al. 2000, Tremaine et al. 2002), our sample of galaxies can be expected to harbor the most massive black holes, and thus have large cores with large amounts of mass ejection.  This paper is organized as follows.  Section~\\ref{s_data} describes the observations and pre-processing steps we undertook to prepare the images for model-fitting and profile analysis.  Section~\\ref{s_model} presents global properties of these galaxies from fitting parametric models to the surface brightness profiles.  Central structure and core properties are studied in Sections~\\ref{s_center} and~\\ref{s_binary_bh}, and photometric evidence for rotation in some of these galaxies is presented in Section~\\ref{s_rotation}.  Section~\\ref{s_dust} studies the difference images to see if they show evidence for dust, and a final section summarizes our results. ",
        "conclusions": "We presented our methodology for obtaining parameters which describe the global and inner structure of high velocity dispersion galaxies ($\\sigma > 350$~km~s$^{-1}$). We performed model-fits to obtain effective radii, luminosities, eliipticities, bulge fractions, and Sersic indices which describe the global structure of the galaxies. We showed that the deVaucouleurs quantities obtained from HST imaging agreed with the SDSS values from which we drew the scaling relations presented in Paper II (Figure~\\ref{f_SDSS_HST}). We studied the galaxies' isophotal deviations from true ellipticity using the $a4$ parameter. This parameter, in combination with the bulge-disk decompositions, ellipticity, and inner profile shape, established the fainter galaxies in our sample as almost certainly fast-rotating ellipticals (Section~\\ref{s_rotation}), meaning that our measured velocity dispersions are likely to be over-estimates.  We also studied the central profile shapes of the galaxies, fitting core-Sersic and Nuker models to the surface brightness profiles to estimate core radii (Section~\\ref{s_center}).  The location of the most luminous of our core-galaxies on the luminosity$-$core-size plane (Figure~\\ref{f_core_radius}) coincides with that of BCGs of Laine et al. (2003). This adds further evidence to the argument presented in Paper II, that the brighter galaxies in our sample are BCGs --- although they are amongst the densest BCGs for their luminosities.  These objects are interesting for the following reason:  when half-light radius is plotted vs luminosity, BCGs are known to have larger radii than the bulk of the early-type galaxy population (Lauer et al. 2007, Bernardi et al. 2007), and this has been used to argue for merger-dominated formation histories.  That they also have core-profiles, and that the cores have {\\em smaller} radii (than expected by extrapolating the core-$L$ relation of the bulk of the population to large $L$) suggests that these objects indeed had formation histories with substantial merger/accretion activity.  The fact that they do not have power-law profiles suggests that in situ star formation is unlikely to be the primary reason for their large central densities.  Either these objects had dense cores around which more stars have since built up, or mergers and accretion events have driven the dense stellar cores of smaller galaxies into the center, erasing a cusp (if there was one).  Such objects are expected to host massive black holes which may have consumed or ejected substantial amounts of mass from the center.  The mass deficit associated with extrapolating the outer Sersic fit inwards, and subtracting the light in the actual best-fit core-Sersic were found to be on average, twice as much as $M_\\bullet$ estimated from the $M_\\bullet-\\sigma$ relation, and equal to $M_\\bullet$ estimated from the $M_\\bullet-L$ relation (Figure~\\ref{f_deficit}) in agreement with previous work (e.g., Graham 2004, Ferrarese et al. 2006a).  Cimatti et al. (2008) studied a sample of superdense z$\\sim$1.5 passive galaxies. They found, when comparing with our sample of high-velocity dispersion galaxies (their figure 19), that the lower-luminosity galaxies in our sample populate a similar locus in the size, mass, surface density plane as their superdense z$\\sim$1.5 passive galaxies. Although rotation may somewhat complicate the picture, it is possible that our low-redshift high-density galaxies are the rare examples of the high-redshift superdense galaxies which have not undergone any dry merging. This scenario is supported by the fact that the low luminosity galaxies in our sample are in low-density environments (Paper II) and have intact power-law centers. It would be interesting to check if the superdense z$\\sim$1.5 galaxies are `fast-rotators' as well.  Finally we found that 35\\% of the objects in this sample (8/23) had nuclear dust, with the majority of these being power-law rather than core galaxies. This also agrees well with previous work (e.g. Laine et al. 2003; Lauer et al. 2005)."
    },
    "0809/0809.5029_arXiv.txt": {
        "abstract": "We present multicolour photometry and modelling of the active eclipsing binary star V405 And. The components of 0.2 and 0.5 solar masses are just below and above the theoretical limit of the full convection, that is thought to be around 0.3 solar mass. The light curves are compositions of constant and variable features: the distorted shape of the components (about 25\\%), a small eclipse, and mainly of spots (about 75\\%) and flares. ",
        "introduction": "The manifestations of stellar activity appearing as spots, plages, flares, activity cycles etc., are consequences of the strong magnetic fields. The origin of this magnetic field of stars is some kind of a dynamo mechanism, which can be the so-called $\\alpha\\Omega$ dynamo, based on the amplification of the fields by differential rotation in the tachocline, which is a thin interface layer between the convection zone and the radiative core. Less massive stars are thought to be fully convective, wherein the $\\alpha^2$ dynamo generates strong, long-lasting, axisymmetric magnetic fields. The mass limit for full convection is thought to be around ~0.35 solar mass. The late type binary V405~And is an excellent target to study the effects of the two possible dynamos, since one component is below, the other is above the mass limit of the full convection \\cite{1997A&A...326..228C}. ",
        "conclusions": ""
    },
    "0809/0809.0921_arXiv.txt": {
        "abstract": "We are searching for new He atmosphere white dwarf pulsators (DBVs) based on the newly found white dwarf stars from the spectra obtained by the Sloan Digital Sky Survey. DBVs pulsate at hotter temperature ranges than their better known cousins, the H atmosphere white dwarf pulsators (DAVs or ZZ Ceti stars). Since the evolution of white dwarf stars is characterized by cooling, asteroseismological studies of DBVs give us opportunities to study white dwarf structure at a different evolutionary stage than the DAVs.  The hottest DBVs are thought to have neutrino luminosities exceeding their photon luminosities (Winget et al. 2004), a quantity measurable through asteroseismology.  Therefore, they can also be used to study neutrino physics in the stellar interior.  So far we have discovered nine new DBVs, doubling the number of previously known DBVs. Here we report the new pulsators' lightcurves and power spectra. ",
        "introduction": "White dwarf stars (WDs) are the endpoints of evolution for most stars.  Their internal structures provide key clues into their complex pre-WD evolution.  As WDs, their subsequent evolution is dominated by cooling. The older they are, the cooler they become. Why then, does there exist a range of temperatures within which we hardly see any He atmosphere WDs (DBs) while we see both the H atmosphere WDs (DAs) and non-DAs (He atmosphere DOs and DBs) at both hotter and cooler temperature than this? This paradox is the so-called ``DB gap'' (Fontaine \\& Wesemael 1987).  Recently, Sloan Digital Sky Survey (SDSS) data have shown us that the DB gap is not completely void of DBs, but rather deficient in the number of DBs (Eisenstein et al.  2006a).  The current best explanation for this effect is based on WDs having specific layer masses (the large gravity in a WD makes it compositionally stratified) which mix and settle at certain temperatures, causing the surface ``flavor'' of a WD to change with time and temperature (Fontaine \\& Wesemael 1987). This explanation demands a thin H layer in at least a substantial fraction of DAs.  However, there have been several works (Fontaine et al. 1992; Clemens 1994; Fontaine et al. 1994; Robinson et al. 1995; Kleinman et al.  1998; Benvenuto et al. 2002) suggesting that perhaps all DAs have thick H layers and if so, spectral evolution by the current model cannot happen.  Once a WD cools past the onset of its instability strip (at a temperature primarily determined by its atmospheric composition and total mass), it begins pulsating in a series of non-radial g-modes, allowing us to study its interior via the technique of asteroseismology. Asteroseismology, the study of stellar pulsations, is an important way to directly measure quantities of the stellar interior. And understanding the interior structure of the DBVs is one very important way to address some of the mysteries of DB evolution.  Among the 9 DBVs known prior to our work, the first DBV discovered (Winget et al. 1982), GD\\,358, is by far the best studied WD pulsator. It has had its internal structure substantially explored by asteroseismology (Winget et al.  1994, Bradley \\& Winget 1994; Vuille et al. 2000; Metcalfe Salaris \\& Winget 2002; Metcalfe 2003;  Kepler et al. 2005; Metcalfe et al. 2005).  The results from the asteroseismological investigations of GD\\,358 (Winget et al., 1994) are impressive: total mass of $0.61\\pm0.03 M_{\\sun}$, He layer mass of log$M_{He}/M_{\\star} = -5.7(+0.18, -0.30)$, $R_{\\star}/R{\\sun}=0.0127\\pm0.0004$, He to C transition zone thickness of about 8 pressure scale heights, absolute luminosity log$L_{\\star}/L{\\sun}=-1.30 (+0.09, -0.12)$ hence a distance of $42\\pm3pc$, weak magnetic field of $1300\\pm300$G and the measurements of radial differential rotation.  More recent detailed model fitting techniques using genetic algorithms along with improvements to the models have been successful in revealing even more information.  We now have a measurement of the oxygen mass fraction in the core which places constraints on both the nuclear burning rate $^{12}C(\\alpha, \\gamma)^{16}O$ and even more detailed structure information, such as the extent of the He/C envelope beneath the pure He envelope (Metcalfe, Salaris \\& Winget 2002; Metcalfe 2003; Metcalfe et al. 2005).  Except for one other DBV, the rest of the class have not been so forthcoming in revealing their internal structures, primarily due to their lack of the abundance of pulsation modes compared to GD\\,358's over 100 detected frequencies. CBS\\,114 is a DBV which showed promise for successful asteroseismological analysis by exhibiting a rich pulsation spectrum, but earlier observational comparisons to the models produced a $\\mathrm{C}(\\alpha,\\gamma)\\mathrm{O}$ nuclear burning rate which was at odds with that obtained from GD\\,358 (Handler, Metcalfe \\& Wood 2002). After several years of additional observations of CBS\\,114, which lead to identifying eleven independent pulsation modes (four of which were new) along with improvements in pulsation models and fitting techniques, Metcalfe et al.(2005) have achieved new asteroseismological results for both stars which are now in agreement with each other.  The one thing CBS\\,114 did not show and which GD\\,358 did were the many fine structure splittings of the pulsation modes caused predominantly by stellar rotation. Our understanding of the pulsation amplitude determining mechanism on these stars is incomplete and we cannot explain why we see significant fine-structure splitting in GD\\,358 and not much in CBS\\,114. We certainly do not believe it is due to lack of rotation on CBS\\,114's part though it could be due to the star being observed near pole on. So the search goes on for a third solvable pulsator to try and distinguish modes, models, fits and reality in these objects.    Another important reason to study DBVs is that they are great cosmic laboratories for high energy physics.  Winget et al. (2004) predict that hot DBs should have significant plasmon neutrino production. Their DB models suggest that 30,000K, $0.6M_{\\sun}$ DBs have a neutrino luminosity that is 1.8 times higher than their photon luminosity. On the cool end, 22,000K,  $0.6M_{\\sun}$ DBV models have a neutrino luminosity less than half of their photon luminosity. Thus the hottest DBVs should be losing energy and cooling significantly faster than the cooler ones.  Since a pulsation mode's period is a function of temperature, we can directly measure a star's cooling rate by measuring a mode's rate of period change (e.g. Kepler et al.  2005b). And thus, the DBVs may be quite revealing laboratories for neutrino physics.   Finally, an increase in the number of known DBVs will help us understand their properties as a group.  Clemens (1994) and Kleinman (1995, 1998) found that the DA pulsators break down nicely into two distinct classes, each subclass exhibiting common class properties which they have used to investigate the dynamics of the pulsation mechanism in these stars.  By increasing the number of known DBVs, we can search for possible subclass distinctions.  Nather, Robinson \\& Stover (1981) noted that the interacting binary white dwarf stars will each eventually form a single DB at the end of their evolution. This means that there may be more than one evolutionary channel leading to the DBs.  Perhaps we will find two distinct classes, each of them retaining the evidence of their evolutionary paths in their pulsation structures.   SDSS is a photometric and spectroscopic survey of the sky covering about 10,000 square degrees around the Northern Galactic cap (York et al. 1996; Stoughton et al. 2002; Gunn et al. 1998; Gunn et al.  2000). In SDSS's Sixth Data Release  (Adelman-McCarthy, et al. 2008), there are photometry of close to 10,000 square degrees in five filters (Fukugita et al. 1996) and 1.27 million spectra. Although the survey's main goal was to produce a 3D map of the large scale structure of the universe, it also contains data on many galactic stellar objects, including WDs.  SDSS data provide the perfect basis set for finding new DBVs which will eventually help solve the DB Gap mystery, measure the neutrino production rates inside the DBs, as well as  answer some other questions about WD structure and evolution.  Kleinman et al. (2004) published the first WD catalogue based on the spectra obtained by SDSS.  and doubled the number of then known WDs. The newest WD catalogue from the SDSS (Eisenstein et al.  2006b, DR4 WD catalogue hereafter) has almost quadrupled the number of WDs. Among the new WDs are DBs whose physical parameters determined from model fitting suggest they are inside the instability strip.  Therefore, we started a project to search for new DBVs using our spectroscopic fits to SDSS spectra, originally from Kleinman et al.  (2004) and later using the DR4 WD catalogue, to identify likely DBV candidates and follow them up with time-series photometry. This survey is the counterpart to the search for new SDSS DAVs reported by Mukadam et al. (2004), Mullally et al.(2005), Kepler et al.  (2005a) and Castanheira et al. (2006a, 2007). ",
        "conclusions": "From the DR4 WD catalogue, we have about 70 DBV candidates brighter than  $g=20$mag.  To date, we have observed 29 of them and found nine new DBVs, doubling the number of known DBVs.  We seek an increased number of DBVs to help us understand their group properties, better determine the location of the instability strip, and perhaps find hot DBVs we can use to measure their cooling rates and place a limit on the neutrino production rate in their interiors. Based on these statistics, we can expect at least another 12 new DBVs from the DR4 sample and 20 more from DR6. These are probably lower limits though,  since we suspect additional observations of our 29 currently observed objects will probably reveal new low amplitude pulsators as well."
    },
    "0809/0809.2786_arXiv.txt": {
        "abstract": "We use a non-equilibrium chemical network to revisit and study the effect of $H_{2}$, $HD$ and $LiH$ molecular cooling on a primordial element of gas. We paid special attention in the variation of $HD$ abundance. We solve both the thermal and chemical equations for a gas element with an initial temperature $T = 1000K$ and a gas number density in the range $n_{tot}=1-10^{3} cm^{-3}$. These are typical propierties of the first halos which formed stars. At low densities, $n_{tot}<10^{2} cm^{-3}$, the gas reaches temperatures $\\sim 100K$ and the main coolant is $H_{2}$, but at higher densities, $n_{tot}>10^{2} cm^{-3}$, the $HD$ molecule dominates the gas temperature evolution and the gas reaches temperatures well below $100K$. The effect of $LiH$ is negligible in all cases. We studied the effect of $HD$ abundance on the gas cooling. The $HD$ abundance was set initially to be in the range $n_{HD}/n_{H}=10^{-7}-10^{-5}$. The simulations show that at $n_{tot}>10^{2} cm^{-3}$ the HD cooling dominates the temperature evolution for $HD$ abundances greater than $10^{-6}n_{H}$. This number decrease at higher densities. Furthermore, we studied the effect of electrons and ionized particules on the gas temperature. We followed the gas temperature evolution with $n_{H_{+}}/n_{H}=10^{-4}-10^{-2}$. The gas temperature reached lower values at high ionization degree because electrons, $H^{+}$ and $D^{+}$ are catalizers in the formation paths of the $H_{2}$ and $HD$ molecules. Finaly, we studied the effect of an OB star, with $T_{eff}=4\\times 10^{4}K$, would have on gas cooling. It is very difficult for a gas with $n_{tot}$ in the range between $1-100 cm^{-3}$ to drop its temperature if the star is at a distance less than $100 pc$. ",
        "introduction": "In a $\\Lambda CDM$ Universe, the first luminous objetcs were formed due to fall-in of gas inside the dark matter potential wells (for a review see  \\citet{Barkana 2001}). In order to let the fall-in of gas inside the dark matter potential wells, the gas thermal energy should have been radiated away by some physical mechanism. Since the primordial molecular clouds would have zero metallicity, the collisional excitation of the existing molecules is the most plausible mechanism for the cooling of baryonic matter in this environment because the collisional excitation cooling of both $H$ and $He$ atoms are inefficient at temperatures lower than $8000K$, which is higher than the temperature of star formation clouds.  \\citet{Tegmark 1997} showed that the first lumious objects may have formed at $z\\sim 30$ inside a $10^{6}M_{\\odot}$ halo ($T_{vir}\\sim 1000K$), which have recently been confirmed by \\citet{O'Shea 2007} and \\citet{Gao 2007}. \\citet{Tegmark 1997} also showed that the stars formed in this environment could be as massives as $\\sim 100M_{\\odot}$ \\citep{Abel2002}.  After the recombination era, the most abundant molecule in the Universe is molecular hydrogen $H_{2}$. Despite of its low primordial abundance, $\\sim 10^{-3}-10^{-4}n_{H}$ \\citep{Palla et al. 1983}, this molecule has a fundamental role in the gas cooling at temperatures less than $8000K$. The first authors to highlight the role of $H_{2}$ in this context were \\citet{Saslaw} and \\citet{Peebles 1968}. Saslaw \\& Zipoy (1967) showed the importance of the charge transfer reaction between $H_{2}^{+}$ and $H$ to form $H_{2}$, and Peebles \\& Dicke (1968) suggested a mechanism to form $H_{2}$ from $H^{-}$.  The $H_{2}$ molecule forms by the $H^{-}$ and $H_{2}^{+}$ channels mainly. The reactions  \\begin{eqnarray} H+H^{+}\\rightarrow H_{2}^{+}+\\gamma \\\\ H+e\\rightarrow H^{-}+\\gamma \\end{eqnarray} are followed by \\begin{eqnarray} H_{2}^{+}+H\\rightarrow H_{2}+H^{+}.\\\\ H^{-}+H\\rightarrow H_{2}+e. \\end{eqnarray}  Due to its zero dipolar moment only quadrupolar rotational transitions are allowed, $J\\rightarrow J \\pm 2$, where $J$ is the quantum number for angular momentum; \\citet{Abgrall}. Furthermore, due to its small moment of inertia (the smallest one between all molecules) the energy gap between its rotational quantum states, $\\Delta E$, is large compared with other molecules ($\\Delta E_{J\\rightarrow J\\pm 2}\\propto 1/I$, where $I$ is the moment of inertia). The smallest energy gap is $\\Delta E_{2\\rightarrow 0}\\approx 500 K$. With this energy difference it is very diffucult to reach temperatures below $\\sim 100K$, (see \\citet{Palla 1999} and references therein).  The $HD$ molecule forms through $D^{+}$ and $D$ channel mainly (see \\citet{Dalgarno 1973}; \\citet{Galli 2002}) \\begin{eqnarray} D^{+}+H_{2} \\rightarrow H^{+}+HD,\\\\ D+H_{2} \\rightarrow H+HD. \\end{eqnarray} $HD$ has a moment of inertia greater than the $H_{2}$ one. Furthermore, it has a finite dipolar moment which allows internal dipolar trasitions (transitions of the kind $J\\rightarrow J\\pm 1$) and internal transition rates greater than in $H_{2}$, \\citet{Abgrall}. Due to its small moment of inertia, the differences between its energy states are smaller than the energy differences of the $H_{2}$ molecule. All these properties make the $HD$ molecule an efficient cooler at low temperatures, below $\\sim 100K$, (see \\citet{NagakuraOmukai}; \\citet{Ripamonti 2007}; \\citet{McGreerBryan 2008}; \\citet{Palla 1999} and references therein).\\\\ The $LiH$ molecule is formed mainly by radiative association of $Li$ and $H$ and associative detachment of $Li^{-}$ and $H$ \\citep{Stancil 1996}: \\begin{eqnarray} Li+H & \\rightarrow & LiH+\\gamma.\\\\ Li^{-}+H & \\rightarrow & LiH+e^{-}. \\end{eqnarray} Moreover, this molecule have both a dipolar moment and a moment of inertia larger than the ones of the $HD$ molecule. These characteristics could make the $LiH$ molecule an efficient cooler at low temperatures, but the cooling functions depend on the number density of the specie, so if the abundance of $LiH$ is too low as expected in primordial environments \\citep{Stancil 1996} its cooling effect will be negligible. For a review of $Li$ chemestry see \\citet{Bodo 2001} and \\citet{Bodo 2003}.  In the work of \\citet{Galli 1998} the effect of $HD$ and $LiH$ was included on the gas cooling. To calculate the photo-destruction ratescoefficient, for photoionization, photodetachment, and photodissociation, they assumed detailed balance with CMB photons. But, to study the effect of the first stars on the primordial gas we need the cross section for each photo-destruction process, in the spirit of \\citet{Glover}. These cross sections are described below. Our current work includes the photo-destruction cross sections of both $Li$ and $Li^{-}$ and the photodisociation of $LiH$ (in its rovibrational ground state) in contrast to previous work.  We improve over recent works \\citep{Glover,GA08} that have studied primordial cooling by exploring the effect of the $HD$ abundance and the inclusion of a stellar radiation field. This paper is organized as follow. In \\S2 we describe both the thermal and chemical model required to follow the evolution of the gas temperature. In \\S3 we present results and discussion. It includes the gas temperature evolution as a function of gas density and molecular coolers; the gas temperature evolution as a function of $HD$ abundance; the temperature evolution as a function of the ionization degree and finaly we show the effect of a star radiation field on the gas temperature. In \\S4 we present the conclusions. ",
        "conclusions": "In this work we have developed a model for the temperature evolution of a primordial gas including 21 different chemical species including reaction rates and cross sections available in the literature. We have paid careful attention to explore the space parameter in abundance, specially $HD$, and to include the $LiH$ to study in detail at what abundances the molecular coolants are relevant. The main results are the following:  \\begin{enumerate} \\item The $HD$ molecule dominates the gas cooling at temperatures below $\\sim 100-200K$ in the range of densities $10^{2}-10^{3} cm^{-3}$ for $HD$ abudances over $10^{-6}n_{H}$. The $HD$ effect is more evident at higher densities. \\item The $LiH$ molecule does not have a clear effect on the gas cooling. The gas would need an abundance at least ten orders of magnitude higher to be an efficient cooler, so $LiH$ is ruled out as an important cooler in primordial gas (\\citet{Mizusawa}). \\item A gas with high ionization degree can drop its temperature more than a neutral gas because both the ionized $H$ and $D$ and electrons are catalizers in the formation of $H_{2}$ and $HD$ cooling molcules. These ionization conditions could be presents in post-shock waves zones or relic $HII$ regions, \\citet{Johnson06}. \\item Is very difficult for a gas to drop its temperature in the presence of an OB star located closer than $100 pc$. So, in order to form more than one star in a primordial halo the formation of seed clumps should almost be instantaneous, otherwise the radiation feedbak of first stars will suppress the star formation conditions (\\citet{OmukaiNishi}). \\end{enumerate}  This work, as previos ones,  suggests the importance of studing the effect firts stars would have on their sorrounding gas in the formation of more than one star within primordial halos (see e.g. \\citet{JimenezHaiman}). For a more accurate study, we need to follow the star formation in a hydrodinamical model, which will we present in forthcoming papers."
    },
    "0809/0809.0260_arXiv.txt": {
        "abstract": "We study the degree to which non--radiative gas dynamics affects the merger histories of haloes along with subsequent predictions from a semi--analytic model (SAM) of galaxy formation. To this aim, we use a sample of dark matter only and non--radiative SPH simulations of four massive clusters. The presence of gas--dynamical processes (e.g. ram-pressure from the hot intra--cluster atmosphere) makes haloes more fragile in the runs which include gas. This results in a 25 per cent decrease in the total number of subhaloes at $z=0$. The impact on the galaxy population predicted by SAMs is complicated by the presence of `orphan' galaxies, i.e. galaxies whose parent substructures are reduced below the resolution limit of the simulation. In the model employed in our study, these galaxies survive (unaffected by the tidal stripping process) for a residual merging time that is computed using a variation of the Chandrasekhar formula. Due to ram--pressure stripping, haloes in gas simulations tend to be less massive than their counterparts in the dark matter simulations. The resulting merging times for satellite galaxies are then longer in these simulations. On the other hand, the presence of gas influences the orbits of haloes making them on average more circular and therefore reducing the estimated merging times with respect to the dark matter only simulation.  This effect is particularly significant for the most massive satellites and is (at least in part) responsible for the fact that brightest cluster galaxies in runs with gas have stellar masses which are about 25 per cent larger than those obtained from dark matter only simulations. Our results show that gas-dynamics has only a marginal impact on the statistical properties of the galaxy population, but that its impact on the orbits and merging times of haloes strongly influences the assembly of the most massive galaxies. ",
        "introduction": "\\label{sec:intro}  During the last decade, a number of observational tests of the standard cosmological model have ushered in a new era of `precision cosmology'. Precise measurements of angular structure in the Cosmic Microwave Background (CMB), combined with other geometrical and dynamical cosmological tests have constrained cosmological parameters tightly \\citep[][and references therein]{Komatsu2008} confirming the hierarchical cold dark matter model (CDM) as the `standard' model for structure formation. While the cosmological paradigm is well established, our understanding of the physical processes regulating the interplay between different baryonic components is still far from complete, and galaxy formation and evolution remains one of the most outstanding questions of modern astrophysics.  Different approaches have been developed in order to link the observed properties of luminous galaxies to those of the dark matter haloes in which they reside. Among these, semi--analytic models (SAMs) of galaxy formation have developed into a flexible and widely used tool that allows a fast exploration of the parameter space, and an efficient investigation of the influence of different physical assumptions. Computational costs are therefore reduced with respect to hydrodynamical simulations, but this is done at the expense of an explicit description of the gas dynamics \\citep[for a recent review on SAMs, see][]{2006RPPh...69.3101B}. Although recent work has started analysing the properties of the galaxy populations in hydrodynamical simulations \\cite[e.g.][]{1996ApJ...472..460F,1999ApJ...521L..99P,2005ApJ...618...23N,2005ApJ...618..557N,2006MNRAS.373..397S,2006MNRAS.373.1265O}, the computational time is still prohibitive for simulations of galaxies in large cosmological volumes.  In addition, the uncertainties inherent in the physical processes at play obviously place strong limits on the accuracy with which galaxies can be simulated. As a consequence, these numerical studies also require an adequate handling of `sub-grid' physics either because the resolution of the simulation becomes inadequate to resolve the scale of the physical process considered, or because we do not have a ``complete theory'' of that particular physical process (which is almost always true). It is therefore to be expected that SAMs will remain a valid method to study galaxy formation for the foreseeable future.  In their first renditions, SAMs took advantage of Monte Carlo techniques coupled to merging probabilities derived from the extended Press-Schechter theory to construct merging history trees of dark matter haloes \\citep{1993MNRAS.264..201K,1994MNRAS.271..781C}. An important advance of later years has been the coupling of semi-analytic techniques with direct $N$-body simulations \\citep{Kauffmann_etal_1999,Benson_etal_2000}. Since dark matter only simulations can handle large numbers of particles, such `hybrid' models can access a very large dynamic range of mass and spatial resolution offering, at the same time, the possibility to model the spatial distribution of galaxies within dark matter haloes. It is also interesting to note that there have been a number of recent studies showing that the extended Press-Schechter formalism does not provide a faithful description of the merger trees extracted directly from N-body simulations \\citep{Benson_et_al_2005,Li_et_al_2007,Cole_et_al_2008}.  This might have important consequences on the predicted properties of model galaxies, although a detailed investigation of the influence of analytical versus numerical merger trees on the predicted properties of model galaxies has not been carried out yet.  A related question is whether the inclusion of the baryonic component alters the halo dynamics with respect to a purely dark matter (DM) simulation.  Processes like ram--pressure stripping and gas viscosity are expected to produce a significant segregation between the collisional and collisionless components \\citep{2001ApJ...561..708V}. These effects are likely more important in environments characterised by high densities and large velocity dispersions (like galaxy clusters), and are expected to change the dynamics and the timing of halo mergers. As the merger history of model galaxies in a SAM is essentially driven by the merger history of its parent halo, any physical process that affects halo mergers will influence model predictions in some measure. We note that recent work has used merger trees from non--radiative hydrodynamic simulations \\citep[e.g.][]{cora08} to study the chemical enrichment of the intra--cluster medium (ICM). This approach offers the advantage of providing a three-dimensional picture of the ICM, while keeping the advantage of exploring different physical choices with sensibly reduced computational times with respect to hydrodynamical simulations. The question of how SAM predictions are affected by using merger trees from different types of simulations (e.g. DM and hydrodynamical simulations) has, however, not been addressed.  The purpose of this paper is to quantify the effects of the presence of gas on the merger histories of haloes, and on predictions from a galaxy formation model.  To this aim, we have used a sample of DM-only and non--radiative hydrodynamical simulations of four massive galaxy clusters (see Sec.~\\ref{sec:sims}).  The merger trees constructed from these simulations have been used as input for a SAM (see Sec.~\\ref{sec:SAM}), and results have been used to carry out a careful comparison of the statistical properties of the galaxy populations and of the formation history of the brightest cluster galaxies (BCGs) from the two sets of simulations.  The use of non--radiative hydrodynamics is only a first step towards a detailed comparison between SAMs and hydrodynamic simulations. A more realistic comparison should include also gas cooling and processes related to compact object physics, such as star formation, supernovae feedback and supermassive black holes production and evolution. We will present this analysis in a future work. We note that previous work has already compared results of smooth particle hydrodynamics (SPH) simulations and SAMs to calculate the evolution of cooling gas during galaxy formation (\\citealt{Benson_et_al_2001,Yoshida_et_al_2002,Helly_et_al_2003}; see also \\citealt{Cattaneo_et_al_2007}), but a detailed comparison is still lacking.  The plan of the paper is as follows. In Sec.~\\ref{sec:sims} we describe the cluster simulations used in this study, and describe the method used for the construction of the galaxy merger trees. In Sec.~\\ref{sec:SAM} we provide a brief description of the SAM adopted, and in Sec.~\\ref{sec:Res} we present the results of our analysis. Finally, in Sec.~\\ref{sec:Concl}, we summarise our findings and give our conclusions. ",
        "conclusions": "\\label{sec:Concl}  In this paper we have used numerical simulations to analyse how the presence of non--radiative gas dynamics affects the predictions of semi-analytic models of galaxy formation for the properties of cluster galaxies. The main results of our work can be summarised as follows.  \\begin{enumerate}  \\item The stellar mass function of galaxies from DM-only runs is in quite good agreement with that obtained from non--radiative hydrodynamical runs. This result is a combination of two different and opposite effects.  \\begin{itemize} \\item Due to a reduced number of subhaloes in the GAS runs (see Figure \\ref{fi:MF_haloes}), these simulations result in a galaxy population with a reduced number of Type-0 and Type-1 galaxies (i.e. central galaxies of a halo, either the main halo or a proper substructure). \\item Due to a systematic increase of the residual merging times assigned to Type-2 galaxies (those associated with haloes disrupted below the resolution limit of the simulation), the cluster galaxy population in the GAS runs contains a larger number of Type-2 galaxies than the DM runs. \\end{itemize}  \\item The longer merging times assigned on average to Type-2 galaxies in the GAS runs are due to ram--pressure stripping, which removes gas from the merging subhaloes and makes them more fragile. The effect of ram--pressure is more important at lower redshift, when the cluster has already assembled in a dominant structure with a high--pressure atmosphere that can efficiently remove gas from substructures. When considering the entire satellite population, we find a systematic difference between the DM and GAS runs in the sense that merging substructures are less massive in the runs with gas. This trend, however, is reversed when concentrating on the most massive satellites (see item iv below). \\vspace{0.1cm}  \\item Type-2 galaxies dominate the radial density profile of cluster galaxies particularly in the inner regions, in agreement with results by \\citet{Gao_etal_2004}. Galaxies associated with distinct dark matter substructures (Type-1 galaxies) exhibit a flatter distribution and their contribution to the inner regions of galaxy clusters is negligible. We did not find any significant difference, in terms of spatial distribution, between the DM and the GAS runs.\\vspace{0.1cm}  \\item Although a statistical comparison between galaxy populations from the two sets of runs results in a quite nice agreement, a one-to-one comparison for the brightest central galaxies shows that these galaxies tend to have larger stellar masses in runs with gas. The difference varies from cluster to cluster and it is generally due to single merging events of relatively massive satellites which get assigned lower merging times in the GAS runs (see the example shown in Figure \\ref{fi:traj}). The final difference in stellar mass is then due primarily to a different accretion history of satellite galaxies in the two sets of runs, and not to intrinsic differences in the star formation rates in the main progenitor. \\end{enumerate}  Our results demonstrate that predictions of semi-analytic models of galaxy formation are not significantly affected when non-radiative hydrodynamic simulations are used to construct the halo merger trees which provide the skeleton of the model. This statement is, however, correct only in a statistical sense. The presence of the gas induces significant differences in the timing of the halo mergers, and affects significantly the halo orbits making them more circular, on average. Although these effects might be over-estimated in our non--radiative runs, our results suggest that an accurate treatment of merging times is crucial for predicted quantities like the mass accretion history of model brightest cluster galaxies. As subhaloes are fragile systems that are rapidly reduced below the resolution limit of the simulation \\citep{2004MNRAS.348..333D,2004MNRAS.355..819G}, the treatment of satellite mergers in semi-analytic models requires the use of analytic formulations (e.g.  the Chandrasekhar formula). Recent work \\citep{2008MNRAS.383...93B,2008ApJ...675.1095J} has shown the limits of the formulation usually adopted in semi-analytic models. This recent work, however, does not provide consistent alternative formulations. Additional work is therefore needed in order to obtain a more realistic and detailed description of the merging process, which represents a crucial ingredient of semi-analytic models of galaxy formation."
    },
    "0809/0809.1785_arXiv.txt": {
        "abstract": "{} {We have studied the bulge and the disk kinematics of the giant low surface brightness galaxy ESO 323-G064 in order to investigate its dynamical properties and the radial mass profile of the dark matter (DM) halo.} {We observed the galaxy with integral field spectroscopy (VLT/VIMOS, in IFU configuration), measured the positions of the ionized gas by fitting Gaussian functions to the \\oiiipg\\ and \\hb\\ emission lines, and fit stellar templates to the galaxy spectra to determine velocity and velocity dispersions. We modeled the stellar kinematics in the bulge with spherical isotropic Jeans models and explored the implications of self consistent and dark matter scenarios for NFW and pseudo isothermal halos.} {In the bulge-dominated region, $r<5''$, the emission lines show multi-peaked profiles. The disk dominated region of the galaxy, $13''<r<30''$, exhibits regular rotation, with a flat rotation curve that reaches $248 \\pm 6$ \\kms. From this we estimate the total barionic mass to be $M_{bar} \\sim 1.9 \\cdot 10^{11}$ \\msun\\ and the total DM halo mass to be $M_{DM} \\sim 4.8 \\cdot 10^{12}$ \\msun. The stellar velocity and velocity dispersion have been measured only in the innermost $\\approx 5''$ of the bulge, and reveal a regular rotation with an observed amplitude of 140 \\kms\\ and a central dispersion of $\\sigma=180$ \\kms. Our simple Jeans modeling shows that dark matter is needed in the central $5''$ to explain the kinematics of the bulge, for which we estimate a mass of $\\approx (7 \\pm 3) \\cdot 10^{10}$ \\msun. However, we are not able to disentangle different DM scenarios. The computed central mass density of the bulge of ESO 323-G064 resembles the central mass density of some high surface brightness galaxies, rather than that of low surface brightness galaxies.} {} ",
        "introduction": "Galaxy properties define a continuum, in size, luminosity, mass, and surface brightness. Low surface brightness (LSB) galaxies have significantly extended the range in surface brightness over which galaxies can be studied, and thus they have received a great deal of attention in studies of dark matter (e.g. \\citealt{Pfenniger+94, deBlok+97, McGaugh+01, Swaters+00, Swaters+03, Kuzio+08}), stellar content (e.g. \\citealt{deBlok+95, Bothun+97, Boissier+08}) and gas content (e.g. \\citealt{ONeil+03, Pizzella+08b}). Because of their unique properties, LSB galaxies also play a significant role in our understanding of the universe. Their distribution in the field and in galaxy clusters will help us to understand the bright and dark matter distribution in the local neighborhood and in the universe; their number and density distribution may be relevant to the study of damped Ly-$\\alpha$ absorption seen against quasars (e.g. \\citealt{Jimenez+99, Oneil02}).  Like high surface brightness galaxies, LSB galaxies also span a range in properties, such as size, mass, and bulge size. For example, \\citet{Bothun+90} discovered a class of giant LSB galaxies, and \\citet{Beijersbergen+90} studied a sample of bulge-dominated LSB galaxies. To date, very little is known of the properties of giant-sized or bulge-dominated LSB galaxies even though they can enhance our understanding of galaxy formation and evolution by providing an opportunity to sample galaxies in a hitherto almost unexplored regime of galaxy properties. Several scenarios have been proposed to explain their formation: density peak in voids \\citep{Hoffman+92}, bar instability \\citep{Noguchi01,Mayer+04}, or secular evolution from ring galaxies \\citep{Mapelli+08}.  To fill the gap between regular and giant LSBs, Swaters, Rubin, \\& McGaugh (SRM in preparation) have recently completed a study of the kinematics of a sample of bulge dominated LSB galaxies. Their aim is to study the dark matter properties and to determine whether bulge-dominated LSB galaxies are dark matter dominated like other LSB galaxies (\\citealt{deBlok+97, Swaters+03}).  In this paper we show additional results obtained for the giant LSB galaxy ESO 323-G064, selected from the SRM project. ESO 323-G064 has an heliocentric velocity of 14830 \\kms\\ at a distance of  194 Mpc (assuming $H_0=75$ \\kmsmpc\\ and a galactocentric correction of $-278$ \\kms).  On its digital sky survey and 2MASS images, ESO 323-G064 has a bright, compact bulge, surrounded by a low surface brightness disk. There is some evidence of a weak bar component. The long slit spectra of ESO 323-G064 obtained (SRM, in preparation) with the 6.5m Baade Telescope at LCO, show strong, double-peaked emission from H$\\alpha$, \\nii, [SII], [OI] in the nuclear region. The large extent (47 kpc) of the H$\\alpha$ and the large peak-to-peak rotation velocity of 445 \\kms\\ (on the plane of the sky) mark ESO 323-G064 as a likely giant LSB.  In order to determine the kinematics of the stars in the bulge-dominated region, and to compare these motions with those of the gas in both the nucleus and at larger radii, we decided to obtain two dimensional data using an integral field unit. This will give a better picture of the complex gaseous kinematics especially the double peaked emission revealed by the long slit data. Optical two-dimensional observations in LSB galaxies have been obtained in the past (i.e. \\citealt{Swaters+03,Kuzio+06, Kuzio+08, Pizzella+08b}) but, so far, never for the stellar component of a giant LSB.  The paper is organized as follow: in Section \\ref{sec:obs} we present the integral field observations and discuss the data reduction process; in Sections 3 and 4 we present the kinematical results for the ionized gas and stellar components respectevely; in Section 5 we present the Jeans model to the stellar kinematics and in Section 6 we discuss the results. ",
        "conclusions": "We presented the two-dimensional velocity and velocity dispersion fields for the gaseous and stellar components of the LSB galaxy ESO 323-G064. The gas emission lines show a very bright and complex structure within the central $5''$, characterized by 3 peaks, which we interpret as due to the presence of 3 spatially unresolved emitting regions. At radii out to $30''$, in the region dominated by the galaxy disk, a weak \\hb\\ emission is detected showing a regular velocity field with a maximal amplitude of $248\\pm6$ \\kms.  The stellar absorption lines are detectable only in the very innermost regions of the bulge.  The stellar kinematics shows a regular rotation with an amplitude of $\\pm 140$ \\kms\\ and a central velocity dispersion of $\\approx 180$ \\kms. The intrinsic $(V/\\sigma)_{INTR}$ ratio for the bulge is $0.56 \\pm 0.02$, which is consistent with an isotropic rotator in the $0.1 < \\epsilon_{OBS} <0.2$ observed ellipticity range, considering an inclination of 62 degres. These values of $V/\\sigma$ and $\\epsilon$ place the bulge of ESO 323-G064 among fast rotator bulges \\citep{Cappellari+06}. The value of $V/\\sigma $ is considerably small when compared to the average value determined from the sample of 6 bulge dominated LSB galaxies of Pizzella et al. 2008 ($<V/\\sigma> > 1$).  The circular velocity, $248 \\pm 6$ \\kms\\ measured from the gaseous disk, places ESO 323-G064 in good agreement with the location of LSB galaxies in the $V_C -\\sigma_c$ plane \\citep{Courteau+07b, Courteau+07}. On the other hand, this value is lower then the value of 320 \\kms predicted by \\citet{Pizzella+05}.  The intrinsic bulge ellipticity value for ESO 323-G064 ($\\epsilon_{INTR}=0.20 \\pm 0.07 $) is consistent with the mean value of bulges of the high surface brightness disk galaxies $<\\epsilon_{HSB}>=0.15$ (as determined by \\citealt{Mendez-Abreu+08}) while it is lower than the average value for bulge dominated LSB galaxies $<\\epsilon_{LSB}>=0.45$ (as found by \\citealt{Pizzella+08} in a small sample of 6 bulge dominated LSB).   The amplitude of the gaseous rotation curve ($248\\pm6$ \\kms) leads to an estimate of the total baryonic mass in the galaxy of $M_{bar}=(1.9\\pm0.2)\\cdot 10^{11}$ \\msun using the empirical baryonic Tully Fisher, as done in \\citet{McGaugh05}. Moreover, under the hypothesis of a $\\Lambda$CDM universe (see Section \\ref{sec:largeRadiiGas} for details on the assumed parameters), we estimate the total mass for the dark matter halo $M_{DM} \\sim 5 \\cdot 10^{12}$ M$_{\\odot}$.   We produce spherical isotropic Jeans models for the stellar kinematics in the bulge, exploring the self consistent, NFW and pseudo isothermal scenarios. Even though the data are to be taken with some caveats (due to the lack of good photometry, limited spatial extension and resolution of the stellar kinematics) with this simple analysis we show that dark matter scenarios fit the data better than the self consistent model. The derived total bulge mass is $(7 \\pm 3)\\cdot 10^{10}$ \\msun\\ but we are not able to disentangle between the two different dark matter models.  The derived central bulge mass density (see Table \\ref{tab.model.result}) is $\\rho =15^{+9}_{-5}$ [M$_{\\odot}$ pc$^{-3}$] in the NFW scenario, and $\\rho =5^{+2}_{-1}$ [M$_{\\odot}$ pc$^{-3}$] in the pseudo isothermal scenario. Typical values of central mass density range from few $10^{-3}$ M$_{\\odot}$ pc$^{-3}$ to few $10^{-2}$ M$_{\\odot}$ pc$^{-3}$ for regular low surface brightness (see for example \\citealt{Kuzio+06, deBlok+01}) and giant low surface brightness galaxies (\\citealt{Pickering+97}). On the contrary, a much wider range of values is measured for regular galaxies, from few $10^{-3}$ M$_{\\odot}$ pc$^{-3}$ to several $10^3$ M$_{\\odot}$ pc$^{-3}$ (i.e. \\citealt{Salucci+01, Noordermeer+07}). Therefore, in this picture, the bulge of ESO 323-G064 resembles more the central mass density of regular bulges than those measured in low surface brightness galaxies. This is consistent also with the fact that bulges of giant LSB galaxies are photometrically similar to those of regular high surface brigthtness galaxies \\citep{McGaugh+95, Beijersbergen+99}     \\begin{figure} \\psfig{file=figure20.ps,width=8.0cm,clip=} \\caption{Fraction of dark matter mass compared to total mass, as a function of radial distance, for the inner $5''$ of ESO 323-G064. NFW ({\\it solid line}); pseudo isothermal ({\\it dashed line}).} \\label{fig:dm_ratio} \\end{figure}  \\begin{figure} \\psfig{file=figure21.ps,width=8.0cm,clip=} \\caption{{\\it Filled circles:} mass density values derived from the observed velocity and velocity dispersion (Equation \\ref{eqn.mass.density.derived}) compared to the best fit NFW ({\\it continuous line}), pseudo isothermal ({\\it dashed line}) and self consistent ({\\it dotted line}) model predictions.} \\label{fig.density.profile} \\end{figure}  \\noindent {\\bf Acknowledgments.} The authors wish to thank the referee A. Bosma for useful suggestions which improved the paper content and the discussion."
    },
    "0809/0809.3113_arXiv.txt": {
        "abstract": "We study the dynamics of the scalar field FLRW flat cosmological models within the framework of the {\\it Unified Dark Matter} (UDM) scenario. In this model we find that the main cosmological functions such as the scale factor of the Universe, the scalar field, the Hubble flow and the equation of state parameter are defined in terms of hyperbolic functions. These analytical solutions can accommodate an accelerated expansion, equivalent to either the {\\em dark energy} or the standard $\\Lambda$ models. Performing a joint likelihood analysis of the recent supernovae type Ia data and the Baryonic Acoustic Oscillations traced by the Sloan Digital Sky Survey (SDSS) galaxies, we place tight constraints on the main cosmological parameters of the UDM cosmological scenario. Finally, we compare the UDM scenario with various dark energy models namely $\\Lambda$ cosmology, parametric dark energy model and variable Chaplygin gas. We find that the UDM scalar field model provides a large and small scale dynamics which are in fair agreement with the predictions by the above dark energy models although there are some differences especially at high redshifts. ",
        "introduction": "The detailed analysis of the available high quality cosmological data (Type Ia supernovae \\cite{Riess07}, \\cite{essence}; CMB \\cite{Spergel07}, \\cite{Komatsu08}, etc.) leads to the conclusion that we live in a flat and accelerating universe. In order to investigate the cosmic history of the observed universe, we have to introduce a general cosmological model which contains cold dark matter to explain the large scale structure clustering and an extra component with negative pressure, the vacuum energy (or in a more general setting the ``dark energy''), to explain the observed accelerated cosmic expansion (Refs. \\cite{Riess07,essence,Spergel07,Komatsu08} and references therein). The nature of the dark energy is one of the most fundamental and difficult problems in physics and cosmology. There are many theoretical speculations regarding the physics of the above exotic dark energy, such as a cosmological constant (vacuum), quintessence, $k-$essence, vector fields, phantom, tachyons, Chaplygin gas and the list goes on (see \\cite{Ratra88,Weinberg89,Wetterich:1994bg,Caldwell98,KAM,Caldwell,Peebles03, Brax:1999gp,fein02,chime04,Brookfield:2005td,Copel06,Boehmer:2007qa,Friem08} and references therein).  Such studies are based on the general assumption that the real scalar field $\\phi$ rolls down the potential $V(\\phi)$ and therefore it could resemble the dark energy \\cite{Ozer87,Peebles88, Weinberg89,Turner97,Caldwell98,Padm03,Peebles03}. This is very important because scalar fields could provide possible solutions to the cosmological coincidence problem. In this framework, the corresponding stress-energy tensor takes the form of a perfect fluid, with density $\\rho_{\\phi}={\\dot \\phi}^{2}/2+V(\\phi)$ and pressure $P_{\\phi}={\\dot \\phi}^{2}/2-V(\\phi)$. From a cosmological point of view, if the scalar field varies slowly with time, so that ${\\dot \\phi}^{2}/2V \\ll 1$, then ${\\rm w}\\equiv P_{\\phi}/\\rho_{\\phi}\\approx -1$, which means that the scalar field evolves like a vacuum energy. Of course in order to investigate the overall dynamics we need to define the functional form of the potential energy. The simplest example found in the literature is a scalar field with $V(\\phi)\\propto \\phi^{2}$ (see for review \\cite{Peebles03}, \\cite{Dolgov90}) and it has been shown that the time evolution of this scalar field is dominated by oscillations around $\\phi=0$. Of course, the issue of the potential energy has a long history in scalar field cosmology (see \\cite{sahni00,santi00,sen02,kehagias04,Gorini05} and references therein) and indeed several parameterizations have been proposed (exponential, power law, hyperbolic etc).   The aim of the present work is to investigate the observational consequences of the overall dynamics of a family of flat cosmological models by using a hyperbolic scalar field potential which appears to act both as dark matter and dark energy \\cite{Bertacca07}. To do so, we use the traditional Hamiltonian approach. In fact, the idea to build cosmological models in which the dark energy component is somehow linked with the dark matter is not new in this kind of studies. Recently, alternative approaches to the unification of dark energy and dark matter have been proposed in the framework of the generalized Chaplygin gas \\cite{Kam01,Bili02} and in the context of supersymmetry \\cite{Taka06}.  The structure of the paper is as follows. The basic theoretical elements of the problem are presented in section 2 by solving analytically [for spatially flat Unified Dark Matter (UDM) scalar field models] the equations of motion. In section 3, we present the functional forms of the basic cosmological functions [$a(t)$, $\\phi(t)$ and $H(t)$]. In section 4 we place constraints on the main parameters of our model by performing a joint likelihood analysis utilizing the SNIa data \\cite{essence} and the observed Baryonic Acoustic Oscillations (BAO) \\cite{Eis05} and \\cite{Pad07}. In particular, we find that the matter density at the present time is $\\Omega_{m} \\simeq 0.25$ while the corresponding scalar field is $\\phi_{0} \\simeq 0.42$ in geometrical units (0.084 in Planck units). Section 5 outlines the evolution of matter perturbations in the UDM model. Also we compare the theoretical predictions provided by the UDM scenario with those found by three different type of dark energy models namely $\\Lambda$ cosmology, parametric dark energy model and variable Chaplygin gas. We verify that at late times (after the inflection point) the dynamics of the UDM scalar model is in a good agreement, with those predicted by the above dark energy models although there are some differences especially at early epochs: (i) the UDM equation of state parameter takes positive values at large redshifts, (ii) it behaves well with respect to the cosmic coincidence problem, and (iii) before the inflection point the cosmic expansion in the UDM model is much more decelerated than in the other three dark energy models implies that the large scale structures (such as galaxy clusters) are more bound systems with respect to those cosmic structures which produced by the other three dark energy models. Finally, we draw our conclusions in section 6. ",
        "conclusions": "In this work we investigate analytically and numerically the large and small scale dynamics of the scalar field FLRW flat cosmologies in the framework of the so called {\\it Unified Dark Matter} scenario. In particular using a Hamiltonian formulation we find that the time evolution of the basic cosmological functions are described in terms of hyperbolic functions. This theoretical approach yields analytical solutions which can accommodate a late time accelerated expansion, equivalent to either the dark energy or the standard $\\Lambda$ models. Furthermore, based on a joint likelihood analysis using the SNIa data and the Baryonic Acoustic Oscillations, we put tight constraints on the main cosmological parameters of the UDM cosmological model. In particular we find $\\Omega_{m}\\simeq 0.25$ and the scalar field at the present time is $\\phi_{0}\\simeq 0.42$ or 0.084 (in Planck units). Also, we compare the UDM scenario with various dark energy models namely $\\Lambda$ cosmology, parametric dark energy model and variable Chaplygin gas. We find that the cosmological behavior of the UDM scalar field model is in a good agreement, especially after the inflection point, with those predicted by the above dark energy models although there are some differences especially at early epochs. In particular, we reveal that the UDM scalar field cosmology has three important differences over the other three dark energy models considered: \\begin{itemize}  \\item It can pick up positive values of the equation of state parameter at large redshifts ($z>0.8$). Also, it behaves relatively well with respect to the cosmic coincidence problem.  \\item At early enough epochs ($a \\sim 0.15$ or $z\\sim 5.5$) the cosmic expansion in the UDM model is much more decelerated than in the other three dark energy models. In order to investigate whether the expansion of the observed universe has the above property, we need a visible distance indicator (better observations) at high redshifts ($2\\le z \\le 6$).  \\item Close to the cluster formation epoch, its collapse factor $\\lambda_{UDM}$ is less than $12\\%$ of the corresponding factor of the other three dark energy models. This feature points to the direction that perhaps the $\\lambda$ parameter can be used as a cosmological tool.  \\end{itemize}"
    },
    "0809/0809.4505_arXiv.txt": {
        "abstract": "The low number density of the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxies (LRGs) suggests that LRGs occupying the same dark matter halo can be separated from pairs occupying distinct dark matter halos with high fidelity.  We present a new technique, Counts-in-Cylinders (CiC), to constrain the parameters of the satellite contribution to the LRG Halo-Occupation Distribution (HOD).  For a fiber collision-corrected SDSS spectroscopic LRG subsample at $0.16 < z < 0.36$, we find the CiC multiplicity function is fit by a halo model where the average number of satellites in a halo of mass $M$ is $\\left<N_{sat}(M)\\right> = ((M - M_{cut})/M_1)^{\\alpha}$  with $M_{cut} = 5.0^{+1.5}_{-1.3} (^{+2.9}_{-2.6}) \\times 10^{13} M_{\\sun}$, $M_1 = 4.95^{+0.37}_{-0.26} (^{+0.79}_{-0.53}) \\times 10^{14} M_{\\sun}$, and $\\alpha = 1.035^{+0.10}_{-0.17} (^{+0.24}_{-0.31})$ at the 68\\% and 95\\% confidence levels using a WMAP3 cosmology and $z=0.2$ halo catalog.  Our method tightly constrains the fraction of LRGs that are satellite galaxies, $6.36^{+0.38}_{-0.39}$\\%, and the combination $M_{cut}/10^{14} M_{\\sun} + \\alpha = 1.53^{+0.08}_{-0.09}$ at the 95\\% confidence level. We also find that mocks based on a halo catalog produced by a spherical overdensity (SO) finder reproduce both the measured CiC multiplicity function and the projected correlation function, while mocks based on a Friends-of-Friends (FoF) halo catalog has a deficit of close pairs at $\\sim 1$ Mpc/$h$ separations.  Because the CiC method relies on higher order statistics of close pairs, it is robust to the choice of halo finder.  In a companion paper we will apply this technique to optimize Finger-of-God (FOG) compression to eliminate the 1-halo contribution to the LRG power spectrum. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS; \\citet{york/etal:2000}) has recorded the largest sample of Luminous Red Galaxies (LRGs), probing a volume of $\\sim 1 \\; (h^{-1}$ Gpc$)^3$ out to $z \\sim 0.5$ \\citep{eisenstein/etal:2001} and making it ideal for studying large scale structure.  Understanding the small-scale relationship between the galaxy and dark matter density fields is essential to extracting the linear matter power spectrum from the galaxy power spectrum, even on very large scales \\citep{schulz/white:2006, sanchez/cole:2008}.\\\\  The Halo Occupation Distribution (HOD) is a popular and useful description of this relationship \\citep{seljak:2000,peacock/smith:2000,cooray/sheth:2002}, and can be used to constrain the rate of merging, disruption or evolution in well-defined galaxy populations.  \\citet{conroy/ho/white:2007}, \\citet{white/etal:2007}, and \\citet{wake/etal:2008} have used this framework to constrain the rate at which LRGs merge or are disrupted in clusters.  \\citet{brown/etal:2008} combine HOD constraints with luminosity function measurements of red galaxies to deduce that stellar mass build-up in clusters occurs primarily in the satellite galaxies or intracluster light between $z=1$ and $z=0$, while the central galaxies grow only modestly, with $L_{cen} \\sim M^{1/3}$.  \\citet{zheng/coil/zehavi:2007} constrain stellar mass growth between DEEP2 ($z\\sim1$) and SDSS ($z\\sim0$) galaxies using the HOD description, and \\citet{conroy/etal:2007} employ HOD modeling to illuminate the fate of $z \\sim 2$ star-forming galaxies.  Others researchers, e.g., \\citet{chen:2007} and \\citet{ho/etal:2007}, use the HOD description to study the spatial distribution of satellite galaxies.  \\citet{wake/etal:2008} also argue that the small-scale clustering at different redshifts constrains the scatter in halo merger histories which can be compared with predictions of hierarchical models.\\\\  Several groups have used two and three point statistics \\citep{blake/collister/lahav:2007, kulkarni/etal:2007, white/etal:2007, zheng/etal:2008, wake/etal:2008, padmanabhan/etal:2008} as well as galaxy-galaxy lensing \\citep{mandelbaum/etal:2006} to constrain the HOD of LRGs.  \\citet{ho/etal:2007} have taken a more direct approach and used X-ray determined cluster masses to measure $N_{LRG}(M)$.  Though these analyses were performed on samples with different luminosity and redshift ranges, they offer seemingly conflicting results on both the slope $\\alpha$ at the high mass limit of the satellite term $N_{sat} \\sim M^{\\alpha}$ and the fraction of LRGs that are satellite galaxies.  \\citet{ho/etal:2007} find $\\alpha \\approx 0.6$ when fitting the {\\em total} LRG number $N_{tot}(M) \\propto M^{\\alpha}$, \\citet{kulkarni/etal:2007} find $\\alpha = 1.4$, \\citet{blake/collister/lahav:2007} find $\\alpha \\sim 2.1 - 2.6$, and \\citet{zheng/etal:2008} find $\\sim 1.8$ for $\\sigma_8 = 0.8$; \\citet{kulkarni/etal:2007} report a satellite fraction of $\\sim 17\\%$, while \\citet{blake/collister/lahav:2007}'s redshift slices span 3-8\\%, and \\citet{zheng/etal:2008} find $5-6\\%$ for the LRG subsample studied in this paper.  The most luminous elliptical galaxies in \\citet{mandelbaum/etal:2006} have a satellite fraction of $\\lesssim 10\\%$.\\\\  The low number density of the SDSS Luminous Red Galaxy (LRG) sample {\\em suggests} that LRG pairs occupying the same dark matter halo can be separated from pairs occupying distinct dark matter halos with high fidelity.  In this paper we explore that intuition, and show that one-halo pairs can be identified with $\\sim 75\\%$ completeness and $\\lesssim 27\\%$ contamination by simple cuts in the transverse separation $\\Delta r_{\\perp}$ and LOS separation $\\Delta r_{\\parallel}$.  Furthermore, these pairs can be grouped together using a Friends-of-Friends (FoF) algorithm to estimate the LRG group multiplicity function.  We apply this technique to a sample of LRGs from SDSS to constrain their HOD.  We find that both the high values of $\\alpha \\sim 2$ and high satellite fractions reported in previous papers are inconsistent with the $0.16 < z < 0.36$ SDSS LRG group multiplicity function measured here.  In contrast to previous methods which rely on 2 and 3 point statistics to constrain the HOD (as in \\citet{kulkarni/etal:2007}), our method probes the HOD more directly by estimating the group multiplicity function from the higher order statistics in the LRG density field in the one-halo dominant regime.\\\\  We present an overview of the CiC method in \\S~\\ref{overview} and apply it to an approximately volume limited subsample of SDSS LRGs in \\S~\\ref{data}, addressing the complications of fiber collisions, incompleteness, and complex angular masks. The CiC technique developed here requires calibration on mock galaxy catalogs.  We summarize our $N$-body simulation parameters in \\S~\\ref{sims}.  \\S~\\ref{hodmodel} presents the HOD model we employ throughout this analysis and details how we populate our simulations with galaxies.  \\S~\\ref{CiCtechnique} describes the CiC technique to measure the LRG group multiplicity function and its calibration with simulations.  The HOD parameters are fit using a maximum likelihood analysis explicated in \\S~\\ref{maximumL}. In \\S~\\ref{results} we present the CiC multiplicity function of our SDSS LRG subsample and describe the relation between the CiC and true group multiplicity functions. We present the constraints on the HOD parameters and their implications for the fraction of LRGs that are satellites, as well as the mass distribution of halos hosting LRG groups with $n_{sat}$ satellites.  Mock catalogs produced using the CiC maximum likelihood HOD and a spherical overdensity (SO) halo catalog agree with the \\citet{masjedi/etal:2006} measurement of $w_p(r_p)$ for this sample when the large scale bias is adjusted with a single parameter for the central galaxy HOD.  In \\S~\\ref{fofkulksec} we compare these results with a mock LRG catalog based on a FoF halo catalog and with other HOD measurements in the literature.  We show that while the FoF and SO catalogs can both match the observed CiC multiplicity function, the FoF catalog produces mock catalogs with a deficit of halos at 1 Mpc/$h$ that is evident in the projected correlation function.  In \\S~\\ref{assessCiC} we comment on the strengths and weaknesses of the CiC method and summarize our conclusions in \\S~\\ref{conc}.\\\\  Throughout this paper we adopt the \\citet{spergel/etal:2007} cosmological parameters used in our simulations to convert redshifts to distances: ($\\Omega_m, \\Omega_b, \\Omega_{\\Lambda}, n_s, \\sigma_8, h$) = (0.26, 0.044, 0.74, 0.95, 0.77, 0.72).  All distances and separations are in comoving coordinates. ",
        "conclusions": "\\label{conc} The low number density of SDSS LRGs allows us to partially separate LRG pairs occupying the same dark matter halo from pairs occupying distinct dark matter halos.  Candidate one-halo pairs are identified using simple cuts in the transverse and LOS separations.  We group these pairs using the FoF algorithm to compute the CiC group multiplicity function.  We measure the CiC group multiplicity function for the subsample of SDSS LRGs satisfying $-23.2 < M_{g} < -21.2$ and $ 0.16 < z < 0.36$, carefully accounting for the effects of fiber collisions and survey boundaries, holes, and incompleteness.\\\\  In order to derive HOD constraints from our measurement, we calibrated the relation between the CiC and true one-halo group multiplicity functions using mock LRG catalogs.  The variance about the mean relation is comparable to the Poisson sampling variance of the CiC multiplicity function and must be properly accounted for in the maximum likelihood parameter estimation.\\\\  The CiC group multiplicity function places strong constraints on the satellite LRG HOD, $N_{sat}(M)$.  When we fix $\\sigma_{log M} = 0.7$ and $\\bar{n}_{LRG} = 10^{-4}$ (Mpc/$h)^{-3}$, the maximum likelihood HOD parameters and marginalized one-dimensional 68\\% and 95\\% confidence intervals are $M_{cut} = 5.0^{+1.5}_{-1.3} (^{+2.9}_{-2.6}) \\times 10^{13} M_{\\sun}$, $M_1 = 4.95^{+0.37}_{-0.26} (^{+0.79}_{-0.53}) \\times 10^{14} M_{\\sun}$, and $\\alpha = 1.035^{+0.10}_{-0.17} (^{+0.24}_{-0.31})$, with $M_{min} = 8.05 \\times 10^{13} M_{\\sun}$ at the maximum likelihood point.  We tightly constrain the satellite fraction to $f_{sat} = 0.0636^{+0.0019}_{-0.0020} (^{+0.0038}_{-0.0039})$.  The projected correlation function $w_p(r_p)$ of mock catalogs derived from an SO halo catalog is an excellent agreement with the measurements of \\citet{masjedi/etal:2006} and \\citet{zehavi/etal:2005a} when the large scale clustering is used to fix $\\sigma_{log M}$.  Fig.~\\ref{fig:fofcompare} shows that FoF halo catalogs have a severe deficit of pairs at $\\sim 1$ Mpc/$h$.  In \\S~\\ref{fofkulksec} we point out that methods using $w_p(r_p)$ with an analytic estimate of the two-halo term calibrated using FoF halos will severely overestimate the number of satellites and the maximum expected one-halo group size.  Our measured CiC group multiplicity function rules out the best fit HOD from \\citet{kulkarni/etal:2007}.\\\\  Despite the increased complexity of our approach and necessary calibration using simulations, we have produced high quality mock catalogs which reproduce both higher order statistics in the density field and the features of the projected correlation function.  These mock catalogs will be used in a forthcoming paper to study the large scale structure statistics of our CiC groups.\\\\"
    },
    "0809/0809.1428_arXiv.txt": {
        "abstract": "{}% { We study the coronal magnetic field structure inside active regions and its temporal evolution. We attempt to compare the magnetic configuration of an active region in a very quiet period with that for the same region during a flare. } { Probably for the first time, we use vector magnetograph data from the Synoptic Optical Long-term Investigations of the Sun survey (SOLIS) to model the coronal magnetic field as a sequence of nonlinear force-free equilibria. We study the active region NOAA 10960 observed on 2007 June 7 with three snapshots taken during a small C1.0 flare of time cadence 10 minutes and six snapshots during a quiet period. } { The total magnetic energy in the active region was approximately $3 \\times 10^{25}$ J. Before the flare the free magnetic energy was about 5~\\% of the potential field energy. A part of this excess  energy was released during the flare, producing almost a potential configuration at the beginning of the quiet period. } { During the investigated period, the coronal magnetic energy was only a few percent higher than that of the potential field and consequently only a small C1.0 flare occurred. This was compared with an earlier investigated active region 10540, where the free magnetic energy was about 60~\\% higher than that of the potential field producing two M-class flares. However, the free magnetic energy accumulates before and is released during the flare which appears to be the case for both large and small flares. } ",
        "introduction": "Methods have been developed to extrapolate the observed photospheric magnetic field vector into the corona. Using the fact that the magnetic field is dominant in solar active regions (ARs), we are able to neglect non-magnetic forces and to assume that the coronal magnetic field is force-free. Different instruments provide photospheric vector magnetograph data, which are used as input to the extrapolation methods. These data have, however had a rather low time cadence. Data of high time cadence are required to investigate in detail, for example the different evolutionary stages of solar flares. Suitable data include magnetic field observations of the Sun provided by the SOLIS Vector-SpectroMagnetograph. With a time cadence of $\\approx$ 10 minutes, the instrument is designed to measure multiple area scans of ARs, which enables us for the first time to investigate the evolution of the coronal magnetic field energy with a high time cadence. Many existing studies deal with the extrapolation based on vector magnetograph data. For instance, \\cite{reg_07b} dealt with the photospheric vector magnetic field provided by the Mees Solar Observatory Imaging Vector Magnetograph, \\cite{wie_05} used spectropolarimetric data recorded with the Tenerife Infrared Polarimeter of the German Vacuum Tower Telescope, and \\cite{tha_08} performed extrapolations of Solar Flare Telescope Vector Magnetograph data. In all of these studies, only one snapshot was however used or, as in the last of the aforementioned studies, a sequence of vector magnetograms with a low time cadence of one magnetogram per day. Therefore, an improvement is achieved by applying our extrapolation technique to the high time cadence SOLIS/VSM data as described in the present study. ",
        "conclusions": "\\begin{figure}[ht]\\centering \\centerline{\\includegraphics[width=0.3\\textwidth]{0235fig4}} \\caption{Panels (a), (b), and (c) show the magnetic field configuration during the C1.0 flare. Panel (d) shows the minimum energy configuration. Shown are field lines of the potential (gray) and NLFF (black) field. For improved visibility, the $z$-axis is drawn elongated.} \\label{fig:fig4} \\end{figure} We have investigated the coronal magnetic field associated with the NOAA AR 10960 on 2007 June 7 by analyzing SOLIS/VSM data. Three vector magnetograms with a time cadence of  $\\approx$ 10 minutes were available to investigate the magnetic energy content of the coronal field during a C1.0 flare, and six further snapshots were acquired to analyze a very quiet time about three hours after the flare. Before as well as after the small flare, the magnetic field energy was $E_{nlff} \\approx 3 \\times 10^{25}$~J. The NLFF field had a free energy of $E_{free} \\approx 1.5 \\times 10^{24}$ J before the flare. As a consequence of the flare/CME, this free magnetic energy reduced by almost a factor of 10 and produced an almost potential configuration. Six snapshots acquired within a time period of about $70$ minutes, during a quiet period of 3 -- 4 hours after the flare, showed again an increase in the free magnetic energy. Since the estimated free magnetic energy remained only about 5~\\% of the total energy content, no large eruption was produced by AR 10960.\\\\ This is clearly different from the flaring of AR 10540 observed on 2004 January 18 -- 21, which was analyzed in a previous work with the help of vector magnetograph data from the Solar Flare Telescope in Japan of time cadence of about 1 day. In this AR, the free energy was $E_{free}\\approx$ 66~\\% of the total energy, which was sufficiently high to power a M6.1 flare \\citep[for details see][]{tha_08}. The activity of AR 10540 investigated earlier was significantly higher than for the data analyzed in the current paper, as was the total magnetic energy. However, despite these differences, we also found some common features. Magnetic energy accumulates before the flare and a significant part of the excess energy is released during the flare. The high amount of free magnetic energy available in AR 10540 produced M-class flares, while the relatively small amount of free energy in AR 10960 powered only a small C-class flare. In both cases, all three components of the vector magnetogram changed during the flare, but the energy decrease in the NLFF field was always higher than that of the potential field, i.e. the energy release was more related to the change in the transverse magnetic field components -- which correspond to the field aligned electric currents in the corona -- than to that of the longitudinal component."
    },
    "0809/0809.1875_arXiv.txt": {
        "abstract": "{Optical interferometry is a powerful tool for observing the intensity structure and angular diameter of stars. When combined with spectroscopy and/or spectrophotometry, interferometry provides a powerful constraint for model stellar atmospheres.} {The purpose of this work is to test the robustness of the spherically symmetric version of the \\textsc{Atlas} stellar atmosphere program, \\textsc{SAtlas}, using interferometric and spectrophotometric observations.} {Cubes (three dimensional grids) of model stellar atmospheres, with dimensions of luminosity, mass, and radius,  are computed to fit observations for three evolved giant stars, $\\psi$ Phoenicis, $\\gamma$ Sagittae, and $\\alpha$ Ceti. The best--fit parameters are compared with previous results.} {The best--fit angular diameters and values of $\\chi^2$ are consistent with predictions using \\textsc{Phoenix} and plane--parallel \\textsc{Atlas} models.   The predicted effective temperatures, using \\textsc{SAtlas},  are about $100$ to $200$ $\\rm{K}$ lower, and the predicted luminosities are also lower due to the differences in effective temperatures.} {It is shown that the \\textsc{SAtlas} program is a robust tool for computing models of extended stellar atmospheres that are consistent with observations.  The best--fit parameters are consistent with predictions using \\textsc{Phoenix} models, and the fit to the interferometric data for $\\psi$ Phe differs slightly, although both agree within the uncertainty of the interferometric observations.} ",
        "introduction": "Optical interferometry accurately measures the combination of the angular diameter and the structure of the intensity distribution of a stellar disk.  The combination of interferometry with spectroscopy and/or spectrophotometry is a powerful tool for determining effective temperatures and other properties of stars, and the growth of optical interferometry is testing theoretical models of stellar atmospheres in new ways.  In a series of articles, \\cite{Wittkowski2004, Wittkowski2006b, Wittkowski2006a} fit Very Large Telescope Interferometer (VLTI) observations of cool giants using the VINCI instrument with spherically symmetric stellar atmosphere models in local thermodynamic equilibrium (LTE) computed with the \\textsc{Phoenix} program \\citep{Hauschildt1999}, and with plane parallel models from \\textsc{Phoenix} and \\textsc{Atlas} \\citep{Kurucz1970, Kurucz1993}.  The authors demonstrated that optical interferometry can detect the wavelength--dependent limb--darkening of cool giants and that the limb--darkening can be used to constrain model stellar atmospheres.  These works are based on observations of three stars: $\\psi$ Phoenicis (M4 III), $\\gamma$ Sagittae (M0 III), and $\\alpha$ Ceti (M1.5 III) also called Menkar.  The purpose of this note is to model these interferometric observations using the new spherically symmetric version of the \\textsc{Atlas} program (\\textsc{SAtlas}) developed by \\cite{Lester2008} and to determine fundamental parameters of the three stars.  There are two versions of the \\textsc{SAtlas} program, one using opacity distribution functions  while the other uses opacity sampling; in this work we use the first version.  The radiation field is computed using the method suggested by \\cite{Rybicki1971}, which is a reorganization of the \\cite{Feautrier1964} method.  In this work we compute cubes (three dimensional grids) of spherically symmetric stellar atmospheres with dimensions of luminosity, mass, and radius to model interferometric visibilities to fit the VLTI/VINCI observations and broadband spectrophotometry from \\cite{Johnson1975}.  This will provide a robust test of the program and a comparison to the results predicted using spherically symmetric \\textsc{Phoenix} models.  In the next section, we outline the method for computing synthetic visibilities and predicting radii, effective temperatures, luminosities, masses, and gravities.  We present the results in the third section and the discussion in the fourth section. ",
        "conclusions": "The purpose of this work is to test the spherical version of the \\textsc{Atlas} program. By fitting interferometric and spectrophotometric observations using model stellar atmospheres, we derive the properties of $\\psi$ Phe, $\\gamma$ Sge, and $\\alpha$ Cet.  The derived parameters are compared to earlier results.  The  angular diameters that are determined by fitting model stellar atmospheres to interferometric observations are the same, except for $\\alpha$ Cet, for which we predict an angular diameter that is $0.1$ $\\rm{mas}$ smaller. The differences in predicted angular diameter are likely due to differences in limb--darkening of the model atmospheres generated using \\textsc{SAtlas} and \\textsc{Phoenix}.  Also our minimum $\\chi^2$ value of the fit of the limb--darkened angular diameter for $\\psi$ Phe, $1.66$, is smaller than the predicted value of $\\chi^2$ from the spherically symmetric \\textsc{Phoenix} models, $1.8$.  The \\textsc{SAtlas} models for $\\psi$ Phe that best fit the \\emph{interferometric} data have effective temperatures in the range of $4000$ to $4500$ $\\rm{K}$, while the \\textsc{Phoenix} models from \\cite{Wittkowski2004} are all $3550$ and $3600$ $\\rm{K}$.  The fit of the \\textsc{SAtlas} models to \\emph{spectrophotometric} observations predict an effective temperature of $3415$ $\\rm{K}$ conflicting with the prediction of a higher temperature. For the effective temperature range of $3400$ to $3800$ $\\rm{K}$, the minimum $\\chi^2$ values predicted by fitting \\textsc{SAtlas} models are $1.70$ to $1.72$, similar to the values found using plane--parallel \\textsc{Atlas} models.  The fit of the \\textsc{SAtlas} models to broadband spectrophotometric observations are used to determine the effective temperatures of the three stars because spectrophotometry is much more sensitive to the effective temperature than interferometry. The predicted effective temperatures are smaller than those found in the previous works, but this difference is due to the methods used for determining the effective temperature, \\emph{not} differences in the stellar atmosphere programs. If we use the method from \\cite{Wittkowski2004} where the bolometric flux $f_{\\rm{bol}}$ is calculated by integrating the spectrophotometric data and then $T_{\\rm{eff}}^4 = 4f_{\\rm{bol}}/(\\sigma \\theta_{\\rm{Ross}}^2)$, we would predict similar effective temperatures.  Any variation in that case would be due to differences in the calculation of the bolometric flux and differences in the angular diameter.  This suggests the effective temperatures here would differ by at most $1\\%$.   Our results also differ because of the spectrophotometric data used.  The  difference is smallest for $\\psi$ Phe because the \\textsc{Phoenix} and \\textsc{SAtlas} fits both use the same spectrophotometric data for the fitting and agree within the uncertainty.  For $\\gamma$ Sge, \\cite{Wittkowski2006a} complement the \\cite{Johnson1975} data with narrow--band data from \\cite{Alekseeva1997} while for $\\alpha$ Cet, \\cite{Wittkowski2006b} use optical \\citep{Glushneva1998b, Glushneva1998a} and infrared \\citep{Cohen1996} spectrophotometry.  The effective temperature difference for the remaining two stars are about $150$ to $200$ $\\rm{K}$.  Because the angular diameters are similar, the predicted radii are also similar for both fits with \\textsc{SAtlas} and \\textsc{Phoenix}.  Therefore the differences between predicted luminosities are due to differences in effective temperatures and thus due to the different methods for determining the effective temperatures.  The lower effective temperatures and luminosities imply lower masses when compared to \\cite{Girardi2000} evolutionary tracks and smaller gravities.  The \\textsc{SAtlas} models are consistent with previous results. Fitting the models to interferometric and spectrophotometric observations have provided a robust test of the spherically symmetric version of the \\textsc{Atlas} program and have shown that the \\textsc{SAtlas} program is a powerful tool for studies of stellar atmospheres."
    },
    "0809/0809.3088_arXiv.txt": {
        "abstract": "Based on a sample of 355 quasars with significant optical polarization, we found that quasar polarization vectors are not randomly oriented over the sky as naturally expected. The probability that the observed distribution of polarization angles is due to chance is lower than 0.1\\%. The polarization vectors of the light from quasars are aligned although the sources span huge regions of the sky ($\\sim$ 1~Gpc). Groups of quasars located along similar lines of sight but at different redshifts (typically z $\\sim$ 0.5 and z $\\sim$ 1.5) are characterized by different preferred directions of polarization. These characteristics make the observed alignment effect difficult to explain in terms of a local contamination by interstellar polarization in our Galaxy. Interpreted in terms of a cosmological-size effect, we show that the dichroism and birefringence predicted by a mixing between photons and very light pseudoscalar particles within a magnetic field can qualitatively reproduce the observations. We find that circular polarization measurements could help constrain this mechanism. ",
        "introduction": "Large-scale alignments of quasar polarization vectors were first uncovered looking at a sample of 170 quasars selected from the literature (Hutsem\\'ekers 1998, hereafter Paper I). The presence of such alignments was confirmed later on a larger sample (Hutsem\\'ekers \\& Lamy 2001, hereafter Paper II). The departure from random orientations was found at significance levels small enough to merit further investigation. Moreover, these alignments seemed to come from high redshift regions, implying that the underlying mechanism might cover physical distances of gigaparsecs. A large survey of linear polarization was then started, with the goal to characterize better the polarization properties of quasars and to investigate the reality of the alignments. A final sample of 355 quasars with reliable polarization measurements was then built on the basis of the new polarization measurements and of a comprehensive compilation from the literature. A detailed analysis of this sample was carried out by Hutsem\\'ekers et al. (2005, hereafter Paper III). The main results are reviewed here and a possible interpretation based on photon-pseudoscalar mixing is presented. ",
        "conclusions": ""
    },
    "0809/0809.0568_arXiv.txt": {
        "abstract": "Distances of galaxies in the Hubble Space Telescope Key Project are based on the Cepheid period-luminosity relation. An alternative basis is the tip of the red giant branch. Using archival HST data, we calibrate the infrared Tully-Fisher relation using 14 galaxies with tip of the red giant branch measurements. Compared with the Key Project, a higher value of the Hubble Constant by 10\\% $\\pm$ 7\\% is inferred. Within the errors the two distance scales are therefore consistent. We describe the additional data required for a conclusive tip of the red giant branch measurement of H$_0$. ",
        "introduction": "The extragalactic distance scale based on the Cepheid period-luminosity (PL) relation and secondary distance indicators, such as the Tully-Fisher relation, the supernova standard candle \\citep{gib00}, surface brightness fluctuations, and the fundamental plane \\citep{{fr01},{mo00}} has been criticized recently \\citep{{str08},{str07}} on the grounds that the PL relation may not be unique. Indeed, the finite width of the Cepheid instability strip in the HR diagram implies that nuisance parameters such as metallity and star formation history may play a role in determining the PL relation. Metallicity was considered as a second parameter by \\cite{fr01}, \\cite{sa04}, and \\cite{ma06}. \\cite{rom} have reviewed the situation and concluded that the Cepheid PL relation is not universal.  It is of interest, therefore, to see how well the distance scale can be measured without reference to Cepheids at all. In this Letter we use the tip of the red giant branch (TRGB) distance indicator to calibrate the Tully-Fisher relation. The TRGB is a good standard candle because it results from the helium flash on the red giant branch, which theory suggests is relatively immune to metallicity effects in old stellar populations. ",
        "conclusions": "\\cite{sa00} obtained H$_0$ = 71 $\\pm$ 4 $\\pm$ 7 km s$^{-1}$ Mpc$^{-1}$ from their multiwavelength Cepheid-based Tully-Fisher calibration. The largest term in the 7 km s$^{-1}$ Mpc$^{-1}$ systematic error is due to the distance of the Large Magellanic Cloud. The largest term in the absolute calibration of the TRGB (population II) distance scale is the uncertainty in M$_{I,TRGB}$ = --4.05 $\\pm$ 0.02 mag \\citep{ri07}, associated with the absolute magnitude of the horizontal branch.  Our principal finding is that, within the 1.5$\\sigma$ uncertainty, the mean difference of the distance moduli derived from Cepheids and from the TRGB magnitude for our sample of 14 galaxies is consistent with zero.  In addition, we conclude that the further steps to a more accurate Cepheid-independent value of H$_0$ are (i) a larger sample of TRGB distances to galaxies which calibrate secondary distance indicators, (ii) multiwavelength photometry of these galaxies, and (iii) TRGB calibration of type Ia supernovae \\citep{tsr08}, surface brightness fluctuations, and the fundamental plane."
    },
    "0809/0809.0797_arXiv.txt": {
        "abstract": "The plasmoid-induced-reconnection model explaining solar flares based on bursty reconnection produced by an ejecting plasmoid suggests a possible relation between the ejection velocity of a plasmoid and the rate of magnetic reconnection. In this study, we focus on the quantitative description of this relation. We performed magnetohydrodynamic (MHD) simulations of solar flares by changing the values of resistivity and the plasmoid velocity. The plasmoid velocity has been changed by applying an additional force to the plasmoid to see how the plasmoid velocity affects the reconnection rate. An important result is that the reconnection rate has a positive correlation with the plasmoid velocity, which is consistent with the plasmoid-induced-reconnection model for solar flares. We also discuss an observational result supporting this positive correlation. ",
        "introduction": "Magnetic reconnection is a process in which the magnetic energy is converted into the kinetic and thermal energy \\citep{1958IAUS....6..123S, 1963ApJS....8..177P, 1964psf..conf..425P}, and it has been widely believed to play a fundamental role in causing solar flares. A model for flares based on magnetic reconnection has been developed since 1960s by \\citet{1964psf..conf..451C}, \\citet{1966Natur.211..697S}, \\citet{1974SoPh...34..323H}, and \\citet{1976SoPh...50...85K}, so this model has been called the CSHKP model. The observations supporting this model has been reported, such as cusps \\citep{1992PASJ...44L..63T}, arcades \\citep{1992PASJ...44L.211T, 1992PASJ...44L.205M, 2002GeoRL..29u..10I}, loop top hard X-ray (HXR) sources \\citep{1994Natur.371..495M,2003ApJ...596L.251S}, X-ray jets \\citep{1992PASJ...44L.173S, 1996PASJ...48..123S}, and so on. Furthermore Yohkoh \\citep{1991SoPh..136....1O} discovered the common property of flares with different appearances, such as Long Duration Events (LDE flares) and impulsive flares (see \\citealt{1999Ap&SS.264..129S} and \\citealt{2002SSRv..101....1A} for review).  So far the dynamic process caused by magnetic reconnection in flares has been widely studied by MHD simulations \\citep[e.g.][]{1983SoPh...84..169F, 1990JGR....9511919F, 1991SoPh..135..361F, 1996ApJ...466.1054M, 1996PhPl....3.4172U, 2000JGR...105.2375L, 2001ApJ...549.1160Y}. Most of these works have been focused on MHD processes producing apparent features of flares such as cusp shaped loops, ejecting plasmoid (blob of plasma), inflows and loop-top HXR sources. However, these works have not clearly explained the fundamental physical process, that is what determines the energy release rate in flares. This question is related to the rate of magnetic reconnection, so to identify the condition under which fast magnetic reconnection \\citep{1977JPlPh..17..337U} operates is important.    Observationally, solar flares are often associated with plasmoid ejections. \\citet{1995ApJ...451L..83S} found that 8 impulsive flares on the limb (Masuda-type; \\citealt{1995PASJ...47..677M}) were associated with plasmoid ejections. \\citet{1997PASJ...49..249O, 1998ApJ...499..934O} carefully analyzed plasmoid ejections in impulsive flares and found that plasmoids undergo strong acceleration during the impulsive phase of these flares. \\citet{2004ApJ...616..578S} found the tiny two-ribbon flare driven by emerging flux accompanying the miniature filament eruption (=plasmoid). This feature was also found in other observations \\citep{2001ApJ...559..452Z,2004ApJ...613..592T,2005ApJ...630.1148S} and numerical simulations \\citep{1997ApJ...487..437M}. It is also found that there is a positive correlation between the plasmoid velocity and the reconnection rate \\citep{1995ApJ...451L..83S, 2005ApJ...634L.121Q, shimizu2008}. There are other literatures discussing this topic \\citep{1999ApJ...525L..57N,2000JGR...10523153F, Klimchuk2001, 2002A&ARv..10..313P,2003NewAR..47...53L,2005ApJ...635.1291K}. \\citet{2000JGR...105.2375L} derived an analytic relation between the acceleration of coronal mass ejections (CMEs) and reconnection rate.  The observations above show an important suggestion that plasmoid ejection plays a key role in causing fast magnetic reconnection. Based on these observational results, \\citet{1996AdSpR..17....9S, 1997cswn.conf..103S} extended the classical CSHKP model and proposed the plasmoid-induced reconnection model. In this model, a plasmoid (or flux rope in a 3D situation) is created in the anti-parallel magnetic field by the magnetic reconnection (Figure \\ref{figure:pir}a, b, c). Then the plasmoid situates in a current sheet inhibits inflows into the sheet, so reconnection is inefficient and magnetic energy is stored (Figure \\ref{figure:pir}d). Then the plasmoid starts to move at the velocity $v_{plasmoid}$, inflows toward the X-point ($v_{inflow}$) are induced following mass conservation and reconnection starts (Figure \\ref{figure:pir}e). If we assume the incompressibility, the mass flux into the reconnection region is given by $\\sim v_{inflow} L_{inflow}$ and the mass flux ejected by the plasmoid  is $\\sim v_{plasmoid} W_{plasmoid}$, and these are balanced, where $W_{plasmoid}$ is the typical width of the plasmoid, and $L_{inflow}$ $(\\ge W_{plasmoid})$ is the typical length of the inflow region. Consequently the induced inflow speed can be estimated as follows: \\begin{equation} v_{inflow} \\sim v_{plasmoid} W_{plasmoid} / L_{inflow}. \\end{equation} Since the reconnection rate is determined by the speed of the inflows, fast reconnection becomes possible when plasmoid ejection is fast. Moreover the jet from a reconnection point accelerate the plasmoid \\footnote{Note that in addition to the acceleration by the reconnection jet, the magnetic pressure gradient force can also accelerate the plasmoid if there is a global magnetic pressure gradient around the plasmoid. See more detailed discussion in Appendix.}, so the fast reconnection further drives fast plasmoid ejection. This suggests a positive correlation between the plasmoid velocity and the reconnection rate. The merit of this model is to provide us with a unified view for understanding various types of flares with different spatial sizes and timescales, and it can naturally cause fast reconnection \\citep{1986JGR....91.5579P} in a current sheet where many plasmoids with different sizes are created by tearing instability \\citep{2000A&A...360..715K,2001EP&S...53..473S,2004A&A...417..325K}.   \\begin{figure} \\plotone{f1a.eps}  \\caption{Schematic picture of the plasmoid-induced reconnection model. Solid lines indicate magnetic field lines. Panel (a), (b) and (c) shows the process creating the plasmoid in the anti-parallel magnetic field by the magnetic reconnection. Panel (d) shows how the plasmoid in the current sheet inhibits the magnetic reconnection  in the reconnection region shown by a light gray area in panel (e). Thick black arrows indicate inflows to a reconnection point, and thick white arrows indicate reconnection jets. Panel (e) explains how the plasmoid ejection can induce strong inflows into the reconnection region.} \\label{figure:pir} \\end{figure}  \\setcounter{figure}{0} \\begin{figure} \\plotone{f1b.eps} \\caption{cont.} \\end{figure}  So far there have been no MHD simulations performed to examine the plasmoid-induced-reconnection model, which should reproduce the correlation between the plasmoid velocity and the reconnection rate. In this study, we performed a series of these MHD simulations by changing the parameters related to resistivity and plasmoid velocity to investigate the relation between these two quantities.  The model and the numerical method are described in section \\ref{section:model}, and the numerical results are presented in section \\ref{section:results}. The discussion is given in section \\ref{section:discussion} and the conclusions are given in section \\ref{section:conclusion}. ",
        "conclusions": "\\label{section:conclusion}  In this paper, we performed several MHD simulations to examine the basic physical relation between the plasmoid velocity and reconnection rate in the context of the plasmoid-induced-reconnection model for solar flares. The initial magnetic configuration, resistivity, and the additional force are parameters in these simulations. When we changed the amplitude of resistivity (case A), the reconnection rate and the plasmoid velocity changed, showing a positive correlation. We showed that the reconnection rate (i.e. inflow speed) and the plasmoid velocity are closely related to each other. This result is consistent with observations \\citep{1995ApJ...451L..83S, 2005ApJ...634L.121Q, shimizu2008} supporting the plasmoid-induced-reconnection model of impulsive flares."
    },
    "0809/0809.2337_arXiv.txt": {
        "abstract": "{Using the IRAM 30m telescope, we have detected the $^{12}$CO $J=2-1$, $4-3$, $5-4$, and $6-5$ emission lines in the millimeter-bright, blank-field selected AGN COSMOS J100038+020822 at redshift $z=1.8275$. The sub-local thermodynamic equilibrium (LTE) excitation of the $J=4$ level implies that the gas is less excited than that in typical nearby starburst galaxies such as NGC253, and in the high-redshift quasars studied to date, such as J1148+5251 or BR1202-0725. Large velocity gradient (LVG) modeling of the CO line spectral energy distribution (CO SED; flux density vs. rotational quantum number) yields H$_{2}$ densities in the range $10^{3.5}-10^{4.0}$ cm$^{-3}$, and kinetic temperatures between 50 K and 200 K. The H$_{2}$ mass of $(3.6 - 5.4) \\times 10^{10}$ M$_{\\sun}$ implied by the line intensities compares well with our estimate of the dynamical mass within the inner 1.5 kpc of the object. Fitting a two-component gray body spectrum, we find a dust mass of $1.2 \\times 10^{9}$ M$_{\\sun}$, and cold and hot dust temperatures of 42$\\pm$5 K and 160$\\pm$25 K, respectively. The broad MgII line allows us to estimate the mass of the central black hole as $1.7 \\times 10^{9}$ M$_{\\sun}$. Although the optical spectrum and multi-wavelength SED matches those of an average QSO, the molecular gas content and dust properties resemble those of known submillimeter galaxies (SMGs).   The optical morphology of this source shows tidal tails that suggest a recent interaction or merger. Since it shares properties of both starburst and AGN, this object appears to be in a transition from a strongly starforming submillimeter galaxy to a QSO. }{} ",
        "introduction": "\\label{introduction} Submillimeter blank field surveys have discovered a population of dust enshrouded high-redshift galaxies \\citep{Smail1997,Hughes1998,Eales1999}, which are massive systems with huge molecular gas reservoirs and star formation at high rates \\citep{Neri2003,Greve2005,Solomon2005}. Their faint X-ray emission \\citep{Alexander2005} suggests that most of the submillimeter output is not powered by active galactic nuclei (AGN), but by massive star formation. These starbursting submillimeter galaxies (SMGs) account for a substancial fraction of the far-infrared (IR) background, and current models suggest this population may represent the formation of massive spheroidals at high-redshift \\citep{Dunlop2001}.  In the local Universe, an evolutionary connection between starbursting ultraluminous infrared galaxies (ULIRGs) and QSOs has been suggested \\citep{Sanders1988a} and discussed controversially \\citep[e.g.][]{Sanders1988a,Sanders1988b, Genzel1998, Tacconi2002}. This evolutionary cycle has found support from hydrodynamical simulations of galaxy formation \\citep[e.g.][]{DiMatteo2005, Hopkins2005, Hopkins2006} and observations \\citep{Sanders1988a,Page2004,Stevens2005}.  The ubiquitous presence of AGN activity in most SMGs \\citep{Alexander2005} suggests a close link between the AGN and starburst activity at redshifts $z \\sim 1-3$. If a QSO is found to be far-IR luminous, it means that large amounts of gas and dust should still be present. Such a source may, in fact, be a good candidate for an object in the transition from a starburst to a QSO, in particular when it shows absorbed X-ray emission \\citep{Page2004, Stevens2005}. Yet only a few high-redshift composite starburst/AGN have been studied in its molecular and multi-wavelength continuum emission \\citep{RowanRobinson2000, leFloch2007, Coppin2008}.  Observations of carbon monoxide (CO) in galaxies are important probes of the physical conditions of the cold and warm molecular gas in the galactic nuclei and disks. They provide estimates of the total amount of gas available to fuel starburst and/or AGN activity, and the CO line profile and intensity can be used to obtain important information about the galaxy kinematics, such as dynamical mass or size of the emitting region \\citep{Solomon1997,Solomon2005}.  Because of their diagnostic value, great efforts have been made to observe CO emission lines in SMGs. These studies have benefited from deep radio continuum imaging (e.g. VLA 1.4 GHz) to locate the SMG accurately \\citep{Ivison2002,Ivison2005, Ivison2007}, and from the determination of optical spectroscopic redshifts \\citep[e.g.][]{Chapman2005}. Due the small spectroscopic bandwidths of current (sub)millimeter telescopes and interferometers, this was necessary to permit the proper frequency tuning for line observations (e.g. of CO, [CI], [CII]). Yet only 19 SMGs have been reported in CO so far \\citep[][]{Frayer1998, Frayer1999, Andreani2000, Sheth2004, Hainline2004, Neri2003, Greve2005, Frayer2008, Coppin2008, Tacconi2008} and only a few have been observed in multiple molecular transitions to study the excitation conditions of their molecular gas reservoir \\citep{Solomon2005}. Spatial structure and dynamics were studied for some SMGs through high resolution CO imaging \\citep{Tacconi2006, Tacconi2008}. Overall, only $\\sim 50$ high-redshift ($z>1$) objects have been detected in CO, most of which are luminous, optically selected QSOs \\citep[][]{Omont1996, Guilloteau1997, Guilloteau1999, Carilli2002, Walter2003, Bertoldi2003, Beelen2004, Riechers2006, Carilli2007} and high-redshift radio galaxies \\citep[HzRG;][]{DeBreuck2003, DeBreuck2005, Klamer2005, Papadopoulos2005}.   \\begin{figure*}[!ht] \\centering \\includegraphics[width=15cm]{all_lines.ps} \\caption{Observed spectra of the CO $J=2-1, 4-3, 5-4$ and $6-5$ emission lines in J100038+020822, obtained with the IRAM 30m telescope. Gaussian fits to the spectra are shown as dotted lines. The velocity reference is at $z=1.8275$.} \\label{fig:lines} \\end{figure*}  Here we report the detection of CO $2-1$, $4-3$, $5-4$, and $6-5$ line emission from a millimeter selected QSO, J$100038.01+020822.4$. Its CO line intensities and flux density ratios allow us to estimate its molecular gas content and its excitation conditions. Its spectral properties suggest that this object may be evolving from a starburst to a QSO.  After a description of the observations in Section 2, we present the molecular gas and dust properties of our source in Sections \\ref{sect:lvg} to \\ref{sect:co_prop2}. We study its morphology and multi-wavelength properties in Sections \\ref{sect:optical_morph}, \\ref{sect:sed} and \\ref{sect:bhmass}. We discuss the results in Section 4 and give a brief summary in Section 5. Throughout, we use a $\\Lambda$CDM cosmology, $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{\\mathrm{M}}=0.3$. ",
        "conclusions": "\\label{sect:discussion} \\subsection{Comparison of Excitation Conditions}  The observations of the CO lines in J100038 show that its molecular gas is somewhat less excited than observed in local starburst or high-redshift QSOs. We find that the turn-over of the CO line SED occurs between the CO\\ $5-4$ and CO\\ $6-5$ transitions whereas in most local starbursts/AGNs, or high-redshift QSOs studied to date, the peak of the CO line SED is typically located between the CO\\ $6-5$ and CO\\ $7-6$ transitions. Examples of such higher excitation CO emission in the local Universe are NGC253 \\citep{Guesten2006}, M82 \\citep{Weiss2005a}, and at high-redshift, J1148+5251 \\citep{Bertoldi2003,Walter2003}, BR1202-0725 \\citep{Carilli2002, Riechers2006} and APM08279+5255 \\citep{Weiss2007}. The only SMG for which the CO SED has been traced over its peak is J16359+6612 \\citep{Weiss2005b}, which interestingly, shows a similar CO excitation to J100038. Other examples of lower excitation are found toward the centers of the local starburst/AGN of Circinus and NGC4945 \\citep{Hitschfeld2008}, or in the main starburst region of the Antennae galaxy \\citep{Zhu2003}. The lower excitation of the molecular gas observed in these cases, and in particular towards J100038, is likely produced by the relatively low H$_{2}$ density values ($n(\\mathrm{H}_{2})<10^{4}$ cm$^{3}$). Furthermore, the best solution derived from the LVG analysis indicates a moderate kinetic temperature ($\\sim95$ K), suggesting that the AGN does not contribute strongly to the heating of the molecular gas. Conversely, this may imply that most of the heating is produced by star formation, as it is also suggested by the similarity of molecular gas conditions ($T_{\\mathrm{kin}}$, $n$(H$_{2}$)) observed between J16359+6612 and J100038.   \\begin{table} \\caption{Summary of derived properties of J100038+020822}             % \\label{table:3}      % \\centering% \\begin{tabular}{l l c}        % \\hline\\hline Property & Name & Value \\\\% inserts double horizontal lines \\hline $M(\\mathrm{H}_{2})$ &Molecular Gas Mass ($10^{10}\\ M_{\\sun})$ &  $3.6 - 5.4$  \\\\ $M_{\\mathrm{dyn}}$  &Dynamical Mass ($10^{10}\\ \\mathrm{sin}^{-2}(i)\\ M_{\\sun}$)& $3.0$  \\\\ $M_{\\mathrm{dust}}$             &Dust Mass ($10^{9}\\ M_{\\sun}$)  &  $1.2$\\\\ SFR                 &Star Formation Rate ($10^{3}\\ M_{\\sun}$ yr$^{-1}$) & $1.7 $  \\\\ SFE                 &Star Formation Efficiency ($L_{\\sun}\\ M_{\\sun}^{-1}$) & $190 $    \\\\ % $\\tau_{\\mathrm{SF}}$&Gas Depletion Lifetime (Myr) & 30 \\\\ $M_{\\mathrm{V}}$             &Optical Absolute Magnitude &  -24.2 \\\\ $M_{\\mathrm{BH}}$   &Black Hole Mass ($10^{9}\\ M_{\\sun}$)& $1.7$\\\\ \\hline                                   %   \\end{tabular} \\end{table}   \\subsection{J100038+020822: A Starburst-QSO Composite} The general picture for QSO and stellar spheroid formation is based on the merger of two gas-rich disk galaxies. The large amounts of gas and dust involved in these mergers provide the fuel for infrared luminous starbursts to occur, and the gas inflow into the inner regions of the galaxy likely feeds the central massive black hole \\citep{Mihos1994,Barnes1996}. Observations of local ULIRGs indicate that the most luminous phase occurs close to the final stage of the merger, when both galaxy disks overlap \\citep{Sanders1988a,Sanders1988b,Veilleux1999}.  When the central massive black hole has reached a sufficient size and luminosity, feedback from an active nucleus phase dissipates the dust \\citep{DiMatteo2005} and an optically luminous QSO emerges. As the gas is being expelled, the feeding of the QSO ceases, stopping its activity, and the system relaxes into a spheroidal galaxy hosting a (quiescent) supermassive black hole at its center \\citep{Hopkins2006}.  One important question in galaxy evolution is whether this scenario connecting gas-rich starburst galaxies and optically bright QSOs constitutes a rule for most of these systems or whether it is a phenomenon that applies only to the most luminous and massive objects. This evolutionary connection has been studied with vast supporting evidence in the local Universe \\citep[e.g.][]{Sanders1988a}. However, for the crucial epoch of star formation and QSO activity ($1<z<3$) only a few studies have so far been undertaken. The study of submillimetre \\citep{Page2004, Stevens2005} and CO line \\citep{Coppin2008} emission of absorbed and unabsorbed high-redshift QSOs indicates that the QSO phase could be preceded by a SMG starburst phase and suggests that submillimetre selected QSOs represent ideal objects for studying these transitional cases at $z>1$.  J100038 appears to be in a transition between the starburst and QSO phases, as a comparison of its properties to those of a sample of typical SMGs \\citep{Greve2005} and QSOs \\citep[][]{Solomon2005} shows. We start by describing its QSO properties followed by the features that classify this source as a starburst galaxy.  J100038 shows an optical spectrum typical of a Type-1 AGN (Figure \\ref{fig:spectra}) with a prominent MgII broad emission line. It is relatively luminous at optical wavelengths ($M_{\\mathrm{V}}\\sim-24$ magnitudes, extinction corrected), shows mild optical extinction with $A_{\\mathrm{V}}\\sim1$ and its X-ray emission indicates it is a heavily absorbed QSO (log$N(\\mathrm{H})=22-23$ cm$^{-2}$). Moreover, its optical to mid-IR SED resembles that of typical QSOs (Figure \\ref{fig:sed}), well in agreement with \\citet{Elvis1994} templates. In addition, we derived a diameter of $\\sim1.5$ kpc for the CO line emitting region. This size is consistent with values observed in high-redshift QSOs which range between 1 and 3 kpc \\citep{Walter2004, Solomon2005, Maiolino2007} and similar to those found in local ULIRGs  \\citep[$\\sim 0.5$ kpc;][]{Downes1998, Soifer2000}, but smaller than the diameters for the CO emitting region in SMGs \\citep[$\\lesssim4$ kpc;][]{Tacconi2006}.  On the other hand, J100038 shows distinctive features of on-going starburst activity. The optical morphology of J100038 seen in the HST imaging is suggestive of a recent merger event (Figure \\ref{fig:closeup_i}). While most of the emission is concentrated in the point-like central source ($C$), the faint emission from the eastern source ($E$) may be indicative of a tidal tail, hinting at a past interaction or merger, although it could also imply the influence of gravitational lensing. The low number density of HST $I$-band sources in the surroundings of J100038 , $\\rho(I<25.5)=62$ arcmin$^{-2}$, implies that the probability of chance association between $E$ and $C$ is only 4.1\\%. If both sources ($E$ and $C$) are physically related, the projected distance between them would be 7.5 kpc. A large fraction of the optical emission from the host galaxy in J100038 could still be absorbed by surrounding dust (Jahnke et al, in prep.), and could be hiding the actual link to the eastern component. We note that the optical emission of the host galaxy is consistent with the obscured SED of the prototypical local starburst galaxy Arp220. Indeed, the optical morphology of this system is strikingly similar to that observed in local ULIRG/PG-QSOs \\citep[IRAS 05189-2524;][]{Sanders1988b, Surace1998}. Furthermore, the far-IR to radio SED of J100038 agrees very well with that of Arp220 (Figure \\ref{fig:sed}), and does not follow the typical behavior of radio-loud or radio-quiet QSOs at these wavelengths. In fact, its radio to far-IR spectral index, when used as a redshift indicator results in $z=1.9$ \\citep{Bertoldi2007}, a strong indication that the far-IR and radio emission are both produced by star formation. As mentioned before, the analysis of the molecular gas physical conditions suggest that the AGN plays a moderate role in the gas heating and hints that the heating may actually be dominated by starforming regions.  All the evidence exposed permits to well fit J100038 in the Sanders et al. scenario. Studies of these transitional cases in the local Universe have found that about $30-50\\%$ of the most luminous ULIRGs ($L_{\\mathrm{FIR}}>10^{12.3}\\ L_{\\sun}$) show broad emission lines \\citep{Veilleux1999} and the presence of compact nuclei in 40\\% of all ULIRGs \\citep{Scoville2000,Soifer2000} as well as large molecular gas reservoirs \\citep{Evans2002, Evans2005}, even in the late stages of this evolutionary sequence. All these properties apply to J100038. Furthermore, the presence of both starburst and QSO attributes imprinted in the galaxy SEDs have largely been observed in nearby objects, constituting the basis of this evolutionary scenario \\citep{Sanders1988a, Sanders1988b, Sanders1989}. These templates compare favorably with the SED observed for J100038, strongly suggesting that this object is evolving from a starburst to a QSO. In fact, it is possible to classify it in a stage between the ``warm ULIRG'' and the ``infrared excess'' QSO phases following the mentioned scenario \\citep{Sanders2004}."
    },
    "0809/0809.1664_arXiv.txt": {
        "abstract": "We present ground-based observations of the transiting Neptune-mass planet Gl 436b obtained with the 3.5-meter telescope at Apache Point Observatory and other supporting telescopes. Included in this is an observed transit in early 2005, over two years before the earliest reported transit detection. We have compiled all available transit data to date and perform a uniform modeling using the JKTEBOP code. We do not detect any transit timing variations of amplitude greater than $\\sim$1 minute over the $\\sim$3.3 year baseline. We do however find possible evidence for a self-consistent trend of increasing orbital inclination, transit width, and transit depth, which supports the supposition that Gl 436b is being perturbed by another planet of $\\lesssim$ 12 M$_{\\earth}$ in a non-resonant orbit. ",
        "introduction": "Gliese 436 is an M-dwarf (M2.5V) with a mass of 0.45 M$_{\\sun}$ and hosts the extrasolar planet Gl 436b, which is currently the least massive transiting planet with a mass of 23.17 M$_{\\earth}$ \\citep{Torres07}, and the only planet known to transit an M dwarf. Gl 436b was first discovered via radial-velocity (RV) variations by \\citet{Butler04}, who also searched for a photometric transit, but failed to detect any signal greater than 0.4\\%. It was thus a surprise when \\citet{Gillon07b} reported the detection of a transit with a depth of 0.7\\%, implying a planetary radius of 4.22 R$_{\\earth}$ \\citep{Torres07} and thus a composition similar to Uranus and Neptune.  In addition, both \\citet{Deming07} and \\citet{Maness07} calculated that the significant eccentricity of the orbit, e = 0.15, coupled with its short period of $\\sim$2.6 days, should result in circularization timescales of $\\sim$10$^{8}$ years, which contrasts with the old age of the system at $\\ga$6$\\times$10$^{9}$ years. The existence of one or more additional planets in the system could be responsible for perturbations to Gl 436b's orbit, and thus result in the observed peculiarities. We considered this possibility right after the initial publication of \\citet{Gillon07b}, and began an intensive campaign to observe the photometric transits of Gl 436b in order to search for variations indicative of orbital perturbations \\citep{Stringfellow08}.  Early this year, \\citet{Ribas08a} reported the possible detection of a $\\sim$5 M$_{\\earth}$ companion in the Gl 436 system located near the outer 2:1 resonance of Gl 436b via analysis of all the RV data compiled to date. Theoretically this planet would be perturbing Gl 436b so as to increase its orbital inclination at a rate of $\\sim$0.1 deg yr$^{-1}$, and thus its transit depth and length, so that the non-detection by \\citet{Butler04} and the observed transit of \\citet{Gillon07b} were compatible. Since the RV detection of this second planet had a significant false-alarm probability of $\\sim$20\\%, \\citet{Ribas08a} proposed that confirmation could be achieved through 2008 observations of Gl 436b's transits, which would show a lengthening of transit duration by $\\sim$ 2 minutes compared to the \\citet{Gillon07b} data. As well, transit-timing variations (TTVs) of several minutes should also be detectable by observing a significant number of transits.  Recently, \\citet{Alonso08} reported a lack of observed inclination changes and TTV evidence for the second planet, based on a comparison of a single H band light curve obtained in March 2008 to 8$\\micron$ data taken with Spitzer 254 days earlier \\citep{Gillon07a,Deming07}. This result, combined with additional radial velocity measurements \\citep{Howard08,Bonfils08} that contradicted the proposed period of the second planet, drove \\citet{Ribas08b} to retract their claim of the companion at IAU Symposium 253. However, very recently \\citet{Shporer08} presented multiple light curves obtained in May 2007, and could not rule out TTVs on the order of a minute. While the planet specifically proposed by \\citet{Ribas08a} most likely does not exist,  \\citet{Ribas08b} makes a strong case that a second planet is still needed to explain the peculiarities of Gl 436b, and most likely exists in a non-resonant configuration where no strong TTVs are induced. Amateur astronomers have been diligent in observing Gl 436b since it's initial transit discovery, and thus along with this data, published data, and our own data, we are able to present a thorough analysis of the TTVs, inclination, duration, and depth of the transit changes in the Gliese 436 system. We present our observations in \\S2, our modeling and derivation of parameters in \\S3, and explore the observed TTVs and parameters of the system over time in \\S4.\\\\ \\\\ ",
        "conclusions": "We have presented a total of ten new transit light curves of Gl 436b, three of which come from the 3.5-meter telescope at APO, and one of which is from the NMSU 1-meter in January 2005. We have collected and uniformly modeled all available professional and amateur light curves, and searched for any trends in transit timing, width of transit, and depth of transit variations. We find statistically significant, self-consistent trends that are compatible with the perturbation of Gl 436b by a planet with mass $\\lesssim$ 12 M$_{\\earth}$ in a non-resonant orbit with semi-major axis $\\lesssim$ 0.08 AU. This conclusion is based on the numerical simulations of \\citet[][see Fig. 1]{Ribas08a} who constrain the mass and semi-major axis of the theoretical second planet by examining which configurations could produce the observed orbital perturbations while still remaining undetected by the existing radial-velocity data. From our analysis, we infer a non-resonant orbit based on a lack of detected TTVs with amplitude $\\gtrsim$ 1 minute.  We stress that our measured trends are moderately dependent on our 2005 data, and thus subsequent high-precision observations over the next few years need to be carried out to confirm or refute this trend. If confirmed, it would be strong evidence for the first extrasolar planet discovered via orbital perturbations to a transiting planet. Also, we would like to note that although \\citet{Alonso08} had previously limited the rate of inclination change to 0.03$\\pm$0.05 deg/yr, they did so only by measuring the change in width between the 2007 Spitzer observations and their own 2008 H-band data, which they found to be 0.5$\\pm$1.2 minutes. Via Table 1, we find the difference in transit width between the two observations to be 1.5$\\pm$1.4 minutes, which is in agreement with our derived inclination and width values, and is a more reliable result due to using full model fits with proper limb-darkening coefficients. With respect to the the amateur observations, although they are numerous, the very small depth of the transit makes it a challenge for most small aperture systems, resulting in very large uncertainties in i and T$_{0}$. Also, while amateur observers are aware of the importance of precision timing, we of course cannot examine each of their observing set-ups, and thus one must be aware of the possibility, although small, of systemic time offsets on a given night when interpreting their data."
    },
    "0809/0809.4041_arXiv.txt": {
        "abstract": "The bulk flow, \\ie the dipole moment of the peculiar velocity field, is a sensitive probe of matter density fluctuations on very large scales.   However, the peculiar velocity surveys for which the bulk flow has been calculated have non-uniform spatial distributions of tracers, so that the bulk flow estimated does not correspond to that of a simple volume such as a sphere.  Thus bulk flow estimates are generally not strictly comparable between surveys, even those whose effective depths are similar.   In addition, the sparseness of typical surveys can lead to aliasing of small scale power into what is meant to be a probe of the largest scales.   Here we introduce a new method of calculating bulk flow moments where velocities are weighted to give an optimal estimate of the bulk flow of an idealized survey, with the variance of the difference between the estimate and the actual flow being minimized.    These ``minimum variance'' estimates can be designed to estimate the bulk flow on a particular scale with minimal sensitivity to small scale power, and are comparable between surveys. We compile all major peculiar velocity surveys and apply this new method to them.  We find that most surveys we studied are highly consistent with each other. Taken together the data suggest that the bulk flow within a Gaussian window of radius 50 \\hmpc\\ is $407\\pm81$ \\kms\\ toward  \\lb{287\\arcdeg\\pm9}{8\\arcdeg\\pm6}.  The large-scale bulk motion is consistent with predictions from the local density field. This indicates that there are significant density fluctuations on very large scales.  A flow of this amplitude on such a large scale is not expected in the WMAP5-normalized $\\Lambda$CDM cosmology, for which the predicted one-dimensional r.m.s.\\ velocity is $\\sim 110$ \\kms. The large amplitude of the observed bulk flow favors the upper values of the WMAP5 $\\Om h^2$-$\\sigma_8$ error-ellipse, but even the point at the top of the WMAP5 95\\% confidence ellipse predicts a bulk flow which is too low compared to that observed at $>98$\\% confidence level. ",
        "introduction": "\\label{sec:intro}  A long-standing question in cosmography is the origin of the $\\sim 600$ km/s peculiar velocity of the Local Group (LG) with respect to the Cosmic Microwave Background (CMB).  The motion of the LG with respect to the ``Local Sheet'' in which it is embedded is only $\\sim 60$ km/s \\citep{TulShaKar08}, thus most of the LG's motion is due to structures on scales larger than the Local Sheet, i.e.\\ beyond  5 \\hmpc\\ (where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$). In the gravitational instability paradigm \\citep{FelFriFry01,ScoFelFri01,VerHeaPer02}, this motion is due to the gravity of structures on larger scales. For a galaxy at position {\\bf{r}}, the peculiar velocity {\\bf{v}} is given by \\citep{PPC} \\begin{equation} {\\bf{v}}\\left( {\\bf{r}} \\right) = \\frac{\\Omega\\sbr{m}^{0.55}}{4\\pi }\\int d^3 {\\bf{r}}^{\\prime } \\delta\\sbr{m}\\left( \\mathbf{r}^{\\prime }\\right) \\frac{\\left( \\mathbf{r}^{\\prime }-\\mathbf{r}\\right) }{\\left| \\mathbf{r}^{\\prime }-\\mathbf{r}\\right| ^3}. \\label{eq:peculiar} \\end{equation} where $\\delta\\sbr{m} \\left( \\mathbf{r}\\right) = ({\\rho -\\overline{\\rho}})/{\\overline{\\rho}}$, and $\\overline{\\rho}$ is the average density of the Universe,  $\\Om$ is the matter density parameter, and we have used 0.55 instead of 0.6 for the power of $\\Omega_m$ in the pre-factor to improve accuracy for models with dark energy \\citep{Lin05}.  The issue of the LG's motion and that of other nearby galaxies has important cosmological and cosmographical implications. Specifically, as shown by Eq. \\ref{eq:peculiar}, the peculiar velocities of individual galaxies are sensitive to the matter power spectrum over a wide range of scales. Indeed, apart from the Integrated Sachs-Wolfe (ISW) effect \\citep{SacWol67}, peculiar velocities are the only probe of the matter density fluctuations on scales of $\\sim 100 \\hmpc$ and clearly the only dynamical probe in the low-redshift universe. A given power spectrum predicts the r.m.s.\\ of the components of a galaxy's peculiar motion. For models with more power, i.e. a higher normalization, one predicts a larger r.m.s.\\ velocity.  For a single galaxy, the contributions to its motion arise from a range of scales: from the local ($\\sim$ 5 \\hmpc) to the very large ($\\ga 100 \\hmpc$) scales. One may reduce the effects of small scale density fluctuations by studying the peculiar velocity of a larger volume using a sample of peculiar velocity tracers such as galaxies, clusters, or Type Ia supernovae. Beginning with the work of \\cite{RubRobTho76}, a number of such surveys have been undertaken over the last couple of decades \\citep{ DreFabBur87, LynFabBur88, AarBotCor89, Wil90, Cou92, HanMou92, MatForBuc92b, LauPos94, HudLucSmi97, WilCouFab97, GioHayFre98, HudSmiLuc99, DalGioHay99, DalGioHay99b, Wil99b, CouWilStr00, ColSagBur01, HudSmiLuc04, HauHanTho06, SprMasHay08}.  The simplest statistic that can be derived from a sample of peculiar velocities is the dipole moment of the sample, also known as its bulk flow. It was quickly realized that the bulk flow was closely related to the amplitude of fluctuations on large scales, and could be used to test cosmological models \\citep{CluPee81, VitJusDav86}. At face value, however, the surveys cited above yield apparently conflicting results: the measured bulk flow ranges from 0 to $\\sim 1000$ km/s.  Note, however, that many of the above-mentioned surveys are sparsely-sampled, and that while authors quote the bulk flow of the sample, this sample bulk flow is often mis-interpreted as the coherent bulk flow of the whole volume occupied by the survey.  The issue of sparse sampling, small-scale aliasing and their effects on statistics such as the bulk flow were first analyzed by \\cite{Kai88}, and later \\citet{WatFel95} and others \\citep{pairwise00,Hud03,pairwise03,SarFelWat07,WatFel07,FelWat08}. These studies addressed the issue of comparing sparse surveys both to each other (to check for consistency between different sparse surveys) and the comparison of sparse peculiar velocity samples with expectations from cosmological models. One lesson from this work is that both sparse sampling and aliasing present an important effect that must be accounted for in interpreting the results, particularly those from sparse surveys such as clusters or SNe.  Bulk flow estimates are essentially weighted averages of the individual velocities in a survey.   Previous work has focused on a weighting scheme that produces a maximum likelihood estimate (MLE) of the bulk flow of a survey, an estimate that minimizes the uncertainties due to measurement noise but does not make any correction for the survey geometry .  Thus the MLE bulk flow is obviously dependent on a given survey's particular geometry and statistical properties.  In this paper, we instead address the question of how peculiar velocity data can be used to estimate a more general statistic: the bulk flow of an ideal, densely-sampled survey with a given depth.   Our approach will be to calculate \\emph{optimal} weights which produce the best possible estimate of this statistic.  An approach related to this question is that of Zaroubi and collaborators, who used Wiener filtering \\citep{ZarHofDek99} and variants \\citep{Zar02} to reconstruct the matter density field directly from peculiar velocities.  That work, however, was focussed more on the mapping of the density field and the measurement of $\\beta = f(\\Om)/b$,  where $b$ is the bias parameter, than on the bulk flow \\cite[but see][]{HofEldZar01}. In this paper, our aim is somewhat different: to construct dipole moments that probe the largest scales.  The goal of this paper is to make the cleanest measurement of the large-scale bulk flow using the best peculiar velocity data available. We discuss the peculiar velocity surveys used in this analysis in Section~\\ref{sec:surveys}.  In Section~\\ref{sec:mom},  we describe the construction of the velocity moments,  the power spectrum model,  and the optimal weighting scheme used to estimate bulk flow components free of small scale noise. In  Section~\\ref{sec:res} we apply these optimal weights to the data. In Section~\\ref{sec:consist}, we assess whether the optimally-weighted bulk flow results from different surveys are mutually consistent. In Section~\\ref{sec:implic}, we discuss the cosmographic implications of our results. In Section~\\ref{sec:compare}, we compare the measured bulk flow with expectations from cosmological models. We discuss our resuts in Section~\\ref{sec:discuss} and conclude in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have calculated optimal ``minimum variance'' weights designed to measure bulk flows with minimal sensitivity to small scale power, and have applied these weights to a number of recent peculiar velocity surveys.  We find that all of the surveys we studied are consistent with each other, with the possible exception of the \\cite{LauPos94} BCG survey. Taken together the data suggest that the bulk flow within a Gaussian window of radius 50 \\hmpc\\ is $407\\pm81$ \\kms\\ toward  \\lb{287\\arcdeg\\pm9}{8\\arcdeg\\pm6}.  This motion is not due to nearby sources, such as the Great Attractor (at a distance of $\\sim40$\\hmpc), but rather to sources at greater depths that have yet to be fully identified.  A flow of this amplitude on such a large scale is not expected in the WMAP5-normalized $\\Lambda$CDM cosmology. The observed bulk flow favors the upper values of the WMAP5 $\\Om h^2$-$\\sigma_8$ error-ellipse, but even the point at the top of the WMAP5 95\\% confidence ellipse predicts a bulk flow which is too small compared to that observed at a confidence level $> 98$\\%.  There are several possible explanations for the discrepancy we have observed.   There is the possibility that we happen to live in a volume with a statistically unlikely ($\\lesssim 2$\\%) bulk flow magnitude.   If this is the case, then the structures that cause this flow should be eventually identified as the depth and sky coverage of redshift surveys increase.    Alternatively, it is possible that the large observed flow is the result of a systematic error in the data, although the independence of the distance indicators (TF, FP and SN Ia) and methodology of the various surveys, as well as the agreement between different surveys makes this unlikely.  The bulk flow in the nearby ($d \\lesssim 60 \\hmpc$) Universe is no longer noise-limited but rather cosmic variance limited, so that increasing the quantity of nearby peculiar velocity data will not alter the significance of this result. At very large depths ($d > 100 \\hmpc$), however, the bulk flow measurement is still quite noisy.  Future peculiar velocity surveys, such as the NOAO Fundamental Plane Survey \\citep{SmiHudNel04}, as well as nearby supernovae surveys \\citep{FilLiTre01,WooAldLee04,KelSchBes07, FriBasBec08}, are expected to yield more precise measurements of the amplitude of the bulk motion on these very large scales, and thus have the potential to strengthen the cosmological constraints therefrom. In order to measure the bulk flow variance directly, one must measure the bulk flow in independent (i.e. distant) regions. For the standard distance estimators used in traditional peculiar velocity work, the errors grow proportional to distance and hence become infeasible at large distances. So other techniques, such as those based on the kinetic Sunyaev-Zel'dovich effect \\citep{SunZel72,RepLah91,HaeTeg96, RuhAdeCar04, Kos06, ZhaFelJus08, KasAtrKoc08a} will be needed to access independent volumes.  To reiterate, the results presented in this paper pose a challenge to the standard $\\Lambda$CDM model with the WMAP5 parameters. As this study shows, the implications to the standard scenario should be explored with as many independent cosmological observations as we can muster, with particular attention paid to clues from probes at low redshift and on the largest scales.  \\vspace{0.2cm}  \\noindent{\\bf Acknowlegements:} We thank Niayesh Afshordi and Russell Smith for interesting discussions and helpful comments on earlier versions of this paper. HAF would like to thank Danny Marfatia for helpful comments and Brad Klee for programming help. HAF has been supported in part by a grant from the Research Corporation, by an NSF grant AST-0807326 and by the National Science Foundation through TeraGrid resources provided by the NCSA. MJH is supported by NSERC."
    },
    "0809/0809.2983_arXiv.txt": {
        "abstract": "{We present the main results of the first long-term spectrophotometric monitoring of the ``Einstein cross'' Q2237+0305 and of the single-epoch spectra of the lensed quasar J1131-1231. \\\\ From October 2004 to December 2006, we find that two prominent microlensing events affect images A \\& B in  Q2237+0305 while images C \\& D remain grossly unaffected by microlensing on a time scale of a few months. Microlensing in A \\& B goes with chromatic variations of the quasar continuum. We observe stronger micro-amplification in the blue than in the red part of the spectrum, as expected for continuum emission arising from a standard accretion disk. Microlensing induced variations of the \\ciii\\ emission are observed both in the integrated line intensity and profile. Finally, we also find that images C \\& D are about 0.1-0.3 mag redder than images A \\& B. \\\\ The spectra of images A-B-C in J1131-1231 reveal that, in April 2003, microlensing was at work in images A and C. We find that microlensing de-amplifies the continuum emission and the Broad Line Region (BLR) in these images. Contrary to the case of Q2237+0305, we do not find evidence for chromatic microlensing of the continuum emission. On the other hand, we observe that the Balmer and \\MgII\\ broad line profiles are deformed by microlensing. These deformations imply an anti-correlation between the width of the emission line and the size of the corresponding emitting region. Finally, the differential microlensing of the \\FeII\\ emission suggests that the bulk of \\FeII\\ is emitted in the outer parts of the BLR while another fraction of \\FeII\\ is produced in a compact region. }  \\FullConference{The Manchester Microlensing Conference: The 12th International Conference and ANGLES Microlensing Workshop\\\\ January 21-25 2008\\\\ Manchester, UK}  \\begin{document} ",
        "introduction": "So far, the effect of microlensing on the continuum and on the Broad Line Region (BLR) of quasars has been investigated mainly theoretically (e.g. Schneider \\& Wambsganss 1990, Wambsganss \\& Paczynski 1991, Lewis et al. 1998, Abajas et al. 2002, Lewis \\& Ibata 2004). In this contribution, we report the main results from a spectrophotometric study of microlensing occurring in two gravitationally lensed quasars: the ``Einstein cross'' Q2237+0305 and J1131-1231 (Fig.~\\ref{fig:aquisition}). The Einstein cross consists of a $z_s= 1.695$ quasar gravitationally lensed into four images in a cross-like pattern about the bulge of a nearby Sab galaxy ($z_l= 0.0394$). Due to the low redshift of the lensing galaxy, to the high density of stars on the line of sight to the lensed images and to the negligible time delay, this system is among the best ones to study microlensing. We present hereafter some results of the first long-term regular spectro-photometric monitoring campaign of this object. The second target, J1131-1231, is one of the nearest gravitationally lensed quasar. The lensing galaxy at $z_l = 0.295$ splits the light rays from the source at $z_s = 0.658$ into four macro-images: three bright images (A-B-C) separated on the sky by typically 1'' and a fainter component (D) located at 3.6'' from A (Sluse et al. 2003). We study microlensing between the three brightest images which also happen to have negligible differential time delays.   \\begin{figure}[!ht] \\begin{centering} \\includegraphics[width=6.5cm]{Q2237_slits3.eps} \\includegraphics[width=6.5cm]{J1131Manchester.eps} \\caption{{\\it Left:} FORS1 R-band acquisition image of Q2237+0305 taken on 12 September 2005. Two different slits (0.7'' width) placed on images A \\& D (mask 1) C \\& D (mask 2) have been used. {\\it Right:} FORS2 R-band image of J1131-1231 taken on 26 April 2003. The slit (1'' width) used is shown. The seeing on both images is of the order of 0.6''.} \\label{fig:aquisition} \\end{centering} \\end{figure} ",
        "conclusions": "We have presented the main results of the first long term spectro-photometric monitoring of the Einstein Cross Q2237+0305 and of the single-epoch spectroscopic observations of the lensed quasar J1131-1231. The microlensing spectral signatures observed in the lensed images of these two systems confirm that quasar microlensing is a powerful tool for studying the continuum and the broad line emitting region of lensed quasars.  We find that all images of the Q2237+0305 are affected by some microlensing during the observing period. Image D seems to be the less affected. Our observations have monitored two important microlensing brightening episodes in image B and A on resp. May 2005-Dec 2005 (HJD 3500-3710) and May 2006-Dec 2006 (HJD 3880-4100). For each of these events, the continuum becomes bluer when the image gets brighter as expected from microlensing magnification of an accretion disk. In addition to the microlensing induced chromatic variations of the quasar continuum we have also observed time-independent reddening of images C and D with respect to A and B. We estimate that the amount of differential extinction between pairs of quasar images is in the range 0.1-0.3 mag, with images C and D being the most reddened. We also report microlensing induced variations of the BELs affecting both the integrated line intensities and profiles. The microlensing of the BLR is the strongest in image A where we observe that the broad component of the broad lines is more microlensed than the core. This is compatible with an anti-correlation between the line width and the size of the corresponding emitting region. Finally, we observe that the continuum of images C and A are significantly more magnified than the BELs during the whole observing period, indicating that long-term microlensing is at work in those two images.  The analysis of the single-epoch spectra of images A-B-C of J1131-1231 has revealed microlensing de-amplification of the quasar images A and C. In these two cases, the continuum and the BELs are microlensed. However, a larger fraction of the BELs is microlensed in image C, indicating that the Einstein radius of the microlens is larger for that image. Microlensing of the BELs confirms that the size of the emission line region is anti-correlated with the FWHM of the corresponding line component. Contrary to the case of Q2237+0305, we do not find evidence for chromatic microlensing of the continuum. The most interesting result concerns the microlensing of the \\FeII\\ emission. On one hand, we find that a large fraction of the near UV and optical \\FeII\\, emission takes place in a region similar to the outer parts of the Balmer Line Emitting Region. On the other hand we find that the \\FeII\\ emitted in the rest-frame ranges 4630-4800 \\& 3080-3540\\AA\\ likely arises from a more compact region possibly similar in size to the region emitting the very broad \\MgII\\ emission.  Our observations have demonstrated that the BLR of the quasars Q2237+0305 and J1131-1231 is small enough to be significantly microlensed. Because of the different source redshifts, these two targets allow to study very different emitting regions in the optical range. Microlensing in Q2237+0305 allows to study the near UV continuum emission, the \\FeII$_{\\rm UV}$, and the high ionisation \\ciii, \\civ\\ broad emission lines while microlensing in J1131-1231 allows to probe the rest-frame continuum emission in the range 2500-6000 \\AA, the \\FeII$_{\\rm opt}$, the low ionisation \\MgII, and the Balmer broad emission lines. The spectro-photometric monitoring of Q2237+0305 demonstrates the great asset of the time information in the study of quasar microlensing. We have began a 2 year spectrophotometric monitoring of J1131-1231 with the VLT to track the microlensing induced deformation of the BELs in this system. The spectro-photometric monitoring data gathered for Q2237+0305 and soon for J1131-1231 will allow to put sharp upper limits on the size of the continuum emission region and on the geometry of the BLR (for different ionization levels) using state-of-the-art inverse ray-shooting simulations.  {\\large{\\bf {Acknowledgements:}}} This work is supported by the Swiss National Science Foundation, by ESA PRODEX under contract PEA C90194HST and by the Belgian Federal Science Policy Office."
    },
    "0809/0809.3917_arXiv.txt": {
        "abstract": "{The Mini-Calorimeter (MCAL) instrument on-board the AGILE satellite is a non-imaging gamma-ray scintillation detector sensitive in the 300~keV--100~MeV energy range with a total on-axis geometrical area of $1400~\\mathrm{cm^2}$. Gamma-Ray Bursts (GRBs) are one of the main scientific targets of the AGILE mission and the MCAL design as an independent self-triggering detector makes it a valuable all-sky monitor for GRBs. Furthermore MCAL is one of the very few operative instruments with microsecond timing capabilities in the MeV range.} {In this paper the results of GRB detections with MCAL after one year of operation in space are presented and discussed.} {A flexible trigger logic implemented in the AGILE payload data-handling unit allows the on-board detection of GRBs. For triggered events, energy and timing information are sent to telemetry on a photon-by-photon basis, so that energy and time binning are limited by counting statistics only. When the trigger logic is not active, GRBs can be detected offline in ratemeter data, although with worse energy and time resolution.} {Between the end of June 2007 and June 2008 MCAL detected 51 GRBs, with a detection rate of about 1~GRB/week, plus several other events at a few milliseconds timescales. Since February 2008 the on-board trigger logic has been fully active. Comparison of MCAL detected events and data provided by other space instruments confirms the sensitivity and effective area estimations. MCAL also joined the $3^{rd}$ Inter-Planetary Network, to contribute to GRB localization by means of triangulation. } {} ",
        "introduction": "The AGILE satellite \\citep{Tavani2008,Tavani2008b}, the Italian space mission dedicated to gamma-ray and hard-X astrophysics, has the study of GRBs among its main scientific targets.  The Gamma-Ray Imaging Detector (GRID), composed of a tungsten-silicon tracker \\citep{Prest2003} and a CsI(Tl) Mini-Calorimeter, has a wide field of view that makes it a valuable instrument for GRB detection in the poorly explored 30~MeV-50~GeV energy band. SuperAGILE \\citep{Feroci2007}, the hard X-ray imager on-board AGILE operating in the 18-60~keV energy band, is equipped with an on-board trigger logic and localization algorithm providing few arcmin position accuracy, allowing rapid dissemination of the coordinates \\citep{DelMonte2007c}. The Mini-Calorimeter, despite being a subsystem of the GRID, is also equipped with a self-triggering operative mode and on-board logic making it an all-sky monitor in the 300~keV-100~MeV energy range. A simultaneous GRB detection with GRID, MCAL and SuperAGILE would allow spectral coverage over six orders of magnitude. In this paper the status of the GRB detection with MCAL, one year after the AGILE launch, is reviewed and discussed. ",
        "conclusions": "MCAL is tailored to the detection of medium-bright GRBs with peak energies above a few hundred keV, as expected from sensitivity calculations and experimental evidence reported in section \\ref{sensitivity}. The main characteristics of the instrument are its spectroscopic capabilities in the MeV range and the microsecond timing accuracy. Between $22^{nd}$ June 2007 and $30^{th}$ June 2008 MCAL detected 51 GRBs, with an average detection rate of about 1~GRB/week. A preliminary flux calibration is in good agreement with expectations. Since the beginning of February 2008 the on-board trigger logic has been active and calibration activities are in progress."
    },
    "0809/0809.1065_arXiv.txt": {
        "abstract": "We present the first results of a survey of 14 low redshift galaxy clusters using {\\sl Suzaku}. Although luminous ($L_x> 1\\times 10^{43}$ erg s$^{-1}$ (0.1-2.4 keV)), these clusters have no prior pointed X-ray data. Together with 47 other systems they form a flux limited sample ($f_x >1.0 \\times 10^{-11}$ erg s$^{-1}$cm$^{-2}$ (0.1-2.4 keV)) with $z\\leq 0.1$ in the northern celestial hemisphere. Using this total sample we evaluate the local $L-T$ relationship and the local cluster temperature function. {\\sl Suzaku} temperature measurements appear to be in accord with those of other missions. General agreement is found with other published estimates of the low redshift cluster temperature function, however the sample used here exhibits slightly lower space densities at gas temperatures below 4-5 keV. We find a corresponding deficit in the number of clusters with temperatures between approximately 4.5 and 5.5 keV. Although at low significance, a similar feature exists in previous low redshift cluster datasets. We suggest that low-redshift cluster samples, while crucial for calibrating precision cosmology measurements, must be used with caution due to their limited size and susceptibility to the effects of cosmic variance. ",
        "introduction": "\\label{sec:introduction}  The nature and evolution of galaxy clusters offers an independent and complementary probe of cosmology to those of the cosmic microwave background \\citep{hinshaw08} or supernova (e.g. Permlmutter \\& Schmidt 2003). Theoretical work (e.g. Haiman, Mohr \\& Holder 2001) has shown that  large surveys of clusters to reasonably high redshift ($z\\sim 1$) can not only provide these much needed constraints, but can, through the use of information such as the spatial distribution of clusters, be made ``self-calibrating'' - {\\em if} the evolution of cluster structure can be parameterized (Hu 2003, Majumdar \\& Mohr 2003), and {\\em if} the scatter between cluster mass and observables (luminosity, temperature) is understood for clusters {\\em at the survey detection threshold}.  Future X-ray surveys or Sunyaev-Zel'dovich Effect (SZE) surveys, such as the South Pole Telescope (SPT), and the Atacama Cosmology Telescope (ACT), will be designed to detect $\\sim 10^4$ clusters to high redshift within a limited sky area and will not sample very many local clusters, or obtain very high physical resolution data. However, both of these characteristics are vital in order to establish the zero-redshift baseline from which we can then evaluate the global evolution of cluster scaling laws and structure, and calibrate possible systematic effects induced by potentially complex astrophysics.  In this paper we describe the results of an effort to utilize existing X-ray cluster data, and new observations by the {\\sl Suzaku} (JAXA/NASA)  mission to add to the existing body of work on the low redshift population of clusters (e.g. Henry 2000; Reiprich \\& Bohringer 2002; Ikebe et al. 2002). Specifically, we present the results of a {\\sl Suzaku} survey of 14 nearby, X-ray luminous, clusters that have otherwise {\\em never} been observed in a direct, pointed X-ray observation with earlier, spectrally sensitive instruments. Our motivation is to increase the robustness of the local cluster temperature function at {\\em lower} masses, which are typical of the mass threshold of the majority of likely future X-ray and SZE cluster surveys. The range of masses probed by this data combined with pre-existing data (a total, complete, sample of 61 objects) is approximately $8\\times 10^{13}$---$10^{15}$ M$_{\\odot}$ ($H_0=70$). We therefore bracket the expected low mass limit of future surveys, all of which are designed to probe to $\\sim 2\\times 10^{14}$ M$_{\\odot}$ (locally for the X-ray, and across essentially all redshifts for SPT/ACT). For modern cosmological tests based on the form of $\\frac{dN(>M)}{dz}$ (the number of clusters per unit redshift above a mass $M$, c.f. Haiman, Mohr \\& Holder 2001), the {\\em only} theoretical requirement is a precise understanding of the scatter in mass around the detection threshold (e.g. X-ray flux limit).    The rest of this paper is organized as follows. In \\S~\\ref{sec:sample}, we compile the local, flux-limited, complete sample. In \\S~\\ref{sec:data}, we describe the Suzaku data, and the spectral analysis method. We also present spectral fitting result for the 14 clusters.  In \\S~\\ref{sec:t_l}, we fit the $L-T$ relation in two different ways. In \\S~\\ref{sec:xtf}, we describe our procedures for computing the XTF and present our results. We also make a comparison with previous results. In \\S~\\ref{sec:conclude}, we briefly summarize our conclusions and the implications of this work. Throughout this paper, we adopt a spatially flat universe dominated by a cosmological constant and cold dark matter (CDM), with the following set of cosmological parameters: $\\Omega_m=0.27$, $\\Omega_{\\Lambda}=0.73$ and $H_0=70~{\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "\\label{sec:conclude}  We have conducted an X-ray galaxy cluster survey of 14 systems using {\\sl Suzaku}. These clusters are part of a  statistically complete sample of 61 clusters in the local ($z<0.1$) Universe with $f_x>1.0 \\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ (0.1-2.4 keV), in the northern celestial hemisphere. None of these 14 clusters have been observed previously in pointed X-ray observations, despite being luminous systems ($L_x>1 \\times 10^{43}$ erg s$^{-1}$ (0.1-2.4 keV)). The {\\sl Suzaku} data is shown to be of high quality and enables robust imaging and spectral analyses to be made. A comparison of spectral analyses of {\\sl Suzaku} data on 3 well known clusters (not in our sample) with those from earlier missions indicates no obvious biases either in the {\\sl Suzaku} calibration or the analysis approach that we use.  We present the first results of a spectral analysis of these 14 clusters with the aim of deriving gas temperatures that reflect the true virial temperature. To do this we employed a spectral modeling scheme that determines whether a one or two plasma temperature model is appropriate. Two temperature models allow us to separate out the contribution of cooling cores in 7 of the clusters. Combining the new {\\sl Suzaku} data with preexisting temperature data for the flux limited sample we obtain a dataset of 57 objects with spectroscopically confirmed gas temperatures and 4 objects for which pointed X-ray data is either not yet public or has not yet been analyzed. For these 4 systems we simply use estimated temperatures based on the $L-T$ relation, and note that their inclusion or exclusion has a negligible impact on our results.  We have determined both the $L-T$ relation and the cumulative XTF for the full sample. In fitting the $L-T$ relationship with a power law we argue that the often used BCES method can produce a power law index that is dependent on the actual measurement errors in $L$ and $T$, owing to the weighting scheme employed. We have therefore also employed a modified $\\chi^2$ fit that takes intrinsic scatter into account in an attempt to mitigate any bias in the fit. We find that  this modified $\\chi^2$ and the BCES  method are in quite close agreement, with power-law indices differing by $\\sim 4$\\% and normalizations by $\\sim 0.1$ \\%. The modified $\\chi^2$ yields $L_{0.1-2.4\\rm{keV}}= 10^{42.31} T^{2.45}$.  The measured XTF is in close agreement with that of previous studies. Obviously there is significant overlap in the actual clusters contained in all such works. However, our lower flux limit (by a factor of 2 from, for example, I02) and the fact that 22\\% of our sample has never been studied in the X-ray before, helps reduce this problem.  There is some evidence that our XTF is slightly lower at temperatures less than $\\sim 4.5$ keV than the earlier measurements. This appears to be partially a consequence of a deficit of clusters in the range of 4.5-5.5 keV. We have considered several possible causes for this. These include; a systematic bias in our spectral analyses, a statistical fluke, or a consequence of large-scale structure and cosmic variance in this very local sample. Since a similar deficit is also apparent in two earlier XTF measurements (I02, Henry (2000)) it may be that a statistical or cosmic variance effect is the most likely.  If local galaxy cluster samples are to be used to construct robust calibrations for future high-precision cosmological tests then it will be crucial to ensure that we understand the finite volume effects (e.g. cosmic variance) that may produce systematic biases. One solution will be to push measurements to slightly lower mass scales where clusters are more numerous. We have demonstrated that {\\sl Suzaku} is an excellent instrument for doing just that, by exploring the population of luminous, nearby, clusters that has nonetheless not yet been fully studied.   \\vspace{-0.5\\baselineskip}"
    },
    "0809/0809.0255_arXiv.txt": {
        "abstract": "Starting from a complete sample of type I AGN observed by \\emph{INTEGRAL} in the 20-40 keV band, we have selected a set of 8 AGN which can be classified as radio loud objects according to their 1.4 GHz power density,  radio to hard X-ray flux flux density ratio and  radio morphology. The sample contains 6 Broad Line Radio Galaxies and 2 candidate ones. Most of the objects in our sample display a double lobe morphology, both on small and large scales. For all the objects, we present broad-band (1-110 keV) spectral analysis using \\emph{INTEGRAL} observations together with archival \\emph{XMM-Newton}, \\emph{Chandra}, \\emph{Swift/XRT} and \\emph{Swift/BAT} data. We constrain the primary continuum (photon index and cut-off energy), intrinsic absorption and reprocessing features (iron line and reflection) in  most of the objects. The sources analysed here show remarkable similarities to radio quiet type I AGN with respect to most of the parameters analysed; we only find marginal evidence for weaker reprocessing features in our objects compared to their radio quiet counterparts. Similarly we do not find any correlation between the spectral parameters studied and the source core dominance or radio to 20-100 keV flux density ratios, suggesting that what makes our objects  radio loud has no effect on their high energy characteristics. ",
        "introduction": "Most of what is known about Active Galactic Nuclei (AGN) is essentially based on studies of radio quiet sources, which make up almost 90\\% of the entire AGN population. It is now widely accepted that active galaxies are powered by accretion onto a supermassive black hole: the observed radiation, spanning the entire electromagnetic spectrum, is produced by a cold accretion disk and by a hot corona, as proposed in the so-called two-phase model \\citep{b15}. In the X-ray domain, the emission of radio quiet Broad Line AGN (Seyfert 1) is, to the first order, well described by a power law of photon index 1.8-2.0, extending from a few keV to over 100 keV; at higher energies there is evidence for an exponential cut-off, the exact value of which is still uncertain \\citep{b22,b23}. Secondary features such as the Fe K$\\alpha$ line and the Compton reflection component are also commonly observed; they are considered to be the effects of reprocessing of the primary continuum and are relatively well understood \\citep{b19}. Studies performed on samples of Broad Line Radio Galaxies (BLRG from now on; \\citealt{b24,b7}) have shown that these objects show optical/UV continuum and emission line characteristics similar to their radio quiet counterparts, but  display some fundamental differences in their X-ray behaviour. BLRG, in fact, exhibit flatter/harder power law slopes than radio quiet Seyfert 1 galaxies and are also known to have weaker reprocessing features  (e.g. \\citealt{b24} and \\citealt{b7}). The origin of these differences is however far from being understood. A possible cause for the observed diversity could be ascribed to a different disk geometry and/or accretion flow efficiencies, or to the presence of jets and beaming effects that contaminate/dilute the AGN component and the reprocessing features; ionised reflection, which naturally produces weaker reflection features, has also been considered in the literature \\citep{b60}. Now that a large sample of AGN detected above 20 keV by \\emph{INTEGRAL} is available, it is possible and important to perform a comparison between different classes of AGN, by measuring the shape of the primary continuum together with the high energy cut-off and the reflection component. In the present work we focus on the broad-band (1-110 keV) spectral analysis of a sample of eight radio loud type 1 AGN, combining \\emph{XMM-Newton}, \\emph{Chandra}, \\emph{Swift/XRT} and \\emph{BAT} data together with \\emph{INTEGRAL/ISGRI} measurements. We also compare our results to a sample of radio quiet Seyfert 1s detected by \\emph{INTEGRAL} and recently studied by \\citet{b21} and finally discuss the implications of our findings. ",
        "conclusions": "In the following we discuss the distribution of the fit parameters in our sample of AGN (see table 6) in comparison to a similar study made on a set of 9 radio quiet Seyfert 1 \\citep{b21} also detected by \\emph{INTEGRAL}. The power law slopes are found to be similar in both samples, i.e. around 1.5-2.0; only one of the radio bright AGN (the one requiring complex absorption) shows a flat power law continuum, while low values of $\\Gamma$ are more numerous in the sample of  \\citet{b21}. The intrinsic absorption measured in our sources is generally small or absent ($\\lesssim$2$\\times$10$^{21}$ atoms cm$^{-2}$), except for 3C 111 and IGR J21247+5058, and also similar to the values found for radio quiet  type 1 AGN; in 3C 111 the extra absorption is probably due to intervening material between us and the source, while in the case of IGR J21247+5058 complex absorption is strongly required by the data. It is interesting to note that similar complexity, i.e. one or more layers of cold material partially covering the central nucleus, is also observed in 2 of the 9 objects studied by \\citet{b21}. We also point out that, up to very recently, the only other BLRG known to have a similar spectral complexity regarding the absorption was 3C 445 \\citep{b40}, for which three layers of cold material were required to model the spectrum. Our analysis suggests that  objects with complex absorption are present within the population of broad line AGN, independent of them being radio quiet or radio loud. \\citet{b7} found evidence in their data for a correlation between absorption and radio core dominance, with more absorbed objects showing lower \\emph{CD} values. We do not find such a trend; furthermore we do not have any indication of a correlation between photon index or absorption and R$_{HX}$ using both our sources and the radio quiet Seyferts of \\citet{b21}.  From the spectral analysis presented here, it is also evident that the high energy cut-off (although not required by all the sources in the sample) spans a wide range of values from $\\sim$40 keV up to more than 300 keV, as already found for radio quiet AGN (see for instance \\citealt{b35} and \\citealt{b21}). We also explored the possibility of a correlation between the high energy cut-off and the photon index but found nothing; a trend of increasing cut-off energy with higher $\\Gamma$ has been reported in the literature (see for instance \\citealt{b22}), but due to the strong dependency between these two parameters in the fitting procedure, it is difficult to discriminate between any true correlation and induced effects.  As far as the reflection fraction is concerned, we find in all but two cases (namely B3 0309+411 and 4C 74.26) very low values for this component, typically below 1; as expected in these sources, the EW of the iron line  is $\\lesssim$100 eV. Indeed if the iron line emission is associated with the optically thick material of the disk, one would expect the line EW to correlate with the reflection fraction as observed. However, in B3 0309+411 and 4C 74.26 the observed iron line EW is too small for the reflection measured: a possible explanation resides in the scatter expected in the EW values due to a variation in the iron abundance, or to a possible anisotropy of the source seed photons which might affect the observed spectrum \\citep{b42,b43}. For a given value of \\emph{R}, the EW could also differ according to the value of the power law photon index: up to $\\Gamma$=2 the EW decreases as the spectrum steepens, while above $\\Gamma$=2 the trend is reversed \\citep{b22, b41}. B3 0309+411 and 4C 74.26 both have steep spectra and so a lower EW than expected on the basis of the measured \\emph{R} value is not a surprise.  Overall we can conclude that, as already observed by a number of authors (i.e. \\citealt{b7} and \\citealt{b24}), the reprocessing features of ttthe radio loud AGN analysed here tend to be, on average, quite weak. For a better comparison with their radio quiet counterparts, we have again used our data in combination with those of \\citet{b21}. Figure \\ref{r_ew} shows a plot of \\emph{R} vs. EW  for the entire sample. It is evident that radio loud AGN are more confined to a region of the plot characterized by low values of  EW and to a lesser extent of \\emph{R}, whereas radio quiet objects tend to be more spread over both axes. This evidence does not seem to be related to the radio core dominance, suggesting that any dilution of the reprocessing features, at least for the sources reported here, is not caused by the presence of a jet. Since the observed difference is not striking, an accretion flow origin for the X/gamma-ray emission is a likely explanation for the production of the reprocessing features in our radio loud AGN; this agrees well with some of them being classified as FR II sources, i.e. the BLRG most closely resembling radio quiet AGN \\citep{b44}. Within this scenario, it is still possible that a different geometry and/or accretion flow efficiency involving the disk provides the condition for weaker reflection components in BLRG. Alternatively weaker  reflection features might be the result of reprocessing in an ionised accretion disk, as suggested by \\citet{b60}; this would alleviate the need for a change in the accretion disk geometry and provide a more similar enviroment for both radio loud and radio quiet AGN.  Our sample is of course very small, limited to sources which have a small dynamic range of parameters (redshift, loudness, core dominance, etc.), and contains two objects (i.e. QSO B0241+62 and IGR J13109-5552) which deserve further radio studies and are not immediately confirmed as BLRG. We hope in the near future to improve the statistics, adding newly discovered AGN selected in the hard X-ray band in order to verify our findings as well as our novel approach of selecting radio loud type 1 sources.   \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 3}} \\vspace{0.2cm} \\begin{tabular}{lcccccccc} \\multicolumn{9}{c}{{\\bf Spectral Fit Results: \\texttt{wa$_g$*wa*(po+zga)}}}\\\\ \\hline {\\bf Name}    &  {\\bf N$_H$}               & {\\bf $\\Gamma$}&{\\bf E$_{line}$}  & {\\bf $\\sigma$} &   {\\bf EW}   & {\\bf C$_1^A$} &{\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}\\\\ &{\\bf (10$^{22}$ cm$^{-2}$)}&                &    {\\bf (keV)}   & {\\bf (eV)}     & {\\bf (eV)}   &                &              & \\\\ \\hline QSO B0241+62   &0.23$^{+0.08}_{-0.08}$&1.65$^{+0.07}_{-0.07}$&6.40$^{+0.06}_{-0.05}$&    10f        &85$^{+43}_{-45}$ &1.02$^{+0.22}_{-0.18}$&1.32$^{+0.32}_{-0.26}$& 227.1 (234) \\\\ B3 0309+411    &        -           &1.82$^{+0.03}_{-0.03}$&6.34$^{+0.07}_{-0.09}$&    10f        &100$^{+43}_{-50}$&      -             &4.65$^{+2.11}_{-2.00}$& 574.6 (603)\\\\ 3C 111         &0.43$^{+0.02}_{-0.02}$&1.69$^{+0.02}_{-0.02}$&6.40$^{+0.05}_{-0.05}$&    10f        &  $<$30         &0.44$^{+0.05}_{-0.04}$&0.54$^{+0.11}_{-0.10}$& 1245.7 (1498) \\\\ IGR J13109-5552& 0.27$^{+0.23}_{-0.27}$ &1.70$^{+0.24}_{-0.22}$&         -          &     -         &       -        &     -              &2.39$^{+1.84}_{-0.99}$& 41.2 (55)\\\\ 3C 390.3       &         -          &1.74$^{+0.01}_{-0.01}$&6.44$^{+0.02}_{-0.04}$&    10f        &  71$^{+8}_{-18}$&0.90$^{+0.07}_{-0.07}$&0.94$^{+0.10}_{-0.11}$& 2200.4 (2234)\\\\ 4C 74.26       &0.12$^{+0.01}_{-0.01}$&1.73$^{+0.02}_{-0.01}$&6.44$^{+0.05}_{-0.06}$&183$^{+84}_{-65}$&103$^{+32}_{-25}$&0.54$^{+0.09}_{-0.09}$&0.87$^{+0.25}_{-0.25}$& 1992.0 (2060)\\\\ S5 2116+81     &    $<$0.19         &1.97$^{+0.12}_{-0.12}$&     -             &      -         &        -       &1.52$^{+0.60}_{-0.44}$&2.76$^{+1.34}_{-0.95}$& 89.1 (128) \\\\ IGR J21247+5058&0.69$^{+0.02}_{-0.02}$&1.33$^{+0.02}_{-0.01}$&6.39$^{+0.05}_{-0.08}$&    10f        &   $<$30        &0.32$^{+0.02}_{-0.02}$&0.36$^{+0.02}_{-0.02}$& 2692.3 (2559)\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 3A}} \\vspace{0.2cm} \\begin{tabular}{lccccccc} \\multicolumn{8}{c}{{\\bf 3C 390.3 Spectral Fits Results: \\texttt{wa$_g$*(bb+po+zga)}}}\\\\ \\hline {\\bf N$_H$}               &{\\bf $\\Gamma$} & {\\bf kT}     &{\\bf E$_{line}^{\\dagger}$} & {\\bf EW}  &{\\bf C$_1^A$}& {\\bf C$_2^B$}& {\\bf $\\chi^2$ (dof)}\\\\ {\\bf (10$^{22}$cm$^{-2}$)}&                & {\\bf (keV)} &    {\\bf (keV)}            & {\\bf (eV)}&              &               &  \\\\ \\hline -            &1.89$^{+0.02}_{-0.02}$&2.41$^{+0.22}_{-0.17}$ &6.42$^{+0.04}_{-0.03}$& 49$^{+9}_{-15}$&1.37$^{+0.16}_{-0.15}$&1.58$^{+0.24}_{-0.21}$& 1970.4 (2232)\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\dagger}$: line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 3B}} \\scriptsize \\vspace{0.2cm} \\begin{tabular}{lccccccccc} \\multicolumn{10}{c}{{\\bf IGR J21247+5058 Spectral Fit Results: \\texttt{wa$_g$*pcfabs*pcfabs*(po+zga)}}}\\\\ \\hline {\\bf $\\Gamma$}&{\\bf N$_H^1$}                      &{\\bf cf$_1$}& {\\bf N$_H^2$}                     &{\\bf cf$_2$}&{\\bf E$_{line}^{\\dagger}$}&{\\bf EW}&{\\bf C$_1^A$}&{\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}\\\\ &{\\bf (10$^{22}$cm$^{-2}$)}&                   &{\\bf (10$^{22}$ cm$^{-2}$)}&                    &    {\\bf (keV)}    &{\\bf (eV)} &          & & \\\\ \\hline 1.73$^{+0.05}_{-0.04}$&10.81$^{+1.88}_{-1.49}$&0.38$^{+0.03}_{-0.03}$&1.08$^{+0.16}_{-0.15}$&0.86$^{+0.04}_{-0.03}$&6.39$^{+0.07}_{-0.08}$&$<$30&0.62$^{+0.06}_{-0.05}$&0.79$^{+0.08}_{-0.06}$& 2356.2 (2556)\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\dagger}$: line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 4}} \\vspace{0.2cm} \\begin{tabular}{lcccccccccc} \\multicolumn{11}{c}{{\\bf Spectral Fit Results: \\texttt{wa$_g$*wa*(cutoffpl+zga)}}}\\\\ \\hline {\\bf Name} &       {\\bf N$_H$}        & {\\bf $\\Gamma$}& {\\bf E$_{cut}$}      &{\\bf E$_{line}$}         & {\\bf $\\sigma$}                 &   {\\bf EW}       & {\\bf C$_1^A$} &{\\bf C$_1^B$}&{\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ &{\\bf (10$^{22}$ cm$^{-2}$)}&     &  {\\bf (keV)}  &    {\\bf (keV)}     & {\\bf (eV)}                  & {\\bf (eV)}       &  && &\\\\ \\hline QSO B0241+61    &0.18$^{+0.09}_{-0.09}$&1.58$^{+0.09}_{-0.09}$&$>$78&6.40$^{+0.07}_{-0.05}$&10f           &80$^{+41}_{-45}$&1.14$^{+0.23}_{-0.21}$&1.54$^{+0.39}_{-0.34}$& 222.9 (233)& 96\\%\\\\ B3 0309+411     &-                   &1.80$^{+0.05}_{-0.05}$&   $>$43         &6.34$^{+0.07}_{-0.09}$&10f           &99$^{+46}_{-47}$&-&6.75$^{+3.44}_{-3.81}$&573.9 (602)& 91\\%\\\\ 3C 111          &0.42$^{+0.02}_{-0.02}$&1.65$^{+0.03}_{-0.03}$&110$^{+118}_{-40}$&6.40$^{+0.05}_{-0.05}$&10f           &$<$30           &0.56$^{+0.09}_{-0.08}$&0.74$^{+0.19}_{-0.17}$&1234.2 (1497)&99.9\\% \\\\ IGR J13109-5552 &$<$0.46            &1.55$^{+0.31}_{-0.26}$&$>$58  &             -      &   -            &       -       &-&2.29$^{+1.84}_{-0.99}$& 39.8 (54)& 83\\%\\\\  4C 74.26        &0.11$^{+0.01}_{-0.01}$&1.70$^{+0.03}_{-0.03}$& $>$78&6.44$^{+0.05}_{-0.06}$&184$^{+78}_{-66}$&104$^{+31}_{-26}$&0.63$^{+0.16}_{-0.14}$&1.09$^{+0.44}_{-0.37}$&1989.9 (2059)& 86\\%\\\\ S5 2116+81      &$<$0.18 &1.95$^{+0.14}_{-0.14}$&$>$81&-&- &-&1.54$^{+0.64}_{-0.44}$&2.87$^{+1.59}_{-1.04}$&89.2 (127) & -\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 4A}} \\vspace{0.2cm} \\begin{tabular}{lccccccccc} \\multicolumn{10}{c}{{\\bf 3C390.3 Spectral Fits Results: \\texttt{wa$_g$*(bb+cutoffpl+zga)}}}\\\\ \\hline {\\bf N$_H$} & {\\bf $\\Gamma$}     & {\\bf kT}             & {\\bf E$_{cut}$}& {\\bf E$_{line}^{\\ddagger}$} & {\\bf EW} & {\\bf C$_1^A$}  & {\\bf C$_2^B$}& {\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ {\\bf (10$^{22}$cm$^{-2}$)}&                     &{\\bf (keV)}          & {\\bf (keV)}    &  {\\bf (keV)}               & {\\bf (eV)}&           &       &     \\\\ \\hline -          &1.89$^{+0.02}_{-0.01}$ & 2.41$^{+0.21}_{-0.16}$ &  NC        &6.43$^{+0.03}_{-0.04}$         & 48$^{+9}_{-14}$ & 1.37$^{+0.16}_{-0.15}$&1.58$^{+0.23}_{-0.11}$ & 1970.2 (2231)& 37\\%\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\ddagger}$: line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3A. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 4B}} \\tiny \\vspace{0.2cm} \\begin{tabular}{lccccccccccc} \\multicolumn{12}{c}{{\\bf IGR J21247+5058 Spectral Fit Results: \\texttt{wa$_g$*pcfabs*pcfabs*(cutoffpl+zga)}}}\\\\ \\hline {\\bf $\\Gamma$}&{\\bf N$_H^1$}                      &{\\bf cf$_1$}& {\\bf N$_H^2$}                     &{\\bf cf$_2$}&{\\bf E$_c$}&{\\bf E$_{line}^{\\ddagger}$}&{\\bf EW}&{\\bf C$_1^A$}&{\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ &{\\bf (10$^{22}$cm$^{-2}$)}&                   &{\\bf (10$^{22}$ cm$^{-2}$)}&                    & {\\bf (keV)}&    {\\bf (keV)}    &{\\bf (eV)} &          & & &\\\\ \\hline 1.49$^{+0.05}_{-0.07}$&0.80$^{+0.15}_{-0.17}$&0.88$^{+0.11}_{-0.05}$&8.23$^{+1.70}_{-2.05}$&0.27$^{+0.04}_{-0.05}$&80$^{+100}_{-63}$&6.39$^{+0.07}_{-0.08}$&$<$30&0.65$^{+0.05}_{-0.06}$&0.86$^{+0.07}_{-0.07}$& 2276.8 (2555)& $>$99.9\\%\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\ddagger}$ line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3B. \\end{table*}   \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 5}} \\vspace{0.2cm} \\begin{tabular}{lcccccccccc} \\multicolumn{11}{c}{{\\bf  Spectral Fit Results: \\texttt{wa$_g$*wa*(pexrav+zga)}, 0$\\leq$R$\\leq$2, E$_c$=10000}}\\\\ \\hline {\\bf Name}      &  {\\bf N$_H$}                        &  {\\bf $\\Gamma$}          &  {\\bf R} & {\\bf E$_{line}$} & {\\bf $\\sigma_{line}$} & {\\bf EW}   & {\\bf C$_1^A$} &{\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ &{\\bf (10$^{22}$ cm$^{-2}$)}&                                     &              &    {\\bf (keV)}     & {\\bf (eV)}                  & {\\bf (eV)}  &                        &                      & &\\\\ \\hline QSO B0241+61    & 0.31$^{+0.04}_{-0.10}$&1.83$^{+0.07}_{-0.15}$&$>$0.37&6.40$^{+0.08}_{-0.03}$&10f&67$^{+40}_{-42}$&0.59$^{+0.28}_{-0.11}$&0.74$^{+0.36}_{-0.16}$& 224.1 (233)& 92\\%\\\\ B3 0309+411     &     -               &1.92$^{+0.03}_{-0.04}$&$>$1.20&6.33$^{+0.09}_{-0.10}$&10f & 70$^{+41}_{-43}$ &-&2.42$^{+1.06}_{-0.97}$&560.4 (602)& 99.99\\%\\\\               3C 111          &0.48$^{+0.03}_{-0.03}$&1.78$^{+0.05}_{-0.06}$&1.06$^{+0.63}_{-0.68}$&6.40$^{+0.07}_{-0.08}$&10f&$<$30&0.30$^{+0.07}_{-0.04}$&0.35$^{+0.11}_{-0.07}$& 1238.5 (1497)&99.7\\%\\\\ IGR J13109-5552 &0.26$^{+0.25}_{-0.25}$ &1.75$^{+0.23}_{-0.25}$&$>$0.1               &     -              &   -  &   -          &-&1.30$^{+2.69}_{-0.74}$& 40.5 (54)& 66\\%\\\\ 4C 74.26        &0.16$^{+0.02}_{-0.02}$&1.84$^{+0.05}_{-0.04}$&1.30$^{+0.67}_{-0.51}$&6.45$^{+0.04}_{-0.06}$&$<$186&51$^{+37}_{-16}$&0.33$^{+0.07}_{-0.06}$&0.54$^{+0.16}_{-0.17}$& 1977.0 (2059)&99.99\\%\\\\ S5 2116+81      &$<$0.22 &2.03$^{+0.20}_{-0.16}$ &$>$0.1&-&- & -&1.10$^{+0.73}_{-0.43}$&1.97$^{+1.51}_{-0.87}$& 88.5 (127)&64\\%\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 5A}} \\vspace{0.2cm} \\begin{tabular}{lccccccccc} \\multicolumn{10}{c}{{\\bf 3C 390.3 Spectral Fits Results: \\texttt{wa$_g$*wa*(bb+pexrav+zga)}, 0$\\leq$R$\\leq$2, E$_c$=10000}}\\\\ \\hline {\\bf N$_H$} & {\\bf $\\Gamma$}     & {\\bf kT}             &  {\\bf R}            &  {\\bf E$_{line}^{\\ddagger}$} & {\\bf EW}            & {\\bf C$_1^A$}  &{\\bf C$_2^B$}&  {\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ {\\bf (10$^{22}$cm$^{-2}$)}&             & {\\bf (keV)}         &                         &  {\\bf (keV)}                         & {\\bf (eV)}            &            &                &\\\\ \\hline -     &1.89$^{+0.03}_{-0.02}$ & 2.38$^{+0.27}_{-0.22}$ & 0.70$^{+0.48}_{-0.59}$ & 6.43$^{+0.03}_{-0.04}$ & 44$^{+10}_{-13}$ & 0.88$^{+0.42}_{-0.13}$ & 0.93$^{+0.55}_{-0.22}$&1965.4 (2231)& 98\\%\\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\ddagger}$: line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3A. \\end{table*}  \\begin{table*} \\scriptsize \\begin{center} \\centerline{{\\bf Table 5B}} \\tiny \\vspace{0.2cm} \\begin{tabular}{lccccccccccc} \\multicolumn{12}{c}{{\\bf IGR J21247+5058 Spectral Fit Results: \\texttt{wa$_g$*pcfabs*pcfabs*(pexrav+zga)}, 0$\\leq$R$\\leq$2, E$_c$=10000}}\\\\ \\hline {\\bf $\\Gamma$}&{\\bf N$_H^1$}                      &{\\bf cf$_1$}& {\\bf N$_H^2$}                     &{\\bf cf$_2$}& {\\bf R}& {\\bf E$_{line}^{\\ddagger}$}&{\\bf EW}&{\\bf C$_1^A$}& {\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}&{\\bf Prob.$^{\\dagger}$}\\\\ &{\\bf (10$^{22}$cm$^{-2}$)}&                   &{\\bf (10$^{22}$ cm$^{-2}$)}&                    &               & {\\bf (keV)}    &{\\bf (eV)} &        &   &\\\\ \\hline 1.77$^{+0.05}_{-0.05}$&8.14$^{+1.81}_{-1.69}$&0.36$^{+0.03}_{-0.04}$&0.99$^{+0.16}_{-0.18}$&0.88$^{+0.05}_{-0.03}$&0.93$^{+0.59}_{-0.45}$ &6.38$^{+0.08}_{-0.08}$&$<$30&0.41$^{+0.08}_{-0.06}$&0.50$^{+0.10}_{-0.08}$& 2341.7 (2555)& 99.99\\% \\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\ddagger}$: line width fixed to 10 eV.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}.\\\\ $^{\\dagger}$: fit improvement with respect to table 3. \\end{table*}   \\begin{table*} \\scriptsize \\centering \\begin{center} \\centerline{{\\bf Table 6}} \\vspace{0.2cm} \\begin{tabular}{lcccccccc} \\multicolumn{9}{c}{{\\bf  Spectral Fit Results: \\texttt{wa$_g$*wa*(pexrav+zga)}}}\\\\ \\hline {\\bf Name}     &{\\bf N$_H$}                      &{\\bf $\\Gamma$}&{\\bf R}&{\\bf E$_c$}&{\\bf EW$^{\\dagger}$}  &{\\bf C$_1^A$}&{\\bf C$_2^B$}&{\\bf $\\chi^2$ (dof)}\\\\ & {\\bf (10$^{22}$ cm$^{-2}$}&                       &            &{\\bf (keV)}&{\\bf (eV)}                     &                     &                       &           \\\\ \\hline QSO B0241+61    &0.21$^{+0.15}_{-0.10}$&1.64$^{+0.22}_{-0.14}$& 0.56$^{+0.82}_{-0.41}$ &    $>$86       &73$^{+44}_{-42}$& 0.85$^{+0.44}_{-0.32}$&1.12$^{+0.65}_{-0.48}$&222.2 (233)\\\\ B3 0309+411     &  -    & 1.90$^{+0.08}_{-0.08}$&3.48$^{+2.24}_{-1.58}$&35$^{+91}_{-17}$&59$^{+42}_{-43}$&-&6.89$^{+9.39}_{-3.90}$&554.2 (601)\\\\          3C 111          &0.46$^{+0.03}_{-0.03}$&1.73$^{+0.06}_{-0.06}$&0.85$^{+0.57}_{-0.58}$&126$^{+193}_{-50}$&    $<$30     &0.40$^{+0.11}_{-0.08}$&0.52$^{+0.19}_{-0.13}$&1227.9 (1497)\\\\ 3C 390.3&    -                                       &1.89$^{+0.03}_{-0.02}$ &0.60$^{+0.60}_{-0.44}$ &$>$300&41$^{+12}_{-11}$&0.95$^{+0.17}_{-0.25}$&1.01$^{+0.39}_{-0.30}$& 1966.3 (2231)\\\\ 4C 74.26        &0.14$^{+0.02}_{-0.03}$&1.79$^{+0.06}_{-0.07}$&1.22$^{+0.69}_{-0.70}$&100$^{+680}_{-52}$&88$^{+23}_{-20}$&0.95$^{+0.17}_{-0.25}$&1.01$^{+0.39}_{-0.30}$ & 1976.0 (2060)\\\\ IGR J21247+5058& complex$^{\\star}$&1.48$^{+0.06}_{-0.06}$&$<$0.21&79$^{+23}_{-15}$&$<$30&0.65$^{+0.05}_{-0.08}$&0.86$^{+0.08}_{-0.12}$& 2276.7 (2555) \\\\ \\hline \\end{tabular} \\end{center} \\scriptsize $^{\\dagger}$: line parameters fixed at value obtained in table 3.\\\\ \\emph{A}: cross-calibration constant between \\emph{XMM} and \\emph{BAT}; \\emph{B}: cross-calibration constant between \\emph{XMM} and \\emph{INTEGRAL/ISGRI}\\\\ $^{\\star}$: N$_H^1$=7.86$^{+2.02}_{-1.66}$, cf$_1$=0.27$^{+0.04}_{-0.05}$; N$_H^2$=0.77$^{+0.18}_{-0.13}$, cf$_2$=0.89$^{+0.10}_{-0.06}$. \\end{table*}   \\begin{small} \\begin{figure*} \\centering \\includegraphics[scale=0.3,angle=-90]{qso0241_pex_bat_uf.ps} \\hspace{0.5cm} \\includegraphics[scale=0.3,angle=-90]{b0309_pex_uf.ps}\\\\ \\caption{Table 6 model for QSO B0241+62 (left panel) and B3 0309+411 (right panel). The model is a cut-off power law absorbed both by Galactic and intrinsic column density  reflected by neutral material plus a narrow Gaussian component describing the iron line.} \\label{qso_b3} \\end{figure*} \\end{small}  \\begin{small} \\begin{figure*} \\centering \\includegraphics[scale=0.3,angle=-90]{3c111_pex_bat_uf.ps} \\hspace{0.5cm} \\includegraphics[scale=0.3,angle=-90]{3c390_pex_bat_uf.ps}\\\\ \\caption{Table 6 model for 3C 111 (left panel) and 3C390.3 (right panel). The model is a cut-off power law absorbed both by Galactic and intrinsic column density reflected by neutral material plus a narrow Gaussian component describing the iron line.} \\label{3c111_3c390} \\end{figure*} \\end{small}  \\begin{small} \\begin{figure*} \\centering \\includegraphics[scale=0.3,angle=-90]{4c74_pex_bat_uf.ps} \\hspace{0.5cm} \\includegraphics[scale=0.3,angle=-90]{21247_pex_bat_uf.ps}\\\\ \\caption{Table 6 model for 4C 74.26 (right panel) and  IGR J21247+5058 (left panel). The model is a cut-off power law absorbed both by Galactic and intrinsic column density  reflected by neutral material plus a narrow Gaussian component describing the iron line.} \\label{4c_21247} \\end{figure*} \\end{small}  \\begin{small} \\begin{figure*} \\centering \\includegraphics[scale=0.3,angle=-90]{s2116_po_bat_uf.ps} \\hspace{0.5cm} \\includegraphics[scale=0.3,angle=-90]{13109_po_uf.ps}\\\\ \\caption{Table 3 model for S5 2116+81 (left panel) and IGR J13109-5552 (right panel). The model is a simple power law, absorbed both by Galactic and intrinsic column densities.} \\label{s5_13109} \\end{figure*} \\end{small}  \\begin{small} \\begin{figure*} \\centering \\includegraphics[scale=0.4,angle=-90]{ew_r.ps} \\caption{Reflection fraction vs. EW for the sample analysed here and the sources presented in \\citet{b21}. Blue squares are radio quiet sources, while radio loud sources are represented by magenta circles. Arrows represent upper or lower limits on the parameter values.} \\label{r_ew} \\end{figure*} \\end{small}"
    },
    "0809/0809.2066_arXiv.txt": {
        "abstract": "{ We present an update to the photometric calibration of the COMBO-17 catalogue on the Extended Chandra Deep Field South, which is now consistent with the GaBoDS and MUSYC catalogues. As a result, photometric redshifts become slightly more accurate, with $<0.01$ rms and little bias in the $\\delta z/(1+z)$ of galaxies with $R<21$ and of QSOs with $R<24$. With increasing photon noise the rms of galaxies reaches 0.02 for $R<23$ and 0.035 at $R\\approx 23.5$. Consequences for the rest-frame colours of galaxies at $z<1$ are discussed. ",
        "introduction": "Almost five years ago, a catalogue of the Extended Chandra Deep Field South (ECDFS) was published by the COMBO-17 survey project \\citep{Wolf04}. It comprised photometry in 17 bands, viz. the broad-band filters UBVRI and 12 medium-band filters covering the wavelength range from 400 to 930~nm. On the basis of the SEDs they derived and published photometric classifications into star, galaxy, QSO or white dwarf, as well as photometric redshifts for galaxies and QSOs. While the A901 and S11 fields of the COMBO-17 survey had straightforward calibrations, it was known that the ECDFS calibration was ambiguous. In spite of this, photometric redshifts in the ECDFS were of very good quality with $\\sigma_z/(1+z)\\approx 0.02$ for galaxies with $R<23$ and for QSOs (=type-1 AGN) brighter than $M_B<-22$.  Recently, it became clear that the photometric calibration was indeed in error with an almost monotonic drift across wavelength. Since the publishing of the original COMBO-17 ECDFS catalogue, two separate projects -- MUSYC  \\citep{MUS} and GaBoDS \\citep{Gab} -- have constructed broadband photometric catalogues of the ECDFS. The GaBoDS consortium retrieved and combined all existing WFI imaging (up to Dec 2005), including raw COMBO-17 data, and calibrated them from nightly standard star imaging. Both the GaBoDS and MUSYC projects use the same WFI data; the MUSYC catalogue adds NIR data to the GaBoDS--reduced optical data. It was through comparison of the catalogue photometry between them and COMBO-17 that the calibration issue was first noticed.  Broad-band photo-z's are susceptible to calibration offsets, but photo-z's from medium-band surveys are more robust against calibration errors. This applies to both random calibration offsets, whose impact is diminished by a large number of bands, and to systematic calibration drifts, since spectral features are reliably located by medium-band filters even when the SED is globally wrong. Hence, the presence of the calibration drift could not be deduced from the fully satisfying photo-z performance.  In this paper, we publish a refined calibration alongside an updated catalogue and explain the origin of the calibration error (see Sect.~2), the details of which might be useful for future multi-band surveys. We then compare the new calibration to the two other ground-based surveys that quantified the calibration difference in the broad-bands originally (see Sect.~3). In Sect.~4 we discuss subtle consequences for the photo-z accuracy and the more important update of the rest-frame colours. ",
        "conclusions": "The calibration of the CDFS field of COMBO-17 is updated over the original publication of the data in 2004. Each COMBO-17 field is calibrated by two spectrophotometric standard stars in each of its fields. While the calibration of the other COMBO-17 fields was straightforward, the two stars on the CDFS suggested calibrations that were inconsistent in colour at the 0.15~mag level from B to I. Both were marginally consistent with the colours of the Pickles atlas, so the choice was unconstrained. Wolf et al. (2004) ended up choosing the wrong star and introduced a colour bias to the blue. Here, we have changed the calibration to follow the other star, and it is now consistent with both the GaBoDS and MUSYC photometry. The consequences of the calibration change for the photometric redshifts is little when all 17 filters are used, but larger when only broad bands are used. Broad-band photo-z's hinge more on colours than on spectral features that are traced in medium-band photo-z's. However, a small positive change in the photo-z's is observed. Galaxies at $R<21$ and QSOs at $R<24$ have photo-z's accurate to within $<0.01$ rms of $\\delta z/1(+z)$ after exclusion of outliers. The new catalogue version is accessible from the COMBO-17 or CDS websites."
    },
    "0809/0809.5193_arXiv.txt": {
        "abstract": "We propose an experimental approach to {\\it macro}scopically test the Kochen-Specker theorem (KST) with superconducting qubits. This theorem, which has been experimentally tested with single photons or neutrons, concerns the conflict between the contextuality of quantum mechnaics (QM) and the noncontextuality of hidden-variable theories (HVTs). We first show that two Josephson charge qubits can be controllably coupled by using a two-level data bus produced by a Josephson phase qubit. Next, by introducing an approach to perform the expected joint quantum measurements of two separated Josephson qubits, we show that the proposed quantum circuits could demonstrate quantum contextuality by testing the KST at a macroscopic level.  PACS number(s): 03.65.Ta, % 03.67.Lx, % 85.25.Dq. % ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0852_arXiv.txt": {
        "abstract": "The past decade has witnessed impressive progress in our understanding of the physical properties of massive stars in the Magellanic Clouds, and how they compare to their cousins in the Galaxy.  I summarise new results in this field, including evidence for reduced mass-loss rates and faster stellar rotational velocities in the Clouds, and their present-day compositions.  I also discuss the stellar temperature scale, emphasizing its dependence on metallicity across the entire upper-part of the Hertzsprung-Russell diagram. ",
        "introduction": "The prime motivation for studies of early-type stars in the Magellanic Clouds over the past decade has been to quantify the effect of metallicity ($Z$) on their evolution.  The intense out-flowing winds in massive stars are thought to be driven by momentum transferred from the radiation field to metallic ions (principally iron) in their atmospheres; a logical consequence of this mechanism is that the wind intensities should vary with $Z$ (Kudritzki et al., 1987).  Monte Carlo simulations predict that, for stars with T$_{\\rm eff}\\,>$\\,25,000\\,K, the wind mass-loss rates should scale with metallicity as $Z^{0.69}$ (Vink et al., 2000; 2001).  This has a dramatic impact on their subsequent evolution.  For example, an O-type star in the SMC should lose significantly less mass over its lifetime than a star in the Galaxy, thus retaining greater angular momentum.  This could then lead to different late-phases of evolution such as the type of core-collapse supernova (SN), and offers a potential channel for long duration, gamma-ray bursts at low $Z$.   The $Z$-dependence of the {\\em initial} rotational velocity distributions of massive stars and the importance of rotationally-induced mixing have also been active areas of research. For instance, \\cite{mgm99} noted that the relative fraction of Be- to B-type stars increases with decreasing metallicity\\footnote{The sample of Maeder et al. comprised only one SMC cluster, NGC\\,330, long known to have a significant Be-fraction (e.g. Grebel et al., 1992) and sometimes suggested as a `pathological' case.  However, new results from Martayan et al. (these proceedings) also find similarly large fractions for other SMC clusters.}, suggesting that this might arise from faster rotational velocities at lower $Z$.  The recent generation of evolutionary models has explored the effects of rotational mixing (e.g. Heger \\& Langer, 2000; Meynet \\& Meynet, 2000), with the prediction of larger relative surface-nitrogen enhancements at faster rotation rates, and at lower Z (Maeder \\& Meynet, 2001).  To date, we have lacked sufficient observations to explore the effects of metallicity thoroughly.  Robust empirical results were needed with which to confront both the stellar wind and evolutionary models for early-type stars -- here I summarise recent observations and quantitative analyses toward this objective.  \\newpage ",
        "conclusions": ""
    },
    "0809/0809.0313_arXiv.txt": {
        "abstract": "We have derived a model of the Kuiper belt luminosity function exhibited by a broken power-law size distribution. This model allows direct comparison of the observed luminosity function to the underlying size distribution. We discuss the importance of the radial distribution model in determining the break diameter. We determine a best-fit break-diameter of the Kuiper belt size-distribution of $30<D_b<90$ km via a maximum-likelihood fit of our model to the observed luminosity function. We also confirm that the observed luminosity function for $m(R)\\sim21-28$ is consistent with a broken power-law size distribution, and exhibits a break at $m(R)=26.0^{+0.7}_{-1.8}$. ",
        "introduction": "} The size distribution (SD) of a belt of planetesimals can provide information on the physical and dynamical properties of those objects. Collisions between belt members causes the SD to evolve over time. Low velocity collisions can cause objects to accrete into larger bodies, while high velocity collisions smash bodies into smaller pieces. A generic outcome of collisional evolution simulations of a belt of planetesimals is a power-law SD with a steep slope for large objects, breaking to a shallower slope at some object diameter. Waves or other shapes can be induced in a size distribution, but the general behavior is that of a broken power-law \\citep{Obrien2003}. The location of the transition between the steep and shallow SD slopes (break), as well as the small and large object slopes provides information on the strengths of the colliding bodies, and constrains time-scales of accretion and total mass in the belt region \\citep{Safronov1969,Dohnanyi1969,Kenyon2001,Kenyon2002,Kenyon2004,Bottke2005}  The SD of the Kuiper belt has been inferred from a measurement of the belt's luminosity function (LF) \\citep[see for example][(F08) and references therein]{Fraser2008}. F08 show that for the largest objects in the belt (diameter $D\\gtrsim 100$km), the SD is well represented by a power-law $N(D)\\propto D^{-q}$ with a slope $q\\sim4.3$. %  Kuiper belt surveys designed to measure the Kuiper belt object (KBO) faint-end LF ($m(R)\\gtrsim 26$), have found evidence to suggest that the LF deviates from a power-law \\citep{Bernstein2004}. \\citet{Bernstein2004} found that a distribution with a steep slope for bright objects, that ``rolled over\" to a shallow slope for faint objects was necessary to describe the substantial dirth of objects with $m(R)\\gtrsim 27$. They concluded that the underlying KBO SD transitions (``breaks'') to shallower slopes as predicted from collisional evolution models. Using archival Subaru Suprimecam images, \\citet{Fuentes2008} provided a substantial increase in the sample of faint ($m(R)\\sim26$) KBOs and concluded that the luminosity function breaks from a steep to shallow slope at $m(R)\\sim24.3$.  \\citet{Bernstein2004} and \\citet{Fuentes2008} use analytic functions to represent the shape of the observed LF. These functional forms were not derived from physical or theoretical properties, making infering the underlying KBO SD based on the reported LF difficult. We present here a new functional form of the LF that results from the observation of a belt of planetesimals exhibiting a broken SD. This function provides a direct interpretation of the observed LF in terms of  the size and radial distributions.  In Section~\\ref{sec:derivation}, we derive the new functional form. We fit the function to the available  observations of the Kuiper belt in Section~\\ref{sec:fit}, describe what future efforts are necessary to constrain the SD in Section~\\ref{sec:errors}, briefly discuss implications of the Kuiper belt SD in Section~\\ref{sec:discussion}, and make concluding remarks in Section~\\ref{sec:conclusions}.. ",
        "conclusions": "} We have developed a model of an LF exhibited from a planetesimal belt with a broken power-law SD. This model allows a direct interpretation of the observed LF in terms of the underlying SD assuming a broken power-law shape. We fit this model to the observed Kuiper belt LF and the break first observed by \\citet{Bernstein2004}. The best-fit parameters are presented in Table~\\ref{tab:fit}. We find that the break in the SD must occur at diameters $27-87$ km ($m(R)\\sim25.8$) for reasonable choices of the radial distribution. We have demonstrated that the inference of the SD from the LF requires an accurate model of the radial distribution. Additional detections of KBOs faintward of $m(R)=26$, and a proper treatment of the KBO orbital distribution are required to provide an accurate measure of the break size and faint-end slope."
    },
    "0809/0809.0125_arXiv.txt": {
        "abstract": "It is widely written and believed that Edwin Hubble introduced the terms `early' and `late types' to suggest an evolutionary sequence for galaxies. This is incorrect. Hubble took these terms from spectral classification of stars to signify a sequence related to complexity of appearance, albeit based on images rather than spectra. The temporal connotations of the terms had been abandoned prior to his 1926 paper on classification of galaxies. ",
        "introduction": "The terms `early' and `late type' in astrophysics have been applied to both stars and galaxies.  Spectral classification of stars follows an early-to-late sequence O-B-A-F-G-K-M with recent additions of L-T \\citep{kirkpatrick99}. This classification closely relates to a sequence in temperature from hot to cool stars. Morphological classification of galaxies is based on a number of factors, including ellipticity, the size of the nuclear region relative to the spiral arms, and the smoothness of the image. A commonly used classification is the revised and extended Hubble system \\citep{deVaucouleurs59,Sandage61,Sandage75} that is based on Hubble's (\\citeyear{Hubble26}) original scheme for `extragalactic nebulae', and Reynold's (\\citeyear{Reynolds20}) earlier ideas.  It follows an early-to-late sequence, ellipticals-lenticulars-spirals-irregulars, E-S0-Sa-Sb-Sc-Sd-Sm-Im (ignoring the barred/unbarred characteristic). \\citet{Sandage05} has reviewed the history of this development.  At first glance, there appears to be no relation between the early to late type sequences of stars and galaxies other than the terminology. Before the implications of $E=mc^2$ \\citep{Einstein05} were realised, and to explain the Hertzsprung-Russell diagram, it was natural to suppose that stars cooled from early to late spectral types because there was no established mechanism for Myr stability of stellar atmospheres (cf.\\ cooling of brown dwarfs, \\citealt{burrows97}). There is now a wide-spread but mistaken belief that Hubble chose this terminology because he thought that the morphological sequence was also a temporal sequence.  For example, \\citet{BM98book} wrote ``Hubble suggested that galaxies evolved from the left-hand end of this sequence to the right. This now discredited speculation lives on in the convention ... early-type ... late-type galaxies.''  While, \\citet{CL02book} noted ``Although it is now not thought this evolutionary sequence is correct, Hubble's nomenclature, in which ellipticals are `early' type and spirals and irregulars `late', is still commonly used.''  Similar explanations can be found in other textbooks (e.g., \\citealt{Tayler93book,Shore03book,CO06book}). The main aim of this article is to show that these explanations for the origin of the terminology are incorrect, and to illuminate the correct explanation. ",
        "conclusions": ""
    },
    "0809/0809.2120_arXiv.txt": {
        "abstract": "We study the kilohertz quasi-periodic oscillations (kHz QPOs) and the band-limited noise (BLN) in the 0.5--16 Hz range observed simultaneously on the horizontal branch (HB) and on the upper normal branch (NB) of the brightest neutron star Low-mass X-ray Binary (LMXB) Scorpius X--1 with the observations performed with the {\\it Rossi X-Ray Timing Explorer (RXTE)}.  We find that the twin kHz QPO frequencies are positively correlated with the flux variations taking place on the BLN time scales on the HB, in contrast to the anti-correlation held on the time scale of the normal branch oscillation (NBO) on the NB reported previously, suggesting that although they occur in sequence along the color-color tracks, the BLN and the NBO are of different origins. We also show the evidence that the frequency separation between the twin kHz QPOs decreases with the flux by $2\\sim~3$ Hz on the BLN time scales, which is consistent with the trend on the longer time scale that the Z source traces the HB. This further suggests that the flux variation associated with the BLN originates from the mass accretion rate variation in the disk accretion flow. We discuss the implications of these results for our understanding of the BLN. ",
        "introduction": "The brightest neutron star Low-mass X-ray Binaries (LMXBs) were divided into two main types, namely,  the ``Z\" sources and the ``atoll\" sources, based on their spectral and timing behavior studied twenty years ago (Hasinger \\& van der Klis 1989). The brighter type is the ``Z\" sources, because they trace out roughly a Z-shape in X-ray color-color diagram (CD), which consists of three branches, called the horizontal branch (HB), the normal branch (NB), and the flaring branch (FB), respectively. The picture is yet not well understood. Recent observations of the transient Z source XTE J1701-462 have brought new clues to the formation of the tracks (Homan et al. 2007), while observations in the past decade revealed some weak neutron star LMXBs accreting at lower mass accretion rates, some of which show similarities to the atoll sources (see the review van der Klis 2006).  The power spectra of neutron star LMXBs usually contain several variability components in a wide range of frequency up to kilohertz. Aperiodic variability includes Quasi-periodic Oscillation (QPO, narrow feature) or noise (broad structure). Band-limited noise (BLN) component, formerly called high-frequency noise (HFN) in atoll sources and low-frequency noise (LFN) in Z sources, has a flat shape up to a break frequency. For the Z sources, the  BLN is observed only on the HB and the very upper part of the NB. In addition to the BLN, various QPOs are observed in different branches of the Z sources. The characteristic frequencies of the variability components often vary monotonically along the tracks in the CD.  The brightest persistent neutron star LMXB \\mbox{Sco X$-$1} is a Z source. In \\mbox{Sco X$-$1} there are three distinct types of QPOs observed, i.e., the near 6 Hz normal-branch oscillation (NBO), the near 45 Hz horizontal-branch oscillation (HBO) with a harmonic near 90 Hz, and the kilohertz quasi-periodic oscillations (kHz QPOs). The kHz QPOs and the harmonic of the HBO, were first seen in the persistent flux of \\mbox{Sco X$-$1} from the very early RXTE observations (van der Klis et al. 1996a; 1997).  The characteristic frequencies of the X-ray variability in X-ray binaries are probably related to some characteristic time scales in the system, some of which may be used to constrain the mass and the spin of the compact object. Many authors have tried to explain these  frequencies and their properties. Models for the kHz QPOs, which are based on orbital and epicyclical motions, include relativistic precession model (Stella \\& Vietri 1999) and sonic-point and spin-resonance model (Miller et al. 1998; Lamb \\& Miller 2004). Some other models for the kHz QPOs are based on disk oscillations or modes in the accretion flow, e.g., nonlinear resonance (Abramowicz \\& Kluzniak 2001; Abramowicsz 2005), Alfve\\'{n} wave oscillation (Zhang 2004), and kink mode (Li \\& Zhang 2005). Models for the HBOs include magnetospheric beat-frequency models (Alpar \\& Shaham 1985; Lamb et al. 1985), the Lense-Thirring precession model (Stella \\& Vietri 1998), and the magnetically warped precessing-disk model (Lai 1999; Shirakawa \\& Lai 2002). For the NBO or the FBO, disk oscillation (e.g. Alpar et al. 1992), radiation-force feedback instability (Fortner et al. 1989), and the oscillation of a spherical shell (Titarchuk 2001) have been proposed. Related to \\mbox{Sco X$-$1}, a possible explanation for the NBO and the coupling between the twin kHz QPOs and the NBO was discussed by Hor$\\acute{a}$k et al. (2004). In addition, a two-oscillator (TO) model has been developed to explain several  characteristic frequencies in neutron star X-ray binaries (Titarchuk \\& Osherovich 1999; Titarchuk et al. 1999).  The frequency correlation between these QPOs and noises has been found to occur in different source types and over a wide range of luminosities (e.g., Psaltis, Belloni, and van der Klis 1999; Belloni, Psaltis, and van der Klis 2002; van der Klis 2006). The correlations between the frequencies of the QPOs and the X-ray count rate on long-term and short-term time scales were observed as well (note: past X-ray timing measurements we discussed below were obtained with two proportional counter instruments, i.e., the Medium Energy (ME) Proportional Counter on board the European Space Agency's X-ray Observatory (EXOSAT) and the Proportional Counter Array (PCA) on board the RXTE. Therefore the correlations between the frequencies of various QPOs and the X-ray count rate as seen with the EXOSAT and the RXTE are generally consistent). In the atoll sources on the time scales of hours to days, the kHz QPO frequency is found correlated with the X-ray count rate, but on time scales longer than days the correlation does not hold and we see the ``parallel track'' phenomenon (e.g., \\mbox{4U~0614+091}, Ford et al. 1997; \\mbox{4U 1608$-$52}, Yu et al. 1997; M\\'{e}ndez et al. 1999). An anti-correlation between the kHz QPO frequency and the count rate associated with the mHz QPO, which might correspond to a thermonuclear burning mode on neutron star, was also found in 4U 1608-52 (Yu \\& van der Klis 2002). In the Z sources on the time scales of hours to days, the kHz QPO frequencies and the X-ray count rate appear to positively correlate on the HB and sometimes anti-correlate on the NB as inferred from observations reported in the literature. In \\mbox{GX 5$-$1} and \\mbox{GX 17+2}, the kHz QPO frequencies increase from the HB to the upper NB (van der Klis et al. 1996b), while the X-ray count rate increases from the HB to the HB/NB vertex and then decreases along the upper NB (Kuulkers et al. 1994; Wijnands et al. 1997; Homan et al. 2002), therefore a positive correlation on the HB and an anti-correlation on the upper NB can be inferred. In \\mbox{GX 340+0} and \\mbox{Cyg X$-$2}, the kHz QPO frequencies and the count rate increase along the HB to the HB/NB vertex (Jonker et al. 1998; Wijnands et al. 1998), so there is a positive correlation on the HB. In \\mbox{Sco X$-$1}, the kHz QPO frequencies increase on the NB (van der Klis et al 1996a) and the X-ray count rate decreases from the HB to the NB (Dieters \\& van der Klis 2000), so there is an anti-correlatation on the NB (Yu et al. 2001). Similar to those of the kHz QPOs, the frequency of the HBO is correlated with the X-ray count rate as well. In general, the HBO frequency and the count rate appear to positively correlate on the HB. For example, in \\mbox{GX 5$-$1}, the HBO frequency is positively correlated with the source count rate on the HB (Lewin et al.1992; Kuulkers et al. 1994). But the situation is complex on the NB. In \\mbox{Cyg X$-$2}, the HBO frequency increases from the HB to the HB/NB vertex and then remains approximately constant on the upper NB, while the count rate increases from the HB to the very upper NB (Wijnands et al. 1998). In \\mbox{GX 340+0}, the HBO frequency increases from the HB to the HB/NB vertex and then remains approximately constant on the upper NB, while the count rate increases from the HB to the HB/NB vertex and decreases from the NB to the NB/FB vertex (Kuulkers \\& van der Klis 1996; Jonker et al. 1998). In \\mbox{GX 17+2}, the HBO frequency increases from the HB to the middle NB and then decreases from the middle NB to the NB/FB vertex, while the count rate increases from the HB to the HB/NB vertex and decreases from the NB to the NB/FB vertex (Wijnands et al. 1996; Homan et al. 2002). Therefore on the upper NB, previous observations suggest that there is no significant correlation between the HBO frequency and the X-ray count rate in \\mbox{Cyg X$-$2} and \\mbox{GX 340+0}, while there is an anti-correlation in GX 17+2. On the lower NB, a positive correlation was seen in \\mbox{GX 17+2}. Similar to the atoll sources, on longer time scales, there is no correlation between the QPO frequencies and the X-ray count rate.  Previous studies show the BLN on the HB and the NBO on the NB in \\mbox{Sco X$-$1} occur in time and frequency sequence (see Figure 2 of van der Klis et al. 1997). One may wonder if they are of similar origins. In this paper, we investigated the dependence of the kHz QPOs on the flux variation due to the BLN in \\mbox{Sco X$-$1}. We found that on the time scales of the BLN, there is a positive correlation between the kHz QPO frequencies and the X-ray flux , and the frequency separation of the kHz QPOs deceases with the X-ray flux, which is different from what was seen for the NBO. ",
        "conclusions": "We have performed a study of f the band-limited noise (BLN) in \\mbox{Sco X$-$1}. We conclude \\begin{itemize} \\item Both the lower and the upper kHz QPO frequencies are positively correlated with the flux variation taking place on the BLN time scales, roughly quantitatively consistent with the kHz QPO frequency vs. the X-ray flux correlation hold along the Z track. \\item The frequency separation between the two kHz QPOs decreases with increasing kHz QPO frequencies on the BLN time scales by a few hertz, consistent with the overall trend held on longer time scales at the same kHz QPO frequency range. \\item Based on the above properties of the BLN which is different from those of the NBO, we conclude that the BLN likely corresponds to the variation in the mass accretion rate of the disk accretion flow. \\item The results obtained for the BLN and the NBO in \\mbox{Sco X$-$1} suggest that further observational studies of $\\Delta\\nu$ at highest kHz QPO frequencies are crucial for testing kHz QPO models. \\end{itemize}"
    },
    "0809/0809.4583_arXiv.txt": {
        "abstract": "We present a preliminary report on the broadband optical photometry of the 2006 outburst of the recurrent nova RS Ophiuchi.  These data were obtained using the robotic 2m Liverpool Telescope and cover the outburst from day 27 through day 548. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1645_arXiv.txt": {
        "abstract": "We present metallicities of 2690 RR0 Lyrae stars observed toward the MACHO Survey fields in the Galactic Bulge. These \\feh values are based upon an empirically calibrated relationship that uses the Fourier coefficients of the light curve and are accurate to $\\pm$ 0.2 dex.  The majority of the RR0 Lyrae stars in our sample are located in the Galactic Bulge, but 255 RR0 stars are associated with Sagittarius dwarf galaxy. We find that the RR0 Lyrae stars that belong to the Galactic Bulge have average metallicities $\\feh = -$1.25, with a broad metallicity range from $\\feh = -$2.26 to $-0.15$. The RR0 stars from Sagittarius dwarf galaxy have lower average metallicity of $\\feh = -1.55 \\pm 0.02$, with an intrinsic dispersion of 0.25 dex, similar to that in the bulge.  A correlation between metallicity and galactocentric distance is found, in a sense that for the metal-poor RR0 Lyrae stars (\\feh $< -$1.5 dex), the closer a star is to the Galactic center, on average, the more metal rich it is.  However, for the metal-rich RR0 Lyrae stars (\\feh $> -$1.2 dex), this trend is reversed. Using mean magnitudes of MACHO RR Lyrae stars, we searched for the evidence of the Galactic bar, and found marginal evidence of a bar. The absence of a strong bar indicates that the RR Lyrae in the bulge represent a different population than the majority of the bulge stars, which are metal rich and are part of a bar. ",
        "introduction": "The bulge RR Lyrae variables are likely to be among the oldest and most metal poor stars in the bulge, so their metal abundances are of considerable importance in determining the mix of populations in the Galactic Bulge, and vital to our understanding of the nature of the bulge itself \\citep{walkThisTurn}.  Studying abundances of RR Lyrae stars lets us probe early chemical evolution of the Milky Way and allows the chemical history of the the oldest population of the bulge to be traced out.  Direct studies of the bulge are difficult due to the severe crowding toward the central regions of the Galaxy and the large, patchy, reddening along the line of sight. Therefore, astronomical literature contains limited spectroscopy of RR Lyrae stars toward the bulge and most spectroscopic studies focus on Baade's window (BW) centered roughly on the globular cluster NGC 6522 at $(l, b) = (1 \\hbox{$.\\!\\!^\\circ$} 0, -3 \\hbox{$.\\!\\!^\\circ$} 9)$.  \\citet{but76} measured $\\Delta S$ abundances \\citep{preston59} for nine RR Lyrae stars in BW and derived a mean metallicity of $<\\feh>$ $= -0.82 \\pm 0.14 $ from a broad abundance distribution.  \\citet{gratzi} measured $\\Delta S$ values for 17 BW variables and found $<\\feh>$ $= -0.76 \\pm 0.12$. Again the results suggested a wide range of abundance. \\citet{walkThisTurn} found a somewhat lower abundance of $<\\feh>$ $= -1.0$. from 59 RR Lyrae variables in BW. In contrast to other studies, they concluded  that the RR Lyrae stars in BW had a relatively narrow range in abundance.  In this paper, we determine iron metallicities of bulge RR0\\footnote{RR0 stars have been traditionally called RRab stars and are simply fundamental pulsators. See \\citet{alc00} for further explanation on the more intuitive nomenclature.}  Lyrae stars determined indirectly from pulsational properties of the variables.  A number of studies have shown that Fourier parameters of light curves of RR Lyrae stars are related to their physics properties, including the metallicity \\citep{sc93}.  \\citet{jk96} used spectroscopic and photometric observations of field RR0 stars to calibrate a relationship between the \\feh of RR0 Lyrae stars and Fourier parameters. Because of the accurate fit of the Fourier formula to the observed metallicities of the RR0 Lyrae field stars, it appears to be very attractive to use this on various large databases. Recently, \\citet{kinemuchi06} applied the \\citet{jk96} formula to find photometric metallicities of 433 of 1188 RR0 Lyrae stars in the Northern Sky Variability Survey.  Although this \\feh method has been used frequently and appears to be reasonably reliable, there has been some question as to the validity of the \\citet{jk96} method \\citep[see discussions in][]{diFab05, grat04}. This is investigated in detail in Appendix A where it is demonstrated that the \\feh derived from MACHO LMC RR0 Lyrae lightcurves using \\citet{jk96} method agrees well with spectroscopic determinations from \\citet{grat04} and \\citet{bori06}.  We utilize the Fourier coefficients for 2690 RR0 Lyrae stars in the MACHO bulge survey to derive \\feh abundances of these variables.  This paper first addresses the observational data and Fourier Coefficients in \\S2.  Next, the derivation of \\feh, along with various tests, is described in \\S3. The presentation of bulge and Sgr metallicities and implications regarding these metallicities follow in \\S4.  Using reddenings determined from the minimum V-band light of the RR0 Lyrae stars, in \\S5, trends coinciding with distance are elucidated. A summary is presented in \\S6. ",
        "conclusions": "Metallicity measurements of 2690 RR0 Lyrae stars are presented from the MACHO bulge database using the Fourier coefficients of their light curves, with an error of 0.20 dex. Parameters derived for individual stars are available electronically in Table~\\ref{tab1} for the bulge stars and Table~\\ref{tab2} for Sgr stars. We find that RR0 stars in the Galactic Bulge span a broad metallicity range from $\\rm \\feh = -2.26$ to $-0.15$ dex, and have an average metallicity of $\\rm <\\feh>$ $= -$1.25.  This compares favorably to other bulge metallicity studies.  Using mean magnitudes of MACHO RR Lyrae stars we searched for the evidence of the Galactic bar, and found a slight signature of a Galactic bar with Galactic latitude $|b| < 3.5^\\circ$. The most straightforward interpretation of the absence of a strong bar is that the metal-rich RR Lyrae in the bulge represent a different population than the metal-poor majority. As we also found that the average metallicity in this region is considerably less varied and diverse, we conclude that there is evidence indicating the presence of a population belonging to the inner extension of the halo, which is relatively metal-poor and which could be among the oldest known stars in our Galaxy.  The RR0 Lyrae stars believed to be in the northernmost extension of the Sgr galaxy have a $\\rm <\\feh>$ $= -$1.53 which is on average more metal poor than the RR0 Lyrae stars from the bulge.  This value lies within the range of metallicity of \\citet{marconi98}, and agrees well with the metallicity found for the $V$-magnitude data of RR0 stars given by \\citet{mateo}.  Based on metallicities of RR0 Lyrae stars, we estimate their absolute magnitude and distances to all stars.  The distance to the Galactic Center is 8.0 $\\pm$ 0.3 kpc, where the error takes into account the estimated error in the MACHO photometry calibration as well as the zero-point in the RR0 Lyrae absolute magnitude relation. A correlation between metallicity and galactocentric distance is found. For metal-poor RR Lyrae stars the closer a star is to the Galactic center, on average, the more metal rich it is.  However for the metal-rich RR0 Lyrae stars, we find that this is not the case; instead, on average, the closer a star is to the Galactic center, the more metal poor it is. The result that the metal-rich and metal-poor RR0 Lyrae stars show a different trend in galactocentric distance needs to be addressed when modeling the chemical evolution of the Milky Way.  The Sgr RR0 Lyrae stars do not show a strong metallicity gradient. However, the metal-poor Sgr RR0 Lyrae stars do indicate a metallicity gradient of $\\sim$ +0.15 dex from $\\sim$ 2 kpc to 10 kpc of the galaxy.  (The further the RR0 Lyrae star from the center of Sgr, the more metal rich.)  This metallicity gradient is in the opposite direction as that found using giant stars by \\citet{ald01, marconi98, bellazzini99}."
    },
    "0809/0809.2470_arXiv.txt": {
        "abstract": "We used merger trees realizations, predicted by the extended Press-Schechter theory, in order to study  the growth of angular momentum of dark matter haloes. Our results showed that:\\\\ 1) The spin parameter  $\\lambda'$  resulting from the above method, is an increasing function of the present day mass of the halo.  The mean value of $\\lambda'$ varies from 0.0343 to 0.0484 for haloes with present day masses in the range of $ 10^9\\mathrm{h}^{-1}M _{\\odot}$ to $10^{14}\\mathrm{h}^{-1}M_{\\odot}$.\\\\ 2)The distribution of $\\lambda'$ is close to a  log-normal , but, as it is already found in the results of N-body simulations, the match is not satisfactory at the tails of the distribution. A new analytical formula that approximates the results much more satisfactorily is presented.\\\\ 3) The distribution of the values of $\\lambda'$ depends only weakly on the redshift. \\\\ 4) The spin parameter of an halo depends on the number of recent major mergers. Specifically the spin parameter is an increasing function of this number. ",
        "introduction": "There are two more likely pictures regarding the growth of angular momentum in dark matter  haloes.\\\\ The first is that galactic haloes acquired their angular momentum from the tidal torques of the surrounding matter. This is an old idea of Hoyle (1949) that has been investigated in a large number of studies (e.g. Peebles 1969; Doroshkevich 1970;  Efstathiou \\& Jones 1979, Barnes \\& Efstathiou 1987, White 1984; Voglis \\& Hiotelis 1989, Warren et al. 1992, Steinmetz \\& Bartelmann 1995). The results of the above studies, analytical and numerical, show that the spin parameter $\\lambda$, introduced by Peebles (1969) and defined by the relation $\\lambda\\equiv J \\sqrt{\\mid E\\mid}/GM^{5/2}$, has an average value of about 0.05-0.07, where $J$ is the modulus of the spin of halo, $M,E $ are the mass and the energy respectively and $G$ is the gravitational constant. According to the above picture, haloes acquire their angular momentum during their linear stage of their evolution because during this stage they have large linear sizes and thus the environment is capable to affect their evolution by tidal torques. Since their expansion is decelerating their relative linear size becomes smaller and the affection by the environment becoms less significant. The moment of the maximum expansion is in practice the end of the epoch of growth of angular momentum. At latter times, the halo evolves as a dynamically isolated system. Steinmetz \\& Bartelmann  (1995) showed that the dependence of the probability distribution of $\\lambda$ on the density parameter of the model Universe as well as on the variance of the density contrast field is very weak. Only a marginal tendency for $\\lambda$ is found to be larger for late-forming objects in an open Universe.\\\\ The second picture is closely related to the hierarchical clustering scenario of cold dark matter (CDM; Blumenthal et al. 1986). According to this scenario, structures in the Universe grow from small initially Gaussian density perturbations that progressively detach from the general expansion, reach a maximum radius and then collapse to form bound objects. Larger haloes are formed hierarchically by mergers between smaller ones, called progenitors. The buildup of angular momentum is a random walk process associated with the mass assembly history of the halo's major progenitor. The main role of tidal torques in this picture is to produce the random tangential velocities of merging progenitors.\\\\ The above two pictures of formation are usually studied by two different kinds of methods. The first kind is the N-body simulations that are able to follow the evolution of a large number of particles under the influence of the mutual gravity, from initial conditions to the present epoch. The second kind consists of semi-analytical methods. Among them, the Press-Schechter (PS) approach and its extensions (EPS) are of great interest since they allow to compute  mass functions (Press \\& Schechter  1974; Bond et al. 1991) to approximate merging histories (Lacey \\& Cole 1993,  Bower 1991, Sheth \\& Lemson 1999b) and to estimate the spatial clustering of dark matter haloes (Mo \\& White 1996; Catelan et al. 1998, Sheth \\& Lemson 1999a).\\\\ Recently large cosmological N-body simulations have been performed in order to study the angular momentum of dark matter haloes  in $\\Lambda CDM$ models of the Universe (e.g. Bullock 2001, Kasun \\& Evrard 2005, Bailin \\& Steinmetz 2005, Avila-Reese et al. 2005, Gottl\\\"{o}ber  \\& Turchaninov 2006).\\\\ Additionally, semi-analytical methods like merging histories resulting from EPS methods, have been used for similar studies (Vitvitska et. al 2002, Maller et. al 2002). \\\\ In this paper, we use such merging histories based on EPS approximations to study the distribution of spins.\\\\ The paper is organized as follows: In Sect.2, basic equations are summarized. In Sect.3, we present our results while a discussion is given in Sect.4. \\begin{figure}[t] \\includegraphics[width=8cm]{hiotelis-fig1.eps} \\caption{The figure shows  the conditions at the onset of the merger of two haloes with masses $m_1$ and $m_2$ and virial radii $R_1$ and $R_2$ respectively with $m_1 > m_2$ and $R_1 > R_2$. The centers of their masses are indicated by $CM_1$ and $CM_2$ and their distance is $r\\equiv max(R_1,R_2)=R_1$. The position vectors of their centers of mass are $\\mathbf{r_1}$ and $\\mathbf{r_2}$ respectively so their  relative position is $\\mathbf{CM_{12}}=\\mathbf{r_2}-\\mathbf{r_1}$. The vectors of the velocities of their center of mass are $\\mathbf{v_1}$ and $\\mathbf{v_2}$ and the vector of their relative velocity is $\\mathbf{v_{rel}}=\\mathbf{v_2}-\\mathbf{v_1}$. Haloes merge if they approach each other, that is the condition $\\mathbf{v_{rel}}\\cdot \\mathbf{CM_{12}} <0$ is fulfilled. See text for more details. }\\label{fg1-eps} \\end{figure} \\begin{figure}[t] \\includegraphics[width=8cm]{hiotelis-fig2.eps} \\caption{Spin parameter  $\\lambda'$ for the most massive progenitor at scale factor $\\alpha$ for the cases $1, 2$ and $3$. Solid, dashed and dotted lines correspond to the cases 1, 2 and 3 respectively. The evolution of $\\lambda'$ is characterized by sharp increases and decreases due to mergers.}\\label{fg2-eps} \\end{figure} ",
        "conclusions": "This study describes a picture for the growth of the angular momentum of dark matter haloes in terms of a hierarchical clustering scenario. The results presented above are, in general, in good agreement with the results that were already known in the literature. Comparing our results with those of large N-body simulations, we have found satisfactory agreement in the following points:\\\\ 1) The values of  spin parameter are in the range of  0.0343 to 0.0484 for haloes with present day masses in the range of $ 10^9\\mathrm{h}^{-1}M _{\\odot}$ to $10^{14}\\mathrm{h}^{-1}M_{\\odot}$.\\\\ 2) A log-normal distribution approximates satisfactorily the distributions of the values of the spin parameter but it fails to describe accurately the tails of the resulting distributions. A new, more satisfactory formula, is presented.\\\\ 3) The role of recent major mergers is very important. The distribution  of the spin parameter is appreciably affected by the number of recent major merger. The present day value of the spin parameter of a halo is an increasing function of the number of the recent major mergers.\\\\ 4) The distributions  of the spin parameter do not depend significantly on the redshift.\\\\ 5) The value of the spin parameter is a function of the present day mass of the halo. The form of this function depends, in N-body simulations, on the halo-finding algorithm but in general seems that spin parameter is a decreasing function of mass. Instead, in our results, $<\\lambda'>$ is an increasing function of mass, approximately very closely a power-law form.\\\\ Our results give rise to some  questions, as for example:  Why semi-analytical methods are not able to predict the correct relation between the spin parameter and the virial mass of the halo?  Does this disagreement reflects the null role of tidal fields in the orbital-merger picture or it arises from other problems associated with the nature of merger-trees?  Are merger-trees able to give the correct relation for better description of the velocity field during the merge? Is any way improvements on both analytical  and numerical methods are required in order to help us answering  some of the above questions and to advance our understanding about the physical processes that created the structures we observe."
    },
    "0809/0809.2193_arXiv.txt": {
        "abstract": "We present multi-colour photometry of the M8.5V ultracool dwarf ``pulsar'' TVLM 513-46546 (hereafter TVLM 513) obtained with the triple-beam photometer {\\sc ultracam}. Data were obtained simultaneously in the Sloan-$g'$ and Sloan-$i'$ bands. The previously reported sinusoidal variability, with a period of 2-hrs, is recovered here. However, the Sloan-$g'$ and Sloan-$i'$ lightcurves are anti-correlated, a fact which is incompatible with the currently proposed starspot explanation for the optical variability. The anti-correlated nature and relative amplitudes of the optical lightcurves are consistent with the effects of persistent dust clouds rotating on the surface of the star. In the absence of other plausible explanations for the optical variability of TVLM 513, it seems likely that dust cloud coverage combined with the rapid rotation of TVLM 513 is responsible for the optical variability in this object. However, crude modelling of a photosphere with partial dust cloud coverage shows that the anti-correlation can only be reproduced using cooler models than the literature temperature of TVLM 513. We suggest this discrepancy can be removed if more dust is present within the photosphere of TVLM 513 than theoretical model atmospheres predict, though a definitive statement on this matter will require the development of self-consistent models of partially dusty atmospheres. ",
        "introduction": "\\label{sec:introduction} \\alph{footnote} \\protect\\footnotetext[1]{Based on observations made at the European Southern Observatory, Paranal, Chile (ESO program 079.C-0686)} Brown dwarfs and very-low mass stars (collectively known as ultracool dwarfs, or UCDs) are strongly affected by the presence of dust in their photospheres. Dust absorbs elements from the gas phase, changing the opacity and metallicity of the atmosphere. Dust begins to form at temperatures corresponding to the transition between M and L spectral types, and becomes more prominent as the atmosphere cools; the presence of dust thus defines the L-dwarfs. Finally, when the dust clouds ``rain-out'' at lower temperatures still, they are responsible for the L-T transition. Thus understanding the formation, chemistry and atmospheric dynamics of dust is the central challenge facing theories of ultracool dwarf atmospheres \\citep{burrows06}.  The presence of dust is also thought to affect the magnetic properties of ultracool dwarfs. A combination of an increasingly neutral atmosphere, and frequent collisions between charged particles and dust grains in the dense atmosphere results in the atmospheres of ultracool dwarfs having a high electrical resistivity \\citep{mohanty02}. Thus, whilst strong quiescent and flaring radio emission reveals strong magnetic fields amongst the ultracool dwarfs \\citep[e.g][]{berger01,burgasser05,hallinan07}, the predominantly neutral atmospheres may explain the relatively low levels of other activity indicators, such as H$\\alpha$ emission and X-rays \\citep[e.g][]{gizis00,west04}.  TVLM 513 (M8.5V) is an ideal object in which to study the interplay of magnetic fields with a cool, dense atmosphere. It is close by \\citep[${\\rm d}=10.6$\\,pc; ][]{dahn02}, has known radio and H$\\alpha$ activity, and is rapidly rotating \\citep{mohanty03}. The radio emission of TVLM 513 is fascinating; observations from 1.4 to 8.5 GHz show the radio emission can switch states between highly polarised pulses with a 2-hr period \\citep{hallinan07} to a fainter, quiescent state interrupted by stochastic flares \\citep{berger07}. The periodic nature of the radio emission has led to this object being dubbed an ultracool dwarf ``pulsar''. Strong H$\\alpha$ emission, modulated on the same 2-hr period, and a detection of X-ray emission \\citep{berger07}, show that TVLM 513 is a magnetically active star, which is capable of supporting a chromosphere and a corona. This is despite its low effective temperature of $\\sim$2300\\,K \\citep{dahn02}, roughly the temperature at which dust is believed to start forming within the photosphere \\citep[see][and references within]{burrows06}.  Optical I-band photometry of TVLM 513 revealed the same 2-hr periodicity as seen in the radio and H$\\alpha$ emission; this was assumed to be due to starspots \\citep{lane07}. However, optical variability in ultracool dwarfs can also be caused by dust clouds within the photosphere \\citep[see][for a review]{bailer-jones02}. Simultaneous optical photometry in two or more bands can distinguish between variability caused by starspots \\citep{rockenfeller06a} and dust clouds \\citep{littlefair06}. We therefore obtained simultaneous, multi-colour photometry  using the fast photometer {\\sc ultracam} \\citep{dhillon07} to determine the cause of the optical variability in TVLM 513. The observations are described in Section~\\ref{sec:obs}, the results presented in Section~\\ref{sec:results} and discussed in Section~\\ref{sec:disc}, whilst in Section~\\ref{sec:conc} we draw our conclusions. ",
        "conclusions": "\\label{sec:conc}  We present simultaneous monitoring of the M8.5V star TVLM 513 in the Sloan-$g'$ and $i'$ bands. Both bands show sinusoidal variability on the $\\sim$2-hr rotation period. The $g'$ and $i'$ band lightcurves are in anti-phase, a fact which is incompatible with the optical variability being due to starspots. Without a plausible alternative explanation for this behaviour, we conclude that the optical variability in TVLM 513 is caused by the presence of a dust cloud in a predominantly dust-free photosphere. The optical variability is consistent with a dust cloud covering a significant fraction of the photosphere, but only if cooler stellar models than appropriate are used. This could be a consequence of the crudity of our modelling, or it may imply there is more dust in the atmosphere of TVLM 513 than the stellar atmosphere models of \\cite{allard01} predict. The $i'$-band variability reported here is a factor of four smaller than previously reported by \\cite{lane07}. Since such a high amplitude of variability is difficult to reconcile with dust cloud models, we suggest that the optical variability of TVLM 513-46546 can change in nature. It is unclear if this is related to the observed changes in radio emission from this object."
    },
    "0809/0809.2464_arXiv.txt": {
        "abstract": "We compile a large sample of broad absorption lines (BAL) quasars with X-ray observations from the \\xmm\\ archive data and Sloan Digital Sky Survey Data Release 5. The sample consists of 41 BAL QSOs. Among 26 BAL quasars detected in X-ray, spectral analysis is possible for twelve objects. X-ray absorption is detected in all of them. Complementary to that of \\citet{gall06} (thereafter G06), our sample spans wide ranges of both BALnicity Index (BI) and maximum outflow velocity (\\vmax\\ ). Combining our sample with G06's, we find very significant correlations between the intrinsic X-ray weakness with both BALnicity Index (BI) and the maximum velocity of absorption trough. We do not confirm the previous claimed correlation between absorption column density and broad absorption line parameters. We tentatively interpret this as that X-ray absorption is necessary to the production of the BAL outflow, but the properties of the outflow are largely determined by intrinsic SED of the quasars. ",
        "introduction": "About 10\\%-30\\% of optically selected QSOs show broad absorption lines (BAL) in their UV spectra, indicative of outflows with velocities up to 0.1c\\citep{hewett03,reichard03}. The similarity in the UV continuum and emission lines between BAL and non-BAL QSOs suggests that BAL QSOs are otherwise normal QSOs viewed in the direction covered by the outflow \\citep[e.g.][]{weymann91}. One exception to these similarities is that BAL QSOs are soft X-ray faint compared to non-BAL QSOs \\citep[e.g.][]{green95,brinkmann99}. The weakness in X-rays is interpreted as due to strong absorption rather than intrinsic difference. Evidence for this has been accumulated now from detailed studies of X-ray spectra of a few bright BAL quasars, which display X-ray absorption with column densities from $10^{22}$ to $\\geq10^{24}$\\cmsq\\ \\citep{wang99,gall99, gall02}.  Giving the ubiquity of X-ray absorption in BAL quasars, it is natural to ask whether and how the X-ray absorbing gas is connected to the UV BAL phenomenon. It has been known for quite long time that BAL gas should be either confined into small clumps or shielded from the intense soft X-rays in order to match the observed profile. \\citet{murray95} proposed that the highly ionized gas at the base of disk wind (shielding gas) can naturally filter the soft X-ray radiation to prevent the gas to be over-ionized so that the radiative acceleration is effective \\citep[See also][]{proga00}. As both UV and X-ray absorbers are part of the continuous outflow, the column densities of the two are expected to be correlated. Indeed, \\citet{brandt00} identified a correlation between the equivalent width of \\CIV\\ absorption line and the soft X-ray weakness in a sample of bright quasars, including half a dozen BAL QSOs.  \\citet{wang05,wang07} found that electron scattering of the shielding gas can explain the distribution of continuum polarization in quasars, and the resonant scattering of BAL outflow can explain the observed polarized spectrum of BAL. They further noted that certain special features should appear in the polarized spectrum if the size of the shielding gas is comparable with that of the BAL outflows. As these features are only found in several low-ionization BAL (LoBAL) QSOs  \\citep[see their paper for details and also][]{ogle99}, the shielding gas is likely well inside the BAL outflow except in LoBAL QSOs. A similar conclusion has been reached by studying the X-ray spectra of BAL QSOs \\citep[e.g.][]{gall04,gall06}. On the other hand, as pointed by \\citet{wang00}, in order to keep sufficient opacity in the soft X-ray between $0.2-0.3$~keV, the absorber must have large column density also of Li-like ions because those ions are responsible for both the soft X-ray absorption between $0.2-0.3$~keV and the high ionization UV BALs. Though this band is notoriously difficult to be studied, they argued that at least in three bright low redshift BAL QSOs, the X-ray absorption opacity around $0.2-0.3$ ~keV is large, suggesting very large column density of Li-like ions. However, a relatively small fraction of X-ray absorbing gas at moderate ionization level will be sufficient to suppress the soft X-ray flux.  If the X-ray shielding is critical to the ionization balance in the BAL outflow, which in turn affects the radiative accelerating force on the outflow, one would expect that kinematic properties and column density of BAL outflow will somehow correlate with the properties of the X-ray absorber. In a sample of BAL quasars observed by \\chandra\\ , G06 found a weak correlation between the maximum outflow velocity(\\vmax\\ ) of BAL and the indicator (\\daox\\ ) of X-ray absorption. Their finding agrees with the qualitative analysis that the strong soft X-ray absorption leads to more Li-like ions, thus more efficiently radiative acceleration by UV photons. However just as G06 pointed out that they had only four sources at the low \\vmax, and their sample of BAL QSOs is biased towards strongly absorbed sources comparing to the BI distribution of SDSS EDR BAL QSOs \\citep{reichard03}. Giving the importance of this question, more study based on a uniform sample is clearly required.  X-ray absorption is not the only factor that affects the ionization equilibrium of BAL gas. \\citet{steffen06} showed that the X-ray luminosity of non-BAL QSOs has a large scatter for a given optical luminosity. According to the current popular scenario that BAL and non-BAL are only a matter of whether our line of sight passes through BAL region or not, the intrinsic spectral energy distribution (SED) between UV and X-ray of BAL QSOs should be also diverse. Therefore, it would be interesting to study how the wind properties depend on the intrinsic SED because ionization equilibrium is also closely related to the intrinsic SED. If such a relation does exist, it may offer insight into the driver of the outflows.  In this paper, we present a study of BAL QSOs from SDSS Data Release 5 (DR5) \\citep{adelman07} that observed by \\xmm\\ satellite in X-ray in order to explore the relations between UV and X-ray absorbers as well as the relations between BAL properties and the intrinsic UV to X-ray spectra. In \\S2 we describe the selection of our \\CIV\\  BAL QSOs sample and the data analysis in \\S3. We show our results and discuss the underlying physics in \\S4. Finally, we summarize our results in \\S5. Throughout the paper, we assume the cosmological parameters $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\mathrm{M}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "\\label{sec_res}  \\subsection{X-ray Properties Of BAL QSOs}\\label{sec xrayprop}  To investigate the general X-ray properties of our SDSS/\\xmm\\ BAL sample (this paper), we try to measure the intrinsic absorption adopting a simple neutral absorption. However, only 12 of 41 sources can be fitted directly to give the intrinsic absorption column densities \\nh\\ in the range $\\sim 4\\times10^{21}$ to $\\sim 2\\times10^{23}$\\cmsq. For 14 sources, upper/lower limits can be placed by the hardness ratios HR. The final sample spans a wide range of intrinsic absorption column densities from $<10^{20}$\\cmsq\\ to $\\sim 10^{24}$\\cmsq\\ (Table 3). The lowest limit is obtained for the LoBAL QSO, SDSS J092238.43+512121.2. We have rechecked the optical spectrum, the identification of this quasar as LoBAL might be questionable because of the presence of narrow absorption lines. Notably, five LoBAL QSOs do not show stronger absorption than HiBAL QSOs.  Following G06,we measure \\aox\\ or place upper/lower limits on it, which ranges from $-1.36$ to $-2.26$ with an average $-1.86$ (Table 3). Similar to G06, we show the \\daox, to account for the luminosity dependence of \\aox(Table 3,see also the dot-dashed line in Fig. \\ref{fig_xdis}). The average value of \\daox\\ is $-0.25$, suggesting that 2~keV X-ray luminosities (at rest frame) of our SDSS/\\xmm\\ BAL sample are roughly three times fainter than the SDSS/\\rosat\\ non-BAL sample\\citep{stra05}. And in Table 3 we also present the \\aoxcorr\\ and \\daoxcorr\\ as a surrogate of the intrinsic X-ray properties of BAL QSOs. \\aoxcorr\\ is calculated by assuming $\\Gamma=2.0$ and using the hard-band counts rate to normalize the X-ray continuum and \\daoxcorr\\ =\\aoxcorr\\ $-$ \\aoxl\\ (see G06 or \\S3.1 for the definition). We find \\daoxcorr\\ in the range from $-0.36$ to $0.29$ with an average value of 0.11, which indicates our SDSS/\\xmm\\ BAL sample is slightly X-ray brighter, relative to the average quasars at that UV luminosity, than the SDSS/\\rosat\\ non-BAL sample\\citep[see Fig. \\ref{fig_xdis}]{stra05}.The X-ray brighter of our sample may be due to the relative shallower detection threshold of \\xmm\\ relative to \\chandra . Comparing the \\daoxcorr\\ distribution of the LBQS/\\chandra\\ BAL sample(G06) with the SDSS/\\rosat\\ non-BAL sample\\citep[figure 2 of G06]{stra05} one can find the LBQS/\\chandra\\ BAL sample(G06) is slightly intrinsic X-ray weaker with a median \\daoxcorr$=-0.14$ than the normal QSOs. Alternately, even hard X-rays are absorbed in the LBQS/\\chandra\\ BAL sample(G06) so that the simple assumption is broken down (See G06 or \\S3.1 for detail). We have carried out simulations to test this effect. Using \\xspec, we simulate the dependence of the \\daoxcorr\\ on the varying neutral hydrogen column density \\nh. We assume that \\daoxcorr\\ equals to zero for a single $\\Gamma=2$ power-law with \\nh$=10^{20}$\\cmsq\\ at the redshift 2. We find that an absorption column density of $3\\times10^{23}$\\cmsq\\ is required in order to account for the mean offset, about $-0.14$, of \\daoxcorr\\ of the LBQS/\\chandra\\ BAL sample(G06) relative to the SDSS/\\rosat\\ non-BAL sample\\citep{stra05}. If this is the main cause, most of X-ray weak sources in the LBQS/\\chandra\\ BAL sample(G06) will have a column density at least order of this. Future X-ray observation is certainly needed to assess this.   Either intrinsic X-ray weak or large column density of absorber may indicate that the LBQS/\\chandra\\ BAL sample(G06) is biased in X-ray properties. Their sample obviously has larger values of BI and \\vmax\\ than the SDSS BAL QSOs\\citep{reichard03} and is not uniform on UV properties too. Note that our SDSS/\\xmm\\ BAL sample is more uniform, especially on UV properties, and can be used as a complement to the LBQS/\\chandra\\ BAL sample(G06). For direct comparison, we show the distribution of \\daox\\ for the SDSS/\\rosat\\ non-BAL sample\\citep{stra05} and \\daoxcorr\\ for our SDSS/\\xmm\\ BAL sample(this paper) and the LBQS/\\chandra\\ BAL QSO sample(G06) in Fig. \\ref{fig_xdis}. In the following we used the combined sample of ours 41 and G06's 35 sources to study the relations between X-ray and UV properties. Note again, the UV properties of the LBQS/\\chandra\\ BAL QSO sample(G06) used in this paper are obtained by using our procedures so that we can use the consistent definition of BI and \\vmax. We also use the hardness ratios presented in G06 to calculate \\nh\\ of these G06 QSOs following the same approach as we have done for \\xmm\\ sources.  \\subsection{X-ray And UV Absorptions}\\label{sec_ab}  One of the purposes of this paper is to study the relationship between the UV and X-ray absorbers. It is generally believed that the X-ray absorber shields the disk winds from soft X-rays and makes line driving more efficient. A naive deduction is that the properties of UV and X-ray absorptions are correlated. Basing on this idea, G06 presented a correlation analysis between the X-ray absorption using \\daox\\ as an indicator and the UV absorption properties such as BI, DI, \\vmax\\ and $f_{\\rm deep}$. They found only a weak correlation between \\daox\\ and \\vmax. We will carry out a similar analysis using a larger sample covering more uniformly the whole BI range. In the following analysis, we will use Kendall-$\\tau$ test to quantify the significance of a correlation.  First, we check whether \\daox\\ is a good indicator of X-ray absorption. In the left panel of the Fig. \\ref{fig_nhaox}, we show \\nh\\ versus \\daox\\ for the combined subsample of 51 BAL QSOs that \\nh\\ is obtained either from spectral fit or from HR analysis. Similar to G06, we find a clear correlation between the two quantities. The probability for null hypothesis is less than 0.01\\%~ using non-parametric Kendall $\\tau$-test (See Table 4). Then we compare the redshift distributions of sources with \\daox\\ $> -0.2$ and \\daox\\ $< -0.2$.The result is their distributions are very similar,which indicates that the correlation between \\nh\\ and \\daox\\ is not from selection effect of redshift.These suggest that \\daox\\ can be used as a measure for X-ray absorption indeed.  Next, we examine the correlations between BAL properties and the X-ray absorption column density \\nh\\ (Fig. \\ref{fig_nhuv}) measured through X-ray spectral fit or HR analysis. We do not find any correlation with a probability of null hypothesis less than 1\\% (Table 4). However, a weak correlation between BI and \\nh\\ cannot be rejected because the large uncertainty in the \\nh\\ measurement may reduce the significance of a weak correlation to the measured level (2\\%).  We show \\daox\\ versus BI and \\daox\\ versus \\vmax\\ in Fig. \\ref{fig_daouv}. Two lensed BAL QSOs are excluded from following analysis because their UV and X-ray light may have been differntly amplified. There appears a correlation between \\daox\\ and BI with the probability for null hypothesis of only 0.05\\% (Table 4). The correlation appears not linear, rather there is an upper envelope. Since LoBAL QSOs may be different from the HiBAL QSOs \\citep{boroson92,wang07}, we also make Kendall test for 45 HiBAL QSOs only. The correlation is marginally significant with a null probability of 1\\%. The decrease in significance is caused by reducing the sample size. However, we do not find any significant correlation between \\vmax\\ and \\daox, which was seen in G06, in neither the whole sample nor in the HiBAL subsample with a null probability of 5\\% and 68\\%, respectively (Table 4). Comparison with the LBQS/\\chandra\\ BAL sample(G06), our sample has a handful BAL QSOs on the upper right of the figure. These BAL QSOs destroy the weak correlation trend of \\vmax\\ vs \\daox\\ in the LBQS/\\chandra\\ BAL sample(G06). These correlations are more or less similar to the correlations using \\nh\\ with an exception of higher significance. It is worthwhile to note that \\daox\\ reflects a combination of the X-ray absorption and the intrinsic deviation to the average quasar SED. Therefore, one must be careful as using it as an indicator of X-ray absorption. We will discuss below the implication of these results.   \\subsection{Intrinsic X-ray Properties And The Outflow}  Previous studies have shown that UV properties, such as the blueshift and the equivalent width of \\CIV\\ emission line, are correlated with X-ray to optical flux ratio for non-BAL QSOs (e.g. \\citet{baskin04,richards06}). It would be interesting to explore whether the BAL properties are correlated with the intrinsic \\aox. Unlike the correlation with X-ray absorption, such correlation should give information for the primary driver of the outflow. Here we use a corrected \\aox\\, i.e. \\aoxcorr\\ to represent the intrinsic \\aox\\, and investigate its relation with the UV absorption line properties. We also study the correlations between UV properties and \\daoxcorr\\ so that we can compare the results with those in previous section and G06. We note that the range of UV luminosity, $l_{2500}$, for the combined sample is only 2 dex, which introduce a scatter in \\aoxcorr\\, through \\aox$\\sim l_{2500}$ correlation, of less than 0.27, a factor of about 2.5 smaller than the dynamic range of \\aoxcorr. Therefore, the difference between \\daoxcorr\\ and \\aoxcorr\\ should be small in correlation analysis. This is verified below that the relationships between \\daoxcorr\\ and UV properties have a very similar behavior as those between \\aoxcorr\\ and UV properties.  Before exploring UV and X-ray connection, we first check whether \\aoxcorr\\ and \\daoxcorr\\ are affected by the X-ray absorption or not. We plot \\nh\\ versus \\daoxcorr\\ on the right panel of Fig. \\ref{fig_nhaox}. There is no apparent correlation between the two quantities. Kendall test gives a probability of chance coincidence of 36\\% (Table 4). Therefore, we can conclude that there is no evidence that \\aoxcorr\\ and \\daoxcorr\\ are significantly affected by X-ray absorption.  We then explore the correlations between \\aoxcorr\\ and BI or \\vmax\\ with Kendall and Spearman tests. We find that \\aoxcorr\\ is significantly correlated with both BI and \\vmax\\ with a Null probability of less than 0.1\\% for either test(See Fig 5; also Table 4). The correlation is still very significant ($P<0.1\\%$) for HiBAL QSO subsample (45 QSOs). For clarity we show only QSOs detected in the hard X-ray band in Fig.\\ref{fig_aocuvd}. Furthermore, these QSOs are more important for the Kendall and Spearman tests than the rest objects, and can give us a clear trend about these correlations. We note that one LoBAL QSOs, SDSS J133553.61+514744.1, which appears largely discrepant with the main sample in Fig. \\ref{fig_aocuv} and \\ref{fig_aocuvd}. This quasar has very steep X-ray photon index $\\Gamma\\sim2.56$ so that our `absorption-correction'  underestimate the intrinsic X-ray luminosity. If we set \\aox$=-1.82$ as the lower limit of \\aoxcorr (it is reasonable since \\aoxcorr\\ is larger than \\aox), this quasar would move rightward and is consistent with other quasars. The correlations between \\daoxcorr\\ and UV properties are very similar to above correlations using \\aoxcorr\\ except the latter appear slightly more significant (Table 4).  This may indicate that the dependence of UV properties on \\aoxcorr\\ is more fundamental than on \\daoxcorr.  Is it possible that these correlations are introduced by some selection effect in the sample? If it is the case this effect would tend to miss the objects which occupy the bottom-left (X-ray weak and high UV absorption) and top-right (X-ray strong and low UV absorption) corners of Fig. \\ref{fig_aocuv} and \\ref{fig_aocuvd}. Note that the sample selections of ours and G06's are both based on the optical luminosity and have nothing to do with the X-ray properties. If the relative X-ray luminosity is uncorrelated with BAL properties it is hard to understand why G06 and we select the objects at the bottom-right(top-left) corner but miss those at the bottom-left(top-right) corner. Since objects at top-right (bottom-left) corners, if they really exist, should have the same optical properties as these at top-left (bottom-right). All of these analysis indicate the correlations between the intrinsic \\aox\\ and the properties of UV absorber are real. We discuss the implications of the correlations in details in the next subsection.   \\subsection{Discussion}\\label{sec_dis}  Using a larger and more uniform sample, we reexamine the correlations between BAL properties and X-ray absorption presented in G06. Two indicators of X-ray absorption, \\nh\\ and \\daox, are used in the work. We identify the correlation between \\daox\\ and $BI$ as the only significant one. In particular, we do not find the correlation between \\daox\\ and \\vmax\\ claimed in G06, and any correlation between \\nh\\ and the UV absorption properties. Although we can not rule out a weak correlation between \\nh\\ and the UV absorption line properties due to relative large error bar of \\nh\\, our results clearly suggest that X-ray absorption is {\\em not} the major factor that determines UV absorption properties. Given the fact that almost all BAL QSOs show strong absorption in X-ray, it seems that X-ray absorption is a necessary condition for launching of the BAL winds, but the properties of the wind depend on other factors. As shown above, the observed correlation between \\daox\\ and BI may be the secondary effect of the correlation between BI and \\aoxcorr\\ or \\daoxcorr, as \\daox\\ is composed of the contributions of absorption and of \\daoxcorr.  In passing, we note that lack of correlations between the X-ray absorption column density and UV properties does not necessarily contradict with the scenario of radiatively accelerated wind as naively thought. For locally optically thin material, the ratio of the radiation force to the gravitational force is a function of Eddington ratio and the cross-section ratio of effective absorption to Thomson scattering. If resonant scattering is responsible for the absorption opacity, the cross-section will be determined by the fraction Li-like ions. According to the equatorial wind model\\citep{murray95}, a clump of highly ionized gas (shielding gas), which accounts for most X-ray opacity, blocks the soft X-ray interior to the wind. The transmitted flux of soft X-rays that ionize Li-like ions in the wind depends strongly on the X-ray column density, \\nh. If \\nh\\ along the direction is very small, high-velocity wind cannot be launched because of the reduction of the radiation force caused by the over-ionization. On the other hand, if \\nh\\ is very large, the wind will end up with a turbulent flow due to the blocking of thick very-low-ionized gas behind\\citep[see the figure 4 of ][]{proga00}. Thus, high-velocity wind can only be launched when the radial column density \\nh\\ is moderate as the fraction of Li-like ions, such as \\CIV\\ and \\NV, is large enough. As far as the X-ray absorption column density is in the right range, the fraction of Li-like ions should be the dominant species. The flow properties are then determined self-consistently by the launching radius, the gas density at the launching radius and the radiation intensity. If X-ray absorber is well separated from the UV absorber, then we would not expect any correlation between the X-ray absorption column density and the flow properties for BAL QSOs apart.  On the other hand, the X-ray absorber may be the 'hitchhiking' gas just located at the inner edge of the wind, and its properties may have a close connection with the boundary conditions of the disk wind\\citep{murray95}. In that case, we should consider globally the structure of gas along a line of sight. The wind starts at a radius where the radiation force is substantially larger than the gravitational force. As far as the gas density is high enough, such a region can certainly exist. \\citet{murray95} has worked out a consistent line acceleration model, and they found that gas column density and final velocity are correlated for a constant Eddington ratio and at a given launch radius. However, if the launching radius is not exactly scaled with luminosity as $L^{1/2}$ and there are a range of Eddington ratio, as they assumed, the correlation can be smeared out.  More interesting results of our work are the strong correlations between the parameters of outflow and intrinsic \\aox. We argue that these correlations are essential rather than due to some selection effect or the secondary effect of other correlations (see previous subsection for details). In fact our results are consistent with \\citet{richards06} who found that the QSOs with large blueshifts of \\CIV\\ emission line, i.e. the parent population of BAL QSOs as suggested by \\citet{richards06}, tend to have lower X-ray luminosity for given optical luminosity (their figure 5). It is also upheld by \\citet{laor02} who presented significant correlation between the equivalent width of \\CIV\\ absorption and \\aox\\ in a sample of non-BAL QSOs. This correlation is actually predicted by \\citet{murray95}, in which they found that quasars with a large X-ray to UV ratio can only produce weak low velocity winds while quasars with a small X-ray to UV ratio can produce strong and large velocity winds. This is exactly what we have found here. As we discussed above, their model also predicted a correlation between the X-ray absorption column density and the maximum velocity of the flow, which is not observed in this sample. Lack of such correlation may be due to two important factors that (1) variation in the Eddington ratio and launching radius; (2) the large uncertainties in the measurement of absorption column density.  We do not fully understand why the variation in the Eddington ratio and launching does not completely smeared out the correlation with \\aoxcorr. There seems one reason for this. \\citet{wang04} found that the 2-10kev luminosities to bolometric luminosities ratio tightly anti-correlated with Eddington ratio for a sample of broad-line and narrow-line Seyfert 1 AGNs. If this correlation holds up for BAL QSOs, one would expect that quasars with high Eddington ratio would have larger radiative acceleration force, or large terminal velocity, and at the same time X-ray weaker. \\citet{ganguly07} find that \\vmax\\ as a function of Eddington ratio has an upper envelope in the SDSS3 BAL catalog, exactly as expected. Since there is no clear correlation of BI with UV luminosities\\citep[cf.][]{laor02}, Eddington rate and black hole mass\\citep{ganguly07},it is very likely that the BAL properties are more likely determined by the SED of quasars rather than Eddington ratio."
    },
    "0809/0809.4242_arXiv.txt": {
        "abstract": "The first metal enrichment in the universe was made by supernova (SN) explosions of population (Pop) III stars. The trace remains in abundance patterns of extremely metal-poor (EMP) stars. We investigate the properties of nucleosynthesis in Pop III SNe by means of comparing their yields with the abundance patterns of the EMP stars. We focus on (1) jet-induced SNe with various energy deposition rates [$\\Ed=(0.3-1500)\\times10^{51}$\\ergs], and (2) SNe of stars with various main-sequence masses ($\\Mms=13-50\\Msun$) and explosion energies [$E=(1-40)\\times10^{51}$ergs]. The varieties of Pop III SNe can explain varieties of the EMP stars: (1) higher [C/Fe] for lower [Fe/H] and (2) trends of abundance ratios [X/Fe] against [Fe/H]. ",
        "introduction": "\\label{sec:intro}  Long-duration $\\gamma$-ray bursts (GRBs) have been found to be accompanied by luminous and energetic Type Ic supernovae [SNe Ic, called hypernovae (HNe)] (\\eg \\cite[Galama \\etal\\ 1998]{gal98}). Although the explosion mechanism is still under debate, photometric observations (a ``jet break'', \\eg \\cite[Frail \\etal\\ 2001]{fra01}) and spectroscopic observations (a nebular spectrum, \\eg \\cite[Maeda \\etal\\ 2002]{mae02}) indicate that they are aspherical explosions with jet(s).  The aspherical explosions are also indirectly suggested from the abundance patterns of extremely metal-poor (EMP) stars with [Fe/H] $<-3$.\\footnote{Here [A/B] $\\equiv\\log_{10}(N_{\\rm A}/N_{\\rm B})-\\log_{10}(N_{\\rm A}/N_{\\rm B})_\\odot$, where the subscript $\\odot$ refers to the solar value and $N_{\\rm A}$ and $N_{\\rm B}$ are the abundances of elements A and B, respectively.} The EMP stars are suggested to show nucleosynthesis yields of a single core-collapse SN (\\eg \\cite[Beers \\& Christlieb 2005]{bee05}). Particularly, C-enhanced EMP (CEMP) stars have been well explained by the faint SNe with large fallback (\\cite[Umeda \\& Nomoto 2005; Iwamoto \\etal\\ 2005; Nomoto \\etal\\ 2006; Tominaga \\etal\\ 2007b]{ume05,iwa05,nom06,tom07b}). On the other hand, some CEMP stars show enhancement of Co and Zn (\\eg \\cite[Depagne \\etal\\ 2002]{dep02}) that requires explosive nucleosynthesis under high entropy. In a {\\sl spherical} model, however, a high entropy explosion is equivalent to a high energy explosion that inevitably synthesizes a large amount of \\Nifs, \\ie leads a bright SN (\\eg \\cite[Woosley \\& Weaver 1995]{woo95}). This incompatibility will be solved if a faint SN is associated with a narrow jet within which a high entropy region is confined (\\cite[Umeda \\& Nomoto 2005]{ume05}). ",
        "conclusions": "We focus on two interesting properties observed in the abundance patterns of the metal-poor stars: (1) the higher [C/Fe] for lower [Fe/H] and (2) the trends of [X/Fe] against [Fe/H]. The variations of the metal-poor stars are explained by the variations of SNe that contribute the metal enrichment of the early universe. Especially, (1) the variation of the energy deposition rates explains the tendency of [C/Fe] against [Fe/H] and (2) the variations of $\\Mms$ and $E$ explain the trends of [X/Fe] against [Fe/H]. We propose that the abundance patterns of the metal-poor stars will provide additional constraints on the explosion mechanism of GRBs and SNe other than the direct observations of present GRBs and SNe.  \\vspace{.5cm}  \\noindent This work has been supported in part by WPI Initiative, MEXT, Japan."
    },
    "0809/0809.1651_arXiv.txt": {
        "abstract": "Existing exoplanet radial velocity surveys are complete in the planetary mass-semimajor axis ($M_p-a$) plane over the range 0.1 AU $< a <$ 2.0 AU where $M_p \\gtrsim 100~M_{\\oplus}$.  We marginalize over mass in this complete domain of parameter space and demonstrate that the observed $a$ distribution is inconsistent with models of planet formation that use the full Type I migration rate derived from a linear theory and that do not include the effect of the ice line on the disk surface density profile.  However, the efficiency of Type I migration can be suppressed by both nonlinear feedback and the barriers introduced by local maxima in the disk pressure distribution, and we confirm that the synthesized $M_p-a$ distribution is compatible with the observed data if we account for both retention of protoplanetary embryos near the ice line and an order-of-magnitude reduction in the efficiency of Type I migration. The validity of these assumption can be checked because they also predict a population of short-period rocky planets with a range of masses comparable to that of the Earth as well as a ``desert\" in the $M_p-a$ distribution centered around $M_p \\sim 30-50~M_{\\oplus}$ and $a <1$ AU. We show that the expected ``desert\" in the $M_p-a$ plane  will be discernible by a radial velocity survey with 1 m s$^{-1}$ precision and $n \\sim 700$ radial velocity observations of program stars. ",
        "introduction": "Over 200 planets with reliable mass ($M_p$) and semimajor axis ($a$) measurements have been discovered around nearby FGK stars in the past decade. At the same time, attempts to build a comprehensive deterministic theory of planet formation have lead to the development of population synthesis models based on the sequential accretion scenario. In \\citet{ida04}, two of us studied the growth of planetesimals into dynamically isolated embryos as well as their tidal interactions with their parent disks.  Using the observed ranges of disk mass, size, and accretion rate we showed that a fraction of embryos evolve into cores with more than a few Earth masses ($M_\\oplus$), accrete massive envelopes, open up gaps near their orbits, and attain asymptotic masses comparable to that of Jupiter. In some massive and persistent disks, the newly formed gas giant planets may migrate toward the proximity of their host stars.  In the end these simulations produced the distribution of dynamical and structural properties of planets.  Presently, the observed sample of extrasolar planetary properties has become large enough to enable direct comparisons between the theoretically predicted and observed $M_p-a$ distributions that not only delineate the dominant physical mechanisms at work in planet formation, but also provide quantitative constraints on the efficiencies of those processes.  In the latest update of the planet formation models two of us have incorporated the effect of Type I migration \\citep{ida08a}. This process is a direct consequence of a protoplanetary core's tidal interaction with its parent protoplanetary disk.  The efficiency of Type I migration was first determined by a linear theory \\citep{gol80,war86,tan02} that neglected the embryo's perturbation on the surface density distribution of its parent protoplanetary disk. In the environment of a minimum mass nebula though, this efficiency factor would imply that a protoplanetary embryo with mass a fraction of an Earth mass would migrate from $\\sim 1$ AU into its host star before the severe depletion of the disk gas that is known to occur over a time scale of several Myr.  Although this critical mass at which Type I migration causes AU-scale migration on the disk depletion time increases with $a$, it is still difficult to retain sufficiently massive cores for the onset of dynamical accretion of gas.  This argument implies that gas giants should be very rare \\citep{ida08a} and this paradox has led to many in-depth analyses of the Type I migration process.  Numerical nonlinear simulations of Type I migration were reviewed by \\citet{pap06} and many potential explanations for slow Type I migration are in the literature: intrinsic turbulence in the disk \\citep{lau04,nel04}, self-induced unstable flow \\citep{kol04,li05}, nonlinear radiative and hydrodynamic feedback \\citep{mas06a}, and variation in surface density and temperature gradients \\citep{mas06b}.  \\citet{dob07} has shown that in some situations the nonlinear Type I migration rate can be less than 10\\% of the linear prediction.  \\defcitealias{ida08b}{IL} Another issue studied by two of us \\citep[IL hereafter]{ida08b} is the critical embryo mass $M_{crit} >$ at least a few $M_{\\oplus}$ required by current models for runaway gas accretion.  The embryo is limited by its isolation mass $M_{iso}$, and the solid surface density profile of the minimum mass solar nebula (MMSN) $\\Sigma_d \\propto a^{-3/2}$ requires that the $M_{iso}$ scales like $a^{3/4}$.  On the other hand, the timescale for growth $\\tau_{c,acc}$ scales with $a^{27/10}$.  As a result, after the characteristic gas depletion time $\\tau_{dep}$ the most massive embryos near the ice line have masses $M_c \\sim M_{iso} < M_{crit}$. However, since we observe many exoplanets with Jupiter masses at $a \\sim 1$ AU, there must be some physical process that is neglected in this simple analysis.  \\citet{kre07} outlined one possible solution to this problem in which solids are trapped near regions of the disk where the local pressure requires the gas to rotate with super-Keplarian velocities.  When the combined contribution of both Lindblad and corotation resonances are taken into account \\citep{mas06b}, the migration of protoplanetary cores may be suppressed as well \\citepalias{ida08b}.  In this paper, we utilize the observed data to calibrate the population synthesis models.  In \\S2 we quantitatively show that the existing distribution of exoplanets cannot be explained by models of planet formation that apply the full Type I migration rate predicted from linear theory and that do not include the effects of the ice line.  We also point out that the existing observed and synthesized $M_p-a$ distributions are in agreement with each other if we take into account the effect of an ice line barrier and assume a reduction in the magnitude of Type I migration.  In addition, we describe the parameters of a radial velocity survey capable of verifying the existence of ``desert\" in the $M_p-a$ diagram predicted by \\citetalias{ida08b}.  In \\S3 we consider the implications of these models and suggest methods to test our assumptions.  In \\S4 we summarize our findings. ",
        "conclusions": "We used the fact that existing exoplanet radial velocity surveys are complete in the planetary mass-semimajor axis ($M_p-a$) plane where 0.1 AU $< a <$ 2.0 AU and $M_p$ is in the range specified by Equation~(\\ref{eq2}) to show that the observed semimajor axis distribution in the complete region cannot be explained by models of planet formation that use the full Type I migration rate predicted by linear theory and that do not include the effects of the ice line.  Moreover, we also demonstrated that the expected ``desert\" in the $M_p-a$ plane at about $M_p \\sim 30~M_{\\oplus}$ and $a <1$ AU predicted by \\citetalias{ida08b} will be discernible by a radial velocity survey with 1 m s$^{-1}$ precision and $n \\sim 700$ radial velocity observations of program stars.  Such an observational campaign will also verify the predicted inner boundary of the ``desert\" where we expect a large population of super-Earths have migrated to and halted in the proximity of their host stars."
    },
    "0809/0809.4597_arXiv.txt": {
        "abstract": "{ We report the discovery of a sub-Jupiter mass exoplanet transiting a magnitude $\\mathrm{V}=11.7$ host star 1SWASP~J030928.54+304024.7. A simultaneous fit to the transit photometry and radial-velocity measurements yield a planet mass $\\rmsub{M}{p}=0.53\\pm 0.07\\,\\mathrm{M_J}$, radius $\\rmsub{R}{p}=0.91^{+0.06}_{-0.03}\\,\\mathrm{R_J}$ and an orbital period of $3.722465^{+0.000006}_{-0.000008}$\\,days. The host star is of spectral type K3V, with a spectral analysis yielding an effective temperature of $4800\\pm 100\\,\\mathrm{K}$ and $\\log g=4.45\\pm 0.2$. It is amongst the smallest, least massive and lowest luminosity stars known to harbour a transiting exoplanet. WASP-11b is the third least strongly irradiated transiting exoplanet discovered to date, experiencing an incident flux $\\rmsub{F}{p}=1.9\\times10^{8}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$ and having an equilibrium temperature $\\rmsub{T}{eql}=960\\pm 70$\\,K.  } ",
        "introduction": "Observations of planets that transit their host star represent the current best opportunity to test models of the internal structure of exoplanets and of their formation and evolution. Since the first detection of an exoplanetary transit signature \\citep{charbonneau00, henry00} over fifty transiting planetary systems have been identified. A number of wide-field surveys are in progress with the goal of detecting transiting exoplanets, for example OGLE \\citep{udalski02}, XO \\citep{mccullough05}, HAT \\citep{bakos04}, TrES \\citep{odonovan06} and WASP \\citep{pollacco06}.  The WASP project operates two identical instruments, at La Palma in the Northern hemisphere, and at Sutherland in South Africa in the Southern hemisphere. Each telescope has a field of view of just under 500 square degrees. The WASP survey is sensitive to planetary transit signatures in the light-curves of hosts in the magnitude range V\\,$\\sim$9--13. A detailed description of the telescope hardware, observing strategy and pipeline data analysis is given in \\citet{pollacco06}.  In this paper we report the discovery of WASP-11b, a sub-Jupiter mass gas giant planet in orbit about the host star 1SWASP~J030928.54+304024.7. We present the WASP discovery photometry plus higher precision optical follow-up and radial velocity measurements which taken together confirm the planetary nature of WASP-11b. ",
        "conclusions": "\\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{wasp11_evol_baraffe.ps}} \\resizebox{\\hsize}{!}{\\includegraphics{wasp11_evol_girardi.ps}} \\caption[]{The position of WASP-11 in the $R/M^{1/3}-\\rmsub{T}{eff}$ plane. Evolutionary tracks for a solar metallicity star from \\citet{baraffe98} (upper panel) and \\citet{girardi00} (lower panel) are plotted along with isochrones for ages 10\\,Myr (solid), 1\\,Gyr (dashed), 5\\,Gyr (dot-dashed), 10\\,Gyr (dotted). Evolutionary mass tracks are shown for 0.7, 0.8, 0.9 and 1.0\\,M$_\\odot$.} \\label{starevol} \\end{center} \\end{figure}  The system parameters derived here place WASP-11b towards the lower end of the mass range of known transiting planets, falling approximately mid-way between the masses of Jupiter and Saturn. The host star WASP-11 is also amongst the smallest and lowest luminosity stars known to host a transiting planet, however it is relatively nearby and thus quite bright ($\\mathrm{V}=11.7$). WASP-11b is irradiated by a stellar flux $\\rmsub{F}{p}=1.9\\times 10^8$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ at the sub-stellar point making it the third least heavily irradiated transiting planet after GJ436b and HD17156b. We compute an equilibrium temperature for WASP-11b of $\\rmsub{T}{eql}(A=0; f=1)=960\\pm70$\\,K, which makes it more typical of the bulk of known exoplanets than of the ``hot Jupiter'' class most commonly found by the transit method.  Theoretical models of the atmospheres of hot giant exoplanets \\citep{fortney06, burrows07} have shown that heavy irradiation can lead to the development of a temperature inversion and a hot stratosphere.  This is due to the absorption of stellar flux by an atmospheric absorber, possibly TiO and VO. In both sets of models the magnitude of the incident stellar flux is the key controlling variable determining whether a given extra-solar giant planet (EGP) will possess a hot stratosphere. Recent observations by \\citet{machalek08} of secondary transits of XO-1b using the {\\it Spitzer Space Telescope} suggest the presence of a temperature inversion in the atmosphere of that exoplanet.  On the other hand analogous observations of HD189733b \\citep{charbonneau08} show no evidence for an inversion, despite the irradiating fluxes of XO-1b and HD189733b being almost identical ($\\rmsub{F}{p}=0.49\\times 10^9$ and $\\rmsub{F}{p}=0.47\\times 10^9$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ respectively). This strongly suggests that the incident stellar flux is not the sole controlling parameter determining the presence of the inversion, a likelihood which the authors of the atmosphere models readily point out themselves. Further observations of planets particularly in the low-irradiation regime are required to help parameterise the thermal inversion. WASP-11b is amongst the nearest and brightest low-irradiation EGPs making it a good candidate for such studies. Moreover we note that the orbital eccentricity of WASP-11b is much lower than the other two bright low-irradiation transiting exoplanets, GJ436b and HD17156b ($e=0.15$ and $e=0.67$ respectively). As a consequence the secular variation in irradiation around the orbit will be correspondingly lower in WASP-11b, removing a potentially complicating factor when comparing follow-up observations with predictions from atmospheric models developed assuming steady-state irradiation.  To estimate the age of the WASP-11 we compared the observed stellar density and temperature against the evolutionary models of low- and intermediate-mass stars of \\citet{girardi00} and \\citet{baraffe98}. In Figure~\\ref{starevol} we plot the position of WASP-11 in the $R/M^{1/3}$ versus $\\rmsub{T}{eff}$ plane atop isochrones of different ages from the two models. For such a cool star, the isochrones are closely spaced in this parameter plane due to the slow post-main-sequence evolution of late-type stars. The sets of isochrones from the two models overlap in this regime, and both models suggest the same mass and age for the host star. WASP-11 falls above the 10\\,Gyr isochrone for both models, though it is consistent with this age within the errors. The very low lithium abundance also points toward WASP-11 being $\\ga 1$--2\\,Gyr old \\citep{sestito05}. We investigated using gyrochronology to age the host star, following \\citet{barnes07}, however we were unable to measure a definite rotational period. No rotation modulation was detected in the lightcurve to an amplitude limit of a few milli-magnitudes. The spectral analysis furnishes only an upper-limit to $v\\sin i$, so no rotational period can be determined in that way. Taken together these factors are all consistent with WASP-11 being an old star, older than maybe 1\\,Gyr, however it is not possible to be more definite than that with the available data.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{fort_mr.ps}} \\caption[]{Planetary mass-radius relations as a function of core mass and system age, interpolated from the models of \\protect\\citet{fortney07}.} \\label{coremass} \\end{center} \\end{figure}  \\citet{fortney07} present models of the evolution of planetary radius over a range of planetary masses and orbital distances, and under the assumption of the presence of a dense core of various masses up to 100\\,M$_\\oplus$.  To compare our results with the Fortney et al. models we plotted the modelled mass-radius relation as a function of core mass in Figure~\\ref{coremass}. To account for the lower-than-Solar luminosity of the host star WASP-11 we calculated the orbital distance $a_\\odot=a(M_\\star/M_\\odot)^{-3.5/2}$ at which a planet in orbit about the Sun would receive the same incident stellar flux as WASP-11b does from its host. We then interpolated the models of Fortney et al. to this effective orbital distance ($a_\\odot=0.068$ for WASP-11b). As the age of the WASP-11 system is poorly constrained we compare our results with the modelled mass-radius relation at 300\\,Myr, 1\\,Gyr and 4.5\\,Gyr. We find that the radius of WASP-11b is consistent with the presence of a dense core with a mass in the range $\\rmsub{M}{core}\\sim$ 42--77\\,M$_{\\oplus}$ for a system age of 300\\,Myr, $\\rmsub{M}{core}\\sim$33--67\\,M$_{\\oplus}$ at 1\\,Gyr, and $\\rmsub{M}{core}\\sim$22--56\\,M$_{\\oplus}$ at 4.5\\,Gyr."
    },
    "0809/0809.3462_arXiv.txt": {
        "abstract": "While there is plentiful evidence in all fronts of experimental cosmology for the existence of a non-vanishing dark energy (DE) density $\\rD$ in the Universe, we are still far away from having a fundamental understanding of its ultimate nature and of its current value, not even of the puzzling fact that $\\rD$ is so close to the matter energy density $\\rM$ at the present time (i.e. the so-called ``cosmic coincidence'' problem). The resolution of some of these cosmic conundrums suggests that the DE must have some (mild) dynamical behavior at the present time. In this paper, we examine some general properties of the simultaneous set of matter and DE perturbations $(\\delta\\rho_M, \\delta\\rho_D)$ for a multicomponent DE fluid. Next we put these properties to the test within the context of a non-trivial model of dynamical DE (the $\\CC$XCDM model) which has been previously studied in the literature. By requiring that the coupled system of perturbation equations for $\\delta\\rM$ and $\\delta\\rD$ has a smooth solution throughout the entire cosmological evolution, that the matter power spectrum is consistent with the data on structure formation and that the ``coincidence ratio'' $r=\\rD/\\rM$ stays bounded and not unnaturally high, we are able to determine a well-defined region of the parameter space where the model can solve the cosmic coincidence problem in full compatibility with all known cosmological data. ",
        "introduction": "\\label{Introduction}  Undoubtedly the most prominent accomplishment of modern cosmology to date has been to provide strong indirect support for the existence of both dark matter (DM) and dark energy (DE) from independent data sets derived from the observation of distant supernovae\\,\\cite{Supernovae}, the anisotropies of the CMB\\,\\cite{WMAP3}, the lensing effects on the propagation of light through weak gravitational fields\\,\\cite{Lensing}, and the inventory of cosmic matter from the large scale structures (LSS) of the Universe\\,\\cite{Cole05,Lahav02}. But, in spite of these outstanding achievements, modern cosmology still fails to understand the ultimate physical nature of the components that build up the mysterious dark side of the Universe, most conspicuously the DE component of which the first significant experimental evidence was reported $10$ years ago from supernovae observations. The current estimates of the DE energy density yield $\\rD^{\\rm exp}\\simeq (2.4\\times 10^{-3}\\,eV)^4$ and it is believed that it constitutes roughly $70\\%$ of the total energy density budget for an essentially flat Universe. The big question now is: what is it from the point of view of fundamental physics? One possibility is that it is the ground state energy density associated to the quantum field theory (QFT) vacuum and, in this case, it is traditional to associate $\\rD$ to $\\rL=\\CC/8\\pi\\,G$, where $\\CC$ is the cosmological constant (CC) term in Einstein's equations. The problem, however, is that the typical value of the (renormalized) vacuum energy in all known realistic QFT's is much bigger than the experimental value. For example, the energy density associated to the Higgs potential of the Standard Model (SM) of electroweak interactions is more than fifty orders of magnitude larger than the measured value of $\\rD$.  Another generic proposal (with many ramifications) is the possibility that the DE stands for the current value of the energy density of some slowly evolving, homogeneous and isotropic scalar field (or collection of them). Scalar fields appeared first as dynamical adjustment mechanisms for the CC\\,\\cite{Dolgov,PSW} and later gave rise to the notion of quintessence\\,\\cite{quintessence}. While this idea has its own merits (specially concerning the dynamical character that confers to the DE) it has also its own drawbacks. The most obvious one (often completely ignored) is that the vacuum energy of the SM is still there and, therefore, the quintessence field just adds up more trouble to the whole fine-tuning CC problem\\,\\cite{weinRMP,CCRev}!  Next-to-leading is the ``cosmological coincidence problem'', or the problem of understanding why the presently measured value of the DE is so close to the matter density. One expects that this problem can be alleviated by assuming that $\\rD$ is actually a dynamical quantity. While quintessence is the traditionally explored option, in this paper we entertain the possibility that such dynamics could be the result of the so-called cosmological ``constants'' (like $\\CC$, $G$,...) being actually variable. It has been proven in \\,\\cite{SS12} that this possibility can perfectly mimic quintessence. It means that we stay with the $\\CC$ parameter and make it ``running'', for example through quantum effects\\,\\cite{cosm,JHEPCC1,Babic}\\,\\footnote{ For a general discussion, see \\,\\cite{SSIRGAC06,ShS08}.}. However, in \\,\\cite{LXCDM12} it was shown that, in order to have an impact on the coincidence problem, the total DE in this context should be conceived as a composite fluid made out of a running $\\CC$ and another entity $X$, with some effective equation of state (EOS) parameter $\\wX$, such that the total DE density and pressure read $\\rD=\\rL+\\rX$ and $\\pD=-\\rL+\\wX\\rX$, respectively. We call this system the $\\CC$XCDM model\\,\\cite{LXCDM12}. Let us emphasize that $X$ (called ``the cosmon'') is not necessarily a fundamental entity; in particular, it need not be an elementary scalar field. As remarked in \\,\\cite{LXCDM12}, $X$ could represent the effective behavior of higher order terms in the effective action (including non-local ones). This is conceivable, since the Bianchi identity enforces a relation between all dynamical components that enter the effective structure of the energy-momentum tensor in Einstein's equations, in particular between the evolving $\\CC$ and other terms that could emerge after we embed General Relativity in a more general framework\\,\\cite{Gruni,fossil}. Therefore, at this level, we do not impose a microscopic description for $X$ and in this way the treatment becomes more general\\,\\footnote{See e.g. \\cite{Ratra08ab,Percival07,Zhang08a} for recent constraints on $\\CC$CDM, XCDM and quintessence-like models. The margin for the energy densities and EOS parameter $\\wX$ is still quite high. In the $\\CC$XCDM case, the fact that $\\CC$ is running provides an even wider range of phenomenological possibilities.}. The only condition defining $X$ is the DE conservation law, namely we assume that $\\rD=\\rL+\\rX$ is the covariantly self-conserved total DE density.   In this paper, we analyze the combined dynamics of DE and matter density perturbations for such conserved DE density $\\rD$. The present study goes beyond the approximate treatment presented in \\cite{GOPS}, where we neglected the DE perturbations and estimated the matter perturbations of the $\\CC$XCDM model using an effective (variable) EOS $\\we$ for the composite fluid $(\\rL,\\rX)$. The main result was that a sizeable portion of parameter space was still compatible with a possible solution of the cosmic coincidence problem. The ``effective approach'' that we employed in \\cite{GOPS} was based on three essential ingredients: i) the use of the effective EOS representation of cosmologies with variable cosmological parameters\\,\\cite{SS12}; ii) the calculation of the growth of matter density fluctuations using the effective EOS of the DE\\,\\cite{LinderJenkins03}; and iii) the application of the, so-called, ``F-test'' to compare the model with the LSS data, i.e. the condition that the linear bias parameter, $b^2(z)= P_{GG}/P_{MM}$ does not deviate from the $\\CC$CDM value by more than $10\\%$ at $z=0$, where $P_{MM}\\propto (\\delta\\rho_M/\\rho_M)^2$ is the matter power spectrum and $P_{GG}$ is the galaxy fluctuation power spectrum\\,\\cite{Cole05,Lahav02} -- see \\cite{GOPS,Ftest} for details. This three-step methodology turned out to be an efficient streamlined strategy to further constrain the region of the original parameter space\\,\\cite{LXCDM12}. However, there remained to perform a full fledged analysis of the system of cosmological perturbations in which the DE and matter fluctuations are coupled in a dynamical way. This kind of analysis is presented here.   The structure of the paper is as follows. In the next section, we outline the meaning of the cosmic coincidence problem within the general setting of the cosmological constant problem. In section \\ref{sect:perturbations}, the basic equations for cosmological perturbations of a multicomponent fluid in the linear regime are introduced. In section \\ref{sect:effEOS}, we describe the general framework for addressing cosmological perturbations of a composite DE fluid with an effective equation of state (EOS). In section \\ref{sect:genericfeatures}, we describe some generic features of the cosmological perturbations for the dark energy component. The particular setup of the $\\CC$XCDM model is focused in section \\ref{sect:LXCDM}. In sections \\ref{sect:perturbLXCDM} and \\ref{numerical}, we put the $\\CC$XCDM model to the stringent test of cosmological perturbations and show that the corresponding region of parameter space becomes further reduced. Most important, in this region the model is compatible with all known observational data and, therefore, the $\\CC$XCDM proposal can be finally presented as a robust candidate model for solving the cosmic coincidence problem. In section \\ref{DEfluct}, we offer a deeper insight into the correlation of matter and DE perturbations. In the last section, we present the final discussion and deliver our conclusions.    \\noindent ",
        "conclusions": "\\label{sect:conclusions} In this paper, we have addressed the impact of the cosmological perturbations on the coincidence problem. In contrast to the previous study\\,\\cite{GOPS}, where this problem was examined in a simplified ``effective approach'' in which the dark energy (DE) perturbations were neglected, in the present work we have taken them into account in a full-fledged manner. We find that the results of the previous analysis were reasonable because the DE perturbations generally tend to smooth at scales below the sound horizon. However, the inclusion of the DE perturbations proved extremely useful to pin down the physical region of the parameter space and also to put the effective approach within a much larger perspective and to set its limitations.  First of all, we have performed a thorough discussion on the coupled set of matter and DE perturbations for a general multicomponent fluid. This has prepared the ground to treat models in which the DE is a composite medium with a variable equation of state (EOS). We have concentrated on those cases in which the DE, despite its composite nature, is described by a self-conserved density $\\rD$. Notice that if matter is covariantly conserved, the covariant conservation of the DE is mandatory. In particular, this is the situation for the standard $\\CC$CDM model, although in this case the self-conservation of the DE appears through a trivial cosmological constant term, $\\rL=\\rL^0$, which remains imperturbable throughout the entire history of the Universe. One may nevertheless entertain generalized frameworks where the DE is not only self-conserved, but is non-trivial and dynamical. This is not a mere academic exercise; for instance, in quantum field theory in curved space-time we generally expect that the vacuum energy should be a running quantity\\,\\cite{JHEPCC1,SSIRGAC06,ShS08}. Therefore, in such cases, the CC density becomes an effective parameter that may evolve typically with the expansion rate, $\\rL=\\rL(H)$, and constitutes a part of the full (dynamical) DE of the composite cosmological system with variable EOS. In these circumstances, if the gravitational coupling $G$ is constant, the running CC density $\\rL=\\rL(H)$ cannot be covariantly conserved unless other terms in the effective action of this system compensate for the CC variation. We have called the effective entity that produces such compensation ``$X$'' or ``cosmon'', and denoted with $\\rX$ its energy density. Therefore, $\\rD=\\rL+\\rX$ is the self-conserved total DE density in this context, which must be dealt with together with the ordinary density of matter $\\rM$. A generic model of this kind is what we have called the $\\CC$XCDM model\\,\\cite{LXCDM12,GOPS}.  Furthermore, from general considerations  based on the covariance of the effective action of QFT in curved space-time\\,\\cite{JHEPCC1,SSIRGAC06,ShS08}, we expect that the running CC density $\\rL=\\rL(H)$ should be an affine quadratic law of the expansion rate $H$, see Eq.\\,(\\ref{runlamb}). Using this guiding principle and the ansatz of self-conservation of the DE, we find that the evolution of $\\rX$, and hence of $\\rD$, becomes completely determined, even though its ultimate nature remains unknown. In particular, $X$ is \\textit{not} a scalar field in general.  The $\\CC$XCDM model was first studied in \\cite{LXCDM12} as a promising solution to the cosmic coincidence problem, in the sense that the coincidence ratio $r=\\rD/\\rM$ can stay relatively constant, meaning that it does not vary in more than one order of magnitude for many Hubble times. The main aim of the present paper was to make a further step to consolidate such possible solution of the coincidence problem, specifically from the analysis of the coupled system of matter and DE perturbations. Let us remark that this has been a rather non-trivial test for the $\\CC$XCDM model. Indeed, after intersecting the region where the DE perturbations of this model can be consistently defined, with the region where the coincidence problem can be solved\\,\\cite{LXCDM12,GOPS}, we end up with a significantly more reduced domain of parameter space where the model can exist in full compatibility with all known cosmological data. The main conclusion of this study is that the predictivity of the model has substantially increased. Therefore, it can be better put to the test in the next generation of precision cosmological observations, which include the promising DES, SNAP and PLANCK projects\\,\\cite{SNAP}.  Interestingly enough, we have found that the final region of the parameter space is a naturalness region which is more accessible to the aforementioned precision experiments. For example, we have obtained the bound $0\\leqslant\\nu\\lesssim \\nu_0\\sim 10^{-2}$ for the parameter that determines the running of the cosmological term. This bound is perfectly compatible with the physical interpretation of $\\nu$ from its definition (\\ref{nu}). Moreover, our analysis indicates that the cosmon entity $X$ behaves as ``phantom matter''\\cite{LXCDM12}, i.e. it satisfies $\\wX<-1$ with negative energy density. This result is a clear symptom (actually an expected one) from its effective nature. It is also a welcome feature; let us recall\\,\\cite{LXCDM12} that ``phantom matter'', in contrast to the ``standard'' phantom energy, prevents the Universe from reaching the Big Rip singularity. Finally, perhaps the most noticeable (and experimentally accessible) feature that we have uncovered from the analysis of the DE perturbations in the $\\CC$XCDM model, is that the overall EOS parameter $\\we$ associated to the total DE density $\\rD$ behaves effectively as quintessence ($\\we\\gtrsim -1$) in precisely the region of parameter space where the cosmic coincidence problem can be solved. In other words, quintessence is mimicked by the $\\CC$XCDM model in that relevant region, despite that there is no fundamental quintessence field in the present framework. A detailed confrontation of the various predictions of the $\\CC$XCDM model (in particular, the kind of dependence $\\we=\\we(z)$) with the future accurate experimental data\\,\\cite{SNAP}, may eventually reveal these features and even allow to distinguish this model from alternative DE proposals based on fundamental quintessence fields.  To summarize: we have demonstrated that the set of cosmological models characterized by a composite, and covariantly conserved, DE density $\\rD$ in which the vacuum energy $\\rL$ is a dynamical component (specifically, one that evolves quadratically with the expansion rate, see Eq.\\,(\\ref{runlamb})), proves to be a distinguished class of models that may provide a consistent explanation of why $\\rD$ is near $\\rM$, in full compatibility with the theory of cosmological perturbations and the rest of the cosmological data. Remarkably, such class of models is suggested by the above mentioned renormalization group approach to cosmology. We conclude that the $\\CC$XCDM model can be looked upon as a rather predictive framework that may offer a robust, and theoretically motivated, dynamical solution to the cosmic coincidence problem.   \\vspace{1cm}    {\\bf Acknowledgments.}\\ The authors are grateful to Julio Fabris for useful discussions. We have been supported in part by MEC and FEDER under project FPA2007-66665 and also by DURSI Generalitat de Catalunya under project 2005SGR00564. The work of J.G. is also financed by MEC under BES-2005-7803. We acknowledge the support from the Spanish Consolider-Ingenio 2010 program CPAN CSD2007-00042.        \\newcommand{\\JHEP}[3]{ {JHEP} {#1} (#2)  {#3}} \\newcommand{\\NPB}[3]{{\\sl Nucl. Phys. } {\\bf B#1} (#2)  {#3}} \\newcommand{\\NPPS}[3]{{\\sl Nucl. Phys. Proc. Supp. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRD}[3]{{\\sl Phys. Rev. } {\\bf D#1} (#2)   {#3}} \\newcommand{\\PLB}[3]{{\\sl Phys. Lett. } {\\bf B#1} (#2)  {#3}} \\newcommand{\\EPJ}[3]{{\\sl Eur. Phys. J } {\\bf C#1} (#2)  {#3}} \\newcommand{\\PR}[3]{{\\sl Phys. Rep. } {\\bf #1} (#2)  {#3}} \\newcommand{\\RMP}[3]{{\\sl Rev. Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\IJMP}[3]{{\\sl Int. J. of Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRL}[3]{{\\sl Phys. Rev. Lett. } {\\bf #1} (#2) {#3}} \\newcommand{\\ZFP}[3]{{\\sl Zeitsch. f. Physik } {\\bf C#1} (#2)  {#3}} \\newcommand{\\MPLA}[3]{{\\sl Mod. Phys. Lett. } {\\bf A#1} (#2) {#3}} \\newcommand{\\CQG}[3]{{\\sl Class. Quant. Grav. } {\\bf #1} (#2) {#3}} \\newcommand{\\JCAP}[3]{{ JCAP} {\\bf#1} (#2)  {#3}} \\newcommand{\\APJ}[3]{{\\sl Astrophys. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AMJ}[3]{{\\sl Astronom. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\APP}[3]{{\\sl Astropart. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AAP}[3]{{\\sl Astron. Astrophys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\MNRAS}[3]{{\\sl Mon. Not. Roy. Astron. Soc.} {\\bf #1} (#2)  {#3}} \\newcommand{\\JPA}[3]{{\\sl J. Phys. A: Math. Theor.} {\\bf #1} (#2)  {#3}} \\newcommand{\\ProgS}[3]{{\\sl Prog. Theor. Phys. Supp.} {\\bf #1} (#2)  {#3}} \\newcommand{\\APJS}[3]{{\\sl Astrophys. J. Suppl.} {\\bf #1} (#2)  {#3}}   \\newcommand{\\Prog}[3]{{\\sl Prog. Theor. Phys.} {\\bf #1}  (#2) {#3}} \\newcommand{\\IJMPA}[3]{{\\sl Int. J. of Mod. Phys. A} {\\bf #1}  {(#2)} {#3}} \\newcommand{\\IJMPD}[3]{{\\sl Int. J. of Mod. Phys. D} {\\bf #1}  {(#2)} {#3}} \\newcommand{\\GRG}[3]{{\\sl Gen. Rel. Grav.} {\\bf #1}  {(#2)} {#3}}"
    },
    "0809/0809.0307_arXiv.txt": {
        "abstract": "{To obtain an accurate description of broad-band photometric star cluster evolution, certain effects should be accounted for. {{Energy equipartition}} leads to {{mass segregation and}} the preferential loss of low-mass stars, while stellar remnants severely influence cluster mass-to-light ratios. Moreover, the stellar initial mass function and cluster metallicity affect photometry as well. Due to the continuous production of stellar remnants, their impact on cluster photometry is strongest for old systems like globular clusters. This, in combination with their low metallicities, evidence for mass segregation, and a possibly deviating stellar initial mass function, makes globular clusters interesting test cases for cluster models.} {In this paper we describe cluster models that include the effects of {{the preferential loss of low-mass stars}}, stellar remnants, choice of initial mass function and metallicity. The photometric evolution of clusters is predicted, and the results are applied to Galactic globular clusters.} {The cluster models presented in this paper represent an analytical description of the evolution of the underlying stellar mass function due to stellar evolution and dynamical cluster dissolution. Stellar remnants are included by using initial-remnant mass relations, while cluster photometry is computed from the Padova 1999 isochrones.} {Our study shows that the preferential loss of low-mass stars strongly affects cluster magnitude, colour and mass-to-light ratio evolution, as it increases cluster magnitudes and strongly decreases mass-to-light ratios. The effects of stellar remnants are prominent in the evolution of cluster mass, magnitude and mass-to-light ratio, while variations of the initial mass function induce similar, but smaller changes. Metallicity plays an important role for cluster magnitude, colour and mass-to-light ratio evolution. The different effects can be clearly separated with our models. We apply the models to the Galactic globular cluster population. Its properties like the magnitude, colour and mass-to-light ratio ranges are well reproduced with our models, provided that {{the preferential loss of low-mass stars}} and stellar remnants are included. We also show that the mass-to-light ratios of clusters of similar ages and metallicities cannot be assumed to be constant for all cluster luminosities. Instead, mass-to-light ratio increases with cluster luminosity and mass.} {These models underline the importance of more detailed cluster models when considering cluster photometry. By including the preferential loss of low-mass stars and the presence of stellar remnants, the magnitude, colour and mass-to-light ratio ranges of modelled globular clusters are significantly altered. With the analytic framework provided in this paper, observed cluster properties can be interpreted in a more complete perspective.} ",
        "introduction": "\\label{sec:intro} In recent studies, the photometric evolution of star clusters has been extensively treated from various approaches \\citep[e.g.,][]{andersfritze03,lamers06,vonhippel06,fagiolini07}. Because cluster photometry is used for a broad range of applications, like age-dating galaxies and tracking their formation history, it is crucial to obtain an accurate description of the photometric evolution of clusters. While {\\it Simple Stellar Population} (SSP) models \\citep[e.g.,][]{leitherer99,bruzual03,andersfritze03,maraston05} only consider the changing photometric properties due to stellar evolution, other models that also use the dynamical input of $N$-body simulations can predict the photometric evolution of clusters under a wider variety of conditions \\citep[e.g.,][]{lamers06,fagiolini07,borch07}. In reality, not only stellar evolution but also the dynamical interaction of a cluster with its environment causes it to lose stars \\citep[e.g.,][]{baumgardt03}. This process, which is called dissolution, occurs due to internal two-body relaxation and external effects like tidal perturbation, spiral arm passages or encouters with Giant Molecular Clouds \\citep[e.g.,][]{baumgardt03,gieles06,gieles07}. It can change the shape of the stellar mass function, and will also affect photometric cluster evolution. This is the case as a cluster {evolves towards energy equipartition}, causing it to preferentially lose low-mass stars \\citep[e.g.,][]{portegieszwart01,baumgardt03,hurley05}.  {The physical driving force of the preferential loss of low-mass bodies is subject to debate. On the one hand, energy equipartition between the bodies constituting a cluster increases the velocities of low-mass objects and thereby gives rise to the preferential loss of low-mass stars. On the other hand, it has been proposed that mass segregation, a phenomenon in which due to energy equipartition the more massive stars sink towards the cluster centre and low-mass objects move outwards, leads to the same effect since bodies in the cluster outskirts are more loosely bound than objects in the cluster centre and are thus more easily lost \\citep[e.g.,][]{leon00,portegieszwart01,lamers06}. This line of reasoning is not compatible with \\citet{king66}, where it is shown that the escape rate of stars from a cluster does not vary with radius. In that scenario, the preferential loss of low-mass stars and mass segregation are both the effects of energy equipartition, but do not necessarily share any causal connection \\citep[e.g.,][]{delafuentemarcos00}. Regardless of its specific nature, in the remainder of this paper we consider energy equipartition to be the fundamental cause of the preferential loss of low-mass stars: either directly, via mass segregation, or a combination of the two. In any case, mass segregation can serve as an indicator for clusters that have undergone a strong preferential loss of low-mass stars \\citep[e.g.,][]{portegieszwart01,baumgardt03} and will therefore be used in that respect.}  There have been many observations of Galactic open and globular clusters in which evidence of mass segregation was found \\citep[e.g.,][]{anderson96,hillenbrand98,zoccali98,albrow02,richer04,koch04,pasquali04}. These clusters can be expected to exhibit non-canonical photometric evolution. Furthermore, for a number of clusters overall mass-to-light ratios are observed that strongly deviate from the mean value presented in \\citet{mclaughlin00} \\citep[e.g.,][]{baumgardt03b,vandeven06}, which suggests a range of scenarios for photometric cluster evolution. This, in combination with the high mass-to-light ratio {\\it in the centre} of some globular clusters \\citep[e.g.,][]{pasquali04,vandenbosch06} and the consequent invocation of intermediate mass black holes (IMBHs) \\citep[e.g.,][]{portegieszwart02,gurkan04}, asks for a cluster model that can explain the observed range of mass-to-light ratios. While some studies suggest IMBHs to explain the high mass-to-light ratio in the centres of globular clusters \\citep[e.g.,][]{gebhardt05,noyola06}, others show that these are not required and central concentrations of stellar remnants also provide a solution \\citep[e.g.,][]{baumgardt03a,baumgardt03b,hurley07}. Therefore, it is important to investigate to what extent either model can be used to explain the observed range of mass-to-light ratios, colours and magnitudes. A model describing cluster mass-to-light ratios may also be able to provide insight in the connection between globular clusters and ultra-compact dwarf galaxies (UCDs), the latter having a different mass-to-light ratio range than globular clusters \\citep[e.g.,][]{hasegan05,evstigneeva07,rejkuba07,mieske08a}.  In this paper we present cluster models that are based on stellar isochrones like all SSP models, but analytically incorporate dynamical effects on cluster photometry by following the results from $N$-body simulations \\citep{baumgardt03}. The resulting speed and applicability to a large parameter space makes it very suitable for studying the effect of a range of parameters on cluster evolution. We show how photometric properties of clusters like their magnitude, colour and overall mass-to-light ratio are affected by {the preferential loss of low-mass stars}, the inclusion of stellar remnants, the stellar initial mass function (IMF) and metallicity.  The structure of the paper is as follows. Our cluster evolution models are presented in Sect.~\\ref{sec:clevo}. In Sect.~\\ref{sec:evo} it is shown how stellar evolution affects the cluster content, including the production of stellar remnants. We derive the equations describing dynamical effects of cluster evolution on a multi-component powerlaw IMF \\citep[e.g.,][]{kroupa01} in Sect.~\\ref{sec:diss}. The effects of {the preferential loss of low-mass stars} and the dynamical loss of remnants are included. In that section, we also provide the final set of equations to describe cluster evolution with our models. The computation of photometric evolution is treated in Sect.~\\ref{sec:phot}. The results are presented in Sect.~\\ref{sec:results}, where also the influences of {the preferential loss of low-mass stars}, stellar remnants, IMF and metallicity are investigated, and the results are compared to previous studies. In Sect~\\ref{sec:appl}, the models are applied to Galactic globular clusters. Section~\\ref{sec:disc} contains a discussion of the results, while our conclusions are provided in Sect.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} We have treated the influence of {the preferential loss of low-mass stars}, stellar IMF, metallicity and the inclusion of stellar remnants on cluster mass, magnitude, colour and mass-to-light ratio evolution. We presented analytical models that describe the evolution of cluster content and photometry, based on stellar evolution from the Padova 1999 isochrones and on simplified dynamical dissolution models as first presented in \\citet{lamers05}. The latter, in turn, is based on the $N$-body simulations by \\citet{baumgardt03}.  {The models represent the cluster evolution part of our new cluster population synthesis code {\\it SPACE}.} We considered Kroupa and Salpeter IMFs and metallicities in the range $Z=0.0004$---0.05. The obtained data are publicly available in electronic form at the CDS and also at \\texttt{http://www.astro.uu.nl/\\~{}kruijs}. The results from our models are as follows. \\begin{itemize} \\item[(1)] {\\it The preferential loss of low-mass stars} slightly decreases the total disruption time of a cluster by a few percent. However, the most significant changes are effected in cluster photometry. The effect of fading is decreased as clusters {including mass loss in the preferential mode} can stay more than 1.5 $V$-band magnitudes brighter than clusters losing mass in the canonical mode, because most of the dynamical mass loss occurs in the form of low-mass stars that contribute little to cluster luminosity. Initially, clusters {exhibiting the preferential loss of low-mass stars} are bluer than standard ones, but they become redder during the last $\\sim 10\\%$ of cluster lifetime. The cluster mass-to-light ratio is severely decreased due to {the preferential loss of low-mass stars}. The decrease typically ranges from 2---4$~\\msun~{\\rm L}_\\odot^{-1}$ (i.e., up to 0.6~dex) near total cluster disruption {for total disruption times $\\tdis>12$~Gyr}. {If the upturn of the $M/L_V$ evolution that is much more prominent in \\citet{baumgardt03} than in our models (see Sect.~\\ref{sec:assump}, point (5)) is accounted for, this range of $M/L_V$ decrease is at most 0.5$~\\msun~{\\rm L}_\\odot^{-1}$ smaller.} \\item[(2)] {Including the mass of {\\it stellar remnants} obviously yields an increase} in the total cluster mass and consequently also in total disruption time with respect to cluster evolution without remnants. The extended lifespan also implies that cluster luminosity less rapidly decreases. The mass-to-light ratio is enhanced by almost 2 $~\\msun~{\\rm L}_\\odot^{-1}$ at its maximum, close to total disruption. \\item[(3)] We compared the evolution of clusters with {\\it Salpeter and Kroupa IMFs}, which can be considered to {favour high stellar masses (Kroupa, or a `top-oriented' IMF) or low masses (Salpeter, or a `bottom-oriented' IMF) alternatives due to the bend in the Kroupa IMF and a slight slope difference}. As can be expected, clusters with a bottom-{oriented} IMF retain more mass due to stellar evolution, which eventually causes these clusters to become brighter than clusters with a top-{oriented} IMF. However, they start out being slightly fainter since a top-{oriented} IMF favours massive stars and is thus brighter than a bottom-{oriented} one. Similarly, clusters with a top-{oriented} IMF are bluer and have smaller mass-to-light ratios than clusters with a bottom-{oriented} IMF. {For the Kroupa and Salpeter IMFs}, the latter change can amount up to several $~\\msun~{\\rm L}_\\odot^{-1}$. \\item[(4)] {\\it Metallicity} variations hardly influence the total mass evolution of clusters. In accordance with stellar studies \\citep{hurley04} low-metallicity clusters are brighter and also much bluer than high-metallicity ones. Consequently, the mass-to-light ratio is a strongly increasing function of metallicity. \\item[(5)] When applying our results to {\\it Galactic globular clusters}, it is evident that {the preferential loss of low-mass stars} is required to explain their low observed {\\it mass-to-light ratios}, especially if stellar remnants are accounted for. Low metallicity is insufficient to serve as an explanation. Another important implication of our study is that the mass-to-light ratio can {\\it not} be assumed to be constant over varying luminosity, as it is strongly affected by the dynamical history of clusters. \\item[(6)] The fact that clusters of high masses may not have reached energy equipartition yet suggests that the effects of {the preferential loss of low-mass stars} disappear with increasing cluster mass. Because clusters {exhibiting the preferential loss of low-mass stars} have much lower mass-to-light ratios than clusters that lose their mass in the canonical mode, clusters with high masses would then have much higher mass-to-light ratios than ones with lower masses. This effect may have been found by \\citet{rejkuba07}. The above interpretation and its application to the observations of \\citet{rejkuba07} is treated more extensively in \\citet{kruijssen08b}. \\item[(7)] The typical {\\it colour range of globular clusters} is covered by our models. When considering the colour-metallicity relation as reported by \\citet{smith07}, from an order-of-magnitude comparison we suggest that the observed colour scatter at fixed metallicity could be the effect of {the preferential loss of low-mass stars}. \\item[(8)] {Only when adopting a Salpeter IMF down to $\\mmini=0.1$~\\msun, the mass-to-light ratios of UCDs are reproduced by our models. While UCDs could represent a natural continuation of the trend of increasing mass-to-light ratio with (globular) cluster mass \\citep{rejkuba07}, this is not expected to be of a dynamical nature, since more massive UCDs are not expected to have reached energy equipartition within a Hubble time.} \\end{itemize}  The retain of remnants and the existence of {the preferential loss of low-mass stars} are found in $N$-body simulations of clusters and in observations, while metallicity and IMF variations are observed among real clusters. Therefore the effects described in this paper should be considered when observing clusters, and observed cluster properties have to be interpreted very carefully\\footnote{Predictions for specific models can be made by the first author upon request.}."
    },
    "0809/0809.0594_arXiv.txt": {
        "abstract": "We observed the southwestern region of the Cygnus Loop in two pointings with \\textit{XMM-Newton}. The region observed is called the ``blow-out'' region that is extended further in the south. The origin of the ``blow-out'' is not well understood while it is suggested that there is another supernova remnant here in radio observation. To investigate the detail structure of this region in X-ray, we divided our fields of view into 33 box regions. The spectra are well fitted by a two-component nonequilibrium ionization model. The emission measure distributions of heavy elements decrease from the inner region to the outer region of the Loop. Then, we also divided our fields of view into 26 annular sectors to examine the radial plasma structure. Judging from metal abundances obtained, it is consistent with that the X-ray emission is the Cygnus Loop origin and we concluded that high-$kT_{e}$ component ($\\sim$0.4\\,keV) originates from the ejecta while low-$kT_{e}$ component ($\\sim$0.2\\,keV) is derived from the swept-up interstellar medium. The flux of low-$kT_{e}$ component is much less than that of high-$kT_{e}$ component, suggesting the ISM component is very thin. Also, the relative abundances in the ejecta component shows similar values to those obtained  from previous observations of the Cygnus Loop. We find no evidence in X-ray that the nature of the ``blow-out'' region originated from the extra supernova remnant. From the ejecta component, we calculated the masses for various metals and estimated the origin of the Cygnus Loop as the core-collapse explosion rather than the Type Ia supernova. ",
        "introduction": "\\label{sec:intro} The Cygnus Loop is one of the brightest Supernova Remnant (SNR) in the X-ray sky. Its age is estimated to be $\\sim$\\,10,000 yrs \\cite{Blair05}. Since the distance is comparatively close to us (540\\,pc; Blair et al. 2005), the apparent size is quite large ($2.5^\\circ\\times3.5^\\circ$; Levenson et al. 1997), which enables us to study the plasma structure of the Loop in detail.  Although the \\object{Cygnus Loop} is an evolved SNR, a hot plasma is still confined inside the Loop \\cite{Hatsukade90}. Miyata et al. (1998) observed the Loop with the \\textit{Advanced Satellite for Cosmology and Astrophysics} (ASCA), and detected the strong  highly-ionized Si-K, S-K line and Fe-L line near the center of the Cygnus Loop. They concluded that the hot plasma, a ``fossil'' of the supernova explosion, left in the core of the Loop. Tsunemi et al. (2007) (hereafter TKNM07) observed the Cygnus Loop along the diameter from the northeast (NE) to the southwest (SW) with \\textit{XMM-Newton} and studied the radial plasma structure. From the spectral analysis, they showed that the Cygnus Loop consists of two component plasma. They concluded that the low-$kT_e$ component originating from the interstellar medium (ISM) surrounds the high-$kT_e$ component originating from the ejecta. In addition, they measured the metal abundances of the high-$kT_e$ component and showed the metal distribution of the ejecta. The results indicate that the abundances are relatively high ($\\sim$5 times solar) and each element is nonuniformly distributed: Si, S and Fe are concentrated in the inner region while the other elements such as O, Ne and Mg are abundant in the outer region. They also estimated the progenitor star's mass to be 15 M$_\\odot$.  The Cygnus Loop is a typical shell-like SNR; this structure is thought to be generated by the cavity explosion \\cite{Levenson97}. The Cygnus Loop is almost circular in shape, however, we can see some breakout in the SW. It is called the ``blow-out'' region \\cite{Aschenbach99}. The origin of the ``blow-out'' is not well understood. Aschenbach \\& Leahy (1999) have explained this extended structure as a breakout into a lower density ISM. On the other hand, Uyaniker et al. (2002) suggested the existence of a secondary SNR (named G72.9-9.0) in the south from a radio observation and some other radio observations support this conclusion (Uyaniker et al. 2004; Sun et al. 2006).  Our observations were performed in a direction from the Cygnus Loop center toward the south ``blow-out'' region. In this paper, we report the result of the spectral analysis and discuss about the plasma structure of this region. ",
        "conclusions": "We observed the SW region of the Cygnus Loop with \\textit{XMM-Newton}. To examine the plasma structure, we divided our FOV in two different ways: 33 box sectors and 26 annular sectors. We fitted the spectrum extracted from each region with two-$kT_{e}$ VNEI model. The plasma structure of the low-$kT_{e}$ component and that of the high-$kT_{e}$ component are quite different from each other: each temperature is $\\sim$0.2\\,keV and $\\sim$0.4\\,keV for the former and the latter, respectively.  The EM distribution of the low-$kT_{e}$ component suggest the rim brightening structure, while that of the high-$kT_{e}$ component monotonously decreases from the center of the Loop to the outside.  In the high-$kT_{e}$ component, the abundances of Si and Fe are relatively high compared to those of Ne and Mg. The distributions of EMs as well as the relative abundances in the high-$kT_{e}$ component match the view that the low- and high-$kT_{e}$ components, respectively, originate from the ISM and the ejecta of the Cygnus Loop, which was derived by earlier observations such as TKNM07 or Katsuda et al. (2008b). We found that the emission from this ISM component is relatively weak. This suggests that the thickness of the shell is thin in Pos-8 and 9. We also calculated the relative abundances of Ne, Mg, Si, and Fe to O in the ejecta component for the entire FOV, and estimated the origin of the Cygnus Loop as the core-collapse explosion rather than the Type Ia supernova. We found no evidence in X-ray that the nature of the ``blow-out'' region originated from the extra SNR."
    },
    "0809/0809.2591_arXiv.txt": {
        "abstract": "We present infrared photometry of SN~1999em, plus optical photometry, infrared photometry, and optical spectroscopy of SN~2003hn.  Both objects were Type II-P supernovae.  The $V-[RIJHK]$ color curves of these supernovae evolved in a very similar fashion until the end of plateau phase. This allows us to determine how much more extinction the light of SN~2003hn suffered compared to SN~1999em.  Since we have an estimate of the total extinction suffered by SN~1999em from model fits of ground-based and space-based spectra as well as photometry of SN~1999em, we can estimate the total extinction and absolute magnitudes of SN~2003hn with reasonable accuracy.  Since the host galaxy of SN~2003hn also produced the Type Ia SN~2001el, we can directly compare the absolute magnitudes of these two SNe of different types. ",
        "introduction": "Supernovae (SNe) come in two basic models: exploding white dwarfs that are members of close binary systems, and single massive stars that develop iron cores.  Type Ia SNe are generally regarded to be exploding carbon-oxygen white dwarf stars that have reached the Chandrasekhar limit due to mass transfer from a nearby non-degenerate donor star \\citep[][and references therein]{Liv00}. Therefore, Type Ia SNe are explosions constrained by a uniform energy budget.  This leads to a high degree of uniformity of their light curves.  The objects with more rapidly declining light curves are fainter at maximum light, and the slow decliners are brighter \\citep{Phi93, Ham_etal96, Rie_etal96, Per_etal97, Phi_etal99, Jha_etal07}. In the near-infrared, however, Type Ia SNe are even better.  They are standard candles \\citep{Kri_etal04a, Kri_etal04b, Woo_etal07}. Except for the fastest declining objects, there are essentially no ``decline rate relations'' in the IR.  As we have shown in a recent series of papers \\citep[see][and references therein]{Kri_etal07,Wan_etal08}, a combination of optical and near-IR photometry allows the accurate determination of host galaxy extinction of Type Ia SNe, even if the dust is different than ``standard'' Galactic dust with R$_V$ = 3.1.  Type II SNe are thought to be single stars born with 8 M$_{\\odot}$ or more which, after a few million years of evolution, undergo the collapse of their iron cores and the subsequent ejection of their hydrogen-rich envelopes \\citep{Heg_etal03}. Observationally, Type II-P SNe are distinguished by prominent hydrogen lines in their spectra. When the star explodes with a significant fraction of its H-rich envelope, in theory it should display a light curve characterized by a phase of $\\sim$ 100 days of nearly constant luminosity followed by a sudden drop of 2-3 mag \\citep{Lit_etal83}. More than half of all Type II SNe belong to this class of ``Plateau'' SNe whose progenitors are attributed to stars born with less than 25 M$_{\\odot}$ \\citep{Hen_etal06,Li_etal07}. Type II-P SNe show a wide mass range, a considerable spread in explosive power, absolute magnitudes, and ejecta velocities \\citep{Ham_03}. \\citet{Ham_Pin02} have shown that there is a correlation between the expansion velocities of Type II-P SNe and their bolometric luminosities during the plateau phase.  Thus, Type II-P SNe can be useful as standardizable candles in their own right.  In this paper we address a simple question.  Can Type II-P SNe be found that exhibit similar enough color curves such that we may attribute the systematic differences of their colors to different amounts of dust extinction along the light of sight?  If the answer to this question is Yes, then in principle photometry can be used to obtain absolute magnitudes of Type II-P SNe which have minimal systematic errors owing to dust extinction along the line of sight. Accurate distances give us accurate luminosities, which are critical for constraining hydrodynamic models of Type II-P SNe as well as for cosmological studies. ",
        "conclusions": "\\citet{Ham_etal01} give an EPM distance to SN~1999em of 7.5 $\\pm$ 0.5 Mpc. A similar analysis by \\citet{Leo_etal02} yields 8.2 $\\pm$ 0.6 Mpc. Using improved photospheric velocities \\cite{Jon_etal08} obtain a distance of 9.3 $\\pm$ 0.5 Mpc. All these EPM distances were obtained using the same set of dilution factors calculated by \\citet{Eas_etal96}. Recently a new set of atmosphere models was calculated by \\citet{Des05} whose dilution factors are systematically higher than those of Eastman et al. and which lead to greater distances. Based on these new models \\citet{Des06} and \\citet{Jon_etal08} obtain 11.5 $\\pm$ 1.0 and 13.9 $\\pm$ 1.4 Mpc, respectively. \\citet{Leo_etal03} give a Cepheid-based distance of 11.7 $\\pm$ 1.0 Mpc, in good agreement with the EPM results derived from the new models of \\citet{Des05}, which are considerably greater than those obtained with the \\citet{Eas_etal96} models. If we adopt the Cepheid-based distance, the maximum observed $V$-band magnitude of 13.79 and total $V$-band extinction of 0.34 mag, it follows that M$_V$ = $-$16.89 $\\pm$ 0.24 mag for SN~1999em.  Using the \\dmm\\ method of \\citet{Phi_etal99}, the distance to NGC 1448, the host of SNe 2001el and 2003hn, is 17.9 $\\pm$ 0.8 Mpc \\citep{Kri_etal03}. Adopting the total $V$-band extinction of 0.586 $\\pm$ 0.050 mag for SN~2001el \\citep{Kri_etal07}, it follows that M$_V$(max) = $-$19.12 $\\pm$ 0.11.  A check on the distance of NGC 1448 can be obtained by assuming that the $JHK$ absolute magnitudes at maximum of SN~2001el equal the mean values given in Table 17 of \\citet{Kri_etal04b}.  \\citet{Kri_etal07} also found that R$_V$ = 2.15 is the most appropriate value for the host galaxy dust associated with SN~2001el. From the near-IR maxima we obtain a distance of 18.1 $\\pm$ 0.4 Mpc.  For comparison, the EPM analysis of SN~2003hn by \\citet{Jon_etal08} yields 16.9 $\\pm$ 2 and 26.3 $\\pm$ 7 Mpc, using dilution factors from \\citet{Eas_etal96} and \\citet{Des05}, respectively. The ``Standardized Candle Method'' (SCM) applied to SN~2003hn yields a distance of 17.8 $\\pm$ 1 Mpc \\citep{Oli_etal08}.  Under the assumption that the early-time photometric behavior of SN~2003hn was the same as that of SN~1999em, we can extrapolate that SN~2003hn was 0.056 mag brighter than our earliest $V$-band measurement.  Adopting V$_{max}$ = 14.41 $\\pm$ 0.03, A$_V$(total) = 0.58 $\\pm$ 0.14, and $d$ = 17.9 $\\pm$ 0.8 Mpc, it follows that M$_V$(max) = $-$17.44 $\\pm$ 0.17. Taken at face value, SN~2003hn was 0.55 $\\pm$ 0.30 mag brighter than its ``cousin'' SN~1999em.  Since SN~2003hn and 2001el occurred in the same galaxy, a comparison of their absolute magnitude differences involves no uncertainties in distance. Corrected for extinction, at the time of maximum light SN~2003hn was 1.68 mag fainter in $V$ than SN~2001el.  This confirms the notion that a typical Type II-P SN is significantly fainter than a Type Ia SN for the optical maxima.  In the near-IR, SNe 2001el and 2003hn are not so dissimilar in brightness at maximum. For SN~2001el the observed $K_{max}$ = 12.83 $\\pm$ 0.04 \\citep{Kri_etal03}.  Its total $K$-band extinction was about 0.057 mag.  The absolute magnitude M$_K$(max) $\\approx -18.49$.  For SN~2003hn, $K_{max}$ = 13.27 $\\pm$ 0.03 (from Table \\ref{03hn_yjhk}), A$_K \\approx 0.064$ mag, so M$_K$(max) $\\approx -18.06$.  This is only 0.43 mag fainter than SN~2001el.  The implication is that a wide angle SN survey carried out at 2.2 microns would find Type II-P SNe almost as easily as Type Ia SNe.  Since Type II SNe are single massive stars that have very short main sequence lifetimes, they end their lives very close to where they were born, in regions of significant levels of star formation.  As a result, the light curves of most (or all) Type II SNe would be affected by interstellar extinction along the line of sight.  We have found two Type II-P SNe whose color curves vary in a similar manner until the end of the plateau phase. The similarities of these color curves and increasing color excesses as we proceed from $V-R$ through $V-K$ imply that the observed color differences are simply due to differing amounts of dust along the line of sight. The implication is that some fraction of Type II-P SNe may exhibit sufficiently uniform color curves that they may be used for a determination of the {\\em relative} amounts of dust extinction that they suffer.  Once we are confident we have data on Type II-P SNe which are minimally reddened in their host galaxies, we can obtain accurate total extinctions for all the SNe with good optical and IR light curves which show those similar color curves. This will lead to a more accurate distance calibration for Type II-P SNe.  As future projects such as Pan-STARRS and LSST discover large numbers of SNe, we should be able to use Type Ia and Type II-P SNe for cosmology."
    },
    "0809/0809.3673.txt": {
        "abstract": "We investigate the connection between dark energy and fourth order gravity by analyzing the behavior of scalar perturbations around a Friedmann-Robertson-Walker  background. The evolution equations for scalar perturbation are derived using the covariant and gauge invariant  approach and applied to two widely studied $f(R)$ gravity models. The structure of the general {\\it fourth order} perturbation equations and the analysis of scalar perturbations lead to the discovery of a characteristic signature of fourth order gravity in the matter power spectrum, the details of which have not seen before in other works in this area. This could provide a crucial test for fourth order gravity on cosmological scales. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In spite of all the efforts made so far, the problem of the nature of Dark Energy (DE) is still far from a completely satisfactory resolution. Among the many theoretical frameworks proposed, the idea of a geometrical origin of Dark Energy has recently received a great deal of attention. The main reason for this popularity can be found in the fact that these type of theories of gravity, which are suggested by the low energy limit of very fundamental schemes \\cite{stringhe,birrell}, lead to cosmologies which admit naturally a Dark Energy era \\cite{SalvSolo,revnostra,Odintsov, Carroll,star2007} (and possibly even an inflationary one \\cite{kerner1,star80,Cognola1}) without the introduction of any additional cosmological fields.  Much of the investigation performed up to now on the idea of Geometric Dark Energy has been focused on the so-called fourth order theories of gravity. In these theories the Hilbert-Einstein action is modified with terms that are at most of order four in the metric tensor. The features of fourth order gravity have been analyzed with different methods \\cite{OurDynSys, BarrowDyn,Barrow Hervik,OdintRec,CognolaDyn} and it has been shown that their cosmologies can give rise to a phase of accelerated expansion which is considered the footprint of Dark Energy.  Although these results are very encouraging there are still some important open problems to be addressed. One of them is the analysis of  the evolution of the linear perturbations and their comparison with observations. Over the past year this problem has been studied by a number of authors,  by (1) considering different ways of parameterizing the non-Einstein modifications of gravity or (2) by simplifying the underlying fourth-order perturbation equations using a quasi-static approximation or a combination of (1) and (2) \\cite{Li:2008ai,HuSawicki,Bertschinger:2008zb,otherperts}.  In a number of recent papers \\cite{SantePert, K1} we derived the evolution equations for scalar and tensor perturbations of a subclass of fourth order theories of gravity characterized by an action which is a general analytic function of the Ricci scalar. In our work we study the dynamics of linear scalar perturbations using the covariant and gauge invariant approach developed for General Relativity (GR) in \\cite{EllisCovariant,EB,EBH,BDE,DBE,BED,DBBE}. This approach has the  advantage of using perturbation variables with a clear geometrical and physical interpretation. Furthermore, we use a specific recasting of the field equations that will make the development of the cosmological perturbation theory even more transparent, allowing one to integrate the perturbation equations exactly for a given $f(R)$ model without making any additional approximations.  The preliminary results obtained in \\cite{SantePert} showed some interesting features. First of all the evolution of scalar perturbations is determined by a fourth order differential equation rather than a second order one. This implies that the evolution of the density fluctuations contains, in general, four modes rather that two and can give rise to a more complex evolution than the one of General Relativity (GR).  Secondly, the perturbations are found to depend on the scale for any equation of state for standard matter (while in GR the evolution of the dust perturbations are not scale dependent). This means that, for example, in this framework the evolution of super-horizon and sub-horizon  perturbations is different. Third, and more surprisingly, we found  that growth of large density fluctuations can occur also in backgrounds in which the expansion rate is increasing in time. This is in striking contrast with what one finds in GR and what one would naively expect, but at the same time suggests new ways to tackle the DE problem.  The  features mentioned above imply that the evolution of perturbations in this framework can be completely different from the one we are familiar with. Yet this does not necessarily mean that they are incompatible with observations. Rather, they are a sign of the fact that in dealing with these models one has to resist the temptation of using assumptions which work well in GR. In this paper, following in this spirit, we will analyze further what was found in \\cite{SantePert} with the aim of achieving  a clearer understanding of the physics of the matter dominated era in fourth order gravity. In order to do this, we will rewrite the perturbation equations in a more physically meaningful way and will develop a series of tools which will make the analysis of the evolution of  density perturbations easier to understand and to compare with GR.  The paper is organized as follows. In section II we will give some  basic equations and we will present briefly the covariant gauge invariant formalism we use to develop the perturbation theory. In section III, we give the background and the perturbation equations. In section IV we rewrite these equation in an interesting form allowing us to discuss their general structure. In section V  we propose some useful tools to understand the behavior of the perturbations and compare it with what one obtains in General Relativity.  In section VI we apply these tools to some simple specific examples. Section VII is dedicated to the conclusions.  Unless otherwise specified, natural units ($\\hbar=c=k_{B}=8\\pi G=1$) will be used throughout this paper, Latin indices run from 0 to 3. The symbol $\\nabla$ represents the usual covariant derivative and $\\partial$ corresponds to partial differentiation. We use the $-,+,+,+$ signature and the Riemann tensor is defined by \\begin{equation} R^{a}{}_{bcd}=W^a{}_{bd,c}-W^a{}_{bc,d}+ W^e{}_{bd}W^a{}_{ce}- W^f{}_{bc}W^a{}_{df}\\;, \\end{equation} where the $W^a{}_{bd}$ are the Christoffel symbols (i.e. symmetric in the lower indices), defined by \\begin{equation} W^a_{bd}=\\frac{1}{2}g^{ae} \\left(g_{be,d}+g_{ed,b}-g_{bd,e}\\right)\\;. \\end{equation} The Ricci tensor is obtained by contracting the {\\em first} and the {\\em third} indices \\begin{equation}\\label{Ricci} R_{ab}=g^{cd}R_{acbd}\\;. \\end{equation} Finally the Hilbert--Einstein action in the presence of matter is given by \\begin{equation} {\\cal A}=\\int d x^{4} \\sqrt{-g}\\left[\\frac{1}{2}R+ L_{m}\\right]\\;. \\end{equation}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In this paper we have discussed the properties of scalar perturbations in fourth order gravity and described some useful methods for extracting physical information from the complex equations that govern their evolution. These tools are devised in  such a way to gain as clear an understanding as possible of the behavior of scalar perturbations on all scales and to facilitate a direct comparison of these results with the corresponding ones in GR. Two simple models:  $f(R)=R^{n}$ and $f(R)=R+\\alpha R^{n}$ were analyzed in detail because of their simplicity and because  we have a relatively good understanding of their background via the dynamical systems approach.  The results obtained show profound differences in the dynamics of the perturbations compared to what occurs in GR: in these models we find the growth rate of the perturbations is in general different from the GR one and always dependent on the scale. This implies, for example, that depending on the value of the parameters the perturbation can grow or dissipate (or both) at different rates, with obvious consequences for the global cosmic history.  The fourth order system of differential equations that governs the behavior of the scalar perturbations of these models, in particular when written in terms of   $\\Delta_m$ and $\\Delta_R$, has a structure that resembles closely one of a two fluids GR model. Although the $\\Delta_m$ and $\\Delta_R$  system of equations is enormously more complicated than (\\ref{EqPerIIOrd1}-\\ref{EqPerIIOrd2}), one can still use it to get an idea of the nature of the interaction between the non-Einstein part of the gravitational interaction and standard matter.  In section VII we found that the wavenumber structure of these coefficients is such that  for very large or very small $k$ they become scale invariant.  This implies, in turn,  that the matter power spectrum is  scale invariant for $k\\rightarrow 0,\\infty$ and can present some characteristic features on scales that depend on the different parameters of the model. Another way of interpreting the form of the spectrum without necessarily using the curvature fluid idea is to interpret $\\mathcal{R}$ and  $\\Re$ as being associated with the propagation of the scalar degree of freedom of the theory, or a scalar gravitational mode, whose interaction with matter is able to emulate the effect of a relativistic component, which, like photons in GR, induces power loss in the oscillations that appear in the spectrum.  The picture that emerges from our results is that in our examples, fourth order gravity influences deeply the structure formation process, but the modifications are only detectable around a specific value of $k$. Everywhere else in $k$ space the results are very close (although dynamically different) to GR. This is particularly interesting since it means that for a suitable choice of values for the parameters, the oscillations can be positioned beyond the observational boundary of currently available data. Hence our results seems to imply that not only could these models be compatible with the observed matter power spectrum, but that we also have a systematic way of constraining their parameters using data coming from  a range of scales, for example by combining CMB and LSS data \\cite{WMAP, SDSS}.  It is also worth commenting briefly about the compatibility of our analysis with the existing literature. For example in  \\cite{Li:2008ai} a class of models which is very similar to the one we analyzed in Section \\ref{SecRRn} is considered. Using a specific background the authors showed that the matter power spectrum  is characterized by an excess of power at small scales when the theory is very close to $\\Lambda$CDM. A similar result is found in \\cite{HuSawicki}  in which classes of models are considered which  contain additive corrections to the Hilbert -Einstein action (i.e. they have the form  $R+g(R)$ ). These general corrections are parameterized by  a quantity $B$, which measures the deviation from GR. Instead in \\cite{Bertschinger:2008zb}, generic modifications are considered (both scale-dependent and scale-independent), using a more flexible parameterization. In the sub-case of scale dependent modifications (like in the $f(R)$ case) examples are given in which one finds  once again excess of power on small scales in the power spectrum. Also one should note that in the same sub-case examples were given when one finds a deficit in power.  Interestingly, both the examples considered in this paper exhibit  excess power on small scales  for $n\\rightarrow 1^{+}$ in the case of $R^{n}$-gravity and for $n\\rightarrow 1^{+}$  and $\\alpha\\ll 1$ in the case $R+\\alpha R^n$. This indicates that, for small corrections  to GR, excess power at large scale  seems to be a generic feature in $f(R)$-gravity. However there are situations in which one might want to analyze theories which are not necessarily close to GR. For example in \\cite{Salv06,Salv07,Salucci} a fit of  $R^n$-gravity  with the data coming from the rotation curves of galaxies and supernovae type Ia leads to values of $n$ in the range $[1.7, 3.5]$. For these values of $n$, both our examples indicate a loss of power  at small scales, which means that we are provided with an opportunity to rule out this model by testing it against available data. Although the presence of an excess or deficit of power on small scales seems closely related  to the value of the specific parameters of the model itself, there are  some indications that other features of the spectrum  found in our examples  are indeed general. For example, the $k$ structure of the general equations (\\ref{EqPerIIOrdParGen1}-\\ref{EqPerIIOrdParGen2}) suggests that the evolution of perturbations in a generic $f(R)$ theory  presents at least three different regimes. Also looking at the derivation of the  $(\\Delta_m, \\Delta_R)$ equations (which can be performed in general, provided a sufficiently large amount of paper and time) one realizes that the $k$ dependence of the dissipation and source coefficients we have found in our examples is expected to be common to any $f(R)$ Lagrangian because it originates from  the $\\3nab^2\\3nab^2$ terms in the perturbation equations.  We end by commenting that these results together with the dynamical systems analysis of the background cosmological history presented in other papers \\cite{OurDynSys,SanteGenDynSys} provides a unified and consistent approach to the combined study of FLRW observational constraints and a complete analysis of linear structure growth in the context of $f(R)$ gravity.  \\vfill {\\noindent{\\bf Acknowledgements:}\\\\ The authors wish to thank Dr J. Larena for useful discussion and suggestions. SC wish to thanks J Donkers for useful discussion and support during the development of this paper. KNA and SC are supported by Claude Leon Foundation fellowships. This work was supported by the National Research Foundation (South Africa) and the {\\it Ministrero degli Affari Esteri - DIG per la Promozione e Cooperazione Culturale} (Italy) under the joint Italy/South Africa science and technology agreement. \\newpage \\appendix %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.2448_arXiv.txt": {
        "abstract": "An observation of the West lobe of radio galaxy Fornax~A (NGC~1316) with Suzaku is reported. Since \\citet{feigelson95} and \\citet{kaneda95} discovered the cosmic microwave background boosted inverse-Comptonized (IC) X-rays from the radio lobe, the magnetic field and electron energy density in the lobes have been estimated under the assumption that a single component of the relativistic electrons generates both the IC X-rays and the synchrotron radio emission. However, electrons generating the observed IC X-rays in the 1 -- 10 keV band do not possess sufficient energy to radiate the observed synchrotron radio emission under the estimated magnetic field of a few $\\mu$G. On the basis of observations made with Suzaku, we show in the present paper that a 0.7 -- 20 keV spectrum is well described by a single power-law model with an energy index of 0.68 and a flux density of $0.12\\pm0.01$ nJy at 1 keV from the West lobe. The derived multiwavelength spectrum strongly suggests that a single electron energy distribution over a Lorentz factor $\\gamma = 300 - 90000$ is responsible for generating both the X-ray and radio emissions. The derived physical quantities are not only consistent with those reported for the West lobe, but are also in very good agreement with those reported for the East lobe. ",
        "introduction": "\\label{sec:intro} % Cosmic jets of active galactic nuclei (AGNs) are significant kinetic energy channels from supermassive blackholes to intergalactic space. Some bulk flows release their entire kinetic energy into expanding radio lobes, which can be regarded as calorimeters for AGN outflows. The injected kinetic energy is thought to be turned into internal energy of plasma and induced magnetic field energy in the lobes. Although the energy densities of electrons and magnetic fields are often estimated assuming equipartition between the electron and magnetic fields, recent X-ray and radio imaging spectroscopy techniques have provided a way to resolve the two physical quantities without relying on this assumption.  From observations of the radio galaxy Fornax~A (NGC~1316) with ASCA and ROSAT, it was first discovered that the radio lobes generate inverse-Compton (IC) X-rays as well as a synchrotron radio emission. It was concluded that the IC X-rays were generated from the cosmic microwave background (CMB) photons (\\cite{kaneda95}; \\cite{feigelson95}), as has long been suspected (e.g. \\cite{harris79}). Following this discovery, through X-ray imaging spectroscopic measurements carried out by satellites in the past decade, including ASCA, ROSAT, Beppo-SAX, Chandra, and XMM-Newton, measurements of the electron energies of the lobes and the magnetic fields of a new level of precision have been achieved, in cooperation with measurements made using radio interferometers (e.g., \\cite{brunetti01}; \\cite{isobe02}; \\cite{isobe05} hereafter I05; \\cite{hardcastle02}; \\cite{grandi03}; \\cite{comastri03}; \\cite{croston04}).  The derived electron energy densities often exceeded those of the magnetic fields by a factor of a few to several dozen in a number of the observed radio lobe objects (e.g. \\cite{isobe02}; I05; \\cite{croston05}). However, we note that there is still room for argument against this observational claim. These evaluations are based on another assumption that the synchrotron and IC emissions are generated by relativistic electrons with a power law shaped energy spectrum; however, the validity of this law has not been strictly proven yet. In the case of typical radio lobes, the radio synchrotron radiation is generated by electrons with $\\gamma \\sim (1-6) \\times10^4$, whereas the observed IC X-rays in the range of 1 -- 10 keV are generated by electrons with $\\gamma \\sim (1-3) \\times 10^3$. In this regard, we need to examine the electron energy distribution in the range of $\\gamma > 3 \\times 10^3$ by observing IC X-rays with energies of over 10~keV.  We, therefore, used Suzaku to observe the hard X-ray extension of IC X-rays with energies of up to 20~keV to cover this observational gap. We expected to observe that the electrons producing the $\\sim 20$ keV IC X-rays, would also be responsible for generating sub-100MHz synchrotron radio radiation. Suzaku, which is equipped not only with a high-throughput X-ray telescope (XRT; \\cite{xrt07}) and CCD cameras (XISs; \\cite{xis07}) but also with the most sensitive hard X-ray spectrometer (HXD; \\cite{hxd07}) currently available, is one of the most suitable instruments for the purpose of measuring the spectrum over 10~keV, the region corresponding to the electrons.  The structure of this paper is as follows. We describe the Suzaku observation and the results in \\S~2 and \\S~3. Following the presentation of the results, we evaluate the newly obtained X-ray spectra from the viewpoint of a multiwavelength spectrum and derive the physical quantities with respect to the lobe field in \\S~4. Finally, we summarize these results in the last section. In the Appendix, we present our analysis and results previously obtained by XMM-Newton in order to evaluate contaminant point sources in the West radio lobe region. Throughout the present paper, we assume the distance to Fornax~A to be 18.6 Mpc (\\cite{madore99}), on the basis of measurements made by the Hubble Space Telescope of Cepheid variables in NGC~1365, a galaxy that, along with Fornax~A, belongs to the Fornax Cluster. This distance gives an angle-to-physical size conversion ratio of 5.41 kpc~arcmin$^{-1}$ and a red shift of the rest frame  $z_{\\rm rest}=0.00465$, while the measured red shift, including the proper motion, is reported to be $z_{\\rm obs} = 0.005871$ (\\cite{longhetti98}). ",
        "conclusions": "% We succeeded in observing non-thermal X-rays from the Fornax~A West lobe with the HXD/PIN instrument onboard Suzaku. Having assumed a hard X-ray emission extending over the lobe with a radius of $12'$, we confirmed that the measured hard X-ray spectral slope and flux are in good agreement with the extrapolation of the diffuse non-thermal component observed with XIS below 10~keV. The non-thermal X-rays can be described sufficiently well by a single power law model over the entire range from 0.5 up to 16 keV (with the nominal NXB) or to 20 (with the rescaled NXB) keV, and the spectral slope is consistent with that observed in the radio spectrum. Employing the derived power law component for the diffuse hard X-rays, we plotted the spectral energy distribution (SED) together with the radio spectra in figure~\\ref{fig:sed}.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig4.eps} \\end{center} \\caption{(A) Spectral energy distribution from the West lobe (squares), East lobe (smaller diamonds with crosses), and total (larger diamonds). Derived models for the synchrotron emission (dashed line) and the CMB-boosted IC emissions (solid line) are shown with lines (see text). The 29.9 and 100 MHz data points are taken from \\citet{finlay73}, the 408 MHz data points from \\citet{robertson73}, the 843 MHz data points from \\citet{jones92}, the 1.4 GHz data points from \\citet{ekers83}, the 2.7 GHz data points from \\citet{ekers69}, \\citet{shimmins79} and \\citet{kuehr81}, and the 5.0~GHz data point from \\citet{kuehr81}.}\\label{fig:sed} \\end{figure}%  An X-ray spectrum following a single power law distribution, which does not smoothly connect to the synchrotron component, requires another origin. The IC process is the most promising origin for the observed hard X-rays reported so far. Taking the size of and the distance from the host galaxy into account, we found that the photon energy density of the CMB at the source rest frame ($4.3\\times 10^{-13}$erg~cm$^{-3}$) exceeds both emissions from the host galaxy emission ($\\sim$ a few $\\times 10^{-15}$erg~cm$^{-3}$) and synchrotron emission from the lobe ($\\sim 6 \\times 10^{-17}$erg~cm$^{-3}$) by two orders of magnitude (e.g. \\cite{rband07} and references in the caption of figure~\\ref{fig:sed}).  In accordance with \\citet{harris79}, we derived the estimated energy density of the electrons and the magnetic field energy density to be $u_{\\rm e} = (5.1 \\pm 1.0) \\times 10^{-13}$erg~cm$^{-3}$ and $u_{\\rm m} = (0.67 \\pm 0.08)\\eta^{-1} \\times 10^{-13}$erg~cm$^{-3}$, respectively, where $\\eta$ is the filling factor of magnetic field against the electron spatial distribution. Here, we assumed an electron Lorentz factor of $\\gamma = 300 - 90000$, as we describe below, and included the systematic errors arising from the ambiguity of the reported radio flux and slope according to \\S~4 in I05. Note that the difference in brightness profiles between radio and X-rays is discussed with ASCA (T01). The displacement is thought to originate from partial displacement between the magnetic field and the electron distribution. T01 examined it and showed the magnetic field filling factor could be $\\eta = 0.95$ or less, with the electrons filling the entire volume homogeneously.  Adopting the above values, we calculated the CMB boosted IC spectrum and plotted the results in the obtained SED in figure~\\ref{fig:sed}. Here, we assumed an electron differential energy spectrum with a single power law distribution of \\begin{displaymath} N_{\\gamma} = 2.0 \\times 10^{-6} \\gamma^{-2.36} \\;\\; {\\rm electrons~cm}^{-3}, \\end{displaymath} Here, we employed the lower boundary electron energy distribution of $\\gamma = 300$ to describe the lower energy end of the detected X-rays, although we have no evidence of a lower energy cutoff of the electron energy distribution. Similarly, we set the upper boundary of $\\gamma=90000$ in accordance with the higher end of the radio observation of the total lobe observations. The calculation program was written to use the SSC equations presented in \\citet{kataoka00}, to which we added the CMB/IC spectrum, and it can be seen that it provides a good description of the obtained multiwavelength data.  The electron Lorentz factor of $\\gamma = 300 - 4000$ (or 4500) is required by the CMB/IC X-ray spectrum in the range of 0.5 -- 16 (or 20) keV, where we employed the nominal (or rescaled) NXB model for the HXD/PIN. At the same time, adopting the derived magnetic field, the observed synchrotron radiation spectrum requires a Lorentz factor of $\\gamma > 4200$. Thus, we conclude that a single electron energy distribution can naturally explain the two independent measurements in radio and X-ray bands. The derived parameters are summarized in table~\\ref{tab:param}. We also note that the derived electron and magnetic field energy densities are not only consistent with those determined in previous reports (T01) but are also in very good agreement with those determined for the East lobe (I05).  \\begin{table} \\caption{Derived physical quantities}\\label{tab:param} \\begin{tabular}{cccc} \\hline \\hline % & West Lobe    & East Lobe\\footnotemark[$\\dagger$] & unit  \\\\ \\hline % (1) $B_{\\rm eq}$     & 1.59          & 1.55                   & $\\mu$G \\\\ (2) $B_{\\rm IC}$     & $1.3\\pm0.1$   & $1.23^{+0.08}_{-0.06}$ & $\\mu$G \\\\ (3) $u_{\\rm m}$      & $0.67\\pm0.08$ & $0.60^{+0.08}_{-0.06}$ & $10^{-13}$ erg~cm$^{-3}$  \\\\ (4) $u_{\\rm e}$      & $5.0\\pm1.0$   & $5.0^{+1.1}_{-1.0}$    & $10^{-13}$ erg~cm$^{-3}$   \\\\ $u_{\\rm e}/u_{\\rm m}$ &  $7.5 \\pm 0.6$ & $5.0^{+1.1}_{-1.0}$  &  \\\\ \\hline % \\end{tabular} \\footnotemark[$\\dagger$]: Taken from I05 for reference.\\\\ (1) Estimated magnetic field assuming electron-magnetic field energy equipartition.\\\\ (2)---(4) Magnetic field strength, magnetic field energy density, and electron energy density as derived from the observed CMB/IC X-rays and the synchrotron radio waves, respectively (see the text). \\end{table}"
    },
    "0809/0809.1741_arXiv.txt": {
        "abstract": "{} % {We aim to identify the stellar populations (mostly red giants and young stars) detected in the ISOGAL survey at 7 and 15$\\mu$m towards a field (LN45) in the direction $\\ell=-45, b=0.0$.} {The sources detected in the survey of the Galactic plane by the Infrared Space Observatory are characterized based on colour-colour and colour-magnitude diagrams.  We combine the ISOGAL catalog with the data from surveys such as 2MASS and GLIMPSE.  Interstellar extinction and distance are estimated using the red clump  stars detected by 2MASS in combination with the isochrones for the AGB/RGB branch.  Absolute magnitudes are thus derived and the stellar populations are identified based on their absolute magnitudes and their infrared excess.  } {A standard approach to the analysis of ISOGAL disk observations has been established.    We identify several hundred RGB/AGB stars and 22 candidate young stellar objects in the direction of this field in an area of 0.16 deg$^2$.  An over-density of stellar sources is found at distances corresponding to the distance of the Scutum-Crux spiral arm.  In addition, we determine mass-loss rates of AGB-stars using dust radiative transfer models from the literature.} {} ",
        "introduction": "%  ISOGAL survey data at 7 and 15~$\\mu$m of about 16deg$^2$ in selected fields of the inner Galactic disk have allowed the detection of about 1~$\\times$~10$^5$ point sources down to $\\sim$10~mJy at 15~$\\mu$m and 7~$\\mu$m. In conjunction with ground based near-IR surveys (DENIS, 2MASS), they offer the possibility to investigate the different populations of infrared stars in the Galactic disk up to 15~$\\mu$m. The technical characteristics of the five wavelength ISOGAL-DENIS point source catalogue are presented and discussed in detail by \\citet{Schuller2003}, while \\citet{Omont2003} have reviewed the scientific capabilities and the main outcome of ISOGAL.  The best detected stellar class is that of AGB stars which are almost completely detected above the RGB tip at least at 7~$\\mu$m up to the Galactic Center \\citep{Glass1999,Omont1999}. The infrared color ($K_S$-[15])$_\\mathrm{0}$ is a very good measure of their mass-loss rate \\citep{Ojha2003}. The mass-loss rate can also be derived from the pure ISOGAL color [7]-[15], which is practically independent of extinction. Combined with variability data from MACHO\\citep{Alcock1997} or EROS\\citep{Palanque1998}, ISOGAL data of bulge fields have shown that practically all sources detected at 15~$\\mu$m are long period variables with interesting correlations between the mass-loss rate traced by the 15~$\\mu$m excess and variability intensity and period \\citep{Alard2001}. The 15~$\\mu$m data from the fields observed (0.29 deg$^2$ in total) in the Galactic bulge ($\\vert b \\vert \\geq 1^\\circ$) have been used by \\citet{Ojha2003} to infer global properties of the mass returned to the interstellar medium by AGB stars in the bulge. However, similar work has not yet been carried out in the disk ISOGAL fields because of the difficulty in properly estimating distances along the line of sight.  Red giants are by far the most numerous population of luminous bright stars in the near and mid-infrared.  They are detectable through out the Galaxy with modern surveys such as 2MASS, DENIS, GLIMPSE, ISOGAL and MSX. The reddening of various infrared colours may be used to trace the interstellar extinction, A$_\\mathrm{V}$. However, the near-IR colours $J-K_S$ or $H-K_S$ are generally the most useful because of greater sensitivity to extinction with A$_\\mathrm{J}$-A$_\\mathrm{Ks}$~$\\sim$~0.17\\,A$_\\mathrm{V}$ and A$_\\mathrm{H}$-A$_\\mathrm{Ks}$~$\\sim$~0.06\\,A$_\\mathrm{V}$ \\citep{Glass1999book}. Such methods have been used for systematic studies of the extinction in large areas from DENIS \\citep{Schultheis1999} and 2MASS \\citep{Dutra2003}, and also for modeling  of Galactic stellar populations from these surveys \\citep{Marshall2006}.  When the luminosities of the red giants are known, e.g. those of the 'red clump' or the RGB tip, one may infer both extinction and the distance from colour-magnitude diagrams (CMD) such as $J$ vs $J-K_S$. The numerous giants of the red clump are known to be, by far, the best way to make 3D estimates of the extinction along Galactic lines of sight by determining distance scales from their well defined luminosities and intrinsic colours. \\citet{Lopez-Corredoira2002}, \\citet{Drimmel2003} and \\citet{Indebetouw2005} have exploited the red clump stars of DENIS and 2MASS to explore the extinction along various lines of sight.  The 6-8~$\\mu$m range is not as good an indicator of AGB mass-loss as 15~$\\mu$m. However, the sensitivity of ISOGAL at 7~$\\mu$m allows the detection of less luminous giants below the RGB tip at the distance of the bulge \\citep{Glass1999}. The large number of such stars, $\\sim$10$^5$, detected by ISOGAL in the Galactic disk and inner bulge may also be used to study the mid-infrared extinction law at 7~$\\mu$m by comparing $K_S$-[7] with $J-K_S$.  \\citet{Jiang2003} have used this approach along one line of sight. This method may even be extended to the derivation of extinction at 15~$\\mu$m from the ratio $K_S$-[15]/$J-K_S$. It has recently been applied to all the exploitable ISOGAL lines of sight (more than 120 directions at both wavelengths of ISO) in the Galactic disk and inner bulge \\citep{Jiang2006}.  Young stars with dusty disks or cocoons are the other class of objects to be addressed using ISOGAL's ability to detect 15 $\\mu$m excess. \\citet{Felli2002} have thus identified 715 candidate young stars from relatively bright ISOGAL sources with a very red [7]-[15] color.  \\citet{Schuller2002} has proposed another criteria for identifying such luminous young stellar objects based on the non-point source like behaviour of the 15~$\\mu$m emission. However, various reasons have limited the full exploitation of ISOGAL for young star studies; e.g.  lack of complementary data at longer or shorter wavelengths which makes it difficult to have a good diagnostic of the nature of the objects, their luminosity and mass; lack of angular resolution which may preclude deblending of nearby sources; limited quality of the data, especially in the regions of active star formation with high diffuse infrared background.  Other infrared surveys, IRAS and MSX \\citep{Price2001}, have covered the complete Galactic disk, including bands at longer wavelengths, but with a very limited sensitivity, especially in the range 12--20~$\\mu$m where ISOGAL is three or four magnitudes deeper than MSX. However, the much increased panoramic capabilities of {\\it Spitzer Space Observatory} have now made available much deeper data in the four IRAC bands (3.6, 4.5, 5.8 and 8.0$~\\mu$m) in the main part of the whole Galactic disk from the GLIMPSE {\\it Spitzer Legacy Project} \\citep{Benjamin2003,Benjamin2005}, extended to the whole inner disk with GLIMPSE II. GLIMPSE is about one order of magnitude deeper than ISOGAL in the range 6-8$\\mu$m, and it has a better angular resolution, but it lacks extension at longer wavelength as provided by ISOGAL at 15~$\\mu$m. We note that the MIPSGAL project with {\\it Spitzer} will provide longer wavelength (MIPS bands at 24$\\mu$m and 70$\\mu$m) coverage in the near future \\citep{Carey2005}.  The main purpose of the present paper is to begin a reassessment of ISOGAL data in conjunction with the availability of GLIMPSE data. We have chosen a standard ISOGAL field with good quality data and covering a relatively large area and latitude range, allowing a valuable statistical study. We have made a complete analysis of the ISOGAL data and validated their quality using the GLIMPSE data. We discuss their various science outputs, especially at 15~$\\mu$m, including complementary information in the line of sight, especially from GLIMPSE/2MASS and \\element{C}\\element{O} millimetre observations. Our main goal from such a case study is to validate general methods for a subsequent complete exploitation of 15~$\\mu$m data in all ISOGAL fields, complemented by near-IR and GLIMPSE data, especially for systematic studies of AGB stars and their mass-loss and dusty young stars of intermediate mass.  The paper is organized as follows: Sect. \\ref{sec_dataset} recalls the general properties of ISOGAL data and describes the associations of ISOGAL point sources in this field with GLIMPSE, 2MASS, and MSX sources. Section \\ref{sec_stats} discusses corresponding statistics and the validation of ISOGAL quality using GLIMPSE data. Section \\ref{sec_ext_dist} is devoted to interstellar extinction in this direction, the mid-IR extinction law and stellar dereddening, and comparison with \\element{C}\\element{O} emission across the field, in order to infer a rough three dimensional picture of extinction and source distribution. Section \\ref{sec_nature} deals with the nature of the sources: the AGB population, its mass-loss, luminosity function and relation with the RGB population, identification of the relatively few young stars in this direction and their properties and relationship with various other indicators of star formation. ",
        "conclusions": "%  We have studied a large ISOGAL field towards the Galactic disk and established a standard mode of studying the large set of ISOGAL observations of the disk.  We show that the  stars of  the  red clump can be traced all along the line of sight using the 2MASS $J$ and $K_S$ data up to the start of the  Scutum--Crux  arm  towards  the  $\\ell=-45\\degr$ direction.  We find  that the  locally accepted  value of  the extinction per unit distance, $c_\\mathrm{J}$, is sufficient to  fit well the red  clump locus up to 2.5kpc in this direction. However,  beyond 2.5kpc,  $c_\\mathrm{J}$ varies  with  galactic  latitude  and increases with distance.  We use the red clump  locus to obtain the distance and  extinction towards  individual stars  assuming them  to be  red giants  and thus  following the  RGB/AGB isochrone.  The  distribution of  stellar  density  rises as one hits the spiral arm   at  4kpc.  The 2MASS  data are  not deep enough to  detect    the stars  of  the  red clump at distances larger than 4kpc at the lowest galactic latitudes.   We do see more luminous AGB  stars to  larger distances.  Most of the red giants brighter  than  the stars of the red  clump are detected by  ISOGAL at 7$\\mu$m.  The ($K_S$-[15])$_0$ colour provides the mass loss  rates for the AGB stars.  There are not many AGBs  with large  mass-loss rate in  this  direction.  From the mid infrared colour excess we identify a total of 22 YSO candidates in this field.  We provide a catalog of the sources detected by ISOGAL with the estimated extinction and distance.  We tabulate also the mass-loss rates for the several hundred red giants towards this field.  \\medskip {\\it Acknowledgements.}  This work is based on observations made with the {\\em Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. We are grateful to the GLIMPSE project for providing access to the data. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This research made use of data products from the Midcourse Space Experiment.  Processing of the data was funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science.  This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. We acknowledge use of the TOPCAT software from Starlink and the VOPLOT virtual observatory software from VO-India (IUCAA) in this work. We thank the referee, Sean  Carey, for constructive comments and criticisms which have improved the paper. S. Ganesh was supported by Marie-Curie EARA fellowship to work at the Institut d'Astrophysique de Paris. We are grateful to the Indo-French Astronomy Network for providing travel support to enable M. Schultheis to visit PRL. Work at the Physical Research Laboratory is supported by the Department of Space, Govt. of India."
    },
    "0809/0809.1094_arXiv.txt": {
        "abstract": "The effective extinction law for supernovae surrounded by circumstellar dust is examined by Monte-Carlo simulations. Grains with light scattering properties as for interstellar dust in the Milky-Way (MW) or the Large Magellanic Clouds (LMC), but surrounding the explosion site would cause a semi-diffusive propagation of light up to the edge of the dust shell. Multiple scattering of photons predominantly attenuates photons with shorter wavelengths, thus steepening the effective extinction law as compared to the case of single scattering in the interstellar medium. Our simulations yield typical values for the total to selective extinction ratio $R_V\\sim 1.5-2.5$, as seen in recent studies of Type Ia supernova colors, with further stiffening differential extinction toward the near-UV. ",
        "introduction": "The uncertainties in the brightness corrections of Type Ia supernovae (SNIa) for color excess is among the largest systematic uncertainties in the use of Type Ia supernovae to measure cosmological distances \\citep{smock}. The standard interpretation of color excess being due to extinction by interstellar dust in the supernova host galaxy has recently been challenged by the empirically deduced color-brightness relation for SNIa. While studies of differential extinction of quasars shining through foreground galaxies yield values of $R_V= A_V/E(B-V)$ compatible with the average MW value \\citep{QSO}\\footnote{It should be noted that low $R_V$ values to individual QSO systems have been found in e.g. \\citep{QSO} and \\citep{Wang04}}, the SNIa Hubble diagram scatter is minimized for values significantly smaller than  $R_V=3.1$ \\citep{SNLS,scp08}. Furthermore, the use of SNIa for cosmology benefits significantly from the understanding the extinction law for the full optical and near-IR wavelength range. E.g. the wavelength dependence of extinction toward the Magellanic Clouds differ from the Milky-Way extinction law in \\citep{CCM}, even for very similar values of $R_V$. For quasar sight-lines, \\citet{QSO} tested both Milky-Way like extinction as well as Small Magellanic Cloud (SMC) extinction law, both giving comparable goodness of fit. A preference for SMC dust for extinction of AGNs has been suggested by \\citet{li}.   Recently, the detection of circumstellar (CS) matter in the local environment surrounding the Type Ia supernova SN2006X in the nearby galaxy M100 has been reported by \\citet{Patat07}. A shell within a few $10^{16}$ cm ($\\sim 0.01$ pc) of the center of the explosion has been suggested to explain the time-variable Na I D lines in the SN spectrum.  \\citet{Wang08a} report $R_V=1.48 \\pm 0.06$ and $E(B-V) = 1.42 \\pm 0.04$ mag for SN2006X and a light echo in the lightcurve was found by \\citet{Wang08b} consistent with dust illuminated at a distance of 27-170 pc from the site of the explosion. Even if the local environment around this supernova may not be very common among SNIa, similar values for the total to selective extinction ratio have been reported for several SNIa with good wavelength coverage. E.g. \\citet{Krisciunas07} found $R_V=1.55 \\pm 0.08$ for SN 1999cl; \\citet{Elias-Rosa06,Elias-Rosa08} report $R_V=1.80 \\pm 0.19$ and $R_V=1.59 \\pm 0.07$ for SN2003cg and SN 2002cv respectively. Furthermore, a statistical study of optical colors of a sample including 80 near-by SNIa, \\citet{Nobili&Goobar} found an average value of $\\bar R_V=1.75 \\pm 0.27$ for SNIa with $E(B-V)<$0.7, and even lower for a subsample of low-reddening SNIa.  Next, we examine the possibility that low values of $R_V$ stem from the semi-diffusive propagation of photons in the neighborhood of the site of the supernova explosion. ",
        "conclusions": "Simple simulations show that circumstellar material, detected in at least one Type Ia supernova, SN2006X \\citep{Patat07,Wang08b}, could potentially explain the empirically determined extinction law for low redshift SNIa, especially if the circumstellar material resembles LMC dust grains. Adopting the CS shell size of \\citet{Patat07}, $R_{CS}\\sim 10^{16}$ cm, we find that for $\\tau\\sim1$, the required mass in dust around the supernova is $M_{dust}\\sim 4 \\pi R_{CS}^2/(\\sigma_a/m_{dust})\\sim 10^{-4} M_\\odot$ when inserting typical values of the absorption cross-section from Table~\\ref{tbl-1}.  A simple power-law expression is found to fit very well the effective extinction law for dust in the CS environment of the supernova produced by Monte-Carlo simulations.  Depending on the thickness of the CS shell, shifts in the time of lightcurve maximum may be expected for different bands since the amount of quasi random-walk will differ. In particular, photons in redder bands will suffer less scattering and thereby less time delay. This effect should correlate with the measured reddening, $E(B-V)$. The assumption of a uniform density is not expected to be critical for the results at first order. However, a second order effect may be expected since a large scale of $R_{CS}$ would result in a longer time for photons being ``trapped'' in the scattering sphere. As the the intrinsic colors of Type Ia change on a time scale of days \\citep{Nobili&Goobar}, time delays of photons of that time scale would affect the measured colors as a function of time.  In a forthcoming paper, potential direct observables from interaction between photons and dust will be investigated. Also, the sensitivity of the effective extinction law to the dust grain sizes and density profile in CS medium and the combination of both scattering in the circumstellar material and the interstellar medium needs to be further investigated. If the presence of CS material is indeed the source of the color-brightness relation found in SNIa, the case for restframe near-IR observations is further strengthened: the peak magnitude corrections, and their model dependence, are smaller than at optical wavelengths."
    },
    "0809/0809.4491.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) {This work aims to provide a theoretical formulation of Surface Brightness Fluctuations (SBF) in the framework of probabilistic synthesis models, in which there are no deterministic relations between the different stellar components of a population but only relations on average, and to distinguish between the different distributions involved in the SBF definition.} % methods heading (mandatory) {By applying the probabilistic theory of stellar population synthesis models, we estimate the shape (mean, variance, skewness, and kurtosis) of the distribution of fluctuations across resolution elements, and examine the implications for SBF determination, definition and application.} % results heading (mandatory) {We propose three definitions of SBF: (i) stellar population SBF, which can be computed from synthesis models and provide an intrinsic metric of fit for stellar population studies; (ii) theoretical SBF, which include the stellar population SBF plus an additional term that takes into account the distribution of the number of stars per resolution element $\\psi(N)$; theoretical SBF coincide with Tonry \\& Schneider (1998) definition in the very particular case that $\\psi(N)$ is assumed to be a Poisson distribution. However, the Poisson contribution to theoretical SBF is around 0.1\\% of the contribution due to the stellar population SBF, so there is no justification to include any reference to Poisson statistics in the SBF definition; (iii) observational SBF, which are those obtained in observations that are distributed around the theoretical SBF. Finally, we show alternative ways to compute SBF and extend the application of stellar population SBF to defining a metric of fitting for standard stellar population studies. } %conclusions heading (optional), leave it empty if necessary {We demostrate that SBF are observational evidence of a probabilistic paradigm in population synthesis, where integrated luminosities have an intrinsic distributed nature, and they rule out the commonly assumed deterministic paradigm  of stellar population modeling. } ",
        "introduction": "Stellar population synthesis studies aim to decompose the integrated light, $L_{\\mathrm{tot};\\lambda}(N^*)$, of a sample of $N^*$ stars into a combination of particular stellar classes or populations, $n_i$. These studies can be developed by following two different methods. The first method \\cite[][ as examples]{ST71,Fab72} completes a direct decomposition of the integrated light into principal components. The resulting solutions are {\\it a posteriori} constrained by additional assumptions to obtain  the  description of the entire stellar population. In the second method  \\cite[evolutionary stellar population synthesis, hereafter EPS; see][for an example]{Tin68}  the stellar populations are modeled following {\\it a priori} rules related to stellar evolution and the relative contribution of each stellar population. Model results are compared with observations, and properties of the stellar populations (and their evolutionary status) of the targeted source are inferred. In both types of methods, degenerate solutions are present. The EPS method, however, is based on stellar evolution theory and provides both a tool for both studying the integrated light of galaxies and comparing its properties with stellar model predictions.  By construction, EPS needs to assume the existence of a probability distribution function that describes the mean value of the relative contribution of different populations to the integrated light, $\\mu_1'(n_i)/N^*$. In addition, it assumes that each population can be characterized by a mean luminosity, $\\mu_1'(l_{i;\\lambda})$. EPS studies therefore refer to the {\\it generic} emission of the integrated light of a set of systems rather than the {\\it particular} emission of a given system. In a work inspired by \\cite{GGS04}, \\cite{CLCL06}, demonstrate that EPS provides a characterization of the probability distributions that describe {\\it all} the possible integrated luminosities of an ensemble of stars with given evolutionary conditions\\footnote{Throughout the paper, evolutionary conditions mean the age, metallicity, and star formation history of the stellar ensemble.}. In general, this characterization is expressed in terms of the {\\it mean}, $\\mu_1'(\\ell_\\lambda)$, of the distribution of the integrated luminosity of a system with a reference number of stars, $N^*_\\mathrm{norm}$, that allows us to obtain the mean value of a distribution for a different number of stars, $N^*$, by simple transformations\\footnote{In the following, we assume $N_\\mathrm{norm}=1$ and we do not include the subindex $\\lambda$ in the luminosities to simplify the notation. We also omit the subindex $\\lambda$ in the following equations.}, $\\mu_1'(L_{\\mathrm{tot};\\lambda};N^*) = \\frac{N^*}{N^*_\\mathrm{norm}} \\, \\mu_1'(\\ell_\\lambda)$. However, it is also possible to provide a more accurate characterization of such distributions using additional parameters of the distribution, such as the second raw moment, $\\mu_2'(\\ell_\\lambda)$, the variance, $\\kappa_2(\\ell_\\lambda)$, the skewness, $\\gamma_1$, and the kurtosis, $\\gamma_2$,  or even provide the distribution itself \\citep{CLCL06}. This natural description of EPS in probabilistic terms, in which it is possible to establish the distribution of the integrated luminosity of a stellar ensemble, but not the precise position of this ensemble in the distribution, is not commonly considered in EPS literature, nor its application. A deterministic paradigm is usually applied:  $\\mu_1'(n_i)/N^*$ is not a mean value but the actual fraction of stars of a given class present in the stellar system, regardless of the number of stars in the system. In the deterministic paradigm, given the total number of stars in the system and its evolutionary conditions, there is one and only one possible value of the integrated luminosity.  However, the probabilistic description of stellar populations,  although not recognized and studied in detail before, has been present in the literature since the late eighties.  \\cite{TS88} presented the first description of the distribution of the integrated luminosities of stellar populations making use of the mean and the variance in the associated theoretical distributions, and defined theoretical Surface Brightness Fluctuations (SBF, $\\bar{L}_{\\mathrm{tot}}(N^*)$) as the variance of the integrated luminosity of a stellar population divided by its mean: $\\bar{L}_{\\mathrm{tot}}(N^*) = \\kappa_2[{L}_{\\mathrm{tot}}(N^*)]/\\mu_1'[{L}_{\\mathrm{tot}}(N^*)]$. Interestingly the dependence of $\\kappa_2[{L}_{\\mathrm{tot}}(N^*)]$ and $\\mu_1'[{L}_{\\mathrm{tot}}(N^*)]$ on the total number of stars cancels out numerically and $\\bar{L}_{\\mathrm{tot}}(N^*) = \\bar{\\ell}$ is independent of the number of the stars in the system. {\\it Assuming a Poisson distribution} for the number of stars in each given stellar population, SBF can also be expressed as the ratio of the second to the first raw moments of the distribution of integrated luminosities, and in this case, it can be redefined to be the luminosity-weighted luminosity of the population.  \\cite{Buzz89} presented an independent probabilistic description  of EPS, also based on the variance of the corresponding distributions, and defined the effective number of stars, ${\\cal{N}}(\\ell)$, to be the ratio  $\\mu_1'(\\ell)^2/\\kappa_2(\\ell)$.  The  value of ${\\cal{N}}[L_\\mathrm{tot}(N^*)]$ for a system with $N^*$ stars scales directly with $\\cal{N}(\\ell)$ and it is an useful parameter for estimating sampling effects produced by the discreteness of real stellar populations. In his formulation, it was assumed that the relative number of stars in a given population fluctuates around the theoretical value  way described by a Poisson distribution \\cite[but see][ and below]{CLCL06}. The connection between $\\cal N$ and SBF was defined by \\cite{Buzz93}, who related both quantities to the statistical entropy of stellar systems \\cite[see also][ for a recent review]{Buzz07}.  As shown by \\cite{TS88}, SBF is not only a theoretical concept, but a real observable that provides a relatively direct technique for determining extragalactic distances. The amplitude of the relative fluctuations in the observed flux with respect to a suitable mean is related directly to the theoretical SBF luminosity, and therefore to distance. In the observational domain, its principal asset is being able to disentangle the fluctuations that provide physical information about the system from the noise. Since they are related to evolutionary conditions, SBF provide information about stellar evolution and stellar population studies \\cite[e.g.][]{CRBC03,GLLB07}.  Most work performed on theoretical SBF was related to its calibration as a distance indicator \\cite[e.g.][]{TAL90}, although current calibrations are performed empirically. Further examples of SBF in stellar populations studies are \\cite{W94,Con97,BCC98,BVA01}, and \\cite{LCG00}. In the SBF definition, it is customary to assume that the number of stars in a given stellar population follows a Poisson distribution. Unfortunately, no complete study exists on theoretical SBF themselves and their implication in the use of EPS: the very observation of SBF implies that EPS models cannot be interpreted in a deterministic way, but as a description of a probability distribution that defines naturally an intrinsic {\\it metric of fit} when physical properties are inferred from comparisons with model results \\cite[see][ for more details]{CL07a,CL07b}.  On the other hand, the distributed nature described by EPS models was established clearly in theoretical studies of systems with a low number of stars, where the most probable number of stars in a given evolutionary stage $n_i/N^*$ differs significantly from the expected value i.e. the mean $\\mu_1'(n_i)/N^*$. The characterization of integrated light in probabilistic terms began more recently. Partial solutions to the problems were presented by \\cite{CLC00,BruTuc02,Gi02,LM00,CVGLMH02,Ceretal01,CVG03,CL04}, and \\cite{FRDI07} among others, and a comprehensive solution to the problem was presented by \\cite{CLCL06}, who introduced a probabilistic paradigm of EPS that applies to both under-sampled and well sampled systems.  In this paper, we establish a framework for theoretical SBF based on the probability distribution theory of stellar populations. The formalism allow us to describe the behavior of SBF for systems with a low number of stars per pixel, which was demonstrated by \\cite{AT94} to be a limitation of the method. Our main objective is to establish a robust definition of SBF. In particular, we show that the Poisson assumption is not essential to the SBF definition. We examine the implications of this robust definition and explore new ways to obtain observational SBF. Paraphrasing \\citet{KP99}, ``once the definitions are laid out, the theory tends to fall into place, and understanding it is comparatively easy''.  The present paper is organized as follows. In Sect. \\ref{sec:probabilityetc}, we summarize the theoretical probability distributions involved in SBF studies and provide an introduction to the probabilistic description of stellar populations. In Sect.~\\ref{sec:theoreticalSBF}, we present a detailed and quantitative description of SBF, establish the origin of SBF in probabilistic terms, and provide robust definitions of SBF according to the distributions involved. A discussion about the implications of the present result for SBF studies and additional SBF analysis methods are presented in Sect. \\ref{sec:discussion}. Finally, our conclusions are given in Sect. \\ref{sec:conclusion}.  In a companion paper (Cervi\\~no et al. in preparation, hereafter Paper II), we will theoretically evaluate the observational error budget of the different ways to obtain SBF. ",
        "conclusions": "\\label{sec:conclusion}   In this work we have investigated the nature of SBF. We have introduced a distinction between stellar population SBF, theoretical SBF, and observational SBF. Stellar population SBF depend only on the intrinsic properties of stellar populations; theoretical SBF are determined by both the population and the contribution from the distribution of the total number of stars in the used resolution element; and the observational SBF represent the observational counterpart of theoretical SBF. Stellar population SBF and theoretical SBF should be considered in probabilistic terms. On the other hand, observed SBF are statistically distributed and should therefore be interpreted in statistical terms.  In this paper, we have applied the probabilistic description of synthesis models to SBF. We have found that SBF are simply a manifestation of the way in which Nature samples either the integrated luminosity distribution function (pLDF) in integrated systems across different resolution elements, or the stellar luminosity distribution function (sLDF) of resolved systems. We have also found that this definition does not require any assumption about the distribution of stars across different stellar types. Stellar population SBF are therefore simply defined to be the ratio  between the {\\it variance} and the mean of the luminosity {\\it distribution} function in linear fluxes:  \\begin{equation} \\bar{L}_{\\mathrm{SP}} = \\frac{\\kappa_{2}}{\\mu_{1}'} = \\frac{\\kappa_{2}(L_{\\mathrm{tot};N^*})}{\\mu_{1}'(L_{\\mathrm{tot};N^*})}. \\end{equation}   In addition to the definition of stellar population SBF, we have introduced a definition for theoretical SBF, which consists in general of a theoretical description of an observational method. We have shown that the effect of the distribution in the number of stars per resolution element, $N$, is included in theoretical SBF in two ways. First, this distribution affects the computation of the mean flux needed to define the amplitude of renormalized fluctuations for each resolution element. Second, the distribution of mean fluxes also affects  theoretical SBF. Both effects introduce a term that is added to the population contribution. Only for flat galaxy profiles can this term be a simple Poisson contribution; in the case of a different profile, the term differs and must be evaluated explicitly. As a result, theoretical SBF are defined to be:  \\begin{equation} \\bar{L}_{\\mathrm{theo}} = \\frac{\\kappa_{2}}{\\mu_{1}'} + \\mu_{1}'\\,\\frac{\\kappa_2(N)}{\\mu_{1}'(N)} = \\bar{L}_{\\mathrm{SP}} + \\mu_{1}'\\,\\frac{\\kappa_2(N)}{\\mu_{1}'(N)}. \\end{equation}   Observational SBF are an {\\it estimation} of theoretical SBF and must be corrected for the contribution of the variation in the total number of stars per pixel, which depends on both the observational strategy (by means of the way in which the local mean is computed) and the particular galaxy (by means of the galaxy profile), before being compared with stellar population SBF.  Assuming, as a particular case, that the total number of stars per resolution element is Poisson-distributed, we have shown that the corresponding contribution is negligible compared with the contribution of  the stellar population in the break-up of theoretical SBF. Therefore, the common theoretical definition of SBF as \"mean, luminosity-weighted luminosity of the stellar population'' resulting from a Poisson assumption is a restrictive definition that does not necessary apply to the observed SBF, nor any definitions based on an interpretation of SBF as Poisson fluctuations in the number of stars.  We have also characterized the distribution of renormalized fluctuations for different cases in terms of skewness and kurtosis, showing their relation with the skewness and kurtosis of the sLDF. We have shown that the shape of the renormalized fluctuation distributions for a system with $N^*$ stars per resolution element is similar to the shape of the integrated luminosity distribution of clusters with $N^*$ stars, the population luminosity distribution function (pLDF). Taking advantage of the results by \\cite{CLCL06}, we have shown that the renormalized fluctuations of flux across galaxies are not Gaussian, as had been commonly assumed.  Additionally, we have shown alternative ways of compute SBF and that stellar population SBF define an intrinsic metric of fit for the comparison of single observations with synthesis models results. In fact,  {\\it SBF are observationaling evidence that the results of evolutionary synthesis models cannot be interpreted in a deterministic way, but must be interpreted using a probabilistic paradigm.}   Finally, as a caution to theoretical computations of SBF, it is fundamental to ensure that the stellar types considered in the computations are defined in so that the luminosity can be assumed to be constant within each stellar class considered. SBF computed with isochrones that include a detailed treatment of the post-AGB evolution \\citep{Cioni06a,Cioni06b} yield only lower limits to real SBF, since the post-AGB sequence provided by these isochrones are the {\\it average} luminosity during the post-AGB evolution. The variance in this average luminosity must also be known before realistic SBF computations can be obtained based on these isochrones."
    },
    "0809/0809.2432_arXiv.txt": {
        "abstract": "{} {We theoretically and phenomenologically investigate the question whether the \\gr emission from the remnants of the type Ia supernovae SN 1006, Tycho's SN and Kepler's SN can be the result of electron acceleration alone.} {The observed synchrotron spectra of the three remnants are used to determine the average momentum distribution of nonthermal electrons as a function of the assumed magnetic field strength. Then the inverse Compton emission spectrum in the Cosmic Microwave Background photon field is calculated and compared with the existing upper limits for the very high energy \\gr flux from these sources.} {It is shown that the expected interstellar magnetic fields substantially overpredict even these \\gr upper limits. Only rather strongly amplified magnetic fields could be compatible with such low \\gr fluxes. However this would require a strong component of accelerated nuclear particles whose energy density substantially exceeds that of the synchrotron electrons, compatible with existing theoretical acceleration models for nuclear particles and electrons.} {Even though the quantitative arguments are simplistic, they appear to eliminate simplistic phenomenological claims in favor of a inverse Compton \\gr scenario for these sources.} ",
        "introduction": "The question, whether the very high energy (VHE) ($E_\\gamma>100$~GeV) \\gr emission of the Galactic supernova remnants (SNRs) implies a sufficiently large production of nuclear cosmic rays (CRs) -- of the same order as that required to replenish the Galactic CRs -- is one of the key problems addressed by \\gr astronomy. There are two ways to deal with this question in the investigation of an individual SNR.  The first approach is a theoretical one. It uses a nonlinear combination of gas dynamics (or eventually magnetohydrodynamics) for the thermal gas/plasma and kinetic transport theory for the collisionless, nonthermal relativistic particle component that is coupled with the plasma physics of the electromagnetic field fluctuations which scatter these particles. In the environment of the collisionless shock wave of a supernova explosion this allows the description of diffusive shock acceleration of the energetic particles which are originally extracted from the thermal gas and thus {\\it injected} into the acceleration process. Since the field fluctuations are excited by the accelerating particles themselves and since the pressure of these particles (typically comparable with the thermal pressure) is reacting back on the thermal plasma, this strongly coupled system becomes a complex problem of nonlinear dynamics, not only for the charged particle components but also for the electromagnetic field and its fluctuations. Models suggest that a sizeable fraction of the entire hydrodynamic explosion energy will be transformed into energetic particle energy. This suggests that the SNRs are indeed the sources of the Galactic CRs.  Both nuclear charged particles and electrons can be accelerated to achieve nonthermal momentum distributions. The energetic electrons show their presence through synchrotron emission from radio frequencies to hard X-ray energies. They may also interact with diffuse interstellar radiation field photons, like the Cosmic Microwave Background, to produce high energy \\grs in inverse Compton (IC) collisions. The injection of electrons into the acceleration process is however poorly understood quantitatively. Even assuming that the electron momentum distribution at high particle energies is only proportional to the total number of electrons injected per unit time at the shock, the amplitude factor of the electron momentum distribution is therefore not known from theory. It is typically inferred from the measured synchrotron spectrum produced by the accelerated electrons by assuming a mean strength of the magnetic field. This will be a key question in the discussion of this paper.  The injection of nuclear particles from the suprathermal tail of the momentum distribution produced in the dissipative shock transition is much better understood, because the same mechanism that produces scattering fluctuations for higher-energy nuclei -- essentially a beam instability from the accelerated nuclei that is so strong that asymptotically the particles scatter along the mean field direction already after one gyro-period -- also works at injection energies. In addition, where heavy ion injection takes place \\citep{vbk03}, it is to be expected that these ions dominate the nonthermal energy density behind a strong shock like in a young SNR. Then the main nonlinear shock modification consists in a weakening of the quasi-discontinuous part of the shock structure, associated with a broad shock precursor. Low energy particles -- ions and electrons -- in the accelerated spectrum are then only accelerated at this weaker subshock and this implies a significantly softer momentum spectrum at low energies than at high energies. This physical effect is visible in the radio part of the electron synchrotron spectrum and therefore a quantitative indication of the degree of shock modification. It provides a means to determine the injection rate of nuclear particles, de facto of protons, from the radio synchrotron spectrum. In all cases, where the synchrotron spectrum of SNRs was measured, this softening was observed. Together with the nonlinear theory of acceleration, and in the strong scattering limit, this determines the nonthermal pressure $P_\\mathrm{c}$, which turns out to be comparable with the kinetic pressure $\\rho V_\\mathrm{s}^2$ of the gas. Here $ V_\\mathrm{s}$ and $\\rho$ denote the shock velocity and the the upstream mass density, respectively. Using therefore the synchrotron measurement, the nonthermal quantities can be determined from theory.  The exception is at first sight the mean magnetic field strength. However, it needs to be consistent with the {\\it overall form} of the synchrotron spectrum, from radio to X-rays, and with the X-ray synchrotron {\\it morphology} that depends on the effective strength of the magnetic field. In this way an interior effective magnetic field strength is determined. It is typically an order of magnitude larger than the MHD-compressed upstream field strength. This amplification of the magnetic field is a characteristic of the effective acceleration of CR nuclei in a SNR, because it can only be the result of strong acceleration of nuclear particles. The pressure of the accelerated electrons -- also for the cases discussed below -- is more than two orders of magnitude below $\\rho V_\\mathrm{s}^2$.  The question is then, whether the observed \\gr emission is also dominated by nuclear particles through their inelastic, $\\pi^0$ - producing collisions with thermal gas nuclei. This need not be the case if the target density of the thermal gas is very low, despite the fact that the energy density of the accelerated nuclear particle component is very high, in fact comparable to the thermal energy density.  Using this theoretical approach \\citep[for reviews, see e.g.][]{Malkov01,vbk04,ber05,ber08} the investigation of half a dozen of {\\it young} Galactic SNRs has shown that the nuclear CR production is in all cases so high that the Galactic SNRs are viable candidates for the Galactic CR population up to particle energies $\\sim 10^{17}$ eV, well above the so-called knee in the spectrum \\citep{bv07}. Even though important details are open to debate because the time-dependent evolution of a point explosion can only be calculated numerically \\citep{byk96}, we believe that this result is quite a robust one.  However, from a strictly observational point of view, the hadronic nature of most of the SNR \\gr sources is not proven this way. This might ultimately be possible in a direct way with a very sensitive neutrino detector. The remaining question whether the Galactic CR population has a SNR origin then still requires the consistency of the observational result and the theoretical picture.  As far as \\gr observations are concerned, there is also a different approach, basically phenomenological. It considers the question, whether and to which extent the hadronic or leptonic origin of the measured \\gr emission can be decided by favoring either one mechanism at the expense of the other directly from the data. For example, it can ask the question whether the necessarily limited dynamical range of the observed \\gr emission allows a distinction between a hadronic and a leptonic scenario. Or it can ask whether observations in other wavelength ranges tend to empirically contradict the theoretically favored scenario of a predominantly nuclear energetic particle energy density. A possible topic consists in the interpretation of spatial correlations in resolved \\gr SNRs, like those noted in RX~J1713.7-3946 \\citep{aha06} and RX~J0852.0-4622 (Vela Jr.) \\citep{aha07}. The correlation of the hard X-ray synchrotron emission with the VHE \\gr emission features might be considered to favor energetic electrons to produce both emissions.  Discussions of the above and similar issues have recently been given for instance in \\citet{aha06,port07,aha07,Katz08,Plaga08}, and \\citet{bv08}. However, the complexity of the configurations that characterize these extended sources introduces severe uncertainties. They arise from the poorly known structure of the circumstellar medium, which could be partly due to the strong winds expected from the progenitor stars or could be partly pre-existing in the form of neighboring interstellar clouds, affected by the progenitor and its subsequent explosion.  We shall add in this paper such a phenomenological argument. It concerns the spatially integrated synchrotron emission spectrum for the simplest available objects, the remnants of the three young type Ia SNe, observed in VHE $\\gamma$-rays. Even though only upper limits exist from the HEGRA, \\hess \\, and CANGAROO experiments for SN 1006 \\citep{aha05}, Tycho's SNR \\citep{aha01}, and Kepler's SNR \\citep{enomoto08,aha08}, they can nevertheless be used to estimate lower limits to the effective mean magnetic field strengths in the SNR that are consistent with the observed spatially-integrated synchrotron spectra. These somewhat naively estimated magnetic fields are then compared to the expectations for these types of SN explosions. The large discrepancies found disfavor leptonic scenarios for these objects. ",
        "conclusions": "Simple one box approximations indicate that a leptonic scenario for the \\gr emission from the three known Galactic type Ia SNRs SN 1006, Tycho's SNR and Kepler's SNR significantly overpredicts the \\gr flux, even when compared to the existing upper limits from observations. The calculation makes direct use of the observed synchrotron emission spectra. Even though the arguments are simplistic, they appear to eliminate equally simplistic phenomenological arguments in favor of such a scenario. Any positive argument in favor of a purely leptonic scenario would therefore have to be based on a full solution of the governing nonlinear equations. From our results, however, we believe that such a positive argument can not be made."
    },
    "0809/0809.2604_arXiv.txt": {
        "abstract": "We construct domain walls and instantons in a class of models with coupled scalar fields, determining, in agreement with previous studies, that many such solutions contain naked timelike singularities. Vacuum bubble solutions of this type do not contain a region of true vacuum, obstructing the ability of eternal inflation to populate other vacua. We determine a criterion that potentials must satisfy to avoid the existence of such singularities, and show that many domain wall solutions in Type IIB string theory are singular. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0601_arXiv.txt": {
        "abstract": "We present an analysis of the chemical abundances of the star Tycho G in the direction of the remnant of supernova (SN) 1572, based on Keck high-resolution optical spectra. The stellar parameters of this star are found to be those of a G-type subgiant with $T_{\\mathrm{eff}} = 5900 \\pm 100$ K, \\loggl\\ $ = 3.85 \\pm 0.30$ dex, and $\\mathrm{[Fe/H]} = -0.05 \\pm 0.09$.  This determination agrees with the stellar parameters derived for the star in a previous survey for the possible companion star of SN 1572 (Ruiz-Lapuente et al. 2004).  The chemical abundances follow the Galactic trends, except for Ni, which is overabundant relative to Fe, $[{\\rm Ni/Fe}] $ $=$ 0.16 $\\pm$ 0.04. Co is slightly overabundant (at a low significance level). These enhancements in Fe-peak elements could have originated from pollution by the supernova ejecta. We find a surprisingly high Li abundance for a star that has evolved away from the main sequence. We discuss these findings in the context of companion stars of supernovae. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are the best known cosmological distance indicators at high redshifts. Their use led to the discovery of the currently accelerating expansion of the Universe (Riess et al. 1998; Perlmutter et al. 1999); see Filippenko (2005) for a review. They were also used to reveal an early era of deceleration, up through about 9 billion years after the big bang (Riess et al. 2004, 2007). Larger, higher-quality samples of SNe~Ia, together with other data, are now providing increasingly accurate and precise measurements of the dark energy equation-of-state parameter, $w = P/(\\rho c2)$ (e.g., Astier et al. 2006; Riess et al. 2007; Wood-Vasey et al. 2007; Kowalski et al. 2008).  Though the increase in the empirical knowledge of SNe~Ia has led to an enormous advance in their cosmological use, the understanding of the explosion mechanism still requires careful evaluation (e.g., Hillebrandt \\& Niemeyer 2000; Wheeler 2007).  While in Type II SNe we have the advantage that the explosion leaves a compact star to which surviving companions (if they exist) often remain bound, thus enabling a large number of studies (e.g., Mart{\\'\\i}n et al. 1992; Israelian et al. 1999), in SNe~Ia the explosion almost certainly does not produce any compact object.  To investigate how the explosion takes place, we examine the rates of SNe~Ia at high redshifts, or we can take a more direct approach and survey the field of historical SNe Ia. The latter strategy has been followed since 1997 in a collaboration that used the observatories at La Palma, Lick, and Keck (Ruiz-Lapuente et al. 2004). The stars appearing within the 15$\\%$ innermost area of the remnant of SN 1572 ($0.65'$) were observed both photometrically and spectroscopically at multiple epochs over seven years. The proper motions of these stars were measured as well using images obtained with the WFPC2 on board the {\\it Hubble Space Telescope (HST)} (GO--9729). The results revealed that many properties of any surviving companion star of SN 1572 were unlike those predicted by hydrodynamical models. For example, the surviving companion star (if present) could not be an overluminous object nor a blue star, as there were none in that field. Red giants and He stars were also discarded as possible companions.  A subgiant star (G2~IV) with metallicity close to solar and moving at high speed for its distance was suggested as the likely surviving companion of the exploding white dwarf (WD) that produced SN 1572 (Ruiz-Lapuente et al. 2004). This star, denoted Tycho G, has coordinates $\\alpha = 00^{\\rm h}25^{\\rm m}23.7^{\\rm s}$ and $\\delta = +64^\\circ08'02''$ (J2000.0).  Comparisons with a Galactic model showed that the probability of finding a rapidly moving subgiant with solar metallicity at a location compatible with the distance of SN 1572 was very low. A viable scenario for the Tycho SN 1572 progenitor would be a system resembling the recurrent nova U~Scorpii. This system contains a  white dwarf of $M_{\\rm WD} = 1.55 \\pm 0.24$~\\Msun and a secondary star with $M_2 = 0.88 \\pm 0.17$~\\Msun orbiting with a period $P_{\\rm orb} \\approx 1.23$~d (Thoroughgood et al. 2001).  On the other hand, based on analysis of low-resolution spectra, Ihara et al. (2007) recently claimed that the spectral type of Tycho G is not G2~IV as found by Ruiz-Lapuente et al. (2004), but rather F8~V. In addition, Ihara et al. (2007) suggested the star Tycho~E as the companion star of the Tycho SN 1572, but we consider their conclusion to be unjustified. They base it on a single \\ion{Fe}{1} feature at 3720~{\\AA} in a spectrum having a signal-to-noise ratio (SNR) of only $\\sim 13$. The short spectral range of their data (3600--4200~{\\AA}) and the low resolution ($\\lambda/\\Delta\\lambda \\approx 400$, with a dispersion of 5~{\\AA}/pixel) add to our concern.  The present work resolves questions regarding the metallicity and spectral classification of Tycho G. We concentrate on providing both the chemical composition of the star and the stellar parameters as derived from new data. In addition, we present a detailed chemical analysis of Tycho G aimed at examining the possible pollution of the companion star by the supernova. Although more data are still required to definitively settle this issue, the overall properties of this star remain consistent with being the surviving companion of SN 1572. ",
        "conclusions": "\\subsection{Has Tycho G been Polluted by the SN Ia?\\label{polluted}}  The measured ratio [Ni/Fe] $= 0.16 \\pm 0.04$ in Tycho G is 4--5$\\sigma$ above the average value of this ratio in the Galaxy ([Ni/Fe] $\\approx -0.05$). It indicates an overabundance of Ni with respect to the Galactic trend. Ni may have been trapped by the star as slowly moving material from deep layers of the SN ejecta orbited around it. While the rapidly moving ejecta consisting mainly of intermediate-mass elements would have escaped away from the site of the explosion, the very slowly moving tail made of heavy elements like Ni could have been captured by the companion.  To be consistent with the evolutionary path suggested above (and also by Ruiz-Lapuente et al. 2004), we use $a = 19.28$ R$_{\\odot}$ for the separation between the white dwarf and its companion, and $R_{*}= R_{L} = 6.75$ R$_{\\odot}$ for the radius of the companion (just before the SN explosion).  The mass that could be trapped by the companion is given by the subtended angle,  \\begin{equation} m_{\\rm t} = \\Delta M  \\ (\\pi R_{2}^{2}/4\\pi \\ a^{2})  \\ f, \\end{equation}  \\noindent where $\\Delta M$ is  the ejected mass and $f$ is the fraction of retained material.  \\noindent  We try to model the contamination of the companion star by the nucleosynthetic products of different SN~Ia models from Iwamoto et al. (1999). In Table~\\ref{tbl4} we show the expected abundances of Tycho G after having captured a significant amount of the SN ejecta. In these model computations, we adopt the yields from SN~Ia models with different kinetic explosion energies in the range (1.30--$1.44)\\,\\times\\,10^{51}$ ergs, different central densities (in units of $109$ g~cm$^{-3}$) of $\\rho_9=1.37$ (C) and 2.12 (W), and different deflagration velocities with fast deflagration in the models W7 and W70 (Nomoto, Thielemann, \\& Yokoi 1984) and slow deflagration in the others. We assumed as initial abundances the average abundances of disk stars with a metallicity of ${\\rm [Fe/H]} = -0.20\\pm 0.09$, and a capture efficiency factor $f=0.03$. The matter that is captured by the companion has a much larger mean molecular weight than the composition of its atmosphere and then is completely mixed with the whole star due to thermohaline mixing, in a short timescale (Podsiadlowski et al. 2002).  Badenes et al. (2006) found, by comparing a grid of X-ray synthetic spectra based on hydrodynamical models, that the fundamental properties of X-ray emission in Tycho SN 1572 are well reproduced with one-dimensional delayed detonation models having a kinetic energy of $\\sim1.2\\,\\times\\,10^{51}$ ergs. In Table~\\ref{tbl4}, the models called WDD1,2,3 and CDD1,2 include transitions from slow deflagration to detonation, which are equivalent to delayed detonation models.  In Figure~\\ref{snfig}, we display the observed abundances of Tycho G in comparison with the expected abundances from several SN~Ia models. Except for Na and Al, we find reasonable agreement between the observed abundances of Tycho G and the expected abundances after contamination from the nucleosynthetic products of the SN~Ia. Surprinsingly, Gonz\\'alez Hern\\'andez et al. (2008a) measured the chemical abundances of the secondary star of the black hole binary Nova Scorpii 1994 and found several $\\alpha$ elements significantly enhanced in the companion star while Al and Na were not enhanced. The enhanced $\\alpha$ elements in that star were explained by comparing the observed abundances with supernova/hypernova explosion models. However, these models also predict that Al and Na should be enhanced at a comparable level as oxygen, and they did not find an explanation for this discrepancy. Therefore, the model of contamination of Tycho G by using SN~Ia models could provide an explanation for the relatively high observed [Ni/Fe] with respect to the Galactic trend of this element.  \\begin{figure}[!ht] \\centering \\includegraphics[width=8cm,angle=0]{f11_color.eps} \\caption{Expected abundances in Tycho G after contamination from the nucleosynthetic products in SN~Ia models (Iwamoto et al. 1999), in comparison with the observed abundances.} \\label{snfig} \\end{figure}  One could find stable isotopes of Co as well among the slowly moving material in a SN~Ia. The ratio [Co/Fe] $ = 0.18 \\pm 0.08$ in Tycho G is consistent with contamination from the SN ejecta. The study by Reddy et al. (2006) finds that [Co/Fe] shows flat behavior with [Fe/H] for thin-disk stars ([Co/Fe] $ = -0.05 \\pm 0.02$). For thick-disk stars, a weak trend of decreasing [Co/Fe] with increasing [Fe/H] is apparent, falling to [Co/Fe] $ = 0.00 \\pm 0.04$ at the metallicity of Tycho G. The effect of a larger [Co/Fe] is $>1\\sigma$, a lower significance level compared to the ratio [Ni/Fe] (see also Fig. 6).  The ratios of the $\\alpha$ elements (O, Si, Ca, and Ti) to iron do not show enhancements over Galactic values.  \\begin{figure*}[!ht] \\centering \\includegraphics[width=15cm,angle=0]{f12.eps} \\caption{Li abundance of Tycho G in comparison with that of subgiant stars from Randich et al. (1999) having similar stellar parameters, metallicity, and rotational velocity.} \\label{lifig} \\end{figure*}  A study of the SN 1572 remnant in X-rays (Hamilton et al. 1985) suggests that it contains $\\sim$0.66 \\Msuno, including 0.20 \\Msun of Fe and 0.02 \\Msun of Ni in an inner layer as well as $\\sim$0.44 \\Msun of Fe plus Ni unshocked by the reverse shock and consequently moving in free expansion. The energy of this supernova is estimated from the ejecta to be $E \\approx$ (4--5) $\\times 10^{50}$ ergs (Hughes 2000).  \\subsection{Li Abundance\\label{licap}}  As shown in Figure~\\ref{fig3}, the Li line at 6708~\\AA\\ is pronounced in Tycho G.  The abundance of this element, $A$(Li) = $2.50 \\pm 0.09$, raises the question of how much Li could have survived in this star that has already evolved off the main sequence. Similar abundances have been found in some turn-off stars (with $6300 < T_{\\rm eff} < 6800$~K) in the sample of do Nascimento et al. (2000, 2003), and in a few subgiant stars with stellar parameters similar to those of Tycho G. In the sample of Randich et al. (1999), there are also a few subgiants with similar stellar parameters and Li abundance as Tycho G; see Figure~\\ref{lifig}. Although the Li abundance of Tycho G is somewhat high for its relatively low effective temperature and surface gravity, the upper-right panel of Figure~\\ref{lifig} shows two stars (with \\teff $\\approx 5500$ K and 6000~K, and \\loggl\\ $\\approx 3.3$ dex) that are even more anomalous than Tycho G.  Canto Martins et al. (2006) found a Li-rich subgiant star (S1242) in the open cluster M67, with \\teff $=5800$~K, \\loggl\\ $=3.9$ dex, [Fe/H] $= -0.05$, and $A$(Li) = 2.7, all similar to the properties of Tycho G. This star is in a wide eccentric binary system, and these authors proposed that the Li has been preserved due to tidal effects. There exist other subgiant stars of similar type showing high Li abundances and belonging to binary systems. Again, this could be linked to being in a close binary system where tidal interaction might inhibit Li destruction.  An alternative is that Li was created as a result of energetic processes. Mart{\\'\\i}n et al. (1992, 1994a,b) suggest that in soft X-ray transients (SXTs) one is observing freshly synthesized Li, and that the Li could come from $\\alpha \\alpha$ and/or other spallation reactions during the outbursts of these binary systems that contain black holes or neutron stars.  Spallation reactions are produced by protons and $\\alpha$ particles with energies above a few MeV when hitting C, N, and O nuclei, and also by $\\alpha + \\alpha$ collisions at similar energies.  In addition, in the case where the compact object is a white dwarf, Mart{\\'\\i}n et al. (1995) have shown that their companion stars do not show high Li abundances. They claimed that this is consistent with the Li production scenario in SXTs, because the weaker gravitational potential well of the white dwarfs does not allow particles to be accelerated to the required high energies (Mart{\\'\\i}n, Spruit, \\& van Paradijs 1994b). However, other authors have shown that Li production is possible during nova explosions (Starrfield et al. 1978; Boffin, Paulus, \\& Arnould 1993).  The energy arguments could also work for the enhancement of Li in the companion of a carbon-oxygen white dwarf. The accretion lasts for a sufficiently long time to build enough Li nuclei in the convective envelope of the secondary star. At the onset of the explosion, the supernova is able to unbind part of the material from the convective envelope. If part of the convective envelope were to survive, it would display a high Li abundance.  However, this scenario of Li production during outbursts in SXTs also predicts Li isotopic ratios in the range $0.1 < N(^6$Li)/$N(^7$Li$) < 10$; see Casares et al. (2007), and references therein. These authors conclude that the preservation scenario by tidal effects is the most likely mechanism to explain the high Li abundances of SXTs.  There are two other alternatives for enrichment of Li in the companions of SNe~Ia. The first is that the high-energy particles accelerated by the shock wave in the outermost layers of the exploding white dwarf induce spallation reactions in the surface of the companion. A fraction of the irradiated material could mix with the layers now making the surface of the companion, and Li enrichment would show.  The second alternative is that the surface of the companion, after the explosion, can be bombarded by high-energy particles accelerated in the SN remnant, producing spallation reactions. Steady bombardment by those high-energy particles has the restriction that while it enriches Li, the energy deposited would increase the star's luminosity.  However, if Li enrichment is limited to a thin layer extending not much below the photosphere, the increase in luminosity would be negligible."
    },
    "0809/0809.0615_arXiv.txt": {
        "abstract": "We present a  measure of the inclination of the  velocity ellipsoid at 1~kpc below the  Galactic plane using a sample  of red clump giants  from the RAVE DR2  release.  We find  that the  velocity ellipsoid  is tilted  towards the Galactic plane with an inclination of $7.3\\pm1.8^{\\circ}$. \\noindent We compare this value to computed inclinations for two mass models of the Milky  Way.  We find that our measurement is  consistent with a short scale  length of the  stellar disc  ($R_d\\simeq2$~kpc) if  the dark  halo is oblate or  with a long scale  length ($R_d\\simeq3$~kpc) if the  dark halo is prolate. \\noindent Once  combined with independent  constraints on the  flattening of the halo,  our measurement suggests  that the scale length  is approximately halfway  between  these  two  extreme   values  ,  with  a  preferred  range [$2.5$-$2.7$]~kpc for a nearly spherical halo. Nevertheless, no model can be clearly ruled  out.  With the  continuation of the  RAVE survey, it  will be possible  to provide a  strong constraint  on the  mass distribution  of the Milky  Way using  refined measurements  of the  orientation of  the velocity ellipsoid. ",
        "introduction": "Our  understanding  of Galactic  stellar  populations  and kinematics  makes regular  progress with  the advent  of  new large  Galactic stellar  surveys providing distances,  photometry, radial velocities or  proper motions.  Our Galaxy is  at the present  the only  place where we  can probe the  6D phase space of  stellar positions and  velocities.  For instance, the  Galactic 3D potential  can  be probed  through  the  orbits  of the  Sagittarius  stream \\citep{ib01,rm05,newberg02,fel06,helmi04}   or   Palomar   5   tidal   tails \\citep{od03,gd06,gj06},   or   through   the   kinematics  of   halo   stars \\citep{ba05}.  At smaller scales, the potential can also be analysed through the force perpendicular  to the galactic plane (Oort  1960; Cr\\'ez\\'e et al. 1998; Kuijken \\&  Gilmore 1989a,b,c,1991; Siebert et al.   2003; Holmberg \\& Flynn 2004) or through the coupling between the 3 components of the velocity in the solar neighbourhood \\citep{bien99}.\\\\  Here,  we concentrate on  the question  of the  orientation of  the velocity ellipsoid that is  known to be tightly related to the  shape and symmetry of the galactic potential \\citep{oll62,hor63,lyn62,ac91}.  In spite of the long interest in this problem, measuring observationally the orientation  of the  velocity ellipsoid  outside of  the galactic  plane has proven to be  very difficult. This is due mainly to  the absence of reliable distances away  from the Solar neighbourhood.  Despite  this limitation, the first stellar  stream detected within the  Milky Way halo  towards the north Galactic pole  by \\citet{maj96} shows  a velocity tilt, the  ellipsoid being inclined  towards  the Galactic  plane.  This  tilt  could result  from  the expected  velocity correlation induced  by a  spheroidal potential  if these stars had similar  integrals of motion \\citep{bien98}. However  we note that this stream is not detected locally in the RAVE data \\citep{seabroke08}.  Building  realistic Galactic  potentials shows  that  the main  axis of  the velocity  ellipsoid, at  1\\,kpc  above  the Galactic  plane,  points in  the direction of the $z$-axis of symmetry  of the Galaxy towards a point located at 5 to 8 kpc behind  the Galactic centre: for instance from numerical orbit computations  \\citep{bin83,kg89a} or applying  to the  \\citet{ci87} Galactic potential  the  \\citet{ac91}  formulae.   Such  estimates  of  the  velocity ellipsoid tilt are necessary for an accurate determination of the asymmetric drift\\footnote{The asymmetric drift is the tendency of a population of stars to lag behind the local standard  of rest for its rotational velocity, the lag increasing  as a  function of age.}   \\citep{bt87}, and for  a correct measurement   of   the   force   perpendicular   to   the   Galactic   plane \\citep{sta89}.\\\\   In this  paper, we study the  2D velocity distribution  perpendicular to the Galactic  plane  for a  sample  of  red clump  stars  from  the RAVE  survey \\citep{dr1,dr2}.   These stars  are  selected between  500\\,pc and  1500\\,pc below  the Galactic  plane and  provide  a measurement  of the  tilt of  the velocity ellipsoid at  $\\simeq$1~kpc.  In Section~\\ref{s:sample}, we present the  selection  of the  sample  while  Section~\\ref{s:tilt}  focuses on  the measurement   of  the   inclination  and   possible  biases.    Finally,  in Section~\\ref{s:model}  we compare our  measurement to  computed inclinations for two extreme  classes of mass models and we  discuss possible outcomes of this measurement. ",
        "conclusions": "We measured  the tilt of the  velocity ellipsoid at  $\\simeq$1~kpc below the Galactic  plane  using a  sample  of  red clump  giants  from  the RAVE  DR2 catalogue. We find its inclination to be 7.3$\\pm$1.8$^{\\circ}$. Estimates of the effect of contamination by  foreground stars and substructures have been shown to be small and their effect on our measured value can be neglected.  We compared this value to predictions  from two extreme cases of mass models for the  Milky Way proposed by  BT08. In the case  of a massive  disc with a small  scale length  ($R_d=2$~kpc), the  inclination is  compatible  with an oblate halo whose  minor--to--major axis ratio $c/a_\\rho$ is  lower than 0.9 at the \\mbox{1--$\\sigma$} level. On the other hand, in the case of a massive halo  with large disc  scale length  ($R_d\\simeq3$~kpc), prolate  haloes are prefered  with $c/a_\\rho \\geq  1$. When  a direct  measurement is  used, low values  for  $c/a_{\\rho}$  can   be  marginally  rejected,  indicating  that $c/a_{\\rho}>0.7$.  When further  independent constraints from previous  studies are considered, we   find  that   an  intermediate   value   for  the   disc  scale   length $R_d\\simeq[2.5-2.7]$~kpc is  preferred for a  nearly spherical halo,  but no extreme model  can be clearly ruled out,  due to our large  error bars. This range is in good agreement with other studies relying on star count analysis and deep  photometric surveys. Nevertheless  these results have  large error bars of  the same  order as our  measurement and  cannot be used  to further constrain the mass distribution. \\\\  RAVE continues  to acquire spectra and  this work relies on  the second data release of the survey. So far  RAVE has collected more than 200~000 spectra, 4 times the  size of the sample used here. With  the current observing rate, we can expect to multiply by ten  the size of our sample in the coming years which will allow  us to significantly reduce our error bars.   By the end of the survey, we  will be able to provide  a new mass model for  the Milky Way galaxy with a constrained scale length of the disc and minor--to--major axis ratio of the dark halo."
    },
    "0809/0809.0109_arXiv.txt": {
        "abstract": "Suzaku observation of the edge-on spiral galaxy NGC 4631 confirmed its X-ray halo extending out to about 10 kpc from the galactic disk. The XIS spectra yielded the temperature and metal abundance for the disk and the halo regions. The observed abundance pattern for O, Ne, Mg, Si and Fe is consistent with the metal yield from type II supernovae, with an O mass of about $10^{6} M_\\odot$ contained in the halo. These features imply that metal-rich gas produced by type II supernova is brought into the halo region very effectively, most likely through a galactic wind. Temperature and metal abundance may be affected by charge exchange and dust. An upper limit for the hard X-ray flux was obtained, corresponding to a magnetic field higher than $0.5 \\mu$G\\@. ",
        "introduction": "For the understanding of the chemical evolution of galaxies and clusters of galaxies, precise knowledge about the metal production from Type Ia (SN Ia) and Type II (SN II) supernovae (SNe) is of vital importance. X-ray imaging spectroscopy of supernova remnants (SNRs), galaxies and intra-cluster medium  (ICM) has shown that metal abundances generally vary from source to source. However, \\citet{sato07a} determined the abundances of O, Ne, Mg, Si, and Fe in several galaxy clusters based on  Suzaku observations, and found  that the abundance patterns are commonly represented by a combination of type Ia and II supernovae  products with an occurrence number ratio of $1:3.5$, based on the theoretical yields from SN II and SN Ia. In order to further look into the past history of the two types of supernova, we need more precise knowledge about the metal yields from the different SNe types. It is, however, quite difficult to extract pure supernova products, in particular for the X-ray emitting hot gas. When we observe young supernova remnants (SNR), there is always a mixture of supernova ejecta and surrounding ISM and line intensities also strongly depend on the ionization condition. One way to provide a good constraint is to observe X-ray halos around starburst galaxies, which are considered to be maintained by the enhanced SNe II activity in the recent time.  Recent X-ray observations of starforming galaxies showed that metals are contained in the extended hot halo of the galaxies. For M82 and NGC 253, RGS spectra for several sliced regions along the outflow axis showed   lines from highly ionized O, Ne, Mg, Si and Fe \\citep{read02,bauer07}. The observed intensity ratios of the lines indicates that the gas is  cooling as it travels outward from the galaxy disk and that the gas around NGC 253 could be partly out of ionization equilibrium. Suzaku observations of a ``cap'' region of M82, which is 11.6 kpc north of the galaxy and is possibly a termination region of the hot-gas outflow, showed a spectrum consisting of emission lines from O through Fe \\citep{tsuru07}. These spectral features strongly suggest that fresh metal-rich gas produced in the starforming region is flowing out mainly along the minor axis of galaxies. However, metal abundances in the halo gas has so far not been well-constrained. RGS data are  limited in statistics, and both the EPIC and the ACIS instruments do  not have sufficient energy resolution below 1 keV\\@ (e.g.\\ \\cite{tuellmann06}). Suzaku offers a good opportunity for measuring the metallicity of  the outflowing hot gas.   NGC~4631 is a nearby Sc/SBd galaxy with  an edge-on morphology with the distance estimated to be about 7.5 Mpc, where $1'$ corresponds to 2.2~kpc. Estimated mass by the Tully-Fisher relation is $2.6\\times 10^{10} M_{\\odot}$ \\citep{strickland04}. The inclination and position angle are $81^\\circ$ and $356^\\circ$ respectively. This galaxy is suitable for Suzaku observations of  the X-ray halo maintained by its SN activity. With its radio halo \\citep{hummel90} and warm IR ratio, it is classified as a mild disk-wide starburst galaxy \\citep{golla94b}. An extended X-ray halo was  discovered by  ROSAT  \\citep{wang95}, and it has been well studied with  Chandra \\citep{wang01,oshima03,strickland04} and XMM-Newton \\citep{tuellmann06}. The size of the halo is several arcmin  and no X-ray counterpart for the central AGN has been  detected \\citep{strickland04}. The association of an H$\\alpha$ filament with the X-ray emission has been   discovered \\citep{wang01}, and FUSE also detected  O\\emissiontype{VI} lines from a region at $2'$ above the disk \\citep{otte03}. These strongly suggest that a halo around NGC 4631 is the site of galactic outflow or fountain, where the gas is floating up from the disk by the SN energy input and possibly cooling.  Another important feature of the halo is the extended synchrotron radio emission, which is  observed from several star-forming galaxies (e.g.\\ \\cite{veilleux05}). These radio halos are populated with relativistic electrons together with magnetic fields of an order of $10 \\mu$G\\@. The radial orientation of the magnetic field lines suggests that the field has been carried into the halo region by the outflow of hot gas \\citep{golla94a}. In some cases, the hot gas may be confined in the halo when the magnetic pressure exceeds the thermal pressure  \\citep{wang95}. To constrain how the non-thermal energy is distributed in the halo, it is important to look into the possibility of  non-thermal emission in the hard X-ray band.   Throughout this paper we adopt the Galactic hydrogen column density of $N_{\\rm H} = 1.27\\times10^{20}$ cm$^{-2}$ \\citep{dickey90} in the direction of NGC~4631\\@. The solar abundance table is given by \\citet{anders89}, and the errors are the 90\\% confidence limits for a single interesting parameter. ",
        "conclusions": "We determined the temperature and metal abundance of the X-ray emitting halo gas around NGC 4631. The total energy, mass, and metals in the halo can by supplied by SNe with the currently estimated SFR, if the outflow efficiently carries metal-rich gas from the starforming regions into the halo. The effect of neutral material and dust should be taken into account to understand the plasma properties in the halo.  \\begin{figure} \\FigureFile(\\columnwidth,\\columnwidth){figure4.eps} \\caption{Number ratios of Ne, Mg, Si, S, and Fe to O for disk and halo regions.  Solid, dotted, dashed, and dot-dashed lines correspond to the number ratios of metals to O for  abundance patterns  of SNe II yield of \\citet{nomoto06}, solar abundance by \\citet{anders89}, cluster average in \\citet{sato07a}, and  SNe Ia yield of \\citet{iwamoto99}, respectively.} \\label{fig:ratio} \\end{figure}"
    },
    "0809/0809.5030_arXiv.txt": {
        "abstract": "\\noindent The gravitino is a promising supersymmetric dark matter candidate which does not require exact $R$-parity conservation. In fact, even with some small $R$-parity breaking, gravitinos are sufficiently long-lived to constitute the dark matter of the Universe, while yielding a cosmological scenario consistent with primordial nucleosynthesis and the high reheating temperature required for thermal leptogenesis. In this paper, we compute the neutrino flux from direct gravitino decay and gauge boson fragmentation in a simple scenario with bilinear $R$-parity breaking. Our choice of parameters is motivated by a proposed interpretation of anomalies in the extragalactic gamma-ray spectrum and the positron fraction in terms of gravitino dark matter decay. We find that the generated neutrino flux is compatible with present measurements. We also discuss the possibility of detecting these neutrinos in present and future experiments and conclude that it is a challenging task. However, if detected, this distinctive signal might bring significant support to the scenario of gravitinos as decaying dark matter. ",
        "introduction": "The question of the nature of dark matter is still one of the unsolved mysteries in modern cosmology. Many particle candidates have been put forward, but until now the only dark matter evidence we have is based on the gravitational interaction. In the context of supersymmetry with conserved $R$-parity, one naturally encounters one of the most favoured solutions to the dark matter problem. There, the lightest supersymmetric particle (LSP) is stable and can be a successful dark matter candidate if it is neutral and weakly interacting like the neutralino. The neutralino is the most thoroughly studied dark matter candidate and will be tested  in the near future in accelerator, direct detection and indirect detection experiments~\\cite{Bertone:2004pz}.  On the other hand, it is also possible that the dark matter interacts only gravitationally, and supersymmetry also offers candidates of this type. A prominent example is the gravitino, the superpartner of the graviton, which was the first supersymmetric dark matter candidate proposed~\\cite{Pagels-Primack}. It is one of the most elusive dark matter candidates due to its extremely weak interactions. In fact, as part of the gravity multiplet, all gravitino interactions are suppressed either by the Planck scale (for the spin-3/2 component) or by the supersymmetry-breaking scale (for the Goldstino component).  Usually, having an LSP with such extremely weak interactions poses a severe problem to Big Bang nucleosynthesis, as it makes the next-to-lightest supersymmetric particle (NLSP) so long-lived that it decays during or after the formation of the primordial nuclei, typically spoiling the successful predictions of the standard scenario~\\cite{BBN-gravitino}. Moreover, if the NLSP is charged, the formation of a bound state with $^4$He catalyzes the production of $^6$Li~\\cite{Pospelov:2006sc} leading to an overproduction of $^6$Li by a factor 300-600~\\cite{Hamaguchi:2007mp}. One way to avoid these constraints is to lower the scale of supersymmetry breaking, thus enhancing the Goldstino interactions. However, the reheating temperature of the Universe has to be lowered accordingly in order to avoid overclosure~\\cite{gravitino-TH} and is then in conflict with the minimal value required by thermal leptogenesis in order to explain the baryon asymmetry of the Universe~\\cite{lepto-TH}.  It was recently proposed in~\\cite{bchiy} that these problems are automatically solved if a small breaking of $R$-parity is introduced in the model. Even though in the presence of $R$-parity violation the neutralino LSP is too short-lived to play the role of dark matter, the gravitino LSP can still have a sufficiently long lifetime, which is typically many orders of magnitude greater than the age of the Universe due to the suppression of the decay rate by the Planck scale and the small $R$-parity violating couplings~\\cite{ty00}. In this scenario, the NLSP population in the early Universe quickly decays into Standard Model particles via $R$-parity violating interactions. Apart from the presence of a rather inert population of gravitinos, produced thermally from superpartner scatterings at reheating, the cosmology then reduces to the non-supersymmetric case well before the synthesis of primordial nuclei.  An intriguing feature of the scenario with $R$-parity breaking is that gravitino dark matter is not necessarily invisible anymore, since it will decay into Standard Model particles at a very slow rate. Since the huge number of gravitinos in our own Galaxy, as well as in nearby galaxies and clusters, may compensate for the  highly suppressed decay rate, this opens up the possibility of observing the dark matter decay products as an anomalous contribution to the diffuse gamma-ray flux~\\cite{bbci07,it07} or the cosmic-ray antimatter fluxes~\\cite{it08,imm08}. Indeed, anomalous excesses have been observed both in the diffuse extragalactic gamma-ray spectrum and in the positron fraction in similar energy ranges. It has been pointed out that the decay of gravitino dark matter with a lifetime of $\\sim 10^{26} \\usk \\unit{s}$ and a mass of $\\sim 150 \\usk \\unit{GeV}$ can account for both of these excesses at the same time~\\cite{it08,imm08}. This motivates our study of the corresponding neutrino spectrum for the same choice of parameters, both as a consistency check and to find out whether in this scenario an anomalous contribution to the neutrino flux may be expected in present and future neutrino experiments.  This paper is organised as follows: in the next section we will briefly review bilinear $R$-parity violating models and discuss the resulting gravitino decay modes. In Section~\\ref{Spectrum} we will present the neutrino spectrum from dark matter decay and describe its main features. In Section~\\ref{Fluxes} we will then give the neutrino flux as a function of the gravitino lifetime and mass both for neutrinos from our own Galaxy and from diffuse extragalactic sources and consider the effect of neutrino oscillations on the signal expected at the Earth. In Section~\\ref{Backgrounds} we will discuss the different neutrino backgrounds in the energy range we are interested in and compare them to our signal. In Section~\\ref{Detection} we will propose strategies to disentangle the signal from the background, compare the result to present neutrino data from Super-Kamiokande. We will then discuss the feasibility of detection in future detectors in Section~\\ref{Future} and conclude in Section~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  We have examined the neutrino spectrum from the decay of unstable gravitino dark matter in a scenario with bilinear $R$-parity violation. It has been pointed out in the recent literature that the decay of gravitino dark matter particles with a lifetime of $\\sim 10^{26}$\\usk s and a mass of $\\sim 150 \\usk \\unit{GeV}$ into massive gauge bosons may account for the anomalies observed in the diffuse extragalactic gamma-ray spectrum as measured by EGRET as well as in the positron fraction as measured by HEAT.\\footnote{The existence of a positron excess seems to be supported by preliminary results from PAMELA~\\cite{pamela08}.} Motivated by this observation, we have computed the neutrino spectrum for the same choice of parameters as a consistency check of this scenario. We find that this spectrum is compatible with results from neutrino experiments.  We have also examined the detectability of this exotic component of the neutrino flux to find an independent way to test this scenario. While the signal in the neutrino spectrum with two or more distinct peaks, resulting from two-body gravitino decays into gauge/Higgs boson and neutrino, is very characteristic, it will be challenging to detect these features in neutrino experiments. On one side, present neutrino detectors do not achieve a sufficiently high energy resolution to resolve the subdominant peaks, and on the other side, the event rate is expected to be so small that the background of atmospheric neutrinos overwhelms the signal in all flavours. The most promising signal-to-background ratio is found in the tau neutrino flavour, especially when analysing only the flux from the upper hemisphere since there the atmospheric tau neutrino flux is vastly reduced. However, tau neutrinos are difficult to identify in Cherenkov detectors and probably only an event-by-event identification procedure could allow the signal to be seen with such extremely limited statistics. At present, therefore, it is not possible to detect this contribution due to technological limitations.  The ideal detector for testing the present scenario would be one of megaton mass with the ability to identify and measure tau neutrinos event by event. Should such a detector ever become available, it could be worthwhile to look for this component of the neutrino flux by employing strategies for background reduction such as the ones discussed here, especially if the anomalous signatures in the positron fraction and the diffuse extragalactic gamma-ray spectrum are confirmed by PAMELA and FGST, respectively. The detection of a signal in neutrinos compatible with signals in the other indirect detection channels would in fact bring significant support to the scenario of decaying dark matter, possibly consisting of gravitinos that are unstable due to bilinear $R$-parity violation."
    },
    "0809/0809.4493_arXiv.txt": {
        "abstract": "One of the most sought-after signatures of reionization is a rapid increase in the ionizing background (usually measured through the \\lya optical depth toward distant quasars).  Conventional wisdom associates this with the ``overlap\" phase when ionized bubbles merge, allowing each source to affect a much larger volume.  We argue that this picture fails to describe the transition to the post-overlap Universe, where Lyman-limit systems absorb ionizing photons over moderate lengthscales ($\\la 20$--$100 \\Mpc$).  Using an analytic model, we compute the probability distribution of the amplitude of the ionizing background throughout reionization, including both discrete ionized bubbles and Lyman-limit systems (parameterized by an attenuation length).  We show that overlap does \\emph{not} by itself cause a rapid increase in the ionizing background or a rapid decrease in the mean \\lya transmission toward distant quasars.  More detailed semi-numeric models support these conclusions.  We argue that rapid changes should instead be interpreted as evolution in the attenuation length itself, which may or may not be directly related to overlap. ",
        "introduction": "\\label{intro}  In the last several years, the cosmological community has made an enormous effort to measure the reionization history of the intergalactic medium (IGM).  But the picture remains murky: some evidence (principally the cosmic microwave background, or CMB) points toward a mostly-ionized universe at $z \\ga 9$ \\citep{page07, komatsu08, dunkley08}, but other observations (principally from quasar absorption) are usually taken to imply that reionization ends only at $z \\approx 6$ \\citep{fan02, mesinger04, fan06, mesinger07-prox}.  Many other techniques cannot yet distinguish between early and late reionization scenarios (e.g., \\citealt{totani06, kashikawa06, ota07, mcquinn07-lya, mcquinn08-damp}; see also \\citealt{fan06-review} for a review).  While there is not necessarily a contradiction between the CMB and quasar measurements -- reionization may simply be relatively extended, which would hardly be surprising from a theoretical standpoint (e.g., \\citealt{cen03-letter, wyithe03-letter, furl05-double, iliev07-selfreg}) -- the tension does call for a critical examination of the existing evidence.  The most widely-recognized point in favor of late reionization is the apparent rapid decrease in the mean \\lya transmission toward quasars at $z \\ga 6$ \\citep{fan01, fan02, white03, fan06}.  There are two aspects to this.  First, a complete \\citet{gunn65} absorption trough appears toward some quasars, although the enormous optical depth of that line means that this still only requires a small neutral fraction.  Second, there appears to be a substantial steepening of the amount of absorption beyond $z \\sim 6$ \\citep{fan06}.  However, the latter conclusion has been challenged on empirical grounds \\citep{songaila02, songaila04, oh05, becker07}; unfortunately, known lines of sight are sparse enough that no clear resolution has emerged (e.g., \\citealt{lidz06}).  Here, we take a complementary approach and examine the theoretical underpinnings of the conclusion that such a rapid change implies a detection of ``reionization.\"  The crux of this argument is that the moment of ``overlap\" (i.e., the point at which ionized bubbles merge into much larger units, usually considered to be the ``end\" of reionization) must be accompanied by a sudden, rapid increase in the amplitude $\\Gamma$ of the ionization rate.  Qualitatively, this expectation comes from percolation models of the reionization process:  when ionized bubbles (presumed to be transparent to ionizing photons) overlap, many more sources illuminate any given patch, so $\\Gamma$ should increase rapidly.  Quantitatively, the first self-consistent simulation of reionization showed precisely such a jump \\citep{gnedin00}.  In that simulation, $\\Gamma$ evolved much more sluggishly both before overlap (when ionized regions grew slowly because the photons had to ionize fresh material) and afterwards (when the ionizing photons were already able to propagate large distances), so overlap appeared to be clearly defined.  However, this picture does not match properly onto the well-understood post-reionization Universe, where dense ``Lyman-limit systems\" (LLSs) absorb ionizing photons over relatively small distances (e.g., $\\sim 110 \\Mpc$ at $z=4$; \\citealt{storrie94, miralda03}, probably falling to $\\la 30 \\Mpc$ by $z=6$; \\citealt{lidz07}).  Once ionized regions grew beyond the mean separation of these systems, LLSs (rather than the edges of ionized bubbles) controlled the mean free path of ionizing photons \\citep{furl05-rec, gnedin06} and by extension the effective horizon to which sources could be seen.  Because reionization is so inhomogeneous, many ionized regions will reach this size ($\\ga 20 \\Mpc$) well before complete overlap \\citep{barkana04, furl04-bub}; there must in fact be a gradual transition from the ``bubble-dominated\" ionization topology characteristic of reionization to the ``web-dominated\" topology characteristic of the post-reionization Universe \\citep{furl05-rec}.  Here we ask whether overlap \\emph{must} be accompanied by a rapid increase in $\\Gamma$ and, conversely, whether such an increase is a robust signal of overlap.  We also examine the implications of recent reionization models for the \\lya forest observations conventionally used to argue for late reionization.  Unfortunately, numerical simulations do not yet have the dynamic range to sample the large ($\\sim 100 \\Mpc$) scales necessary for reionization and simultaneously predict detailed properties of the \\lya forest (though see \\citealt{gnedin06} for a detailed study of LLSs and the \\lya forest during reionization in a $8 h^{-1} \\Mpc$ box).  Thus we will use a simple analytic model that incorporates both the discrete ionized bubbles and intervening absorption by LLSs to show that, during the end stages of reionization, $\\Gamma$ is primarily controlled by the mean separation of LLSs, which may not evolve rapidly at the moment of overlap -- and, more importantly, may continue evolving long afterwards.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H_0=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{dunkley08,komatsu08}.   Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  We have shown that, because attenuation due to LLSs must become important before overlap, the late stages of reionization need not be accompanied by a rapid increase in $\\VEV{\\Gamma}$ (or $\\tau_{\\rm eff}$).  Thus, the conventional wisdom that a sharp increase in the ionizing background is a robust indicator of overlap must be modified:  namely, such a feature indicates only that the attenuation length $r_0$ is evolving rapidly.  Of course, it may be that $r_0$ does evolve quickly at the end of reionization (and some theoretical calculations suggest that this does occur; \\citealt{wyithe08-prox, choudhury08}).  For example, $\\VEV{\\Gamma}$ must (at least to some degree) control the ionization structure of LLSs and hence $r_0$ \\citep{miralda05, schaye06}.  As $\\VEV{\\Gamma}$ increases, LLSs will shrink, increasing $r_0$ and hence $\\VEV{\\Gamma}$, creating a positive feedback loop.  Thus even a slow initial increase in the emissivity and/or mean free path at the end of reionization may spiral into a relatively rapid change in $\\VEV{\\Gamma}$.  Unfortunately, our understanding of LLSs is not yet advanced enough to determine the effectiveness of such a feedback loop; for now, we can only say that effective feedback requires that the LLSs have quite steep density profiles (see, e.g., the Appendix to \\citealt{furl05-rec}).  More importantly, there is no obvious reason that the mean free path can \\emph{only} change during overlap:  so long as (i) the emissivity continues to increase and (ii) $\\VEV{\\Gamma}$ controls the properties of LLSs, the feedback loop will continue -- and with the rapid increase in the collapsed fraction at $z \\ga 6$, such a scenario seems entirely plausible.  Thus, even if $\\VEV{\\Gamma}$ is evolving strongly at $z \\sim 6$ \\citep{fan02, fan06}, we do not have a robust indicator of overlap.  Indeed, it may be better to regard the final stage of reionization as the disappearance of LLSs \\emph{after} overlap.  This ``post-overlap\" phase, which consumes just a few percent of the neutral gas, is typically thought to be indistinguishable from the cosmic-web dominated Universe at later times.  We have shown instead that $\\VEV{\\Gamma}$ may evolve slowly during overlap but rapidly afterwards, making quasar absorption spectra a useful probe of this tail end of reionization (which matches smoothly onto the post-reionization Universe) but not necessarily of overlap itself.  This contrasts with the conventional picture in which $\\VEV{\\Gamma}$ evolves rapidly \\emph{only} during overlap; that intuition came from reionization simulations that were too small to include the full inhomogeneity of reionization and so exaggerated the importance of overlap (e.g., \\citealt{gnedin00}).  The recent semi-numeric calculations of \\citet{choudhury08}, which included an approximate prescription for self-shielding gas, present a similar picture to ours, in which these LLSs are burned off over the (rather extended) final stages of reionization that follow overlap.  Beyond an increase in $\\Gamma$, there are other reasons that $r_0$ may evolve throughout (and after) reionization.  For example, photoheating can evaporate loosely-bound structures, substantially modifying the IGM gas distribution by eliminating dense systems \\citep{haiman01, shapiro04, pawlik08}.  This will increase $r_0$ as well; however, the process should happen gradually throughout reionization as more and more of the volume is illuminated (and so will be particularly slow if reionization is extended).  Evaporation also occurs over the sound-crossing time, which is relatively long for the moderate overdensities of most interest at these high redshifts \\citep{pawlik08}, so that LLSs probably continue to evolve past the ``end\" of reionization.  Given the challenges we have raised to the conventional interpretation, it is worth asking two further questions about the data.  First, are there other effects that can cause a rapid increase in $\\tau_{\\rm eff}$ \\emph{without} a significant change in $\\Gamma$?  One possibility is the density distribution itself:  at least according to current models, it is extremely steep in the low-density tail that is responsible for \\lya forest transmission, so it is possible that a small increase in the ionizing background renders a relatively large fraction of the IGM visible \\citep{oh05}.  This must await more detailed calculations of the evolving density field at the end of reionization.  Second, how certain can we be that overlap has actually occurred by $z \\sim 6$?  Is it possible that, although most of the IGM is highly ionized before that point, a small fraction far from ionizing sources could still be completely neutral \\citep{lidz07}?  In the conventional picture, in which $\\VEV{\\Gamma}$ evolves rapidly at overlap, such a conclusion can be easily dismissed.  However, our results suggest that overlap itself may be buried inside the smoothly evolving transmission at $z \\la 6$ -- or it may have occurred at much higher redshifts.  Finally, we expect our conclusions to hold with even more force during helium reionization:  the clumpy IGM and enhanced recombination rate during that era makes attenuation more important \\citep{furl08-helium, mcquinn08, bolton08} and high-energy photons can more easily create a nearly-uniform ionizing background, so the transition from bubble-domination to web-domination will be even smoother.  In summary, with our present knowledge of high-$z$ LLSs there is no compelling reason to associate evolution in $r_0$ or \\lya forest transmission exclusively with overlap.  If indeed the quasar data shows a more rapid evolution in $\\tau_{\\rm eff}$ at $z > 6$ \\citep{fan06}, this may indicate that we are seeing a rapid increase in $r_0$ that followed or even preceded overlap (note, however, that this measurement is itself controversial; \\citealt{songaila02, songaila04, becker07}).   More detailed studies of the coupling between the ionizing background and the dense, neutral blobs that trap ionizing photons are needed before the next step -- associating such a jump with overlap itself -- can be taken securely.  We thank A. Lidz and S.~P. Oh for comments that greatly improved the manuscript.  This research was partially supported by the NSF through grant AST-0607470 to SF.  Support for this work was also provided by NASA through Hubble Fellowship grant \\#HF-01222.01 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."
    },
    "0809/0809.4941_arXiv.txt": {
        "abstract": "We study brane inflation in a warped deformed conifold background that includes general possible corrections to the throat geometry sourced by coupling to the bulk of a compact Calabi-Yau space. We focus specifically, on the perturbation by chiral operator of dimension 3/2 in the CFT.  We find that the effective potential in this case can give rise to required number of e-foldings and the spectral index $n_S$ consistent with observation. The tensor to scalar ratio of perturbations is generally very low in this scenario. The COBE normalization, however,  poses certain difficulties which can be circumvented provided model parameters are properly fine tuned. We find the numerical values of parameters which can give rise to enough inflation, observationally consistent values of density perturbations, scalar to tensor ratio of perturbations and the spectral index $n_S$. ",
        "introduction": "It is well known that the standard model for Big bang cosmology is plagued with some intrinsic problems like horizon and flatness problems. Such problems could be cured if one postulates that the early universe went through a brief period of accelerated expansion, otherwise called the cosmological inflation \\cite{inflation}. The inflationary scenario not only explains the large-scale homogeneity of our universe but also proposes a mechanism, of quantum nature, to generate the primordial inhomogeneities which is the seed for understanding the structure formation in the universe. Such inhomogeities have been observed as anisotropies in the temperature of the cosmic microwave background, thus the theoretical ingredients of inflation are subject to observational constraints. Moreover, the paradigm of inflation has stood the test of theoretical and observational challenges in the past two decades \\cite{Spergel1,Spergel2}. However, in spite of its cosmological successes, it still remains a paradigm in search of a viable theoretical model that can be embedded in a fundamental theory of gravity . In this context, enormous amount of efforts are underway to derive inflationary models from string theory, a consistent quantum field theory around the Planck's scale. Progress in the understanding of the nonperturbative dualities in string theory, to be precise, discovery of D-branes, has further provided an important framework to build and test inflationary models of cosmology.  In past few years, many inflationary models have been constructed in the context of D-brane cosmology. It includes inflation due to tachyon condensation on a non-BPS brane, inflation due to the motion of a D3-brane towards an anti-D3-brane \\cite{sen,linde,kallosh,lindeD}, inflation due to geometric tachyon arising from the motion of a probe brane in the background of a stack of either NS5-branes or the dual D5-branes \\cite{GTach} . However, these models are based on effective field theory and assume an underlying mechanism for the stabilization of various moduli fields. Thus, they do not take into account the details of compactification and the effects of moduli stabilization and hence any predictions from these models are questionable.  Progress in search of a realistic inflationary model in string theory was made when it was learnt that background fluxes can stabilize all the complex structure moduli fields. In fact it has been shown in Ref.~\\cite{GKP} that the fluxes in a warped compactification, using a Klebanov-Strassler (KS) throat \\cite{KS}, can stabilize the axio-dilaton and the complex structure moduli of type IIB string theory compactified on an orientifold of a Calabi-Yau threefold. Further important progress was achieved when it was shown in Ref.~\\cite{KKLT} that the K\\\"ahler moduli fields also can be stabilized by a combination of fluxes and nonperturbative effects. The nonperturbative effects, in this context, arise, via gauge dynamics of either an Euclidean D3-brane or from a stack of   D7-branes wrapping super-symmetrically a four cycle in the warped throat. The warped volume of the four cycle controls the magnitude of the nonperturbative effect since it affects the gauge coupling on the D7-branes wrapping this four cycle.  Armed with these results, an inflationary model \\cite{KKLMMT} has been built taking into account of the compactification data (see also Refs.~\\cite{dbpapers}). The inflaton potential is obtained by performing string theoretic computations involving the details of the compactification scheme. In this setup inflation is realized by the motion of a D3-brane towards a distant static anti-D3-brane, placed at the tip of the throat and the radial separation between the two is considered to be the inflaton field. The effect of the moduli stabilization resulted in a mass term to the inflaton field which is computed in \\cite{KKLMMT}. It turned out, unfortunately that the mass is large and of the order of hubble parameter and hence spoils the inflation. To avoid this disappointment (see, however, \\cite{IT}), one needs to search for various sources of correction to the inflaton potential. For example, in Ref.~\\cite{Bau1}, the embedding of the D7-branes as given in \\cite{Kup} was considered. Assuming that at least one of the four-cycles carrying the nonperturbative effects descend down a finite distance into the warped throat so that the brane is constrained to move only inside the throat, it was found that such a configuration leads to a perturbation to the warp factor affecting a correction to the warped four cycle volume. Further, this correction depends on the position of the D3-brane and thus the superpotential for the nonperturbative effect gets corrected by an overall position-dependent factor. Thus the full potential on the brane is the sum of the potential (F-term) coming from the superpotential and the usual D-term potential contributed by the interaction between the D3-brane and the anti-D3-brane. Taking these corrections into account, the volume modulus stabilization has been re-analyzed in Refs.~\\cite{Bau2,Bau3,KP} and  the viability of inflation was investigated in this modified scenario. The stabilization of the volume modulus puts severe constraint which is difficult to solve analytically without invoking approximations. Moreover, the model needs extreme fine tuning.  The above model has been reexamined in Ref.~\\cite{PSS} where inflation, involving both the volume modulus and the radial distance between the brane and the anti-brane participate in the dynamics. Making a rotation in the trajectory space, a linear combination of the two fields, which become independent of time and hence can be stabilized, is identified with the volume modulus. The orthogonal combination then becomes the inflaton field. The model again needs severe fine tuning and further it has been observed that when the spectral index of scalar perturbation reaches the scale invariant value, the amplitude tends to be larger than the COBE normalized value by about three order of magnitude, making the model seemingly unrealistic. However, a recent analysis using Monte Carlo method for searching the parameter space shows that COBE normalization as well as the requirement of nearly flat  spectrum can be satisfied at the same time \\cite{HC}.  In this paper, we analyse the possibility of a realistic model for brane inflation remaining within the large volume compactification scheme but incorporating a different correction to the inflaton potential. Infact, the authors of Ref.~\\cite{BDKKM} have recently reported that there can be corrections to the inflaton potential that arise from ultra violet deformations corresponding to perturbation by the lowest-dimension operators in the dual conformal field theory. These contributions to the potential have sensitive dependence on the details of moduli stabilization and the gluing of the throat into the compact Calabi-Yau space. In the next section we briefly review the origin and form of these contributions to the potential. In section 3, we analyse the inflationary dynamics based on the full potential. Section 4 is devoted to the summery of our analysis and conclusions. ",
        "conclusions": "In this paper we have analysed the possibilities of inflation in a warped background with an effective potential (\\ref{Spot}). The model includes three parameters, ${\\cal D}$, $\\alpha=M_{UV}/M_P$ and $C_{3/2}$. The fact that $D3$ brane moves towards the tip of the throat and can reach close to it and not vise-versa imposes constraints on the model parameters, namely, $C_{3/2} \\lesssim 0.1$ for a viable range of $\\alpha$. The COBE normalization demands that typically, ${\\cal D}^{1/4}\\sim 10^{-4}$ which makes $C_{3/2}$ much smaller than one. Inflation becomes possible near the origin where potential can be made sufficiently flat by appropriately choosing the parameters. Numerical values of model parameters can easily be set to obtain enough inflation and observationally consistent value of $n_S$. For instance, in case of, $\\alpha^{-1}=2.000, C_{3/2}=0.00982,{\\cal D}=1.000\\times 10^{-15}$, we find that $N\\simeq 60$ and $n_S \\simeq 0.96$. The tensor to scalar ratio is very low in this case. The problem is caused by COBE normalization; the amplitude of density perturbations is large in this case, $\\delta_H^2 \\simeq 6\\times 10^{-8}$ . This is related to the fact that the constant ${\\cal D}$ does not give the over all scale of the potential but crucially effects the slow roll parameters and the spectral index. Infact, for several other choices of ${\\cal D}$, we can obtain the required values of $N$ and $n_S$ keeping the low value of tensor to scalar ratio of perturbations. However, it is tricky to satisfy  the COBE normalization. Our search led to the viable numerical values of the parameters and we have shown that for ${\\cal D}=1.210\\times 10^{-17},~~\\alpha^{-1}=2.11991,~~C_{3/2}=0.0062284$; the total number of e-foldings after horizon crossing  $N$ is around $60$ and $n_S\\simeq 0.96$ at the cosmologically relevant scales. For this choice of parameters, the COBE normalization can be set correctly and we find that $\\delta_H^2 \\simeq 2.4\\times 10^{-9}$. It is really interesting that in this case the tensor to scalar ratio is very low, $r=16\\epsilon \\simeq 10^{-11}$, at the horizon crossing. It is possible to vary the parameters around the given values and still satisfy the observations constraints provided that we fine tune the parameters ${\\cal D}$, $\\alpha$ and $C_{3/2}$. Our search in the parameter space was carried out manually which allowed us to demonstrate that the scenario under consideration can be made consistent with the findings of WMAP5.  Finally, we should comment on the phenomenological aspects of model discussed in Ref.[18] and compare it with the analysis presented here. It was demonstrated in Ref.[20] that the single field models based upon the scenario of Ref.[18] can give rise to only few e-folds given the constrains on the model parameters. However, the two field model (i.e. the volume modulus also participates in the inflation dynamics) performs much better in this case. But it is difficult to satisfy all the observational constraints simultaneously [20], specially constraints of the spectral index and the COBE normalization are problematic. In the scenario of Ref.[22](where the microscopic theory is different from Ref.[18]) analyzed here, the constraints on model parameters are relaxed and hence the phenomenology of single field inflation dynamics becomes possible. For such a  single field model, we have demonstrated that it is possible to satisfy all the observational constraints."
    },
    "0809/0809.4988_arXiv.txt": {
        "abstract": "We present the characterization of additional properties of the night-sky at the Calar Alto observatory, following the study started in S\\'anchez et al. (2007), hereafter Paper I. We focus here on the night sky-brightness at the near-infrared, the telescope seeing, and the fraction of useful time at the observatory. For this study we have collected a large dataset comprising 7311 near-infrared images taken regularly along the last four years for the ALHAMBRA survey ($J$, $H$ and $Ks$-bands), together with a more reduced dataset of additional near-infrared images taken for the current study. In addition we collected the information derived by the meteorological station at the observatory during the last 10 years, together with the results from the cloud sensor for the last $\\sim$2 years. We analyze the dependency of the near-infrared night sky-brightness with the airmass and the seasons, studying its origins and proposing a zenithal correction. A strong correlation is found between the night sky-brightness in the $Ks$-band and the air temperature, with a gradient of $\\sim -$0.08 mag per 1 $\\degree$ C. The typical (darkest) night sky-brightness in the $J$, $H$ and $Ks$-band are 15.95 mag (16.95 mag), 13.99 mag (14.98 mag) and 12.39 mag (13.55 mag), respectively. These values have been derived for the first time for this observatory, showing that Calar Alto is as dark in the near-infrared as most of the other astronomical astronomical sites in the world that we could compare with. Only Mauna Kea is clearly darker in the $Ks$-band, but not only compared to Calar Alto but to any other observatory in the world. The typical telescope seeing and its distribution was derived on the basis of the FWHM of the stars detected in the considered near-infrared images. This value, $\\sim$1.0$\\arcsec$ when converted to the V-band, is only slightly larger than the atmospheric seeing measured at the same time by the seeing monitor, $\\sim$0.9$\\arcsec$. Therefore, the effects different than the atmosphere produce a reduced degradation on the telescope seeing, of the order of $\\sim$10\\%. Finally we estimate the fraction of useful time based on the relative humidity, gust wind speed and presence of clouds. This fraction, $\\sim$72\\%, is very similar to the one derived in Paper I, based on the fraction of time when the extinction monitor is working. ",
        "introduction": "The night sky brightness in the optical and near-infrared, the number of clear nights, the seeing, transparency and photometric stability are some of the most important parameters that qualify a site for front-line ground-based astronomy. There is limited control over all these parameters, and only in the case of the sky brightness is it possible to keep it at its natural level by preventing light pollution from the immediate vicinity of the observatory. Previous to the installation of any observatory, extensive tests of these parameters are carried out in order to find the best locations, maximizing then the efficiency of these expensive infrastructures.  However, most of these parameters are not constant along the time.  We have started a program to determine the actual values of the main observing conditions for the Calar Alto observatory.  The Calar Alto observatory is located at 2168m height above the sea-level, in the Sierra de los Filabres (Almeria-Spain) at $\\sim$45 km from the Mediterranean sea. It is the second largest European astronomical site in the north hemisphere, behind the Observatorio del Roque de los Muchachos (La Palma), and the most important in the continental Europe. Currently there are six telescopes located in the complex, three of them operated by the Centro Astronomico Hispano Aleman A.I.E. (CSIC-MPG), including the 3.5m, the largest telescope in the continental Europe.  Along its 26 years of operations there has been different attempts to characterize some of the main properties described before: (i) Birkle et al. (1976) presented the first measurements of the seeing at Calar Alto, being the basis of the site testing; (ii) Leinert et al. (1995) determined the optical sky brightness corresponding to the year 1990; (iii) Hopp \\& Fernandez (2002) studied the extinction curve corresponding to the years 1986-2000; (iv) Ziad et al. (2005) estimated the median site seeing in the observatory from a single campaign in may 2002. A consistent study of all these properties was presented in S\\'anchez et al. (2007), including the first results of our program.  In that study, hereafter Paper I, we focused on the optical properties of the night sky, presenting (i) the first optical night-sky spectrum, identifying the natural and light pollution emission lines and their strength, (ii) the moon-less night-sky brightness in different optical bands, (iv) the extinction and its yearly evolution and (v) the atmospheric seeing and its yearly evolution. It was found that most of these properties were similar to those of major astronomical sites. In particular Calar Alto is among the darkest places in the world in the optical range.  In the current study we focus on the near infrared (NIR) properties of the sky at Calar Alto. In particular we characterize here the typical and darkest sky brightness in the $J$, $H$ and $K$-bands, comparing them with similar properties in other observatories world-wide. Their dependencies with positioning in the sky and seasons are also analyzed. The typical seeing measured in real observations (telescope seeing) is derived and compared with the atmospheric seeing estimated by the seeing monitor (out of the telescope). Finally we estimate the fraction of useful time in the observatory based on the data obtained by the meteorological station along the last ten years (relative humidity and gust wind) and the presence of clouds based on the cloud sensor.  The structure of this article is as follows: in Section \\ref{data} we describe the dataset collected for the current study, including a description of data and the data reduction; in Section \\ref{ana} we show the analysis and results performed over the different types of data and the results derived for each one; in Section \\ref{conc} we summarize the main results and present the conclusions. All the magnitudes listed in this article are in the Vega system. ",
        "conclusions": "\\label{conc}  In this article we have continued the characterization of the main properties of the night-sky at the Calar Alto observatory that we started in Paper I. These properties were compared with similar ones of other different astronomical sites. The main results of this article can be summarized in the following points:  \\begin{itemize}  \\item The night sky brightness at the near-infrared presents a strong airmass dependency in the $J$ and $H$ band, and a mild one in the $Ks$-band. These dependencies can be modelled and corrected (or predicted) for any night and airmass by assuming that they are due to airmass dependency of the OH lines emission. Obviously, $J$ and $H$-band observations are recomended to be performed near the zenith, while at the $Ks$-band this is less critical. This result should be taken into account when defining observing strategies in the future.  \\item There is a seasonal pattern in the $Ks$-band sky brightness, which shows an average difference of $\\sim$0.8 mag from summer, when is brighter, to winter, when is darker. This pattern is due to the strong dependecy of the sky-brightness in this band with the atmospheric temperature, which may indicate that the contribution from the thermal background emission is the dominant component in this band.  \\item The brightness of the night sky at Calar Alto is similar or darker than any in other major astronomical site in the $J$ and $H$ bands, including both Paranal and Mauna Kea. In the $Ks$-band all the astronomical sites are very similar apart from Mauna Kea, which is clearly darker (by $\\sim$1 mag). Indeed, Calar Alto has a similar sky brightness than Paranal or La Palma at this wavelength range.   \\item The seeing on the detector is only slightly larger than the site seeing measured by the DIMM, outside the dome (at least for the 3.5m telescope). The effects different than the purely atmospheric ones produce a statistical increase of a $\\sim$10\\% in the measured seeing on the detector compared with the site one.  \\item The fraction of useful time at the observatory is of the order of $\\sim$70\\%, with $\\sim$40\\% of the nights 100\\% completely useful. These fractions are better during Summer and worst during Winter.  \\end{itemize}  We finally conclude that our additional analysis of the astronomical conditions at Calar Alto agree with our previous results presented in Paper I. The new data strength our previous conclusion that it is a good astronomical site, similar in many aspects to places where there are 10m-like telescopes under operation or construction.   For both reasons we consider that this observatory is a good candidate for the location of future large aperture optical/NIR telescopes."
    },
    "0809/0809.2389_arXiv.txt": {
        "abstract": "We present spectra of a sample of Herbig Ae and Be (HAeBe) stars obtained with the Infrared Spectrograph on the \\textit{Spitzer Space Telescope}. All but one of the Herbig stars show emission from polycyclic aromatic hydrocarbons (PAHs) and seven of the spectra show PAH emission, but no silicate emission at 10~\\mum. The central wavelengths of the 6.2, 7.7--8.2, and 11.3~\\mum\\ emission features decrease with stellar temperature, indicating that the PAHs are less photo-processed in cooler radiation fields. The apparent low level of photo processing in HAeBe stars, relative to other PAH emission sources, implies that the PAHs are newly exposed to the UV-optical radiation fields from their host stars. HAeBe stars show a variety of PAH emission intensities and ionization fractions, but a narrow range of PAH spectral classifications based on positions of major PAH feature centers. This may indicate that, regardless of their locations relative to the stars, the PAH molecules are altered by the same physical processes in the proto-planetary disks of intermediate-mass stars. Analysis of the mid-IR spectral energy distributions indicates that our sample likely includes both radially flared and more flattened/settled disk systems, but we do not see the expected correlation of overall PAH emission with disk geometry. We suggest that the strength of PAH emission from HAeBe stars may depend not only on the degree of radial flaring, but also on the abundance of PAHs in illuminated regions of the disks and possibly on the vertical structure of the inner disk as well. ",
        "introduction": "Circumstellar accretion disks that contain dust and gas are common around intermediate mass stars. Studying the material in these disks and their evolution can aid us in understanding the very latest stages of star formation and the early stages of planet formation. Herbig Ae/Be (HAeBe) stars are intermediate-mass (M$\\sim$2-10 M$_{\\sun}$) pre-main-sequence stars of spectral class A, B, or early F \\citep{her60, str72, the94, mal98}. Though often associated with nebulosity, HAeBe photospheres are directly observable at visible wavelengths. Unlike normal A and B stars, HAeBe stars have strong emission lines (e.g. H$\\alpha$ and Br$\\gamma$) in their spectra, they can be highly variable at visible wavelengths, and their infrared fluxes are strongly in excess of purely photospheric emission. Current theory predicts that intermediate-mass stars should form more quickly than lower-mass (T Tauri) stars, but they will not reach their stellar birth lines heavily enshrouded in dense circumstellar envelopes as high-mass stars do. This means that we can directly view the stellar photospheres and residual accretion disks of HAeBe stars, with a relatively unobscured view of both. Unlike high-mass stars, while some HAeBe stars are associated with star forming regions they are not common in the extreme environments of HII regions near OB associations. Thus the stars themselves, and not their environments, drive their disk structure, chemistry, and evolution.  Chemical and dynamical studies of Galactic HAeBe stars have shown holes and gaps in their inner disks indicating the possible dynamical influence of newly formed planets \\citep[e.g.][]{gra05,bri07}, as well as evidence of different disk geometries from those that flare with increasing radius and have dust with a large abundance of small grains, to flattened disks composed of larger dust grains \\citep[e.g.][]{mee01,aa04}. These characteristics are similar to those of the lower-mass T Tauri systems.  Solid-state features and thermal radiation from warm dust, as well as emission from atomic and molecular gas all produce measurable signatures in the mid-IR spectra of young stars with circumstellar disks. Mid-IR imaging and spectroscopy, much of it from space telescopes, have therefore played a major role in measuring and characterizing the disks orbiting T Tauri and HAeBe stars. Unlike T Tauri stars, of which only $\\sim$8\\% show emission from polycyclic aromatic hydrocarbons \\citep[PAHs,][]{gee06}, nearly 50\\% of HAeBe stars have strong PAH emission \\citep[e.g.][]{mee01,aa04}, which requires a direct line of sight from the emitting material---presumably located on the surfaces and/or inner rims of disks---to the photospheric optical and UV radiation fields. Furthermore, spectra from the Infrared Space Observatory (\\iso) and recent \\textit{Spitzer} observations have revealed that many HAeBe stars in the Galaxy have large abundances of crystalline silicate grains in their disks indicating significant thermal dust processing subsequent to the formation of the star-disk system. These latter characteristics are not as common in T Tauri disks. Moreover, an evolutionary link from HAeBe proto-planetary disks to Vega-like debris disks is becoming evident in studies of large samples of HAeBe stars \\citep[e.g.][]{her06,gra07}, but the details have yet to be completely worked out. Though details of time scale and evolution remain elusive, recent modeling efforts have succeeded in predicting spectral energy distributions from varying physical characteristics in the disks like dust grain size and composition, properties of the stellar radiation field, disk heating and cooling, and disk geometry \\citep[e.g.][]{dul04,dal06,dul07,rob07}.  It is remarkable that such a large fraction of HAeBe stars are PAH emission sources. \\cite{mee01} suggested that PAH emission may change measurably with the structure and physical properties of the disks and perhaps with disk evolution. PAHs have been observed in many different astrophysical contexts: the Galactic and extra-galactic ISM, dust envelopes surrounding post-AGB (asymptotic giant branch) stars, planetary nebulae (PNe), and young stellar objects (YSOs). The molecular structure of PAHs can differ significantly depending on the physical characteristics of the objects or environments where they are found. The bottom line is that PAHs are abundant in interstellar and circumstellar environments, and they can be important constituents in the energy balance of those environments as sources of photoelectrons that heat the gas component \\citep{kam04}.  Emission from PAH molecules may also trace proto-planetary disk geometry \\citep{mee01,aa04,gra05}. Exactly what the PAH emission can tell us about the disk environments will depend upon a thorough and detailed understanding of how the molecules are energized and processed by the stellar radiation fields and how they are distributed radially and vertically in the disks. Throughout this paper when we refer to \"processing\" of PAHs, we mean photo-processing by which the optical and UV stellar radiation break molecular bonds producing smaller PAH molecules or completely destroying them. We have begun a detailed analysis of PAH emission from HAeBe stars using very high S/N spectra that allow much higher precision in analyzing PAH feature strengths and wavelengths, and thus the molecular structure that produces them, than has been possible before.  In a previous paper (Sloan et. al. 2005, hereafter Paper 1), we analyzed four HAeBe stars in the 5-14~\\mum\\ range that have strong PAH emission but no silicate peak at 10~\\mum. We identified a trend towards decreasing PAH molecule size with increasing PAH ionization fraction and noted that the center wavelengths of the PAH features shift systematically with increasing UV field strength from the stars. We suggested that HAeBe stars may have distinctive PAH spectra relative to other PAH sources (e.g. interstellar photon-dominated regions, PNe, etc).  We present 5-36~\\mum\\ spectra of 18 HAeBe stars, two intermediate-mass T Tauri stars, HD 97300, and 51 Oph for a total of 22 sources observed with \\textit{Spitzer}. This extends our original sample of HAeBe stars in Paper 1 by more than a factor of four, and we now analyze HAeBe stars that do have the 10~\\mum\\ silicate emission feature. We present a detailed analysis of the PAH feature strengths and their ratios in an attempt to infer their physical environments and to test for trends in PAH emission when compared with other diagnostics of proto-planetary disk evolution that have recently been helpful in understanding the evolution of T Tauri disks \\citep[e.g.][]{fur06}.  The major trends suggested in Paper 1 persist in the present, larger sample. \\cite{aa04} noted in their study of 46 HAeBe stars that those with strong mid-IR excesses relative to near-IR excesses have stronger PAH emission than those with weaker mid-IR excesses. This result supports a relation between the detailed structure of PAH emission and the evolution of the disks. Recent studies by \\cite{gra05} suggest that the strength of (UV-optical) photon-dominated spectral features should correlate with overall disk geometry. Disks that are flat and whose outer disks are therefore shadowed or illuminated at grazing incidence should have weaker H$_{2}$ line and PAH features relative to disks that are flared and therefore well-illuminated by the photospheres of their host stars. These results imply a strong correlation between PAH strength and disk structure, which we can now test with our data. ",
        "conclusions": "  \\begin{itemize}  \\item PAHs exposed to stronger stellar radiation fields may become shattered and develop more jagged edges so that the 11.3~\\mum\\ mono C-H bending mode is attenuated with respect to the 12.7~\\mum\\ trio C--H bending mode. PAH ionization fraction, inferred from the PAH F$_{7.7-8.2}$/F$_{11.3}$ ratio, is related to the structure (and possibly the size) of the PAH molecules.  \\item The center wavelengths of the 6.2, 7.7--8.2, 8.6, and 11.3~\\mum\\ emission features change as the stellar effective temperature changes. We confirm our previous finding that the center wavelengths of the PAH emission features increase with decreasing effective temperature of the host star, now with a much larger sample of stars. We conclude that hotter stars photo-process the PAHs more than cooler stars do and that the degree of processing alters the PAH spectrum measurably.  \\item Our sample shows spectral characteristics generally consistent with the boundary between Class B or C PAH emission sources defined by \\cite{pee02}, viz. a 6.2~\\mum\\ feature shifted to 6.3~\\mum\\ and 7.7--8.2~\\mum\\ emission dominated by the component at 7.9~\\mum. However, in some of our sources, the 7.7--8.2~\\mum\\ feature center is shifted closer to 8.0~\\mum, as in the spectrum of HD 100546, a HAeBe star observed by the SWS on \\iso\\ and included in our sample. It is remarkable that HAeBe stars as a class have such a variety of PAH luminosities and ionization fractions, but such a narrow range of PAH spectral classifications. This may indicate that the PAH molecules are altered by the same physical processes in the disks, perhaps unique to HAeBe stars. In any case we propose that the HAeBe stars comprise a distinct PAH spectral group of their own. HAeBe PAH spectra indicate very little processing of the PAH molecules by the stellar radiation field, relative to other types of PAH source (e.g. HII regions, PDRs, and PNe).  \\item The other Peeters Class B objects are post-AGB and that implies that the PAH are less processed and therefore relatively new to the scene in Herbig Ae/Be disks. This is consistent with the suggestion that the PAHs are introduced at a stage of disk evolution when icy mantles evaporate from large dust grain surfaces. If this is the case, then the PAHs near HAeBe stars are being constantly destroyed and replenished by interactions with the stellar radiation field since we do not observe HAeBe systems with the (apparently more processed) Class A spectra.  \\item We have detected the 17~\\mum\\ PAH complex in one of our sample HAeBe stars, HD 97300, which stands apart form the rest of our sample as the only class A (e.g. highly photo-processed) PAH spectrum and is likely not a Herbig Ae/Be star. The 17~\\mum\\ feature is commonly observed in the Galactic ISM and in other galaxies where strong UV fields excite extensive PDRs. HD 97300 is also a Meeus Group IIb source having no warm small grain component in its SED and no silicate features at 10~\\mum\\ or longward of 20~\\mum. We suggest that, whatever the classification of HD 97300, it's spectrum is dominated by PAHs in the surrounding interstellar gas and not in a circumstellar disk.  \\item Our sample includes both flared (Meeus et al. 2001, Group I) and flattened/settled (Group II) systems, but the overall PAH emission does not decrease monotonically with indicators of depletion of small dust grains and disk settling, as one would expect if the PAH emission only originates in an optically and geometrically thin surface layer of a radially flared disk.  \\item Taking into account estimates of disk inclination for the sources in our sample does not uncover a direct correlation of weak PAH emission with indicators of nearly edge-on disk geometry. Nor can we explain the anomaly by invoking strong stellar radiation fields to destroy or process the PAHs so that they emit less light in some disks.  \\item Differences in PAH abundance in the illuminated parts of HAeBe disks and/or evolution of the inner disk scale height appear to be possibilities for explaining highly flared disks that have relatively weak PAH emission. These sources may be intermediate-mass versions of the transitional T Tauri disks.  \\end{itemize}  A larger sample of HAeBe stars will allow us to better determine what a \"normal\" HAeBe star looks like. We have begun the next step, a thorough modeling effort including the addition of detailed PAH spectral synthesis and careful application of mineralogical models for an even larger sample of stars."
    },
    "0809/0809.1175_arXiv.txt": {
        "abstract": "Using measured radial velocity data of nine double lined spectroscopic binary systems NSV 223, AB And, V2082 Cyg, HS Her, V918 Her, BV Dra, BW Dra, V2357 Oph, and YZ Cas, we find corresponding orbital and spectroscopic elements via the method introduced by Karami \\& Mohebi (2007a) and Karami \\& Teimoorinia (2007). Our numerical results are in good agreement with those obtained by others using more traditional methods. ",
        "introduction": "\\label{intro} Determining the orbital elements of binary stars helps us to obtain fundamental information, such as the masses and radii of individual stars, that has an important role in understanding the present state and evolution of many interesting stellar objects. Analysis of both light and radial velocity (hereafter RV) curves, derived from photometric and spectroscopic observations, respectively, yields a complete set of basic absolute parameters. One historically well-known method to analyze the RV curve is that of Lehmann-Filh\\'{e}s (cf. Smart, 1990).   In the present paper we use the method introduced by Karami $\\&$ Mohebi (2007a) (=KM2007a) and Karami $\\&$ Teimoorinia (2007) (=KT2007), to obtain orbital parameters of the nine double-lined spectroscopic binary systems: NSV 223, AB And, V2082 Cyg, HS Her, V918 Her, BV Dra, BW Dra, V2357 Oph, and YZ Cas. Our aim is to show the validity of our new method to a wide range of different types of binary.  The NSV 223 is a contact system of the A type and the mass ratio is believed to be small (Rucinski et al. 2003a,b). The large semiamplitudes ,$K_i$, suggest that the orbital inclination is close to $90^\\circ$. The spectral type is $F7V$ and the period is $0.366128$ days (Rucinski et al. 2003a,b). The AB And is a contact binary. The spectral type is $G8V$. The period is $0.3318919$ days. It is suggested that observed period variability may be a result of the orbital motion in a wide triple system. The third body should then have to be a white dwarf (Pych et al. 2003,2004). V2082 Cyg is most probably an A-type contact binary with a period of $0.714084$ days. The spectral type is $F2V$ (Pych et al. 2003,2004). In HS Her, the effective temperatures were found to be $T_1=15200\\pm750K$ and $T_2=7600\\pm400K$ for the primary and secondary stars, respectively (Cakirli et al. 2007). The secondary component appears to rotate more slowly. The presence of a third body physically bound to the eclipsing pair has been suggested by many investigators. The two component are located near the zero-age main sequence, with age of about $32$ Myr (Cakirli et al. 2007). It is classified as an Algol-type eclipsing binary and single-lined spectroscopic binary. The spectral type of more massive primary component is $B4.5V$. The effective temperature is about $15200\\pm750K$ for the primary component and $7600\\pm400K$ for the secondary component. The period is $1.6374316$ days (Cakirli et al. 2007). The V918 Her is an A-type contact binary. The spectral classification is A7V. The period of this binary star is $0.57481$ days (Pych et al. 2003,2004). The BV Dra  and BW Dra have a circular orbit. From spectrophotometry the components of BV Dar are classified as F9 and F8 while the components of BW Dra are G3 and G0. The period of BV Dra is $0.350066$ days and for BW Dra is $0.292166$ days (Batten \\& Wenxian 1986). The V2357 Oph was classified as a pulsating star with a period of $0.208$ days. The spectral type is G5V (Rucinski et al. 2003a,b). The spectral type of primary component of YZ Cas is A1V and for the secondary is F7V. The period of this binary is $4.4672235$ days (Lacy 1981).  This paper is organized as follows. In Sect. 2, we give a brief review of the method of KM2007a and KT2007. In Sect. 3, the numerical results are reported, while the conclusions are given in Section 4. ",
        "conclusions": "Using the method introduced by KM2007a and KT2007, we obtain both the orbital elements and the combined spectroscopic parameters of nine double lined spectroscopic binary systems. Our results are in good agreement with the those obtained by others using more traditional methods. There are some awareness about very close or contact binaries which show some anomalies from more general (detached) binaries. Systematic differences are sometimes found from standard Keplerian models. Even so, the results which are obtained appear comparable to those found by other authors using other methods. This does not mean that all these results are correct. It just implies that where systematic differences appear they affect the results found by different methods of analysis in similar ways. There are also some awareness about the complication binaries which do not behave in an ideal way. For example, active RS CVn type binaries with spots and the non-uniform surface flux distribution should affect the radial velocities.\\\\\\\\ \\noindent{{\\bf Acknowledgements}} This work has been supported financially by Research Institute for Astronomy $\\&$ Astrophysics of Maragha (RIAAM), Maragha, Iran."
    },
    "0809/0809.4420_arXiv.txt": {
        "abstract": "This document presents the results from our spectroscopic survey of H$\\alpha$ emitters in galactic and SMC open clusters with the ESO Wide Field Imager in its slitless spectroscopic mode. First of all, for the galactic open cluster NGC6611, in which, the number and the nature of emission line stars is still the object of debates, we show that the number of true circumstellar emission line stars is small. Second, at low metallicity, typically in the Small Magellanic Cloud, B-type stars rotate faster than in the Milky Way and thus it is expected a larger number of Be stars. However, till now, search for Be stars was only performed in a very small number of open clusters in the Magellanic Clouds. Using the ESO/WFI in its slitless spectroscopic mode, we performed a H$\\alpha$ survey of the Small Magellanic Cloud. 3 million low-resolution spectra centered on H$\\alpha$ were obtained in the whole SMC. Here, we present the method to exploit the data and first results for 84 open clusters in the SMC about the ratios of Be stars to B stars. ",
        "introduction": "Observations were performed in 2002, at the 2.2m of the ESO at la Silla equiped with the Wide Field Imager in its slitless  spectroscopic mode (Baade et al. 1999). This kind of instrumentation is not sensitive to the ambient diffuse nebula and displays only emission lines, which come from circumstellar matter like in the case of Be stars. Be stars are very fast rotating stars, which are surrounded by a  circumstellar decretion disk. This instrumentation allowed Martayan et al. (2008a) to find true cirumstellar emission line stars in the  Eagle Nebula and NGC6611 open cluster located in the Milky Way, while slit-spectroscopic observations show strong nebular lines. Only a small number of true emission line stars (less than 10) was found.  In the Milky Way, we used broad bandpass filter centered in H$\\alpha$, but in the Magellanic Clouds due to the crowding of the fields, we used a narrow bandpass filter also centered in H$\\alpha$. The exposure times range from 120 to 600s, and the resolving power is low ($\\sim$100). In the SMC, 14 images were obtained, $\\sim$8000 spectra were treated in 84 SMC open clusters among the 3 million obtained for the whole SMC, and in NGC6611 $\\sim$10000 spectra were treated. In the LMC, 5 million spectra were obtained.  The data-reduction was performed using IRAF tasks and the spectra extraction with SExtractor (Bertin \\& Arnouts 1996). The analysis of spectra and emission line stars detection were done using lecspec and ALBUM codes by Martayan et al. (2008a,b). To classify the stars in SMC open clusters, we cross-correlated our WFI catalogues with OGLE ones (Udalski et al. 1998) to obtain the photometry (B, V, I) for each star and various information for each open cluster (E[B-V], age, reddening). More than 4000 stars of SMC open clusters were classified. An example of colour-magnitude digrams is shown for different open clusters in the SMC in Fig.~\\ref{fig1}.   \\begin{figure}[ht] \\begin{center} \\resizebox{8cm}{!}{\\includegraphics[angle=-90]  {martayan1_fig1.ps}} \\resizebox{8cm}{!}{\\includegraphics[angle=-90]  {martayan1_fig2.ps}} \\caption{{\\bf Left}: SMC open cluster OGLE-SMC107 (NGC330).  {\\bf Right}: SMC open cluster OGLE-SMC069. Black dots correspond to the absorption stars, red full circles to definite emission line stars, and red empty circles to candidate emission line stars.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "We conducted a large slitless spectroscopic survey in the Magellanic Clouds and in 2 open clusters in the Milky Way with the ESO/WFI in its slitless spectroscopic mode. In the open cluster NGC6611 and the Eagle Nebula (M16), we show that there is only a small number of true emission line stars. With the stars from 83 open clusters in the SMC, we show that there are twice to four times more Be stars in the SMC than in the MW open clusters. The exploitation of the spectra in the SMC field and in the whole LMC (field and open clusters) is currently ongoing."
    },
    "0809/0809.3036_arXiv.txt": {
        "abstract": "This article investigates the full Boltzmann equation up to second order in the cosmological perturbations. Describing the distribution of polarized radiation by a tensor valued distribution function, we study the gauge dependence of the distribution function and summarize the construction of the gauge-invariant distribution function. The Liouville operator which describes the free streaming of electrons, and the collision term which describes the scattering of photons on free electrons are computed up to second order. Finally, the remaining dependence in the direction of the photon momentum is handled by expanding in projected symmetric trace-free multipoles and also in the more commonly used normal modes components. The results obtained remain to be used for computing numerically the contribution in the cosmic microwave background bispectrum which arises from the evolution of second order perturbations, in order to disentangle the primordial non-Gaussianity from the one generated by the subsequent non-linear evolution. ",
        "introduction": "\\label{Sec_Genform}  \\subsection{Momentum and tetrad}  In the kinetic description of radiation, the momentum of photons is usually decomposed onto an orthonormal basis, that is using a tetrad field. Though not compulsory, this facilitates the separation between the magnitude of the momentum and its direction represented by a spacelike unit vector. The tetrad vectors $\\gr{e}_\\aT$ and their corresponding tetrad forms $\\gr{e}^\\aT$ satisfy the orthonormality conditions \\be\\label{deftetrad} \\gr{e}_\\aT.\\gr{e}_\\bT\\equiv e_{\\aT}^{\\,\\,\\mu} e_{\\bT}^{\\,\\,\\nu} g_{\\mu \\nu}=\\eta_{\\aT \\bT},\\qquad\\gr{e}^\\aT.\\gr{e}^\\bT\\equiv e^{\\aT}_{\\,\\,\\mu} e^{\\bT}_{\\,\\,\\nu} g^{\\mu \\nu}=\\eta^{\\aT \\bT}, \\ee where $g_{\\mu\\nu}$ is the space-time metric, $g^{\\mu\\nu}$ its inverse and $\\eta^{\\aT \\bT}=\\eta_{\\aT \\bT}$ the Minkowski metric. In the previous expressions and throughout this paper, we use Greek indices ($\\mu,\\nu,\\rho,\\sigma\\dots$) for abstract indices and the beginning of the Latin alphabet ($\\aT,\\bT,\\cT,\\dT\\dots$) for tetrad labels. Since the tetrad labels run from $0$ to $3$, we also use Latin indices starting from the letter $\\iT$ (that is $\\iT,\\jT,\\kT,\\lT\\dots$) with values ranging from $1$ to $3$ to label the spacelike vectors or forms of a tetrad. We then reserve the label $\\zT$ for the timelike vector and form in a tetrad. A momentum can then be decomposed as \\be \\gr{p}=p^\\aT \\gr{e}_\\aT=p^\\zT \\gr{e}_\\zT +p^\\iT \\gr{e}_\\iT\\,, \\ee where the components can be extracted as \\be p^\\aT=\\gr{p}.\\gr{e}^\\aT \\equiv e^{\\aT}_{\\,\\mu}p^\\mu\\,. \\ee It can be decomposed into a magnitude $p^\\zT$ and the a direction vector $\\gr{n}$ according to \\be p^\\mu=p^\\zT (e_{\\zT}^{\\,\\mu} +n^\\mu)\\,,\\quad \\gr{n}.\\gr{e}_\\zT \\equiv n_\\mu e_{\\zT}^{\\,\\mu}=0\\,, \\quad n^\\mu n_\\mu\\equiv\\gr{n}.\\gr{n}=1\\,. \\ee This decomposition can be used to define a screen projector \\be\\label{DefScreen} {S_{\\gr{e_\\zT}}}_{\\mu\\nu}(\\gr{p})=g_{\\mu \\nu}+e^\\zT_{\\,\\mu} e^\\zT_{\\,\\nu}-n_\\mu n_\\nu\\,, \\ee which projects on a space which is both orthogonal to $\\gr{e}_\\zT$ and orthogonal to the direction $\\gr{n}$, since \\be \\flamoi{S_{\\gr{e}_\\zT}}_{\\mu \\nu} e_\\zT^{\\,\\mu}=0,\\qquad {S_{\\gr{e}_\\zT}}_{\\mu \\nu}p^\\mu=0,\\qquad {S_{\\gr{e}_\\zT}}_{\\mu \\nu}n^\\nu=0,\\qquad {S_{\\gr{e}_\\zT}}_{\\mu}^{\\,\\,\\nu}{S_{\\gr{e}_\\zT}}_{\\nu\\sigma}={S_{\\gr{e}_\\zT}}_{\\mu\\sigma}\\,. \\ee In this paper, a projected tensorial quantity is orthogonal to $\\gr{e}_\\zT$ only and a screen projected quantity is orthogonal to both $\\gr{n}$ and $\\gr{e}_\\zT$. If there is no risk of confusion about the choice of the observer in this screen projector, then we will use the notation $S_{\\mu \\nu}$ instead of ${S_{\\gr{e}_\\zT}}_{\\mu\\nu}$. If we also omit the dependence of the screen projector in the photon momentum, then we abbreviate $S_{\\mu \\nu}(\\gr{p})$ in $S_{\\mu \\nu}$, and $S_{\\mu \\nu}(\\gr{p}')$ in $S'_{\\mu \\nu}$. The polarization of a photon is represented by its polarization vector $\\mybepsilon(\\gr{p})$ which is a complex spacelike unit vector ($\\epsilon_\\mu^\\star (\\gr{p})\\epsilon^\\mu(\\gr{p})=1$) and taken in the Lorentz gauge ($\\epsilon_\\mu(\\gr{p}) p^\\mu =0$). Since there is a residual gauge freedom (of electromagnetism) in the choice of the polarization vector, we will work with the screen projected polarization vector $S^\\mu_{\\,\\,\\nu}\\epsilon^\\nu(\\gr{p})$ which is unique.  \\subsection{The (screen-projected) polarization tensor}  The radiation is represented by a Hermitian tensor valued distribution function (which is thus complex valued), also called polarization tensor~\\cite{Bildhauer1989,Kosowsky1996,Tsagas2007,Stebbins2007} satisfying \\be F_{\\mu \\nu}(x^\\aB,p^\\aT)\\qquad p^\\mu F_{\\mu\\nu}(x^\\aB,p^\\aT)=0\\,, \\ee from which we can form the distribution function of photons in a given state of polarization $\\mybepsilon$ by \\be F_{\\mu \\nu}(x^\\aB,p^\\aT){\\epsilon^\\star}^\\mu(\\gr{p}) \\epsilon^\\nu(\\gr{p})\\,. \\ee Here the $x^\\aB$ are the coordinates used to label points on the space-time manifold. Throughout this paper, indices which refer to these coordinates are $\\aB,\\bB,\\cB,\\dots$ if they run from $0$ to $3$. Furthermore, indices which are $\\iB,\\jB,\\kB,\\dots$ run from $1$ to $3$, and the time component index is $\\zB$.\\\\ For a given electromagnetic plane wave with potential vector amplitude $A^\\mu=A \\epsilon^\\mu$ and wave vector $k^\\mu$ in the geometric optics limit, this polarization tensor can be defined~\\cite{MTW} by \\be\\label{Def_buildpolartensor} \\dirac{1}(\\gr{p}.\\gr{p})F_{\\mu \\nu}(p^\\aT) \\equiv \\frac12 (2\\pi)^3 \\dirac{4} (\\gr{k} -\\gr{p})  A^2 \\epsilon_\\mu(\\gr{k}) \\epsilon^\\star_\\nu(\\gr{k}) \\,, \\ee where $\\dirac{n}$ is the Dirac function of dimension $n$. This tensor $F_{\\mu\\nu}$ should not be confused with the Faraday tensor. Since the remaining gauge freedom of electromagnetism also affects the polarization tensor, we can define the screen-projected distribution function by \\be f_{\\mu\\nu}(x^\\aB,p^\\aT)=S_\\mu^\\rho S_\\nu^\\sigma F_{\\rho \\sigma}(x^\\aB,p^\\aT)\\,. \\ee This tensor is no more dependent on the residual gauge freedom and encodes all the polarization properties of radiation as seen by an observer having a velocity $\\gr{e}_\\zT$. Similarly to the definition of the screen projector, if the context requires it, we will use an index notation of the type $f_{\\gr{e}_\\zT}^{\\mu\\nu}$ to remind with which velocity, and thus with which screen projector, it is defined. The screen projected tensor has four degrees of freedom which can be split according to \\be f_{\\mu \\nu}(x^\\aB,p^\\aT)\\equiv\\frac{1}{2} I(x^\\aB,p^\\aT)S_{\\mu \\nu}+P_{\\mu \\nu}(x^\\aB,p^\\aT)+\\frac{\\ii}{2} V(x^\\aB,p^\\aT) \\epsilon_{\\mu \\nu \\sigma} n^\\sigma, \\ee with $\\epsilon_{\\mu \\nu \\sigma}\\equiv e_\\zT^\\rho\\epsilon_{\\rho \\mu \\nu \\sigma}$, and  where $\\epsilon_{\\rho  \\mu \\nu \\sigma}$ is the completely antisymmetric tensor. $P_{\\mu \\nu}$, which encodes the degree of linear polarization, is real symmetric and trace free, as well as orthogonal to $\\gr{e}_\\zT$ and $\\gr{n}$, and thus encodes two degrees of freedom, which can be described by the Stokes functions $Q$ and $U$~\\cite{Tsagas2007}. $I(x^\\aB,p^\\aT)$ and $V(x^\\aB,p^\\aT)$ are respectively the intensity (or distribution function) for both polarizations and the degree of circular polarization. We also define the normalized (screen-projected) polarization tensor by \\be U_{\\mu \\nu}\\equiv \\frac{f_{\\mu\\nu}}{f^{\\alpha}_{\\,\\,\\alpha}}=\\frac{f_{\\mu\\nu}}{I}\\,\\,. \\ee  \\subsection{Stress-energy tensor and energy-integrated functions}  For a distribution of photons, a stress-energy tensor can be defined by \\be T^{\\mu \\nu}(x^\\aB)\\equiv e_\\aT^{\\,\\mu} e_\\bT^{\\,\\nu} \\left(\\int \\dirac{1}(\\gr{p}.\\gr{p})I(x^\\aB,p^\\cT) p^\\aT p^\\bT \\frac{\\dd p^\\zT \\dd^3 p^{\\iT}}{(2 \\pi)^3}\\right). \\ee Performing the integral over $p^\\zT$ (we choose the convention in which we integrate on the two mass hyperboloides), we eliminate the Dirac function and it leads to \\be T^{\\mu \\nu}(x^\\aB)\\equiv e_\\aT^{\\,\\mu} e_\\bT^{\\,\\nu} \\left(\\int I(x^\\aB,p^\\iT) p^\\aT p^\\bT \\frac{\\dd^3 p^{\\iT}}{p^\\zT (2\\pi)^3}\\right)\\,, \\ee where now the intensity distribution function has to be considered as a function of $p^\\iT$, and $p^\\zT$ is positive and taken on the mass shell, that is $p^\\zT=\\sqrt{p_\\iT p^\\iT}$, and is thus identified with the energy of the photon. Splitting the integral over $\\dd^3 p^\\iT$ into magnitude and angular direction leads to \\be\\label{Def_Tmunu3} T^{\\mu \\nu}(x^\\aB)\\equiv \\frac{1}{(2\\pi)^3}e_\\aT^{\\,\\mu} e_\\bT^{\\,\\nu} \\left(\\int I(x^\\aB,p^\\iT) \\left(p^\\zT\\right)^3 N^\\aT N^\\bT \\dd p^\\zT \\dd^2 \\Omega\\right)\\,, \\ee where $\\dd^2 \\Omega$ is the differential solid angle associated with the unit vector $n^\\iT$, and where $N^\\iT \\equiv n^\\iT$ and $N^\\zT\\equiv 1$. This motivates the definition of the (energy-) integrated counterparts of $I$, $V$ and $P_{\\mu \\nu}$ which are \\bea\\label{Brightnessdef} {\\cal I}(x^\\aB,n^\\iT)&\\equiv &\\frac{4 \\pi}{(2 \\pi)^3} \\int I(x^\\aB,p^\\zT,n^\\iT)(p^\\zT)^3 \\dd p^\\zT\\,,\\\\* {\\cal V}(x^\\aB,n^\\iT)&\\equiv& \\frac{4 \\pi}{(2 \\pi)^3} \\int V(x^\\aB,p^\\zT,n^\\iT)(p^\\zT)^3 \\dd p^\\zT\\,,\\\\* {\\cal P}_{\\mu \\nu}(x^\\aB,n^\\iT)&\\equiv& \\frac{4 \\pi}{(2 \\pi)^3} \\int P_{\\mu \\nu}(x^\\aB,p^\\zT,n^\\iT)(p^\\zT)^3 \\dd p^\\zT\\,. \\eea $\\cal I$ is the brightness, ${\\cal V}$ is the degree of linear polarization in units of ${\\cal I}$, and ${\\cal P}_{\\mu \\nu}$ is the tensor of linear polarization in units of ${\\cal I}$.  \\subsection{Description of massive particles}  For massive particles such as electrons ($\\ie$), protons ($\\ip$) or cold dark matter ($\\ic$), we do not need to describe polarization and thus we can rely solely on a the distribution functions $\\gel$, $g_{\\mathrm p}$ and $g_{\\mathrm c}$ (chosen to describe the two helicities). Additionally, following common practice in cosmology, we will refer to electrons and protons together as baryons though electrons are leptons. This is motivated by the fact that most of the mass is carried by protons which are baryons, and the Compton interaction between protons and electrons makes these two components highly interdependent. For a particle with impulsion $q^\\aT$, we use the following notation \\be \\flamoi n^\\iT\\equiv \\frac{q^\\iT}{q^\\zT}\\,,\\quad \\beta=\\sqrt{n^\\iT n_\\iT}\\,,\\quad\\lambda \\equiv \\beta q^\\zT\\,, \\quad \\hat n^\\iT= n^\\iT/\\beta\\,, \\quad \\gamma= \\left(1 -\\beta^2 \\right)^{-1/2}\\,, \\ee where now we have to distinguish between the unit vector $\\hat n^\\iT$ and the velocity vector $n^\\iT$ because of the mass $m$ of the particles. The stress-energy tensor can be defined in a similar manner to equation~(\\ref{Def_Tmunu3}) by \\be T^{\\mu \\nu}(x^\\aB)\\equiv e_\\aT^{\\,\\mu} e_\\bT^{\\,\\nu} \\left(\\int \\dirac{1}(\\gr{q}.\\gr{q}-m^2) g(x^\\aB,q^\\cT) q^\\aT q^\\bT \\frac{\\dd q^\\zT \\dd^3 q^{\\iT}}{(2 \\pi)^3}\\right). \\ee Integrating over $q^\\zT$, this leads to \\be\\label{Def_Tmunu4} T^{\\mu \\nu}(x^\\aB)\\equiv \\frac{1}{(2\\pi)^3}e_\\aT^{\\,\\mu} e_\\bT^{\\,\\nu} \\left(\\int g(x^\\aB,q^\\iT) q^\\zT N^\\aT N^\\bT \\dd^3 q^\\iT\\right)\\,, \\ee where we recall that $N^\\aT \\equiv (1,n^\\iT)$, and $q^\\zT$ is taken on the mass shell ($q^\\zT=\\sqrt{q_\\iT q^\\iT+m^2}$).\\\\  \\subsection{The fluid limit}\\label{Fluidlimit}  The stress-energy tensor of radiation or matter is equivalent to the one of an imperfect fluid with stress-energy tensor \\begin{equation} T_{\\mu\\nu} = \\rho u_\\mu u_\\nu + P\\left( g_{\\mu\\nu} + u_\\mu u_\\nu\\right) +\\Pi_{\\mu \\nu}\\,. \\label{defstressenergytensor} \\end{equation} In this decomposition $\\rho$ is the energy density, $P$ is the pressure, $u^\\mu$ is the fluid velocity and $\\Pi^{\\mu \\nu}$, which satisfies $\\Pi^{\\mu}_{\\,\\,\\mu}=u^\\mu \\Pi_{\\mu \\nu}=0$, is the anisotropic stress. For an isotropic distribution of radiation, that is where ${\\cal I}(x^\\aB,p^\\iT)$ depends only on the magnitude of $p^\\iT$, it is straightforward to show by comparison of the expressions~(\\ref{Def_Tmunu3}) and (\\ref{defstressenergytensor}) that $P=\\rho/3$. Similarly, for a set of heavy particles (or non-relativistic particles), that is with $\\sqrt{q^\\iT q_\\iT} \\ll m$, the pressure satisfies $P\\ll \\rho$ and the anisotropic stress tensor is also similarly small. For cold dark matter, it is assumed that the mass of particles is large enough so that we can use this approximation.\\\\ In the case of electrons and protons, that is baryons, the Coulomb interaction ensures that the distribution of momenta follows a Fermi-Dirac distribution in the reference frame where they have no bulk velocity~\\cite{Bernstein1988}, at least as long as the baryonic matter is ionized. This distribution is isotropic in this adapted frame and depends only on $\\lambda$ \\be g_{\\mathrm{fd}}(\\lambda)=\\left(\\exp\\left[\\frac{\\sqrt{(\\lambda^2+m^2)}-\\mu}{T}\\right]+1\\right)^{-1}\\,\\,\\,. \\ee As a consequence, the anisotropic stress vanishes, and in this adapted frame the baryons are then ideally described by the energy density and the pressure \\bea \\rho_\\ib&\\equiv& \\frac{4\\pi}{(2\\pi)^3}\\int g_{\\mathrm{fd}}(\\lambda)\\sqrt{\\lambda^2+m^2}\\lambda^2 \\dd \\lambda\\,,\\\\* P_\\ib&\\equiv& \\frac{4\\pi}{3(2\\pi)^3}\\int g_{\\mathrm{fd}}(\\lambda)\\frac{\\lambda^4}{\\sqrt{\\lambda^2+m^2}} \\dd \\lambda\\,. \\eea If we can neglect the chemical potential $\\mu$, which is the case in the cosmological context, then for non-relativistic particles we obtain \\bea \\rho_\\ib&\\simeq&m\\left(1+\\frac{3T}{2m}\\right) n_\\ib\\,,\\\\* P_\\ib&\\simeq&n_\\ib T\\,. \\eea The baryonic matter is ionized roughly until recombination where the temperature of photons is of order $T_{\\mathrm{LSS}} \\simeq 0.25\\,\\, \\mathrm{eV}$~\\cite{Komatsu2008}. For electrons, the thermal correction is of order $T_{\\mathrm{LSS}}/\\me\\simeq 0.25/511000\\simeq 0.5\\,\\, 10^{-6}$, and this ratio is even $m_{\\ip}/m_{\\ie}\\simeq 1836$ times smaller for protons. We thus deduce that the baryons either have no anisotropic stress because of Coulomb interaction, or have a very tiny pressure and anisotropic stress after recombination has occurred. However, as discussed in section~\\ref{Sec_pertscheme} this thermal correction is still of order of the metric perturbations and should not {\\it in principle} be ignored for second order computations. Nevertheless, the thermal corrections are not relevant for computing the bispectrum and it is for this reason that we will, from now on, drop terms in $T/m$ and describe baryons as cold matter (but not dark since it can interact with radiation). It should be mentionned that for neutrinos these conclusions are not valid anymore since they are very light~\\cite{Lesgourgues2006}. We will here assume that they are light enough to be treated as collisionless radiation, and the equations which govern their evolution can be found by setting $\\st=0$ in the equations for photons, where $\\st$ is the Thomson cross section.  \\subsection{Multipole expansion for radiation}  Functions of $p^\\aT$ can be viewed as functions of $(p^\\zT,n^\\aT)$ and we can separate the dependence into the energy and the direction of the momentum. The dependence in the direction can be further expanded in multipoles using projected symmetric trace-free (PSTF) tensors, where projected means that they are orthogonal to $\\gr{e}_\\zT$. For instance, $I$ can be expanded as \\be\\label{Def_multipoleS} I(x^\\aB,p^\\zT,n^\\aT)=\\sum_{\\ell=0}^\\infty I_{\\uline{\\aT_\\ell}}(x^\\aB,p^\\zT)n^{\\uline{\\aT_\\ell}}\\,, \\ee where the $I_{\\uline{\\aT_\\ell}}\\equiv I_{\\aT_1\\dots\\aT_\\ell}$ are PSTF. For the lowest multipole, i.e. the one corresponding to $\\ell=0$, we use the notation $I_{\\emptyset}$. Note that we have defined the notation $n^{\\uline{\\aT_\\ell}}\\equiv n^{\\aT_1}\\dots n^{\\aT_\\ell}$. We also remind that $n^\\aT\\equiv \\gr{n}.\\gr{e}^\\aT\\equiv n^\\mu e^\\aT_{\\,\\mu}$. Since $\\gr{n}$ is projected, $n^\\zT=0$ and thus if any of the indices in $I_{\\uline{\\aT_\\ell}}$ is $\\zT$, then the multipoles is chosen to vanish. These multipoles can be obtained by performing the following integrals \\be \\flamoi I_{\\uline{\\aT_\\ell}}(x^\\aB,p^\\zT)=\\Delta_\\ell^{-1}\\int  I(x^\\aB,p^\\zT,n^\\aT) n_{\\langle \\uline{\\aT_\\ell}\\rangle}\\dd^2 \\Omega\\,,\\qquad \\Delta_\\ell\\equiv 4 \\pi \\frac{\\ell!}{(2\\ell+1)!!}\\,\\,\\,, \\ee where $\\langle \\dots\\rangle$ means the symmetric trace-free part. A similar expansion can be performed on $V$, by replacing $I$ by $V$ in the above expressions, as well as for their energy integrated counterparts. Note that in particular \\bea {\\cal I}^{\\emptyset}&=&T^{\\zT \\zT}\\,,\\\\* {\\cal I}^\\iT&=&4 \\pi \\Delta_1^{-1} T^{\\zT \\iT}\\,,\\\\* {\\cal I}^{\\iT \\jT}&=&4 \\pi  \\Delta_2^{-1} T^{\\langle \\iT \\jT \\rangle}\\,. \\eea As for $P_{\\aT \\bT}\\equiv P_{\\mu \\nu} e_\\aT^{\\,\\mu}e_\\bT^{\\,\\nu}$, it can be expanded in electric and magnetic type components according to~\\cite{Challinor2000,Tsagas2007,Dautcourt1978,Thorne1980} \\be\\label{Def_multipoleP} \\flamoi P_{\\aT \\bT}(x^\\aB,p^\\aT)=\\sum_{\\ell=2}^\\infty \\left[E_{\\aT \\bT \\uline{\\cT_{\\ell-2}}}(x^\\aB,p^\\zT)n^{ \\uline{\\cT_{\\ell-2}}} \\,- n_{\\cT}\\epsilon^{\\cT \\dT}_{\\,\\,\\,\\,\\,(\\aT}B_{\\bT)\\dT \\uline{\\cT_{\\ell-2}}}(x^\\aB,p^\\zT)n^{\\uline{\\cT_{\\ell-2}}}\\right]^{\\mathrm{TT}}, \\ee where the notation $\\mathrm{TT}$ denotes the transverse (to $\\gr{n}$) symmetric trace-free part, which for a second rank tensor is \\be \\left[X_{\\aT\\bT}\\right]^{\\mathrm{TT}}\\equiv S_{(\\aT}^{\\,\\,\\cT} S_{\\bT)}^{\\,\\,\\dT}X_{\\cT \\dT}-\\frac{1}{2}S_{\\aT\\bT} S^{\\cT\\dT}X_{\\cT \\dT}\\,. \\ee The electric and magnetic multipoles can be obtained by performing the integrals \\bea E_{\\uline{\\aT_\\ell}}(x^\\aB,p^\\zT)&=&M_\\ell^2\\Delta_\\ell^{-1}\\int \\,n_{\\langle \\uline{\\aT_{\\ell-2}}}P_{\\aT_{\\ell-1}\\aT_{\\ell}\\rangle}(x^\\aB,p^\\zT,n^\\aT)\\dd^2 \\Omega\\,\\,,\\\\* B_{\\uline{\\aT_\\ell}}(x^\\aB,p^\\zT)&=&M_\\ell^2\\Delta_\\ell^{-1}\\int \\,n_\\bT \\epsilon^{\\bT \\dT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}n_{\\uline{\\aT_{\\ell-2}}}P_{\\aT_{\\ell-1}\\rangle \\dT}(x^\\aB,p^\\zT,n^\\aT)\\dd^2 \\Omega\\,\\,, \\eea where \\be M_\\ell=\\sqrt{\\frac{2 \\ell (\\ell-1)}{(\\ell+1)(\\ell+2)}}\\,\\,. \\ee \\subsection{Transformation rules under a change of frame}\\label{Sec_Changeframe}  \\subsubsection{The photon momentum}  So far, everything was defined with respect to an observer having a velocity $\\gr{e}_\\zT$. How does all this machinery transform when the radiation is observed by an observer with a different velocity ${\\tilde{\\gr{e}}}_{\\zT}$? Since two velocities can be related by a Lorentz transformation, there exists a vector $\\gr{v}$ such that \\be {\\tilde{\\gr{e}}}_{\\zT}=\\gamma(\\gr{e}_\\zT+\\gr{v}),\\qquad\\gamma\\equiv\\frac{1}{\\sqrt{1-\\gr{v}.\\gr{v}}}\\,\\qquad v^\\zT\\equiv\\gr{v}.\\gr{e}^\\zT=0\\,\\,. \\ee We deduce immediately that the magnitude and the direction unit vector of the photon momentum transform as \\bea\\label{Eq_Tmomentum1} p^{\\zTt}&\\equiv& \\gr{p}.\\tilde{\\gr{e}}^{\\zT}=\\gamma p^\\zT\\left(1-\\gr{n}.\\gr{v} \\right)\\,,\\\\* {\\tilde{\\gr{n}}}&=&\\frac{1}{\\gamma(1- \\gr{v}.\\gr{n})}(\\gr{e}_\\zT+\\gr{n})-\\gamma(\\gr{e}_\\zT + \\gr{v})\\,. \\eea We remind that the direction, as observed by the transformed observer, is given by the decomposition $ \\gr{p}=p^{\\zTt} ({\\tilde{\\gr{e}}}_{\\zT} +{\\tilde{\\gr{n}}})$. These rules imply the following transformation rule for the screen projector \\be \\tilde S_{\\mu \\nu}=S_{\\mu \\nu}+\\frac{2 \\gamma}{p^{\\zTt}}p_{(\\mu}S_{\\nu)\\rho}v^\\rho+\\left(\\frac{\\gamma}{p^{\\zTt}}\\right)^2 p_\\mu p_\\nu S_{\\alpha \\beta}v^\\alpha v^\\beta\\,, \\ee which implies directly the following useful relations \\bea\\label{TruleforS} \\tilde S_{\\mu \\nu}=\\tilde S_\\mu^{\\,\\,\\rho} \\tilde S_\\nu^{\\,\\,\\sigma} S_{\\rho \\sigma}\\,,\\qquad \\tilde n^\\mu \\tilde \\epsilon_{\\mu \\rho \\sigma}=n^\\mu \\epsilon_{\\mu \\alpha \\beta} \\tilde S^\\alpha_{\\,\\,\\rho}\\tilde S^\\beta_{\\,\\,\\sigma}\\,. \\eea The last relation can also be demonstrated easily by noting that $n^\\mu \\epsilon_{\\mu \\alpha \\beta} \\tilde S^\\alpha_{\\,\\,\\rho}\\tilde S^\\beta_{\\,\\,\\sigma}$ is by construction orthogonal to $\\tilde{\\gr{e}}_\\zT$ and ${\\tilde{\\gr{n}}}$, and since it is also obviously antisymmetric in its two free indices, it has to be proportional to $\\tilde n^\\mu \\tilde \\epsilon_{\\mu \\rho \\sigma}$. By contracting both expressions with themselves we obtain that they are indeed equal.\\\\ The rest of the tetrad can be transformed without rotation, that is with a pure boost, by \\be \\tilde e^{\\aT} = \\Lambda^\\aT_{\\,\\,\\bT}e^\\bT,\\qquad\\tilde e_\\aT = e_\\bT \\Lambda^\\bT_{\\,\\,\\aT}\\,, \\ee and the components of this transformation are given by \\be \\Lambda^\\zT_{\\,\\zT}=\\gamma,\\qquad \\Lambda^\\zT_{\\,\\iT}=-\\gamma v_\\iT,\\qquad \\Lambda^\\iT_{\\,\\jT}=\\delta^\\iT_\\jT+\\frac{\\gamma^2}{1+\\gamma}v^\\iT v_\\jT, \\ee where we remind that $v^\\iT$ is the component of $\\gr{v}$ along the tetrad $\\gr{e}^\\iT$, that is $v^\\iT=\\gr{v}.\\gr{e}^\\iT $ . The transformation rule for the photon direction when expressed along tetrads is thus \\be \\tilde n^{\\iTt}\\equiv{\\tilde{\\gr{n}}}. {\\tilde{\\gr{e}}}^\\iT=\\frac{1}{\\gamma(1-\\gr{n}.\\gr{v})}\\left[n^\\iT+\\frac{\\gamma^2}{(1+\\gamma)}\\gr{n}.\\gr{v} \\,v^\\iT-\\gamma v^\\iT \\right]. \\ee \\subsubsection{The radiation multipoles}\\label{Transfomultipoles}  It can be easily checked that for a vector orthogonal to $\\gr{p}$, such as the polarization vector, $\\tilde S^\\mu_{\\,\\,\\nu}S^\\nu_{\\,\\,\\sigma}\\epsilon^\\sigma=\\tilde S^\\mu_{\\,\\,\\sigma}\\epsilon^\\sigma$. As an immediate consequence, we deduce from equation~(\\ref{Def_buildpolartensor}) that the screen-projected polarization tensor transforms according to \\be\\label{Eq_Tpropertydebase} \\tilde f_{\\mu \\nu}(x^\\aB,p^{\\zTt},\\tilde n^\\aTt)=\\tilde S^\\mu_{\\,\\,\\alpha} \\tilde S^\\nu_{\\,\\,\\beta} f_{\\alpha \\beta}(x^\\aB,p^{\\zT},n^\\aT)\\,. \\ee We deduce from equation~(\\ref{TruleforS}) that $P_{\\mu \\nu}$ transforms following the same rule, whereas $I$ and $V$ transform as scalars \\be \\tilde I(p^{\\zTt},\\tilde n^\\aTt)=I(p^{\\zT},n^\\aT)\\,,\\qquad\\tilde V(p^{\\zTt},\\tilde n^\\aTt)=V(p^{\\zT},n^\\aT)\\,. \\ee Here and in the rest of this paper, we omit the dependence in the position $x^\\aB$ to simplify the notation. We can deduce from equation~(\\ref{Eq_Tmomentum1}) that the differential solid angle transforms according to (this can also be deduced from using the transformation rule of $p^\\zT$ and the fact that $\\dd^3 p^\\iT/p^\\zT =p^\\zT \\dd p^\\zT \\dd^2 \\Omega$ is a scalar) \\be \\dd \\tilde \\Omega =\\left[\\frac{1}{\\gamma(1-\\gr{v}.\\gr{n})}\\right]^2\\dd \\Omega\\,, \\ee and this can be used to deduce the transformation rules of the multipoles \\be\\label{Eqinttochangeframe} \\flamoi\\tilde I_{\\uline{\\tilde{\\aT}_\\ell}}(p^{\\zTt})=\\Delta_\\ell^{-1}\\int \\dd \\Omega \\left[\\gamma(1-\\gr{v}.\\gr{n})\\right]^{-2}\\sum_{\\ell'=0}^\\infty I_{\\uline{\\bT_{\\ell'}}}[p^{\\zTt}\\gamma^{-1}(1-n^\\cT v_\\cT)^{-1}]n^{\\uline{\\bT_{\\ell'}}} \\tilde n_{\\langle \\uline{\\tilde{\\aT}_\\ell}\\rangle}\\,. \\ee In the previous integral, $I_{\\uline{\\bT_{\\ell'}}}[p^{\\zTt}\\gamma^{-1}(1-n^\\cT v_\\cT)^{-1}]$ has to be considered as a function of the direction $n^\\aT$. It is thus necessary to Taylor expand it as \\be I_{\\uline{\\bT_\\ell}}[p^{\\zTt}\\gamma^{-1}(1-n^\\cT v_\\cT)^{-1}]=\\sum_{n=0}^\\infty\\frac{1}{n!}\\left( \\frac{\\gamma^{-1}n^\\cT v_\\cT}{1-\\gamma^{-1}n^\\cT v_\\cT}\\right)^n I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^{\\zTt}\\gamma^{-1}]\\,, \\ee where we define \\be I^{\\{n\\}}_{\\uline{\\bT_\\ell}}(p^\\zT)\\equiv (p^\\zT)^n\\frac{\\partial^n I_{\\uline{\\bT_\\ell}}(p^\\zT)}{\\partial (p^\\zT)^n}\\,, \\ee with the conventions $I'_{\\uline{\\bT_\\ell}} \\equiv I^{\\{1\\}}_{\\uline{\\bT_\\ell}}$ and $I''_{\\uline{\\bT_\\ell}} \\equiv I^{\\{2\\}}_{\\uline{\\bT_\\ell}}$. Under this form, it is then possible to perform the integration using the following well known integrals~\\cite{Uzan1998} \\ifcqg \\be\\label{Intonn} \\flamoi \\int n^{\\iT_1}\\dots n^{\\iT_k}\\frac{\\dd^2 \\Omega}{4 \\pi} = \\cases{0\\,, &if\\;\\;$k= 2p + 1$,\\\\\\frac{1}{k+1}\\left[\\delta^{(\\iT_1 \\iT_2}\\dots\\delta^{\\iT_{(k-1)}\\iT_k)}\\right]\\,, &if\\;\\;$k= 2p$.} \\ee \\else \\be\\label{Intonn} \\flamoi \\int n^{\\iT_1}\\dots n^{\\iT_k}\\frac{\\dd^2 \\Omega}{4 \\pi} = \\begin{cases}0\\,, &if\\;\\;$k= 2p + 1$,\\\\\\frac{1}{k+1}\\left[\\delta^{(\\iT_1 \\iT_2}\\dots\\delta^{\\iT_{(k-1)}\\iT_k)}\\right]\\,, &if\\;\\;$k= 2p$.\\end{cases} \\ee \\fi Note that we have used the standard notation $(\\dots)$ for the symmetrization of indices which leaves unchanged an symmetric tensor and we will also use the notation $[\\dots]$ for the antisymmetrization of indices which leaves unchanged an antisymmetric tensor. The integrals~(\\ref{Intonn}) used in the transformation rule~(\\ref{Eqinttochangeframe}) are ideally suited for a tensor calculus package, and we used \\emph{xAct}~\\cite{xAct} to compute them. \\\\ Since the result of equation~(\\ref{Eqinttochangeframe}) will involve terms like $I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^{\\zTt}\\gamma^{-1}]=I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^\\zT(1-n^\\cT v_\\cT)]$, it will require an additional Taylor expansion in order to have the result expressed only in function of the $I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^\\zTt]$ or $I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^\\zT]$. Note that the expression of $\\tilde I_{\\uline{\\tilde{\\aT}_\\ell}}(p^{\\zTt})$ in function of the $I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^\\zTt]$, which is the choice that we make in the expressions that we report below, does not depend on the direction $n^\\iT$, whereas it does when expressed in function of the $I^{\\{n\\}}_{\\uline{\\bT_\\ell}}[p^\\zT]$ since $p^\\zT$ is unambiguously defined from $p^\\zTt$ only once a direction $n^\\iT$ is specified.  A similar method, with similar definitions can be used to determine the transformation rules of the electric and magnetic multipoles~\\cite{Tsagas2007}. At first order in the velocity $\\gr{v}$, we obtain the following transformation rules \\bea \\flamoi\\tilde I_{ \\uline{\\aTt_\\ell}}(p^{\\zTt})&=&I_{\\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell+2)(\\ell+1)}{(2\\ell+3)}v^\\bT I_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell+1)}{(2\\ell+3)}v^\\bT I'_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})\\nonumber\\\\* \\flamoi&&-(\\ell-1)v_{\\langle \\aT_\\ell}I_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt}) +v_{\\langle \\aT_\\ell}I'_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt})\\,, \\eea  \\bea \\flamoi\\tilde E_{ \\uline{\\aTt_\\ell}}(p^{\\zTt})&=&E_{\\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell-1)(\\ell+2)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT E_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell-1)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT E'_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})\\nonumber\\\\* \\flamoi&&-(\\ell-1)v_{\\langle \\aT_\\ell}E_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt}) +v_{\\langle \\aT_\\ell}E'_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt})\\nonumber\\\\* \\flamoi&&-\\frac{2}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}B_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}(p^{\\zTt})-\\frac{2}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}B'_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}(p^{\\zTt})\\,, \\eea  \\bea \\flamoi\\tilde B_{\\uline{\\aTt_\\ell}}(p^{\\zTt})&=&B_{\\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell-1)(\\ell+2)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT B_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})+\\frac{(\\ell-1)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT B'_{\\bT \\uline{\\aT_\\ell}}(p^{\\zTt})\\nonumber\\\\* \\flamoi&&-(\\ell-1)v_{\\langle \\aT_\\ell}B_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt}) +v_{\\langle  \\aT_\\ell}B'_{\\uline{\\aT_{\\ell-1}}\\rangle}(p^{\\zTt})\\nonumber\\\\* \\flamoi&&+\\frac{2}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}E_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}(p^{\\zTt})+\\frac{2}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}E'_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}(p^{\\zTt})\\,. \\eea As for the energy-integrated counterparts, they transform at first order in $\\gr{v}$ as \\be \\flamoi\\tilde {\\cal I}_{\\uline{\\aTt_\\ell}}={\\cal I}_{\\uline{\\aT_\\ell}}+\\frac{(\\ell-2)(\\ell+1)}{(2\\ell+3)}v^\\bT {\\cal I}_{\\bT \\uline{\\aT_\\ell}}-(\\ell+3)v_{\\langle \\aT_\\ell}{\\cal I}_{\\uline{\\aT_{\\ell-1}}\\rangle}\\,, \\ee  \\be \\flamoi\\tilde {\\cal E}_{ \\uline{\\aTt_\\ell}}={\\cal E}_{\\uline{\\aT_\\ell}}-\\frac{(\\ell-1)(\\ell-2)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT {\\cal E}_{\\bT \\uline{\\aT_\\ell}}-(\\ell+3)v_{\\langle \\aT_\\ell}{\\cal E}_{\\uline{\\aT_{\\ell-1}}\\rangle}-\\frac{6}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}{\\cal B}_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}\\,, \\ee and \\be \\flamoi\\tilde {\\cal B}_{\\uline{\\aTt_\\ell}}={\\cal B}_{\\uline{\\aT_\\ell}}-\\frac{(\\ell-1)(\\ell-2)(\\ell+3)}{(\\ell+1)(2\\ell+3)}v^\\bT {\\cal B}_{\\bT \\uline{\\aT_\\ell}}-(\\ell+3)v_{\\langle \\aT_\\ell}{\\cal B}_{\\uline{\\aT_{\\ell-1}}\\rangle}+\\frac{6}{(\\ell+1)}v_\\bT \\epsilon^{\\bT \\cT}_{\\,\\,\\,\\,\\,\\langle \\aT_\\ell}{\\cal E}_{\\uline{\\aT_{\\ell-1}}\\rangle \\cT}\\,. \\ee We report in~\\ref{AppTrule2} the transformation rules up to second order in $\\gr{v}$ for multipoles of further interest.  \\subsection{The Liouville equation}  In this section, we present the equation which governs the free-streaming of photons directly in the tetrad basis though nothing prevents it from being expressed with formal indices. The evolution of the polarization tensor is dictated by the Boltzmann equation \\be\\label{EqBoltzmann} L[f_{\\aT \\bT}(x^\\aB,p^\\zT,n^\\iT)]=C_{\\aT \\bT}(x^\\aB,p^\\zT,n^\\iT)\\,. \\ee $L$ is the Liouville operator whose action on TT tensors like the screen-projected polarization tensor $f_{\\aT \\bT}$ is given by~\\cite{Tsagas2007} \\be \\flamoi L[f_{\\aT \\bT}(x^\\aB,p^\\aT)]\\equiv S_\\aT^{\\,\\cT} S_\\bT^{\\,\\dT}\\left[p^\\hT \\nabla_\\hT f_{\\cT \\dT}(x^\\aB,p^\\aT)+ \\frac{\\partial f_{\\cT \\dT}(x^\\aB,p^\\aT)}{\\partial p^\\hT}\\frac{\\dd p^\\hT}{\\dd s} \\right]\\,, \\ee where $s$ is the affine parameter along the particle geodesic. $\\nabla$ is the covariant derivative which in the tetrad basis is related to the partial derivative by \\be \\flamoi p^\\hT \\nabla_\\hT f_{\\cT \\dT}(x^\\aB,p^\\aT)\\equiv p^\\hT \\partial_\\hT f_{\\cT \\dT}(x^\\aB,p^\\aT)+p^\\hT\\omega_{\\hT\\cT}^{\\,\\,\\,\\,\\,\\,\\bT}f_{\\bT \\dT}(x^\\aB,p^\\aT)+p^\\hT\\omega_{\\hT\\dT}^{\\,\\,\\,\\,\\,\\,\\bT}f_{\\cT \\bT}(x^\\aB,p^\\aT)\\,, \\ee where the $\\omega_{\\aT\\bT\\cT}$ are the Ricci rotation coefficients (see \\ref{app_TRS} and \\cite{Wald1984} for details). $C_{\\aT \\bT}$ is the collision tensor whose expression will be detailed in section~\\ref{Sec_Collision}.\\\\ In the case where the collision tensor can be ignored, that is when the collision of photons with electrons or protons can be neglected [the latter type of collision can be always ignored compared to the former since its cross section is reduced by a factor $(m_\\ie/m_\\ip)^2$], this reduces to the Liouville equation. The Liouville equation arises from the fact that the polarization vector $\\mybepsilon$ of a photon is parallel transported and thus satisfies $p^\\mu \\nabla_\\mu \\epsilon^\\nu=0$, and then it follows from the construction~(\\ref{Def_buildpolartensor}) of the (not screen projected) polarization tensor $F_{\\mu \\nu}$ that \\be p^\\hT \\nabla_\\hT F_{\\cT \\dT}(x^\\aB,p^\\aT)+ \\frac{\\partial F_{\\cT \\dT}(x^\\aB,p^\\aT)}{\\partial p^\\hT}\\frac{\\dd p^\\hT}{\\dd s}=0\\,. \\ee By using the expression~(\\ref{DefScreen}) for $S^{\\mu \\nu}$ and the property $F_{\\mu \\nu}(x^\\aB,p^\\aT)p^\\mu=0$, we obtain directly from the previous equation that the screen-projected tensor satisfies $L[f_{\\aT \\bT}(x^\\aB,p^\\aT)]=0$. It can also be shown that the Liouville operator preserves the decomposition of $f_{\\mu \\nu}$ in an antisymmetric part ($V$), a trace ($I$) and a symmetric traceless part ($P_{\\mu \\nu}$), that is \\be \\flamoi L[f_{\\aT \\bT}(x^\\aB,p^\\aT)]=\\frac{1}{2}L[I(x^\\aB,p^\\dT)]S_{\\aT \\bT}+L[P_{\\aT \\bT}(x^\\aB,p^\\aT)]+\\frac{\\ii}{2}L[V(x^\\aB,p^\\dT)]n^\\cT \\epsilon_{\\cT \\aT \\bT}\\,, \\ee with the Liouville operator acting on a scalar valued function like $I$ or $V$ according to \\be L[I(x^\\aB,p^\\aT)]\\equiv p^\\hT \\partial_\\hT I(x^\\aB,p^\\aT)+ \\frac{\\partial I(x^\\aB,p^\\aT)}{\\partial p^\\hT}\\frac{\\dd p^\\hT}{\\dd s}\\,\\,. \\ee To see this, we need only to use the property $\\omega_{\\aT [\\bT \\cT]}=\\omega_{\\aT \\bT \\cT}$ and remark that \\be S^\\cT_{(\\aT}S_{\\,\\,\\bT)}^{\\dT} p^\\hT \\omega_{\\hT\\cT \\fT}S^\\fT_{\\,\\dT}= p^\\hT \\omega_{\\hT\\cT \\fT}S^\\cT_{(\\aT}S^\\fT_{\\,\\bT)}=0\\,, \\ee \\be S^\\cT_{[\\aT}S_{\\,\\,\\bT]}^{\\dT} p^\\hT \\omega_{\\hT\\cT \\fT} \\epsilon^{\\fT}_{\\,\\,\\dT}\\propto S^\\cT_{[\\aT}S_{\\,\\,\\bT]}^{\\dT} \\epsilon_{\\cT \\fT} \\epsilon^{\\fT}_{\\,\\,\\dT}=\\epsilon_{\\fT[\\aT} \\epsilon^{\\,\\,\\fT}_{\\bT]}=0\\,, \\ee where we have used the definition $\\epsilon_{\\mu \\nu} \\equiv n^\\alpha \\epsilon_{\\alpha \\mu\\nu}$.\\\\ Additionally, since the affine parameter $s$ is a scalar, the transformation properties of the Liouville operator and the collision tensor under a local change of frame is the same as the transformation property of $f_{\\mu \\nu}$ given in equation~(\\ref{Eq_Tpropertydebase})~\\cite{Tsagas2007}. ",
        "conclusions": ""
    },
    "0809/0809.1519_arXiv.txt": {
        "abstract": "Dark stars powered by dark matter annihilation have been proposed as the first luminous sources in the universe. These stars are believed to form {in the central dark matter cusp} of low-mass minihalos. {Recent calculations indicate stellar masses up to $\\sim1000\\,M_\\odot$ and/or have very long lifetimes. The UV photons from these objects could therefore contribute significantly to cosmic reionization.} Here we show that such dark star models would require a somewhat artificial reionization history, based on a double-reionization phase and a late star-burst near redshift $z\\sim6$, in order to fulfill the WMAP constraint on the optical depth as well as the Gunn-Peterson constraint at $z\\sim6$. This suggests that, if dark stars were common in the early universe, then models are preferred which predict a number of UV photons similar to conventional Pop.~III stars. This excludes $800\\ M_\\odot$ dark stars that enter a main-sequence phase and other models that lead to a strong increase in the number of UV photons.\\\\ We also derive constraints {for massive as well as light dark matter candidates} from the observed X-ray, gamma-ray and neutrino background, considering dark matter profiles which have been steepened during the formation of dark stars. This increases the clumping factor at high redshift and gives rise to a higher dark matter annihilation rate in the early universe. {{We furthermore estimate the potential contribution from the annihilation products in the remnants of dark stars, which may provide a promising path to constrain such models further, but which is currently still uncertain.}} ",
        "introduction": "Growing astrophysical evidence suggests that dark matter in the universe is self-annihilating. X-ray observations from the center of our Galaxy find bright $511$ keV emission which cannot be attributed to single sources \\citep{Jean06, Weidenspointner06}, but can be well-described assuming dark matter annihilation \\citep{BoehmHooper04}. {Further observations indicate also an excess of GeV photons \\citep{deBoer05}, of microwave photons \\citep{Hooper07}, and of positrons \\citep{Cirelli08}. A common feature of these observations is that the emission seems isotropic and not correlated to the Galactic disk. However, there is usually some discrepancy between the model predictions and the amount of observed radiation, which may be due to uncertainties in the dark matter distribution, astrophysical processes and uncertainties in the model for dark matter annihilation \\citep{deBoer08}.      }  {It is well-known that weakly-interacting massive dark matter particles may provide a natural explanation of the observed dark matter abundance \\citep{Drees93, Kolb90}.} Calculations by \\citet{Ahn05b} indicated that the extragalactic gamma-ray background cannot be explained from astrophysical sources alone, but that also a contribution from dark matter annihilation is needed at energies between $1$-$20$ GeV. {It is currently unclear whether this is in fact the case or if a sufficient amount of non-thermal electrons in active galactic nuclei (AGN) is available to explain this background radiation \\citep{Inoue08}. Future observations with the FERMI satellite \\footnote{http://www.nasa.gov/mission\\_pages/GLAST/science/index.html} will shed more light on such questions and may even distinguish between such scenarios due to specific signatures in the anisotropic distribution of this radiation \\citep{Ando07}. }  {The first stars have been suggested to have high masses of the order $\\sim100\\ M_\\odot$, thus providing powerful ionizing sources in the early universe \\citep{Abel02, Bromm04}. The effect of dark matter annihilation on the first stars has been explored recently in different studies. \\citet{Spolyar08} showed that an equilibrium between cooling and energy deposition from dark matter annihilation can always be found during the collapse of the proto-stellar cloud. This has been explored further by \\citet{Iocco08} and \\citet{FreeseSpolyar08}, who considered the effect of scattering between baryons and dark matter particles, increasing the dark matter abundance in the star. \\citet{IoccoBressan08} considered dark star masses in the range $5\\leq M_*\\leq600\\ M_\\odot$ and calculated the evolution of the pre-main-sequence phase, finding that the dark star phase where the {energy input} from dark matter annihilation dominates may last up to $10^4$~yr. \\citet{FreeseBodenheimer08} examined the formation process of the star in more detail, considering polytropic equilibria and {additional mass accretion} until the total Jeans mass of $\\sim800\\ M_\\odot$ is reached. They find that this process lasts for $\\sim5\\times10^5$~yr. They suggest that dark stars are even more massive than {what is} typically assumed for the first stars, and may be the progenitors for the first supermassive black holes at high redshift. {\\citet{IoccoBressan08}, \\citet{Taoso08} and \\citet{Yoon08} have calculated the stellar evolution for the case in which the dark matter density inside the star is enhanced by the capture of addition WIMPs via off-scattering from stellar baryons. \\citet{IoccoBressan08} followed the stellar evolution until the end of \\HeI burning, \\citet{Yoon08} until the end of oxygen burning and \\citet{Taoso08} until the end of \\HI burning. \\citet{Yoon08} also took the effects of rotation into account.} The calculations found a potentially very long lifetime of dark stars and correspondingly a strong increase in the number of UV photons that may contribute to reionization. Dark stars in the Galactic center have been discussed by \\citet{Scott08a, Scott08b}. }  {Such models for the stellar population in the early universe imply} that the first luminous sources produce much more ionizing photons, and reionization starts earlier than for a population of conventional Pop.~III stars. In fact, we recently demonstrated that reionization based on massive Pop.~III can well reproduce the observed reionization optical depth \\citep{SchleicherBanerjee08a}. Increasing the number of ionizing photons per stellar baryon may thus reionize the universe too early and produce a too large reionization optical depth. This can only be avoided by introducing a transition to a stellar population which produces less ionizing photons, such that the universe can recombine after the first reionization phase. We therefore consider a double-reionization scenario in order to re-obtain the required optical depth. We discuss such models in \\S \\ref{reionization} and demonstrate that some models of dark stars require considerable fine-tuning in reionization models in order to be compatible with the reionization optical depth from the WMAP \\footnote{http://lambda.gsfc.nasa.gov/} 5-year data \\citep{Nolta08, Komatsu08} and to complete reionization at redshift $z\\sim6$ \\citep{Becker01}. In \\S \\ref{21cm}, we show how such scenarios can be tested via $21$ cm measurements.  A further consequence of the formation of dark stars is the steepening of the density profiles in minihalos \\citep{FreeseBodenheimer08, Iocco08}, thus increasing the dark matter clumping factor with respect to standard NFW models. In \\S \\ref{xray}, we estimate the increase in the clumping factor during the formation of dark stars and compare the calculation with our expectation for conventional NFW profiles {and heavy dark matter candidates. In \\S \\ref{light}, we perform similar calculations for the light dark matter scenario}. Further discussion and outlook is provided in \\S \\ref{outlook}. ",
        "conclusions": "\\label{outlook} In this work, we have examined whether the suggestion of dark star formation in the early universe is consistent with currently available observations. We use these observations to obtain constraints on dark star models and dark matter properties. From considering cosmic reionization, we obtain the following results: \\begin{itemize} \\item Dark stars with masses of the order $800\\ M_\\odot$ as suggested by \\citet{FreeseBodenheimer08} can only be reconciled with observations if somewhat artificial double-reionization scenarios are constructed. They consist of a phase of dark star formation followed by a phase of weak Pop.~II star formation and a final star burst to reionize the universe until redshift $6$. \\item The same is true for dark stars in which the number of UV photons is significantly increased due to dark matter capture, as suggested by \\citet{IoccoBressan08}. \\item It appears more reasonable to require that dark stars, if they were common, should have similar properties as conventional Pop.~III stars. For MS-dominated models, this requires that typical dark star masses are of order $100\\ M_\\odot$ or below. For CD models it requires a dark matter density {above} $10^{11}-10^{12}\\ \\mathrm{GeV}\\ \\mathrm{cm}^{-3}$ if a spin-dependent elastic scattering cross section of $t\\times10^{-39}$~cm$^2$ is assumed \\citep{Yoon08}. \\item {Alternatively, it may imply that the elastic scattering cross section is smaller than the current upper limits, that the dark matter cusp is destroyed by mergers or friction with the gas or that the star is displaced from the center of the cusp.} \\item A further interpretation is that dark stars are very rare. This would require some mechanism to prevent dark star formation in most minihalos. \\item However, if the double-reionization models are actually true, it would indicate that dark stars form only at redshifts beyond $14$, which makes direct observations difficult. \\item We also note that $21$ cm observations may either confirm or rule out double-reionization models. \\end{itemize} We have also examined whether the formation of dark stars and the corresponding enhancement of dark matter density in dark matter halos due to adiabatic contraction may increase the observed X-ray, gamma-ray and neutrino background. Here we found the following results: \\begin{itemize} \\item {For massive dark matter particles, direct annihilation into gamma-rays provides significant constraints for masses less than $30$~GeV.} \\item {For massive dark matter particles, the contribution from direct annihilation into neutrinos is well below the observed background.} \\item {In light dark matter scenarios, the $511$ keV emission is significantly enhanced below frequencies of $100$ keV in the observers restframe. For a certain range of parameters, this emission may even form a significant contribution of the total X-ray background. In this case, we derive a lower limit of $10$ MeV for the dark matter particle mass (while we find $7$ MeV for standard NFW profiles). } \\item {In light dark matter scenarios, the background radiation due to internal bremsstrahlung is not affected significantly from adiabatic contraction at early times, as the main contribution comes from low redshift. \\\\} \\item {Both for light and massive dark matter particles, the annihilation products in the remnants of dark stars may provide significant contributions {that may be used to constrain such models in more detail. However, whether this contribution can be reached is highly model-dependent and relevant questions regarding the death of dark stars has not been explored in the literature.}} \\end{itemize} Future observations may provide further constraints on this exciting suggestion. Small-scale $21$ cm observations may directly probe the HII regions of the first stars and provide a further test of the luminous sources at high redshift, and extremely bright stars might even be observed with the James-Webb telescope, if they form sufficiently late. With this work, we would like to initiate a discussion on observational tests and constraints on dark stars, which may tighten theoretical dark star models and provide a new link between astronomy and particle physics."
    },
    "0809/0809.3807_arXiv.txt": {
        "abstract": "We present the results of a deep ($15 \\la r \\la 23$), 20 night survey for transiting planets in the intermediate age open cluster M37 (NGC 2099) using the Megacam wide-field mosaic CCD camera on the 6.5m MMT. We do not detect any transiting planets among the $\\sim 1450$ observed cluster members. We do, however, identify a $\\sim 1~R_{J}$ candidate planet transiting a $\\sim 0.8~M_{\\odot}$ Galactic field star with a period of $0.77~{\\rm days}$. The source is faint ($V = 19.85~{\\rm mag}$) and has an expected velocity semi-amplitude of $K \\sim 220~{\\rm m/s}~(M/M_{J})$. We conduct Monte Carlo transit injection and recovery simulations to calculate the $95\\%$ confidence upper limit on the fraction of cluster members and field stars with planets as a function of planetary radius and orbital period. Assuming a uniform logarithmic distribution in orbital period, we find that $< 1.1\\%$, $< 2.7\\%$ and $< 8.3\\%$ of cluster members have $1.0~R_{J}$ planets within Extremely Hot Jupiter (EHJ, $0.4 < P < 1.0~{\\rm day}$), Very Hot Jupiter (VHJ, $1.0 < P < 3.0~{\\rm day}$) and Hot Jupiter (HJ, $3.0 < P < 5.0~{\\rm day}$) period ranges respectively. For $0.5~R_{J}$ planets the limits are $< 3.2\\%$, and $< 21\\%$ for EHJ and VHJ period ranges, while for $0.35~R_{J}$ planets we can only place an upper limit of $< 25\\%$ on the EHJ period range. For a sample of $7814$ Galactic field stars, consisting primarily of FGKM dwarfs, we place $95\\%$ upper limits of $< 0.3\\%$, $< 0.8\\%$ and $< 2.7\\%$ on the fraction of stars with $1.0~R_{J}$ EHJ, VHJ and HJ assuming the candidate planet is not genuine. If the candidate is genuine, the frequency of $\\sim 1.0 R_{J}$ planets in the EHJ period range is $0.002\\% < f_{EHJ} < 0.5\\%$ with $95\\%$ confidence. We place limits of $< 1.4\\%$, $< 8.8\\%$ and $< 47\\%$ for $0.5~R_{J}$ planets, and a limit of $< 16\\%$ on $0.3~R_{J}$ planets in the EHJ period range. This is the first transit survey to place limits on the fraction of stars with planets as small as Neptune. ",
        "introduction": "\\label{sec:intro}  The discovery by \\citet{Mayor.95} of a planet with half the mass of Jupiter orbiting the solar-like star 51 Pegasi with a period of only $4.23~{\\rm days}$ shocked the astronomical community. The existence of such a ``hot Jupiter'' (HJ) defied the prevailing theories of planet formation which had been tailored to explain the architecture of the Solar System. Since then, radial velocity (RV) surveys for planets orbiting nearby F, G and K main-sequence stars have determined that $1.2\\% \\pm 0.2\\%$ of these stars host a HJ \\citep[][where a HJ is defined as a planet roughly the size of Jupiter that orbits within $0.1~{\\rm AU}$ of its star]{Marcy.05}, with indications that this frequency depends on the metallicity of the host star \\citep[such that the frequency is roughly proportional to $10^{2{[}Fe/H{]}}$,][]{Fischer.05}.  Over the last decade there have been numerous surveys for extra-solar planets following a variety of techniques \\citep[e.g.][]{Butler.06} with the goal of determining the planet occurrence rate in new regions of parameter space. A successful technique has been to conduct photometric searches for planets that transit their host stars. This technique is particularly sensitive to planets on close-in orbits. To date more than 50 planets have been discovered by this technique\\footnote{http://exoplanet.eu}, including numerous very hot Jupiters (VHJ) with orbital periods between 1 and 3 days. \\citet{Gaudi.05} used the four transiting planets discovered at the time by the OGLE collaboration \\citep{Udalski.02a, Konacki.03, Bouchy.04, Konacki.04, Pont.04} to determine that only $0.1 - 0.2\\%$ of FGK stars host a VHJ. \\citet{Gould.06} conducted a thorough analysis of the OGLE survey to determine that the frequency of VHJs is $f_{VHJ} = (1/710)(1^{+1.10}_{-0.54})$ and $f_{HJ} = (1/320)(1^{+1.37}_{-0.59})$, while \\citet{Fressin.07} found $f_{VHJ} = (1/560)$ and $f_{HJ} = (1/320)$. The SWEEPS survey for transiting planets in the Galactic bulge conducted with the Hubble Space Telescope \\citep{Sahu.06} identified a putative class of ultra-short-period planets, or Extremely Hot Jupiters (EHJ) with periods less than $1.0~{\\rm day}$ orbiting stars lighter than $0.88 M_{\\odot}$. They conclude that $\\sim 0.4\\%$ of bulge stars more massive than $\\sim 0.44 M_{\\odot}$ are orbited by a Jupiter-sized planet with a period less than $4.2~{\\rm days}$, though they estimate that this fraction is uncertain by a factor of 2. Note that due to their faintness the majority of the SWEEPS candidates are unconfirmed with RV follow-up. In addition to these two surveys, the TrES \\citep[e.g][]{Alonso.04}, HAT \\citep[e.g.][]{Bakos.07}, XO \\citep[e.g.][]{McCullough.06}, and WASP \\citep[e.g.][]{CollierCameron.06} surveys have all discovered planets orbiting relatively bright stars in the Galactic field, though to date these surveys have not been used to calculate the planet occurrence frequency.  While photometric surveys of Galactic field stars have been quite successful at finding transiting planets over the last few years, it is generally difficult to measure the planet occurrence frequency with these surveys \\citep[for discussions of how this can be done despite the difficulties see][]{Gould.06,Fressin.07,Gaudi.07,Beatty.08}. The difficulty arises from the uncertainty in the parameters (mass, radius, metallicity) of the surveyed stars. Moreover, typical field surveys yield numerous false positives that are often culled in part by eye, these culling procedures are generally difficult to model in determining the detection efficiency of the survey. In contrast to field surveys, surveys of globular and open star clusters observe a population of stars with parameters that are relatively easy to determine en masse, moreover many of the false positive scenarios are less common for this type of survey. There has been significant work invested in developing optimum strategies to search for planets in stellar clusters \\citep{Janes.96,vonBraun.05,Pepper.05}. A number of groups have completed transit surveys of open clusters, including the UStAPS \\citep{Street.03,Bramich.05,Hood.05}, EXPLORE-OC \\citep{vonBraun.05}, PISCES \\citep{Mochejska.05, Mochejska.06}, STEPSS \\citep[][hereafter B06]{Burke.06} and MONITOR \\citep{Aigrain.07} projects and a survey by \\citet{Montalto.07}. There have also been several surveys of globular clusters \\citep{Gilliland.00, Weldrake.05, Weldrake.08}.  While to date no confirmed transiting planet has been found in a stellar cluster, many of these surveys have placed limits on the frequency of hot transiting planets, typically as functions of planetary radius as well as period. The globular cluster surveys have placed the most stringent constraints; the null result for the core of 47~Tucanae by \\citet{Gilliland.00} suggests that the frequency of HJ in this environment is at least an order of magnitude less than for the solar neighborhood, while the null result for the outskirts of the same cluster by \\citet{Weldrake.05} is inconsistent with the planet frequency in the solar neighborhood at the $3.3\\sigma$ level and suggests that the dearth of planets in this globular cluster is due to low metallicity rather than crowding effects. The open cluster surveys, on the other hand, have typically placed limits on the occurrence frequency that are well above the $1.2\\%$ measured by the RV surveys. Notably B06 conducted a thorough Monte Carlo simulation of their transit survey of the open cluster NGC 1245 to limit the frequency of EHJ, VHJ and HJ with radii of $1.5~R_{J}$ to $< 1.5\\%$, $< 6.4\\%$ and $< 52\\%$ respectively. The fundamental limit on the ability of open cluster surveys to place meaningful limits on the occurrence frequency of Jupiter-sized planets appears to be due to the relatively small number of stars in an open cluster. B06 find that for their survey strategy, $\\sim 7400$ dwarf stars would have to be observed for at least a month to put a limit of less than $2\\%$ on the planet frequency, which is significantly larger than the typical size of an open cluster. One open cluster that has been a popular target is NGC 6791. This cluster is old \\citep[$t \\sim 8~{\\rm Gyr}$,][]{Carraro.06, Kalirai.07}, metal rich \\citep[${[}M/H{]} \\sim +0.4$,][]{Gratton.06, Origlia.06} and contains a large number of stars \\citep[$M > 4000 M_{\\odot}$,][]{Kaluzny.92}, though it is also very distant \\citep[$(m-M)_{0} \\sim 12.8$,][]{Stetson.03} so that lower main sequence stars in the cluster are quite faint. It has been the target of three transit searches \\citep{Bruntt.03,Mochejska.05,Montalto.07}, the most recent of which found that their null result is inconsistent at the $\\sim 95\\%$ level with the RV HJ frequency at high metallicity.  The paucity of stars in open clusters appears to limit their usefulness as probes of the HJ frequency \\citep[excluding, perhaps, the result from][]{Montalto.07}. They may, however, be useful for probing smaller planet radii \\citep[see][]{Pepper.06a}. In the last several years RV surveys have discovered a number of Neptune and super-Earth-mass planets \\citep[HN, $M < 0.1 M_{J}$;][]{Butler.04,Endl.07,Fischer.07,Lovis.06,Melo.07,Rivera.05,Santos.04a,Udry.05,Udry.07,Vogt.05}. One of these planets, GJ 436 b, has been discovered to transit its host star \\citep{Gillon.07}. Little, however, is known about the frequency of these planets. Determining, or placing meaningful limits on this frequency would provide a powerful test of planet formation models. The theoretical predictions of the frequency of these objects run the gamut from a steep decline in the frequency of HN relative to HJ \\citep{Ida.04}, except perhaps for M-dwarfs \\citep{Ida.05}, to HN being ubiquitous \\citep{Brunini.05}.  In this paper, the fourth and final in a series, we present the results of a survey for transiting hot planets as small as Neptune in the intermediate age open cluster M37 (NGC 2099) using the MMT. We were motivated to conduct this transit survey by \\citet{Pepper.05,Pepper.06a} who suggested that it may be possible to find Neptune-sized planets transiting solar-like stars by surveying an open cluster with a large telescope. The Megacam mosaic imager on the MMT \\citep{McLeod.00} is an ideal facility for conducting such a survey due to its wide field of view and deep pixel wells that oversample the stellar point spread function (PSF). Preliminary observations of NGC 6791 suggested that finding Neptune-sized planets was indeed technically feasible using this facility \\citep{Hartman.05}. Using the formalism developed by \\citet{Pepper.05} we found that M37 is the optimum target for MMT/Megacam to maximize the number of stars to which we would be sensitive to Neptune-sized planets. We note that a drawback of this type of survey is that any identified candidates will be quite faint making follow-up RV confirmation difficult. For planets significantly smaller than $1.0 R_{J}$, false positives where the transiting object is a small star or brown dwarf are no longer applicable. Given the depth of the survey very few giants will be included in the sample, and those that are, can easily be rejected based on their colors. Low-precision spectroscopy may be sufficient to rule out various blended binary scenarios. Therefore, it is reasonable to suppose that for a given small-radius candidate one could make a strong argument that the object is a real planet without obtaining a RV determination of its mass. Also note that similar difficulties will be faced by the \\emph{Corot} and \\emph{Kepler} space missions (albeit for even smaller planets), so our experiment may provide a useful analogy to these missions.  We conducted the survey over twenty nights between December, 2005 and January, 2006, accumulating more than 4000 quality images of the cluster. This is easily the largest telescope ever utilized for such a survey. In the first paper in the series \\citep[Paper I]{Hartman.08a} we describe the observations and data reduction, combine photometric and spectroscopic data to refine estimates for the cluster fundamental parameters ($t = 550 \\pm 30~{\\rm Myr}$ with overshooting, $[M/H] = +0.045 \\pm 0.044$, $(m-M)_{V} = 11.57 \\pm 0.13~{\\rm mag}$ and $E(B-V) = 0.227 \\pm 0.038~{\\rm mag}$), and determine the cluster mass function down to $0.3 M_{\\odot}$. In the second paper \\citep[Paper II]{Hartman.08b} we analyze the light curves of $\\sim 23000$ stars observed by this survey to discover 1430 variable stars. In the third paper \\citep[Paper III]{Hartman.08c} we use the light curves to measure the rotation periods of 575 probable cluster members. This is the largest sample of rotation periods for a cluster older than a few hundred Myr, and thus provides a unique window on the late time rotation evolution of low-mass main sequence stars.  In the following section we will summarize our observations and data reduction. In \\S~\\ref{sec:transselect} we discuss the pipeline used to remove systematic variations from light curves and identify candidate transiting planets. In \\S~\\ref{sec:transcand} we describe the candidate transiting planets identified by this survey, finding no candidates that are probable cluster members. In \\S~\\ref{sec:deteff} we conduct Monte Carlo simulations to determine the transit detection efficiency of our survey. In \\S~\\ref{sec:results} we present our results on the limit of stars with planets for various planetary radii and orbital periods. Finally, we conclude in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:sumobs}  The observations and data reduction procedure were described in detail in Papers I and II, we provide a brief overview here. The photometric observations consist of $gri$ photometry for $\\sim 16,000$ stars, $gri$ photometry for stars in a field located two degrees from the primary M37 field and at the same Galactic latitude, and $r$ time-series photometry for $\\sim 23,000$ stars, all obtained with the Megacam instrument \\citep{McLeod.00} on the 6.5 m MMT. Megacam is a $24\\arcmin \\times 24\\arcmin$ mosaic imager consisting of 36 2k$\\times$4k, thinned, backside-illuminated CCDs that are each read out by two amplifiers. The mosaic has an unbinned pixel scale of $0\\farcs 08$ which allows for a well sampled point-spread-function (PSF) even under the best seeing conditions. To decrease the read-out time we used $2 \\times 2$ binning with the gain set so that the pixel sensitivity became non-linear before the analog-to-digital conversion threshold of 65,536 counts. Because of the fine sampling and the relatively deep pixel wells, one can collect $2\\times 10^7$ photons in $1\\arcsec$ seeing from a single star prior to saturation, setting the photon limit on the precision in a single exposure to $\\sim$ 0.25 mmag.  The primary time-series photometric observations consist of $\\sim 4000$ high quality images obtained over twenty four nights (including eight half nights) between December 21, 2005 and January 21, 2006. We obtained light curves for stars with $14.5 \\la r \\la 23$ using a reduction pipeline based on the image subtraction technique and software due to \\citet{Alard.98} and \\citet{Alard.00}. The resulting light curves were passed through the processing and transit detection pipeline that we describe in the following section. We used the {\\scshape Daophot} 2 and {\\scshape Allstar} PSF fitting programs and the {\\scshape Daogrow} program \\citep{Stetson.87,Stetson.90,Stetson.92} to obtain the $g$, $r$, and $i$ single-epoch photometry.  As described in Paper I we also take $BV$ photometry for stars in the field of this cluster from \\citet{Kalirai.01}, $K_{S}$ photometry from 2MASS \\citep{Skrutskie.06} and we transform our $ri$ photometry to $I_{C}$ using a transformation based on the $I_{C}$ photometry from \\citet{Nilakshi.02}.  In addition to the photometry, we also obtained high-resolution spectroscopy of 127 stars using the Hectochelle multi-fiber, high-dispersion spectrograph \\citep{Szentgyorgyi.98} on the MMT. The spectra were obtained on four separate nights between February 23, 2007 and March 12, 2007 and were used to measure $T_{eff}$, $[Fe/H]$, $v\\sin i$ and the radial velocity (RV) via cross-correlation against a grid of model stellar spectra computed using ATLAS 9 and SYNTHE \\citep{Kurucz.93}. The classification procedure was developed by Meibom et al. (2008, in preparation), and made use of the \\emph{xcsao} routine in the {\\scshape Iraf}\\footnote{{\\scshape Iraf} is distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy, Inc., under agreement with the National Science Foundation.} \\emph{rvsao} package \\citep{Kurtz.98} to perform cross-correlation. We use these spectra to provide stellar parameters and radial velocity measurements for several of the host stars to the candidate transiting planets.  \\label{sec:discussion}  We have presented the results of a deep $\\sim 20$ night survey for transiting hot planets in the open cluster M37. This survey stands out from previous ground-based transit surveys both in terms of the size of the telescope used, and the photometric precision attained. We observed $\\sim 1450$ cluster members with masses between $0.3~M_{\\odot} \\la M \\la 1.3~M_{\\odot}$ as well as 7814 Galactic field stars with masses between $0.1~M_{\\odot} \\la M \\la 2.1~M_{\\odot}$. While no candidate planets were found among the cluster members, we did identify one candidate extremely hot Jupiter with a period of $0.77~{\\rm days}$ transiting a Galactic field star. However, the follow-up spectroscopic observations needed to confirm the planetary nature of this candidate would be difficult (although perhaps not impossible) to obtain with current technology, given the faintness of the source. We note that if this candidate is real, then we conclude that $0.09^{+0.2}_{-0.077}\\%$ of FGKM stars have Jupiter-sized planets with periods between 0.5 and 1.0 days. This result would be consistent with the results from the SWEEPS survey, and would confirm this new class of ultra short period planets. We also note that this planet frequency is small enough that these planets could have escaped detection in most other radial velocity and transit surveys.  The primary result of this survey is an upper limit on the frequency of planets smaller than $1.0R_{J}$. For cluster members, we find that at $95\\%$ confidence $< 25\\%$ of stars have planets with radii as small as $0.35 R_{J}$ and periods shorter than $1.0~{\\rm day}$, $< 44\\%$ of stars have planets with radii as small as $0.45 R_{J}$ and periods between $1.0$ and $3.0~{\\rm days}$, and $< 49\\%$ of stars have planets with radii as small as $0.6 R_{J}$ and periods between $3.0$ and $5.0~{\\rm days}$. The upper limits on the smallest planets may be as much as a factor of $50\\%$ higher if all the variable stars near the cluster main sequence are cluster members. For the field stars we are able to place $95\\%$ confidence upper limits of $16\\%$ on the fraction of stars with planets as small as $0.3 R_{J}$ with periods less than $1.0~{\\rm day}$, $8.8\\%$ on the fraction of stars with planets as small as $0.5 R_{J}$ with periods between $1.0$ and $3.0~{\\rm days}$ and $47\\%$ on the fraction of stars with planets as small as $0.5 R_{J}$ with periods between $3.0$ and $5.0~{\\rm days}$. We estimate that these upper limits may be higher by at most a factor of $\\sim 11\\%$ due to binarity. While these limits do not approach the observed frequency of Jupiter-sized planets with similar periods, they do represent the first limits on the frequency of planets as small as Neptune. We can now state empirically that extremely hot Neptunes (periods shorter than $1~{\\rm day}$) are not ubiquitous, nor are very hot planets with radii intermediate between Neptune and Saturn.  The limits that we place on Jupiter-sized planets are more stringent than previous open cluster transit surveys, but are still above the frequencies measured by RV and Galactic field transit surveys. The primary limitation on open cluster transit surveys appears to be the paucity of stars in these systems. To place a limit on the frequency of HJ that is less than $2\\%$ with the same set of observations, M37 would have to have been $\\sim 4$ times richer than it is. We also note that for a relatively young cluster like M37, variability may reduce the detectability of Neptune-sized planets by as much as $\\sim 50\\%$."
    },
    "0809/0809.0919_arXiv.txt": {
        "abstract": "We have discovered two dusty intervening Mg~{\\sc ii} absorption systems at $z\\sim1.3$ in the Sloan Digital Sky Survey (SDSS) database. The overall spectra of both QSOs are red (u-K$>$4.5 mag) and are well modelled by the composite QSO spectrum reddened by the extinction curve from the Large Magellanic Cloud(LMC2) Supershell redshifted to the rest-frame of the Mg~{\\sc ii} systems. In particular, we detect clearly the presence of the UV extinction bump at $\\lambda_{\\rm rest}\\sim 2175$~\\AA. Absorption lines of weak transitions like Si~{\\sc ii}$\\lambda$1808, Cr~{\\sc ii}$\\lambda$2056, Cr~{\\sc ii}+Zn~{\\sc ii}$\\lambda$2062, Mn~{\\sc ii}$\\lambda$2594, Ca~{\\sc ii}$\\lambda$3934 and Ti~{\\sc ii}$\\lambda$1910 from these systems are detected even in the low signal-to-noise ratio and low resolution SDSS spectra, suggesting high column densities of these species. The depletion pattern inferred from these absorption lines is consistent with that seen in the cold neutral medium of the LMC. Using the LMC $A_V$ vs. $N$(H~{\\sc i}) relationship we derive $N$(H~{\\sc i})$\\sim 6\\times 10^{21}$ cm$^{-2}$ in both systems. Metallicities are close to solar. Giant Metrewave Radio Telescope (GMRT) observations of these two relatively weak radio loud QSOs (f$_\\nu$ $\\sim$ 50 mJy) resulted in the detection of 21-cm absorption in both cases. We show that the spin temperature of the gas is of the order of or smaller than $500$~K. These systems provide a unique opportunity to search for molecules and diffuse interstellar bands at $z>1$. ",
        "introduction": "Studying the physical conditions in the interstellar medium (ISM) at high redshift and the processes that maintain these conditions is important for our understanding of how galaxies form and evolve. The presence of dust influences the physical state of the gas through photo-electric heating,  UV shielding, and formation of molecules on the surface of grains. However, we know very little about dust properties in the ISM at high redshifts. Properties of the dust can be derived from extinction curves observed in different astrophysical objects. Recently, Noll et al. (2007) found evidence for the presence of a moderate UV bump in at least part of the population of massive galaxies at $z>1$. Studies of dust extinction in the circum-burst environment of Gamma Ray Bursts indicate that in most cases, a LMC-like extinction curve is preferred (e.g. Heng et al. 2008).  The depletion of Cr with respect to Zn in intervening damped Lyman-$\\alpha$ systems (DLAs)  shows that dust is indeed an important component of the high density gas (Pettini, Smith \\& Hunstead 1994). A correlation is observed in DLAs between metallicity and dust-depletion (Ledoux, Petitjean \\& Srianand 2003) which is confirmed by the higher detection rate of H$_2$ in DLAs with higher metallicities (Petitjean et al. 2006; Noterdaeme et al. 2008). The corresponding gas is cold (T$\\sim$150 K; Srianand et al. 2005) as expected from multiphase ISM models in which the gas with high metallicity and dust content has lower kinetic temperature than the gas with lower metallicity and dust content (Wolfire et al. 1995). However, even in the highest metallicity DLAs typical dust signatures like high extinction (i.e 0.16$\\le$ E(B$-$V)$\\le$0.40), 2175~\\AA~absorption bump or the diffuse interstellar bands (DIBs) are not seen.    Signatures of dust are seen from few intermediate and low redshift absorption systems. Both DIBs and 2175\\AA~absorption bump have been detected in the $z_{\\rm abs}$~=~0.524 system toward AO~0235+164 (Wolfe \\& Wills 1977; York et al. 2006; Junkkarinen et al. 2004; Kulkarni et al. 2007). This system also shows strong 21-cm absorption (Wolfe \\& Wills 1977) and has $N$(H~{\\sc i}) = 5$\\times10^{21}$ cm$^{-2}$, E(B-V) = 0.23 (Junkkarinen et al. 2004) and $T_{\\rm s}$ = 220$\\pm$60 K. Wang et al. (2004) have reported the detection of the 2175\\AA~absorption bump in 3 intermediate redshift ($z\\sim1.3$) Mg~{\\sc ii} systems. Recently, Ellison et al. (2008) have detected DIBs in the \\zabs = 0.1556 Ca~{\\sc ii} systems towards SDSS J001342.44-002412.6. Composite spectra have been obtained for different samples of absorbers from the Sloan Digital Sky Survey (SDSS). They show in a statistical way that dust is present in strong Mg~{\\sc ii} (York et al. 2006a)  and Ca~{\\sc ii} systems (Wild, Hewett \\& Pettini 2006). In the former case the mean extinction curve is similar to the SMC curve with a rising ultraviolet extinction below 2175~\\AA~ with E(B$-$V)$\\le$0.08 and with no evidence of an UV bump. In the latter case, evidence for the UV bump is marginal and LMC extinction curve provides E(B-V)$\\le$0.103 for different sub-samples.  An efficient way to reveal cold and dusty gas is to search for 21-cm absorption as shown by the high detection rate of 21-cm absorption towards red QSOs (Carilli et al. 1998; Ishwara-Chandra, Dwarakanath \\& Anantharamaiah 2003; Curran et al. 2006). Multiple lines of sight toward lensed QSOs passing through regions producing high extinction also show 21-cm and molecular absorption lines (see Wiklind \\& Combes 1996).  We have recently completed a systematic search for 21-cm absorption in a complete sample of 38 Mg~{\\sc ii} systems, drawn from the SDSS DR5, with redshifts in the range $1.10\\le z\\le 1.45$ corresponding to the frequency coverage of 610~MHz feed at GMRT (see Gupta et al. 2007 for early results). Using an automatic procedure we have identified 5 systems in this sample showing strong absorption lines from Si~{\\sc ii}, Zn~{\\sc ii}, Cr~{\\sc ii}, Fe~{\\sc ii}, Mg~{\\sc ii} and Mg~{\\sc i} in front of radio-loud QSOs with flux density $\\ge50$ mJy. Here, we report the detection of 21-cm absorption from two of these systems at \\zabs$\\sim$1.3 towards red (u-K$\\ge$4.5 mag) QSOs. The optical spectra of these two quasars possess 2175\\AA~dust absorption features. ",
        "conclusions": "\\begin{figure} \\psfig{figure=J0850dusta.ps,height=7.0cm,width=8.cm,angle=0}  \\psfig{figure=J0852dusta.ps,height=7.cm,width=8.cm,angle=0} \\caption[]{The SDSS spectrum of J0850+5159 (top) and J0852+3435(bottom) are fitted with SDSS composite spectrum  corrected using average extinction curves from the Milky Way (dotted), the LMC2 Supershell (solid) and the SMC (dashed). Rest wavelengths at \\zabs are indicated at the top of the figure. The arrow marks the location of 2175 {\\AA} feature at \\zabs. For presentation purpose the observed spectrum is boxcor smoothed by 10 pixels. } \\label{fig1} \\end{figure}  \\subsection {The $z_{\\rm abs} = 1.3265$ absorption system towards J085042.21$+$515911.7}  The spectrum of this QSO (with  u-K $\\sim$ 4.8 mag) shows a curvature around 2175~\\AA~ in the rest frame of the Mg~{\\sc ii} system (see  top panel of Fig.~\\ref{fig1}). The Milky Way extinction curve that fits the 2175 {\\AA} feature over predicts the QSO flux below 4500~\\AA. This is also the case (not shown in Fig.~\\ref{fig1}) when we use the average LMC extinction curve. The SMC extinction curve that fits the observed spectrum around 5500~{\\AA} under predicts flux below 4500~{\\AA}. The LMC2 extinction curve that has a shallow UV-bump and relatively high extinction at rest wavelength $\\lambda_r < 2100$ {\\AA} compared to that of the Milky Way (and of the LMC) reproduces the data better. We also fitted the spectra of $\\sim$ 200 non-BAL QSOs with \\zem =1.892$\\pm$0.005 using SMC extinction curve  and find the probability of the 2175~{\\AA} feature being produced by QSO-to-QSO spectral variation to be $\\le$0.03. For the range of composite spectra considered here, the mean $A_{V_a}$ is 0.74$\\pm$0.07 (and E(B$-$V) = 0.27). Our best fit model prediction J = 16.7, H = 15.8 and K = 14.7 mag, with a typical uncertainty of 0.2 mag, agrees well with the observed  values, J = 16.99$\\pm$0.23, H = 15.72$\\pm$0.16 and K = 15.27$\\pm$0.15, from Strutskie et al. (2006). The excess flux predicted in K band can be mainly attributed to the known host galaxy contribution ($\\sim$0.3 mag) in the SDSS composite spectrum (See Vanden Berk et al. 2001) at $\\lambda_r>$ 6000\\AA.  From Table~2 of Gordon et al. (2003) we have, \\begin{equation} N({\\rm HI})\\sim {A_{V_a}\\over \\kappa} (6.97\\pm0.67)\\times 10^{21} {\\rm cm}^{-2}. \\end{equation} Here $\\kappa$ is the ratio of dust-to-gas ratio in the absorption system to that in  LMC. We can use the equivalent widths of weak transitions (Table~\\ref{sdssuv}) to estimate the column densities of species and derive the depletion factors onto dust-grains assuming Zn is not depleted: $-$0.5,$-$0.9, $-$1.0, $-$0.9, and $-$1.0 for [Si/Zn], [Cr/Zn], [Fe/Zn], [Ti/Zn],  and [Mn/Zn] respectively. Note that Zn~{\\sc ii}$\\lambda2062$ is blended with a Cr~{\\sc ii} absorption but we can remove the contribution from Cr~{\\sc ii} by using the average $N$(Cr~{\\sc ii}) derived from unblended lines. Depletion is higher than what is typically seen in high-$z$ absorption systems and are intermediate between what is seen in the Cold and Warm phases in the Galactic ISM (Welty et al. 1999). Clearly more than 90\\% of the refractory metals are in dust as in the LMC. If we assume $\\kappa\\sim 0.9$ then we get $N$(H~{\\sc i}) = (5.73$\\pm$1.10)$\\times10^{21}$ cm$^{-2}$.  Our GMRT spectrum shows very strong absorption that can be modelled with two Gaussian components (see Table~\\ref{gmrtres} and upper panel in Fig.~\\ref{gmrtdet}). The radio spectrum of this source is flat and VLBA maps at 5 and 15 GHz show that more than 90\\% of the flux density is in the unresolved core (Taylor et al. 2005). Using $f_{\\rm c} =0.9$ in Eq.~1, we derive $N$(H~{\\sc i}) = 3.02$\\pm0.84\\times10^{19}$ $T_{\\rm s}$ cm$^{-2}$. From the total $N$(H~{\\sc i}) inferred from extinction we can estimate the spin temperature, $T_{\\rm s}\\sim190^{+124}_{-69}$ K. Note that this value is an upper limit as the dust-to-gas ratio may be larger than 1. All this is consistent with the system being associated with a cold neutral medium with dust properties similar to that seen in LMC2 supershell sample of Gordon et al. (2003). \\par\\noindent  \\subsection{The $z_{\\rm abs} = 1.3095$ absorption system towards J085244.74$+$343540.5}  This QSO has u-K$\\sim$5.6 mag and shows a curvature around 2175~\\AA~in the rest frame of the Mg~{\\sc ii} system (see bottom panel of Fig.~\\ref{fig1}) that is stronger than for the previous quasar. The Mg~{\\sc ii} system at \\zabs = 1.3095 shows absorption lines that are also stronger than that of the previous system. The SDSS spectrum of this quasar is also well reproduced when applying a LMC2 like extinction curve to the QSO composite spectrum (with $A_{V_a}$ = 1.00$\\pm$0.09 and E(B$-$V) = 0.36). { Using spectra of $\\sim$ 330 non-BAL QSOs with \\zem=1.655$\\pm$0.005, we find the probability of the 2175~{\\AA} feature being produced by QSO-to-QSO spectral variation to be $\\le$0.001. } {Our best fit model prediction J = 16.5, H = 15.2 and K = 14.2 mag agrees reasonably with the observed  values, J = 16.75$\\pm$0.14, H = 15.57$\\pm$0.11 and K = 15.10$\\pm$0.12, considering the possible uncertainties in the QSO composite spectrum discussed above. } We derive lower limits on the column densities using the weaker metal lines. The absence of Cr~{\\sc ii}$\\lambda$2056 and Fe~{\\sc ii}$\\lambda$2249 features together with the high equivalent width of the Zn~{\\sc ii}+Cr~{\\sc ii} blend at $\\lambda_r = 2062$~\\AA~ is consistent with higher depletion factor and higher metallicity in this system. The detection of Ca~{\\sc ii} absorption strongly supports this conclusion. From the SED fit and assuming $\\kappa = 1$ we derive ${\\sl N}$(H~{\\sc i}) =  6.97$\\pm1.30\\times10^{21}$ cm$^{-2}$. This may well be an upper limit as $\\kappa$ could be larger than 1.  This system, despite having $N$(H~{\\sc i})$\\kappa$ similar to that in the \\zabs = 1.3265 towards J0850+5159, has a factor of two lower total integrated $N$(H~{\\sc i}) from 21-cm absorption. This source has a flat radio spectrum and is unresolved in the VLA A-array 8.4 GHz image\\footnote{from CLASS (Cosmic Lens All-Sky Survey) and JVAS (Jodrell/VLA Astrometric Survey) archive website\\\\ http://www.jb.man.ac.uk/research/gravlens/class/gmodlist.html} having a resolution of 0''.255$\\times$0''.236. However, high resolution (few mas scale) VLBA map is not available for this source. Therefore, while $f_{\\rm c}$ could be close to 1 its exact value is still uncertain. We get a constraint $N$(H~{\\sc i})$f_{\\rm c}$/$T_{\\rm s}$ = $1.31\\pm0.26\\times 10^{19}$ cm$^{-2}$ K$^{-1}$. By comparing the two $N$(H~{\\sc i}) estimates we derive $T_{\\rm s}/f_{\\rm c} = 536^{+234}_{-88}$ K. As $f_{\\rm c}$ can be less than unity and $\\kappa$ is probably larger than one, the above value is an upper limit for the gas kinetic temperature. This is consistent with the gas being part of a cold neutral medium."
    },
    "0809/0809.0896_arXiv.txt": {
        "abstract": "We present a sample of 42 high-mass low-metallicity outliers from the mass--metallicity relation of star-forming galaxies.  These galaxies have stellar masses that span $\\log(M_{\\star}/\\mbox{M}_{\\odot}) \\sim 9.4$ to 11.1 and are offset from the mass--metallicity relation by $-0.3$ to $-0.85$\\,dex in \\tlogoh.  In general, they are extremely blue, have high star formation rates for their masses, and are morphologically disturbed.  Tidal interactions are expected to induce large-scale gas inflow to the galaxies' central regions, and we find that these galaxies' gas-phase oxygen abundances are consistent with large quantities of low-metallicity gas from large galactocentric radii diluting the central metal-rich gas.  We conclude with implications for deducing gas-phase metallicities of individual galaxies based solely on their luminosities, specifically in the case of long gamma-ray burst host galaxies. ",
        "introduction": "\\label{sec:intro}  Star-forming galaxies fall on a luminosity--metallicity relation such that more luminous galaxies tend to have higher gas-phase metallicities (the proxy for which is typically the oxygen abundance in units of {$12+\\log[\\mbox{O}/\\mbox{H}]$).  This relation is observed to hold---shifted to lower metallicities---at redshifts as high as $z\\sim 3.5$ \\citep{erb06a,maiolino08}, and its scatter decreases to $\\sim 0.15$\\,dex at $z\\sim 0$ when galaxy stellar mass replaces luminosity in the relation \\citep{tremonti04}.  Most commonly-accepted theories as to the origin of the mass--metallicity relation have as a central proposition that low-mass galaxies are metal-deficient, rather than that high-mass galaxies are metal-enhanced \\citep{larson74, dalcanton07,finlator08}.  In fact, there is some evidence that the mass--metallicity relation may flatten at large stellar masses, $\\log(M_{\\star}/\\mbox{M}_{\\odot}) \\gtrsim 10.5$ \\citep[e.g.,][]{tremonti04}, though it is unclear to what extent this flattening is due to a saturation of the metallicity indicator used at high oxygen abundances \\citep[see e.g.,][]{kewley02,bresolin06,kewley08}.  Because most star-forming galaxies {\\em do} lie on a mass--metallicity locus, we can also learn about the gas-phase metallicity evolution of galaxies by studying the properties of galaxies that do {\\em not} fall on the relation.  In \\citet[][hereafter Paper~I]{peeples08}, we explored the population of low-mass high-metallicity outliers, and postulated that these metal-rich dwarf galaxies have low gas fractions and are therefore nearing the end of substantial epochs of star formation.  Here we investigate the other corner of the mass--metallicity plane by asking what we can learn about the evolution of massive galaxies from the properties of the high-mass low-metallicity outliers.  As shown in Figure~\\ref{fig:ohmstar}, we find 42 low-metallicity high-mass galaxies with masses ranging from $\\log(M_{\\star}/\\mbox{M}_{\\odot}) \\sim 9.4$ to 11.1 and offsets from the central mass--metallicity relation of $-0.3$ to $-0.85$\\,dex.  We describe in \\S\\,\\ref{sec:sample} how we selected this sample and verified the galaxies' outlier status.  In \\S\\,\\ref{sec:disc}, we describe the physical properties of these galaxies and discuss possible origins for their low oxygen abundances.  Specifically, we find that they have highly disturbed morphologies strongly suggestive of merging or post-merging systems, are extremely blue, and have high specific star formation rates.  We summarize our conclusions and state some implications of these findings in \\S\\,\\ref{sec:conc}.  \\begin{figure*} \\plotone{OHmstar_grey_95.ps} \\caption{\\label{fig:ohmstar}Plot of \\tlogoh\\ v.\\ $\\log(\\mbox{M}_{\\star}/\\mbox{M}_{\\odot})$. The small grey points represent galaxies from the main \\citet{tremonti04} SDSS sample described in \\S\\,\\ref{sec:sample}, and the step-curve is the median of these points in bins of $0.1$\\,dex of total stellar mass.  The high-mass low-metallicity massive galaxies in our sample are plotted as the large grey circles in the lower-right region of the diagram; the only two clearly spiral galaxies in the sample are denoted with ``\\S'' symbols.  For reference we also plot the high-metallicity outliers from Paper~I as large open points ({\\em upper left}).  } \\end{figure*} ",
        "conclusions": "\\label{sec:conc} We have identified a sample of 42 low-metallicity high-mass outliers from the mass--metallicity relation.  As a population, these galaxies have disturbed morphologies and high specific star formation rates, implying that they are undergoing a merger-induced starburst.  We propose that their observed low oxygen abundances are due the tidal interaction inducing large-scale gas inflow which subsequently dilutes the central interstellar medium in these galaxies.  While there have been several observational and theoretical suggestions that interacting galaxies will result in a decreased observed gas-phase metallicity, this is the first work that has shown that the vast majority of severe low-metallicity outliers from the mass--metallicity relation are morphologically disturbed.  Finally, we note that the cuts outlined in \\S\\,\\ref{sec:sel} provide an effective means of using {\\em only} colors and spectra to identify a rather pure (though not necessarily complete) sample of tidally disturbed galaxies.  One striking implication of these results is that, while it is safe to assume that the metallicities for {\\em populations} of galaxies will fall within a particular range of values given their luminosities at redshifts for which the luminosity--metallicity relation has been measured, one should not assign metallicities to {\\em individual} galaxies based solely on their luminosities.  For example, while studies of the host galaxies of long gamma-ray bursts (GRBs) suggest that GRBs are only found in low-metallicity environments \\citep{stanek06,kewley07}, several recent GRBs have been associated with very luminous hosts, such as GRB~070306 and its $M_B \\sim -22.3$ host galaxy \\citep{jaunsen08}.  As we have shown in this paper and Paper~I, assuming an oxygen abundance for a galaxy from its luminosity alone can result in mis-estimating \\tlogoh\\ in luminous galaxies by a whole dex, and thus one should not assume that, e.g., the host of GRB~070306 necessarily has a high metallicity.  For cases in which \\logoh\\ has been measured via emission line spectra for GRB hosts, it is generally low, even if the galaxy is luminous.  For example, the host galaxy of GRB~031203 has an absolute $B$-band magnitude similar to that of the Milky Way, yet \\citet{margutti07} find it to have a low metallicity ($12+\\log[\\mbox{O}/\\mbox{H}] = 8.12$) as well as a high star formation rate ($\\sim 13\\, \\mbox{M}_{\\odot}$\\,yr$^{-1}$).  \\citet{prochaska04} interpret this offset in metallicity as a sign that GRB~031203 went off in a ``very young star-forming region,'' which is in line with our interpretation in \\S\\,\\ref{sec:merge}.  Futhermore, the brighter GRB hosts studied by \\citet{fruchter06} are morphologically very similar to our massive low-metallicity galaxies, and our results strongly suggest that the oxygen abundances inferred from luminosity alone are uncertain by as much as 1\\,dex. Specifically, if a luminous galaxy is morphologically disturbed, has a high specific star formation rate, and is extremely blue, then it should not be assumed to lie within the luminosity--metallicity locus of star-forming galaxies."
    },
    "0809/0809.0291_arXiv.txt": {
        "abstract": "We have studied the sensitivity of \\textsl{s}-process nucleosynthesis in massive stars to $\\pm {2\\sigma}$ variations in the rates of the triple-$\\alpha$ and $^{12}$C($\\alpha,\\gamma$)$^{16}$O reactions.  We simulated the evolution of massive stars from H-burning through Fe-core collapse, followed by a supernova explosion. We found that: the production factors of \\textsl{s}-process nuclides between $^{58}$Fe and $^{96}$Zr change strongly with changes in the He burning reaction rates; using the \\citet{lod03} solar abundances rather than those of \\citet{and89} reduces \\textsl{s}-process nucleosynthesis; later burning phases beyond core He burning and shell C burning have a significant effect on post-explosive production factors.  We also discuss the implications of the uncertainties in the helium burning rates for evidence of a new primary neutron capture process (LEPP) in massive stars. ",
        "introduction": "Half of the elements between Fe and Bi are produced by the slow (\\textsl{s}) neutron capture process and most of the remainder by the rapid (\\textsl{r}) neutron capture process.  About 35 additional neutron deficient stable nuclides above $^{56}$Fe, the \\textsl{p}-process nuclei, are produced in explosive processes (\\citealt{pra90a}).  Two components, the main and the weak \\textsl{s}-processes, are required to explain the  isotopic distributions of \\textsl{s}-process nuclides. The main \\textsl{s}-process occurs in low-mass ($\\lesssim4\\,\\Msun$) asymptotic giant branch (AGB) stars and contributes mainly to the production of heavier elements, with  a smaller contribution to $A \\leq 90$. The weak \\textsl{s}-process occurs during the late evolutionary stages of massive stars ($\\gtrsim 10\\,\\Msun$) and produces nuclides up to $A \\simeq 90$.  Recently, \\cite{tra04} summed the contributions of the \\textsl{r}-process, the main and weak \\textsl{s}-process, and the \\textsl{p}-process to the abundances of Sr, Y, and Zr.  The summed contributions were smaller than the observed solar abundances by 8\\%, 18\\%, and 18\\%, respectively. They concluded that an additional light element primary \\textsl{s}-process contribution from massive stars (LEPP) is needed to explain this difference; the nature and site of the LEPP are unknown.  This LEPP has also been invoked by \\cite{mon07} to explain the abundances of a larger group of light \\textsl{r}-process elements.  Since the LEPP effects are relatively small, and since some of the LEPP elements are produced with relatively large abundance in the weak \\textsl{s}-process, it is important to establish whether the uncertainties in the weak \\textsl{s}-process are sufficiently small that the claim of LEPP contributions is robust.  Nuclide production in the weak \\textsl{s}-process also depends on the rate of the neutron source $^{22}$Ne($\\alpha,n$)$^{25}$Mg and on the capture cross sections for the neutron poisons (medium-weight isotopes up to Fe, including $^{12}$C, $^{16}$O, $^{20}$Ne, $^{22}$Ne, $^{23}$Na, $^{24}$Mg, and $^{25}$Mg). Neither the source strength nor the neutron capture cross sections for the poisons are known with sufficient accuracy (\\citealt{hei08}). We shall not deal with these issues in this paper, but rather with the more indirect effects of uncertainties in the rates ($R_{3\\alpha}$ and $R_{\\alpha,12}$) of the triple alpha and $^{12}$C($\\alpha,\\gamma$)$^{16}$O reactions, and in the initial stellar composition. For example, we have shown in a previous paper (\\citealt{tur07}; see also \\citealt{wea93}; \\citealt{woo03}; \\citealt{woo07}), that the amount of the above neutron poisons present during the weak \\textsl{s}-process in massive stars depends sensitively on these rates and the initial stellar composition.  Most earlier studies of the weak \\textsl{s}-process focused on production toward the end of core He burning by neutrons from the $^{22}$Ne($\\alpha,n$)$^{25}$Mg reaction (\\citealt{cou74}; \\citealt{lam77}; \\citealt{arn85}; \\citealt{bus85}; \\citealt{lan89}; \\citealt{pra90b}; \\citealt{the00}; \\citealt{rai91a}; \\citealt{bar92}). Later papers considered also a second exposure at higher temperatures and neutron densities peaking during shell C burning (\\citealt{rai91b}; \\citealt{rai92}; \\citealt{rai93}; \\citealt{the07}).  Explosive processing in the supernova explosion was not considered. Recent calculations of \\textsl{s}-process yields have been extended to consider the entire evolutionary history of the star, including  explosive burning (\\citealt{hof01}; \\citealt{rau02}; \\citealt{lim03}).  In this paper we simulate the evolution of massive stars from H-burning through Fe-core collapse, followed by a supernova explosion. Our principal purposes are (1) to establish the magnitude of weak \\textsl{s}-process production in a self-consistent model; (2) to study the uncertainties in weak \\textsl{s}-process nucleosynthesis arising from uncertainties in $R_{3\\alpha}$ and $R_{\\alpha,12}$; (3) to study the effects of different stellar abundances, specifically those of \\cite{lod03} and \\cite{and89}, hereafter L03 and AG89; (4) to delineate the stages of stellar evolution during which weak \\textsl{s}-process elements are produced; and (5) to assess the bearing of these uncertainties on the robustness of the LEPP process.  In Section~2, we describe the stellar model and the input physics relevant to the treatment of the weak \\textsl{s}-process. Section~3 gives our results for the dependence of the post-explosive weak-s process production factors on changes in the rates of the helium burning reactions and on different initial stellar abundances. In Section~4, we show the contribution to the weak \\textsl{s}-process abundances of the various stellar burning stages prior to the supernova explosion.  In Section~5, we investigate the range of weak \\textsl{s}-process production of Sr, Y, and Zr allowed by the uncertainties in the helium burning reactions.   \\begin{figure*} \\centering \\includegraphics[angle=90,width=0.475\\textwidth]{AA_15Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{LA_15Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{AA_20Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{LA_20Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{AA_25Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{LA_25Msun_pfVsRate_final.eps} \\caption{Post-explosive production factors as a function of $R_{\\alpha,12}$. \\textbf{(a)} The AA series, $15\\,\\Msun$. \\textbf{(b)} The LA series, $15\\,\\Msun$. \\textbf{(c)} The AA series, $20\\,\\Msun$.  \\textbf{(d)} The LA series, $20\\,\\Msun$. \\textbf{(e)} The AA series, $25\\,\\Msun$. \\textbf{(f)} The LA series, $25\\,\\Msun$.} \\label{pfVs12cag} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_AA15Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_LA15Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_AA20Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_LA20Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_AA25Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR12Cag_LA25Msun.eps} \\caption{Post-explosive production factors for s-only isotopes as a function of $R_{\\alpha,12}$. \\textbf{(a)} The AA series, $15\\,\\Msun$. \\textbf{(b)} The LA series, $15\\,\\Msun$. \\textbf{(c)} The AA series, $20\\,\\Msun$.  \\textbf{(d)} The LA series, $20\\,\\Msun$. \\textbf{(e)} The AA series, $25\\,\\Msun$. \\textbf{(f)} The LA series, $25\\,\\Msun$.} \\label{sOnlyPfVs12cag} \\end{figure*}    \\begin{figure*} \\centering \\includegraphics[angle=0,width=\\columnwidth]{3a-c12-ne22-fe56-t9c.eps}  \\caption{Central temperature at the end of central helium burning ($1\\,\\%$ helium mass fraction - ``$Y=0.01$'') in units of $10^9\\,$K ($T_{9,\\mathrm{c}}$), central abundances of $^{12}$C, $^{22}$Ne, and $^{56}$Fe after core helium depletion (``$Y=0.00$'') for $25\\,\\Msun$ stars with L03 initial abundance as a function of $R_{3\\alpha}$. }\\label{hedep} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[angle=90,width=0.475\\textwidth]{LC_25Msun_pfVsRate_final.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{LB_25Msun_pfVsRate_final.eps} \\caption{Post-explosive production factors for a 25 \\Msun star as a function of $R_{3\\alpha}$. \\textbf{(a)} The LC series. \\textbf{(b)} The LB series.  The ``reaction rate multiplier'' is the factor that multiplies both the standard rates.} \\label{pfVs3alpha} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsR3alpha_LC25Msun.eps} \\hfill \\includegraphics[angle=90,width=0.475\\textwidth]{sOnlyPfVsRate_LB25Msun.eps} \\caption{Post-explosive production factors for s-only isotopes for a 25 \\Msun star as a function of $R_{3\\alpha}$. \\textbf{(a)} The LC series. \\textbf{(b)} The LB series.  The ``reaction rate multiplier'' is the factor that multiplies both the standard rates.} \\label{sOnlyPfVs3alpha} \\end{figure*} ",
        "conclusions": "We followed the entire nucleosynthesis throughout the life of massive stars, from H-burning through Fe-core collapse, followed by a supernova explosion.  We observe a strong sensitivity of the PFs for weak-\\textsl{s} process isotopes to $\\pm 2\\sigma$ variations of the rates of the triple-$\\alpha$ and of the $^{12}$C($\\alpha,\\gamma$)$^{16}$O reactions.  This can be explained by the variations in the PFs of medium-weight neutron poisons found by \\cite{tur07}, the changes in the amount of carbon left at the end of central helium burning, and the amount of $^{22}$Ne burnt in central helium burning.  In most cases, our simulations yield lower PFs for the abundances of L03 than for the abundances of AG89; the lower CNO content of the L03 abundance set is responsible for the reduced efficiency of the \\textsl{s}-processing.  This tendency is not always followed, however, as we see for the \\textsl{s}-only nuclei with the central rates, in Table~\\ref{sonlySumm}.  The production of the weak-\\textsl{s} nuclei is highly sensitive to the rates of the helium burning reactions.  In some cases (see the discussion of the results of Table~\\ref{sonlySumm}), we find that a $15\\,\\%$ change in these rates may change the nucleosynthesis rates by more than a factor of two.  We find that one must follow the entire evolution of the star to evaluate accurately the contribution of massive stars to \\textsl{s}-process abundances.  Most earlier studies took into account only the contribution of core He burning and shell C burning.  We show that significant production takes place in later burning phases as detailed for the \\textsl{s}-only isotopes in Figure~\\ref{stages}.  The burning phases beyond shell C burning mostly destroy isotopes in the core (especially near the mass cut), but overall lead to an increased production of some isotopes.  The passage of the shock wave further modifies those PFs, usually reducing them slightly.  We have examined the uncertainties in the weak-\\textsl{s} process production for Sr, Y, Zr owing to uncertainties in the helium-burning reaction rates, and find that they are in the $3\\,\\%-7\\,\\%$ range. This is smaller than the observed deficiencies of $8\\,\\%-18\\,\\%$ in the known production mechanisms that led to the introduction of the LEPP process.  On the other hand, we surely underestimate the total uncertainties.  Uncertainties in the reaction rates of the neutron producing and capture reactions are significant.  In addition, uncertainties in the stellar models can have significant effects (e.g., the treatment of hydrodynamics, convection, overshoot, etc.; see \\citealt{cos07} for a study of the effects of overshooting).  We refer the reader to \\cite{tur07} for a more complete discussion of those approximations and physics uncertainties.  Combining all these uncertainties may render the evidence for the LEPP process less convincing.  We have shown that many aspects of nucleosynthesis in supernovae, the production of the medium weight nuclei, as discussed in \\cite{tur07}, and that of the weak-\\textsl{s}-process nuclei described here, are highly sensitive to variations of the helium burning reaction rates, within their experimental uncertainties.  This further emphasizes the need for better values of the helium burning reaction rates."
    },
    "0809/0809.4151_arXiv.txt": {
        "abstract": "We have imaged the disc of the young star HL Tau using the VLA at 1.3~cm, with $0.08''$ resolution (as small as the orbit of Jupiter). The disc is around half the stellar mass, assuming a canonical gas-mass conversion from the measured mass in large dust grains. A simulation shows that such discs are gravitationally unstable, and can fragment at radii of a few tens of AU to form planets. The VLA image shows a compact feature in the disc at 65~AU radius (confirming the `nebulosity' of \\citet{welch}), which is interpreted as a localised surface density enhancement representing a candidate proto-planet in its earliest accretion phase. If correct, this is the first image of a low-mass companion object seen together with the parent disc material out of which it is forming. The object has an inferred gas plus dust mass of $\\approx 14$~M$_{\\rm Jupiter}$, similar to the mass of a proto-planet formed in the simulation. The disc instability may have been enhanced by a stellar flyby: the proper motion of the nearby star XZ Tau shows it could have recently passed the HL Tau disc as close as $\\sim 600$~AU. ",
        "introduction": "The mechanisms by which giant planets form are uncertain. Core-accretion models \\citep[e.g.]{pollack,hubickyj} have successfully linked high abundances of rocky elements in the star to higher planet-probability \\citep[e.g.]{fischer,santos}, and may explain planets with substantial rocky cores \\citep{sato} -- but have theoretical difficulties with slow planetary cooling that limits mass accretion rates and with rapid inwards migration leading to loss of cores into the star. Also, the time of around 6~Myr to complete Jupiter may conflict with the infrared detection rate of discs that declines close to zero by 6~Myr \\citep{haisch}, and with the latest ages of $\\sim 15$~Myr when gas is detected \\citep{dent} since planet completion takes longer in low-mass discs. The alternative model of gravitational instability can create a proto-planet very rapidly, on dynamical (orbital) timescales \\citep{boss,rice05}, provided that the disc cooling time is similarly short \\citep{gammie01,rafikov}. Only relatively rare discs of $\\ga 0.1 \\times M_{star}$ will be unstable, but recent radio studies \\citep[e.g]{rodmann} that account for mass in large dust grains may boost more of the total disc masses into this category.  Since metre-sized particles are gathered up into the gas fragments \\citep{rice06}, this model may also account for solid cores to giant planets.  \\begin{figure*} \\label{fig1} \\includegraphics[width=145mm,angle=0]{NATUNANSUB.CPS} \\caption{VLA 1.3 cm images towards HL Tau. Main image: natural weighting with a beam (small inset) of $0.11''$ for maximum sensitivity; contours increase by $\\sqrt{2}$ and are at (4.0, 5.7, 8.0, 11.4, 16.0) $\\times 17 \\umu$Jy/beam (1$\\sigma$). The arrow indicates the jet axes and the ellipse shows the approximate extent of the inclined disc in a previous $0.6''$ resolution image at 2.7~mm \\citep{looney}. The compact object lies to the upper-right. Upper inset: same but with central peak of 263 $\\umu$Jy subtracted, to highlight the jet bases and two features at $\\approx 20$~AU at the ends of the disc major axis. Contours from the unsubtracted image are overlaid. Lower inset: uniform weighted image at higher resolution of $0.08''$, with contours at (3.0, 4.2, 6.0, 8.5, 12.0) $\\times 21 \\umu$Jy/beam (1$\\sigma$), highlighting the compact object. } \\end{figure*}  Imaging a planet within its birth-disc would illuminate the processes involved. Companions of several Jupiter masses upwards have been imaged, but large separations from the primary and in some cases small mass-ratios of the two components suggest that these objects may have formed like binary stars \\citep{luhmann}. No discs within these systems have been imaged, so the formation processes remain obscure -- the most direct observational evidence for a forming object is a cleared cavity within the disc of AB~Aur \\citep{oppenheimer}.  In this letter, we present the first candidate for a low mass companion imaged in the accretion stage and within the parent disc. Such an object would be expected to appear as a cool condensation of dust and gas, possibly without a distinct dense core as sedimentation timescales for large dust grains can be a few $10^4$ years \\citep{helled}. Here, we use radio-wavelength data to trace the thermal emission from large dust particles in the disc around HL Tau. This pre-main-sequence Class~I (remnant envelope) object has been modelled by \\citet{tom} at around 0.33~M$_{\\odot}$ and 5~L$_{\\odot}$, seen at $< 10^5$~years old. The HL~Tau disc was selected as one of the brightest known at millimetre wavelengths, with estimates for gas plus dust mass of up to 0.1~M$_{\\odot}$ \\citep{beckwith}, and thus within the disc-to-star mass regime where instability could occur. Millimetre interferometry (resolving out the envelope) has shown emission from the dust-disc extending out to at least $\\sim 100$~AU radius \\citep{wilner,mundy,lay,looney,rodmann}. ",
        "conclusions": "Including the large grains now detected, the HL~Tau disc mass is $\\approx 0.13$~M$_{\\odot}$. \\citet{tom} find good fits to the stellar mass for 0.2--1~M$_{\\odot}$; for their best-fit value of 0.33~M$_{\\odot}$ the disc is around $0.4 M_{\\rm star}$. This proportionally massive disc should be gravitationally unstable, and a simulation at the higher end of the $M_{\\rm disc, star}$ ranges confirms that planetary objects could form at a few tens of AU. The VLA data show such a flux peak in the parent disc material, interpreted here as a surface density enhancement. (The clump is three times brighter than the local disc flux, while warming of the gas by gravitational collapse should only contribute marginally to higher emission; simulation results suggest the beam-averaged dust temperature is raised by $\\approx 50$~\\%.) This clump at 65 AU from HL Tau lies in the appropriate unstable region, and is compact as expected for a low-mass object accreting from the disc.  The simulated disc is unstable for the adopted parameters, but external forces could have increased the real disc's tendency to fragment. Notably, another cluster member, XZ Tau, appears close-by, which is unusual within the diffuse Taurus association, and the relative motions suggest a possible recent encounter of the two stars.  Their line-of-sight distances are unknown, but the similar radial velocities \\citep{folha} suggest they are not located in very different parts of the association, and in 2-D the stars are presently diverging. XZ Tau lies $23''$ east of HL Tau and the proper motions \\citep{ducourant} are (+11,--19) and (--3,--21) milliarcsec/year respectively (errors of 2-5 mas/yr). Around 1600 years earlier, the stars could thus have passed within $\\sim 600$~AU (in 2-D projection).  Such an event would have been dynamically recent, given that the compact object has an orbital period of 900~years for $M_{star}$ of 0.33~$M_{\\odot}$.  The final mass of this still-forming companion may increase, by absorbing more of the disc, but our estimate of 14~M$_{\\rm Jupiter}$ for the condensation is well down into the sub-stellar regime. If all this material is accreted, the final object would be around the brown dwarf / planet boundary by the definition of short-lived deuterium-burning capability, which occurs at $\\ga 12$--13~M$_{\\rm Jupiter}$. A more recently developed definition of a planet is a low-mass object that formed in the disc of a star. This `origins' definition sidesteps the deuterium-burning issue, which as \\citet{chabrier} point out is irrelevant for the evolution of brown dwarfs. In the case of HL~Tau `b', imaging the object within the parent disc marks it as a candiate proto-planet by this origins definition."
    },
    "0809/0809.1859_arXiv.txt": {
        "abstract": "{\\vspace{5pt}\\par   The impact of a kination-dominated phase generated by a quintessential exponen-tial model on the thermal abundance of gravitinos and axinos is investigated. We find that their abundances become proportional to the transition temperature from the kination to the radiation era; since this temperature is significantly lower than the initial (``reheating\") temperature, the abundances decrease with respect to their values in the standard cosmology. For values of the quintessential energy-density parameter close to its upper bound, on the eve of nucleosynthesis, we find the following: (i) for unstable gravitinos, the gravitino constraint is totally evaded; (ii) If the gravitino is stable, its thermal abundance is not sufficient to account for the cold dark matter of the universe; (iii) the thermal abundance of axinos can satisfy the cold dark matter constraint for values of the initial temperature well above those required in the standard cosmology. A novel calculation of the axino production rate by scatterings at low temperature is also presented.} ",
        "introduction": "\\setcounter{equation}{0}  A plethora of recent data \\cite{wmap, snae} indicates \\cite{wmapl} that the two major components of the present universe are \\emph{Cold Dark Matter} (CDM) and \\emph{Dark Energy} (DE) with density parameters \\cite{wmap} \\beq {\\sf (a)}~~\\Omega_{\\rm CDM}=0.214\\pm0.027~~\\mbox{and}~~{\\sf (b)}~~\\Omega_{\\rm DE}=0.742\\pm0.03 \\label{cdmba}\\eeq at $95\\%$ \\emph{confidence level} (c.l.). Identifying the nature of these two unknown substances, is one of the major challenges in contemporary cosmo-particle theories.  The DE component can be explained by modifying the \\emph{standard cosmology} (SC) via the introduction of a slowly evolving scalar field called quintessence \\cite{early} (for reviews, see Ref.~\\cite{der}). An open possibility in this scenario is the existence of an early \\emph{kination dominated} (KD) era \\cite{kination}, where the universe is dominated by the kinetic energy of the quintessence field; this period is an indispensable ingredient of quintessential inflationary scenaria \\cite{qinf, dimopoulos, chung}. During this era, the expansion rate of the universe is larger compared to its value during the usual \\emph{radiation domination} (RD) epoch. This implies that the relic abundance of the \\emph{weakly interacting massive particles} (WIMPs) can be significantly enhanced \\emph{with respect to} (w.r.t) its value in the SC \\cite{Kam, salati, prof, jcapa}, provided that they decouple from the thermal bath during the KD era. Further phenomenological implications of this effect for the future collider or astrophysics experiments have also been studied \\cite{kinpheno}.  WIMPs are the most natural candidates \\cite{candidates} to account for the second  major component of the present universe, the CDM. Among them, the most popular is the lightest neutralino \\cite{goldberg, lkk} which turns out to be the \\emph{lightest supersymmetric particle} (LSP) in a sizeable fraction of the parameter space of \\emph{sypersymmetric} (SUSY) models and therefore, stable under the assumption of the conservation of $R$-parity. However, SUSY theories predict the existence of even more weakly interacting massive particles, known as \\emph{e}-WIMPs \\cite{ewimps} which can naturally play the role of LSP. These are the gravitino, $\\Gr$, and the axino, $\\ax$ ($\\Gr$ is the spin-3/2 fermionic SUSY partner of the graviton, and $\\ax$ the spin-$1/2$ fermionic SUSY partner of the axion which arises in SUSY extensions \\cite{goto} of the \\emph{Peccei-Quinn} (PQ) solution \\cite{pq} to the strong CP problem). As their name indicates, the interaction rates of gravitinos and axinos are \\emph{extremely} weak, since they are respectively suppressed by the reduced Planck scale, $m_{\\rm P}=M_{\\rm P}/\\sqrt{8\\pi}$ ($M_{\\rm P}=1.22\\times10^{19}~{\\rm GeV}$ being the Planck mass) and by the axion decay constant, $f_a\\sim(10^{10}-10^{12})~{\\rm GeV}$ (for a review, see Ref.~\\cite{kim}).  Due to the weakness of their interactions, \\emph{e}-WIMPs depart from chemical equilibrium very early ($\\Gr$ at a energy scale close to $m_{\\rm P}$ and $\\tilde a$ close to $f_a$) and we expect that their relic density (created due to this early decoupling) is diluted by the primordial inflation. However, they can be reproduced in the following ways: (i) in the thermal bath, through scatterings \\cite{gravitinoc, moroi1, Bolz, steffen, axino, steffenaxino} and decays \\cite{axino, small, strumia} involving superpartners, and (ii) non-thermally \\cite{gravitinont, axinont}, from the out-of-equilibrium decay of the \\emph{next-to-LSP} (NLSP). In this paper we do not consider the possible non-thermal production of \\emph{e}-WIMPs, since this mechanism is highly model dependent (i.e., it is sensitive to the type and decay products of the NLSP). As a consequence, we do not consider either the out-of-equilibrium decay of the one \\emph{e}-WIMP to the other (as in the case of \\cref{asaka}, where $\\Gr$ is the NLSP and $\\ax$ the LSP). In all these cases, extra restrictions have to be imposed in order not to jeopardize the success of the standard Big Bang \\emph{Nucleosynthensis} (NS).  The latter requirement  has to be satisfied also for unstable $\\Gr$. This restriction imposes  a tight upper bound on the initial (``reheating\") temperature, $\\Ti$, of the universe in the SC \\cite{gravitinoc, moroi1,kohri, kohri2, oliveg}.  On the other hand, if one of the \\emph{e}-WIMPs is a stable LSP, it has to obey the CDM constraint. In particular, its relic density $\\Omega_{\\chia}h^2$ has to be confined in the region \\cite{wmap} \\beq\\label{cdmb} {\\sf (a)}~0.097\\lesssim \\Omega_{\\chia}h^2 \\lesssim 0.12~~\\mbox{for}~~{\\sf (b)}~10~\\keV\\leq m_\\chia\\leq m_{\\rm NLSP} \\eeq where $m_{\\rm NLSP}$ is the mass of the NLSP. Let us note, in passing, that the lower bound of \\sEref{cdmb}{a} is valid under the assumption that CDM is entirely composed by $X$'s and the abundance of non-thermaly produced $X$'s is negligible. The lower bound on $m_{\\chia}$ arises from the fact that smaller $m_\\chia$ cannot explain \\cite{sformation} the observed early reionization \\cite{wmap}. For about $10\\leq m_\\chia/\\keV\\leq 100$, $X$'s may constitute warm dark matter (the mass limits above are to be considered only as indicative).  In this paper we reconsider the creation of a KD era in the context of the exponential quintessential model \\cite{wet, expo}, taking into account restrictions arising from NS, the inflationary scale, the acceleration of the universe and the DE density parameter. Although this model does not possess a tracker-type solution \\cite{salati, attr} in the allowed range of its parameters, it can produce a viable present-day cosmology in conjunction with the domination of an early KD era, for a reasonable region of initial conditions \\cite{brazil, cline, silogi, german}. We then investigate the impact of KD on the thermal production of \\emph{e}-WIMPs,  solving the relevant equations both numerically and semianalytically. We find that the abundance of \\emph{e}-WIMPs becomes proportional to the transition temperature from KD to RD era, $\\Tkr$, and decreases w.r.t its value in the SC since  $\\Tkr$ can be much lower than $\\Ti$. In particular, we consider two cases, depending on whether $\\Tkr$ is higher -- \\emph{high $T$ regime} (HTR) -- or lower -- \\emph{low $T$ regime} (LTR) -- than a threshold $\\Tc\\simeq10~\\TeV$, below which the thermal production of $\\ax$ via the decay of the superpartners becomes important; for this low temperature region, a novel formulae for the production of $\\ax$'s through scatterings is presented. It turns out that, in this part of the parameter space, $\\ax$ becomes an attractive CDM candidate within the \\emph{quintessential kination scenario} (QKS).  Modifications to the thermal production of \\emph{e}-WIMPs have also been investigated in the context of extra dimensional theories \\cite{seto, panot}, where the expansion rate of the universe can be also enhanced w.r.t its value in the SC, due to the presence of an extra term including the brane-tension. However, this increase is more drastic than in the QKS. In addition, the production of $\\ax$ via scatterings  at low temperature and via the decay \\cite{axino} of the SUSY particles has not been taken into account \\cite{panot}.  We start our analysis by reviewing the basic features of the exponential quintessential model in Sec.~\\ref{sec:quint}. We then present our numerical and semi-analytical calculations of the thermal abundance of \\emph{e}-WIMPs in Sec.~\\ref{sec:boltz} and study the parameter space allowed by several requirements for $\\Gr$ (Sec.~\\ref{sec:grv}) and for $\\ax$ (Sec.~\\ref{sec:axn}). Our conclusions are summarized in Sec.~\\ref{sec:con}. Computational issues on the low temperature $\\ax$-production are discussed in Appendix A.  Throughout the text, brackets are used by applying disjunctive correspondence, natural units ($\\hbar=c=k_{\\rm B}=1$) are assumed, the subscript or superscript $0$ refers to present-day values (except in the coefficient $V_0$) and $\\log~[\\ln]$ stands for logarithm with basis $10~[e]$. Moreover, we assume that the domain wall number \\cite{kim} is equal to 1. ",
        "conclusions": "\\label{sec:con}  We presented an exponential quintessential model which generates a period dominated by the kinetic energy of the quintessence field. The parameters of the quintessential model ($\\lambda, \\Ti, \\Omqns$) were confined so that $0.5\\leq\\Omega_q(\\Ti)\\leq1$, and were constrained by current observational data originating from NS, the acceleration of the universe, the inflationary scale and the DE density parameter. We found $0<\\lambda<0.9$ and studied the allowed region in the ($\\Ti, \\Omqns$)-plane.  We proceeded to examine the impact of this KD epoch to the thermal abundance of $\\Gr$ and $\\tilde a$. We solved the problem {\\sf (i)} semi-analytically, producing approximate relations for the cosmological evolution before and after the transition from KD to RD, and solving the appropriately re-formulated Boltzmann equation that governs the evolution of the $\\chia$-number density and {\\sf (ii)} numerically, integrating the relevant system of the differential equations.  Although we did not succeed to achieve general analytical solutions in all cases, we consider as a significant development the derivation of a result by solving numerically just one equation, instead of the whole system above. Moreover, for typical values of $m_i$'s in \\Eref{mi}, empirical formulas that reproduce quite successfully our numerical results were derived. For unstable $\\Gr$, the  $\\Gr$-constraint poses a lower bound on $\\Omqns$, which turns out to be almost independent of $\\Ti$. The CDM constraint can be satisfied by the thermal abundance of $\\Gr$ for extremely low values $(10^{-23}-10^{-24})$ of $\\Omqns$. On the contrary, the former constraint can be fulfilled by the thermal abundance of $\\ax$ with values of $\\Omqns$ close to the upper bound posed by the requirement for successful NS.   Let us also comment here on three minor subtleties of our calculation which, do not alter the basic features of our conclusions (although could potentially create some quantitative modifications to our results). In particular:  \\begin{itemize}  \\item Throughout our investigation we did not identify the nature of NLSP. Therefore the upper bound, shown in Fig. 5-a [Fig. 7], on $m_{\\Gr}$ [$m_{\\ax}$] derived from the requirement [$m_{\\Gr}\\leq m_{\\tilde B }$] $m_{\\ax}\\leq m_{\\tilde B }$ could be modified if there is another SUSY particle lighter than $\\tilde B$. Moreover, we did not consider the NS constraints concerning the late decays of the NLSP into $X$'s. These constraints \\cite{axino,gravitinont, axinont} depend very much on the properties of the NLSP, i.e. its composition, its mass relative to the $m_X$, and its coupling to $X$'s. Consequently, additional bounds on the $m_X$ might arise. As the $\\ax$ interactions are not as strongly suppressed as the $\\Gr$ interactions, the $\\ax$ LSP anyhow is far less problematic than the $\\Gr$ LSP w.r.t these constraints.   \\item In the case of $\\Gr$, we did not incorporate contributions to $\\Omgr$ from the process of reheating. Indeed, these extra contributions can be  a fraction of the result shown in \\Eref{YhT} \\cite{kohri2,sahu}, in the case of the usual reheating realized by the coherent oscillations of a massive particle \\cite{turner}. However, in the QKS several reheating processes have been proposed \\cite{sami} and therefore, any  safe comparison between the SC and the QKS has to be performed for $T<\\Ti$. In other words, to keep our investigation as general as possible, we preferred to study the evolution of the universe after the start of the RD [KD] era in the SC [QKS] (we simply assumed the existence of an earlier inflationary epoch).  \\item  In the case of the $\\ax$-CDM, we used throughout our investigation some representative masses for the superpartners, given in \\Eref{mi}. Variation in these values (especially in $m_{\\gl}$ and $m_{\\sq}$) has an impact on $\\Omax$ in the LTR (e.g., for $m_{\\gl}\\simeq m_{\\sq}$ the contribution of the process $\\sq^*q$ to $C_{\\ax}^{\\rm LT}$ can be enhanced). In addition, a further uncertainty in our calculation arises from the determination of $\\Ts$ below which $C_{\\ax}^{\\rm HT}$ is replaced by $C_{\\ax}^{\\rm LT}$ in the integration of the relevant equations. Note however, that the aim of the present paper is to demonstrate the change of $\\Omx$ due to the presence of the KD era and not a full scan of the SUSY parameter space. \\end{itemize}   Although our results have been derived in the context of an exponential quintessential model, their applicability can be extended to every model that generates a KD phase, even without \\cite{kination} quintessential consequences. It is worth mentioning that in the presence of kination we can obtain simultaneous compatibility of both the gravitino and the CDM constraint (i.e., the lower bound on $\\Omqns$ from the gravitino constraint is compatible with the $\\Omqns$'s needed in order to have the correct amount of $\\ax$ CDM). It would be probably interesting to check if we can obtain a simultaneous compatibility of these two constraints with additional bounds, arising e.g. from leptogenesis and neutrino masses \\cite{kinlept, Bento} or the quintessino abundance \\cite{quintessino}.   \\ack We would like to thank K.Y. Choi, R. Ruiz de Austri and L. Roszkowski for helpful discussions. The research of S.L was funded by the FP6  Marie Curie Excellence Grant MEXT-CT-2004-014297. The work of M.E.G, C.P and J.R.Q was supported by the Spanish MEC project FPA2006-13825 and the project P07FQM02962 funded by the ``Junta de Andalucia''.   \\appendix"
    },
    "0809/0809.1890_arXiv.txt": {
        "abstract": "In an earlier investigation, we proposed population boundaries for both inspiralling and mass-transferring double white dwarf (DWD) systems in the distance independent ``absolute'' amplitude-frequency domain of the proposed space-based gravitational-wave (GW) detector, {\\it LISA}. The degenerate zero temperature mass-radius (M-R) relationship of individual white dwarf stars that we assumed, in combination with the constraints imposed by Roche geometries, permits us to identify five key population boundaries for DWD systems in various phases of evolution. Here we use the  non-zero entropy donor M-R relations of \\cite{DB2003} to modify these boundaries for both DWD and neutron star-white dwarf (NSWD) binary systems. We find that the mass-transferring systems occupy a larger fraction of space in ``absolute'' amplitude-frequency domain compared to the simpler $T=0$ donor model. We also discuss how these boundaries are modified with the  new evolutionary phases found by \\cite{Deloyeetal2007}. In the initial contact phase, we find that the contact boundaries, which are the result of end of inspiral evolution,  would have some width, as opposed to an abrupt cut-off described in our earlier $T=0$ model. This will cause an overlap between a DWDs $\\&$ NSWDs evolutionary trajectories, making them indistinguishable with only LISA observations within this region.  In the cooling phase of the donor, which follows after the adiabatic donor evolution, the radius contracts, mass-transfer rate drops and slows down the orbital period evolution. Depending upon the entropy of the donor, these systems may then lie inside the fully degenerate $T=0$ boundaries, but LISA may be unable to detect these systems as they might be below the sensitivity limit or within the unresolved DWD background noise. We assess the limits and applicability of our theoretical population boundaries with respect to observations and find that a measurement of $\\dot f$ by LISA at high frequencies (Log $[f] \\geq 2$) would likely distinguish between DWD/NSWD binary. For low frequency sources, GW observations alone would unlikely tell us about the binary components, without the help of electromagnetic observations. ",
        "introduction": "\\label{sec1} The proposed space-based Gravitational-wave (GW) detector, {\\it LISA}\\footnote{http://lisa.nasa.gov} ({\\it Laser Interferometer Space Antenna}) \\citep{FB84,EIS87,Bender98}, is sensitive to GWs in the $10^{-4} - 1$ Hz frequency range. Within this band, one of the most promising sources are double white dwarf (DWD) binary systems, as it is expected that a large fraction of main-sequence binaries end their lives as close DWDs \\citep{IT84, IT86}. For this reason, the GWs emitted by these systems in our Galaxy may form a background noise in the low frequency ($\\leq 3 \\times 10^{-3}$ Hz) band of {\\it LISA}. The population of DWDs in our Galaxy is expected to be dominated by systems that undergo two distinct, long-lived phases of evolution: an ``inspiral'' phase, where both the stars are detached from their Roche lobes and the loss of angular momentum in the form of GW emission causes the two stars to slowly spiral in towards each other; and a ``stable mass transfer'' phase, where the less massive star fills its Roche lobe initially and starts transferring mass steadily to its companion. An example of the stable mass transferring systems are the AM CVn type systems, of which 18 \\citep{Nelemans2005} are known through electromagnetic observations\\footnote{Two controversial candidate systems, RX J0806+15 and V407 Vul may change their number to $16$. See \\cite{Cropper1998, Wu2002, MS2002} for more details.}.  Apart from DWDs, neutron star white dwarf (NSWD) binary systems are also  one of the promising sources of GWs for LISA. Various authors \\citep{Kim2004, cooray2004, Nelemans2001} have estimated the GW background from these systems and concluded that the number of NSWD systems detectable with LISA is 1-2 orders of magnitude less than DWD systems. Similar to DWDs, NSWD systems also undergo inspiral and stable mass transfer phases. Specifically, several studies \\citep[see for example:][]{Nelson1986, BT2004a, Nelemans2006} have suggested that one of the possible formation scenarios of the so-called ultra compact x-ray binary (UCXB) systems, with orbital periods  $\\le 80$ minutes, is that a low mass white-dwarf donor ($\\le 0.1 M_\\mathrm{\\odot}$) transferring mass to an accreting neutron star (NS) primary in a short orbit. In this scenario, a detached NSWD binary system initially evolves to a minimum orbital period as  angular momentum is lost from the system due to GW radiation. At this minimum orbital period, the companion white dwarf (WD) star starts filling its Roche lobe and transfers mass to the NS  and the system evolves to longer orbital periods. At present, there are $12$ known UCXB systems with measured orbital periods, see Table\\ref{table1}. Some of these systems have accreting millisecond pulsars (XTE J1807-294, XTE J1751-305 \\& XTE J0929-314: see \\cite{Mark2002}, for example) and several of them are found in globular clusters as the stellar density and close encounters are more common \\citep{clark1975,ivanova2007}. The total population of field UCXBs may be low, $\\sim 10$ \\citep{BT2004a, cooray2004} and theoretical studies by \\cite{BT2004b} indicates that even at the Galactic center, accreting NS systems do not contribute much to the faint x-ray population.  In general, the capabilities of {\\it LISA} as a GW detector are usually discussed in the context of the log $(h)$-- log $(f)$ domain. \\cite{KT2007} proposed  that an analogy can be drawn between the astronomy community's familiar color-magnitude (CM) diagram and {\\it LISA}'s amplitude-frequency  diagram. For LISA sources, an analogous quantity to absolute magnitude $M$ is $\\log(rh)$, where $r$ is the distance to the source. The underlying physical properties of compact binary  systems such as DWDs and NSWDs, their evolution, and their relationship to one another in the context of stellar populations can be ascertained only if the observational properties of such systems are displayed in a $\\log(rh) - \\log(f)$ diagram, rather than in a plot of $\\log(h)$ versus $\\log(f)$. \\cite{KT2007} discussed the DWD binary systems in this context of ``absolute'' amplitude-frequency domain assuming that the donors in these systems follow zero temperature mass-radius relation. Here, we will consider the effect of ``warm'' donors as proposed in \\cite{DB2003, Deloyeetal2005, Deloyeetal2007} and extend the discussion to NSWD binary systems and compare the population boundaries between both of them. We also discuss the limits and applicability of these population boundaries in the log$(rh)$ - log$(f)$ space within the context of observed AM CVn and UCXB systems. ",
        "conclusions": "\\label{sec3} The population boundaries for DWD $\\&$ NSWD binary systems discussed in previous sections were generated using \\cite{DB2003} non-zero entropy donor models,  which assume that the donors are fully convective and undergo adiabatic evolution throughout  the mass-loss phase. Recently, \\cite{Deloyeetal2007} showed that these assumptions may not properly estimate the donor's orbital period evolution and specifically the donor's adiabatic evolution may not hold true for whole mass-transfer phase. Instead they identified three distinct evolutionary phases and here we discuss the implications of their new findings on our population boundaries.  During the first phase, which happens during the mass-transfer ``turn-on'' phase (when the donor comes into contact initially), the radius of the donor decreases , the mass-transfer rate $\\dot M_\\mathrm{d}$ increases and the orbital period $P_\\mathrm{orb}$ continues to decrease until the donor radius reaches it's minimum value ($\\dot M_\\mathrm{d}$ becomes maximum) and starts expanding again. \\cite{Deloyeetal2007} calculated that this turn on phase lasts up to $\\sim  10^{6}$ years. In the second phase, the donor responds to the mass loss adiabatically and starts expanding, which is considered to be the normal AM CVn phase. But this phase of adiabatic expansion ends and a third phase of evolution begins at around $P_\\mathrm{orb} \\sim 45$ min, when the mass-transfer rate (and the donors thermal time) drop enough for the donor to cool and start contracting to a fully degenerate configuration, stalling the $P_\\mathrm{orb}$ evolution.  In the case of DWDs, the initial turn-on phase will result in the $q=1$ contact boundary (green curve with crosses in Fig. \\ref{fig1}) to have a width, instead of a sharp boundary as discussed in \\cite{KT2007}. This is because it is calculated assuming that once the system comes into contact, it would evolve towards lower amplitudes and frequencies, whereas $P_\\mathrm{orb}$ decreases in the turn on phase, to a minimum even after the initial contact. Accordingly, there will be a slight overlap of lower contact boundaries for He, C $\\&$ O donors. Once the system evolves off this initial contact phase, the donor expands adiabatically in response to the mass-transfer and the system follows a typical AM CVn evolution,  where the GW amplitude $\\&$ frequency keeps decreasing as $P_\\mathrm{orb}$ increases. Assuming an adiabatic evolution means that the cooling time of the donor is longer than the mass-transfer time-scale ($M_\\mathrm{d} / \\dot M_\\mathrm{d}$), it will in turn affects the orbital evolution time-scale. Accordingly, the donor follows a trajectory where it passes through different isotherms of decreasing $q$, after the initial contact. As \\cite{KT2007} illustrate, the lower right curve in Fig. \\ref{fig1} shows the contact boundary for $q=1$ systems and similar contact boundaries for lower $q$'s would lie to the left of it, but the contact boundaries will shift towards lower amplitudes and frequencies. This phase of adiabatic evolution comes to an end between $P_\\mathrm{orb} \\approx 40-55$ min, when the donor starts to cool and contract eventually towards a fully degenerate  star. This will drastically (almost an order of magnitude, see \\cite{Deloyeetal2007} Fig. 15) reduce $\\dot M_\\mathrm{d}$ at these long orbital periods. Hence, after this range of $P_\\mathrm{orb}$, the systems GW frequency evolution slows down in accordance with the drop in $P_\\mathrm{orb}$ evolution. If the donor has cooled enough to approximate it as a $T=0$ degenerate model, then it may lie close to one of the $T=0$ contact boundaries, depending upon its $q$. But these $T=0$ boundaries are bounded within the region constrained by $M_\\mathrm{a} = M_\\mathrm{ch}$ (light blue dot-dashed) curve and the $q=1$ contact boundary (green curve with crosses in Fig. \\ref{fig1}), because these boundaries are drawn assuming that the donor is a fully degenerate star with $T=0$ and such a system can not exist beyond these curves. Therefore, we have drawn a vertical (brown) dot-dashed line in Fig. \\ref{fig1} at $P_\\mathrm{orb} = 55$ min (Log$f = 3.21$) beyond which we would expect these systems to be within the $T=0$ contact boundaries. Note that this vertical line is not a sharp boundary: some systems may not reach a fully degenerate configuration by this $P_\\mathrm{orb}$. Rather, a system's $P_\\mathrm{orb}$ evolution slows down at a particular GW frequency once they start cooling towards a degenerate configuration at the above mentioned $P_\\mathrm{orb}$. Fig. \\ref{fig1} also shows the observed AM CVn type systems, for which the masses and $P_\\mathrm{orb}$ are taken from \\cite{Deloyeetal2005} $\\&$ \\cite{Roelofs2007}. Couple of them are clearly outside $T=0$ region; for some of them, the system's mass function limits the minimum donor mass to a value above that of a Roche-filling $T=0$ donor at the same $P_\\mathrm{orb}$.  It is very difficult to know their exact temperature and/or composition purely from GW observations, as there is a good overlap of systems with these characteristics\\footnote{Out of these systems, CE 315 (the system with lowest GW amplitude and frequency in Fig. \\ref{fig1}) has $P_\\mathrm{orb} = 65$ min and lies to the left of the $P_\\mathrm{orb} = 55$ line and below $T=0, M_\\mathrm{a} = M_\\mathrm{ch}$ boundary. This may seem to imply that this system may be cooling off and trying to reach a $T=0$ degenerate configuration. But in Fig 1. of \\cite{Deloyeetal2005}, they give a temperature range of Log $T = 4.0-6.5$ and \\cite{Bildsten2006} noted that this system may have a hot donor.  If that is the case, \\cite{Bildsten2006} note that this donor may have evolved with constant entropy that it was born with or may have been heated by either the disk or the accretor. However, if irradiation is the cause, then \\cite{Deloyeetal2007} note that it will delay the onset of donor's cooling and  increases the temperature of the donor. So though this system lies below the  $T=0, M_\\mathrm{a} = M_\\mathrm{ch}$ boundary, it does not necessarily mean that the donor can be approximated as a zero-temperature object. }.  For NSWDs, similar to DWDs, the  contact boundaries  with He composition drawn  assuming NS mass is $1.2 M_\\odot$ (green curve with crosses ) and $3.0 M_\\odot$ (green dashed curve) will also have a width due to decrease in $P_\\mathrm{orb}$ even after the contact. Furthermore, the lower contact boundary lies {\\it below} the upper contact boundary for DWDs ($M_\\mathrm{a} = 1.44 M_\\mathrm{\\odot}$). This will make the contact boundaries to ``overflow'' into the DWD region and LISA observations may not be able to distinguish these two types of systems in this region. Here also the donor undergoes adiabatic evolution after $\\dot M_\\mathrm{d}$ reaches a maximum and consequently enters the cooling phase at $P_\\mathrm{orb}\\approx 55$ min. Accordingly, in Fig. \\ref{fig2}, the brown line shows the orbital period of this transition phase. Beyond this line, the donors should start cooling and GW frequency slows down accordingly. Fig. \\ref{fig2} also shows the currently observed UCXBs, which we assume to be mass-transferring NSWD binary systems. Table \\ref{table1} gives the values of the masses of donors and orbital periods that we used. Some of the UCXBs shown in Fig. \\ref{fig2} lie {\\it below} the lower contact boundary (green cross curve), indicating that probably the donors are not of He composition. Although, this contact boundary has a width, and hence they may have He composition, it is unlikely that we can observe them during this relatively short lived phase. Moreover, these systems are plotted assuming the minimum mass of the donors mass range derived from observations, so this provides additional uncertainty in determining the composition of donors. It is unlikely that LISA will be able to observe  cooling donors in either DWD or NSWD systems because of the instrumental and/or DWD background noise. In Fig.\\ref{fig3}, we plot the known UCXBs and AM CVns on top of LISA's sensitivity curve\\footnote{http://www.srl.caltech.edu/~shane/sensitivity/MakeCurve.html} (SNR = 1) to assess the detectability of these systems. In the case of known UCXBs, it is clear that only one system (4U 1820-30) has enough signal strength to be visible to LISA, whereas some of the known AM CVns emit GWs above the instrumental and DWD background noise.     Recalling from \\S\\ref{sec1},  in order to transform from log$(h_\\mathrm{norm})$ to log$(rh_\\mathrm{norm})$  space, we need to know the distance $r$ to the binary system. The relation between the unknown binary parameters $r$, $M_\\mathrm{tot}$ and $q$ and the observables $h_\\mathrm{norm}$, $f$ and $\\dot f$ can be written as \\citep{KT2007} \\begin{eqnarray} \\label{mass_r_relation} \\frac{M_\\mathrm{tot}^{5}}{r^{3}}\\biggl[\\frac{q}{(1+q)^{2}}\\biggr]^{3} &=& \\frac{c^{12}} {2^{6}\\pi^{2}G^{5}}\\frac{h_\\mathrm{norm}^{3}}{f^{2}} , \\\\ r(1 - 2 g) &=& \\frac{5 c}{24 \\pi^{2}} \\frac{\\dot f}{h_\\mathrm{norm} f^{3}} \\label{r_fdot_relation} \\end{eqnarray} where $g=0$ in the inspiral phase of evolution and hence, it is easy to determine $r$ from  $h_\\mathrm{norm}$, $f$ and $\\dot f$ through Eq.(\\ref{r_fdot_relation}). For mass-transferring systems $g$  is a function of $M_\\mathrm{tot}$ and $q$, and they can be related to the observable $f$ by the requirement that in the mass transfer phase, $R_\\mathrm{d} = R_\\mathrm{L}$.   The determination of $r$ and/or the masses of the stars in DWD/NSWD binary system depends on the determination of $\\dot f$.  If an $\\dot f$ can not be measured for a system, then there is no way to tell whether that system is a NSWD or DWD system based only on LISA observations. If an $\\dot f$ can be measured, and if it turns out to be negative, then it is possible that particular system is a mass-transferring system (DWD or NSWD). But as shown in Figs. \\ref{fig1} $\\&$ \\ref{fig2}, it will still not be possible to know the type of the system, at least for low frequency sources (Log$[f] \\lessapprox -2$). There is a fairly good overlap in $\\dot f$ between DWD $\\&$ NSWD systems in this region because of the non-zero entropy nature of the donors and also due to the lower limit on the mass of the NS ($1.2 M_\\odot$). But for high frequency mass-transferring sources (Log$[f] \\gtrapprox -2$), it {\\it may} still be possible to know the type of the system, as the overlap region reduces\\footnote{There still will be some uncertainty for systems with NS mass lower than $M_\\mathrm{ch}$, but for NS masses higher than  $M_\\mathrm{ch}$, LISA should be able to distinguish both types of systems through the measurement of $\\dot f$. }. The same thing can be said about inspiralling systems because there is a large  overlap of DWD and NSWD inspirals at low frequencies  and even a measurement of positive $\\dot f$ would unlikely be able distinguish these two types of inspiralling systems."
    },
    "0809/0809.0966.txt": {
        "abstract": "An abundance analysis is presented of 60 metal-poor stars drawn from catalogues of nearby stars provided by Ariyanto et al. (2005) and Schuster et al. (2006). In an attempt to isolate a sample of metal-weak thick disc stars, we applied the kinematic criteria $V_{\\rm rot} \\geq 100$ km s$^{-1}$, $|U_{LSR}| \\leq 140$ km s$^{-1}$, and $|W_{LSR}| \\leq 100$ km s$^{-1}$. Fourteen  stars  satisfying these criteria and having [Fe/H] $\\leq -1.0$ are included in the sample of 60 stars. Eight of the 14 have [Fe/H] $\\geq -1.3$ and may be simply thick disc stars of slightly lower than average [Fe/H]. The other six have [Fe/H] from $-1.3$ to $-2.3$ and are either metal-weak thick disc stars or halo stars with kinematics mimicking those of the thick disc. The sample of 60 stars is completed by eight thick disc stars, 20 stars of a hybrid nature (halo or thick disc stars), and 18 stars with kinematics distinctive of the halo. ",
        "introduction": "Stars in the the Sun's immediate neighbourhood belong in the main to the Galactic disc with a sprinkling of halo stars. The majority of the disc stars are from the thin disc with a  minority from the thick disc. Introduction of the concept of the thick disc is broadly  attributed  to Gilmore \\& Reid (1983) who from star counts found that a double-exponential offered a fit to the space density perpendicular to the Galactic plane toward the south Galactic pole. The dominant component with a scale height of 300 pc was the familiar disc, now referred to as the thin disc. The second component with a scale height of 1350 pc was dubbed the thick disc. At the Galactic plane, the thick disc stars comprise no more than a few per cent of the thin disc population. Obviously, the fractional population represented by the thick disc increases with height above the plane.  After its introduction, the thick disc was a controversial innovation for quite some years but, today, many properties of the thick disc in the solar neighbourhood are rather well determined. Many studies have found thick disc stars to be old with ages in the range of 8-13 Gyr (Fuhrmann 1998; Reddy et al. 2006) but other studies have proposed ages as young as 2 Gyr for some thick disc stars (Bensby et al. 2007). Their velocity dispersion perpendicular to the Galactic plane is about 40 km s$^{-1}$, a value to be compared with about 20 km s$^{-1}$ for the old thin disc, and 90 km s$^{-1}$ for the halo. Thick disc stars lag behind thin disc stars in rotation about the centre of the Galaxy by about 50 km s$^{-1}$. The mean metallicity of the thick disc is about  [Fe/H]  $= -0.6$ with most stars falling in the interval [Fe/H] of $-0.3$ to $-1.0$. Differences in relative abundances [X/Fe] between thin and thick disc stars are now well established over this [Fe/H] range (Bensby et al. 2005; Reddy, Lambert, \\& Allende Prieto 2006).  Remaining significant controversies surround the upper and lower bounds to the [Fe/H] distribution function for thick disc stars. In this paper, we are concerned with the lower bound for thick disc stars in the solar neighbourhood. Stars of the thick disc with [Fe/H] less than about $-1$ are referred to as `metal-weak thick disc' stars (here, metal-weak thick disk stars). The metal-weak thick disk has now been the subject of many investigations. But, although many studies of putative metal-weak thick disk stars exist, the results as far as their composition are concerned  have been limited to a measurement of metallicity, usually [Fe/H]. In broad terms, two kinds of samples have been discussed: stars in the Sun's immediate neighbourhood and  giant stars beyond the solar neighbourhood. The significant advantage provided by local stars is that the sample has  well determined kinematics. The  advantage of a sample composed of stars at greater distances from the Galactic plane is that the relative fraction of thick to thin disc stars will be higher than locally but the obvious disadvantage is, in general, that knowledge of stellar kinematics is compromised by the lack of accurate distance and the small amplitude of the proper motions. In addition, halo stars provide an increasingly severe contamination of stellar samples as distance from the Galactic plane increases.  The persisting  controversy about the existence of metal-weak thick disk stars may be traceable to reconsideration of the metallicities obtained  by Morrison et al. (1990) from DDO photometry for their sample of giants with disclike kinematics. Ryan \\& Lambert (1995) undertook a high-resolution spectroscopic analysis of a subset of Morrison et al.'s stars and showed that the most of the stars identified as belonging to the metal-weak thick disk have metallicities [Fe/H] $>$ $-1$ and are thus `normal' thick disc stars. Twarog \\& Anthony-Twarog (1994, 1996) provided a recalibration of the DDO photometry that also weakened Morrison et al.'s evidence for existence of metal-weak thick disk stars. These critiques of the pioneering paper on metal-weak thick disk stars have percolated through the literature. Beers et al. (2002), ardent advocates for metal-weak thick disk stars as an important component of Galactic structure, wrote that  `acceptance of their [metal-weak thick disk stars] presence has been cast in doubt because of incorrectly assigned metallicities'. Oddly, major studies of  candidate metal-weak thick disk stars have rarely obtained the metallicity and never elemental abundances from high-resolution high-signal-to-noise ratio spectra.  In this paper, we  search for metal-weak thick disk stars among two recent  catalogues of local main sequence stars with reliable kinematics and metallicities. A comprehensive abundance analysis is undertaken for metal-weak thick disk candidates and halo stars in order to search for definitive differences in composition between candidate metal-weak thick disk stars and halo stars of similar metallicity. ",
        "conclusions": "Our collection of 14 stars in Table 1 are offered as metal-weak thick disk candidates but by no means as certain members of the elusive metal-weak thick disk population. We have been unable to identify a conclusive signature distinguishing a metal-weak thick disc star from a halo star. In terms of age, present knowledge and precision of age determinations do not provide a discriminant between thick disc and halo. At a given [Fe/H], the relative abundances [X/Fe] of metal-poor stars appear to show a dispersion no larger than the measurement uncertainties. Relative to the uncertainties, the halo stars and our metal-weak thick disk candidates of similar [Fe/H] are not distinguishable. Our sole proposed discriminant involves a combination of the LSR-velocities.  This discriminant is imperfect because kinematics of the thick disc and halo involve overlapping distribution functions. The overlap is essentially complete for the $U$ and $W$ components: the thick disc velocity distributions are wholly contained within the broader halo distributions. The overlap for the $V$ distributions is dependent on the Fe-dependence of the (unknown) velocity distribution of the halo stars (see the difference between Figures~2a and 2b) and on the extrapolation of the thick disc velocity distribution into the regime of [Fe/H] $\\leq -1$, as discussed above. The $V_{\\rm rot}$ of the metal-weak thick disk candidates is comparable to that of Schuster et al's second group of thick disc stars. While our velocity criteria are intended to favour selection of thick disc stars, contamination of the metal-weak thick disk candidates by halo stars remains a possibility. The eccentricity is no more than an intriguing discriminator. Selection of metal-weak thick disk candidate stars by their $V_{\\rm rot}$ ensures that they have a lower eccentricity than the stars of much lower $V_{\\rm rot}$ in Table 2. %The distribution %functions for the LSR-velocities admit of stars with velocities %satisfying the criteria we impose on the metal-weak thick disk candidates.  Selection effects influencing the two catalogues certainly play a role. This is strongly suggested, as noted above, by the rather different distributions of metal-poor ([M/H] $< -1$) stars in the Arifyanto et al. catalogue (Figure~2a) and the Schuster et al. catalogue (Figure~2b). In Figure~2a, the stars are rather symmetrically distributed about $V_{\\rm rot}$ = 0, the boundary between stars on prograde and retrograde orbits. In contrast, Figure~2b shows a preponderance of retrograde orbits and a mean $V_{\\rm rot}$ decreasing with decreasing [M/H]. Setting aside the influence of selection effects, very different models of the halo $V_{\\rm rot}$ distribution versus [M/H] result from the two catalogues. In turn, these different models would result in different conclusions about the extension of the thick disc to [M/H] $< -1$ that result from subtraction of the halo distribution from the observed distribution. Most probably, neither represents the true halo distribution.  The small sample of metal-weak thick disk candidates have the composition established for the thick disc and the halo at their interface. Obviously, the case for a metal-weak thick disk population needs refinement beginning with the isolation of  a larger sample with reliable metallicities assuring that stars are indeed metal-weak and reliable kinematics from accurate distances, radial velocities, and proper motions to establish that the stars have disc-like motions. The tools with which to meet these goals are on the horizon.  Inevitably, the tools will enable stars of the halo, the  thick disc and its metal-weak tail to be examined at up to several kpc from the Sun. Examination of the populations characteristics as a function height from the Galactic plane and distance from the Galactic centre will provide much needed information for digestion by theorists seeking to unravel the history of the Galaxy. Results are now beginning to appear from analyses of the SDSS database - see, for example, Allende Prieto et al. (2006) and Ivezi\\'{c} et al. (2007) who use the SDSS data to analyse spectra of large numbers of F and G dwarfs at kpc distances from the Sun. It will now be interesting to obtain detailed abundances (i.e., [X/Fe]) for samples of these faint stars.  Two considerations strongly imply that stars different in [X/Fe] at a given [Fe/H] should be uncovered by examining such samples. First, stars in dwarf spheroidal galaxies (dSphs) and the Magellanic Clouds have lower [$\\alpha$/Fe] at a given [Fe/H] than the Galactic halo and thick disc stars; a compilation by Koch et al. (2008) shows that across the range of [Fe/H] from about $-0.5$ to $-3$ the [$\\alpha$/Fe] of the dSphs  is 0.2 to 0.3~dex smaller than for the Galactic halo and thick disc stars (see their Figure~3). Second, the Galaxy is known to be accreting dSphs. Therefore, accretion of stellar systems akin to the present day dSphs is expected to add stars of lower than typical [$\\alpha$/Fe] to the halo populations. Conversely, the seeming rarity of such stars in the general local stellar population implies that systems like the dSphs did not play a major role in the merger history of the halo and thick disc.   This research has been supported in part by the Robert A. Welch Foundation of Houston, Texas. We thank anonymous referee for the detailed review which helped to improve the paper.  \\newpage"
    },
    "0809/0809.2227_arXiv.txt": {
        "abstract": "{} {The BL Lac object S5\\,0716+71 was observed in a global multi-frequency campaign to search for rapid and correlated flux density variability and signatures of an inverse-Compton (IC) catastrophe during the states of extreme apparent brightness temperatures.} {The observing campaign involved simultaneous ground-based monitoring at radio to IR/optical wavelengths and was centered around a 500-ks pointing with the INTEGRAL satellite (November 10--17, 2003). Here, we present the combined analysis and results of the radio observations, covering the cm- to sub-mm bands. This facilitates a detailed study of the variability characteristics of an inter- to intra-day variable IDV source from cm- to the short mm-bands. We further aim to constrain the variability brightness temperatures ($T_{B}$) and Doppler factors ($\\delta$) comparing the radio-bands with the hard X-ray emission, as seen by INTEGRAL at 3--200\\,keV.} {0716+714 was in an exceptionally high state and different (slower) phase of short-term variability, when compared to the past, most likely due to a pronounced outburst shortly before the campaign. The flux density variability in the cm- to mm-bands is dominated by a $\\sim 4$\\,day time scale amplitude increase of up to $\\sim\\,35$\\,\\%, systematically more pronounced towards shorter wavelengths. The cross-correlation analysis reveals systematic time-lags with the higher frequencies varying earlier, similar to canonical variability on longer time-scales. The increase of the variability amplitudes with frequency contradicts expectations from standard interstellar scintillation (ISS) and suggests a source-intrinsic origin for the observed inter-day variability. We find an inverted synchrotron spectrum peaking near 90\\,GHz, with the peak flux increasing during the first 4 days. The lower limits to $T_{B}$ derived from the inter-day variations exceed the $10^{12}$\\,K IC-limit by up to 3--4 orders of magnitude. Assuming relativistic boosting, our different estimates of $\\delta$ yield robust and self-consistent lower limits of $\\delta \\geq 5 - 33$ -- in good agreement with $\\delta_{VLBI}$ obtained from VLBI studies and the IC-Doppler factors $\\delta_{IC}\\,>$\\,14--16 obtained from the INTEGRAL data.} {The non-detection of S5\\,0716+714 with INTEGRAL in this campaign excludes an excessively high X-ray flux associated with a simultaneous IC catastrophe. Since a strong contribution from ISS can be excluded, we conclude that relativistic Doppler boosting naturally explains the apparent violation of the theoretical limits. All derived Doppler factors are internally consistent, agree with the results from different observations and can be explained within the framework of standard synchrotron-self-Compton (SSC) jet models of AGN.} ",
        "introduction": "Rapid intensity and polarization variability on time scales of hours to days is frequently observed in compact blazar cores at centimeter wavelengths. Since its discovery in 1985 such IntraDay variability (IDV, Witzel et al. 1986, Heeschen et al. 1987) has been found in a significant fraction ($\\sim$\\,20--30\\%) of flat-spectrum quasars and BL\\,Lac objects (Quirrenbach et al. 1992, Kedziora-Chudczer et al. 2001, Lovell et al. 2003). The `classical' IDV sources \\citep[of type II,][]{1987AJ.....94.1493H} show flux density variations in the radio bands on time scales of $\\lesssim$\\,0.5\\,--2\\,days and variability amplitudes ranging from a few up to 30\\,\\%. Stronger and often faster variability is seen in both linear polarization \\citep[e.g.][]{1989A&A...226L...1Q,2003A&A...401..161K} and more recently also in circular polarization \\citep[][]{2000ApJ...538..623M}.  When the observed IDV is interpreted as being source-intrinsic, the apparent variability brightness temperatures of $T_{B}\\,\\geq\\,10^{18}$\\,K, largely exceed the inverse Compton (IC) limit of $10^{12}$\\,K \\citep[][]{1969ApJ...155L..71K,1994ApJ...426...51R} by several orders of magnitude \\citep[see also][]{1995ARA&A..33..163W}. Consequently, extreme relativistic Doppler-boosting factors ($\\delta\\,>$\\,100) would be required to bring these values down to the IC-limit. As such high Doppler-factors (and related bulk Lorentz factors) are not observed in AGN jets with VLBI, alternate jet (and shock-in-jet) models use non-spherical geometries, which allow for additional relativistic corrections \\citep[e.g.][]{1999ApJ...511..136S,1991A&A...241...15Q}. It is also possible that the brightness temperatures are intrinsically high and coherent emission processes are involved \\citep[][]{1992ApJ...391L..59B,1998MNRAS.301..414B}. Another possibility to explain intrinsic brightness temperatures of $>\\,10^{12}$\\,K is a more homogeneous ordering of the magnetic field \\citep[][]{2006ChJAA...6..530Q}.  Owing to the small source sizes in compact flat-spectrum AGN, interstellar scintillation (ISS) is an unavoidable process. For a small number of more rapid, intra-hour variable sources, seasonal cycles and/or variability pattern arrival time delays are observed. This is used as convincing argument that such rapid variability is caused by ISS \\citep[PKS\\,0405$-$385, PKS\\,1257$-$326 and J1819+3845, e.g.][] {2000aprs.conf..147J,2003ApJ...585..653B,2002Natur.415...57D}. The situation for the slower and `classical' (type II) IDV sources is less clear. The much longer and complex variability time scales and the existence of more than one characteristic variability time-scale, did not yet allow to unambiguously establish seasonal cycles nor arrival time measurements and by one of this proof that IDV of type II is solely due to ISS. In fact, many of these sources show an overall more complex variability behavior, \\citep[e.g. QSO 0917+624, BL 0954+658, J\\,1128+5925;][]{2001ApJ...550L..11R,2001A&A...370L...9J, 2002PASA...19...64F,2004PhDT........38K,2006astro.ph.10795B,2007A&A...470...83G}, which points towards a more complicated and time-variable blending between propagation induced source extrinsic effects and source intrinsic variability.  In order to shed more light on the physical origin of IDV in such type II sources, we performed a coordinated broad-band variability campaign on one of the most prominent IDV sources. Since the late 1980's, the compact blazar S5\\,0716+71 (hereafter 0716+714) showed strong and fast IDV of type II whenever it was observed. In the optical band, 0716+714 is identified as a BL\\,Lac, but has unknown redshift. A redshift of $z\\,>$\\,0.3 was deduced using the limits on the surface brightness of the host galaxy \\citep[][]{1991ApJ...372L..71Q,1993A&AS...97..483S,1996AJ....111.2187W,2005ApJ...635..173S,2008arXiv0807.0203N}. In this paper we will use $z\\,\\geq$\\,0.3, which is within the measurement uncertainty of the previous redshift estimates. Consequently, only lower limits to any linear size and brightness temperature measurement can be obtained. The observed IDV time scales imply -- via the light travel time argument -- source sizes of only light-hours (corresponding to micro-arcsecond scales) leading to lower limits of $T_{B} \\geq 10^{18}$\\,K. Direct measurement of the nucleus with space-VLBI at 5 GHz gives a robust lower limit of $T_{B} \\ge\\!2\\times 10^{12}$\\,K, implying a minimum equipartition Doppler factor of $\\delta \\!\\ge\\!4$ \\citep[][]{2006A&A...452...83B}. The VLBI kinematics leads to estimates of the bulk jet Lorentz factor $\\gamma_{\\rm min}\\,\\geq$\\,16--21, based on the measured jet speeds \\citep[][]{2001ApJS..134..181J,2005A&A...433..815B}.  This is is unusually high for a BL\\,Lac object.  0716+714 is one of the best studied blazars in the sky \\citep[e.g.][] {1996AJ....111.2187W,1998A&A...333..445G,2000A&AS..141..221Q,2003A&A...401..161K, 2005A&A...433..815B,2006A&A...452...83B}. It has been observed repeatedly in various multi-frequency campaigns \\citep[e.g.][]{1990A&A...235L...1W,1996AJ....111.2187W,1999hxra.conf..170O, 1999A&A...351...59G,2003A&A...400..477T,2003A&A...402..151R} and is well known to be extremely variable ($\\lesssim$\\,hours to months) at radio to X-ray bands \\citep[e.g.][and references therein]{1996AJ....111.2187W,2003A&A...402..151R}. 0716+714 is so far the only source in which a correlation between the radio/optical variability was observed \\citep[see][]{1991ApJ...372L..71Q,1995ARA&A..33..163W} indicating that the radio/optical emission has a common and source-intrinsic origin, during the observations in 1990 \\citep[][]{1996ChA&A..20...15Q,2002ChJAA...2..325Q}. Further, the detection of IDV at mm-wavelengths and the observed frequency dependence of the IDV variability amplitudes \\citep[][]{2002PASA...19...14K,2003A&A...401..161K} make it difficult to interpret the IDV in this source solely by standard ISS.  In a source-intrinsic interpretation of IDV it is unclear if and for how long the IC-limit can be violated \\citep[][]{1992ApJ...391..453S,2002PASA...19...77K}. The quasi-periodic and persistent violation of this limit (on time scales of hours to days) would imply subsequent Compton catastrophes, each leading to outbursts of IC scattered radiation. The efficient IC-cooling would rather quickly restore the local $T_{B}$ in the source, and the scattered radiation should be observable as enhanced emission in the X-/$\\gamma$\\,-ray bands \\citep[][]{1969ApJ...155L..71K,1996ApJ...461..657B,2006A&A...451..797O}. In order to search for multi-frequency signatures of such short-term IC-flashes, 0716+714 was the target of a global multi-frequency campaign carried out in November 2003, and centered around a 500-ks INTEGRAL\\footnote{INTErnational Gamma-Ray Astrophysics Laboratory} pointing on November 10\\,--\\,17, 2003 (`core-campaign'). In order to detect contemporaneous IC-limit violations at lower energies, quasi-simultaneous and densely time-sampled flux density, polarization and VLBI monitoring observations were organized, covering the radio, millimeter, sub-millimeter, IR and optical\\footnote{The IR/optical data were collected in collaboration with the WEBT; http://www.to.astro.it/blazars/webt} wavelengths.  Early results of this campaign including radio data at two frequencies, optical and INTEGRAL soft $\\gamma$-ray data were presented by \\citet{2006A&A...451..797O}, who also showed the simultaneous spectral energy distribution. The results from the mm-observations (3\\,mm, 1.3\\,mm) performed with the IRAM 30\\,m radio-telescope were presented by \\citet{2006A&A...456..117A}. The results of the VLBI observations will be presented in a forthcoming paper. In this paper III we present the intensity and polarization data obtained with the Effelsberg 100\\,m radio telescope (5, 10.5, \\& 32\\,GHz) during the time of the core-campaign. We then combine these data with the (sub-) millimeter and optical data in order to determine the broad-band variability and spectral characteristics of 0716+714, with special regard to (i) the intra- to inter-day cm-/mm-variability, (ii) the brightness temperatures and IC-limit, and (iii) the Doppler factors combining the radio and high energy INTEGRAL data from this campaign.  \\begin{table*} \\begin{center} \\caption{ The participating radio observatories, their observing wavelengths/frequency, dates and total observing time $T_{Obs}$.} \\begin{tabular}{l|lllll} \\hline \\hline Radio telescope \\& Institute  & Location              & $\\lambda_{obs}$ [mm] & $\\nu_{obs}$ [GHz] & dates & $T_{obs}$ [hrs]\\\\ \\hline WSRT (14x25\\,m), ASTRON       & Westerbork, NL        & 210, 180           & 1.4, 2.3        & Nov. 10--11 & 8   \\\\ Effelsberg (100\\,m), MPIfR    & Effelsberg, D         & 60, 28, 9          & 4.85, 10.45, 32 & Nov. 11--18 & 164 \\\\ Pico Veleta (30\\,m), IRAM     & Granada, E            & 3.5, 1.3           & 86, 229         & Nov. 10--16 & 135 \\\\ Mets\\\"ahovi (14\\,m), MRO      & Mets\\\"ahovi, Finland  & 8                  & 37              & Nov. 08--19 & 263 \\\\ Kitt Peak (12\\,m), ARO        & Kitt Peak, AZ, USA    & 3                  & 90              & Nov. 14--18 & 91  \\\\ SMT/HHT  (10\\,m), ARO         & Mt. Graham, AZ, USA   & 0.87               & 345             & Nov. 14--17 & 77  \\\\ JCMT (15\\,m), JAC             & Mauna Kea, HI, USA    & 0.85, 0.45         & 345, 666        & Nov. 09--13 & 99 \\\\ \\hline \\hline \\end{tabular} \\label{observatories} \\end{center} \\end{table*} ",
        "conclusions": "We presented the results of a combined variability analysis from an intensive multi-frequency campaign on the BL Lac 0716+714 performed during seven days in November 2003. From seven participating radio observatories we obtain a frequency coverage of 1.4--666\\,GHz. Densely sampled IntraDay Variability (IDV) light curves obtained at wavelengths of 60, 28, 9, 3, and 1.3\\,mm allowed for the first time a detailed analysis of the source's intra- to inter-day variability behavior over the full radio- to short mm-band. In this observing campaign 0716+714 was found to be in a particular slow mode of variability, when compared to all previous IDV observations of this source. While in total intensity a component of faster variability was observed only at $\\lambda$\\,60\\,mm, the source's flux density in the cm- to mm-regime was dominated by a nearly monotonic increase on inter-day time scales and with variability amplitudes strongly increasing towards higher frequencies. Here, our CCF analysis confirms that the flux density variations are correlated across the observing bands, with variability at shorter wavelengths leading. This and the observed frequency dependence, which cannot be explained by a model for weak scintillation, strongly suggest that the observed inter-day variability has to be considered as source-intrinsic rather than being induced by ISS. Only at 60\\,mm wavelength a component of faster variability is seen and implies an unusually high apparent brightness temperature $T_{B}$. Hence, ISS might be present at this frequency.  The non-detection of the `classical', more rapid (type II) IDV behavior of 0716+714 in the cm-bands is most likely caused by opacity effects and can be related to the overall flaring activity of the source shortly before and during this campaign. Since episodic IDV behavior is also observed in other sources \\citep[e.g.][]{2002PASA...19...64F,2006MNRAS.369..449K}, it appears likely that the variability pattern in `classical' type\\,II IDV sources is strongly affected by the evolution of their intrinsic complex structure on time scales of weeks to months. The observed complicated variability patterns in total intensity and polarization, and their correlations, indicate the existence of a multi-component structure with individually varying and polarized sub-components of different size.  From daily averages, the spectral evolution of the highly inverted radio-to-sub-mm spectrum of 0716+714 could be studied. During the seven observing days the spectrum always peaked near $\\nu_{m}\\!\\sim\\!86$\\,GHz. Significant changes of the peak flux mainly occurred during the first 5 days with a continuous {\\bf rise} of the peak flux density from 4 to 5\\,Jy. Together with possibly small changes in the daily spectral slopes and the peak frequency $\\nu_{m}$, the observed variations follow the `canonical' behavior and indicate time-variable synchrotron self-absorption and adiabatic expansion of a shock or a flaring component as described by standard models \\citep[e.g.][]{1985ApJ...298..114M}.  The apparent brightness temperatures $T_{B}$ obtained from the inter-day variations exceed theoretical limits by several orders of magnitudes. Although $T_{B}$ decreases towards the mm-bands, the $10^{12}$\\,K IC-limit is always violated. Assuming relativistic boosting of the radiation, the source must always be strongly Doppler boosted. We obtain lower limits to the Doppler factor of the source using different methods, including (i) inverse Compton ($\\delta_{var,IC}$) and equipartition ($\\delta_{var,eq}$) estimates using the variability brightness temperatures, (ii) an estimate  $\\delta_{eq}$ using calculations of the synchrotron and equipartition magnetic field, and (iii) an inverse Compton Doppler factor $\\delta_{IC}$ using the data from the simultaneous INTEGRAL observations. These methods reveal robust and self-consistent lower limits to the Doppler factor with $\\delta_{var,IC}>$\\,5, $\\delta_{var,eq}>$\\,8 and $\\delta_{eq}>$\\,12. These limits are in good agreement with estimates based on recent kinematical VLBI studies of the source and the IC Doppler factor $\\delta_{IC}>$\\,14 obtained from the upper limits to the high energy emission in the 3--200\\,keV bands. The non-detection of the source in the soft $\\gamma$-ray bands implies the absence of a simultaneous strong IC catastrophe during the period of our IDV observations. Since a strong contribution of interstellar scintillation to the observed inter-day variability can be excluded, we conclude that relativistic Doppler boosting appears to naturally explain the observed apparent violation of the theoretical brightness temperature limits."
    },
    "0809/0809.2157_arXiv.txt": {
        "abstract": "{} {X-ray observations of gravitationally lensed quasars may allow us to probe the inner structure of the central engine of a quasar. Observations of Q2237+0305 (Einstein Cross) in X-rays may be used to constrain the inner structure of the X-ray emitting source.} {Here we analyze the XMM-$\\it{Newton}$ observation of the quasar in the gravitational lens system Q2237+0305 taken during 2002. Combined spectra of the four images of the quasar in this system were extracted and modelled with a power-law model. Statistical analysis was used to test the variability of the total flux.} {The total X-ray flux from all the images of this quadruple gravitational lens system is $6\\cdot10^{-13}$ erg/(cm$^2\\cdot$sec) in the range 0.2--10~keV, showing no significant X-ray spectral variability during almost 42~ks of the observation time. Fitting of the cleaned source spectrum yields a photon power-law index of $\\Gamma=1.82_{ - 0.08}^{ + 0.07}$. The X-ray lightcurves obtained after background subtraction are compatible with the hypothesis of a stationary flux from the source.} {}  \\keywords {gravitational lensing -- X-rays: general -- quasars: general} ",
        "introduction": "Since its discovery, the gravitational lens system (GLS) Q2237+0305 (Einstein Cross; (Huchra et al., \\cite{huchra}) has attracted much attention as a unique laboratory to study gravitational lensing effects. This GLS consists of a quadruply-imaged quasar and a lensing galaxy that is the nearest among all known gravitational lens systems (the quasar redshift is $z_Q=1.695$, while the lens redshift is $z_G=0.0395$). Due to its very convenient configuration, it is one of the best investigated GLSs. It has been continuously monitored by a number of groups from 1992--2005 (Rix et al., \\cite{rix}; Falco et al., \\cite{vla}; Oestensen et al., \\cite{not}; Blanton et al., \\cite{hubble}; Bliokh et al., \\cite{maidanak}; Nadeau et al., \\cite{monica}; Wozniak et al., \\cite{ogle}; Alcalde et al., \\cite{glitp}; Schmidt et al., \\cite{apo}). Significant microlensing-induced brightness peaks on lightcurves of quasar images were detected in this system (see, e.g., Wozniak et al., \\cite{ogle}; Alcalde et al. \\cite{glitp}; Moreau et al., \\cite{moreau}). The OGLE group is continuing observations of the Einstein Cross and the database $http://bulge.princeton.edu/\\sim ogle$ is being constantly renewed.  Observations of the Einstein Cross in X-rays are likely to provide important data that can be used to constrain the source inner structure. It is hoped that the detection of a variable X-ray flux from this system will lead to estimates of the relative time-delays $\\Delta\\tau$ between the images.  The X-ray emission of Q2237+0305 was first detected during {\\it ROSAT} observations in 1997 (Wambsganss et al., \\cite{rosat}). From the analysis of Wambsganss et al. (\\cite{rosat}), a 0.1-2.4 keV count rate of 0.006 count/sec was obtained and, assuming a $\\Gamma=1.5$ power-law model and Galactic absorption (hydrogen column density $N_{H}=5.5\\cdot 10^{20}$~cm$^{-2}$; Dickey \\& Lockman, \\cite{HI}), a flux of $2.2\\cdot 10^{-13}$ erg/(cm$^{2}\\cdot$ s) was derived. The {\\it ROSAT} observation of Q2237+0305 did not show any significant flux variability, and the spatial resolution of the {\\it ROSAT}/HRI detector is not sufficient to resolve the different quasar images.   After {\\it ROSAT}, GLS Q2237+0305 was observed several times with the Advanced CCD Imaging spectrometer onboard the $\\it{Chandra\\ X-ray\\ Observatory}$. Results from a 30~ks and a 10~ks  observation of Q2237+0305, carried out on September 2000 and on December 2001, respectively, were published by Dai et al. (\\cite{dai}). For the former of these $\\it{Chandra}$ observations, the X-ray flux was $4.6 \\cdot 10^{-13}$ erg/(cm$^{2}\\cdot$s) in the energy range 0.4--8.0~keV, and the lensed luminosity was $1.0 \\cdot 10^{46}$ erg/s; for the latter observation, the flux was $3.7 \\cdot 10^{-13}$ erg/(cm$^{2}\\cdot$ s) and the lensed luminosity $8.3 \\cdot 10^{45}$ erg/s (in the same energy range).   An important point concerns the variability of the X-ray signal from the source that carries useful information about the innermost source structure. In the case of Q2237+0305, the detection of X-ray variability is especially important because it may lead to a direct estimate of the gravitational time-delays that are a significant characteristic of lensed systems. The model prediction of $\\Delta\\tau$ between different Einstein Cross images is of the order of hours (Schmidt et al. \\cite{schmidt}), whereas the optical data may provide time-delay values that are accurate to only 1-2 days (Vakulik et al., \\cite{vakulik}). The only observational estimate of the shortest time-delay is obtained by cross-correlating the X-ray lightcurves from different images of Q2237+0305 (Dai et al., \\cite{dai}) of the $\\it{Chandra}$ observations: here all four X-ray images of this quasar have been resolved and sufficient variability has been reported (the difference between total fluxes during the first and the second $\\it{Chandra}$ observations is close to 20$\\%$ of the averaged value). The latter observation enabled Dai and collaborators to determine the time delay of $\\Delta\\tau_{BA}= 2.7_{-0.9}^{+0.5}$~hours between two of the four GLS images. The delays of the other images of Einstein Cross are still unknown.   The XMM-$\\it{Newton}$ receivers cannot separate different images of the Einstein Cross. Nevertheless, one may hope to obtain some useful information about the source from the autocorrelation function of the total X-ray flux from all the images. The applicability of this method depends on the presence of sufficient flux variability and may be complicated by the presence of possible microlensing.    In the present paper we analyze the XMM-$\\it{Newton}$ observation of the GLS Q2237+0305 taken in 2002. We describe the data reduction and present the parameters of the X-ray spectra obtained by EPIC cameras and RGS spectrometers onboard XMM-$\\it{Newton}$. Also, we present the results on variability in the range 0.2--10~keV. Then we discuss an estimate of the source inhomogeneity scale in view of high-magnification microlensing events in this GLS.     \\begin{figure} \\resizebox{\\hsize } {!} { \\includegraphics{Q2237fig1.ps} } \\caption{EPIC pn (upper left), MOS1 (upper right) and  MOS2 (below) images of the Q2237+0305. Black circles define areas collecting source signal, while white-bordered areas show the regions chosen for estimates of the background.} \\label{fig1} \\end{figure} ",
        "conclusions": "We analysed the  XMM-$\\it{Newton}$ data (May, 2002) on the total X-ray flux from all four images of the GLS Q2237+0305 ``Einstein Cross\". The flux value in the 0.2--10~keV range is $5.9\\cdot 10^{-13}$ erg/(cm$^2\\cdot$ sec), corresponding to a lensed luminosity of $1.25\\cdot 10^{46}$ erg/sec. In order to obtain the true quasar luminosity, this value must be corrected, taking into account the lens magnification. The most recent model estimate of the magnification is $16^{+5}_{-4}$ (Schmidt et al., \\cite{schmidt}). Fitting the X-ray spectra in the 0.2--10~keV energy range by an absorbed power-law yields the photon index $\\Gamma=1.88\\pm0.10$ and hydrogen column density in the lensing galaxy $N_H=(7\\pm$2)$\\cdot 10^{20}$~cm$^{-2}$. This is in satisfactory agreement with the results of Dai et al. (\\cite{dai}).    The XMM-$\\it{Newton}$ lightcurves obtained after the background subtraction  are compatible with the hypothesis of a constant flux from the source. An additional argument in favor of no variability is the absence of a significant correlation between the light-curves from different cameras. This does not completely rule out source variability: (i) it may be obscured by the Poisson variance; (ii) it may be smeared out in the total X-ray flux from all the images that have different time delays; (iii) there may be periods of different variability amplitudes. These periods may correspond to the lifetimes of X-ray emitting blobs orbiting the black hole and therefore we might expect variability periods of the order of the revolution time around the black hole. In any case, the net variability during the XMM-$\\it{Newton}$ observation cannot be large. As for concern (iii), we note that the {\\it ROSAT} (1997) observations  did not reveal significant variability. On the other hand, the variability was considerable during the $\\it{Chandra}$ (2000, 2001) observations (Dai et al., \\cite{dai}). This suggests the presence of slow variability changes with characteristic times much longer than the inner variability timescale $\\tau_v\\sim 1$~hour following from the results of Dai et al. (\\cite{dai}).   Light-travel arguments imply that the one hour time-scale variability corresponds to a size-scale of the X-ray emitting region responsible for this variation of $R_X\\sim\\tau_v c/(1+z_s)=4\\cdot10^{13}$ cm. An interesting question is whether $R_X$ represents an inhomogeneity in the accretion disk or the entire X-ray emitting region. This question might be answered, provided that the black hole mass in QSO 2237+0305 is known. Kochanek (\\cite{kochan}) estimated this mass, implying a radius of the innermost stable circular orbit of $3r_g\\approx (4\\div 22.5)\\cdot10^{14}$ cm  (see, e.g., Chandrasekhar, \\cite{chandrasekhar}) in the case of a non-rotating black hole (where $r_g=2GM/c^2$ is the gravitational radius). The inner radius of a real accretion disk must be even larger. However, the mass estimate of Kochanek (\\cite{kochan}) is model dependent and it is not sufficient to make a rigorous judgement concerning a comparison with the order-of-magnitude estimate of $R_X$.   One may ask how this inhomogeneous structure (if it exists) would appear in the presence of microlensing. The theory of high-amplification events in the case of a small continuous source is well known (see, e.g., Grieger et al., \\cite{grieger}; Shalyapin et al., \\cite{shalyapin}; Popovic et al., \\cite{popovic} and references in these papers). In the case of one emitting spot, the microlensing amplification scales as $\\sim R_X^{-1/2}$. Obviously, in the case of a large number of small inhomogeneities, the total amplification may be smeared out over the source. However, different regions of the time-dependent inhomogeneous structure will be differently amplified, resulting in an increase of the variability amplitude in microlensed image. Note that independent contributions of this microlensing effect may reduce the correlation between the lightcurves in different images. On the other hand, any prominent peak in the lightcurve without repetitions after the expected time delays can be considered as a probable sign of microlensing.   The idea of inhomogeneity of the quasar core in GLS has been repeatedly discussed from different viewpoints (see, e.g., Wyithe \\& Turner, \\cite{wyither1}; Wyithe \\& Loeb, \\cite{wyither2}; Schechter et al., \\cite{schechter}). Some evidence in favor of inhomogeneous structure may be found in optical observations of different GLSs, where the rapid flux variations are observed on the background of slower changes (Colley \\& Schild, \\cite{colley}; Paraficz et al., \\cite{paraficz}). Note that there is an alternative explanation for these variations that invoke a planetary-mass object in the lensing galaxy (Schild \\& Vakulik, \\cite{vakshild}, Colley \\& Schild, \\cite{colley}). It might be possible to distinguish between these two models by long-term optical and X-ray monitoring of microlensing events in AGN.   Observations of the resolved X-ray images with $\\it{Chandra}$ are more informative for the Q2237+0305 variability studies and determination of the time delays. Nevertheless, observations of the total X-ray flux with XMM-$\\it{Newton}$ may provide valuable complementary information on this issue. In the case of XMM-$\\it{Newton}$, sufficiently long observations are needed to obtain a good result for the autocorrelation function. To avoid uncorrelated contributions from different images to the autocorrelation function, the monitoring period must be larger than the differential time delays in this system. In the case of Q2237+0305, the available measurement interval of almost 42~ks is clearly not long enough. The model prediction (Schmidt et al., \\cite{schmidt}) of maximal delay for Q2237+0305 is 17~hours. Large delays ($>35$~ks) are also predicted by earlier models of Q2237+0305 (Rix et al., \\cite{rix}; Schneider et al., \\cite{schneidtau}). This is consistent with the analysis of the optical lightcurves of Q2237+0305 (Vakulik et al., \\cite{vakulik}), which yields an upper limit for the delays of about $\\sim 2$~days.    While waiting for new observations one may  ask whether in principle it is possible to determine the variable X-ray signal $x(t)$ from the source if the total flux from all the GLS images is available. Typically, this problem has an infinite number of solutions. However, if we know $x(t)$ over a sufficiently large initial interval, then  $x(t)$ may be recovered uniquely from the total flux lightcurve (Zhdanov et al., \\cite{zh}). This may be of interest if one needs to combine results of long GLS monitoring of the total flux (cf. XMM-$\\it{Newton}$) with shorter observations of separated different X-ray images (cf. $\\it{Chandra}$)."
    },
    "0809/0809.4292_arXiv.txt": {
        "abstract": "We investigate the effect of a finite source on the  planetary-lensing signals of high-magnification events.  From this, we find that the dependency of the finite-source effect on the caustic shape is weak and perturbations survive even when the source is substantially bigger than the caustic. Specifically, we find that perturbations with fractional magnification excess $\\geq 5\\%$ survive when the source star is roughly 4 times bigger than the caustic.  We also find characteristic features that commonly appear in the perturbation patterns of planetary lens systems affected by severe finite-source effect and thus can be used for the diagnosis of the existence of a companion. These features form in and around a circle with its center located at the caustic center and a radius corresponding to that of the source star. The light curve of an event where the source crosses these features will exhibit a distinctive signal that is characterized by short-duration perturbations of either positive or negative excess and a flat residual region between these short-duration perturbations. ",
        "introduction": "Since the first discovery in 2004, eight planets have been detected by using the microlensing method \\citep{bond04, udalski05, beaulieu06, gould06, gaudi08, bennett08, dong08}. The microlensing method is important in various aspects of exoplanet studies.  First, due to its high sensitivity to planets located in the outer region of planetary systems beyond the snow line, the microlensing method is complementary to other methods such as the radial velocity and the transit methods that are sensitive to planets orbiting close to their host stars.  In addition, the sensitivity of the microlensing method extends to very low-mass planets and Earth-mass planets can be detected from ground-based observations.  Furthermore, it is the only proposed method that can detect free-floating planets \\citep{bennett02, han06b} that are thought to be kicked out from planetary systems.  The detection rate of microlensing planets is rapidly increasing and at least five additional planet candidates were detected during the 2007 and 2008 observation seasons (A.\\ Gould 2008, private communication).   The microlensing signal of a planet is a perturbation to a smooth standard light curve of a primary-induced lensing event occurring on a background star \\citep{mao91, gould92}. The duration of the perturbation is short: several hours for an Earth-mass planet and several days even for a gas-giant planet.  Thus it is difficult to detect planetary signals from microlensing survey observations where stars are monitored on a roughly nightly basis.  Currently, the observational frequency required for planet detections is achieved by employing an early-warning system and follow-up observations, where the early-warning system (OGLE: Udalski et al.\\ 1994; MOA: Bond et al.\\ 2001) enables to issue alerts of ongoing events by analyzing data from survey observations real time and follow-up observations (PLANET: Kubas et al.\\ 2008; Micro-FUN: Dong et al.\\ 2006) are focused on these alerted events.  However, the limited number of telescopes available for follow-up observations restricts the number of events that can be followed up at any given time and thus priority is given to events which will maximize the planet detection probability.  Currently, the highest priority is given to high-magnification events \\citep{bond02, yoo04}.  These events have high intrinsic planet detection efficiency because the source trajectories always pass close to the perturbation region around the central caustic induced by the planet \\citep{griest98}.   Despite the high chance of perturbation, however, it is often thought that detecting low-mass planets through the channel of high-magnification events would be difficult. This thought is based on the fact that the central caustic induced by a low-mass planet is usually smaller than the source star and thus the planetary signal would be greatly washed out by severe finite-source effect \\citep{bennett96}. However, it might be that perturbations persist despite the finite-source effect and could still be detected thanks to the high photometry precision achieved by the enhanced brightness of the highly magnified source star.  In this paper, we test this possibility by investigating how the pattern of central planetary perturbations is affected by the finite-source effect.   The paper is organized as follows.  In \\S\\ 2, we briefly describe the physical properties of central caustics. In \\S\\ 3, we investigate the effect of a source size on the perturbation pattern.   For this, we construct maps of perturbation pattern for planetary systems with various caustic/source size ratios and caustic shapes. Based on these maps, we search for characteristic features that may be used to identify the existence of planets.  We summarize the results and conclude in \\S\\ 4. ",
        "conclusions": "We investigated the effect of a finite source on the central perturbation pattern. From this, we found that the dependency of the finite-source effect on the caustic shape is weak and perturbations survive even when the source is substantially bigger than the caustic.  Specifically, we found that perturbations with fractional magnification excess $\\geq 5\\%$ survive when the source star is roughly 4 times bigger than the caustic.  We also found characteristic features that commonly appear in the perturbation patterns of lens systems affected by severe finite-source effect and thus can be used for the diagnosis of the existence of a companion.  These features form in and around a circle with its center located at the caustic center and a radius corresponding to that of the source star.  The light curve of an event where the source crosses these features will exhibit a distinctive signal that is characterized by short-duration perturbations of either positive or negative excess and a flat residual region between these short-duration perturbations."
    },
    "0809/0809.1648_arXiv.txt": {
        "abstract": "We present observations of SCP 06F6, an unusual optical transient discovered during the \\emph{Hubble Space Telescope} Cluster Supernova Survey. The transient brightened over a period of $\\sim$100 days, reached a peak magnitude of $\\sim$21.0 in both $i_{775}$ and $z_{850}$, and then declined over a similar timescale. There is no host galaxy or progenitor star detected at the location of the transient to a $3\\sigma$ upper limit of $i_{775} \\ge 26.4$ and $z_{850} \\ge 26.1$, giving a corresponding lower limit on the flux increase of a factor of $\\sim$120. Multiple spectra show five broad absorption bands between $4100$~\\AA\\ and $6500$~\\AA\\ and a mostly featureless continuum longward of $6500$~\\AA. The shape of the lightcurve is inconsistent with microlensing. The transient's spectrum, in addition to being inconsistent with all known supernova types, is not matched to any spectrum in the Sloan Digital Sky Survey (SDSS) database. We suggest that the transient may be one of a new class. ",
        "introduction": "Supernova (SN) surveys are designed to detect the brightening of supernovae over timescales of days to weeks. They often cover large areas at high sensitivity. As a result, they are able to discover unusual and rare transients with similar timescales. For example, in 2006 the Lick Observatory Supernova Search (LOSS) discovered an optical transient in the galaxy M85 \\citep{kulkarni07a,rau07b,ofek08a} with a lightcurve plateau of $\\sim$80 days. It is suggested that the origin of this transient is a stellar merger and that an entire class of similar transients, \\emph{luminous red novae}, exists. Other recent discoveries of rare objects include a Type Ia SN with a super-Chandrasekhar mass progenitor \\citep{howell06a} from the Supernova Legacy Survey (SNLS) and SN 2005ap, the most luminous SN ever observed \\citep{quimby07a} from the Texas Supernova Search (TSS).  Here we report the observations of the optical transient SCP 06F6 discovered during the course of the 2005-2006 \\emph{Hubble Space Telescope} (\\emph{HST}) Cluster SN Survey (PI: Perlmutter; Dawson et al. 2008, in preparation). The discovery was originally reported in a June 2006 IAU circular \\citep{dawson06a}. Its lightcurve rise time of $\\sim$100 days is inconsistent with all known SN types, and its spectroscopic attributes are not readily matched to any known variable. We present photometry in \\S2 and spectroscopy in \\S3. In \\S4, we discuss constraints and summarize. ",
        "conclusions": "We have presented photometric and spectroscopic data of an unusual optical transient discovered during the \\emph{HST} Cluster SN Survey. Its key features are as follows: a rise time of $\\sim$100 days with a roughly symmetric lightcurve; small but statistically significant color variations across the lightcurve; no detected host galaxy or progenitor; broad spectral features in the blue, with a red continuum, and some evidence for spectral evolution. Below, we first discuss constraints on the distance to the source. Next we consider the possibility that the transient is the result of microlensing, finding this to be unlikely. Lastly, we search for similar objects in the Sloan Digital Sky Survey (SDSS) spectral database, finding no convincing matches.  Any detection of proper motion or parallax would be strong evidence of a galactic source. We tested for this by fitting the position of the transient in each of the six ACS detection epochs using a 2-d Gaussian (Fig.~\\ref{fig:propmotion}). In all epochs the position uncertainty is dominated by image coalignment errors caused by residual distortion, rather than statistical error in the fit. Note that this uncertainty in relative position between epochs is distinct from the uncertainty in the absolute position of the transient discussed in \\S2. For these images and this position, we estimate this error to be ~0.1 pixel ($0.''005$) in each coordinate. The most discrepant positions differ by approximately 0.25~pixels ($0.''0125$). As a whole, the positions are consistent with no proper motion or parallax and give little indication of either. The upper limit on proper motion is $62~\\mathrm{mas}~\\mathrm{yr}^{-1}$. The upper limit on parallax is $\\sim$25~mas, which gives a lower limit on distance of $\\sim$40~pc. This limit excludes virtually any possible solar system object as the source of the brightening.    \\begin{figure} \\begin{center} \\epsscale{1.0} \\plotone{f4.eps} \\end{center} \\caption{Position of transient in each of the six ACS detection epochs with respect to the overall best fit position, $\\alpha = 14^\\mathrm{h} 32^\\mathrm{m} 27^\\mathrm{s}.395$, $\\delta = +33^\\circ 32' 24''.83$ (J2000.0). The \\emph{top} and \\emph{bottom} panels are the position in Right Ascension ($\\alpha$) and Declination ($\\delta$), respectively. 5 mas = 0.1 pixels. \\label{fig:propmotion}} \\end{figure}   We can derive a more significant constraint on distance from reference image magnitude limits, assuming that the transient is an outburst and a progenitor star exists. The dimmest stars (aside from neutron stars) known to undergo outbursts are white dwarfs (WDs). WDs range in absolute magnitude from approximately 10~mag ($T_\\mathrm{eff}\\sim25000$~K) to approximately 16~mag ($T_\\mathrm{eff}\\sim3000$~K) and dimmer. If we assume the progenitor is a WD with absolute magnitude $i_{775} = 16$, the reference image $3\\sigma$ upper limit of $i_{775} > 26.4$ gives a distance $3\\sigma$ lower limit of $\\sim$1.2~kpc. If the progenitor is a source dimmer than $i_{775} = 16$ (e.g., a cooler WD or a neutron star), the constraints on distance are weaker. Because the source is at high galactic latitude ($b = 67.3^\\circ$), a limit of 1.2~kpc places the WD outside the thin disk of the galaxy, as the thin disk has a WD scale height of approximately 300~pc \\citep{boyle89a}. However, there is a significant population of relatively cool WDs residing in the thick disk and stellar halo. Galaxy models predict between tens \\citep{castellani02a} and hundreds \\citep{robin03a} of WDs far dimmer than our reference image detection limit in a single ACS field at this galactic latitude. Despite the large range (which reflects the uncertainty in the density of ancient WDs in the stellar halo) it is clear that the density is high enough to make a galactic WD a possible progenitor.  Although the symmetry of the lightcurve (Fig.~\\ref{fig:lightcurve}) suggests that the transient is a microlensing event, this interpretation is unlikely. The lightcurve is dramatically broader than the theoretical lightcurve for microlensing of a point source by a single lens \\citep{paczynski86a}. The typical lightcurve FWHM of high-magnification (peak magnification $\\ge$ 300) microlensing events is on the order of a few hours \\citep[e.g.,][]{abe04a,dong06a} whereas the transient's lightcurve FWHM is $\\sim$100 days with a peak magnification $3\\sigma$ lower limit of $\\sim$120. Also, the color evolves a small but significant amount over the lightcurve, particularly between epochs eight and nine. Some of these difficulties can be overcome if we assume the source is resolved; this can both change the shape of the lightcurve and allow for color variation as different source regions are differentially magnified. However, this typically results in a lower peak magnification. Finally, microlensing would still not explain the mysterious spectrum.  In an effort to identify objects with similar spectra, we cross-correlated the broad features of the spectrum with the SDSS spectral database. Each SDSS spectrum was warped with a polynomial function to best fit the Keck spectrum, based on a least squares fit. The value of the root mean square of the difference between the spectra was used to determine the correlation. An allowance for relative redshift was made, with the requirement that the spectra overlapped in the range of the strongest features (3500 to 6200~\\AA). No convincing matches were found. Changing the warping function between linear and quadratic and varying the wavelength range used in the fit altered which SDSS objects had the highest correlation, but did not result in a more convincing match. The SDSS objects with the highest correlation were broad absorption line quasars (BAL QSOs) at various redshifts and carbon (DQ) WDs. Although BAL QSOs do have similarly broad features, they don't exhibit the spacing or rounded profiles of those of the transient. Also, BAL QSOs typically include emission features. The DQ WDs most similar to the transient are known as DQp WDs. Like the transient, DQp WDs have broad, rounded absorption features between 4000 and 6000~\\AA\\ with a red continuum \\citep[see, e.g.,][]{hall08a}. However, the positions and spacing of the absorption features shortward of 5000~\\AA\\ differ greatly from those of the transient spectrum. In addition, DQp WDs show increased emission toward the UV, which is not seen in the transient.  The absence of similar spectra in the SDSS database implies that such variables are either very rare or typically fainter than the SDSS detection threshold, or both. If they are typically faint, this would seem to argue for an extragalactic origin, though a galactic origin is of course still possible. If this transient does indeed represent a new class of either galactic or extragalactic transients, such objects will be of great interest for future extensive surveys of the time-variable sky."
    },
    "0809/0809.3151_arXiv.txt": {
        "abstract": "{The Mn to Cr mass ratio in supernova ejecta has recently been proposed as a tracer of Type Ia SN progenitor metallicity. We review the advantages and problems of this observable quantity, and discuss them in the framework of two Galactic supernova remnants: the well known Tycho SNR and W49B, an older object that has been tentatively classified as Type Ia. The fluxes of the Mn and Cr K$\\alpha$ lines in the X-ray spectra of these SNRs observed by the \\textit{Suzaku} and \\textit{ASCA} satellites suggest progenitors of supersolar metallicity for both objects.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "In \\cite{badenes08:mntocr}, the $M_{Mn}/M_{Cr}$ ratio was introduced as a tracer of SN Ia progenitor metallicity. It might be more appropriate to say that it is a tracer of \\textit{neutron excess} in the explosive Si burning region of Type Ia SNe. This is interesting in its own right, because it also opens a window into the extent of the convective core of pre-explosion WDs. We hope that more and better observations will let us disentangle the contributions of metallicity and C simmering to the neutron excess in Type Ia SNe, and that this can help us to understand the lingering mystery of Type Ia progenitors."
    },
    "0809/0809.4637_arXiv.txt": {
        "abstract": "\\noindent We present a new method of incorporating radiative transfer into Smoothed Particle Hydrodynamics (SPH).  There have been many recent attempts at radiative transfer in SPH \\cite{Stam_2005_1,Stam_2005_2,Mayer,WB_1}, however these are becoming increasingly complex, with some methods requiring the photosphere to be mapped (which is often of non-trivial geometric shape), and extra conditions to be applied there (matching atmospheres as in \\cite{Mejia_4}, or specifying cooling at the photosphere as in \\cite{Mayer}).  The method of identifying the photosphere is usually a significant addition to the total simulation runtime, and often requires extra free parameters, the changing of which will affect the final results. Our method is not affected by such concerns, as the photosphere is constructed implicitly by the algorithm without the need for extra free parameters.  The algorithm used is a synergy of two current formalisms for radiative effects: a) the polytropic cooling formalism proposed by \\cite{Stam_2007}, and b) flux-limited diffusion, used by many authors to simulate radiation transport in the optically thick regime (e.g. \\cite{Mayer}).  We present several tests of this method:  \\begin{itemize} \\item The evolution of a \\(0.07\\, M_{\\odot}\\) protoplanetary disc around a \\(0.5\\,M_{\\odot}\\) star (\\cite{Mejia_1,Mejia_2,Mejia_3,Mejia_4}) \\item The collapse of a non-rotating \\(1\\,M_{\\odot}\\) molecular cloud (\\cite{Masunaga_1,Stam_2007}) \\item The thermal relaxation of temperature fluctuations in an static homogeneous sphere (\\cite{Masu_98,Spiegel,Stam_2007}) \\end{itemize} ",
        "introduction": "\\subsection{The Equation of State}  \\noindent SPH alone only evolves the density and internal energy of each particle: radiative transfer requires extra variables such as temperature and opacity.  Therefore, some kind of prescription is required to calculate all required variables using only the density and internal energy.  The prescription used in this work is split into two expressions: the \\emph{equation of state} which calculates temperatures and mean molecular weights, and the \\emph{opacity law} which calculates opacities.  The equation of state is the same as used by \\cite{Stam_2007}: it accounts for the vibrational and rotational energy states of hydrogen and helium, as well as their various dissociation and ionisation states.  \\subsection{Polytropic Cooling}  \\noindent This approximation uses an SPH particle's density \\(\\rho_i\\), temperature \\(T_i\\), and gravitational potential \\(\\psi_i\\) to estimate a mean optical depth for the particle \\cite{Stam_2007}.  This relies on calculating a mass averaged column density, and a mass averaged opacity.  This averaging process allows us to account for the particle's surroundings: hot particles may be surrounded by a cooler, denser environment, increasing the effective opacity.  The formalism guarantees optimum cooling at the photosphere - however, this formalism cannot model detailed energy exchange between particles.  \\subsection{Flux Limited Diffusion}  \\noindent This approximation handles the exchange of energy between pairs of neighbouring particles, using the diffusion approximation.  This approximation is only valid in the optically thick regime, so a flux limiter is used to allow the approximation to be expanded to less optically thick situations.  Energy flows along temperature gradients in accordance with the laws of thermodynamics.  However, energy exchange terms are \\emph{pairwise}: hence there is no energy loss from the system.  \\subsection{The Hybrid Method}  \\noindent Comparing the limitations of the above two methods, it should be clear that a union of these two procedures should be complementary: polytropic cooling handles the important energy loss from the system (which flux-limited diffusion cannot), and flux-limited diffusion handles the detailed exchange of heat between neighbouring fluid elements (which polytropic cooling cannot).  Indeed, polytropic cooling's inability to model the detailed exchange of heat between neighbouring fluid elements - and flux-limited diffusion's inability to model energy loss - allow the two methods to work together correctly, modelling all aspects of the system's energy budget without encroaching on each other. ",
        "conclusions": "\\noindent We have created a new radiative transfer algorithm by merging two currently existing methods.  By doing so, the algorithm can model in detail frequency-averaged radiative transfer.  Several tests of the algorithm against both analytic and numerical benchmarks have shown it to be successful in capturing the necessary physics.  \\begin{theacknowledgments}  \\noindent The authors would like to acknowledge D. Price for the use of SPLASH \\citep{Price}.  All simulations were performed using the supa64 machine at the University of St. Andrews.  DS and AW gratefully acknowledge the support of an STFC rolling grant (PP/E000967/1) and a Marie Curie Research Training Network (MRTN-CT2006-035890).  \\end{theacknowledgments}"
    },
    "0809/0809.4415_arXiv.txt": {
        "abstract": "{IGR~J18483$-$0311 is a high-mass X-ray binary recently discovered by \\textit{INTEGRAL}. Its periodic fast X-ray transient activity and its position in the Corbet diagram - although ambiguous - led to the conclusion that the source was a likely Be/X-ray binary (BeXB), even if a supergiant fast X-ray transient (SFXT) nature could not be excluded.} {We aimed at identifying the companion star of IGR~J18483$-$0311 to discriminate between the BeXB and the SFXT nature of the source.} {Optical and near-infrared photometry, as well as near-infrared spectroscopy of the companion star were performed to identify its spectral type. We also assembled and fitted its broad-band spectral energy distribution to derive its physical parameters.} {We show that the companion star of IGR~J18483$-$0311 is an early-B supergiant, likely a B0.5Ia, and that its distance is about 3-4~kpc.} {The early-B supergiant nature of its companion star, as well as its fast X-ray transient activity point towards an SFXT nature of IGR~J18483$-$0311. Nevertheless, the long duration and the periodicity of its outbursts, as well as its high level of quiescence, are consistent with IGR~J18483$-$0311 being an intermediate SFXT, in between classical supergiant X-ray binaries (SGXBs) characterised by small and circular orbits, and classical SFXTs with large and eccentric orbits.} ",
        "introduction": "High-mass X-ray binaries (HMXBs) are X-ray sources for which high-energy emission stems from accretion onto a compact object of material coming from a massive companion star. Until recently, the majority of known HMXBs were Be/X-ray binaries (BeXBs), and just a few were supergiant X-ray binaries (SGXBs). The launch of the \\textit{INTErnational Gamma-Ray Astrophysics Laboratory} \\citep[\\textit{INTEGRAL}, ][]{2003Winkler} in October 2002 completely changed the situation, as many more HMXBs whose companion stars are supergiants were discovered during the monitoring of the Galactic centre and the Galactic plane using the onboard IBIS/ISGRI instruments \\citep{2003Ubertini, 2003Lebrun}. Most of these sources are reported in \\citet{2007Bird} and \\citet{2007Bodaghee}, and their studies reveal two main features which were not present on previously known SGXBs: first, many of them exhibit a considerable intrinsic absorption, far above the interstellar level; second, some of these new sources revealed a transitory nature and occasionally exhibit a fast X-ray transient activity lasting a few hours. It then appears that the \\textit{INTEGRAL} supergiant HMXBs can be classified in two classes: one class of considerably obscured persistent sources \\citep[see e.g. ][]{2006Chaty} and another of transitory sources called supergiant fast X-ray transients \\citep[SFXTs,][]{2006Negueruelaa}. \\newline  IGR~J18483$-$0311 was discovered during observations performed with \\textit{INTEGRAL} in 2003 April 23-28 \\citep{2003Chernyakova}, and it was found to have average fluxes of about 10~mCrab and 5~mCrab in the 15$-$40~keV and 40$-$100~keV bands, respectively. \\citet{2004Molkov} also detected the source in 2003 March-May during a survey of the Sagittarius arm tangent with \\textit{INTEGRAL}, and gave average fluxes of about 4.3~mCrab and 3.9~mCrab in the 18$-$60~keV and the 60$-$120~keV bands, respectively. Moreover, analysing archival data from the \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}) All-Sky Monitor (ASM), \\citet{2006Levine} reported a 18.55~days periodicity of the lightcurve of IGR~J18483$-$0311, likely corresponding to its orbital period.  With archival data from observations performed with \\textit{INTEGRAL} from 2003 May to 2006 April, \\citet{2007Sguera} found five new outbursts from IGR~J18483$-$0311 whose activity lasted no more than a few days, and characterised by fast flares lasting a few hours. They fitted its 3$-$50~keV spectrum with an absorbed cut-off power law, and derived $N_{\\rm H}\\sim9\\times10^{22}$~cm$^{-{2}}$ (well above the galactic column density in the line of sight of about $1.6\\times10^{22}$~cm$^{-{2}}$), $\\Gamma\\sim1.4$, and $\\textrm{E}_\\textrm{c}\\sim22$~keV. They also showed that all the outbursts are well-fitted by an absorbed bremsstrahlung whose parameters are $N_{\\rm H}\\sim7.5\\times10^{22}$~cm$^{-{2}}$ and $\\textrm{kT}\\sim21.5$~keV. They reported a periodicity in the long term ISGRI light curve of about 18.52~days - confirming the result of \\citet{2006Levine} - and argued it is most likely an orbital period. They also derived a periodicity of about 21.05~s in the JEM-X light curve of the first outburst, which is likely a neutron star pulse period. Finally, from archival \\textit{Swift} observations, the authors obtained a very accurate position of IGR~J18483$-$0311, which allowed them to pinpoint its USNO~B1.0~0868$-$0478906 and 2MASS~J18481720$-$0310168 counterparts. They found that the magnitudes were typical of an absorbed massive late O/early B star which, along with the position of the source in the Corbet diagram as well as the periodicity of its high-energy activity, led them to conclude that IGR~J18483$-$0311 was likely a BeXB, without excluding the possibility of an SFXT nature.  \\citet{2008Chaty} assembled the broad-band spectral energy distribution (SED) of the companion star of IGR~J18483$-$0311 from 0.7~$\\mu$m to 8~$\\mu$m, with near-infrared (NIR) photometric observations performed at the ESO/NTT using the SofI instrument, as well as archival optical data from the USNO-A.2 catalogue, and mid-infrared (MIR) data from the GLIMPSE survey. They fitted it with a combination of two absorbed black bodies and showed that the best fit was consistent with the companion star being a heavily absorbed B star, confirming the HMXB nature of IGR~J18483$-$0311. At last, \\citet{2008Masetti} performed optical spectroscopy of the companion star of IGR~J18483$-$0311, and concluded that it was an O/B giant star because of the large equivalent width of the H$_{\\rm \\alpha}$ emission line. \\newline  In this paper, we report optical and NIR observations of IGR~J18483$-$0311 performed in service mode at the ESO/NTT through our Target of Opportunity program ID~078.D-0268 (P.I. S. Chaty). These observations aimed at constraining the nature of its companion star and of the binary system. In Sect. 2, we describe the optical/NIR photometric and spectroscopic observations, as well as their reduction. In Sect. 3, we present the results and show the broad-band SED of the companion star of IGR~J18483$-$0311. We finally discuss the outcomes in Sect. 4. ",
        "conclusions": "The identification of its companion star as a B0.5Ia supergiant, as well as its fast X-ray transient activity, are consistent with IGR~J18483$-$0311 being an SFXT. Nevertheless, as already pointed out in \\citet{2007Sguera}, the IGR~J18483$-$0311 X-ray behaviour appears unusual. Indeed, whereas outbursts last a few hours and $L_{\\rm max}$/$L_{\\rm min}\\sim10^4$ for typical SFXTs, the outbursts of IGR~J18483$-$0311 last a few days and its $L_{\\rm max}$/$L_{\\rm min}$ ratio is $\\sim\\,10^3$ (meaning that its quiescence is higher). Moreover, \\citet{2006Levine} and \\citet{2007Sguera} reported a 18.52~days orbital period, which is quite low compared to what is expected for classical SFXTs, with large and eccentric orbits.  Recently, several authors pointed out the importance of clumpy winds to explain the SFXT behaviour \\citep{2005Zand, 2007Leyder, 2007Walter, 2008Negueruela}. They argue that the flares are created by the interaction of the compact object with dense clumps (created in the stellar wind of the companion star), the frequency of the flares depending mainly on the geometry of the system. On the contrary, the quiescent emission would be due to the accretion onto the compact object of diluted inter-clumps medium, which would explain the very low quiescence exhibited by classical SFXTs. Moreover, by classifying twelve SFXTs in function of the duration and the frequency of their outbursts, as well as their $L_{\\rm max}$/$L_{\\rm min}$ ratio, \\citet{2007Walter} showed the existence of intermediate SFXTs, characterised by a lower ratio and longer flares. Finally, \\citet{2008Negueruela} proposed a general scheme to explain the emission of both SGXBs and SFXTs. The authors argued about the existence of a zone around the companion supergiant star (radius $\\le\\,2R_{\\rm \\ast}$), in which the clump density is very high, and another one (radius $\\ge\\,3R_{\\rm \\ast}$) in which it is low. In function of the orbital parameters of the compact object, the system could be a classical SGXB (small and circular orbit inside the zone of high clump density), a classical SFXT (large and eccentric orbit), or an intermediate SFXT (small orbits, circular or eccentric) with possible periodic outbursts \\citep[see also][for an alternative model]{2007Sidoli}.  Considering the longer duration of its flares as well as their 18.52~days periodicity, and its lower $L_{\\rm max}$/$L_{\\rm min}$ ratio, IGR~J18483$-$0311 is likely an intermediate SFXT. Moreover, to figure out - in the framework of the model proposed by \\citet{2008Negueruela} - whether its orbit is circular or elliptic, we can first use the Kepler's third law for a B0.5Ia supergiant mass and radius $M_{\\rm \\ast}=33{\\textrm{M}_{\\rm \\odot}}$ and $R_{\\rm \\ast}=33.8{\\textrm{R}_{\\rm \\odot}}$ \\citep{2008Searle}, and a neutron star mass of $1.4{\\textrm{M}_{\\rm \\odot}}$ to derive the semi-major axis and we find $a\\,\\sim\\,2.83R_{\\rm \\ast}$. Then, assuming that IGR~J18483$-$0311 only gets into activity when orbiting within the zone of high clump density, we find that an eccentricity $0.43 \\le e \\le 0.68$ is needed to explain the average three days bursting activity reported in \\citet{2007Sguera}. We then conclude that IGR~J18483$-$0311 is an intermediate SFXT with a small and eccentric orbit."
    },
    "0809/0809.1006_arXiv.txt": {
        "abstract": "We consider the evaporation of rotating micro black holes produced in highly energetic particle collisions, taking into account the polarization due to the coupling between the spin of the emitted particles and the angular momentum of the black hole. The effect of rotation shows up in the helicity dependent angular distribution significantly. By using this effect, there is a possibility to determine the axis of rotation for each black hole formed, suggesting a way to improve the statistics. Deviation from thermal spectrum is also a signature of rotation. This deviation is due to the fact that rapidly rotating holes have an effective temperature $T_{\\rm eff}$ significantly higher than the Hawking temperature $T_H$. The deformation of the spectral shape becomes evident only for very rapidly rotating cases. We show that, since the spectrum follows a blackbody profile with an effective temperature, it is difficult to determine both the number of extra-dimensions and the rotation parameter from the energy spectrum alone. We argue that the helicity dependent angular distribution may provide a way to resolve this degeneracy. We illustrate the above results for the case of fermions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1230_arXiv.txt": {
        "abstract": "{ \\grb\\   is the first  {gamma ray burst (GRB)}, since the time of EGRET, for which individual photons of energy above several tens of MeV have been detected with a pair-conversion tracker telescope. This burst was discovered with the Italian \\textit{AGILE} gamma-ray satellite. The GRB was localized with a cooperation by \\textit{AGILE} and  {the interplanetary network (IPN)}. The gamma-ray imager (GRID) estimate of the position, obtained before the SuperAGILE-IPN localization, is found to be consistent with the burst position. {The hard X-ray emission observed by SuperAGILE lasted about 7 s, while there is evidence that the emission above 30 MeV extends for a longer duration (at least ~13 s).} Similar behavior was seen in the past from a few other GRBs observed with EGRET. However, the latter measurements were affected, during the brightest phases, by instrumental dead time effects, resulting in only lower limits to the burst intensity. Thanks to the small dead time of the \\textit{AGILE}/GRID we could assess that in the case of \\grb\\ the gamma-ray to X--ray flux ratio changes significantly between the prompt and extended emission phase. ",
        "introduction": "Only a relatively small number of gamma-ray bursts (GRBs) have been detected so far at energies above the MeV region.  The largest sample of high-energy observations has been collected using the large calorimeter of the EGRET instrument on board the  {Compton Gamma Ray Observatory} (\\textit{CGRO}). This detector provided GRB light curves and spectra in the 1--200 MeV energy range, thus extending the spectral coverage of the brightest bursts discovered at lower energy with the  {Burst And Transient Source Experiment} \\citep[BATSE][]{kaneko2008}. Only for a handful of GRBs were high-energy photons detected also with the EGRET spark chamber, {sensitive in the energy range 30 MeV - 10 GeV} \\citep[see, e.g.][and references therein]{dingus2001}. These observations showed the existence of a few bursts with very hard spectra, with peak energy up to \\gsim170 MeV,  and possibly of GRBs with  high energy ``excesses'', {i.e. an emission above 100 MeV larger than the extrapolation of their X-rays spectra}. It is particularly important for theoretical models to understand whether the high-energy emission is a separate spectral component, possibly with a different time evolution from than of the prompt emission. This possibility is supported by the slow time decay of the high-energy flux in GRB 941017 \\cite{gonzalez2003}, as well as by the EGRET detection of delayed photons from GRB 940217 \\cite{hurley1994}, including one with an energy of 18 GeV after 1.3 hours. Two factors that affected the study of GRBs at high energy are intrinsic to gamma-ray imagers of the old generation, based on spark chamber trackers: the relatively small field of view, limiting the sample of observed GRBs,  and the significant instrumental dead time, reducing the number of detected photons in the brightest parts of the bursts, thus preventing a reliable measure of their peak flux and fluence.  Both limitations are overcome by new gamma-ray satellites using self-triggering trackers based on silicon microstrip technology, like \\textit{AGILE} \\cite{tavani}. The \\textit{AGILE} Gamma-Ray Imaging Detector (GRID), operating in the  30 MeV -- 50 GeV energy range, has a field of view as large as one fifth of the whole sky and a dead time of only $\\sim$200 $\\mu$s, which makes it particularly suited for the observation of GRBs. {For comparison, EGRET had a field of view of $\\sim$1 sr and a dead time of $\\sim$200 ms}. Furthermore, \\textit{AGILE} carries other detectors that allow the study of GRB emission in different energy ranges: SuperAGILE provides two one-dimensional images (along orthogonal directions), light curves and spectra in the 17--50 keV range \\cite{Feroci2007}. The MiniCalorimeter (MCAL) \\cite{Labanti2008}, besides being used as part of the GRID, can be used to autonomously detect and study GRBs in the 0.35--100 MeV range. Finally, GRB light curves in the hard X-ray band can be obtained from the GRID Anti Coincidence scintillator panels \\citep[ACS,][]{perotti}. The SuperAGILE capabilities in terms of rapid and accurate localizations of GRBs have been well demonstrated already with the discovery of its first GRB, 070724B \\citep{Del_Monte_et_al_2008} and continue with a rate of about 1 GRB per 1--2 months. Instead, MCAL detects approximately 1 GRB per week, as described in \\citep{Marisaldi2008}.  Here we present the observation of the first GRB for which a positive detection has been obtained by the GRID. Interestingly, the AGILE data give strong evidence that the GRB duration at energies above 30 MeV is at least double than that observed at lower energies, suggesting that different emission processes might be at work in the gamma and X-ray range. \\begin{figure}[ht!]% \\centering \\includegraphics[angle=90, width=9. cm]{IPN-SA_position.eps} \\caption{ Combined SuperAGILE and IPN localization of GRB 080514B. The star ($*$) marks the position of the afterglow found in X-rays by the Swift-XRT \\citep[see][]{Page08_GCN7723} and in optical by IAB80 telescope \\citep[see][]{GCN_7719, GCN_7720}, the solid lines indicate the SuperAGILE error box and the dashed lines represent the IPN triangulation based on SuperAGILE and Mars Odyssey data. } \\label{fig:SuperAGILE-IPN_position}  \\vspace{-0.4 cm}  \\end{figure} ",
        "conclusions": "We have reported evidence for high-energy ($>$25 MeV) emission from \\grb\\ , a burst belonging to the class of long duration GRBs, discovered thanks to the complement of detectors on board the \\textit{AGILE} satellite. The highest-energy photon consistent with the direction of the burst has an energy of about 300 MeV. Most of the detected GRB photons are at lower energy, in the range $\\sim$25-50 MeV, consistent with a power law spectrum with photon index $2.5$.  Previous detections of GRBs at these high energies were obtained more than 10 years ago with EGRET, before the discovery of GRB afterglows confirmed their cosmological origin. {X-ray and optical/NIR observations indicate that the afterglow properties of \\grb\\ are similar to those of other bursts \\cite{rossi2008}}.  The most striking feature of \\grb\\ concerns the fact that the arrival times of the high energy photons detected with the \\textit{AGILE}/GRID do not coincide with the brightest peaks seen at hard X-rays. Three photons are concentrated within 2 s at the beginning of the burst, while the next ones arrive only when the  X-ray emission has returned to a level consistent with the background (7 s after the beginning of the burst). It is important to note that the number of photons in the initial 2 s is not limited by instrumental dead time effects, as was the case for some of the EGRET GRBs for which only a flux lower limit could be measured in correspondence of the brightest peaks. This implies a rapid time evolution of the gamma-ray to X-ray flux ratio, although a quantitative assessment of this variability is hampered by the small statistics.  {Previous EGRET observations indicated a similar behavior for the bursts GRB 941017 \\cite{gonzalez2003} and GRB 940217 \\cite{hurley1994}, suggesting the possibility that the high-energy emission, at least in some cases, is not a simple extension of the main component, but originates from a different emission mechanism and/or region. Theoretical models predict that Inverse Compton (IC) plays a significant role in the high-energy emission from GRBs.% The ratio of IC ($\\sim$MeV-GeV) to synchrotron ($\\sim$keV) fluences depends on the energy of the relativistic electrons accelerated in the shocks and on the ratio between magnetic field and photon densities. Detailed time resolved spectroscopy covering the whole energy range, as now possible with further \\textit{AGILE} and \\textit{GLAST} observations (complemented by low energy data from one of the X and soft gamma-ray experiments currently on orbit) is required to disentangle the different components and obtain some constraints on the physical parameters of the sources.}"
    },
    "0809/0809.1189_arXiv.txt": {
        "abstract": "We define a volume limited sample of over 14,000 early-type galaxies (ETGs) selected from data release six of the Sloan Digital Sky Survey. The density of environment of each galaxy is robustly measured. By comparing narrow band spectral line indices with recent models of simple stellar populations (SSPs) we investigate trends in the star formation history as a function of galaxy mass (velocity dispersion), density of environment and galactic radius. We find that age, metallicity and $\\alpha$-enhancement all increase with galaxy mass and that field ETGs are younger than their cluster counterparts by $\\sim 2\\;\\rm Gyr$. We find negative radial metallicity gradients for all masses and environments, and positive radial age gradients for ETGs with velocity dispersion over $180\\;\\rm km\\,s^{-1}$. Our results are qualitatively consistent with a relatively simple picture for ETG evolution in which the low-mass halos accreted by a proto-ETG contained not only gas but also a stellar population. This fossil population is preferentially found at large radii in massive ETGs because the stellar accretions were dissipationless. We estimate that the typical, massive ETG should have been assembled at $z \\lesssim 3.5$. The process is similar in the cluster and the field but occurred earlier in dense environments. ",
        "introduction": "Observational determinations of the history of star formation in Early-Type Galaxies (hereafter ETGs) are of great importance because hierarchical models of galaxy formation make firm predictions for the relation between age, metallicity and $\\alpha$-enhancement as a function of mass. These scaling relations, plotted against the observational proxy for mass, the velocity dispersion (hereafter $\\sigma$), have been the focus of many recent studies of ETGs (Annibali et al. 2007, Bernardi et al., 2006, de la Rosa et al. 2007, Gallazzi et al. 2006, Jimenez et al. 2007, Kuntschner et al. 2001/2, Lucey et al. 2007, Mateus et al. 2007, Nelan et al. 2005, Proctor et al. 2004/8, S{\\'a}nchez-Bl{\\'a}zquez et al. 2006, Smith et al. 2007, Terlevich \\& Forbes 2002, Thomas et al. 2005).  These studies clearly show that the simple picture of early formation of low mass galaxies, which then merge to form more massive systems is incorrect. The oldest stellar populations are found in the most massive galaxies -- one aspect of so-called ``downsizing''. However, the observational constraint is sensitive to the epoch at which star formation ceased, not when it started, so that low mass ETGs can still be ``old'', as dynamically bound objects, but have a \\emph{mean} stellar age that is much younger. However, this is not sufficient to save the simple hierarchical picture because the $\\alpha$-enhancement is also seen to increase with mass - implying more rapid formation for massive objects. This is a prediction of monolithic collapse models. It is important to remember that both the star formation history and mass assembly history of ETGs determine their evolution.  The most common method of determining the age, metallicity and $\\alpha$-enhancement of ETGs is by comparison of the narrow band absorption line indices with simple stellar population (SSP) models. It is well-known that stellar population parameters based on these indices are also sensitive to minor episodes of recent star formation. Luminosity-weighted, SSP equivalent stellar population parameters, such as those discussed here, do not therefore, distinguish between a genuinely young galaxy and an old galaxy that has experienced a ``rejuvenation'' event.  Some recent studies, based on the colour-magnitude relation of ETGs using colours that are extremely sensitive to recent star formation, in fact paint a surprising picture of ETG evolution. Schawinski et al. (2007) have used GALEX ultraviolet imaging to show that 30\\% of massive ETGs show \\emph{ongoing} star formation and that this fraction is higher in low-density environments. A similar picture is given by mid-infrared Spitzer data. Both Clemens et al. (2008) and Bressan et al. (2006) find that $\\sim 30\\%$ of ETGs in the Coma and Virgo clusters have experienced some star formation in the recent past.  A recent study by Rogers et al. (2007) has combined SDSS spectra and GALEX data to conclude that ``weak episodes of recent star formation'' are a phenomenon more commonly associated with ETGs in the \\emph{cluster} environment, a result seemingly, but not necessarily, inconsistent with several studies that find older SSP-equivalent ages in denser environments.   Here we repeat the analysis carried out in Clemens et al. (2006, hereafter Paper~I), which used 3614 objects selected from data release 3 (DR3) of the SDSS. Applying the same selection criteria to DR6 we define a sample of 14353 ETGs, four times as many objects. ",
        "conclusions": "We find positive correlations between age, metallicity and $\\alpha$-enhancement and the velocity dispersion, $\\sigma$, in ETGs. Galaxies in dense environments are $\\sim 20\\%$ older than those in low density environments for all $\\sigma$ ($\\sim 2\\;\\rm Gyr$ for an age of $10\\;\\rm Gyr$). The trend with age flattens above $\\sigma \\sim 200\\;\\rm km\\,s^{-1}$, especially for galaxies in low density environments. We find a marginally significant trend towards higher metallicities in low density environments but the environment has no effect on the $\\alpha$-enhancement.   Apart from the marginal metallicity difference between field and cluster the results are very similar to those of Paper~I. There we concluded that an anti-hierarchical scenario, in which star formation lasts longer but with lower efficiency in lower mass objects (see Granato et al. 2004) was consistent with the data. Here the additional determination of SSP parameters as a function of galactic radius places additional constraints on the evolutionary scenario.  Massive ETGs ($\\sigma > 180\\;\\rm km\\,s^{-1}$) have positive radial age gradients, negative metallicity gradients and marginally significant negative $\\alpha$-enhancement gradients. When a massive halo becomes non-linear it accretes smaller halos which started to collapse at earlier times. The radial trends suggest that these halos do not contain only gas, but also pristine stars. The gaseous component falls dissipatively into the potential well of the massive (proto-)spheroid, fueling rapid star formation. The increase in mass increases the rate and efficiency of star formation, driving the main correlations with galaxy mass. The pristine stellar component of each sub-halo, however, being dissipationless, is deposited at a radius consistent with the angular momentum of the encounter. These stars, which are slightly older, more metal poor and have moderate $\\alpha$-enhancement will therefore be spread over larger radii. At early times, rapid gas-rich mergers lead to an almost monolithic formation, at later times mergers become increasingly ``dry''. Very similar scenarios have been proposed to explain both the bi- modal metallicity distribution of globular clusters and the greater radial scale length of metal-poor relative to metal-rich globular clusters in elliptical galaxies (C\\^ot\\'e, Marzke \\& West, 1998, Bekki et al. 2008). Our results imply that the metal-poor globular cluster population in ETGs should be older, less metal-rich and slightly less $\\alpha$-enhanced than the metal-rich clusters.   Because the age difference at larger radii is actually the luminosity weighted SSP equivalent age in a larger aperture (not an annulus), the value of 0.05 dex, $\\sim 1\\;\\rm Gyr$, is a lower limit to the real age difference at larger radii. This time difference limits the assembly redshift of massive ETGs simply due to the lack of time to accommodate the formation of stars in the lower mass halos. Our limit translates into an upper limit to the assembly redshift of massive ETGs of $z \\lesssim 3.5$; in which case the stars in low mass halos formed at $z\\sim 10$, for a standard cosmology ($H_0=70$, $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$). This also provides an estimate of the star formation rate in the assembled spheroid. If the final stellar mass were $10^{12}\\;\\rm M_{\\odot}$ then stars must have formed at a rate $\\sim 10^{12}\\;\\rm M_{\\odot}/1\\;\\rm Gyr \\sim 10^3\\;\\rm M_{\\odot}\\,yr^{-1}$.  We conclude by stressing the statistical nature of our results. Because a galaxy's velocity dispersion is a function of both the halo mass and virialization redshift, variations in these parameters may render small samples insensitive to the trends we find.  The catalogue on which this article is based can be found at, {\\tt www.mrao.cam.ac.uk/$\\sim$bn204/galevol/clemensetal08.html}."
    },
    "0809/0809.4765_arXiv.txt": {
        "abstract": "High resolution observations of solar filaments suggest the presence of groups of prominence threads, i.e. the fine-structures of prominences, which oscillate coherently (in phase). In addition, mass flows along threads have been often observed. Here, we investigate the effect of mass flows on the collective fast and slow nonadiabatic magnetoacoustic wave modes supported by systems of prominence threads. Prominence fine-structures are modeled as parallel, homogeneous and infinite cylinders embedded in a coronal environment. The magnetic field is uniform and parallel to the axis of threads. Configurations of identical and nonidentical threads are both explored.  We apply the $T$-matrix theory of acoustic scattering to obtain the oscillatory frequency and the eigenfunctions of linear magnetosonic disturbances. We find that the existence of wave modes with a collective dynamics, i.e. those that produce significant perturbations in all threads, is only possible when the Doppler-shifted individual frequencies of threads are very similar. This can be only achieved for very particular values of the plasma physical conditions and flow velocities within threads. ",
        "introduction": "Prominences/filaments are fascinating coronal magnetic structures, whose dynamics and properties are not well-understood yet. The long life of the so-called quiescent prominences (several weeks) suggests that the cool and dense prominence material is maintained against gravity and thermally shielded from the much hotter and much rarer solar corona by means of some not well-known processes. However, it is believed that the magnetic field must play a crucial role in both the support and isolation of prominences. High-resolution observations of solar filaments reveal that they are formed by a myriad of horizontal structures called threads \\citep[e.g.,][]{lin2005}, which have been observed in the spines and barbs of both active region and quiescent filaments \\citep{lin2008}. The width of these fine-structures is typically in the range 0.2~arcsec --  0.6~arcsec, which is close to the resolution of present-day telescopes, whereas their lengths are between 5~arcsec and 20~arcsec \\citep{lin2004}. Threads are assumed to be the basic substructures of filaments and to be aligned along magnetic field lines. From the point of view of theoretical modeling, prominence threads are interpreted as large coronal magnetic flux tubes, with the denser and cooler (prominence) region located at magnetic field dips that correspond to the observed threads. Although some theoretical works have attempted to model such structures \\citep[e.g.,][]{ballesterpriest, schmitt, rempel, heinzel06}, there are some concerns about their formation and stability that have not been resolved yet.  Small amplitude oscillations, propagating waves and mass flows are some phenomena usually observed in prominences and prominence threads \\citep[see some recent reviews by][]{oliverballester02, ballester, banerjee}. Periods of small-amplitude prominence oscillations cover a wide range from less than a minute to several hours, and they are usually attenuated in a few periods \\citep{molowny, terradasobs}. Focusing on prominence threads, some works have detected oscillations and waves in such fine-structures \\citep[e.g.,][]{yi1, yi2, lin2004, okamoto,lin2007}. In particular, \\citet{yi1} and \\citet{lin2007} suggested the presence of groups of near threads that moved in phase, which may be a signature of collective oscillations. On the other hand, mass flows along magnetic field lines have been also detected \\citep{zirker94, zirker, lin2003, lin2005}, with typical flow velocities of less than 30~km~s$^{-1}$ in quiescent prominences, although larger values have been detected in active region prominences \\citep{okamoto}. Regarding the presence of flows, a phenomenon which deserves special attention is the existence of the so-called counter-streaming flows, i.e., opposite flows within adjacent threads \\citep{zirker, lin2003}.  Motivated by the observational evidence, some authors have broached the theoretical investigation of prominence thread oscillations by means of the magnetohydrodynamic (MHD) theory in the $\\beta = 0$ approximation. First, some works \\citep{joarder,diaz2001,diaz2003} focused on the study of the ideal MHD oscillatory modes supported by individual nonuniform threads in Cartesian geometry. Later, \\citet{diaz2002} considered a more representative cylindrical thread and obtained more realistic results with respect to the spatial structure of perturbations and the behavior of trapped modes. Subsequently, the attention of authors turned to the study of collective oscillations of groups of threads, and the Cartesian geometry was adopted again for simplicity. Hence, \\citet{diaz2005} investigated the collective fast modes of systems of nonidentical threads and found that the only nonleaky mode corresponds to that in which all threads oscillate in spatial phase. Later, \\citet{diazroberts} considered the limit of a periodic array of threads and obtained a similar conclusion. Therefore, these results seem to indicate that all threads within the prominence should oscillate coherently, even if they have different physical properties. However, one must bear in mind that the Cartesian geometry provides quite an unrealistic confinement of perturbations, and so systems of more realistic cylindrical threads might not show such a clear collective behavior.  The next obvious step is therefore the investigation of oscillatory modes of systems of cylindrical threads. The first approach to a similar problem was done by \\cite{luna1}, who considered a system of two identical, homogeneous cylinders embedded in an unlimited corona. Although \\cite{luna1} applied their results to coronal loops, they are also applicable to prominence threads. These authors numerically found that the system supports four trapped kink-like collective modes. These results have been analytically re-obtained by \\citet{tom}, by considering the thin tube approximation and bicylindrical coordinates. Subsequently, \\citet{luna2} made use of the $T$-matrix theory of acoustic scattering to study the collective oscillations of arbitrary systems of non-identical cylinders. Although the scattering theory has been previously applied in the solar context \\citep[e.g.,][]{bogdan87,keppens}, the first application to the study of normal modes of magnetic coronal structures has been performed by \\citet{luna2}. They concluded that, contrary to the Cartesian case of \\citet{diaz2005}, the collective behavior of the oscillations diminishes when cylinders with nonidentical densities are considered, the oscillatory modes behaving in practice like individual modes if cylinders with mildly different densities are assumed.  The present study is based on \\citet{luna2} and applies their technique to the investigation of MHD waves in systems of cylindrical prominence threads. Moreover, we extend their model by considering some effects neglected by them. Here, the more general $\\beta \\neq 0$ case is considered, allowing us to describe both slow and fast magnetoacoustic modes. In addition, the adiabatic assumption is removed and, following previous papers \\citep{soler1,solerapj}, the effect of radiative losses, thermal conduction and plasma heating is taken into account. The detection of mass flows in prominences has motivated us to include this effect in our study, and so the presence of flows along magnetic field lines is also considered here. Therefore, the present work extends our recent investigation \\citep[][hereafter Paper~I]{solerapj}, which was focused on individual thread oscillations, to the study of collective MHD modes in prominence multi-thread configurations with mass flows. On the other hand, the longitudinal structure of threads \\citep[e.g.,][]{diaz2002} is neglected in the present investigation. For this reason, the effect of including a longitudinal variation of the plasma physical conditions within threads should be investigated in a future work. Finally, the prominence multi-thread model developed here could be an useful tool for future seismological applications \\citep[similar to that of][]{hinode}.  This paper is organized as follows. The description of the model configuration and the mathematical method are given in \\S~\\ref{sec:math}. Then, the results are presented in \\S~\\ref{sec:results}. First, the case of two identical prominence threads is investigated in \\S~\\ref{sec:2treads}. Later, this study is extended to a configuration of two different threads in \\S~\\ref{sec:ntreads}. Finally, our conclusion is given in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we have assessed the effect of mass flows on the collective behavior of slow and fast kink magnetosonic wave modes in systems of prominence threads. We have seen that the relation between the individual Alfv\\'en (sound) speed of threads is the relevant parameter which determines whether the behavior of kink (slow) modes is collective or individual. In the absence of flows and when the Alfv\\'en speeds of threads are similar, kink modes are of collective type. On the contrary, perturbations are confined within an individual thread if the Alfv\\'en speeds differ. In the case of slow modes, the conclusion is equivalent but replacing the Alfv\\'en speeds by the sound speeds of threads. On the other hand, when flows are present in the equilibrium, one can find again collective motions even in systems of non-identical threads by considering appropriate flow velocities. These velocities are within the observed values if threads with not too different temperatures and densities are assumed. However, since the flow velocities required for collective oscillations must take very particular values, such a special situation may rarely occur in real prominences.  Therefore, if coherent oscillations of groups of threads are observed in prominences  \\citep[e.g.,][]{lin2007}, we conclude that either the physical properties and flow velocities of all oscillating threads are quite similar or, if they have different properties, the flow velocities within threads are the appropriate ones to allow collective motions. From our point of view, the first option is the most probable one since the flow velocities required in the second case correspond to a very peculiar situation. This conclusion has important repercussions for future prominence seismological applications, in the sense that if collective oscillations are observed in large areas of a prominence, threads in such regions should possess very similar temperatures, densities, and magnetic field strengths  Here, we have only considered two-thread systems, but the method can be applied to an arbitrary multi-thread configuration. So, the model developed here could be used to perform seismological studies of large ensembles of prominence threads if future observations provide with positions and physical parameters of such systems."
    },
    "0809/0809.0693_arXiv.txt": {
        "abstract": "We present the first systematic derivation of the one-loop correction to the large scale matter power spectrum in a mixed cold+hot dark matter cosmology with subdominant massive neutrino hot dark matter.  Starting with the equations of motion for the density and velocity fields, we derive perturbative solutions to these quantities and construct recursion relations for the interaction kernels, noting and justifying all approximations along the way. We find interaction kernels similar to those for a cold dark matter-only universe, but with additional dependences on the neutrino energy density fraction $f_\\nu$ and the linear growth functions of the incoming wavevectors.  Compared with the $f_\\nu=0$ case, the one-loop corrected matter power spectrum for a mixed dark matter cosmology exhibits a decrease in small scale power exceeding the canonical $\\sim 8 f_\\nu$ suppression predicted by linear theory, a feature also seen in multi-component $N$-body simulations. ",
        "introduction": "Standard big bang theory predicts a background of relic neutrinos, permeating the universe at an average of 112 neutrinos per cubic cm per neutrino flavour. This enormous abundance means that even a sub-eV to eV neutrino mass $m_\\nu$  will render these otherwise elusive particles a significant dark matter component, $\\Omega_\\nu h^2 = \\sum m_\\nu/(93 \\ {\\rm eV})$. Experimentally, a lower limit of $\\sum m_\\nu \\gwig 0.05 \\ {\\rm eV}$ on the sum of the neutrino masses has been established by neutrino oscillation experiments~(e.g., \\cite{Amsler:2008zz}). On the other hand, tritium $\\beta$-decay end-point spectrum measurements point to an upper limit of $\\sum m_\\nu \\lwig 6 \\ {\\rm eV}$~(e.g., \\cite{Amsler:2008zz}). The corresponding energy densities, $0.001 \\lwig \\Omega_\\nu \\lwig 0.12$, make for an interesting prediction to be tested against cosmological observations.  Importantly, neutrino dark matter is of the hot variety because of their inherent thermal velocity, which prevents them from clustering gravitationally on scales smaller than their free-streaming length. This free-streaming effect then feedbacks into the evolution of the dominant cold dark matter (CDM) component via the gravitational source term, leading to a suppressed rate in the formation of structures  on small length scales that is in principle manifest in the power spectrum of the large scale structure distribution~\\cite{Bond:1980ha,Doroshkevich:1980zs,bib:khlopov,Shafi:1984ek,Schaefer:1989ua}. At present, observations of the cosmic microwave background anisotropies, galaxy clustering, and type Ia supernovae together engender a conservative upper limit on the contribution of neutrino dark matter, or equivalently, on the neutrino masses, of $\\sum m_\\nu \\lwig 1 \\ {\\rm eV}$ within the $\\Lambda$CDM framework.  See~\\cite{Lesgourgues:2006nd,Hannestad:2006zg} for recent reviews.  Many future cosmological probes will continue to improve on this limit, or perhaps even detect neutrino dark matter (e.g., \\cite{Lesgourgues:2006nd,Abdalla:2007ut,% Hannestad:2007cp,Hannestad:2006as,% Kitching:2008dp,Lesgourgues:2005yv,Gratton:2007tb,Ichikawa:2005hi,% Lesgourgues:2007ix,Pritchard:2008wy}). To this end, it is important to note that many of these probes, particularly high redshift galaxy surveys and weak gravitational lensing, will derive most of their constraining power at wavenumbers $k \\sim 0.1 \\to 1 \\ h \\ {\\rm Mpc}^{-1}$, where the evolution of density perturbations has become weakly nonlinear. Coincidentally, these are also the scales at which neutrino free-streaming is expected to produce the largest effect on the large scale matter power spectrum. Thus, it is crucial that we understand the nonlinear evolution of density perturbations at these scales, in order to maximise our gain in detecting/constraining neutrino dark matter.  A number of recent works have attempted to model the effects of massive neutrinos on the matter power spectrum at nonlinear scales using various different techniques. These include multi-component $N$-body simulations~\\cite{Brandbyge:2008rv} and semi-analytic halo models~\\cite{Abazajian:2004zh,Hannestad:2005bt,Hannestad:2006as}. However, with the exception of $N$-body simulations, these methods all contain elements that require prior calibration against simulation results, and not surprisingly, all calibrations to date have been performed in a $\\Lambda$CDM setting. This renders these nonlinear models at present not completely satisfactory for use with cosmologies containing massive neutrinos. Recalibration against simulations in the appropriate cosmology will, however, enhance/restore our confidence in these methods.  Recently, Saito {\\it et al.}~\\cite{Saito:2008bp} proposed an alternative method based on higher order cosmological perturbation theory.  Since its inception in the 1980s~\\cite{bib:juszkie,bib:vishniac,Fry:1983cj,Goroff:1986ep}, higher order cosmological perturbation theory has found numerous applications ranging from computation of the weakly nonlinear power spectrum and gravity-induced bispectrum, to the exploration of  nonlinear galaxy bias. See reference~\\cite{Bernardeau:2001qr} for a review. In the perturbative  approach, one envisages nonlinear evolution as the outcome of interactions of a collection of linear waves (perturbations). The equations of motion of the system define the interaction kernels.   Saito {\\it et al.}'s recipe is simple:  assume the neutrino density perturbations  remain linear at all times, and apply nonlinear modelling only to the CDM+baryon component.  This basic scheme has also been adopted in some earlier nonlinear models including massive neutrinos~\\cite{Hannestad:2006as}. For $\\sum m_\\nu \\lwig 0.6 \\ {\\rm eV}$, multi-component $N$-body simulations have confirmed that a linear approximation for the neutrino density contrast is sound~\\cite{Brandbyge:2008rv}.  For the CDM+baryon component, Saito {\\it et al.}~calculated the one-loop correction to the CDM+baryon auto-correlation power spectrum, using the correctly computed linear waves (perturbations), but interaction kernels that have been developed for a CDM-only universe. This amounts to ignoring additional mode-coupling effects between the linear growth functions at different wavenumbers.  Although, as we shall show, this approximation does lead to a considerable simplification in the final form of the nonlinear  power spectrum, it is nonetheless reminiscent of the mismatch between the cosmology to which we apply and that on which we calibrate some of the semi-analytic methods discussed above, and therefore calls for closer scrutiny.  In this paper we present a systematic derivation of the one-loop correction to the matter power spectrum in the presence of massive neutrinos from first principles. As in~\\cite{Hannestad:2006as,Saito:2008bp}, we assume linearity for the neutrino component. For CDM+baryons, however, we begin with the relevant equations of motion, and solve them (approximately) perturbatively in an Einstein--de Sitter $\\Omega_m=1$ background at high redshifts ($z \\gwig 0.5$). From these solutions we derive interaction kernels that capture also the physics of mode-coupling between linear growth functions at different wavenumbers.  With these kernels we compute the correct nonlinear power spectrum.  In section~\\ref{sec:eom} we give the relevant equations of motion.  We discuss  briefly the linear order solutions in section~\\ref{sec:linear}, before formulating our higher order perturbative approach in section~\\ref{sec:beyond}.  In section~\\ref{sec:approx} we outline our scheme to obtain approximate solutions to the equations of motion for higher order perturbations, while in section~\\ref{sec:nthorder} we derive recursion relations for the interaction kernels and thus generalise our approximate solutions to arbitrary orders in perturbative expansion. Section~\\ref{sec:spectra} deals with the calculation of the one-loop correction to the matter power spectrum, which we evaluate numerically for realistic cosmological models and discuss in detail  in section~\\ref{sec:discuss}. We conclude in section~\\ref{sec:conclusions}. ",
        "conclusions": "}  In this paper we have presented the first rigorous and systematic derivation of the one-loop correction to the large scale matter power spectrum in a mixed dark matter cosmology with subdominant massive neutrino hot dark matter.   Beginning with the relevant equations of motion, we find that by invoking an ``adiabatic'' approximation, accurate to better than $\\sim 0.1 f_\\nu$, higher order corrections to the CDM+baryon density contrast and velocity field can be rendered into a form nearly identical to that for a pure-CDM cosmology. The interaction kernels and their recursion relations also exhibit striking similarities to their standard CDM-only counterparts, but contain additional dependences on the neutrino energy density fraction $f_\\nu$ and the linear growth functions of the incoming wavevectors.  These results, generalised to $n$th order in perturbative expansion, are summarised in equations~(\\ref{eq:solution}) to (\\ref{eq:sigmastuff}).  Using these approximate solutions we compute the usual ``22'' and ``13'' one-loop correction terms to the matter power spectrum. As in the standard CDM-only case, these correction terms take the form of integrals over the wavevector ${\\bm q}$ of the linear power spectrum $P^L(q,\\tau)$ multiplied by the interaction kernels. In addition to the corrections to the CDM+baryon auto-correlation, we also find a one-loop correction term for the cross-correlation between the CDM+baryon and the neutrino components which was previously neglected. These correction terms appear in their evaluated and most simplified form in equations~(\\ref{eq:22explicit}) to (\\ref{eq:13explicit}), (\\ref{eq:xintegral}) and (\\ref{eq:t}).  Evaluating these expressions numerically, we find that nonlinear corrections to the large scale matter power spectrum can enhance the suppression  of small scale power due to neutrino free-streaming relative to the $f_\\nu=0$ case to beyond the canonical linear suppression factor of $\\sim 8 f_\\nu$. This enhanced suppression has been observed in multi-component $N$-body simulations~\\cite{Brandbyge:2008rv}.   As said, the interaction kernels contain hitherto unaccounted dependences on $f_\\nu$ and the linear growth functions.  Neglecting these dependences in principle generates large deviations in the one-loop corrections.  However, since these deviations occur at wavenumbers at which the linear contribution dominates over the correction terms, their net effect on the total matter power spectrum never exceeds 1\\%. Future cosmological probes will require an accuracy of $\\sim 1$\\% in the matter power spectrum in order not to bias parameter estimation. We have thus verified the validity of the approach of~\\cite{Saito:2008bp}.   An important assumption in our present treatment is that the neutrino density perturbations have been taken to remain linear at all times.  Although for realistic values of $f_\\nu$ this can be justified by $N$-body simulation results~\\cite{Brandbyge:2008rv}, a truly complete analysis of higher order corrections to the clustering statistics of the large scale structure distribution in the presence of massive neutrinos should include also a proper account of nonlinear neutrino evolution. We defer this investigation to a future publication.  Finally, as noted in reference~\\cite{Saito:2008bp}, although higher order perturbation theory appears at first glance to have a limited range of validity---we expect our one-loop corrections to improve on linear theory to better than 1\\% accuracy in the region $0.1 \\lwig k/(h \\ {\\rm Mpc}^{-1}) \\lwig 0.4$ at $z=3$, it does enable an approximate factor of four increase in the maximum usable wavenumber in a data set. This is equivalent to a factor of 64 gain in the number of independent Fourier modes, or an eight-fold gain in statistical power for a fixed survey volume. Such an improvement is no small feat, and may very well be just what we need to detect neutrino dark matter.   \\ack  ${\\rm Y}^3$W thanks Jacob Brandbyge, Steen Hannestad and Georg Raffelt for useful discussions and/or comments on the manuscript."
    },
    "0809/0809.0236_arXiv.txt": {
        "abstract": "We present observational constraints on the nature of the different core-collapse supernova types through an investigation of the association of their explosion sites with recent star formation, as traced by \\ha +\\NII\\ line emission. We discuss results on the analysed data of the positions of 168 core-collapse supernovae with respect to the \\ha\\ emission within their host galaxies.\\\\ From our analysis we find that overall the type II progenitor population does not trace the underlying star formation. Our results are consistent with a significant fraction of SNII arising from progenitor stars of less than 10\\msun . We find that the supernovae of type Ib show a higher degree of association with HII regions than those of type II (without accurately tracing the emission), while the type Ic population accurately traces the \\ha\\ emission. This implies that the main core-collapse supernova types form a sequence of increasing progenitor mass, from the type II, to Ib and finally Ic. We find that the type IIn sub-class display a similar degree of association with the line emission to the overall SNII population, implying that at least the majority of these SNe do not arise from the most massive stars. We also find that the small number of SN `impostors' within our sample do not trace the star formation of their host galaxies, a result that would not be expected if these events arise from massive Luminous Blue Variable star progenitors. ",
        "introduction": "\\label{intro} Despite years of observational and theoretical research on the nature of supernova (SN) explosions and the properties of their progenitors there remain substantial gaps in our knowledge of all SN types. Although there are many different theoretical predictions as to the nature of SN progenitors, the observational evidence to discriminate between various progenitor scenarios remains sparse. SNe can be split into two theoretical classes; SNIa which are thought to arise from the thermonuclear explosion of an accreting white dwarf, and core-collapse (CC) SNe which are believed to signal the collapse of the cores of massive stars at the end points in their stellar evolution. \\\\ Results from the first paper in this series (\\citealt{jam06}, JA06 henceforth) suggested that the SNIb/c arise from higher mass progenitors than SNII (albeit with small statistics; only 8 SNIb/c). We test these initial results with an increased sample size enabling us to distinguish between the various CC sub-types, and present results from a combined sample of 100 SNII (that can be further separated into 37 IIP, 8 IIL, 4 IIb, 12 IIn), 62 SNIb/c (22 Ib, 34 Ic and 6 that only have Ib/c classification), and 6 SN `impostors'. We will present results and discussion of SNIa within the context of the methods used in this paper elsewhere. We will also present further research on the radial positions of SNe within galaxies, and on correlations between CC SN type and local metallicity in future publications. Here we concentrate on results on the progenitor masses of the different CC SNe.  \\subsection{Core-collapse supernovae} \\label{cc} CC SNe are thought to be the final stage in the stellar evolution of stars with initial masses $>$8-10\\msun , when fusion ceases in the cores of their progenitors and they can no longer support themselves against gravitational collapse. The different types of CC SNe are classified according to the presence/absence of spectral lines in their early time spectra, plus the shape of their light curves. The first major classification comes from the presence of strong hydrogen (H) emission in the SNII. SNIb and Ic lack any detectable H emission, while the SNIc also lack the helium lines seen in SNIb. SNII can also be separated in various sub-types. SNIIP and IIL are classified in terms of the decline shape of their light curves (\\citealt{bar79}; plateau in the former and linear in the latter), thought to indicate different masses of their H envelopes prior to SN, while SNIIn show narrow emission lines within their spectra \\citep{sch90}, thought to arise from interaction of the SN ejecta with a slow-moving circumstellar medium (e.g. \\citealt{chug94}). SNIIb are thought to be intermediate objects between the SNII and Ib as at early times their spectra are similar to SNII (prominent H lines), while at later times they appear similar to SNIb (\\citealt{fil93}).\\\\ Strong evidence has been presented to support the belief that SNII and SNIb/c arise from massive progenitors, through their absence in early type galaxies \\citep{bergh02}, and the direct detection of a small sample of progenitors on pre-explosion images (\\citealt{sma04,maun05,hen06,li06,gal07,li07,crock08}). However, it is unclear how differences in the nature of their progenitors produce the different SNe we see. It is clear that there must be some process by which the progenitors of the different SNe lose part (or almost all in the case of SNIb and Ic) of their envelopes prior to explosion. The differences in efficiency of this mass loss process could be dependent primarily on progenitor mass, with higher mass progenitors having higher mass loss rates due to stronger stellar winds, and losing more of their envelopes. In this picture a sequence of SNe types emerges from SNIIP and IIL to SNIIb, SNIb and finally Ic having successively higher initial masses. There are also other factors that probably play an important role. Initial chemical abundance will also affect the progenitor mass loss, with higher metallicity producing stronger radiatively driven winds ({e.g. \\citealt{pul96,kud00,mok07}). It has also been proposed \\citep{pod92} that massive binaries could produce a significant fraction of CC SNe, with mass transfer ejecting matter and leading to some of the various CC sub-types.\\\\ Since the theoretical separation of SNe into two distinct explosion classes by \\cite{hoy60}, there have been many predictions as to how the different CC SN types emerge from different progenitors. There are two main theoretical routes to achieving the observed different SN types. The first attempts to describe the full range of SNe from a single star progenitor scenario. \\cite{heg03} and \\cite{eld04} produced SN progenitor maps showing how variations in initial mass and metallicity produced the different CC SN types. These models both predict that single stars of up to $\\sim$25-30\\msun\\ will produce SNIIP, with stars of slightly higher mass producing SNIIL and IIb, and those of $>$30\\msun\\ ending their lives as SNIb/c (both authors also predict that these initial mass ranges will shift to higher values with decreasing metallicity). In both of these models no attempt was made to differentiate between the SNIb and the SNIc, but one would presume that within this single star scenario SNIc would arise from higher mass progenitors than the SNIb as they have lost even more of their stellar envelopes. Alternatively it could be that massive binaries produce the majority of CC SNe other than SNIIP (with these SNe still arising from single star progenitors). The initial mass of the stars producing SNIb/c, SNIIL and SNIIb would then be similar to those of SNIIP (12-20\\msun , e.g. \\citealt{shi90}) but would arise from binary evolutionary processes. There is also a growing number of SNe that show evidence of binarity (e.g. SN 1987A; \\citealt{pod90} and SN 1993J; \\citealt{nom93,pod93,maun04}). Recent comparisons of the observed ratio of SNIb/c rates to those of SNII also argue that binaries are playing the dominant role in producing SNIb/c \\citep{kob07}, while \\cite{eld08} predict a SNIb/c rate produced by a combination of single and binary progenitors that best produces the observed SN rate. Again one should note that these binary models group SNIb and Ic together and do not attempt to predict what differences in progenitor produce these two types.\\\\ Given the different predictions for the origin of the CC SN types described above, observations are needed to discriminate between these models and thus firmly tie down the progenitors of the different SN types. However, apart from a small number of direct detections of progenitors (\\citealt{sma04,maun05,hen06,li06,gal07,li07,crock08}) this observational evidence remains sparse. Therefore here we present results to test the above predictions and constrain differences in progenitor mass of the different CC SN sub-types by investigating the nature of their parent stellar populations within host galaxies.  \\subsection{Progenitor constraints from parent stellar populations} \\label{pops} The most obvious way to determine the nature of SN progenitors is to investigate the properties of their stars on pre-explosion images. This has had some success although it is only possible for events in very nearby galaxies and therefore the statistics remain low. Another way is to investigate how the rates of the various SN types vary with different parameters, such as redshift or host galaxy properties. Our approach is intermediate to these methods as we attempt to constrain the nature of SN progenitors through investigating the environments and stellar populations at the positions of historical SNe. Here we concentrate on the association of the different CC SNe types with recent star formation (SF) as traced by \\ha\\ emission.\\\\ \\cite{ken98} states in a review paper on \\ha\\ imaging techniques that: ``only stars with masses $>$10\\msun\\ and lifetimes of $<$20 Myr contribute significantly to the ionising flux''. Thus, if our understanding of this line emission is correct, we can use this assumption as a starting point to constrain the relative stellar lifetimes and therefore the relative masses of the various SN progenitors, through investigating how accurately the different SN types trace the emission. In JA06 we presented a statistic to quantitatively measure the association of individual SNe with the \\ha\\ emission of their host galaxies, and presented results from an initial galaxy sample (\\ha GS, discussed in \\S \\ref{data}). It was found that overall the SNII progenitor population did not trace the underlying SF of their host galaxies, with a significant fraction lying on regions of low or zero emission line flux which were ascribed to a `Runaway' fraction of progenitor stars (however, this assumed that SNII arise from progenitors of $>$10\\msun ). The SNIb/c did appear to follow the emission implying that these progenitors come from higher mass stars than the SNII, although the statistics on this class were small (only 8 SNe for SNIb and Ic combined). This SN/galaxy sample has now been significantly increased, enabling the full parameter space of CC SN progenitors to be investigated and results from this increased sample are presented here.\\\\ The paper is arranged as follows: in the next section we present the data and discuss the reduction techniques employed, in \\S \\ref{ncr} we summarise the statistic introduced in JA06 and used throughout this paper, in \\S \\ref{results} we present the results for the different CC SN types, in \\S \\ref{diss} we discuss possible explanations for these results and their implications for the relative masses of the SN progenitors, and finally in \\S \\ref{con} we draw our conclusions. ",
        "conclusions": "\\label{con} We find that there is a significant fraction of the SNII population that do not show any association to the distribution of \\ha\\ line emission. This excess of $\\sim$35\\%\\ of SNII falling on sites of little or zero \\ha\\ flux, compared to what would be expected if they accurately traced the underlying SF, suggests that a large fraction of SNII arise from progenitor stars of less than 10\\msun . Our results also imply that the different CC SN types can be separated into a sequence of increasing progenitor mass running from the SNII through the Ib, with finally the SNIc arising from the highest mass progenitors. We now summarise our findings on the possible relative mass ranges of the progenitors of the different CC SN types. \\begin{itemize} \\item Assuming that only stars of 10\\msun\\ and above significantly contribute to the ionising flux that produces \\ha\\ emission within galaxies, we calculate a lower mass limit for producing SNII of 7.8\\msun . \\item We confirm the results of JA06, that the SNIb/c trace the SF of their host galaxies more accurately than the SNII, implying that they arise from a higher mass progenitor population than the SNII. \\item SNIc accurately trace the underlying SF within their host galaxies and therefore probably arise from the highest mass progenitors of all SNe. \\item SNIIL show a similar degree of association to \\ha\\ emission as the overall SNII population implying that they arise from stars of similar mass to those of SNIIP, with metallicity or binarity probably playing an important role in removing part of their envelopes and thus changing the shape of their light curves. \\item Our results suggest that SNIIb arise from more massive stars than the overall SNII population. \\item Although some SNIIn may arise from very massive stars, our results suggest that the majority come from the low end of the CC mass spectrum. \\item SN `impostors' do not seem to trace the high mass SF within host galaxies. \\end{itemize}"
    },
    "0809/0809.4003_arXiv.txt": {
        "abstract": "We perform an autocorrelation study of the Auger data with the aim to constrain the number density $n_s$  of ultrahigh energy cosmic ray (UHECR) sources, estimating at the same time the effect on $n_s$ of the systematic energy scale uncertainty and of the distribution of UHECR. The use of global analysis has the advantage that no biases are introduced, either in $n_s$ or in the related error bar, by the \\emph{a priori} choice of a single angular scale. The case of continuous, uniformly distributed sources is nominally disfavored at 99\\% C.L. and the fit improves if the sources follow the large-scale structure of matter in the universe. The best fit values obtained for the number density of proton sources are within a factor $\\sim$2 around $n_s\\simeq 1\\times 10^{-4}/$Mpc$^3$ and depend mainly on the Auger energy calibration scale, with lower densities being preferred if the current scale is correct. The data show no significant small-scale clustering on scales smaller than a few degrees. This might be interpreted as a signature of magnetic smearing of comparable size, comparable with the indication of a $\\ap 3^\\circ$ magnetic deflection coming from cross-correlation results. The effects on the above results of some approximations done is also discussed. ",
        "introduction": "Evidence is now emerging that ultrahigh energy cosmic rays (UHECRs) have an astrophysical origin, as opposed to being generated in exotic top-down models: The detection of a spectral suppression consistent with the Greisen-Zatsepin-Kuzmin (GZK) effect \\cite{Greisen:1966jv,Zatsepin:1966jv} by both HiRes~\\cite{Abbasi:2007sv} and Auger~\\cite{Abraham:2008ru} collaborations, together with the Auger bounds on the UHE neutrino flux~\\cite{Abraham:2007rj}, on the photon fraction~\\cite{Aglietta:2007yx} and on the anisotropy towards the Galactic Center~\\cite{Aglietta:2006ur,Aloisio:2007bh} are all consistent with this scenario. The next step is clearly the identification of the sources of UHECRs, an arena where anisotropy studies play a crucial role. Yet, the limited angular resolution of extensive air shower detectors and especially the deflections that charged particles suffer in astrophysical magnetic fields make the task highly non trivial. This is especially troublesome given that the UHECRs chemical composition is unknown, that we lack a detailed knowledge of the Galactic magnetic field structure and, above all, of the very magnitude and structure of extragalactic magnetic fields (EGMF) outside of cluster cores. These limitations---together with the small statistics available at present---suggest that, at least in an initial phase, charged particle astronomy may be limited to the inference on the statistical properties of UHECR sources, rather than a detailed study of single accelerators.  In Ref.~\\cite{I}, we found that a global comparison of the two-point auto-correlation function of the data with the one of catalogues of potential sources is a powerful diagnostic tool: This observable is less sensitive to unknown deflections in magnetic fields than the cross-correlation function, while keeping a strong discriminating power among source candidates. In particular,  the autocorrelation function of (sub-) classes of galaxies have different biases with respect to the large-scale structure (LSS) of matter. As a result, the best fit value for the density $n_s$  of different source classes may differ, especially if only one or a small range of angular scales is considered. Although the bias of different source classes differs maximally at small angular scales, we showed that the statistically most significant differences are at intermediate angular scales, where both the larger number of cosmic ray pairs (CR) and of galaxy pairs leads to relatively smaller error bars. Moreover, the autocorrelation function on larger angular scales becomes less dependent on possible deflections in the Galactic and extragalactic fields.\\\\ \\indent In this article we derive the number density of UHECR sources using the recently published arrival directions and energies of the 27 Auger events~\\cite{PAO} with estimated energy $E\\geq 57$\\,EeV, thereby complementing the study~\\cite{I} with a concrete example for a comparison of the {\\em global\\/} cumulative autocorrelation function of sources and UHECRs. Note that we showed in Ref.~\\cite{I} that, even in an idealized case where systematics play no major role, roughly three times the number of ``useful'' events that can be extracted from Ref.~\\cite{PAO} are required to start distinguishing between different sub-classes of sources. Thus a study of the kind envisaged in Ref.~\\cite{I} is unrealistic at present. Here, we restrict ourselves to more modest goals: i) To compare the data against predictions of two toy model cases of uniformly distributed sources and of ``normal galaxies''  (i.e.\\ sources that have the same distribution as the PSCz catalogue~\\cite{Saunders:2000af}) which we shall refer to with the two values for the label $\\kappa=\\{{\\rm uni},{\\rm LSS}\\}$, respectively. ii) To study the effect on the allowed range of $n_s$ of a systematic error on the energy scale of the UHECR experiment. Note that preliminary results of the clustering of the Auger events has been presented in \\cite{Mollerach:2007vb}, but astrophysical implications have not been discussed there.  We review the statistical method we use in Sec.~\\ref{analysis}, and apply it in  Sec.~\\ref{int} to the Auger data, providing some interpretation of the results. In Sec.~\\ref{disc} we discuss some limitations and caveats of the analysis. Finally, we summarize in Sec.~\\ref{sum}. ",
        "conclusions": "i) that a correlation between the two optimal cuts in energy is indeed present; ii) that a slightly different cut, in the range $40\\lsim E_{\\rm cut}/{\\rm EeV}\\lsim 60$, should still lead to an appreciable clustering of the events.  We checked that this is indeed the case by performing the same analysis  as before, but now adding a further scan  over the range of accessible energy-cuts, i.e. any value $E>57\\,$EeV. We plot the results in Fig.~\\ref{fig:energyscan}\\footnote{Some ripples visible in Fig.~\\ref{fig:energyscan} are due to the relatively low number of Montecarlo: the scan to account for the further penalty factor is computationally quite expensive and given the partial nature of the answer there is no motivation to refine the results further.}. One can note that in all cases the best fit improves and the constraints on $n_s$ worsen, as expected given the further penalty due to the energy scan. Yet, most qualitative features described in the previous section stay the same: for example, the LSS model is still preferred over a uniform one. It should be also said that if the true minimum of the chance probability is below $E=57$ EeV and thus not included in the scan, then these constraints are ``over-penalized'' and thus looser than necessary. At the moment, it is impossible to draw more quantitative conclusions,   since our scan suggests that it is likely that the optimal cut for the autocorrelation function is below $E=57$ EeV, a range for which the events are not publicly available.   \\begin{figure*}[!htbp] \\begin{center} \\begin{tabular}{cc} \\epsfig{file=penalty_like_Uniform60EeV_EnergyScan.eps,width=0.47\\textwidth} & \\epsfig{file=penalty_like_LSS60EeV_EnergyScan.eps,width=0.47\\textwidth}\\\\ \\end{tabular} \\begin{tabular}{cc} \\epsfig{file=penalty_like_Uniform80EeV_EnergyScan.eps,width=0.47\\textwidth} & \\epsfig{file=penalty_like_LSS80EeV_EnergyScan.eps,width=0.47\\textwidth}\\\\ \\end{tabular} \\end{center} \\vspace{0pc} \\caption{ Penalized chance probabilities $p_{-}(n_s)$, $p_{+}(n_s)$ and $p(n_s)$, for  $\\Ecut=60\\,$EeV (top panels) and $\\Ecut=80\\,$EeV (bottom panels) taking into account the effect of the energy scan. The left column reports the case for uniformly distributed sources, the right panel for sources following LSS with the bias of the PSCz Galaxy catalogue. Also shown is the 95\\% and 99\\% confidence level.}\\label{fig:energyscan} \\end{figure*}     \\subsection{Assumptions on the chemical composition}\\label{chemadd} In this article, we assumed dominant proton primary as a basic working hypothesis. It is worth noting, however,  that the experimental situation on the chemical composition at UHE is far from  settled:  while anisotropy data point to a relatively light composition, the results of the fluorescence detector of the Pierre Auger Collaboration favor a significant fraction of heavy nuclei~\\cite{Unger:2007mc}. Yet, for the interpretation of these results one must rely on simulations employing hadronic interaction models. These are not based on a first-principle theory, rather on models calibrated on ``low-energy''  collider data, then {\\it extrapolated} about two orders of magnitude beyond the center of mass energies experimentally probed.  A proper quantitative assessment on how our conclusions vary in  a mixed composition scenario goes beyond the purpose of the present paper. Yet, at a qualitative level, we can note that several effects would come into play. First of all,  in the unrealistic case where one could forget about magnetic deflections, the major effect would be a reduction of the energy-loss horizon (but for iron, whose horizon is similar to the proton one). This should {\\it enhance} the anisotropy pattern, due to the prominence of nearby accelerators.  When including (the poorly known) magnetic fields, two additional effects are relevant: i) for a given energy, the higher the charge the larger the deflection and hence the  loss of information at {\\it small} angular separations. Quantifying how large is this scale is a difficult task. We note that protons of these energies in the sole Galactic field likely suffer a few degrees absolute deflections, see~\\cite{Kachelriess:2005qm}, implying a degree-scale smoothing in the relative deflections important for the 2pcf. This is comparable with the angular resolution of the PAO: in this optimal case the whole information in the 2pcf starting at $\\theta_{\\rm min}\\sim 1^\\circ$ could be used. However, a different Galactic field model (especially towards the Galactic Center), the presence of heavy nuclei, and/or significant extragalactic magnetic fields can easily lift $\\theta_{\\rm min}$ by one order of magnitude or more. ii) the other effect is that the real path-length of the nucleus would be longer than the distance to the source: thus, the distance of accessible sources would be even shorter than estimated from energy-loss considerations. This is particularly relevant if the nucleus spends a lot of time in a magnetized region surrounding the accelerator (e.g. a magnetized cluster in which it is immersed) before escaping in the Intergalactic Medium.  Finally, if the maximal energy of acceleration of different  species of nuclei fall by unfortunate coincidence in the same region expected for the GZK feature, slightly different energy cuts in the data (as well as statistical and systematic errors on the energy scale) might significantly change the expected pattern of anisotropies. The same happens if the proportions of different nuclei accelerated at the source change as a function of the energy. This is in principle a possibility, especially if different classes of objects contribute to the events at slightly different energies.  The general pessimistic conclusion is that if several of  the above effects are relevant (or perhaps a single one is {\\it dominant}) the capability of performing UHECR astronomy would be greatly reduced. While one still expects indications for anisotropies, inverting the problem and  inferring the source/propagation medium properties would require a much larger statistics (especially in the trans-GZK region): disentangling the different effects is in fact a formidable task.  To provide a glimpse of how some of the  above effects alter the reconstruction of $n_s$ (our main topic here), in Table~\\ref{tab:2} we report how the constraint on $n_s$ degrades as a function of a ``smoothing angle'' $\\theta_{\\rm min}$, below which we assume that the 2pcf information is completely lost.  The main trend is that, if the smallest angular scales are neglected, it is easier to find parameter configuration fitting the data and, correspondingly, the allowed range for $n_s$ widens. This had to be expected, given the shape of the correlation functions shown in Figs. 2-3.   In particular we find that with $\\theta_{\\rm min}=3^\\circ$ we obtain almost the same results as in the global case. The case $\\theta_{\\rm min}=10^\\circ$ still places useful constraints especially for the $\\Ecut=80\\,$EeV case, while finally using only the information above $\\theta_{\\rm min}=30^\\circ$ basically no constraints on $n_s$ are obtained. Note however that relative deflection angles of that sort would imply overall deflections even larger, seriously questioning the perspectives of present instruments to perform some form of UHECR astronomy.  \\begin{table*}[!t] \\begin{center} \\footnotesize{ \\begin{tabular}{|c|c|c|c||c|} \\hline $n_s/10^{-4}$Mpc$^{-3}$, \\,  $\\theta_{\\rm min}$= & $3^\\circ$ & $10^\\circ$ & $30^\\circ$ & global \\\\ \\hline \\hline LSS (80)&   $\\; 1.0^{+100}_{-0.8}  \\;$      &  $\\; 0.5^{+30}_{-0.4}    \\;$&  $\\; 0.5^{+\\infty}_{-0.4} \\; $ & $\\; 1.3^{+100}_{-0.8} \\;$\\\\ \\hline Uniform (80)& $\\;0.8^{+1.5}_{-0.6} \\;$      &  $\\; 0.3^{+1.7}_{-0.2}   \\;$&  $\\; 0.3^{+\\infty}_{-0.2} \\; $ & $\\; 1.4^{+1.4}_{-0.7} \\;$\\\\ \\hline LSS (60) &   $\\; 0.3^{+20}_{-0.28}   \\;$    &  $\\; 0.1^{+5}_{-0.09}    \\;$&   n.c.                   & $\\; 0.8^{+19}_{-0.6} \\;$\\\\ \\hline Uniform (60) & $\\; 0.2^{+0.8}_{-0.12} \\;$   &  $\\; 0.1^{+0.9}_{-0.085} \\;$&   n.c.                   & $\\; 0.5^{+0.5}_{-0.2} \\;$\\\\ \\hline \\end{tabular}  } \\end{center} \\caption{\\label{tab:2}  The estimated number density of sources (at 95\\% confidence level) under different assumptions on the minimum angle above which the 2pcf information is preserved, $\\theta_{\\rm min}$. n.c. stands for ``no constraints''.} \\end{table*}"
    },
    "0809/0809.0766_arXiv.txt": {
        "abstract": "We present Hubble Space Telescope imaging and spectroscopic observations of three Brightest Cluster Galaxies, Abell 1836-BCG, Abell 2052-BCG, and Abell~3565-BCG, obtained with the Wide Field and Planetary Camera 2, the Advanced Camera for Surveys and the Space Telescope Imaging Spectrograph. The data provide detailed information on the structure and mass profile of the stellar component, the dust optical depth, and the spatial distribution and kinematics of the ionized gas within the innermost region of each galaxy. Dynamical models, which account for the observed stellar mass profile and include the contribution of a central supermassive black hole (SBH), are constructed to reproduce the kinematics derived from the H$\\alpha$ and [N II]$\\lambda\\lambda$6548,6583 emission lines. Secure SBH detection with $M_{\\bullet} = 3.61^{+0.41}_{-0.50} \\times 10^9$ M$_{\\odot}$ and $M_{\\bullet} = 1.34^{+0.21}_{-0.19} \\times 10^9$ M$_\\odot$, respectively, are obtained for Abell 1836-BCG and Abell 3565-BCG, which show regular rotation curves and strong central velocity gradients. In the case of Abell~2052-BCG, the lack of an orderly rotational motion prevents a secure determination, although an upper limit of $M_{\\bullet} \\lesssim 4.60 \\times 10^9$ M$_{\\odot}$ can be placed on the mass of the central SBH. These measurements represent an important step forward in the characterization of the high-mass end of the SBH mass function. ",
        "introduction": "\\label{sec:intro}  Within the past decade, the focus in supermassive black holes (SBHs) studies has moved from the dynamical measurement of SBH masses, $M_{\\bullet}$, in nearby galaxies (see review by Ferrarese \\& Ford 2005), to the characterization of scaling relations connecting $M_{\\bullet}$ to the large scale properties of their hosts (Kormendy \\& Richstone 1995; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Graham et al. 2001; Ferrarese 2002; Marconi \\& Hunt 2003). Such relations, combined with the knowledge of the galaxy luminosity or velocity dispersion functions, lead to a direct determination of the local SBH mass function and, by comparison with the energetics of high redshift AGNs, accretion history (e.g. Marconi et al. 2004; Shankar et al. 2004; Benson et al. 2007; Tundo et al. 2007; Graham et al. 2007; Lauer et al. 2007). Furthermore, the tightness of the relations linking galactic properties to $M_{\\bullet}$ is indicative of a formation/evolutionary history in which SBHs and galaxies are causally connected. Indeed, feedback from AGN activity is believed to play an important, perhaps even fundamental role in the evolution of galaxies (e.g. Binney \\& Tabor 1995; Suchkov et al. 1996; Ciotti \\& Ostriker 2001; Schawinski et al.2006; Springel et al. 2005; Hopkins et al. 2007; Di Matteo et al. 2007).   In this context, Brightest Cluster Galaxies (BCGs) are of particular interest. Their privileged location at the center of a massive cluster implies that they have undergone a particularly extensive merging history (Khochfar \\& Silk 2006) and are likely to host the most massive SBHs in the local Universe (Yoo et al. 2007). The latter conclusion is also supported by all scaling relations which are known to be obeyed by local SBHs, that predict the most massive SBHs to reside in most massive galaxies. As such, BCGs constitute an excellent laboratory to search for the local relics of the most powerful high-redshift quasars (Willott et al. 2003; Vestergaard 2004), and to investigate the role of mergers in the black hole mass function.   Unfortunately, direct dynamical measurements of SBHs masses in BCGs are exceedingly difficult - with only two such measurements made to-date (in M87 and NGC 1399, Harms et al. 1994; Macchetto et al. 1997; Houghton et al. 2006). The reason for this is simple: SBH mass measurements based on resolved kinematic tracers (gas or stars) need to be carried out within the SBH ``sphere of influence\", i.e. the region of space within which the SBH dominates the overall gravitational potential. For a $10^9$ $M_{\\odot}$ SBH, the sphere of influence becomes unresolved for optical measurements with the Hubble Space Telescope at distances beyond $\\sim 100$ Mpc (see Figure 44 of Ferrarese \\& Ford 2005). Few BCGs are located within this limit.  To compound the problem, bright ellipticals are characterized by shallow stellar density profiles and faint central surface brightnesses (e.g., Ferrarese et al. 1994; Lauer et al. 1995; Rest et al. 2001; Ferrarese et al. 2006) making the detection of stellar absorption lines with {\\it HST} prohibitively expensive.  The latter difficulty can be overcome with the use of gas dynamics. Emission lines (mainly \\ha\\ and \\nii) from the gas are bright and easily measured. If the gas is confined in a disk, there is little ambiguity in the velocity distribution (Ho et al. 2002), and since the method was first used (Harms et al. 1994; Ferrarese et al. 1996; Ferrarese \\& Ford 1999) an increasing amount of attention has been devoted to the theoretical aspects of the dynamical modeling (Maciejewski \\& Binney 2001; Barth et al.  2001; Cappellari et al. 2002; Coccato et al. 2006).  The lack of a secure characterization of the SBH mass function above the $10^9$ $M_{\\odot}$ mark is troublesome.  Lauer et al. (2007) suggest that the relations between $M_{\\bullet}$ and host bulge luminosity, $L_{bulge}$, (Kormendy \\& Richstone 1995) and velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000) would predict significantly different $M_{\\bullet}$ if extrapolated above $10^9$ $M_{\\odot}$. In particular, the $M_{\\bullet} - \\sigma$ relation would predict less massive SBHs in BCGs than the $M_{\\bullet} - L_{bulge}$ relation, due to the slower increase of $\\sigma$ with galaxy luminosity observed for BCGs compared to the bulk of the Early-Type galaxy population (Bernardi et al. 2007). The difference in the predicted mass is such to significantly affect (by an order of magnitude) the high-mass end of the local SBH mass function.  von der Linden et al. (2007) argue that the shallower dependence of $\\sigma$ on $L$ applies to BCG but not to comparably massive non-BCG galaxies, which instead follow the canonical Faber-Jackson relation defined by less massive systems. This implies that BCGs and non-BCG galaxies of comparable mass must occupy a different locus in either, or both, the $M_{\\bullet} - \\sigma$ and $M_{\\bullet} - L_{bulge}$. This result is in contrast with the findings of Batcheldor et al. (2007), who argue that SBHs masses predicted from NIR luminosities (Marconi \\& Hunt 2003) are consistent with those predicted from $\\sigma$. They attributed the discrepancy noted by Lauer et al. (2007) (who used $V-$band luminosities) to a bias introduced by the inclusion, when computing $L_{bulge}$, of extended blue envelopes around BCGs. At present these ambiguities prevent us from testing any theoretical prediction on the high-mass end of the black hole mass function based on the observed quasar abundances and merger histories.  In this paper, we analyze the kinematics of the ionized gas in the nuclear region of three BCGs -- Abell~1836-BCG, Abell~2052-BCG, and Abell~3565-BCG -- in order to constrain the mass of the central SBHs. The data were obtained using the Space Telescope Imaging Spectrograph (STIS) on board the Hubble Space Telescope (HST), supplemented with imaging data from the Advanced Camera for Surveys (ACS) and the Wide Field and Planetary Camera 2 (WFPC2).  Furthermore, the central stellar velocity dispersion of Abell 1836-BCG was obtained from ground based spectroscopy.  The paper is organized as follows. The criteria of galaxy selection are presented in \\S \\ref{sec:sample}. ACS and STIS observations are described and analyzed in \\S\\S \\ref{sec:imaging} and \\ref{sec:spectroscopy}, respectively. Ground based spectroscopic observations of Abell 1836-BCG are described and analyzed in \\S \\ref{sec:groundspectrospcopy}.  The SBH masses of Abell~1836-BCG and Abell~3565-BCG, and an upper limit for the SBH mass of Abell~2052-BCG are derived in \\S \\ref{sec:model}. In \\S \\ref{sec:conclusions} results are compared to the predictions of the SBHs scaling laws and conclusions are given. ",
        "conclusions": "\\label{sec:conclusions}   We have presented a dynamical analysis aimed at constraining the masses of the SBHs in three brightest cluster galaxies: Abell 1836-BCG, at a distance of 147.2 Mpc, Abell 2052-BCG, at a distance of 141.0 Mpc, and Abell 3565-BCG, at a distance of 50.8 Mpc. The models are based on data obtained with the Hubble Space Telescope. Broad-band WFPC2 and broad and narrow-band ACS images were used to constrain the luminosity profile and the distribution of the ionized gas, as traced by the H$\\alpha+$[N~{\\small II}] emission, while STIS was employed to measure the rotation velocity and velocity dispersion profiles, from the [N~{\\small II}]$ \\,\\lambda6583$ emission, along three parallel slits positions, the first aligned along the photometric major axis, and the others adjacent to the first on either side of the nucleus.  In the case of Abell~1836-BCG and Abell~3565-BCG, the regular morphology and kinematics observed for the ionized gas led to secure black hole mass detections of \\mbh$=$ $3.61^{+0.41}_{-0.50}\\times 10^9$ \\msun\\ and $1.34^{+0.21}_{-0.19}\\times 10^9$ \\msun\\ (where the uncertainties represent 1$\\sigma$ errors\\footnote{Barth et al. 2001 show that the total error budget for gas-dynamical modelling is likely to be twice as large as such formal error estimates.}), respectively. For Abell-2052-BCG, which displays irregular kinematics, an upper limit of \\mbh$\\leq 4.60 \\times10^9$ \\msun\\ was derived under the conservative assumption of a negligible stellar contribution to the gravitational potential, and an inclination angle for the gas disk of $33\\dg$. At face value, the black hole in Abell~1836-BCG is the most massive to have been dynamically measured to-date.   In our modeling we have neglected the potential impact on our SBH mass determination of assigning a dynamical origin to the additional kinematic broadening that we have used in our models to match the observed profile for the gas velocity dispersion. Although it has been argued that this additional turbulence does not affect the bulk motion of the gas (e.g. van der Marel \\& van den Bosch 1998), it is important to notice that ignoring the possibility that the gas might be supported by dynamical pressure will lead to {\\it under-}estimate the SBH mass. For instance, Barth et al. (2001) have shown that including a classical asymmetric drift correction led to an 12\\% increase of their best \\mbh\\ value. Unfortunately, we could not repeat this analysis in our models since the observed line-broadening exceeds by far the circular velocity, breaking down the epicyclic approximation on which the asymmetric drift correction is based. Nonetheless, our best SBH mass measurements are still likely to be underestimations of the real SBH mass.   Figure~\\ref{scal_rel} shows the location of Abell~1836-BCG, Abell~2052-BCG, and Abell~3565-BCG in the \\mbh$-$\\sigmac\\ (Ferrarese \\& Ford 2005; Tremaine et al. 2002) and near-infrared \\mbh$-$\\lbulge\\ (Graham 2007; Marconi \\& Hunt 2003) planes. $K$-band Two-Micron All-Sky Survey (2MASS) magnitudes were retrieved from the NASA/IPAC Infrared Science Archive, and corrected for Galactic extinction following Schlegel et al. (1998).  Stellar velocity dispersions are from this paper (see \\S \\ref{sec:groundspectrospcopy}) for Abell 1836-BCG, from Smith et al. (2000) in the case of Abell~3565-BCG (we adopt his high S/N measurement, $\\sigma=335\\pm 12$ \\kms), and Tonry (1985) for Abell~2052-BCG. We note that published estimates of $\\sigma$ for Abell 2052-BCG range between $250$ \\kms\\ and $370$ \\kms; the high end of this spread is probably due to contamination from a companion galaxy located only $8''$ away in the North-East direction. The contribution from this galaxy was explicitly accounted for by Tonry (1985), whose value, $\\sigma=253\\pm 12 $ \\kms, we have adopted. Correcting the adopted values of $\\sigma$ to a circular aperture of size $1/8 r_e$ (Table~\\ref{tab:sample}), following Jorgensen et al. (1995), we find $\\sigma_c=288 \\pm 9$ \\kms, $\\sigma_c=233 \\pm 11$ \\kms, and $\\sigma_c=322\\pm12$ \\kms\\ for Abell~1836-BCG, Abell~2052-BCG and Abell~3565-BCG, respectively.   The SBH mass detected in Abell~3565-BCG, \\mbh = $1.34^{+0.21}_{-0.19}\\times 10^9$ \\msun\\ , is consistent both with the \\mbh$-$\\sigmac\\ (Ferrarese \\& Ford 2005) and the $K$-band \\mbh$-$\\lbulge\\ (Graham 2007) relations, which predict \\mbh = $1.7^{+2.0}_{-0.9}\\times 10^9$ \\msun\\ (the error is computed adopting a 0.34 dex scatter in \\mbh) and \\mbh = $1.1^{+1.1}_{-0.6}\\times 10^9$ \\msun\\ (adopting a 0.30 dex scatter), respectively.  For Abell~2052-BCG, although the conservative upper limit on the black hole mass obtained for reasonable assumptions for \\mlstar\\ and $i$ (\\S 5.3), is consistent with both the \\mbh$-$\\sigmac\\ (Ferrarese \\& Ford 2005) and the $K$-band \\mbh$-$\\lbulge\\ (Graham 2007) relations, these predict significantly different masses, $3.5^{+4.1}_{-1.9}\\times 10^8$ \\msun\\ (adopting a 0.34 dex scatter) and $1.3^{+1.3}_{-0.6}\\times10^9$ \\msun\\ (adopting a 0.30 dex scatter) respectively (but note the above caveat regarding the measurement of $\\sigma$ for this galaxy). Although the sense of the discrepancy is consistent with the one argued for by Lauer et al. (2007), the lack of a precise determination of the SBH mass in this galaxy prevents us from establishing which, if either, of the two relations better predicts $M_{\\bullet}$ in this case.  The only two BCGs in the literature with a dynamically measured SBH mass are M87 (Macchetto et al. 1997) and NGC 1399 (Houghton et al. 2006). M87 appears to have a more massive black hole than expected based on its $K-$band luminosity while it obeys the \\mbh$-$\\sigmac\\ relation. NGC 1399 is consistent with both the relations within their scatter.  Abell 1836-BCG appears to be an outlier in all SBH scaling relations. The dynamically measured SBH mass reported in this paper, $M_\\bullet=3.61^{+0.41}_{-0.50}\\times 10^9 $ \\msun, is larger, at the 3$\\sigma$ level, than the value of \\mbh = $9.6^{+9.5}_{-4.8}\\times 10^8$ \\msun\\ (adopting a 0.30 dex scatter), predicted by the $K-$band \\mbh$-$\\lbulge\\ relation (Graham 2007). It is also larger, at the 3$\\sigma$ level, than the value of \\mbh = $9.77^{+11.6}_{-5.30}\\times 10^8$ \\msun\\ (adopting a 0.34 dex scatter), predicted by the \\mbh$-$\\sigmac\\ relation of Ferrarese \\& Ford (2005). Finally, the SBH mass of Abell 1836-BCG is not consistent with the value of $M_\\bullet=6.3^{+11.5}_{-4.1}\\times 10^8 $ \\msun\\ predicted by the fundamental plane relation for SBHs by Hopkins et al. (2007), adopting an effective radius $r_e=13\\Sec1=9.3$ kpc with a scatter of 0.45 dex in \\mbh.  The existence of extremely massive SBHs as outliers to the empirical \\mbh$-$\\sigmac\\ relation would not be surprising. As shown by Lauer et al. (2007), the \\mbh$-$\\lbulge\\ relation predicts a greater abundance of the most massive SBHs, given the luminosities of the BCGs and the lack of galaxies with velocity dispersions larger than $\\sim 350$ \\kms. At the same time, the observed luminosity function of AGN implies the presence of SBHs as massive as $\\sim 5\\times 10^9$~\\msun\\ within the distance of the clusters we have observed, and SBHs mergers can increase the highest SBH masses in these massive clusters up to $\\sim 10^{10}\\,$\\msun\\ (Yoo et al. 2007). A larger sample of SBH masses measured in massive clusters would be required to test models of the effect of mergers in increasing the largest SBH masses in the universe.  In conclusion, both the \\mbh$-$\\sigmac\\ and \\mbh$-$\\lbulge\\ relations appear at their high-\\mbh\\ end to be consistent with the SBH mass measured for one BCG, Abell 3565-BCG, but inconsistent with another one, Abell 1836-BCG. For the remaining target, Abell 2052-BCG, although the ionized-gas kinematics allowed us only to set an upper-limit on \\mbh\\ it seems unlikely that this galaxy could obey both relations simultaneously. The fact that Abell 1836-BCG is an outliers in both the \\mbh$-$\\sigmac\\ and \\mbh$-$\\lbulge\\ relations would appear to weaken the claim of Lauer et al (2007) that \\lbulge\\ is more reliable predictor of \\mbh\\ for BCGs. Overall, our results might indicate that the scatter of SBH scaling relations increases at the high end, although additional data are necessary to test this claim."
    },
    "0809/0809.5233_arXiv.txt": {
        "abstract": "Measurements by dust detectors on interplanetary spacecraft appear to indicate a substantial flux of interstellar particles with masses $> 10^{-12}\\gram$. The reported abundance of these massive grains cannot be typical of interstellar gas: it is incompatible with both interstellar elemental abundances {\\it and} the observed extinction properties of the interstellar dust population. We discuss the likelihood that the Solar System is by chance located near an unusual concentration of massive grains and conclude that \\newtext{this} is unlikely, unless dynamical processes in the ISM are responsible for such concentrations.  Radiation pressure might conceivably drive large grains into ``magnetic valleys''. If the influx direction of interstellar gas and dust is varying on a $\\sim10$~yr timescale, \\newtext{as suggested by some observations,} \\newtext{this would} have dramatic implications for the small-scale structure of the interstellar medium. ",
        "introduction": "Introduction}  The interstellar medium (ISM) consists of a partially-ionized, magnetized gas mixed with solid particles of dust. The ionization state and molecular fraction of the gas depend primarily on the gas density and the local intensity of ultraviolet radiation that can photodissociate molecules and photoionize molecules and atoms. The dust content is determined by the prior history of the gas, including injection of newly-formed dust in stellar winds and supernova explosions, grain destruction in violent events such as supernova blast waves, and grain growth in the interstellar medium by both vapor deposition and coagulation in dense regions.  While we do not know the properties of interstellar dust with precision, they are strongly-constrained by a variety of observations. The observed wavelength dependence of interstellar extinction -- the so-called ``reddening curve'' (reviewed in \\S~\\ref{sec:ism dust}) -- provides strong constraints on both the composition and size distribution of interstellar dust. In the local regions of the Milky Way, interstellar dust is abundant, containing a large fraction of the elements (such as Mg, Si, and Fe) that can be incorporated into refractory solids. As discussed in \\S\\ref{sec:abundances}, interstellar abundances therefore provide a strong constraint on grain models.  The size distribution of interstellar grains can be inferred from the observed average reddening curve together with interstellar abundance constraints.  Microparticle impacts on detectors on Ulysses and Galileo have been interpreted as showing a flux of solid particles entering the Solar system from the local interstellar medium \\citep{Grun+Zook+Baguhl+etal_1993}. In \\S~\\ref{sec:abundances} we show that the population of large grains inferred from dust impact detectors on Ulysses and Galileo \\citep{Landgraf+Baggaley+Grun+etal_2000, Kruger+Landgraf+Altobelli+Grun_2007, Kruger+Grun_2008} is incompatible with average elemental abundances in the ISM. In \\S~\\ref{sec:extinction}, we show that such a large grain population would result in wavelength-dependent extinction very different from what is observed.  The Ulysses and Galileo data, if correctly interpreted, imply that the Solar System is, by chance, located in a very atypical spot in the ISM, with an overabundance of very large grains. The likelihood of such a scenario is discussed in \\S~\\ref{sec:atypical}. In \\S~\\ref{sec:velocity structure} we comment on \\newtext{suggestions} that the interstellar dust inflow vector \\newtext{might have changed} appreciably over only $\\sim$5~yrs. Our conclusions are summarized in \\S~\\ref{sec:summary}. ",
        "conclusions": "Summary}  The size distribution of interstellar grains entering the heliosphere, as inferred from observations by Ulysses and Galileo \\citep{Landgraf+Baggaley+Grun+etal_2000,Kruger+Landgraf+Altobelli+Grun_2007} cannot be typical of the general interstellar medium, as can be demonstrated by two independent arguments: \\begin{enumerate} \\item The required abundance of elements in grains would substantially exceed what is available in the interstellar medium. \\item If such a size distribution were generally present, it would produce an interstellar ``reddening law'' very different from what is observed. \\end{enumerate} Therefore, if the size distribution of local interstellar dust does have the large grain population reported by \\citet{Landgraf+Baggaley+Grun+etal_2000}, the dust grain/gas ratio in the interstellar medium must be quite nonuniform.  The length scale characterizing these nonuniformities is not known.  If the velocity vector of the incoming dust flow is \\newtext{actually} changing over time scales of only years \\newtext{ -- one possible explanation for the variations in the directions of impacting particles reported by} \\citet{Kruger+Landgraf+Altobelli+Grun_2007} \\newtext{--} this would require that the dust velocity vary over lengthscales of only tens of AU. Such small scale structure was not expected.  Mechanisms that might account for such nonuniformity are considered. It seems extremely unlikely that the Sun is passing through a region that has recently been enriched with dust from a stellar source. The least unlikely scenario may involve concentration of dust in certain regions, and removal of dust from other regions, by dynamical processes.  One possible mechanism involving anisotropic starlight driving dust grains along deformed magnetic field lines is outlined. Whether this can compete with the diffusive effects of turbulent mixing is far from clear, however.  It is important to carry out additional observations to confirm the enhanced grain size distribution, and to confirm the time-dependence of the density and velocity vector of the inflowing dust and gas. If the reported density of large grains, and the time-dependence of the inflow, are confirmed, this may require revision of our understanding of the small-scale structure of the ISM. Absorption line studies seem to suggest that, by coincidence, the heliosphere is just now passing through the transition zone -- possibly a shock transition -- between the ``Local Interstellar Cloud'' and ``Cloud G''.  If so, the flow into the heliosphere offers the opportunity to study the small-scale structure in this transition zone. The Ulysses observations indicate that this region is heavily enriched with large dust particles, although why this should be so remains unclear."
    },
    "0809/0809.2902_arXiv.txt": {
        "abstract": "We calculate relativistic Fermi liquid parameters (RFLPs) for the description of the properties of dense nuclear matter (DNM) using Effective Chiral Model. Analytical expressions of Fermi liquid parameters (FLPs) are presented both for the direct and exchange contributions. We present a comparative study of perturbative calculation with mean field (MF) results. Moreover we go beyond the MF so as to estimate the pionic contribution to the FLPs. Finally, we use these parameters to estimate some of the bulk quantities like incompressibility, sound velocity, symmetry energy etc. for DNM interacting via exchange of $\\s$, $\\o$ and $\\p$ meson. In addition, we also calculate the energy densities and the binding energy curve for the nuclear matter. Results for the latter have been found to be consistent with two loop calculations reported recently within the same model. ",
        "introduction": "One of the most exciting field of contemporary nuclear research has been the studies of the properties of dense nuclear matter (DNM). Such studies are important both in the context of laboratory experiments and nuclear astrophysics. Therefore, several attempts have been made in recent years to ascertain the properties of nuclear system at densities higher than the normal nuclear matter densities \\cite{chin76,speth80,horowitz83,friman96,friman99,song01,holt07}.   The suitable description of nuclear matter at such high densities is provided by Quantum Hadrodynamics (QHD) \\cite{serot92}. Historically, QHD was developed by Walecka \\cite{walecka74,chin74,vol16} to study the properties of neutron star where the nucleons are assumed to interact via the exchange of $\\sigma$ and $\\omega$ mesons. In this model starting with interacting Lagrangian the relativistic field equations are solved by making MF approximation where the meson fields are replaced by their vacuum expectation values. Subsequently, starting from the same model Chin developed a full diagrammatic scheme and showed that the MF results can be obtained by making Hartree approximation {\\em i.e.} by retaining only the direct terms in a relativistic field theoretic approach \\cite{chin77}. In the same work, it was also shown how exchange corrections can be made and analytical expressions can be found for the energy density and related quantities by making some long range approximation for the $n-n$ interaction. Since then the QHD has undergone a series of developments which we do not discuss here but refer the reader to ref.\\cite{serot79,kapusta81, serot82,furnstahl87,furnstahl93,furnstahl95,furnstahl96}.  The most recent model which we use here for the description of dense nuclear system is provided by the Chiral Effective Field theory (chEFT) \\cite{furnstahl89,furnstahl97}. It might be recalled here, that, in such a framework, the explicit calculation of the Dirac vacuum is not required, rather, on the contrary, here, the short distance dynamics are absorbed into the parameters of the theory adjusted phenomenologically by fitting empirical data. For detail discussion refer the reader to ref. \\cite{furnstahl97,serot97,hu07,biswas08}. Recently this model has been applied \\cite{hu07} to calculate the exchange corrections by evaluating nucleon loops involving $\\sigma$, $\\omega$ and $\\pi$ as intermediate states, which we address here partly.  Our approach here is to study the dense nuclear system in terms of relativistic Fermi liquid parameters (RFLPs). Such an extention of the Fermi Liquid theory \\cite{pines_book,baym_book} was first made by Baym and Chin in ref.\\cite{baym76}. It should, however, be noted that the calculation presented in ref.\\cite{baym76} were performed perturbatively where the original QHD model was used. It should, however, be noted that the first application of Fermi liquid theory (FLT) to study the nuclear system was due to Migdal \\cite{migdal78} who used FLT to investigate the properties of unbound nuclear matter and finite nuclei \\cite{mig_book}. FLT also provides theoretical foundation for the nuclear shell model \\cite{mig_book} as well as nuclear dynamics of low energy excitations \\cite{baym_book,krewald88}. The connection between Landau, Brueckner-Bethe and Migdal theories was discussed in ref.\\cite{brown71}. While these are all non-relativistic calculations, the relativistic calculations involving RFLT are rather limited.  After the original work of \\cite{baym76},  the relativistic problem was revisited in \\cite{matsui81} where one starts from the expression of energy density in presence of scalar and vector meson MF and takes functional derivatives to determine the FLPs. The results are found to be qualitatively different than the perturbative results \\cite{baym76,matsui81}. Moreover, besides $\\sigma$ and $\\omega$ meson, ref.\\cite{matsui81} also includes the $\\rho$ and $\\pi$ meson and the model adopted was originally proposed by Serot that incorporates pion into the Walecka model. The latter, however, do not contribute to the parameters presented in \\cite{matsui81} as the calculation was restricted only to the MF level where pion fails to contribute. On the other hand, the Migdal parameters using one-boson-exchange models of the nuclear force calculated in ref.\\cite{celenza_book,anastasio83}, in which a comparison of relativistic and non-relativistic results have also been studied.  In the present work, we use a model, where we have pions and we extend the calculation beyond MF to include the pionic contributions into the FLPs. Furthermore, we evaluate and compare the perturbative results with MF approximated results within the framework of the present model. In addition we also calculate various physical quantities like incompressibility, sound velocity and symmetry energy etc. Moreover, the results are compared whenever possible with the previous calculations by taking suitable limits. For instance, the exchange energy, we compare results calculated within the present scheme with a more direct evaluation of the loop diagrams like in ref.\\cite{hu07}.   This paper is organized as follows. In Sec.II, we will depict brief outline of the formalism of FLT. We find the analytic expressions for the FLPs both for direct and exchange contributions in Sec.III. Subsequently, we determine chemical potential, energy density and various other thermodynamic quantities like incompressibility and sound velocity. Sec.IV, is devoted to calculate isovector LPs to which involves the $\\pi$ meson contribution, and used to express the symmetry energy.       \\vskip 0.4in \\section {Formalism}   In FLT total energy $E$ of an interacting system is the functional of occupation number $n_{p}$ of the quasi-particle states of momentum $p$. The excitation of the system is equivalent to the change of occupation number by an amount $\\delta n_{p}$. The corresponding energy of the system is given by \\cite{baym_book,baym76},   \\beq{\\label {total_energy}} E&=&E^{0}+\\sum_{s}\\int\\frac{d^3{p}}{(2\\pi)^3} \\varepsilon_{ps}^{0}\\delta n_{ps} +\\frac{1}{2}\\sum_{ss'}\\int\\frac{d^3{p}}{(2\\pi)^3}\\frac{d^3{p'}}{(2\\pi)^3} f_{ps,p's'} \\delta n_{ps}\\delta n_{p's'}, \\eeq  where $E^0$ is the ground state energy and $s$ is the spin index, and the quasi-particle energy can be written as,  \\beq\\label{quasi_energy} \\veps_{ps}=\\veps_{ps}^{0}+\\sum_{s'}\\int\\frac{d^3{p'}}{(2\\pi)^3}f_{ps,p's'} \\delta n_{p's'}, \\eeq  where $\\veps_{ps}^{0}$ is the non-interacting single particle energy. The interaction between quasi-particles is given by $f_{ps,p's'}$, which is defined to be the second derivative of the energy functional with respect to occupation functions,  \\beq\\label{quasi_interac} f_{ps,p's'}=\\frac{\\delta^{2}E}{\\delta{n}_{ps}~\\delta{n}_{p's'}}. \\eeq   Since quasi-particles are well defined only near the Fermi surface, one assumes  \\beq \\left.\\begin{array}{lll} &\\veps_{p}&=\\mu+v_{f}(p-p_{f})\\\\ {\\rm and~~~} & p&\\simeq p'\\simeq p_{f}. \\end{array} \\right\\} \\eeq  Then LPs $f_{l}$s are defined by the Legendre expansion of $f_{ps,p's'}$ as \\cite{baym_book,baym76},  \\beq\\label{landau_para} f_l=\\frac{2l+1}{4}\\sum_{ss'}\\int\\frac{d\\Omega}{4\\pi}P_{l}(\\cos\\theta)f_{ps,p's'}, \\eeq  where $\\theta$ is the angle between $p$ and $p'$, both taken to be on the Fermi surface, and the integration is over all directions of $p$ \\cite{baym76}. We restrict ourselves for $l\\le 1$ i.e. $f_{0}$ and $f_{1}$, since higher $l$ contribution decreases rapidly.  Now the Landau Fermi liquid interaction $f_{ps,p's'}$ is related to the two particle forward scattering amplitude via \\cite{baym_book,baym76},  \\beq f_{ps,p's'}&=&\\frac{M}{\\veps_{p}^0}\\frac{M}{\\veps_{p'}^0} {\\cal M}_{ps,p's'}, \\eeq  where $M$ is the mass of the nucleon and the Lorentz invariant matrix ${\\cal M}_{ps,p's'}$ consists of the usual direct and exchange amplitude, which may be evaluated directly from the relevant diagrams as shown in Fig.~\\ref{dir_jax} and Fig.~\\ref{ex_jax}.    \\vskip 0.2in  \\begin{figure}[htb] \\begin{center} \\includegraphics[scale=0.65,angle=0]{dir_jax.eps} \\caption{Schematic diagrams of direct contribution to forward scattering amplitude. Nucleons are represented by solid lines. $\\sigma$, $\\omega$ and $\\pi$ mesons are denoted by dotted, wavy and dashed lines respectively.} \\label{dir_jax} \\end{center} \\end{figure}    \\vskip 0.2in  \\begin{figure}[htb] \\begin{center} \\includegraphics[scale=0.65,angle=0]{ex_jax.eps} \\caption{Schematic diagrams of exchange contribution to forward scattering amplitude. Nucleons are represented by solid lines. $\\sigma$, $\\omega$ and $\\pi$ mesons are denoted by dotted, wavy and dashed lines respectively.} \\label{ex_jax} \\end{center} \\end{figure}     The spin averaged scattering amplitude  $(f_{pp'})$ is given by \\cite{baym76},  \\beq\\label{fermi_inter} f_{pp'}&=&\\frac{1}{4}\\sum_{ss'}\\frac{M}{\\veps_{p}^0} \\frac{M}{\\veps_{p'}^0}{\\cal M}_{ps,p's'}. \\eeq  The dimensionless LPs are $F_{l}=N(0)f_{l}$, where $N(0)$ is the density of states at the Fermi surface defined as,  \\beq\\label{dens_of_state} N(0)&=&\\sum_{s}\\int\\frac{\\rm d^3{p}}{(2\\pi)^3} \\delta(\\veps_{ps}-\\mu)\\nn\\\\ &=&\\frac{g_{s}g_{I}p_{f}^2}{2\\pi^2}\\left(\\frac{\\del p}{\\del\\veps_{p}} \\right)_{p=p_{f}}\\nn\\\\ &\\simeq&\\frac{g_{s}g_{I}p_{f}\\veps_{f}}{2\\pi^2}. \\eeq  Here $g_{s},g_{I}$ are the spin and isospin degeneracy factor respectively.  In the above expression $(\\del p/\\del\\veps_{p})_{p=p_{f}}$ is the inverse Fermi velocity $(v_{f}^{-1})$ related to the FL parameter $F_{1}$, \\beq\\label{inv_vel} v_{f}^{-1}=(\\del p/\\del\\veps_{p})_{p=p_{f}}=(\\mu/p_{f})(1+F_{1}/3). \\eeq  With the help of Eq.(\\ref{dens_of_state}) and Eq.(\\ref{inv_vel}) one writes {\\cite{matsui81}}  \\beq\\label{relative} \\veps_{f}=\\mu(1+\\frac{1}{3}F_{1}). \\eeq  To compare Eq.(\\ref{inv_vel}) and Eq.(\\ref{relative}) with the well known non-relativistic expressions one has to put $\\veps_{f}=M^*$ and $\\mu=M$.    \\vskip 0.4in \\section {Chiral Lagrangian and Landau parameters}  We adopt the non-linear chiral model to calculate the FLPs and consequently estimate various quantities of physical interest like effective chemical potential, sound velocity, incompressibility, symmetry energy etc. Here all the fields are treated relativistically \\cite{baym76}. By retaining only the lowest order terms in the pion fields, one obtains the following Lagrangian from the chirally invariant Lagrangian \\cite{furnstahl97,hu07}:  \\beq \\cl &=& \\bar{\\Psi}\\left[\\gm(i\\del_{\\mu}-g_{\\o}\\o_{\\mu})-i\\frac{g_{A}}{f_{\\pi}} \\gm\\gamma_{5}\\del_{\\mu}\\underline{\\pi}-(M-g_{\\s}\\Phi_{\\s})\\right]\\Psi\\nn\\\\&& +\\frac{1}{2}\\del^{\\mu}\\Phi_{\\s}\\del_{\\mu}\\phi_{\\s}-\\frac{1}{2}m_{\\s}^2\\Phi_{\\s}^2 -\\frac{1}{4}\\o^{\\mn}\\o_{\\mn}+\\frac{1}{2}m_{\\o}^2\\o^{\\m}\\o_{\\m} +\\frac{1}{2}\\del^{\\m}{\\vec\\Phi_{\\p}}\\cdot\\del_{\\m}{\\vec\\Phi_{\\p}} -\\frac{1}{2}m_{\\p}^2{\\vec\\Phi_{\\p}^2}\\nn\\\\&& +\\cl_{NL}+\\delta\\cl, \\eeq  where $\\o_{\\mn}=\\del_{\\m}\\o_{\\n}-\\del_{\\n}\\o_{\\m}$, $\\underline{\\p}=\\frac{1}{2}({\\vec\\tau}\\cdot{\\vec\\Phi_{\\p}})$ and ${\\vec\\tau}$ is the isospin index. Here $\\Psi$ is the nucleon field and $\\Phi_{\\s}$, $\\o_{\\m}$ and ${\\vec\\Phi_{\\p}}$ are the meson fields (isoscalar-scalar, isoscalar-vector and isovector-pseudoscalar respectively). The terms $\\delta\\cl$ and $\\cl_{NL}$ contain the non-linear and counterterms respectively (for explicit expression see{\\cite{hu07}}). Note that in this work, the convention of {\\cite {kapusta81}} is used.    \\subsection{Perturbative calculation}  Let us calculate the LPs perturbatively due to the exchange of scalar and vector mesons between the nucleons \\cite{baym76}. The direct contribution ({\\em see } Fig.~\\ref{dir_jax}) is given by \\cite{baym76}  \\beq\\label{dir_int} \\left.\\begin{array}{ll} f_{pp'}^{dir,\\s}&=-\\frac{g_{\\sigma}^2}{m_{\\sigma}^2}\\frac{M^2} {\\veps_{p}^0\\veps_{p'}^0}\\\\ f_{pp'}^{dir,\\o}&=\\frac{g_{\\omega}^2}{m_{\\omega}^2}\\frac{P.P'} {\\veps_{p}^0\\veps_{p'}^0}, \\end{array} \\right\\} \\eeq  where  $\\veps_{p}^0=\\sqrt{p^2+M^2}$. Now with the help of Eq.(\\ref{landau_para}) and Eq.(\\ref{dir_int}), the LPs become  \\beq\\label{f0_dir} \\left.\\begin{array}{ll} f_{0}^{dir,\\s}&=-\\f{g_{\\s}^2}{m_{\\s}^2}\\f{M^{2}}{\\veps_{f}^2}\\\\ f_{0}^{dir,\\o}&=\\f{g_{\\o}^2}{m_{\\o}^2}, \\end{array} \\right\\} \\eeq  and  \\beq\\label{f1_dir} \\left.\\begin{array}{ll} f_{1}^{dir,\\s}&=0\\\\ f_{1}^{dir,\\o}&=-\\f{g_{\\o}^2}{m_{\\o}^2}\\frac{p_{f}^2}{\\veps_{f}^2}. \\end{array} \\right\\} \\eeq  One may neglect the contribution of $f_{1}^{dir,\\o}$ as discussed in ref.\\cite{speth80}. A better approach was developed by Matsui \\cite{matsui81} where the magnetic interaction is included which reduces the value of $f_{1}^{dir,\\o}$.  One may now, for the direct contribution plug in $f_{pp'}^{dir,\\s}$ and $f_{pp'}^{dir,\\o}$ in Eq.(\\ref{total_energy}) and Eq.(\\ref{quasi_energy}) to obtain the energy density and the SPE spectrum, respectively. The SPE spectrum is given by \\cite{baym76}  \\beq\\label{dir_single} \\veps_{p}^{dir}&=&\\veps_{p}^{0}+\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho -\\frac{g_{\\s}^2}{m_{\\s}^2}\\frac{M}{\\veps_{p}^0}n_{s}. \\eeq   Here $\\rho $ and $n_{s}$ are the baryon and scalar density given by  \\beq\\label{baryon_den} \\rho&=&g_{s}g_{I}\\frac{p_{f}^3}{6\\p^2}, \\eeq  and  \\beq\\label{scalar_den} n_{s}&=&g_{s}g_{I}\\frac{M}{4\\p^2} \\left[p_{f}\\veps_{f}-M^2\\ln\\left(\\frac{p_{f}+\\veps_{f}}{M}\\right)\\right]. \\eeq  The energy density for direct contribution is \\cite{baym76}  \\beq\\label{dir_totalE} E^{dir}&=&E^0+\\frac{1}{2}\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho^2 -\\frac{1}{2}\\frac{g_{\\s}^2}{m_{\\s}^2}n_{s}^2. \\eeq  The chemical potential is  \\beq\\label{mu_dir1} \\mu^{dir}&=&\\frac{\\del E^{dir}}{\\del \\rho}\\nn\\\\ &=&\\veps_{f}+\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho-\\frac{g_{\\s}^2}{m_{\\s}^2}n_{s} \\frac{\\del n_{s}}{\\del \\rho}\\nn\\\\ &=&\\veps_{f}+\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho -\\frac{g_{\\s}^2}{m_{\\s}^2}\\frac{M}{\\veps_{f}}n_{s}. \\eeq  One can derive the same result directly from Eq.(\\ref{dir_single}) as $\\mu = \\veps_p{\\Big |_{p=p_f}}$.  \\subsection{FLPs in mean field model}  It is well known that in the MF approximation, one replaces the mesonic fields by their vacuum expectation values viz. $\\sigma \\ra <\\sigma> = \\sigma_0$, $\\omega \\ra <\\omega> = \\delta_{\\mu 0} \\omega^\\mu$. The pion, however, fails to contribute at the MF level as  $ <\\pi> = 0$.  In the MF approximation the energy density can be written as \\cite{vol16}  \\beq\\label{mft_totalE} E^{MFT}&=&\\frac{1}{2}\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho^2 +\\frac{1}{2}\\frac{g_{\\s}^2}{m_{\\s}^2}n_{s}^2+\\sum_{i}n_{i}\\sqrt{p_{i}^2+M^{*2}}. \\eeq  In the above equation $M^*$ denotes the effective nucleon mass to be determined self consistently \\cite{matsui81,hu07}. With the help of Eq.(\\ref{quasi_interac}), the interaction parameter takes the following form \\cite{matsui81}  \\beq\\label{mft_int} f_{pp'}^{MFT}&=&\\frac{g_{\\o}^2}{m_{\\o}^2}-\\frac{g_{\\s}^2}{m_{\\s}^2}\\frac{M^{*2}} {\\veps_{p}^0\\veps_{p'}^0}\\left[1+\\zeta(M^*)\\right]^{-1}. \\eeq  Here $\\veps_{p}^0=\\sqrt{p^2+M^{*2}}$ and  \\beq\\label{brk_trm} \\zeta(M^*)&=&\\frac{g_{\\s}^2}{m_{\\s}^2}{\\sum_{i}}n_{i} \\frac{p_{i}^2}{(p_{i}^2+M^{*2})^{3/2}}\\nn\\\\ &=&\\frac{g_{\\s}^2}{m_{\\s}^2}M^{*2}\\frac{g_{s}g_{I}v_{f}}{2\\p^2} \\left(1+\\frac{1}{2(1-v_{f}^2)}-\\frac{3}{4v_{f}} \\ln\\left\\vert\\frac{1+v_{f}}{1-v_{f}}\\right\\vert\\right), \\eeq  where $v_{f}={p_{f}}/{(p_{f}^2+M^{*2})^{1/2}}$, is the relativistic Fermi velocity. The inverse part of Eq.(\\ref{mft_int}) reduces the magnitude of interaction parameter compared to what is obtained in absence of the MF Eq.(\\ref{dir_int}) \\cite{matsui81,anastasio83}   The LPs as defined in ref.\\cite{matsui81} are  \\beq\\label{f0_mft} \\left.\\begin{array}{ll} f_{0}^{MFT,\\s}&=-\\f{g_{\\s}^2}{m_{\\s}^2}\\f{M^{*2}}{\\veps_{f}^2} \\left[1+\\frac{g_{\\s}^2}{m_{\\s}^2} {\\sum_{i}}n_{i}\\frac{p_{i}^2}{(p_{i}^2+M^{*2})^{3/2}}\\right]^{-1}\\\\ f_{0}^{MFT,\\o}&=\\f{g_{\\o}^2}{m_{\\o}^2}. \\end{array} \\right\\} \\eeq  When we evaluated the above Eq.(\\ref{f0_mft}), we neglect the ``magnetic interaction'' between the quasiparticles which is induced by the microscopic currents. In presence of current density \\cite{matsui81}  \\beq\\label{f1_mft} f_{1}^{MFT,\\o}&=&-\\f{g_{\\o}^2}{m_{\\o}^2}\\f{p_{f}^2}{\\veps_{f}^2} \\left[1+\\frac{g_{\\o}^2}{m_{\\o}^2} {\\sum_{i}}n_{i}\\frac{\\f{2}{3}p_{i}^2+M^{*2}}{(p_{i}^2+M^{*2})^{3/2}}\\right]^{-1}. \\eeq  Clearly, the current contribution reduces the value of $f_{1}^{MFT,\\o}$. Previously we showed in Eq.(\\ref{dir_single}) and Eq.(\\ref{dir_totalE}) the SPE spectrum and energy density in absence of MF, but in presence of MF, SPE is given by \\cite{matsui81,chin76,vol16,chin74}  \\beq\\label{mft_single} \\veps_{p}^{MFT}&=&\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho+\\sqrt{p^2+M^{*2}}. \\eeq  Therefore,  \\beq\\label{mu_mft} \\mu^{MFT}&=&\\frac{g_{\\o}^2}{m_{\\o}^2}\\rho+\\sqrt{p_{f}^2+M^{*2}}. \\eeq  In the low density limit, Eq.(\\ref{mft_single}) reduces to Eq.(\\ref{dir_single}) as $M^*= M-\\frac{g_{\\s}^2}{m_{\\s}^2} {\\sum_{i}}n_{i}\\frac{M}{(p_{i}^2+M^{2})^{1/2}}$. It is to be noted that in the MF approximation scalar meson contribution is absorbed in the effective mass does not appear explicitly as in Eq.(\\ref{dir_single}). Another interesting difference is also noticed in the expressions for the total energy densities given by Eqs.(\\ref{dir_totalE}) and (\\ref{mft_totalE}). Note that, our MF result is consistent with ref.\\cite{matsui81} but differs with that of ref.\\cite{chin74}.    \\begin{figure}[htb] \\vskip 0.2in \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{mu_mat_chin.eps}} \\caption{Chemical potential for direct contribution with $\\sigma$ and $\\omega$ meson exchange in symmetric nuclear matter. The dashed and solid curve represent the perturbative and MF results, respectively.} \\label{fig1} \\end{center} \\end{figure}      \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{e_mf_per.eps}} \\caption{ Total energy from direct contribution with $\\sigma$ and $\\omega$ meson exchange in symmetric nuclear matter. The dashed and solid curve represent the perturbative and MF results, respectively.} \\label{fig1a} \\end{center} \\end{figure}    In Fig(\\ref{fig1}) we present the comparative study of the chemical potential obtained perturbatively and with MF calculation. At low density they tend to merge, while at higher density MF results differ significantly from the perturbative result. Numerically $\\mu^{MFT}$ and $\\mu^{per}$ are given by 861.07 MeV and 832.64 MeV respectively at normal matter density ($\\rho_{0}=0.148{\\rm fm^{-3}}$). For the numerical estimate we adopt the coupling parameter set as designated by ${\\bf M0A}$ in ref.\\cite{hu07}. In Fig(\\ref{fig1a}) we compare the results for total energy obtained from perturbative and MF calculation. This also shows at low density they tend to merge, while at higher density MF results become larger than the  perturbative results. This is easily understood from Eq.(\\ref{dir_totalE}) and (\\ref{mft_totalE}). At saturation density numerical values are given by 886.43 MeV and 893.31 MeV for perturbative and MF calculation respectively.  Now we consider the exchange modification over the MF. Evaluating the exchange diagrams (Fig.~\\ref{ex_jax}), we obtain the interaction parameter as \\cite{baym76}  \\beq \\left.\\begin{array}{ll} f_{pp'}^{ex,\\s}&=-\\frac{g_{\\sigma}^2}{4\\veps_{p}^0\\veps_{p'}^0}\\frac{P.P'+M^{*2}} {(P-P')^2-m_{\\sigma}^2}\\\\ f_{pp'}^{ex,\\o}&=-\\frac{g_{\\omega}^2}{2\\veps_{p}^0\\veps_{p'}^0}\\frac{P.P'-2M^{*2}} {(P-P')^2-m_{\\omega}^2}. \\end{array} \\right\\} \\eeq  With the help of Eq.(\\ref{landau_para}), the LPs for scalar meson exchange reads as  \\beq\\label{f0sig} f_{0}^{ex,\\sigma}&=&\\frac{g_{\\s}^2}{8\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)+2M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\s}^2} {\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\frac{g_{\\sigma}^2}{8\\veps_{f}^2}\\left[1-\\left( \\frac{m_{\\sigma}^2-4M^{*2}}{4p_{f}^2}\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\sigma}^2} \\right)\\right], \\eeq  and  \\beq\\label{f1sig} \\frac{1}{3}f_{1}^{ex,\\sigma}&=&\\frac{g_{\\s}^2}{8\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)+2M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\s}^2}(\\cos\\theta) {\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\frac{g_{\\sigma}^2}{8\\veps_{f}^2}\\left[\\left(\\frac {m_{\\sigma}^2-4M^{*2}}{2p_{f}^2}\\right)\\left\\{1-\\left(\\frac{m_{\\sigma}^2+2p_{f}^2} {4p_{f}^2}\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\sigma}^2}\\right)\\right\\}\\right]. \\eeq  Using Eqs.(\\ref{f0sig}) and (\\ref{f1sig}) we have  \\beq\\label{f0f1sig} f_{0}^{ex,\\sigma}-\\frac{1}{3}f_{1}^{ex,\\sigma}&=& \\frac{g_{\\s}^2}{8\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)+2M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\s}^2} (1-\\cos\\theta){\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\frac{g_{\\sigma}^2}{8\\veps_{f}^2} \\left[1-\\left(\\frac{m_{\\sigma}^2-4M^{*2}}{2p_{f}^2}\\right)\\left\\{1- \\frac{m_{\\sigma}^2}{4p_{f}^2}\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\sigma}^2}\\right) \\right\\}\\right]. \\eeq  It is this combination {\\em i.e.} $f_{0}-\\frac{1}{3}f_{1}$, which appears in the calculation of chemical potential and other relevant quantities. For massless scalar meson interaction, the above Eq.(\\ref{f0f1sig}) turns out to be finite,  \\beq\\label{f0f1sig1} \\left(f_{0}^{ex,\\sigma}-\\frac{1}{3}f_{1}^{ex,\\sigma}\\right)_{m_{\\s}\\ra 0} &=&\\frac{g_{\\s}^2}{8p_{f}^2}\\left(1+\\frac{M^{*2}}{\\veps_{f}^2}\\right). \\eeq  Note that in the limit $m_{\\s}\\ra 0$ both $f_{0}$ and $f_{1}$  are individually diverge because of the presence of $(1-\\cos\\theta)$ term in the denominator. In the massless limit such divergences are also contained in Eq.(\\ref{f0sig}) and Eq.(\\ref{f1sig}).  The dimensionless LPs $F_{0}$ and $F_{1}$ are defined as $F_{0}=N(0)f_{0}$ and $F_{1}=N(0)f_{1}$, where $N(0)$ is the density of states at the Fermi surface defined in Eq.(\\ref{dens_of_state}).  \\beq F_{0}^{ex,\\s}&=&g_{s}g_{I}\\frac{g_{\\sigma}^2p_{f}}{16\\pi^2\\veps_{f}}\\left[1-\\left( \\frac{m_{\\sigma}^2-4M^{*2}}{4p_{f}^2}\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\sigma}^2} \\right)\\right], \\eeq  and  \\beq \\f{1}{3}F_{1}^{ex,\\s}&=&g_{s}g_{I}\\frac{g_{\\sigma}^2p_{f}}{16\\pi^2\\veps_{f}} \\left[\\left(\\frac {m_{\\sigma}^2-4M^{*2}}{2p_{f}^2}\\right)\\left\\{1-\\left(\\frac{m_{\\sigma}^2+2p_{f}^2} {4p_{f}^2}\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\sigma}^2}\\right)\\right\\}\\right]. \\eeq  Similarly, for vector meson exchange we have  \\beq\\label{f0omega} f_{0}^{ex,\\omega}&=&\\frac{g_{\\omega}^2}{4\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)-M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\o}^2} {\\rm d(\\cos\\theta)}\\nn\\\\ &=&-\\frac{g_{\\omega}^2}{8\\veps_{f}^2}\\left[-2+ \\frac{(m_{\\omega}^2+2M^{*2})}{2p_{f}^2}\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\omega}^2} \\right)\\right], \\eeq  and  \\beq\\label{f1omega} \\frac{1}{3}f_{1}^{ex,\\omega}&=&\\frac{g_{\\omega}^2}{4\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)-M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\o}^2} (\\cos\\theta){\\rm d(\\cos\\theta)}\\nn\\\\ &=&-\\frac{g_{\\omega}^2(m_{\\omega}^2+2M^{*2})} {16\\veps_{f}^2p_{f}^2}\\left[-2+\\left(1+\\frac{m_{\\omega}^2}{2p_{f}^2}\\right) \\ln\\left(1+\\frac{4p_{f}^2}{m_{\\omega}^2}\\right)\\right]. \\eeq  Using Eq.(\\ref{f0omega}) and Eq.(\\ref{f1omega}) we have  \\beq\\label{f0f1ome} f_{0}^{ex,\\omega}-\\frac{1}{3}f_{1}^{ex,\\omega}&=& \\frac{g_{\\omega}^2}{4\\veps_{f}^2}\\int_{-1}^{1} \\frac{p_{f}^2(1-\\cos\\theta)-M^{*2}}{2p_{f}^2(1-\\cos\\theta)+m_{\\o}^2} (1-\\cos\\theta){\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\frac{g_{\\omega}^2}{4\\veps_{f}^2} \\left[1+\\frac{(m_{\\omega}^2+2M^{*2})}{4p_{f}^2}\\left\\{-2+\\frac{m_{\\omega}^2} {2p_{f}^2}\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\omega}^2}\\right)\\right\\}\\right]. \\eeq  In the limit $m_{\\o}\\ra 0$ the above Eq.(\\ref{f0f1ome}) turns into  \\beq\\label{f0f1ome1} \\left(f_{0}^{ex,\\o}-\\frac{1}{3}f_{1}^{ex,\\o}\\right)_{m_{\\o}\\ra 0} &=&\\frac{g_{\\o}^2}{4p_{f}^2}\\left(1-\\frac{2M^{*2}}{\\veps_{f}}\\right). \\eeq  This expression agrees with the previous calculation by Baym and Chin \\cite{baym76} who arrived at this result by direct evaluation of the integral by putting $m_{\\o}=0$ in Eq.(\\ref{f0f1ome}). Here also to be noted, in the limit $m_{\\o}\\ra 0$ both $f_{0}$ and $f_{1}$ are individually divergent, but the combination $f_{0}-\\frac{1}{3}f_{1}$ is finite as observed in the case of scalar ($\\s$) meson exchange.  The dimensionless LPs for vector meson exchange reads:  \\beq F_{0}^{ex,\\o}&=&-g_{s}g_{I}\\frac{g_{\\omega}^2p_{f}}{16\\pi^2\\veps_{f}}\\left[-2+ \\frac{(m_{\\omega}^2+2M^{*2})}{2p_{f}^2}\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\omega}^2} \\right)\\right], \\eeq  and  \\beq \\f{1}{3}F_{1}^{ex,\\o}&=&-g_{s}g_{I}\\frac{g_{\\omega}^2(m_{\\omega}^2+2M^{*2})} {32\\pi^2p_{f}\\veps_{f}}\\left[-2+\\left(1+\\frac{m_{\\omega}^2}{2p_{f}^2}\\right) \\ln\\left(1+\\frac{4p_{f}^2}{m_{\\omega}^2}\\right)\\right]. \\eeq   \\vskip 0.2in \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F_sca.eps}} \\caption{Dimensionless LPs in symmetric nuclear matter for $\\sigma$ meson exchange in relativistic theory. $F_{0}^{MF}$, $F_{0}^{ex}$, $F_{0}^{tot}$ and $F_{1}^{ex}$ are denoted by dot-dashed, dashed, solid and dotted line respectively.} \\label{fig2} \\end{center} \\end{figure}     \\begin{table} \\caption{Parameter sets used in this work.} \\label{table-1} \\begin{tabular}{ccccc} \\hline \\hline Meson    &~~~~    & Mass               &~~~~ & Coupling \\\\  \\hline $\\s$    &~~~~    & $m_{\\s}/M$=0.54     &~~~~ & $g_{\\s}/{4\\p}$=0.7936\\\\ $\\o$    &~~~~    & $m_{\\o}/M$=0.8328   &~~~~ & $g_{\\o}/{4\\p}$=0.9681\\\\ \\hline \\hline \\end{tabular} \\end{table}   \\vskip 0.2in \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F_vec.eps}} \\caption{Dimensionless LPs for symmetric nuclear matter for $\\omega$ meson exchange in relativistic theory. $F_{0}^{MF}$, $F_{0}^{ex}$, $F_{0}^{tot}$ and $F_{1}^{ex}$ are denoted by dashed, dot-dashed, solid and dotted line respectively.} \\label{fig3} \\end{center} \\end{figure}     \\begin{table} \\caption{Dimensionless LPs for $\\s$ and $\\o$ exchange at $\\rho=\\rho_{0}$.} \\label{table-2} \\begin{tabular}{ccccccccc}\\hline\\hline Meson &~~~~&$F_{0}^{MF}$&~~~~&$F_{0}^{ex}$&~~~~&$F_{0}^{tot}$&~~~~&$F_{1}^{ex}$ \\\\ \\hline $\\s$  &~~~~& -8.65 &~~~~ &  3.61 &~~~~ &  -5.04 &~~~~  &  0.875   \\\\ $\\o$  &~~~~&  7.35 &~~~~ & -1.91 &~~~~ &   5.44 &~~~~  &  -0.932   \\\\ \\hline \\hline \\end{tabular} \\end{table}      In Fig.(\\ref{fig2}) and Fig.(\\ref{fig3}) we present $F_{0}^{MF}$, $F_{0}^{ex}$, $F_{0}^{tot}$ and $F_{1}^{ex}$ as a function of baryon density for symmetric nuclear matter due to $\\sigma$ and $\\omega$ meson interaction respectively. It is to be noted, that $F_{0}^{MF}$ and $F_{0}^{ex}$ contribute in opposite sign for both $\\s$ and $\\o$ meson exchange as it is seen from Table~(\\ref{table-2}). We quote few numerical values of $F^{\\s}$ and $F^{\\o}$ in Table~(\\ref{table-2}) at normal matter density ($\\rho_{0}=0.148 {\\rm fm^{-3}}$). It is to be noted that the numerical estimation of $F_0$ with MF in our case, differs from ref.\\cite{matsui81}. This is due to different coupling parameters in these two models.   We now proceed to calculate chemical potential due to the exchange terms denoted by $\\mu^{ex}$. As in ref.{\\cite{baym76}} we have  \\beq\\label{delmu_deln} \\frac{\\del\\mu}{\\del \\rho}&=&\\frac{2\\pi^2}{g_{s}g_{I}p_{f}^2} \\left(\\frac{\\del\\veps_{p}}{\\del p}\\right)_{p=p_{f}}+f_{0} \\nn \\\\ &=&\\frac{2\\pi^2}{\\mu g_{s}g_{I}p_{f}}-\\frac{1}{3}f_{1}+f_{0}, \\eeq  and  \\beq\\label{fermi_vel} \\left(\\frac{\\del\\veps_{p}}{\\del p}\\right)_{p=p_{f}}=v_{f}&=& \\frac{p_{f}}{\\mu}-\\frac{g_{s}g_{I}p_{f}^2}{2\\pi^2}\\frac{f_{1}}{3}. \\eeq  Now from Eq.(\\ref{delmu_deln}) one gets  \\beq\\label{mudmu} \\mu{\\rm d\\mu}&=&\\left[p_{f}+g_{s}g_{I}\\frac{\\mu p_{f}^2}{2\\p^2} (f_{0}-\\frac{1}{3}f_{1})\\right]dp_{f}. \\eeq  To calculate $\\mu$, it is sufficient to let $\\mu=\\veps_{f}$ in the right hand side of Eq.(\\ref{mudmu}). With the constant of integration adjusted so that at high density $ p_{f}\\simeq \\veps_{f}$, Eq.(\\ref{mudmu}) upon integration together with Eq.(\\ref{f0f1sig}) yield  \\beq\\label{mu_sigma} \\mu^{ex}_{\\sigma}&=&\\veps_{f}+\\frac{g_{s}g_{I}g_{\\sigma}^2}{128\\pi^2\\veps_{f}}M^{*2} \\left[-2y_{\\sigma}(4-y_{\\sigma}^2)^{3/2}\\tan^{-1}\\left (\\frac{x\\sqrt{4-y_{\\sigma}^2}}{y_{\\sigma}\\sqrt{1+x^2}}\\right)\\right.\\nn\\\\&&\\left. ~~~~~~~~~~~~~~~~~~~~~~~~~~~+\\frac{y_{\\sigma}^2(4-y_{\\sigma}^2)\\sqrt{1+x^2}}{x} \\ln\\left(1+\\frac{4x^2} {y_{\\sigma}^2}\\right)+4x\\sqrt{1+x^2} \\right.\\nn\\\\&&\\left. ~~~~~~~~~~~~~~~~~~~~~~~~~+2(y_{\\sigma}^4-6y_{\\sigma}^2+6)\\ln(x+\\sqrt{1+x^2}) \\right], \\eeq  where $x=p_{f}/M^*$ and $y_{\\s}=m_{\\s}/M^*$.  \\vskip 0.2in  \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{muex_svp.eps}} \\caption{Density dependent of exchange chemical potential in symmetric nuclear matter. $\\s$,$\\o$ and $\\p$ mesons are denoted by solid, dotted and dashed line respectively.} \\label{fig4} \\end{center} \\end{figure}  For massless meson limit {\\em i.e.} at $m_{\\s}\\ra 0$ implies $y_{\\s}\\ra 0$ we have  \\beq \\mu^{ex}_{\\s}\\left\\vert\\right._{m_{\\s}\\ra 0}&=&\\veps_{f}+g_{s}g_{I} \\frac{g_{\\s}^2M^{*2}}{32\\p^2\\veps_{f}}\\left[x\\sqrt{1+x^2}+3\\ln(x+\\sqrt{1+x^2}) \\right]. \\eeq  Similarly for vector meson interaction, using Eq.(\\ref{f0f1ome}) and Eq.(\\ref{mudmu}) one obtains  \\beq \\mu_{\\omega}^{ex}&=&\\veps_{f}-\\frac{g_{s}g_{I}g_{\\omega}^2M^{*2}}{64\\pi^2\\veps_{f}} \\left[ \\frac{2y_{\\omega}(y_{\\omega}^4-2y_{\\omega}^2-8)}{\\sqrt{-y_{\\omega}^2+4}}\\tan^{-1} \\left(\\frac{x\\sqrt{-y_{\\omega}^2+4}}{y_{\\omega}\\sqrt{1+x^2}}\\right)\\right.\\nn\\\\ &&\\left.~~~~~~~~~~~~~~~~~~~~~~~+\\frac{y_{\\omega}^2(y_{\\omega}^2+2)\\sqrt{1+x^2}}{x} \\ln\\left(1+\\frac{4x^2}{y_{\\omega}^2}\\right)-4x\\sqrt{1+x^2}\\right.\\nn\\\\&&\\left. ~~~~~~~~~~~~~~~~~~~~~~~~+2(6-y_{\\omega}^4)\\ln(x+\\sqrt{1+x^2})\\right]. \\eeq  Here $y_{\\o}=m_{\\o}/M^*$. For massless limit of vector meson the expression for chemical potential reads as  \\beq \\mu^{ex}_{\\o}&=&\\veps_{f}+g_{s}g_{I}\\frac{g_{\\o}^2M^{*2}}{16\\p^2\\veps_{f}} \\left[x\\sqrt{1+x^2}-3\\ln(x+\\sqrt{1+x^2})\\right]. \\eeq  In low density limit ($M^*\\ra M$) for the massless meson exchange we reproduce the expression derived earlier {\\cite{baym76,chin77}}.   \\vskip 0.2in \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{mu_direx_svp.eps}} \\caption{Comparison of mean field and exchange results of chemical potential in symmetric nuclear matter. Direct contributions are plotted by solid curve and exchange contributions by dashed curve.} \\label{fig5} \\end{center} \\end{figure}     \\vskip 0.2in  \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{mu_HF_MFT.eps}} \\caption{Comparison of chemical potential with MFT and HF case in symmetric nuclear matter. MF and HF results denoted by dashed and solid line respectively.} \\label{fig11} \\end{center} \\end{figure}     Once the $\\mu^{ex}$ is determined, one can readily calculate its contribution to the energy density. For scalar meson interaction it is given by {\\cite{baym76,chin77}},  \\beq\\label{xe_sig} E^{ex}_{\\s}&=&\\frac{1}{2}\\int{\\rm d\\rho}(\\mu^{ex}_{\\s}-\\veps_{f})\\nn\\\\ &=&g_{s}^2g_{I}^2\\frac{g_{\\s}^2M^{*4}}{512\\p^4} \\left[y_{\\s}^2(4-y_{\\s}^2)\\left\\{-\\frac{x^2}{2}+\\frac{y_{\\s}^2+4x^2}{8} \\ln\\left(1+\\f{4x^2}{y_{\\s}^2}\\right)\\right\\}+x^4\\right.\\nn\\\\ &&\\left.~~~~~~~~~~~~~~~~~ -\\f{1}{2}(y_{\\s}^4-6y_{\\s}^2+6)(\\{x\\sqrt{1+x^2}-\\ln(x+\\sqrt{1+x^2})\\}^2-x^4) +I_{\\s}\\right], \\eeq  where  \\beq I_{\\s}&=&-2y_{\\sigma}(4-y_{\\sigma}^2)^{3/2}\\int\\f{x^2}{\\sqrt{1+x^2}} \\tan^{-1}\\left(\\frac{x\\sqrt{4-y_{\\sigma}^2}}{y_{\\sigma}\\sqrt{1+x^2}}\\right) {\\rm d}x. \\eeq  Similarly, for vector meson exchange we obtain  \\beq\\label{xe_ome} E^{ex}_{\\o}&=&\\frac{1}{2}\\int{\\rm d\\rho}(\\mu^{ex}_{\\o}-\\veps_{f})\\nn\\\\ &=&-g_{s}^2g_{I}^2\\frac{g_{\\o}^2M^{*4}}{256\\p^4} \\left[y_{\\o}^2(2+y_{\\o}^2)\\left\\{-\\frac{x^2}{2}+\\frac{y_{\\o}^2+4x^2}{8} \\ln\\left(1+\\f{4x^2}{y_{\\o}^2}\\right)\\right\\}-x^4\\right.\\nn\\\\ &&\\left.~~~~~~~~~~~~~~~~~ +\\left(\\f{y_{\\o}^4}{2}-3\\right)\\left(\\{x\\sqrt{1+x^2}-\\ln(x+\\sqrt{1+x^2})\\}^2-x^4 \\right)+I_{\\o}\\right], \\eeq  where  \\beq I_{\\o}&=&\\frac{2y_{\\o}(y_{\\o}^4-2y_{\\o}^2-8)}{\\sqrt{4-y_{\\o}^2}} \\int\\f{x^2}{\\sqrt{1+x^2}} \\tan^{-1}\\left(\\frac{x\\sqrt{4-y_{\\o}^2}}{y_{\\o}\\sqrt{1+x^2}}\\right) {\\rm d}x. \\eeq  Thus for the case of massless mesons {\\cite{baym76,chin77}}, we have  \\beq\\label{sig_eng_mless} E^{ex}_{\\s}|_{m_{\\s}\\ra 0}&=&\\f{g_{\\s}^2M^{*4}}{8\\p^4} [x^4-\\f{3}{4}\\{x\\sqrt{1+x^2}-\\ln(x+\\sqrt{1+x^2})\\}^2], \\eeq  and  \\beq\\label{omg_eng_mless} E^{ex}_{\\o}|_{m_{\\o}\\ra 0}&=&-\\f{g_{\\o}^2M^{*4}}{8\\p^4} [x^4-\\f{3}{2}\\{x\\sqrt{1+x^2}-\\ln(x+\\sqrt{1+x^2})\\}^2]. \\eeq  Due to presence of pion fields in the chiral Lagrangian we have component in the interaction which acts on the isospin fluctuation. One can derive the quasiparticle interaction with isospin dependency by the same procedure as for $\\sigma$ and $\\omega$ meson. Pion, being a pseudoscalar, fails to contribute at the MF level forcing us to go beyond MFT so as to include pionic contribution to the FLPs. It is to be noted that, in exchange diagram pion have both isoscalar and isovector contribution to FLPs. Detailed calculation for isoscalar contribution to FLPs is similar as $\\sigma$ and $\\omega$ meson. For brevity, we present only dimensionless FLPs and their contribution to energy density. We also quote their numerical values.   The dimensionless LPs due to $\\pi$ exchange are  \\beq\\label{dless_F0_pi} F^{ex,\\p}_{0}=-g_{s}g_{I}\\f{3g_{A}^2p_{f}M^{*2}}{64\\p^2f_{\\p}^2\\veps_{f}} \\lt[-2+\\frac{m_{\\pi}^2}{2p_{f}^2} \\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\rt], \\eeq  and  \\beq\\label{dless_F1_pi} \\f{1}{3}F^{ex,\\p}_{1}=-g_{s}g_{I}\\f{3g_{A}^2m_{\\p}^2M^{*2}} {128\\p^2f_{\\p}^2p_{f}\\veps_{f}} \\lt[-2+\\left(\\frac{m_{\\pi}^2} {2p_{f}^2}+1\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\rt]. \\eeq       \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F_pi.eps}} \\caption{Dimensionless LPs in symmetric nuclear matter for pion exchange in relativistic theory. Solid and dashed line represent $F_{0}^{\\pi}$ and $F_{1}^{\\pi}$ respectively.} \\label{fig8} \\end{center} \\end{figure}     \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F0_svp.eps}} \\caption{Dimensionless relativistic LP $F_{0}$ in symmetric nuclear matter. $\\s$, $\\o$, $\\p$ and total contribution are denoted by dashed, dot-dashed, dotted and solid line respectively.} \\label{fig9} \\end{center} \\end{figure}     \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F1_svp.eps}} \\caption{Dimensionless relativistic LP $F_{1}$ in symmetric nuclear matter. $\\s$, $\\o$, $\\p$ and total contribution are denoted by dashed, dot-dashed, dotted and solid line respectively.} \\label{fig10} \\end{center} \\end{figure}   In the non-relativistic limit $\\veps_{f}\\ra M^*$, one obtains the same expression of $F_{1}^{\\pi}$ as reported in \\cite{friman96,rho80}. In Fig.(\\ref{fig8}) we show the density dependence of $F_{0}$ and $F_{1}$ due to pionic interaction. Numerically at nuclear saturation density $(\\rho_{0}=0.148 {\\rm fm^{-3}})$, $F^{ex,\\p}_{0}=0.68$ and $F^{ex,\\p}_{1}=-0.2$.  In Fig.(\\ref{fig9}) and Fig.(\\ref{fig10}) we plot separate and total contribution of $F_0$ and $F_1$ due to $\\s$, $\\o$ and $\\p$ exchange respectively. Interestingly, individual contribution to LPs of $\\s$ and $\\o$ meson are large while sum of their contribution to $F_{0}^{tot}$ is small due to the sensitive cancellation of $F_{0}^{\\s}$ and $F_{0}^{\\o}$ as can be seen from Fig.(\\ref{fig9}). Such a cancellation is responsible for the nuclear saturation dynamics  \\cite{celenza_book,anastasio83}. Numerically, $F_{0}^{\\s+\\o}$ is approximately $3/2$ times smaller than $F_{0}^\\p$ as can be seen from Table (\\ref{table-3}).        \\begin{figure}[!tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{xe.eps}} \\caption{Comparison of separate and total exchange energy obtained from FLT in symmetric nuclear matter.$\\s$, $\\o$, $\\p$ and total contribution are denoted by dashed, dot-dashed, dotted and solid line respectively.} \\label{fig12} \\end{center} \\end{figure}     \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{dir_xe.eps}} \\caption{Comparison of mean field energy and exchange energy obtained from FLT in symmetric nuclear matter. $E_{\\s}^{MF}$, $E_{\\o}^{MF}$, $E_{\\s}^{ex}$, $E_{\\o}^{ex}$ and $E_{\\p}^{ex}$ are denoted by solid, dashed, dot-dashed, dash-dashed and dotted line respectively.} \\label{fig13} \\end{center} \\end{figure}     The exchange energy density is given by  \\beq\\label{xe_pi} E^{ex}_{\\p}&=&-g_{s}^2\\frac{3g_{A}^2M^{*6}}{128f_{\\p}^2\\p^4} \\left[I_{\\p}+y_{\\p}^4\\left\\{-\\frac{x^2}{2}+\\frac{y_{\\p}^2+4x^2}{8} \\ln\\left(1+\\f{4x^2}{y_{\\p}^2}\\right)\\right\\}-x^4\\right.\\nn\\\\ &&\\left. +\\left(\\f{y_{\\p}^4}{2}-y_{\\p}^2-1\\right) \\left(\\{x\\eta-\\ln(x+\\eta)\\}^2-x^4\\right)\\right], \\eeq  where $y_{\\p}=m_{\\p}/M^*$ and  \\beq I_{\\p}=-2y_{\\p}^3\\sqrt{4-y_{\\p}^2} \\int\\f{x^2}{\\eta} \\tan^{-1}\\left(\\frac{x\\sqrt{4-y_{\\p}^2}}{y_{\\p}\\eta}\\right) {\\rm d}x. \\eeq  For the massless pion this reads as  \\beq\\label{pi_eng_mless} E^{ex}_{\\p}{\\Big\\vert}_{m_{\\p}\\ra 0}&=&\\f{3g_{A}^2M^{*6}}{32f_{\\p}^2\\p^4} [x\\eta-\\ln(x+\\eta)]^2. \\eeq  In Fig.(\\ref{fig12}) and Fig.(\\ref{fig13}) we show the density dependence of energy due $\\s$, $\\o$ and $\\p$ meson exchanges. Numerical values are quoted in Table(\\ref{table-5}).  It might be mentioned here that in the massless meson limit, Eq.(\\ref{sig_eng_mless}), (\\ref{omg_eng_mless}) and (\\ref{pi_eng_mless}) can be evaluated analytically from two loop ring diagrams of ref.\\cite{hu07} using Eqs.(54), (55) and (56). We have checked and the expression for the energies are found to be consistent with each other. With massive meson results are compared numerically. Numerical estimation of exchange energy from loop diagram and RFLT are found to be few percent limit.     \\begin{table} \\caption{Dimensionless Landau parameters and chemical potential at $\\rho=\\rho_{0}$. Note that, $F_{0}$, $F_{1}$ are the dimensionless isoscalar LPs.} \\label{table-3} \\begin{center} \\begin{tabular}{ccccccc} \\hline \\hline Meson   & $F_{0}$ & ~~~~~~ & $F_{1}$  &  ~~~~~~  & $\\mu^{ex}(MeV)$\\\\ \\hline $\\s$   &  -5.04  & ~~~~~~ & 0.875    &  ~~~~~~  & 731.89 \\\\ $\\o$   &   5.44  & ~~~~~~ & -0.93    &  ~~~~~~  & 501.82 \\\\ $\\p$   &   0.68   & ~~~~~~ & -0.20   &  ~~~~~~  & 609.88  \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}       \\begin{table} \\caption{Chemical potential in MeV from FLT at $\\rho=\\rho_{0}$.} \\label{table-4} \\begin{tabular}{ccccccc} \\hline \\hline Meson    &~~~~   & $\\mu^{ex}$ &~~~~  & $\\mu^{MF}$ &~~~~ &$\\mu^{HF}$\\\\ \\hline $\\s$    &~~~~   &  731.89   &~~~~   &    -      &~~~~ &    -     \\\\ $\\o$    &~~~~   &  501.82   &~~~~   &    -      &~~~~ &    -      \\\\ $\\p$    &~~~~   &  609.88   &~~~~   &    -      &~~~~ &    -      \\\\ $\\s+\\o+\\p$ &~~~~ & 675.63   &~~~~   &  861.07   &~~~~ &  952.73    \\\\ \\hline\\hline \\end{tabular} \\end{table}    Finally we reproduce saturation property of nuclear matter {\\em i.e.} $E/{\\rho}-M=-16.12$ MeV at $p_{f}=1.3 {\\rm fm^{-1}}$ with those energy calculated from RFLPs.    \\begin{table} \\caption{MF and Exchange energy in MeV from FLT at $\\rho=\\rho_{0}$.} \\label{table-5} \\begin{tabular}{ccccc} \\hline \\hline Meson    &~~~~    & $E^{MF}$      &~~~~    & $E^{ex}$    \\\\  \\hline $\\s$    &~~~~    &  193.86      &~~~~    &  40.48      \\\\ $\\o$    &~~~~    &  138.39      &~~~~    & -23.41       \\\\ $\\p$    &~~~~    &    -         &~~~~    &  12.49       \\\\    \\hline\\hline \\end{tabular} \\end{table}      \\begin{figure}[ht] \\vskip 0.15in \\begin{center} \\resizebox{8.0cm}{6.0cm}{\\includegraphics[]{be.eps}} \\caption{Binding energy graph from FLT for symmetric nuclear matter.} \\label{fig15} \\end{center} \\end{figure}       \\vskip 0.4in \\subsection{Incompressibility and First Sound Velocity}   In nuclear matter several important relationships exist between nuclear observables and the FLPs. The thermodynamical parameters can be expressed in terms of few LPs. For example we present the incompressibility ($K$) and first sound velocity ($c_{1}$)\\cite{holt07,matsui81}.  \\vskip 0.2in \\begin{figure}[htb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{K_svp.eps}} \\caption{Compressibility $K$ in symmetric nuclear matter. } \\label{fig6} \\end{center} \\end{figure}   Incompressibility of the Fermi liquid may be derived as in the non-relativistic theory by the second derivative of energy density ($E$) with respect to the number density $\\rho$ {\\cite{holt07,matsui81}};  \\beq K&\\equiv&\\rho\\frac{\\del^2 E}{\\del \\rho^2}\\nn\\\\ &=&\\rho\\frac{\\del\\mu}{\\del \\rho}. \\eeq  If energy density $E$ is given in terms of number density $\\rho$, then the expression for incompressibility or compression modulus in terms of LPs is given by,  \\beq\\label{incompressibility} K&=&\\frac{3p_{f}^2}{\\veps_{f}}(1+F_{0}). \\eeq  Now consider the effect of quasiparticle collision on the collective modes of a neutral Fermi liquid. Suppose the frequency of the mode is $\\omega$, while the quasiparticle collision frequency is $\\nu$. For the limit $\\omega<<\\nu$, many quasiparticle collision takes place during time interval $\\omega^{-1}$. Then the region is collision-dominated, or {\\it  hydrodynamic regime} {\\cite{matsui81}}. Under this circumstances, organized density fluctuation is possible and  hydrodynamic or first sound waves are generated. Hydrodynamic sound propagates only when the system attains the local thermodynamic equilibrium in a time much shorter than the time interval of the sound oscillation.    \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{c1_svp.eps}} \\caption{First sound velocity $c_{1}$ in symmetric nuclear matter. } \\label{fig7} \\end{center} \\end{figure}   The first sound velocity is given by {\\cite{baym76}}  \\beq\\label{first_sound} c_{1}^2~=~\\frac{\\del P}{\\del E}&=&\\frac{\\del P}{\\del \\mu}\\frac{\\del \\mu} {\\del \\rho}\\frac{\\del \\rho}{\\del E}\\nn\\\\ &=&\\frac{\\rho}{\\mu}\\frac{\\del \\mu}{\\del \\rho}, \\eeq  With the help of Eq.(\\ref{delmu_deln}), Eq.(\\ref{first_sound}) yields  \\beq\\label{first_sound_1} c_{1}^2&=&\\frac{1}{3}\\frac{p_{f}^2}{\\mu^2}\\frac{1+F_{0}}{1+\\frac{1}{3}F_{1}}\\nn\\\\ &=&\\frac{1}{3}\\frac{p_{f}^2}{\\mu^2}\\left[1+\\frac{g_{s}g_{I}\\mu p_{f}}{2\\pi^2} (f_{0}-\\frac{1}{3}f_{1})\\right]. \\eeq  Corresponding values of the incompressibility and the first sound velocity are plotted in Fig.(\\ref{fig6}) and Fig.(\\ref{fig7}) separately with $\\s+\\o$ and $\\s+\\o+\\p$ contribution. It is observed that for combined $\\s$ and $\\o$ meson at $\\rho<0.75\\rho_{0}$, $F_{0}<-1$ and the resulting compressibility turns out to be negative. While for $\\pi$ meson together with $\\s$ and $\\o$ meson the same conclusion can be drawn at $\\rho<0.55\\rho_{0}$ . This is the region where the attractive interaction due to the exchange of scalar mesons overwhelms the repulsive force coming from vector meson exchange, and consequently the system becomes unstable {\\cite{matsui81}}. On the other hand, as the density increases, the attractive scalar meson exchange force tends to be suppressed by the relativistic effect and the net quasiparticle interaction become repulsive. At nuclear saturation density ($\\rho_{0}=0.148 {\\rm fm^{-3}}$) we have $K=476.04$ MeV and $705.84$ MeV for combined $\\s+\\o$ and $\\s+\\o+\\p$ mesons respectively. The small effective mass is responsible for large incompressibility. The first sound velocity $c_{1}=0.19$ for $\\s+\\o$ and $0.23$ for $\\s+\\o+\\p$ at normal nuclear matter density and at all densities $c_{1}$ is smaller than the velocity of light, which is consistent with causality {\\cite{matsui81}}.    \\vskip 0.4in \\section {Isovector Landau parameter and symmetry energy}  In this section we proceed to calculate isovector LPs due to pion exchange. The isovector contribution to interaction parameter is  \\beq\\label{pion_interaction} f_{pp'}^{\\prime}{\\Big\\vert}_{p\\simeq p'=p_{f}}&=& -\\frac{1}{2}\\f{g_{A}^2M^{*2}}{4f_{\\p}^2\\veps_{f}^2} \\lt\\{\\f{p_{f}^2(1-\\cos\\th)}{2p_{f}^2(1-\\cos\\th)+m_{\\p}^2}\\rt\\}. \\eeq  where $g_{A}^2=1.5876$ , $f_{\\p}=93 MeV$  and $m_{\\p}=139$ MeV {\\cite{hu07}}. Here $-1/2$ is isospin factor in isovector channel {\\cite{friman99}}.  Using Eq.(\\ref{landau_para}) and (\\ref{pion_interaction}) we derive isovector LPs $f'_{0}$ and $f'_{1}$,  \\beq\\label{f0_pi} f^{\\prime}_{0}&=&-\\f{g_{A}^2M^{*2}p_{f}^2}{16f_{\\p}^2\\veps_{f}^2} \\int_{-1}^{1}\\f{(1-\\cos\\theta)}{2p_{f}^2(1-\\cos\\theta)+m_{\\p}^2} {\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\f{g_{A}^2M^{*2}}{32f_{\\p}^2\\veps_{f}^2} \\left[-2+\\frac{m_{\\pi}^2}{2p_{f}^2} \\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\right], \\eeq  and  \\beq\\label{f1_pi} \\frac{1}{3}f^{\\prime}_{1}&=&-\\f{g_{A}^2M^{*2}p_{f}^2}{16f_{\\p}^2\\veps_{f}^2} \\int_{-1}^{1}\\f{\\cos\\theta(1-\\cos\\theta)}{2p_{f}^2(1-\\cos\\theta)+m_{\\p}^2} {\\rm d(\\cos\\theta)}\\nn\\\\ &=&\\f{g_{A}^2M^{*2}m_{\\p}^2} {64f_{\\p}^2\\veps_{f}^2p_{f}^2} \\left[-2+\\left(\\frac{m_{\\pi}^2} {2p_{f}^2}+1\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\right]. \\eeq   The dimensionless LPs $F_{0}^{\\prime}=N(0)f_{0}^{\\prime}$ and $F_{1}^{\\prime}=N(0)f_{1}^{\\prime}$ are given below. Here, $N(0)$ is the density of states at the Fermi surface defined in Eq.(\\ref{dens_of_state}).  \\beq F^{\\prime}_{0}&=&g_{s}g_{I}\\f{g_{A}^2p_{f}M^{*2}}{64\\p^2f_{\\p}^2\\veps_{f}} \\lt[-2+\\frac{m_{\\pi}^2}{2p_{f}^2} \\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\rt], \\eeq  and  \\beq \\f{1}{3}F^{\\prime}_{1}&=&g_{s}g_{I}\\f{g_{A}^2m_{\\p}^2M^{*2}} {128\\p^2f_{\\p}^2p_{f}\\veps_{f}} \\lt[-2+\\left(\\frac{m_{\\pi}^2} {2p_{f}^2}+1\\right)\\ln\\left(1+\\frac{4p_{f}^2}{m_{\\pi}^2}\\right)\\rt]. \\eeq   \\begin{figure}[tb] \\begin{center} \\resizebox{8cm}{6.0cm}{\\includegraphics[]{F_prm.eps}} \\caption{Density dependent of isovector LPs in symmetric nuclear matter.} \\label{fig14} \\end{center} \\end{figure}     Knowing the isovector LPs, to which here only the pion contributes, one can calculate nuclear symmetry energy. The symmetry energy is defined as the difference of energy between the neutron matter and symmetric nuclear matter which we denote as $\\beta$ \\cite{matsui81}.  Analytically, the symmetry energy is defined from the expansion of the energy per nucleon $E'(\\rho,\\alpha)$ in terms of the asymmetry parameter $\\alpha$ defined as {\\cite{greco03}}  \\beq \\alpha \\equiv-\\frac{\\rho_{3}}{\\rho}=\\frac{\\rho_{n}-\\rho_{p}}{\\rho}=\\frac{N-Z}{A}. \\eeq  We have  \\beq E'(\\rho,\\alpha)\\equiv\\frac{E(\\rho,\\alpha)}{\\rho}=E'(\\rho)+E'_{sym}(\\rho)\\alpha^2 +{\\cal{O}}(\\alpha^4)+..... \\eeq  and so, in general,  \\beq E'_{sym}\\equiv\\beta&=&\\frac{1}{2}\\frac{\\del^{2}E'(\\rho,\\alpha)}{\\del\\alpha^2} {\\Big|}_{\\alpha=0}\\nn\\\\ &=&\\frac{1}{2}\\rho\\frac{\\del^{2}E}{\\del\\rho_{3}^2}{\\Big|}_{\\rho_{3}=0}. \\eeq  In terms of LPs, the symmetry energy can be expressed as  \\beq\\label{symm_energy} \\beta = \\frac{p_{f}^2}{6\\veps_{f}}(1+F'_{0}), \\eeq  where $F'_{0}$ is the isospin dependent LP. Numerically at saturation density $(\\rho=\\rho_0)$ we obtain $\\beta=14.57$ MeV. So relatively small contribution to $\\beta$ comes from one-pion exchange diagram {\\cite{die03}}.   Similar to the hydrodynamic sound corresponding to the total baryon density oscillation, we consider hydrodynamic isospin sound which accompany the out-of-phase oscillations of proton and neutron density. Isospin sound velocity $c_{1}^{\\prime}$ is given by {\\cite{matsui81}}  \\beq\\label{iso_sound} c_{1}^{\\prime}&=& v_{f}\\sqrt{\\frac{\\veps_{f}}{3\\mu}(1+F_{0}^{\\prime})} \\eeq  Numerically at saturation density $(\\rho=\\rho_0)$ we obtain $c_1^{\\prime}=0.17$.     \\vskip 0.4in ",
        "conclusions": ""
    },
    "0809/0809.3245_arXiv.txt": {
        "abstract": "This article investigates the full Boltzmann equation up to second order in the cosmological perturbations. Describing the distribution of polarized radiation by using a tensor valued distribution function, the second order Boltzmann equation, including polarization, is derived without relying on the Stokes parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.4805.txt": {
        "abstract": "{ XTE J1810$-$197 and 1E 1547.0$-$5408 are two transient AXPs exhibiting radio emission with unusual properties. In addition, their spin down rates during outburst show opposite trends, which so far has no explanation. Here, we extend our quark-nova model for AXPs to include transient AXPs, in which the outbursts are caused by transient accretion events from a Keplerian (iron-rich) degenerate ring. For a ring with inner and outer radii of $23.5$ km and $26.5$ km, respectively, our model gives a good fit  to the observed X-ray outburst from XTE J1810$-$197 and the behavior of temperature, luminosity, and area of the two X-ray blackbodies with time. The two blackbodies in our model are related to a heat front (i.e. Bohm diffusion front)  propagating along the ring's surface and an accretion hot spot on the quark star surface. Radio pulsations in our model are caused  by dissipation at the light cylinder of magnetic bubbles, produced near the ring during the X-ray outburst. The delay between X-ray peak emission and radio emission in our model is related to the propagation time of these bubbles to the light cylinder and  scale with the period  as $t_{\\rm prop.}\\propto P^{\\frac{7}{2}-\\frac{\\alpha}{2}}/\\dot{P}^{1/2}$ where $\\alpha$ defines the radial dependence of matter density in the magnetosphere ($\\propto r^{-\\alpha}$);  for an equatorial wind, $\\alpha=1$, we predict  a $\\sim 1$ year and $\\sim 1$ month delay for XTE J1810$-$197 and 1E 1547.0$-$5408, respectively.  The observed flat spectrum, erratic pulse profile, and the pulse duration are all explained in our model as a result of X-point reconnection events induced by the  dissipation of the bubbles at the light cylinder.  The spin down rate of the central quark star can either increase or decrease depending on how the radial drift velocity of the magnetic islands changes with distance from the central star.  We suggest an evolutionary connection between transient AXPs and typical AXPs in our model. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) are magnetars with rotation period of 2-12 seconds and inferred surface magnetic field strength $B\\sim 10^{14-15}$ G (e.g. Woods \\& Thompson 2006; Kaspi 2007).  In this work we focus on 2 AXPs, XTE J1810$-$197 and 1E 1547.0$-$5408, which are the only magnetars known to emit in the radio (Camilo et al. 2006). Both are demonstrably transient radio sources, having not been detected in previous surveys of adequate sensitivity. XTE J1810$-$197 is a transient AXP\\footnote{In the sense that in quiescence their surface temperature are as low as those of some ordinary  young neutron stars} first detected when its X-ray flux increased $\\sim 100$-fold compared to a quiescent level it maintained for at least 24 years (Ibrahim et al. 2004). Discovered with the {\\it Einstein X-ray} satellite in 1980, 1E 1547.0$-$5408 was eventually identified as a magnetar candidate (Gelfand \\& Gaensler 2007) with spectral characteristics of an AXP. In this paper we look at these sources in the Quark-Nova context (hereafter QN; Ouyed et al. 2002) building on three previous papers where we explore its application to Soft Gamma-ray Repeaters (SGRs) (Ouyed, Leahy, \\& Niebergal 2007a; OLNI), to AXPs (Ouyed, Leahy, \\& Niebergal 2007b; OLNII), and to Rotating Radio Transients (RRATs) (Ouyed et al. 2009; OLNIII), and superluminous supernovae (Leahy\\&Ouyed 2008). But first, we briefly describe their observed X-ray and radio properties, during quiescence and bursting phases.     \\subsection{The X-ray emission}  In the pre-burst era, XTE J1810$-$197's ROSAT spectrum showed a single blackbody (BB) with temperature $T = 0.18$ keV, an emitting area of $\\sim 520\\ {\\rm km}^2  (d/3.3\\ {\\rm kpc})^2$, and a luminosity of $L_{\\rm BB}\\sim 5.6\\times 10^{33}\\ {\\rm erg\\ s}^{-1} (d/3.3\\ {\\rm kpc})^2$. During its bursting phase,   XTE J1810$-$197  showed a hot blackbody ($T\\sim 0.65$ keV) with an exponential decay in X-ray luminosity of $\\sim 280$ days, as well as a warm blackbody ($T\\sim 0.3$ keV) decaying at a rate of $\\sim 870$ days (Gotthelf \\& Halpern 2007). For the case of 1E 1547.0$-$5408, after its radio detection (Camilo et al. 2007a), an X-ray outburst was confirmed (Halpern et al.  2008) with a record high luminosity of $\\sim 1.7\\times 10^{35}~(d/9.9\\ {\\rm kpc})^2$ erg s$^{-1}$ and with a total outburst energy of $10^{42}\\ {\\rm erg}  < E_{\\rm b} < 10^{43}$ erg.     \\subsection{The radio emission}    For XTE J1810$-$197, the radio emission began within 1 yr of  its only known X-ray outburst (Camilo et al. 2006 and references therein). At its observed peak more than 3 yr after the X-ray outburst, the radio flux density was more than 50 times the pre-burst upper limit. The X-ray flux has since returned to its quiescent level nearly 4 yrs after the burst. 1E 1547.0$-$5408, although not as well sampled as XTE J1810$-$197, exhibits similar variations in flux density and was reported with a factor of 16 times the pre-burst upper limit (Camilo et al 2007a).  Trends in radio emission between the 2 sources can be summarized as follows:   \\begin{itemize} \\item  Both are very highly linearly polarized showing a  flat spectrum over a wide range of frequencies. Their striking spectra (i.e. spectral index $> -0.5$) clearly distinguishable from ordinary radio pulsars (with a spectral index $\\sim -1.6$; Camilo et al. 2007a\\&b).  \\item At their peak, both magnetars are very luminous in radio with luminosity at 1.4 GHz $L_{1.4}\\ge 100 {\\rm mJy}\\  {\\rm (d/kpc)}^2$, which is larger than the $L_{1.4}\\le 2 {\\rm mJy}\\ {\\rm (d/kpc)}^2$ of most any ordinary young pulsar (e.g. Camilo et al. 2006).   For XTE  J1810$-$197, its assumed isotropic radio luminosity up to 42 GHz is about $2\\times 10^{30}$ erg s$^{-1}$ (Camilo et al. 2006).  \\item   Both have variable pulse profiles (exhibiting sudden changes in radio pulse shape) and radio flux densities. The flux changes at all frequencies. At a given frequency there is no stable average pulse profile. Different pulse components change in relative intensity and new components sometimes appear. Sub-pulses with typical width approximately $<$ 10 ms are observed (Camilo et al. 2007a\\&b).  \\item For XTE J1810$-$197, the torque was decreasing, at a time when the star was returning to quiescence years after the large outburst. As the torque decreased, so did the radio flux (Camilo et al. 2007b).  \\item In 1E 1547.0$-$5408, in contrast, the torque has been increasing, at a time when the X-ray flux has been gradually decreasing (Camilo et al. 2007a).   \\end{itemize}  In this paper we extend our existing quark star model for AXPs to account for the observed behavior of these two transients. This paper is structured as follows: Section 2 gives the basic elements of the model. Section 3 describes the quiescent phase. Section 4, the bursting phase. Section 5, the radio emission. Model predictions are highlighted in Section 6 before we conclude in Section 7. ",
        "conclusions": "There are two fundamental components in our model for AXPs and transient AXPs namely, the QS and the Keplerian ring. In quiescence, vortex annihilation on the QS gives rise to thermal and non-thermal X-ray emission. The ring reprocesses the emission to give a second cooler BB emission component. Outburst is triggered by accretion of a small inner part of the ring (i.e. the wall). The two main consequences are production of light ($Z \\sim 13$) nuclei and triggering MHD accretion onto the QS (yielding the HS). The interplay between the Bohm diffusion (i.e. $R_{\\rm BF}$ term) and depletion of light nuclei (i.e. $\\mu$)   gives rise to a rich behavior, necessary in order to  account for the observed   behavior of XTE J1810$-$197. Finally, one can ask if such a small Keplerian degenerate iron-rich ring  could form around a neutron star. Ring formation when the neutron star is born appears implausible since a proto-neutron star  is large compared to the ring size. After formation, there is no obvious mechanism to eject degenerate material unless a violent change of state, like a QN occurs.       %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.3073_arXiv.txt": {
        "abstract": "We present the first results of $\\textit{AKARI}$ Infrared Camera near-infrared spectroscopic survey of the Large Magellanic Cloud (LMC). We detected absorption features of the H$_2$O ice 3.05$\\mu$m and the CO$_2$ ice 4.27$\\mu$m stretching mode toward seven massive young stellar objects (YSOs). These samples are for the first time spectroscopically confirmed to be YSOs. We used a curve-of-growth method to evaluate the column densities of the ices and derived the CO$_2$/H$_2$O ratio to be 0.45$\\pm$0.17. This is clearly higher than that seen in Galactic massive YSOs (0.17$\\pm$0.03). We suggest that the strong ultraviolet radiation field and/or the high dust temperature in the LMC may be responsible for the observed high CO$_2$ ice abundance. ",
        "introduction": "Properties of extragalactic young stellar objects (YSOs) provide us important information on the understanding of the diversity of YSOs in different galactic environments. The Large Magellanic Cloud (LMC), the nearest irregular galaxy to our Galaxy \\citep[$\\sim$50kpc;][]{Alv04}, offers an ideal environment for this study since it holds a unique metal-poor environment \\citep{Luc98}. Because of its proximity and nearly face-on geometry, various types of surveys have been performed toward the LMC \\citep[e.g., ][ and references therein]{Zar04,Mei06,Kat07}.  An infrared spectrum of YSOs shows absorption features of various ices which are thought to be an important reservoir of heavy elements and complex molecules in a cold environment such as a dense molecular cloud or an envelope of a YSO \\citep[e.g., ][]{Chi98,Num01,Whit07,Boo08}. These ices are thought to be taken into planets and comets as a result of subsequent planetary formation activity \\citep{Ehr00}. Studying the compositions of ices as functions of physical environments is crucial to understand the chemical evolution in circumstellar environments of YSOs and is a key topic of astrophysics. H$_2$O and CO$_2$ ices are ubiquitous and are major components of interstellar ices \\citep{vDB98,BE04}. Since the absorption profile of the ices is sensitive to a chemical composition of icy grain mantles and a thermal history of local environments, the ices are important tracers to investigate the properties of YSOs. However, our knowledge about the ices around extragalactic YSOs is limited because few observations have been performed toward extragalactic YSOs. Therefore infrared spectroscopic observations toward YSOs in the LMC are important if we are to improve our understanding of the influence of galactic environments on the properties of YSOs and ices.  $\\textit{AKARI}$ is the first Japanese satellite dedicated to an infrared astronomy launched in February 2006 \\citep{Mur07}. We have performed a near infrared spectroscopic survey of the LMC using a powerful spectroscopic survey capability of Infrared Camera \\citep[IRC;][]{TON07} on board $\\textit{AKARI}$. In this letter, we present 2.5--5$\\mu$m spectra of newly confirmed YSOs in the LMC with our survey, and discuss the abundance of H$_2$O and CO$_2$ ice. ",
        "conclusions": "The obtained column densities of H$_2$O and CO$_2$ ices are plotted in Fig 2. The error bars become larger for the larger column density due to the saturation effect of the curve-of-growth. A linear fit to the data points indicates that the CO$_2$/H$_2$O ice column density ratio in the LMC is 0.45 $\\pm$ 0.17. The large uncertainty mainly comes from the errors in the curve-of-growth analysis. For comparison, column densities of Galactic massive YSOs taken from \\citet{Gib04} and their CO$_2$/H$_2$O ice column density ratio of 0.17 $\\pm$ 0.03 \\citep{Ger99} are also plotted in Fig 2. A similar CO$_2$/H$_2$O ratio of 0.18 $\\pm$ 0.04 is also observed toward a Galactic quiescent dark cloud \\citep{Whit07}, while a relatively high CO$_2$/H$_2$O ratio of 0.32 $\\pm$ 0.02 is observed toward Galactic low-- and intermediate-- mass YSOs, and some of them reach $\\sim$0.4 \\citep{Pon08}. Although the uncertainty is large, it is clear from the present results that the CO$_2$/H$_2$O ice ratio in the LMC is higher than the typical ratios of the Galactic objects. Since the distribution range of the H$_2$O ice column density in the LMC is comparable to that of the massive Galactic YSOs, it can be concluded that the abundance of the CO$_2$ ice is higher in the LMC. The present results suggest that the different galactic environment of the LMC is responsible for the high CO$_2$ abundance.   The formation mechanism of CO$_2$ ice in circumstellar environments of YSOs is not understood, however a number of scenarios have been proposed. Several laboratory experiments indicate that the CO$_2$ ice is efficiently produced by UV photon irradiation to H$_2$O-CO binary ice mixtures \\citep[e.g., ][]{Wat07}. The LMC has an order-of-magnitude stronger UV radiation field than our Galaxy due to its active massive star formation \\citep{Isr86}, which could lead to the higher CO$_2$/H$_2$O ratio in the LMC. The high CO$_2$/H$_2$O ratio toward a YSO in the LMC is also reported in \\citet{vanL05}, and they suggest that a different radiation environment in the LMC is one of the reasons for the high CO$_2$ abundance. On the other hand, the model of diffusive surface chemistry suggests that high abundance of CO$_2$ ice can be produced at relatively high dust temperatures \\citep{Ber99,Ruf01} . Several studies have reported that the dust temperature in the LMC is generally higher than in our Galaxy based on far-infrared to submillimeter observations of diffuse emission \\citep[e.g., ][]{Agu03, Sak06}. Therefore the high dust temperature may also have an effect on the high CO$_2$ ice abundance in the LMC.  It is difficult to separate the effect of the UV radiation field and the dust temperature on the high abundance of CO$_2$ ice in the LMC by our low-resolution NIR spectra. The 4.62 $\\mu$m XCN feature is known to be indicative of strong UV irradiation \\citep{Ber00,Spo03}. On the other hand, detailed profile analysis of the 3.05 $\\mu$m H$_2$O ice stretching mode and the 15.2 $\\mu$m CO$_2$ ice bending mode should reveal the temperature and compositions of the ices \\citep{Ehr96,Obe07}. Future observations of the XCN feature and the H$_2$O and CO$_2$ ice features with a sufficient wavelength resolution will be useful to investigate this problem."
    },
    "0809/0809.3559_arXiv.txt": {
        "abstract": "We present results from the Suzaku observations of the dwarf nova SS~Cyg in quiescence and outburst in 2005 November. High sensitivity of the HXD PIN and high spectral resolution of the XIS enable us to determine plasma parameters with unprecedented precision. The maximum temperature of the plasma in quiescence $20.4^{+4.0}_{-2.6}\\,\\mbox{(stat.)}\\pm 3.0\\,\\mbox{(sys.)}$~keV is significantly higher than that in outburst $6.0^{+0.2}_{-1.3}$~keV. The elemental abundances are close to the solar ones for the medium-Z elements (Si, S, Ar) whereas they decline both in lighter and heavier elements, except for that of carbon which is 2 solar at least. The solid angle of the reflector subtending over an optically thin thermal plasma is $\\Omega^{\\rm Q}/2\\pi = 1.7\\pm 0.2\\,\\mbox{(stat.)}\\pm 0.1\\,\\mbox{(sys.)}$ in quiescence. A 6.4~keV iron {\\ka} line is resolved into narrow and broad components. These facts indicate that both the white dwarf and the accretion disk contribute to the reflection. We consider the standard optically thin boundary layer as the most plausible picture for the plasma configuration in quiescence. The solid angle of the reflector in outburst $\\Omega^{\\rm O}/2\\pi = 0.9^{+0.5}_{-0.4}$ and a broad 6.4~keV iron line indicate that the reflection in outburst originates from the accretion disk and an equatorial accretion belt. The broad 6.4~keV line suggests that the optically thin thermal plasma is distributed on the accretion disk like solar coronae. ",
        "introduction": "Dwarf novae (DNe) are non-magnetic cataclysmic variables (CVs; binaries between a white dwarf primary and a mass-donating late-type star) which show optical outbursts typically with $\\Delta m_V =$ 2--5 lasting 2--20~d with intervals of $\\sim$10~d to tens of years \\citep{1995CAS....28.....W}. These outbursts can be explained as a result of a sudden increase of mass-transfer rate within the accretion disk surrounding the white dwarf due to a thermal-viscous instability \\citep{1974PASJ...26..429O,1981A&A...104L..10M,1982MNRAS.199..267B,1984PASP...96....5S,1993ApJ...419..318C,1996PASP..108...39O}. A boundary layer (hereafter abbreviated as BL) is formed between the inner edge of the accretion disk and the white dwarf where matter transferred through the disk releases its Keplerian motion energy and settles onto the white dwarf. BL is a target of EUV and X-ray observations since its temperature becomes $T\\simeq 10^5$--10$^8$~K. \\citet{1979MNRAS.187..777P} and \\citet{1985ApJ...292..535P} have predicted that radiation from the BL starts to shift from hard X-ray to EVU when the high $\\dot{M}$ front arrives at the inner edge of the disk, because BL becomes optically thick to its own radiation. This prediction has been verified by a number of multi-waveband coordinated observations \\citep{1979MNRAS.186..233R,1992MNRAS.257..633J,2003MNRAS.345...49W}.  The region around the inner edge of the disk is filled with a lot of intriguing but still unresolved issues. Within the framework of the standard accretion disk \\citep{1973A&A....24..337S}, half of the gravitational energy is released in the accretion disk, and hence, the other half is released in BL. The observations in extreme-ultraviolet band of VW~Hyi and SS~Cyg, however, revealed that the fractional energy radiated from BL is only $<10$\\% of the disk luminosity \\citep{1991ApJ...372..659M,1995ApJ...446..842M}.  According to the classical theory, the temperature of BL in outburst is predicted to be 2--5$\\times 10^5$~K \\citep{1979MNRAS.187..777P}, whereas the temperature estimated by ultraviolet and optical emission lines is constrained to a significantly lower range 5--10$\\times 10^4$~K \\citep{1991MNRAS.249..452H}. These discrepancies may be resolved if we assume that BL is terminated not on the static white dwarf surface but on a rapidly rotating accretion belt on the equatorial surface of the white dwarf \\citep{1978necb.conf...89P,1978A&A....63..265K}. Suggestions of the accretion belt, rotating at a speed close to the local Keplerian velocity, have been reported from a few DNe in outburst \\citep{1993ApJ...405..327L,1996ApJ...458..355H,1996ApJ...471L..41S,1997AJ....114.1165C,1998ApJ...497..928S}. Mechanism has not been understood yet to drive a dwarf nova oscillation (DNO) which is a highly coherent oscillation of soft X-ray and optical intensities with a period of 3--40~s \\citep{1978ApJ...219..168R,1980ApJ...235..163C,1984ApJ...278..739C,1986A&A...158..233S,1998MNRAS.299..921M,1998PASP..110..403P,2001ApJ...562..508M}. \\citet{2002MNRAS.335...84W} try to understand DNO by assuming a magnetically driven accretion onto the accretion belt by enhancing a magnetic field the belt through a dynamo mechanism.  One of the unresolved outstanding issues may be the origin of a hard X-ray optically thin thermal emission in outburst, since BL is believed to be optically thick. In order to identify its emission site, and to obtain some new insight on BL in quiescence as well, we planned to observe SS~Cyg both in quiescence and outburst with the X-ray observatory Suzaku \\citep{2007PASJ...59S...1M}. SS~Cyg is a dwarf nova in which the $1.19\\pm 0.02\\MO$ white dwarf and the $0.704\\pm 0.002 \\MO$ secondary star \\citep{1990MNRAS.246..654F} are revolving in an orbit of $i = 37^\\circ\\pm 5^\\circ$ \\citep{1983PhDT........14S} with a period of 6.6 hours. The distance to SS~Cyg is measured to be $166\\pm12$~pc using HST/FGS parallax \\citep{1999ApJ...515L..93H}. SS~Cyg shows an optical outburst roughly every 50 days, in which $m_V$ changes from 12th to 8th magnitude. The optically thin to thick transition of BL has clearly been detected with coordinated observations of optical, EUVE, and RXTE \\citep{2003MNRAS.345...49W}. In \\S~2, we show an observation log and procedure of data reduction. In \\S~3, detail of our spectral analysis is explained. Owing to a high spectral resolution of the X-ray Imaging Spectrometer (XIS; \\cite{2007PASJ...59S..23K}) and a high sensitivity of the Hard X-ray Detector (HXD; \\cite{2007PASJ...59S..35T}; \\cite{2007PASJ...59S..53K}) over 10~keV enable us to determine spectral parameters of hard X-ray emission of SS~Cyg with unprecedented precision. In \\S~4, we discuss on the emission site and its spatial extension both in quiescence and outburst by utilizing the spectral parameters, a 6.4~keV iron line parameters in particular. We summarize our results and discussions in \\S~5. ",
        "conclusions": "We have presented results of the Suzaku observations on the dwarf nova SS~Cyg in quiescence and outburst in 2005 November. The X-ray spectra of SS~Cyg are composed of a multi-temperature optically thin thermal plasma model with a maximum temperature of a few tens of keV, its reflection from the white dwarf surface and/or the accretion disk, and a 6.4~keV neutral iron {\\ka} line from the reflectors via fluorescence. High sensitivity of the HXD PIN detector and the high spectral resolution of the XIS enable us to disentangle degeneracy between the maximum temperature and the reflection parameters, and to determine the emission parameters with unprecedented precision. The maximum temperature of the plasma in quiescence $kT^{\\rm Q}_{\\rm max} = 20.4^{+4.0}_{-2.6}\\,\\mbox{(stat.)}\\pm 3.0\\,\\mbox{(sys.)}$~keV is significantly higher than that in outburst $kT^{\\rm O}_{\\rm max} = 6.0^{+0.2}_{-1.3}$~keV. The elemental abundances of the plasma are close to the solar ones for the medium-Z elements (Si, S, Ar) whereas they declines both in lighter and heavier elements. Those of oxygen and iron are 0.46$^{+0.04}_{-0.03}\\,\\mbox{(stat.)}\\pm 0.01\\,\\mbox{(sys.)}Z_\\odot$ and 0.37$^{+0.01}_{-0.03}\\,\\mbox{(stat.)}\\pm 0.01\\,\\mbox{(sys.)}Z_\\odot$. The exception is carbon whose abundance is at least $2Z_\\odot$ even if we take into account all possible systematic errors. These trends are similar to other dwarf novae observed with XMM-Newton \\citep{2005ApJ...626..396P}.  The solid angle of the reflector subtending over the optically thin thermal plasma is $\\Omega^{\\rm Q}/2\\pi = 1.7\\pm 0.2\\,\\mbox{(stat.)}\\pm 0.1\\,\\mbox{(sys.)}$ in quiescence. Since even an infinite slab can subtend a solid angle of $\\Omega/2\\pi = 1$ over a radiation source above it, this large solid angle can be achieved only if the plasma views both the white dwarf and the accretion disk with substantial solid angles. Thanks to high energy resolution of the XIS, we have resolved a 6.4~keV iron {\\ka} line into a narrow and broad components (significance of the broad component is $\\sim$99\\%), which also indicate contributions from both the white dwarf and the accretion disk to the reflected continuum spectra. The equivalent widths of them are both $\\sim$50~eV. From all these results, we consider the standard optically thin BL formed between the inner edge of the accretion disk and the white dwarf surface \\citep{1985ApJ...292..535P} as the most plausible model to explain the observed large solid angle. From the equivalent width of the narrow 6.4~keV component, the height of the BL from the white dwarf surface is $h < 0.12R_{\\rm WD}$. The total equivalent width of the 6.4~keV line ($\\sim$100~eV) is consistent with that expected from $\\Omega^{\\rm Q}/2\\pi$, the iron abundance, and the incident illuminating continuum spectrum.  The solid angle of the reflector in outburst $\\Omega^{\\rm O}/2\\pi = 0.9^{+0.5}_{-0.4}$, on the other hand, is significantly smaller than that in quiescence, and is consistent with an infinite slab. Since the 6.4~keV iron emission line is broad with no narrow component ($\\lesssim$20\\% of the broad component), the reflection originates from the accretion disk. The accretion belt can also contribute to the reflection. The 6.4~keV line from the accretion belt is expected to be broad, which is consistent with the absence of the narrow 6.4~keV component. The EW of the 6.4~keV line is so large that it cannot be interpreted within a simple scheme of reflection from the disk. Even if Compton down-scattering of the observed He-like {\\ka} line is taken into account, we can only find a solution which marginally reconciles the large EW with the solid angle of the reflector. We consider the optically thin thermal plasma in outburst as being distributed on the accretion disk. The Chandra HETG observation in outburst revealed that the He-like and H-like emission lines from O, Ne, Mg, and Si are broad and their widths ($\\sim$2000~km~s$^{-1}$) are consistent with those expected from the Keplerian velocity of the accretion disk \\citep{2008ApJ...680..695O}. This fact suggests that the optically thin thermal plasma is anchored to the accretion disk and the accretion belt by magnetic field, for example, like solar coronae.  \\appendix"
    },
    "0809/0809.4186_arXiv.txt": {
        "abstract": "The Cryogenic Dark Matter Search experiment (CDMS) employs low-temperature Ge and Si detectors to detect WIMPs via their elastic scattering interaction with the target nuclei. The current analysis of 397.8\\,kg-days Ge exposure resulted in zero observed candidate events, setting an upper limit on the spin-independent WIMP-nucleon cross-section of 6.6\\,$\\times$\\,$10^{-44}$\\,cm$^2$ (4.6\\,$\\times$\\,$10^{-44}$\\,cm$^2$, when previous CDMS Soudan data is included) at the 90\\% confidence level for a WIMP mass of 60\\,GeV. To increase the sensitivity, new one inch thick detectors have been developed which will be used in the SuperCDMS phase. SuperCDMS 25kg will be operated at SNOLAB with an expected sensitivity on the spin-independent WIMP-nucleon elastic scattering cross-section of 1\\,$\\times$\\,$10^{-45}$\\,cm$^2$. ",
        "introduction": "The Cryogenic Dark Matter Search (CDMS) experiment operates 19 Ge (250\\,g each) and 11 Si (100\\,g each) detectors at the Soudan underground laboratory (MN, USA) to search for non-luminous, non-baryonic Weakly Interacting Massive Particles (WIMPs), that could form the majority of the matter in the universe \\cite{Spergel,Jungman} . Each detector is a disk 7.6\\,cm in diameter and 1\\,cm thick. The detectors are operated at cryogenic temperatures $\\sim$ 40 mK to collect the athermal phonons created upon an interaction in the crystal in four independent sensors. In addition, the electron hole pairs created by a recoil are drifted in a field of 3\\,V/cm (Ge), 4\\,V/cm (Si) towards two concentric electrodes lithographically patterned on one flat side of the crystals \\cite{zipdetectors}. In the analysis events from the outer part of the detectors are removed by a fiducial volume cut based on the partitioning of energy between the two concentric charge electrodes. The simultaneous measurement of the phonon and ionization recoil energy of an interaction in the crystals not only allows an accurate measurement of the recoil energy independent of recoil type (nuclear/electron recoil), but also allows the discrimination between nuclear and electron recoils by the so called ionization yield parameter, which is the ratio of the ionization and phonon energy, providing a rejection factor of $>$\\,$10^4$. Nuclear recoils produce fewer charge pairs, and hence less ionization energy than do electron recoils of the same energy. The ionization yield for electron and nuclear recoils is determined from $^{133}$Ba and $^{252}$Cf calibrations respectively, providing the bands shown in Fig.\\ref{fig:yieldtiming}. ",
        "conclusions": "The CDMS-II experiment has maintained high dark matter discovery potential by limiting expected backgrounds to less than one event in the signal region. The current data sets the world's most stringent upper limit on the spin-independent WIMP-nucleon cross-section for WIMP masses above 42\\,GeV/c$^2$ with a minimum of 4.6\\,$\\times $\\,$10^{-44}$\\,cm$^2$ for a WIMP mass of 60\\,GeV/c$^2$. Ongoing runs aim to accumulate roughly 2000\\, kg-days of WIMP search exposure until the end of 2008. By this the CDMS-II experiment is expected to reach a sensitivity of 1\\,$\\times$\\,$10^{-44}$\\,cm$^{2}$. \\begin{wrapfigure}{l}{0.5\\textwidth} \\centering \\includegraphics[scale=0.7]{bruch_tobias.fig4.eps} \\caption{\\small{Upper limits on the spin independent WIMP-nucleon cross-section from the current analysis (red dashed) and the combined limit by including previous CDMS data (red solid) \\cite{r123analysis}. Also shown are the limit from the XENON10 experiment \\cite{xenon10} (black solid) and expected sensitivities of the CDMS-II setup until the end of 2008 (light gray); two super towers operated at Soudan (dark gray/solid) and the SuperCDMS 25 kg stage (dark gray/dashed). Filled regions indicate CMSSM models \\cite{baltz,ruiz}.}} \\label{fig:exlimits} \\end{wrapfigure} The first two super towers with new 1\\,inch thick detectors will be installed at the Soudan site by 2009 demonstrating the improved discrimination capabilities. The next upgrade of the CDMS experiment to SuperCDMS 25\\,kg operating seven super towers will be installed at SNOLAB, increasing the sensitivity by one order of magnitude.        \\begin{footnotesize}"
    },
    "0809/0809.4523_arXiv.txt": {
        "abstract": "A  recent paper by   Stanimirovi{\\'c} \\etal \\ (2008) presents quit interesting results  from H$_{\\rm I}$ observations of the Magellanic Stream (MS) tip. The high spatial resolution of the  data    reveals   rich and complex  morphological and   kinematic structures; notably four coherent $H_I$ substreams   extending over angular size of about $20^o$ were found.  We suggest to use the data to search for the existence of an underlying turbulence in the residuals of velocity fields. If existent, a turbulence  would provide a {\\it dynamical} evidence ,that the sub streams are cohorent structures. The characteristics of the turbulence could yield information about the energy source, as well as about the physical parameters of the gas in these streams.  We use the position-velocity images of  Stanimirovi{\\'c} \\etal \\ (2008) to derive   spatial power spectra for the velocity residuals. These, indicate the presence of a large scale turbulence with  size comparable   to  that of the streams themselves. The turbulent velocity on the largest scale is estimated to be about  $15 \\ {\\rm km/s}$. Adopting, a distance of $120 \\  {\\rm kpc}$, implies a turbulent largest scale of $\\sim 40 \\  {\\rm kpc}$ and timescale for decay of about  $3 {\\ \\rm Gyr}$.  For a turbulence with scale that large, the natural  energy source   is the tidal interaction between the Magellanic Clouds, and between them and the Milky Way galaxy. The  estimated turbulent timescale for decay is consistent with this mechanism. Such a mechanism has been suggested for the turbulence in the ISM of the SMC by Goldman (2000, 2007). In effect, the turbulence is a fossil from the era of the streams formation.  The shape of derived turbulence spectrum is used here  to  obtain constraints on the inclination of the streams and on the density of the emitting neutral hydrogen. ",
        "introduction": "In a recent paper    Stanimirovi{\\'c} \\etal \\ (2008) present  quit interesting   results  from H$_{\\rm I}$ observations of the tip of the Magellanic Stream (MS) . The high spatial resolution of the data obtained by   Stanimirovi{\\'c} \\etal \\ (2008), reveals   rich and complex  morphological and kinematic structures. The authors find four coherent $H_I$ substreams in the tip of the MS extending over projected angular size of about $20^o$. Three of the these streams (S2, S3, S4) originate from about the same location and are clumpy. The remaining, S1 stream, seems more diffuse and doesn't share the common location of the former streams. In all streams, the kinematic data show  large scale velocity gradients   of $\\sim  (5-10){\\rm km\\ s^{-1}\\ deg^{-1}}$.  By comparing the observations to the simulations of Connors \\etal \\ (2006),  Stanimirovi{\\'c} \\etal \\ (2008), interpret the three former streams to be the result of the tidal splitting of the main MS  by   tidal interaction of the LMC and the MS about  1.05 Gyr  and  0.55 Gyr ago. In this picture, the MS stream itself was formed  about 1.5  Gyr ago by close tidal encounter of the SMC, LMC and the Milky Way (MW).  The   S1 sub-stream, is interpreted to have formed much more recently, about  0.2 Gyr ago, and consists of gas drawn from the Magellanic Bridge. Contrary to the former three streams, it had not enough  time to cool and fragment. ",
        "conclusions": "The main results of the present work are: \\renewcommand{\\labelenumi}{\\arabic{enumi}.} \\begin{enumerate} \\item Each of the streams exhibits a large scale turbulence   comparable to the size of the stream. This provides {\\it dynamical} proof that the streams are coherent structures, as stated in  Stanimirovi{\\'c} \\etal \\ (2008). \\item For an assumed distance of 120 kpc to the MS tip, the projected size of the stream and the largest turbulence scale is $\\sim 40$ kpc. \\item The turbulent r.m.s velocity on the largest scale is about $15\\ {\\rm km s^{-1}}$, namely mildly supersonic. This is in line with the inertial range index of the turbulent spectral energy function  being $-2$ instead of the Kolmogorov value $-5/3$, appropriate for subsonic turbulence. \\item The resulting timescale for decay of the turbulence is about 3 Gyr and is longer than the age of the MS, thus providing \"fossil evidence\". \\item The fact that a break in the power spectrum is evident in S1, and even more so in S3 indicates that the inclination angle of these stream, with respect to the line of sight is not large; they are viewed almost face-on. \\item The absence of a clear break in the power spectra of S2 and S4, suggest that their inclination is larger. Taking the absence of the break to imply that the  projected depth is larger than about $8^0$ and assuming that the true depth is $4^0$, as deduced for S2 and S4, results in an inclination of $\\sim 60 ^0$. From Fig. 6 of Connors \\etal \\ (2006) one can deduce a similar inclination at the tip of the MS. \\item A depth of  $4^0$ at a distance of 120 kpc corresponds to a physical depth of a 8 kpc. For   column densities in the range  $(1-10) \\times 10^{19} {\\rm cm^{-2}}$ this implies a depth-averaged   density of the warm neutral medium: $n_{WNM} = (0.3 -3) \\times 10^{-3} {\\rm cm^{-3}}$. \\end{enumerate}  \\begin{figure}[h] \\centerline{ \\includegraphics[scale=0.75]{fig.pdf}}  \\caption{ {\\it Dots}: The power spectrum of the  velocity residuals in arbitrary units as function of the normalized wavenumber. Wavenumber 1 corresponds to a largest angular scale in each stream--  S1: $15^0$, S2: $11.9^0$, S3: $17^0$, S4: $15.9^0$.  {\\it Lines}:  $k^{-2}, k^{-3}$ for S1 and S3; $k^{-3}$ for S2 and S4.} \\end{figure}"
    },
    "0809/0809.3135_arXiv.txt": {
        "abstract": "Non-monotonic features of rotation curves, and also the related gravitational effects typical of thin disks -- like backward-reaction or amplification of rotation by negative surface density gradients -- which are characteristic imprints of disk-like mass distributions, are discussed in the axisymmetric thin disk model. The influence of the data cutoff in rotational velocity measurements on the determination of the mass distribution in flattened galaxies is studied.  It has also been found that the baryonic matter distribution in the spiral galaxy NGC 5475, obtained in the axisymmetric thin disk approximation, accounts for the rotation curve of the galaxy. To obtain these results, the iteration method developed recently by the authors has been applied. ",
        "introduction": "The analysis of rotation curves provides the most reliable means for ascertaining at least the gross distribution of gravitating matter within spiral galaxies, as Alar Toomre pointed out \\citep{bib:toomre_ad_disk_eq}. In the same paper Toomre formulated a complete mathematical model of an axisymmetric and infinitely thin disk rotating under its own gravity. The model offers a tool for determining the equilibrium mass distribution directly from the rotation law of a highly flattened system, such as a spiral galaxy. It is assumed in the framework of the model, that orbits of stars, gas, etc., are circular, and that gravitational forces are balanced solely by centrifugal forces; the effect of pressure is thus ignored. A few years earlier, following the work reported by \\citep{bib:prender}, Brandt discussed less general situation of a very flattened system regarded as an assembly of osculating homoeoids compressed to a circular disk \\citep{bib:brandt}. All these papers followed many other pioneering ones referred to in \\citep{bib:brandt}.   The customary parametric few-component  models relate the mass distribution of baryonic matter mainly to luminosity measurements. The obtained amount of luminous mass is usually insufficient to account for the observed rotation of galaxies and a massive spherical dark halo is introduced as a remedy for the missing mass. The thin disk model with the surface mass density reconstructed mainly from rotation curves performs very well with strongly non-monotonic rotation curves, whereas the customary models have difficulties in explaining them. The thin disk model can easily account for high local gradients of rotational velocity, typically giving a lower total mass. In an extreme example, the rotational velocity of an outer galactic region, falling off faster than Keplerian, cannot be explained by the presence of a massive spherically symmetric halo. However, such a feature may be explained using a flattened mass distribution with a suitably changing mass profile.  The major difficulty in the practical use of the disk model lies with unambiguous mass density reconstruction. Unlike for spherical symmetry, the density depends on the assumed extrapolation of the rotation curve beyond the last measurement point (referred to in this paper as the 'cutoff radius'). This is a consequence of the nonlocal relation between the surface density and the rotational velocity of a disk-like system. For a reliable reconstruction of the disk surface density, the rotation curve has to be known globally (which, for observational reasons, is impossible) or at least out to its Keplerian falloff.\\footnote{\\cite{bib:22} report on galaxies with Keplerian rotation curves at large radii} Therefore, the velocity measurements must be supplemented with independent data in order to constrain the sought mass distribution. This may be achieved, for example, by taking into account the amount of hydrogen measured in the outermost galactic regions as was done in \\citep{bib:bratek}.  Mass estimates of spiral galaxies are strongly model-dependent. Rotational velocities, when interpreted in a spherical dark halo model, may greatly overestimate galactic masses as compared to the disk model. It is thus important to consider a phenomenologically acceptable model that gives a lower limit to the mass. The thin disk model may serve as such a reference model.  NGC 4736 is an example of a spiral galaxy with a rotation curve of which outer parts cannot be satisfactorily reconstructed when a spherical halo is assumed. By a simple examination one finds that in the circular orbit approximation, the rotation curve of this galaxy cannot be created by a spherical matter distribution. By applying the thin disk model, with no constraining assumptions about the mass-to-light ratio profile, one can find a mass profile in this galaxy that perfectly conforms  with its rotation curve, and agrees with the amount of hydrogen observed at large radii, beyond the cutoff radius \\citep{bib:bratek}. Only insignificant (if any) amount of dark matter is required, whereas the customary models predict the galaxy to be dark matter dominated. In this paper we report also on another spiral galaxy, NGC 5457, in which the amount of baryonic matter found in the disk model accounts for this galaxy rotation.  \\subsection{The global thin disk model as a means of relating the rotation law to a mass distribution} The main objection against the use of the axisymmetric thin disk as a model of a spiral galaxy, is its instability with too short time scale \\citep{bib:toomre_stab, bib:stab_numer}. The simplest way to stabilize such a system is to add a sufficiently strong spherically symmetric potential. This is usually done by introducing a massive spherical halo of dark matter, which is believed to surround spiral galaxies. However, it should be noted that the galactic interior, irrespectively of its structure, produces at radii sufficiently large almost spherically symmetric potential that may stabilize the external circular orbits.  Putting the stability issue aside, the disk model is still useful for approximate description of the gravitational field of a flattened galaxy. Suppose that the rotation curve of such a galaxy represents the velocity of the streaming motion of matter rotating on roughly circular orbits in the galactic disk. The disk model associates with this rotation curve a formal surface mass density that may be considered as an approximation of the column density of matter in this galaxy. To ensure its applicability, we use the model only for describing galaxies with rotation curves breaking the condition of spherical mass distribution at large distances from their centers. We assume that this non-sphericity can be attributed to flattening of these galaxies. In these regions we therefore expect the disk model to better reflect the oblateness of such galaxies than the models with massive spherical halo. What's more, the customary rotation curve modeling already assumes a resultant thin disk surface density as a superposition of the stellar disk component's surface density and of the column density of the spherical bulge component projected onto the thin disk's plane. Based on these premises the use of the single global thin disk for modeling the whole mass distribution in a flattened galaxy seems justified. The only new qualitative thing is that the  effective disk becomes extended further out to the outer part in the case when the sphericity condition is broken there, that is, when the massive spherical halo cannot be introduced. ",
        "conclusions": "For the accuracy of determination of the mass distribution and its features in the disk component, local non-monotonic properties of rotation curves are important. They are not less important than the value of the rotational velocity. In particular, a local increase in the velocity field may be caused simply by a local decrease of matter density in the disk component and not by the increase of density in the spherical component of a galaxy. The important information about mass distribution carried by rotation curves may be simply overlooked if one assumes that the local mass distribution is proportional to the local brightness. In some cases the dark matter halo is introduced simply because the amount of luminous matter with the constant mass-to-light ratio found by the least square fitting method, cannot account for the rotation at large radii. The example of the spiral galaxy NGC 4736 discussed in \\citep{bib:bratek} very well illustrates this observation. This shows also that determination  of mass distribution in spiral galaxies is model dependent.  For more reliable predictions it is thus important to have a generic model, more flexible and general than the parametric models with smooth profiles and constant mass-to-light ratios. It seems that the global thin disk model provides such a generic model, at least for flattened galaxies with rotation curves breaking at large radii the sphericity condition.  \\medskip Due to observational cutoff in rotation data, the surface mass density reconstruction in the galactic disk (and consequently in the whole galaxy), is subject to high uncertainties, which increase with the distance from the center. Only in particular situations of rotation curves with almost Keplerian tails, the cutoff errors can be eliminated and taken into account. The observationally determined part of a given rotation curve may be explained by various internal disk mass distributions depending on how the rotation law has been extrapolated beyond the cutoff radius. This uncertainty can be removed if additional constraints on the mass distribution in the outer parts, not covered with rotation measurements, are taken into account. In this respect the hydrogen distribution in  the outer regions can be used. We have established for galaxies NGC4736 and NGC5457 that the observed amount of baryonic matter accounts for rotation of these galaxies in the approximation of the global disk model."
    },
    "0809/0809.3373.txt": {
        "abstract": "We develop and demonstrate a probabilistic method for classifying rare objects in surveys with the particular goal of building very pure samples. It works by modifying the output probabilities from a classifier so as to accommodate our expectation (priors) concerning the relative frequencies of different classes of objects.  We demonstrate our method using the {\\em Discrete Source Classifier}, a supervised classifier currently based on Support Vector Machines, which we are developing in preparation for the Gaia data analysis. DSC classifies objects using their very low resolution optical spectra.  We look in detail at the problem of quasar classification, because identification of a pure quasar sample is necessary to define the Gaia astrometric reference frame. By varying a posterior probability threshold in DSC we can trade off sample completeness and contamination. We show, using our simulated data, that it is possible to achieve a pure sample of quasars (upper limit on contamination of 1 in 40\\,000) with a completeness of 65\\% at magnitudes of G=18.5, and 50\\% at G=20.0, even when quasars have a frequency of only 1 in every 2000 objects.  The star sample completeness is simultaneously 99\\% with a contamination of 0.7\\%. Including parallax and proper motion in the classifier barely changes the results. We further show that not accounting for class priors in the target population leads to serious misclassifications and poor predictions for sample completeness and contamination. We discuss how a classification model prior may, or may not, be influenced by the class distribution in the training data. Our method controls this prior and so allows a single model to be applied to any target population without having to tune the training data and retrain the model. ",
        "introduction": "} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %Finding rare objects is hard, because we expect to have to look at many objects %before we encounter one and even then we may not recognise it. Finding rare objects is hard, for two reasons. First, we expect to have to look at many objects before we encounter one, and second, we may not even recognise it even when we do. The reason for this is that the prior probability that any one object is of the rare class, $P(C_{\\rm rare})$, is very small. So even if it has very characteristic features, i.e.\\ the likelihood of the data given the rare object, $P({\\rm Data} \\,|\\, C_{\\rm rare})$, is high, the posterior probability, $P(C_{\\rm rare} \\,|\\, {\\rm Data}$), could still be low.  A survey for rare objects obviously requires a very discriminating classifier, but we could assist it by modifying the class probabilities. If we raised the prior, $P(C_{\\rm rare})$, we are more likely to find the rare objects, but it is inevitable that we will then incorrectly classify other objects as being of the rare class.  Depending on our goals, there may be a satisfactory balance between maximizing the expected number of true positives (sample completeness) and minimizing the expected number of false positives (sample contamination).  Here we present a method for achieving an optimal balance and a correct prediction of this balance which, although simple, is not trivial and has important implications.  We illustrate our method in the context of the problem of detecting quasars in the Gaia survey based on their low resolution ($R \\simeq 30$) optical spectra.  Gaia is an all-sky astrometric and spectrophotometric survey complete to $G=20$, expecting to observe some $10^9$ stars, a few million galaxies and half a million quasars. Its primary mission is to study Galactic structure by measuring the 3D spatial distribution and 2D kinematic distribution of stars throughout the Galaxy and correlating these with stellar properties (abundances, ages etc.) derived from the spectra. With astrometric accuracies as good as 10\\,$\\mu$as, Gaia cannot be externally calibrated with an existing catalogue. Instead it must observe a large number of quasars over the whole sky with which to define its own reference frame.  This quasar sample must be very clean (low contamination).  (The quasar sample is also, of course, of intrinsic interest.)  We present our Gaia classification model, the {\\em Discrete Source Classifier} (DSC), and report its classification performance using simulated data.  We will show how, using our probabilty modification approach, we can use this to build pure quasar samples, at the (acceptable) loss of sample completeness.  {\\bf Related work}. One of the most comprehensive search for quasars to date is that done with the SDSS. Richards et al.~\\cite{richards02} defined a colour locus in the {\\it ugriz} space to identify objects for spectroscopic follow up and estimated the contamination rate of their photometric selection to be 34\\%.  Using the spectroscopy of all point sources taken in SDSS stripe 82, Vanden Berk et al.~\\cite{vandenberk05} assessed the completeness of the Richards et al.\\ selection at 95.7\\% for sources brighter than $i=19.1$. Later, Richards et al.~\\cite{richards04} trained a photometric classifier (based on kernel density estimation) on a set of 22,000 spectroscopically identified quasars and used this to build a sample of 100\\,000 quasars brighter than $g=21.0$.  For the unresolved UV excess quasars in this sample, they estimated the completeness to be 94.7\\% down to $g$=19.5, with a contamination of just 5\\% down to $g$=21.0. Suchkov et al.~\\cite{suchkov05} and Ball et al.~\\cite{ball06} both use decision trees to classify objects in the SDSS photometric catalogue, by training on objects with spectroscopic classifications from SDSS. Gao et al.~\\cite{gao08} used support vector machines and Kd-tree nearest neighbours to classify spectroscopically-confirmed SDSS and 2MASS objects as quasars or stars using just the photometry, with stars and quasars present in equal proportions. With both methods they obtained contamination rates of a few percent (averaged across both classes).  % Wolf et al.~\\cite{wolf01} discussed methods for photometric % classification of objects into quasar, galaxy or star classes and % recovery of redshifts for the extragalactic classes.  Cabanac et % al.~\\cite{cabanac02} developed a photometric classifier based on % principal component analysis using simulated data in anticipation of % several large surveys.  {\\bf Outline}. We start in section~\\ref{theory} by presenting our general prior modification method. In section~\\ref{model} we present the Gaia-DSC classifier and the data used to train and test it. The performance of this and the application of our method are given in section~\\ref{results} and discussed in section~\\ref{discussion}, where we also discuss the interpretation of the probabilities and some of the limitations of this work. We conclude in section~\\ref{conclusions}.  {\\bf Definitions and notation}.  We use the term {\\em class fractions} to mean the relative numbers of objects of each class in a data set. The {\\em nominal model} refers to the classifier ``as is'', i.e.\\ the output probabilities without any modifications. In contrast, in the {\\em modified model} we have modified these outputs to be appropriate to a situation in which there would be very different class fractions, e.g.\\ quasars very rare (section~\\ref{probmod}).  The subscript $i$ refers to true classes and the subscript $j$ to model output classes.  $f^{train}_i$, $f^{test}_i$ and $f^{mod}_i$ refer to the class fractions of class $i$ for the training data, test data and effective test data respectively. They are normalized, $\\sum_i f^{train}_i = \\sum_i f^{test}_i = \\sum_i f^{mod}_i = 1$. This {\\em effective test data set} is one which we don't actually have, but reflects a target population with different class fractions, e.g.\\ quasars very rare. We make predictions of the performance on this by using the actual test data but by modifying the performance equations (section~\\ref{modcalcs}).  In the figures especially, lower case class names, e.g.\\ {\\tt quasar}, denote estimated classes and upper case names, e.g.\\ {\\tt STAR}, true classes. Thus $P({\\tt quasar} | {\\tt STAR})$ means ``the DSC-assigned quasar probability given that the source is truly a STAR''. This refers to DSC outputs for objects with known classes (e.g.\\ in a test set). This is not to be confused with the notation for the DSC outputs in the general case, $P(C_j | x_n, \\theta)$, which means ``the probability which DSC model $\\theta$ assigns to class $C_j$ for object $x_n$''. ``C\\&C'' stands for ``completeness and contamination''.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We have introduced a method of probability modification which can be used to build clean samples of rare objects for which the completeness and contamination can be reliably predicted.  We have demonstrated this using a support vector machine classifier, although it may, of course, be used in conjunction with any classifier which gives probabilities. The main conclusions of our work are as follows.  \\begin{itemize}  \\item To construct a pure sample of objects we should use a probabilistic classifier and only select objects with high probabilities. By varying this probability threshold we can trade off sample completeness and contamination.  \\item To achieve pure samples of rare objects, we must take into account the expected class fractions in the target population, which act as a prior probability on the classifier. We use these to modify the nominal classifier outputs to give the {\\em modified model}.  \\item We applied our modified model to a three class problem in which quasars are simulated to be 1000 times rarer than stars and galaxies.  We can achieve a pure quasar sample (zero contamination) yet still reach a sample completeness of 50--65\\% for magnitudes down to G=20.0. Although the test set is finite in size, this correponds to an upper limit in the contamination of 1 in 39\\,000.  This is more than adequate for establishing an astrometric reference frame for Gaia: If Gaia observes half a million quasars, we can build a quasar sample of 250\\,000 with no more than 13 contaminants.  \\item While achieving this pure quasar sample, we simultanesouly achieve very complete galaxy and star samples (both 99\\%) with low contamination (both 0.7\\%) (figures for G=18.5).  \\item These results were achieved after removing quasars with low equivalent width emission lines from the training sample (defined somewhat arbitrarily as 5000\\,\\AA). Including these precluded establishing a low contamination sample, because it resulted in cool (4000--8000\\,K) highly reddened (\\av\\,=\\,8--10) stars being confused with them. After removing these quasars, the first stars to contaminate a quasar sample (if we set a low threshold) are these reddened cool stars as well as hot stars (\\teff\\,$>$\\,40\\,000\\,K).  \\item Including parallax and proper motion (either as additional SVM inputs, or in a separate mixture model classifier) hardly changes the performance.  This is not surprising since the majority of objects with astrometry consistent with zero are actually stars.  \\item Extending the training and testing sets to include quasars with a full range of interstellar extinction does not significantly alter the results (completeness slightly lower, but contamination unaffected)  \\item All classification models have a prior, but the prior is often not explicit and are sometimes implicitly influenced by the training data distribution. We have introduced a simple method for calculating the implicit priors in a classification model, which we call the {\\em model-based priors}.  In many cases we have experimented with, these priors are close to the true class fractions in the training data (nominal model) or modified class fractions (modified model).  \\item We recommend that a classifier be trained on roughly equal numbers of objects in each class so that it can properly learn the class distributions or boundaries. By determining the model-based priors and replacing them with something more appropriate to the target population (e.g.\\ quasars being rare), we can produce a modified model with superior performance. In particular, this is far better at producing large, pure samples of the rare class.  \\item With our approach we can apply the model to any new target population by specifying the appropriate class fractions (priors) without having to change the training data distribution or re-train the model.  \\end{itemize}    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.0716_arXiv.txt": {
        "abstract": "We consider the early stages of cosmic hydrogen or helium reionization, when ionizing sources were still rare. We show that Poisson fluctuations in the galaxy distribution substantially affected the early bubble size distribution, although galaxy clustering was also an essential factor even at the earliest times. We find that even at high redshifts, a significant fraction of the ionized volume resided in bubbles containing multiple sources, regardless of the ionizing efficiency of sources or of the reionization redshift. In particular, for helium reionization by quasars, one-source bubbles last dominated (i.e., contained $90\\%$ of the ionized volume) at some redshift above $z=7.3$, and hydrogen reionization by stars achieved this milestone at $z>23$. For the early generations of atomic-cooling halos or molecular-hydrogen-cooling halos, one-source ionized regions dominated the ionized volume only at $z>31$ and $z>48$, respectively. To arrive at these results we develop a statistical model for the effect of density correlations and discrete sources on reionization and solve it with a Monte Carlo method. ",
        "introduction": "The earliest generations of stars are thought to have transformed the universe from darkness to light and to have reionized and heated the intergalactic medium. Knowing how the reionization process happened is a primary goal of cosmologists, because this would tell us when the early stars formed and in what kinds of galaxies. The strong fluctuations in the number density of galaxies, driven by large-scale density fluctuations in the dark matter, imply that the dense regions reionize first, producing on large scales an inside-out reionization topology \\citep{BLflucts}. This basic picture has been studied and confirmed with detailed analytical models \\citep{fzh04}, semi-numerical methods \\citep{mesinger}, and by a variety of large numerical simulations \\citep{mellema, zahn, cen} that solve gravity plus radiative transfer. The distribution of neutral hydrogen during reionization can in principle be measured from maps of 21-cm emission by neutral hydrogen \\citep{Madau}, although upcoming experiments such as the Murchison Widefield Array (MWA)\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} and the Low Frequency Array (LOFAR)\\footnote{http://www.lofar.org/} are expected to be able to detect ionization fluctuations only statistically \\citep[for reviews see, e.g.,][]{fob06,bl07}.  The infancy of cosmic reionization, when only a small fraction of the volume of the universe was ionized, is of interest for a number of reasons. First, when ionizing sources were rare at early times, they are expected to have formed separate H~II bubbles which if observed can be used to study directly the properties of individual sources and their surroundings \\citep{cen2}, without the complications of later times, when overlapping bubbles imply that galaxy clustering dominates the ionization distribution and the 21-cm power spectrum.  Second, when ionization fluctuations disappear over much of the universe, it becomes possible to use the 21-cm technique for other applications including those of fundamental cosmology, without the complications of ionization fluctuations which are intrinsically non-linear (since the ionization fraction varies from 0 to 1). Major such applications include measurements of the density power spectrum \\citep{rees1,rees2}, of fluctuations in the \\Lya radiation emitted by the first galaxies \\citep{BL05b,Jonathan06a,Shapiro}, and of fluctuations in the rate of heating from early X-rays \\citep{Jonathan07}. If ionization fluctuations are negligible then the angular anisotropy of the 21-cm power spectrum makes it possible to measure separately various fluctuation sources, including in particular the cosmologically-interesting baryonic density power spectrum \\citep{BL05a}. On small scales, the existence of H~II bubbles (even when rare) affects the fluctuations in \\Lya and X-ray radiation, producing a small-scale cutoff in the 21-cm power spectrum that can be used to detect and study the population of galaxies that formed just 200 million years after the Big Bang \\citep{NB08}.  While analytical models and numerical simulations exist that can be used to study the later epochs of reionization, the early times are very difficult to investigate. Simulations, which in general must overcome the huge disparity between the large characteristic scales of galaxy clustering at high redshift and the small scales of individual galaxies \\citep{BLflucts}, are stretched even further at early times, when ionizing sources become very rare and even larger cosmological volumes are required in order to assemble a reasonable statistical sample. As discussed in detail below, current analytical models based on the model of \\citet{fzh04} account for galaxy clustering but are based on continuous variables and cannot account for the fact that galaxies are discrete sources. This discreteness becomes a crucial factor in the early stages of reionization, when the number of ionizing sources per bubble is small. In this limit, Poisson fluctuations also become substantial, weakening the correlation between the galaxy distribution and the underlying large-scale density fluctuations in the dark matter. Discreteness can also play a significant role during the central stages of reionization, particularly in the case of He reionization by quasars, which are rare sources believed to form only in massive halos that correspond to many-$\\sigma$ density fluctuations at high redshift. These various aspects of discrete sources are not accounted for in current analytical models. \\citet{FurOh} considered helium reionization and showed that the continuous models break down when discreteness is important. They suggested to instead use a pure stochastic Poisson model, without halo correlations, when He is less than $\\sim 50\\%$ ionized globally.  In this paper we develop a model that accounts for discrete sources as well as density correlations. We solve the model with a Monte Carlo method and use it to show that galaxy correlations play a major role even in the infancy of cosmic reionization. Isolated one-source bubbles do dominate at sufficiently high redshifts, but the pure stochastic Poisson model is essentially never a good description of the bubble size distribution. In the next section we first review previous models (section~\\ref{s:prev}), then develop ours (section~\\ref{s:full}) and summarize all the various models whose results we later compare (section~\\ref{s:sum}). We illustrate our results during the infancy of reionization (section~\\ref{s:1perc}) and then develop an approximate calculation that allows us to scan through a wide parameter space of possible reionization scenarios (section~\\ref{s:aprx}). Finally, we illustrate our results during later stages of reionization (section~\\ref{s:late}) and summarize our conclusions (section~\\ref{s:conc}). We assume a standard $\\Lambda$CDM universe with cosmological parameters that match the five-year WMAP data and other large scale structure observations \\citep{wmap}, namely $\\Omega_m=0.28$ (dark matter plus baryons), $\\Omega_\\Lambda=0.72$ (cosmological constant), $\\Omega_b=0.046$ (baryons), $h=0.7$ (Hubble constant), $n=0.96$ (power spectrum index) and $\\sigma_8=0.82$ (power spectrum normalization). ",
        "conclusions": "\\label{s:conc}  We have developed a model of reionization that adds discrete ionizing sources and Poisson fluctuations to the continuous model of \\citet{fzh04}. We have shown how to obtain the distribution of ionized bubbles, versus both bubble size and number of ionizing sources, with a two-step Monte Carlo method that accounts for both density and Poisson correlations among regions of various sizes surrounding a given random point in the universe. The bubble size distribution we obtained differs substantially from previous models, but if the continuous barrier model is cut off below $\\Vb$ (the minimum bubble volume corresponding to a single halo of mass $M_{\\rm min}$) then it yields a reasonable rough estimate to the true bubble size distribution. More specifically, this estimate is generally accurate for H reionization even as early as a mean ionized fraction $\\bar{x}^i = 1\\%$, while for He reionization it works best for small volumes and at later times, and at $\\bar{x}^i = 1\\%$ is accurate only up to $V \\sim 3 \\Vb$. Note that with the cutoff at $\\Vb$, the linear barrier approximation (which can be calculated analytically) gives nearly identical results to the exact continuous barrier.  Our full model yields a bubble distribution by number $N$ that drops more rapidly with $N$ than does the volume distribution drop with $V$, but still, multi-source bubbles are always far more abundant than a pure stochastic Poisson model would suggest. This is due to the fact that density fluctuations are strongly correlated with ionization even when Poisson fluctuations are large. Thus, the density of ionized regions is strongly biased high compared to unconstrained regions, but on the other hand, Poisson fluctuations allow regions to fully ionize themselves even if their density is not as high as would be needed in the continuous barrier model.  The main parameters controlling the relative dominance of single-source bubbles are the effective efficiency $\\zeta$ and the effective slope $n$ of the power spectrum on the scale of a one-source bubble. The ratio of how much harder (in terms of number of $\\sigma$ of the fluctuation) it is to ionize large bubbles compared to small ones, is approximately proportional to $\\zeta^{1+(n/3)}-1$. Reionization by rare sources that are massive and bright corresponds to having a high $\\zeta$ and to a high minimum bubble size, which brings larger scales into play, making the effective $n$ less negative and thus making it harder to produce multi-source bubbles.  We have developed a quick, $15\\%$ accuracy approximate calculation of the ratio $N_{1/2}$ between the total ionized volume and that in multi-source bubbles. This allowed us to sweep through the full parameter space of possible halo masses and efficiencies of the ionizing sources, and to show that sources with a given minimum circular velocity $V_c$ can only achieve a dominance of one-source bubbles at high redshift, regardless of their efficiency or of the reionization redshift. In particular, for He reionization by quasars, one-source bubbles can dominate (i.e., contain $90\\%$ of the ionized volume) only at $z>7.3$, and fill half the ionized volume at $z>4.9$, while H reionization by stars can achieve these milestones only at $z>23$ and $z>18$, respectively (assuming $10 < \\zeta < 1000$). The generation of atomic-cooling halos can place $90\\%$ of the ionized volume in isolated bubbles only at $z>31$ and $50\\%$ at $z>24$, while the earliest generation of molecular-hydrogen-cooling halos can achieve the same only at $z>48$ and $z>36$, respectively.  We note that reality likely includes even more fluctuations than included in our Poisson model, since we have still assumed that the number of ionizing photons emitted from a galactic halo is proportional to its mass. In reality, variations in the ionizing efficiency (through spatial or temporal fluctuations in the star formation efficiency and in the escape fraction of ionizing photons), and in the merger histories of halos of a given mass (even within a given environment, as measured by the average density of a surrounding region) will increase the role of (now generalized) Poisson fluctuations compared to that of galaxy bias due to the underlying large-scale density fluctuations. Simple forms of such variability can be included in a model of the type that we presented, since the ionizing photon outputs from sources are added as individual units (which could be generated from additional distributions for a given halo mass). In general, the model we developed can be used to investigate helium reionization and observational prospects for 21-cm observations during the infancy of hydrogen reionization."
    },
    "0809/0809.2180_arXiv.txt": {
        "abstract": "IGR J17091--3624 and IGR J17098--3628 are two X-ray transients discovered by {\\it INTEGRAL} and classified as possible black hole candidates (BHCs). We present here the results obtained from the analysis of multi-wavelength data sets collected by different instruments from 2005 until the end of 2007 on both sources. IGR J17098--3628 has been regularly detected by {\\it INTEGRAL} and {\\it RXTE} over the entire period of the observational campaign; it was also observed with pointed observations by {\\it XMM} and {\\it Swift}/XRT in 2005 and 2006 and exhibited flux variations not linked with the change of any particular spectral features. IGR J17091--3624 was initially in quiescence (after a period of activity between 2003 April and 2004 April) and it was then detected again in outburst in the XRT field of view during a {\\it Swift} observation of IGR J17098--3628 on 2007 July 9. The observations during quiescence provide an upper limit to the 0.2-10 keV luminosity, while the observations in outburst cover the transition from the hard to the soft state. Moreover, we obtain a refined X-ray position for IGR J17091--3624 from the {\\it Swift}/XRT observations during the outburst in 2007. The new position is inconsistent with the previously proposed radio counterpart. We identify in VLA archive data a compact radio source consistent with the new X-ray position and propose it as the radio counterpart of the X-ray transient. ",
        "introduction": "Among the sources included in the {\\it INTEGRAL} IBIS/ISGRI survey catalog~\\citep{Bird}, the two X-ray transients IGR J17091--3624 and IGR J17098--3628 are remarkable because of their proximity, being only 9\\arcmin.6 away from each other.  The two sources have been detected in different periods of time from April 2003 until the beginning of 2008. In particular, IGR J17091--3624 was detected by IBIS in 2003 and remained detectable for one year and, after a period of quiescence, it was again in outburst in July 2007. IGR J17098--3628 was detected for the first time by {\\it INTEGRAL} in 2005 (while IGR J17091--3624 was not visible) and it has then remained detectable with variable flux in the soft X-ray energy range (E $<$ 20 keV) up to now. On the basis of their spectral behavior, both IGR J17091--3624 and IGR J17098--3628 are classified as probable BHCs \\citep{Luto03,Greb2007}.  In this paper, we present the results of a long term monitoring, primarily at high energy, of these two transients, and discuss the identification in archival data of a radio counterpart for IGR J17091--3624. The date, the exposure time, the instruments used and the sources detection of the different observations are summarized in Table~\\ref{log}. The paper is organized as follows: in \\S\\ref{intro1} and \\S\\ref{intro2}, we introduce the two BHCs; in \\S\\ref{obs}, we describe our observations; in \\S\\ref{res}, we present the results; finally, we give a discussion and present our conclusions in \\S\\ref{disc}.  \\begin{table} \\begin{center} \\caption{IGR J17098-3628 and IGR J17091-3624 observations log from 2003 until September 2007} \\label{log} \\begin{tabular}{lcccc} \\hline \\hline  Start date & End date & Satellite & Detected sources & Exposure \\\\ 2003-04-12 & 2003-04-21 & {\\it INTEGRAL} & IGR J17091-3624\\tablenotemark{a} & 15 ks \\\\ 2003-04-20 & \\nodata & {\\it RXTE} & IGR J17091-3624\\tablenotemark{b} & 4.5 ks  \\\\ 2003-04-23 & 2003-05-09 & VLA\\tablenotemark{c}& IGR J17091/IGR J17098 & 500 s \\\\ 2003-08-10 & 2004-04-02 & {\\it INTEGRAL} & IGR J17091-3624\\tablenotemark{a} & 154 ks  \\\\ 2005-02-20 & 2005-05-01 & {\\it INTEGRAL} & IGR J17098-3628\\tablenotemark{d} & 300 ks \\\\ 2005-03-29 & \\nodata & {\\it RXTE} & IGR J17098-3628\\tablenotemark{d}& 2.1 ks \\\\ 2005-05-01 & 2007-02-28 & {\\it INTEGRAL} & IGR J17098-3628 &  450 ks\\\\ 2005-05-01 & 2005-09-14 & {\\it Swift} & IGR J17098-3628 & 8.5 ks \\\\ 2006-08-25 & \\nodata & {\\it XMM} & IGR J17098-3628  & 8 ks \\\\ 2007-02-19 & \\nodata  & {\\it XMM} & IGR J17098-3628  & 16 ks \\\\ 2007-07-19 & 2007-08-07 & {\\it Swift} & IGR J17091/IGR J17098 & 11 ks \\\\ 2007-08-25 & 2007-09-30 & {\\it INTEGRAL} & IGR J17091-3624 & 50 ks  \\\\ \\hline \\end{tabular} \\tablenotetext{a}{Capitanio et al. 2006} \\tablenotetext{b}{$ $Lutovinov \\& Revnivtsev 2003} \\tablenotetext{c}{Rupen et al. 2003} \\tablenotetext{d}{Grebenev et al. 2007} \\end{center} \\end{table}   \\subsection{IGR J17091--3624} \\label{intro1}  IGR J17091--3624 was discovered by {\\it INTEGRAL}/IBIS during a Galactic Center observation on 2003 April 14--15~\\citep{Kuulk}.  Initially, the flux was $\\sim$20 mCrab in the 40--100 keV energy band exhibiting a hard spectrum, while it was not detected in the 15--40 keV band, with an upper limit of $\\sim$10 mCrab. During subsequent observations of the Galactic Center Deep Exposure (GCDE) on 2003 April 15--16, the source flux increased to $\\sim$40 mCrab in the 40--100 keV band and to 25 mCrab in the 15--40 keV (the IBIS flux statistical error is less then 10\\%).  Immediately after the {\\it INTEGRAL} discovery, providing the position of IGR J17091--3624, an {\\it RXTE\\/} observation was performed and the source was then searched in the X-ray catalogs. IGR J17091-3624 was found in the archival data of both the TTM telescope on board the KVANT module of the {\\it Mir\\/} orbital station \\citep{Atel2}, and in the {\\it BeppoSAX} WFC \\citep{Atel4}. A first study of the IBIS/ISGRI spectral evolution of the source~\\citep{Luto03, Luto05} showed a source hardening with a photon index changing from $\\Gamma = 2.2 \\pm 0.1$ to $\\Gamma=1.6\\pm 0.1$ from 2003 April to 2003 August. A subsequent detailed analysis of the IBIS, JEM-X and {\\it RXTE}/PCA data of the entire outburst duration~\\citep{Cap2006} revealed an indication of an hysteresis like behavior and the presence of a hot disc black body emission component during the source softening.  From the investigations reported above, IGR J17091--3624 appears as a moderately bright variable transient source, with a flaring activity in 1994 October ({\\it Mir}/KVANT/TTM), 1996 September ({\\it BeppoSAX}/WFC), 2001 September \\citep[{\\it BeppoSAX}/WFC,][]{Atel4}, 2003 April \\citep[{\\it INTEGRAL}/IBIS,][]{Kuulk}, and 2007 July. In this paper, we report on the last episode of activity, as well as on limits on the quiescent state.   \\subsection{IGR J17098--3628}  \\label{intro2}  IGR J17098--3628 was detected for the first time with {\\it INTEGRAL}/IBIS 9.4' off IGR J17091--3624 \\citep{Atel444} during deep Open Program observations of the Galactic Center region on 2005 March 24.  The average fluxes were 28.2 $\\pm$ 1.4 and 38.7 $\\pm$ 2.8 mCrab in the 18--45 and 45--80 keV bands, respectively.  Further analysis \\citep{Atel447} reported that the source was evolving in both brightness and spectral shape with an indication of softening.  From 2005 March 29 to April 4, an observational campaign was performed with {\\it RXTE}/PCA. The spectral shape given by {\\it INTEGRAL}/IBIS and {\\it RXTE}/PCA varied throughout the observations and it is modelled by a soft black body emission component plus a hard tail and an absorption consistent with the Galactic one~\\citep{Greb2007}. The spectral variations suggested that this source was an X-ray nova going in outburst and probably a BHC \\citep{Greb2007}.  The source was then observed with the {\\it Swift} satellite \\citep{Atel476} with an exposure time of 2.8 ks on 2005 May 1; it was quite bright, with an estimated flux of 1.3 $\\times$ $ 10^{-9}$ ergs s$^{-1}$ cm$^{-2}$ in the 0.5--10 keV energy band.  The analysis of the {\\it Swift}/XRT data for IGR J17098--3628 refined the source coordinates as follows: RA 17$^h$ 09$^m$ 45.9$^s$, Dec --36$^\\circ$ 27$\\arcmin$ 57$\\arcsec$ (J2000), with an uncertainty radius of about 5$\\arcsec$ \\citep{Atel476}. This position is 30$\\arcsec$ from the {\\it INTEGRAL} position reported by \\cite{Atel444}.    Following the soft X-ray detection of IGR J17098--3628 with {\\it Swift}/XRT, a probable radio counterpart has also been found \\citep{Atel490}. In particular, the first data set of four consecutive Very Large Array (VLA) radio observations, made on 2005 March 31, April 5, April 12, and May 4 at 4.86 GHz, showed only one significant radio source within the 2$\\arcmin$ {\\it INTEGRAL} error circle, located at RA 17$^h$ 09$^m$ 45.934$^s$ $\\pm$ 0.011s, Dec --36$^\\circ$ 27$\\arcmin$ 57.30$\\arcsec~\\pm$ 0.55$\\arcsec$ (J2000) \\citep{Atel490}.  Thanks to the radio observations, a probable optical/infrared identification has been found within the 2MASS All-Sky Catalog and the SuperCOSMOS Sky Survey~\\citep{Atel478,Atel479,Atel490,Atel494}.  On the base of the optical identification, \\citet{Greb2007} estimated the source distance as ${\\it d}$= 10.5 kpc and an upper limit on the inclination angle: ${\\it }\\geq$ 77$^{\\circ}$. ",
        "conclusions": "\\label{disc} The huge amount of data collected and the broad energy coverage allowed us to perform a detailed spectral analysis of these two sources, following their evolution during more than two years and fixing upper limits when one of the sources was not visible. In the following a detailed summary of the results for both sources can be found.  \\subsection{IGR J17098--3628} Our analysis, spanning from May 2005 to September 2007, proves that the source spectrum shows a soft black body component with an internal temperature of about $\\sim$ 1 keV and an internal radius comparable to the last stable orbit of the accretion disc. The luminosity due to soft component \\citep[assuming $d=10.5$ kpc,][]{Greb2007} varied during the two years of our observational campaign from $\\sim 2 \\times 10^{37}$ to $\\sim 5 \\times 10^{36}$ ergs~s$^{-1}$. The flux variation range and the spectral parameters are comparable with the ones reported by \\citet{Greb2007} during the first phases of the outburst on April 2005.  In our data set each flux variation seems not to be correlated with any relevant spectral features, as confirmed by the lack of evident variation of the {\\it RXTE}/ASM hardness-intensity diagram (Figure~\\ref{2005Camp}).  However respect to the previous April 2005 observations, we did not detect any power law emission at higher {\\it INTEGRAL}/IBIS energy band. This indicates that the spectral shape was evolved from the first phase of the outburst being dominated only by a thermal disc component coming from the accretion disc around the compact object.  We can conclude that this source  seems  to have spent 2.5 years in the high soft state with a disc black body component substantially identical to the one previously observed at the beginning of the outburst~\\citep{Greb2007}. Curiously the only difference is the lack of any high energy emission. In fact the hard component fell below the detection limit of IBIS after about three months from the beginning of the outburst. Hence the geometry and the temperature of the accretion disc have not shown any significant variation up to now, on the other end the power law emission quenched probably because of the electron temperature of the corona fallen below the disc seed photons temperature making the inverse Compton scattering processes inefficient.   \\subsection{IGR J17091--3624}  Previous studies have demonstrated that the X-ray luminosity of the BHCs in quiescence is lower than that of neutron star X-ray binaries and falls below $\\sim$ 10$^{32}$-10$^{33}$ ergs~s$^{-1}$ \\citep{Campana}. In the quiescent state, neutron star X-ray binaries are normally detected and hence the {\\it XMM} upper limit of IGR J17091--3624 is a good indication that the source is a BHC.  The {\\it Swift}  ToO caught the source at the beginning of its outburst when it was still in a low/hard state. The best fit of the first observation  (2007-07-09/16) is a power law spectrum  with $\\Gamma$=1.4$\\pm$ 0.1. The source spectrum substantially softened and the black body component became dominant with an increasing temperature and a steeper power law. {\\it INTEGRAL} continued to monitor the source with IBIS  after the end of the XRT observational campaign, up to the end of September 2007. These observations confirm the softening of the source. A power law component without an energy cutoff provided the best fit of the IBIS spectra between 20 and 80 keV, probably due to jets or corona reprocessing of disc seed photons. The power law photon index varied during two month period from about 1.4 (beginning of July) to about 3 (end of September). Unfortunately the relatively low source flux at high energies (7 $\\times$ 10$^{-10}$ ergs~cm$^{-2}$ s$^{-1}$ between 20-100 keV) and the variation of the spectral shape did not provide sufficient statistics to extend the spectra up to 80 keV and to verify the presence or not of an high energy cutoff in the spectrum.  Thanks to the XRT refined position, we found the radio counterpart of the source. As soon as 9 days after the detection by IBIS in 2003, a radio source at the sub-mJy level was detected at 5 GHz. The source has been detected also at 8.4 GHz, and it showed an increase in flux over the subsequent two weeks. The spectrum is inverted, characteristic of self-absorbed synchrotron radiation from a compact jet. This behavior is typical of BHC in low hard states. Although it is difficult to guess the spectral shape at higher frequencies, it is reasonable to estimate that the total radiative luminosity of the compact jet is of the order of 10$^{31}$ erg sec$^{-1}$, assuming the distance to the Galactic Center. No information is available after the 2007 outburst, and it will be clearly valuable to obtain new radio observations after future episodes of activity.  Chaty et al. 2008, on the basis of ESO NTT observations, report on two possible infrared counterparts of IGR J17091--3624 consistent with the {\\it Swift}/XRT error box. Our refined position, based on the radio observations, could support the identification of the real infrared counterpart of this source.  Little is known about nature, duration and recurrence of outbursts in transient X-ray binaries. It is probable that the outbursts are due to randomly acting factors such as the mass transfer variations or truncation of the inner disc radius. However for IGR J17091--3624 five outbursts are known from 1994 to 2007. Thus, it appears that this source goes in outburst every three or five years. Generally the flux varies from 5 mCrab to about 20 mCrab in the range 2-10 keV over a period of few months.  The 2003 outburst was the first detected  at higher energy range (20-150 keV) for one year period. Comparing the two hard states of respectively the 2003 and 2007 outbursts, the last one seems to be the hardest. In fact it was possible to determine the high energy power law cutoff of the 2003 hard state \\citep[$\\sim$ 49 keV,][]{Cap2006}  while in the 2007 hard state the cutoff is not detectable up to 80 keV. Concerning the soft state the 2003 outburst has an higher temperature black body emission~\\citep{Cap2006} that does not appear to be present in the 2007 outburst. However the XRT monitoring campaign observed the source only at the beginning of its transition to the soft state and unfortunately in the subsequent {\\it INTEGRAL} monitoring the source was outside the JEM-X field of view and so it was not possible to follow the entire evolution of the black body temperature."
    },
    "0809/0809.2463_arXiv.txt": {
        "abstract": "In May 2007 the compact radio source Sgr\\,A* was observed in a global multi-frequency monitoring campaign, from radio to X-ray bands. Here we present and discuss first and preliminary results from polarization sensitive VLBA observations, which took place during May 14-25, 2007. Here, Sgr\\,A* was observed in dual polarization on 10 consecutive days at 22, 43, and 86\\,GHz. We describe the VLBI experiments, our data analysis, monitoring program and show preliminary images obtained at the various frequencies. We discuss the data with special regard also to the short term  variability. ",
        "introduction": "There is overwhelming evidence that  Sagittarius\\,A* (Sgr\\,A*), the extremely compact radio source at the center of our Galaxy, is associated with a super-massive black hole.  It shows sudden bursts of radiation (called flares) a few times per day, which are thought to be related to the guzzling of gas and dust from its environment and which can be detected in X-ray and near infra-red (NIR) regime (e.g., \\cite{eckart08} and references therein). To search for correlated short term variability and to investigate the flares in as many wavebands as possible, a world wide radio--sub-mm--NIR--X-ray observing campaign on Sgr\\,A* has been carried out in May 2007, making use of many telescopes including in the radio bands the VLBA (Very Long Baseline Array), GMVA (The Global Millimeter VLBI Array), CARMA (Combined Array for Research in Millimeter-Wave Astronomy), ATCA (Australia Telescope Compact Array), IRAM (Institut de Radio Astronomie Millimetrique), and Effelsberg.  The VLBA observed Sgr\\,A* at 22, 43, and 86\\,GHz on 10 consecutive days. These data are used to measure the source size and to search for possible deviations from point-like (or symmetric) structure,  and variations of the source on time scales of $\\sim$ 1 week. Since the intrinsic source size and shape is masked by the interstellar scattering at centimeter wavelengths, the combination of our 3 observing frequencies is of great importance for the size determination and is crucial for discriminating interstellar broadening from the source intrinsic structure of Sgr\\,A* by the known scattering law (e.g., \\cite{bower06}).  Here we show and discuss preliminary images obtained with the VLBA at the 3 frequencies and focus our discussion on the results mainly obtained at 43\\,GHz. Other results regarding to the total flux density measurements in this global campaign are presented in other papers (see Kunneriath \\textit{et al}, Eckart \\textit{et al}, this conference). ",
        "conclusions": "We have shown images and parameters from 10 epochs of high frequency VLBI observations of Sgr\\,A* during a multi-frequency campaign in May 2007. In particular, we have measured the flux density variations and shown the size measurement at 43\\,GHz. Our result indicate that the measured size at 43\\,GHz is stable on a daily basis. Future studies will show if there is a correlation between the NIR flaring activities and the radio flux density variations. The radio (cm-mm-wavelengths) flux of Sgr\\,A* is known to be variable on timescales of months to hours, with variability being more pronounced at shorter wavelengths. The time scales of the variations indicate that the higher frequency emission originates from a very compact region (e.g.,~\\cite{Miya}, ~\\cite{Mau}). The mm-variability naturally raises the question of whether or not the flux density variations are accompanied by the structural change, which could be detected with VLBI at 86\\,GHz. Previous size measurements at 86\\,GHz may have already indicated possible structural variability, where the increase in size seems to correlate with a higher flux density level~\\cite{kri06}. However, these variations of the source size is at present uncertain and still requires confirmation. As such, the 86\\,GHz VLBI measurements from this campaign will be very crucial for the possible structural variability.  \\ack % R-S Lu and D Kunneriath are members of the International Max Planck Research School (IMPRS) for Astronomy and Astrophysics. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."
    },
    "0809/0809.0650_arXiv.txt": {
        "abstract": "{N-body simulations have shown that the dynamical decay of the young ($\\sim 1$ Myr) Orion Nebula cluster could be responsible for the loss of at least half of its initial content of OB stars. This result suggests that other young stellar systems could also lose a significant fraction of their massive stars at the very beginning of their evolution. To confirm this expectation, we used the Mid-Infrared Galactic Plane Survey (completed by the {\\it Midcourse Space Experiment} satellite) to search for bow shocks around a number of young ($\\la$ several Myr) clusters and OB associations. We discovered dozens of bow shocks generated by OB stars running away from these stellar systems, supporting the idea of significant dynamical loss of OB stars. In this paper, we report the discovery of three bow shocks produced by O-type stars ejected from the open cluster NGC\\,6611 (M16). One of the bow shocks is associated with the O9.5Iab star HD165319, which was suggested to be one of ``the best examples for isolated Galactic high-mass star formation\" (de Wit et al. 2005). Possible implications of our results for the origin of field OB stars are discussed.} ",
        "introduction": "Dynamical three- and four-body encounters between stars in the dense cores of young star clusters impart to some stars velocities exceeding the escape velocity from the cluster's potential well (Poveda et al. \\cite{po67}; Leonard \\& Duncan \\cite{le90}). The concentration of massive stars towards the center of young clusters (e.g. Hillenbrand \\& Hartmann \\cite{hi98}; Gouliermis et al. \\cite{go04}; Chen et al. \\cite{ch07}) implies that the dynamical evolution of cluster cores is dominated by massive stars and that the majority of escaping stars should be massive. The escaping massive stars constitute a population of OB field stars (Gies \\cite{gi87}). The escape velocity ranges from several ${\\rm km} \\, {\\rm s}^{-1}$ (for loose and low-mass clusters) to several tens of ${\\rm km} \\, {\\rm s}^{-1}$ (for compact and massive ones), so that the typical space velocity of OB field stars is $\\sim 10 \\, {\\rm km} \\, {\\rm s}^{-1}$ (e.g. Gies \\cite{gi87}). The stars with peculiar (transverse) velocities exceeding $30 \\, {\\rm km} \\, {\\rm s}^{-1}$ are called runaway stars (Blaauw \\cite{bl61}). It should be realized, however, that any star unbound from the parent cluster should be considered as a `runaway', independently of its peculiar velocity (cf.  de Wit et al. \\cite{de05}). Although the peculiar velocity of some runaway OB stars could be as high as several $100 \\, {\\rm km} \\, {\\rm s}^{-1}$  (e.g. Ramspeck et al. \\cite{ra01}) or even $\\sim 500-700 \\, {\\rm km} \\, {\\rm s}^{-1}$ (Edelmann et al. \\cite{ed05}; Heber et al. \\cite{he08}), the majority have velocities of several tens of ${\\rm km} \\, {\\rm s}^{-1}$. Therefore, most runaway OB field stars should be located not far from their birth clusters.  A (massive) star moving supersonically with respect to the ambient medium generates a bow shock (Baranov et al. \\cite{ba71}; Van Buren \\& McCray \\cite{va88}) visible in infrared (Van Buren et al. \\cite{va95}; Noriega-Crespo et al. \\cite{no97}; France et al. \\cite{fr07}) and/or ${\\rm H}_{\\alpha}$ (Kaper et al. \\cite{ka97}; Brown \\& Bomans \\cite{br05}). Observations of runaway OB stars showed that only a small fraction ($\\la 20$ per cent) produce bow shocks (Van Buren et al. \\cite{va95}; Noriega-Crespo et al. \\cite{no97}; Huthoff \\& Kaper \\cite{hu02}). The main reason for this is that many runaway stars move through the hot interstellar gas and their peculiar velocities are lower than the sound speed in the ambient medium (Huthoff \\& Kaper \\cite{hu02}). The geometry of a bow shock allows us to infer the direction of motion of the associated star and thereby trace its trajectory backwards to the parent cluster in those cases in which the proper motion of the star has not been measured directly.  N-body simulations by Pflamm-Altenburg \\& Kroupa (\\cite{pf06}) showed that the dynamical decay of the young ($\\sim 1$ Myr) Orion Nebula cluster could be responsible for the loss of at least half of its initial content of OB stars (see also Kroupa \\cite{kr04}; cf. Aarseth \\& Hills \\cite{aa72}). This result suggests that other young stellar systems could also lose a significant fraction of their massive stars at the very beginning of their evolution. We therefore expect to observe numerous signatures of interaction between ejected stars and the dense environment of young star clusters. To confirm this expectation, we searched for bow shocks around a number of young ($\\la$ several Myr) clusters and OB associations using the Mid-Infrared Galactic Plane Survey. This survey, carried out with the Spatial Infrared Imaging Telescope onboard the {\\it Midcourse Space Experiment (MSX)} satellite (Price et al. \\cite{pr01}), covers the entire Galactic plane within $|b|<\\pm 5\\degr$ and provides images at $18\\arcsec$ resolution in four mid-infrared spectral bands centered at 8.3~$\\mu$m (band A), 12.1~$\\mu$m (band C), 14.7~$\\mu$m (band D), and 21.3~$\\mu$m (band E). To search for possible optical counterparts to the detected bow shocks, we used the SuperCOSMOS H-alpha Survey (SHS; Parker et al. \\cite{pa05}) and the Digitized Sky Survey (McLean et al. \\cite{mc00}). Our choice of young stellar systems implies that the majority of ejected stars should not have traveled far from their birth places (so that the parent clusters could be identified with more confidence) and that supernova explosions in massive binaries do not significantly contribute to the production of runaway stars (Blaauw \\cite{bl61}).  We discovered dozens of bow shocks. Three are discussed in this paper, while the results of study of other objects are presented elsewhere. The geometry of the three bow shocks suggests that they are driven by stars expelled from the young open star cluster \\object{NGC\\,6611}. In Sect.\\,2, we review the relevant data for this cluster. The results of the search for bow shocks around NGC\\,6611 are presented in Sect.\\,3 and discussed in Sect.\\,4. Conclusions are summarized in Sect.\\,5. ",
        "conclusions": "With the detection of three bow-shock-producing stars in the vicinity of NGC\\,6611, we have found strong evidence that the proposed loss of massive stars, due to dynamical processes in the early evolution of young clusters, is indeed operating. While existing astrometric data for the three stars do not allow us to determine precisely the timing of the ejection or study the interaction of the runaway stars with the ISM in detail, the morphology of the bow shocks and the proper motion of the cluster itself allows us to argue that these stars originated in NGC\\,6611.  Clearly, our finding opens up a number of possibilities for follow-up investigations of the NGC\\,6611 bow shocks and the ejected stars. More accurate proper motion measurements will allow us to check if the stars were ejected continuously or during a short period of cluster evolution. The most important point nevertheless is that the idea proposed by Pflamm-Altenburg \\& Kroupa (\\cite{pf06}) is supported by direct observation, at least in the case of NGC\\,6611. The results for several other young clusters will be presented in forthcoming papers."
    },
    "0809/0809.1183_arXiv.txt": {
        "abstract": "The strong equivalence principle is extended in application to averaged dynamical fields in cosmology to include the role of the average density in the determination of inertial frames. The resulting cosmological equivalence principle is applied to the problem of synchronisation of clocks in the observed universe. Once density perturbations grow to give density contrasts of order one on scales of tens of megaparsecs, the integrated deceleration of the local background regions of voids relative to galaxies must be accounted for in the relative synchronisation of clocks of ideal observers who measure an isotropic cosmic microwave background. The relative deceleration of the background can be expected to represent a scale in which weak--field Newtonian dynamics should be modified to account for dynamical gradients in the Ricci scalar curvature of space. This acceleration scale is estimated using the best--fit nonlinear bubble model of the universe with backreaction. At redshifts $z\\lsim0.25$ the scale is found to coincide with the empirical acceleration scale of modified Newtonian dynamics. At larger redshifts the scale varies in a manner which is likely to be important for understanding dynamics of galaxy clusters, and structure formation. Although the relative deceleration, typically of order $10^{-10}$ms$^{-2}$, is small, when integrated over the lifetime of the universe it amounts to an accumulated relative difference of 38\\% in the rate of average clocks in galaxies as compared to volume--average clocks in the emptiness of voids. A number of foundational aspects of the cosmological equivalence principle are also discussed, including its relation to Mach's principle, the Weyl curvature hypothesis and the initial conditions of the universe. ",
        "introduction": "The strong equivalence principle (SEP) stands at the conceptual core of general relativity, as a physical principle which limits the choice of our physical theory of gravitation among all possible metric theories of gravitation one can construct. In this paper I will argue that the ramifications of this principle have not been fully explored, and that its physical interpretation requires further clarification to deal with the dynamical properties of spacetime inherent in Einstein's theory when the nonequilibrium situation is considered. In particular, the problem of how to synchronise clocks in the absence of a spacetime background with specific symmetries does not have a general solution in general relativity. In this paper I will show that at least for universes which began with a great deal of symmetry, as ours did, the reasoning of the equivalence principle can be extended: the average regional density provides a clock in expanding regions. This has particular consequences for cosmological models and the definition of gravitational energy. It underlies the author's proposal that dark energy is a misidentification of cosmological gravitational energy gradients in an inhomogeneous void--dominated universe \\cite{clocks,sol,essay}. Broader foundational consequences may also follow.  To set the scene, it pays to recall that historically the equivalence principle\\cite{eep}, and indeed general relativity\\cite{GR}, was formulated at a time before the dynamical properties of spacetime were understood. The conceptual route that Einstein took began with the {\\em weak equivalence principle} or the {\\em principle of uniqueness of free fall}, known since the experiments of Galileo, that all bodies (subject to no forces other than gravity) will follow the same paths given the same initial positions and velocities. Realising that this observational statement embodies a feature of universality of the gravitational interaction, Einstein created a theory in which gravity is a property of spacetime itself.  Einstein's identification of what the true gravitation field should be began 101 years ago with first identifying what it is not, based on thought experiments with elevators, concerning what motions of particles cannot be distinguished operationally. His 1907 principle of equivalence \\cite{eep} may be translated as follows: {\\em All motions in an external static homogeneous gravitational field are identical to those in no gravitational field if referred to a uniformly accelerated coordinate system}. A uniformly accelerated reference frame may be found operationally in empty Minkowski spacetime by firing rockets; if matter is the source of the true gravitational field then such choices of frame cannot represent gravity.  More generally, since special relativity with nongravitational forces appears to always be valid in small regions, one should always be able to get rid of gravity near a point. This is embodied in the SEP: {\\em At any event, always and everywhere, it is possible to choose a local inertial frame such that in a sufficiently small spacetime neighbourhood all nongravitational laws of nature take on their familiar forms appropriate to the absence of gravity, namely the laws of special relativity}. This means that gravity is made to be universal, as it is contained in spacetime structure. The true gravitational field strength is encoded in the Riemann curvature tensor, and is determined regionally by the tidal effects of geodesic deviation.  One of the most profound and difficult consequences of the SEP is that gravitational energy and momentum cannot be described by a local density, and so are not local quantities. General relativity overcomes the nonlocality problem of Newtonian gravity: there is no action at a distance. General relativity is an entirely local theory in the sense of {\\em propagation} of the gravitational interaction, which is causal. However, the background on which the interaction propagates may contain its own energy and momentum, when integrated over sufficiently large regions, and this has to be understood in the calibration of local rods and clocks at widely separated events.  Whereas the calibration of rods and clocks is mathematically determined by invariants of the {\\em local} metric, and the spacetime connection which relates local invariants at widely separated events, in practice we cannot analytically solve Einstein's equations for the most general distribution of matter to unambiguously determine the metric and its connection. A slicing of spacetime into hypersurfaces, and a threading of these hypersurfaces by timelike worldlines of observers or by null geodesics, is inevitable for any operational description of spacetime in terms of rods and clocks. Such splittings of space and time, together with additional symmetry assumptions, are also necessary for analytically solving Einstein's equations in particular cases, or more generally for numerical modelling.  The definition of quasilocal gravitational energy and momentum \\cite{quasi_rev} then turns out to depend on the choice of slicing, associated surfaces of integration, and the identification of observers that thread the slices. These procedures are inherently noncovariant and nonunique, and many questions of naturalness of any particular definition inevitably arise. (See Ref.\\ \\cite{NSV} for a recent discussion.) We have the dilemma that a spacetime split inevitably breaks any given particle motion into a motion of the background and a motion with respect to the background; and this may involve a degree of arbitrariness. The viewpoint that will be adopted here is that since quasilocal gravitational energy gradients have their origin in the equivalence principle, the primary criterion for making canonical identifications of physically relevant classes of observer frames is that the equivalence principle itself must be properly formulated, and respected, when making macroscopic cosmological averages.  If the SEP is applied to macroscopic objects then strictly speaking we can only apply it to systems in which the gravitational interaction is ignored. Yet we implicitly apply the SEP to scales at least as large as galaxies which are treated as particles of dust in an expanding fluid -- with an expansion rate given by a Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) model -- subject to additional Newtonian gravitational interactions. The Newtonian approximation, which is made in present--day numerical simulations of structure formation must, by the rules of general relativity, presuppose that a sufficiently large static Minkowski frame can be found.  I shall argue in this paper that to correctly derive a Newtonian limit, without prior assumptions about the cosmological background geometry, we must first correctly apply the SEP. If galaxies are to be treated as particles of dust we must address the following question: given that the background is not static, what is the largest scale on which the SEP can be applied? An attempt to answer this question, which is unavoidable if we are to consistently apply the principles of general relativity to cosmology, means that we have to deal with the relationship of inertial frames to averages of matter fields and motions. In other words we must address Mach's principle, which may be stated as \\cite{Bondi,BKL} follows: {\\em``Local inertial frames are determined through the distributions of energy and momentum in the universe by some weighted average of the apparent motions''}. The leading questions in this statement are what is to be understood by ``local'', and what is the ``suitable weighted average''?  The problem with any process of averaging in general relativity is that by coarse graining we can lose information about the calibration of local rods and clocks within the coarse grained cells relative to average quantities. Rods and clocks are related to invariants of the local metric, can vary greatly within an averaging cell, and will not in general coincide with some average calibration of rods and clocks once averaging cells become sufficiently large in the nonlinear regime of general relativity. Galaxies, for example, contain supermassive black holes, in whose local neighbourhood the determination of proper lengths and times for typical observers differs extremely from typical observers in the outskirts of a galaxy.  It is commonly believed that as long as we are ``in the weak--field limit'' we do not have to worry about complications such as the extreme ones posed by black holes. However, the weak--field limit is always taken about a {\\em background}, and once inhomogeneities develop in the universe there are no exact symmetries to describe the background. One set of uniformly calibrated rods and clocks is no longer sufficient to describe the background itself. Given an initially homogeneous isotropic universe with small scale--invariant density perturbations, once one is within the nonlinear regime of structure formation, homogeneity and isotropy are only defined in a statistically average sense.  Within a cell of statistical homogeneity -- which can be taken to be \\cite{clocks} of the same order as the baryon acoustic oscillation (BAO) scale, $100h^{-1}$Mpc, $h$ being the dimensionless parameter related to the Hubble constant by $H\\Z0=100h\\kmsMpc$ -- there are density contrasts of order unity over scales of tens of megaparsecs. Since the universe is inhomogeneous over these scales {\\em not every observer is the same average observer}. Different classes of canonical observers are required to interpret cosmological parameters \\cite{clocks}. In particular, we and the objects we see are typically in galaxies in locally nonexpanding regions which formed from density perturbations which were greater than the critical density. In such regions local spatial curvature can differ markedly from the volume--average location in voids, where space is freely expanding. Small differences of spatial curvature and rates of clocks can accumulate to give large differences over the lifetime of the universe \\cite{clocks}. Dynamical gradients in the curvature of space are a physical reality which must be properly understood.  In this paper I will estimate the scale of the small relative deceleration of the background in regions of different density, by proposing an extension of the SEP to cosmology. I will demand an equivalence between the particular example of geodesic flow of a congruence of particles ``at rest'' in a dynamically expanding universe, whose average volume expansion is governed by an average over all masses and motions of the particles, and the equivalent ``volume--expanding'' motion of point particles in a Minkowski space. In other words, the local Ricci scalar curvature due to the volume average of pressureless dust can always be ``renormalised away'' on sufficiently small scales, but in a way which may lead to relative recalibrations of local rods and clocks between different spacetime regions.  Equivalently, for volume expansion we cannot {\\em locally} distinguish between particles at rest in a dynamic spacetime, and particles moving in a static spacetime: the two situations are equivalent in a sense which deserves the designation {\\em cosmological principle of equivalence}. This might be viewed as a further Machian style refinement of the notion of inertia. Although we measure geodesic deviation, in terms of the scalar curvature part of the Riemann tensor and volume expansion, we are unable to distinguish whether the geodesics are deviating because of local accelerations of particles in a static space, or whether the particles are ``at rest'' in an expanding space which is decelerating due to the gravitational attraction of the average density of matter.  Historically, it might be said that although Einstein was conceptually guided by Mach's principle, he never quite succeeded in fully implementing it in general relativity, because when he first formulated the theory he did not fully appreciate the dynamical nature of spacetime. His first attempt to study cosmology indicated that for any usual source of matter the theory was not stable, but intrinsically dynamical \\cite{esu}. Famously, he invoked the cosmological constant -- his {\\em``gr\\\"o{\\ss}te Eselei''} -- to avoid the issue. This paper attempts to lay the conceptual groundwork for an alternative first principles route to the physical interpretation of cosmological general relativity, taking account of observational evidence that is immeasurably better now than it was in 1917.  The plan of the paper is as follows. Preliminary definitions, motivations, and a statement of the cosmological equivalence principle are presented in Sec.\\ \\ref{aep}. In Sec.\\ \\ref{gedanken} the thought experiments introduced in Sec.\\ \\ref{aep} are generalised to the case of regions of different density, and an estimate of clock rate variance is given based on presently observed density contrasts. The definition of the cosmic rest frame, and its relation to other frames used in cosmological averaging, is clarified in Sec.\\ \\ref{frame}. In Sec.\\ \\ref{alpha} a numerical estimate of the time--varying relative deceleration of voids relative to walls, where galaxies are located, is computed over the lifetime of the universe, and its cosmological implications discussed. The role of initial conditions and a possible conceptual relationship to Penrose's Weyl curvature hypothesis are discussed in Sec.\\ \\ref{Weyl}. The paper concludes with a summary discussion, including some further speculations, in Sec.\\ \\ref{conclude}. ",
        "conclusions": "}  In this paper I have extended the strong equivalence principle to account for the average effect of the density of matter in the definition and relative calibration of clocks in inertial frames on cosmological scales. Since the resulting cosmological equivalence principle relates the single scalar degree of freedom of Newtonian gravity to the framework of general relativity, it may provide a means to better understand the calibration of cosmological weak fields once density perturbations have grown large to form a universe that is very inhomogeneous on scales of tens of megaparsecs. It should thereby give a setting for better understanding the Newtonian limit in the dynamical situation of cosmology. The numerical estimate of the relative deceleration between observers in the walls around galaxy clusters and volume-average observers in voids, typically of order $10^{-10}$ms$^{-2}$, is acceptably small for weak field scales and yet leads cumulatively to the present epoch clock rate variance of 38\\% found in Refs.\\ \\cite{clocks,LNW}. Intriguingly, at redshifts $z\\lsim0.25$ the relative deceleration required by the CEP coincides precisely with the empirical acceleration scale of MOND \\cite{mond,McGaugh}.  At a conceptual level I have attempted to present a framework for the consistent definition of average inertial frames in relation to average dynamically--varying matter densities in cosmological general relativity. The hope is that the cosmological equivalence principle is thereby a key step to the incorporation of Mach's principle into general relativity, in the way that Einstein intended but never quite realised. Mach's principle is most commonly invoked in distinguishing inertial frames from rotating frames \\cite{BKL,BKL2,Schmid,wave}. What is studied here is a different aspect of Mach's principle: the role of the average volume deceleration of the local geometry in defining the standard of time of inertial frames. In the absence of a timelike Killing vector, the evolution of the average density provides a relevant clock. To fully incorporate Mach's principle in general relativity, it is of course necessary to deal with the other dynamical gravitational degrees of freedom which can affect the distinction of inertial frames from rotating ones, such as gravitational waves \\cite{wave}.  In this paper I have expounded the view that to deal with the volume--contracting average dynamical effects of matter density, a reduction to a frame (\\ref{cif}) is the relevant step in the normalisation of gravitational energy before the final step to a static Minkowski space. It is quite possible that when average degrees of freedom in addition to the scalar Ricci curvature are considered, in order to deal with gravitational waves and spinning matter fluids, there are other steps in the relevant relative calibrations of rods and clocks. After all, energy, momenta, and angular momenta are only defined with respect to a frame. In the dynamical regime of general relativity the question arises as to which collective average frames have physical utility in the absence of exact symmetries described by Killing vectors. My view is that a truly deep understanding of quasilocal gravitational energy and momentum is still to be found, but the path to such enlightenment requires a better conceptual understanding of the equivalence principle in application to collective dynamical degrees of freedom of matter fields.  Historically speaking, in the early stages of the development of general relativity Einstein did not fully appreciate the dynamical importance of the energy and momentum of spacetime itself. Spacetime is inevitably dynamical for matter obeying the strong energy condition. Einstein's first journey through the conceptual landscape of cosmological general relativity had him worrying about boundary conditions at spatial infinity \\cite{esu}, as he overlooked the possibility that the universe had a beginning.  Since general relativity is causal the geometry at any event can only depend on events within its past light cone, and is independent of what lies beyond the particle horizon. Thus boundary conditions at spatial infinity beyond the particle horizon are physically irrelevant if the universe had a beginning, a possibility that Einstein did not consider when he first formulated his static universe \\cite{esu}. For a universe like ours which had a beginning, the initial conditions are of vital importance in determining the relevant weighted average of the apparent motions, as I have discussed in Sec.\\ \\ref{Weyl}.  The conceptual journey discussed in this paper arose in an effort to model the universe more realistically \\cite{clocks,sol,LNW}, to account for the structure we actually observe, by realising that the quasilocal gravitational energy of a dynamical spacetime geometry -- which has real effects on the calibration of clocks -- should be an essential feature of a universe with large dynamical density gradients. If successful, this will eliminate the need for a cosmological constant or other fluidlike vacuum energy as the source of ``dark energy'', but it still leaves the other cosmological problem, why $\\Lambda=0$, unsolved. If we take the strong equivalence principle literally then $\\Lambda$ must be zero since otherwise we could not have a vacuum Minkowski spacetime for our local inertial frames. My own personal view is that quantum field theoretic calculations based in a flat spacetime which suggest that $\\Lambda\\goesas M\\ns{Planck}^4$ miss the mark, because the spacetime vacuum cannot be understood without accounting for the intrinsically dynamic nature of spacetime. It is not a problem for flat space quantum field theory. While the cosmological constant problem is no doubt a problem for quantum gravity, I believe that quantum gravity research might benefit from a more physical understanding of dynamical gravitational energy and the equivalence principle.  \\medskip {\\bf Acknowledgement} I wish to thank many people for discussions and correspondence including in particular Thomas Buchert, Syksy R\\\"as\\\"anen, Frederic Hessman, Herbert Balasin, Helmut Rumpf, Peter Aichelburg, Lars Andersson and Stacy McGaugh. I warmly thank Prof.\\ Remo Ruffini and ICRANet for support and hospitality while the paper was completed."
    },
    "0809/0809.4590_arXiv.txt": {
        "abstract": "We are currently involved in a multifaceted campaign to study extragalactic classical novae in the Local Group and beyond.  Here we report on-going results from the exploitation of the POINT-AGAPE M31 dataset; initial results from our Local Group imaging, and spectroscopic CNe follow-up campaign and introduce the Liverpool Extragalactic Nova Survey. ",
        "introduction": "Although much has been learnt from the study of Galactic classical novae (CNe), it is clear that Galactic data are not ideal for establishing the population characteristics of novae because these data are often heavily biased by selection effects.  To-date no RNe have been identified outside the Milky Way and its companions, but CNe have been studied in about a dozen galaxies.  To gain further insight into the population of novae, and specifically to explore further the question of whether there exist two distinct nova populations \\citep[see e.g.][]{1992A&A...266..232D}, we are involved in a number of campaigns to study a large sample of novae in the Local Group and beyond.  The following sections briefly describe the current status of the three parts of our extragalactic CN work: the POINT-AGAPE survey; our Local Group CN follow-up project, and the Liverpool Extragalactic Nova Survey. ",
        "conclusions": ""
    },
    "0809/0809.3465_arXiv.txt": {
        "abstract": "The precision study of dark matter using weak lensing by large scale structure is strongly constrained by the accuracy with which one can measure galaxy shapes. Several methods have been devised but none have demonstrated the ability to reach the level of precision required by future weak lensing surveys. In this paper we explore new avenues to the existing {\\it Shapelets} approach, combining a priori knowledge of the galaxy profile with the power of orthogonal basis function decomposition. This paper discusses the new issues raised by this matched filter approach and proposes promising alternatives to shape measurement techniques. In particular it appears that the use of a matched filter (e.g. \\sersic\\ profile) restricted to elliptical radial fitting functions resolves several well known Shapelet issues. ",
        "introduction": "Galaxy shapes provide the unique signature of gravitational lensing by large scale structure, which has been recognized as a key to the study of dark matter and dark energy \\citep{munshi2008}. A limiting factor is the accuracy with which one can measure shapes \\citep{CH2006,M2007,HJ2008}. Among the different existing methods, one particularly interesting approach is the decomposition of galaxy images using basis functions e.g. \\citet{bj2002} or {\\it Shapelets} \\citep{ref1,rb2003,mr1}. The strengths of this approach rely on the fact that the shape measurement is analytical and therefore time efficient as it involves rather small matrix multiplications. {\\it Shapelets} decompose an image into a linear combination of orthonormal components up to some truncation order, and the shape parameters are extracted from a least-squares best fit using the recomposed (noise free) model. However, the Shapelet type approach suffers from a few difficulties:  (i) The choice of the decomposition truncation order is arbitrary. In \tpractice, different lensing groups use radically different ``optimal'' truncation orders. Some prefer low \\citep{KK2006}, while others prefer high \\citep{BPR2008}, although the $\\chi^2$ values for different truncation order could be very different. Therefore a constant $\\chi^2$ criterion to measure the shape does not appear to be a robust guarantee of unbiased shape measurement.  (ii) ``Easy cases'' such as large and bright elliptical galaxies are poorly fitted. This suggests that a good fit for low signal-to-noise galaxies does not necessarily mean that the shape has been correctly measured, since it could just be buried in the sky noise. This is the overfitting problem.  (iii) Basis decomposition has too many degrees of freedom for shear measurement, since ideally we are only interested in two numbers (or six if we include the flexion). This is where galaxy morphology and shear measurement are clearly two different problems.  All of those problems have one common origin, namely the choice of the zeroth order weight function -- Gaussian functions for both Cartesian \\citep{ref1} and Polar Shapelets \\citep{mr1}, as well as \\citet{bj2002}. Unfortunately, Gaussian functions are poor matches to real galaxy profiles. Ideally, we would like the zeroth order to be as close as possible to the real profile, and leave to the basis decomposition the task to fit departures from this ``typical'' profile.  Currently the most promising shape measurement method uses a bayesian model fitting approach \\citep{MKH2007,K2008}. This method does not suffer from the same issues as Shapelets, but is limited by the strong galaxy profile prior.  In this paper, we investigate how the change of the weight function affects the basis decomposition method, and how it leads naturally to a hybrid method which combines Shapelets and fitting techniques. We choose to focus on the \\sersic\\ profile (hence the term Sersiclets), but our discussion can be extended to any profile\\footnote{While working on the concepts discussed in this paper, we became aware of similar investigations using exponential and hyperbolic sech functions (Kuijken \\& van Uitert in prep).}, e.g. Moffat profile for ground based point spread function (PSF). \\sect{methodology} introduces the notation and gives a technical description of the new fitting functions. \\sect{experiment} shows the impact of those fitting functions on shape fitting and decomposition. Finally, \\sect{conclusion} summarizes our work so far and future possibilities.  Note that in this paper we choose not to discuss the PSF deconvolution. Indeed, the problems we mentioned earlier affect equally the measurement of galaxy shapes whether or not the galaxies are convolved with a PSF, and Shapelets are a popular approach because their Gaussian properties allow for very efficient PSF treatment. The PSF deconvolution issue goes beyond this work because it depends on how the PSF is measured and interpolated between stars. Moreover the approach developed here could as well be applied to the PSF profile measurement separately, and then used later to address the deconvolution step through a forward convolution model fitting method. See for example \\citet{K2008}. ",
        "conclusions": "\\label{sec:conclusion} We presented an extension of Shapelets by using an arbitrary weight function in place of the Gaussian function. As galaxies' light profiles follow the \\sersic\\ profile on average, we used the \\sersic\\ function as our weight function. This allowed us to fit cuspy galaxies at lower orders than Shapelets could.  Because the \\sersic\\ function lacks analytical properties, we used the Gram-Schmidt process to generate the orthonormal polynomials as radial components for the fitting functions, where the integrals in the process must be evaluated numerically. As the \\sersic\\ profile has poor local support, the integration limit must be truncated to a finite limit.  We found that the full set of fitting functions for Sersiclets cannot decompose an arbitrary image even at high orders, as the fitting functions do not form a complete set. Instead of modeling objects using all fitting functions, we reduced the fitting set to only the circularly symmetric components ($m=0$). The model was then sheared by $e_1$ and $e_2$ to render elliptical shapes. The reduced set of fitting functions defines a hybrid method which combines the most interesting features of the basis decomposition and the fitting technique.  Our experiments so far only focused on idealized images simulated using known profiles and noise. Both the full and the reduced Sersiclets outperformed Shapelets, as we expected. The Shapelet matched filter's true performance will be tested in a future paper on image simulations such as those for GREAT08 \\citep{Br2008}. The C code to evaluate Sersiclet models is publicly available on request."
    },
    "0809/0809.2092.txt": {
        "abstract": "We present the Observations of Redshift Evolution in Large Scale Environments (ORELSE) survey, a systematic search for structure on scales greater than 10 $h_{70}^{-1}$ Mpc around 20 well-known clusters at redshifts of $0.6 < z < 1.3$. The goal of the survey is to examine a statistical sample of dynamically active clusters and large scale structures in order to quantify galaxy properties over the full range of local and global environments. We describe the survey design, the cluster sample, and our extensive observational data covering at least 25$'$ around each target cluster. We use adaptively-smoothed red galaxy density maps from our wide-field optical imaging to identify candidate groups/clusters and intermediate-density large scale filaments/walls in each cluster field. Because photometric techniques (such as photometric redshifts, statistical overdensities, and richness estimates) can be highly uncertain, the crucial component of this survey is the unprecedented amount of spectroscopic coverage. We are using the wide-field, multi-object spectroscopic capabilities of the DEep Multi-Object Imaging Spectrograph to obtain 100-200+ confirmed cluster members in each field. Our survey has already discovered the Cl 1604 supercluster at $z \\approx 0.9$, a structure which contains at least eight groups and clusters and spans 13 Mpc $\\times$ 100 Mpc. Here, we present the results on the large scale environments of two additional clusters, Cl 0023+0423 at $z = 0.84$ and RX J1821.6+6827 at $z = 0.82$, which highlight the diversity of global properties at these redshifts. The optically--selected Cl 0023+0423 is a four-way group-group merger with constituent groups having measured velocity dispersions between 206--479 km s$^{-1}$. The galaxy population is dominated by blue, star-forming galaxies, with 80\\% of the confirmed members showing [OII] emission. The strength of the H$\\delta$ line in a composite spectrum of 138 members indicates a substantial contribution from recent starbursts to the overall galaxy population.  In contrast, the X-ray--selected RX J1821.6+6827 is a largely-isolated, massive cluster with a measured velocity dispersion of $926 \\pm 77$ km s$^{-1}$. The cluster exhibits a well defined red sequence with a large quiescent galaxy population. The results from these two targets, along with preliminary findings on other ORELSE clusters, suggest that optical selection may be more effective than X-ray surveys at detecting less-evolved, dynamically-active systems at these redshifts. ",
        "introduction": "Clusters of galaxies are accurate tracers of the large scale structure in the local Universe. Redshift surveys at $z < 0.1$ (e.g., Geller \\& Huchra 1989; Einasto et al.\\ 1997; Colless et al.\\ 2001) and numerical simulations (e.g., Colberg et al.\\ 2000; Evrard et al. 2002; Dolag et al.\\ 2006) reveal the filamentary structure of the Universe, stretching between clusters and superclusters of galaxies.  Clusters form at the nodes of these filaments, growing through the continuous accretion of individual galaxies and groups from the surrounding field (e.g., Frenk et al\\ 1996; Eke et al.\\ 1998).  The precursors of the filaments should be present around distant clusters, containing many of the galaxies which will eventually infall into the virialized core and form the cluster population observed today. Since galaxy environment should change dramatically during the course of vigorous cluster assembly, the large scale structure present around high-redshift clusters offers us the unique opportunity to probe, over the full range of local environments, the physical effects on galaxies as they assemble into denser regions.  As a result, we are undertaking the Observations of Redshift Evolution in Large Scale Environments (ORELSE) survey, a systematic search for structure on scales greater than 10 $h_{70}^{-1}$ Mpc around 20 known clusters at $z > 0.6$.  The survey covers 5 square degrees, all targeted at overdense cluster %intermediate-to-high density regions, making it complementary and comparable to field surveys such as DEEP2 (Davis et al.\\ 2003) and COSMOS (Scoville et al.\\ 2007; Koekemoer et al.\\ 2007).  While superclusters have been studied at low redshift (Davis et al.\\ 1980; Postman, Geller \\& Huchra 1988; Quintana et al.\\ 1995; Small et al.\\ 1998; Rines et al.\\ 2002; Fadda et al.\\ 2008; Porter et al.\\ 2008) and at $z \\approx 0.2 - 0.6$ (Kaiser et al.\\ 1999; Gray et al.\\ 2002; Kodama et al.\\ 2001, 2003, 2005; Ebeling et al.\\ 2004; Tanaka et al.\\ 2007a; Kartaltepe et al.\\ 2008), large scale structures at high redshift ($z \\ge 0.6$) are just now being explored (Kodama et al.\\ 2005; Nakata et al.\\ 2005; Tanaka et al.\\ 2006, 2007b; Swinbank et al.\\ 2007; Gal et al.\\ 2008; Fassbender et al.\\ 2008; Gilbank et al.\\ 2008; Patel et al.\\ 2008). Previous studies at these redshifts have largely focused on the central regions ($< 2~h_{70}^{-1}$ Mpc) of massive clusters. These studies have revealed strong evolution in the cluster galaxy population which includes (1) an increase in the fraction of blue, star-forming, late-type galaxies with redshift, implying that early-types are forming out of the excess of late-types over the last $\\sim 7$ Gyr (e.g., Butcher \\& Oemler 1984; Ellingson et al.\\ 2001; Dressler et al.\\ 1997; van Dokkum et al.\\ 2000, 2001; Lubin et al.\\ 2002); (2) a larger fraction of post-starburst (``K+A'') galaxies in the cluster versus field environment, indicating strong star formation activity in the recent past (e.g., Dressler et al.\\ 1999, 2004; Balogh et al.\\ 1999; Tran et al.\\ 2003); (3) a deficit of faint, passive, red galaxies, suggesting that a large fraction of these galaxies in present-day clusters has moved onto the red sequence relatively recently as a result of a truncation of star formation (e.g., Smail et al.\\ 1998; van Dokkum \\& Franx 2001; Kodama et al.\\ 2004; De Lucia et al.\\ 2004, 2007; Tanaka et al.\\ 2005, 2007b; Koyama et al.\\ 2007); and (4) an increase in the overdensities of X-ray and radio sources with cluster redshift, indicating enhanced starburst/nuclear activity in the past (e.g., Best 2003; Cappelluti et al.\\ 2005; Eastman et al.\\ 2007; Kocevski et al.\\ 2008a).  All of these findings imply significantly increased star-forming, starburst, and nuclear emission in the past and raise the question -- what physical mechanisms associated with the cluster environment are responsible for the suppression of star formation and nuclear activity and the transformation of gaseous, disk galaxies into passive spheroids?  Several mechanisms have been suggested, including galaxy harassment (Moore et al.\\ 1998), ram pressure stripping (Gunn \\& Gott 1972), starvation (Larson et al.\\ 1980), and merging (Mihos 1995, 1999). Most of these processes are associated {\\it not} with the densest cluster regions, but rather with the infall regions and lower-density environments far from the cluster cores.  Studies at low-to-moderate redshift suggest that the non-cluster processes play a pivotal role in driving galaxy evolution well before the galaxies reach the cluster cores. Data from the Sloan Digital Sky Survey (SDSS) and the 2dF Survey indicate a sharp transition between galaxies with field-like star formation rates (SFRs) and galaxies with low SFRs comparable to that of galaxies within the cluster cores (Lewis et al.\\ 2002; Gomez et al.\\ 2003; Goto et al.\\ 2003a,b; Balogh et al.\\ 2004). The density at which this transition occurs (log $\\Sigma \\sim 1$ Mpc$^{-2}$) corresponds to the cluster virial radius, well outside the core region on which most studies have been focused. The fact that low SFRs and passive (gas-less) spiral galaxies are observed well beyond the virial radius rule out severe processes, such as ram pressure stripping or galaxy-galaxy merging, as being completely responsible for the variations in galaxy properties with environment (Lewis et al.\\ 2002; Goto et al.\\ 2003b). Similarly, studies on large scales (clustocentric distances of $> 4~h_{70}^{-1}$ Mpc) at $z \\sim 0.4$ indicate that galaxy colors change sharply at relatively low densities and that the morphological mix, at a given density, is independent of cluster radius, implying that galaxy transformation occurs outside the cluster core and is physically associated with the infalling filaments and the chains of smaller galaxy groups within them (Kodama et al.\\ 2001; Treu et al.\\ 2003). Consequently, these groups serve as a pre-processing phase in the evolution of cluster galaxies (Kodama et al.\\ 2001, 2003, 2004; Gray et al.\\ 2002; Treu et al.\\ 2003; Bower \\& Balogh 2004; Mihos 2004).  These trends continue to redshifts approaching $z \\sim 1$ where there are indications of ``downsizing'' with galaxy evolution accelerated in high-density and high-mass (cluster) regions, compared to lower-density and lower-mass (group) regions. Studies of large scale structures at $z \\sim 0.5-0.8$ suggest that the galaxy population in clusters is more evolved, with the faint-end of the red sequence being more developed and the red-sequence galaxies showing no signs of current or recent star formation. In contrast, groups and lower-density environments have a stronger deficit of the faint red-sequence galaxies and clear signs of star-formation activity (weak [OII] emission and/or strong H$\\delta$ absorption) in the existing red-galaxy population (De Lucia et al.\\ 2004, 2007; Kodama et al.\\ 2004; Tanaka et al.\\ 2005, 2005; Koyama et al.\\ 2007). As such, these results confirm that, at these epochs, much of the activity is taking place in lower-density, lower-mass environments.  This environment--activity level connection also extends to the extremes of active galactic nuclei (AGN) and starburst galaxies. From studies covering a wide redshift range ($0.2 < z < 1.2$), excesses of 24 $\\mu$m and X-ray sources are observed preferentially in intermediate-density, not high-density, regimes, consisting of groups, cluster infall regions, and filaments (Cappelluti et al.\\ 2005; Gilmour et al.\\ 2007; Marcillac et al.\\ 2007, 2008; Kocevski et al. 2008b). These results, as well as those described above, imply that the environmental processes which induce nuclear activity, truncate star formation, and change galaxy properties are not necessarily driven by cluster specific mechanisms, like ram pressure stripping by the hot intracluster medium or harassment (i.e., truncation of the galaxy halo) by the cluster tidal field. Thus, it is essential to look well beyond the regions of traditional study (the cluster cores) to find answers to our questions concerning galaxy evolution. The ORELSE survey is designed specifically to target these largely-unexplored regions and answer these questions by correlating multi-wavelength (radio, optical, infrared, and X-ray) photometric data with galaxy kinematics and spectral features in a statistical sample of large scale structures at redshifts approaching unity.  The first large scale structure detected in the ORELSE survey was the Cl 1604 supercluster at $z \\approx 0.9$ which includes two massive clusters Cl 1604+4304 at $z = 0.90$ and Cl 1604+4321 at $z = 0.92$ originally detected by Gunn, Hoessel \\& Oke (1986) and further studied by Oke, Postman \\& Lubin (1998). Wide-field imaging and spectroscopy as part of the ORELSE survey has revealed a more complex, massive structure, containing at least eight groups and clusters and spanning 13 Mpc $\\times$ 100 Mpc (Lubin et al.\\ 2000; Gal \\& Lubin 2004; Gal, Lubin \\& Squires 2005; Gal et al.\\ 2008). Our extensive spectroscopic data on the Cl 1604 supercluster (over 400 confirmed members) demonstrate that comprehensive redshift surveys, like ORELSE, are essential for understanding galaxy and cluster evolution. Specifically, (1) the entire structure size in redshift space is equivalent to typical photometric redshift errors (Margoniner \\& Wittman 2008); (2) superpositions of groups/clusters mean that mass measures based on weak lensing signal, richness, or X-ray luminosity are highly uncertain; (3) cluster velocity dispersions based even on traditionally large numbers of galaxies (i.e.\\ 20-40) can be substantially overestimated due to outliers (Gal et al.\\ 2008); and (4) the expected overdensities of radio and X-ray sources are small, making the identification and study of individual active galaxies impossible (Kocevski et al.\\ 2008a).  In this paper, we build on our studies of the Cl 1604 supercluster. We present the experimental design of the full survey (\\S 2) and the photometric and spectroscopic results on two additional target clusters, Cl 0023+0423 (hereafter Cl 0023) at $z = 0.84$ and RX J1821.6+6827 (hereafter RX J1821) at $z = 0.82$, which highlight the significant diversity of structure properties at these redshifts (\\S 4 and 5). Throughout the paper we use a cosmology with $H_0=70$ $h_{70}^{-1}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "In this paper, we have presented the motivation, design, and implementation of the Observations of Redshift Evolution in Large Scale Environments (ORELSE) survey, a systematic search for structure on scales of 10 $h^{-1}_{70}$ Mpc around 20 well-known clusters at $0.6 < z < 1.3$. The survey covers 5 square degrees, all targeted at intermediate-to-high-density regions, making it complementary and comparable to field surveys such as DEEP2 and COSMOS.  The program utilizes optical/near-infrared imaging at the UKIRT 3.8-m, KPNO 4-m, Palomar 5-m, and Subaru 8-m covering at least 25$'$ around each target cluster. Following the successful application in the Cl 1604 supercluster at $z \\approx 0.9$, we use an adaptively smoothed red-galaxy density map to visually identify associated groups/clusters and larger scale filaments/walls. Guided largely by actual observations, we use the extensive sample of 400+ confirmed members in Cl 1604 supercluster to adapt our color and magnitude cuts, quantify significant density peaks, and estimate false contamination rates for other target fields at lower and higher redshifts.  We find that this technique is highly efficient at detecting systems even down to group masses ($\\sigma \\sim 300$ km s$^{-1}$), as well as extended structures covering significant portions of the imaging field.  The crucial component of the ORELSE survey, and what distinguishes it from similar surveys, is the unprecedented amount of spectroscopic coverage. Utilizing the wide field, multi-object spectrograph DEIMOS on the Keck 10-m, we are obtaining high quality spectra for 100--200+ confirmed members per system, allowing us to measure properties on a galaxy-to-galaxy basis. Targeting galaxies down to $i' = 24$, our system of color selection provides a spectroscopic efficiency for cluster members of up to 45\\%.  Based on the first results from the survey, we already see significant diversity of large-scale and galaxy-scale properties at a given redshift. Although at similar redshifts, Cl 0023 and RX J1821 differ in their evolutionary state, global dynamics, activity level, and galaxy population. The optically-selected Cl 0023 shows multiple high-density peaks in the red-galaxy density map and a complex velocity structure, suggesting a four-way group-group merger. The measured velocity dispersions of the member groups range from 206--479 km s$^{-1}$. The overall galaxy population is exceptionally active, with a high fraction of blue galaxies and [OII] emitters. Conversely, the X-ray selected RX J1821 is a relatively isolated, centrally-condensed massive cluster. Only one significant overdensity and one redshift peak are observed in this field. The measured velocity dispersion is 926 km s$^{-1}$, indicative of a rich cluster. However, there are obvious signs of continuing dynamical evolution, including an elongated galaxy distribution and significant velocity substructure. The galaxy population is substantially more quiescent than that of Cl 0023; however, it still has a larger fraction of blue, star-forming galaxies than local clusters, consistent with the bluing of the cluster population with redshift. Composite spectra from both systems suggest some contribution from galaxies which have undergone a recent starburst (strong H$\\delta$ absorption), with the contribution in Cl 0023 being substantial.  These two systems indicate that global cluster properties (i.e., mass) play a role in driving galaxy evolution (see also Discussion of Poggianti et al.\\ 2006). Cl 0023 is comprised of group-sized halos, has a small quiescent population, and subsequently contains more active galaxies. Cl 0023 will likely be a massive cluster, with a velocity dispersion on the order of RX J1821, in $\\sim 1$ Gyr, although noticeable dynamical evolution will still be taking place after $\\sim 3$ Gyr, i.e.\\ until $z \\lesssim 0.3$ (Lubin, Postman \\& Oke 1998). If this dynamical evolution does take place, the star formation in a significant fraction of the galaxy population must be quenched during this period in order to be consistent with average star-formation rates and blue fractions in moderate-redshift clusters of similar mass (e.g., Balogh et al.\\ 1997; Nakata et al.\\ 2005). The current activity level in Cl 0023 and this evolutionary time frame implies that the group and/or merger environment is efficient at inducing and quenching star formation.  Conversely, RX J1821 is already a cluster-sized halo, with a galaxy population dominated by passive red galaxies (from both early-epoch formation and later-time quenching). Although both systems will end up as massive clusters by the present day, they have clearly different evolutionary histories. We note that these results, along with preliminary findings from other ORELSE targets, suggest that X-ray--selection may be biased, at least with respect to optical surveys, toward more evolved systems at a given redshift. On the other hand, due to projection effects, optical surveys detect multi-group systems and systems embedded in large scale structures that are not traditionally part of ``cluster'' surveys. As a result, the cluster detection technique may have a strong impact on any conclusions concerning the presence and effect of large scale structure, the relative importance of global versus local environment, and the timescales for galaxy and cluster evolution.  The scientific goals of the ORELSE survey are to identify and examine a statistical sample of dynamically active clusters and large scale structures during an active period in their history. Based on our completed (including spectroscopic) observations so far, such as those presented here and in Gal et al.\\ (2008), we have found two superclusters (containing 7+ groups/clusters), a four-way group-group merger, and two largely-isolated, massive X-ray--selected clusters. When the survey is complete, we will have a significant sample with which to constrain large scale cluster dynamics, cluster formation mechanisms, and the effect of environment on galaxy evolution.  Combining our ground-based databases with follow-up radio, X-ray, and/or mid-infrared observations, the program will (1) determine the relation between global cluster properties (such as richness, mass, luminosity, and dynamical state) and the nature of the surrounding large scale structures; (2) quantify the dependence of galaxy multi-wavelength and spectral properties on position and density to determine when, where, and why star formation and nuclear activity are triggered (or truncated) in and near the cluster environment; and (3) chart the process of structure formation and its effect on galaxy evolution over this active 3 Gyr period."
    },
    "0809/0809.4135_arXiv.txt": {
        "abstract": "A dynamical classification of the cosmic web is proposed. The  large scale environment is classified into four web types: voids, sheets, filaments and knots. The classification is based on the evaluation of the deformation tensor, i.e. the Hessian of the gravitational potential, on a grid. The classification is based on counting the number of eigenvalues above a certain threshold, $\\lambda_\\mathrm{th}$, at each grid point, where the case of zero, one, two or three such eigenvalues corresponds to void, sheet, filament or a knot grid point. The collection of neighboring grid points, friends-of-friends, of the same web attribute   constitutes  voids, sheets, filaments and knots as web objects.  A simple dynamical consideration suggests that $\\lambda_\\mathrm{th}$ should be approximately unity, upon an appropriate scaling of the deformation tensor.  The algorithm has been applied and tested against a suite of (dark matter only) cosmological N-body simulations. In particular, the dependence of the volume and mass filling fractions on $\\lambda_\\mathrm{th}$ and on the resolution has been calculated for the four web types. Also, the   percolation properties of voids and filaments have been studied.   Our main findings are: (a) Already at $\\lambda_\\mathrm{th}=0.1$ the resulting web classification reproduces  the visual impression of the cosmic web. (b)  Between $0.2 \\lesssim \\lambda_\\mathrm{th} \\lesssim 0.4 $, a system of percolated voids coexists with a net of interconected filaments. This suggests a reasonable choice for $\\lambda_\\mathrm{th}$ as the parameter that defines the cosmic web. (c) The dynamical nature of the suggested classification provides a robust framework for incorporating environmental information into galaxy formation models, and in particular the semi-analytical ones. ",
        "introduction": "\\label{sec:intro}  The large scale structure of the Universe, as depicted from galaxy surveys, weak lensing mapping and numerical simulations, shows a web like  three dimensional structure. There are three features that can be generally observed. First,  most of the volume resides in underdense regions; second, most of the volume is permeated by filaments; third, the densest clumps are located at the intersection of filaments \\citep{1996Natur.380..603B}.  This motivates a classification of the cosmic web into at least three categories: voids (underdense regions), filaments and knots (densest clumps).  There are  clear evidences for the  correlations of the observed  properties of galaxies with the environment.  We have for instance the morphology density correlation,  where elliptical galaxies are found preferentially in crowded  environments, and spiral galaxies are found in the field \\citep{1980ApJ...236..351D}. The same kind of correlation can be found in terms of the colors of the galaxies \\citep{2005ApJ...629..143B}.   According to the current paradigm of  structure formation  galaxies form and evolve in  dark matter (DM) halos \\citep{1978MNRAS.183..341W}. It follows that the study of such environmental dependence should commence  with the effort to understand the formation of DM halos in the context of the cosmic web \\citep{2005MNRAS.363L..66G,2005ApJ...634...51A,2007ApJ...654...53M}. This motivates us to search for a robust and meaningful  method   to classify  the different environments in numerical simulations. Such a classification should provide the framework for studying  the environmental dependence of  galaxy formation.   Translating the visual impression into an algorithm that classifies the local geometry into different environments is not a trivial task. A somewhat less challenging, yet very closely related, task is that of identifying just the voids out of the cosmic web. A thorough review and comparison of different algorithms of void finders  has been recently presented in \\cite{2008MNRAS.387..933C}.  The   void finders can be classified according to the method employed. Most are based on the point distribution of galaxies or dark matter (DM) halos and some on the smoothed density or potential fields. Some of the finders are based on spherical filters  while others assume no inherent symmetry \\citep{2007MNRAS.375..184B,2005MNRAS.360..216C,2003MNRAS.344..715G, 2002MNRAS.332..205A, 2008MNRAS.386.2101N, 2007MNRAS.380..551P, 2002MNRAS.330..399P, 2006MNRAS.367.1629S}.  It should be emphasized that an environment finder should be evaluated by its merits and it cannot be labeled as correct or wrong. A good algorithm should provide a quantitative classification which agrees with the visual impression and it should be based on a robust and well defined numerical scheme. Yet, it is desirable for an algorithm to be based on an analytical prescription so that its outcome can be calculated analytically. Also, simplicity is always  very highly desired.   A variety of approaches have been employed in the classification of the cosmic environment into its basic elements. The simplest way is based on the association of the environment with the local density, evaluated with a top-hat filter, say, of some width \\citep{1999MNRAS.302..111L}.  The density field can be analyzed in a much more sophisticated and elaborated way. This is the case of the web classification based on the multi-scale analysis of the Hessian matrix of the density field  \\citep{2007A&A...474..315A} or the skeleton analysis of the density field \\citep{2006MNRAS.366.1201N,2008MNRAS.383.1655S}. Both methods classify the cosmic web by pure geometrical tools applied to the density field. A very different approach is done within a dynamical framework in which the analysis of the gravitational potential is used to classify the web. This has been inspired by the seminal work of \\cite{1970A&A.....5...84Z} that led  to the \"Russian school of structure formation\" (e.g. \\cite{1982GApFD..20..111A, 1983MNRAS.204..891K}). The quasi-linear theory of  the Zeldovich approximation predicts the existence of an infinitely  connected web of pancakes (i.e. sheets), filaments and knots. This morphological classification is based on the study of the eigenvalues of the deformation tensor, namely the Hessian matrix of the linear gravitational field.   A recent application of the Zeldovich-based classification has been provided by \\cite{2008arXiv0801.1558L} who used a Wiener filter linear reconstruction of the local density field and evaluated the linear deformation, and hence also the shear, tensor on a grid. The cosmic web has been classified according the structure of the shear tensor.  A different approach has been followed by \\cite{2007MNRAS.375..489H}   who suggested that the full non-linear  gravitational potential should be used for the  geometrical classification. Apart from the difference between the linear and the  non-linear potential, both \\cite{2007MNRAS.375..489H} and \\cite{2008arXiv0801.1558L}  are using the same classification. Namely, the Hessian of the gravitational  potential is evaluated on a grid and its eigenvalues are examined locally. A  grid point is classified as a void, sheet, filament or knot point if the  number of eigenvalues greater that a null threshold is zero, one, two or three.   The algorithm presented here is an extension and improvement of the one suggested by \\cite{2007MNRAS.375..489H}. The extension is represented in the selection of a new free parameter with a dynamical interpretation. As such it provides a classification of local environment. Namely each spatial point is flagged as belonging to a either a void, sheet, filament or a knot point. It is the collective classification of all points in space which gives rise to the geometrical construction we call the cosmic web. This opens the door for defining voids, sheets, filaments and knots as individual objects. Each object is defined as a collecion of connected points having the same environmental attribute. We use a friends-of-friends (FoF) algorithm to find connected sets of points in simulations. Having defined the web objects the statistical and dynamical properties of these can be readily studied. Here we shall focus on analyzing the statistical properties of the voids and filaments, aiming to better constrain the value for the free parameter we introduce. We perform as well a statistical study to asses effect of the cosmic variance on our conclusions.    This paper is organized as follows. The web classification scheme is described in \\S \\ref{sec:web}. The N-body simulation used in the paper and the numerical implementation of the web classification are presented in \\S \\ref{sec:n-body}.  \\S \\ref{sec:res-web} describes the main properties of the cosmic web and in particular its dependence on the smoothing scale and the free parameter of our classification scheme. \\S  \\ref{sec:res-void} concentrates on the properties of the voids sector of the cosmic web. \\S \\ref{sec:frag-filament} studies the fragmentation of filaments in order to give a confidence interval to the free parameter that was introduced. In \\S \\ref{sec:cosmic-variance} we revisit the sections \\S \\ref{sec:res-web} and \\S \\ref{sec:res-void} to study the effect of cosmic variance. The paper concludes with a general discussion and a summary of the main results of the paper  (\\S \\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}      This paper presents an improved method to identify large scale environment in dark matter simulations. Our scheme is based on the analysis of the Hessian of the gravitational potential generated by the dark matter distribution. The algorithm presented here constitutes an improvement on the scheme of  \\cite{2007MNRAS.375..489H}, involving a pertinent reinterpretation of the dynamics in the problem.  The change is done through the addition of a free parameter $\\lambda_\\mathrm{th}$ related to the dynamical time associated with the collapse. This important improvement allows a more realistic treatment of the cosmic web.   Inspection of the different plots of Figure \\ref{fig:factors} reveals the striking difference in the way the cosmic web, and in particular the voids, respond to the changes in the Gaussian smoothing and the $\\lambda_\\mathrm{th}$'s threshold. Keeping a null $\\lambda_\\mathrm{th}$  and changing $R_\\mathrm{s}$ we see that the non-linear evolution does not change the ranking of the VFF and MFF found in the deep linear regime (i.e., $R_\\mathrm{s}\\approx 12.5 \\hMpc$). Namely, for the null threshold the sheets have the highest VFF and the filaments the highest MFF, independent of $R_\\mathrm{s}$. Considering the case of a fixed $R_\\mathrm{s}=1.95\\hMpc$, the void VFF grows strongly with $\\lambda_\\mathrm{th}$ and above   $\\lambda_\\mathrm{th} \\sim 0.5$  the voids have the highest  VFF.  The results on the VFF are extremely robust respect to the simulation size and cosmic variance effects. In the case of void MFF, it changes  from the lowest one at  $\\lambda_\\mathrm{th}=0$ to the second highest  at $\\lambda_\\mathrm{th} =1.0$.    The nature of the web changes dramatically with the threshold and it becomes volume dominated by the voids as $\\lambda_\\mathrm{th}$ increases. The MFF shows a larger cosmic variance, compared with the VFF. Also, it is  equally sensitive to changes in the smoothing scale and the threshold value.   The web classification provides  a set of flagged points on a grid and the collection of neighboring grid points of a given environmental type, connected by a FoF algorithm,  forms objects we call voids, sheets, etc. The statistical properties of the system of voids and their dependence on $\\lambda_\\mathrm{th}$ have been explored here.  In particular the number of isolated voids and their sizes have been analyzed, finding that the number of isolated voids roughly decreases exponentially with $\\lambda_\\mathrm{th}$. In the $1 \\hGpc$ simulation the system of voids percolates between  $0.1 \\lesssim \\lambda_\\mathrm{th} \\lesssim 0.2$, at which  the largest void jumps from having less than $10\\%$ to $90\\%$ of the volume occupied by voids.   The percolation dynamics is also seen on average in the different smaller simulations, with a wide spread on the way the percolation is observed, but keeping the same threshold interval for the transition.   The association of the eigenvectors of the deformation tensor with the collapse time, and hence with the age of the universe, enables in principle a theoretical determination of $\\lambda_\\mathrm{th}$. Using the spherical collapse model, calculated within the WMAP3 cosmology, the threshold value is $\\lambda_\\mathrm{th}=3.21$. However, at that high value the web looks very fragmented, in particular the network of filaments.  The application of the ellipsoidal collapse model \\citep{2002MNRAS.329...61S} might provide a better theoretical estimate.  Short of a ``first principle'' determination of $\\lambda_\\mathrm{th}$ we resort to a heuristic approach. We look for the range of $\\lambda_\\mathrm{th}$ over which the percolated networks of the voids and filaments coexist for a fixed smoothing scale of $R_\\mathrm{s}=1.95 \\hMpc$ in the $1 \\hGpc$ simulation. The voids and filaments behave in an opposite way in terms of the dependence of the percolation on the threshold value. At low $\\lambda_\\mathrm{th}$ the voids are isolated and the filaments percolated and at high $\\lambda_\\mathrm{th}$ the voids percolate and the filaments are fragmented. Adopting a  $95\\%$ in the fractional volume as defining the percolation transition we find the web to be defined by a threshold in the range of  $0.20\\lesssim\\lambda_\\mathrm{th}\\lesssim 0.40$. This  range stands in good agreement with the visual impression obtained from the simulations.     The notion of the cosmic web is not new. The filamentary structure has been extensively studied, mostly within the context of the Zeldovich pancakes \\citep{1970A&A.....5...84Z}. The role of voids has also been heavily studied and many algorithms for voids  finding have been suggested (see \\cite{2008MNRAS.387..933C}). We have been motivated  by the computational simplicity and the elegance of the \\cite{2007MNRAS.375..489H}  approach and have modified it in a way that reproduces the web as it emerges  from observations and simulations. Our main drive is to provide a simple, fast  and precise tool for classifying the environmental properties of each point in space. Using  the non-zero thresholding of the eigenvalues of the Hessian   of the potential indeed provides a very efficient tool that can be easily applied to simulations.  The same analysis can be performed at different redshifts in the simulation, allowing the classification of environment as a function of time.  The dynamical nature of the web classification implies that the web type might affects the dynamical evolution of DM halos and of galaxy formation. This might  have profound consequences   on the star formation timescale.... \\citep{2008arXiv0805.2191G}.  It follows that the web  classification can be introduced into semi-analytical modeling, as a dynamical  tag that together with the mass of the DM halos dictate the the mode gas accretion onto galaxies.   A major challenge that is still to be addressed is the application of the method to the distribution of galaxies. Short of that, the algorithm remains in the theoretical realm of simulations and semi-analytical modeling of galaxy formation."
    },
    "0809/0809.3838_arXiv.txt": {
        "abstract": "Although the orbits of comparable mass, spinning black holes seem to defy simple decoding, we find a means to decipher all such orbits. The dynamics is complicated by extreme perihelion precession compounded by spin-induced precession.  We are able to quantitatively define and describe the fully three dimensional motion of comparable mass binaries with one black hole spinning and expose an underlying simplicity. To do so, we untangle the dynamics by capturing the motion in the orbital plane and explicitly separate out the precession of the plane itself. Our system is defined by the 3PN Hamiltonian plus spin-orbit coupling for one spinning black hole with a non-spinning companion. Our results are twofold: (1) We derive highly simplified equations of motion in a non-orthogonal orbital basis, and (2) we define a complete taxonomy for fully three-dimensional orbits.  More than just a naming system, the taxonomy provides unambiguous and quantitative descriptions of the orbits, including a determination of the zoom-whirliness of any given orbit. Through a correspondence with the rationals, we are able to show that zoom-whirl behavior is prevalent in comparable mass binaries in the strong-field regime, as it is for extreme-mass-ratio binaries in the strong-field. A first significant conclusion that can be drawn from this analysis is that all generic orbits in the final stages of inspiral under gravitational radiation losses are characterized by precessing clovers with few leaves and that no orbit will behave like the tightly precessing ellipse of Mercury. The gravitational waveform produced by these low-leaf clovers will reflect the natural harmonics of the orbital basis -- harmonics that, importantly, depend only on radius. The significance for gravitational wave astronomy will depend on the number of windings the pair executes in the strong-field regime and could be more conspicuous for intermediate mass pairs than for stellar mass pairs. The 3PN system studied provides an example of a general method that can be applied to any effective description of black hole pairs. ",
        "introduction": "To provide the reader with a road map through intermediate results accumulated on our way to the periodic taxonomy, we preview some highlights here.  Our method can be broken into two main steps:  \\begin{figure} \\hspace{-155pt} \\centering \\hfill \\includegraphics[width=90mm]{xplots/orbplane1.eps} \\hspace{0pt} \\includegraphics[width=60mm]{xplots/orbplane2.eps} \\hfill \\caption{Top: The orbital plane precesses around the $\\bhj=\\bhk$ axis through the angle $\\Psi$. Bottom: The orbital plane can be spanned by the vectors $(\\bhx,\\bhy)$ or the vectors $(\\bn,\\bhPhi)$. \\label{orbplane}} \\end{figure}   {\\bf 1. Simplified Equations of Motion.} Since the perihelion precesses and the orbital plane precesses, the motion around a spinning black hole depends on angles as well as on the radius. In usual spherical coordinates, the equations of motion are quite complicated. By working in a non-orthogonal orbital basis, we show explicitly that the equations of motion are independent of (non-orthogonal) angular variables.  Physically, we exploit the observation\\footnote{The orbital plane is also emphasized in applications of PN dynamics to pulsar timing \\cite{{schafer1993},{wex1998},{gong2004},{konigsdorffer2005}}.} that the orbit lies in the plane spanned by the coordinate $\\br$ and its canonical momentum $\\bp$; that is, the orbital plane is perpendicular to the orbital angular momentum, $\\bl=\\br \\times \\bp$. The plane itself then precesses around the constant total angular momentum, $\\bj$. The importance of the orbital plane was clear in some of the earliest papers on spin-precession \\cite{apostolatos1994}, although that early work generally imposed a quasi-circular restriction on the orbits. We decompose all motion into precession of the perihelion within the orbital plane with a precession of the entire plane superimposed, with no restrictions or approximations. A preview of the explicit construction is shown in Fig.\\ \\ref{orbplane}. A fully precessing orbit is shown on the top in Fig.\\ \\ref{4leaf} while on the bottom the orbital plane traps a much simplified orbit, reminiscent of the equatorial orbits of Kerr black holes \\cite{levin2008}.  The simplified Hamilton's equations immediately inform us that all eccentric orbits have constant aphelia and perihelia.\\footnote{The constancy of periastron and apastron for every orbit might must have been implicitly understood in Refs.\\ \\cite{{hartl2005},{damoureob2001}}.} When the aphelia and perihelia are one and the same, we have non-equatorial constant radius orbits, also known as spherical orbits (as previously found in \\cite{{damoureob2001},{buonanno2006}}). The spherical orbits are not necessarily periodic; they fill out a band on the surface of a sphere. They are nonetheless significant in our campaign to fully dissect the dynamics and are treated in a companion paper \\cite{companion}.  \\begin{figure} \\centering \\includegraphics[width=70mm]{xplots/4leaf.93dlong.eps} \\hspace{+15pt} \\includegraphics[width=45mm]{xplots/4leaf.9.eps} \\hfill \\caption{ Top: Fully three-dimensional orbit. Bottom: The trajectory as captured in the orbital plane.} \\label{4leaf} \\end{figure}  The simplified Hamilton's equations show that zoom-whirl patterns will be symmetric from one radial cycle to another when viewed in the orbital plane, as can be seen on the right of Fig.\\ \\ref{4leaf}. A related subtle feature is that the three coordinate velocities in the orbital basis depend only on radius and are therefore periodic as an orbit executes a radial cycle from apastron to apastron. Taken together these symmetries in the orbital plane are intriguing for gravitational wave analysis. The waveforms must be decomposable into the orbital basis and therefore Fourier decomposable into the three fundamental frequencies that are the time average over one radial cycle of the instantaneous velocities. It remains to be seen how advantageous this might be for gravitational wave astronomy.   We restrict ourselves to completing the dynamical picture in this article since it is the dynamics that shapes the gravitational waves. In this spirit, step 1 above allows us to proceed to step 2:  {\\bf 2. Taxonomy of Fully 3D Orbits.} We offer a method to completely taxonomize the dynamics with the restriction that only one of the black holes spins. Our approach includes {\\it all} fully three-dimensional orbits described by the third-order Post-Newtonian (3PN) Hamiltonian plus spin-orbit couplings.  Our taxonomy extends the periodic tables for Kerr equatorial orbits \\cite{levin2008} to fully non-equatorial orbits of comparable mass black hole binaries. In Ref.\\ \\cite{levin2008}, we introduced a taxonomy for equatorial Kerr motion with the following salient features. Each entry in the Kerr periodic tables of Ref.\\ \\cite{levin2008} is a perfectly closed equatorial orbit identified by a rational number \\begin{equation} q=w+\\frac{v}{z} \\end{equation} where $w$ counts the number of whirls, $z$ counts the number of leaves, and $v$ indicates the order in which the leaves are traced out. Since the rationals are dense on the number line, the periodics are dense in phase space. Consequently, {\\it any} generic equatorial orbit can be {\\it arbitrarily} well-approximated by a nearby periodic orbit. In this way, any generic orbit is approximately equivalent to a high-leaf orbit (high $z$). Additionally, any generic orbit can be approximated as a precession around a low-leaf orbit, a technique that might ultimately benefit signal extraction.  Our ambition in this paper is both to extend the taxonomy to comparable mass binaries and to resolve the non-equatorial motion of spinning binaries. Truly periodic three-dimensional motion follows when the trajectory closes in the orbital plane {\\it and} the precession of the entire plane closes simultaneously. Fully closed motion requires two rationals, each representing a ratio of fundamental frequencies. And although in principle there must exist orbits that are fully periodic in the three-dimensional motion -- as Poincar\\' e argued \\cite{poincare1892} -- our taxonomy of bound orbits needs only the weaker condition of periodicity in the orbital plane. Not every orbit that is closed in the orbital plane will be closed in the full three-dimensional space. In other words, for the less restrictive condition of orbital plane periodicity we only need one rational ratio of frequencies. As will be explained in detail in \\S \\ref{closed}, the aperiodic orbit of Fig.\\ \\ref{4leaf} can be approximated as a precession around a 4-leaf clover in the orbital plane.  Although the PN approximation is poor in the strong-field, the qualitative results should survive a full relativistic treatment. Spin-spin couplings will impose additional modulations on the orbital plane picture but since spin-spin couplings are higher-order in the PN expansion, the expectation has been that their effect can be treated as a perturbation \\cite{buonanno2006}. In a companion paper we will argue that although a small perturbation, the spin-spin couplings are responsible for the emergence of chaos around unstable orbits \\cite{companion}.  We emphasize that for real spinning astrophysical black holes, the orbits we resolve in this paper are not an exotic subset of orbits, but rather are descriptive of {\\it all} bound orbits -- {\\it all} non-circular orbits are captured in the spectrum of rationals. There is a long-standing argument that black hole binaries will circularize by the time they enter the bandwidth of the gravitational wave observatories. However this is not possible for spinning black holes. Circular orbits {\\it do not exist} for misaligned spins. Although spherical orbits do exist, they are destroyed by the spin-spin effects \\cite{companion}. What's more, black hole pairs formed in dense clusters are not expected to circularize by the time of merger and are expected to be plentiful sources for advanced LIGO \\cite{{wen2003}} While we restrict ourselves to one spinning black hole and one non-spinning in this paper, the scenario is both astrophysically possible in its own right and theoretically important to lay the foundation for the two spinning case with spin-spin included, a task we return to in a companion paper \\cite{companion}  We express Hamilton's equations in a non-orthogonal orbital basis in \\S \\ref{simple}. We discuss the closed orbit taxonomy in \\S \\ref{closed}. In \\S \\ref{periodic}, we show periodic tables for two different black hole binaries, a comparable mass binary and a non-spinning extreme mass ratio pair. Appendix \\S \\ref{orbitalapp} details the projection of Hamilton's equations onto our orbital basis. In the conclusions, \\S \\ref{conc},  we discuss the modulations predicted from spin-spin couplings and those imposed by spinning both black holes. ",
        "conclusions": "\\label{conc}  To recap in bullet format, for comparable mass binaries with one spinning black hole and one non-spinning black hole as approximated by the 3PN Hamiltonian plus spin-orbit coupling, our main results are:  {\\bf 1. Simplified Equations of Motion in an Orbital Basis}\\par From which we find \\par $\\bullet $ constant aphelia and perihelia for non-equatorial eccentric orbits, and \\par $\\bullet $ three fundamental frequencies that depend only on radius.\\par  \\noindent and  {\\bf 2. Taxonomy of Fully Three-Dimensional Orbits}\\par For which we find\\par $\\bullet $ there exists a spectrum of closed orbits in the orbital plane corresponding to a subset of the rationals;\\par $\\bullet $ one rational, not two is required for an orbital plane taxonomy of constant angular momentum slices;\\par $\\bullet ${\\it all} orbits can be approximated as near an orbit that is perfectly closed in the orbital plane; and\\par $\\bullet $ zoom-whirl behavior is ubiquitous in comparable mass binary dynamics and entirely quantifiable through the spectrum of rationals.  The first discovery we made in Ref.\\ \\cite{levin2008} with our periodic taxonomy for equatorial Kerr orbits is that precessing elliptical orbits such as Mercury's are excluded in the strong-field Kerr regime -- just as Keplerian ellipses are excluded in any relativistic regime. Instead, at close separations, {\\it all } equatorial Kerr eccentric orbits trace out precessions of patterns best described as multi-leaf clovers, whirling around the central black hole before zooming out quasi-elliptically \\cite{{barack2004},{glampedakis2002}}. In this paper we have found that this same conclusion applies in the strong-field regime of our comparable mass black hole binaries, even out of the equatorial plane. Our periodic tables in the orbital plane show zoom-whirl behavior as the norm in the strong-field regime and not as the exception.  The further importance of the orbital dynamics lies in its direct imprint in the gravitational waveform \\cite{{levino2000},{drasco2006}}. The waveform will necessarily reflect the features above. For instance, an equatorial circular orbit (neglecting radiation reaction) is described by essentially one frequency. By contrast, {\\it all other orbits} in the strong-field regime generate highly modulated waveforms naturally described by harmonics of the 3 orbital basis frequencies, which in turn directly correspond to the natural frequencies of a nearby periodic orbit.  Naturally, we should ask about the astrophysical likeliness of detecting any such orbits with either the LIGO or LISA observatories. Although estimates vary \\cite{belczynski2007}, stellar mass black hole pairs are currently the favored source for advanced ground-based dectectors and intermediate mass black hole pairs are considered an important objective for space-based LISA science. It is challenging to definitively assess the spins and eccentricities of black hole/black hole binaries given the absence of observational constraints \\cite{oshaughnessy2005}. Still, one can guess that long-lived stellar binaries that might collapses to a pair of bound black holes would circularized by the time the pair enters the strong-field due to angular momentum lost in the form of gravitational radiation.\\footnote{However, when spin-spin coupling is included, there are no circular or even spherical orbits \\cite{companion}.} By contrast, for shorter-lived black hole binaries formed in globular clusters, the astrophysical likeliness of eccentric orbits sliding in the LIGO bandwidth is assessed to be $\\gtrsim 30\\%$ for eccentricities $>0.1$ in Ref.\\ \\cite{wen2003}. All such binaries would necessarily transit near the periodic set on inspiral. Even if the inspiral happens too quickly to witness multiple executions near a low-leaf clover, the orbit can still be sewn together as a skip from a piece of one periodic to a piece of another. Finally, the spins and eccentricities of intermediate black hole binaries detectable by LISA are most difficult to predict although we should expect them to spend a more generous allotment of windings on eccentric orbits in the strong-field.  Before closing, we have to mention the effects of spin-spin coupling on the orbital basis picture. The spin-spin correction introduces additional precessions of the spin and this destroys the constancy of the angle between $\\bs\\cdot \\bl$ and generally introduces explicit angular dependence in the equations of motion. Interpreted as a small perturbation to the system here, the spin-spin couplings cause additional wobbling of the precessional motion.  When both objects spin, the impact of spin-spin coupling can be particularly destructive. It is by now well documented that two spinning black holes in comparable mass binaries exhibit chaotic motion in the conservative system \\cite{{suzuki1997},{levin2000},{levin2003},{cornish2002},{cornish2003},{hartl2005},{wu2007},{wu2008}}. There are not enough constants of the motion to ensure regular behavior. Even the restricted spin-orbit scenario of this paper is clearly vulnerable to chaos even perturbed as it admits a homoclinic orbit. Under perturbation, homoclinic orbits tend to disrupt into a homoclinic tangle, an infinite intersection of stable and unstable manifolds. At root, chaos emerges as periodic orbits proliferate in a bounded region of space. Our clean taxonomy of periodic orbits corresponding to a simple set of rationals would give way to a glut of periodic orbits corresponding to some fractal set, as in systems with a strange repeller -- the Hamiltonian analog to the strange attractor. The complex motion may be damped by energy and angular momentum loses to gravitational waves but, at the least, chaos signals the breakdown of the simple set of periodic orbits. The onset of chaos can be directly identified with the spin-spin couplings, and we leave this task to a forthcoming companion paper \\cite{companion}.  \\bigskip \\bigskip \\bigskip \\bigskip \\bigskip  \\noindent{*Acknowledgments*}   We are especially thankful to Gabe Perez-Giz for his valuable and generous contributions to this work and to Jamie Rollins for his careful reading of the manuscript. We also thank Szabi Marka for important discussions. JL and BG gratefully acknowledge the support of a Columbia University ISE grant. This material is based in part upon work supported under a National Science Foundation Graduate Research Fellowship.   \\vfill\\eject \\appendix"
    },
    "0809/0809.4904_arXiv.txt": {
        "abstract": "At the end of inflation, dynamical instability can rapidly deposit the energy of homogeneous cold inflaton into excitations of other fields. This process, known as preheating, is rather violent, inhomogeneous and non-linear, and has to be studied numerically. This paper presents a new code for simulating scalar field dynamics in expanding universe written for that purpose. Compared to available alternatives, it significantly improves both the speed and the accuracy of calculations, and is fully instrumented for 3D visualization. We reproduce previously published results on preheating in simple chaotic inflation models, and further investigate non-linear dynamics of the inflaton decay. Surprisingly, we find that the fields do not want to thermalize quite the way one would think. Instead of directly reaching equilibrium, the evolution appears to be stuck in a rather simple but quite inhomogeneous state. In particular, one-point distribution function of total energy density appears to be {\\em universal} among various two-field preheating models, and is exceedingly well described by a lognormal distribution. It is tempting to attribute this state to scalar field turbulence. ",
        "introduction": "\\label{sec:intro}  The idea of inflation is a cornerstone of the modern theory of the early universe. According to inflationary paradigm, universe at early times undergoes a period of rapid (quasi-exponential) expansion, which wipes the initial state of the universe clean while seeding the primordial inhomogeneities with quantum fluctuations generated during expansion \\cite{Linde:2005dd, Linde:2005ht}. While universe is inflating, all of its energy sits in the homogeneous scalar field or condensate (known as inflaton), which is in a vacuum-like state with little entropy or particle excitations. But eventually inflation ends, and this energy has to be deposited into excitations of other matter fields, starting the thermal history of the universe with a hot big bang.  Decay of the inflaton can be very efficient if the fields experience dynamical instability at the end of inflation; such a stage became known as preheating. In most chaotic inflation models, oscillations of the inflaton field can cause instability via parametric resonance \\cite{Dolgov:1989us, Traschen:1990sw, Kofman:1994rk, Shtanov:1994ce, Kofman:1995fi}. Although linear development of this instability can be understood analytically \\cite{Kofman:1997yn, Greene:1997fu}, it might be chaotic \\cite{Podolsky:2002qv}, and one needs to resort to numerical simulations to investigate non-linear dynamics that soon takes over \\cite{Khlebnikov:1996mc,Podolsky:2005bw, Dufaux:2006ee, Felder:2006cc}. In hybrid inflation models \\cite{Linde:1993cn}, in addition to parametric resonance \\cite{GarciaBellido:1997wm}, one also has a tachyonic instability associated with the symmetry breaking \\cite{Felder:1998vq}, dynamics of which has been explored in \\cite{Felder:2000hj, Felder:2001kt, GarciaBellido:2002aj}.  Non-equilibrium dynamics of preheating can lead to a multitude of interesting phenomena. Some of the topics discussed in the literature are formation of topological defects \\cite{Tkachev:1998dc, Battye:1998xe}, production of various particles (with applications to baryo- and leptogenesis) \\cite{GarciaBellido:1999sv, GarciaBellido:2000dc, GarciaBellido:2001cb, GarciaBellido:2003wd}, possibility of primordial black hole formation \\cite{Suyama:2004mz, Suyama:2006sr}, generation of primordial magnetic fields \\cite{DiazGil:2007dy, DiazGil:2008tf}, and production of stochastic gravitational wave background \\cite{Easther:2006gt, Easther:2006vd, GarciaBellido:2007dg, GarciaBellido:2007af, Dufaux:2007pt, Caprini:2007xq}. Due to difficulties of dealing with non-linear evolution equations, most of these studies rely on numerical simulations.  This paper describes a new code for simulating non-linear scalar field dynamics in expanding universe developed to study preheating, called DEFROST, and the first results obtained with it. There are other codes available for this purpose, most notably LATTICEEASY by Gary Felder and Igor Tkachev \\cite{Felder:2000hq}, and its parallel version CLUSTEREASY \\cite{Felder:2007nz}. Through the use of more advanced algorithms and careful optimization, the new code significantly improves both accuracy and performance achievable in simulations of preheating. An important design goal has been the ease of visualization and analysis of the results, which is all too important for understanding dynamics of complex systems.  We present results on preheating in a simple two-field chaotic inflation model with massive inflaton decaying into another scalar field via quartic coupling \\cite{Felder:2006cc}. Through our simulations, a new and somewhat simpler picture of the late stages of preheating emerges. After initial transient when instability develops, bubbles form and then break-up, the matter distribution soon arranges itself in a clumpy state which persists with little changes for a long time. One-point probability distribution function of total energy density in this state appears to be {\\em universally lognormal} among various two-field preheating models. It is tempting to attribute this to relativistic turbulence \\cite{Nordlund:1998wj, Micha:2002ey, Micha:2004bv}. We also see some evidence that the structure formed during preheating continues to grow in size on a much longer timescale. Somehow, this picture reminds one of large scale structure formation \\cite{Bardeen:1985tr, Bernardeau:1994aq, Bond:1995yt}.  This paper is organized in the following way: In Section~\\ref{sec:eom}, we introduce scalar field models of preheating, derive equations of motion, and discuss physical approximations we use. Section~\\ref{sec:solver} describes the detailed implementation of numerical evolution algorithm. Initial conditions including quantum fluctuations of the fields produced during inflation are discussed in Section~\\ref{sec:ic}, with particular attention paid to implementation of Gaussian random field generator. We briefly recount the theory of preheating via broad parametric resonance at the end of chaotic inflation in Section~\\ref{sec:chaotic}, and present our simulations in Section~\\ref{sec:results}. We conclude by summarizing our main results in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  This paper presents a new numerical code I developed for simulating preheating of the Universe after the end of the inflation, which I call DEFROST. It is small (about 600 lines of Fortran code), fast, easy to modify, and is fully instrumented for 3D visualizations (using \\href{http://wci.llnl.gov/codes/visit/}{LLNL's VisIt}, for example). The source code is available for download online at \\url{http://www.sfu.ca/physics/cosmology/defrost/} and is distributed under the terms of GNU Public License. While the main design goal of DEFROST has been the accuracy of the simulations, performance of the solver has also been significantly improved compared to LATTICEEASY~\\cite{Felder:2000hq}, which is the most mature and widely used reheating code publicly available today.  As a result of all the optimizations (and a bit of black magic), DEFROST outperforms LATTICEEASY by about a factor of four in raw PDE solver speed (for 2 fields on a $256^3$ grid in double precision on a dual Xeon 5160 machine) while using more accurate (and more expensive) discretization. If one takes into account the time spent on analysis of the results, the difference is even larger, as FFTW libraries used by DEFROST are vastly faster than FFT routines shipped with LATTICEEASY (especially on multi-processor machines). The speed-up offered by DEFROST is so significant that the studies done a few years ago on a big parallel cluster \\cite{Podolsky:2005bw, Dufaux:2006ee} using MPI version of LATTICEEASY \\cite{Felder:2007nz} can now be carried out on a single fast workstation. The planned MPI version of DEFROST should be able to push the accessible simulation size over the $1024^3$ barrier, provided the code scales well.  The code was tested on a number of chaotic inflation models which end via parametric resonance. In this paper, we report the simulations of the simplest two-field preheating model with massive inflaton and quartic coupling to decay field (\\ref{eq:model}). We reproduce the previously published numerical results for this model \\cite{Podolsky:2005bw, Felder:2006cc} (and the ones for trilinear coupling \\cite{Dufaux:2006ee}, which we will not discuss here, although our simulation data and results for that model are available \\href{http://www.sfu.ca/physics/cosmology/defrost/M2S12L4}{online} as well). We further investigate the dynamics of the scalar field evolution in these preheating models, taking advantage of advanced visualization and analysis capabilities DEFROST offers. In particular, we study the behaviour of energy density distribution and scalar gravitational potential during preheating, something which has not been looked at closely before. Our main science results are summarized by two observations, both novel and quite surprising.  First, the evolving scalar fields quickly end up in a simple state, which, although highly inhomogeneous, appears to have a certain universality to it. In this state, the one-point distribution function of total energy density is nearly stationary (apart for the overall dilution due to expansion), and is described by a lognormal distribution for {\\em all two-field parametric resonance preheating models we tried so far}, namely the ones described by interaction potentials \\begin{itemize} \\item $V(\\phi,\\psi) = \\frac{1}{2}\\, m^2 \\phi^2 + \\frac{1}{2}\\, g^2 \\phi^2 \\psi^2$, \\item $V(\\phi,\\psi) = \\frac{1}{2}\\, m^2 \\phi^2 + \\frac{1}{2}\\, \\sigma \\phi \\psi^2 + \\frac{1}{4}\\, \\lambda \\psi^4$, \\item $V(\\phi,\\psi) = \\frac{1}{4}\\, \\lambda \\phi^4 + \\frac{1}{2}\\, g^2 \\phi^2 \\psi^2$, \\item $V(\\phi,\\psi) = \\frac{1}{4}\\, \\lambda (\\phi^2 + \\psi^2)^2$. \\end{itemize} This is true even if distributions of field values or other correlators might still be evolving, and appears to be a very general statement about random scalar fields one encounters in preheating. It is tempting to attribute this state to scalar field turbulence \\cite{Micha:2002ey, Micha:2004bv}, especially since lognormal density distributions are known to occur in supersonic isothermal turbulence in hydrodynamics \\cite{Nordlund:1998wj}. We do not see obvious signs of thermalization, even if the simulations are run for a time much longer than the dynamical timescale of the problem (the longest done so far for massive inflaton is $2^{12}/m$ corresponding to five $e$-folds since the end of inflation; this is limited mainly by my patience rather than the code stability).  Second, less general but still amusing, is the observation that the small-scale structure in the gravitational potential can grow faster than comoving box expands. It is not quite clear whether the reason for it happening is kinematic or dynamical in nature. As we neglected gravitational interactions in our simulations, the only thing that can cause the structure to grow is the interaction between scalar fields themselves. In our preheating model it is repulsive; yet the structure still grows! Although one might suspect that any inhomogeneity in gravitational potential on sub-horizon scales would probably get washed away by subsequent evolution (and is too small to form primordial black holes), this effect still might have some interesting cosmological consequences.  All in all, we find that the picture of preheating dynamics is simpler in real space than what it looks like in particle representation. The final stage of preheating, with growing structure and lognormal density distribution, eerily reminds one of large scale structure formation in later cosmology (although of course it occurs on vastly smaller scales and is driven by completely different physics). Perhaps the analytical methods developed for the latter \\cite{Bardeen:1985tr, Bernardeau:1994aq, Bond:1995yt} could be fruitfully applied to preheating as well. This is what we intend to explore next."
    },
    "0809/0809.1829_arXiv.txt": {
        "abstract": "The CRESST cryogenic direct dark matter search at  Gran Sasso, searching for WIMPs via nuclear recoil, has been upgraded to CRESST-II by  several changes and improvements. The upgrade includes a new detector support structure capable of   accommodating 33 modules,  the associated multichannel readout with 66 SQUID channels,  a neutron shield, a calibration source lift, and the installation of a muon veto. We present the results of a commissioning run carried out in 2007.  The basic element  of CRESST-II  is a  detector module consisting of a large ($\\sim$300 g) $\\mathrm{CaWO_4}$ crystal and a very sensitive smaller ($\\sim 2$ g) light  detector to detect the scintillation light  from the $\\mathrm{CaWO_4}$. The large crystal gives an  accurate total energy measurement.  The  light  detector permits a determination of the light  yield for an event,  allowing an effective separation of nuclear recoils from electron-photon backgrounds. Furthermore, information from light-quenching factor studies allows the definition of a region of the energy-light yield plane which  corresponds to tungsten recoils. A neutron test is reported which supports the principle of using the light yield to identify the recoiling nucleus.     Data obtained with two detector modules for a total exposure of 48 kg-days are presented. Judging by the rate of events in the ``all nuclear recoils'' acceptance region the apparatus  shows a  factor $\\sim$ten improvement with respect to  previous results, which we attribute  principally to the presence of the neutron shield. In the ``tungsten recoils''  acceptance region  three events are found, corresponding to a rate of 0.063 per kg-day. Standard assumptions on the dark matter flux,  coherent or spin independent interactions, then yield  a limit  for WIMP-nucleon scattering of $4.8 \\times 10^{-7}\\uu{pb}$, at $M_{\\uu{WIMP}}\\sim50\\uu{GeV}$. ",
        "introduction": "Evidence continues to grow that the majority of the mass of our galaxy and in the universe is made up of non-baryonic dark matter \\cite{bertone2005}. Precision measurements of the cosmic microwave background have given accurate figures for the density of matter and energy in the universe as a whole, suggesting that about a fourth of the energy density of the universe is in the form of  dark matter \\cite{komatsu2008}. Measurements of the rotation curves of other spiral galaxies, indicate large amounts of dark matter, which would  suggest it also  dominates our own galaxy. Gravitational  lensing of light by galactic clusters indicates large amounts of dark matter, and dark matter is necessary for a reasonable description of the formation of galaxies.  Although some of these observations might also be explained by alternative theories such as modified Newtonian dynamics (MOND) \\cite{sanders2002}, recent studies of the Bullet Cluster appear to favour the dark matter hypothesis \\cite{clowe2006}. However  the need  for direct detection of  dark matter evidently remains strong.  A plausible origin for dark matter comes from particle physics in the form of WIMPs (weakly interacting massive particles). These would be stable massive particles with an interaction cross section typical for weak interaction processes. A relic density sufficient to make a significant contribution to the energy density of the universe arises naturally  in this way. Supersymmetry provides  a well motivated candidate in the form of the  neutralino and this has been extensively studied theoretically.  According to  the WIMP hypothesis  a galaxy has  a large gravitationally bound  halo of such particles, making up most of the mass.  The   local mass density may be   estimated  from halo models and is found to be  on the order of  $0.3\\uu{GeVcm^{-3}}$. The range of the coherent or spin-independent WIMP-nucleon scattering cross-sections predicted by minimal supersymmetric models extends from below $10^{-10}$ to above $10^{-7}\\uu{pb}$ \\cite{trotta2007,trotta2007b}; this makes the direct detection of such particles a difficult but in  principle achievable task.  The CRESST  project is a dark matter search aiming to detect the nuclear scattering of WIMPs by the use of  cryogenic detectors. Since only small recoil energies ($\\sim$ keV's)  are   anticipated in  WIMP-nucleus scattering, cryogenic detectors with their high sensitivity are well suited to the problem.  CRESST-II is an upgrade of the original CRESST appartus where the detectors are arranged in modules, each  consisting of two detectors, a large detector comprising the target mass  and a similar but much smaller light detector measuring the light yield. While the large detector gives  an accurate  total  energy measurement, the light yield  measurement   permits an effective rejection of  events which are not  nuclear recoils.  This arises from the property of  electron-photon events,  constituting the dominant background, to give a much higher  light output than nuclear recoils.   The large detector consists of a $\\mathrm{CaWO_4}$ crystal of  $\\sim300\\uu{g}$ with a tungsten superconducting-to-normal thermometer deposited on the surface. The energy of an event  is measured in this detector. The energy deposited in the crystal is quickly converted into a gas of phonons  which are then absorbed in the superconducting thermometer. The  separate  light detector, based on the same principle, measures the simultaneously emitted   scintillation light. A comparison of these two signals then allows  an effective background discrimination.     Results from an earlier   prototype run with two modules were presented in Ref.\\,\\cite{angloher2005}, where taking  the better of the two modules and assuming  coherent WIMP scattering on the tungsten led to the  limit $1.6\\times10^{-6}\\uu{pb}$. The improvements discussed here are aimed at ultimately  increasing this sensitivity by  up to two orders of magnitude. Such an increase would result, for example, with background free data and a threshold of 12 keV from 10 kg for 100 days giving 1000 kg-days. Or if the  threshold could be reduced to 5 keV, with 300 kg days.   During 2007 an extended commissioning run of the new  setup was carried out. Although this was primarily for optimization purposes, the results also represent an interesting  further improvement. We report on them and the  modifications for CRESST-II  in the present paper. ",
        "conclusions": "\\label{summary}  CRESST-II successfully completed its commissioning run in 2007. New elements of the apparatus  were successfully installed and operated. A neutron test demonstrated the ability to detect nuclear recoils with a light yield consistent with that found from quenching factor studies, supporting the principle of identifying the recoil nucleus via the light yield.  Data were taken with two detector modules  for a total of 48 kg-days. Three candidate events of uncertain origin are present in the acceptance region for  tungsten recoils, yielding  a rate of 0.063 per kg-day. A factor $\\sim 10$ improved performance is found with respect to previous work for the ``all nuclear recoils'' acceptance region.  A limit on  coherent WIMP-nucleon scattering, is obtained, which at its lowest  value, for $M_{\\uu{WIMP}}\\approx 50 $GeV, is $4.8\\times10^{-7}\\uu{pb}$."
    },
    "0809/0809.4537_arXiv.txt": {
        "abstract": "The primordial curvature fluctuation spectrum is reconstructed by the cosmic inversion method using the five-year Wilkinson Microwave Anisotropy Probe data of the cosmic microwave background temperature anisotropy. We apply the covariance matrix analysis and decompose the reconstructed spectrum into statistically independent band-powers. The statistically significant deviation from a simple power-law spectrum suggested by the analysis of the first-year data is not found in the five-year data except possibly at one point near the border of the wavenumber domain where accurate reconstruction is possible. ",
        "introduction": "Inflationary cosmology \\cite{inf1,inf2,inf3} explains the origin of cosmic structure on various scales in an unified way. The expansion history during the inflationary stage of the early universe is recorded in the spectrum of seed structure which can be revealed by modern cosmological observations such as the Wilkinson Microwave Anisotropy Probe (WMAP) observation of the cosmic microwave background (CMB) \\cite{WMAPBASIC,WMAPTEMP,WMAPCOSMO,WMAPINF,WMAPTTTE,loglike,WMAP3TEMP,WMAP5,WMAP5ONLY,WMAP5COSMO}. The WMAP mission evaluated every multipole moment of the CMB anisotropy spectrum in a wide range of scales which is potentially a record of the inflationary expansion history with {\\it high time resolution}. It is a feasible challenge to probe the highly time-resolved behavior of the inflaton field (s) which drives inflation.  Since the initial release of the WMAP results \\cite{WMAPBASIC,WMAPTEMP,WMAPCOSMO,WMAPINF,WMAPTTTE}, it has been argued that the CMB temperature anisotropy spectrum has nontrivial features such as running of the spectral index, oscillatory behaviors on intermediate scales, and lack of power on large scales \\cite{INFRA,JF94,COBE,JY99}, which cannot be explained by a  power-law primordial spectrum that is a generic prediction of simplest inflation models. These features may provide clues to unnoticed physics of inflation. Some of these anomalous structures disappeared on the three-year spectrum, however several anomalies are still existing \\cite{WMAP3TEMP,COSPA07}. To explain these features, a number of inflation models have been proposed \\cite{CPKL03,CCL03,FZ03,KT03,HM04,LMMR04,KYY03,YY03,KS03,YY04,MR04A,MR04C,KTT04,HS04,HS07,CHMSS06}.  As an observational approach to these possible nontrivial features, which can be an alternative to model fitting, there have been several attempts to reconstruct the primordial spectrum using CMB anisotropy data without any prior assumptions about the shape of the primordial spectrum \\cite{WSS99,BLWE03,MW03A,MW03B,MW03C,MW03D,SH01,SH04,BLH07,SVJ05,WMAP3COSMO,LSV08,VP08}. One such attempt to reconstruct the primordial spectrum is a filtering method where the primordial spectrum is characterized by amplitudes on a few number of representative scales. While such methods have an advantage for reconstructing the global structure of the primordial spectrum, they may miss possible fine structures if their scale width is narrower than the filtering scale, which has been chosen rather arbitrarily so far.  On the other hand, there exist other methods which can reconstruct the primordial spectrum as a continuous function without any {\\it ad hoc} filtering scale to investigate detailed features such as the cosmic inversion method \\cite{MSY02,MSY03,KMSY04,KSY04,KSY05}, the Richardson-Lucy method \\cite{SS03,SS06,SS07}, or a nonparametric method \\cite{TDS04,THS05}. The cosmic inversion method has proved its ability of reconstructing the modulations off a power-law spectrum quite well by the analysis of mock anisotropy data. In the analysis of the first-year WMAP data, we pointed out the possibility of nontrivial structures in the primordial spectrum around the scales of $2.4\\times10^2 {\\rm Mpc}$ and $4.2\\times10^2 {\\rm Mpc}$.  Theoretically, according to the standard inflation paradigm, each wavenumber ($k$-) mode of the power spectrum is mutually independent. On the other hand, each $k$-mode of the power spectrum reconstructed from the observed CMB anisotropy has a strong correlation with neighboring modes because each multipole of CMB anisotropy, $C_\\ell$, depends on the $k$-modes in the wide range around $k=\\ell/d$ where $d$ is the distance to the last scattering surface. In order to extract real features, therefore, it is important to decompose the reconstructed spectrum to mutually independent band-powers keeping resolution as high as possible.  The purpose of this work is to apply the cosmic inversion method to the five-year WMAP temperature anisotropy spectrum \\cite{WMAP5}, update the reconstructed primordial curvature spectrum, and perform the above-mentioned band-power analysis by diagonalizing the covariance matrix. Then, we revisit the possibility of fine structures in the primordial spectrum. Because of the arbitrariness of the primordial spectrum, we inevitably incorporate infinite degree of freedom to our analysis, which results in degeneracy among spectral shape and cosmological parameters \\cite{KNN01,SBKET}. In this paper, we consider the concordance adiabatic $\\Lambda$CDM model, where the cosmological parameters (except for the ones characterizing the primordial spectrum) are those of the WMAP team's best-fit power-law model \\cite{WMAP5ONLY,WMAP5COSMO}, and instead focus on the detailed shape of the primordial spectrum. Note that, as shown in \\cite{KMSY04}, different choices of cosmological parameters affect only the overall shape. The fine structures of the reconstructed power spectrum remain intact in the relatively small wavenumber range we probe.  This paper is organized as follows: In Sec. II, the overview of our analysis is described. In Sec. III, we show the reconstructed primordial power spectrum from the five-year WMAP data and discuss its implication. Finally, Sec. IV is devoted to the conclusion. ",
        "conclusions": "Using the cosmic inversion method, we have reconstructed the primordial power spectrum of curvature perturbations assuming the absence of isocurvature modes and the best-fit values of cosmological parameters for the power-law $\\Lambda$CDM model. While the range of accurate reconstruction is rather narrow, $120 \\lesssim kd \\lesssim 380$, we can reproduce the fine structures of modulations off a simple power-law with which we can recover the highly oscillatory features observed in $C_\\ell$.  The statistical significance of the oscillatory structures in the reconstructed power spectrum is difficult to quantify due to the strong correlation among the neighboring $k$-modes.  We have therefore performed the covariance matrix analysis to calculate the window matrix, $W$, which diagonalizes the covariance matrix into statistically independent modes.  We have chosen the large enough number of band-powers, $N=40$, to probe the feature of the primordial spectrum as precisely as possible while keeping the overlap of the window functions $W_{ij}$ small enough.  As a result we have found all the independent modes are consistent with the best-fit power-law $\\Lambda$CDM model for the five-year WMAP data except possibly for a point around $kd\\simeq 120$. Whereas, the possible deviation from a simple power-law around $kd \\simeq 180,~220$ and $ 350$ reported in the analysis of the first-year WMAP data \\cite{KMSY04} has disappeared. This difference is due to the fact that the observed $C_\\ell$ in the five-year WMAP data has become much smoother than the first-year counterpart around the first peak.  On the other hand, there still remain some nontrivial deviations from the expected  $C_\\ell$ of a simple power-law spectrum around $\\ell\\approx kd \\approx 40$ even in the five-year data.  Unfortunately, the cosmic inversion method cannot probe the primordial spectrum in this region. So we need to develop a different method which can probe the power spectrum for smaller wavenumber region $kd \\leq 120$ precisely.  From the covariance matrix analysis, we have shown that statistically independent bands of the primordial spectrum have an effective width $\\Delta kd \\approx 10$.  This means that it is difficult to probe the possible fine structures on $A(k)$ predicted by, say, trans-Planckian processes \\cite{Martin:2000xs}, if the scale width of their characteristic modulation is narrower than $\\Delta kd \\approx 10$. In other words, even if our result is consistent with a simple power-law spectrum, this does not rule out such possibility of narrow modulation of $\\Delta kd \\ll 10$."
    },
    "0809/0809.1124_arXiv.txt": {
        "abstract": "In a previous work (Pichardo \\etal 2005), we studied stable configurations for circumstellar discs in eccentric binary systems. We searched for ``invariant loops'': closed curves (analogous to stable periodic orbits in time-independent potentials) that change shape with the binary orbital phase, as test particles in them move under the influence of the binary potential. This approach allows us to identify stable configurations when pressure forces are unimportant, and dissipation acts only to prevent gas clouds from colliding with one another.  We now extend this work to study the main geometrical properties of circumbinary discs. We have studied more than 100 cases with a range in eccentricity $0 \\le e \\le 0.9$, and mass ratio $0.1 \\le q \\le 0.9$. Although gas dynamics may impose further restrictions, our study sets lower stable bounds for the size of the central hole in a simple and computationally cheap way, with a relation that depends on the eccentricity and mass ratio of the central binary. We extend our previous studies and focus on an important component of these systems: circumbinary discs. The radii for stable orbits that can host gas in circumbinary discs are sharply constrained as a function of the binary's eccentricity.  The circumbinary disc configurations are almost circular, with eccentricity $e_d < 0.15$, but if the mass ratio is unequal the disk is offset from the center of mass of the system.  We compare our results with other models, and with observations of specific systems like GG Tauri A, UY Aurigae, HD 98800 B, and Fomalhaut, restricting the plausible parameters for the binary. ",
        "introduction": "\\label{Intro}  It is currently believed that fragmentation is the most probable mechanism for star formation, and the main product of fragmentation are multiple stellar systems with preference for wide eccentric binaries with separations $\\ge 10 AU$ (Bonnell \\& Bastien 1992; Bate 1997; Bate \\& Bonnell 1997; Bodenheimer \\etal 2000). Even in isolated stars, there is evidence that the majority of Sun-like stars formed in clusters (Carpenter 2000; Lada \\& Lada 2003), including the Sun (Looney, Tobin \\& Fields 2006).  In the last decade the interest in binary systems has increased. This is in part because of the discovery that many T-Tauri and other pre-main sequence binary stars possess circumstellar and circumbinary discs as inferred from observations of excess radiation at infrared to millimeter wavelengths, polarization, and both Balmer and forbidden emission lines (Mathieu \\etal 2000, Itoh \\etal 2002, for a review see Mathieu 1994). On the other hand, recent observations of binary star systems, using the Spitzer Space Telescope, show evidence of debris discs in these environments (Trilling \\etal 2007) and planets (Fischer \\etal 2008). In their studies they find that $60\\%$ of the observed close binary systems (separations smaller than 3 $AU$) have excess in their thermal emission, implying on-going collisions in their planetesimal regions.  Over 150 extrasolar planets have been identified in surveys using the Doppler technique. Of the first 131 extrasolar planetary systems that have been confirmed, at least 40 are in binary or multiple systems (for an up-to-date list see Haghighipour 2006). Approximately 30 of them are on S-type orbits (around one of the components: circumstellar discs) with wide stellar separations (between 250 and 6500 $AU$), including at least 3 that orbit one member of a triple star (Raghavan \\etal 2006). Although most of these binaries are very wide, a few have separations smaller than 20 $AU$ (Els \\etal 2001; Hatzes \\etal 2003), challenging standard ideas of Jovian planet formation. Some interesting ideas try to explain the formation of Jovian planets within close binaries, but they could only explain few cases (Pfahl \\& Muterspaugh 2006), if more are discovered soon, these theories would not be sufficient. Although close binaries are not included in precise Doppler radial velocity search programs because of their complex and varying spectra, at least one planet with a minimum of $~2.5$ Jupiter masses has been detected in a P-type orbit (around both components: circumbinary discs), with a distance from the center of mass of $~23$ $AU$. The source is a radio pulsar binary comprised by a neutron star and a white dwarf in a $~191$ day stellar orbit (Lyne \\etal 1988; Sigurdsson and Phinney, 1993; Sigurdsson \\etal 2003). An example of accretion in P-type orbits about close binaries is given by Quintana \\& Lissauer (2006) who note the observation of the two small moons orbiting in nearly circular/planar orbits about the binary system Pluto-Charon (Weaver \\etal 2006). Specifically regarding to circumbinary disc material, millimeter and mid-infrared excess emission has been detected around several spectroscopic pre-main sequence binary star systems including GW Ori (Mathieu \\etal 1995), UZ Tau E (Jensen \\etal 1996), DQ Tau (Mathieu \\etal 1997).  In this work we have followed the same steps as in Pichardo, Sparke \\& Aguilar (2005, hereafter Paper~I), where we opted for a simpler approach, analogous to using the structure of periodic orbits in a circular binary, to predict the gas flow. The path followed by a gas parcel in a stable disk around a star must not intersect itself, or the path of a neighboring parcel (unlike the case of planets, where the paths may cross). In our work we follow Rudak \\& Paczynski (1981) and explore those non-crossing orbits of test particles that could be interpreted as gas particles in the low pressure regime, or as protoplanets or planets. An important issue in Celestial Mechanics is to determine the regions around a stellar binary system where accretion discs can form. Important theoretical effort carried out to answer this question is reviewed in Paper~I, where we studied circumstellar and circumbinary discs in binaries of arbitrary eccentricity and mass ratio.  In this work we extend those studies, which were based on identifying families of stable {\\it invariant loops}, a concept introduced by Maciejewski \\& Sparke (1997, 2000) in studies of nested galactic bars. We focus this time specifically on the geometry of circumbinary discs. We employ for this approach a test particle method probing the orbital structure of binaries of various eccentricities and mass ratios.  In Section \\ref{method} we briefly review the concept of an {\\it invariant loop}, describe the method to solve the motion equations and the strategy used to find invariant loops. The geometry of the circumbinary discs including a fit for the inner radii of the circumbinary disc and a fit for the lopsidedness, are presented in Section \\ref{geom}. In Sections \\ref{theory} and \\ref{observations} we apply this study to compare with theoretical work and observations of some well known systems, respectively. Our conclusions are presented in section \\ref{conclusions}. ",
        "conclusions": "We have extended the studies started in Paper~I to a more detailed analysis of circumbinary discs in eccentric binary systems from the geometrical and physical point of view.  The disks are defined by a family of stable {\\it invariant loops}, the analogs to stable periodic orbits around a circular binary, which do not cross each other or themselves.  Thus they define paths that can be followed by clouds of gas, which dissipate energy when they run into each other.  Just as with a binary in circular orbit, the circumbinary disk is truncated at its inner edge when there are no longer any stable non-crossing orbits for the gas to follow.  We have used this property of the invariant loops to define the inner edge of the circumbinary disk.  We showed already in Paper~I that the inner radius of a circumbinary disk depends strongly on eccentricity, opening wider gaps for higher eccentricities.  The size of the inner hole depends only slightly on the mass ratio.  The geometric centre of the circumbinary disc is off-centre with respect to the centre of mass of the binary system. The disc is closer to being symmetrical around the whole orbit of the secondary star.  Here, we have explored the range of parameters to quantify both the off-centering of the circumbinary disc with respect to the centre of mass of the system, and the average inner radius of circumbinary discs, as a function of mass ratio and eccentricity.  We compare our results with the work of HW99 who searched for initially-circular orbits that survive more than 10,000 periods of the binary.  When the eccentricity is small this procedure gives similar results to ours, but at larger eccentricity HW99 find fewer stable orbits close to the binary, and hence larger inner gaps.  This could be related to the fact that the larger the eccentricity, the smaller the available phase space of orbits that are trapped so that they must remain close to a stable invariant loop.  If the properties of a circumbinary disc are observable, it is possible to constrain the binary system properties by comparing the predictions based on invariant loops with what is observed. Likewise, if the orbital parameters of a binary are known, the geometry of the circumbinary and circumstellar regions permitted for stable orbits are readily obtained.  For the well known circumbinary disc of the binary system GG Tau A, we use the inner radius of the circumbinary disk to restrict the possibilities for the binary parameters.  Since two stars have nearly equal masses, we would expect the ring to be nearly symmetrical about the center of mass of the stars.  The observed offset is substantial and much larger than can be explained by orbital dynamics; effects such as the finite ring thickness may be important (e.g. Duch\\^ene \\etal 2004).  In the case of UY Aurigae, we show that the observed radius of the circumbinary disk requires that the binary orbit be noncircular, with eccentricity $ e > 0.1$.  The center of the disk should be offset from the mass center of the stars in the direction towards the center of the secondary star's orbit by $R_{sh}\\ga 0.05a$.  We have modeled the system HD 98800 B with the parameters given by Boden \\etal (2005): the shift of the disc with respect to the centre of mass of the system should be about $0.1~AU$.  Although the fits we provide here are valid for values of $q\\ge 0.1$, our technique is also applicable to extreme cases of $q \\le 0.001$. For the disk around Fomalhaut, we have used invariant loops to propose plausible solutions for the orbital parameters of a planet that explains the morphology of the debris disc of this system.  We reach similar results to Quillen (2006), but for the larger planetary mass proposed by Deller \\& Maddison (2005) we find that the planet's orbit must be less eccentric."
    },
    "0809/0809.3987.txt": {
        "abstract": "MOST (Microvariability \\& Oscillations of STars) and ASAS (All Sky Automated Survey) observations have been used to characterize photometric variability of TW~Hya on time scales from a fraction of a day to 7.5 weeks and from a few days to 8 years, respectively. The two data sets have very different uncertainties and temporal coverage properties and cannot be directly combined, nevertheless, they suggests a global variability spectrum with ``flicker noise'' properties, i.e.\\ with amplitudes $a \\propto 1/\\sqrt{f}$, over $>4$ decades in frequency, in the range $f = 0.0003$ to 10 cycles per day (c/d). A 3.7 d period is clearly present in the continuous 11 day, 0.07 d time resolution, observations by MOST in 2007. Brightness extrema coincide with zero-velocity crossings in periodic (3.56 d) radial velocity variability detected in contemporaneous spectroscopic observations of \\citet{Set2008} and interpreted as caused by a planet. The 3.56/3.7~d periodicity was entirely absent in the second, four times longer MOST run in 2008, casting doubt on the planetary explanation. Instead, a spectrum of unstable single periods within the range of 2 -- 9 days was observed; the tendency of the periods to progressively shorten was well traced using the wavelet analysis. The evolving periodicities and the overall flicker-noise characteristics of the TW~Hya variability suggest a combination of several mechanisms, with the dominant ones probably related to the accretion processes from the disk around the star. ",
        "introduction": "\\label{intro}  Photometric variability of the very young, T~Tauri-type stars is still puzzling and remains an object of very active research; for the most recent literature, see \\citet{Percy2006} and \\citet{ROTOR-I}. The variability comes in part from accretion and matter ejection phenomena, in part from photospheric spots coming and going into view, and in part from the inner accretion disk and then the accretion region as matter is channelled by magnetic fields into the photosphere. \\citet{Herbst1994} identified 3 basic types of variability which may coexist in a given T~Tauri star: Type~I, photospheric dark spots; Type~II, variable accretion with some rotation modulation component; Type~III, totally random variations, mostly due to variable obscuration. Quasi-periodic variations, intertwined with chaotic changes are the most natural outcome of several different mechanisms contributing simultaneously.  Typical time scales of T~Tauri stars are of the order of hours to days, but they are very difficult to characterize from the ground because of diurnal observation breaks and discontinuous temporal coverage. We are not aware of any attempt to obtain a continuous record of T~Tauri variability from a satellite or through inter-observatory coordination of efforts.  In this paper we present new, single broad-band (located between $V$ and $R$ bands), continuous photometric observations of the T~Tauri star TW~Hya obtained by the MOST satellite mission over 11 days in 2007 and 46 days in 2008. Although the high frequency coverage (above 10 cycles per day) was inadequate, the temporal coverage at the once-per-orbit satellite sampling period of 101.4 minutes was practically uninterrupted permitting a study of the stellar variability on time scales of a fraction of a day to a few tens of days. For still longer time scales, we used the ASAS project data obtained over 8 yearly seasons from 2001 to 2008. These data sampled the brightness changes of TW~Hya in the $V$ and $I$ bands at intervals of a day to a few days.  In this paper, after a brief introduction of TW~Hya itself (Section~\\ref{target}) and of its previous photometric variability studies (Section~\\ref{prev}), we discuss the results of the analysis of the 2007 and 2008 MOST data (Sections~\\ref{most7} and \\ref{most8}). The TW~Hya variability at very low frequencies is analyzed on the basis of the ASAS data (Section~\\ref{asas}). We conclude (Sections~\\ref{disc} and \\ref{concl}) that TW~Hya shows a flicker-noise variability spectrum (the special type of a ``red noise'' spectrum) over a wide range of the time scale from hours to years. ",
        "conclusions": "\\label{concl}  Following our analysis of the MOST and ASAS very different but partly complementary observational data sets, and considering conclusions of the important paper by \\citet{Press1978} on the common occurrence of flicker noise (with Fourier amplitudes scaling as $a \\propto 1/\\sqrt{f}$), we cannot resist a comparison of the TW~Hya photometric variability with an orchestra. As pointed out by Press, the spectrum of the orchestra sound can be usually described by the prosaic flicker noise. Thus, although we hear the music defined by an entire orchestra where everybody plays a different piece, sometimes the melody (tune) produced by an individual instrument becomes noticeable against the general flicker-noise level. We observed such distinct, time-variable ``tones'' in the period range of 2 to 9 days; they may be manifestations of the accretion process as the matter spirals into the star.  The flicker-noise spectrum of TW~Hya photometric variations appears to extend from very low frequency of 0.0003 c/d (accessible from 8 years of the ASAS data) to $\\simeq 10$ c/d (accessible to our MOST observations). The lowest frequencies require long monitoring time, so that good photometric stability and consistency of calibrations are essential, even at relatively moderate accuracy (0.01 -- 0.02 mag) of the ASAS project. At frequencies corresponding to days to weeks, the uniform time coverage and high accuracy of the MOST mission (0.003 -- 0.005 mag, even with the included intrinsic variability of the star in times scales shorter than an hour) permitted us to study the difficult (from the ground) frequency range of $\\simeq 0.025$ c/d to 10 c/d. In this range, we observed very clear, well defined, varying ``tones'' in the 0.1 -- 1.0 c/d range only (specifically, the periods 2 -- 9 days). It would be tempting to identify those changing periods as signatures of the orbital decay as the accreting material spirals into the star. The strong periodicity observed by MOST in 2007 with the period of 3.7 d, in phase with the 3.56 d radial-velocity variations of \\citet{Set2008}, appears to be one such ``tone''. It does not seem to be related to a possible planet orbiting TW~Hya because it disappeared within one year. However -- contrary to the clear period changes observed in 1.5 month of the MOST observations in 2008 -- it was relatively stable through the 3 month of the 2007 radial velocity observations.  The wide range of variability frequencies suggests a multitude of mechanisms. While we suspect that the main mechanism in the range of time scales accessible to MOST is accretion within the innermost disk at distances of 2 -- 15 $R_\\odot$ from the star, it is hard to imagine that accretion would produce photometric variability on time scales of years at the implied radii of several astronomical units. Thus, in the symphony orchestra analogy, more instruments (mechanisms) must be contributing to create the extended, well defined flicker-noise variability spectrum of TW~Hya."
    },
    "0809/0809.2899_arXiv.txt": {
        "abstract": "We present a novel solver for an analogue to Poisson's equation in the framework of modified Newtonian dynamics (MOND). This equation is highly non-linear and hence standard codes based upon tree structures and/or FFT's in general are not applicable; one needs to defer to multi-grid relaxation techniques. After a detailed description of the necessary modifications to the cosmological \\nbody\\ code \\amiga\\ (formerly known as \\mlapm) we utilize the new code to revisit the issue of cosmic structure formation under MOND. We find that the proper (numerical) integration of a MONDian Poisson's equation has some noticable effects on the final results when compared against simulations of the same kind but based upon rather ad-hoc assumptions about the properties of the MONDian force field. Namely, we find that the large-scale structure evolution is faster in our revised MOND model leading to an even stronger clustering of galaxies, especially when compared to the standard \\LCDM\\ paradigm. ",
        "introduction": "\\label{sec:introduction}  Modified Newtonian dynamics (MOND) was proposed by \\citet{Milgrom83} as an alternative to Newtonian gravity to explain galactic dynamics without the need for dark matter. Although current cosmological observations point to the existence of vast amounts of non-baryonic dark matter in the Universe \\citep[e.g.][]{Komatsu08}, it remains interesting to explore other alternatives, especially as not all of the features of CDM models appear to match observational data (e.g., the ``missing satellite problem'' \\citep{Klypin99, Moore99} and the so-called ``cusp-core crisis'' \\citep[e.g.][]{deBlok03,Swaters03}). In that regards it appears important to look for tests able to discriminate between MOND and Newtonian gravity, especially in the context of cosmology now that there exist various relativistic formulation of the MOND theory \\citep{Bekenstein04,Zhao07,Zhao08}. However, progress in that field has been hampered by the fact that MOND is a non-linear theory making any analytical predictions as well as numerical simulations a tedious task. To date, only a few codes exist that actually solve the MONDian analogue to Poisson equation in a non-cosmological context \\citep{Brada99, Nipoti07, Tiret07}, none of which is publically available. We are augmenting this list by making available a new solver for the MONDian Poisson's equation primarily designed to work in a cosmological context but readily adjusted to allow for simulations of isolated galaxies.  Until recently MOND was merely a heuristic theory tailored to fit rotation curves with little (if any) predictive power for cosmological structure formation. One of the most severe problems for the general appreciation and acknowledgment of MOND as a \"real\" theory (and a conceivable replacement for dark matter) was the lack of success to formulate the theory in a general relativistic manner.  This situation though changed during the last couple of years and at present there are a number of covariant theories (e.g. \\citep{Bekenstein04, Zhao07, Zhao08}) whose non-relativistic weak acceleration limit accords with MOND while its non-relativistic strong acceleration regime is Newtonian. However, there remains a lot to be done with regards to the relativistic formulations of MOND in order to a get completly acceptable theory. But nevertheless has MOND reached a stage of development in wich we can think in doing cosmology with reducing the need of unjustifiable assumptions.  Furthermore, the study of cosmological structure formation in the non-linear regime will provide new constrains on such generalizations of MOND.  The first valiant attempts at simulating cosmic structure formation under the influence of MOND were done by \\citet{Nusser02} and \\citet[][KG04 from now on]{Knebe04}. While their studies provided great insights into the non-linear clustering behaviour of MONDian objects (we dare to call them galaxies as none of these simulations included the physics of baryons) they were still based upon some rather unjustifyable assumptions. The main objective of this paper now is to refine the implementation of MOND in the \\nbody\\ code \\amiga\\ (successor of \\texttt{MLAPM}, \\citet{Knebe01}), i.e. we modified the code to numerically integrate a MONDian analogue to Poisson's equation. This equation is a highly non-linear partial differential equation whose solution is non-trivial to obtain. Only sophisticated multi-grid relaxation techniques (on adaptive meshes of arbitrary geometry) are capable of tackling this task. We argue that only when MOND has been thoroughly and \"properly\" studied and tested against \\LCDM\\ can we safely either rule it out for once and always or confirm this rather venturesome theory. The theory has become a valid competitor to dark matter and it therefore only appears natural -- if not mandatory -- to (re-)consider its implications. We developed a tool and the subsequently necessary analysis apparatus allowing to test and discriminate cosmological structure formation in a MONDian Universe from the standard dark matter paradigm and used it to study a particular model of MOND theory. ",
        "conclusions": "\\label{sec:conclusions}   We presented a novel solver for the analogue to Poisson's equation taking into account the effects of modified Newtonian dynamics (MOND). This equation is a highly non-linear partial differential equation for which standard solvers based upon Fourier transformation techniques and/or tree structures are not applicable anymore; one has to defer to multi-grid relaxation techniques \\citep[e.g.,][]{Brandt77, Press92, Knebe01} in order to numerically solve it.  The major part of this paper hence deals with the necessary (non-trivial) adaptions to the existing multi-grid relaxation solver \\amiga\\ (successor to \\mlapm\\ introduced by \\citet{Knebe01}). We show that the accuracy of our MOND solver is at a credible level for a static problem with known MONDian solution and that it reproduces correct results in the Newtonian limit for cosmological simulations.  In the case that we don't want to adhere to any (undjustifyable and hence ad-hoc) assumptions as, for instance, done by KG04, we are facing the problem that there is still a lot of freedom from the covariant point of view to describe the original phenomenological MOND theory.  To be able to actually perform MONDian cosmological simulations we decided to choose one particular model for MOND and to leave for later studies the analysis of the effects that differents theories producces on the structure formation.  One of the suppositions KG04 made relates to the relation between the MONDian and the Newtonian force vector. They use the rather simple relation as given by setting $\\mathbf C=\\mathbf 0$ in \\Eq{eq:curl} making the right-hand side simply the Newtonian force. This implicit definition for $\\vecg_M$ could be inverted and used instead of $\\vecg_N$ when updating the particles' velocities during the time integration of \\Eq{eq:eom}. As noted in \\Eq{eq:curl} they neglected the so-called MONDian curl-field $\\mathbf C=\\nabla \\times \\mathbf h$. Others have shown that this field decreases like $O(r^{-3})$ and vanishes for any kind of symmetry \\citep{Bekenstein84}, it yet remains unclear whether it will leave any imprint on inhomogeneous strucure formation as found in cosmological simulations.  In the second part of this paper we therefore presented a series of cosmological simulations set out to quantify both structure formation under MOND as well as the effects of the curl-field.  Surprisingly, we found that our results are consistent with KG04 even if we take in account that they use a different MOND model and initial conditions generated with standard linear theory in a MONDian regimen.  We revised some of their findings (mainly the discrepancy between the abundance of galaxies at galaxies at high redshift between \\LCDM\\ and OCBMond) and noted that the curl-field leaves marginal yet noticable effects. The major result of this study though is that the curl-field appears to drive structure formation! Our results obtained with the new solver (and hence including the effects of the curl-field) can be summarized as follows  \\begin{itemize} \\item the curl-field drives structure formation, \\item the curl-field leads to more objects at $z=0$ where \\item cross-identified objects are more massive in OCBMond2 than in OCBMond, and \\item the OCBMond2 model shows a stronger two-point correlation function. \\end{itemize}  We acknowledge that there are still a lot of puzzling results to be investigated in greater detail. However, we postpone this to future studies. The aim of this paper was primarily to describe the novel gravity solver that is freely available for download.\\footnote{\\amiga\\ can be downloaded from the following web page \\texttt{http://www.aip.de/People/aknebe/AMIGA}. The new MOND solver can be switched on via the compilation flag \\texttt{-DMOND2}. The code is able to run cosmological simulations, ie with expansion and periodic boundary conditions and to work also in isolation, wich implies no expansion and fixed boundary conditions.  It can be also use without temporal evolution in order to get 3D potentials from density distributions given by particles or analitic formulas. Please contact the authors for more details.}"
    },
    "0809/0809.0314_arXiv.txt": {
        "abstract": "We present results from deep radio observations taken with the VLA at a center frequency of 1400 MHz covering a region of the SWIRE Spitzer Legacy survey, centered at 10$^h$46$^m$00$^s$, 59\\arcdeg01\\arcmin00\\arcsec (J2000). The reduction and cataloging of radio sources are described. An electronic catalog of the sources detected above 5 sigma is also presented.  The survey presented is the deepest so far in terms of the radio source density on the sky. Perhaps surprisingly, the sources down to the bottom of the catalog appear to have median angular sizes still $> 1$ arcsecond, like their cousins 10-100 times stronger. The shape of the differential log N - log S counts also seems to require a correction for the finite sizes of the sources to be self-consistent. If the log N - log S normalization remains constant at the lowest flux densities, there are about 6 sources per square arcminute at 15 $\\mu$Jy at 20cm. Given the finite source size this implies that we may reach the natural confusion limit near 1 $\\mu$Jy. ",
        "introduction": "Understanding the growth and evolution of galaxies is one of the key goals of modern astronomy. A new generation of multi-wavelength surveys from X-ray to radio wavelengths are providing the starting point for new leaps of understanding of the subject. The key has proved to be deep surveys of the same fields at every wavelength possible, so that as complete a SED as possible can be determined for a large number of objects. At radio wavelengths, the VLA at 20cm is the workhorse instrument. The large field-of-view, high sensitivity and good spatial resolution allow images with thousands of sources per pointing and $\\sim$1.5 arcsecond resolution. New computing algorithms combined with the current generation of fast computers have unlocked this potential of the VLA, which has been there for 25 years.  The $\\mu$Jy radio sky gives unique information about the evolution of faint, distant galaxies from the synchrotron emission observed from both star-forming and AGN-dominated systems. This complements data from other wavelength bands. The ratio of radio to FIR luminosity has the potential to tell us whether star-formation or AGN are dominant \\citep[e.g.,][]{Yun01}. For star-forming systems the physical size of the emitting regions and their relation to brightness distributions in other bands can tell us about the extent of the star-forming regions independent of dust extinction and perhaps can tell us if a wind is required. Radio AGN are dominantly  radio jets and thus tell us about the mechanical energy flowing away from the black hole, instead of the radiation dominated emission seen at other wavelengths. Radiation and mechanical energy dominated AGN may result from different phases of the black hole growth \\citep{ch05,si07} and thus it is important to study both types of AGN. The sensitivity of this survey reaches a surface density of several sources per square arcminute and thus allows us to study a large sample of distant galaxies with only one pointing of the VLA. In this paper we report the results of the radio survey only. Other related papers will discuss the observations at other wavelengths and what we can learn by combining all the data on the objects in this field. ",
        "conclusions": "We have presented the deepest image so far of the radio sky in terms of source density and provided an electronic catalog for sources with peak signal-to-noise ratios $> 5$.  The median size distribution for these sources continues to be $\\sim1.2$ arcseconds or a bit larger down to the bottom of the catalog. After a correction for incompleteness due to source size, we find that the log N - log S normalization factor remains approximately flat down to 15$\\mu$Jy, corresponding to about six sources per square arcminute. If this trend continues the 20cm natural confusion limit may be reached near one $\\mu$Jy.        \\clearpage"
    },
    "0809/0809.0913_arXiv.txt": {
        "abstract": "We present the discovery of a new quadruply lensed quasar. The lens system, SDSS J1330+1810 at $z_s=1.393$, was identified as a lens candidate from the spectroscopic sample of the Sloan Digital Sky Survey. Optical and near-infrared images clearly show four quasar images with a maximum image separation of $1\\farcs76$, as well as a bright lensing galaxy. We measure a redshift of the lensing galaxy of $z_l=0.373$ from absorption features in the spectrum. We find a foreground group of galaxies at $z=0.31$ centred $\\sim 120''$ southwest of the lens system. Simple mass models fit the data quite well, including the flux ratios between images, although the lens galaxy appears to be $\\sim 1$~mag brighter than expected by the Faber-Jackson relation. Our mass modelling suggests that shear from nearby structure is affecting the lens potential. ",
        "introduction": "\\label{sec:intro}  \\begin{figure*} \\begin{center} \\includegraphics[width=0.75\\hsize]{f1.eps} \\end{center} \\caption{The SDSS $i$-band image of the SDSS~J1330+1810 field. The pixel scale of the image is $0\\farcs396$~${\\rm pixel^{-1}}$, and the seeing was $1\\farcs0$. The inset at lower left shows an  expanded view of the lens system. Several bright galaxies in the field are indicated by G1-G3. The $i$-band Petrosian magnitudes of these galaxies (without correcting for Galactic extinction) are $17.72$ (G1), $18.20$ (G2), and $18.59$ (G3). \\label{fig:fc1330}} \\end{figure*}  \\begin{figure} \\begin{center} \\includegraphics[width=0.95\\hsize]{f2.eps} \\end{center} \\caption{The SDSS spectra of SDSS~J1330+1810 ({\\it upper}) and the nearby galaxy G2 ({\\it lower}). See Figure \\ref{fig:fc1330} for the relative positions. The spectral resolution is $R\\sim 2000$, but the spectra are smoothed with a 5-pixel boxcar to suppress noise. Thick vertical line segments in the upper panel mark MgII and Ca absorption lines, which suggest the lens redshift of $z_l=0.373$. A Mg/Fe absorption system at $z=1.054$ is also shown by thin lines. The feature at $\\sim 5600${\\AA} in the spectrum of G2 is due to poor subtraction of a strong sky emission line ($5575${\\AA}). \\label{fig:spec1330}} \\end{figure}  Thus far about 100 gravitationally lensed quasars are known, of which $\\sim30$ are quadruple (four-image) lenses. The number ratio of quadruple lenses to double (two-image) lenses contains information on both the shapes of lensing galaxies and the luminosity function of source quasars \\citep{rusin01,chae03,huterer05,oguri07b,madelbaum08}. In addition, quadruple lenses allow more detailed mass modelling of individual lenses. For instance, the larger number of images provides more constraints on the lens potential, which is essential in probing the effects of external perturbations on primary lenses \\citep{keeton97} and constraining the Hubble constant from time delay measurements \\citep[e.g.,][]{suyu06}. Magnifications of merging image pairs in quadruple lenses satisfy distinct relations if the lens potential is smooth, but small-scale structures near the image can violate the relations. Thus flux ratios of quadruple lens images serve as unique probes of substructure or microlensing in lens galaxies \\citep{mao98,metcalf01,chiba02,dalal02,schechter02}.  In this paper, we present the discovery of a new gravitationally lensed quasar with four lensed images, SDSS~J133018.65+181032.1 (SDSS~J1330+1810). It was discovered as part of the Sloan Digital Sky Survey Quasar Lens Search \\citep[SQLS;][]{oguri06,oguri08,inada08}, which takes advantage of the large spectroscopic quasar catalog \\citep[see][]{schneider07} of the Sloan Digital Sky Survey \\citep[SDSS;][]{york00} to locate new lensed quasars. We place particular emphasis on mass modeling and investigation of the structure around the lens.  The outline of this paper is as follows. We describe the SDSS and follow-up data in Sections \\ref{sec:sdss} and \\ref{sec:followup}, respectively. The environment of the lens is discussed in Section \\ref{sec:env}. Section \\ref{sec:model} is devoted to mass modelling. Our results are summarised in Section \\ref{sec:sum}. Throughout the paper, we adopt the standard Lambda-dominated flat universe cosmology with matter density $\\Omega_M=0.26$ and the Hubble constant $h=0.72$ \\citep{dunkely08}. ",
        "conclusions": "\\label{sec:sum}  We have presented the discovery of a new four-image lensed quasar SDSS J1330+1810 ($z_s=1.393$). This source was selected as a lens candidate by the SQLS due to its extended morphology. Our observations in optical and near-infrared indicate that it is a typical fold-type quadruple lens, with a maximum separation between images of $1\\farcs76$.  From the spectrum we measured a lens redshift of $z_l=0.373$. Standard simple elliptical mass models fit the data well, including the flux ratios, implying no evidence for substructure. The mass modelling suggests an important contribution of the external shear, probably from a nearby bright galaxy. There is also a foreground group of galaxies whose centre is $\\sim120''$ from the lens. The lens galaxy is $\\sim 1$~mag brighter than predicted by mass modelling.  Thus far the SQLS has discovered 31 lensed quasars including SDSS J1330+1810\\footnote{The current status of the survey is summarised at http://www-utap.phys.s.u-tokyo.ac.jp/\\~{}sdss/sqls/}, of which 5 are quadruple lenses \\citep[e.g.,][]{inada03,kayo07}. The low fraction of quadruple lenses may imply that the faint end slope of the quasar optical luminosity function is shallow, although it is marginally consistent with standard theoretical expectations \\citep{oguri07b}. The SQLS has completed $\\sim 2/3$ of its survey, implying that a few more quadruple lenses will be discovered by the completion of the survey."
    },
    "0809/0809.1725_arXiv.txt": {
        "abstract": "Using new \\ubvri\\halpha CCD photometric observations and the archival infrared and X-ray data, we have carried out a multi-wavelength study of a Perseus arm young galactic star cluster NGC 7419. An age of $22.5\\pm3.0$ Myr and a distance of $3230^{+330}_{-430}$ pc are derived for the cluster. Our photometric data indicates a higher value of color excess ratio $E(U-B)/E(B-V)$ than the normal one. There is an evidence for mass segregation in this dynamically relaxed cluster and in the range $1.4-8.6M\\subsun$, the mass function slope is in agreement with the Salpeter value. Excess emissions in near-infrared and \\halpha support the existence of a young ($\\le 2$ Myr) stellar population of Herbig Ae/Be stars ($\\geq 3.0 M_\\odot$) indicating a second episode of star formation in the cluster region. Using XMM-Newton observations, we found several X-ray sources in the cluster region but none of the Herbig Ae/Be stars is detected in X-rays. We compare the distribution of upper limits for Herbig Ae/Be stars with the X-ray distribution functions of the T-Tauri and the Herbig Ae/Be stars from previous studies, and found that the X-ray emission level of these Herbig Ae/Be stars is not more than $L_X \\sim 5.2\\times 10^{30}$ \\egs,  which is not significantly higher than for the T-Tauri stars. Therefore, X-ray emission from Herbig Ae/Be stars could be the result of either unresolved companion stars or a process similar to T-Tauri stars. We report an extended X-ray emission from the cluster region NGC 7419, with a total X-ray luminosity estimate of $\\rm{\\sim 1.8\\times10^{31}~erg~s^{-1}~arcmin^{-2}}$. If the extended emission is due to unresolved emission from the point sources then we estimate $\\sim$288 T-Tauri stars in the cluster region each having X-ray luminosity $\\rm{\\sim 1.0\\times10^{30}~erg~s^{-1}}$. Investigation of dust attenuation and ${}^{12}$CO emission map of a square degree region around the cluster indicates the presence of a foreground dust cloud which is most likely associated with the local arm star forming region (Sh2-154). This cloud harbors uniformly distributed pre-main-sequence stars ($0.1-2.0M\\subsun$), with no obvious trend of their distribution with either $(H-K)$ excess or \\av. This suggests that the star formation in this cloud depend mostly upon the primordial fragmentation. ",
        "introduction": "\\label{sec:int}  NGC 7419 ($\\rm RA_{J2000}=22^h 54^m 20^s$, ${\\rm DEC_{J2000}}=+60\\degr 48\\arcmin 54\\arcsec$; $l=109\\fdg13$, $b=1\\fdg12$), is a moderately populated young and heavily reddened galactic star cluster in Cepheus with a large number of Be stars. The cluster contains high mass ($\\geq10M\\subsun$); intermediate mass ($2-10M\\subsun$) and low-mass ($\\leq2M\\subsun$) stars. It is therefore an ideal laboratory for the study of initial stellar mass distribution as well as duration of star formation process in a molecular cloud. Presence of statistically significant number of Herbig Ae/Be stars in the cluster makes it very attractive for understanding the formation of these stars and origin of various atmospheric activities like \\halpha emission and X-ray emission in them. However, to address these questions in detail, one would like to know accurate distance and age parameters of the cluster NGC 7419, which is lacking despite a number of photometric and spectroscopic studies. This is mainly because of the fact that the cluster is heavily reddened in comparison to the nearby clusters situated at the similar distances, and suffers from variable reddening.  In order to determine cluster reddening reliably, accurate $UBV$ broadband photometry of early type stars is essential. A comparison of the photometries available in the literature indicates that $UB$ data, many have systematic calibration error. For example, \\citet{beauchamp94} have mentioned that their color may have an offset of $\\sim0.2$ mag due to the calibration problems in $U$ band. Their photometric observations have been carried out in the poor seeing ($2\\farcs5-4\\farcs0$) conditions. This will affect cluster photometric data particularly in the crowded regions.  The main goals of present study are to determine the distance, age and its spread, and mass function (MF) of the cluster as accurate as possible. This will help us to understand the star formation history of the cluster, and to investigate the X-ray emission properties of Herbig Ae/Be stars. Deep optical \\ubvri observations ($V\\sim22.0$ mag), narrow band \\halpha photometric observations along with the Two Micron All Sky Survey (2MASS), Infrared All Sky Survey (IRAS), Midcourse Space experiment (MSX) and XMM-Newton archival data are used to understand the X-ray emission properties of Herbig Ae/Be stars, and the global scenario of star formation in the cluster NGC 7419 and its surrounding region.  \\citet{blanco55} has reported the distance of this cluster as $\\sim$ 6 kpc based on the $RI$ photometric observations. A similar value for cluster distance has also been obtained by \\citet{moffat73}. However, \\citet{hulst54} has obtained a significantly smaller  distance of 3.3 kpc. Using CCD data, a distance of 2.0 kpc and 2.3 kpc was estimated by \\citet{bhatt93} and \\citet{beauchamp94}, respectively. The age estimated by \\citet{bhatt93} is $\\sim$ 40 Myr while \\citet{beauchamp94} have estimated a much younger age of $\\sim$ 14 Myr. Recent CCD observations reported by \\citet{subramaniam06} estimated its distance as 2.9 kpc and an age of 20--25 Myr.  The paper describes optical observations and the derivation of cluster parameters in \\S\\ref{sec:phoDat} and \\S\\ref{sec:proClu}. The near-infrared (NIR) data are dealt in \\S\\ref{sec:nirDat}, while distribution of young stellar objects (YSOs), MF and mass segregation are given in  \\S\\ref{sec:spaDis}, \\S\\ref{sec:masFun} and \\S\\ref{sec:masSeg}. Finally, the X-ray data and its analysis (for the first time) are described in \\S\\ref{sec:xraDat}, followed by the summary and conclusions in \\S\\ref{sec:sumCon}. ",
        "conclusions": "\\label{sec:sumCon}  A deep optical $UBVRI$ and narrow band \\halpha observations along with  multi-wavelength archival data from the surveys such as 2MASS, MSX, IRAS and XMM-Newton are used to understand the global scenario of star formation and the basic parameters of the cluster NGC 7419. XMM-Newton archival data have also been used to study the X-ray emission mechanism from the cluster.  The radius of the cluster NGC 7419 has been found to be 4\\farcm0$\\pm$0\\farcm5 using radial density profile. The reddening law in the direction of the cluster is found to be normal at longer wavelengths but anomalous at shorter wavelengths. Reddening, $E(B-V)$, is found to be varying between 1.5 to 1.9 mag with a mean value $\\sim 1.7\\pm0.2$ mag. The turn-off age and the distance of the cluster are estimated to be $22.5\\pm2.5$  Myr and $\\rm{3230^{+330}_{-430}}$, respectively. The MF for the main-sequence stars in the cluster is estimated as having $ \\Gamma =-1.10 \\pm 0.19 $ in the mass range $8.6 M\\subsun < M \\leq 1.4 M\\subsun$, which is a similar to the \\citet{salpeter55} value. Effect of mass segregation is found in the main-sequence stars which may be the result of dynamical evolution.  Using the NIR color-color diagram and narrow band $\\halpha$ observations, we have identified 21 Herbig Ae/Be in the cluster region with the masses lying between 3 to $7 M\\subsun$. The ages of these Herbig Ae/Be  stars are found to be in the range of $\\sim$ 0.3 to 2.0 Myr. The significant difference between  turn-off age and  turn-on age of the cluster represents a second episode of star formation in the cluster. We have found 90 YSOs having masses in the range from 0.1 to $2.0 M\\subsun$ with the help of NIR color-color diagram around the cluster. The presence of  such a large number of  NIR excess sources (T-Tauri stars) shows a recent star formation episode in the surroundings of the cluster region. Using extinction, dust and $^{12}$CO maps, we found that these YSOs are probably associated with a foreground star forming region Sh-154 and not related with the cluster region. We found no obvious trend in spatial distribution of YSOs with $A_V$ and $(H-K)_{\\rm excess}$. The dispersion in $A_V$ and $(H-K)_{\\rm excess}$ is also very low which indicates that a majority of the YSOs are born at the same time in the same environment. Therefore, it is possible that primordial fragmentation of the cloud may be responsible for the formation of the low mass stars within this cloud.   We have detected 66 X-ray sources in the observed field by XMM-Newton observatory using archival X-ray data. Out of these sources, 23 are knowm to be the probable members of the cluster based on analysis of their X-ray colors. Fifteen X-ray sources are without any optical or NIR counterparts. These may be the young embedded sources which need to be investigated further. We have derived the detection limits of X-ray observations based on the position of the X-ray source in the field of view and on its energy spectrum. Thus, the median value for the detection limit for the 21 Hebig Ae/Be stars in the field is $\\rm{L_X \\sim}$$\\rm{5.2\\times 10^{30}~erg~s^{-1}}$. We have compared the XLF for the cluster members with the T-Tauri stars and Herbig Ae/Be stars. Because of insufficient exposure and resolution, the sensitivity was not enough to reach the level of the median X-ray luminosity observed in T-Tauri stars in Taurus-Auriga region and Herbig Ae/Be stars in \\citet{stelzer06}. Therefore a conclusive comparison of X-ray properties of the stars cannot be made. However, the comparison indicates that Herbig Ae/Be stars in NGC 7419 tend to be less X-ray luminous than in the sample of \\citet{hamaguchi05}, which shows that X-ray activity level of the Herbig Ae/Be stars is not more than in the T-Tauri stars. Therefore, we can support the binary T-Tauri companion hypothesis for the generation of  X-rays  in Herbig Ae/Be stars. It is also possible that a Herbig Ae/Be star is itself emitting X-rays but the level of the X-ray emission is similar to that of the T-Tauri stars. The cluster region shows an extended X-ray emission with a total luminosity estimated to be $\\rm{ L_X \\approx 1.8\\times10^{31}~erg~s^{-1}~arcmin^{-2}} $. This diffuse emission might be the result of X-ray emission from T-Tauri type stars which could not be resolved by XMM-Newton observations. It requires $\\sim$288 T-Tauri stars  each having $\\rm{L_X}$ $\\rm{\\sim 1.0\\times10^{30}~erg~s^{-1}}$, if it originates from such stars. High resolution deep observations such as from CHANDRA are required for a detailed analysis of this cluster region."
    },
    "0809/0809.1349_arXiv.txt": {
        "abstract": "We present the {\\it AEGIS-X} survey, a series of deep \\chandra\\ ACIS-I observations of the Extended Groth Strip. The survey comprises pointings at 8 separate positions, each with nominal exposure 200ks, covering a total area of approximately 0.67 deg$^{2}$ in a strip of length 2 degrees. We describe in detail an updated version of our data reduction and point source detection algorithms used to analyze these data. A total of 1325 band-merged sources have been found to a Poisson probability limit of $4 \\times 10^{-6}$, with limiting fluxes of $5.3 \\times 10^{-17}$ erg cm$^{2}$ s$^{-1}$ in the soft (0.5--2 keV) band and $3.8 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ in the hard (2--10 keV) band. We present simulations verifying the validity of our source detection procedure and showing a very small, $<1.5$\\%, contamination rate from spurious sources. Optical/NIR counterparts have been identified from the DEEP2, CFHTLS, and Spitzer/IRAC surveys of the same region. Using a likelihood ratio method, we find optical counterparts for 76\\% of our sources, complete to $R_{AB}$=24.1, and, of the 66\\% of the sources that have IRAC coverage, 94\\% have a counterpart to a limit of 0.9 $\\mu$Jy at 3.6 $\\mu m$ ($m_{AB}$=23.8). After accounting for (small) positional offsets in the 8 \\chandra\\ fields, the astrometric accuracy of the \\chandra\\ positions is found to be $0\\farcs8$ RMS, however this number depends both on the off-axis angle and the number of detected counts for a given source. All the data products described in this paper are made available via a public website. ",
        "introduction": "\\label{sec:intro}  There is intense current interest in the formation of galaxies, the history of star formation and accretion power in the universe, and the inter-relations between these phenomena. This has motivated the investment of major observational resources in obtaining images and spectroscopy in various areas of the sky. These multi-wavelength surveys cover a wide range of areas, depths, angular resolutions and wavebands, ranging from very deep ``pencil\" beam surveys in small areas of the sky  to wider surveys at shallower depths.  Surveys at X-ray wavelengths provide an important component of the multiwavelength arsenal, primarily because of the efficiency and sensitivity of X-ray emission in selecting Active Galactic Nuclei (AGN). There is apparently a strong relationship between AGN and their host galaxy bulges \\citep[e.g.][]{fm00,ge00}, so studying the growth of supermassive black holes in the context of the evolution of their host galaxies is likely to be a fruitful area of exploration. X-ray surveys also enable an estimate of the history of accretion power in the universe and the origin of the X-ray background \\citep[e.g.][]{bh05}.  The deepest X-ray surveys in existence are the \\chandra\\ Deep Fields North and South (hereafter the CDF-N and CDF-S, respectively), which each cover an area of $\\sim 0.1$~deg$^{2}$ to nominal depths of $\\sim 2-3 \\times 10^{-17}$~erg cm$^{-2}$ s$^{-1}$ \\citep{al03,lu08}. The widest surveys to resolve a significant fraction of the X-ray background cover of order 10 deg$^{2}$, but to much shallower (by a factor $\\sim 100$) depths (e.g. XBootes: \\citealt{mu05}; \\citealt{ke05}; XMM-LSS: \\citealt{pi04}). The major advantage of ultradeep surveys is that they are able to probe into the distant universe, to detect ``typical\" objects at high redshift. On the other hand, larger areas are required to sample significant large scale structures in the universe and hence determine the relationship between galaxy evolution and local environmental density. Moreover, larger area surveys are able to sample and find unusual, rare objects. For a complete picture it is clearly also necessary to explore the parameter space in between the ultradeep and ultrawide surveys.  Motivated by this, we have obtained deep X-ray observations using the \\chandra\\ X-ray observatory in a region of sky of area $\\sim 0.5$~deg$^{2}$ known as the Extended Groth Strip (EGS), covering the energy range 0.5--7~keV. Extensive multiwavelength data in this region have been obtained as part of the ``AEGIS\" project \\citep{da07} making it one of the premier datasets to study the co-evolution of black hole accretion and galaxy formation. The purpose of the present paper is to describe the X-ray dataset and reduction, and present a catalog of point sources derived from the \\chandra\\ X-ray survey, which we designate \\ax.  The structure of the paper is as follows. In \\S\\ref{sec:dr} we describe the data and the reduction, including the method used to calculate the \\chandra\\ point spread function. Detailed descriptions of our source detection and photometry procedures are given in \\S\\ref{sec:detphot}. In \\S\\ref{sec:results} the results of the source detection and analysis are presented. In \\S\\ref{sec:counterpart} the optical and infrared counterparts to the X-ray sources are provided and in \\S\\ref{sec:conc} the conclusions are given. ",
        "conclusions": ""
    },
    "0809/0809.4270_arXiv.txt": {
        "abstract": "Photometric calibration to $\\sim5\\%$ accuracy is frequently needed at arbitrary celestial locations; however, existing all-sky astronomical catalogs do not reach this accuracy and time consuming photometric calibration procedures are required. I fit the Hipparcos $B_{T}$, and $V_{T}$ magnitudes, along with the 2MASS $J$, $H$, and $K$ magnitudes of Tycho-2 catalog-stars with stellar spectral templates. From the best fit spectral template derived for each star, I calculate the synthetic SDSS $griz$ magnitudes and constructed an all-sky catalog of $griz$ magnitudes for bright stars ($V\\ltorder12$). Testing this method on SDSS photometric telescope observations, I find that the photometric accuracy, for a single star, is usually about $0.12$, $0.12$, $0.10$ and $0.08$\\,mag (1\\,$\\sigma$), for the $g$, $r$, $i$, and $z$-bands, respectively. However, by using $\\sim10$ such stars, the typical errors per calibrated field (systematic + statistical) can be reduced to about 0.04, 0.03, 0.02, and 0.02\\,mag, in the $g$, $r$, $i$, and $z$-bands, respectively. Therefore, in cases for which several calibration stars can be observed in the field of view of an instrument, accurate photometric calibration is possible. ",
        "introduction": "\\label{Introduction}  Often in astronomical research, absolute photometric accuracy better than $10\\%$, is required. In many cases, the method of choice is to observe photometric standards (e.g., Landolt 1992). However, this requires photometric observing conditions, and additional observations. The Sloan Digital Sky Survey (SDSS; York et al. 2000) provides an excellent photometric calibration in $ugriz$ bands, with accuracy better than $2\\%$ (Adelman-McCarthy et al. 2008). However, this is available only for about a quarter of the sky. Other all-sky visible-light catalogue, like the USNO-B1 (Monet et al. 2003) and the US Naval Observatory CCD astrograph catalogue (Zacharias et al. 2004) provide relatively poor photometric accuracy. The magnitudes of individual stars in these catalogue are accurate to about $0.3$\\,mag, and even with a large number of stars there is still a considerable field-to-field systematic errors.  In this paper, I calculate the SDSS $griz$ magnitudes of Tycho-2 stars over the entire sky. In case $\\gtorder10$ of these stars are visible in a camera field of view, these stars can be used to photometrically calibrate an astronomical image to accuracy of better than $0.04$\\,mag. The only overhead is that typically a shorter exposure in which the Tycho-2 catalog stars will not be saturated is required. ",
        "conclusions": "\\label{Disc}   In order to test the accuracy of the derived synthetic magnitudes, I constructed a catalog of all the photometric measurements available from the SDSS photometric telescope ``secondary patches'' fields\\footnote{Available from: http://das.sdss.org/PT/} (Tucker et al. 2006). From this catalog I selected all the non-saturated stars brighter than 12th magnitude in both the $g$- and $r$-bands. Davenport et al. (2007) analyzed the systematic offset between the SDSS magnitude system and the SDSS photometric telescope bands. They found that the magnitudes of very red stars are different in the two systems. Using the transformations given by Davenport et al. (2007) I converted all the SDSS photometric-telescope magnitudes to the SDSS system.  Next I cross-correlated this list with the catalog of $griz$ synthetic magnitudes presented in \\S\\ref{Cat}, and selected stars which have $JHK$ magnitudes fainter than 6 (i.e., not saturated); $B_{T}$ and $V_{T}$ magnitude errors smaller than 0.15; and RMS of the best RMS-fit template of less than 0.15\\,mag. For each of the 3714 stars satisfying these criteria, I compared the corrected magnitudes (i.e., Davenport et al. 2007) as measured by the SDSS photometric telescope with its best-fit synthetic magnitude. In Figure~\\ref{fig:Mag_SDSS_vs_Fit}, I show histograms of the fitted synthetic magnitudes minus the SDSS magnitudes for the $griz$-bands. The median value, along with half the range containing $68\\%$ of the stars, are indicated in each panel. \\begin{figure*} \\centerline{\\includegraphics[width=16.0cm]{f2.eps}} \\caption{Histograms of the actual SDSS magnitude minus the best fit magnitude for a set of 3714 stars which magnitude was measured by the SDSS photometric telescope. The histograms are shown for the $g$, $r$, $i$, and $z$-bands (from left to right), respectively. The upper row is for the best-fit magnitude obtained by minimizing the RMS, while the lower row is for the $\\chi^{2}$ fit, which is somewhat better. The median value is the median of the histogram, while std is half of the range containing $68\\%$ of the values in the histogram (rather than the usual definition). \\label{fig:Mag_SDSS_vs_Fit}} \\end{figure*} The upper row is for the best-fit magnitude by minimizing the RMS, while the lower row is for the $\\chi^{2}$ fit. This plot suggests that the $\\chi^{2}$ fit is marginally better. Therefore, it should be preferred over the RMS-fit. In all the bands there are small offsets, listed in Fig.~\\ref{fig:Mag_SDSS_vs_Fit}, between the SDSS photometric telescope magnitudes and the derived synthetic magnitudes. The magnitudes in the catalog (Table~\\ref{Tab-CatDesc}) are corrected for these small offsets. I note that I repeated this test using 42 SDSS photometric standards\\footnote{Electronic version available from: http://home.fnal.gov/$\\sim$dtucker/ugriz/tab08.dat} (Smith et al. 2002), and found similar results.    The field of view of large format cameras may contain only several suitable Tycho-2 stars (see Fig.~\\ref{fig:Tyc2_StarDen}). To estimate the uncertainty in magnitude calibration when using several stars, I have carried out the following simulations: I randomly selected $N$ stars, for any $N$ between 2 and 100, out of the 3714 SDSS standard stars that I used for the comparison of the derived magnitudes with the actual SDSS magnitudes. For each set of randomly selected $N$ stars I calculated the median of differences between the $\\chi^{2}$-fitted synthetic magnitudes and the SDSS magnitudes. Next, I repeated this procedure 10,000 times (for each $N$), and calculated the standard deviation of the 10,000 median of differences. \\begin{figure} \\centerline{\\includegraphics[width=8.5cm]{f3.eps}} \\caption{The expected error in photometric calibration as a function of the number of stars used in the calibration process (see text). The thick-solid line, thick-dashed line, thin-dashed line, and thin-solid line, are for the $g$, $r$, $i$, and $z$-bands, respectively. \\label{fig:Error_vs_Nstars}} \\end{figure} In Figure~\\ref{fig:Error_vs_Nstars} I show the standard deviation of the distribution of the median of differences between the synthetic magnitudes and the measured SDSS magnitudes, as a function of $N$. The plot suggests that, for example, by using five stars the errors in photometry reduce to 0.07\\,mag, 0.04\\,mag, 0.03\\,mag and 0.03\\,mag in the $g$, $r$, $i$ and $z$-bands, respectively. When using ten stars the errors reduce to about 0.04\\,mag, 0.03\\,mag, 0.02\\,mag and 0.02\\,mag in the $g$, $r$, $i$ and $z$-bands, respectively. In order to remove outliers it is important to use the median of differences (i.e., and not mean of differences).    To summarize, I suggest an alternative method for photometric calibration that may work to accuracy of about several percents. This method relies on selected stars in the Tycho-2 catalog for which I fitted spectral templates to their $B_{T}V_{T}JHK$ magnitudes. Two major limitation of this method is that several Tycho-2 stars are needed in the field of view, and that shorter exposures, in which these stars are not saturated, are required. Given the catalog star density as a function of Galactic latitude (Fig.~\\ref{fig:Tyc2_StarDen}), this method is applicable for instruments with large field of view or for low Galactic latitude observations."
    },
    "0809/0809.3982_arXiv.txt": {
        "abstract": "The LMC clusters with similar ages to the Milky Way open clusters are in general more metal-poor and more populous than the latter, being located close enough to allow their stellar content to be well resolved. Therefore, they are unique templates of simple stellar population (SSP), being crucial to calibrate models describing the integral light as well as to test the stellar evolution theory. With this in mind we analyzed HST/WFPC2 (V, B--V) colour-magnitude diagrams (CMDs) of 15 populous LMC clusters with ages between $\\sim$0.3 Gyr and $\\sim$4 Gyr using different stellar evolutionary models. Following the approach described by \\cite[Kerber, Santiago \\& Brocato (2007)] {KSB07}, we determined accurate and self-consistent physical parameters (age, metallicity, distance modulus and reddening) for each cluster by comparing the observed CMDs with synthetic ones generated using isochrones from the PEL and BaSTI libraries. These determinations were made by means of simultaneous statistical comparison of the main-sequence fiducial line and the red clump position, offering objective and robust criteria to select the best models. We compared these results with the ones obtained by \\cite[Kerber, Santiago \\& Brocato~(2007)]{KSB07} using the Padova isochrones. This revealed that there are significant trends in the physical parameters due to the choice of stellar evolutionary model and treatment of convective core overshooting. In general, models  that incorporate overshooting presented more reliable results than those that do not. Furthermore, the Padova models fitted better the data than the PEL and BaSTI models. Comparisons with the results found in the literature demonstrated that our derived metallicities are in good agreement with the ones from the spectroscopy of red giants. We also confirmed that, independent of the adopted stellar evolutionary library, the recovered 3D distribution for these clusters is consistent with a thick disk roughly aligned with the LMC disk as defined by field stars. Finally, we also provide new estimates of distance modulus to the LMC center, that are marginally consistent with the canonical value of 18.50. ",
        "introduction": "The LMC contains a rich system of stellar clusters, with more than three thousand cataloged objects (\\cite[Bica \\etal\\ 2008]{Bica_etal08}) and covering ages from few Myr to about 13 Gyr. There are about one hundred that can be considered as populous ones ($> 10^{5}$ stars), which offer the opportunity to recover the age-metallicity relation for the LMC by means of accurate age (from CMD analysis) and metallicity (from spectroscopy analysis) determinations. Furthermore, the objects with ages between $\\sim$ 0.3 and 4 Gyr -- the intermediate-age LMC clusters (IACs) -- are more metal poor than the open clusters in the Milky Way, being therefore fundamental pieces in the local universe to calibrate integrated light models as well as to test the evolutionary models in the sub-solar metallicity regime.  Taking the advantage of the superior photometric quality of the HST and the intrinsic large stellar statistics for the populous IACs we have applied a statistical method to recover accurate physical parameters -- not only age, but also metallicity, distance and reddening -- for these objects in a self-consistent way. By self-consistency we mean the ability to simultaneously infer these parameters from the same data-set without any prior assumptions about any of them. The method, presented by \\cite[Kerber, Santiago \\& Brocato~(2007)~(hereafter KSB07)]{KSB07} using the Padova isochrones (\\cite[Girardi \\etal\\ 2002]{Girardi_etal02}), joins CMD modelling and statistical analysis to objectively determine which are the synthetic CMDs that best reproduce the observed ones. In the present work we expanded this analysis to two other stellar evolutionary libraries: PEL (or Pisa, \\cite[Castellani \\etal\\ 2003]{Castellani_etal03}) and BaSTI (or Teramo, \\cite[Pietrinferni \\etal\\ 2004]{Pietrinferni_etal04}). This allowed us to quantify how the recovered physical parameters depend on the adoption of different stellar evolutionary libraries, including the treatment for the convective core overshooting. Furthermore, since we are determining the individual distance to each cluster, we could also probe the three dimensional distribution of these clusters, which seems to be roughly aligned with LMC disk (\\cite[KSB07]{KSB07}; \\cite[Grocholski \\etal\\ 2007]{Grocholski_etal07}), and to obtain new determinations of distance modulus to the LMC centre. ",
        "conclusions": ""
    },
    "0809/0809.5276_arXiv.txt": {
        "abstract": "{} {This paper reports on a search for new classical nova candidates in the M81 galaxy based on archival, as well as recent, new images.} {We used images from 1999--2007 to search for optical transients in M81. The positions of the identified classical nova candidates were used to study their spatial distribution. Kolmogorov - Smirnov test (KS) and bottom-to-top (BTR) ratio diagnostic were used to analyze the nova candidate distribution and differentiate between the disk and the bulge populations.} {In total, 49 classical nova candidates were discovered. In this study, we present the precise positions and photometry of these objects, plus the photometry of an additional 9 classical nova candidates found by Neill \\& Shara (2004). With our large sample, we find a different spatial distribution of classical nova candidates when compared to the results of earlier studies. Also, an extraordinarily bright nova was found and studied in detail.} {} \\keywords {galaxies: individual: M81 -- binaries: close -- novae} ",
        "introduction": "Novae are important objects for the study of close binary evolution, but our location in the Milky Way prevents us from getting an unbiased sample locally. Studying novae in nearby galaxies can provide a more homogeneous sample of these objects. For galaxies several Mpc away, long-term monitoring coupled with a rapid cadence using a relatively large telescope is necessary. Significant amounts of telescope time are difficult to obtain  so searching for nova candidates using archival images (already obtained for a variety of different science studies) have the possibility for getting useful results. This method has several disadvantages,including a lack of control over cadence, bandpasses, and exposure depth.  The outburst of classical novae (CNe) are caused by explosive hydrogen burning on the white dwarf (WD) surface of a close binary system with material transfer from the companion star onto the WD surface. During the thermonuclear runaway, a fraction of the envelope is ejected, while a part of it remains in steady nuclear burning on the WD surface (Jos\\'e \\& Hernanz 1998; Prialnik \\& Kovetz 1995). This powers a supersoft X-ray source (SSS). The duration of the SSS phase is inversely related to the WD mass (Pietsch et al. 2006). Since WD envelope models also show that the duration of the SSS phase depends on the metalicity of the envelope, monitoring of the SSS phase of CNe also provides important information about the chemical composition of the post-outburst envelope (Pietsch et al. 2006). Results of recent work aimed at X-ray monitoring optical novae in M31 (Pietsch et al. 2006) bring new important results, while showing the necessity of having a good catalog of optical novae available for such studies.  Nearby galaxies with high annual nova rates are the best targets for statistically conclusive studies of the properties of extragalactic novae at optical, as well as X-ray, wavelengths. Besides M31, the M81 galaxy is another nearby large spiral galaxy. Only two recent papers aimed at the study of CNe in M81 have been published up to now -- Shara et al. (1999) and Neill \\& Shara (2004). Here, we take advantage of M81 being a relatively common target for optical imaging and, in order to search for classical nova (CN) candidates in this galaxy, we analyze available archival CCD images, together with our recent images. ",
        "conclusions": "We discovered and classified 49 transient objects as M81 CN candidates. Together with previously known CN candidates (35 objects in total; published by Shara et al., 1999 and by Neill \\& Shara, 2004), the number of all CN candidates is more than doubled. These results are important for future studies concerning the identification of optical counterparts of supersoft X-ray sources in the M81 galaxy and for possible identification of recurrent novae.  {The relatively large number of CN candidates in this sample provides a more accurate view of the CN spatial distribution in M81 than we had before this study. There is no strong evidence of asymmetry in the distribution of our CN candidates across the major axis of M81. We cannot sustain the claim of a very high bulge-to-disk nova ratio in M81. Our results from the BTR diagnostic have give a low bulge-to-disk nova ratio and thus support the existence of a significant disk nova population.  The KS test applied to the radial distribution of our CN candidates and the M81 light distributions indicate a good match between the CN candidates and the total light of the parent galaxy with a high confidence level. This indicates that both CN populations, i.e., the disk and the bulge, are certainly present in the galaxy.  However, since we cannot rule out the possibility that we have missed a few novae in the outer part of the galaxy due to sparse time coverage with respect to the central part of the galaxy (and thus we tested conformity of distribution of galaxy light and CN candidates only up to radius of 323\\arcs), the conclusions from this method are probably of lower significance than the BTR statistics.  We are not able to refine the annual nova rate due to time gaps in the coverage and a generally low cadence of images. However, the number of CN candidates detected is consistent with the previous result of 30 yr$^{-1}$ (Neill \\& Shara, 2004) within relatively large uncertainties.  We have shown that relatively inhomogeneous material, as well as use of archival data originally obtained for different purposes, can be used to obtain significant results. However, only a long-period (three years running at least), deep (going down to $H_{\\alpha}$ = 21 mag) and comprehensive survey covering the whole galaxy with minimum  detection incompleteness is able to provide an accurate bulge-to-disk nova ratio. Such a survey would significantly refine the characteristics of both bulge and disk nova populations as well as the total annual nova rate.  \\medskip"
    },
    "0809/0809.0515_arXiv.txt": {
        "abstract": "{ We simulate anisotropic outflows of AGN, and investigate the large-scale impact of the cosmological population of AGN outflows over the Hubble time by performing N-body $\\Lambda$CDM simulations. % Using the observed quasar luminosity function to get the redshift and luminosity distribution, and analytical models for the outflow expansion, AGNs are allowed to evolve in a cosmological volume. By the present epoch, $13 - 25\\%$ of the total volume % is found to be pervaded by AGN outflows, with $10^{-9}$ G magnetic field. ",
        "introduction": "\\label{sec-intro}  Outflows from AGN are observed in a wide variety of forms: radio galaxies, broad absorption line quasars, Seyfert galaxies exhibiting intrinsic absorption in the UV, broad emission lines, warm absorbers and absorption lines in X-rays \\citep[e.g.,][]{crenshaw03, everett07}. There have been studies on the cosmological impact of quasar outflows in large scales (\\citealt{FL01}, hereafter FL01; \\citealt{so04}, hereafter SO04; \\citealt{lg05}, hereafter LG05). \\citet{barai08} investigated the cosmological influence of radio galaxies over the Hubble time. All these studies considered spherically expanding outflows.  On cosmological scales an outflow is expected to move away from the high density regions of large-scale structures, with the outflowing matter getting channelled into low-density regions of the Universe \\citep{martel01}. We implement such anisotropic AGN outflows within a cosmological volume. The simulation methodology is given in \\S\\ref{sec-numerical}, and the results are discussed in \\S\\ref{sec-results}. ",
        "conclusions": "\\label{sec-results}   \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{f1.eps}} \\caption{ \\footnotesize Volume of the simulation box filled by AGN outflows ($N_{\\rm AGN}$) as a fraction of the total volume {\\it (solid)}, and as a fraction of volumes of various overdensities: $N (\\rho > \\overline{\\rho})$ {\\it (dash dot)}, $N (\\rho > 2 \\overline{\\rho})$ {\\it (dashed)}, $N (\\rho > 3 \\overline{\\rho})$ {\\it (dotted)}, $N (\\rho > 5 \\overline{\\rho})$ {\\it (dash dot dot dot)}, $N (\\rho > 7 \\overline{\\rho})$ {\\it (long dashes)}. } \\label{fig1} \\end{figure}  At each timestep the total volume occupied by the AGN outflows is computed by counting the contributions of all the sources born by then, both the active ones and those in the anisotropic phase. We performed 4 simulations with opening angles of $\\alpha = 60^{\\rm o}, 90^{\\rm o}, 120^{\\rm o}, 180^{\\rm o}$, all with $\\epsilon_K = 0.1$, and one with $\\alpha = 120^{\\rm o}$ and $\\epsilon_K = 0.05$, whose results are shown in the figures.  We count the grid cells in the simulation box which occur inside the volume of one or more AGN outflows. The total number of these filled cells, $N_{\\rm AGN}$, gives the total volume of the box occupied by outflows. We express the total volume filled as a fraction of volumes of various overdensities in the box, $N_{\\rho} = N (\\rho > {\\cal C} \\overline{\\rho})$, where $\\overline{\\rho} = (1+z)^3 \\Omega_M 3H_0^2 / (8 \\pi G)$ is the mean matter density of a spatially flat Universe (the box) at an epoch $z$. So $N_{\\rho}$ gives the number of cells which are at a density ${\\cal C}$ times the mean density. We find $N_{\\rho}$ for ${\\cal C} = 0, 1, 2, 3, 5, 7$; ${\\cal C} = 0$ gives the total volume of the box.  Figure~\\ref{fig1} shows the redshift evolution of the volume filling factors for our 5 simulation runs. With $10\\%$ kinetic efficiency, $0.13$ of the entire Universe is filled at present % by AGN outflows with an opening angle of $60^{\\rm o}$; the fraction increases to $0.17$ with $90^{\\rm o}$, $0.21$ with $120^{\\rm o}$, and $0.25$ with $180^{\\rm o}$. A $5\\%$ kinetic efficiency and $\\alpha = 120^{\\rm o}$ fills $0.13$ of the volume. In all our runs, the outflows fill up all of the regions with $\\rho > 2 \\overline{\\rho}$ % by $z=0.3$. % With $\\epsilon_K = 0.1$ and $\\alpha = 90^{\\rm o}$ or higher, the outflows permeate all the overdense regions ($\\rho > \\overline{\\rho}$) by $z = 0.1$.   \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{f2.eps}} \\caption{ \\footnotesize Volume fraction filled by AGN outflows ({\\it top}), the volume weighted average of the total energy density inside outflow volumes $\\langle u_E \\rangle$ ({\\it middle}), and the equipartition magnetic field within the filled volumes $\\langle B_0 \\rangle$ ({\\it bottom}). } \\label{fig2} \\end{figure}  It is the overdense regions of the Universe which gravitationally collapse to form stars and galaxies. So evidently the AGN outflows have a profound cosmological impact on the protogalactic regions. We note that the volumes obtained by LG05 (100\\% filling by $z\\sim1$) are higher than ours.  We perform preliminary estimates of the energy density and magnetic field in the volumes of the Universe filled by the AGN outflows. The energy density inside the outflow behaves similar to the outflow pressure evolving adiabatically (\\S\\ref{sec-num-aniso}), $u_E = 3 p_0$. Assuming equipartition of energy between magnetic field of strength $B_0$ and relativistic particles inside the outflow, the magnetic energy density is $u_B = u_E / 2 = B_0^2 / (8 \\pi)$. We define the volume weighted average of a physical quantity ${\\cal A}$ as $\\langle {\\cal A} \\rangle (z) \\equiv \\sum ({\\cal A} V_0) / \\sum V_0$, where the summation is over all outflows existing in the simulation box at that epoch.  Figure~\\ref{fig2} shows the redshift evolution of the total volume filling fraction, $\\langle u_E \\rangle$ and $\\langle B_0 \\rangle$. The energy density and magnetic field decrease with redshift as larger volumes are filled. At $z = 0$, a magnetic field of $\\sim 10^{-9}$ G permeates the filled overdense volumes, consistent with the results of \\citet{ryu08}. At a given redshift, the energy density and magnetic field are larger for smaller opening angles of the anisotropic outflows.   We conclude that, using our N-body simulations, the cosmological population of AGN outflows pervade $13-25\\%$ of the volume of the Universe by the present. A magnetic field of $\\sim 10^{-9}$ G is infused in the filled volumes at $z = 0$."
    },
    "0809/0809.0209_arXiv.txt": {
        "abstract": "Since that very memorable day at the Beijing 2008 Olympics, a big question on every sports commentator's mind has been ``What would the 100 meter dash world record have been, had Usain Bolt not celebrated at the end of his race?'' Glen Mills, Bolt's coach suggested at a recent press conference that the time could have been 9.52 seconds or better. We revisit this question by measuring Bolt's position as a function of time using footage of the run, and then extrapolate into the last two seconds based on two different assumptions. First, we conservatively assume that Bolt could have maintained Richard Thompson's, the runner-up, acceleration during the end of the race. Second, based on the race development prior to the celebration, we assume that he could also have kept an acceleration of 0.5 m/s$^2$ higher than Thompson. In these two cases, we find that the new world record would have been $9.61\\pm0.04$ and $9.55\\pm0.04$ seconds, respectively, where the uncertainties denote 95\\% statistical errors. ",
        "introduction": "On Saturday, August 16th 2008, Usain Bolt shattered the world record of 100 meter dash in the Bird's Nest at the Beijing Olympics 2008. In a spectacular run dubbed ``the greatest 100 meter performance in the history of the event'' by Michael Johnson, Bolt finished at 9.69 seconds, improving his own previous world record from earlier this year by 0.03 seconds. However, the most impressive fact about this run was the way in which he did it: After accelerating away from the rest of the field, he looked to his sides when two seconds and 20 meters remained, and when that noting he was completely alone, he started celebrating! He extended his arms, and appeared to almost dance along the track.  \\begin{figure*} \\begin{center} \\mbox{\\epsfig{file=screenshot2.eps,width=0.8\\linewidth,clip=}} \\end{center} \\caption{Example screen shot used to estimate the runners' position as a function of time.} \\label{fig:screenshot} \\end{figure*}  Despite this, he broke the world record by 0.03 seconds. But, needless to say, this celebration left spectators and commentators all over the world wondering about one big question: What would the world record have been if he had \\emph{not} celebrated the last 20 meters? Bolt's coach, Glen Mills, recently suggested at a press conference of the Golden League tournament in Z{\\\"u}rich, that the record could have been 9.52 seconds, or even better.  We wanted to check this for ourselves, by attempting to measure Bolt's position as a function of time, and extrapolate from the dynamics before the celebration began, into the last two seconds of the race. Based on (hopefully) reasonable assumptions, we could then obtain an estimate of the new world record.  In this paper we analyze footage of the run obtained from various web sites and the Norwegian Broadcasting Corporation (NRK), with the goal of estimating this ``hypothetical'' world record. The main technical difficulty in performing this analysis lies in obtaining accurate distance measurements as a function of time for each runner. Fortunately, this task is made considerably easier by the presence of a moving camera mounted to a rail along the track. This rail is bolted to the ground at regular intervals, and thereby provides the required standard ruler. Using the methods detailed in the following sections, and properly taking into account all major sources of statistical uncertainty, we believe that our measurements are sufficiently accurate and robust to support interesting conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Glen Mills, Usain Bolt's coach, suggested that the world record could have been 9.52 seconds if Bolt had not danced along the track in Beijing for the last 20 meters. According to our calculations, that seems like an good, but perhaps slightly optimistic, estimate: Depending on assumptions about Bolt's acceleration at the end of the race, we find that his time would have been somewhere between 9.55 and 9.61 seconds, with a 95\\% statistical error of $\\pm0.04$ seconds. Clearly, the uncertainties due to the assumptions about the acceleration are comparable to or larger than the statistical uncertainties. Therefore, 9.52 seconds does by no means seem to be out of reach.  In Figure \\ref{fig:manipulated} we show an illustration of how such a record would compare to the actual world record of 9.69 seconds, relative to the rest of the field: The left version of Bolt shows his actual position at $\\sim9.5$ seconds, while the right version indicates his position in the new scenarios.  Of course, there are potential several systematics errors involved in these calculations. For instance, it is impossible to know for sure whether Usain might have been tired at the end, which of course would increase the world record beyond our estimates. On the other hand, judging from his facial expressions as he crossed the finishing line, this doesn't immediately strike us as a very plausible hypothesis.  \\begin{figure} \\mbox{\\epsfig{file=bolt_manip.eps,width=\\linewidth,clip=}} \\caption{Photo montage showing Bolt's position relative to his competitors for real (left Bolt) and projected (right Bolt) world records.} \\label{fig:manipulated} \\end{figure}  Another issue to consider is the wind. It is generally agreed that a tail wind speed of 1 m/s improves a 100 meter time by 0.05 seconds \\citep{mureika:2000}. Further, for IAAF (International Association of Athletics Federations) to acknowledge a given run as a record attempt, the wind speed must be less than +2 m/s. When Bolt ran in Beijing, there was no measurable wind speed at all, and one can therefore safely assume that the world record could have been further decreased, perhaps by as much as 0.1 seconds, under more favorable wind conditions.  A corollary of this study is that a new world record of less than 9.5 seconds is within reach for Usain Bolt in the near future."
    },
    "0809/0809.1619_arXiv.txt": {
        "abstract": "Astronomical mid-IR spectra show two minor PAH features at 5.25 and 5.7 $\\mu$m (1905 and 1754 cm$^{\\rm - 1}$) that hitherto have been little studied, but contain information about the astronomical PAH population that complements that of the major emission bands. Here we report a study involving both laboratory and theoretical analysis of the fundamentals of PAH spectroscopy that produce features in this region and use these to analyze the astronomical spectra. The ISO SWS spectra of fifteen objects showing these PAH features were considered for this study, however only four (HD 44179; NGC 7027; Orion Bar, 2 positions) have sufficient signal-to-noise between 5 and 6 $\\mu$m to allow for an in-depth analysis. All four astronomical spectra show similar peak positions and profiles. The 5.25 $\\mu$m feature is peaked and asymmetric, with a FWHM of about 0.12 $\\pm$ 0.01 $\\mu$m ($\\sim$40 $\\pm$ 6.5 cm$^{\\rm -1}$), while the 5.7 $\\mu$m feature is broader and flatter, with a FWHM of about 0.17 $\\pm$ 0.02 $\\mu$m (50 $\\pm$ 5.6 cm$^{\\rm -1}$). Detailed analysis of the laboratory spectra and quantum chemical calculations show that the astronomical 5.25 and 5.7 $\\mu$m bands are a blend of combination, difference and overtone bands primarily involving CH stretching and CH in-plane and CH out-of-plane bending fundamental vibrations. The experimental and computational spectra show that, of all the hydrogen adjacency classes possible on PAHs, solo and duo hydrogens consistently produce prominent bands at the observed positions whereas quartet hydrogens do not. In all, this study supports the picture that astronomical PAHs are large with compact, regular structures. From the coupling with primarily strong CH out-of-plane bending modes one might surmise that the 5.25 and 5.7 $\\mu$m bands track the neutral PAH population. However, theory suggests the role of charge in these astronomical bands might also be important. ",
        "introduction": "\\label{sec:introduction}  Some thirty years of observations, combined with computational and laboratory studies, have shown that the mid-IR astronomical emission features, formerly referred to as the Unidentified Infrared (UIR) bands, are produced by mixtures of highly vibrationally excited Polycyclic Aromatic Hydrocarbons (PAHs) and closely related species. Detected in many Galactic and extra-galactic objects, including several with significant redshift \\citep[e.g.][]{2005ApJ...628..604Y}, the astronomical infrared emission features present an important and unique probe of astrochemical and astrophysical conditions across the universe. Recent reviews and papers of the observational and laboratory work \\citep[e.g.][]{2004ApJ...617L..65P, 2004ASPC..309..665H, 2004ARA&A..42..119V, 2007ApJ...659.1338S, 2007ApJ...656..770S, 2008ARA&A..46..289T} and work on theoretical models \\citep[e.g.][]{2001A&A...372..981V, 2001ApJ...556..501B, 2001ApJ...554..778L, 2002A&A...388..639P, 2006A&A...460..519R, 2007ApJ...657..810D} can be found elsewhere.  The major features at 3.3, 6.2, `7.7', 8.6 and the complex of bands between 11 and 20 $\\mu$m have been studied in great detail \\citep[e.g.][and references therein]{2000A&A...357.1013V, 2001A&A...370.1030H, 2002A&A...390.1089P, 2004ApJ...611..928V}, and the fundamental spectroscopic information is now available with which one can analyze the strongest astronomical features. However, there are several components of the astronomical PAH emission spectra that have been widely overlooked. Many of these contain valuable, sometimes subtle, information which is equally important to that revealed by the more well-known features. This study focuses on just such features, namely the weak bands that fall between 5 and 6 $\\mu$m (2000 and 1667 cm$^{\\rm -1}$).  One of the early predictions of the PAH hypothesis was the expectation of a weak emission feature near 5.25 $\\mu$m in all objects showing the major PAH bands. Its detection in 1989 (\\citeauthor{1989ApJ...345L..59A}) was an early confirmation of the PAH hypothesis. Since that time, although evident in many spectra showing the major PAH features, little has been published on this feature and its companion near 5.7 $\\mu$m.  The weak 5.25 and 5.7 $\\mu$m PAH features do not correspond to fundamental vibrational frequencies ($\\nu_{i}, \\nu_{j}, \\cdots$), but are produced by overtones ($n\\times\\nu_{i}$), combinations ($\\nu_{i} + \\nu_{j}$), and difference ($\\nu_{i} - \\nu_{j}$)bands of these fundamental vibrations \\citep{1989ApJ...345L..59A}. For example, the strong CH out-of-plane (CH$_{\\rm oop}$) fundamental bending vibration ($\\nu_{\\rm oop}$) for solo hydrogens produces the well known band at 11.2 $\\mu$m (893 cm$^{\\rm -1}$).  The overtone of this vibration, $2\\times\\nu_{\\rm oop}$, is expected to produce a much weaker feature near 5.6 $\\mu$m (1786 cm$^{\\rm -1}$) and is seen to contribute to the blue side of the broad 5.7 $\\mu$m interstellar feature.  Likewise, $\\nu_{\\rm oop}$ could combine with a CC stretching vibration ($\\nu_{\\rm oop} \\pm \\nu_{\\rm CC}$), resulting in other weak features.  Here we present high quality ISO SWS \\citep{1996A&A...315L..49D} spectra from four astronomical sources that show these features. The 5.25 and 5.7 $\\mu$m bands are analyzedin terms of overtone, combination and difference frequencies using experimental and theoretical PAH spectra. This manuscript is structured as follows. The astronomical observations are presented in Sect. \\ref{sec:data}, PAH spectroscopy is described in Sect. \\ref{sec:spectroscopy}, astrophysical implications are drawn in Sect. \\ref{sec:astrophysical} and a summary with conclusions is given in Sect. \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  This paper presents a study of the two minor PAH features in the 5 - 6 $\\mu$m (2000 - 1667 cm$^{\\rm -1}$) region, centered at positions near 5.25 (1905 cm$^{\\rm -1}$) and 5.7 $\\mu$m (1739 cm$^{\\rm -1}$). These contain information about the interstellar PAH population and conditions in the emission regions that both complement and extend the information revealed by the major bands.  Fifteen high quality ISO SWS spectra have been investigated for emission in the this wavelength region, with four spectra having sufficient signal-to-noise to allow for an in-depth analysis. Combined with a spectral database comprised of laboratory studies and dedicated quantum-chemical calculations, these spectra allow us to probe the main characteristics of the carriers of the astronomical PAH features. After continuum and emission line removal, all four astronomical spectra show similar, almost universal profiles. However, the signal-to-noise level can be improved and there are hints of subtle, but interesting, variations.  The absence of bands between 5 - 6 $\\mu$m in laboratory spectra of deuterated polycyclic aromatic molecules, as well as the absence of fundamentals in the quantum-chemical calculations in this region along with the strong correlation between the 5.25 and 5.75 $\\mu$m band strength with the 11.2 $\\mu$m band strength in the astronomical spectra, substantiates the involvement of CH$_{\\rm ip}$ and CH$_{\\rm oop}$ bending vibrations. In-depth analysis of the laboratory spectra and quantum-chemical calculations show that the astronomical 5.25 and 5.75 $\\mu$m bands are a blend of combination, difference, and overtone bands, involving CH$_{\\rm ip}$ and CH$_{\\rm oop}$ bending and stretching modes and, likely for the larger ionized PAHs, CC$_{\\rm oop}$ modes. When it becomes possible to compute the intensities of overtone and combination bands, this work should be extended.  Turning to the hydrogen adjacency classes, PAHs with solo and duo hydrogens consistently produce prominent bands in the appropriate wavelength regions, whereas PAHs with higher adjacency hydrogens show far richer spectra. These produce bands in-between and beyond the 5.25 and 5.7 $\\mu$m bands, ruling such species out as important members of the emitting population. The 5.7 $\\mu$m feature in itself - through its profile - contains adjacency class information that might be more easily accessible than through the 10 - 15 $\\mu$m CH$_{\\rm oop}$ region, where it is difficult to separate the duo and trio hydrogen modes.  The data suggest that the emitting astronomical PAHs are mostly large (50 $\\simeq N_{\\rm C}\\simeq$ 100 ), compact, and not fully deuterated. Furthermore, both the quantum-chemical calculations and the absence of a correlation between the 5.25/6.2 and 5.7/6.2 $\\mu$m band strengths with the 11.2/6.2 $\\mu$m band strength ratio suggest that the 5.25 and 5.7 $\\mu$m PAH band do not trace ionization and are carried predominately by neutrals. This point is reinforced by the lack of connection with class A and class B PAH band behavior. This suggests that high quality spectra from 5 - 10 $\\mu$m provide insight into the neutral as well as the cation and anion members of the emitting astronomical PAH family. Even so, a bigger collection of spectra of large compact PAHs, with mixtures of solo, duo and trio hydrogens, is still needed to firm up these conclusions."
    },
    "0809/0809.1333_arXiv.txt": {
        "abstract": "{A longstanding problem in astrochemistry is how molecules can be maintained in the gas phase in dense inter- and circumstellar regions at temperatures well below their thermal desorption values. Photodesorption is a non-thermal desorption mechanism, which may explain the small amounts of observed cold gas in cloud cores and disk mid-planes. } {This study aims to determine the UV photodesorption yields and to constrain the photodesorption mechanisms of three astrochemically relevant ices: CO, N$_2$ and CO$_2$. In addition, the possibility of co-desorption in mixed and layered CO:N$_2$ ices is explored.} {The UV photodesorption of ices is studied experimentally under ultra high vacuum conditions and at astrochemically relevant temperatures (15 -- 60~K) using a hydrogen discharge lamp (7--10.5~eV). The ice desorption is monitored by reflection absorption infrared spectroscopy of the ice and simultaneous mass spectrometry of the desorbed molecules.} {Both the UV photodesorption yield per incident photon and the photodesorption mechanism are highly molecule specific. The CO photodesorbs without dissociation from the surface layer of the ice, and N$_2$, which lacks a dipole allowed electronic transition in the wavelength range of the lamp, has a photodesorption yield that is more than an order of magnitude lower. This yield increases significantly due to co-desorption when N$_2$ is mixed in with, or layered on top of, CO ice. CO$_2$ photodesorbs through dissociation and subsequent recombination from the top 10 layers of the ice. At low temperatures (15 -- 18~K), the derived photodesorption yields are $2.7(\\pm1.3)\\times10^{-3}$ and $<2\\times10^{-4}$ molecules photon$^{-1}$ for pure CO and N$_2$, respectively. The CO$_2$ photodesorption yield is $1.2(\\pm0.7)\\times10^{-3}\\times (1-e^{-(x/2.9(\\pm1.1) \\rm)})+1.1(\\pm0.7)\\times10^{-3}\\times(1- e^{-(x/4.6(\\pm2.2)} \\rm))$ molecules photon$^{-1}$, where $x$ is the ice thickness in monolayers and the two parts of the expression represent a CO$_2$ and a CO photodesorption pathway, respectively. At higher temperatures, the CO ice photodesorption yield decreases, while that of CO$_2$ increases.} {} ",
        "introduction": "In dark clouds molecules and atoms collide with and stick to cold submicron-sized dust particles, resulting in icy mantles \\citep{Leger85, Boogert04}. The ices are subsequently processed by atom or light interactions to form more complex species \\citep{Tielens82, Watanabe03,Ioppolo08}. Observations show that H$_2$O, CO and CO$_2$ are the main ice constituents, with abundances up to $10^{-4}$ with respect to the total hydrogen density. These molecules are key constituents in the formation of more complex species \\citep{Tielens97}, and their partitioning between the grain and gas phase therefore strongly affects the chemical evolution in star- and planet-forming regions \\citep{Vandishoeck06b}.  Whether formed on the grains or frozen out from the gas phase, chemical models of cloud cores show that all molecules except for H$_2$ are removed from the gas phase within $\\sim 10^9 / n_{\\rm H}$ years, where $n_{\\rm H}$ is the total hydrogen number density \\citep{Willacy98}. For a typical cloud core density of $10^4$ cm$^{-3}$, this time scale is much shorter than the estimated age of such regions and thus molecules like CO and CO$_2$ should be completely frozen out. Yet gas-phase molecules, like CO, are detected in these clouds \\citep{Bergin01, Bergin02}. Cold CO gas is also detected in the midplanes of protoplanetary disks \\citep{Dartois03, Pietu07} where the densities are higher and the freeze-out time scales are even shorter, suggesting the existence of either efficient non-thermal desorption or an efficient mixing process in the disks. Similarly \\citet{Sakai08} have detected cold HCO$_2^+$, tracing gas phase CO$_2$, toward the embedded low-mass protostar IRAS 04368+2557 in L 1527 also referred to as L 1527 IRS. From the high column density and the thin line profile they conclude that the observed CO$_2$ cannot originate from thermal evaporation of ices in the hot inner regions of the envelope. They instead suggest gas phase formation of CO$_2$ to explain their observations, but do not consider non-thermal desorption in the cold envelope as an alternative. HCO$_2^+$ is also detected by \\citet{Turner99} toward several small translucent molecular clouds. They conclude that the observed HCO$_2^+$ can only form through gas phase chemistry for very specific C/O ratios and time spans and that the source of gas phase HCO$_2^+$ may instead be desorbed CO$_2$ ice. Both the CO and CO$_2$ observations may thus be explained by non-thermal desorption of ices, but this has not been quantified to date.  In dense clouds and in outer disks and disk midplanes, desorption must occur non-thermally since the grain temperature is low enough, around 10~K, that thermal desorption is negligible. Suggested non-thermal desorption pathways include photon and cosmic ray induced processes and desorption following the release of chemical energy \\citep{Shen04, Roberts07}.The importance of these processes depend both on the intrinsic desorption yields and on the local environment, especially the UV and cosmic ray fluxes. External UV photons from the interstellar radiation field can penetrate into the outer regions of dense clouds and disks and this UV field may be enhanced by orders of magnitude in disks through irradiation by the young star. In addition to direct interaction with ices, cosmic rays and X-rays also produce a UV field inside of the clouds through interaction with H$_2$.  UV photodesorption is therefore possible in most dense astrophysical environments, but it has been proposed as an important desorption pathway of ices mainly in protoplanetary disks and other astrophysical regions with dense clumps of material and excess UV photons \\citep{Willacy00,Dominik05}. There is however a lack of experimentally determined photodesorption yields for most astrophysically relevant molecules. This has prevented progress in the field and in most models UV photodesorption is simply neglected. Recently we showed that CO photodesorption is an efficient process with a yield of $3(\\pm1)\\times10^{-3}$ photon$^{-1}$ \\citep{Oberg07b}. This is of the same order as H$_2$O photodesorption, investigated by \\citet{Westley95a,Westley95b}, though the dependence of the H$_2$O yield on different parameters remains unclear. The photodesorption of H$_2$O and benzene in a H$_2$O dominated ice has also been investigated by \\citet{Thrower08} who only find substrate and matrix mediated desorption processes.  In this study we determine the photodesorption yield of CO$_2$ and its dependence on ice thickness, temperature, morphology, UV flux and integrated UV flux or fluence as well as UV irradiation time. In addition, we extend the previously reported investigation of CO and N$_2$ photodesorption to include different temperatures and ice morphologies. From the deduced yield dependencies we constrain the different desorption mechanisms and discuss the astrophysical implications. ",
        "conclusions": "\\begin{enumerate} \\item The CO photodesorption yield is temperature dependent between 15 and 27 K, which is described empirically by $2.7\\times10^{-3}-(T-15)\\times1.7\\times10^{-4}$ molecules photon$^{-1}$. The anti-correlation between yield and temperature is probably due to ice re-structuring into a more compact configuration -- the observed linearity may be coincidental, however. For most astrophysical applications the yield measured at 15 K is appropriate to use. \\item The CO photodesorption is initially reduced by more than 80\\% when the CO ice is covered by 1 ML of N$_2$ ice and decreases with UV fluence when mixed with N$_2$ due to surface build-up of N$_2$ ice, confirming that CO only desorbs from the ice surface. \\item N$_2$ co-desorbs with CO at 16 K in an ice mixture and in a layered ice with a yield of $3\\times10^{-4}$ molecules photon$^{-1}$. \\item A CO$_2$ photodesorption event starts with the photodissociation of a CO$_2$ molecule into CO and O. The fragments either desorb directly or react and recombine to form CO$_2$ and CO$_3$ before desorbing. The CO$_3$ yield is however less than 1\\% and the two main desorption products are CO and CO$_2$. \\item The CO$_2$ photodesorption yield is thickness dependent at all temperatures between 18 and 60 K. At 18--30 K the yield is well described by $1.2\\times10^{-3}\\times(1-e^{-x/2.9})+1.1\\times10^{-3}\\times(1-e^{-x/4.6})$, and at 40--60 K by $2.2\\times10^{-3}\\times(1-e^{-x/5.8})+0.22\\times10^{-3}\\times x$ molecules photon$^{-1}$, where $x$ is the ice thickness in monolayers. The first part in each yield equation is due to desorbing CO$_2$ molecules and the second part to desorbing CO molecules. \\item The thickness dependence of CO$_2$ photodesorption is understood from a mean-free-path perspective, where the different excited fragments travel a different average distance through the ice before being stopped. At higher temperatures, this mean free path increases due to increased mobility of molecules in the ice. \\item A simple model of an envelope using the observed CO$_2$ abundance in L 1527 IRS shows that CO$_2$ photodesorption can maintain CO$_2$ fractional abundances up to $1\\times10^{-6}$ in the gas phase at A$_V\\sim10$ mag after moderate grain growth and $2-3\\times10^{-8}$ using small ISM grains. At lower extinctions the photodesorption is higher due to the external irradiation field and a high fraction of the total CO$_2$ ice abundance is maintained in the gas phase. \\end{enumerate}"
    },
    "0809/0809.3931_arXiv.txt": {
        "abstract": "\\noindent We present observations of continuum ($\\lambda$ = 0.7, 1.3, 3.6 and 18 cm) and OH maser ($\\lambda$ = 18 cm) emission toward the young planetary nebula IRAS~17347$-$3139, which is one of the three planetary nebulae that are known to harbor water maser emission. From the continuum observations we show that the ionized shell of IRAS~17347$-$3139 consists of two main structures: one extended (size $\\sim$1$\\rlap{.}^{\\prime\\prime}$5) with bipolar morphology along PA=$-$30$^{\\circ}$, elongated in the same direction as the lobes observed in the near-infrared images, and a central compact structure (size $\\sim$0$\\rlap{.}^{\\prime\\prime}$25) elongated in the direction perpendicular to the bipolar axis, coinciding with the equatorial dark lane observed in the near-infrared images. Our image at 1.3 cm suggests the presence of dense walls in the ionized bipolar lobes. We estimate for the central compact structure a value of the electron density at least $\\sim$5 times higher than in the lobes. A high resolution image of this structure at 0.7 cm shows two peaks separated by about 0$\\rlap{.}^{\\prime\\prime}$13 (corresponding to 100-780 AU, using a distance range of 0.8$-$6 kpc). This emission is interpreted as originating in an ionized equatorial torus-like structure, from whose edges the water maser emission might be arising. We have detected weak OH~1612~MHz maser emission at $V_{\\rm LSR}$~$\\sim$~$-$70~km~s$^{-1}$ associated with IRAS 17347$-$3139. We derive a 3$\\sigma$ upper limit of $<$ 35\\% for the percentage of circularly polarized emission.  Within our primary beam, we detected additional OH~1612~MHz maser emission in the LSR velocity ranges $-$5 to $-24$ and $-$90 to $-$123~km~s$^{-1}$, associated with the sources 2MASS J17380406$-$3138387 and OH 356.65$-$0.15, respectively. ",
        "introduction": "The study of transition objects from the asymptotic giant branch (AGB) to the planetary nebula (PN) phase is very important to understand the processes by which low and intermediate mass stars evolve. It has been observed that planetary nebulae (PNe) display a large variety of morphologies, including bipolar or multipolar structures (Balick 1987; Schwarz, Corradi, \\& Melnick 1992; Manchado et al. 1996). However, it is not well understood how they develop such morphologies. Given that the transition phase occurs in a very short time scale of $\\sim$1000 years (Kwok 1993), only a few objects are expected to be in this evolutionary stage, making the observational study of the physical conditions under which they evolve a difficult task. A simple model which in general explains the development of bipolar morphologies in PNe is the generalized interacting stellar winds (GISW) model (Kahn \\& West 1985; Balick 1987; Icke 1988, Mellema et al. 1991). This model assumes that the ``superwind'', expelled during the AGB phase, produces a circumstellar envelope (CSE) which has a latitude-dependent density profile, with an enhancement in the equatorial region and decreasing monotonically toward the poles. Subsequently, the slow massive wind is replaced by a fast tenuous wind; the latter interacts hydrodynamically with the former, resulting in the creation of the bipolar lobes (Mellema et al. 1991, Frank et al. 1993, Garc\\'\\i a-Segura et al. 1999, Balick \\& Frank 2002).  High angular resolution and sensitive images of PNe obtained with the Hubble Space Telescope (HST) have revealed collimated structures whose formation cannot be explained by the GISW model (Miranda \\& Solf 1992; Sahai \\& Trauger 1998). The presence of a companion, collimated outflows (e.g.  Sahai \\& Trauger 1998; Soker \\& Rappaport 2000; Vel\\'azquez et al. 2007) or magnetic fields (e.g. Garc\\'\\i a-Segura et al. 1999), are required in most cases to explain the formation of such collimated structures. Nonetheless, the existence and study of a disk or an equatorial density enhancement in the CSE is considered a key ingredient to understand the processes that form bipolar lobes. A detailed study of particular PNe can provide crucial information about the physical conditions under which they develop their morphologies, and help to determine the relevance of the different shaping mechanisms proposed.   IRAS~17347$-$3139 is a young PN with a clear bipolar morphology, as revealed by the near-infrared images (de~Gregorio-Monsalvo~et~al. 2004 [hereafter dGM04], S\\'anchez-Contreras et al. 2006; Sahai et al. 2007). The lobes show an extent of $\\sim$4$^{\\prime\\prime}$, separated by a dark lane, which probably is tracing a dense dusty equatorial region. S\\'anchez-Contreras et al. (2006) suggest that the limb brightened appearance of the lobes could be indicating the presence of bubble-like structures with dense walls and tenuous interiors, presumably excavated by jet-like winds.  dGM04 detected water masers arising from this young planetary nebula. Up to now, only other two PNe are known to exhibit water maser emission (Miranda et al. 2001; G\\'omez et al. 2008). Since the water maser emission is expected to last for a very short period after the intense mass-loss rate stops, at the end of the AGB phase ($\\sim$ 100 yr, G\\'omez, Moran \\& Rodr\\'\\i guez 1990), the detection of this emission suggests that these stars have entered the PN phase only some decades ago, making this objects good candidates to study the early stages of PN formation. On the other hand, the nature of the radio continuum emission in IRAS~17347$-$3139 has been discussed by dGM04 and G\\'omez et al. 2005 (hereafter G05). These authors showed that the flux density of IRAS~17347$-$3139 rises with frequency, deriving a spectral index $\\alpha $~$\\simeq$~0.7 ($S_{\\nu}\\propto \\nu^{\\alpha}$), between 4.9 and 22 GHz, which was interpreted in terms of free-free emission from an ionized nebula. Moreover, G05 found that the radio continuum flux density is increasing rapidly with time. They estimated a dynamic time scale for the ionized envelope of $\\sim$ 100 yr, supporting the idea that this star entered the PN phase only some decades ago.  OH maser emission at 1612 MHz toward IRAS~17347$-$3139 was first reported by Zijlstra et al. (1989). Recently, Szymczak and G\\'erard (2004) presented single-dish polarimetric observations of OH masers toward this source. However, the association of this emission with the PN is uncertain due to the low angular resolution of their observations.  In order to clarify some questions originated in previous works, and to further investigate the PN IRAS~17347$-$3139, we have carried out high sensitivity and angular resolution continuum and OH maser observations with the Very Large Array (VLA). This work is structured as follows: In \\S~2 we describe the new observations that allowed us to image with higher sensitivity and higher angular resolution the ionized envelope of this source. The results are presented in \\S~3, the analysis of the data is discussed in  \\S~4, and the conclusions are given in \\S~5. ",
        "conclusions": "We have carried out sensitive high angular resolution VLA observations of the young planetary nebula IRAS 17374$-$3139. We present the first images of its ionized structure at cm wavelengths. The radio continuum images revealed the presence of a bright central structure, and an extended more tenuous bipolar component.  A double Gaussian fit shows that the extended component is elongated in the same direction as the bipolar lobes observed in the near-infrared images, while the central structure shows an elongation in the perpendicular direction, parallel to the dark lane observed in the IR images. We interpret that the radio continuum emission is arising in two extended ionized lobes, and in an equatorial ionized torus. The electron density at the base of the lobes is 2-6$\\times$10$^{5}$~cm$^{-3}$, for a distance range from 6 to 0.8 kpc. This relatively high density in the lobes supports the idea that this source is a very young PN. Given the subtle point-symmetric morphology of the extended component, we suggest the possible presence of a collimated ionized wind in this source. On the other hand, we derive a lower limit for the electron density for the equatorial torus of $n_{e} \\geq 1$-$3\\times10^6$~cm$^{-3}$, given the same distance range as above. A high resolution image at 0.7 cm reveals the presence of a double peak structure in the central component, supporting the interpretation of the equatorial torus. We compared the distribution of the water maser emission with our high resolution radio continuum images; the relative positions of the maser and the continuum emission suggest that the water masers arise from the outer parts of the ionized torus.  We detected OH maser emission at 1612~MHz toward IRAS~17347$-$3139. The spectrum shows only one weak feature at $V_{\\rm LSR}=$~$-$70~km~s$^{-1}$ which coincides spatially with the continuum emission. We derived a 3-$\\sigma$ upper limit of $<$ 35\\% for the percentage of circularly polarized emission (m$_{c}$=V/I). We also report the detection of OH 1612 MHz maser emission coming from two other sources, J17380406-3138387 and OH 356.65-015, located within our primary beam."
    },
    "0809/0809.1105_arXiv.txt": {
        "abstract": "We present a catalog of 9017 X-ray sources identified in \\chandra\\ observations of a 2$\\times$0.8\\degree\\ field around the Galactic center. This enlarges the number of known X-ray sources in the region by a factor of 2.5. The catalog incorporates all of the ACIS-I observations as of 2007 August, which total 2.25 Msec of exposure. At the distance to the Galactic center (8~kpc), we are sensitive to sources with luminosities of $4\\times10^{32}$~\\ergsec\\ (0.5--8.0 keV; 90\\% confidence) over an area of one square degree, and up to an order of magnitude more sensitive in the deepest exposure (1.0~Msec) around \\sgrastar. The positions of 60\\% of our sources are accurate to $<$1\\arcsec (95\\% confidence), and 20\\% have positions accurate to $<$0\\farcs5. We search for variable sources, and find that 3\\% exhibit flux variations within an observation, 10\\% exhibit variations from observation-to-observation. We also find one source, CXOUGC J174622.7--285218, with a periodic 1745~s signal (1.4\\% chance probability), which is probably a magnetically-accreting cataclysmic variable. We compare the spatial distribution of X-ray sources to a model for the stellar distribution, and find 2.8$\\sigma$ evidence for excesses in the numbers of X-ray sources in the region of recent star formation encompassed by the Arches, Quintuplet, and Galactic center star clusters. These excess sources are also seen in the luminosity distribution of the X-ray sources, which is flatter near the Arches and Quintuplet than elsewhere in the field. These excess point sources, along with a similar longitudinal asymmetry in the distribution of diffuse iron emission that has been reported by other authors, probably have their origin in the young stars that are prominent at $l$$\\approx$0.1\\degree. ",
        "introduction": "Stars are detectable as X-ray sources at several important stages of their lives. Pre-main sequence stars are X-ray sources because of their enhanced magnetic activity \\citep{pf05}. Massive OB and Wolf-Rayet stars produce X-rays through shocks in their stellar winds \\citep{ber97,gagne05}, and possibly from magnetically-confined plasma close to their stellar surfaces \\citep{wc07}. Neutron stars are bright X-ray sources if they are young and still have latent heat from what was once the stellar core \\citep{wwn96}, if they accelerate particles in rotating, moderate-strength ($B$$\\sim$$10^{12}$ G) fields \\citep{gs06}, or if they have extremely strong fields ($B$$\\sim$$10^{14}$ G) that decay and accelerate particles \\citep{wt06}. White dwarfs, neutron stars, and black holes are bright X-ray sources if they are accreting matter from a binary companion \\citep{war95,psa06}, or in principle from the interstellar medium \\citep[see, e.g.,][]{per03}. Therefore, X-ray surveys can be used to study the life cycles of stars, particularly their start and end points.  Here we present a catalog of X-ray sources detected in \\chandra\\ observations toward the inner 2\\degree\\ by 0.8\\degree\\ of the Galaxy. The region encompasses about 1\\% of the Galactic stellar mass \\citep{lzm02}, and possibly up to 10\\% of the Galactic population of young, massive stars \\citep{mp79,fig04}. Therefore, these data provide a statistically meaningful sample of the Galactic stellar population.  Previous catalogs based on \\chandra\\ data on the Galactic center have been published by \\citet{m-cat} using 630 ks of data taken through 2002 June on the central 17\\arcmin$\\times$17\\arcmin\\ around \\sgrastar, and by \\citet{m-wide} using observations taken through 2005 June on the inner 2\\degree$\\times$0.8\\degree\\ of the Galaxy. However, since the publication of these catalogs, a large amount of new data have been obtained. These data increase the number of point sources identified by a factor of 2.5. They also provide much better astrometry for individual X-ray sources. The improvement in astrometry enables the identification of rare objects such as Wolf-Rayet stars, X-ray binaries, and rotation-powered pulsars, through comparisons of our X-ray catalog with radio and infrared data sets (e.g., Mauerhan, J. \\etal, in prep). Therefore, we provide here an updated catalog of point sources, which incorporates and supercedes the previous catalogs. We also describe the spatial and luminosity distributions of the X-ray sources.  Throughout this paper, we adopt a distance to the Galactic center of $D$=8~kpc \\citep{reid93,mcn00}, and an average absorption column of $N_{\\rm H}$=$6\\times10^{22}$~cm$^{-2}$ \\citep{bag03}. ",
        "conclusions": "We have presented a catalog of 9017 X-ray sources located in the inner 2\\degree\\ by 0.8\\degree\\ around the Galactic center. This increases the number of sources known in the region by a factor of 2.5. For all of the sources, we provide tables listing their positions (Table~\\ref{tab:positions}), photometry, and colors (Table~\\ref{tab:phot}). Of these sources, 6760 have hard colors that are consistent with high absorptions columns $N_{\\rm H}$$\\ga$$4\\times10^{22}$~cm$^{-2}$, which indicates that they lie at or beyond the Galactic center. In addition, the positions of the X-ray sources in this catalog are more accurate than earlier versions. This catalog contains 2029 sources with $<$0.5\\arcsec\\ uncertainties (90\\% confidence), and another 3981 with uncertainties between 0.5\\arcsec\\ and 1\\arcsec. This catalog will be excellent for comparisons with multi-wavelength ones, in order to search for young stars, high-mass X-ray binaries, and pulsars \\citep[e.g.,][]{wll02b,lwl03,mik06,m-ys,mau07}.  The luminosity range that we cover, from $10^{31}$ to $10^{34}$~\\ergsec\\ (0.5--8.0 keV; assuming a $\\Gamma$=1.5 power law, $N_{\\rm H}$=$6\\times10^{22}$ cm$^{-2}$, and $D$=8 kpc), is at least an order of magnitude fainter than studies of Local Group galaxies \\citep[e.g.,][]{tp04,kil05,plu08}. Consequently, the natures of the sources that we study are also very different. Whereas the detectable stellar population of external galaxies in X-rays is dominated by accreting black holes and neutron stars, most of our sources are probably cataclysmic variables \\citep[e.g.,][]{m-wide}. The hardness of the X-ray colors (Fig.~\\ref{fig:hit}) suggests that the sources are specifically magnetically-accreting white dwarfs \\citep{ei99,m-wide}. Therefore, the X-ray population probably represents old stars. Indeed, the spatial distribution of sources brighter than $2\\times10^{-6}$~\\phcms\\ (2--8 keV) traces that of the old stellar population (Fig.~\\ref{fig:londist}). This makes the population of X-ray sources in the Galactic center similar to those seen in globular clusters \\citep[e.g.,][]{ver97,hei06}.  Although the distribution of the majority of the X-ray sources traces that of the old stellar population, we have found 2.8$\\sigma$ evidence for an excess of sources in two regions where young, massive stars are forming: in the inner few arcminutes around \\sgrastar, and in the region where the Arches and Quintuplet star clusters lie. The excess of sources near these young star clusters also appears in the number of sources as a function of limiting flux, in which relatively more bright X-ray sources are found near the Arches and Quintuplet (Fig.~\\ref{fig:lognlogs} and Table~\\ref{tab:lognlogs}). In total, these two regions contain a couple dozen more bright sources than our stellar mass model predicts. We suggest that these excess X-ray sources are part of the young stellar population in these region \\citep{mik06,m-ys,mau07}. In the near future, we will publish additional OB and  Wolf-Rayet stars that have been identified through infrared spectroscopy of counterparts to X-ray sources (J. Mauerhan \\etal, in prep).  \\begin{deluxetable*}{llcccl}[htp] \\tabletypesize{\\scriptsize} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Luminous X-ray Binaries Covered by Our Observations\\label{tab:transients}} \\tablehead{ \\colhead{\\chandra\\ name} & \\colhead{Common Name} & \\colhead{RA} & \\colhead{DEC} & \\colhead{uncertainty} & \\colhead{Reference} \\\\ \\colhead{(CXOUGC J)} & \\colhead{} & \\multicolumn{2}{c}{(Degrees, J2000)} & \\colhead{(arcsec)} & colhead{} } \\startdata 174354.8-294441 & 1E 1740.7-2942 & 265.97864 & $-29.74499$ &  0.5 & \\citet{sid99} \\\\ 174417.2-293943 & AX J1744.3-2940 & 266.07190 & $-29.66234$ &  0.5 & \\citet{sid01} \\\\ 174433.0-284427 & Bursting Pulsar & 266.13788 & $-28.74096$ &  0.5 & \\citet{ww02} \\\\ 174451.6-292042 & KS 1741-293 & 266.21515 & $-29.34522$ &  0.5 & \\citet{int97} \\\\ 174457.4-285021 & XMM J174457-2850.3 & 266.23944 & $-28.83917$ &  0.3 & \\citet{sak05} \\\\ [2pt] 174502.3-285449 & Granat 1741.9-2853 & 266.25983 & $-28.91397$ &  0.4 & \\citet{m-grs} \\\\ 174535.6-290133 & AX J1745.6-2901 & 266.39853 & $-29.02612$ &  0.4  & \\citet{mae96} \\\\ 174535.5-290124 & \\nodata & 266.39822 & $-29.02337$ &  0.3  & \\citet{m-trans}\\\\ 174537.1-290104 & 1A 1742-289 & 266.40494 & $-29.01796$ &  0.4 & \\citet{dav76} \\\\ 174538.0-290022 & \\nodata & 266.40863 & $-29.00623$ &  0.3  & \\citet{m-trans} \\\\  [2pt] 174540.0-290005 & \\nodata & 266.41699 & $-29.00160$ &  0.4 & \\citet{m-trans} \\\\ 174540.0-290030 & \\nodata & 266.41684 & $-29.00859$ &  0.3  & \\citet{m-trans} \\\\ 174540.9-290014 & \\nodata & 266.42078 & $-29.00398$ &  0.4  & \\citet{m-trans} \\\\ 174553.9-290346 & SWIFT J174553.9-290347 & 266.47467 & $-29.06305$ &  0.4 & \\nodata \\\\ 174554.4-285455 & XMM J174554.4-285456 & 266.47690 & $-28.91533$ &  0.4 & \\citet{por05} \\\\ [2pt] 174621.0-284342 & 1E 1743.1-2843 & 266.58768 & $-28.72868$ &  0.4 & \\citet{por03} \\\\ 174702.5-285259 & SAX J1747.0-2853 & 266.76080 & $-28.88307$ &  0.4 & \\citet{wmw02} \\\\ \\nodata & XTE J1748-288 & 267.02108 & --28.47383 & 0.6 & \\citet{hrm98} \\\\ \\nodata & XMM J174544-2913.0 & 266.43546 & --29.21683 & 4.0 & \\citet{sak05} \\\\ \\enddata \\end{deluxetable*}  A small fraction of the X-ray sources should be accreting black holes and neutron stars. Around 300 such X-ray binaries are known in the Galaxy, about half of which contain low-mass donors that over fill their Roche lobe, and half of which contain high-mass (OB and Wolf-Rayet) stars that donate mass through a stellar wind \\citep{liu06,liu07}. These X-ray binaries are most-easily identified when they are bright and variable \\citep{m-trans}. In total, over the history of X-ray astronomy, 19 X-ray sources in our survey field have been observed to be $>$$10^{34}$~\\ergsec\\ in X-rays, and have varied by at least an order of magnitude in X-ray flux (Table~\\ref{tab:transients}). Fifteen of these transient X-ray sources were bright during the time span of our \\chandra\\ observations (1A 1742-289 and XTE J1748-288 never entered outburst). Half of them have been discovered in the last 9 years using \\chandra, \\xmm\\, or {\\it Swift} \\citep[e.g.,][]{sak05,por05,m-trans,wij06,ken06}. Surprisingly, despite having obtained 600 ks of new data in 2006 and 2007, we did not detect any new, bright ($>$$10^{34}$~\\ergsec), transient X-ray sources. This suggests that we have identified all of the X-ray binaries that are active on time scales of a decade.  As mentioned in \\S\\label{ref:obs}, the tables from this work will be available in the electronic edition of this journal, and additional products will be made available from the authors' web site.\\footnote{{\\tt http://www.srl.caltech.edu/gc\\_project/xray.html}} The data available from the authors' site includes FITS images of all of the images presented in this paper, as well as the averaged event lists, snapshot images, spectra, and calibration files for each source in the catalog. Combined with an increasing amount of multi-wavelength data, this data set can be used to better understand the interactions between stars and interstellar media in the Galactic center, and the population of X-ray emitting objects in general."
    },
    "0809/0809.3100_arXiv.txt": {
        "abstract": "We present deep radio images at 1.4 GHz of a large and complete sample of BL Lacertae objects (BL Lacs) selected from the Deep X-ray Radio Blazar Survey (DXRBS). We have observed 24 northern ($\\delta \\ga -30^{\\circ}$) sources with the Very Large Array (VLA) in both its A and C configurations and 15 southern sources with the Australia Telescope Compact Array (ATCA) in its largest configuration. We find that in the DXRBS, as in the 1-Jy survey, which has a radio flux limit roughly ten times higher than the DXRBS, a considerable number (about a third) of BL Lacs can be identified with the relativistically beamed counterparts of Fanaroff-Riley type II (FR II) radio galaxies. We attribute the existence of FR II-BL Lacs, which is not accounted for by current unified schemes, to an inconsistency in our classification scheme for radio-loud active galactic nuclei (AGN). Taking the extended radio power as a suitable measure of intrinsic jet power, we find similar average values for low- (LBL) and high-energy peaked BL Lacs (HBL), contrary to the predictions of the blazar sequence. ",
        "introduction": "The basic nature of BL Lacertae objects (BL Lacs) is believed to be understood within the unified schemes for radio-loud active galactic nuclei (AGN). These sources are radio galaxies, which have their relativistic jets oriented close to the observer's line of sight. For objects with such a preferred jet alignment, which are commonly referred to as 'blazars', the relativistic beaming effect can explain most of the observed properties, such as, e.g., non-thermal continuum emission from radio up to $\\gamma$-ray frequencies, high core luminosities, core-dominated radio morphologies, irregular and rapid variability, and strong radio and optical polarization (see review by \\citet{Urry95} and references therein).  Current unified schemes identify the parent population of BL Lacs with the low-luminosity Fanaroff-Riley type I \\citep[FR I;][]{Fan74} radio galaxies, whereas the high-luminosity Fanaroff-Riley type II (FR II) radio galaxies are assumed to appear as radio quasars when their jets are viewed at relatively small angles. However, in recent years evidence has accumulated that some BL Lacs might in fact be beamed FR II radio galaxies. Their extended radio powers are higher than those of known FR Is and their extended radio morphologies are consistent with those of FR IIs once orientation effects are accounted for \\citep[][]{Kol92, Cas99, Rec01}.  If indeed a considerable number of FR IIs are part of the BL Lac parent population, it will have important implications. Estimates of relativistic beaming parameters such as, e.g., the jet bulk Lorentz factor or viewing angle, which are usually derived from a comparison of the radio luminosity function \\citep[e.g.,][]{Urry91, Urry95} or the distribution of radio core-dominance values \\citep[e.g.,][]{Kol92, Per93} of the parent population with that of the corresponding blazar class, will have to be revisited. Furthermore, since, by definition, the broad emission line region, which is assumed to trace directly the accretion disk power, is absent or only very weak in BL Lacs, the existence of a considerable number of FR II-like jet powers among this blazar class could mean that accretion disk and jet luminosities are not as closely linked as current jet formation models suggest \\citep[e.g.,][]{Bla82, Raw91, Mar03}.  A further implication will concern the so-called 'blazar sequence' \\citep{Fos98, Ghi98}. This model posits that Compton-cooling determines the frequency of the jet synchrotron emission peak, in particular that the higher the (intrinsic) jet power, the stronger the cooling, and the lower the synchrotron emission peak frequency. Therefore, finding in particular high-energy peaked BL Lacs with FR II-like extended radio powers could present a severe challenge for this model \\citep[see also][]{Pad07}.  Our current knowledge of the radio properties of BL Lacs is based mainly on deep radio observations of two complete samples selected at widely different frequencies, namely, the radio-selected 1-Jy \\citep{Sti91, Rec01} and the X-ray-selected {\\it Einstein} Medium Sensitivity Survey \\citep[EMSS;][]{Mor91, Rec00} BL Lac samples. However, these two surveys sample the extreme ends of the radio flux distribution of BL Lacs and, therefore, have presented a strongly biased view of BL Lac physics. In order to rectify this situation we have obtained deep radio images of a complete sample of BL Lacs with intermediate radio properties.  The paper is structured as follows. In Section 2 we discuss the selection of the BL Lac sample, for which the radio images have been obtained and analyzed as detailed in Section 3. In Section 4 we address the question of the parent population of BL Lacs. In Section 5 we investigate if some BL Lacs with featureless optical spectra are at high redshifts rather than strongly relativistically beamed. The radio properties of low- and high-energy peaked BL Lacs are compared in Section 6. Finally, Section 7 summarizes of our main results and presents our conclusions. For consistency with previous work we have assumed throughout this paper cosmological parameters $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0$. Energy spectral indices have been defined as $f_\\nu \\propto \\nu^{-\\alpha}$. ",
        "conclusions": "Our knowledge of the radio properties of BL Lacs is based mainly on the 1-Jy and EMSS samples. However, these surveys have presented a biased view of BL Lac physics. Therefore, we have obtained deep radio images of a complete sample of 44 BL Lacs selected from the Deep X-ray Radio Blazar Survey (DXRBS). We have observed the northern sources with the VLA in both its A and C configurations and the southern sources with the ATCA in its largest configuration. Our main results can be summarized as follows.  (i) Current unified schemes identify the parent population of BL Lacs with FR I radio galaxies, however, in recent years evidence has accumulated that some BL Lacs might in fact be relativistically beamed FR II radio galaxies. We find that also in the DXRBS, based on both the extended radio powers as well as the radio morphologies, (at least) a third of the BL Lac sample can be identified with relativistically beamed FR IIs.  (ii) We discuss an inconsistency in the current classification scheme for radio-loud AGN, which explains why FR II-BL Lacs are in fact {\\it expected} to exist. This inconsistency emerges since we separate radio galaxies and blazars into their subclasses based on different criteria, namely, radio morphology (and so radio power) and emission line strength, respectively. However, these two criteria are not equivalent.  (iii) The so-called 'blazar sequence' posits that the amount of Compton cooling of the jet particles determines the frequency of the synchrotron emission peak \\citep{Fos98, Ghi98}. In particular, it expects low-energy peaked BL Lacs (LBL) to have on average higher {\\it intrinsic} jet powers than high-energy peaked BL Lacs (HBL). We compare the extended radio powers of the DXRBS LBL (23 objects) and HBL (21 objects) and find that, contrary to the expectations of the blazar sequence, their average values are similar ($\\log L_{\\rm ext}=25.35\\pm0.23$ and $25.28\\pm0.20$, respectively).  We are currently extending our radio observation program for the DXRBS BL Lac sample to investigate smaller-scale source structure. In a future paper we will present results from VLBI."
    },
    "0809/0809.4516_arXiv.txt": {
        "abstract": "I discuss how the chemical abundance distributions, kinematics and age distributions of stars in the thin and thick disks of the Galaxy can be used to decipher the merger history of the Milky Way, a typical large galaxy. The observational evidence points to a rather quiescent past merging history, unusual in the context of the `consensus' cold-dark-matter cosmology favoured from observations of structure on  scales larger than individual galaxies. ",
        "introduction": "CDM} Dissipational collapse of a gas-rich system is an important ingredient in establishing the thin disks so prevalent today.  In the context of hierarchical clustering scenarios, such as {$\\Lambda$}CDM, gaseous mergers are required to produce disks (Zurek, Quinn \\& Salmon 1988; Robertson et al.~2006).  Centrifugally supported, extended disks also are only produced if angular momentum is largely conserved during collapse within the dark halo (Fall \\& Efstathiou 1980; Mo, Mao \\& White 1998).  However, angular-momentum transport and evolution, for example due to gravitational torques and tidal effects, are natural during the mergers inherent in {$\\Lambda$}CDM, leading to low angular momentum, very concentrated disks (Navarro, Frenk \\& White 1995).  The typical merger history (of dark haloes) is fixed by the dark matter power spectrum, so with {$\\Lambda$}CDM additional baryonic processes are introduced to implement `feedback', to both suppress early star formation and prevent dissipation, delaying disk formation until after the epoch of most active merging (cf.~Simon White's talk).  Later mergers into the disk will heat the thin disk, with minor mergers producing a thick stellar disk (some thin stellar disk component persists, e.g. Kazantzidis et al.~2007) and also driving gas into the central regions to build up a bulge (Mihos \\& Hernquist 1996).  During a merger, orbital energy is absorbed into the internal degrees of freedom of the merging systems, and orbital angular momentum is redistributed, some absorbed by the larger system, and some being lost to the system. The evolution of a satellite, and of its orbit, as it merges with, and is assimilated by, a larger system depends on the relative masses (the dynamical friction timescale on which the satellite sinks to the center scales like the mass ratio, being shorter for more massive satellites), on the relative density profiles of the satellite and host (denser satellites can survive tidal effects, leading to mass loss and disruption, closer to the center of the larger system), and on the initial orbital parameters (e.g. sense of rotation, inclination angle to the plane of the host, periGalacticon and orbital eccentricity).  The effect of the satellite(s) on a pre-existing stellar disk is also a sensitive function of the satellite's properties and orbit.  The mechanism by which a thin disk is heated by the merging process is a combination of local deposition of orbital energy from the satellite(s), local scattering (in which azimuthal streaming motions are transformed into random motions), and resonant excitation of modes in the disk, providing heating on a more global scale (e.g.~Quinn \\& Goodman 1986; T\\'oth \\& Ostriker 1992; Sellwood, Nelson \\& Tremaine 1998).   Early simulations focussed on the impact of one minor merger (e.g.~Quinn, Hernquist \\& Fullagar 1993; Velazquez \\& White 1999). These produced a plausible thick disk, similar in structure to that of the Milky Way (Gilmore \\& Reid 1982; see also Gilmore \\& Wyse 1985 for kinematics and an order-of-magnitude estimation of the mass of satellite needed), from a merger of a robust (dense) satellite with a mass ratio to the stellar {\\it disk\\/} (not to the total mass) of 10--20\\%, for a range of initial orbital parameters.  Simulations of the merging of cosmologically motivated ensembles of satellites are more relevant, and also find significant heating of a pre-existing stellar disk, over an extended period of time (e.g.~Hayashi \\& Chiba 2006; Kazantzidis et al.~2007, also this volume).  The satellites in these later simulations have a orbital distribution similar to that of subhaloes identified in dissipationless $\\Lambda$CDM cosmological simulations, with a mean ratio of initial apocenter to pericenter distance of around 6:1 (e.g. Ghigna et al.~2000; Diemand, K\\\"uhlen \\& Madau 2007). The distribution of pericenter distances is also important, since satellites with pericenters that are significantly beyond the disk (larger than say 10 disk scale lengths, see Fig.~2 of Hayashi \\& Chiba 2006) do not couple effectively to the disk. In addition to realistic orbits, the internal density profiles of the satellites are critical, since fluffy satellites are disrupted early and provide little heating (Huang \\& Carlberg 1997).  High-resolution, N-body simulations of the formation of the dark halo of a `Milky Way galaxy' with the CDM power spectrum have demonstrated that a significant population of substructure is indeed dense enough to persist and survive many pericenter passages.  The shapes of the mass- and velocity-functions of subsystems are reasonably independent of redshift and at $z=0$ are well-established as power laws (see convergence tests in Reed et al.~2005; S.~White's talk at this conference). The amplitudes depend on resolution, with the number of satellite dark haloes still increasing with increased resolution (S.~White, these proceedings); `overmerging', particularly in the central regions of the larger host galaxy halo, can artificially reduce substructure.  Indeed, the population of subhaloes within the analogue of the solar neighborhood is not yet well-established, even in pure dark-matter simulations (the addition of baryons will increase central densities). However, the present generation of (baryon-free) simulations imply that robust satellites penetrate the region of the disk (e.g.~Diemand, K\\\"uhlen \\& Madau 2007).  Simulations which model gas physics and star formation also find that thick (and thin) disks are produced. For example, the bulk of the younger stars (ages less than 8~Gyr) in the thick disk in the simulation of Abadi et al.~(2003) are produced by heating of the pre-existing stellar disk (we return to the older stars in section~2 below).    The most massive satellite provides the greatest heating, with the scaling such that the increase in scaleheight (or, equivalently, in the square of the vertical velocity dispersion) is proportional to the square of the mass of the satellite (Hayashi \\& Chiba 2006). For an ensemble of satellites with the differential mass function seen in CDM simulations, namely a power-law with slope close to $-2$ (e.g.~Diemand, K\\\"uhlen \\& Madau 2007), the cumulative heating is also dominated by the most massive systems (see also White 2000). The substructure distribution at earlier times contains more satellites of higher mass, since these are more affected by dynamical friction, which brings them into the central regions where they are more effectively stripped of mass from their outer parts (e.g.~Zentner et al.~2005).  These massive satellites, after this shriking of their orbits, can be more effective at heating the thin disk, prior to their demise.  Thus at early times the satellite distribution is more concentrated than is the host dark matter halo, while at the present day it has evolved to be less concentrated.  One must allow for this evolution of the mass function and orbital parameters, rather than simply adopting a surviving satellite retinue from the end-point of a simulation, which minimizes the predicted overall heating (as found by e.g.~Font et al.~2001).  Indeed following the full merging history is preferable. Such analyses  imply that late (after redshifts of unity) heating of thin stellar disks seems to be inevitable in $\\Lambda$CDM (e.g.~Abadi et al.~2003; Stewart et al.~2008; Kazantzidis, this volume).  The current models (Stewart et al.~2008) show that over the last 10Gyr, fully 95\\% of galaxy haloes of present total mass $10^{12}$~M$_\\odot$ have accreted a system of mass equal to that of the present-disk ($5 \\times 10^{10}$M$_\\odot$ -- their simulation resolution limit is $10^{10}$~M$_\\odot$); the mass ratio of the substructure to that of the stellar disk at the epoch of accretion is the more important ratio for disk heating and such a large mass is highly likely to cause severe heating. More numerous, lower-mass mergers are also expected, continuing to late epochs.      Gas physics can also play a role in the formation of the thick disk, in terms of slow settling to the disk plane (Burkert, Truran \\& Hensler 1992) or a starburst in a rapidly changing potential such as a gas-rich merger (Robertson et al.~2006; Brook et al.~2007). The latter mechanism has some observational support at high redshift (see Elmegreen's and Genzler's contributions to this volume).  One must of course also take account of adiabatic compression and heating of an existing thick disk by subsequent slow accretion of gas to buildup the thin disk (cf.~T\\'oth \\& Ostriker 1992; Elmegreen \\& Elmegreen 2006). Here I will focus on the -- apparently inevitable -- late heating of thin stellar disks, as a probe of $\\Lambda$CDM.  The important issues are the predicted chemical abundance and age distributions of stars in the thin and thick disks, given a typical merger history, and how they compare with the observations. ",
        "conclusions": ""
    },
    "0809/0809.1275_arXiv.txt": {
        "abstract": "The Doppler technique measures the reflex radial motion of a star induced by the presence of companions and is the most successful method to detect exoplanets. If several planets are present, their signals will appear combined in the radial motion of the star, leading to potential misinterpretations of the data. Specifically, two planets in 2:1 resonant orbits can mimic the signal of a single planet in an eccentric orbit. We quantify the implications of this statistical degeneracy for a representative sample of the reported single exoplanets with available datasets, finding that 1) around $35\\%$ percent of the published eccentric one-planet solutions are statistically indistinguishible from planetary systems in 2:1 orbital resonance, 2) another $40\\%$ cannot be statistically distinguished from a circular orbital solution and 3) planets with masses comparable to Earth could be hidden in known orbital solutions of eccentric super-Earths and Neptune mass planets. ",
        "introduction": "\\label{sec:math}  The solution degeneracy between a single planet eccentric orbit and two planets in circular resonant orbits is a direct consequence of the well known Fourier expansion of the Kepler equation into powers of the eccentricity \\citep[see][as an example]{moulton:1914}. In \\cite{konacki:1996}, the method of frequency analysis was first applied to an extrasolar planetary system and \\cite{konacki:1999} adapted it to Doppler measurements. The potential confusion between eccentric orbits and resonant systems has been briefly mentioned in \\citep{marcy:2001} and \\cite{ford:2006}, but this issue has not been specifically considered until now in a broad statistical sense.  Mathematically, the degeneracy between the resonant and the eccentric solutions comes from the fact that their equations of Keplerian trajectories are identical up to first order in the eccentricity. Detailed analytical expressions in terms of the Bessel functions can be found elsewhere \\citep[eg.][]{konacki:1999}. The relevant terms up to the 7th power in the eccentricity can be found in \\citet{lucy:2005}. Here we only discuss the first order term, which is the one relevant to the degeneracy under discussion  In the case of a single eccentric planet, the reflex radial velocity motion of the star is \\begin{eqnarray} v^e_r = v_{r0} &+& K \\cos \\left[ W\\left(t-\\tau_0\\right) \\right] \\, \\label{eq:velecc} \\\\ &+& Ke\\cos \\left[ 2 W\\left(t-\\tau_0\\right) - \\omega\\right] \\nonumber \\\\ &+& O({Ke^2})\\,, \\nonumber \\end{eqnarray} \\noindent where $v_{r0}$ is the linear radial velocity of the barycenter of the system, $K$ is the semi-amplitude of the radial velocity variations induced by the planet on the star, $\\omega$ is the argument of the periastron (angle between the periastron of the orbit and the ascending node), $\\tau_0$ is the time of crossing of the ascending node, and $W=2\\pi/P$ is the orbital frequency, where $P$ is the orbital period (see Fig.~\\ref{fig:dibu}a). The term proportional to $Ke$ is called the \\textit{first eccentric harmonic}, while the term $O({Ke^2})$ contains all the higher order contributions.  If instead we have a two--planet system, both in circular orbits and the inner planet having an orbital period half of the outer one, i.e. $W_2=2W$ (see Fig.~\\ref{fig:dibu}b), the expression for the radial velocity of the star is \\begin{eqnarray} v^R_r = v_{r0} &+& k_1 \\cos \\left[ W\\left(t-\\tau_0\\right) \\right] \\label{eq:velres}\\\\ &+& k_2 \\cos \\left[ 2 W\\left(t-\\tau_0\\right) + \\phi_0\\right]\\nonumber\\, \\\\ &+& O({k_1e_1,k_2e_2, Ke^2}),\\nonumber \\end{eqnarray} \\noindent where $k_1$ and $k_2$ are the radial velocity semi--amplitudes of the outer and the inner planet. $W$ and $\\tau_0$ are the orbital frequency and the time of crossing of the ascending node of the outer planet, and the angle $\\phi_0$ is the relative phase between the two planets at $\\tau_0$. Higher order terms, summarized here as $O({k_1e_1,k_2e_2, Ke^2})$, become significant if the orbits are allowed to be eccentric.  To a first order approximation, $v^e_r$ and $v^R_r$ are formally identical if $k_1 = K$, $k_2 = Ke$, and $\\phi_0 = -\\omega$. This implies that the signal $k_2$ of an inner lower-mass planet will be indistinguishable from the \\textit{first eccentric harmonic} $Ke$ unless the observations are precise enough to resolve the second order term in the harmonic expansion. The amplitude of that second order term is $9/8\\,Ke^2\\sim Ke^2$. \\citep[see Appendix A in][]{lucy:2005}.  The similarity of both solutions is illustrated in Figure \\ref{fig:resvsecc}, which shows how the Doppler radial velocity curves would look like in each case (one-planet in an eccentric orbit versus two resonant planets in circular orbits), for different values of $\\omega$ and $e$.  The two configurations can be easily confused, especially when e $<$ 0.3 in the single planet case (or equivalently, when the inner planet is significantly less massive than the outer planet,i.e. $k_2 << k_1$). Confusion is also possible for larger values of $e$ if the uncertainties are large and the radial velocity curves are sparsely sampled, which is the case for several  published Doppler velocity curves.  As an example estimate, if the detected semi-amplitude and eccentricity are $K\\sim 100~ms^{-1}$ and $e\\sim0.1$, the amplitude of the second harmonic will be $Ke^2\\sim 1~ms^{-1}$ and both orbital solutions are indistinguishable at the 3--$\\sigma$ level unless the precision of the data is better than $0.3~ms^{-1}$. This is a problem, since only recently have planet hunting groups started to achieve that level of precision \\citep{mayor:2008,fischer:2008}. All these statements will be made more precise in Section \\ref{sec:analysis}. There is also the accuracy limitation imposed by the stellar jitter\\footnote{intrinsic noise associated to the stellar activity}, which has typical amplitudes of 3--5 $ms^{-1}$ \\citep{cumming:2008}. The optimal strategies and the limitations of the Doppler technique to disentangle this degeneracy are discussed in Sec.~\\ref{sec:breaking}. ",
        "conclusions": "We show that the Doppler signal of a single eccentric planet can mimic the signal of a two-planet system in a 2:1 circular or near--circular resonant orbit. This degeneracy is affecting a large fraction of the known exoplanets, this is, around $30-40\\%$ of the published single planet sytems. We also find strong evidence of at least one case (HD 125612) where the resonant solutions is significantly better than the published eccentric one by \\citet{fischer:2007}. The analysis described in this paper can also be applied to multi-planet systems with eccentric candidates, where the degree of degeneracy is expected to be similar or even stronger due to the mixed signals of the different planets involved. The detailed analysis of multiplanetary systems is more complex and usually involves dynamical stability consideration. Therefore, a case by case study is imposed. A remarkable example is the planetary system around GJ 581 \\citep{mayor:2009} were $2-3$ earth mass planet could be hidden in the habitable zone of the system.  The only techniques currently able to distinguish between two resonant planets and single eccentric planet systems are limited to the Doppler and photometric approaches described above. In the future, other methods such as astrometry and direct imaging, will also provide ways to uncover \\textit{eccentric imposters} \\citep[see][]{moorhead:2009}. Astrometry will be the first one to become sensitive enough, once the upcoming space astrometric missions Gaia/ESA \\citep[][to be launched in 2011]{lindegren:2008} and SIM/NASA \\citep[][to be launched after 2015]{unwin:2008}, go on-line. High precision astrometry will be most helpful if the resonant orbits are not coplanar. Otherwise, it will suffer from the same degeneracies as the Doppler technique \\cite{konacki:2002}. The ultimate test will be direct imaging, which will make possible to measure whether the orbit of the detected planet is indeed eccentric. This will have to wait until spaceborne missions such as a Darwin/TPF launch.  The conclusions of this work make it worth reconsidering some published orbital solutions and motivate the follow-up of some interesting systems. We find that future announcements of eccentric planets should be carefully tested before publication since it is relatively simple to check if there is any improvement using a resonant configuration. Also, we find that several reported radial velocity curves \\citep{mayor:2008, mayor:2009} may contain hidden signals of rocky planets.  \\vspace{0.2in}  \\textit{Acknowledgements.} GAE thanks A. Boss \\& A. Weinberger for financial support.  MLM acknowledges support provided by NASA through Hubble Fellowship grant HF-01210.01-A awarded by the STScI, which is operated by the AURA, Inc., for NASA, under contract NAS5-26555. JEC would like to thank NASA's Origins of Solar Systems Program for support. We also thank A. Bonanos and all the other members of the Astronomy group at CIW-DTM for fruitful discussions. The authors also recognize the comments and suggestions made by L.Lucy and an anonymous referee which helped to improve the manuscript significantly."
    },
    "0809/0809.1796_arXiv.txt": {
        "abstract": "Preliminary results of our analysis on the extended emission of short/medium duration GRBs observed with Swift/BAT are presented. The Bayesian blocks algorithm is used to analyze the burst durations and the temporal structure of the lightcurves in different energy bands. We show here the results of three bursts (GRBs 050724, 061006 and 070714B) that have a prominent soft extended emission component in our sample. The extended emission of these bursts is a continuous, flickering-liked component, lasting $\\sim 100$ seconds post the GRB trigger at 15-25 keV bands. Without considering this component, the three bursts are classified as short GRBs, with $T_{90}=2\\sim 3$ seconds. GRB 060614 has an emission component similar to the extended emission, but this component has pulse-liked structure, possibly indicating that this emission component is different from that observed in GRBs 050724, 061006, and 070714B. Further analysis on the spectral evolution behavior of the extended emission component is on going. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2089_arXiv.txt": {
        "abstract": "\\label{abstract}  We present global, three-dimensional numerical simulations of HD 189733b and HD 209458b that couple the atmospheric dynamics to a realistic representation of non-gray cloud-free radiative transfer.  The model, which we call the Substellar and Planetary Atmospheric Radiation and Circulation (SPARC) model, adopts the MITgcm for the dynamics and uses the radiative model of McKay, Marley, Fortney, and collaborators for the radiation.   Like earlier work with simplified forcing, our simulations develop a broad eastward equatorial jet, mean westward flow at higher latitudes, and substantial flow over the poles at low pressure.  For HD 189733b, our simulations without TiO and VO opacity can explain the broad features of the observed 8 and 24-$\\mu$m light curves, including the modest day-night flux variation and the fact that the planet/star flux ratio peaks before the secondary eclipse.  Our simulations also provide reasonable matches to the {\\it Spitzer} secondary-eclipse depths at 4.5, 5.8, 8, 16, and $24\\,\\mu$m and the groundbased upper limit at $2.2\\,\\mu$m.  However, we substantially underpredict the  $3.6\\,\\mu$m secondary-eclipse depth, suggesting that our simulations are too cold in the 0.1--1 bar region.  Predicted temporal variability in secondary-eclipse depths is $\\sim1$\\% at {\\it Spitzer} bandpasses, consistent with recent observational upper limits at $8\\,\\mu$m.  We also show that nonsynchronous rotation can significantly alter the jet structure. For HD 209458b, we include TiO and VO opacity; these simulations develop a hot ($>2000\\,$K) dayside stratosphere whose horizontal dimensions are small at depth but widen with altitude.  Despite this stratosphere, we do not reproduce current {\\it Spitzer} photometry of this planet. Light curves in {\\it Spitzer} bandpasses show modest phase variation and satisfy the observational upper limit on day-night phase variation at $8\\,\\mu$m. ",
        "introduction": "\\label{Introduction}  Blasted by starlight $10^3$--$10^5$ times stronger than that received by Jupiter and experiencing modest rotation rates due to their presumed tidal locking \\citep{guillot-etal-1996}, hot Jupiters occupy a fascinating meteorological regime that does not exist in our Solar System \\citep[for an extensive review see][] {showman-etal-2008b}. Despite the wide range of transiting exoplanets that have been discovered, HD 189733b and HD 209458b remain the best characterized hot Jupiters and represent important test cases for our understanding of these objects generally.  A variety of observations now exist that constrain the 3D temperature structure, composition, hazes, and albedo for these two planets.  There is now hope that, by comparing these observations with detailed models, insights into this novel atmospheric regime can be achieved.   The strongest observational evidence for atmospheric motions comes from infrared light curves obtained with the {\\it Spitzer Space Telescope}. Continuous light curves of HD 189733b over half an orbit have now been obtained at both 8 and $24\\,\\mu$m, constraining this planet's day-night heat transport \\citep{knutson-etal-2007b, knutson-etal-2009a}. The inferred dayside and nightside brightness temperatures are $\\sim1250\\,$K and $\\sim1000\\,$K, respectively.\\footnote{At $8\\,\\mu$m, the dayside and nightside brightness temperatures are $1258\\pm11\\,$K and $1011\\pm51\\,$K.  At $24\\,\\mu$m, the dayside and nightside brightness temperatures are $1220\\pm47\\,$K and $984\\pm48\\,$K.} Because the nightside temperatures would be extremely cold ($\\sim200\\,$K) in the absence of winds, these observations imply the existence of a vigorous atmospheric circulation that efficiently transports heat from dayside to nightside. Further evidence for winds is the fact that the peak flux in both light curves occurs {\\it before} secondary eclipse, implying that the hottest region lies 20--$30^{\\circ}$ of longitude east of the substellar point \\citep{knutson-etal-2007b, knutson-etal-2009a}. For HD 209458b, current data contain insufficient temporal sampling to determine whether similar offsets exist but nevertheless demonstrate that the 8-$\\mu$m day-night brightness temperature difference is also modest \\citep{cowan-etal-2007}.   For both planets we now also have dayside photometry at all Spitzer broadband channels (centered at 3.6, 4.5, 5.8, 8.0, 16, and $24\\,\\mu$m), placing constraints on dayside composition and vertical temperature profile.  These data were obtained by differencing the photometry of star$+$planet taken just before/after secondary eclipse from photometry taken during secondary eclipse, when only the star is visible. When compared with one-dimensional (1D) atmosphere models, these data suggest that, near the photosphere pressures, the dayside temperature of HD 189733b decreases with altitude \\citep{charbonneau-etal-2008, barman-2008, knutson-etal-2009a}, whereas HD 209458b instead contains a thermal inversion (a hot stratosphere) \\citep{knutson-etal-2008a, burrows-etal-2007b}.   Following pioneering work by \\citet{hubeny-etal-2003}, \\citet{fortney-etal-2008} and \\citet{burrows-etal-2008} suggested that hot Jupiters subdivide into two classes depending on whether or not their atmospheres contain highly absorbing substances such as gaseous TiO and VO. For a Sun-like primary, solar-composition planetary atmospheres inward of 0.04--$0.05\\,$AU are hot enough to contain TiO and VO; because of the extreme visible-wavelength opacity of these compounds, such planets absorb starlight at low pressure ($\\sim$mbar) and naturally exhibit dayside stratospheres.  For planets outward of $\\sim0.05\\,$AU, temperatures drop sufficiently for TiO and VO to condense in the deep atmospheres; these planets absorb starlight deeper in the atmosphere, lack stratospheres, and show spectral bands in absorption. Based on simple comparisons of radiative and advective timescales, \\citet{fortney-etal-2008} further suggested that the former category (dubbed ``pM'' class planets) would exhibit large day-night temperature differences whereas the latter category (``pL'' class) would exhibit more modest temperature contrasts.  Their calculations suggest that HD 189733b is a pL-class planet while HD 209458b is a pM-class planet.  This dichotomy makes these two planets a particularly interesting pair for comparison.  Given these developments, there is a pressing need for 3D calculations of the atmospheric circulation on hot Jupiters.  Several groups have investigated a range of 2D and 3D models \\citep[]{showman-guillot-2002, cho-etal-2003, cho-etal-2008, cooper-showman-2005, cooper-showman-2006, showman-etal-2008a, showman-etal-2008b, langton-laughlin-2007, langton-laughlin-2008, dobbs-dixon-lin-2008}. However, to date, all published models have adopted severe approximations to the radiative transfer or excluded radiative heating/cooling entirely. While such simplified approaches are invaluable for investigating the underlying dynamics, a detailed attempt to explain wavelength-dependent photometry and light curves must include a realistic coupling of the radiative transfer to the dynamics.  Here, we present new numerical simulations that couple a realistic representation of non-gray cloud-free radiative transfer to the dynamics. Surprisingly, coupling of 3D dynamics to nongray radiative transfer has never previously been done for any giant planet, not even Jupiter, Saturn, Uranus, and Neptune. This is the first 3D dynamical model for any giant planet --- inside our Solar System or out --- to do so.  We dub our model the Substellar and Planetary Atmospheric Radiation and Circulation model, or SPARC model, and to honor its dynamical heritage usually refer to it as SPARC/MITgcm. \\S 2 presents the model. \\S 3 describes the basic circulation regime and resulting light curves and spectra for HD 189733b. \\S 4 describes the effect of a nonsynchronous rotation rate, \\S 5 considers HD 209458b, and \\S 6 concludes. ",
        "conclusions": "We presented global, three-dimensional numerical simulations of the atmospheric circulation of HD 189733b and HD 209458b that couple the dynamics to a realistic representation of cloud-free non-gray radiative transfer.  This new model, which we dub SPARC/MITgcm, is the first 3D dynamical model for any giant planet --- including those in our Solar System --- to incorporate nongray radiative transfer. Our model adopts the MITgcm for the dynamics and an optimized version of the radiative model of McKay, Marley, Fortney, and collaborators for the radiative transfer.  Opacities are calculated assuming solar composition (or some multiple thereof) with equilibrium chemistry that accounts for rainout.  Like earlier work with simplified forcing \\citep{showman-etal-2008a, cooper-showman-2005, showman-guillot-2002}, our simulations develop a broad eastward equatorial jet, mean westward flow at high latitudes, and substantial flow over the poles at low pressure.  The jet structure depends significantly on longitude at pressures $<100\\,$mbar.  For HD 189733b, our simulations that exclude TiO and VO opacity explain the broad features of the observed 8 and 24-$\\mu$m light curves \\citep{knutson-etal-2007b, knutson-etal-2009a}, including the modest day-night flux variation and the fact that the planet/star flux ratio peaks before the secondary eclipse.  In our simulations, the offset results from the eastward displacement of the hot regions from the substellar point. On the other hand, we do not fit the flux minimum seen after transit in the 8-$\\mu$m light curve \\citep{knutson-etal-2007b}. Our simulations also provide reasonable matches to the {\\it Spitzer} secondary-eclipse depths at 4.5, 5.8, 8, 16, and $24\\,\\mu$m \\citep{charbonneau-etal-2008, deming-etal-2006} and the groundbased upper limit at $2.2\\,\\mu$m from \\citet{barnes-etal-2007}.  The temporal variability in these simulations is modest --- of order 1\\% --- and is fully consistent with the upper limit on temporal variability from \\citet{agol-etal-2008}.  The primary HD 189733b observation where we fare poorly is the $3.6\\,\\mu$m secondary-eclipse depth from \\citet{charbonneau-etal-2008}, which we underpredict by about a factor of two.  Because the $3.6\\,\\mu$m channel is expected to sense down to $\\sim0.1$--1 bar on this planet (Fig.~\\ref{contribution-fcns}, {\\it left panel}), this suggests that our simulation is too cold in this region of the atmosphere and/or has the incorrect opacity in this wavelength range.  Previous 1D models of HD 189733b have suffered a similar problem at this wavelength \\citep{barman-2008, knutson-etal-2009a}, as have 1D models of brown dwarfs with similar effective temperatures \\citep{geballe-etal-2009}. A full-orbit light curve of HD 189733b at $3.6\\,\\mu$m, possible with warm {\\it Spitzer}, would provide crucial insights to help resolve this problem.  For HD 209458b, we include gaseous TiO and VO opacity to see whether it allows us to explain the inference of a stratosphere from {\\it Spitzer} photometry \\citep{knutson-etal-2008a, burrows-etal-2007b, fortney-etal-2008}.  As expected, these simulations develop a hot ($>2000\\,$K) dayside stratosphere whose horizontal dimensions are small at depth but widen with altitude until the stratosphere covers most of the dayside at pressures $<0.1\\,$mbar. Interestingly, both branches of the bifurcation discussed by \\citet{hubeny-etal-2003} for different planets occur here on a single planet's dayside.  Using the terminology of \\citet{fortney-etal-2008}, the substellar region is a pM-class planet but the terminators and nightside are a pL-class planet.  It is thus perhaps more proper to talk about pM-class {\\it daysides} rather than pM-class {\\it planets}.  But despite the stratosphere in our simulations of HD 209458b, we do not reproduce current {\\it Spitzer} photometry of this planet, which includes particularly high ($\\sim1700$--$1900\\,$K) brightness temperatures in the 4.5 and 5.8-$\\mu$m channels.  This could mean that our stratosphere has the incorrect properties (e.g., temperature range, altitude range, and vertical thermal gradient). However, a fundamental difficulty in explaining the diverse brightness dayside temperatures (ranging from $\\sim1500$--$1900\\,$K in {\\it Spitzer} bandpasses) is that the range of pressures from which emergent infrared photons originate are very similar for all {\\it Spitzer} bandpasses, at least as calculated by our radiative model with equilibrium chemistry.   This means that {\\it regardless} of the planet's temperature profile, the brightness temperatures in all these bands should be similar. Breaking this degeneracy would require changing the opacities so that the opacities in different {\\it Spitzer} bandpasses differ significantly. Dropping the equilibrium chemistry assumption would make this possible; this provides a clue that disequilibrium chemistry may be important. Disequilibrium chemistry appears to influence infrared spectra of brown dwarfs \\citep[e.g.][]{leggett-etal-2007a, geballe-etal-2009}, lending weight to this possibility.   Our light curves of HD 209458b in {\\it Spitzer} bandpasses exhibit modest day-night variation, and we successfully explain the upper limit on the day-night flux contrast from \\citet{cowan-etal-2007} at $8\\,\\mu$m.  Our ability to meet this constraint results from the fact that much of the dayside 8-$\\mu$m radiation emanates from altitudes near the base of the stratosphere, where temperatures are not too hot, while nightside 8-$\\mu$m radiation emanates from substantially deeper pressures where temperatures are warmer than they are aloft. This dual effect leads to modest day-night flux variations despite large day-night temperature variations (on isobars) at low pressures. Assuming the high inferred 4.5 and 5.8-$\\mu$m brightness temperatures indeed result from a stratosphere, we suggest that the real planet should exhibit large phase variations at 4.5 and $5.8\\,\\mu$m yet more modest phase variations at other {\\it Spitzer} bandpasses and $K$ band.   The task of developing exoplanet GCMs that couple dynamics and radiative transfer has now been espoused by many authors, and the work presented here demonstrates that this approach indeed holds significant promise for explaining atmospheric observations of hot Jupiters.  Our 3D SPARC/MITgcm simulations, especially of HD 189733b, show encouraging resemblance to observations of the real planet.  While some discrepancies remain, and a wider range of parameters need be explored, these simulations support the idea that detailed model/data comparisons can eventually allow robust inferences about the circulation regime of hot Jupiters to be inferred. Given the huge range in properties of currently known transiting planets, with future observations sure to unveil additional surprises, this helps energize the exciting prospect that planetary meteorology can successfully be extended beyond the confines of our Solar System."
    },
    "0809/0809.2040_arXiv.txt": {
        "abstract": "We investigated the influence of environment on cluster galaxies by examining the alignment of the brightest cluster galaxy (BCG) position angle with respect to the host cluster X-ray position angle. The cluster position angles were measured using high spatial resolution X-ray data taken from the Chandra ACIS archive, that significantly improved the determination of the cluster shape compared to the conventional method of using optical images. Meanwhile, those of the BCGs were measured using homogeneous dataset composed of high spatial resolution optical images taken with Suprime-Cam mounted on Subaru 8m telescope.  We found a strong indication of an alignment between the cluster X-ray emission and optical light from BCGs, while we see no clear direct correlation between the degree of ellipticity of X-ray and optical BCG morphologies, despite the apparent alignment of two elliptical structures. We have also investigated possible dependence of the position angle alignment on the X-ray morphology of the clusters, and no clear trends are found. The fact that no trends are evident regarding frequency or degree of the alignment with respect to X-ray morphology may be consistent with an interpretation as a lack of dependence on the dynamical status of clusters. ",
        "introduction": "It is well established that the major axes of galaxy clusters tend to point toward their nearest neighbour \\citep[e.g.][Hashimoto et al. 2007a]{binggeli1982sao,flin1987agc,rhee1987sea,plionis1994paa,west1995scf,plionis2002csa,chambers2000eca,chambers2002nna}.  \\nocite{hashimoto2007icx}  Another alignment effect is that between the orientation of the brightest cluster galaxy (BCG), or cD galaxy,  and that of their parent cluster \\citep[e.g.][]{sastry1968cas,dressler1978csv,carter1980mcg,binggeli1982sao,struble1990pas,plionis2003gap}. Similar alignment is also reported for poor groups of galaxies \\citep{fuller1999adg}. Numerical work \\citep[e.g.][]{west1991fsu,vanhaarlem1993vfa,west1994amh,onuora2000acu,splinter1997eao,faltenbacher2005ima} show that substructure-cluster alignments can occur naturally in hierarchical clustering models of structure formation such as the cold dark matter model.  Unfortunately, all of these previous galaxy-cluster alignment studies are using optically-determined cluster position angles, most of them are  based on the Palomar Observatory Sky Survey (POSS). Despite the importance of these optical investigations, individual galaxies may not be the best tracers of the shape of a cluster. Problems can arise from foreground/background contamination, as well as the fact that galaxies contribute discreteness noise. However, it is believed that the X-ray emitting gas within a cluster traces its gravitational potential \\citep{sarazin1986xec}. X-ray morphology may then be one of the best observable phenomenon for determining the cluster shape and orientation. Indeed, there are several X-ray studies for cluster vs. neighbour-cluster alignment \\citep*[e.g.][Hashimoto et al. 2007a]{ulmer1989mar,chambers2000eca,chambers2002nna}, but there are few galaxy-cluster alignment studies using X-ray morphology, except for \\citet*{porter1991coa} and \\citet*{rhee1991xro}, where they  investigated the BCG-cluster alignment using low spatial resolution X-ray data, as well as traditional cluster shape parameter from apparent galaxy distribution, and reported a significant alignment. Unfortunately, previous X-ray studies are mostly based on $Einstein$ data. These data are important, but the exposure depths are small and the spatial resolution is rather low compared to recently available X-ray data. The low spatial resolution may not significantly affect relatively robust measures such as position angle in a direct way, but it will critically hinder the accurate removal of contaminating point sources and the accurate determination of cluster center, and that may significantly affect the estimate of cluster X-ray morphology including the position angle. Hence, a new investigation using deeper X-ray data with much higher spatial resolution is needed.  Here we report a new investigation of galaxy alignment with respect to its parent cluster, using the cluster position angle and ellipticity determined by high spatial resolution X-ray data taken from the Chandra ACIS archive. Meanwhile, position angle and ellipticity of BCGs are determined from optical images taken with Subaru 8m telescope. This paper is organized as follows. In Sec. 2, we describe our main  sample and X-ray measures, in Sec. 3, details of our optical data are described. Sec. 4 summarizes our results. Throughout the paper, we use $H_{o}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$=0.3, and $\\Omega_{\\Lambda}$=0.7, unless otherwise stated.   \\section[]{The X-ray Sample, X-ray Data Preparation, and X-ray measures}  Here we briefly summarize our main X-ray sample, X-ray data preparation, and X-ray measures. \\nocite{hashimoto2007rqm} More detailed descriptions can be found in Hashimoto et al. (2007b).  Almost all clusters are selected from flux-limited X-ray surveys, and X-ray data are taken from the Chandra ACIS archive. A lower limit of z = 0.05 or 0.1 is placed on the redshift to ensure that a cluster is observed with sufficient field-of-view with ACIS-I or ACIS-S, respectively. The majority of our sample comes from the $ROSAT$ Brightest Cluster Sample \\citep[BCS;][]{ebeling1998rbc} and the Extended $ROSAT$ Brightest Cluster Sample \\citep[EBCS;][]{ebeling2000rbc}. When combined with EBCS, the BCS clusters represent one of the largest and most complete X-ray selected cluster samples, which is currently the most frequently observed by $Chandra$. To extend our sample to higher redshifts, additional high-z clusters are selected from various deep  surveys; 10 of these clusters are selected from the $ROSAT$ Deep Cluster Survey \\citep[RDCS;][]{rosati1998rdc}, 10 from the $Einstein$ Extended Medium Sensitivity Survey \\citep[EMSS;][]{gioia1990eoe,henry1992ems}, 14 from the 160 Square Degrees $ROSAT$  Survey \\citep{vikhlinin1998cgc}, 2 from the Wide Angle $ROSAT$ Pointed Survey \\citep[WARPS;][]{perlman2002wsv}, and 1 from  the North Ecliptic Pole survey \\citep[NEP;][]{henry2006rne}, RXJ1054 was discovered by \\citet{hasinger1998rds}, RXJ1347 was discovered in the $ROSAT$ All Sky Survey \\citep{schindler1995das}, and 3C295 has been mapped with $Einstein$ \\citet{henry1986xrs}.  The resulting sample contains 120 clusters. At the final stage of our data processing, to employ our full analysis, we further applied a selection based on the total counts of cluster emission, eliminating clusters with very low signal-to-noise ratio. Clusters whose center is too close to the edge of the ACISCCD are also removed. The resulting final sample contains 101 clusters with redshifts between 0.05 - 1.26 (median z = 0.226), and bolometric luminosity between 1.0 $\\times$ 10$^{44}$ -- 1.2 $\\times$ 10$^{46}$ erg s$^{-1}$ (median 8.56 $\\times$ 10$^{44}$ erg s$^{-1}$). We reprocessed the level=1 event file retrieved from the archive. The data were filtered to include only the standard event grades 0,2,3,4,6 and status 0, then multiple pointings were merged, if any. We eliminated time intervals of high background count rate by performing a 3 $\\sigma$ clipping of the background level. We corrected the images for exposure variations across the field of view, detector response and telescope vignetting.  We detected point sources using the CIAO routine celldetect with a signal-to-noise threshold for source detection of three. We removed point sources, except for those at the center of the cluster which was mostly the peak of the surface brightness distribution rather than a real point source. The images were then smoothed  with Gaussian $\\sigma$=5\". We decided to use isophotal contours to characterize an object region, instead of a conventional circular aperture, because we did not want to introduce any bias in the shape of an object. To define constant metric scale to all clusters, we adjusted an extracting threshold in such a way that the square root of the detected object area times a constant was 0.5 Mpc, i.e. const$\\sqrt{area}$ = 0.5 Mpc. We chose const =1.5, because the isophotal limit of a detected object was best represented by this value.  The morphology of cluster X-ray emission is then characterized objectively by the position angle, as well as the ellipticity and the asymmetry. The position angle is defined by the orientation of the major axis measured east from north. Ellipticity is simply defined by the ratio of semi-major (A) and semi-minor axis (B) lengths as: \\begin{eqnarray} Elli & = & 1  -B/A \\end{eqnarray} where A and B are defined by the maximum and minimum spatial {\\it rms} of the object profile along any direction and computed from the centered-second moments by the formula: \\\\ \\begin{eqnarray} A^2 & = &\\frac{\\overline{x^2}+\\overline{y^2}}{2}+\\sqrt{\\left(\\frac{\\overline{x^2}-\\overline{y^2}}{2}\\right)^2 + \\overline{xy}^2}  \\\\ B^2 & = &\\frac{\\overline{x^2}+\\overline{y^2}}{2}-\\sqrt{\\left(\\frac{\\overline{x^2}-\\overline{y^2}}{2}\\right)^2 + \\overline{xy}^2} \\end{eqnarray}  The asymmetry is measured by first rotating a cluster image by 180 degrees around the object center, then subtracting the rotated image from the original unrotated one. The residual signals above zero are summed and then normalized. Please see Hashimoto et al. 2007b for more detailed definitions of morphological measures.   \\section[]{Optical data}  To determine the position angle and the ellipticity of BCGs, we used optical broad band images taken with Supreme-Cam \\citep{miyazaki1998cam} on the Subaru telescope. The data were retrieved from Subaru-Mitaka-Okayama-Kiso Archive (SMOKA). Reduction software developed by \\citet{yagi2002lfn} was used for flat-fielding, instrumental distortion correction, differential refraction, sky subtraction, and stacking. The camera covers a 34' $\\times$ 27' field of view with a pixel scale of 0\\farcs202. The photometry is calibrated to Vega system using Landolt standards \\citep{landolt1992ups}. We refine the original astrometry written as WCS keyword in the distributed archival data using using USNO-A2 catalog with positional uncertainties less than  $\\sim$ 0.2 arcsec. The data were taken under various seeing conditions, and we used only images with less than $\\sim$ 1\\farcs2 seeing. The optical data retrieved from SMOKA contains 30 clusters with redshifts between 0.08 - 0.9,  Some clusters have observed through many wavebands, and that allowed us to investigate the possible variations of our measures caused by waveband shifts. We have decided to rely primarily on the R band images  for this alignment study, because we found that the effect of waveband shift is negligible.   The position angle and ellipticity of BCGs are measured exactly the same way as the X-ray cluster emission, namely, the position angle is defined by the orientation of the major axis measured east from north, and the ellipticity is  defined by  the ratio of semi-major and semi-minor axis. Please see Hashimoto et al. 2007b \\citep[see also][]{hashimoto1998ies} for further details. As a precaution, we investigate the effect of superposed small galaxies sometimes lying on top of the extended structure of some BCGs, and we found that these superposed small galaxies have little effect on our robust measures such as, position angle and ellipticity.   \\section[]{Results}  \\subsection[]{Systematics}  One of the haunting, yet unfortunately often lightly treated, problem of any study comparing complex morphological characteristics of astronomical objects is the  possible systematics introduced by various data quality, exposure times and object redshifts. Depending on the sensitivity of measures of characteristics, some susceptible measures may be seriously affected by these systematics, producing the misleading results.  Unfortunately, investigating the systematics on the complex characteristics is not an easy task. To investigate the  systematic effect of, for example, various exposure times, one of the  standard approaches is to simulate an image with a given exposure time by using an exposure-time-scaled and noise-added model image. Unfortunately, we need to approximate the various characteristics of a model to the complicated characteristics of a real object, (and those characteristics are often what we want to investigate) and this is an almost impossible task.  Meanwhile, if we use the real data, instead of the model, we will not have this problem. We can simulate lower signal-to-noise data caused by a shorter integration time by scaling the real data  by the exposure time, and adding Poisson noise taking each pixel value as the mean for a Poisson distribution. However, this simple rescaling and adding-noise process will produce an image containing an excessive amount of Poisson noise for a given exposure time, thus lead us to underestimate the data quality. This inaccurate estimate of noise is caused by  the intrinsic noise already presented in the initial real data. The intrinsic noise is difficult to be removed without sacrificing the fine spatial details of the object.  Similarly, to investigate the effect of dimming and smaller angular size caused by higher  redshifts, in addition to the rest waveband shift effect, simple rescaling and rebinning of the real data  will not work, because these  manipulations will again produce the incorrect amount of noise.  Further difficulty associated with simulation using the real data comes from the fact that exposure and redshift effects are often coupled, because, in the real observation, low redshift objects are usually observed with shorter exposures than high redshift objects. This coupling further poses a serious problem, because simple standard method of simulating an observation with `decreased' exposure time will force high redshift data to get degraded, which greatly reduces signal-to-noise ratio of already low quality high redshift data.  To circumvent all of these challenging problems, we developed a very useful simulating technique employing a series of `adaptive scalings'  accompanied by a noise adding process applied to the real images. This technique allows us to simulated an image of desired fiducial exposure time and redshift with correct signal-to-noise ratio without using a tricky artificial model image, thus we can easily investigate the effect of various image quality and/or  easily change the real data to common fiducial exposure and redshift for easy comparison. Moreover, the technique can provide us with a powerful tool for conducting evolutionary studies, enabling us to compare the local objects  to the high redshift objects  without degrading photon-expensive high redshift data of low signal-to-noise ratio at all. This method is originally developed for the comparison of X-ray image data, but can be  used for almost all kind of imaging data, including optical and NIR images.  Here we briefly describe the method. Please see Hashimoto et al. 2007b for further details. To simulate data with integration t=t1, an original unsmoothed image (including the background) taken with original integration time t0 was at first rescaled by a factor R$_0$/(1-R$_0$), instead of simple R$_0$,  where R$_0$=t1/t0, t0$>$t1. That is, an intermediate scaled image I$_1$ was created from the original unsmoothed image I$_0$ by: \\begin{eqnarray} I_1   &=& I_0\\frac{R_0}{(1-R_0)}. \\end{eqnarray}  Poisson noise was then added to this rescaled image by taking each pixel value as the mean for a Poisson distribution and then randomly selecting a new pixel value from that distribution. This image was then rescaled again by a factor (1-R$_0$) to produce an image whose {\\it signal} is scaled by R$_0$ relative to the original image, but its {\\it noise} is approximately scaled by $\\sqrt{R_0}$, assuming that the intrinsic noise initially present in the real data is Poissonian.  Similarly, to simulate the dimming effect by the redshift, an intermediate scaled image I$_1$ is created from the background subtracted image I$_0$ by a pixel-to-pixel manipulation:  \\begin{eqnarray} I_{1}(x,y) &=& \\frac{I_{0}(x,y)^2R_1^2}{[I_{0}(x,y)R_1+B-R_1^2(I_{0}(x,y)+B)]} \\\\ where & &\\nonumber \\\\ R_1&=&[(1+z0)/(1+z1)]^4 \\end{eqnarray} where z0 and z1 are the original redshift and the  new redshift of the object, respectively, and B is the background.  Finally, to  simulate the angular-size change due to the redshift difference between z0 and z1, the original image will be rebinned by a factor R$_2$, then intermediate scaled image will be created by  rescaling the rebinned image by a factor 1/(R$^2_2$-1), before the addition of the  Poisson noise. For the simulation with `increased' exposure time, this factor can be changed to R$_3$/(R$_2^2$-R$_3$) where  R$_3$ = t2/t0, t2$>$t0, where  t2 is the increased exposure time, and t0 is the original integration time, and (R$_2^2$ -R$_3$) $>$ 0. The maximum length of integration time we can `increase' (t2$_{max}$) is naturally limited by the original  exposure time and how much we increase the redshift for the redshift-effect part, and determined by the relationship, \\begin{eqnarray} R_2^2-R_3 = 0, \\end{eqnarray} which is equivalent to the case when  no Poisson noise is added after the rebinning. Thus, \\begin{eqnarray} t2_{max}=t0R_2^2. \\end{eqnarray} This t2$_{max}$ can be also used as a rough estimate of the effective image depth. The t2$_{max}$ provides an estimate of the image depth much more effectively  than the conventional simple exposure time because t2$_{max}$ is related to a quantity that is affected both by exposure time and redshift, and thus enabling us to quantitatively compare exposure times of observations involving targets at different redshifts (e.g. 100 ksec at z=0.1 and 100 ksec at z=0.9).  \\begin{figure} \\center{ {\\includegraphics[height=7cm,width=5cm,clip,angle=-90]{sampleimage.ps}} } \\caption{ Simulating an image with desired exposure time and redshift using the real data: Even simulating an image with prolonged exposure time is possible with our adaptive scaling method. Here, optical R band images, taken with  Subaru Suprime-Cam, of the BCG at the center of an example cluster (Abell 2219) are shown. Images with original and modified exposure time and redshift are presented with north is up and east is left. (a) Original image: exptime(t)=240s, and redshift(z)=0.228, (b) Simulated shorter exposure image with t=10s (c) Simulated high-z  image with z=0.9, t=240s (d) Simulated prolonged exposure at high-z  with t=1092s, z=0.9. } \\label{FigTemp} \\end{figure}  \\begin{figure} \\center{ {\\includegraphics[height=7cm,width=5cm,clip,angle=-90]{sampleimage_x.ps}} } \\caption{ Similarly with Fig. 1, X-ray images from Chandra ACIS, of Abell 2219 are shown with original and modified exposure time and redshift. North is up and east is left. (a) Original: t=41ks, z=0.228 (b) Simulated shorter exposure: t=10ks (z=0.228) (c) Simulated High-z image: z=0.9 (t=41ks) (d) Prolonged exposure at High-z: t=188ks, z=0.9. } \\label{FigTemp} \\end{figure}  Although we suspected that our ellipticity and position angle were quite robust, as a precaution we investigated  the possible systematics on these measures introduced by  various  exposure times and redshifts, using our scaling technique described above.   In Fig. 1, we demonstrate our technique of simulating desired exposure time and redshift using the real optical image of the BCG at the cluster center taken with  Subaru Suprime-Cam. Original and modified exposure time and redshift  of an example cluster (Abell 2219) are shown with north  up and east left, where (a) original image: exptime(t)=240s, and redshift(z)=0.228, (b) simulated shorter exposure image with t=10s, (c) simulated high-z  image with z=0.9, t=240s, and (d) simulated prolonged exposure at high-z  with t=1092s, z=0.9.  Similarly, in Fig. 2, we use the real X-ray images of Abell 2219 from Chandra ACIS, and simulated various exposures and redshifts, where (a) original image with t=41ks, z=0.228, (b) simulated shorter exposure: t=10ks (z=0.228), (c) simulated High-z image: z=0.9 (t=41ks), and (d) prolonged exposure at High-z: t=188ks, z=0.9.  Using this technique, we simulated datasets with various exposure  times and redshifts, and measured our cluster parameters. We found that our X-ray and optical position angles are robust against various exposure times and redshifts. Similarly, we found that other morphological measures such as, the ellipticity and asymmetry are  quite robust, as well.   \\subsection[]{Analyses}  \\begin{table} \\tiny \\caption{Summary of optical cluster sample} \\label{symbols} \\begin{tabular}{@{}lcccccc} \\hline Cluster & z &   PA\\_X & PA\\_BCG & $\\Delta$PA$^a$ & Elli\\_X & Elli\\_BCG \\\\ & (redshift) &   (degree) & (degree) & (degree)   &          &  \\\\ \\hline \\hline \\scriptsize a2034 & 0.110  &  206.9 & 22.24 & 4.65 & 0.15 & 0.50 \\\\ a2069 & 0.114  &  327.7 & 331.8 & 4.18 & 0.46 & 0.66 \\\\ a750 & 0.163  &  249.0 & 249.6 & 0.66 & 0.14 & 0.31 \\\\ rxj1720 & 0.164  &  355.3 & 32.08 & 36.7 & 0.15 & 0.34 \\\\ a520 & 0.203  &  192.4 & 228.6 & 36.2 & 0.26 & 0.49 \\\\ a963 & 0.206  &  175.9 & 347.0 & 8.85 & 0.15 & 0.42 \\\\ a2261 & 0.224  &  225.8 & 14.48 & 31.3 & 0.14 & 0.13 \\\\ a2219 & 0.228  &  309.5 & -79.5 & 29.0 & 0.38 & 0.42 \\\\ a2390 & 0.233  &  298.6 & -57.6 & 3.76 & 0.30 & 0.32 \\\\ rxj2129 & 0.235  &  246.0 & 65.28 & 0.71 & 0.21 & 0.54 \\\\ a2631 & 0.278  &  258.2 & 84.27 & 6.07 & 0.29 & 0.41 \\\\ a1758 & 0.280  &  308.7 & 85.53 & 43.2 & 0.47 & 0.21 \\\\ a2552 & 0.299  &  201.7 & 216.4 & 14.7 & 0.18 & 0.45 \\\\ a1722 & 0.327  &  204.4 & 357.7 & 26.6 & 0.27 & 0.14 \\\\ zwcl3959 & 0.351  &  333.4 & 346.7 & 13.3 & 0.21 & 0.36 \\\\ a370 & 0.357  &  187.2 & 86.96 & 79.7 & 0.37 & 0.11 \\\\ rxj1532 & 0.361  &  227.7 & 78.24 & 30.5 & 0.18 & 0.37 \\\\ zwcl1953 & 0.373  &  351.2 & 306.3 & 44.8 & 0.24 & 0.51 \\\\ zwcl2661 & 0.382  &  159.5 & 206.9 & 47.3 & 0.12 & 0.43 \\\\ zwcl0024 & 0.390  &  5.200 & 246.6 & 61.4 & 0.03 & 0.48 \\\\ rxj2228 & 0.412  &  263.7 & 56.27 & 27.4 & 0.21 & 0.54 \\\\ rxj1347 & 0.451  &  359.2 & 343.8 & 15.3 & 0.20 & 0.13 \\\\ ms0451 & 0.540  &  279.9 & 344.9 & 65.0 & 0.26 & 0.04 \\\\ cl0016 & 0.541  &  228.7 & 244.6 & 15.9 & 0.19 & 0.31 \\\\ ms2053 & 0.583  &  304.8 & 336.8 & 32.0 & 0.25 & 0.16 \\\\ rxj1350 & 0.810  &  334.2 & 308.9 & 25.2 & 0.20 & 0.47 \\\\ rxj1716 & 0.813  &  236.2 & 255.1 & 18.9 & 0.17 & 0.52 \\\\ ms1054 & 0.830  &  266.1 & 35.96 & 50.1 & 0.40 & 0.51 \\\\ rxj0152 & 0.835  &  221.6 & 231.0 & 9.46 & 0.58 & 0.29 \\\\ wga1226 & 0.890  &  284.5 & 264.6 & 19.8 & 0.09 & 0.15 \\\\ \\hline \\hline \\end{tabular} a: $\\Delta$PA is an acute angle of PA\\_BCG-PA\\_X  \\end{table}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[height=3cm,width=2cm,clip,angle=90]{plot1g.elli_elli_gal_Sa_.degr_BCG01.ps}} \\caption{ Ellipticity of BCGs is plotted against ellipticity of the X-ray morphology of the host clusters. Interestingly, no clear correlation is seen. } \\label{FigTemp} \\end{figure}   Table 1 shows a summary of our optical cluster sample, where $\\Delta$PA is an acute relative angle between the position angle of X-ray (PA\\_X) and BCG (PA\\_BCG), namely, the relative position angle differences greater than 90 degree are `folded' and changed to be acute ranging between 0 and 90 degree. Despite  the robust nature of our measures, we modify, as a precaution, all of the X-ray and optical observations to be equivalent to z=0.9 and t=t2$_{max}$ to eliminate any possible small systematics, but otherwise to maximize the image quality.  In Fig. 3, the ellipticity of cluster X-ray morphology is plotted against the ellipticity of optical morphology of BCGs. Interestingly, in spite of expected alignment of two elliptical structures, there are a large scatter and we found no strong correlation in the relationship between the two ellipticities.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[height=3cm,width=2cm,clip,angle=90]{plot1g.difang2_elli_gal_Sa_.degr_BCG01.ps}} \\resizebox{\\hsize}{!}{\\includegraphics[height=3cm,width=2cm,clip,angle=90]{plot1g.difang2_elli_Sa_.degr_BCG01.ps}} \\caption{ The acute position angle difference plotted versus ellipticity of BCGs (top panel) and ellipticity of cluster X-ray emission. The position angle difference is determined by the difference between  the cluster X-ray position angle and the position angle of BCG galaxy. Figures illustrate that clusters with the position angle difference less than 45 degree tend to be more abundant, particularly for high ellipticity BCGs or clusters, implying that cluster X-ray emission  and optical light from BCG are aligned. } \\label{FigTemp} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[height=2cm,width=3cm,clip,angle=0]{plot2g.Num_difang2_Sa_degr.ps}} \\caption{ The frequency distribution of the  position angle difference. There is a tendency that we have more clusters with an angle difference less than 45 degrees, consistent with the observation made in Fig. 4 implying that cluster X-ray emission  and optical light from BCG are aligned. } \\label{FigTemp} \\end{figure}  Fig. 4 shows the acute relative position angle difference between cluster X-ray morphology and BCG morphology plotted against  the ellipticity of BCGs (top panel) and the ellipticity of cluster X-ray morphology (bottom panel). Fig. 4 shows that the position angle difference tends to be  smaller than 45 degree, implying that BCGs tend to elongated in the same direction of the X-ray distribution of  their host clusters, particularly for clusters exhibiting relatively high ellipticity in their optical BCG morphology and/or in their cluster X-ray  morphology. Meanwhile, for clusters with very low BCG or X-ray ellipticity (ellipticity $<$ 0.1) position angle are generally poorly determined, and thus position angle difference can be inaccurate.  Fig. 5 shows the frequency distribution of the  position angle difference. There is a strong tendency that we have more clusters with an angle difference less than 45 degrees, consistent with the observation made in Fig. 4.  To test this trend more rigorously, we first employed the Kolmogorov-Smirnov (K-S) test. The null hypothesis here is that our sample can be drawn from a parent population of random position angle differences. However, the K-S test detects the  deviation from the parent population (here the population of  random position angle differences), thus it may loose some sensitivity for testing the cluster alignment, where it is likely that position angle difference is systematically lower than the random sample. To increase the sensitivity to an alignment signal, as a second test we employed the Wilcoxon-Mann-Whitney  rank-sum test. The null hypothesis of this test is that the position angle difference is not systematically smaller or larger than the random sample. Therefore the test is insensitive to an excess of angles around the mean (i.e. 45 deg). When applied to our sample, both K-S and rank-sum tests show, not surprisingly  the strong alignment signals, and we find that the null hypothesis can be rejected with 99.93\\% and 99.99\\% confidence, respectively, thus confirming that BCGs are significantly aligned to the X-ray emissions of the host clusters. We have also investigated the alignment of other luminous non-BCG galaxies to the X-ray emissions and we found no significant alignment. The results are summarized in table 2, where LG2 is the second brightest galaxy, LG3 is the third brightest galaxy, and LGn is the n-th brightest galaxy  within a projected distance of 1 Mpc from the X-ray center.   \\begin{table} \\caption{Significance levels of the alignment for various luminous galaxies } \\label{symbols} \\begin{tabular}{@{}lccccc} \\hline Statistics & BCG & LG2 &  LG3 & LG4 & LG5  \\\\ \\hline \\hline K-S      & 99.93  & 81.67 & 49.93 & 4.22 & 8.99 \\\\ Rank Sum & 99.99  & 83.06 & 57.09 & 54.51 & 63.33 \\\\ \\hline \\hline \\end{tabular} \\end{table}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[height=3cm,width=2cm,clip,angle=90]{plot1g.elli_asym_Sa_.degr_BCG01.ps}} \\caption{ X-ray morphology versus BCG alignment: Cluster X-ray asymmetry is plotted against cluster X-ray ellipticity. Large solid circles are clusters showing strong alignment with BCG-to-cluster position angle difference less than 20 degree, while solid triangles are clusters showing the ``modest''alignment, but with position angle difference between 20 and 45 degree. Crosses represent clusters with no sign of the alignment, and small dots are clusters in our main sample without optical Subaru data. No clear trends are visible regarding frequency or degree of the alignment with respect to X-ray morphology, which can be interpreted as lack of dependence on the dynamical status of clusters. } \\label{FigTemp} \\end{figure}  In Fig. 6, we investigated possible dependence of the position angle alignment on the X-ray morphology of the clusters. In Fig. 6, the cluster X-ray asymmetry is plotted against cluster X-ray ellipticity. Large solid circles are clusters showing strong alignment between the cluster and BCG with position angle difference less than 20 degree, while solid triangles are clusters showing the alignment, but with position angle difference between 20 and 45 degree. Crosses represent clusters with no sign of the alignment, and small dots are clusters in our main X-ray sample without optical Subaru data. No clear trends are evident regarding frequency or degree of the alignment with respect to X-ray morphology.  We have also attempted to investigate possible dependence of the alignment on cluster redshifts. We found that for clusters less than  z=0.35, both K-S and rank-sum tests show that the null hypothesis can be rejected with 99.92\\% and 99.99\\% confidence, respectively, while for clusters greater than or equal to  z=0.35, alignment signals are somewhat weaker that the null hypothesis can be rejected with 93.88\\% and 83.69\\% confidence, respectively for K-S and rank-sum tests. Similarly, we have investigated the dependence of the alignment on the cluster X-ray bolometric luminosity (Lbol), and we did not find any significant trend: for clusters with  Lbol greater than or equal to 2 $\\times$10$^{45}$ erg/s, the null hypothesis can be rejected with 93.45\\% and 96.42\\% confidence, while for clusters with Lbol smaller than 2 $\\times$10$^{45}$ erg/s, the null hypothesis can be rejected with 99.78\\% and 99.90\\% confidence, respectively for K-S and rank-sum tests. ",
        "conclusions": "We investigated the influence of environment on cluster galaxies by examining the alignment of the  BCG position angle with respect to the host cluster X-ray position angle. The cluster position angles were measured using high spatial resolution X-ray data taken from the Chandra ACIS archive, that significantly improved the determination of the cluster shape compared to the conventional method of using optical images. Meanwhile, those of the BCGs were measured using high spatial resolution optical images taken with Suprime-Cam mounted on Subaru 8m telescope.  We found a strong indication of an alignment between the cluster X-ray emission and optical light from BCGs, while we see no clear direct correlation between the ellipticity of X-ray morphology and optical BCG morphology despite of the apparent alignment of two elliptical structures. In the hierarchical structure formation models, the alignment effect could be produced  by clustering models of structure formation such as the cold dark matter model \\citep[e.g.][]{salvadorsole1993tie,west1994amh,usami1997ter,onuora2000acu,faltenbacher2002cog,faltenbacher2005ima}. The existence of the alignment effects is also consistent with a cosmic structure formation model such as the hot dark matter model \\citep[e.g.][]{zeldovich1970gia},  where clusters and galaxies form by fragmentation in already flattened sheet- and filament-like structures.  We have also investigated possible dependence of the position angle alignment on the X-ray morphology of the clusters, and no clear trends are found. If the X-ray morphology of clusters reflects dynamical status of clusters \\citep[e.g.][]{hashimoto2004aca}, the fact that no trends are evident regarding frequency or degree of the alignment with respect to X-ray morphology may be consistent with an interpretation as a lack of dependence of alignment on the dynamical status of clusters. Primordial galaxy alignments in clusters can be damped by various mechanisms such as  the exchange of angular momentum in galaxy encounters, violent relaxation, and secondary infall \\citep[e.g.][]{quinn1992gaa,coutts1996sad}  over a Hubble time. Thus, in highly relaxed clusters, we might naively expect to observe weaker primordial galaxy alignment, because there has been sufficient time to mix the phases. The fact that we do not see any significant alignment trend with respect to the X-ray morphology may provide an important constrain on these damping scenarios."
    },
    "0809/0809.2276_arXiv.txt": {
        "abstract": "In the infrared, the heavily reddened LkH$\\alpha$ 101 is one of the brightest young stars in the sky.  Situated just north of the Taurus-Auriga complex in the L1482 dark cloud, it appears to be an early B-type star that has been serendipitously exposed during a rarely observed stage of early evolution, revealing a remarkable spectrum and a directly-imaged circumstellar disk. While detailed studies of this star and its circumstellar environment have become increasingly sophisticated in the 50 years since \\citet{herbig56} first pointed it out, the true nature of the object still remains a mystery.  Recent work has renewed focus on the young cluster of stars surrounding LkH$\\alpha$ 101, and what it can tell us about the enigmatic source at its center (e.g., massive star formation timescales, clustered formation mechanisms).  This latter effort certainly deserves more intensive study.  We describe the current knowledge of this region and point out interesting work that could be done in the future. ",
        "introduction": "In a generalized sense, there are three distinct types of young star clusters: ($a$) high-mass star-forming regions with an associated extensive network of low-mass stars (e.g., Orion); ($b$) quiescent environments that host low-mass star formation exclusively (e.g., Taurus-Auriga); and ($c$) smaller clusters of low-mass stars surrounding one or a few A/B stars.  Naturally, there is a continuum of such types, and the picture is not quite so simple.  However, an important goal in this line of research is to generally understand the differences and commonalities between these cluster types in an effort to better explain the various clustered modes of star formation and their consequences.  The LkH$\\alpha$ 101 cluster is an interesting example of the ($c$) type; a handful of B stars and a hundred or more low-mass stars with a dominant source (LkH$\\alpha$ 101) at the center.  The remarkable central source and apparent young age for the cluster indicate that we have been afforded a fortuitous opportunity to investigate this formation mode at a very early time.  In this chapter, we highlight various studies of the LkH$\\alpha$ 101 region, separated into sections focused on the local interstellar medium (Sect.~2), distance estimates (Sect.~3), the embedded young cluster (Sect.~4), and LkH$\\alpha$ 101 itself (Sect.~5).  We conclude with a brief preview of a new, comprehensive multiwavelength study of the region, and summarize the information with an eye toward future studies (Sect.~6).  \\begin{figure}[ht!] \\begin{center} \\includegraphics[height=4.7in,draft=False]{andrews.fig1.eps} \\end{center} \\caption{Optical $VRI$ image of the NGC 1579 reflection nebula that hosts a young cluster surrounding LkH$\\alpha$ 101 \\citep[from][]{herbig04}.  The image is roughly 7 arcminutes on a side, with north up and east to the left.} \\end{figure} ",
        "conclusions": ""
    },
    "0809/0809.2795_arXiv.txt": {
        "abstract": "We present environmental dependence of dusty star forming activity in and around the cluster RXJ1716.4+6708 at $z= 0.81$ based on wide-field and multi-wavelength observations with the Prime Focus Camera on the Subaru Telescope (Suprime-Cam) and the Infrared Camera onboard the AKARI satellite (IRC). Our optical data shows that the optical colour distribution of galaxies starts to dramatically change from blue to red at the medium-density environment such as cluster outskirts, groups and filaments. By combining with the AKARI infrared data, we find that 15~$ \\mu$m-detected galaxies tend to have optical colours between the red sequence and the blue cloud with a tail into the red sequence, consistent with being dusty star forming galaxies.  The spatial distribution of the 15~$\\mu$m-detected galaxies over $\\sim$ 200 arcmin$^2$ around the cluster reveals that few 15~$\\mu$m galaxies are detected in the cluster central region. This is probably due to the low star forming activity in the cluster core. However, interestingly, the fraction of 15~$\\mu$m-detected galaxies in the medium-density environments is as high as in the low-density field, despite the fact that the optical colours start to change in the medium-density environments. Furthermore, we find that 15~$\\mu$m-detected galaxies which have optically red colours (candidates for dusty red galaxies) and galaxies with high specific star formation rates are also concentrated in the medium-density environment. These results imply that the star forming activity in galaxies in groups and filaments is enhanced due to some environmental effects specific to the medium-density environment (e.g. galaxy-galaxy interaction), and such a phenomenon is probably directly connected to the truncation of star forming activity in galaxies seen as the dramatic change in optical colours in such environment. ",
        "introduction": "\\label{sec:intro}  \\subsection{Galaxy properties as a function of environment} Galaxies live in various environments.  Recent redshift surveys have shown filamentary large scale structures in the local Universe. In the distant Universe, at least up to $z \\lsim 1$, similar filamentary nature of large scale structures is found around clusters through wide field observations of distant clusters (e.g.\\ \\citealt{kod05}). There are also some hints that the large scale structure is present at much higher redshifts up to $z\\sim 6$ (\\citealt{shi03}; \\citealt{ouc05}).  Environment must have played an important role in the history of galaxy evolution since galaxy properties are strongly dependent on environment. This is first noted by \\cite{dre80}, who showed that early-type galaxies dominate in high density regions while late-type galaxies tend to live in low density regions. This trend is called ``morphology--density relation''. After the Dressler's work, this interesting trend was confirmed and extended by many authors (e.g. \\citealt{pos84}; \\citealt{dre97}; \\citealt{got03d}; \\citealt{pos05}). However, it is still unclear what is the key physical process to produce the morphology--density relation or other environmental dependence of galaxy properties.  The morphology--density relation can be understood, at least partly, as a result of morphological transformation of a substantial number of galaxies when they entered high-density environments. Many mechanisms to suppress the star forming activity and to contribute to the morphological transformation have been proposed (see the review by \\citealt{bos06}).  For example, ram-pressure stripping due to the interaction with hot plasma gas filled in the cluster core (e.g. \\citealt{gun72}), is expected to be effective in rich cluster cores.  High-speed encounters between galaxies, which are often called ``galaxy harassment'', should occur in very high-density environments (e.g.  \\citealt{moo96}). In addition, mergers or galaxy-galaxy interactions should also contribute to the galaxy transformation (e.g. \\citealt{too72}). Interactions with the cluster potential may also cause a tidal force when galaxies pass the central region of clusters (e.g. \\citealt{byr90}). Another mechanism that can be effective is the so-called ``strangulation'' (e.g. \\citealt{lar80}), which leads to a slow decline in star formation rate after a galaxy falls into a more massive (i.e. group or cluster) halo. Identification of the key processes behind the galaxy transformation is one of the major remaining issues in galaxy evolution.  It is well known that the fraction of blue star forming galaxies in clusters increases towards higher redshifts (Butcher-Oemler effect; \\citealt{but84}). Therefore, the transition from blue active galaxies to red passive galaxies should be more commonly seen in distant clusters. We can expect to see directly such truncation in action through observations of distant clusters (see also \\citealt{got03a}). However, importantly, it is reported that most of the actions take place in the outskirts of clusters rather than in the cluster core. In fact, some recent studies focus on the galaxy properties in the surrounding regions of clusters and try to identify the environment where the truncation of star formation occurs in accreted galaxies (e.g. \\citealt{abr96}; \\citealt{bal99}; \\citealt{pim02}). In such environment, it is reported that passive spirals (i.e.\\ spiral morphology but no star formation) tend to be found (e.g.\\ \\citealt{got03b}). \\cite{kod01} performed a wide-field imaging of the CL0939 cluster at $z=0.41$ and discovered that the colour distribution changes dramatically at the intermediate density environment which corresponds to groups/filaments. A very similar result was reported by \\cite{tan05} for the surrounding regions of higher redshift clusters, CL0016 at $z=0.55$ and RXJ0152 at $z=0.83$. They suggest that the intermediate density environment such as groups or filaments around clusters are the very sites where the truncation of galaxies is taking place. These pioneer works are really telling us the need for wide-field observation of distant clusters in order to study the physical mechanisms that are responsible for the truncation of galaxies from active phase to passive phase during the course of hierarchical assembly of galaxies to clusters.   \\subsection{Dusty star forming galaxies in the local and distant Universe} It is well known that red galaxies tend to have little star forming activity (passive galaxies), while blue galaxies have on-going star-forming activity (star-forming galaxies) at a given redshift. However, the classification of passive or star-forming galaxies based only on their optical colours is sometimes highly uncertain. In fact, \\cite{hai08} showed that $\\sim 30 \\%$ of field red sequence galaxies selected from optical colour--magnitude diagrams have on-going star formation activity with EW(H$\\alpha$) $>$ 2\\AA, using their local galaxy samples from Sloan Digital Sky Survey (SDSS; \\citealt{yor00}). \\cite{dav06} also showed in their SDSS and SWIRE survey (\\citealt{lon03}) that $\\sim 18 \\%$ of their samples have red optical colours and infrared excess at the same time, which include both AGNs and on-going dusty star-forming galaxies. Similar results are shown for cluster red sequence. For example, \\cite{wol05} studied the Abell 901/902 cluster at $z=0.17$, and showed that $\\sim 22 \\%$ of red sequence galaxies are dusty red galaxies from the SED fitting using the medium-band survey (COMBO17; \\citealt{wol01}).  Instead of using optical colours, optical emission lines (e.g. \\oii ($\\lambda = 3727$\\AA) and H$\\alpha $ ($\\lambda = 6563$ \\AA)) are often and widely used to measure star formation rates (\\citealt{ken98}). However, these lines, especially \\oii { } lines in the rest-frame ultra-violet, are attenuated by inter-stellar dust in the galaxies. Also, \\oii { } strength is affected by AGN contribution, if any, and dependent on metallicity as well. Although \\halpha\\, line is a much better indicator than \\oii { } in this respect (and is in fact one of the best indicators of star formation rates among emission lines), it is redshifted to near-infrared (NIR) regime at $z \\gsim 0.5$, where large format wide-field camera or spectrograph has become available only recently. Given these difficulties in optical and NIR observations, infrared (IR) luminosity of a galaxy that can be sensitively obtained by space telescopes, serves as an ideal measure of star formation rate, and it is also well calibrated with local starburst galaxies (\\citealt{ken98}). Since the total IR luminosity of a galaxy can be estimated through single broad-band imaging at mid-infrared (MIR) (e.g. \\citealt{tak05}), it is critically important to observe clusters in MIR bands as well as in the optical/NIR in order to reveal the hidden star formation activity and hence trace its true star formation history. This is especially true at high redshifts, because the luminous infrared galaxies (LIRGs) which have $10^{11} L_{\\odot} \\le L_{\\textrm{IR}} (8-1000\\mu \\textrm{m}) \\le 10^{12} L_{\\odot}$ and ultraluminous infrared galaxies (ULIRGs) which have $L_{\\textrm{IR}} (8-1000\\mu \\textrm{m}) \\ge 10^{12} L_{\\odot}$ are more commonly seen in the distant Universe than in the local Universe (e.g.\\ \\citealt{san96}; \\citealt{lef05}).   \\subsection{Infrared observation of galaxy clusters}  Taking these situations into account, wide-field MIR study of galaxy clusters covering entire structure around the cluster is required to investigate the ``true'' environmental dependence of star formation activities of galaxies. However, until recently, infrared studies of galaxy clusters have been conducted mainly with the ISO satellite (\\citealt{kes96}), which have been limited to very inner regions of clusters (see the review by \\citealt{met05}). Even though, some of these studies showed the importance of large amount of hidden star formation activity in the cluster environment (e.g. \\citealt{duc02}). The recent advent of the Spitzer MIPS (\\citealt{rie04}) has enabled us to investigate wider fields of distant clusters up to $z\\sim 0.8$. \\cite{gea06} observed very wide field (25$' \\times$ 25$'$) around CL0024 ($z=0.39$) and MS0451 ($z=0.55$) clusters by mosaicing MIPS images.  \\cite{mar07} and \\cite{bai07} observed RXJ0152 and MS1054 (both at $z=0.83$), respectively, but their field coverage is limited only to cluster central regions.  \\cite{bai07} actually imply that for rich clusters at $z \\sim 0.8$ even wider-field infrared studies are needed.  In the local Universe, it is well established that star formation activity depends strongly on environment in the sense that it systematically declines towards higher density regions (e.g.\\ \\citealt{gom03}). Recently, however, surprising results are reported where such a relationship between star formation activity and local galaxy density becomes inverted at $z \\sim 1$ (\\citealt{elb07}; \\citealt{coo07}), i.e., star formation rate {\\it increases} towards higher density environment at $z\\sim1$. This is naively expected that one approaches to the formation epoch of cluster galaxies which is probably systematically skewed to higher redshifts compared to the formation epoch of field galaxies. Therefore, looking back the environmental dependence of star formation activity in galaxies as a function of redshift is a basic but vital step towards understanding the environmentally dependent galaxy formation and evolution.  In this paper, we conduct for the first time a panoramic MIR study of a $z \\sim 0.8$ cluster, over a very large area including surrounding groups, filaments and outer fields. Throughout this paper, we use $\\Omega_M =0.3$, $\\Omega_{\\Lambda} =0.7$, and $H_0 =70$ km s$^{-1}$Mpc$^{-1}$. Magnitudes are all given in the AB system, unless otherwise stated. ",
        "conclusions": "\\label{sec:summary} We have performed a wide-field and multi-wavelength optical and infrared study of the distant galaxy cluster RXJ1716.4+6708 (RXJ1716) at $z=0.81$. A unique wide field coverage both in optical and infrared has enabled us to classify galaxies into three environmental bins, namely, high-density regions (cluster core), medium-density regions (cluster outskirts, groups, filament), and low-density regions (field), and has thus allowed us to investigate galaxy properties as a function of environment along the structures in and around the distant cluster.  We find many of the 15~$\\mu$m cluster members show intermediate optical colours between the red sequence and the blue cloud. This may indicate that these 15~$\\mu$m cluster members are actively star-forming galaxies but attenuated by dust hence showing intermediate optical colours, although relatively fewer detection at 15~$\\mu$m of blue galaxies may be partly because they tend to be optically faint.  We quantified the environment around the cluster using the local projected number density of cluster member galaxies, and confirmed that the optical colour distribution starts to dramatically change at the ``medium'' density environment that corresponds to groups and/or filaments.  We showed that the 15~$\\mu$m members are very rare in the high-density cluster centre. This is probably due to the low star forming activity in such regions. However, interestingly, the fraction of the 15~$\\mu$m-detected cluster members in the medium-density environment is as high as in the low-density fields, despite the fact that optical colours of galaxies start to dramatically change from blue to red in the medium-density environment.  Although the statistic is not very good, the fraction is slightly higher in the medium-density environments even compared with the low-density fields.  We also find that dusty red galaxies (optically red 15~$\\mu$m cluster members) and the galaxies with high specific star formation rates (red in $z'-$ 15~$\\mu$m colour) are both concentrated in the medium-density environment. These results may suggest that the star formation activity in galaxies is once enhanced by some physical processes which are effective in group/filament environment (e.g. galaxy-galaxy interaction), before their star forming activity is eventually truncated and they move on to the red sequence.  We stress that all these new findings are based on the widest field MIR observation of a $z \\sim 0.8$ cluster so far. There is no other study that covers such a wide field around clusters in MIR at $z \\gsim 0.8$.  Since our study is a case study for just one cluster at $z = 0.81$, we are desperately in need for a larger sample of distant clusters viewed in the infrared regime and at the same time covering a wide field of view so that we can witness the galaxy truncation in action in the in-fall regions of distant clusters along the filaments. This is essential in order to confirm this interesting trend and to obtain a general view of galaxy evolution."
    },
    "0809/0809.0429_arXiv.txt": {
        "abstract": "We consider an extreme case of disc accretion onto a gravitating centre when the viscosity in the disc is negligible. The angular momentum and the rotational energy of the accreted matter is carried out by a magnetized wind outflowing from the disc. The outflow of matter from the disc occurs due to the Blandford \\& Payne(1982) centrifugal mechanism. The disc is assumed to be cold. Accretion and outflow are connected by the conservation of the energy, mass and the angular momentum. The basic properties of the outflow, angular momentum flux and energy flux per particle in the wind, do not depend on the details of the structure of the accretion disc. In the case of selfsimilar accretion/outflow, the dependence of the rate of accretion $\\dot M$ in the disc depends on the disc radius $r$ on the law $\\dot M \\sim r^{{1\\over2(\\alpha^2-1)}}$, where $\\alpha$ is a dimensionless Alfvenic radius. In the case of $\\alpha \\gg 1$, the accretion in the disc is provided by very weak matter outflow from the disc and  the outflow predominantly occurs from the very central part of the disc. The solution obtained in the work provides mechanism which transforms the gravitational energy of the accreted matter into the energy of the outflowing wind with efficiency close to $100\\%$. The final velocity can essentially exceed Kepler velocity at the site of the wind launch. This mechanism allows us to understand the nature of the astrophysical objects with low luminosity discs and energetic jet-like outflows. ",
        "introduction": "Conventional theory of accretion discs proposed by  \\cite{shakurasunyaev} was successful in interpretation of observations of the accretion discs in the binary systems. This theory satisfactory predicts the general properties of the discs around compact objects. Nevertheless some important phenomena connected with the accretion appeared incompatible with the conventional theory. The most important among them are the accretion discs with anomalously low luminosity and jet-like outflows from the objects. The best example in this regard is AGN M87.  Advection dominated disc model (ADAF model ) was proposed to explain the under luminous discs, in particular disc around BH located in the Galactic Centre \\citep{adafs,narayan}. This theory, however does not solve the problem of jet-like outflow especially in the cases when the kinetic luminosity of jets is comparable (like in the case of SS433 case) or even exceeds ( like in the case of the jet from M87) the bolometric luminosity of the object. In fact, these problems including low luminosity accretion discs, very efficient outflow from the discs and its collimation into jets could be internally connected.  Observations show that jets are strictly connected with the disc accretion. In all jet detections a signature of a disc accretion was found as well. It is important for understanding of the mechanism of the jet ejection that this phenomena is not connected with the specific nature of the central object. Jets are formed irrespectively to the fact that the central object is black hole ( like in the case of jet from AGNs), neutron stars or protostar as it takes place in all cases of the jets from Young stellar objects. It is reasonable to assume that the jets are directly connected with the accretion mechanism itself rather than with the nature of the central object. However, in this case one observational fact needs to be explained. In all cases when it was possible to observe the base of the jets it appears that the jets are launched from the very central part of the disc which produce impression that central object could be connected some way with the process of the jet ejection. At least the fact that jet is ejected from the very central part of the accretion disc rather then from all the disc surface demands explanation.  Another difficult problem is that in some cases the mechanism of ejection is surprisingly efficient. For example, total bolometric luminosity of M87 does not exceed $10^{42} \\rm ergs/s$ \\citep{bolometry}, while the total kinetic luminosity of the jet from M87 is as high as $10^{44} \\rm ergs/s$ \\citep{lkin,lkin2}. Thus, if to estimate the gravitational energy release on the basis of conventional theory of \\cite{shakurasunyaev}, the kinetic energy luminosity of the jet from M87 exceeds by two orders of magnitudes the gravitational energy released at the accretion. The conventional models of the disc accretion (Shakura \\& Sunyaev type or ADAFs) do not predict the existence of such objects. The example with M87 jet shows that one needs to explore new regimes of accretion.   The most evident modification is incorporation into the model of the magnetic field. It has already been widely recognized that the magnetic field has important impact and may play leading role in the ejection of the plasma from accretion discs. In this regard, two processes are of special interest. First one is related the idea of \\cite{blandford}. They have demonstrated that the magnetic field results into instability of the particles at the Kepler orbit. If the angle between the force line of the magnetic field and the disc plane is less than $60^{\\circ}$ the particles are freely ejected from the disc by centrifugal force. The second process was proposed by \\cite{pelletier}. They argue that the winds from the accretion disc can carry out noticeable part of the angular momentum of the accreted material increasing the accretion rate. Together, these work clearly demonstrate that the magnetic field of the disc results into the outflow from the disc at rather general conditions and this wind can carry out essential part of the accreted material angular momentum.   In this work we consider an extreme case when the angular momentum of the disc is carried out by the wind. The viscous stresses are fully neglected. We solve the problem of the disc accretion under these conditions selfconsistently with the problem of the wind outflow from the disc. Fortunately, to provide selfconsistency of the processes of accretion and outflow one do not need to know information about the detailed structure of the accretion disc. The laws of conservation of mass, energy and angular momentum appeared sufficient to provide the selfconsistency of the processes of accretion and outflow.  The paper is organized as follows. In the first section we discuss the qualitative picture of the outflow due to the \\cite{blandford} centrifugal instability and related structure of the poloidal magnetic field. In the second section the basic equations and connection of the accretion and outflow are discussed based on of the conservation laws. In the third section the selfsimilar solutions of the problem are presented. In the last section we discuss the physical sense of the solution. ",
        "conclusions": "The selfconsistent solution of the problem of the disc accretion and plasma outflow from the disc shows that such puzzling properties observed in real astrophysical objects as high energetic efficiency of jets, launching of jets from very central part of the accretion discs and high velocities of outflows which in some objects well exceeds the Keplerian velocities, are naturally explained in frameworks of general approach assuming that the largest part of the angular momentum from the disc is carried out by the wind rather than due to viscous stresses. Although a significant work should still be done to explore the astrophysical implications of the obtained solution to specific astrophysical objects and generalization of the solution for relativistic objects, we believe that this approach is rather promising for explanation of properties of jets from objects of different nature, starting with jets from protostars up to jets from AGN."
    },
    "0809/0809.1377.txt": {
        "abstract": "We present the results from an analysis of the $\\lambda 8727$ forbidden [\\ion{C}{1}] line in high-resolution Gemini-S/bHROS spectra of three Carbon-Enhanced Metal-Poor (CEMP) stars.  Previous derivations of C abundances in CEMP stars primarily have used the blue bands of CH and the C$_2$ Swan system, features which are suspected to be sensitive to photospheric temperature inhomogeneities (the so-called 3D effects).  We find the [C/Fe] ratios based on the [\\ion{C}{1}] abundances of the two most Fe-rich stars in our sample (\\two: $\\mathrm{[Fe/H]} = -1.42$ and \\one: $\\mathrm{[Fe/H]} = -2.66$) to be in good agreement with previously determined values. For the most Fe-deficient star in our sample (\\three: $\\mathrm{[Fe/H]} = -3.08$), however, the [C/Fe] ratio is  found to be 0.34 dex lower than the published molecular-based value.  We have carried out 3D local thermodynamic equilibrium (LTE) calculations for [\\ion{C}{1}], and the resulting corrections are found to be modest for all three stars, suggesting that the discrepancy between the [\\ion{C}{1}] and molecular-based C abundances of \\three\\ is due to more severe 3D effects on the molecular lines.  Carbon abundances are also derived from \\ion{C}{1} high-excitation lines and are found to be 0.45 -- 0.64 dex higher than the [\\ion{C}{1}]-based abundances.  Previously published non-LTE (NLTE) \\ion{C}{1} abundance corrections bring the [\\ion{C}{1}] and \\ion{C}{1} abundances into better agreement; however, targeted NLTE calculations for CEMP stars are clearly needed. We have also derived the abundances of nitrogen, potassium, and iron for each star.  The Fe abundances agree well with previously derived values, and the K abundances are similar to those of C-normal metal-poor stars. Nitrogen abundances have been derived from resolved lines of the CN red system assuming the C abundances derived from the [\\ion{C}{1}] feature. The abundances are found to be approximately 0.44 dex larger than literature values, which have been derived from CN blue bands near 3880 and 4215 {\\AA}. We discuss evidence that suggests that analyses of the CN blue system bands underestimate the N abundances of metal-poor giants. ",
        "introduction": "The discovery that a significant fraction of very metal-poor (VMP; $\\mathrm{[Fe/H]} \\leq -2.0$) stars have highly enhanced abundances of carbon ($\\mathrm{[C/Fe]} \\geq +1.0$) has spurred vigorous research efforts focused on delineating the nucleosynthetic histories of these objects and their role in the chemical evolution of the Galaxy.  The actual fraction of VMP stars that are carbon-enhanced has yet to be precisely determined, but current estimates range from $\\sim 10\\%$ to $\\sim 25\\%$ \\citep{2005ApJ...633L.109C,2005NuPhA.758..312M,2006ApJ...652.1585F,2006ApJ...652L..37L}. This fraction rises at lower metallicities, reaching $\\sim 40\\%$ of stars with $\\mathrm{[Fe/H]} \\leq -3.5$, and stars with $\\mathrm{[Fe/H]} < -4.0$, only three of which are currently known, are all carbon-enhanced.  The increasing incidence of carbon enhancement at lower and lower metallicities implies that the nucleosynthetic pathway(s) leading to these Carbon-Enhanced Metal-Poor (CEMP) stars is highly efficient at low metallicities and that it played an important role in the nucleosynthetic history of the early Galaxy.  Indeed, the fraction of CEMP stars as a function of metallicity may have critical implications for the initial mass function (IMF) in the early Universe \\citep{2007ApJ...665.1361T}.  The manner in which the majority of VMP stars are discovered, and how CEMP stars are subsequently identified, follows a well established procedure\\footnotemark[6].  Stars are first tagged as VMP candidates based on the strength of the \\ion{Ca}{2} K line in low-resolution spectra from objective-prism surveys, in particular the HK survey (Beers, Preston, \\& Shectman 1985, 1992; Beers 1999) and the Hamburg/ESO survey \\citep[HES;][]{2000A&A...358...77W,2003RvMA...16..191C}.  Medium-resolution spectra of the VMP candidates are then obtained and used to identify {\\it bona fide} VMP stars within the sample.  The medium-resolution spectrum of a given VMP star is used to estimate its metallicity ([Fe/H]) by analyzing the strength of the \\ion{Ca}{2} K line as a function of the broadband colors of the star, and inspection of the CH G-band found at 4300 {\\AA} reveals whether or not the star is enhanced in C \\citep[e.g.,][]{2005AJ....130.2804R}.  Estimates of a star's C abundance can be obtained from the CH G-band; however, in the spectra of many CEMP stars, the G-band feature is sufficiently strong as to be saturated, rendering it incapable of providing accurate abundance determinations.  In these cases, C$_2$ lines of the Swan system can be used instead to estimate C abundances.  Once identified, VMP stars are often slated for high-resolution spectroscopic follow-up studies, with which more accurate abundances of Fe, C, and numerous other elements, as well as important ratios such as $^{12}$C/$^{13}$C, can be derived.  \\footnotetext[6]{Here we describe the identification of metal-poor stars based on the low-resolution spectra of objective-prism surveys, the most prolific sources of VMP stars.  For a description of other methods used to find metal-poor stars, please see \\citet{2005ARA&A..43..531B}.}   Analyses of C abundances in high-resolution spectroscopic follow-up studies of CEMP stars focus on the same molecular features as those analyzed in medium-resolution spectra, namely the blue bands of the CH and C$_2$ Swan systems, because it is the molecular features that stay strong and most easily measurable in the spectra of VMP giants.  Another reason is that atomic \\ion{C}{1} lines result from high-excitation transitions, and if measurable, these lines are not believed to be accurate abundance indicators because of their sensitivity to non-local thermodynamic equilibrium (NLTE) effects \\citep{2005ARA&A..43..481A,2006A&A...458..899F}.  However, the abundances derived from the CH and C$_2$ features in the spectra of VMP stars using standard 1-dimensional (1D) local thermodynamic equilibrium (LTE) analyses may not be accurate either.  Abundances derived from molecular features are generally very sensitive to the temperature structure of stellar atmosphere models, and studies of time-dependent 3-dimensional (3D) hydrodynamical models have shown that the molecular lines are susceptible to temperature inhomogeneities due to photospheric granulation, the so-called 3D effects \\citep{2005ARA&A..43..481A}.  The effects on molecular lines are such that abundances based on 3D models are lower than those derived using 1D models, with differences as large as $\\sim 1.0$ dex at $\\mathrm{[Fe/H]} = -3$ (Asplund \\& Garc{\\'i}a P{\\'e}rez 2001; Collet, Asplund, \\& Trampedach 2007).  The derivation of accurate C abundances of CEMP stars is essential to their proper characterization and subsequently to the correct interpretation of their role in the chemical evolution of the Galaxy.  Here we present the results of our investigation into the accuracy of published C abundances as derived from CH and C$_2$ features by analyzing the forbidden [\\ion{C}{1}] line at 8727.13 {\\AA} in high-resolution Gemini-S/bHROS spectra of three CEMP stars: \\object{HE 0054-2542} (\\object{CS 22942-019}), \\object{HE 0507-1653}, and \\object{HE 1005-1439}.  The $\\lambda 8727$ [\\ion{C}{1}] line is known to be immune to NLTE effects and to be a highly reliable abundance indicator \\citep{1999A&A...342..426G}, and it has been used to derive the C abundance of the Sun (e.g., Lambert 1978; Allende Prieto, Lambert, \\& Asplund 2002; Asplund et al. 2005b), field dwarfs \\citep[e.g.,][]{1999A&A...342..426G,2006MNRAS.367.1181B}, R Coronae Borealis stars \\citep[e.g.,][]{2004MNRAS.353..143P}, and Cepheid variables \\citep[e.g.,][]{1981ApJ...245.1018L}.  For CEMP stars, it provides an excellent benchmark for C abundances derived from molecular features. ",
        "conclusions": "An investigation into the accuracy of published C abundances as derived from the traditionally used CH and C$_2$ Swan bands of CEMP stars has been presented. We have analyzed the $\\lambda 8727$ [\\ion{C}{1}] line, a normally reliable C abundance indicator that is impervious to NLTE effects, in high-quality Gemini-S/bHROS spectra of three CEMP stars and have compared the C abundances derived from this line to previously published values.  The 1D LTE abundances derived from the [\\ion{C}{1}] feature are found to confirm the abundances derived from both CH and C$_2$ molecular features in the two most Fe-abundant stars, \\one\\ and \\two\\ ($[\\mathrm{Fe/H}] = -2.66$ and -1.42, respectively), but the [\\ion{C}{1}]-based abundance of the most Fe-deficient star, \\three\\ ($[\\mathrm{Fe/H}] = -3.08$), is 0.34 dex, or about a $2\\sigma$ deviation, lower than the abundance derived from the C$_2$ Swan (0-0) band at 5170 {\\AA}.  As described in $\\S 5.3.2$, the specific C$_2$ Swan band used to derive C abundances does not appear to be a factor in the difference seen for \\three. Also, no particular component of the two abundance analyses could be identified as the most likely source of the discordance, although internal errors may yet be the underlying cause.  Unidentified systematic errors in the analyses also cannot be ruled out as the source of the difference, but the agreement in the abundances of \\two\\ derived by us and \\citet{2007ApJ...655..492A}, the study that also derived the discordant abundance of \\three, suggests that no such systematic errors are to be expected.  Thus, the difference in the C abundance seen between the two studies may indeed be the result of 3D effects in the analysis of the C$_2$ Swan band.  \\citet{2007A&A...469..687C} investigated the impact of 3D effects on CNO abundances of red giants derived from weak CH, NH, and OH lines and found the 3D C abundances to be 0.5 to 0.8 dex lower (for stars with $\\mathrm{[Fe/H]} = -3.0$) than those derived using traditional 1D analyses.  While these corrections are larger than the difference between the C$_2$ and [\\ion{C}{1}]-based abundances for \\three, it must be pointed out that Collet et al. considered stars with {\\it scaled-solar} metallicities, and the corrections may not be applicable to CEMP stars.  Also, the 3D corrections for CH lines are not necessarily similar to those for C$_2$ lines (e.g., Asplund et al. 2005b; Collet et al. 2006).  To investigate the potential 3D effects on [\\ion{C}{1}]-based abundances of low-metallicity stars, we have carried out 3D LTE line formation calculations for [\\ion{C}{1}] using the same 3D hydrodynamical model atmospheres as described by \\citet{2007A&A...469..687C}.  The 3D correction for \\two\\ is estimated to be slightly positive (+0.03 dex), while the corrections for \\one\\ and \\three\\ are negative (-0.07 and -0.15 dex, respectively).  Thus, the corrections are quite modest but do increase in magnitude at lower metallicities, similar to what is seen for the [\\ion{O}{1}] line \\citep{2005ARA&A..43..481A,2007A&A...469..687C}.  While more accurate [\\ion{C}{1}] 3D corrections for these CEMP stars await C-enhanced 3D models and more certain O abundances, the observations and calculations presented here are qualitatively in line with expectations that the 3D corrections for molecular line-based C abundances become more severe towards lower metallicities.  We have also determined abundances from high-excitation \\ion{C}{1} lines and found them to be systematically higher than the [\\ion{C}{1}]-based abundances by $\\sim 0.5$ dex when assuming LTE. One-dimensional abundance corrections have been taken from \\citet{2006A&A...458..899F}, who performed NLTE calculations for a range of stellar parameters assuming a C abundance enhancement of $\\mathrm{[C/Fe]} = +0.4$, and the NLTE-corrected \\ion{C}{1} abundances are in significantly better agreement with the 3D [\\ion{C}{1}]-based values.  However, there remains a $\\sim 0.25$ dex difference in the two C abundances which can be tentatively attributed to underestimated NLTE effects on the \\ion{C}{1} line because calculations for $\\mathrm{[C/Fe]}>+0.4$, which are appropriate for our sample, have not been investigated, or possibly erroneous \\teff\\ values.  At this time, the complicated 3D and NLTE effects that may impact [\\ion{C}{1}], \\ion{C}{1}, CH, and C$_2$ features in the spectra of CEMP stars have not yet been fully established.  Analyses of the [\\ion{C}{1}] line in high-resolution spectra of additional CEMP stars, particularly those at the lowest metallicities, for which CH and/or C$_2$-based abundances exist or are forthcoming are needed to determine if abundances derived from the forbidden line diverge in a systematic way from those derived from the molecular features.  Additionally, 3D models and NLTE line formation calculations with chemical compositions appropriate for CEMP stars are clearly needed to complement the observations.  These measures are necessary so that the most accurate C abundances possible can be derived for CEMP stars."
    },
    "0809/0809.2112_arXiv.txt": {
        "abstract": "% We want to develop spectral diagnostics of stellar populations in the near-infrared (NIR), for unresolved stellar populations. We created a semi-empirical population model and we compare the model output with the observed spectra of a sample of elliptical and bulge-dominated galaxies that have reliable Lick-indices from literature to test if the correlation between Mg2 and CO 1.62 $\\mu$m remains valid in galaxies and to calibrate it as an abundance indicator. We find that (i) there are no significant correlations between any NIR feature and the optical Mg2; (ii) the CaI, NaI and CO trace the $\\alpha$-enhancement; and (iii) the NIR absorption features are not influenced by the galaxy\u2019s age. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2781_arXiv.txt": {
        "abstract": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  We present Keck/DEIMOS spectroscopy of Segue\\,1, an ultra-low luminosity ($M_V = -1.5^{+0.6}_{-0.8}$) Milky Way satellite companion. While the combined size and luminosity of Segue\\,1 are consistent with either a globular cluster or a dwarf galaxy, we present spectroscopic evidence that this object is a dark matter-dominated dwarf galaxy.  We identify \\nmembers\\ stars as members of Segue\\,1 with a mean heliocentric recession velocity of $206 \\pm 1.3$\\kms. We measure an internal velocity dispersion of $4.3\\pm 1.2$\\kms.  Under the assumption that these stars are gravitationally bound and in dynamical equilibrium, we infer a total mass of $4.5^{+4.7}_{-2.5} \\times 10^5 M_{\\odot}$ in the case where mass-follow-light; using a two-component maximum likelihood model, we determine a similar mass within the stellar radius of 50\\,pc.  This implies a mass-to-light ratio of ln$(M/L_V) = 7.2^{+1.1}_{-1.2}$ or $M/L_V = 1320^{+2680}_{-940}$.  The error distribution of the mass-to-light ratio is nearly log-normal, thus Segue\\,1 is dark matter-dominated at a high significance.  Although Segue\\,1 spatially overlaps the leading arm of the Sagittarius stream, its velocity is 100\\kms\\ different than that predicted for recent Sagittarius tidal debris at this position.  We cannot rule out the possibility that Segue\\,1 has been tidally disrupted, but do not find kinematic evidence supporting tidal effects.  Using spectral synthesis modeling, we derive a metallicity for the single red giant branch star in our sample of [Fe/H] $= -3.3\\pm0.2$\\,dex.  Finally, we discuss the prospects for detecting gamma-rays from annihilation of dark matter particles and show that Segue\\,1 is the most promising satellite for indirect dark matter detection.  We conclude that Segue\\,1 is the least luminous of the ultra-faint galaxies recently discovered around the Milky Way, and is thus the least luminous known galaxy. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  The discovery of ``ultra-faint'' dwarf spheroidal (dSph) galaxies around the Milky Way has revolutionized our understanding of dwarf galaxies and their prevalence in the Universe.  These newly discovered satellites, with total absolute magnitudes fainter than $M_V = -8$, have all been found in the Sloan Digital Sky Survey (SDSS) via slight statistical over-densities of individual stars \\citep{willman05b,willman05a,zucker06a,zucker06b, belokurov06a,belokurov06b,sakamoto06a, irwin07a, walsh07a}.  These objects provide important clues to galaxy formation on the smallest scales \\citep{madau08a,ricotti08a} and substantially alleviate the discrepancy between the observed mass function of Milky Way satellites and that predicted by standard Lambda Cold Dark Matter models \\citep[][hereafter SG07]{tollerud08a,simon07a}.  \\citet{strigari07b} note that the ultra-faint dSphs have high central dark matter densities and are good candidates for indirect dark matter detection via gamma-ray emission by particle annihilation.  Future wide-field surveys that improve on the sky coverage and photometric depth of the SDSS are likely to discover many additional ultra-faint Milky Way satellites in the coming years \\citep{koposov07a,walsh08b}.  While the total luminosities of the ultra-faint satellites are comparable to globular clusters, spectroscopic studies for the majority of the newly discovered objects firmly suggest that these objects are dark matter-dominated dwarf galaxies \\citep[][SG07]{kleyna05a,munoz06a,martin07a}. The mass-to-light ratios for all the ultra-faint dSphs are $M/L_V > 100$~M$_{\\odot}$/L$_{\\odot}$, with several systems approaching 1000~M$_{\\odot}$/L$_{\\odot}$, assuming mass-follows-light. \\citet{strigari08a} loosened this constraint, confirming the high mass-to-light ratios and finding a tight anti-correlation between mass-to-light ratio and luminosity such that all the Milky Way dwarfs are consistent with having a common dark matter mass of $\\sim10^7\\Msun$ within their central 300\\,pc. A theoretical understanding of the physics that sets the mass-luminosity relation will provide insight into the formation of galaxies at the smallest scales.  Further evidence that the ultra-faint satellites are indeed galaxies comes from metallicity measurements.  The ultra-faint satellites are the most metal-poor known stellar systems ([Fe/H] $< -2$) and show internal metallicity spreads up to 0.5~dex in several objects (SG07). This is in contrast to Milky Way globular clusters which are, on average, more metal-rich and show little to no internal metallicity spread \\citep[e.g.][]{Pritzl05a}.  In further contrast to globular clusters, the ultra-faint dwarfs also follow the luminosity-metallicity relationship established by brighter Milky Way dwarf galaxies \\citep{kirby08b}.  Thus, both the kinematics and composition of the ultra-faint satellites strongly argue that these objects are dark matter-dominated galaxies.  The combined size and luminosity of the spectroscopically confirmed dSphs in the Milky Way are well separated from globular clusters: at a given luminosity dwarf galaxies have larger sizes and are thus less compact \\citep{belokurov06b,martin08a}.  However, the three faintest SDSS discoveries, Segue\\,1, Willman~1 and Bootes~II, are all in a region that overlaps with globular clusters.  Studying these extreme systems should provide important insight to dSphs, and the difference between dwarfs and star clusters, at all luminosities.  Of these three objects, only Willman~1 has published kinematics \\citep{martin07a}. Because the systemic velocity of Willman~1 is similar to that of the foreground Milky Way stars, possible contamination in the kinematic sample make it difficult to assess whether this object is a dwarf or globular cluster \\citep[][Willman et al.~in prep]{siegel08a}.  Here, we present the first spectroscopic study of an even lower luminosity system, Segue\\,1.  The systemic velocity of Segue\\,1 is far removed from the Milky Way foreground and thus should be a cleaner object to study the properties of the least luminous ultra-faint systems.  Segue\\,1 was discovered by \\citet{belokurov06b} as an over-density of resolved stars in the SDSS located at ($\\alpha_{2000}, \\delta_{2000})$ = (10:07:03, +16:04:25) = $(151.763^{\\circ}, 16.074^{\\circ}$).  Via isochrone fitting, these authors estimate a distance of $23\\pm2$\\,kpc and an absolute luminosity of $M_V \\sim -3\\pm0.6$. \\citet{martin08a} recently revised the luminosity of Segue\\,1 to $M_V = -1.5^{+0.6}_{-0.8}$ using a more robust method to estimate flux in systems with small numbers of observable stars. While the possibility of tidal tails and/or tidal distortion of Segue\\,1 was found in the initial SDSS analysis, deeper imaging and more thorough simulations suggest that these features can be explained via Poisson scatter of the few bright stars in this system \\citep{martin08a}.  Segue\\,1 has no detected gas content, with an observed HI gas mass limit of less than $13\\Msun$ \\citep{putman08a}. This limit is consistent with other dSphs around the Milky Way in which any gas has been presumably removed via ram pressure stripping or used up via tidally-induced star formation \\citep{mayer06a}.   \\citet{belokurov06b} note that Segue\\,1 is spatially superimposed on the leading arm of the Sagittarius stream.  Because it has a similar luminosity and size as the most diffuse globular cluster, they proposed that Segue\\,1 is a globular cluster formerly associated with the Sagittarius dSph.  Spectroscopy of member stars in Segue\\,1 is required to test this hypothesis and answer the crucial question of whether or not this intrinsically faint stellar system is truly a globular cluster (i.e.~a stellar system with a single stellar population with no dark matter).  Here, we present Keck/DEIMOS multi-object spectroscopy for individual stars in the vicinity of Segue\\,1, identifying \\nmembers\\ stars as members of Segue\\,1.  This paper is organized as follows: in \\S\\,\\ref{sec_data} we discuss target selection and data reduction for our Keck/DEIMOS spectroscopy. In \\S\\,\\ref{sec_kin} we discuss the spectroscopic results including estimates of the velocity dispersion, mass, mass-to-light ratio and metallicity.  In \\S\\,\\ref{sec_sgr}, we examine the spatial and kinematic position of Segue\\,1 relative to the Sagittarius stream. In \\S\\,\\ref{sec_gamma} we note that Segue\\,1 may be a good target for indirect detection of dark matter.  Finally, in \\S\\,\\ref{sec_disc}, we discuss Segue\\,1 in context of the Milky Way dSph population.  Throughout the analysis, we use the photometric properties of Segue\\,1 as derived by \\citet{martin08a} of $M_V = -1.5^{+0.6}_{-0.8}$ (i.e.\\ the 1$\\sigma$ magnitude limits are $M_V = -0.9$ and $-2.3$) and $r_{\\rm eff}=4.4'^{+1.2}_{-0.6} = 29^{+8}_{-5}$\\,pc.  We also assume a fixed reddening to Segue\\,1 based on the \\citet{schlegel98a} value of E(B-V) = 0.032~mag.  We list these and other key parameters in Table~1.  \\begin{figure*}[t!] \\epsscale{1.0} \\plotone{fig1.eps} \\caption{{\\it Left:\\/} Color-magnitude diagram of all stars (small black points) within $30'$ of the center of Segue\\,1 from SDSS DR~6 $g$- and $r$-band photometry.  The larger symbols indicate stars with measured Keck/DEIMOS velocities: solid blue circles fulfill our requirements for membership in Segue\\,1, red asterisks are higher velocity stars and open squares are foreground Milky Way stars.  Two fiducial isochrone are shown shifted to the distance of Segue\\,1: M92 ([Fe/H] =$-2.3$, solid line) and M3 ([Fe/H]=$-1.6$, dashed line).  {\\it Right:\\/} Spatial distribution of stars near Segue\\,1.  The solid ellipse is the half-light radius of Segue\\,1 as measured by \\citet{martin08a}.\\label{fig_cmd}} \\end{figure*}    \\begin{figure*}[t!] \\plotone{fig2.eps} \\caption{{\\it Left:\\/} Keck/DEIMOS velocity histogram for all stars in our sample; velocities are corrected to the heliocentric frame.  We identify Segue\\,1 as the velocity peak near $v=206$\\kms.  Stars with less positive velocities are identified as foreground Milky Way, the four stars with $v\\sim300$\\kms\\ are tentatively associated with the Sagittarius stream as discussed in \\S\\,\\ref{ssec_highv}.  {\\it Right:\\/} Radial distance from the center of Segue\\,1 plotted against heliocentric velocity.  Stars to the East of the galaxy center are plotted as triangles, stars to the West are plotted as squares.  We indicate the effective half-light radius ($r_{\\rm eff}$), the mean systemic velocity of the system (black dashed line) and velocity dispersion (grey shaded region). \\label{fig_velocity}} \\end{figure*} ",
        "conclusions": "\\label{sec_disc}  As seen in Figure~\\ref{fig_corr}, Segue\\,1 lies on an extension of the luminosity-metallicity and luminosity-mass relationships established by brighter Milky Way dSphs.  While the dSphs span nearly five orders of magnitude in luminosity, their mass enclosed within 300\\,pc remains nearly constant at $10^7$\\Msun \\citep{strigari08a}. This common mass scale has been noted in previous studies \\citep{mateo93a, gilmore07a}, but remains a very surprising result given the much larger luminosity range spanned by the present data.  It strongly suggests the existence of a characteristic scale in either galaxy formation processes or dark matter physics.  At the same time, the average metallicities of the dSphs are correlated with luminosity such that stars in the least luminous dSph are the most metal-poor \\citep{kirby08b}. Segue\\,1 is at the extreme end of these relationships: its luminosity is merely $L = 340$\\Lsun, yet its total mass enclosed within 300\\,pc is $10^7$\\Msun (projecting the mass model discussed in \\S\\,\\ref{ssec_mass}), resulting in the highest $M/L_V$ ratio of any known stellar system. The metallicity for the single RGB star in Segue\\,1 is [Fe/H] = $-3.3$\\,dex, one of the most metal-poor stars known in a dSph galaxy. This metallicity is slightly less than that predicted by the Kirby et al.~log-linear relationship, however, we note that the average galactic metallicity may be higher than this single star.   The correlations in Figure~\\ref{fig_corr} are the key to understanding how dSphs form.  While several formation avenues exist to modify the mass-to-light ratio of dSphs, the added constraint of the luminosity-metallicity correlation reduces the number of allowable models.  This correlation rules out a tidal stripping scenario in which lower luminosity systems initially form as more luminous galaxies outside the environment of the Milky Way and are then tidally stripped to their present state as they enter the Milky Way environs.  In this scenario, the metallicity of stars would not be tied to the present luminosity. While ruling out formation scenarios is certainly progress, determining what formation processes can explain the observed correlations will be more challenging \\citep[e.g.][]{bovill08a}.  A key question raised by the Segue\\,1 results is why the Milky Way dwarf dSphs have such remarkably different luminosities, yet appear to have similar total masses.  Why do all these objects have a common mass halo and is this consistent with the mass spectrum of dark matter halos predicted by simulations?  Explaining the mechanism that sets both the mass-luminosity and luminosity-metallicity relationships in the Milky Way will provide insight to the formation of galaxies at all scales."
    },
    "0809/0809.4444_arXiv.txt": {
        "abstract": "Dynamical properties of two-component galaxy models whose stellar density distribution is described by a $\\gamma$-model while the total density distribution has a pure $r^{-2}$ profile, are presented. The orbital structure of the stellar component is described by Osipkov--Merritt anisotropy, while the dark matter halo is isotropic.  After a description of minimum halo models, the positivity of the phase-space density (the model consistency) is investigated, and necessary and sufficient conditions for consistency are obtained analytically as a function of the stellar inner density slope $\\gamma$ and anisotropy radius. The explicit phase-space distribution function is recovered for integer values of $\\gamma$, and it is shown that while models with $\\gamma>4/17$ are consistent when the anisotropy radius is larger than a critical value (dependent on $\\gamma$), the $\\gamma=0$ models are unphysical even in the fully isotropic case.  The Jeans equations for the stellar component are then solved analytically; in addition, the projected velocity dispersion at the center and at large radii are also obtained analytically for generic values of the anisotropy radius, and it is found that they are given by remarkably simple expressions.  The presented models, even though highly idealized, can be useful as starting point for more advanced modeling of the mass distribution of elliptical galaxies in studies combining stellar dynamics and gravitational lensing. ",
        "introduction": "Analysis of stellar kinematics (e.g. Bertin et al. 1994, Rix et al. 1997, Gerhard et al. 2001), as well as several studies combining stellar dynamics and gravitational lensing strongly support the idea that the dark and the stellar matter in elliptical galaxies are distributed so that their {\\it total} mass profile is described by a density distribution proportional to $r^{-2}$ (e.g., see Treu \\& Koopmans 2002, 2004; Rusin et al. 2003; Rusin \\& Kochanek 2005; Koopmans et al. 2006; Czoske et al. 2008; Dye et al. 2008). In particular, Gavazzi et al. (2007), with a gravitational lensing analysis of 22 early-type strong lens galaxies, reported a total $r^{-2}$ density profile in the range 1-100 effective radii.  It is clear that in this field the availability of simple dynamical models of two-component galaxies can be useful as starting point of more sophisticated investigations based on axysimmetric or triaxial galaxy models (e.g., Cappellari et al. 2007, van den Bosch et al. 2008).  A few simple yet interesting models with flat rotation curve have been in fact constructed, such as those in which the stellar mass was described by a power-law in a total $r^{-2}$ mass distribution (e.g. Kochaneck 1994), or those obtained from physical arguments (in case of disk galaxies, see e.g. Naab \\& Ostriker 2007).  Here the family of two-component galaxy models whose {\\it total} mass density is proportional to $r^{-2}$, while the visible (stellar) mass is described by the well-known $\\gamma$ models (Dehnen 1993, Tremaine et al. 1994), is presented.  Some preliminary numerical investigation of these models has been done in Keeton (2001), and they have been used in Nipoti et al. (2008) as diagnostics of the total mass distribution in elliptical galaxies.  In this paper a more systematic study of the dynamical properties of these models is presented.  It is shown that the Jeans equations for the stellar component with Osipkov-Merritt (Osipkov 1979, Merritt 1985, hereafter OM) radial anisotropy can be solved analytically.  Remarkably, the projected velocity dispersion at the center and at large radii can be expressed in terms of the model circular velocity by means of extremely simple formulae for generic values of the anisotropy radius and of the central stellar density slope $\\gamma$.  In principle this feature opens the possibility to obtain preliminary indications about the anisotropy from observations at small and large radii.  The positivity of the phase-space density (the so-called consistency) is investigated, by obtaining analytically the necessary and sufficient conditions for model consistency in terms of $\\gamma$, of the anisotropy radius, and of the dark-to-stellar mass ratio within some prescribed radius.  It is found that the phase-space distribution function (hereafter DF) can be recovered analytically for $\\gamma=0,1,$ and $2$.  In particular, it is shown that $\\gamma =0$ models in a total $r^{-2}$ density profile are unphysical for any value of the anisotropy radius.  These results extend the class of two-component galaxy models with explicit DF and add to the large amount of phase-space information already available about one and two-component $\\gamma$ models (e.g., see Dehnen 1993, Tremaine et al. 1994, Hiotelis 1994, Carollo et al. 1995, Ciotti 1996, 1999; Baes et al. 2005, Buyle et al. 2007, Ciotti \\& Morganti 2008).  The paper is organized as follows. In Section 2 the main structural properties of the models are presented, while in Section 3 an investigation of the phase-space properties of the models is carried out both from the point of view of necessary and sufficient conditions for consistency and of direct recovery of the DF in specific cases. In Section 4 the solution of the Jeans equation with OM radial anisotropy is presented, together with their projection at small and large radii. Finally a short summary of possible use of the present models in observational works is given. ",
        "conclusions": "In this paper a family of spherical, two-component galaxy models with a stellar density profile described by a $\\gamma$ model and a total (stars plus dark matter) density profile $\\propto r^{-2}$ at all radii has been investigated, under the assumption that the internal dynamics of the stellar component is described by Osipkov-Merritt anisotropy. The models are fully determined when the inner density slope $\\gamma$ and the anisotropy radius $\\ra$ of the stellar component are assigned, together with a density scale for the total density profile.  The dark matter halo remains defined as the difference between the total and stellar density profiles. The main results can be summarized as follows.  \\begin{itemize}  \\item After having provided the most common structural quantities of the models that are of interest for observations, limitations on the {\\it total} density scale as a function of $\\gamma$ are analytically determined by requiring the positivity and monotonicity of the dark matter halo distribution.  In particular, the request of positivity limits the range of acceptable stellar density slopes to $0\\leq\\gamma\\leq 2$.  Models corresponding to the minimum total density scale (for given $\\gamma$) are called minimum halo models. The central density profile of the dark matter halo diverges as $r^{-2}$ in general, but in the minimum halo $\\gamma=2$ model (in which the positivity and monotonicity limits coincide), the central dark matter profile is $\\propto r^{-1}$.  \\item The minimum value of anisotropy radius corresponding to a dynamically consistent stellar component has been derived analytically as a function of $\\gamma$ by using the necessary and sufficient conditions of CP92.  As expected, an increase of $\\gamma$ results in a decrease of the minimum value of the anisotropy radius $\\ra$ required by consistency.  It is also proved that models with $\\gamma>4/17$ are certainly consistent for sufficiently isotropic velocity dispersion, while for centrally shallower stellar density profiles the sufficient condition for consistency is never satisfied.  The necessary and sufficient condition for the halo consistency are also analytically obtained, and the minimum halo models corresponding to (isotropic) dark matter halos are derived. In the case $\\gamma=2$, the minimum halo coincides with the minimum halo obtained from the positivity and monotonicity conditions.  \\item The phase--space DF of the stellar component for $\\gamma=0,1,2$ is analytically recovered in terms of Polylogarithms and exponentials.  It is found that the $\\gamma=0$ model is inconsistent no matter how much anisotropy is considered. Instead, the isotropic $\\gamma=1,2$ models have a positive DF, and the true critical anisotropy radius for consistency can be determined directly from their DF.  A comparison with the analogous study of one-component $\\gamma$ models shows that the presence of the halo sligthly reduces the maximum amount of sustainable radial anisotropy. The obtained values of the anisotropy radius are independent of the total density scale $\\MR$.  \\item The Jeans equations for the stellar component are solved explicitely for generic values of $\\gamma$ and $\\ra$ in terms of elementary functions.  The asymptotic expansions of $\\srad$ for $r\\to0$ and $r\\to\\infty$ are obtained, and it is shown that $\\srad$ tends to finite (non-zero) values (except for the divergent central velocity dispersion of $\\gamma=0$ model) which are simply related to the model circular velocity.  The projected velocity dispersion $\\sigp(0)$ cannot be calculated analytically, in general. However, by asymptotic expansion of the projection integral, exact values at large radii and at the center are obtained.  In particular, it is shown that for $\\gamma\\geq 1$, and independently of the value of the anisotropy radius, $\\sigp(0)$ coincides with the central velocity dispersion $\\srad(0)$ in the isotropic case.  Instead, for $0\\leq\\gamma<1$, $\\sigp(0)$ depends also on $\\sa$.  In the anisotropic case $\\sigp(0)$ cannot be obtained analytically; however, a very simple form of the projection integral suitable for numerical integrations is given. In the isotropic case the integral can be evaluated analytically, and $\\sigp(0)$, as expected, coincides with the model circular velocity.  \\item Finally, we have shown that the Velocity Profile of the models can be obtained in a very simple form (Gaussian) near the center (independently of the value of the anisotropy radius and for $1\\leq\\gamma\\leq 2$), and at large radii (in the isotropic case).  \\end{itemize}  We conclude by noting that these models, albeit highly idealized, seem to suggest two interesting remarks of observational character.  The first is that in real galaxies with a total $r^{-2}$ density profile and sufficiently peaked stellar density (i.e. $\\gamma\\geq1$), measures of central velocity dispersion should not be strongly affected by radial anisotropy.  Second, for given central stellar density slope $\\gamma$, measures of the projected velocity dispersion at the center and in the external regions are able, at least in principle, to determine the value of the anisotropy radius under the assumption of Osipkov-Merritt anisotropy."
    },
    "0809/0809.1866_arXiv.txt": {
        "abstract": "We calculate the effects of frame dragging on the Galactic-Center stars. Assuming the stars are only slightly relativistic, we derive an approximation to the Kerr metric, which turns out to be a weak field Schwarzschild metric plus a frame dragging term. By numerically integrating the resulting geodesic equations, we compute the effect on keplerian elements and the kinematics. We find that the kinematic effect at pericenter passage is proportional to $(a(1-e^2))^{-2}$. For known Galactic-center stars it is of order 10 m/s. If observed this would provide a measurement of the spin of the black hole. ",
        "introduction": "The center of the Milky way is a very interesting region. It contains a massive black hole (MBH) of $ \\sim 3 \\times 10^6 M_\\odot$. The central parsec contains thousands of stars . For a small but a growing number of these, due to the relatively close proximity of the MBH and short orbital periods, the orbital parameters have been accurately measured \\citep{2003ApJ...596.1015S, 2005ApJ...620..744G, 2005ApJ...628..246E} . Some of the stars have pericenter velocities as high as a few percent of $c$. Hence, as shown in \\cite{2006ApJ...639L..21Z} the general relativistic effect of O($\\beta^2$) should be observable.  At $O(\\beta^3)$ general relativity predicts a new effect, which is that a spinning black hole drags the surrounding space-time along with it. Under the rotational frame-dragging effect (also known as Lense-Thirring effect), the frame of reference with minimal time dilation is one which is rotating around the object as viewed by a distant observer. If this effect could be observed for GC stars, then in principle the spin of the MBH can be measured. A method based on the orbital dynamics of the GC stars is more direct than the usual approach which requires modeling the effect of spin on the accretion disk \\citep[see for example][]{2008AIPC..968..265N}.  The effect of $O(\\beta^3$) terms on the keplerian elements and on astrometry were discussed by \\cite{1998AcA....48..653J} and \\cite{2000ApJ...542..328F}. \\cite{2008ApJ...674L..25W} goes on to consider $O(\\beta^4$) as well.  The resulting astrometric effects are so small that they can only be observed on stars which are closer to the MBH than the observed GC stars.  In this paper we concentrate on the effects of the $O(\\beta^3)$ terms on the kinematics of the GC stars.  Traditionally the effect of relativistic perturbations on orbital dynamics have been studied either using post-newtonian celestial mechanics \\citep{1972gcpa.book.....W} or pseudo-newtonian equations \\citep{1999A&A...343..325S}. We adopt a different and conceptually simpler approach. We do a low velocity perturbative expansion of the Kerr metric and then numerically integrate the resulting geodesic equations. ",
        "conclusions": "We see that the maximum kinematic effect in known GC stars $\\Delta V\\fd$ is of the order of a few 10's of m/s during the few weeks around the pericenter passage. Although this level of accuracy is difficult to achieve for GC stars, it is not implausible. Extrasolar planet searches regularly reach an accuracy better than 1 m/s \\citep{2006Natur.441..305L} and new technologies for radial velocity measurements may be able to obtain precision as high as 1 cm/s \\citep{2008Natur.452..610L}.  There are also two theoretical problems which remain to be solved.  First, an accurate calculation of the redshift as a function of time is required (as it is the observable quantity), rather than velocity as a function of time as calculated here.  Both the kinematic and gravitational redshifts are involved.  We do not know of any approximate method for calculating the redshift in this case, as our approximate metric is not valid for light.  It may be necessary to calculate null geodesics in the full Kerr metric.  Second, the relativistic effects have to be separated from the newtonian effects of other masses, such as nearby stars, gas and dark-matter clouds.  These could overshadow the frame-dragging contribution, especially since some newtonian dynamical processes in the GC region can be unexpectedly strong because of resonances \\citep{2007MNRAS.379.1083G,2008ApJ...683L.151L}.  However, newtonian perturbations from other masses would not give the distinctive time dependence in the kinematics that frame-dragging does (Fig.~\\ref{kincomp}). Hence, we can be optimistic about dis-entangling frame-dragging from all the newtonian effects.  \\newpage  \\appendix"
    },
    "0809/0809.1782_arXiv.txt": {
        "abstract": "We present an analysis of optical and ultraviolet {\\it Hubble Space Telescope} photometry for evolved stars in the core of the distant massive globular cluster NGC 2419. We characterize the horizontal branch (HB) population in detail including corrections for incompleteness on the long blue tail. The majority of the horizontal branch stars can be identified with two main groups (one slightly bluer than the instability strip, and the other at the extreme end of the HB).  We present a method for removing (to first order) lifetime effects from the distribution of HB stars to facilitate more accurate measurements of helium abundance for clusters with blue HBs and to clarify the distribution of stars reaching the zero-age HB. The population ratio $R = N_{HB} / N_{RGB}$ implies there may be slight helium enrichment among the EHB stars in the cluster, but that it is likely to be small ($\\Delta Y < 0.05$).  An examination of the upper main sequence does not reveal any sign of multiple populations indicative of helium enrichment.  The stellar distribution allows us to follow how the two main types of stars evolve after the HB.  We find that the transition from stars that reach the asymptotic giant branch to stars that remain at high temperatures probably occurs among the extreme horizontal branch stars (EHB) at a larger temperature than predicted by canonical evolution models, but qualitatively consistent with helium-enriched models.  Through comparisons of optical CMDs, we present evidence that the EHB clump in NGC 2419 contains the end of the canonical horizontal branch, and that the boundary between the normal HB stars and blue hook stars shows up as a change in the density of stars in the CMD. This corresponds to a spectroscopically-verified gap in NGC 2808 and an ``edge'' in $\\omega$ Cen.  The more clearly visible HB gap at $V \\sim 23.5$ identified by \\citet{rip} appears to be too bright.  Once corrected for lifetime effects, we find that NGC 2419 is currently converting about 25 -- 31\\% of the first-ascent red giant stars in its core into extreme blue horizontal branch stars --- the largest fraction for any known globular cluster.  A comparison of upper red giant branch with theoretical models indicates there is a slight deficiency of bright red giant stars.  This deficiency occurs far enough below the tip of the red giant branch that it is unlikely to be associated with the production of extreme horizontal branch stars via strong mass loss before the core helium flash. ",
        "introduction": "NGC 2419 is the fourth brightest globular cluster known in the Milky Way ($M_V = -9.58$; \\citealt{har96}).  NGC 2419's large mass would normally make it ideal for studying short phases in the lives of evolved low-mass stars because massive cluster (having more stars) should produce a greater chance of catching {\\it some} stars in those phases. Until recently though, it has been neglected because it resides far in the outer halo of the Milky Way (approximately 90 kpc from the center). Massive clusters can make it possible to observe stars in short-lived evolution phases because there is a greater chance of catching {\\it some} star in such a phase when there are more chances.  Evolved stars have unusually strong influences on the luminosity-integrated properties of a galaxy: in old stellar populations, the overall color and the ultraviolet emission result from them.  The luminous giants at the tip of the red giant branch (TRGB) are frequently used as standard candles in resolved stellar populations. Extremely blue HB (hereafter, EHB) stars are a leading candidate for galactic ultraviolet emission.  These populations are linked at least in an evolutionary sense, but it remains a longstanding problem to explain the details of how HB stars (and more specifically the EHB stars) are produced. There has been much recent interest in multiple stellar populations having different chemical compositions (particularly helium, but also heavier elements) within individual clusters (NGC 2808, \\citealt{dan2808}; $\\omega$ Cen, \\citealt{bedomega}; M13, \\citealt{cadm13}; NGC 6218, \\citealt{cn6218}; NGC 6388, \\citealt{busso}; NGC 6441, \\citealt{cadn6441}) However, even allowing for composition effects, it still seems to be necessary for giant stars to lose most of the mass in their envelopes if EHB stars are to be produced --- the turnoff mass of a helium-enriched population cannot be reduced enough to have a giant star consume its envelope before it reaches the TRGB.  In this paper, we examine deep photometry of the evolved stars in the core of the cluster using high resolution {\\it Hubble Space Telescope (HST)} imagery. We have three main goals: to characterize brief evolutionary phases and their effects on the integrated light of the cluster, to examine the paths (often brief) that lead from one evolutionary phase to another, and to use the stellar populations to constrain chemical evolution within the cluster. EHB and bright red giant branch (RGB) stars are particularly important to the integrated colors of clusters.  EHB stars can be seen in several earlier CMDs for NGC 2419 \\citep{har97,stetn2419}, although the impression was that the population was small due to incompleteness at the faint end. \\citet{dale} recently presented a reduction of a portion of the available archived {\\it HST} data for the cluster that clearly shows the faint end of the blue HB tail, and that the EHB population is a substantial fraction of all HB stars in the cluster. In section \\ref{HB}, we conduct a deeper analysis of the HB population of the cluster.  The details of the transition from the RGB to the HB remain poorly understood. In the case of the massive cluster NGC 2808, \\citet{sm07} found a deficit of bright red giants compared to theoretical predictions that may be linked to the EHB stars there --- if a star loses enough mass on the red giant, it can leave the red giant branch before the helium flash at the TRGB, igniting helium later in a ``hot flash'' that ultimately deposits it at the blue end of the HB. NGC 2419 appears to be more efficient at producing EHB stars than most clusters, so in \\S \\ref{urgb} we examine the upper RGB for signatures of EHB star production.  Integrated luminosity has been linked to a number of unusual groups of stars within globular clusters, including blue stragglers (their relative frequency is anti-correlated with total luminosity), extremely blue horizontal branch stars, and populations with different age and/or chemical composition. To date, NGC 2419 has shown little indication that it contains anything but a stellar population with a single composition ([Fe/H] $= -2.1$; \\citealt{sunt}) and age. However, NGC 2419 is now known to have a strongly double-peaked HB, and so we examine the helium abundance indicator $R$ in \\S \\ref{R}. ",
        "conclusions": "We have examined archival HST imagery in the optical and ultraviolet to study evolved stars in the core of the massive globular cluster NGC 2419.  An important aspect of these populations is the clear separation of the horizontal branch population into two groups: the primary peak just to the blue of the instability strip and a secondary peak of comparable size at the extreme end of the HB. More than 38\\% of the HB stars in this cluster are in the EHB peak, and if their initial compositions were the same, this means that more than 33\\% of its red giants have been diverted into this population (more than 25\\% into blue hook candidates). While our study only examined the core population (and in $\\omega$ Cen, for example, there is some evidence for a decrease in the abundance of EHB stars with distance from the cluster center), this is a feat that is even remotely matched by only {\\it some} of the very most massive clusters in the Galaxy.  For example, $\\omega$ Cen has approximately 27\\% of its HB stars on the EHB \\citep{dale} --- a slightly smaller proportion than NGC 2419. M54 also has a modest population of EHB stars \\citep{rosem54,siem54}, but it is a significantly smaller fraction of the total HB population (we estimate $\\la 15$\\% from the CMD in \\citealt{rosem54}). NGC 2808 also has a modest EHB population that is 12\\% of its total HB sample (\\citealt{castel}, where we have included their EBT3 stars and the HBp stars that are likely to be composed of blends of EBT3 and MS stars). In the very metal-rich cluster NGC 6388 \\citep{dale08}, only about 3.5\\% of the HB stars are on the EHB (2\\% as blue hook stars).  A few other massive clusters (e.g. NGC 6441, M15) have a handful of EHB stars that compose a few percent of their populations at most, and another massive cluster (47 Tuc) has no known EHB stars. $\\omega$ Cen, M54, NGC 6388, NGC 6441, and 47 Tuc are probably the only clusters that are more massive than NGC 2419 \\citep{mcl}, while NGC 2808 appears to have a similar mass and M15 is approximately 50\\% less massive. So while large cluster mass certainly makes the appearance of EHB stars more likely, another factor is influencing their relative abundance. This factor is probably not {\\it past} cluster mass --- NGC 2419 is unlikely to have lost as much mass in its history as clusters like $\\omega$ Cen and M54 because its orbit keeps it far out in the Galactic halo, minimizing the effects of Galactic tides and disk shocking. There is also no evidence of populations of greatly different age in NGC 2419 based on a superficial examination --- NGC 1851 shows two distinct subgiant branches that could result from an age difference of about 1 Gyr \\citep{milo}.  For the clusters above, there is a hint that the fraction of EHB stars anticorrelates with metallicity. M15 is about 0.1 dex more metal poor than NGC 2419, but also by far the least massive of these clusters.  $\\omega$ Cen has a large metallicity spread ($-2.2 \\la $ [Fe/H] $\\la -0.7$), but the most common [Fe/H] $\\sim -1.75$ \\citep{johnomega}. Of the clusters with moderate EHB populations, M54 has [Fe/H] $\\sim -1.6$ \\citep{brownm54} and NGC 2808 has [Fe/H] $\\sim -1.15$ \\citep{carn2808}. The remainder of the massive clusters (NGC 6388, NGC 6441, and 47 Tuc) are very metal rich ([Fe/H] $> -0.7$). This crude correlation may be related to well-known relation between [Fe/H] and HB morphology --- more metal-poor clusters tend to have bluer HB stars.  Even if we compare NGC 2419 to its nearest rival in producing EHB stars ($\\omega$ Cen), we find a lack of a clear lead.  $\\omega$ Cen has an HB morphology similar to NGC 2419, though $\\omega$ Cen has diverted a greater proportion of stars into the blue HB tail between the two peaks, which would imply that multiple populations should overlap to a greater degree in $\\omega$ Cen.  $\\omega$ Cen shows very clear signs of multiple populations on the giant branch and main sequence \\citep{bedomega}, including a blue main sequence that is more metal-rich than the majority of stars, and therefore is probably also helium enriched \\citep{pomega}. The number ratio of chemically enriched group to the total MS star population is similar to the ratio of EHB stars to the total HB population, and thus \\citeauthor{pomega} suggested that the EHB stars are the descendants of the enriched population.  In NGC 2419, there is no evidence to date of multiple populations among red giants or main sequence stars.  Although there is a little evidence consistent with slight helium enrichment in the cluster (based on post-HB evolution and the $R$ population ratio), there is no evidence that it is as strong as in $\\omega$ Cen.  It is potentially possible to hide multiple populations on the main sequence by increasing metallicity and helium abundance in a way to cancel out their opposing color shifts. If this is happening in NGC 2419, it would imply a smaller helium enrichment ($\\Delta Y \\sim 0.03$) than is implied for $\\omega$ Cen ($\\Delta Y = 0.14$). Such a small helium enrichment would be consistent with the indicators presented in this paper, but would be unable to explain the two HB star populations. Spectroscopic observations of cluster giants would still be needed to determine whether a metallicity spread is truly present in the cluster.  In one sense NGC 2419 is like $\\omega$ Cen: neither appears to have undergone strong dynamical relaxation \\citep{dale,ferromega}.  There is little or no evidence of a correlation with concentration $c$ \\citep{mcl} though. The three massive clusters with the largest EHB fractions (NGC 2419, $\\omega$ Cen, and M54) are also largest in $R_h$, although 47 Tuc has similar structural properties to M54 and has no EHB stars.  At this point, we must confess that while NGC 2419 seems to have less complicated stellar populations overall than clusters like $\\omega$ Cen and NGC 2808, its relative simplicity helps to stymie attempts to draw out a coherent picture of the source of multiple populations and extreme horizontal branch stars."
    },
    "0809/0809.3787_arXiv.txt": {
        "abstract": "The Magellanic Clouds were the largest members of a group of dwarf galaxies that entered the Milky Way (MW) halo at late times.  This group, dominated by the LMC, contained $\\sim 4$\\% of the mass of the Milky Way prior to its accretion and tidal disruption, but $\\approx 70\\%$ of the known dwarfs orbiting the MW.  Our theory addresses many outstanding problems in galaxy formation associated with dwarf galaxies.  First, it can explain the planar orbital configuration populated by some dSphs in the MW.  Second, it provides a mechanism for lighting up a subset of dwarf galaxies to reproduce the cumulative circular velocity distribution of the satellites in the MW.  Finally, our model predicts that most dwarfs will be found in association with other dwarfs.  The recent discovery of Leo V (Belokurov et al. 2008), a dwarf spheroidal companion of Leo IV, and the nearby dwarf associations supports our hypothesis. ",
        "introduction": "In the cold dark matter (CDM) model, the dark halos of galaxies like the Milky Way build up hierarchically, through the accretion of less massive halos.  When these sub-systems avoid complete tidal disruption, they can survive in the form of satellite dwarf galaxies. However, the dwarf galaxies in the Local Group exhibit several puzzling features.  Numerical simulations of CDM predict 10 to 30 times more satellites within 500 kpc of the Milky Way and M31 than the modest observed population (e.g. Moore et al. 1999).  This discrepancy between the expected and known numbers of dwarf galaxies has become known as the {\\it missing dwarf problem}.  The newly discovered population of ultra-faint dwarfs around the Milky Way and M31 found in the Sloan Digital Sky Survey increases by a factor of two the number of known satellites (Simon \\& Geha 2007), but goes to even lower circular velocities where a comparable or even greater increase in the number of satellites is expected.  Another peculiarity is that many dwarf galaxies in the Local Group lie in the orbital plane of the Magellanic Clouds and Stream.  These dwarfs have been associated with the Magellanic Clouds and termed the Magellanic Group (Lynden-Bell 1976, Fusi Pecci et al. 1995; Kroupa et al. 2005).  In order to reproduce this planar configuration in the current scenario for structure formation, Libeskind et al. (2005) proposed that subhalos are anisotropically distributed in cosmological CDM simulations and that the most massive satellites tend to be aligned with filaments.  Similarly, Zentner et al. (2005) suggested that the accretion of satellites along filaments in a triaxial potential leads to an anisotropic distribution of satellites.  Systems anisotropically distributed falling into the Galactic halo may not lie in a plane consistent with the orbital and spatial distribution of the MW satellites.  For example, a theoretical bootstrap analysis of the spatial distribution of CDM satellites (taken from a set of CDM simulations) by Metz et al. (2008) finds that even if they are aligned along filaments, they will be consistent with being drawn randomly.  This could mean that alignment of the satellites along filaments may not be sufficient to reproduce the observed planar structures.  As we propose here, the origin of planar distributions is facilitated by concentrating infalling satellites into groups.   Another issue is that the dSphs of the Local Group tend to cluster tightly around the giant spirals.  Proximity to a large central galaxy might prevent dwarf irregulars from accreting material, turning off star formation, and they may then undergo tidal interactions to convert them into dwarf spheroidals.  However, isolated dSphs like Tucana or Cetus found in the outskirts of the Local Group (Grebel et al. 2003) suggest that dSphs might also form at great distances from giant spirals prior to their being accreted.  Clues to the questions raised by these observations may be contained in measurements of the metallicities of a large sample of stars in four nearby dwarf spheroidal galaxies: Sculptor, Sextans, Fornax, and Carina. Work by Helmi et al. (2006) shows that all four lack stars with low metallicity, implying that their metallicity distribution differs significantly from that of the Galactic halo, indicating a non-local origin for these systems. ",
        "conclusions": "We assume a model where the LMC was the largest member of a group of dwarf galaxies that was accreted into the MW halo.  Our picture addresses several questions in galaxy formation: ({\\it i}) It explains the association of some dwarf galaxies in the Local group with the LMC-SMC system.  ({\\it ii}) It provides a mechanism to light up dwarf galaxies. ({\\it iii}) It predicts that other isolated dwarfs will have companions. The recent discovery of Leo V (Belokurov et al. 2008), a dwarf spheroidal companion of Leo IV, and the nearby dwarf associations supports our hypothesis.   E.D is grateful to Jacco van Loon and Joana Oliveira for organizing an interesting meeting. She also would like to thank J. Gallagher, G. Besla, K. Bekki, L. Hernquist,  N. Kallivayalil, C. Mastropietro for fruitful discussions."
    },
    "0809/0809.4358_arXiv.txt": {
        "abstract": "We propose to apply an object point process to automatically delineate filaments of the large-scale structure in redshift catalogues. We illustrate the feasibility of the idea on an example of the recent 2dF Galaxy Redshift Survey, describe the procedure, and characterise the results. ",
        "introduction": "The large-scale structure of the Universe traced by the three-dimensional distribution of galaxies shows intriguing patterns: filaments and sheet-like structures connecting in huge clusters surround nearly empty regions, the so-called voids. Surveys of galaxies are created by measuring for each galaxy in addition to its angular position on the sky its distance, estimated from their recession velocities, within the framework of a cosmological model. However, the measured recession velocities are not only due to the Hubble expansion, they appear contaminated by the line-of-sight contribution of the dynamical velocity of a galaxy, due to the local gravitational effects, so the distances of galaxies are in error. Galaxy surveys based on recession velocities are called 'redshift space' maps, and although they are distorted versions of the three-dimensional distribution of galaxies, distance errors are not as serious as to change the overall picture of the large-scale structure.  An overview of such galaxy maps is given in \\cite{martsaar02}. As an example, we present here a map from a recently completed 2dF Galaxy Redshift Survey (2dFGRS, \\cite{2dFGRS}). This survey measured the redshifts (recession velocities) of galaxies in about 1500 square degrees, up to the distances of about $700\\,h^{-1}\\mbox{Mpc}$\\footnote{ Distances between galaxies are usually measured in megaparsecs (Mpc); $1\\mbox{Mpc}\\approx3\\cdot10^{24}\\mbox{cm}$. The constant $h$ is the dimensionless Hubble parameter; the latest determinations give for its value $h\\approx 0.71$. } (corresponding to a redshift $z=0.2$ for the standard cosmological model). The redshifts were measured in two different regions of the sky; Fig.~\\ref{fig:2dfgrs} shows the positions of galaxies in two $2.6^\\circ$ thick slices from both regions.   \\begin{figure} \\centering \\resizebox{0.8\\textwidth}{!}{\\includegraphics*{2dfgrs.eps}}\\\\ \\caption{Galaxy map for two 2dFGRS slices. The observers (we) are situated at the centre of the figure. Both slices are thin, with the thickness of $2.6^\\circ$. The distances are given in redshifts $z$; approximately, the physical distance $D\\approx 3000\\,h^{-1}\\,z\\,\\mbox{Mpc}$. The numbers along the arcs show the right ascension (in hours). The filamentary network of galaxies is clearly seen; the disappearance of structure with depth (towards the sides of the figure) is caused by luminosity selection. \\label{fig:2dfgrs} } \\end{figure}  The most characteristic feature of all three-dimensional galaxy cartographies are the remarkable network of filaments that can be appreciated when diagrams like the one shown in Fig.~\\ref{fig:2dfgrs} are depicted. All filaments are not equal, they show different size and contrast, but typically relatively empty voids are found between them. Different scales are involved in the filamentary patterns, and it is of great importance to have algorithms to identify them and to characterise their properties. We should mention here that the gradual disappearance of structures with distance observed in Fig.~\\ref{fig:2dfgrs} is a selection effect due to how this surveys are built. They are flux-limited samples, and since the apparent luminosity of a galaxy is fainter if the galaxy lies further away, only the brightest galaxies are seen in the more distant regions of the surveyed volume.  Certainly filaments are prominent and visually they dominate the galaxy maps, however there are still no standard methods to describe the observed filamentary structure. Second-order summary statistics like the two-point correlation function, the $K$-function or the power spectrum (in Fourier space) do not provide morphological information. Minkowski functionals, minimal spanning tree (MST), percolation and shapefinders have been introduced for this purpose (for a review see \\cite{martsaar02}).  The minimal spanning tree was introduced in cosmology by \\cite{barrow85}. It is a unique graph that connects all points of the process without closed loops, but certainly describes mainly the local nearest-neighbour distribution, being unable to provide a whole characterisation of the global and large-scale properties of the filamentary network.  A better method, named {\\it skeleton}, has been recently proposed (\\cite{eriksen04,novikov06}) to describe the possible filamentary structure of continuous density fields. The skeleton is determined by segments parallel to the gradient of the field, connecting saddle points to local maxima. Calculating the skeleton involves interpolation and smoothing the point distribution, which introduces an extra parameter, the band-width of the kernel function used to estimate the density field from the point distribution, typically a Gaussian function. In any case, this method has been recently applied to the Sloan Digital Sky Survey (\\cite{sousbie06}, providing, by means of the length of the skeleton, a good discriminant tool for the analysis of the filamentary structures.  In a previous paper (\\cite{StoiMartMateSaar05}), we proposed to use an automated method to trace filaments for realisations of point processes, that has been shown to work well for detection of road networks in remote sensing situations (\\cite{LacoDescZeru05,StoiDescLiesZeru02,StoiDescZeru04}). This method is based on the Candy model, a marked point process where segments serve as marks. The Candy model can be applied to 2-D filaments, and we tested it on simulated galaxy distributions. The filaments we found delineated well the filaments detected by eye.  Many methods used to automatically detect filaments have been developed so far for two-dimensional maps. This is a natural approach for studying the cosmic microwave sky background (\\cite{eriksen04}), which is two-dimensional, but galaxy maps are three-dimensional. The previous filament studies (e.g., \\cite{bhavsar03, bharadwaj04,pandey05}) have also used two-dimensional galaxy maps, projections for thin slices. The main reason for that is that most of the past large-scale galaxy maps were observed for relatively thin spatial slices; also, in projection filaments seem more prominent. But, of course, both slicing of filaments and projecting the distribution onto a plane distorts the geometry and the properties of the filamentary network.  The study of the three-dimensional filamentary network has just begun. Three-dimensio\\-nal filaments have been extracted from galaxy distribution as a result of special observational projects (\\cite{pimbblet04a}), or by searching for filaments in the 2dFGRS catalogue (\\cite{pimbblet04b}). These filaments have been searched for between galaxy clusters, determining the density distribution and deciding if it is filamentary, individually for every filament. Similar studies have been carried out for N-body simulations (\\cite{colberg05}). A review of these and previous studies of large-scale filaments is given in (\\cite{pimbblet05}).  No automated methods to trace filaments in the three-dimensional galaxy maps have been proposed so far. Based on our previous experience with the Candy process, we generalised the approach for three dimensions. As the interactions between the structure elements are more complex in three dimensions, we had to define a more complex model, the Bisous model (\\cite{StoiGregMate05}). This model gives a general framework for the construction of complex patterns made of simple interacting objects. In our case, it can be seen as a generalisation of the Candy model. We will describe the Bisous model below and will apply it to the samples chosen from the real three-dimensional galaxy distribution, the 2dFGRS catalogue. ",
        "conclusions": "We have applied an object point process -- Bisous model -- to objectively find filaments in galaxy redshift surveys (three-dimensional galaxy maps). For that, we defined the model, fixed some of the interaction parameters and chose priors for the remaining parameters. The definition of the data term is very intuitive and rather a simple test. Much more elaborated methods testing the alignment of the points along the direction of the cylinder against the completely spatial randomness need investigation. The uniform law for the interaction parameters was preferred in order to give the same chances to a wide range of topologies of the filamentary network. Still, if concrete prior information about the topology of the filamentary network is available then this should be integrated in the model.  We have run simulated annealing sequences and select filaments on the basis of coverage probabilities for individual cells of sample volume. The coverage probabilities are to be seen as a way of averaging the shape of the filamentary structure. They have the advantage to allow inference from statistics instead of a single realisation. Their main drawback is that the coverage probabilities are computed locally, for small regions. Still, the visualisation of the visit maps built on these probabilities brings new ideas and hypotheses about the topologies of the cosmic filaments. A global Monte Carlo statistical test was built to test the existence of the filaments in the data. To do this, we have calculated sufficient statistics for the data sets and made a comparison with the sufficient statistics obtained on binomial point fields having the same number of points as the data. Our test was indicating that the filaments we find are defined by the data, not by the chosen model.  The method used in this paper can be extended in different ways. One natural extension is to use different generating elements instead on cylinders, e.g., planar elements or clusters. (\\cite{StoiGregMate05}). Although traces of planar structures are seen in superclusters of galaxies, these have been difficult to quantify, mainly because of their low density contrast. Another interesting application is to search for dynamically bound groups and clusters of galaxies that have a typical 'finger-of-God' signature in redshift space, extended along the line-of-sight. And there remains a question if the method could be extended to inhomogeneous point processes -- this would allow us to use all the observational data, not only volume-limited subsamples."
    },
    "0809/0809.4672_arXiv.txt": {
        "abstract": "We obtain new HI and $^{13}$CO images around Supernova Remnants (SNR) Kes 69 and G21.5-0.9. By comparing HI spectra with $^{13}$CO emission spectra, we significantly revise the kinematic distance for Kes 69 to $\\sim$ 5.5 kpc, which was 11.2 kpc, and refine the kinematic distance for G21.5-0.9 to $\\sim$ 4.8 kpc. For Kes 69, the highest velocity of absorption is $\\sim$ 86 km s$^{-1}$ and a prominent HI emission feature at $\\sim$ 112 km s$^{-1}$ has no respective absorption. These new results suggest that Kes 69 is associated with a newly detected extended 1720 MHz OH maser at velocity of $\\sim$ 85 km s$^{-1}$ that originates from within the bright southern radio shell of Kes 69. For G21.5-0.9, the highest velocity of absorption is $\\sim$ 67 km s$^{-1}$. The HI absorption spectra of the nearby bright source PMN J1832-1035 and of Kes 69 show a common absorption feature at velocity of $\\sim$ 69 km s$^{-1}$, which is not seen for G21.5-0.9. The resulting velocity of $\\sim$ 68 km s$^{-1}$ gives the best distance estimate of $\\sim$ 4.8 kpc for G21.5-0.9 and associated young pulsar J1833-1034. ",
        "introduction": "Determining distances to Galactic objects (HII regions, pulsars (PSR) and supernova remnants (SNR) etc.) may help understand the kinematics of the Milky Way. Also, distance measurement of SNRs is a key to obtain their basic parameters such as the luminosity, size and age. As an energetic class of objects, SNRs are associated with many highly active astrophysical phenomena, e.g. anomalous X-ray pulsars, soft  $\\gamma$-ray repeaters, Pulsar Wind Nebula (PWN), non-thermal X-rays and very-high-energy $\\gamma$-rays (see Yang et al. 2008 for a review). The determination of distance of a SNR may test reality of a SNR/PSR/PWN or SNR/TeV $\\gamma$-ray source association, and help constrain the mass range of the progenitor star and type of supernova responsible for the remnant. It can provide direct evidence whether a SNR is physically associated with molecular clouds, or if strong interaction between a SNR shock and surrounding clouds is possible, which is widely believed to be a source of TeV $\\gamma$-ray and non-thermal X-ray emission in the Milky Way. The SNR/molecular cloud interaction may also potentially lead to new generations of star formation, so constitutes an important part of Galactic ecology.  Comparing  HI absorption spectra toward  Galactic SNRs with respective HI and $^{13}$CO emission spectra along the line of sight, we previously revised distances to five SNRs and several overlapping HII regions (Tian et al. 2007; Tian \\& Leahy 2008; Leahy \\& Tian 2008a;). Some values differ greatly from previously published values, e.g. Kes 75 (Leahy \\& Tian 2008b), PWN G54.1+0.3 (Leahy, Tian \\& Wang 2008), because our methods implement improved background subtraction and spurious emission rejection, and resolve the near/far distance ambiguity in the inner Galaxy, which may have been previously done incorrectly. The distance changes can result in significant changes in interpretation of the SNR (and associated object) properties. As the newest paper of a series, we use the methods to directly re-measure distances of two more SNRs: Kes 69 which has 1720 MHz OH maser emission (Hewitt et al. 2008), and G21.5-0.9 which hosts a young pulsar (Camilo et al. 2006; Bietenholz \\& Bartel 2008) and also has TeV $\\gamma$-ray emission (i.e. HESS J1833-105, Djannati-Atai et al. 2007). 1720 MHz masers are associated with warm, shocked molecular gas and are seen as signposts of SNR-molecular cloud interactions (Wardle \\& Yusef-Zadeh 2002), therefore they likely give reliable distance estimates to SNR/molecular cloud systems.  An OH maser at velocity of 69.3$\\pm$0.7 km s$^{-1}$ (hereafter, we use $\\sim$ 70 km s$^{-1}$to describe the radial velocity of this maser) was detected toward Kes 69 (Green et al. 1997), leading to a distance of $\\sim$ 11.2 kpc for the remnant. However, the masers' site is outside of the northeastern radio and X-ray shells of Kes 69 (Yusef-Zadeh et al. 2003). This spurred us to re-measure its distance by analyzing HI+$^{13}$CO spectra in the direction of Kes 69 in order to test the reality of the claimed association. SNR G21.5-0.9 is an interesting SNR nearby Kes 69, so we also analyze HI and CO spectra in the direction of G21.5-0.9.  In this paper, we significantly revise the distance to Kes 69, and refine the distance estimation to G21.5-0.9 by analyzing HI and CO spectra. We use the Galactic rotation curve model and recent measurements of the parameters for this (i.e. R$_{0} \\sim$ 8 kpc, Eisenhauer et al. 2005; V$_{0}$$\\sim$ 220 km s$^{-1}$, Feast \\& Whitelock 1997, Reid \\& Brunthaler 2004).  The radio data come from 1420 MHz continuum plus HI-line observations of the VLA Galactic Plane Survey (VGPS, Stil et al. 2006) and the $^{13}$CO-line (J = 1-0) observations of the Galactic ring survey (Jackson et al. 2006) of the Five College Radio Astronomy Observatory 14 m telescope. The short-spacing information for the HI spectral line images is from additional observations with the 100 m Green Bank telescope (GBT) of the National Radio Astronomical Observatory (NRAO). ",
        "conclusions": "\\subsection{Distances to Kes 69 and G21.5-0.9}  The OH-maser determined distance should be reasonably consistent with the HI+CO determined distance to the same SNR. We note that distances are derived here from a circular rotation model, and contain errors due to uncertainties in the rotation model and non-circular motions. E.g. using the observed $l-V$ diagram, Weiner \\& Sellwood (1999) derived the radial velocity distribution in the inner galaxy. Applying to Kes 69, the distances are reduced by $\\sim$ 0.4 kpc compared to the circular rotation model. Observed random motion of up to 7 km s$^{-1}$ (Shaveret al. 1982) yields a distance uncertainty of $\\sim$ 0.3 kpc for this case. Using the circular rotation model, our analysis of the HI+CO spectra reveals that Kes 69 has a distance of 5.5 to 7.4 kpc, far smaller than 11.2 kpc determined by the OH maser at 70 km s$^{-1}$. This is strong evidence against the OH maser being associated with Kes 69. This is supported by further evidence: the site of the OH maser is outside the detected radio emission of Kes 69, and no CO emission at $\\sim$ 70 km s$^{-1}$ is seen in the CO spectrum of the OH maser. A collisional pumping model shows that shock-excited OH maser emission at 1720 MHz appearing behind a SNR shock front results from the interaction of an SNR with an adjacent, warm, dense shocked molecular cloud (Wardle \\& Yusef-Zadeh 2002).  Furthermore, an extended 1720 MHz OH maser at velocity of $\\sim$ 85 km s$^{-1}$ has been detected from within the bright southern radio shell of Kes 69 by GBT and VLA observations (Hewitt et al. 2008). The new OH maser fits nicely with the highest velocity of HI absorption in the direction of Kes 69, therefore we conclude that Kes 69 is located at a distance of $\\sim$ 5.5 kpc at the near side kinematic distance for the OH maser at $\\sim$ 85 km s$^{-1}$ and consistent with the highest observed HI absorption velocity of $\\sim$ 86 km s$^{-1}$. A bright $^{13}$CO cloud at $\\sim$ 86 km s$^{-1}$ overlapping Kes 69 has been revealed here (Fig. 1 upper right). A research group from Nan Jing University newly finds $^{12}$CO, $^{13}$ CO and HCO$^{+}$ emissions near 85 km/s$^{-1}$ from the bright southern shell (Zhou et al. 2008), consistent with the previous formaldehyde molecular (H$_{2}$CO) and the 1665/7 MHz OH absorption line observations at velocities of 82 -- 87 km s$^{-1}$ (Wilson 1972, Turner 1970). All these support our conclusion on the distance of 5.5 kpc to Kes 69. So the 1720 MHz maser at 70 km s$^{-1}$ is likely not associated with Kes 69.  G21.5-0.9 is located between a cloud at velocity of $\\sim$ 67 km s$^{-1}$ and that of $\\sim$ 69 km s$^{-1}$ (i.e. d $\\sim$ 4.8 kpc). G21.5-0.9 hosts a young pulsar PWN J1833-1034. New measurement refines the most recent velocity measurement to the SNR/PSR system by Camilo et al. 2006 who analyzed HI spectra data from the VLA and the Leiden/Dwingeloo 25 m telescope observations, and lower resolution $^{12}$CO data from the CfA 1.2 m telescope observations, to obtain lower and upper limits on the velocity of 65 km s$^{-1}$ and 76 km s$^{-1}$ (i.e. d $\\sim$ 4.7$\\pm$ 0.4 kpc).  \\subsection{The Evolutionary State of Kes 69}  The 1420 MHz continuum map shows Kes 69 has a roughly elliptical outline 26$^\\prime$ by 20$^\\prime$ whereas the ROSAT PSPC image (Yusef-Zadeh et al. 2003), also elliptical and oriented in the same way, has smaller dimensions of  20$^\\prime$ by 13$^\\prime$. This is mainly due to absence of X-rays from the northwest radio filament (upper left side in Fig. 1 here) and from the northeast radio wing (lower left in Fig. 1), but also the X-rays do not extend out to the edge of the main radio filament in the south. We use 20$^\\prime$ as the mean angular diameter of Kes 69, which is the mean of radio and X-ray values. At the distance of 5.5 kpc, the diameter is $\\sim$ 32 pc.  We apply a Sedov model (Cox  1972) to estimate the parameters of Kes 69, using the ROSAT PSPC temperature of 1.6 keV and X-ray luminosity of 8.4$\\times 10^{34}$ erg/s (Yusef-Zadeh et al. 2003). The high temperature and moderate X-ray luminosity are indicative of explosion in a low density medium. Application of the Sedov model yields an age of 5000 years, an explosion energy of 0.8$\\times 10^{51}$ erg, and a pre-explosion density of $\\sim$0.1 cm$^{-3}$. Based on the estimated errors in the ROSAT PSPC spectral parameters, the errors in these values are $\\sim$60\\% for age and explosion energy and $\\sim$30\\% for density. These results are consistent with observation of OH maser emission from the main south filament of Kes 69 if the explosion occurred in a moderately low density cavity ($\\sim$0.1 cm$^{-3}$) and has only recently ($<<$ 5000 years ago) run into dense molecular gas at the main southern filament. Supporting this are Spitzer infrared observations of line emissions from shocked molecular gas in Kes 69 (Reach et al. 2006). For the ROSAT temperature of 1.6 keV, the SNR shock velocity is 1250 km s$^{-1}$ in 0.1 cm$^{-3}$ gas, but is slowed down to $\\sim$ 4 km s$^{-1}$ in cold molecular gas with density of $\\sim$ $10^{4}$ cm$^{-3}$. The shock would still be supersonic, since the sound speed is 0.2 to 0.6 km s$^{-1}$ in 10 to 100 K molecular hydrogen, and would result in mild heating of the gas to $\\sim$ 4000 K. This would subsequently result in good conditions for production of the observed OH maser emission (Wardle 1999).  In summary, we significantly revise the distance to Kes 69 and obtain a better distance to G21.5-0.9, by an analysis using HI and $^{13}$CO spectra. The evolutionary state of Kes 69 is evaluated, and it is consistent with a moderate aged SNR, just recently encountering a dense molecular cloud.  \\begin{figure*} \\vspace{150mm} \\begin{picture}(100,100) \\put(-180,360){\\special{psfile=f1a.eps hoffset=0 voffset=0 hscale=40 vscale=35 angle=0}} \\put(-10,590){\\special{psfile=f1b.eps hoffset=0 voffset=0 hscale=48 vscale=48 angle=-90}} \\put(-235,385){\\special{psfile=f1c.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\put(30,385){\\special{psfile=f1d.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\put(-235, 210){\\special{psfile=f1e.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\put(80, -80){\\special{psfile=f1f.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=0}}  \\end{picture} \\caption{First row: the 1420 MHz continuum image (left) and  a channel CO map with velocity of $\\sim$ 86 km s$^{-1}$ (right), centered at [b=-0.5, l=21.75]; Kes 69 is the large extended object near center, G21.5-0.9 is inside box 3, PMN J1832-1035 is inside box 4. Second and third row: HI maps at velocities of $\\sim$ 112 km s$^{-1}$ (close to the tangent point velocity in the direction towards Kes 69, middle left), 86 km s$^{-1}$ (middle right), 69 km s$^{-1}$ (lower left) and 67 km s$^{-1}$ (lower right), respectively. The first panel has superimposed contours (green: 25, 43, 75, 118, 145K) of the 1420 MHz continuum emission to show the SNRs, other panels have superimposed contours (green: 25, 50, 500 K).  In the first panel, the plus sign shows the site of the 1720 MHz OH maser at $\\sim$ 70 km s$^{-1}$, the four boxes mark  regions where  HI and CO spectra are extracted.} \\end{figure*}   \\begin{figure*} \\vspace{150mm} \\begin{picture}(120,120) \\put(-180,420){\\special{psfile=f2a.eps hoffset=0 voffset=0 hscale=40 vscale=25 angle=0}} \\put(60,560){\\special{psfile=f2a1.eps hoffset=0 voffset=0 hscale=30 vscale=25 angle=-90}} \\put(-180,275){\\special{psfile=f2b.eps hoffset=0 voffset=0 hscale=40 vscale=25 angle=0}} \\put(60,415){\\special{psfile=f2b1.eps hoffset=0 voffset=0 hscale=30 vscale=25 angle=-90}} \\put(-180,140){\\special{psfile=f2c.eps hoffset=0 voffset=0 hscale=40 vscale=25 angle=0}} \\put(60,280){\\special{psfile=f2c1.eps hoffset=0 voffset=0 hscale=30 vscale=25 angle=-90}} \\put(-180,0){\\special{psfile=f2d.eps hoffset=0 voffset=0 hscale=40 vscale=25 angle=0}} \\put(60,140){\\special{psfile=f2d1.eps hoffset=0 voffset=0 hscale=30 vscale=25 angle=-90}} \\end{picture} \\caption{Left: four HI emission and absorption spectra (from top to bottom), extracted from boxes 1, 2, 3 and 4 shown in Fig. 1a. Right: four CO emission spectra, extracted from boxes 1, 2, 3 and from OH maser site shown in Fig. 1a.} \\end{figure*}"
    },
    "0809/0809.4391_arXiv.txt": {
        "abstract": "ASTEP South is the first phase of the ASTEP project that aims to determine the quality of Dome C as a site for future photometric searches for transiting exoplanets and discover extrasolar planets from the Concordia base in Antarctica. ASTEP South consists of a front-illuminated 4k x 4k CCD camera, a 10 cm refractor, and a simple mount in a thermalized enclosure. A double-glass window is used to reduce temperature variations and its accompanying turbulence on the optical path. The telescope is fixed and observes a $4\\,^{\\circ}$ x $4\\,^{\\circ}$ field of view centered on the celestial South pole. With this design, A STEP South is very stable and observes with low and constant airmass, both being important issues for photometric precision. We present the project, we show that enough stars are present in our field of view to allow the detection of one to a few transiting giant planets, and that the photometric precision of the instrument should be a few mmag for stars brighter than magnitude 12 and better than 10 mmag for stars of magnitude 14 or less. ",
        "introduction": "The ASTEP project (Antarctic Search for Transiting Extrasolar Planets) aims to  determine the quality of Dome C, Antarctica as a site for future photometric surveys and to detect transiting planets (\\cite[Fressin et al. 2005]{}). The 3 month continuous night as well as a very dry atmosphere should yield to a great improvement of the photometric precision when compared to other sites. The goal of ASTEP South is to qualify the photometry from Dome C and to attempt to detect planets. First, the instrument setup is described. We then present the expected planet detection probability using a simple model. Finally, we use SimPhot to simulate the ASTEP South observations and a first data reduction process. The expected noise is then derived for stars in the ASTEP South field of view. ",
        "conclusions": "ASTEP South is the first transit survey from Dome C, Antarctica. The instrument inside its thermalized box is currently observing towards the celestial South pole. A simple analytical model shew that the observed field of view contains enough stars to enable transit detections with our 10 cm refractor. Simulations made with SimPhot show a rms noise close to the photon noise for stars of magnitude 12 or less, and a noise level better than 10 mmag for stars of magnitude 14 or less. Thus, this first campaign should allow us to test this new observing method, to qualify Dome C for photometric observations and possibly to detect planets."
    },
    "0809/0809.1327_arXiv.txt": {
        "abstract": "{ Since Baade's photographic study of M32 in the  mid 1940s, it has been accepted as an established fact that M32 is a compact dwarf satellite of M31. The purpose of this paper is to report on the findings of our investigation into the nature of the existing evidence. We find that the case for M32 being a satellite of M31 rests upon Hubble Space Telescope (HST) based stellar population studies which have resolved  red-giant branch (RGB) and red clump stars in M32 as well as other nearby galaxies. Taken in isolation, this recent evidence could be considered to be conclusive in favour of the existing view. However, the conventional scenario does not explain M32's anomalously high central velocity dispersion for a dwarf galaxy (several times that of either NGC 147, NGC 185 or NGC 205) or  existing planetary nebula observations (which suggest that M32 is more than twice as distant as M31) and also requires an elaborate physical explanation for M32's inferred compactness. Conversely, we find that the case for M32 being a normal galaxy, of the order of three times as distant as M31, is supported by: (1) a central velocity dispersion typical of intermediate galaxies, (2) the published planetary nebula observations, and (3) known scaling relationships for normal early-type galaxies. However, this novel scenario cannot account for the high apparent luminosities of the RGB stars resolved in the M32 direction by HST observations. We are therefore left with two apparently irreconcilable scenarios, only one of which can be correct, but both of which suffer from potentially fatal evidence to the contrary. This suggests that current understanding of some relevant fields is still very far from adequate. ",
        "introduction": "\\label{Sec1}  As M32 appears projected onto the disc of M31, separation of the two galaxies' integrated light and constituent stars is notoriously difficult. Consequently, hardly any formal distance measurements are available for M32. On the other hand, much effort has gone into attempting to demonstrate the existence of physical associations between M31 and M32 (which would imply physical proximity) and/or that M32 is in front of M31's disc (which would place a hard upper limit on M32's distance). This is because since \\citet{B 1944} it has been generally accepted that M32 must be a satellite of M31. M32's distance has therefore not been regarded as a major issue as it has generally been assumed to be similar to M31's (and therefore to be known at least approximately). Note though that before \\citet{B 1944} there was no consensus on M32's status with respect to M31 or on its distance.   Assuming that M32 is at a similar distance to that of M31 (whose distance we take to be 0.76~Mpc from e.g.\\ \\citealt{Vdb 2000a}) then, on account of its angular size and degree of central concentration, it must be an unusually compact dwarf galaxy--regardless of  whether it is elliptical or lenticular in nature. Note that although M32 has long been presumed to be an elliptical galaxy, the possibility that it may instead be a lenticular needs to be taken seriously \\citep{G 2002}. Since e.g.\\ \\citet{Dv 1961}, it has been accepted that M32 is a red compact galaxy (RCG). Most members of this extremely rare class of galaxy appear to be associated with larger neighbours and since \\citet{K 1962} it has been accepted that their compactness arises due to tidal truncation by their neighbours. Red compact dwarfs (RCDs) are all the more unusual (and rare) because their red intrinsic colours and high degree of central concentration with respect to mass are characteristic of early-type giants and intermediates, not dwarfs. In addition to M32, only five other candidates have been identified \\citep{C+ 2007}. However, as is evident from figures~3 and 4 of \\citet{C+ 2007}, M32 remains the faintest, smallest and most compact RCD [candidate] known, and therefore the most extreme case in terms of its deviation from most known scaling relationships for normal galaxies.   In Appendix~\\ref{AppA}, in a radical departure from current practices, we treat M32 as if it were a normal galaxy in order to investigate what distance it would need to be projected to in order for it to fit as many scaling relationships as possible. We find that if M32 were projected to a distance of 2.3($\\pm$0.8) [$-$0.77$A_B$(M31)~mag$^{-1}$]~Mpc and if $A_B$(M31) ($B$-band absorption due to M31) were small, it would obey the known scaling relationships defined by normal early-types, and a physical explanation for its inferred extreme compactness would not be necessary. With this result in mind, in this paper, we test our radical new hypothesis against the existing evidence, the null hypothesis being that M32 is an RCD satellite of M31.   Unfortunately, in the testing of any controversial new hypothesis a critical approach is often unavoidable, and this paper is no exception. It is not our aim to be unnecessarily critical of the existing results, but as we hope most readers will agree, we believe that it is a necessary exercise to investigate the possibilities for re-interpreting evidence that runs contrary to the new hypothesis. Only this way can robust evidence (that is not open to re-interpretation) be separated from circumstantial arguments (that are). Having said this, our aim is to find the truth and we do not have a vested interest in the outcome. This paper merely documents a thought experiment that we have conducted. However, it may also be of historical interest and even shed some light on how science is done.   In Sections~\\ref{Sec2}, \\ref{Sec3} and \\ref{Sec4}, italic type is used for the existing evidence in order to distinguish it from our own analysis. It was found to be impractical to present the existing evidence in strict chronological order. Instead, we have attempted to introduce readers to the issues in an order that requires minimal prior knowledge of the subject, whilst at the same time minimizing the need for any repetition. With a view to remaining objective throughout this paper, we have adopted wordings that neither imply that M32 is an RCD nor imply that it is a background galaxy. ",
        "conclusions": "\\label{Sec5}   Prior to \\citet{B 1944} it was not known whether M32 was a dwarf satellite of M31 or a background galaxy. However, following \\citet{B 1944} it became established as an unquestionable fact that M32 must be a dwarf satellite of M31. We have demonstrated that contrary to popular belief, the only robust evidence for this case is very recent in nature and ultimately rests entirely on resolved stellar photometry obtained with the HST. In other words, in the absence of these HST observations, the hypothesis that M32 might be a background galaxy would still be irrefutable. The assumptions on which the distance measurements derived from SBFs \\citep{TS 1988,T 1991,LT 1993} and ground-based resolved stellar photometry \\citep{F 1989} were based, have been confirmed by later HST observations\\footnote{As SBF studies amount to unresolved stellar photometry, they are related to (and complementary to) resolved stellar photometry but do not constitute independent evidence.}, but all of the preceding evidence beginning with that of \\citet{B 1944} was found to be circumstantial.   We accept that M32 could well be an RCD satellite of M31. Unfortunately however, we would be left with the well known problem of explaining why M32 (and to a lesser extent five other less extreme RCD candidates) should defy most known scaling relationships defined by normal early-type galaxies. This is regardless of whether M32 is an elliptical or a lenticular. Presumably this might be because RCDs are in the process of being tidally stripped and are therefore stellar systems that are not in equilibrium. More specifically though, a major problem we are left with is why as a dwarf galaxy, M32 should have a central stellar velocity dispersion several times higher than that of M31's other main companions (which are of similar apparent brightness to M32). Could this be because M32 used to be a much larger normal galaxy prior to the commencement of the stripping process? If so, perhaps it could have retained its original central velocity dispersion. However, such a scenario would contradict the findings of \\citet{NP 1987} and \\citet{CGJ 2002} who concluded that the precursor of  M32 was unlikely to have been a [larger] normal intermediate/giant elliptical galaxy. Alternatively, \\citet{B+ 2001} have suggested that the precursor of M32 might have been a late-type spiral bulge. If so, this might help to explain M32's anomalously high central velocity dispersion. Another outstanding issue though is why M32's planetary nebulae appear to be systematically fainter in apparent magnitude than M31's if the two galaxies are at the same distance.   The main arguments in favour of M32 being a normal background galaxy are as follows.   (1) M32's central velocity dispersion is typical of intermediate galaxies (including  the bulges of lenticulars), and not of dwarfs. The homogenized mean value listed by the Hyperleda database is 72~km~s$^{-1}$ cf.\\ 23~km~s$^{-1}$ for NGC~205, 20~km~s$^{-1}$ for NGC~185 and 22~km~s$^{-1}$ for NGC~147 (homogenized mean values from Hyperleda)\\footnote{The Lyon-Meudon Extragalactic Database (Hyperleda) is maintained by the Centre de Recherches Astronomiques de Lyon and available online at http://leda.univ-lyon1.fr}. However, recent high-resolution measurements have yielded even higher values for M32 e.g.\\ 126 $\\pm$10~km~s$^{-1}$ \\citep{Vdm+ 1997} and 130~km~s$^{-1}$ (based on a Gaussian fit) or $\\gs$175~km~s$^{-1}$ (when corrected for the wings) \\citep{J+ 2001}.   (2) Published planetary nebula observations from a variety of sources \\citep{CJFN 1989,HkMH 2004,M+ 2006} are all consistent in suggesting that M32 is at least twice as distant as M31.   (3) M32 would obey known scaling relationships defined by normal intermediate/giant early-type galaxies including the Fundamental plane \\citep{DD 1987} if it were of the order of three times more distant than M31.   (4) A physical explanation for M32's apparent compactness would no longer be needed.   However, this scenario cannot explain the results of space-based resolved stellar photometry of M32, which find its RGB to be very similar to M31's in both general morphology and apparent brightness--thereby requiring that the two galaxies are at almost identical distances. Only in the event of these results becoming open to re-interpretation could the new hypothesis become a viable explanation.   Unfortunately, as is evident from Figures~\\ref{f01}, \\ref{f02} and \\ref{f03}, there does not appear to be any middle ground between these two scenarios. If M32 were behind M31 but close enough to M31 to be interacting with it, explanations for M32's high central velocity dispersion and faint planetary nebulae would still be needed. We therefore find that we are left with two apparently irreconcilable scenarios, only one of which can be correct, but both of which suffer from potentially fatal evidence to the contrary.  This suggests that current understanding of some relevant fields is inadequate and that further study is urgently needed.   As far as further work on M32 is concerned, the optical thicknesses of the relevant regions of M31 are critical quantities that remain to be determined, and more studies on this problem would therefore be welcomed. Another remaining long-standing problem with wide-ranging implications is how the irregular component of the light contribution from M31 should be accounted for. The most rigorous subtraction of light from M31 to date is probably that performed for the surface photometry of \\citet{CGJ 2002}\\footnote{Following the example of \\citet{CGJ 2002} meaningful tests of the viability of models for M31 and M32's light distributions are possible. Synthetic images can be generated for the entire M31$\\cup$M32 field, in which (1) only M32 has been subtracted, (2) only M31 has been subtracted, and (3) both M32 and M31 have been subtracted.}. As these authors noted though, their elliptical isophotal model of M31 was unable to account for fine-scale structure such as dust lanes and spiral structure. With this in mind and on account of the newly-found importance of the HST-based stellar photometry, there is clearly a need for future space-based resolved stellar population studies, to obtain  not just one, but several such M31-only control fields and at more carefully chosen locations. Longer telescope-time allocations will therefore be needed."
    },
    "0809/0809.2504_arXiv.txt": {
        "abstract": "We present a time-dependent cosmic-ray modified shock model for which the calculated \\Ha{} emissivity profile agrees well with the \\Ha{} flux increase ahead of the Balmer-dominated shock at knot g in \\object[Tycho SNR]{Tycho's supernova remnant}, observed by \\citet{Lee-etal2007}. The backreaction of the cosmic ray component on the thermal component is treated in the two-fluid approximation, and we include thermal particle injection and energy transfer due to the acoustic instability in the precursor. The transient state of our model that describes the current state of the shock at knot g, occurs during the evolution from a thermal gas dominated shock to a smooth cosmic-ray dominated shock. Assuming a distance of $2.3\\kpc$ to Tycho's remnant we obtain values for the cosmic ray diffusion coefficient, $\\kappa$, the injection parameter, $\\epsilon$, and the time scale for the energy transfer, $\\tau$, of $\\kappa=2\\times10^{24}\\cms$, $\\epsilon=4.2\\times10^{-3}$, and $\\tau=426\\yr$, respectively. We have also studied the parameter space for fast ($300\\kms\\lesssim\\vs\\lesssim3000\\kms)$, time-asymptotically steady shocks and have identified a branch of solutions, for which the temperature in the cosmic ray precursor typically reaches $2$--$6\\times10^{4}\\Kv$ and the bulk acceleration of the flow through the precursor is less than $10\\kms$. These solutions fall into the low cosmic ray acceleration efficiency regime and are relatively insensitive to shock parameters. This low cosmic ray acceleration efficiency branch of solutions may provide a natural explanation for the line broadening of the \\Ha{} narrow component observed in non-radiative shocks in many supernova remnants. ",
        "introduction": "Balmer-dominated filaments in supernova remnants (SNRs) trace fast, non-radiative shocks that propagate into partially neutral, diffuse media. The filaments are sheets of shocked gas seen edge-on \\citep{Hester1987} and have been observed and studied in many SNRs including \\object[Tycho SNR]{Tycho's SNR} (\\citealp{CKR1980,KWC1987,SKBW1991}; \\citealp[][hereafter G00 and G01, respectively]{GRHB2000,GRSH2001}; \\citealp{Lee-etal2007}, hereafter L07), the Cygnus Loop \\citep[e.g.][]{RBFG1983,HRB1994}, RCW86 \\citep[e.g.][]{SGLS2003}, Kepler's SNR \\citep[e.g.][]{BLV1991,Sankrit-etal2005}, SN1006 \\citep[e.g.][]{GWRL2002}, and four remnants in the LMC \\citep[e.g.][]{SRL1994}. This work focuses on modeling a particular Balmer-dominated shock located at knot g in Tycho's SNR \\citep[after][]{KvB1978} and recently observed by L07.  The \\Ha{} spectral line profile of a Balmer-dominated filament has narrow and broad components, whose widths represent the preshock temperature of neutral H and the postshock temperature of protons, respectively \\citep{CR1978}. The width of the broad component and the ratio of intensities are diagnostics for the shock speed and the degree of electron-ion temperature equilibration behind the shock. Shock models to quantify these diagnostics were first developed by \\citet{CKR1980} and later improved by \\citet{SKBW1991}, G01, \\citet{HMcC2007}, and \\citet{HAMR2007}. Currently, the most advanced non-radiative shock models are those of \\citet{AHMR2008}. Except for the work by \\citet{BC1988}, shock models used to interpret optical observations, to date, have not included modifications of the shock structure due to diffusive shock acceleration (DSA) of cosmic rays (CRs).  Balmer-dominated filaments are usually associated with forward shocks of the expanding SNR bubbles. The only remnant in which Balmer emission from a reverse shock has been observed is SN 1987 \\citep{Michael-etal2003}. Forward shocks in SNRs are also sites of CR electron acceleration, as evidenced by synchrotron radio and X-ray emission \\citep{Koyama-etal1995,Gotthelf-etal2001,Long-etal2003,Bamba-etal2005,CHBD2007}, as well as likely sites of CR ion acceleration \\citep{BE1987,Drury-etal2001SSRv, Warren-etal2005}. In most cases, Balmer-dominated shocks do not show evidence for synchrotron emission, which suggests that particle acceleration in Balmer-dominated shocks is not efficient. The neutral component in the upstream thermal gas that is required for non-radiative shocks to produce Balmer-dominated filaments may be damping the turbulence necessary for efficient cosmic ray acceleration in SNR shocks \\citep{DDK1996}. Conversely, the heating of the upstream medium due to efficient CR acceleration may prevent most neutrals reaching the shock before being ionized \\citep{HRB1994}. Two known exceptions where Balmer emission and synchrotron X-ray emission coincide are knot g in Tycho's remnant and a small portion of the eastern rim of SN 1006 \\citep{Cassam-Chenai-etal2008}. However, a direct connection between the Balmer emission producing shock and the X-ray producing shock cannot be made for either case because the X-ray morphology is not resolved to the level of the optical emission and because the effects of projection are uncertain. Many Balmer-dominated filaments are observed to bound regions of X-ray emission whose spectra are consistent with thermal emission of a shock-heated ambient medium \\citep{HRB1994,Raymond-etal2007}.  The theory of DSA predicts a CR precursor ahead of the gas subshock \\citep{D-V1981,BE1999}, in which the upstream gas is pre-heated and accelerated over a characteristic distance $\\kappa/\\vs$, where $\\kappa$ is a momentum averaged CR diffusion coefficient, and $\\vs$ is the shock speed. The value of $\\kappa$ for CR ions depends on the spectrum of the magnetic wave-field, thought to be generated by the CR streaming instability \\citep{Bell1978a}, and has not been well constrained by observations yet. While in the general ISM $\\kappa\\sim3$--$5\\times10^{28}\\cms$ \\citep{SMP2007CR}, at the forward shocks of SNRs $\\kappa$ is thought to be close to the Bohm diffusion limit $\\kappaB\\approx3\\times10^{22}\\,\\mathrm{cm}^2\\,\\mathrm{s}^{-1}\\,\\left(\\beta\\,B_0/\\mu{}\\mathrm{G}\\right)^{-1}\\left(p/\\mathrm{GeV}\\,\\mathrm{c}^{-1}\\right)$, where $\\beta$ is the ratio of particle speed to the speed of light, $B_0$ is the large scale magnetic field, and $p$ is the CR particle momentum \\citep{Drury1983}. $\\kappaB$ is the lower limit to $\\kappa$ allowed by the standard theory of DSA, and corresponds to saturated field fluctuations, $\\delta{}B/B_0=1$. \\citet{SGLS2003} estimated upper limits for the diffusion coefficient of CR ions in several SNRs in the range $\\kappa\\sim10^{25}\\cms$--$2\\times10^{27}\\cms$ from the condition that neutrals must survive ionization while experiencing the amount of heating implied by the \\Ha{} narrow component linewidth. The expression derived by \\citet{PMBG2006} for the electron diffusion coefficient as a function of the synchrotron X-ray cutoff energy implies that the electron diffusion coefficient for the SNRs studied in their work is within only a factor of a few greater than the Bohm diffusion coefficient.  The spectral shape and the sharply peaked radial profiles of the synchrotron X-ray emission at the forward shocks of young SNRs require a postshock magnetic field strength of the order of $100\\muG$ \\citep{VL2003,VBK2005,Ballet2006}. Since compression alone is insufficient to produce such a gain in field strength, it is thought that perturbations in the preshock medium generated by the CR streaming instability are non-linearly amplified beyond $\\delta{}B/B_0=1$ \\citep{BL2001,Bell2004,PMBG2006}. While some simulations show that there exist rapidly growing nonresonant modes \\citep{ZPV2008}, these are saturated at $\\delta{}B/B_0\\sim1$ in other simulations \\citep{NPS2008}, and the problem of magnetic field amplification in the CR precursor remains an unresolved. Another mechanism for wave generation is the CR driven acoustic instability \\citep{Drury1984,D-F1986,Chalov1988,KJR1992}, which may play a role in prolonging the confinement of high energy CRs in the CR precursor \\citep{Berezhko1986,MD2006,DM2007}.   In general, strong MHD turbulence in the CR precursor leads to energy dissipation by wave damping that will affect the structure of the shock \\citep{MV1982,CBAV2008}. Although theories of wave damping exist \\citep[see e.g.][]{Whang1997}, neither the rate of energy dissipation nor the fractions of the dissipated energy going into internal energy of the various components of the flow (CR electrons and ions, and thermal electrons and ions) are known from observations. In models of DSA, the original treatment of Alfv\\'{e}n wave damping in CR modified shocks by \\citet{VM1981} is commonly adopted. The damping of acoustic waves, though less common in models of DSA, may also substantially heat the ions \\citep{D-F1986} or electrons \\citep{GLR2007} of the thermal component.  A further quantity important for DSA is the efficiency of injection of particles from the thermal population into the CR population. The fraction of swept up thermal protons injected into the acceleration process in SNRs is thought to lie in the range $10^{-4}$--$10^{-2}$ \\citep{VBK2003,EC2005} if particle acceleration is efficient.  If the structure of a non-radiative shock is modified by CRs, several subtle signatures in the optical emission from the shock are expected \\citep{Raymond2001}. One signature would be a FWHM of the narrow \\Ha{} component broader than $20\\kms$, the value expected if there is no CR acceleration for an upstream medium at a temperature of $T_0\\sim10^{4}\\Kv$. Spectra of many Balmer-dominated filaments associated with shocks over a wide range of Mach numbers, show a narrow \\Ha{} component with a FWHM in the range $30$--$50\\kms$ \\citep[see][Table 1]{SGLS2003}, indicative of some common form of preshock heating. Significant bulk acceleration of the upstream flow through the precursor would also be detectable as a Doppler shift of the narrow component centroid with respect to \\Ha{} emission from the upstream gas. With the exception of the filament observed by L07, which also concerns this work, this has not been observed yet. Currently, a CR precursor is the favored mechanism for the inferred preshock heating (\\citealp{HRB1994,SRL1994,SGLS2003}; L07), but no self-consistent model of such a precursor has been compared with data. Here we show that the precursor structures predicted by two-fluid models of CR modified shocks including particle injection at the subshock and energy transfer due to the acoustic instability are consistent with the above features of \\Ha{} spectra.  Recently, L07 obtained high-resolution \\Ha{} echelle spectra of an optical filament at knot g in Tycho's SNR, covering the postshock region and the ionization precursor far upstream. Knot g, located in the eastern rim ($\\alpha=00\\ahour25\\aminute56.5\\asecond$, $\\delta=64^\\circ09^\\prime28\\arcsec$, J2000.0), is the brightest region in \\Ha{} emission in the remnant. The synchrotron X-ray emission is also particularly bright in this region \\citep{Decourchelle-etal2001,Hwang-etal2002}. Radio data and \\HI{} absorption studies suggest that the northeastern rim is decelerating into an inhomogeneous ambient medium, possibly the edges of a molecular cloud \\citep[see][and references therein]{LKT2004}. The observations by L07 have spatially resolved a steep \\Ha{} (narrow component) flux increase ahead of the shock discontinuity, distinct from the photoionization precursor (G00), which L07 attribute to enhanced emission from a CR precursor. They also reported a broadening of the narrow component linewidth by $15\\kms$ and a redward Doppler shift of the narrow component centroid with respect to that of the distant upstream \\Ha{} emission in this region of $5\\kms$.  In this paper, we provide a self-consistent CR modified shock model applied to the observational data from L07. We model the shock structure with a time-dependent hydrodynamic two-fluid code. We adjust model parameters to obtain a best fit for the calculated spatial \\Ha{} profile to the observed profile. The two-fluid equations along with the free parameters and boundary conditions of the shock model are described in Sect.~\\ref{sec:hydro}. The method of calculation for the \\Ha{} emissivity is given in Sect.~\\ref{sec:halpha}. A time dependent transient solution that provides the best fit to the observed spatial \\Ha{} profile is presented in Sect.~\\ref{sec:trans}. In Sect.~\\ref{sec:steady}, we discuss the solution space of steady CR modified shocks, and propose that the branch of solutions for which the CR acceleration efficiency is low, may explain the line broadening of the narrow component of the \\Ha{} line, observed in many SNRs. We discuss our results in Sect.~\\ref{sec:discussion}, and conclude the paper in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In summary, CR acceleration in the forward shocks of SNRs results in the heating and acceleration of the preshock medium which may explain some features of the optical emission of Balmer dominated filaments. We have found a transient state in the evolution of a shock from one that is initially not modified by CRs to one that is CR dominated, for which the calculated \\Ha{} emissivity profile matches the emissivity profile across the Balmer-dominated filament in knot g observed by \\citet{Lee-etal2007}. The values of the parameters for this shock model are an initial shock speed $\\vs=2000\\kms$, a distant upstream CR pressure to thermal gas pressure ratio $\\phi_0=1$, a diffusion coefficient $\\kappa=2\\times10^{24}\\cms$, an energy transfer time-scale due to the acoustic instability $\\tau=426\\yr$, and a lower limit to the injection parameter $\\epsilon=4.2\\times10^{-3}$.  The structure of steady shocks that belong to the low CR acceleration efficiency branch of solutions for fast shocks are relatively insensitive to the values of $\\tau$, $\\epsilon$, $\\kappa$, and $\\vs$ in the range $300\\kms<\\vs<3000\\kms$, provided that the parameters are chosen such that a steady solution in the low CR acceleration efficiency branch exists. The solutions are usually reached as time-asymptotic states within less than $100\\yr$, even if $\\epsilon=0$. The mild heating of the preshock gas up to typically $2$--$6\\times10^4\\Kv$, and the negligible bulk acceleration of the flow in the precursor $(\\Delta{}u\\leq10\\kms)$ may provide a natural explanation for the characteristic broadening of the narrow component linewidth that is observed to lie in the small range $\\mathrm{FWHM}=30$--$50\\kms$ in many SNRs, and the lack of bulk Doppler shift of the narrow component observed for these cases."
    },
    "0809/0809.0392_arXiv.txt": {
        "abstract": "We present Very Large Array (VLA) observations of the water maser emission towards IRAS 16552-3050. The maser emission shows a velocity spread of $\\simeq 170$ km~s$^{-1}$, and a bipolar distribution with a separation between the red and blueshifted groups of $\\simeq 0.08''$. These observations and the likely post-AGB nature of the source indicate that IRAS 16552-3050 can be considered as a member of the ``water fountain\" class of sources (evolved stars showing H$_2$O maser emission with a velocity spread $\\ga 100$ km~s$^{-1}$, probably tracing collimated jets).  The water maser emission in IRAS 16552-3050 does not seem to be associated with with any known optical counterpart. Moreover, this source does not have a near-IR 2MASS counterpart, as it happens in about half of the water fountains known. This suggests that these sources tend to be heavily obscured objects, probably with massive precursors ($\\ga 4-5$ M$_\\odot$). We suggest that the water maser emission in IRAS 16552-3050 could be tracing a rapidly precessing bipolar jet. ",
        "introduction": "Water fountains are evolved Asymptotic Giant Branch (AGB) and young post-AGB stars that show water maser emission with velocity separations in their spectral features $\\ge$100~km~s$^{-1}$ \\citep{likkel88}. Some of them have been observed with radio interferometric techniques and all show bipolar distributions of their H$_2$O maser emission, indicating the presence of jets with extremely short dynamical ages of 5-100 yr \\citep{likkel88,boboltz05,boboltz07,imai02,imai04,imai07iau}. OH maser emission has been detected in most sources catalogued as water fountains and, in general, this emission exhibits lower velocities and is spatially less extended than that of H$_2$O \\citep[see, for example][]{deacon07}. Water fountains are evolutionarily located between two phases of stellar evolution, AGB stars and planetary nebulae (PNe), when mass-loss properties, among others, change dramatically. While mass ejection during the AGB phase is spherical, most PNe exhibit elliptical, bipolar, or multipolar morphologies. These morphologies have been attributed to the action of collimated outflows on the previously expelled spherical AGB shell \\citep{sahai98}. Therefore, water fountains very probably trace the onset of non-spherical and highly collimated mass ejection. Understanding the transformation of an AGB star into a PN relies on a precise knowledge of the properties and evolution of water fountains. Unfortunately, very few water fountains are known at present (11 sources reported as of March 2008, see \\citealt{imai07iau}), so that adding new members to this important class of objects represents a valuable progress.   Recently, in a single-dish search for H$_2$O masers in evolved stars, \\citet{suarez07} found water maser emission toward IRAS\\,16552$-$3050 (hereafter IRAS16552), with a maximum separation among its velocity components of $\\sim$170~km~s$^{-1}$. IRAS16552 was first proposed to be a candidate post-AGB star by \\citet{preite88}, based on its IRAS colors. A possible optical counterpart of the IRAS far-infrared source was classified by \\citet{hu93} and \\citet{suarez06} as a star of spectral type K9III and K0I, respectively.  In this paper, we present Very Large Array (VLA) observations, carried out to confirm the association of this high-velocity water maser emission with the IRAS source, and to determine the spatial distribution of the masers. Our results allow us to include IRAS16552 as a bona fide member of the water fountain class.  In Sec.~\\ref{obs} we describe the performed observations. Sec.  \\ref{results} presents the obtained results, and Sec. \\ref{discussion} a discussion on the properties of IRAS16552 and a comparison of this source with other water fountains. Finally, in Sec. \\ref{conclusions} we summarize the main conclusions of this work. ",
        "conclusions": "\\label{conclusions}  In this paper we have presented VLA observations of H$_2$O masers at 22 GHz towards IRAS 16552$-$3050.  Our main conclusions are as follow:  \\begin{itemize}  \\item We have found that the water maser emission in IRAS 16552$-$3050 span a velocity range of $\\sim$170~km~s$^{-1}$ and shows a bipolar distribution of $\\simeq$ 0.08'' in size. The properties of the water maser emission and the likely post-AGB nature of the central source allow us to confirm IRAS 16552$-$3050 as a member of the class of ``water fountain'' evolved objects.  \\item Our result confirms the trend of bipolarity in all water fountains known up to now, strongly suggesting that water masers in these sources are produced in bipolar, collimated jets.  \\item The water maser emission in IRAS 16552$-$3050 does not seem to be associated with the optical source identified by \\citet{hu93} and \\citet{suarez06}, but with a different evolved object, for which there is no optical nor near-IR counterpart known.  \\item About half of the water fountains known do not have a near-IR counterpart in the 2MASS catalog. This suggest that water fountains are surrounded by thicker envelopes than the rest of AGB and post-AGB stars (which usually have near-IR counterparts), probably implying that they are relatively massive objects ($\\ga 4-5$ M$_\\odot$).  \\item The distribution of the water maser emission in the red- and blueshifted group is almost perpendicular to the proposed jet direction. This might be explained if the maser emission traces a rapidly precessing/rotating bipolar jet. A later evolution could give rise to point-symmetric filaments observed in some PNe, which appear oriented almost perpendicular to the radial direction from the central star.  \\end{itemize}"
    },
    "0809/0809.5054_arXiv.txt": {
        "abstract": "We present spatially resolved near-IR spectroscopic observations of 15 young stars.  Using a grism spectrometer behind the Keck Interferometer, we obtained an angular resolution of a few milli-arcseconds and a spectral resolution of 230, enabling probes of both gas and dust in the inner disks surrounding the target stars. We find that the angular size of the near-IR emission typically increases with wavelength, indicating hot, presumably gaseous material within the dust sublimation radius.  Our data also clearly indicate Br$\\gamma$ emission arising from hot hydrogen gas, and suggest the presence of water vapor and carbon monoxide gas in the inner disks of several objects.  This gaseous emission is more compact than the dust continuum emission in all cases.  We construct simple physical models of the inner disk and fit them to our data to constrain the spatial distribution and temperature of dust and gas emission components. ",
        "introduction": "} Protoplanetary disks provide a reservoir of material from which planets may form, and the abundance and properties of extra-solar planets \\citep[e.g.,][]{MARCY+05}, as well as the architecture of our own solar system, suggest that planets frequently form in or migrate through inner regions of protoplanetary disks.  Furthermore, the innermost disk regions represent the interface between the inwardly accreting disk and the magnetized central star, and it is here that material accretes inward or is launched in winds or outflows \\citep{SHU+94}.  Knowledge of the distribution of material in the inner disk is therefore crucial for understanding  the mass assembly and angular momentum evolution of pre-main-sequence stars.  Modeling of spectral energy distributions \\citep[e.g.,][]{BBB88,LA92} and spectrally resolved gaseous emission lines \\citep[e.g.,][]{NAJITA+96,BB04,NAJITA+06} provide important insights into the structure of protoplanetary disks in the terrestrial planet forming region. However, since these techniques typically rely on spectral information as a proxy for spatial information, they require assumptions about underlying geometric, temperature, or velocity structure.  Recently, the technique of spectro-astrometry, which capitalizes on the fact that emission centroids can be measured more precisely than the available angular resolution, has enabled less ambiguous constraints on disk structure \\citep{PONTOPPIDAN+08}. However, substantial gaps in our understanding remain due to a lack of high angular resolution observations.  Near-IR interferometry, which synthesizes a large aperture using two or more smaller, separated apertures, can achieve orders of magnitude higher angular resolution than conventional telescopes, and can spatially resolve disk terrestrial regions.  For example, the Keck Interferometer, which combines the light from the two 10-m Keck apertures over an 85-m baseline, achieves a resolution approximately an order of magnitude higher than that attained with a single aperture. This means that angular scales of a few milli-arcseconds, corresponding to a few tenths of an AU at typical distances to nearby star-forming regions, can be spatially resolved with near-IR interferometers.  Most interferometric measurements to date have probed only inner disk dust \\citep[e.g.,][and references therein]{MILLAN-GABET+07}, which typically dominates the near-IR emission.  To distinguish gas and dust, spectrally dispersed observations are required. Very few sources have been observed with spectrally dispersed interferometric observations to date \\citep{EISNER+07a,EISNER07,MALBET+07,TATULLI+07,KPO08,ISELLA+08}. The observations showed intriguing evidence that gas and dust are not distributed uniformly in inner disk regions.  Moreover, while observations of stars less massive than a few M$_{\\odot}$ found gaseous emission to be more compact than dust emission \\citep{EISNER+07a,EISNER07}, observations of the Br$\\gamma$ emission line around two more luminous stars found the gas to be more extended than the dust, presumably because the Br$\\gamma$ emission traces outflows from these young systems \\citep{MALBET+07,TATULLI+07}.  A larger sample is required to further investigate such potential trends, and to constrain the general properties of inner disk gas in young stars.  Here we present spectrally dispersed near-IR interferometry observations of a sample of young stars, including four T Tauri stars and 11 Herbig Ae/Be stars.  Our data constrain the relative spatial and temperature distributions of dust and gas in sub-AU-sized regions of the disks around these stars. ",
        "conclusions": "} We presented spatially resolved near-IR spectroscopic observations that probed the gas and dust in the inner disks around 15 young stars. One source, HD 141569, was unresolved at all wavelengths between 2.0 and 2.4 $\\mu$m, indicating a lack of dust or gas in the inner disk regions.  Another target, VV Ser, was over-resolved, indicating that the emission spans angles larger than $\\sim 5$ mas at all observed wavelengths.  The near-IR emission from the remaining targets was resolved, and our data show that the angular size of the near-IR emission increases with wavelength in all cases. This behavior suggests temperature gradients in these inner disks, arising from the combination of warm dust at its sublimation temperature and hotter, presumably gaseous material within the dust sublimation radius.  Our data clearly indicate emission from the Br$\\gamma$ transition of hydrogen in several objects, and suggest that water vapor and carbon monoxide gas are present in the inner disks of some targets.  We constructed simple physical models of the inner disk, including dust and gas emission, and we fitted them to our data to constrain the spatial distribution and temperature of dust and gas emission components.  We considered models including only dust emission; dust, gas continuum, and Br$\\gamma$ emission; and dust, gas continuum, water vapor, and Br$\\gamma$ emission.  Models incorporating only dust emission can not fit the data for any of our sources well.  In contrast, models including dust and gas emission are suitable for explaining our data. The inclusion or exclusion of water vapor in these dust+gas models did not substantially affect the quality of the fits in most cases.  For all sources where Br$\\gamma$ emission is observed, we find it to be compact relative to the continuum emission.  This contrasts with previous findings, which found Br$\\gamma$ emission to be extended relative to the continuum around some high-mass stars.  The results presented here suggest that Br$\\gamma$ commonly traces infalling material around young stars spanning a large range in stellar mass.  CO emission is tentatively observed towards several objects, and we see evidence that this emission has a more compact spatial distribution than the dust around RW Aur.  For other objects, our data are insufficient to place meaningful constraints on the relative spatial distribution of CO and other emission components. We will re-observe these targets in the near future with higher dispersion, to obtain better signal-to-noise for the relatively narrow CO lines and better constrain their spatial distribution.  While models including water vapor opacity often fit our data well, the best-fit models generally also require continuum emission from material that is too hot to be water (since water dissociates at $\\sim 3000$ K). We do not have a ready explanation for the source of this hot continuum emission, but we speculate that it may trace free-free emission from hydroden and/or H$^-$.  The gas densities and fractional ionizations required to produce such emission seem plausible in the inner regions of protoplanetary disks, suggesting that free-free emission from H and H$^-$ is a viable explanation for the compact continuum emission seen in our data.   \\medskip  Data presented herein were obtained at the W. M. Keck Observatory, in part from telescope time allocated to the National Aeronautics and Space Administration through the agency's scientific partnership with the California Institute of Technology and the University of California. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognize and acknowledge the cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. This work has used software from the Michelson Science Center at the California Institute of Technology. The authors thank the entire Keck Interferometer team for making these observations possible.  We also wish to thank the referee, Geoff Blake, for his thoughtful and detailed referee report, which greatly improved the manuscript."
    },
    "0809/0809.2990_arXiv.txt": {
        "abstract": "Annihilation of Dark Matter usually produces together with gamma rays comparable amounts of electrons and positrons. The $e^+e^-$ gyrating in the galactic magnetic field then produce secondary synchrotron radiation which thus provides an indirect mean to constrain the DM signal itself. To this purpose, we calculate the radio emission from the galactic halo as well as from its expected substructures  and we then compare it with the measured diffuse radio background. We employ a multi-frequency approach using data in the relevant frequency range 100 MHz--100 GHz, as well as the WMAP Haze data at 23 GHz. The derived constraints are of the order $\\sigva= 10^{-24}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV sensibly depending however on the astrophysical uncertainties, in particular on the assumption on the galactic magnetic field model. The signal from single bright clumps is instead largely attenuated by diffusion effects and offers only poor detection perspectives. ",
        "introduction": "Cosmology and Astrophysics provide nowadays a compelling evidence of the existence of Dark Matter (DM) \\cite{Komatsu:2008hk,Bertone:2004pz}. Nevertheless, its nature still remains elusive, and Dark Matter constituents have escaped a direct detection in laboratory so far. Promising candidates are DM particles produced in thermal equilibrium in the early universe, the so-called Weakly Interacting Massive Particles (WIMPs). Theoretically, models of WIMPs naturally arise, for example, in SUSY as the Lightest Super-symmetric Particle or as the Lightest Kaluza-Klein Particle in the framework of extra-dimensions. These candidates are self-conjugate and can thus annihilate in couples to produce as final states: neutrinos, photons, electrons, light nuclei (as wells as their antiparticles), etc., which can in principle be detected.  Among the indirect DM detection channels, gamma-ray emission represents one of the most promising opportunity due to the very low attenuation in the interstellar medium, and to its high detection efficiency. See for example Ref.s \\cite{Bertone:2004pz,Jungman:1995df,Bergstrom:2000pn} for a review of this extensively studied issue. The expected neutrino detection rates are generally low although  forthcoming km$^3$ detectors offer some promising prospect \\cite{Bergstrom:1998xh,Barger:2001ur}. Finally, positrons and protons strongly interact with gas, radiation and magnetic field in the galaxy and thus the expected signal sensibly depends on the assumed propagation model \\cite{Baltz:1998xv,Hooper:2004bq,Donato:2003xg}. However, during the process of thermalization in the galactic medium the high energy $e^+$ and $e^-$ release secondary low energy radiation, in particular in the radio and X-ray band, that, hence, can represent a chance to look for DM annihilation. Furthermore, while the astrophysical uncertainties affecting this signal are similar to the case of direct $e^+$, $e^-$ detection, the sensitivities are quite different, and, in particular, the radio band allows for the discrimination of tiny signals even in a background many order of magnitudes more intense.  Indirect detection of DM annihilation through secondary photons has received recently an increasing attention, exploring the expected signature in X-rays \\cite{Bergstrom:2006ny,Regis:2008ij,Jeltema:2008ax}, at radio wavelengths \\cite{Blasi:2002ct,Aloisio:2004hy,Tasitsiomi:2003vw,Zhang:2008rs} , or both \\cite{Colafrancesco:2005ji,Colafrancesco:2006he,Baltz:2004bb}. In the following we will focus our analysis on the radio signal expected from the Milky Way (MW) halo and its substructures. It is worth noticing that the halo signal has been recently discussed in Ref.s \\cite{Hooper:2007kb,Hooper:2008zg,Grajek:2008jb} in connection to the WMAP Haze, which has been interpreted as a signal from DM annihilation. In this concern we will take in the following a more conservative approach, by assuming that the current radio observations are entirely astrophysical in origin, and thus deriving constraints on the possible DM signal. The main point will be the use of further radio observations besides the WMAP ones, in the wide frequency range 100 MHz-100 GHz, and a comparison of the achievable bounds. Furthermore, the model dependence of these constraints on the assumed astrophysical inputs will be analyzed. We will also discuss the detection perspectives of the signal coming from the brightest DM substructures in the forthcoming radio surveys.  The paper is organized as follows: in section \\ref{AstrInput} we will discuss the astrophysical inputs required to derive the DM signal such as the structure of the magnetic field, the DM spatial distribution and the radio data employed to derive the constraints. In section \\ref{synchsignal} we describe in detail the processes producing the DM radio signal either when it is originated from the halo or from the substructures. In section IV we present and discuss our constraints, while in section V we analyze the detection sensitivity to the signal coming from the single DM clump. In section VI we give our conclusions and remarks. ",
        "conclusions": "Using conservative assumptions for the DM distribution in our galaxy we  derive the expected secondary radiation due to synchrotron emission from high energy electrons produced in DM annihilation. The signal from single bright clumps offers only poor sensitivities because of diffusion effects which spread the electrons over large areas diluting the radio signal. The diffuse signal from the halo and the unresolved clumps is instead  relevant and can be compared to the radio astrophysical background to derive constraints on the DM mass and annihilation cross section.  Constraints in the radio band, in particular, are complementary to similar (less stringent but less model dependent) constraints in the X-ray/gamma band \\cite{Mack:2008wu,Kachelriess:2007aj} and from neutrinos \\cite{Yuksel:2007ac}. Radio data, in particular, are more sensitive in the GeV-TeV region while neutrinos provide more stringent bounds for very high DM masses ($\\agt 10$ TeV). Gammas, instead, are more constraining for $m_\\chi \\alt 1$ GeV. The combination of the various observations provides thus interesting constraints over a wide range of masses  pushing the allowed window significantly near the thermal relic possibility.    More into details, we obtain conservative constraints at the level of $\\sigva \\sim 10^{-23}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV from the WMAP Haze at 23 GHz. However, depending on the astrophysical uncertainties, in particular on the assumption on the galactic magnetic field model, constraints as strong as $\\sigva \\sim 10^{-25}$\\,cm$^3$s$^{-1}$ can be achieved. Complementary to other works which employ the WMAP Haze at 23 GHz, we also use  the information in a wide frequency band in the range 100 MHz-100 GHz. Adding this information the constraints become of the order of $\\sigva \\sim 10^{-24}$\\,cm$^3$s$^{-1}$ for a DM mass $m_{\\chi}=100$\\,GeV. The multi-frequency approach thus gives comparable constraints with respect to the WMAP Haze only, or generally better for $m_\\chi\\alt100$ GeV where the best sensitivity is achieved at $\\sim$ GHz frequencies.   The derived constraints are quite conservative because no attempt to model the astrophysical background is made differently from the case of the WMAP Haze. Indeed, the Haze residual map itself should be interpreted with some caution, given that the significance of the feature is at the moment still debated and complementary analyses from different groups (as the WMAP one) miss in finding a clear evidence of the feature. In this respect, the multifrequency approach will be definitely necessary to clarify the nature of controversial DM signals as in the case of the WMAP Haze. Progresses are expected with the forthcoming data at high frequencies from Planck and at low frequencies from LOFAR and, in a more distant future, from SKA. These surveys will help in disentangling the various astrophysical contributions thus assessing the real significance of the Haze feature. Further, the low frequency data in particular, will help to improve our knowledge of the galactic magnetic field. Progresses in these fields will provide a major improvement for the interpretation of the DM-radio connection.   \\vspace{1pc}"
    },
    "0809/0809.4052_arXiv.txt": {
        "abstract": "We propose to use a simple observable, the fractional area of ``hot spots'' in weak gravitational lensing mass maps which are detected with high significance, to determine background cosmological parameters. Because these high-convergence regions are directly related to the physical nonlinear structures of the universe, they derive cosmological information mainly from the nonlinear regime of density fluctuations. We show that in combination with future cosmic microwave background anisotropy measurements, this method can place constraints on cosmological parameters that are comparable to those from the redshift distribution of galaxy cluster abundances.  The main advantage of the statistic proposed in this paper is that projection effects, normally the main source of uncertainty when determining the presence and the mass of a galaxy cluster, here serve as a source of information. ",
        "introduction": "\\label{sec:introduction} Weak gravitational lensing (WL), i.e., the coherent distortion of images of faint distant galaxies by the gravitational tidal field of the intervening matter distribution \\citep{twv90}, has been established as a powerful cosmological tool [see, e.g., \\citet{rev} for a review].  Since this effect is purely gravitational, it directly probes the matter distribution along the line of sight, thus providing a way to understand the nature and the evolutionary history of the universe that is relatively insensitive to how light or baryons trace dark matter.  Since the earliest measurements of WL by galaxy clusters \\citep{for88,twv90}, WL by large-scale structure, known as ``cosmic shear'', has been detected based on optical \\citep{cs00a,cs00b,cs00c,cs00d} and radio observations \\citep{crh04}.  Future WL surveys, such as the {\\it Large Synoptic Survey Telescope} (LSST)\\footnote{www.lsst.org}, will measure the cosmic shear field with great precision over half the sky.  Recent attention has focused on how best to extract information from the shear/convergence field in order to constrain cosmological parameters.  The primary focus so far has been using the standard two-point statistics \\citep{bsbv91,mir91,kai92}, which probes the underlying matter power spectrum in projection.  In this paper, we propose to use a simple, direct observable: the fraction of high signal-to-noise ratio ($S/N$) points detected in WL surveys, as another discriminator of cosmology.  We utilize the results of an N-body simulation to quantify both the theoretical predictions and the observational uncertainties of this statistic.  As is well-known, the common two-point statistics do not contain all the statistical information of the WL convergence field, as the nonlinear gravitational instability induces non-Gaussian signatures in the mass distribution and hence in the WL convergence field as well. Unfortunately, there is no complete statistical analysis in practice for the underlying matter density field in the nonlinear regime.  Previous studies of weak lensing statistics on small angular scales, correspondingly, were restricted to the low-order statistics, using the halo model [see, e.g., \\citet{cs02} for a review], or the ``scaling ansatz'' \\citep{hmkl}, later extended and calibrated by N-body simulations \\citep{jmw95,pd96,smi03}, for the nonlinear evolution of clustering.  These low-order statistics include: the two-point correlation function or equivalently the power spectrum, at small scales \\citep{js97}, the three-point correlation function or the bispectrum \\citep{tj03a,tj03b,tj04}, the third-order moment: skewness \\citep{bvm97,js97,hui99} and the fourth-order moment: kurtosis \\citep{tj02}.  An alternative approach is to use the redshift distribution of nonlinear object abundance \\citep{hmh01}, such as shear-selected galaxy cluster samples \\citep{wk03,us,fh07}, using the Press-Schechter prescription \\citep{ps,bcek} with calibration by N-body simulations \\citep{st99,jen01}.  In addition, the ``ratio-statistic'' \\citep{jt03,bj04,sk04,hj04,zhs05} has been constructed to use the geometrical information from tomographic shear/convergence power spectra on small scales.  These analyses have shown that comparable information may be contained in the linear and in the nonlinear regime.  The statistics mentioned above are by no means a full characterization of the convergence field, yet they all face challenges, either from the theoretical or the observational side, which need to be addressed. Including higher-order statistics might give a substantial increase in information, but they are in practice noisy and computationally intensive.  To take advantage of the synergy between these statistics, one needs to account for their covariance, which is difficult to calculate analytically.  One hope is to run large simulations and directly measure this covariance.  Galaxy cluster samples selected optically, by their X-ray flux or by their Sunyaev-Zel'dovich effect signatures, are limited by the uncertain astrophysics when modeling the mass-observable relations and require ``self-calibration'' \\citep{mm03,us,lh05}, as the mass function is exponentially sensitive to errors of limiting mass.  Shear-selected galaxy clusters, on the other hand, have the advantage that the selection function can be determined ab initio by N-body simulations, since only gravity is involved. However, projection effects result in false detections, missing clusters \\citep{wvm02,hty04,hs05}, and producing significant uncertainty in the cluster mass derived from the shear signals \\citep{mwnl99,dw05}, degrading the cosmological information content.  In this paper, we instead focus on the one-point probability distribution function (PDF) of the WL convergence field.  There are three main motivations.  First, the one-point PDF is a simple yet powerful tool to probe non-Gaussian features, and since the non-Gaussianity in the convergence field is induced by the growth of structure, it holds cosmological information \\citep{rkjs99,jsw00,ks00,vmb05}.  These previous works have shown that the one-point PDF is capable of discriminating cosmologies with different $\\Omega_m$, such as an open cold dark matter (CDM) model, a flat cosmological constant-dominated ($\\Lambda$)-CDM model, and a standard CDM model.  Second, and the primary motivation of this work, is that the fractional area statistic we propose in this paper takes into account projection effects by construction.  Determined by the high-convergence tail of the PDF, it is an analog of Press-Schechter formalism, thus similar to the abundance of galaxy clusters but without contamination due to projection effects.  This statistic also utilizes information mainly from the nonlinear regime and complements the well-established statistics in the linear regime. The goal of our work is to give a more quantitative assessment of the statistical information in the nonlinear regime (particularly focusing on the properties of dark energy) provided by this fractional area statistic and its complementarity to other probes of cosmology, e.g., the cosmic microwave background (CMB) anisotropies.  Third, besides utilizing the cosmic shear field \\citep{ztp05} which we will focus on in this paper, there are several other observational techniques which could be used to map out the convergence PDF.  For example, with forthcoming large samples of high redshift supernovae from LSST-like or Joint Dark Energy Mission (JDEM)-like surveys, one could measure the magnification distribution of these standard candles due to lensing by the large scale structure in the foreground, and construct the convergence PDF \\citep{dv06,chh06}; it is also possible to measure the convergence field through the statistics of cosmic magnification \\citep{j02}, using, for instance, 21cm-emitting galaxies \\citep{zpp05,zpp06}.  The rest of this paper is organized as follows.  Our basic calculational methodology, including a calibration of the (co)variance using simulation outputs, is described in \\S~\\ref{sec:II}.  Results of using the fractional area statistic for an LSST-like WL survey are presented in \\S~\\ref{sec:III}, with discussions of its complementarity to other dark energy probes, as well as of various uncertainties.  Conclusions and implications of this work are given in \\S~\\ref{sec:IV}.  Finally, in a series of appendices, we show the details of our calculations. ",
        "conclusions": "\\label{sec:IV} We have shown that, in future wide field weak gravitational lensing surveys, a simple one-point statistic -- the total fractional area of high $S/N$ points in the convergence field -- is a promising probe of cosmology.  It is sensitive to the total matter content of the universe and the amplitude of the density fluctuations, which helps breaking the intrinsic degeneracies in the growth of structure between cosmological parameters, thus constraining the properties of dark energy.  The main conclusion of this work is that the fractional area statistic provides constraints on cosmological parameters similar to the redshift distribution of galaxy cluster abundance, but without suffering from the projection effects.  Indeed, the statistic is explicitly constructed to take advantage of projections as part of the signal.  We expect the fractional area statistic will help achieve the goal of ``precision cosmology'' and shed light on the mystery of dark energy."
    },
    "0809/0809.1993_arXiv.txt": {
        "abstract": "We report the discovery of four dusty cometary tails around low mass stars in two young clusters belonging to the W5 star forming region. Fits to the observed emission profiles from 24 $\\micron$ observations with the {\\it Spitzer} Space Telescope give tail lifetimes $<$ 30 Myr, but more likely $\\lesssim$ 5 Myr. This result suggests that the cometary phase is a short lived phenomenon, occurring after photoevaporation by a nearby O star has removed gas from the outer disk of a young low mass star \\citep[see also][]{balog06,balog08}. ",
        "introduction": "Recent observations with the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) have discovered dusty nebulae associated with several young stars in three Galactic star forming regions \\citep{balog06}. Appearing as extended, cometary objects pointing away from nearby O stars at 24 $\\micron$, (and sometimes at 8 $\\micron$), they resemble the proplyds seen in Orion with HST \\citep{odell93,odell94}. \\citet{balog06} show that these cometary tails are probably dust swept out of a young circumstellar disk by an O star.  \\citet{balog06} show that the emission profiles at 24 and 8 $\\micron$ can be explained by a model where gas is photoevaporated from a young protoplanetary disk. As the gas leaves, entrained dust also escapes and is then blown away from the disk by radiation pressure from nearby O stars. The dust is swept into a cometary morphology which glows in the mid-IR thermally as it reprocesses the incident UV radiation to longer wavelengths. \\citet{balog06} cite \\citet{johnstone98}, \\citet{richling00}, \\citet{hollenbach04} and \\citet{matsuyama03}, who find that the portion of a protoplanetary disk likely responsible for emission at 24 $\\micron$ ($\\sim$2--50 AU from the star) is removed on timescales $\\sim$3$\\times 10^5$ yr.  We have recently discovered four dusty cometary structures in the W5 star forming region. Here we expand on the work of \\citet{balog06} and place constraints on the lifetimes and physical nature of these interesting objects. ",
        "conclusions": "We have discovered four dusty cometary tails in W5 at 24 $\\micron$ with {\\it Spitzer} MIPS. An apparent deficit of gas suggests an origin in radiation pressure blowout of dust from a young stellar disk by nearby O stars. Future observations with high sensitivity, high spatial resolution mid-infrared and sub-millimeter instruments will help constrain their nature.  {\\it Facilities:} \\facility{2MASS ($JHK_S$)}, \\facility{Spitzer (IRAC, MIPS)}"
    },
    "0809/0809.3819_arXiv.txt": {
        "abstract": "In recent work we presented the first results of global general relativistic magnetohydrodynamic (GRMHD) simulations of tilted (or misaligned) accretion disks around rotating black holes. The simulated tilted disks showed dramatic differences from comparable untilted disks, such as asymmetrical accretion onto the hole through opposing ``plunging streams'' and global precession of the disk powered by a torque provided by the black hole. However, those simulations used a traditional spherical-polar grid that was purposefully underresolved along the pole, which prevented us from assessing the behavior of any jets that may have been associated with the tilted disks. To address this shortcoming we have added a block-structured ``cubed-sphere'' grid option to the Cosmos++ GRMHD code, which will allow us to simultaneously resolve the disk and polar regions. Here we present our implementation of this grid and the results of a small suite of validation tests intended to demonstrate that the new grid performs as expected. The most important test in this work is a comparison of identical tilted disks, one evolved using our spherical-polar grid and the other with the cubed-sphere grid. We also demonstrate an interesting dependence of the early-time evolution of our disks on their orientation with respect to the grid alignment. This dependence arises from the differing treatment of current sheets within the disks, especially whether they are aligned with symmetry planes of the grid or not. ",
        "introduction": "\\label{sec:intro}  We have recently undertaken a series of numerical studies of titled accretion disks around rapidly rotating black holes, first in the hydrodynamic \\citep{fra05} and then in the magnetohydrodynamic (MHD) \\citep{fra07b} limits. All of these simulations have been fully general relativistic, using the Kerr-Schild metric to represent the spacetime of the black hole.  Tilted accretion disks are of particular interest because they are subject to differential warping due to the Lense-Thirring precession of the rotating black hole. For very thin disks, close to the black hole the competition between the differential twisting and ``viscous'' damping causes the angular momenta of the disk and hole to align. Further out in the disk, beyond some warp radius, the disk maintains its misaligned state.  For moderately thin to thick disks, such as the ones we simulated previously, the situation is more complex and interesting. The primary difference is that warping is transported via bending waves rather than diffusively, as for thin disks. One consequence of this is that the midplane of a thick disk does not tend to align with the symmetry plane of the black hole at small radii, as in the thin disk case. In fact, the relative tilt between the black-hole and disk angular momenta can {\\em increase} at small radii. Having the tilted disk penetrate very close to the black-hole has many interesting consequences. For instance, we found that accretion onto the hole occurs predominantly through two opposing ``plunging streams'' that start from high latitudes with respect to both the black-hole and disk midplanes \\citep{fra07a}. There is also a strong epicyclic driving within the disk attributable to the gravitomagnetic torque of the misaligned (tilted) black hole \\citep{fra08b}. The induced motion of the gas can be coherent over the scale of the entire disk. The gas also experiences periodic (twice per orbit) compressions. The compressions occur as the gas orbits past the line-of-nodes between the black-hole symmetry plane and disk midplane. Near the black hole these compressive motions can become supersonic and transform into a pair of quasi-stationary shocks. The shocks act to enhance angular momentum transport and dissipation near the hole, forcing some material to plunge toward the black hole from well outside the innermost stable circular orbit. Finally, because we are simulating disks with finite radial extents and fast sound-crossing times, the torque of the black hole causes the entire disk body to precess globally.  The main shortcoming of our work so far comes from limitations imposed on us by our use of a spherical-polar grid. First, construction of a uniform spherical polar grid in three-dimensions results in very small zones surrounding the two poles, where all of the lines of longitude converge. These very small zones constrain the Courant-limited timestep to be exceedingly small, such that the required CPU cycle count becomes prohibitively large. To avoid this problem, researchers have either excised a small conical section around each pole \\citep[e.g.][]{dev03a} or used a lower grid resolution near the poles \\citep{fra07b}. Although these techniques are reasonable when one is primarily interested in studying the equatorial region (where a disk may form), these are not satisfactory when one is interested in what is happening in the polar regions (where jets may form).  A second problem with the spherical-polar grid is that the poles themselves actually represent coordinate singularities, which present significant challenges for numerical advection and curvature coupling schemes (e.g. solving Riemann curvature source terms).  For these reasons we have added the cubed-sphere grid \\citep{sad72, ron96} as an option within our numerical code, Cosmos++.  The advantage of this grid construction is that its topological properties more closely resemble a Cartesian coordinate system than a spherical-polar system. The cubed-sphere grid uses a more uniform zone spacing than spherical polar, so the timestep can remain reasonably large even in high resolution simulations. Also important, the grid does not contain any coordinate singularities except at the origin, which is not a concern for our intended use since we truncate the grid just inside the event horizon of the black hole. Ours is not the first application of the cubed-sphere grid to problems in computational astrophysics; it has been used previously to study accretion onto rotating stars with inclined magnetic fields \\citep{kol02,rom03} and a few problems in stellar evolution \\citep{dea05,dea06}. However this is the first application of this grid to the study of black-hole accretion disks and their attendant jets.  The paper is organized as follows: \\S \\ref{sec:cubedsphere} describes the cubed-sphere mesh in detail and our particular implementation. In \\S \\ref{sec:gradtest} we discuss results of basic gradient tests on the cubed-sphere grid. In \\S \\ref{sec:disks} we compare two sets of numerical simulations of black-hole accretion disks. In the first set we compare simulations of disks accreting onto a Schwarzschild black hole. We compare different grids, resolutions, and orientations of the disk with respect to the grid. Because these simulations use a Schwarzschild black hole, the orientation should have no physical meaning. However, we show that there are, nevertheless, considerable differences in their evolution at early times. We present the case that the differences have to do with the differing treatments of the midplane current sheets in the disks, which forms from the differential winding of our initial poloidal field loops. Finally we compare simulations of tilted accretion disks around Kerr black holes. We use a tilt of $\\beta_0=15^\\circ$ and a spin of $a/M_\\mathrm{BH}=J_\\mathrm{BH}/M^2_\\mathrm{BH}=0.5$, in geometrized units where $G=c=1$, and $M_\\mathrm{BH}$ and $J_\\mathrm{BH}$ are the mass and angular momentum of the black hole, respectively. One of these simulations is run on a spherical-polar mesh, the others on the cubed-sphere grid. We demonstrate that, for the most part, the simulations agree very well. ",
        "conclusions": "In this paper we have presented our implementation of the cubed-sphere grid within Cosmos++. The cubed-sphere grid has at least three significant advantages over more-traditional grid options: 1) it has topological properties similar to a Cartesian grid, but generally conserves angular momentum much better (and nearly as well as a spherical-polar mesh); 2) it can run at a larger Courant-limited timestep than a spherical-polar mesh at comparable resolution (almost a factor of 30 at the resolution used in this work); and 3) it distributes zones more evenly than a spherical-polar mesh, which is desirable for problems where the symmetry is imperfect, such as in tilted accretion disks around rotating black holes, a problem of particular interest to us.  In Section \\ref{sec:cubedsphere} and Appendix \\ref{app:transform} we gave detailed prescriptions for the construction of the cubed-sphere grid and shared a few ``lessons learned'' in regards to extrapolating fields at inter-block boundaries and applying limiters to the field gradients in advection. After implementing these lessons ourselves, we found we could recover second order convergence over the entire grid, including along the inter-block boundaries.  To specifically demonstrate that the cubed-sphere grid is a viable option for the black-hole accretion disk work we have in mind, we have carried out a series of such simulations on our new cubed-sphere grid, using results from our spherical-polar grid as a reference standard. From these tests we conclude that:  \\begin{itemize} \\item{The cubed-sphere grid conserves angular momentum nearly as well a spherical-polar grid at comparable resolution.} \\item{The angular momentum conservation error on the cubed-sphere grid is only weakly dependent on the tilt of the disk.} \\item{Results on the cubed-sphere grid converge to the same solutions obtained on a spherical-polar grid when the two grids approach comparable resolutions. This is true for both untilted {\\em and} tilted disks.} \\item{Important disk properties such as density, pressure, specific angular momentum, inflow velocity, tilt and twist agree to better than 10-20\\% for simulations carried out on cubed-sphere and spherical-polar grids with roughly (2-3)$\\times10^6$ zones.} \\end{itemize}  During our testing, we made one surprise discovery -- that the early-time evolution was considerably different between our untilted and tilted simulations on the cubed-sphere grid. We found this to be true even for non-rotating Schwarzschild black holes, for which a tilt should have no physical meaning or significance. This is something we had not seen on the spherical-polar grid, but there we had tilted the black-hole, not the disk as we do now. We did not anticipate how important this difference would be for the initial growth and development of the $\\Omega$-dynamo and MRI.  We attribute the disparate early-time behavior to the differing ways in which the strong initial current sheet in our disk is handled numerically when it is tilted with respect to the grid. This is another reminder of the important role that numerical reconnection plays in the evolution of numerically simulated magnetized flows even though this topic is perhaps not given enough emphasis in the literature. The appearance of current sheets is virtually unavoidable in strongly sheared MHD flows such as accretion disks. One possible technique for treating the current sheets more consistently throughout the simulation may be to use an artificial resistivity. This would ensure that the current sheets are always resolved in a similar fashion regardless of their orientation with respect to the grid. However, this technique has only been implemented very recently in relativistic MHD \\citep{kom07}. An alternative, although only partial, solution might be to use a total-energy conserving scheme instead of the internal-energy conserving one used here. This, at least, guarantees that the energetics of the flow remain consistent by recapturing in the form of thermal energy any energy lost through magnetic reconnection. When coupled with a radiative cooling package, this can give a much more physical description of the evolution of the flow \\citep{fra08a}.    Although the numerical treatments of current sheets and reconnection are important to understand and appreciate, it is equally important in the context of this paper to point out that we demonstrated by numerical example that the long-term evolution of our disks is relatively unaffected by whether or not they are tilted with respect to the grid. As expected, only when the tilt is relative to a rotating black hole are there long-term implications within the disk.  We are not surprised to find significant discrepancies between our ``low'' and ``high'' resolution simulations, as previous experience had shown us that $128^3$ was roughly the minimum resolution necessary to follow the evolution of black-hole accretion disks in global general-relativistic MHD simulations such as these. Below that resolution the characteristic MRI wavelength ($\\lambda_\\mathrm{MRI} \\equiv 2\\pi v_A/\\Omega$, where $v_A$ is the Alfv\\'en speed) is not covered by a sufficient number of zones over much of the disk volume. This has nothing to do with the cubed-sphere grid itself, but is rather a universal constraint for these types of problems.  Overall we consider our experimentation with the cubed-sphere grid to be a success. In future work we will present further analysis of tilted disks (and their associated jets) evolved using this new grid option."
    },
    "0809/0809.3734_arXiv.txt": {
        "abstract": "{Peculiar velocities of clusters of galaxies can be measured by studying the fluctuations in the cosmic microwave background (CMB) generated by the scattering of the microwave photons by the hot X-ray emitting gas inside clusters. While for individual clusters such measurements result in large errors, a large statistical sample of clusters allows one to study cumulative quantities dominated by the overall bulk flow of the sample with the statistical errors integrating down. We present results from such a measurement using the largest all-sky X-ray cluster catalog combined to date and the 3-year WMAP CMB data. We find a strong and coherent bulk flow on scales out to at least $\\gsim 300 h^{-1}$Mpc, the limit of our catalog. This flow is difficult to explain by gravitational evolution within the framework of the concordance $\\Lambda$CDM model and may be indicative of the tilt exerted across the entire current horizon by far-away pre-inflationary inhomogeneities.} ",
        "introduction": "If a cluster at angular position $\\vec{y}$ has the line-of-sight velocity $v$ with respect to the CMB, the CMB fluctuation caused by the SZ effect at frequency $\\nu$ at this position will be $\\delta_\\nu(\\vec y)=\\delta_{\\rm TSZ}(\\vec y)G(\\nu)+ \\delta_{\\rm KSZ}(\\vec y)H(\\nu)$, with $ \\delta_{\\rm TSZ}$=$\\tau T_{\\rm X}/T_{\\rm e,ann}$ and $\\delta_{\\rm KSZ}$=$\\tau v/c$. Here $G(\\nu)\\simeq-1.85$ to $-1.25$ and $H(\\nu)\\simeq 1$ over the WMAP frequencies, $\\tau$ is the projected optical depth due to Compton scattering, $T_{\\rm X}$ is the temperature of the intra-cluster gas, and $k_{\\rm B}T_{\\rm e,ann}$=511 keV. Averaged over many isotropically distributed clusters moving at a significant bulk velocity with respect to the CMB, the dipole from the kinematic term will dominate, allowing a measurement of $V_{\\rm bulk}$. Thus KA-B suggested measuring the dipole component of $\\delta_\\nu(\\vec y)$.  We use a normalized notation for the dipole power $C_{1}$, such that a coherent motion at velocity $V_{\\rm bulk}$ leads to $C_{1,{\\rm kin}}= T_{\\rm CMB}^2 \\langle \\tau \\rangle^2 V_{\\rm bulk}^2/c^2$, where $T_{\\rm CMB} =2.725$K. For reference, $\\sqrt{C_{1,{\\rm kin}}}\\simeq 1 (\\langle \\tau \\rangle/10^{-3}) (V_{\\rm bulk}/100{\\rm km/sec}) \\; \\mu$K. When computed from the total of $N_{\\rm cl}$ positions, the dipole will also have positive contributions from 1) instrument noise, 2) the thermal SZ (TSZ) component, 3) the cosmological CMB fluctuation component arising from the last-scattering surface, and 4) the various foreground components within the WMAP frequency range. The last of these contributions can be significant at the lowest WMAP frequencies (channels K \\& Ka) and, hence, we restrict this analysis to the WMAP Channels Q, V \\& W which have negligible foreground contributions. The contributions to the dipole from the above terms can be estimated as $\\langle\\delta_\\nu(\\vec{y})\\cos\\theta\\rangle$ at the $N_{\\rm cl}$ different cluster locations with polar angle $\\theta$. For $N_{\\rm cl}\\gg1$ the dipole of $\\delta_\\nu$ becomes $a_{1m} \\simeq a_{1m}^{\\rm kin} +a_{1m}^{\\rm TSZ} + a_{1m}^{\\rm CMB} + \\frac{\\sigma_{\\rm noise}}{\\sqrt{N_{\\rm cl}}}$. Here $a_{1m}^{\\rm CMB}$ is the residual dipole produced at the cluster locations by the primordial CMB anisotropies. The dipole power is $C_1= \\sum_{m=-1}^{m=1} |a_{1m}|^2$. The notation for $a_{1m}$ is such that $m$=$0,1,-1$ correspond to the $(x,y,z)$ components, with $z$ running perpendicular to the Galactic plane towards the NGP, and $(x,y)$ being the Galactic plane with the $x$-axis passing through the Galactic center. This dipole signal should not be confused with the \"global CMB dipole\" that arises from our {\\it local} motion relative to the CMB. The kinematic signal investigated here does not contribute significantly to the ``global CMB dipole\" arising from only a small number of pixels. When the latter is subtracted from the original CMB maps, only a small fraction, $\\sim (N_{\\rm cl}/N_{\\rm total})\\lsim 10^{-3}$, of the kinematic signal $C_{1,{\\rm kin}}$ is removed.  The TSZ dipole for a random cluster distribution is $a_{1m}^{\\rm TSZ}\\sim(\\langle \\tau T_{\\rm  X}\\rangle/T_{\\rm e,ann}) N_{\\rm cl}^{-1/2}$ decreasing with increasing $N_{\\rm cl}$. This decrease could be altered if clusters are not distributed randomly and there may be some cross-talk between the monopole and dipole terms especially for small/sparse samples \\cite{watkins-feldman}, but the value of the TSZ dipole will be estimated directly from the maps as discussed below and in greater detail in Kashlinsky et al (2008 - KA-BKE). The residual CMB dipole, $C_{1,{\\rm CMB}}$, will exceed $\\sigma_{\\rm CMB}^2/N_{\\rm cl}$ because the intrinsic cosmological CMB anisotropies are correlated. On the smallest angular scales in the WMAP data $\\sigma_{\\rm CMB}\\simeq 80\\mu$K, so these anisotropies could be seen as the largest dipole noise source. However, because the power spectrum of the underlying CMB anisotropies is accurately known, this component can be removed with a filter described next.  To remove the cosmological CMB anisotropies we filtered each channel maps separately with the Wiener filter as follows. With the known power spectrum of the cosmological CMB fluctuations, $C_\\ell^{\\Lambda{\\rm CDM}}$, a filter $F_\\ell$ in $\\ell$-space which minimizes $\\langle (\\delta T - \\delta_{\\rm instrument \\; noise})^2\\rangle$ in the presence of instrument noise is given by $F_\\ell = (C_\\ell - C_\\ell^{\\Lambda{\\rm CDM}})/C_\\ell$, with $C_\\ell$ being the measured power spectrum of each map. Convolving the maps with $F_\\ell$ minimizes the contribution of the cosmological CMB to the dipole. The maps for {\\it each} of the eight WMAP channels were thus processed as follows: 1) for $C_\\ell^{\\Lambda {\\rm CDM}}$ we adopted the best-fit cosmological model for the WMAP data \\cite{hinshaw} available from http://lambda.gsfc.nasa.gov; 2) each map was decomposed into multipoles, $a_{\\ell m}$, using HEALPix \\cite{healpix}; 3) the power spectrum of each map, $C_\\ell$, was then computed and $F_\\ell$ constructed; 4) the $a_{\\ell m}$ maps were multiplied by $F_\\ell$ and Fourier-transformed back into the angular space $(\\theta,\\phi)$. We then removed the intrinsic dipole, quadrupole and octupole. The filtering affects the effective value of $\\tau$ for each cluster and we calculate this amount later.  Here we use an all-sky cluster sample created by combining the ROSAT-ESO Flux Limited X-ray catalog (REFLEX) \\cite{bohringer} in the southern hemisphere, the extended Brightest Cluster Sample (eBCS) \\cite{ebeling1,ebeling2} in the north, and the Clusters in the Zone of Avoidance (CIZA) \\cite{ebeling3,kocevski1} sample along the Galactic plane. These are the most statistically complete X-ray selected cluster catalogs ever compiled in their respective regions of the sky. All three surveys are X-ray selected and X-ray flux limited using RASS data. The creation of the combined all-sky catalogue of 782 clusters is described in detail by \\cite{kocevski2} and KA-BKE.  We started with 3-year ``foreground-cleaned\" WMAP data (http://lambda.gsfc.nasa.gov) in each differencing assembly (DA) of the Q, V, and W bands. Each DA is analyzed separately giving us eight independent maps to process: Q1, Q2, V1, V2, W1,..., W4. The CMB maps are pixelized with the HEALPix parameter $N_{\\rm side}$=512 corresponding to pixels $\\simeq 7^\\prime$ on the side or pixel area $4\\times 10^{-6}$ sr. This resolution is much coarser than that of the X-ray data, which makes our analysis below insensitive to the specifics of the spatial distribution of the cluster gas, such as cooling flows, deviations from spherical symmetry, etc. In the filtered maps for each DA we select all WMAP pixels within the total area defined by the cluster X-ray emission, repeating this exercise for cluster subsamples populating cumulative redshift bins up to a fixed $z$. In order to eliminate the influence of Galactic emission and non-CMB radio sources, the CMB maps are subjected to standard WMAP masking. The results for the different masks are similar and agree well within their statistical uncertainties.  The SZ effect ($\\propto n_e$, the electron density) has larger extent than probed by X-rays (X-ray luminosity $\\propto n_e^2$), which is confirmed by our TSZ study using the same cluster catalogue (AKKE). As shown in AKKE, KA-BKE contributions to the TSZ signal are detected out to $\\gsim 30^\\prime$. What is important in the present context, is that the X-ray emitting gas is distributed as expected from the $\\Lambda$CDM profile \\cite{nfw} scaling as $n_e\\propto r^{-3}$ in outer parts. In order to be in hydrostatic equilibrium such gas must have temperature decreasing with radius \\cite{komatsu}. Indeed, the typical polytropic index for such gas would be $\\gamma \\simeq$1.2, leading to the X-ray temperature decreasing as $T_{\\rm X} \\propto n_e^{\\gamma-1} \\propto r^{-0.6}$ at outer radii. This $T_{\\rm X}$ decrease agrees with simulations of cluster formation within the $\\Lambda$CDM model \\cite{simulations_temp} and with the available data on the X-ray temperature profile \\cite{pratt}. For such gas, the TSZ monopole ($\\propto T_{\\rm X} \\tau$) decreases faster than the KSZ component ($\\propto \\tau$) when averaged over a progressively increasing cluster area. To account for this, we compute the dipole component of the final maps for a range of effective cluster sizes, namely $[1,2,4,6] \\theta_{\\rm X-ray}$ and then the maximal cluster extent is set at $30^\\prime$ to avoid a few large clusters (eg. Coma) bias the dipole determination. We note that at the final extent our clusters effectively have the same angular radius of 0.5$^\\circ$. Our choice of the maximal extent is determined by the fact that the SZ signal is detectable in our sample out to that scale (AKKE), which is $\\sim$(3-4)Mpc at the mean redshift of the sample. Increasing the cluster radius further to 1$^\\circ-3^\\circ$, causes the dipole to start decreasing with the increasing radius, as expected if the pixels outside the clusters are included diluting the KSZ signal (KA-BKE).  To estimate uncertainties in the signal only from the clusters, we use the rest of the map for the distribution and variance of the noise in the measured signal. We use two methods to preserve the geometry defined by the mask and the cluster distribution: 1) $N_{\\rm cl}$ central random pixels are selected outside the mask away from the cluster pixels adding pixels within each cluster's angular extent around these central random pixels, iteratively verifying that the selected areas do not fall within either the mask or any of the known clusters. 2) We also use a slightly modified version of the above procedure in order to test the effects of the anisotropy of the cluster catalog. There the cluster catalog is rotated randomly, ensuring that the overall geometry of the cluster catalog is accurately preserved. Both methods yield very similar uncertainties; for brevity, we present results obtained with the first method. ",
        "conclusions": ""
    },
    "0809/0809.3412_arXiv.txt": {
        "abstract": "Massive black holes are key components of the assembly and evolution of cosmic structures and a number of surveys are currently on-going or planned to probe the demographics of these objects and to gain insight into the relevant physical processes. Pulsar Timing Arrays (PTAs) currently provide the only means to observe gravitational radiation from massive black hole binary systems with masses $\\simgt 10^7\\Ms$. The whole cosmic population produces a stochastic background that could be detectable with upcoming Pulsar Timing Arrays. Sources sufficiently close and/or massive generate gravitational radiation that significantly exceeds the level of the background and could be individually resolved. We consider a wide range of massive black hole binary assembly scenarios, we investigate the distribution of the main physical parameters of the sources, such as masses and redshift, and explore the consequences for Pulsar Timing Arrays observations. Depending on the specific massive black hole population model, we estimate that on average at least one resolvable source produces  timing residuals in the range $\\sim5-50$ ns. Pulsar Timing Arrays, and in particular the future Square Kilometre Array (SKA), can plausibly detect these unique systems, although the events are likely to be rare. These observations would naturally complement on the high-mass end of the massive black hole distribution function future surveys carried out by the Laser Interferometer Space Antenna ({\\it LISA}). ",
        "introduction": "Massive black hole (MBH) binary systems with masses in the range $\\sim 10^4-10^{10}\\msun$ are amongst the primary candidate sources of gravitational waves (GWs) at $\\sim$ nHz - mHz frequencies (see, e.g., Haehnelt 1994; Jaffe \\& Backer 2003; Wyithe \\& Loeb 2003, Sesana et al. 2004, Sesana et al. 2005). The frequency band $\\sim 10^{-5}\\,\\mathrm{Hz} - 1 \\,\\mathrm{Hz}$ will be probed by the {\\it Laser Interferometer Space Antenna} ({\\it LISA}, Bender et al. 1998), a space-borne gravitational wave laser interferometer being developed by ESA and NASA. The observational window $10^{-9}\\,\\mathrm{Hz} - 10^{-6} \\,\\mathrm{Hz}$ is already accessible with Pulsar Timing Arrays (PTAs;  e.g. the Parkes radio-telescope, Manchester 2008). PTAs exploit the effect of GWs on the propagation of radio signals from a pulsar to the Earth (e.g. Sazhin 1978, Detweiler 1979, Bertotti et al. 1983), producing a characteristic signature in the time of arrival (TOA) of radio pulses.  The timing residuals of the fit of the actual TOA of the pulses and the TOA according to a given model, carry the physical information about unmodelled effects, including GWs (e.g. Helling \\& Downs 1983, Jenet et al. 2005). The complete Parkes PTA (Manchester 2008), the European Pulsar Timing Array (Janssen et al. 2008), and NanoGrav \\footnote{http://arecibo.cac.cornell.edu/arecibo-staging/nanograv/} are expected to improve considerably on the capabilities of these surveys and the planned Square Kilometer Array (SKA; {\\it www.skatelescope.org}) will produce a major leap in sensitivity.  Popular scenarios of MBH formation and evolution (e.g. Volonteri, Haardt \\& Madau 2003; Wyithe \\& Loeb 2003, Koushiappas \\& Zentner 2006, Malbon et al. 2007, Yoo et al. 2007) predict the existence of a large number of massive black hole binaries (MBHB) emitting in the frequency range between $\\sim 10^{-9}$ Hz and $10^{-6}$ Hz. PTAs can gain direct access to this population, and address a number of unanswered questions in astrophysics (such as the assembly of galaxies and dynamical processes in galactic nuclei), by detecting gravitational radiation of two forms: (i) the stochastic GW background produced by the incoherent superposition of radiation from the whole cosmic population of MBHBs and (ii) GWs from individual sources that are sufficiently bright (and therefore massive and/or close) so that the gravitational signal stands above the root-mean-square (rms) value of the background. Both classes of signals are of great interest, and the focused effort on PTAs could lead to the discovery of systems difficult to detect with other techniques.  The possible level of the GW background, and the consequences for observations have been explored by several authors (see {\\em e.g.} Rajagopal \\& Romani 1995; Phinney 2001, Jaffe \\& Backer 2003; Jenet et al. 2005; Jenet et al. 2006; Sesana et al. 2008). Recently, Sesana Vecchio \\& Colacino (2008, hereinafter PaperI) studied in details the properties of such a signal and the astrophysical information encoded into it, for a comprehensive range of MBHB formation models. As shown in PaperI, there is over a factor of 10 uncertainty in the characteristic amplitude of the MBHB generated background in the PTA frequency window. However, the most optimistic estimates yield an amplitude just a factor $\\approx 3$ below the upper-bound placed using current data (Jenet et al. 2006), and near-term future observations could either detect such a stochastic signal or start ruling out selected MBHB population scenarios. Based on our current astrophysical understanding of the formation and evolution of MBHBs and the estimates of the sensitivity of SKA, one could argue that this instrument guarantees the detection of this signal in the frequency range $3\\times 10^{-9}\\,\\mathrm{Hz} - 5\\times 10^{-8} \\,\\mathrm{Hz}$ for essentially every assembly scenario that is considered at present.  The background generated by the cosmic population of MBHBs is present across the whole observational window of PTAs (cf. PaperI). The Monte Carlo simulations reported in PaperI show clearly the presence of distinctive strong peaks well above the average level of the stochastic contribution (cf. Figure 1 and 4 in PaperI). This is to be expected, as individual sources can generate gravitational radiation sufficiently strong to stand above the rms value of the stochastic background. These sources are of great interest because they can be individually resolved and likely involve the most massive MBHBs in the Universe. Their observation can offer further insight into the high-mass end of the MBH(B) population, galaxy mergers in the low-redshift Universe and dynamical processes that determine the formation of MBH pairs and the evolution to form close binaries with orbital periods of the order of years.  Some exploratory studies have been carried out about detecting individual signals from MBHBs in PTA data (Jenet et al. 2004, 2005). In this paper we study systematically for a comprehensive range of assembly scenarios the properties, in particular the distribution of masses and redshift, of the sources that give rise to detectable individual events; we compute the induced timing residuals and the expected number of sources at a given timing residual level. To this aim, the modelling of the high-mass end of the MBHB population at relatively low redshift is of crucial importance. We generate a statistically significant sample of merging massive galaxies from the on-line Millennium database  ({\\it http://www.g-vo.org/Millennium}) and populate them with central MBHs according to different prescriptions (Tremaine et al. 2002, Mclure et al. 2006, Lauer et al. 2007, Tundo et al. 2007). The Millennium simulation (Springel et al. 2005) covers a comoving volume of $(500/h_{100})^3$ Mpc$^3$ ($h_{100}=H_0/100$  km s$^{-1}$ Mpc$^{-1}$ is the normalized Hubble parameter), ensuring a number of massive nearby binaries adequate to construct the necessary distribution. For each model we compute the stochastic background, the expected distribution of bright individual sources and the value of the characteristic timing residual $\\delta t_\\mathrm{gw}$, see Equation (\\ref{e:deltatgw}), for an observation time $T$. The signal-to-noise ratio at which a source can then be observed scales as SNR$\\approx \\delta t_\\mathrm{gw}/\\delta t_\\mathrm{rms}$ where $\\delta t_\\mathrm{rms}$ is the root-mean-square level of the timing residuals noise, both coming from the receiver {\\em and} the GW stochastic background contribution. In the following we summarise our main results: \\begin{enumerate} \\item The number of detectable individual sources for different thresholds of the effective induced timing residuals $\\delta t_\\mathrm{gw}$ is shown in Table \\ref{tab:summary}. Depending on the specific MBH population model, we estimate that on average at least one resolvable source produces  timing residuals in the range $\\sim5-50$ns. Future PTAs, and in particular SKA, can plausibly detect these unique systems; the detection is however by no means guaranteed, events will be rare and just above the detection threshold.  \\item As expected, the brightest signals come from very massive systems with ${\\cal M}>5\\times10^8\\msun$. Here ${\\cal M}=M_1^{3/5}M_2^{3/5}/(M_1+M_2)^{1/5}$ is the chirp mass of the binary and $M_1>M_2$ are the two black hole masses. Most of the resolvable sources are located at relatively high redshift ($0.2 < z < 1.5$), and not at $z \\ll 1$ as one would naively expect, giving the opportunity to probe the universe at cosmological distances.  \\item The number of resolvable MBHBs depends on the actual level of the stochastic background generated by the whole population; here we have used the standard simplified assumption that the background level is determined by having more than one source per frequency resolution bin of width $1/T$, where $T$ is the observational time. Using this definition we find that at frequencies less than $10^{-7}$ Hz there are typically a few resolvable sources, considering $T= $ 5 yrs, with residuals in the range $\\sim 1\\,\\mathrm{nHz} - 1\\,\\mu\\mathrm{Hz}$. As the level of the background decreases for increasing frequencies, fainter sources become visible individually.  \\item As a sanity check, we have compared the MBHB populations and stochastic background levels obtained using data from the Millennium simulation (adopted in this paper) with those derived by means of merger tree realisations based on the Extended Press \\& Schechter  (EPS) formalism (considered in PaperI) and have found good agreement. This provides an additional validation of the results of this paper and PaperI. Moreover it supports that EPS merger trees, if handled sensibly, can offer a valuable tool for the study of MBH evolution even at low redshift. \\end{enumerate}  The paper is organised as follows. In Section 2 we describe MBHB population models, in particular the range of scenarios considered in this paper. A short review of the timing residuals produced by GWs generated by an individual binary (in circular orbit) in the data collected by PTAs is provided in Section 3. Section 4 contains the key results of the paper: the expected timing residuals from the estimated population of MBHBs, including detection rates for current and future PTAs. We also provide a comparison between the stochastic background computed according to the prescriptions considered here and the results of PaperI. Summary and conclusions are given in Section 5. ",
        "conclusions": "We have investigated the ability of Pulsar Timing Arrays (PTAs) to resolve individual massive black hole binary systems by detecting gravitational radiation produced during the in-spiral phase through its effect on the residuals of the time of arrival of signals from radio pulsars. We have considered a broad range of assembly scenarios, using the data of the Millennium simulation to evaluate the galactic haloes merger rates, and a total of twelve different models that control the relations between the mass of the central black holes and the galactic haloes, and the evolution of the black hole masses through accretion. These models therefore cover qualitatively (and to large extent quantitatively) the whole spectrum of MBH assembly scenarios currently considered. Regardless of the model, we estimate that at least one resolvable source is expected to produce timing residuals in the range $\\sim5-50$ ns, and therefore future PTAs, and in particular the Square Kilometre Array may be able to observe these systems. A whistle-stop summary of the models and results is contained in Table~\\ref{tab:summary}. The total number of visible events clearly depend on the sensitivity of PTAs and on the astrophysical scenario. As expected, the brightest sources (for PTAs) are very massive binaries with chirp mass ${\\cal M}>5\\times10^8\\msun$. However, (initially) quite surprisingly most of the resolvable sources are located at relatively high redshift ($z>0.2$). In conjunction with the observation of the stochastic GW background from the whole population of MBHBs, the identification of individual MBHBs could provide new constraints on the populations of these objects and the relevant physical processes.  As a by-product of the analysis, we have also estimated the level of the GW stochastic background produced by the different models, finding good agreement with the estimates derived using merger tree realisations based on the Extended Press \\& Schechter formalism considered in PaperI. Such agreement provides a further validation of the results of this paper and PaperI, it shows that we can now have in hand self-consistent predictions for stochastic and deterministic signals from the cosmic population of MBHBs, and suggests that EPS merger trees could provide a valuable approach to the studies of MBH evolution at low-to-medium redshift.  As a final word of caution, we would like to stress that the results of this paper clearly suffer from considerable uncertainties determined by the still poor quantitative information about several parameters that control the models. The spread of the predictions of the expected events is therefore likely dominated by the lack of knowledge of the model parameters, rather than the differences between the assembly scenarios."
    },
    "0809/0809.1417_arXiv.txt": {
        "abstract": "The evolution of marginally bound supercluster-like objects in an accelerating $\\Lambda$CDM Universe is followed, by means of cosmological simulations, from the present time to an expansion factor $a=100$. The objects are identified on the basis of the binding density criterion introduced by \\cite{dunner06}. Superclusters are identified with the ones whose mass ${\\rm M}>10^{15}h^{-1}$M$_{\\odot}$, the most massive one with ${\\rm M}\\sim8\\times10^{15}h^{-1}$M$_{\\odot}$, comparable to the Shapley supercluster. The spatial distribution of the superclusters remains essentially the same after the present epoch, reflecting the halting growth of the Cosmic Web as $\\Lambda$ gets to dominate the expansion of the Universe. The same trend can be seen in the stagnation of the development of the mass function of virialized haloes and bound objects. The situation is considerably different when looking at the internal evolution, quantified in terms of their shape, compactness and density profile, and substructure in terms of their multiplicity function. We find a continuing evolution from a wide range of triaxial shapes at $a=1$ to almost perfect spherical shapes at $a=100$. We also find a systematic trend towards a higher concentration. Meanwhile, we see their substructure gradually disappearing, as the surrounding subclumps fall in and merge to form one coherent, virialized system. ",
        "introduction": "The evidence for an accelerated expansion of the Universe has established the dominant presence of a `dark energy' component. In the present cosmological paradigm, the Universe entered into an accelerating phase at $z\\approx 0.7$. Observational evidence points towards a dark energy component which behaves like Einstein's cosmological constant. As long as the matter density in the Universe dominated over that of dark energy, the gravitational growth of matter concentrations resulted in the emergence of ever larger structures. Once dark energy came to dominate the dynamics of the Universe and the Universe got into accelerated expansion, structure formation came to a halt \\citep{peebles80,heath77}. With the present-day Universe having reached this stage, the largest identifiable objects that will ever populate our Universe may be the ones that we observe in the process of formation at the present cosmological time. While no larger objects will emerge, these sufficiently overdense and bound patches will not be much affected by the global cosmic acceleration. They will remain bound and evolve as if they are {\\it island universes}: they turn into isolated evolving regions \\citep{chiueh02,nagamine03,busha03,dunner06}.  While clusters of galaxies are the most massive and most recently collapsed and virialized structures, the present-day superclusters are arguably the largest bound but not yet fully evolved objects in our Universe. In our accelerating Universe, we may assume they are the objects that ultimately will turn into island universes. A large range of observational studies, mostly based on optically or X-ray selected samples, show that clusters are strongly clustered and group together in large supercluster complexes \\citep[see e.g.][]{oort1983,bahcall1988,einasto1994,einasto2001,quintana2000}. These superclusters, the largest structures identifiable in the present Universe, are enormous structures comprising a few up to dozens of rich clusters of galaxies, a large number of more modestly sized clumps, and thousands of galaxies spread between these density concentrations.  In this study, we aim at contrasting the large-scale evolution of structure in an accelerated Universe with that of the internal evolution of bound objects. In order to infer what will be the largest bound regions in our Universe, the {\\it island universes}, we study the mass function of bound objects. The abundance or mass function of superclusters serves as a good indicator of the growth of structure of a cosmological model. While large-scale structure formation comes to a halt, this will manifest itself in the asymptotic behaviour of the supercluster mass function. Meanwhile, the internal evolution of the superclusters continues as they contract and collapse into the largest virialized entities the Universe will ever contain.  We address three aspects of the continuing internal evolution of bound regions: their shape, density profile and internal substructure in terms of their cluster multiplicity. The shape of supercluster regions is one of the most sensitive probes of their evolutionary stage. We know that superclusters in the present-day Universe are mostly flattened or elongated structures, usually identified with the most prominent filaments and sheets in the galaxy distribution \\citep[e.g.][]{plionis92,sathya1998,basilakos01,sheth2003, einasto2007c}. The Pisces-Perseus supercluster chain is a particularly well-known example of a strongly elongated filament \\citep[see e.g.][]{hayngiov1986}. The distribution of shapes of bound structures is a combination of at least two factors. One is the shape of the proto-supercluster in the initial density field. The second factor is the evolutionary state of the bound structure. We know that the gravitational collapse of cosmic overdensities --- whose progenitors in the primordial density perturbation field will never be spherical \\citep{peacheav1985,bbks} --- proceeds in a distinctly anisotropic fashion via flattened and elongated configurations towards a final more compact triaxial virialized state \\citep[see eg.][]{zeldovich1970,icke73,whitesilk1979, eistloeb1995,bondmyers1996,sathya1996,desjacques2008,weybond2008}.  The collapse of the superclusters will also result in a continuous sharpening of the internal mass distribution, reflected in the steepening of their density profile. While they are in the process of collapse, internal substructure of constituent clusters remains recognizable. While the subclumps merge into an ever more massive central concentration the supercluster substructure gradually fades, resulting in an increasingly uniform mass distribution. The evolving and decreasing level of substructure will be followed in terms of the evolving supercluster multiplicity function, i.e. the number of cluster-sized clumps within the supercluster region.  Representing moderate density enhancements on scale of tens of Mpc, in the present Universe superclusters are still expanding with the Hubble flow, although at a slightly decelerated rate, or have just started contracting. Because these structures have not yet fully formed, virialized and clearly separated from each other, it is difficult to identify them unambiguously. In most studies, superclusters have been defined by more or less arbitrary criteria, mostly on the basis of a grouping and/or percolation algorithm \\citep[see e.g][]{oort1983,bahcall1988,einasto1994,einasto2001,quintana2000}. This introduces the need for a user-specified percolation radius. \\citet[hereafter Paper I]{dunner06} attempted to define a more physically based criterion, identifying superclusters with the biggest gravitationally bound structures that will be able to form in our Universe. On the basis of this, they worked out a lower density limit for gravitationally bound structures. This limit is based on the density contrast that a spherical shell needs to enclose to remain bound to a spherically symmetric overdensity.  We use this spherical density criterion to identify bound structures in a large cosmological box. In this we follow the work of \\cite{chiueh02} and \\cite{dunner06}. \\cite{chiueh02} numerically solved the spherical collapse model equations for self-consistent growing mode perturbations in order to obtain a theoretical criterion for the mean density enclosed in the outer gravitationally bound shell. The resulting density criterion was evaluated by \\cite{dunner06} on the basis of numerical simulations. They generalized it by deriving the analytical solution which also forms the basis of the current study, and in \\cite{dunner07} extended the criterion to limits for bound structures in redshift space.  Various authors have addressed the future evolution of cosmic structure \\citep{chiueh02,busha03,nagamine03,dunner06,dunner07,hoffman07,busha07}. The internal evolution of the density and velocity structures of bound objects was followed by \\cite{busha03}, with \\cite{busha07} focusing on the effects of small-scale structure on the formation of dark matter halos in two different cosmologies. \\cite{nagamine03} and \\cite{hoffman07} specifically focused on the evolution of the Local Universe. \\cite{nagamine03} found that the Local Group will get detached from the rest of the Universe, and that its physical distance to other systems will increase exponentially. \\cite{hoffman07} investigated the dependence on dark matter and dark energy by contrasting $\\Lambda$CDM and OCDM models. They concluded that the evolution of structure in comoving coordinates at long term is determined mainly by the matter density rather than by the dark energy. A key point of attention in \\cite{nagamine03}, \\cite{hoffman07} and \\cite{busha07} was the mass function of objects in their simulations, on which they all agree that it hardly changes after the current cosmic epoch.  This paper is the first in a series addressing the future evolution of structure in FRW Universes. In an accompanying publication we will specifically look at the influence of dark matter and the cosmological constant on the emerging (super)cluster population, based on the work described in \\cite{arayamelo08}. The present study concentrates on the details of this evolution in a standard flat $\\Lambda$CDM Universe.  This paper is organized as follows. In section \\ref{sec:scdef}, we present a review of the spherical collapse model, including a derivation of the critical overdensity for a structure to remain bound. Section \\ref{sec:simulation} describes the simulation and the group finder algorithm that we employ when determining the mass functions. This is followed in sect.~\\ref{sec:evolution} by a a qualitative description of the evolution, including a case study of the evolution of some typical bound mass clumps from the present epoch to $a=100$. The lack of evolution in their spatial distribution in the same time interval is studied in sect.~\\ref{sec:spatial}. Section \\ref{sec:ps_mf} presents the mass functions of the bound structures at $a=1$ and $a=100$ and a comparison with the ones obtained by the Press-Schechter formalism and its variants. The evolution of the shapes of the structures is studied in section \\ref{sec:shapes}. In section \\ref{sec:dens_prof} we look into the mass distribution and density profiles. Section \\ref{sec:multi} presents the stark changes in the supercluster multiplicity function. In sect.~\\ref{sec:shapley} its results are combined with those on the supercluster mass functions obtained in section~\\ref{sec:ps_mf} to relate our findings to the presence and abundance of monster supercluster complexes like the Shapley and Horologium-Reticulum supercluster. Finally, in section~\\ref{sec:conclusions}, we discuss our findings and draw conclusions on various issues addressed by our study. ",
        "conclusions": "\\label{sec:conclusions} In this work, we have followed the evolution of bound objects from the present epoch up to a time in the far future of the Universe, at $a=100$, in a standard $\\Lambda$CDM ($\\Omega_{m,0}=0.3$, $\\Omega_{\\Lambda,0}=0.7$ and $h=0.7$) Universe. We contrasted the external global evolution of the population of bound objects with their vigorous internal evolution, starting from the contention that in a dark energy dominated Universe they have the character of {\\it island universes}. Within such a Universe we expect them to become increasingly isolated objects in which cosmic evolution proceeds to the ultimate equilibrium configuration of a smooth, spherical, virialized and highly concentrated mass clumps. We identify the most massive of these objects with superclusters.  For the external evolution, we investigate the spatial distribution and clustering of bound objects and superclusters, along with the weak change of their mass functions. To assess the internal structure, we have looked into their rapidly changing shape, their evolving density profile and mass concentration and the level of substructure of the superclusters in terms of their multiplicity, i.e. the number of clusters within their realm.  We defined the bound structures by the density criterion derived in Paper I (see Eqn.~\\ref{eq:ratio}), and identified them from a 500$h^{-1}$ Mpc cosmological box with $512^3$ dark matter particles in a $\\Lambda$CDM ($\\Omega_{m,0}=0.3$, $\\Omega_{\\Lambda,0}=0.7$ and $h=0.7$) Universe. We ran the simulation up to $a=100$, which is a time where structures have stopped forming. We used the HOP halo identifier in order to identify independently virialized structures, at each of the five timesteps we have analyzed in detail: $a=1$, 2, 5, 10 and 100.  The main results of the present study can be summarized as follows:  \\begin{itemize}  \\item While the large-scale evolution of bound objects and superclusters comes to a halt as a result of the cosmic acceleration, their internal evolution continues vigorously until they have evolved into single, isolated, almost perfectly spherical, highly concentrated, virialized mass clumps. This development is very strong between $a=1$ and 10, and continues up to $a=100$.  \\item The marginally bound objects that we study resemble the superclusters in the observed Universe. While clusters of galaxies are the most massive, fully collapsed and virialized objects in the Universe, superclusters are the largest bound --but not yet collapsed-- structures in the Universe.  \\item The superclusters are true \\emph{island universes}: as a result of the accelerating expansion of the Universe, no other, more massive and larger, structures will be able to form.  \\item While the superclusters collapse between $a=1$ and $a=100$, their surroundings change radically. While at the present epoch they are solidly embedded within the Cosmic Web, by $a=100$ they have turned into isolated cosmic islands.  \\item The large scale distribution of bound objects and superclusters (in comoving space) does not show any significant evolution in between $a=1$ and $a=100$. The cluster and supercluster correlation functions do not change over this time interval, and retain their near power-law behaviour. Superclusters remain significantly stronger clustered than the average bound object, with a supercluster correlation length of $23 \\pm 5h^{-1}$Mpc compared to $r_0\\approx 11.5h^{-1}$Mpc for the full bound object distribution in our simulation sample.  \\item The mass functions of bound objects and superclusters hardly change from $a=1$ to $a=100$, as we expect on theoretical grounds. The mass functions in the simulations are generally in good agreement with the theoretical predictions of he Press-Schechter, Sheth-Tormen, and Jenkins mass functions. At $a=1$, the Sheth-Tormen prescription provides a better fit. At $a=100$, the pure Press-Schechter function seems to be marginally better. This may tie in with the more anisotropic shape of superclusters at $a=1$ in comparison to their peers at $a=100$.  \\item The change in the internal mass distribution and that in the surroundings is directly reflected by the radial density profile. Without exception towards $a=100$ all objects attain a highly concentrated internal matter distribution, with a concentration index $c=0.2$. In general, the vast majority of objects has evolved into a highly concentrated mass clump after $a=10$.  \\item The mass profile in the outer realms of the supercluster changes radically from $a=1$ to $a=100$. At $a=1$ it is rather irregular, while there are large differences between the individual objects. This is a reflection of the surrounding inhomogeneous mass distribution of the Cosmic Web. In between $a=5$ and $a=10$, nearly all superclusters have developed a smooth, regular and steadily declining mass profile.  \\item The inner density profile steepens substantially when the inner region of the supercluster is still contracting. On the other hand, when at $a=1$ it has already developed a substantial virialized core, the inner density profile hardly changes.  \\item As a result of their collapse, the shapes of the bound objects systematically change from the original triaxial shape at $a=1$ into an almost perfectly spherical configuration at $a=100$. For example, at $a=1$ their mean axis ratios are ($\\langle s_{2}/s_{1} \\rangle$, $\\langle s_{3}/s_{1} \\rangle $)= (0.69,0.48). At $a=100$, they have mean axis ratios of ($\\langle s_{2}/s_{1} \\rangle$,$\\langle s_{3}/s_{1} \\rangle $)= (0.94,0.85).  \\item At the current epoch the superclusters still contain a substantial amount of substructure. Particularly interesting is the amount of cluster-mass virialized objects within its realm, expressed in the so called \\emph{multiplicity function}. Restricting ourselves to superclusters with a mass larger than $5\\times10^{15}h^{-1}$M$_{\\odot}$, of which we have 17 in our simulation sample, we find a multiplicity of 5 to 15 at the current epoch. As time proceeds there is a systematic evolution towards unit multiplicity at $a=100$, following the accretion and merging of all clusters within the supercluster's realms.  \\item In a volume comparable to the Local Universe ($z<0.1$) we find that the most massive supercluster would have a mass of $\\sim8\\times10^{15}h^{-1}$M$_{\\odot}$. This is slightly more massive than the mass of the Shapley Supercluster given in \\cite{dunner08}. When turning towards the multiplicity, we find that the largest superclusters in the Local Universe would host between 10 and 15 members, close to the number found in the bound region of the Shapley supercluster \\citep{dunner08} (which contains $\\sim 1/3$ of the clusters traditionally assigned to this structure, e.g., \\cite{proust2006}). \\end{itemize}  While in this study we have addressed a large number of issues, our study leaves many related studies for further investigation. One of the most pressing issues concerns an assessment of the nature of our supercluster objects. This involves a comparison with other supercluster definitions, in particular in how far our density-based definition fares at the earlier epochs when most objects of similar mass will have distinct anisotropic shapes.  Also, while here we follow the evolution of superclusters in the standard $\\Lambda$CDM Universe, in an accompanying publication we will systematically address the influence of dark matter and dark energy on the emerging supercluster. There we found that dark matter is totally dominant in determining the supercluster's evolution. For a preliminary, detailed report on the analysis of the role of the cosmological constant in the formation and evolution of structures, we refer to \\cite{arayamelo08}."
    },
    "0809/0809.3909_arXiv.txt": {
        "abstract": "We investigate the effects of non-Gaussianity in the primordial density field on the reionization history. We rely on a semi-analytic method to describe the processes acting on the intergalactic medium (IGM), relating the distribution of the ionizing sources to that of dark matter haloes.  Extending previous work in the literature, we consider models in which the primordial non-Gaussianity is measured by the dimensionless non-linearity parameter $f_{\\rm NL}$, using the constraints recently obtained from cosmic microwave background data.  We predict the ionized fraction and the optical depth at different cosmological epochs assuming two different kinds of non-Gaussianity, characterized by a scale-independent and a scale-dependent $f_{\\rm NL}$ and comparing the results to those for the standard Gaussian scenario.  We find that a positive $f_{\\rm NL}$ enhances the formation of high-mass haloes at early epochs, when reionization begins, and, as a consequence, the IGM ionized fraction can grow by a factor up to 5 with respect to the corresponding Gaussian model. The increase of the filling factor has a small impact on the reionization optical depth and is of order $\\sim 10$ per cent if a scale-dependent non-Gaussianity is assumed.  Our predictions for non-Gaussian models are in agreement with the latest WMAP results within the error bars, but a higher precision is required to constrain the scale dependence of non-Gaussianity. ",
        "introduction": "\\label{sect:intro}  Reionization marks a crucial event in the history of the universe, when the first sources of ultra-violet (UV) radiation ionize the neutral Intergalactic Medium (IGM) and affect the subsequent formation of the cosmic structures. When reionization ends, the small amount of left neutral hydrogen is responsible for the absorption lines that we observe today in the spectra of far objects. However, the way in which this complex phenomenon occurs is still not well understood and the most recent observations paint it as a spatially inhomogeneous and not istantaneous process. While the Gunn \\& Peterson trough of the high-$z$ QSO spectra suggests a late epoch of reionization at $z\\approx 6$ \\citep{fan2001,becker2001,white2003b,fan2006}, the very recent analysis of the 5-year WMAP data on the cosmic microwave background (CMB) polarization shows an IGM optical depth $\\tau\\sim 0.084$ which is in better agreement with an earlier reionization redshift, $z\\sim 10.8$ \\citep{komatsu2008}.  On the other hand, a late reionization end at $z\\sim 6$ is also probed by the IGM temperature measured at $z<4$ \\citep{hui2003} and by the lack of evolution in the luminosity function of Lyman-$\\alpha$ galaxies between $z=5.7$ and $z=6.5$ \\citep[][see, however, {\\protect \\cite{ota2008}} for evidences of a decline at high $z$]{malhotra2004}. Overall, the present situation regarding reionization at redshift $z\\sim6$ as probed by QSO spectra is still unclear \\citep{becker2006}.  Many analytic, semi-analytic and numerical models \\citep[see, e.g.][]{ gnedin2000,ciardi2003b,wyithe2003,barkana2004, haiman2003,madau2004,Wyithe2006,choudhury2007,iliev2007,ricotti2008} have been proposed to describe this poorly understood reionization process. They basically relate the statistical properties and morphology of the ionized regions to the hierarchical growth of the ionizing sources, making more or less detailed assumptions to describe the ionization and recombination processes acting on the IGM.  Since the first sources of UV background radiation appear in the firstly formed dark matter haloes, which correspond to the highest peaks of the primordial density field, the reionization process is expected to strongly depend on the main parameters describing the cosmological model and the power spectrum of primordial density fluctuations.  For instance, the possible presence of an evolving component of dark energy can imprint signatures in the resulting morphology of the ionized regions and change the time-scales of the whole process \\citep[see, e.g.][]{maio2006,crociani2008}. Also the nature and the statistical distribution of the primordial matter fluctuations can influence the reionization history. Although the standard scenario for the origin of the structures assumes that the primordial perturbations are adiabatic and have a (almost) Gaussian distribution, small deviations from primordial Gaussianity affect the dark halo counts, in the rare-event tail, thus also in the high peaks of the density field which originated collapsed objects at high $z$.  Aim of this work is to investigate the effects of some level of non-Gaussianity in the primordial density field on the reionization history. We will make use of analytical techniques to describe the processes in action on the IGM. In particular, we will extend previous works in which the considered non-Gaussian models have density fluctuations described by a renormalized $\\chi^{2}$ probability distribution with $\\nu$ degrees of freedom \\citep{avelino2006} or by a modified Poisson distribution with a given expectation value $\\lambda$ \\citep{chen2003}. Here we will adopt a more convenient way to introduce primordial non-Gaussianity, which has now become standard in the literature, based on the parameter $f_{\\rm NL}$ (see the next section for its definition). In particular, we will assume two different kinds of non-Gaussianity, characterized by a scale-independent and a scale-dependent $f_{\\rm NL}$ parameter.  The paper is organised as follows. In Section \\ref{sect:ng} we introduce the main characteristics of the cosmological models with primordial non-Gaussianity considered here. Section \\ref{sect:model} reviews the main assumptions of the analytical model adopted to describe the cosmic reionization process. The main results of our analysis are presented and discussed in Section \\ref{sect:res}. Finally in Section \\ref{sect:conclu} we draw our conclusions. ",
        "conclusions": "\\label{sect:conclu}  The aim of this work was to investigate how primordial non-Gaussianity may alter the reionization history when compared to the standard scenario based on Gaussian statistics.  We have chosen a simple analytic method to describe the physical processes acting on the IGM to make predictions on the evolution of the ionized fraction and the reionization optical depth.  Our work extends previous analyses \\citep{chen2003,avelino2006} based on simplified ways to introduce primordial non-Gaussianity, considering models motivated by inflation. In particular we assume two different hierarchical evolution scenarios for the ionizing sources, characterized by scale-independent and scale-dependent non-Gaussianity.  All scenarios here considered are not violating the constraints coming from the recent analysis of the 5-year WMAP data.  Our main conclusions can be summarized as follows:  \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item non-Gaussianity affects the abundance of the dark matter haloes, since the formation of high-mass collapsed objects is enhanced (reduced) when positive (negative) values for the $f_{\\rm NL}$ parameter are assumed.  This effect is more evident at earlier cosmological epochs, exactly when reionization begins. \\item As a consequence, for positive primordial non-Gaussianity, the collapsed fraction in type Ia and Ib sources is higher than for the Gaussian case at the same epoch, and the difference increases with $z$. The opposite result applies when $f_{\\rm NL}<0$ is assumed. \\item The IGM filling factor is higher and its evolution is faster than in the Gaussian scenario if a positive $f_{\\rm NL}$ is assumed. This effect is enhanced for a scale-dependent non-Gaussianity, that can produce a 5 times higher $F_{\\rm HI\\!I}$ with respect to the Gaussian case, at early cosmological epochs. Viceversa the filling factor is smaller and has a mild redshift evolution for negative $f_{\\rm NL}$. \\item Both local and equilateral non-Gaussianity have a small (less than 10 per cent) impact on the reionization optical depth, but the effect is enhanced assuming a scale-dependent $f_{\\rm NL}$ parameter. \\end{enumerate}  We finally remark that our predictions of the reionization optical depth in non-Gaussian cosmologies are in agreement with that estimated by 5-year WMAP analysis within $1\\sigma$ error bars and a precision higher than that of WMAP is required to constrain non-Gaussianity and its scale dependence. Ideally one would simulate reionization in non-Gaussian cosmological models using hydrodynamical simulations that incorporate all the relevant physical processes in a consistent framework (most importantly radiative transfer effects in the IGM).  However, such approach is very time consuming due to the large box size and the high resolution required to simulate large volumes and, at the same time, the physics of the sources of radiation.  For this reason, some approximate semi-analytic schemes such as the one presentend here are still useful especially when calibrated on the more robust results of the hydrodynamical runs."
    },
    "0809/0809.2287_arXiv.txt": {
        "abstract": "We present new $V$-band differential brightness measurements as well as new radial-velocity measurements of the detached, circular, 0.84-day period, double-lined eclipsing binary system CV~Boo. These data along with other observations from the literature are combined to derive improved absolute dimensions of the stars for the purpose of testing various aspects of theoretical modeling. Despite complications from intrinsic variability we detect in the system, and despite the rapid rotation of the components, we are able to determine the absolute masses and radii to better than 1.3\\% and 2\\%, respectively. We obtain $M_{\\rm A} = 1.032 \\pm 0.013$\\,M$_{\\sun}$ and $R_{\\rm B} = 1.262 \\pm 0.023$\\,R$_{\\sun}$ for the hotter, larger, and more massive primary (star A), and $M_{\\rm B} = 0.968 \\pm 0.012$\\,M$_{\\sun}$ and $R_{\\rm B} = 1.173 \\pm 0.023$\\,R$_{\\sun}$ for the secondary. The estimated effective temperatures are $5760 \\pm 150$\\,K and $5670 \\pm 150$\\,K. The intrinsic variability with a period $\\sim$1\\% shorter than the orbital period is interpreted as being due to modulation by spots on one or both components. This implies that the spotted star(s) must be rotating faster than the synchronous rate, which disagrees with predictions from current tidal evolution models according to which both stars should be synchronized.  We also find that the radius of the secondary is larger than expected from stellar evolution calculations by $\\sim$10\\%, a discrepancy also seen in other (mostly lower-mass and active) eclipsing binaries. We estimate the age of the system to be approximately 9~Gyr. Both components are near the end of their main-sequence phase, and the primary may have started the shell hydrogen-burning stage. ",
        "introduction": "CV~Boo (= BD~+37~2641 = GSC~2570~0843; $\\alpha = 15^{\\rm h}\\,26^{\\rm m}\\,19\\fs54$, $\\delta = +36\\arcdeg\\,58\\arcmin\\,53\\farcs4$, J2000.0; $V \\approx 10.8$, SpT = G3V) was discovered as a possible eclipsing binary star by \\cite{peniche85}. \\cite{busch85} confirmed it as an eclipsing binary of type EA and found its period to be 0.8469935 days. In his last published paper, a study of 4 lower main sequence binaries, \\cite{Popper:00} determined a spectroscopic orbit for CV~Boo.  Popper was pessimistic about the prospects for determining accurate absolute properties of them because ``It appears unlikely that definitive photometry will be obtained for these stars, partly because of intrinsic variability.''  Recently, a light curve and radial velocity study of the system were done by \\cite{Nelson:04b}, resulting in the first estimates of its absolute properties.  The parameters of CV~Boo make it potentially interesting as the most evolved system among the well-studied double-lined eclipsing binaries with components near 1~M$_{\\sun}$ (see Figure~\\ref{fig:other_mr}), a regime where some discrepancies with theoretical models have been pointed out.  We describe in the following our extensive new photometric and spectroscopic observations of the object intended to improve our knowledge of the system. The presence of starspots does in fact limit somewhat our ability to determine highly accurate absolute properties for this binary star, but the results are still accurate enough for meaningful tests of current stellar models.  As we describe here, CV~Boo contributes significantly to the body of evidence concerning the differences with theory mentioned above.  \\begin{figure} \\epsscale{1.35} \\vskip -0.3in {\\hskip -0.2in\\plotone{f1.eps}} \\vskip -0.3in \\caption{Main-sequence eclipsing binaries in mass range of CV~Boo with accurate determinations of their absolute properties (masses and radii good to better than $\\sim$2\\%). Data are taken from \\cite{Andersen:91} and updates from the literature. Primary and secondary components are connected with solid lines. CV~Boo is represented with open circles. The dashed line shows the solar-metallicity zero-age main sequence from the models by \\cite{Yi:01}, for reference.\\label{fig:other_mr}} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  Despite the system's intrinsic variability, the absolute dimensions for the components of CV~Boo have now been established quite precisely. The relative errors are better than 1.3\\% in the masses and 2\\% in the radii. The object can now be counted among the group of eclipsing binaries with well-known parameters. Under different circumstances the large number and high quality of the photometric observations we have collected might have permitted a more detailed study of the limb darkening laws and a comparison with theoretically predicted coefficients, but this possibility was thwarted here by the intrinsic variability. This phenomenon is not itself without interest. If interpreted as due to the presence of spots, as we have done here, it implies that at least one of the stars is rotating about 1\\% more rapidly than the synchronous rate, a result that was unexpected for a close but well detached system such as this.  We conclude that our current understanding of tidal evolution is still incomplete, or that other processes are at play in this system that theory does not account for. One interesting possibility is differential rotation. The interpretation of measurements of the rotation period of a star made by photometric means, as we have implicitly done here, usually relies on the assumption of solid-body rotation. More often than not, spots are located at intermediate latitudes rather than on the equator, or at high latitudes in more active stars, and differential rotation is such that the stellar surface revolves more slowly at higher latitudes, at least in the Sun. This will tend to bias photometric rotation measurements towards \\emph{longer} periods, if differential rotation is significant enough. In CV~Boo we see the opposite: the period is \\emph{shorter} than the equatorial rate, assuming that synchronization holds. Thus differential rotation can only explain the signal we have detected if it is ``anti-solar'', with the polar regions rotating more rapidly. A handful of stars do indeed show evidence of weak anti-solar differential rotation \\citep[e.g., IL~Hya, HD~31933, $\\sigma$~Gem, UZ~Lib;][]{Weber:03, Strassmeier:03, Kovari:07, Vida:07}. They all happen to be very active (some of them with high-latitude spots, as in CV~Boo), although they tend to be giants or subgiants rather than dwarfs. It is thought that this phenomenon may result from fast meridional flows \\citep[see, e.g.,][]{Kitchatinov:04}. Further progress in understanding the rotation of the CV~Boo components could be made with additional differential photometric observations in several passbands, along with simultaneous high-resolution, high signal-to-noise ratio spectroscopy over a full orbital cycle.  Another significant discrepancy we find with theory is in the radius of the secondary, which appears to be some 10\\% too large compared with predictions from stellar evolution models. This difference is in the same direction as seen for a number of other low-mass eclipsing binaries, as mentioned in \\S\\,\\ref{sec:evolution}.  In those cases one of the explanations most often proposed is that the strong magnetic fields associated with activity (which is common in rapidly-rotating K and M dwarfs in close binaries) tend to inhibit convective motions, and the structure of the star adjusts by increasing its size to allow the surface to radiate the same amount of energy.  At the same time, the effective temperature tends to decrease in order to preserve the total luminosity. Spot coverage can produce similar effects. Theoretical and observational evidence for the conservation of the luminosity in these systems has been presented by \\cite{Delfosse:00}, \\cite{Mullan:01}, \\cite{Torres:02}, \\cite{Ribas:06}, \\cite{Torres:06}, \\cite{Chabrier:07}, and others \\citep[see also][]{Morales:08}. We do not see any obvious discrepancy in the temperature of CV~Boo~B compared to models, although our uncertainties are large enough that the effect may be masked.  If we restrict ourselves to well studied double-lined eclipsing binaries in which the mass and radius determinations are the most reliable, deviations from theory such as those described above have usually been seen in stars that are considerably less massive than the Sun, which have deep convective envelopes. However, the recent study by \\cite{Torres:06} pointed out that the problem is not confined to the lower mass stars, but extends to active objects approaching 1~M$_{\\sun}$, such as V1061~Cyg~Ab, with $M = 0.93$~M$_{\\sun}$. CV~Boo~B has an even larger mass of 0.968~M$_{\\sun}$, and also appears to be oversized.  Similarly with the virtually identical active star FL~Lyr~B ($M = 0.960$~M$_{\\sun}$). The convective envelopes of these objects are considerably thinner than in K and M dwarfs and represent only a few percent of the total mass, yet they appear sufficient for magnetic fields to take hold and alter the global properties of the star, if that is the cause of the discrepancies.  These examples show once again that our understanding of stellar evolution theory is incomplete, even for stars near the mass of the Sun."
    },
    "0809/0809.1942.txt": {
        "abstract": "We investigate the formation of GMCs in spiral galaxies through both agglomeration of clouds in the spiral arms, and self gravity. The simulations presented include two-fluid models, which contain both cold and warm gas, although there is no heating or cooling between them. We find agglomeration is predominant when both the warm and cold components of the ISM are effectively stable to gravitational instabilities. In this case, the spacing (and consequently mass) of clouds and spurs along the spiral arms is determined by the orbits of the gas particles and correlates with their epicyclic radii (or equivalently spiral shock strength). Notably GMCs formed primarily by agglomeration tend to be unbound associations of many smaller clouds, which disperse upon leaving the spiral arms. These GMCs are likely to be more massive in galaxies with stronger spiral shocks or higher surface densities. GMCs formed by agglomeration are also found to exhibit both prograde and retrograde rotation, a consequence of the clumpiness of the gas. At higher surface densities, self gravity becomes more important in arranging both the warm and cold gas into clouds and spurs, and determining the properties of the most massive GMCs. These massive GMCs can be distinguished by their higher angular momentum, exhibit prograde rotation and are more bound. For a 20 M$_{\\odot}$ pc$^{-2}$ disc, the spacing between the GMCs fits both the agglomeration and self gravity scenarios, as the maximum unstable wavelength of gravitational perturbations in the warm gas is similar to the spacing found when GMCs form solely by agglomeration. ",
        "introduction": "The accumulation of the ISM into giant molecular clouds (GMCs) represents the earliest stage in star formation. The properties of GMCs, such as turbulence and magnetic field strength regulate how star formation evolves on local scales (e.g. \\citealt{McKee1999,Pudritz2002,Larson2003}), and are intrinsically linked to many of the issues in star formation, such as the time for stellar collapse (e.g. \\citep{Elmegreen2007}). Thus understanding how GMCs form, and how their characteristics depend on the dynamics of galaxies and nature of the interstellar medium (ISM), is essential for progress in star formation. The formation and evolution of GMCs is also interrelated to the global properties of the ISM. For example, GMC formation by coalescence is much more likely in a cold, clumpy medium, compared to a warm diffuse environment \\citep{Dobbs2006}, where instabilities in the ISM are required. In turn stellar feedback controls the ejection of hot gas back into the ISM and generates turbulence (e.g. \\citealt{MacLow2004,Joung2006,deAvillez2007}).  There have been numerous suggestions of how GMCs form (as described in a recent review by \\citealt{McKee2007} and references therein). However they predominantly fall into two categories: either GMCs form by the agglomeration of smaller clouds of gas (e.g. \\citealt{Field1965,Taff1972,Scoville1979,Casoli1982}), or through instabilities in the ISM. The latter include gravitational \\citep{Cowie1981,Elmegreen1983,Balbus1985,LaVigne2006}, magnetic (either Parker instabilities \\citep{Mous1974,Blitz1980} or MRI \\citep{Kim2003}) or thermal instabilities \\citep{F1965,Koyama2000,Stiele2006}. Another recent suggestion is that GMCs form in colliding or turbulent flows \\citep{Vaz1995,Ball1999,Vaz2006,Heitsch2006}. This however requires a mechanism to produce these flows, which is most likely gravity \\citep{Elmegreen2003}, spiral shocks or supernovae \\citep{Koyama2000,Bergin2004}.  Coalescence of smaller clouds was first instigated to explain GMC formation, but the timescales of $10^8$ Myr for formation \\citep{Kwan1979} were presumed too long to account for the observed properties of GMCs \\citep{Blitz1980}. Furthermore star formation may be expected to disrupt the constituent clouds before a more massive GMC is assembled \\citep{McKee2007}. By including spiral density waves, the time for formation by coalescence is reduced to 10's of Myrs \\citep{Casoli1982,Kwan1987,Roberts1987}.  In more recent reviews, gravitational instabilities generally appear the preferred mechanism for GMC formation (e.g. \\citealt{Elmegreen1990,McKee2007}). Theoretical analysis \\citep{Elmegreen1982} and numerical simulations \\citep{Kim2001,KOS2002} suggest that the Parker instability alone produces insufficient density enhancements, and GMCs can form by gravitational instabilities in relatively short timescales \\citep{Elmegreen1990}. Predictions from gravitational analysis have also been compared to the observed masses and sizes of GMCs. In spiral galaxies, GMCs appear to be regularly spaced along the spiral arms. This spacing, and the mass of GMCs  can be interpreted in terms of gravitational perturbations to the gas \\citep{Cowie1981,Elmegreen1983,Balbus1985,LaVigne2006}. For a sound speed of 7 km s$^{-1}$, the estimated mass of a GMC is $10^6-10^7$ M$_{\\odot}$, whilst the characteristic spacing is expected to be around 3 times the width of the spiral arm, corresponding approximately with observations \\citep{Elmegreen1995}.  Finally, GMC formation by gravitational instabilities has further been endorsed by recent numerical simulations \\citep{Kim2002,Kim2006,Shetty2006} which produce spacings and masses in line with the theoretical predictions.  The problem of molecular cloud formation has been reopened in the last decade, with computational resources now enabling numerical simulations on both local \\citep{Kim2002,Glover2007a,Glover2007b,Vaz2006,Heitsch2006} and galactic \\citep{Shetty2006,DBP2006,DB2008,Tasker2008} scales. These simulations show the formation of GMCs by self gravity \\citep{Kim2002,Kim2006,Shetty2006,Glover2007a}, turbulent or colliding flows \\citep{Vaz2006,Glover2007b,Heitsch2006,Henne2008,Heitsch2008}, combined Parker and thermal instabilities \\citep{Kosinski2007}, and Kelvin Helmholtz instabilities \\citep{Wada2004,Wada2008}. In previous results \\citep{Dobbs2006,DBP2006}, we showed that molecular clouds, and inter-arm spurs, can form as a result of cold gas passing through a spiral shock. The formation of this structure occurs without self gravity, and is still present when magnetic fields are included, although magnetic pressure acts to smooth out clumps in the spiral shock \\citep{Dobbs2008}. The formation of clouds in these calculations bears most resemblance to the collisional models. Clumps of gas are forced together by the spiral shock, and agglomerate into larger structures, which become most discernible as interarm spurs extending from the spiral arms. Generally in simulations of grand design spirals, the spiral pattern is assumed to be long-lived (including the present work), although this may not be the case \\citep{Merrifield2006,Shetty2007}.  This paper concentrates on two possibilities, the formation of GMCs by agglomeration, and by gravitational instabilities. In particular, we discuss the relative contribution to GMC formation for different initial conditions. When the surface density is lower, GMCs form by the agglomeration of smaller gas clouds as their orbits converge in spiral shocks, similar to the model by \\citet{Roberts1987}. In this scenario, the spacing and size of clumps depends on the time (or equivalently distance) gas spends in the spiral arm, and interactions between clumps during orbit crossings. We show also the formation of GMCs form by gravitational instabilities for a high surface density disc with solely warm gas, and show that for a high surface density disc containing cold gas, both agglomeration and gravitational instabilities contribute to GMC formation. We assess the differences in structure of the disc, in particular the separation of spurs, by performing MHD calculations with and without self gravity, varying the strength of the potential, disc mass, temperature and magnetic field strength. ",
        "conclusions": "We have investigated GMC formation by agglomeration and self gravity. Agglomeration occurs in spiral shocks providing there is a clumpy constituent of the ISM, assumed to be cold HI or molecular gas. This process occurs regardless of whether self gravity is included, and in low surface density calculations, the formation of GMCs is predominantly due to agglomeration. The converse situation arises when there is only warm gas, and GMC formation occurs by gravitational instabilities, providing the disc is gravitationally unstable. Generally, both agglomeration and self gravity are expected to contribute to GMC formation, with self gravity becoming more important to GMC formation and disc structure as the surface density increases.  The degree to which these processes determine GMC properties will depend on the surface density of the galaxy, the thermal nature of the ISM, and most likely the magnetic field strength. In particular, GMCs with retrograde rotation can be produced when the ISM is clumpy. A caveat to these calculations is that we assume a two fluid model with no heating or cooling. We expect to include the thermodynamics of the ISM in future work. A further caveat is that gravitational collapse should occur in our models and lead to stellar feedback. Cooling of the gas to temperatures $<$ 100~K, the inclusion of non-ideal MHD processes, and possibly higher numerical resolution would induce collapse. However the aim of the current paper is to investigate the structure of the disc without feedback. The aim of future simulations will be to see how this picture changes with stellar feedback."
    },
    "0809/0809.4624_arXiv.txt": {
        "abstract": "We introduce a new CMB temperature likelihood approximation called the Gaussianized Blackwell-Rao (GBR) estimator. This estimator is derived by transforming the observed marginal power spectrum distributions obtained by the CMB Gibbs sampler into standard univariate Gaussians, and then approximate their joint transformed distribution by a multivariate Gaussian. The method is exact for full-sky coverage and uniform noise, and an excellent approximation for sky cuts and scanning patterns relevant for modern satellite experiments such as WMAP and Planck. The result is a stable, accurate and computationally very efficient CMB temperature likelihood representation that allows the user to exploit the unique error propagation capabilities of the Gibbs sampler to high $\\ell$'s. A single evaluation of this estimator between $\\ell =2$ and 200 takes $\\sim 0.2$ CPU milliseconds, while for comparison, a singe pixel space likelihood evaluation between $\\ell = 2$ and 30 for a map with $\\sim2500$ pixels requires $\\sim20$ seconds. We apply this tool to the 5-year WMAP temperature data, and re-estimate the angular temperature power spectrum, $C_{\\ell}$, and likelihood, $\\mathcal{L}(C_{\\ell})$, for $\\ell \\le 200$, and derive new cosmological parameters for the standard six-parameter $\\Lambda$CDM model. Our spectrum is in excellent agreement with the official WMAP spectrum, but we find slight differences in the derived cosmological parameters. Most importantly, the spectral index of scalar perturbations is $n_{\\textrm{s}} = 0.973\\pm 0.014$, $1.9\\sigma$ away from unity and $0.6\\sigma$ higher than the official WMAP result, $n_{\\textrm{s}} = 0.965\\pm 0.014$. This suggests that an exact likelihood treatment is required to higher $\\ell$'s than previously believed, reinforcing and extending our conclusions from the 3-year WMAP analysis. In that case, we found that the sub-optimal likelihood approximation adopted between $\\ell=12$ and 30 by the WMAP team biased $n_{\\textrm{s}}$ low by $0.4\\sigma$, while here we find that the same approximation between $\\ell=30$ and 200 introduces a bias of $0.6\\sigma$ in $n_{\\textrm{s}}$. ",
        "introduction": "\\label{sec:introduction}  Detailed measurements of fluctuations in the cosmic microwave background (CMB) have established cosmology as a high-precision science. One striking illustration of this is the fact that it is today possible to predict a vast number of observables based on six numbers only, with only a few (but nevertheless intriguing) ``glitches'' overall. The key to this success has been making accurate measurements of the CMB power spectrum, perhaps most prominently exemplified by Wilkinson Microwave Anisotropy Probe (WMAP; Bennett et al.\\ 2003; Hinshaw et al.\\ 2007, Hinshaw et al.\\ 2008).  The primary connection between theoretical models and CMB observations is made through the CMB likelihood, $\\mathcal{L}(C_{\\ell}) = P(\\mathbf{d}|C_{\\ell})$. This is a multivariate, non-Gaussian function that quantifies the match between the data and a given power spectrum, $C_{\\ell}$. Unfortunately, it is impossible to evaluate this function explicitly for modern high-resolution data sets, due to the sheer size of the problem, and one therefore instead typically resolves to various approximations.  However, given the importance of the CMB in modern cosmology, it is of critical importance to characterize this likelihood accurately, and all approximations must be thoroughly verified. One example is the approximation of the large angular scale likelihood, where $\\mathcal{L}(C_{\\ell})$ is strongly non-Gaussian. This turned out to be a non-trivial issue after the original analysis of the 3-year WMAP temperature data by \\citet{hinshaw:2007}, in which a Master-based \\citep{hivon:2002, verde:2003} approximation was used at $\\ell > 12$. An exact likelihood analysis \\citep{eriksen:2007b} later demonstrated that this sub-optimal approximation, when applied to harmonic modes between $\\ell=13$ and 30, biased the spectral index of scalar perturbations, $n_{\\textrm{s}}$, low by $0.4\\sigma$.  A second example is that of non-cosmological foregrounds. Unless properly accounted for, such foregrounds bias the observed power spectrum to high values, and can seriously compromise any cosmological conclusions. While important for temperature observations, this is an absolutely crucial issue for polarization observations, as the desired CMB in amplitude is comparable to or weaker than the interfering foregrounds over most of the sky.  In recent years, a new set of statistical methods have been developed that allows the user to address these issues within a single well-defined framework \\citep{jewell:2004, wandelt:2004, eriksen:2004}. The heart of this method is the Gibbs sampling algorithm (see, e.g, Gelfand and Smith 1990), in which samples from a (typically complicated) joint distribution are drawn by alternately sampling from (simpler) conditional distributions. In the CMB setting, this is realized by drawing joint samples from $P(\\mathbf{s}, C_{\\ell}|\\mathbf{d})$, by alternately sampling from $P(C_{\\ell}|\\mathbf{s}, \\mathbf{d})$, where $C_{\\ell}$ is the CMB power spectrum, $\\mathbf{s}$ is the CMB sky signal, and $\\mathbf{d}$ are the observed data. In addition to allow for exact likelihood analysis at reasonable computational cost, an equally important feature of this framework is its unique capability of including additional degrees of freedom, such as non-cosmological foregrounds, into the analysis \\citep{eriksen:2008a, eriksen:2008b}. Further, very recently an additional Metropolis-Hastings MCMC sampling step was introduced by \\citet{jewell:2008}, that effectively resolves the previously described inefficiency of the Gibbs sampler at low signal-to-noise \\citep{eriksen:2004}. The framework has also been extended to handle polarization \\citep{larson:2007, eriksen:2007b} and anisotropic universe models \\citep{groeneboom:2008}.  By now, the CMB Gibbs sampler is well established and demonstrated to sample efficiently from the exact CMB posterior. However, a long-standing issue has been the characterization of the joint likelihood, given a set of such samples. Originally, \\citet{wandelt:2004} proposed to use the so-called Blackwell-Rao (BR) estimator for this purpose, and this approach was later implemented and studied in detail by \\citet{chu:2005}. While highly accurate for the large angular scale and high signal-to-noise temperature likelihood, it suffers from one major drawback: Because it attempts to describe the full $\\ell_{\\textrm{max}}$-dimensional likelihood without any constraints on allowed correlations, the number of samples required for convergence scales exponentially with $\\ell_{\\textrm{max}}$. In practice, this limits the BR estimator to $\\ell \\lesssim 30$ for temperature data, and just $\\ell \\lesssim 3-4$ for low signal-to-noise polarization data.  In this paper, we introduce a new temperature likelihood approximation based on samples drawn from the CMB posterior, by modifying the original BR estimator in a way that restricts the allowed $N$-point functions of $\\mathcal{L}(C_{\\ell})$, but still captures most of the relevant information. Explicitly, this is done through a specific change of variables, such that the observed marginal posterior for each multipole, $P(C_{\\ell}|\\mathbf{d})$, is transformed into a Gaussian. Then, in these new variables the joint distribution is approximated by a multivariate Gaussian. As long as the correlation between any two multipoles is reasonably small, as is the case for nearly full-sky experiments such as WMAP and Planck, we shall see that this provides an excellent approximation to the exact joint likelihood. As a result, the new approach greatly reduces the overall number of samples required for convergence, and allows us to obtain a highly accurate likelihood approximation to arbitrary $\\ell_{\\textrm{max}}$. Generalization to a full polarized likelihood will be discussed in a future paper (Eriksen et al., in preparation).  This paper is organized as follows: In \\S \\ref{sec:review}, we first briefly review the Gibbs sampling algorithm together with the original Blackwell-Rao estimator, and in \\S\\ref{sec:method} we introduce the new Gaussianized Blackwell-Rao estimator. Next, in \\S \\ref{sec:application}, we apply the new estimator to simulated data, and compare results with brute-force likelihood evaluations in pixel space. In \\S \\ref{sec:analysis}, we analyze the 5-year WMAP temperature data, and provide an updated power spectrum and set of cosmological parameters. We summarize and conclude in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a new likelihood approximation to be used within the CMB Gibbs sampling framework. This approximation is defined by Gaussianizing the observed marginal power spectrum posteriors, $P(C_{\\ell}|\\mathbf{d})$, through a specific change-of-variables, and then coupling these univariate posteriors into a joint distribution through a multivariate Gaussian in the new variables. This process is exact, i.e., an identity operation, in the uniform and full-sky coverage case, and it is also an excellent approximation in for the moderate sky cuts relevant to satellite missions such as WMAP and Planck.  Our new approach relies on the previously described CMB Gibbs sampling framework \\citep{jewell:2004,wandelt:2004,eriksen:2004}, and thereby inherits many important advantages from that. First and foremost, this framework allows for seamless propagation of uncertainties from various systematic effects (e.g., foregrounds, beam uncertainties, calibration or noise estimation errors) to the final cosmological parameters. This is not straightforward in the hybrid scheme used by the WMAP code. Second, this new approach corresponds to the exact low-$\\ell$ pixel space likelihood part of the WMAP code, not the approximate high-$\\ell$ MASTER part. Still, our method can handle arbitrary high $\\ell$'s. Third, once the one-time pre-processing step has been completed, the computational expense of our estimator is determined by the cost of $\\ell_{\\textrm{max}}$ spline evaluations, while a pixel space approach requires a matrix inversion, and therefore scales as $\\mathcal{O}(N_{\\textrm{pix}}^3)$. For the cases considered in this paper, the CPU time required for the GBR WMAP estimator up to $\\ell=200$ was $\\sim0.2$ milliseconds, while it was $\\sim20$ seconds for the pixel space approach up to $\\ell=32$, for a map with $~2500$ pixels.  In order to validate our estimator, we applied it to a low-resolution simulated data set, and compared it to slices through the exact joint likelihood as computed by brute-force evaluation in pixel space. The agreement between the two approaches was excellent. We then applied the same estimator to the 5-year WMAP temperature data, and estimated both a new power spectrum and new cosmological parameters within a standard six-parameter $\\Lambda$CDM model.  The results from these calculations are interesting. First, our power spectrum is statistically very similar to the official WMAP spectrum, with no visible biases seen and relative fluctuations within the level predicted by noise and sky cut. Nevertheless, we do find significant differences in terms of cosmological parameters, and most notably in the spectral index of scalar perturbations, $n_{\\textrm{s}}$. Specifically, we find $n_{\\textrm{s}} = 0.973\\pm 0.014$, which is only $1.9\\sigma$ away from unity and $0.6\\sigma$ higher than the official WMAP result, $n_{\\textrm{s}} = 0.965\\pm 0.014$.  This result resembles very much the outcome of a re-analysis we did with the 3-year WMAP temperature data \\citep{eriksen:2007a}, for which we found a bias of $0.4\\sigma$ in $n_{\\textrm{s}}$ compared to the official WMAP results. This bias was due to the sub-optimal MASTER-based likelihood approximation \\citep{hivon:2002,verde:2003} used by the WMAP team between $\\ell=12$ and 30, whereas we used an exact estimator in the same range. This study later prompted the WMAP to change their codes to use an exact likelihood evaluator up to $\\ell=30$.  In the same study, we also tried to increase the $\\ell$-range for our exact estimator to $\\ell=50$, but found small differences. We therefore concluded, perhaps somewhat prematurely, that an exact estimator up to $\\ell=30$ was sufficient for obtaining accurate results. On the contrary, in this paper we find still find significant changes when increasing the exact estimator up to $\\ell=200$.  In retrospect, this should perhaps not come as a complete surprise, when realizing that the impact on a particular cosmological parameter typically depends logarithmically on $\\ell$. For instance, \\citet{hamimeche:2008} considered a simple power spectrum model with a single free amplitude, $C_{\\ell} = q\\,C_{\\ell}^{\\textrm{fid}}$, and found that, for a given likelihood estimator to be ``statistically unbiased'', the systematic errors in that same estimator must fall off faster than $\\sim 1/\\ell$.  \\begin{deluxetable}{cccc} \\tablewidth{0pt} \\tablecaption{Marginal 5-year WMAP cosmological parameters\\label{tab:parameters}} \\tablecomments{Comparison of cosmological parameters obtained with the standard 5-year WMAP likelihood code (second column) and with the new GBR estimator at $\\ell\\le200$ (third column), given in terms of marginal means and standard deviations. The shift between the two in units of $\\sigma$ is listed in the fourth column.} \\tablecolumns{4} \\tablehead{Parameter  & WMAP  & GBR  & Shift in $\\sigma$ } \\startdata $\\Omega_{\\textrm{b}}h^2$ & $0.0228\\pm0.0006$ & $0.0230\\pm0.0006$ &    0.4 \\\\ $\\Omega_{\\textrm{c}}h^2$ & $0.109\\pm0.006$ & $0.0108\\pm0.006$ &    -0.3 \\\\ $\\textrm{log}(10^{10}A_{\\textrm{s}})$  & $3.06\\pm0.04$ & $3.06\\pm0.04$ &    0.0 \\\\ $h$                    & $0.722\\pm0.03$ & $0.732\\pm0.03$ &     0.3 \\\\ $n_{\\textrm{s}}$         & $0.965\\pm0.014$ & $0.973\\pm0.014$ &    0.6 \\\\ $\\tau$                 & $0.090\\pm0.02$ & $0.090\\pm0.02$ &    0.0 \\enddata \\end{deluxetable}  A similar consideration holds for $n_{\\textrm{s}}$. Intuitively, $n_{\\textrm{s}}$ is as much affected by $\\ell=2$ to 10 as it is between $\\ell=20$ and 100. In the previous 3-year WMAP re-analysis paper, we increased the range of the accurate likelihood estimator from $\\ell=12$ to 30, corresponding to a factor of 2.5 in $\\ell$, and removed a bias of $\\sim0.4\\sigma$ in $n_{\\textrm{s}}$. In this paper, we increase the range from $\\ell=30$ to 200, corresponding to a factor of 6.7 in $\\ell$, and find an additional bias of $0.6\\sigma$. However, increasing $\\ell$ from 30 to 50 corresponds only to a factor of 1.7 in $\\ell$, and this appears to be too small to produce a statistically significant result.  The main conclusions from this work are two-fold. First, it seems that an accurate likelihood description is required to higher $\\ell$'s than previously believed, and at least up to $\\ell=200$, in order to obtain unbiased results. By extrapolation, it also does not seem unlikely that even higher multipoles should be included. This issue will be revisited in a future publication.  Our second main conclusion is that we find a spectral index only $1.9\\sigma$ away from unity, namely $n_{\\textrm{s}} = 0.973\\pm0.014$. To us, it therefore seem premature to make strong claims concerning $n_{\\textrm{s}}\\ne 1$; the statistical significance of this is rather low, and there are likely still unknown systematic errors in this number.  In a future publication we will generalize the GBR estimator to polarization. Once completed, this will enable a fully Gibbs-based CMB likelihood analysis at low $\\ell$'s, and remove the need for likelihood techniques based on matrix operations, i.e., inversion and determinant evaluation. The computational cost of a standard cosmological parameter MCMC analysis (e.g., CosmoMC) will then once again be driven by the required Boltzmann codes (e.g., CAMB or CMBFast) and not by the likelihood evaluation. In turn, this will increase the importance of fast interpolation codes such as Pico \\citep{fendt:2007} or COSMONET \\citep{auld:2007}. With such fast algorithms for both spectrum and likelihood evaluations ready at hand, the CPU requirements for cosmological parameter estimation may possibly be reduced by orders of magnitude."
    },
    "0809/0809.1976_arXiv.txt": {
        "abstract": "We explore the possibility improving the $\\Lambda$CDM model at megaparsec scales by introducing a scalar interaction that increases the mutual gravitational attraction of dark matter particles. Using N-body simulations, we study the spatial distribution of dark matter particles and halos. We measure the effect of modifications in the Newton's gravity on properties of the two-point correlation function, the dark matter power spectrum, the cumulative halo mass function and density probability distribution functions. The results look promising: the scalar interactions produce desirable features at megaparsec scales without spoiling the $\\Lambda$CDM successes at larger scales. ",
        "introduction": "Introduction}   In this paper we study cosmological implications of a scalar interaction that produces a long-range fifth force in the dark sector, proposed by G.~Farrar , S.~Gubser and J.~Peebles \\cite{FarrarPeebles,GP1,GP2,Farrar2007,Brookfield}. The physical motivation for this model comes from the string theory \\cite{StringGas}.  The cosmological motivation comes from small-scale difficulties of the $\\Lambda$CDM model, which successfully passes almost all observational tests (see e.g. \\cite{WMAP07}). Difficulties appear at length scales below few megaparsecs: (1) the $\\Lambda$CDM voids are less empty than the real voids; (2) the present accretion rate of intergalactic debris onto thin spiral galaxies poses a problem for current galaxy formation paradigm; (3) merger rates at low redshifts for giant elliptical galaxies suggest violent accretion histories for their haloes at low redshifts which is in contradiction with observations.   The void problem, pointed out by J. Peebles \\cite{PeeblesVoid} is hotly debated in the literature, with arguments supporting his original claim \\cite{TikhonovKlypin1,TikhonovKlypin2} as well as arguments to the contrary \\cite{Patiri,Tinker1}.  The late merging problem appears in simulations which shows that accretion onto giant elliptical galaxies at cluster centers continues until z = 0 \\cite{Gao} while the independence of the color-magnitude relation of SDSS\\footnote{\\textit{Sloan Digital Sky Survey} http://www.sdss.org/} galaxies on their environment \\cite{Hogg2004} and the remarkable stability of the color-magnitude relation at z = 0.7 \\cite{Cassata2007} is consistent with the picture of giant galaxies as island universes, contrary to $\\Lambda$CDM simulations. Likewise, according to J. Peebles, the very existence of large galaxies like the Milky Way, with its spiral structure intact, suggests a lack of major mergers in recent history (see \\cite{Disksurv} for thin disk dominated galaxies survival issues). There is observational evidence that the latest ``major invasion'' happened to our Galaxy 10-12 Gyr ago \\cite{Gilmore2002}. There are also arguments to the contrary, claiming that Milky Way-like galaxies survive late merging in N-body simulations \\cite{LCDMdisks}.  For an excellent discussion of the observational situation and comparison with the $\\Lambda$CDM model, see Refs. \\cite{6puzzles,NGP,PeeblesLambda-DE}, and the references therein.  Theoretical suggestions of Peebles and Gubser were followed by the work of Nusser \\textit{et al.} \\cite{NGP}, exploring the cosmological implications with N-body simulations. Some preliminary results on similar scalar model were also obtained by Rodr{\\'{\\i}}guez-Meza \\textit{et al.} \\cite{Rodriguez-Meza1,Rodriguez-Meza2,Rodriguez-Curves}. The Long-Range Scalar Interaction (LRSI) model started a debate in the literature recently, mainly focused on the weak-equivalence principle violation and it's impact on dynamic of Milky-Way satellites \\cite{Satellites1,Satellites2,Satellites3}.  The work, presented here should be regarded as the next step in this line of research. Like Nusser \\textit{et al.}, we study the two-point correlation function and the statistical properties of the mass distribution of the dark matter halos. To our knowledge, for the first time in the literature, we also study the power spectra for a set of scalar field parameters with comoving screening length. Analogous model was analyzed in great detail by Grawdohl \\& Frieman nearly 20 years ago \\cite{GradwohlFrieman1,GradwohlFrieman2}. However, their model assumed a fixed physical screening length, which is a fundamental difference in comparison to our model. We have used more particles in our simulations compared to those used by Nusser \\textit{et al.}, and our resolution is better than their resolution; we also consider a wider range of scalar model parameters. As a consequence, we resolve the power spectrum near the screening length characteristic scale. Our results show a clear feature in the power spectrum near a wave number,that is the inverse of the screening length. The extra power seen at higher wavenumbers is generated by the gravitational field, enhanced by the scalar interaction. Nusser \\textit{et al.} could not see a similar feature in their correlation function because the screening lengths they considered were too close to the mean interparticle separation in their simulations. Apart from this important difference, we confirm their results. The scalar field generates lower density in voids. It also shifts structure formation to higher redshifts.  This paper is organized as follows: In section \\ref{model} we introduce the effective gravitational potential and modified force law used as an approximation of the scalar field. In Section \\ref{sim} we describe our N-body simulations. Our results are presented in Section \\ref{results}. A brief summary and discussion appears in Section \\ref{CR}. ",
        "conclusions": ""
    },
    "0809/0809.3142_arXiv.txt": {
        "abstract": "We investigate the instability and nonlinear saturation of temperature-stratified Taylor-Couette flows in a finite height cylindrical gap and calculate angular-momentum transport in the nonlinear regime. The model is based on an incompressible fluid in Boussinesq approximation with a positive axial temperature gradient applied. While both ingredients themselfs, the differential rotation as well as the stratification due to the temperature gradient, are stable, together the system becomes subject of the stratorotational instability and nonaxisymmetric flow pattern evolve. This flow configuration transports angular momentum outwards and will therefor be relevant for astrophysical applications. The belonging coefficient of $\\beta$-viscosity is of the order of unity if the results are adapted to the size of an accretion disc. The strength of the stratification, the fluids Prandtl number and the boundary conditions applied in the simulations are well-suited too for a laboratory experiment using water and a small temperature gradient around five Kelvin. With such a set-up the stratorotational instability and its angular momentum transport could be measured in an experiment. ",
        "introduction": "In recent years instabilities in stratified media became of higher interest. Especially in view of astrophysical objects the inclusion of stratification is relevant. One simple model to study stratification effects is the classical Taylor-Couette (TC) system. \\cite{thorpe_1968} found a stabilizing effect of stratification only. In the context of stratorotational instability (SRI) stratification in TC flows is investigated numerically with fixed axial temperature gradient in a linear analysis studying the suppression of the onset of Taylor vortices by \\cite{boub_1996}. With fixed density gradient \\cite{mol_2001}, \\cite{yav_2001} and \\cite{shal_2005} show the onset of a linear instability and the growth of nonaxisymmetric modes. Experiments using artifically enlarged buoyancy due to salt concentration \\cite[see][]{with_1974,boub_1995,lebars_2006} are in very good agreement with the linear results. \\cite{caton_2000} use an artifical diffusivity in the continuity equation and show also results of linear stability analysis for the case of non-rotating outer cylinder. \\cite{umur_2006} analyses SRI analytically, especially the influence of vertically changing buoyancy frequency, in the quasi-hydrostatic semi-geostrophic limit. The author also shows that SRI survives only in the presence of no-slip radial boundary conditions. In the context of accretion disks, \\cite{dub_2005} figure out stability conditions and the influence of viscous dissipation and thermal diffusivity. In particular they find that the Prandtl number dependence of the critical parameters is unimportant. But fully nonlinear three-dimensional simulations do not exist. It is the aim of the first part of this paper to describe the characteristics of SRI in the nonlinear regime. The easiest way to do nonlinear simulations is the application of an axial temperature gradient, where the top of the TC system has higher temperature than the bottom. The absence of a diffusivity in the continuity equation makes nonlinear simulations with explicit density gradient more demanding and favors a temperature gradient. Further on this is interesting also from an astrophysical view point: accretion disks are heated from the central object at their top and bottom.  Thus the TC system heated from above could be seen as a simplified model for half of a disc. Of major interest for accretion or protostellar discs is the problem of angular momentum redistribution. Accretion only works if angular momentum is transported outwards effectively. Angular momentum transport of SRI is the second aspect of this work. ",
        "conclusions": "We have shown fully nonlinear simulations of the SRI with temperature stratification in a cylindrical annulus. The stratification is stable, as well as the differentially rotating flow. Both together lead to unstable nonaxisymmetric modes. Depending on the gap width, these are the $m=1$ or $m=2/3/4/5$ modes in the investigated parameter range of rather weak stratification with Froude numbers around $\\Fr=1$. Weak stratification results in the lowest critical Reynolds numbers for the onset of the SRI. On the other hand the instability is influenced by the Prandtl number of the flow only slightly. Especially the dominating mode $m$ does not depend on $\\Pr$, only on $\\Fr$. The time the SRI needs to evolve and reach a saturated state is of the order of 120 rotations or two times the viscous time scale. Thus the growth rate of the instability is rather slow compared to magnetorotational instability and Tayler instability, where it is of the order of ten rotations. For both instabilities magnetic effects play the essential role. The SRI, even if slower, might be of comparable importance, when magnetic effects are rather unimportant or if stratification suppresses MRI or TI. Thus it might become the most efficient instability mechanism in environments with weak or very strong magnetic fields or in low-conducting environments like protostellar discs. We indeed find that SRI could be a mechanism able to transport angular momentum. The normalized angular momentum transport in terms of the $\\beta$ viscosity is of the order of $10^{-3}$, in terms of the $\\alpha_{\\rm SS}$ for a thin disk around unity. This comparison with a thin disk assumes that SRI still occurs in such a flat disk, which is not possible to answer with our simulations at the moment. Nevertheless is the size of $\\beta$ and its linear growth with $\\Re$ a sign of significant influence of the SRI for angular momentum transport which might dominate over magnetic effects to produce turbulence in low-conducting environments. Beside astrophysical motivation angular momentum transport and turbulent transport coefficients are important also for technical applications. In tradition of the idea of Couette to measure viscosity of a fluid, the presented experimental configuration could be used to measure the transport of angular momentum or the increase of viscosity in the laboratory. With water it needs a TC system with $\\Gamma\\approx 10$ and a gap width of \\mbox{$6$ cm} to observe the SRI with a temperature difference between top and bottom of \\mbox{$5$ K}. For Reynolds numbers around $\\Re=1000$ the time the instability needs to grow and to saturate is around \\mbox{$25$ min} or 120 rotations with \\mbox{$0.5$ Hz}. All these conditions seem to be appropriate for an experiment."
    },
    "0809/0809.4886_arXiv.txt": {
        "abstract": "We discuss how the solar occultations of bright sources of energetic gamma rays can be used to extract non--trivial physical and astrophysical information, including the angular size of the image when it is significantly smaller than the experiment's angular resolution. We analyze the EGRET data and discuss prospects for other instruments. The Fermi Gamma Ray Space Telescope will be able to constrain the size of a possible halo around 3C~279 from observations it makes on the 8th of October each year. ",
        "introduction": "\\label{sec:intro} The brightest source in the sky almost at any wavelength, the Sun is very weak in high-energy ($E \\gtrsim 100$~MeV) gamma rays. This property can be used to study solar occultations of gamma-ray sources.  The width of the point-spread function (PSF) of telescopes detecting photons at these energies is quite large, of order several degrees. The enormous exposure of the Fermi Gamma Ray Space Telescope (the telescope previously known as GLAST) would partially compensate for the poor resolution; however, it would be almost impossible to directly measure the angular size of the image which may be smaller than the PSF width. On the other hand, energetic gamma-ray images of distant sources may indeed have a significant angular size due to the cascading of photons on the background radiation and magnetic deflections of the cascade electrons and positrons. It has long been known that one can obtain the angular size of stars from lunar occultations, we suggest that one may determine the image size of gamma-ray sources screened by the Sun \\footnote{Interestingly, we note that the Moon is much brighter than the Sun in this energy band because of secondary emission from cosmic rays hitting the lunar surface, see e.g.\\ \\citet{EGRET:SunMoon,Brigida:2009}.}.  The current collection of known energetic gamma-ray point sources is scarce ($\\sim 300$ sources detected by EGRET), so only a few are expected to be on the strip on the sky such that they are screened by the Sun. Fortunately, the brightest EGRET source identified with an extragalactic object, 3C~279, has an ecliptic latitude of $0.2^\\circ$ and is screened by the Sun on the 8th of October each year. It is 3C~279 which is the main subject of our discussion because, as we will see in Sec.~\\ref{sec:extended}, it represents a perfect target for this kind of study.  The simplest and most direct effect of an extended image size would be the detection of flux from the source during occultation.  Such a result could also be the signal of the transparency of the Sun to gamma rays possible in several scenarios of new physics \\citep{FRT}; however the parameter space of particular models is strongly disfavoured by results of other experiments (e.g.\\ \\citet{CAST2007}).  In Sec.~\\ref{sec:EGRET}, we review our analysis of archival EGRET data of the 1991 occultation of 3C~279 during which a non-zero flux was indeed observed, although at a very low statistical significance.  With a sensitive enough telescope, a more detailed study of the light curve during ingress and egress would be possible. In Sec.~\\ref{sec:GLAST}, we discuss the potential of Fermi for this kind of a study and mention sources other than 3C~279 while Sec.~\\ref{sec:concl} summarizes our conclusions. ",
        "conclusions": "\\label{sec:concl} It will be interesting to try and measure the angular sizes of images of energetic gamma--ray sources by means of observation of their solar occultations. The best target is 3C~279, whose occultation happens each year on the 8th of October.  EGRET observations made during such a period did not exclude the unsuppressed flux of the quasar when it was screened by the Sun. The sensitivity of the Fermi telescope is high enough that if the flux was unsuppressed during occultation, it could be observed more definitively than with EGRET.  Fermi can also constrain the angular size of the image even in the survey mode, and is capable of  obtaining a light curve by the combination of several observations in the pointing mode. This would help to constrain models of particle acceleration and magnetic fields in and around the quasar. If the flux during the occultation exceeds the non-variable flux of the source (as it is slightly favoured by the EGRET data), it would mean that either the extended image is formed relatively nearby or the Sun is partially transparent for the point-like gamma-ray emission (both options would mean a discovery of some unconventional physical or astrophysical phenomenon). The same method may be applied to refine the coordinates and/or to estimate the angular size of images of other gamma-ray sources screened by the Sun.  We are grateful for discussions with J.~Conrad, V.~Rubakov and M.~Tavani.  This work was supported in part by DFG (Germany) and CONACYT (Mexico) (TR), by the grants RFBR 07-02-00820, RFBR 09-07-00388, NS-1616.2008.2 and by FASI under state contracts 02.740.11.0244 and 02.740.11.5092 (ST). We made use of NED which is operated by the Jet Propulsion Laboratory, CalTech, under contract with NASA."
    },
    "0809/0809.3374_arXiv.txt": {
        "abstract": "We consider the linear growth of matter perturbations on low redshifts in some $f(R)$ dark energy (DE) models. We discuss the definition of dark energy (DE) in these models and show the differences with scalar-tensor DE models. For the $f(R)$ model recently proposed by Starobinsky we show that the growth parameter $\\gamma_0\\equiv \\gamma(z=0)$ takes the value $\\gamma_0\\simeq 0.4$ for $\\Omega_{m,0}=0.32$ and $\\gamma_0\\simeq 0.43$ for $\\Omega_{m,0}=0.23$, allowing for a clear distinction from $\\Lambda$CDM. Though a scale-dependence appears in the growth of perturbations on higher redshifts, we find no dispersion for $\\gamma(z)$ on low redshifts up to $z\\sim 0.3$, $\\gamma(z)$ is also quasi-linear in this interval. At redshift $z=0.5$, the dispersion is still small with  $\\Delta \\gamma\\simeq 0.01$. As for some scalar-tensor models, we find here too a large value for $\\gamma'_0\\equiv \\frac{d\\gamma}{dz}(z=0)$, $\\gamma'_0\\simeq -0.25$ for $\\Omega_{m,0}=0.32$ and $\\gamma'_0\\simeq -0.18$ for $\\Omega_{m,0}=0.23$. These values are largely outside the range found for DE models in General Relativity (GR). This clear signature provides a powerful constraint on these models. ",
        "introduction": "The present stage of accelerated expansion of the universe  \\cite{P97} is a major challenge for cosmology. There are basically two main roads one can take: either to introduce a new (approximately) smooth component or to modify the laws of gravity on cosmic scales. In the first class of models, the new component dubbed dark energy (DE) must have a sufficiently negative pressure in order to induce a late-time accelerated expansion. One typically adds a new isotropic perfect fluid with negative pressure able to accelerate the expansion. In the second class of models one is trying to explain the accelerated expansion by modifying gravity, after all the universe expansion is a large-scale gravitational effect. It is far from clear at the present time which class of DE models will finally prevail and one must study carefully all possibilities. While all models are called DE models \\cite{SS00}, the models of the second class are called more specifically modified gravity DE models.  It is clear that in both classes we have many models able to reproduce a late-time accelerated expansion in agreement with the present data probing the background expansion like luminosity distances from SNIa. Probably this will remain the case even with more accurate data. However the evolution of matter perturbations will be affected differently depending on the class of models we are dealing with. This can be used for a consistency check in order to find out whether a DE model is inside General Relativity (GR) or not \\cite{S98}. Therefore the evolution of matter perturbations seems to be a powerful tool to discriminate between models inside and outside GR \\cite{B06}. In particular, the growth of matter perturbations on small redshifts seems a promising tool and it has been the subject of intensive investigations recently \\cite{HL06}.  A particular class of modified gravity DE models are $f(R)$ models where $R$ is replaced by some function $f(R)$ in the gravitational Lagrangian density \\cite{Cap02} (for a recent review, see \\cite{SF08}). The idea of producing an accelerated stage in such models was first successfully implemented in Starobinsky's inflationary model \\cite{S80}. In some sense it was rather natural to try explaining the late-time accelerated expansion by the same kind of mechanism, a modification of gravity but now at much lower energies. Some problems arising in these $f(R)$ DE models were quickly pointed out, related to solar-system constraints \\cite{Chi03} and instabilities \\cite{Dol03}. It was then found that a very serious problem arises where it was not expected: many of these models are unable to produce the late-time acceleration together with a viable cosmic expansion history \\cite{APT07}. It was shown that many popular $f(R)$ models, like those containing a $R^{-1}$ term \\cite{CDTT04}, are unable to account for a viable expansion history because of the absence of a standard matter-dominated stage $a\\propto t^{\\frac{2}{3}}$ which is replaced instead by the behaviour $a\\propto t^{\\frac{1}{2}}$ \\cite{AGPT07}.  However some interesting cosmological models remain viable like the model recently suggested by Starobinsky \\cite{S07}, (for another interesting viable model see also \\cite{HS07}). It is interesting in the first place because of its non-trivial ability to allow for a standard matter-dominated stage. This goes together with the appearance of large oscillations in the past and the overproduction of massive scalar particles in the very early universe which, as noted already in \\cite{S07}, can be a serious problem of all viable $f(R)$ DE models. Another interesting property is the scale dependence of the growth of matter perturbations deep inside the Hubble radius \\cite{T07},\\cite{B07}. This scale dependence was actually used in \\cite{S07} in order to constrain one of the parameters of the model. Of course this model also satisfies the local gravity constraints for certain parameter values. In the end a window in parameter space remains where the model is in principle physically acceptable.  In this letter we want to assess the viability of this model with respect to the growth of matter perturbations (see e.g. \\cite{DJJNST08} for other possible constraints and approaches). Some viable $f(R)$ DE models have their free parameters constrained in such a way that they cannot be distinguished observationally from $\\Lambda$CDM. The situation is different in this case. Our results show that the growth of matter perturbations on low redshifts $z\\lesssim 1$ provides a discriminating signature of these models, able to clearly differentiate it from $\\Lambda$CDM and also from any other DE model inside GR. These results confirm that the growth of matter perturbations can be used efficiently to track the nature of DE models. ",
        "conclusions": ""
    },
    "0809/0809.4248_arXiv.txt": {
        "abstract": "Data from the power output of the radioisotope thermoelectric generators aboard the Cassini spacecraft are used to test the conjecture that small deviations observed in terrestrial measurements of the exponential radioactive decay law are correlated with the Earth-Sun distance.  No significant deviations from exponential decay are observed over a range of $0.7 - 1.6 A.U.$ A $90\\%$ Cl upper limit of $0.84\\times10^{-4}$ is set on a term in the decay rate of $^{238}Pu$ proportional to $1/R^{2}$ and $0.99\\times10^{-4}$ for a term proportional to $1/R$.  \\vskip 0.5cm {PACS numbers: 23.60.+e, 23.40.-s, 95.55.Pe, 96.60.Vg, 0620.Jr.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.4281_arXiv.txt": {
        "abstract": "We present results from our X-ray data analysis of the supernova remnant (SNR) G330.2+1.0 and its central compact object (CCO), CXOU J160103.1--513353 (J1601 hereafter). Using our {\\it XMM-Newton} and {\\it Chandra} observations, we find that the X-ray spectrum of J1601 can be described by neutron star atmosphere models ($T^{\\infty}$ $\\sim$ 2.5--5.5 MK). Assuming the distance of $d$ $\\sim$ 5 kpc for J1601 as estimated for SNR G330.2+1.0, a small emission region of $R$ $\\sim$ 0.4--2 km is implied. X-ray pulsations previously suggested by {\\it Chandra} are not confirmed by the {\\it XMM-Newton} data, and are likely not real. However, our timing analysis of the {\\it XMM-Newton} data is limited by poor photon statistics, and thus pulsations with a relatively low amplitude (i.e., an intrinsic pulsed-fraction $<$ 40\\%) cannot be ruled out. Our results indicate that J1601 is a CCO similar to that in the Cassiopeia A SNR. X-ray emission from SNR G330.2+1.0 is dominated by power law continuum ($\\Gamma$ $\\sim$ 2.1--2.5) which primarily originates from thin filaments along the boundary shell. This X-ray spectrum implies synchrotron radiation from shock-accelerated electrons with an exponential roll-off frequency $\\nu_{\\rm rolloff}$ $\\sim$ 2--3 $\\times$ 10$^{17}$ Hz. For the measured widths of the X-ray filaments ($D$ $\\sim$ 0.3 pc) and the estimated shock velocity ($v_s$ $\\sim$ a few $\\times$ 10$^3$ km s$^{-1}$), a downstream magnetic field $B$ $\\sim$ 10--50 $\\mu$G is derived. The estimated maximum electron energy $E_{\\rm max}$ $\\sim$ 27--38 TeV suggests that G330.2+1.0 is a candidate TeV $\\gamma$-ray source. We detect faint thermal X-ray emission in G330.2+1.0. We estimate a low preshock density $n_0$ $\\sim$ 0.1 cm$^{-3}$, which suggests a dominant contribution from an inverse Compton mechanism (than the proton-proton collision) to the prospective $\\gamma$-ray emission. Follow-up deep radio, X-ray, and $\\gamma$-ray observations will be essential to reveal the details of the shock parameters and the nature of particle accelerations in this SNR. ",
        "introduction": "OBSERVATIONS \\& DATA REDUCTION}  We observed G330.2+1.0/J1601 with the European Photon Imaging Camera (EPIC) on board {\\it XMM-Newton} Observatory on 2008-03-20 (ObsID 0500300101). The pointing (RA[J2000] = 16$^h$ 01$^m$ 3$\\fsecs$14, Dec[J2000] = --51$^{\\circ}$ 33$^{\\prime}$ 53$\\farcs$6) is to J1601 which is positioned at the center of the nearly circular X-ray shell of SNR G330.2+1.0. We chose the small-window mode (4$\\farcm$4 $\\times$ 4$\\farcm$3 field of view [FOV] and 6 ms time resolution) for the EPIC pn to search for pulsations of J1601. We chose the full-window mode ($\\sim$30$'$ diameter FOV) for the EPIC MOS detectors to study the entire SNR. The medium filter was used for all detectors. We reduced the data using the Science Analysis System (SAS) software package v7.1.0.  Our {\\it XMM-Newton} observations of G330.2+1.0/J1601 were significantly contaminated by flaring particle background. We removed time bins in which the overall count rate is 2$\\sigma$ (the pn) or 3$\\sigma$ (the MOS1 and MOS2) above the mean value for time intervals unaffected by flaring background. Time intervals including a considerable contamination by the flaring background ($\\ga$ 50\\% above the average quiescent rate) were eliminated by these time-filters. After the time filtering, 26, 31, and 33 ks exposures for the pn\\footnote{The $\\sim$30\\% deadtime-corrected exposure for the small window mode of the pn is 18.3 ks.}, MOS1, and MOS2, respectively, are available for further data analysis, which is $\\sim$40--45\\% of the total exposure. We then reduced the data following the standard screening of event pattern (PATTERN $\\leq$ 12 for the MOS1/2 and PATTERN $\\leq$ 4 for the pn) and hot pixels (FLAG = 0). (For the timing analysis of J1601, we used a longer exposure while choosing a smaller aperture and more strict event pattern criteria as described in \\S~\\ref{sec:cco_time}.) There are stable components of instrumental background in the EPIC detectors. The primary components that could affect this work are Al-K ($E$ $\\sim$ 1.5 keV) and Si-K ($E$ $\\sim$ 1.7 keV) fluorescence lines due to the interactions of high energy particles with the structure surrounding the detectors and the detectors themselves\\footnote{{\\it XMM-Newton} Users' Handbook, \\S~3.3.7.2.}. We removed these events from our image analysis by excluding narrow energy bands centered on these lines. Our background-subtracted source spectra show little evidence of these lines. Thus, we believe that the impact of contamination from this instrumental background on our EPIC data analysis is negligible.  Because of the severe contamination by the flaring background, photon statistics of the filtered {\\it XMM-Newton} data are significantly lower than originally intended. Thus, in addition to the {\\it XMM-Newton} data, we use the {\\it Chandra} data (ObsID 6687) for the spectral analysis to improve overall photon statistics. The high angular resolution of {\\it Chandra} data is also essential to measure the widths of the thin X-ray filaments of G330.2+1.0.  We performed the {\\it Chandra} observation of G330.2+1.0 with the I-array of the Advanced CCD Imaging Spectrometer (ACIS) on 2006-05-22 as part of the Guaranteed Time Observations program. The effective exposure after the data screening is $\\sim$50 ks, and thus photon statistics for the SNR and CCO in the ACIS data are similar to those obtained by the EPIC MOS1+MOS2 data. The details of the {\\it Chandra} observation and data reduction are described by Park et al. (2006). ",
        "conclusions": ""
    },
    "0809/0809.0970_arXiv.txt": {
        "abstract": "We present a method of mapping dust column density in dark clouds by using near-infrared scattered light. Our observations of the Lupus 3 dark cloud indicate that there is a well defined relation between (1) the $H-K_s$ color of an individual star behind the cloud, i.e., dust column density, and (2) the surface brightness of scattered light toward the star in each of the $J$, $H$, and $K_s$ bands. In the relation, the surface brightnesses increase at low $H-K_s$ colors, then saturate and decrease with increasing $H-K_s$. Using a simple one-dimensional radiation transfer model, we derive empirical equations which plausibly represent the observed relationship between the surface brightness and the dust column density. By using the empirical equations, we estimate dust column density of the cloud for any directions toward which even no background stars are seen. We obtain a dust column density map with a pixel scale of 2.3 $\\times$ 2.3 arcsec$^2$ and a large dynamic range up to $A_V$ = 50 mag. Compared to the previous studies by Juvela et al., this study is the first to use color excess of the background stars for calibration of the empirical relationship and to apply the empirical relationship beyond the point where surface brightness starts to decrease with increasing column density. ",
        "introduction": "Dark clouds are seen as dark patches against bright star fields, while they are observed as reflection nebulae at near-infrared (NIR) wavelengths in the recent decade (Lehtinen \\& Mattila 1996, Nakajima et al. 2003, Foster \\& Goodman 2006). The images of NIR reflection nebulae give us an insight that darker areas with fewer background stars have larger column densities. In our previous paper (Nakajima et al. 2003),  we showed that the Lupus 3 dark cloud appears as an NIR reflection nebula and suggested that there is a relationship between the NIR surface brightness and dust column density toward the Lupus 3 dark cloud. It is worth investigating whether if we can quantify the relationship between the NIR surface brightness and column density.  Star counting, color excess of background stars, molecular gas and thermal dust emissions have been used for column density mapping of dark clouds. Juvela et al. (2006) reviewed these methods intensively. Each of these methods has limitations. A common limitation to all these methods is that it is hard to obtain a high spatial resolution. Padoan et al. (2006) and Juvela et al. (2006) examined the use of scattered NIR surface brightness as a new high resolution tracer of the interstellar clouds. They derived an analytical formula to convert NIR surface brightness to visual extinction. The formula is linear at low visual extinctions and starts to saturate when visual extinction reaches $\\sim$ 10 mag. Juvela et al. (2006) used numerical simulations and radiative transfer calculations to show that the NIR surface brightness can be converted to column density in the range of $A_V$ $<$ 15 mag with accuracy of better than 20\\%. Juvela et al. (2008) applied the method to a filamentary cloud in Corona Australis. The method using surface brightness and one using color excess of background star agrees below $A_V$ $\\sim$ 15 mag.  In this paper, we propose an empirical method of estimating column density of dark clouds up to $A_V$ $\\sim$ 50 mag by using NIR surface brightness. We re-examine the NIR data of the Lupus 3 dark cloud in Nakajima et al. (2003) and obtain a plausible empirical equations which fit the observed relationship between the surface brightness and column density. Using the relationship, we obtain the dust column density map from the surface brightness.   \\begin{figure} \\begin{center} \\includegraphics{figure1.eps} \\caption{Surface brightness of the Lupus 3 dark cloud in the $J$, $H$, and $K_s$ bands from top to bottom. Saturated stars are masked with circles in black. Bad pixel clusters in the J band are also shown in black. The lowest contour level denotes the 3-sigma limiting flux and the interval of contour is set to the 3-sigma limiting flux value at each band. }\\label{fig:fig1} \\end{center} \\end{figure} ",
        "conclusions": "We constructed an empirical relationship equation between the surface brightness and column density of the Lupus 3 dark cloud for each $J$, $H$, and $K_s$ band. By using the equations, we obtained a column density map with a pixel scale of 2.3 $\\times$ 2.3 arcsec$^2$ and a large dynamic range up to $A_V$ = 50 mag of the cloud from the NIR surface brightness. We expect the empirical method to be one of the new tools of column density mapping of dark clouds."
    },
    "0809/0809.2975_arXiv.txt": {
        "abstract": "We study, in this paper, the non-Gaussian features of the mass density field of neutral hydrogen fluid and the Ly$\\alpha$ transmitted flux of QSO absorption spectrum from the point-of-view of self-similar log-Poisson hierarchy. It has been shown recently that, in the scale range from the onset of nonlinear evolution to dissipation, the velocity and mass density fields of cosmic baryon fluid are extremely well described by the She-Leveque's scaling formula, which is due to the log-Poisson hierarchical cascade. Since the mass density ratio between ionized hydrogen to total hydrogen is not uniform in space, the mass density field of neutral hydrogen component is not given by a similar mapping of total baryon fluid. Nevertheless, we show, with hydrodynamic simulation samples of the concordance $\\Lambda$CDM universe, that the mass density field of neutral hydrogen, is also well described by the log-Poisson hierarchy.  We then investigate the field of Ly$\\alpha$ transmitted flux of QSO absorption spectrum. Due to redshift distortion, Ly$\\alpha$ transmitted flux fluctuations are no longer to show all features of the log-Poisson hierarchy. However, some non-Gaussian features predicted by the log-Poisson hierarchy are not affected by the redshift distortion. We test these predictions with the high resolution and high S/N data of quasars Ly$\\alpha$ absorption spectra. All results given by real data, including $\\beta$-hierarchy, high order moments and scale-scale correlation, are found to be well consistent with the log-Poisson hierarchy. We compare the log-Poisson hierarchy with the popular log-normal model of the Ly$\\alpha$ transmitted flux. The later is found to yield too strong non-Gaussianity at high orders, while the log-Poisson hierarchy is in agreement with observed data. ",
        "introduction": "Baryon matter of the universe is mainly in the form of intergalactic medium (IGM), of which the dynamics can be described as compressible fluid. Luminous objects are formed from baryon matter in the gravitational well of dark matter. Therefore, the dynamical state of baryon fluid in nonlinear regime is crucial to understand the formation and evolution of the large scale structures of the universe. In the linear regime, the baryon fluid follows the mass and velocity fields of collisionless dark matter. In the nonlinear regime, however, the baryon fluid statistically decouples from the underlying dark matter field. The statistical behavior of baryon fluid is no longer described by a similar mapping of the underlying dark matter field (e.g. Pando et al. 2004).  It was first pointed out by Shandarin and Zeldovich (1989) that the dynamical behavior of baryon matter clustering on large scales is similar to turbulence. The expansion of the universe eliminates the gravity of uniformly distributed dark matter. The motion of baryon matter on scales larger than dissipation is like that of matter moving by inertia. In this regime, the evolution of baryon matter is scale-free and dynamically like fully developed turbulence in inertial range. The turbulence of incompressible fluid leads the energy passes from large to the smallest eddies, and finally dissipates into thermal motion. While the clustering of cosmic baryon fluid is also due to the transform of density perturbations on different scales, and finally falls and dissipates into virialized halos of dark matter. Yet, the turbulence of incompressible fluid is rotational (Landau \\& Lifshitz 1987), while the clustering of cosmic matter is irrotational, because vorticities do not grow in an expanding universe (Peebles 1980).  Nevertheless, the turbulence-like behavior of cosmic baryon fluid has been gradually noticed. First, the dynamics of growth modes of the cosmic matter is found to be sketched by a stochastic force driven by Burger's equation (Gurbatov, et al 1989, Berera \\& Fang, 1994). The Burger's equation driven by the random force of the gravity of dark matter can also sketch the evolution of baryon fluid, if cooling and heating are ignored (Jones 1999; Matarrese \\& Mohayaee 2002). Later, the Burger's fluid is found to show turbulence behavior if the Reynolds number is large enough (Polyakov 1995; Lassig 2000; Bec \\& Frisch 2000; Davoudi et al. 2001). The Reynolds number of IGM at nonlinear regime actually is large. Therefore, we may expect that, in the scale free range, the dynamical state of cosmic baryon fluid should be Burger's turbulence. The turbulence of Burger's fluid is different from the turbulence of incompressible fluid. The later consists of vortices on various scales, while the former is a collection of shocks.  With the cosmological hydrodynamic simulation based on Navier-Stokes equations in which heating and cooling processes are properly accounted, it has been found that the velocity field of the IGM consists of an ensemble of shocks, and satisfies some scaling relations predicted by Burger's turbulence (Kim et al. 2005). This result reveals that the turbulence features of cosmic baryon fluid are independent of the details of dissipation (heating and cooling) mechanism if we consider only the scale free range, i.e. from the scale of the onset of nonlinear evolution to the scale of dissipation, say Jeans length.  A new progress is to show that the velocity field of cosmic baryon fluid can extremely well described by She-Leveque's (SL) scaling formula (He et al. 2006). The SL formula is considered to be the basic statistical features of the self-similar evolution of fully developed turbulence. Very recently, the non-Gaussianities of mass density field of the hydrodynamic simulation samples are found to be well consistent with the predictions of the log-Poisson hierarchy, which originates from some hidden symmetry of the Navier-Stokes equations. This hierarchical model gives a unified explanation of non-Gaussian features of baryon fluid, including the intermittence, hierarchical relation, scale-scale correlations etc (Liu \\& Fang 2008). These results strongly indicate that, in the scale free range, dynamical state of cosmic baryon fluid is similar to a fully developed turbulence.  In this paper, we investigate the log-Poisson hierarchy of cosmic baryon fluid with observed data -- the Ly$\\alpha$ transmitted flux of quasar absorption spectrum, which is due to the absorptions of quasar continuum by the diffusely distributed neutral hydrogen (Bi et al 1995, Bi \\& Devidsen 1997, Rauch 1998). These samples offer a unique way to study the non-Gaussian feature of cosmic baryon fluid. It has been known for a long time that the fields of Ly$\\alpha$ forests and transmitted flux are highly non-Gaussian. Observation samples of Ly$\\alpha$ forests and transmitted flux show scale-scale correlation (Pando et al 1998), intermittence (Jamkhedkar et al. 2000; Pando et al. 2002; Feng et al. 2003), non-thermal broadening of \\ion{H}{1} and \\ion{He}{2} Ly$\\alpha$ absorption lines (Zheng et al. 2004; Liu et al. 2006). We will show that the log-Poisson hierarchy provides a crux to understand the non-Gaussian behavior.  The outline of this paper is as follows. \\S 2 gives an introduction of the log-Poisson hierarchy. \\S 3 shows that the neutral hydrogen component of cosmic baryon fluid is of log-Poisson hierarchy. In \\S 4, we study the log-Poisson hierarchy of the field of Ly$\\alpha$ transmitted flux with observed samples of quasar absorption spectra. A comparison between log-Poisson hierarchy and log-normal model is also presented in \\S 4. The conclusion and discussion are given in \\S 5. ",
        "conclusions": "Nonlinear evolution of mass and velocity fields is a central problem of large scale structure of the universe.  The clustering of the cosmic baryon fluid, governed by the Navier-Stokes equation in gravitational field of an expanding universe, has to be self similarly hierarchical in the scale free range in which the dynamical equations and initial perturbations are scale-invariant.  The log-Poisson hierarchical clustering sketches the nonlinear evolution of cosmic baryon fluid in the scale free range. If the initial density perturbations are Gaussian, and their power spectrum is given by power law $P(k)\\propto k^{\\alpha}$, the structure functions initially have to be $S_p(r)\\propto r^{-p\\alpha/2}$. In the regime of linear evolution, the structure functions will keep to be $S_p(r)\\propto r^{\\xi_{l}(p)}$, and the intermittent exponent is $\\xi_{l}(p)=-p\\alpha/2$. According to the log-Poisson hierarchy scenario, the nonlinear evolution leads to the hierarchical transfer of the power of clustering from large scales to small scales. The structure function will become $S_p(r)\\propto r^{\\xi_{l}(p)+\\xi_{nl}(p)}$, where the nonlinear term of the intermittent exponent is $\\xi_{nl}(p)=-\\gamma[p-(1-\\beta^{p})/(1-\\beta)]$, in which the parameters $\\beta$ and $\\gamma$ are dimensionless. $\\beta$ measures the intermittency of the field, and $\\gamma$ measures the singularity of the clustering. For Gaussian field, we have $\\beta=1$, and therefore, $\\xi_{nl}(p)=0$ for all order $p$. Since the onset of nonlinear evolution, the parameter $\\beta$ will gradually decrease, and the field becomes intermittent.  With simulation samples, we found that the parameter $\\beta$ is decreasing with the decrease of redshift $z$. It means that the field is weakly intermittent at earlier time, but strong intermittent at later time (Liu \\& Fang 2008). Although $\\xi_{nl}(p)\\neq  0$, the nonlinear evolution keeps the cosmic baryon fluid to be scale-invariant. We showed that the mass density field of neutral hydrogen fluid in the scale free range is also well described by the log-Poisson hierarchy in spite of the neutral hydrogen fraction of the baryon fluid is not constant in space. This is because the UV ionization photon is assumed to be uniform, and it does not violate the scale invariance of this system. However, the number of $\\beta$ of neutral hydrogen is found to be less than that of total baryon fluid. Therefore, the neutral hydrogen is less intermitted.  The Ly$\\alpha$ transmitted flux of quasars Ly$\\alpha$ absorption spectrum is considered to be effective to probe the mass and velocity fields of neutral hydrogen. However, the redshift distortion of the velocity field leads to the deviation of the field of the Ly$\\alpha$ transmitted flux from the neutral hydrogen field. The transmitted flux does not satisfy all predictions of log-Poisson hierarchy. Fortunately, the effect of radshift distortion is approximately negligible for some log-Poisson hierarchical predicted features. Thus, we can test the log-Poisson hierarchy with quasars Ly$\\alpha$ absorption spectrum.  Using high resolution and high S/N data of quasars Ly$\\alpha$ absorption spectrum, we show that all the non-Gaussian features predicted by the log-Poisson hierarchy, including the hierarchical relation, the intermittent exponent, the ratios of different moments, and the scale-scale correlation, are consistent with observed samples. The observed samples of the transmitted flux yield the same intermittence parameter $\\beta$ as that of neutral hydrogen field produced with hydrodynamic simulation of the concordance $\\Lambda$CDM universe. Our result shows that the intermittent exponent $\\xi(p)$, or parameters $\\beta$ and $\\gamma$, is effective to discriminate among models of nonlinear evolution.  The log-normal model can well fit observed data on lower order statistics, but not good on higher orders. On the other hand, the log-Poisson model gives good fitting on lower order as well as higher order statistics. Therefore, a comparison between the log-Poisson model and log-normal model on lower order statistics will be able to find the relation between parameters of the log-Poisson and log-normal models. This relations would be useful to explain the parameters of log-Poisson model with well-known parameters in cosmology, as the parameters of log-normal model generally are known in cosmology.  Recent studies have shown that the turbulence behavior of baryon gas can be detected by the Doppler-broadened spectral lines (Sunyaev et al. 2003; Lazarian \\& Pogosyan 2006). Although these works focus on the turbulence of baryon gas in clusters, the result is still applicable, at least, for the warm-hot intergalactic medium (WHIM), which is shown to follow the evolution of Burger's fluid on large scales (He et al. 2004, 2005). The last but not least, the polarization of CMB is dependent on the density of electrons, and therefore, the map of CMB polarization would provide a direct test on the non-Gaussian features of ionized gas when the data on small scales become available."
    },
    "0809/0809.0836_arXiv.txt": {
        "abstract": "Circumbinary disks are considered to exist in a wide variety of astrophysical objects, e.g., young binary stars, protoplanetary systems, and massive binary black hole systems in active galactic nuclei (AGNs). However, there is no definite evidence for the circumbinary disk except for some in a few young binary star systems. In this Letter, we study possible oscillation modes in circumbinary disks around eccentric and circular binaries. We find that progarde, nonaxisymmetric waves are induced in the inner part of the circumbinary disk by the tidal potential of the binary. Such waves would cause variabilities in emission line profiles from circumbinary disks. Because of prograde precession of the waves, the distance between each component of the binary and the inner edge of the circumbinary disk varies with the beat period between the precession period of the wave and the binary orbital period. As a result, light curves from the circumbinary disks are also expected to vary with the same period. The current study thus provides a new method to detect circumbinary disks in various astrophysical systems. ",
        "introduction": "\\label{sec:intro}  Astrophysical disks are ubiquitous in the various systems of the universe: star-compact object systems, star-planet systems, active galactic nuclei (AGNs), and so forth. These disks surround the individual objects as a circumobject disk. If the object surrounded by a rotating disk is a binary, the disk is called a circumbinary disk (see Figure~\\ref{fig:system} for a schematic view of the circumbinary disk).  About 60$\\%$ of main sequence stars are considered to be born as { binary or multiple systems \\citep{dm91}. Numerical simulations have confirmed that young binary stars embedded in dense molecular gas have a circumbinary disks \\citep{al96a, bb97, gk1, gk2} and that a circumsteller disk is also formed around each star \\citep{bb97,gk1,gk2}. Indeed, direct imaging of the circumbinary disk was successfully achieved by interferometric observations of a few young binary systems including GG Tau \\citep{dut94} and UY Aur \\citep{duv98}.  In the early stage of planet formation, a planet orbiting a star will be embedded in a rotating disk (hereafter, circumbinary disk) surrounding them. \\cite{kd06} showed, by performing numerical simulations, that the cicrumbinary disk becomes eccentric due to the resonant interaction between the planet and the disk. Such a planet--disk interaction also causes the significant evolution of the orbital elements of the planet, which gives a possible explanation about the observed high orbital eccentricities in extrasolar planetary systems \\citep{gs03}. The direct probing of the circumbinary disk is, therefore, a key to testing this scenario.  Massive black holes are considered to co-evolve with their host galaxies \\citep{mag98,fm00,geb00}. There is inevitably an evolutionary stage as a binary black hole in the course of galaxy merger until the coalescence of two black holes \\citep{bege80,may07,kh08}. \\cite{haya07} found that if a binary black hole is surrounded by a circumbinary disk, mass is transferred from the disk to each black hole. The system then has a triple disk composed of an accretion disk around each black hole and a circumbinary disk as a mass reservoir around the binary \\citep{haya08}. There is, however, still little evidence for the circumbinary disk as well as for the binary black hole itself. Therefore, the detection of the circumbinary disk is also one of the important scientific motivations in probing massive binary black holes with parsec/subparsec separations. Quite recently, \\cite{bt08} and \\cite{do08} proposed the hypothesis that SDSSJO92712.65+294344.0 is a massive binary black hole, by interpreting the observed emission line features as those arising from the mass-transfer stream from the circumbinary disk.  There are many phenomena caused by oscillations in circumstellar/accretion disks. One of the most famous among them is superhumps. A superhump is a periodic luminosity hump on the light curves in an accreting binary system, with a slightly longer period than the orbital period of the binary. The superhump phenomenon was first discovered in the SU Ursae Majoris class of dwarf novae, which consists of a white dwarf and a late-type star with a low mass ratio, less than 0.2 \\citep{pat79,vogt80}. It is attributed to the precession of a deformed accretion disk induced by the tidal potential of the binary \\citep{osaki85}. {\\cite{lu91} showed that the deformation of the disk is due to the growth of an eccentric (i.e., $m=1$) perturbation through nonlinear coupling with the tidal potential.} Hitherto, no phenomena caused by oscillation modes have been detected in circumbinary disks.  In this Letter, we investigate the tidally induced oscillation modes in circumbinary disks. These modes can be driven by resonantly excited modes  at particular resonance radii, which are similar to the ones responsible for superhumps in dwarf novae systems. The Letter is organized as follows. In Section 2, we derive the azimuthally and temporally averaged tidal potential around an eccentric binary and discuss the possible oscillation modes and their frequencies in circumbinary disks. Section~3 is devoted to a summary and discussion.  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics*[width=86mm]{f1.eps} } \\caption{ Schematic diagram of a circumbinary disk surrounding the primary object and the secondary object which are gravitationally bound as a binary. } \\label{fig:system} \\end{figure} ",
        "conclusions": "We have studied possible oscillation modes in a circumbinary disk induced by the tidal potential of a binary system such as young binary stars, protoplanetary systems, and massive binary black holes in AGNs. We have pointed out that observationally interesting waves are those with precession frequency $\\omega_{{\\rm p},m}/m \\lesssim \\Omega(r_{\\rm in})-\\kappa(r_{\\rm in})/m$, where $r_{\\rm in}$ is the inner radius of the circumbinary disk. These waves are trapped between the inner disk radius and the ILR radius, where the pattern speed of the wave is equal to $\\Omega(r)-\\kappa(r)/m$. Among them, only $m=1$ waves have wavelengths comparable to the inner disk radius. {When the $m=1$ mode is dominant, the inner region of the circumbinary disk becomes eccentric and precesses very slowly \\citep[c.f.,][]{kd06,mm08}}. On the other hand, waves with $m \\ne 1$ have much shorter wavelengths and affect only the innermost narrow region of the disk. For example, if the $m=2$ mode is dominant, the disk inner edge is deformed to an elliptical shape, which precesses at the local Keplerian frequency.  It is important to note that there are two types of excitation mechanism for these waves. In circular binaries, nonaxisymmetric perturbations in the disk can grow through the resonant interaction with the tidal potential at particular resonance radii, a mechanism similar to that for superhumps in dwarf novae systems. For example, the growth of an eccentric perturbation is driven by the excitation of an $m=2$ wave at the 1:3 OLR radius. In addition to this mechanism, in eccentric binaries, a one-armed ($m=1$) spiral wave can also be excited through direct driving as a result of a one-armed ($m=1$) potential \\citep{al96b}.  Whatever the excitation mechanism, the deformed inner part of the circumbinary disk precesses at the frequency given by equation~(\\ref{eq:op1}) for the $m=1$ mode and equation (\\ref{eq:op2}) for the $m=2$ mode. Since the velocity field in the disk is also perturbed by the waves, the emission line profiles from the inner part of the circumbinary disk will vary with the precession period. Such a variability has a distinct feature and will easily be observed. Since the $m=1$ mode is a low-frequency, eccentric mode, the relative intensity of the violet (V) and red (R) peaks of double-peaked line profiles varies with a long period, e.g., $\\sim40P_{\\rm{orb}}$ for an equal-mass binary with $e=0.5$. Such a line-profile variability caused by $m=1$ waves has long been known as the V/R variation in Be stars, B-type stars with circumstellar decretion disks \\citep[e.g.,][]{por03}. In contrast, if the $m=2$ mode is dominant, the double-peaked profiles stay symmetric, but their peak separations (and FWHMs) will vary with a short period of $\\sim 2^{3/2} P_{\\rm{orb}}$ irrespective of binary parameters.  In addition to these line-profile variabilities, another type of variability from circumbinary disks is expected. There is the beat between the orbital frequency of the binary and the precession frequency of the circumbinary disk. The beat period is slightly longer than the orbital period for $m=1$ and $m=2$ modes, as seen from equations~(\\ref{eq:op1}), (\\ref{eq:op2}), and (\\ref{eq:beatp}). The radiation emitted from the circumbinary disk through the tidal dissipation is expected to vary periodically with the beat period, because the distance between each component of the binary and the inner edge of the circumbinary disk periodically changes. Moreover, the mass transfer rate from the circumbinary disk to each binary component {is also expected to} vary with the same period, because of angular momentum removal by the tidal torques. Therefore, the circumbinary disk perturbed by a trapped density wave will show the light-curve {modulation} at the beat period.  As for observability, the light-curve modulation caused by an $m=1$ deformation is expected to be detected more easily than those caused by the $m\\neq1$ one. This is because the $m=1$ wave can exist globally in the circumbinary disk, while the $m\\neq1$ waves exist only in a narrow region.  In this Letter, we have shown that variabilities with the beat period and/or the precession period are expected as a natural consequence of the tidal interaction between a binary and a circumbinary disk and that one can identify circumbinary disks in terms of these periodic variabilities. This provides a new method to probe circumbinary disks in various astrophysical systems. In a forthcoming paper, we will perform a more detailed analysis of the mode characteristics, including numerical simulations."
    },
    "0809/0809.2827_arXiv.txt": {
        "abstract": "{We have attributed the elements from Sr through Ag in stars of low metallicities (${\\rm [Fe/H]}\\lesssim -1.5$) to charged-particle reactions (CPR) in neutrino-driven winds, which are associated with neutron star formation in low-mass and normal supernovae (SNe) from progenitors of $\\sim 8$--$11\\,M_\\odot$ and $\\sim 12$--$25\\,M_\\odot$, respectively. Using this rule and attributing all Fe production to normal SNe, we previously developed a phenomenological two-component model, which predicts that ${\\rm [Sr/Fe]}\\geq -0.32$ for all metal-poor stars. This is in direct conflict with the high-resolution data now available, which show that there is a great shortfall of Sr relative to Fe in many stars with ${\\rm [Fe/H]}\\lesssim -3$. The same conflict also exists for the CPR elements Y and Zr. We show that the data require a stellar source leaving behind black holes and that hypernovae (HNe) from progenitors of $\\sim 25$--$50\\,M_\\odot$ are the most plausible candidates. If we expand our previous model to include three components (low-mass and normal SNe and HNe), we find that essentially all of the data are very well described by the new model. The HN yield pattern for the low-$A$ elements from Na through Zn (including Fe) is inferred from the stars deficient in Sr, Y, and Zr. We estimate that HNe contributed $\\sim 24\\%$ of the bulk solar Fe inventory while normal SNe contributed only $\\sim 9\\%$ (not the usually assumed $\\sim 33\\%$). This implies a greatly reduced role of normal SNe in the chemical evolution of the low-$A$ elements. This work was supported in part by US DOE grants DE-FG03-88ER13851 (G.J.W.) and DE-FG02-87ER40328 (Y.Z.Q.). G.J.W. acknowledges NASA's Cosmochemistry Program for research support provided through J. Nuth at the Goddard Space Flight Center. He also appreciates the generosity of the Epsilon Foundation.}  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1, 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "The high-resolution (e.g., \\cite{johnson02,honda04,aoki05,francois07,cohen08} and medium-resolution \\cite{barklem05} observations of elemental abundances in a large number of low-metallicity stars in the Galactic halo (and a single star in a dwarf galaxy \\cite{fulbright04}) now provide a data base for determining the stellar sources contributing to the interstellar/intergalactic medium (ISM/IGM) at metallicities of $-5.5<{\\rm [Fe/H]}<-1.5$. These data taken in conjunction with stellar models appear to define the massive stars active in the early epochs. This changes our views of what may be Population III (Pop III) stars and what stellar types are continuing contributors through the present epoch. A key to our understanding is the recognition that low-mass and normal core-collapse supernovae (SNe) from progenitors of $\\sim 8$--$11\\,M_\\odot$ and $\\sim 12$--$25\\,M_\\odot$, respectively, end their lives as neutron stars and that nucleosynthesis in the neutrino-driven winds of nascent neutron stars produce a large group of elements including Sr, Y,  and Zr via charged-particle reactions (CPR, e.g., \\cite{wh92}). This process is not the true ``$r$-process'' with rapid neutron capture. In considering the contributions of low-mass and normal SNe that are certainly key contributors, certain rules are now established:  1. The true ``$r$-process'' is not connected with any significant production of Fe  group nuclei (or the low-$A$ elements below Zn with mass numbers $A\\sim 23$--70, e.g., \\cite{qw02});  2. The CPR nuclei are not directly coupled to the true ``$r$-process'' elements. While some workers consider that a ``turn-on'' of high neutron densities may be achieved in some unknown way coupled to a neutrino-driven wind, the observational data do not appear to be in support of this;  3. If one assumes that only low-mass and normal SNe were the sources, then [Sr/Fe] must exceed $-0.32$ for all metal-poor stars.  The last rule appears to be well followed for ${\\rm [Fe/H]}>-3$, consistent with a two-component model of chemical evolution \\cite{qw07}. However, below ${\\rm [Fe/H]}\\sim -3$ there is a gross deficiency of Sr (and other CPR elements) relative to Fe. This is shown in Figure~\\ref{fig1}a. This observation requires that there be an early stellar source of Fe that leaves behind a black hole instead of a neutron star with neutrino-driven winds producing the CPR nuclei. It further follows that a third stellar component in addition to low-mass and normal SNe is required to account for all the abundance data.  \\begin{figure} \\vskip -0.75cm \\begin{center} \\includegraphics[angle=270,width=0.5\\textwidth]{fig1a.eps}% \\includegraphics[angle=270,width=0.5\\textwidth]{fig1b.eps} \\caption{(a) High-resolution data on $\\log\\epsilon({\\rm Sr})$ vs. [Fe/H]. The solid line is calculated from the two-component model for a well-mixed ISM/IGM. Note that there is a great deficiency of Sr for many sample stars with ${\\rm [Fe/H]}\\lesssim-3$. It is evident that a source producing Fe and no Sr is required. (b) Evolution of [Sr/Fe] with [Ba/Fe] for the data shown in (a). The curves correspond to different fractions $f_{{\\rm Fe},L}$ of Fe due to the $L$ source (normal SNe). Details for these and the other figures of this contribution can be found in \\cite{qw08}.\\label{fig1}} \\end{center} \\end{figure} ",
        "conclusions": "From the results reviewed above it appears that the whole chemical evolution in the ``juvenile epoch'' of the first Gyr after the Big Bang may be explained by the concurrent contributions of massive stars associated with low-mass and normal SNe and HNe in a standard IMF and that this same relative contribution continues into the present epoch. The efforts to seek Pop III stars that only occur in early epochs and then stop are considered by us to be invalid as were our earlier efforts to find a ``prompt inventory'' in the ISM/IGM. It follows that models for the formation of the ``first'' stars, which has been the focus of intensive studies with due consideration of the complex condensation and cooling processes at zero to low metallicities (e.g., \\cite{abel02,bromm04}), should consider the stellar populations inferred here with HNe ($\\sim 25$--$50\\,M_\\odot$) being the dominant metal source. This source is highly disruptive and certainly can disperse debris through the IGM until halos of substantial mass have formed. The apparent sudden onset of heavy ``$r$-process'' elements, which motivated our earlier search for a ``prompt inventory'', is most plausibly related to the formation of halos of sufficient mass that remain bound following both SNe and HNe \\cite{qw04}. It also follows that the earlier models of Galactic chemical evolution that aimed to provide $\\sim 1/3$ of the solar Fe inventory by normal SNe must now be subject to reinvestigation. The observational evidence for ongoing HNe in the current epoch cannot be ignored. There is further the fact that production of heavy (true) $r$-process nuclei is strongly decoupled from Fe production. The extent to which this presents a further challenge to stellar models is to be resolved."
    },
    "0809/0809.0363_arXiv.txt": {
        "abstract": "The emission from neutral hydrogen (HI) clouds in the post-reionization era ($z \\le 6$), too faint to be individually detected, is present as a diffuse background in all low frequency radio observations below $1420 \\, {\\rm MHz}$.  The angular and frequency fluctuations of this radiation ($\\sim 1 \\, {\\rm mK}$) is an important future probe of the large scale structures in the Universe. We show that such observations are a very effective probe of the background cosmological model and the perturbed Universe. In our study we focus on the possibility of determining the redshift space distortion parameter $\\beta$, coordinate distance $r_{\\nu}$, and its derivative with redshift, $r'_{\\nu}$. Using reasonable estimates for the observational uncertainties and configurations representative of the ongoing and upcoming radio interferometers, we predict parameter estimation at a precision comparable with supernova Ia observations and galaxy redshift surveys, across a wide range in redshift that is only partially accessed by other probes. Future HI observations of the post-reionization era present a new technique, complementing several existing one, to probe the expansion history and to elucidate the nature of the dark energy. ",
        "introduction": "Determining the expansion history of our Universe and parameterizing the constituents of the Universe at a high level of precision, are currently some of the most important goals in cosmology. While high-redshift ($z \\leq 2$) supernova~Ia observations (e.g. \\cite{riess,perlmutter}) and galaxy surveys ($z \\leq 1$ ) (e.g. \\cite{tegmark}) probe the local universe; and CMBR observations (e.g. \\cite{dunkley,komatsu}) probe the recombination era $(z \\sim 1000)$, the expansion history is largely unconstrained across the vast intervening redshift range.  Observations of redshifted $21 \\,{\\rm cm}$ radiation from neutral hydrogen (HI) hold the potential of probing the universe over a large redshift range ($20 \\ge z \\ge 0$): from the dark ages to to the present epoch (eg.  \\cite{BA5,furla}). Such observations can possibly be realized at several redshifts, using the currently functioning GMRT \\footnote{http://www.gmrt.ncra.tifr.res.in/}. Several new telescopes are currently being built with such observations in mind (eg. MWA \\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} \\& LOFAR \\footnote{http://www.lofar.org/}). Such observations will map out the large-scale HI distribution at high redshifts. It has recently been proposed \\cite{wlg,chang} that Baryon Acoustic Oscillations (BAO) in the redshifted $21 \\,{\\rm cm}$ signal from the post-reionization era ($z \\le 6$) is a very sensitive probe of the dark energy. The BAO is a relatively small ($\\sim 10-15$ per cent) feature that sits on the HI large-scale structure (LSS) power spectrum. In this paper we investigate the possibility of probing the expansion history in the post-reionization era using the HI LSS power spectrum without reference to the BAO. Unless otherwise stated we use the parameters $(\\Omega_{m0},\\Omega_{\\Lambda0}, \\Omega_b h^2, h, n_s,\\sigma_8)=(0.3,0.7, 0.024,0.7,1.0,1.0)$ referred to as the LCDM model in our analysis.  At redshifts $z \\le 6$, the bulk of the neutral gas is in clouds that have HI column densities in excess of $2 \\times 10^{20}\\,\\,{\\rm atoms/cm^{2}}$ \\cite{peroux,lombardi,lanzetta}. These high column density clouds are observed as damped Lyman-$\\alpha$ absorption lines seen in quasar spectra. These observations indicate that the ratio of the density $\\rho_{\\rm gas}(z)$ of neutral gas to the present critical density $\\rho_{\\rm crit}$, of the universe has a nearly constant value $\\rho_{\\rm gas}(z)/\\rho_{\\rm crit} \\sim 10^{-3}$, over a large redshift range $0 \\le z \\le 3.5$. This implies that the mean neutral fraction of the hydrogen gas is $ \\bar{x}_{\\HI}=50\\,\\,\\Omega_{\\rm gas} h^2 (0.02/\\Omega_b h^2) =2.45 \\times 10^{-2}$, which we adopt for the entire redshift range $z \\le 6$.  The redshifted $21 \\, {\\rm cm}$ radiation from the HI in this redshift range will be seen in emission. The emission from individual clouds ($ < 10 \\,\\mu{\\rm Jy}$) is too weak to be detected with existing instruments unless the image is significantly magnified by gravitational lensing \\cite{saini}. The collective emission from the undetected clouds appears as a very faint background in all radio observations at frequencies below $1420 \\, {\\rm MHz}$. The fluctuations in this background with angle and frequency is a direct probe of the HI distribution at the redshift $z$ where the radiation originated. It is possible to probe the HI power spectrum at high redshifts by quantifying the the fluctuations in this radiation (\\cite{bns,bs}). ",
        "conclusions": "\\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig1.ps,width=0.45\\textwidth,angle=0}} \\caption{Here we plot $C_l(0)$  at redshifts $z=\\{1.5,3.0,4.5\\}$. The signal decreases monotonically with increasing redshift, so the lowest plot is for the highest redshift. We assume the bias to be $b=1$ throughout.  } \\label{fig:cl} \\end{center}  \\end{figure} \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig2.ps,width=0.45\\textwidth,angle=0}} \\caption{Here we plot the frequency decorrelation function $\\kappa_{\\ell}(\\Delta \\nu)$ as a function of $\\Delta \\nu$, for a fixed redshift $z=3.0$ and $\\ell=\\{100,1000,10000\\}$. The signal declines more sharply for higher value of $\\ell$.} \\label{fig:kappa} \\end{center} \\end{figure}  The expected signal $C_l(\\Delta \\nu)$ from a few representative redshifts, calculated for the LCDM model, is plotted in Figure~\\ref{fig:cl}, and in Figure~\\ref{fig:kappa} we have plotted the frequency decorrelation function $\\kappa_{\\ell}(\\Delta \\nu)$ as a function of $\\Delta \\nu$, for a fixed redshift $z=3.0$ and for $\\ell=100,\\,1000 \\,\\& \\,10000$. The HI signal is smaller than $\\sim 1 \\, {\\rm mK}$, and it decreases with increasing $l$.  The shape or $\\ell$ dependence is decided by the shape of $P(k)$ at all comoving wave-numbers $k \\ge \\ell/r_{\\nu}$.  The signal at two different frequencies $\\nu$ and $\\nu+\\Delta \\nu$ decorrelates rapidly with increasing $\\Delta \\nu$ and $\\kappa_{\\ell}(\\Delta \\nu) < 0.1$ at $\\Delta \\nu > 5 \\, {\\rm MHz}$.  The decorrelation occurs at a smaller $\\Delta \\nu$ for the larger multipoles (Figure \\ref{fig:kappa}). While the HI signal at a frequency separation $\\Delta \\nu>5\\, {\\rm MHz}$ is expected to be uncorrelated, the foregrounds are expected to be highly correlated even at frequency separations larger than this (eg. \\cite{santos}).  This should in principle allow the HI signal to be separated from the foregrounds, which are a few orders of magnitude larger (eg. \\cite{mcquinn,mor2}).  It is clear from eq.~(\\ref{eq:fsa}) that $C_{\\ell}(\\Delta \\nu)$ depends on the background cosmological model through the parameters $(\\beta,r_{\\nu},r^{'}_{\\nu})$.  Assuming that the dark matter power spectrum $P(k)$ is known a priori, observations of $C_{\\ell}(\\Delta\\nu)$ can be used to determine the values of these three parameters. It is convenient to replace $r^{'}_{\\nu}$ with the dimensionless parameter \\cite{ali1} \\begin{equation} p(z)=\\frac{d \\ln\\left [r_{\\nu}(z) \\right]}{d \\ln(z)} \\,. \\end{equation} Figure~\\ref{fig:parm} shows the variation of the three parameters $(\\beta,r_{\\nu},p)$ across the redshift range $z\\le 6$ for the LCDM model.  \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig3.ps,width=0.45\\textwidth,angle=0}} \\end{center} \\caption{Here we plot the parameters $(\\beta,r,p)$ as a function of redshift $z$ for the concordance LCDM model. The parameter $r=r_{\\nu}/(6000\\,\\,{\\rm Mpc)}$.} \\label{fig:parm} \\end{figure}  We separately consider parameter estimation using $C_{\\ell}$ and $\\kappa_{\\ell}(\\Delta \\nu)$. The former does not depend on $p$. The amplitude $A= (\\bar{T} \\bar{x}_{\\HI} b)^2 /\\pi r_{\\nu}^2$ of $C_{\\ell}$ is uncertain, and we consider the joint estimation of three parameters $(A,\\beta,r_{\\nu})$ from observations of $C_{\\ell}$.  The value of $\\kappa_{\\ell}(\\Delta \\nu)$ is insensitive to the amplitude $A$, leaving three parameters $(\\beta,r_{\\nu},p)$ that can be jointly estimated from this.  We use the Fisher matrix (e.g. \\cite{tth}) to determine the accuracy at which these parameters can be estimated.  Parameter estimation depends on two distinct aspects of the observing instrument. The first is the $\\ell$ range {\\it ie.}  $\\ell_{min}$, $\\ell_{max}$, and the sampling interval $\\Delta \\ell$, which corresponds to the smallest $\\ell$ spacing at which we have independent estimates of $C_{\\ell}(\\Delta \\nu)$ . This is determined by the instrument's field of view, and is inversely related to it. The second is the observational uncertainty in $C_{\\ell}(\\Delta \\nu)$. This is a sum, in quadrature, of the instrumental noise and the cosmic variance.  The cosmic variance contribution $\\delta C_{\\ell}/C_{\\ell}=\\sqrt{{2}/{((2 \\ell +1) \\, {\\rm f} \\, \\Delta \\ell})}$ (${\\rm f}$ is the fraction of sky observed) is further reduced because the large frequency bandwidth $\\Delta\\nu_B$ provides several independent estimates of $C_{\\ell}$. We assume that $\\delta C_{\\ell}$ is reduced by a factor we $\\sqrt{\\Delta\\nu_B/(1 \\, {\\rm MHz})}$ because of this. The instrumental uncertainties were estimated using relations \\cite{ali08} between $\\delta C_{\\ell}$ and the noise in the individual visibilities measured in radio-interferometric observations.  For this we assume that the baselines in the radio-interferometric array have a uniform u-v coverage.   We consider three different instrumental configurations  for parameter estimation.   \\begin{itemize} \\item[A.] The currently functional GMRT has too few antennas for cosmological parameter estimation. We consider an enhanced version of the GMRT with a substantially larger number of antennas ($N=120$) , each identical to those of the existing GMRT. The antennas have a relatively small field of view ($\\theta_{\\rm FWHM}\\sim 0.8^{\\circ}$ at $610 \\, {\\rm MHz}$) and the array has relatively large baselines spanning $\\ell_{min}=500$ to $\\ell_{max}=10,000$ with $\\Delta \\ell=100$. \\item[B.] The upcoming MWA will have a large number of small sized antennas.  The antennas have a relatively large field of view ($\\theta_{\\rm FWHM}\\sim 5^{\\circ}$ at $610 \\, {\\rm MHz}$), and the array is expected to be quite compact spanning $\\ell_{min}=100$ to $\\ell_{max}=2000$ with $\\Delta \\ell=20$.  The first version of this array is expected to have $N=500$ antennas which is what we consider. \\item[C.] This is a future, upgraded version of the MWA which is expected to have $N=5000$ antennas. \\end{itemize} For each of these configurations, we assume that 16 simultaneous primary beams can be observed.  We present results for $2$ years of observation for A and B, and $1000$ hours for C.  Throughout we assume frequency channels $0.05 \\, {\\rm MHz}$ wide, a bandwidth $\\Delta \\nu_B =32 \\, {\\rm MHz}$, and that a single field is observed for the entire duration. For parameter estimation we use: $\\delta \\kappa_{\\ell}(\\Delta \\nu)=\\sqrt{2} \\, \\delta C_{\\ell}/C_{\\ell}$.   \\begin{figure*} \\begin{center} \\mbox{\\epsfig{file=fig4.ps,width=0.8\\textwidth,angle=0}} \\end{center} \\caption{Expected one-sigma fractional errors for parameter estimation at different redshifts for the LCDM model. The curves in each panel correspond, from top to bottom, to the cases A, B, and C, respectively. } \\label{fig:par} \\end{figure*}   \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig5.ps,width=0.45\\textwidth,angle=-90}} \\end{center} \\caption{Expected  one-sigma confidence regions for the parameters $\\Omega_{m0}$ and $\\Omega_{k0}$, based on estimated errors for observations of $p$, corresponding to Figure~4, at $z=3$. } \\label{fig:contour} \\end{figure}  We find that observations of $C_{\\ell}$ impose very poor constraints on the parameters $\\beta$ and $r_{\\nu}$, and we do not show these here.  The accuracy is considerable higher for $\\kappa_{\\ell}(\\Delta \\nu)$, which captures the three dimensional clustering of the HI as compared to $C_{\\ell}$, which quantifies only the angular dependence. Figure~\\ref{fig:par} shows the predicted estimates for the parameters $\\beta$, $r_{\\nu}$ and $p$ at various redshifts.  Further, we find that a compact, wide-field array (B,C) is considerably more sensitive to these parameter as compared to case A.  Considering the three parameters individually:  {\\bf Redshift-space distortion parameter: $\\beta$ }. This has traditionally been measured from galaxy redshift surveys \\cite{Peacock,Hawkins,Ross,Guzzo}, with uncertainties in the range $0.1 \\le \\Delta \\beta/\\beta \\le 0.2$. These observations have, till date, been restricted to $z \\le 1$. Future galaxy surveys are expected to achieve higher redshifts and smaller uncertainties. Galaxy surveys have the drawback that at very high redshifts they probe only the most luminous objects, which are expected to be highly biased. HI observations do not have this limitation and could provide high precision $(\\Delta \\beta/\\beta <0.1)$ estimates over a large redshift range.  {\\bf Coordinate distance, $r_\\nu$}: The most direct measurement of the coordinate distance comes from supernova type~Ia observations for $z \\le 2$.  Current Sn~Ia observations give $\\Delta r_\\nu/r_\\nu \\simeq 0.07$ \\citep{sai04} for a single supernova. The statistical error in the coordinate distance can be further reduced by observing a large number of supernovae in a small redshift bin; thus the fundamental limitation of this technique is due to unknown systematics in the supernovae themselves, since it is certainly possible that supernovae at high redshift are different. Figure~4 shows that the HI method might have the potential to enable a precise measurement of the coordinate distance up to much larger redshifts.  Furthermore, such a complimentary probe will also help in ascertaining systematics in the supernova probe.  {\\bf Derivative of coordinate distance, p}: This quantifies the Alcock-Paczynski (AP) effect \\cite{Alcock}, which is well accepted as a means to study the expansion history at high $z$, though such observations have not been possible till date.  Observations of redshifted $21\\, {\\rm cm}$ radiation hold the potential of measuring the AP effect \\cite{Nusser,ali1,Barkana}.  The parameter $p$ is not affected by the overall amplitude $A$ and the bias $b$, and is a sensitive probe of the spatial curvature (Figure \\ref{fig:parm}).  Our estimates indicate that it will be possible to measure $p$ with an accuracy $\\Delta p/p \\sim 0.03$ over a large $z$ range.   The parameters $(\\beta,r_{\\nu}, p)$ chosen for our analysis occur naturally when we interpret $C_{\\ell}(\\Delta \\nu)$ in terms of the three dimensional dark matter power spectrum $P(k)$. Further, these parameters are very general in that they do not refer to any specific model for either the dark energy or the dark matter, and are valid even in models with alternate theories of gravity (eg. \\cite{Carroll,Dvali}). In fact, observations of these three parameters at different redshifts can in principle be used to distinguish between these possibilities.  For the purpose of this paper, we illustrate the cosmological parameter estimation by considering the simplest LCDM model, with two unknown parameters $\\Omega_{m0}$ and $\\Omega_{k0}$, and $\\Omega_{\\Lambda0}=1-\\Omega_{m0}-\\Omega_{k0}$. In Figure~\\ref{fig:contour} we plot the $1\\hbox{--}\\sigma$ confidence interval for the estimation of $\\Omega_{m0}$ and $\\Omega_{k0}$, using a single measurement of $p$ alone, {\\it ie.} only one of the three parameters measured at a single redshift $z=3$.  Note that $p$ is insensitive to $H_0$ and hence it is not considered as an additional parameter here.  It is possible to combine measurements at different $z$ to improve the constraints on cosmological parameters.  We shall undertake a detailed analysis for quantifying the precision that can be achieved by combining different data sets (CMBR, galaxy surveys) for a more complicated dark energy model in a future work.  In conclusion, HI observations of the post-reionization era can, in principle, determine the expansion history at a high level of precision and thereby constrain cosmological models.  Neither the upcoming initial version of the MWA which is planned to have $ 500$ antenna elements nor any conceivable upgradation of the existing GMRT will be in a position to carry out such observations, the observation time needed being too large.  We find that an enhanced version of the MWA, which is planned to have $5000$ antenna elements, would be in a position to meaningfully constrain cosmological models.  By combining different probes, we expect to achieve an unprecedented precision in the determination of cosmological parameters. This will be a step towards pinning down the precise nature of dark energy in the universe."
    },
    "0809/0809.4776_arXiv.txt": {
        "abstract": "Two radiation mechanisms, inverse Compton scattering (ICS) and synchrotron radiation (SR), suffice within the cannonball (CB) model of long gamma ray bursts (LGRBs) and X-ray flashes (XRFs) to provide a very simple and accurate description of their observed prompt emission and afterglows. Simple as they are, the two mechanisms and the burst environment generate the rich structure of the light curves at all frequencies and times. This is demonstrated for 33 selected Swift LGRBs and XRFs, which are well sampled from early until late time and faithfully represent the entire diversity of the broad-band light curves of Swift LGRBs and XRFs. Their prompt gamma-ray and X-ray emission is dominated by ICS of `glory' light. During their fast decline phase, ICS is taken over by SR, which dominates their broad-band afterglow. The pulse shape and spectral evolution of the gamma-ray peaks and the early-time X-ray flares, and even the delayed optical `humps' in XRFs, are correctly predicted. The `canonical' and non-canonical X-ray light curves and the chromatic behaviour of the broad-band afterglows are well reproduced. In particular, in canonical X-ray light curves, the initial fast decline and rapid softening of the prompt emission, the transition to the plateau phase, the subsequent gradual steepening of the plateau to an asymptotic power-law decay, and the transition from chromatic to achromatic behaviour of the light curves agrees well with those predicted by the CB model. The Swift early-time data on XRF 060218 are inconsistent with a black-body emission from a shock break-out through a stellar envelope. Instead, they are well described by ICS of glory light by a jet breaking out from SN2006aj. ",
        "introduction": "Since the launch of the Swift satellite, precise data from its Burst Alert Telescope (BAT) and X-Ray Telescope (XRT) have been obtained on the spectral and temporal behaviour of the X-ray emission of long-duration $\\gamma$-ray bursts (LGRBs) and X-ray flashes (XRFs) from their beginning until late times.  The early data are often complemented by the ultraviolet-optical telescope (UVOT) on board Swift, and by ground-based {\\it UVO} and $NIR$ robotic and conventional telescopes. The ensemble of these data have already been used to test the most-studied theories of long duration GRBs and their afterglows (AGs), the {\\it Fireball} (FB) models (see, e.g.~Zhang \\& M\\'esz\\'aros~2004, Zhang~2007, and references therein) and the {\\it Cannonball} (CB) model [see, e.g.~Dar \\& De R\\'ujula~2004 (hereafter DD2004), Dado, Dar \\& De R\\'ujula (hereafter DDD)~2002a, 2003a, and references therein].  The Swift X-ray light curves of LGRBs roughly divide into two classes, `canonical' and non-canonical (Nousek et al.~2006, O'Brien et al.~2006, Zhang~2007).  When measured early enough, the observed X-ray emission has prompt peaks which coincide with the $\\gamma$-ray peaks of the GRB, and a rapidly declining light curve with a fast spectral softening after the last detectable peak of the GRB. This rapid decline and spectral softening of the prompt emission end within a few hundreds of seconds. In canonical LGRBs the X-ray light curve turns sharply into a much flatter `plateau' with a much harder power-law spectrum, typically lasting thousands to tens of thousands of seconds, and within a time of order one day it steepens into a power-law decay, which lasts until the X-ray AG becomes too dim to be detected (Fig.~\\ref{f1}).  The plateau phase is missing in non canonical GRBs, and the asymptotic power-law decline begins the decay of the prompt emission and lasts until the X-ray become too dim to be detected (Fig.~\\ref{f2}) without any observable break.   In an significant fraction of otherwise canonical GRBs, the rapid decay and fast spectral softening of the prompt emission changes to a slower power-law decay, $\\sim t^{-2.1}$, and a harder spectrum, before it reaches the plateau (Fig.~\\ref{f3}). We shall refer to such light curves as `semi-canonical'.    The Swift X-ray data show a flaring activity in a large fraction of GRBs, both at early and late times. The X-ray peaks during the prompt $\\gamma$- ray emission follow the pattern of the $\\gamma$-ray pulses, they must have a common origin. In many GRBs, superimposed on the early-time fast decaying X-ray light curve, there are X-ray flares, whose peak intensities also decrease with time and whose accompanying $\\gamma$-ray emission is probably below the detection sensitivity of BAT. Yet, their spectral and temporal behaviour is similar to that of the prompt X/$\\gamma$ pulses. Very often the flaring activity continues into the afterglow phase. Late-time flares appear to exhibit different temporal and spectral behaviours than early-time flares.    Neither the general trend, nor the frequently complex structure of the Swift X-ray data were predicted by (or can be easily accommodated within) the standard FB models (see, e.g.~Zhang \\& M\\'esz\\'aros 2004, Piran~2005, for reviews). Much earlier confrontations between predictions of the FB models and the observations also provided severe contradictions, such as the failure to understand the prompt spectrum on grounds of synchrotron radiation (e.g.~Ghisellini, Celotti, \\& Lazzati 2000), or the `energy crisis' in the comparison of the bolometric prompt and AG fluences (e.g.~Piran~1999, 2000). We have discussed elsewhere other problems of FB models (DD2004, Dar~2005 and references therein), including those related to `jet breaks' (e.g.~DDD2002a, Dar~2005, DDD2006), and the a-posteriori explanations of the reported detections (GRB 021206: Coburn and Boggs~2003, see however Wigger et al.~2004 and Rutledge \\& Fox~2004; GRBs 930131 and GRB 960924: Willis et al.~2005; GRB 041219A: Kalemci et al.~2007; McGlynn et al.~2007) of large $\\gamma$-ray polarization (DDD2007b, and references therein).   The Swift data have challenged the prevailing views on GRBs. Kumar et al.~(2007) concluded that the prompt $\\gamma$-ray emission cannot be produced in shocks, internal or external. Zhang, Liang \\& Zhang~(2007) found that the fast decay and rapid spectral softening ending the prompt emission cannot be explained by high latitude emission. The X-ray and optical afterglows of Swift GRBs are very chromatic at early time in contrast with the fireball model expectation. Moreover, Curran et al.~(2006) have carefully examined Swift data and found that X-ray and optical AGs have chromatic breaks which differ significantly from the jet break of the blast-wave model of AGs. Burrows and Racusin~(2007) examined the XRT light curves of the first $\\sim\\! 150$ Swift GRBs and reported that the expected jet breaks are extremely rare. In particular, Liang et al.~(2008) have analyzed the Swift X-ray data for the 179 GRBs detected between January 2005 and January 2007 and the optical AGs of 57 pre- and post-Swift GRBs. They did not find any burst satisfying all the criteria of a jet break.   In spite of the above failures, not all authors are so critical. Some posit that the Swift data require only some modifications of the standard FB models to accommodate the results (e.g.~Panaitescu et al.~2006, Dai et al.~2007, Sato et al.~2007). Others still view the situation with faith (e.g.~Covino et al.~2006, Panaitescu~2008, Dai et al.~2008, Racusin et al.~2008a).  The situation concerning the CB model is different. The model was based on the assumption that LGRBs are produced by highly relativistic jets of plasmoids of ordinary matter (Shaviv \\& Dar~1995) ejected in core-collapse supernova (SN) explosions akin to SN1998bw (Dar \\& Plaga~1999, Dar \\& De R\\'ujula 2000).  It successfully described the broad-band AGs observed before the Swift era (e.g.~DDD2002a, DDD2003a) and exposed the consistent photometric evidence for a LGRB/SN association in all nearby GRBs (DDD2002a, DD2004 and references therein) long before GRB 030329. In the case of GRB 030329 the first $\\sim\\!6$ days of AG data were described by the CB model precisely enough to extrapolate them to predict even the date in which its associated SN would be bright enough to be detected spectroscopically (DDD2003c). General acceptance of the GRB-SN association waited until the spectroscopic discovery of SN2003dh, coincident with GRB 030329 (Hjorth et al.~2003, Stanek et al.~2003), and other spectroscopically-proven associations, e.g.~GRB030213/SN2003lw (Malesani et al.~2004), GRB021211/SN2002lt (Della Valle et al.~2003), XRF060218/SN2006aj (Campana et al.~2006b, Pian et al.~2006, Mazzali et al.~2006) and XRF080109/SN2008D (Malesani et al.~2008, Modjaz et al.~2008, Soderberg et al.~2008).   The CB model (DD2004) has been applied successfully to explain all the main observed properties of long GRBs and XRFs before the Swift era (e.g.~Dar 2005 and references therein). The model is summarized in $\\S$\\ref{CBMODEL}. For detailed accounts see, e.g.,~De R\\'ujula, 2007a,b.  In this report we extend and refine our analysis of the temporal and spectral behaviour of the $\\gamma$-ray, X-ray and optical light curves of GRBs during the prompt emission, the rapid-decay phase, and the afterglow phase. The observed prompt spectrum in the $\\gamma$-ray to X-ray domain is the predicted one, which is Compton-dominated in the CB model  (DD2004). The observed widths of the $\\gamma$-ray and X-ray peaks, as well as lag-times between them and their relative fluences, are in accordance with the model's predictions, if free-free absorption dominates the transparency of the CBs to eV photons in the CBs' rest frame. We investigate whether or not the CB model can describe all the data in terms of only two emission mechanisms: inverse Compton scattering and synchrotron radiation. We shall see that this simple picture, explicitly based on the predictions in DDD2002a and DD2004, gives a straightforward and successful description of the Swift GRB data, at all observed energies and times.  An exploding SN illuminates the progenitor's earlier ejecta, creating a {\\it glory} of scattered and re-emitted light. In the CB model inverse Compton scattering (ICS) of glory photons is the origin of the prompt $\\gamma$/X-ray peaks, as we review in $\\S$\\ref{Inverse}. Each peak is generated by a single CB emitted by the `engine', the accreting compact object resulting from a core-collapse supernova event. We shall see that ICS correctly describes the prompt peaks, extending even into the optical domain in XRFs in which the relevant observations are available, such as XRF 060218. The natural explanation of the early time flares is the same as that of the stronger flares: ICS of glory photons by the electrons of CBs ejected in late accretion episodes of fall-back matter on the newly formed central object. These CB emissions must correspond to a weakening activity of the engine, as the accreting material becomes scarcer.  In the CB model, from the onset of the `plateau' onwards, the X-ray, optical (DDD2002a) and radio (DDD2003a) afterglows are dominated by synchrotron radiation (SR), the CB-model predictions for which are reviewed in $\\S$\\ref{Synchrotron}. On occasion these AGs also have transient rebrightenings (`very late' flares), two notable cases before the Swift era being GRB 970508 (Amati et al.~1999, Galama et al.~1998a) and GRB 030329 (Lipkin et al.~2004). During these episodes, the spectrum continues to coincide with the one predicted on the basis of the synchrotron mechanism that dominates the late AGs. These very late flares are well described by encounters of CBs with density inhomogeneities in the interstellar medium (DDD2002a, DD2004). Very late flares in the {\\it XUVONIR} AG may have this origin as well.  In this article we compare the predictions of the CB model and the observed X-ray and optical light curves of 33 selected GRBs, which are well sampled from very early time until late time, have a relatively long follow-up with good statistics and represent well the entire diversity of Swift GRBs. These include the brightest of the Swift GRBs (080319B), the GRB with the longest measured X-ray emission (060729), a few with canonical X-ray light curves (050315, 060526, 061121 and 080320) with and without superimposed X-ray flares, GRBs with semi-canonical light curves (060211A, 061110A, 070220, 080303, 080307, 051021B) and non-canonical light curves (061007, 061126, 060206), and some of the allegedly most peculiar GRBs (050319, 050820A, 060418, 060607A, 071010A, 061126). We also compare the CB model prediction and the observed X-ray light curve of additional 12 GRBs with the most rapid late-time temporal decay.  In the CB model, LGRBs and XRFs are one and the same, the general distinction being that XRFs are viewed at a larger angle relative to the direction of the approaching jet of CBs or have a relatively small Lorentz factor (DD2004, Dado et al.~2004c). Thus we include a Swift XRF of particular interest in our analysis: 060218. Its X-ray light curve is shown to be the normal X-ray light curve of a GRB viewed far off axis, and not the emission from the break-out of a spherical shock wave through the stellar envelope. Its optical AG at various frequencies shows, before the SN becomes dominant, a series of broad peaks between 30 ks and 60 ks after trigger, which we interpreted as the optical counterparts of the dominant prompt X-ray peak of this XRF. The expressions for an ICS-generated peak at all frequencies allow us to predict the positions, magnitudes and pulse shape of these broad peaks, a gigantic extrapolation in time, radiated energy and frequency.  After submitting for publication a first version of a comparison between the CB-model predictions and Swift observations (DDD2007c), we have compared many more Swift data with the CB-model predictions, in order to further test its ability to predict correctly all the main properties of GRB light curves. These included the rapid spectral evolution observed during the fast decay of the prompt emission in `canonical' GRBs (DDD2008a) and the `missing AG breaks' in the AG of several GRBs (DDD2008b). We have also extended the CB model to describe short hard bursts (SHBs) and confronted it with the entire data on all SHBs with well-measured X-ray and/or optical afterglows (Dado \\& Dar 2008).  Together with the GRBs discussed in this report, we have analyzed and published CB model fits to the light curves of more than 100 LGRBs and SHBs.  The CB model continued to be completely successful in the confrontation of its predictions with the data. ",
        "conclusions": "\\label{outlook}  The rich data on GRBs gathered after the launch of Swift, as interpreted in the CB model and as we have discussed here and in recent papers (e.g.~Dado et al.~2006, 2007, 2008a, 2008b) has taught us several things:  \\begin{itemize} \\item{} Two radiation mechanisms, inverse Compton scattering and synchrotron radiation, suffice within the CB model to provide a very simple and accurate description of long-duration GRBs and XRFs and their afterglows. Simple as they are, these two mechanisms and the bursts' environments generate the rich structure and variety of the light curves at all frequencies and times.  \\item{} The historical distinction between prompt and afterglow phases is replaced by a physical distinction: the relative dominance of the Compton or synchrotron mechanisms at different, frequency-dependent times.  \\item{} The relatively narrow pulses of the $\\gamma$-ray signal, the somewhat wider prompt flares of X-rays, and the much wider humps sometimes seen at $UVOIR$ frequencies in XRFs, have a common origin. They are generated by inverse Compton scattering.  \\item{} The synchrotron radiation component dominates the prompt  optical emission in ordinary GRBs, the broad-band afterglow in GRBs and XRFs and the late-time flares of both types of events.  \\item{}  The early-time XRT and UVOT data on XRF 060218 are inconsistent with a black-body emission from a shock break-out through the stellar envelope. Instead, they support the CB-model interpretation of ICS of glory light by an early jet of CBs from what is later seen as SN2006aj. The start time of the X-ray emission does not constrain the exact time of the core's collapse before the launch of the CBs, nor the possible ejection of other CBs farther off axis.  \\item{} Despite its simplicity and approximate nature, the CB model continues to provide an extremely successful description of long GRBs and XRFs. Its testable predictions, so far, are in complete agreement with the main established properties of their prompt emission and of the afterglow at all times and frequencies.  \\end{itemize}   We re-emphasize that the results presented in this paper are  based on direct applications of our previously published explicit predictions. Our master formulae, Eq.~(\\ref{ICSlc}) for ICS and Eqs.~(\\ref{decel}, \\ref{Fnu}, \\ref{SRP}) for the synchrotron component describe all the data very well. But, could they just be very lucky guesses? The general properties of the data are predictions. But, when fitting cases with many flares, are we not `over-parametrizing' the results? Finally, the $E\\,t^2$ law plays an important role. Could it also be trivially derived in a different theory?     When their collimated radiation points to the observer, GRBs are the brightest sources in the sky. In the context of the CB model and of the simplicity of its underlying physics, GRBs are not persistent mysteries, nor `the biggest of explosions after the Big Bang', nor a constant source of surprises, exceptions and new requirements. Instead, they are well-understood and can be used as cosmological tools, to study the history of the intergalactic medium and of star formation up to large redshifts, and to locate SN explosions at a very early stage. As interpreted in the CB model, GRBs are not `standard candles', their use in `Hubble-like' analises would require further elaboration. The GRB conundra have been reduced to just one: `how does a SN manage to sprout mighty jets?' The increasingly well-studied ejecta of quasars and microquasars, no doubt also fired in catastrophic accretion episodes on compact central objects, provides observational hints with which, so far, theory and simulations cannot compete.   The CB model underlies a unified theory of high energy astrophysical phenomena. The information gathered in our study of GRBs can be used to understand, also in very simple terms, other phenomena. The most notable is (non-solar) cosmic rays. We allege (Dar et al.~1992, Dar \\& Plaga~1999) that they are simply the charged ISM particles scattered by CBs, in complete analogy with the ICS of light by the same CBs. This results in a successful description of the spectra of all primary cosmic-ray nuclei and electrons at all observed energies (Dar and De R\\'ujula~2006a). The CB model also predicts very simply the spectrum of the gamma background radiation and explains its directional properties (Dar \\& De R\\'ujula~2001a, 2006b). Other phenomena understood in simple terms include the properties of cooling core clusters (Colafrancesco, Dar \\& De R\\'ujula~2003) and of intergalactic magnetic fields (Dar \\& De R\\'ujula~2005). The model may even have a say in `astrobiology' (Dar, Laor \\& Shaviv~1998, Dar \\& De R\\'ujula~2001b).     {\\bf Acknowledgment:} A.D. would like to thank the Theory Division of CERN for its hospitality during this work. We would also like to thank S.~Campana, G.~Cusumano, P.~Ferrero, D.~Grupe, D.~A.~Kann, K.~L.~Page and S.~Vaughan, for making available to us the tabulated data of their published X-ray and optical light curves of Swift GRBs and an anonymous referee for an exceptionally constructive and useful report.  \\newpage"
    },
    "0809/0809.4295_arXiv.txt": {
        "abstract": "We report on the discovery of \\hatcurb{}, the lowest mass (\\hatcurPPm\\,\\mjup) transiting extrasolar planet (TEP) discovered to date by transit searches. \\hatcurb\\ orbits the moderately bright V=\\hatcurCCtassmv\\ K dwarf \\hatcurCCgsc, with a period $P=\\hatcurLCP\\,d$, transit epoch $T_c = \\hatcurLCT$ (BJD) and duration \\hatcurLCdur\\,d. \\hatcurb\\ has a radius of \\hatcurPPr\\,\\rjup\\ yielding a mean density of \\hatcurPPrho\\,\\gcmc. Comparing these observations with recent theoretical models we find that \\hatcur\\ is consistent with a $\\sim 4.5\\,{\\rm Gyr}$, coreless, pure hydrogen and helium gas giant planet. With an equilibrium temperature of $T_{eq} = \\hatcurPPteff\\,K$, \\hatcurb\\ is one of the coldest TEPs. Curiously, its Safronov number $\\theta=\\hatcurPPtheta$ falls close to the dividing line between the two suggested TEP populations. ",
        "introduction": "\\label{sec:introduction}  It has become clear in recent years that transiting extrasolar planets (TEPs), especially those around bright stars, are extremely valuable for understanding the physical properties of planetary bodies. The transit itself is a periodic event, which --- together with high precision spectroscopic observations and radial velocity (RV) follow-up -- reveals a number of key parameters, notably the relative radius of the planet with respect to the star, and the true mass of the planet without the inclination ambiguity. These allow determination of the mean density of the planet, and an insight into its basic structural properties. These advantages have been realized early on, and the recent rise in the detection of TEPs is due to a number of dedicated transit searches, such as TrES \\citep{brown:00,dunham:04}, XO \\citep{pmcc:05}, SuperWASP \\citep{pollacco:06}, OGLE \\citep[targeting fainter stars;][]{udalski:08}, and HATNet \\citep{bakos:02,bakos:04}. At the time of this writing, the number of published TEPs with a unique identification is $\\sim 40$, with $\\sim 35$ of these due to systematic searches. The properties of known TEPs already span a wide range, from the hot Neptune GJ436b with a mass of $\\mpl = 0.072\\,\\mjup$ \\citep{butler:04,gillon:07} to XO-3 with $\\mpl = 11.79\\mpl$ \\citep{johns-krull:08}, from short period orbits like OGLE-TR-56b with $P=1.2$\\,days \\citep{udalski:02,konacki:03} to $P=21.2$\\,days of \\hd{17506}b \\citep{barbieri:07}. Although most of these planets have circular orbits, some planets with significant eccentricities, such as HAT-P-2b \\citep{bakos:07a}, have also been reported. TEPs have been discovered in a wide range of environments, from orbiting M dwarfs (GJ436b) up to mid F-dwarfs, such as HAT-P-7b \\citep{pal:08a}.  Theoretical investigations have been thriving during this vigorous discovery era, some focusing on the radius of these planets \\citep[][]{burrows:07,liu:08,chabrier:04,fortney:07}, and others on the atmospheres \\citep[e.g.][]{burrows:06,fortney:07}, to mention two of the key observable properties of transiting planets. When confronting theory with observations, it is also essential to use accurate observational values, along with proper error estimates. The recent compilation of TEP parameters by \\citet{torres:08} represents a step forward in this sense. It was also noted throughout these works that a much larger sample is required for better understanding of the underlying physics, i.e.~more planets are needed to populate the mass--radius (or other) parameter space, to improve the statistical significance of correlations between planetary and stellar parameters, and to reveal any previously undetected correlations that may shed light on the physical processes governing the formation and evolution of TEPs.  The HATNet survey has been a major contributor to TEP discoveries. Operational since 2003, it has covered approximately 7\\% of the Northern sky, searching for TEPs around bright stars ($8\\lesssim I \\lesssim 12$\\,mag). HATNet operates six wide field instruments: four at the Fred Lawrence Whipple Observatory (FLWO) in Arizona, and two on the roof of the Submillimeter Array Hangar (SMA) of SAO\\@. Since 2006, HATNet has announced and published 9 TEPs. In this work we report on the tenth such discovery.\\footnote{After submission of this paper, it was realized that HAT-P-10b and WASP-11b refer to the same object, independently discovered, with WASP-11b submitted 7 days earlier to A\\&A. The two discovery groups agreed on calling it in future papers as WASP-11b/HAT-P-10b, with separate entries on www.exoplanet.eu}. ",
        "conclusions": "\\label{sec:discussion}   It is interesting to compare the properties of \\hatcurb\\ with the other known TEPs so as to place it in a broader context. This planet falls at the low-mass end of the current distribution, as shown in \\figr{exomrsaf} (left panel), where we overplot \\citet{baraffe:03} planetary isochrones, which indicate that the radius of \\hatcur\\ is broadly consistent with these models. As we noted earlier, \\hatcur\\ is formally the smallest mass TEP discovered by transit searches. The even smaller \\hd{149026b} \\citep{sato:05} and GJ436b \\citep{butler:04} were discovered by RV searches, and their transits were found later.  We compared our mass and radius to theoretical estimates of \\cite{liu:08} for a 0.5\\,\\mjup\\ body at various orbital distances from a G2V star. We note that the equivalent semi-major axis (with the same incident flux) of \\hatcurb\\ around a solar type star is $a_{rel} = 0.076$. It is also noteworthy that when a detailed comparison is done, the effects of the environment on the planetary properties are not as simple as scaling the integrated stellar flux, since the detailed spectrum of the star (e.g.~UV flux) may also be important.  Based on the models presented by \\cite{liu:08}, for $\\dot{E_{h}}/\\dot{E}_{ins} = 10^{-6}$ the equilibrium radius of \\hatcurb\\ would be $\\sim$1.1\\,\\rjup, where $\\dot{E_{h}}$ is the energy per unit time due to orbital tidal heating or similar internal heating, and $\\dot{E}_{ins}$ is the energy received via insolation. For larger values of $\\dot{E_{h}}/\\dot{E}_{ins}$ the expected radius is larger, and for smaller values it asympotically converges to 1.1\\,\\rjup.  This makes us conclude that a small core of approx.~20\\mearth\\ is required so that the model values match the observed \\hatcurPPrshort\\,\\rjup\\ radius of \\hatcurb. This would be also consistent with the core-mass---stellar metallicity relation proposed by \\cite{burrows:07}  When comparing with models of \\cite{fortney:08}, we obtain similar results. The current mass, radius and insolation of \\hatcurb\\ are consistent with a 500\\,Myr model with a 25\\,\\mearth\\ core mass, or a 4.5\\,Gyr coreless pure hydrogen and helium model. It is noted that low-mass, core-free planets are hard to model, thus our current finding will hopefully provide a further constraint for theoretical models.  The radiation that \\hatcurb\\ receives from its host star is $\\sim 2.56\\cdot10^{8}\\ergs$. With the definitions of \\cite{fortney:08}, \\hatcurb\\ belongs to the pL class of planets. There is only one transiting planet that has a lower mean incident flux: \\hd{17156}b \\citep{barbieri:07}, but this planet orbits on a highly eccentric orbit, with incident flux increasing to over $10^9\\ergs$ at periastron.  The other planet with a similarly low incident flux is OGLE-TR-111b \\citep{udalski:02,pont:04}, orbiting an $I=15.55$\\,mag star with $\\teffstar = 5040\\,K$ (Santos, 2006). \\hatcurb\\ appears to be a near-by analog of OGLE-TR-111b in many respects, since their parameters are very similar (parentheses show those of OGLE-TR-111); the period is \\hatcurLCPshort\\,d (4.01\\,d), the stellar mass is \\hatcurYYmshort\\,\\msun\\ (0.85\\,\\msun), the stellar radius is \\hatcurYYrshort (0.83\\,\\rsun), the luminosity is \\hatcurYYlumshort\\ (0.4), the metallicity is \\hatcurSMEzfeh\\ ($0.19\\pm0.07$), and the planetary radius is \\hatcurPPrshort\\,\\mjup\\ (1.05). Interestingly, even the impact parameter of their transits is similar. There is a slight difference in their masses, with \\hatcurb\\ being smaller (\\hatcurPPm\\,\\mjup\\ vs.~$0.55\\pm0.1\\mjup$). One crucial difference between the two systems is that \\hatcur\\ is 10 times closer to us, being at \\hatcurXdist\\,pc vs.~1500\\,pc for OGLE-TR-111, and is more than 4 magnitudes brighter, thus enabling more detailed follow-up in the near future.  Another interesting observational fact is that the $\\Theta = \\hatcurPPtheta$ Safronov number of \\hatcurb\\ falls fairly close to the dividing line between the proposed Class I and Class II planets \\citep{hansen:07}. At the low end of the equilibrium temperature range of the plot (excluding GJ436b), \\hatcurb\\ seems to be at a point where the two distributions overlap (see \\figr{exomrsaf}, right panel). Finally, we note that \\hatcurb\\ strengthens the orbital period vs.~surface gravity relation \\citep{southworth07}, falling almost exactly on the linear fit between these two quantities \\citep{torres:08}.   \\clearpage"
    },
    "0809/0809.4589_arXiv.txt": {
        "abstract": "We report new spectroscopic and photometric observations of the parent stars of the recently discovered transiting planets \\mbox{TrES-3} and \\mbox{TrES-4}.  A detailed abundance analysis based on high-resolution spectra yields [Fe/H] $= -0.19\\pm 0.08$, $T_\\mathrm{eff} = 5650\\pm 75$~K, and $\\log g = 4.4\\pm 0.1$ for \\mbox{TrES-3}, and [Fe/H] $= +0.14\\pm 0.09$, $T_\\mathrm{eff} = 6200\\pm 75$~K, and $\\log g = 4.0\\pm 0.1$ for \\mbox{TrES-4}. The accuracy of the effective temperatures is supported by a number of independent consistency checks.  The spectroscopic orbital solution for \\mbox{TrES-3} is improved with our new radial-velocity measurements of that system, as are the light-curve parameters for both systems based on newly acquired photometry for \\mbox{TrES-3} and a reanalysis of existing photometry for \\mbox{TrES-4}. We have redetermined the stellar parameters taking advantage of the strong constraint provided by the light curves in the form of the normalized separation $a/R_\\star$ (related to the stellar density) in conjunction with our new temperatures and metallicities. The masses and radii we derive are $M_\\star = 0.928_{-0.048}^{+0.028}~M_{\\sun}$, $R_\\star = 0.829_{-0.022}^{+0.015}~R_{\\sun}$, and $M_\\star = 1.404_{-0.134}^{+0.066}~M_{\\sun}$, $R_\\star = 1.846_{-0.087}^{+0.096}~R_{\\sun}$ for \\mbox{TrES-3} and \\mbox{TrES-4}, respectively. With these revised stellar parameters we obtain improved values for the planetary masses and radii. We find $M_p = 1.910_{-0.080}^{+0.075}~M_\\mathrm{Jup}$, $R_p = 1.336_{-0.036}^{+0.031}~R_\\mathrm{Jup}$ for \\mbox{TrES-3}, and $M_p = 0.925 \\pm 0.082~M_\\mathrm{Jup}$, $R_p = 1.783_{-0.086}^{+0.093}~R_\\mathrm{Jup}$ for \\mbox{TrES-4}. We confirm \\mbox{TrES-4} as the planet with the largest radius among the currently known transiting hot Jupiters. ",
        "introduction": "\\label{sec:introduction}  The 40 transiting planet systems confirmed as of August 2008\\footnote{For a complete listing see {\\tt http://www.inscience.ch/transits/}, or {\\tt http://exoplanet.eu}~.} show a remarkable diversity of properties, which is indicative of the complexity of planet formation and evolution processes. Many different follow-up studies enabled by the special orientation of these systems \\citep[e.g.,][see Charbonneau et al.\\ 2007 for a review]{queloz00, charbon02, charbonneau05, knutson07, tinetti07} have brought about rapid improvements in evolutionary models of planet interiors and atmospheres \\citep{baraffe04, lecavelier06, guillot06, burrows07}. The increasing predictive power of these models is beginning to drive even more challenging observations. On the other hand, the precision and accuracy with which the most basic planet properties such as the mass and radius can be determined is currently limited by our knowledge of the properties of the host stars.  Significant uncertainties remain in the stellar mass and radius determinations of many systems. In some cases this is due to poorly determined photospheric properties (mainly temperature and metallicity), and in others to a lack of an accurate luminosity estimate. Additionally, the variety of methodologies used for these determinations and the different approaches towards systematic errors have resulted in a rather inhomogeneous set of planet properties, as discussed by \\cite{torres08}. This complicates the interpretation of patterns and correlations that are being proposed \\citep[e.g.,][]{Mazeh:05, guillot06, Hansen:07}. Recent improvements in the analysis techniques have the potential to increase the accuracy of the stellar and planetary parameters significantly, especially for the large majority without a direct distance estimate. In particular, the application of the constraint on the stellar density that comes directly from the transit light curves has been shown to be superior to the use of other indicators of luminosity such as the surface gravity ($\\log g$) determined spectroscopically \\citep{sozzetti07, holman07}. \\cite{torres08} have recently reanalyzed a large subset of the known transiting planets, incorporating these improvements and applying a uniform methodology to all systems.  In the present work we focus on two of the recently discovered transiting systems, \\mbox{TrES-3} \\citep{ODonovan:07} and \\mbox{TrES-4} \\citep{Mandushev:07}, which lack accurate estimates for the photospheric properties of the parent stars and as a result have more uncertain stellar and planetary parameters.  To improve upon these properties, we present new radial-velocity and photometric observations of \\mbox{TrES-3} with which we refine both the light curve solution and the spectroscopic orbit. We also carry out a reanalysis of existing \\mbox{TrES-4} photometry utilizing a technique which treats stellar limb darkening using adjustable parameters. We perform the first detailed spectroscopic determination of the photospheric properties of both stars, and we make use of the constraint on the stellar density mentioned above to infer more accurate values for the stellar and planetary masses and radii. Our paper is organized as follows.  In \\S\\,\\ref{sec:obs} we summarize the observations.  In \\S\\,\\ref{sec:atmpar} we present our effective temperature and metallicity determinations, along with several consistency checks aimed at establishing the accuracy of the temperatures.  In \\S\\,\\ref{sec:orb} we report an updated spectroscopic orbital solution for \\mbox{TrES-3}, and the light curve solutions for both systems are discussed in \\S\\,\\ref{lcanalysis}.  Section\\,\\ref{sec:physics} then describes our determination of the stellar masses and radii, which in turn lead to refined values for the planetary parameters over those reported in the discovery papers. We conclude in \\S\\,\\ref{sec:summ} by providing a summary of our results and by discussing whether the properties of the host stars give any useful clues on the origin of the strongly contrasting densities of their close-in gas giant planets, particularly in comparison with the other known transiting planet systems. ",
        "conclusions": "\\label{sec:summ}  Our detailed spectroscopic analyses of \\mbox{TrES-3} and \\mbox{TrES-4} have yielded accurate values of the atmospheric properties ($T_\\mathrm{eff}$ and [Fe/H]), which are critical for establishing the fundamental properties of the hosts. The accuracy of the temperatures is supported by a number of independent checks (low-resolution spectroscopy, line-depth ratios, H$_\\alpha$ line profiles, color-temperature calibrations) that gives us confidence that the inferred stellar properties are reliable. We find that \\mbox{TrES-3} is a main-sequence G dwarf with a metallicity about 1.5 times lower than the Sun's, not a very common occurence among exoplanets hosts. \\mbox{TrES-4} is a somewhat evolved late F star that is nearing the end of its main-sequence phase, and is slightly enhanced in its iron content with respect to the solar abundance.  The agreement we find between age indicators for \\mbox{TrES-3} and \\mbox{TrES-4} based on measurements of the \\ion{Ca}{2} activity levels, the lithium abundances, and rotation, and the evolutionary age inferred from the models is fair, although, as discussed in \\S\\,\\ref{sec:age}, the reliability of empirical age estimates for stars that are not young ($t\\gtrsim 1$ Gyr) is somewhat questionable. The model estimates themselves are not without their problems. Nevertheless, we conclude the stars are 1--3 Gyr old. Neither star stands out as peculiar when compared with other planet hosts with similar physical properties. In \\mbox{TrES-3} the small $v\\sin i$ value and the fact that we can only place an upper limit on the Li abundance are consistent with the notion that planet hosts with $T_\\mathrm{eff}$ similar to the Sun appear to rotate more slowly and are more Li-depleted than stars without detected planets. As pointed out recently \\citep[][and references therein]{gonzalez08}, this evidence suggests that a planet-forming disk may induce additional rotational braking, leading to enhanced mixing in the stellar envelope, which in turn accelerates the destruction of lithium. \\mbox{TrES-4}, on the other hand, does not seem consistent with the claim by \\cite{gonzalez08} that hotter planet hosts with $T_\\mathrm{eff} \\gtrsim 6100$~K have higher Li abundances, possibly due to self-enrichment processes. It is worth keeping in mind, however, that these discrepancies may not be significant given the large spread in the Li abundance for field stars with the temperature and mass of \\mbox{TrES-4} \\citep[e.g.,][]{lambred04}. New investigations on these issues are clearly needed based on uniform analyses of large samples of planet hosts and statistically significant, well-defined control samples of stars without detected planets.  New radial-velocity measurements for \\mbox{TrES-3} presented here have enabled us to revise the spectroscopic orbit for that system. We detect no indication of any longer-term variations in the radial velocities that might suggest the presence of another body in the system.  However, the small number of observations and their limited time span of only 6 months emphasize the need for continued Doppler monitoring of this and other transiting planet systems to investigate the possibility of additional companions.  In stark contrast to the considerable ground-based and space-based efforts invested in studying in great detail the atmospheric properties of many of these objects, which have undoubtedly led to tremendous insights into their structure, formation, and evolution, the amount of radial-velocity data available for transiting planets is meager, and often does not go beyond the handful of observations published in the discovery papers. Interest in the radial velocities seems to be quickly lost. It should be pointed out that the frequency of close-in giant planets ($P < 10$ day) with additional massive planets in outer orbits \\citep[up to the detection limit of today's Doppler surveys, $\\sim 4$ AU; see, e.g.,][and references therein]{butler06} is about 12\\% (8 out of 70 systems discovered via RV methods), and it would therefore be wise to continue the velocity monitoring of some of the transiting systems. If additional planets in a transiting system were detected, they might also be found to undergo transits, and such a discovery would allow us to constrain structural models for gas giants akin to Jupiter or Saturn, and open exciting opportunities for additional investigations with present ({\\em Spitzer}) and upcoming (JWST) space-borne observatories.  Accurate stellar properties for \\mbox{TrES-3} and \\mbox{TrES-4} have been derived here following the approach described in \\cite{sozzetti07} and \\cite{holman07}, comparing the spectroscopically determined $T_\\mathrm{eff}$ and [Fe/H] values along with the photometrically measured $a/R_\\star$ with current stellar evolution models. These properties have in turn allowed us to refine the determination of the mass and radius of the planets. In particular, we confirm that \\mbox{TrES-4} is the planet with the largest radius among the currently known transiting hot Jupiters.  Recently, Winn et al. (2008a) derived upper limits on the albedo of TrES-3 based on the nondetection of occultations observed at optical wavelengths. Our findings are relevant to this study in two ways.  Firstly, they strengthen the case for a circular orbit, which is important because if the orbit is eccentric then it would be possible that Winn et al. (2008a) did not observe TrES-3 at the actual times of occultations and that their data place no constraint on the albedo.  Secondly, our revised light-curve parameters are relevant because the upper limit on the geometric albedo ($p_\\lambda$) was inferred from the measured upper limit on the planet-to-star flux ratio ($\\epsilon_\\lambda$) according to  \\begin{equation} p_\\lambda = \\epsilon_\\lambda(a/R_p)^2 \\end{equation}  Our revised value of $(a/R_p)$ therefore leads to revised upper limits on the geometric albedo of TrES-3. However, this revision turns out to be minor: the new value of $(a/R_p)$ is only 1.5\\% larger than the value used by Winn et al.~(2008a). The upper limits on the geometric albedo become weaker by about 3\\%. At 99\\% confidence, the revised upper limits are 0.31, 0.64, and 1.10 in $i-$, $z-$, and $R-$band, respectively.  There is a large spread in the observed radii and densities for transiting planets of comparable mass placed at similar orbital distance from stars of very similar properties ($T_\\mathrm{eff}$, $\\log g$, [Fe/H], and age). For example, if we consider \\mbox{TrES-4} along with 5 other transiting planet hosts (excluding \\mbox{HAT-P-2}) with similar characteristics (HD~149026, HD~209458, OGLE-TR-56, OGLE-TR-132, and WASP-1), the nominal masses of the attending planets vary by a factor of $\\sim3.5$, but reported densities vary by a factor of $\\sim7.5$ \\citep[e.g.,][]{torres08}.  Many theoretical mechanisms have been proposed to inflate the radius of a strongly irradiated planet, such as additional sources of internal heating due to stellar insolation \\citep{guillot02}, tidal heating due to non-zero eccentricity caused by gravitational interaction with an outer companion \\citep{bodenheimer01} or by rotational obliquity \\citep{winn05}, elevated interior opacity due to enhanced atmospheric metallicity \\citep{guillot06, burrows07}, or varying core masses \\citep{fortney07}. None of these appear capable of explaining the observed spread in density in a natural way \\citep[see also][]{Fabrycky:07}.  Of the 6 stars just mentioned, all are more metal-rich than the Sun, except for HD~209458, which has [Fe/H] close to solar. Some of the above models \\citep{guillot06, burrows07} predict a positive correlation between the inferred planetary core mass and the host star's metallicity, as in the framework of the core-accretion model of giant planet formation \\citep[e.g.,][]{pollack96, alibert05}. This idea assumes [Fe/H] closely tracks the metallicity of the protoplanetary disk, so that more metal-rich stars should be orbited by more metal-rich planets, with a larger heavy-element content. However, among these 6 systems only one planet (HD~149026b) has an inferred core mass significantly larger than zero, and four of the other planets have measured radii so large that the cores are likely to be insignificant. In fact, even in the absence of a core the observed radii cannot be reproduced by the models, with \\mbox{TrES-4} being the extreme case. Interestingly, over 40\\% of the transiting planets reported in Table~5 of \\cite{torres08} do not appear to require any core at all to explain their sizes \\citep[according to the models of][]{fortney07}.  Taking all this into consideration, the claimed evidence for a core mass--metallicity correlation could indeed be seen as supporting the more widely accepted scenario of formation by core accretion, but from the indications above it may also be that a significant fraction of these objects formed in a different way \\citep[e.g.,][and references therein]{durisen07}.  Given the evidence collected so far, we suggest that simply connecting the host star's characteristics to the structural properties of transiting planets may in fact be an over-simplification. We conclude by stressing the importance of refining our understanding of the complex interplay between the disk environment and a forming giant planet, and its evolutionary history after envelope accretion, which might turn out to be more directly responsible for its final structure and composition than the metal content of the parent star."
    },
    "0809/0809.2400_arXiv.txt": {
        "abstract": "I review the basic observational properties of accreting millisecond pulsars that are important for understanding the physics involved in formation of their pulse profiles. I then discuss main effects responsible for shaping these profiles. Some analytical results that help to understand the results of simulations are presented. Constraints on the pulsar geometry and the neutron star equation of state obtained from the analysis of the pulse profiles are discussed. ",
        "introduction": "First evidences for  rapid rotation of neutron stars (NS) in low-mass X-ray binaries appeared in 1996 with the discoveries of sub-kHz coherent oscillations observed during X-ray bursts by the {\\it Rossi X-ray Timing Explorer (RXTE)} \\cite[see][ for review]{SB06}. Two years later, the first accreting X-ray millisecond pulsar (AMP) was discovered \\cite{WvdK98}, which was a nice confirmation of the theory of production of millisecond radio pulsar in the course of accretion. By 2008, there are at least 8 accreting NS discovered that show coherent oscillations for the extended periods of time during the outbursts \\cite{W06,P06}. These sources are exceptional laboratories to study the physics of accretion onto magnetized stars. They also have a great potential to test the NS structure. These goals can be achieved by studying the pulse profiles.  This, however, requires understanding of the processes responsible for the production of the X-rays at the NS surface  as well as detailed modeling of the propagation of the radiation to the observer.  I review here the main observational data that have direct relation to this topic. These include averaged and phase-resolved broad-band spectra, pulse profiles and their energy dependence, and time lags. Then I describe recent efforts to develop theoretical models devoted to modeling these data and the main results obtained from these studies. ",
        "conclusions": "I have reviewed the basic observational properties of AMP such as spectral energy distribution, pulse profiles, time lags that are important for understanding of the physics of accretion in these sources.  I also discussed the major effects that are responsible for shaping the pulses in AMP: light bending, Doppler boosting, anisotropy of emission, eclipses by the accretion disk, etc. The complicated dependences of the pulses on various parameters trigger the analytical approach which allows to understand better the results of simulations. I discussed how various observables such as amplitudes of fundamental and harmonic depend on the main parameters as well as how the invisibility of the secondary spot  constrains geometrical parameters and the inner disk radius. This in turn can be used to constrain the pulsar magnetic field. Various attempts to measure NS compactness produce contradicting results, which probably reflects our ignorance of the angular distribution of the radiation as well as the spot geometry.  Theoretical studies of these aspects of accretion onto AMP are certainly welcome.      \\begin{theacknowledgments} I acknowledge the support from the Academy of Finland grant 110792. I thank my collaborators on this topic Andrei Beloborodov, Maurizio Falanga, Marek Gierli\\'nski, and Askar Ibragimov. \\end{theacknowledgments}"
    },
    "0809/0809.5199_arXiv.txt": {
        "abstract": "{ Three very massive clusters are known to reside in the Galactic Center region, the Arches cluster, the Quintuplet cluster and the Central Parsec cluster, each of them being rich in young hot stars. With new infrared instruments this region is no longer obscured to the observer. } {For understanding these very massive clusters, it is essential to know their stellar inventory. We provide comprehensive spectroscopic data for the stellar population of the Quintuplet cluster that will form the basis of subsequent spectral analyses. } { Spectroscopic observations of the Quintuplet cluster have been obtained with the Integral Field Spectrograph SINFONI-SPIFFI at the ESO-VLT. The inner part of the Quintuplet cluster was covered by 22 slightly overlapping fields, each of them of $8\\,\\arcsec \\times 8\\,\\arcsec$ in size. The spectral range comprises the near-IR $K$-band from 1.94 to 2.45\\,$\\mu$m. The 3D data cubes of the individual fields were flux-calibrated and combined to one contiguous cube, from which the spectra of all detectable point sources were extracted. } { We present a catalog of 160 stellar sources in the inner part of the Quintuplet cluster. The flux-calibrated $K$-band spectra of 98 early-type stars and 62 late-type stars are provided as Online Material. Based on these spectra, we assign spectral types to all detected sources and derive synthetic $K_\\textrm s$-band magnitudes. Our sample is complete to about the 13th $K$-magnitude. We report the detection of two hitherto unknown Wolf-Rayet stars of late WC type (WC9 or later). Radial velocities are measured and employed to asses the cluster membership. The quantitative analysis of the early-type spectra will be subject to a subsequent paper. } {} ",
        "introduction": "The Galactic Center has long been hidden from optical observation due to high extinction in the visual. {\\changed Today} the advanced instruments for infrared radiation give access to this region. Three surprisingly young and massive stellar clusters were found within {\\changed a projected distance of} 35\\,pc from the central black hole: the Arches, the Quintuplet and the Central cluster. These clusters are found to be rich in massive stars (Arches:\\,\\citealt{Blum+2001, Figer2004, Martins+2008}; Quintuplet:\\,\\citealt{FigerMcLeanNajarro1997, FMM99, Figer2004}; Galactic Center Cluster:\\,\\citealt{Eckart+2004, Figer2004}). The prominent constellation of five infrared-bright stars gave reason to name one of the {\\changed clusters} the ``Quintuplet cluster'' \\,\\citep{Okuda+1987, Okuda+1989}. It has a projected distance of 30\\,pc to the Galactic Center, a cluster radius of about 1\\,pc, and an estimated age of about 4 million years\\,\\citep{Okuda+1990, FMM99}.  First spectroscopic observations of the five ``Quintuplet proper'' stars {\\changed \\citep[later named Q1, Q2, Q3, Q4 and Q9 by][]{FMM99}} showed only a continuum without any features. Some of them have been detected also as far-IR and radio sources \\citep{Lang+1997, Lang+1999, Lang+2005}.  \\citet{Tuthill+2006} resolved two of the original Quintuplet stars (Q2 and Q3) in space and time as colliding wind binaries, ejecting dust in the shape of a rotating ``pinwheel''.  According to further spectroscopic studies \\citep{FMM99, Figer2004}, the cluster contains at least 100 O stars and 16 Wolf-Rayet stars.  Two massive{\\changed ,} evolved stars in their luminous blue variable {\\changed (LBV)} phase, the ``Pistol star'' and one further LBV candidate, have been found in the Quintuplet cluster \\citep{Geballe+2000}.  {\\changed At least two luminous} stars of late spectral type K and M {\\changed were} found in the cluster as well.  The present paper {\\changed (hereafter the ``LHO catalog'')} provides a comprehensive spectral atlas of $K$-band spectra from stellar sources in the Quintuplet cluster. These spectra can form the basis for a subsequent analysis of the luminous stellar population, which is {\\changed a} prerequisite for understanding the formation and evolution of this very massive cluster in its special galactic environment.  The paper is structured as follows: Section\\,\\ref{sec:observations} gives details on the observations and data reduction. Section\\,\\ref{sec:catalog} contains the catalog with position, spectral classification, synthetic K magnitude and radial velocity. The final Sect.\\,\\ref{sec:summary} summarizes the catalog. The spectral atlas comprising 160 stars is available as Online Material. ",
        "conclusions": "\\label{sec:summary} We present {\\changed the first} $K$-band spectral catalog of {\\changed 160} stellar sources in the central region of the Quintuplet cluster. Integral field {\\changed spectra} were obtained with the SINFONI-SPIFFI on UT4 from May to July 2006. Our catalog comprises {\\changed flux-calibrated spectra of} 98 {\\changed early-type} stars (OB and WR) and 62 {\\changed late-type} stars (K to M) {\\changed with a spectral resolution of $R \\approx 4000$.} Table\\,\\ref{tab:number-statistics} lists the number of stars per spectral class.  \\begin{table}[!ht] \\caption[]{Spectral type distribution} \\label{tab:number-statistics} \\begin{center} \\begin{tabular}{lr} \\hline \\hline \\noalign{\\smallskip} Spectral Type & No. of stars \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} WN             & 4      \\\\ WC             & 9  \\\\ {\\it{WR total}}& {\\it{13}} \\\\ \\hline \\noalign{\\smallskip} O              & 60 \\\\ B              & 25 \\\\ {\\it{OB total}}& {\\it{85}} \\\\ \\hline \\noalign{\\smallskip} K              & 43 \\\\ M              & 19  \\\\ {\\it{KM total}}& {\\it{62}} \\\\ \\hline {\\it{total}}   & 160 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} All {\\changed 160} sources are listed in Table\\,\\ref{tab:catalog} with their {\\changed LHO} number that was assigned to the source in our data reduction and identification processes (column 1), the determined coordinates of the source ({\\changed columns} 2 and 3) and the derived synthetic $K_\\textrm s$ magnitude (column 4). The spectral type is given in column 5. {\\changed Radial velocities (RV) are given in column 6, together with an assessment if this RV value is indicating a foreground object (``f''). Finally we give a complete list of alias names and cross-identification with previous catalogs and surveys in colum 7. Note that about 100 objects of our list are cataloged here for the first time.}  Two new WR stars {\\changed of WC spectral subtype} are found in our sample, LHO\\,76 and 79, increasing the number ratio WC:WN in the Quintuplet cluster to {\\changed 10:6} in total (9:4 contained in our field). WC stars are outstandingly frequent and bright in the Quintuplet cluster. For comparison, {\\changed the WC:WN number ratios from \\citet{vdH2006} are 0:17 for the Arches cluster, 12:17 for the Central cluster, and 8:19 for Westerlund~I, respectively. A} difference {\\changed to} the solar neighbourhood {\\changed WC:WN} ratio of 38:25 {\\changed \\citep{vdH2001} can be seen.} A detailed analysis of the WR star sub-sample with the Potsdam models for expanding atmospheres (PoWR) is in preparation.  {\\changed Among the new stars in our catalog there are a number of late-type supergiants. In case that the red supergiants (RSG) are real cluster members, as supported by their radial velocities, the question of age and coeval evolution of the Quintuplet cluster has to be readdressed in this context.}  The spectral atlas of the catalog is only available as Online Material. It comprises for each detected source the flux-calibrated $K$-band spectrum in the wavelength range from 1.95 to 2.45\\,$\\mu$m. The spectra are binned to 4\\,\\AA\\,. Some spectra suffer from atmospheric OH emission lines."
    },
    "0809/0809.3295_arXiv.txt": {
        "abstract": "In this paper, we report results of our near-infrared (NIR) photometric variability studies of the BL Lacertae object S5 0716$+$714. NIR photometric observations spread over 7 nights during our observing run April 2$-$9, 2007 at 1.8 meter telescope equipped with KASINICS (Korea Astronomy and Space Science Institute Near Infrared Camera System) and J, H, and Ks filters at Bohyunsan Optical Astronomy Observatory (BOAO), South Korea. We searched for intra-day variability, short term variability and color variability in the BL Lac object. We have not detected any genuine intra-day variability in any of J, H, and Ks passbands in our observing run. Significant short term variability $\\sim$ 32.6\\%, 20.5\\% and 18.2\\% have been detected in J, H, Ks passbands, respectively, and $\\sim$ 11.9\\% in (J-H) color. ",
        "introduction": "Blazars constitute a small subclass of the most enigmatic class of radio-loud active galactic nuclei (AGNs) consisting of BL Lacertae objects (BL Lacs) and flat spectrum radio quasars (FSRQs). BL Lacs show largely featureless optical continuum. Blazars exhibit strong flux variability at all wavelengths of the complete electromagnetic (EM) spectrum, strong polarization ($> 3\\%$) from radio to optical wavelengths, usually core dominated radio structures, and predominantly nonthermal radiation at all wavelengths. In a unified model of radio-loud AGNs based on the angle between the line of sight and the emitted jet from the source, blazars jets make angle of $<$ 10$^{\\circ}$ from the line of sight (Urry \\& Padovani 1995).  From the study of the spectral energy distributions (SEDs) of blazars, it is found that blazars SEDs have two peaks (Fossati et al. 1998; Ghisellini et al. 1998). The first component peaks at near-infrared (NIR)/optical in low energy peaked blazars (LBLs) and at UV/X-rays in high energy peaked blazars (HBLs). The second component peaks at GeV energies in LBLs and at TeV energies in HBLs. The EM emission is dominated by synchrotron component at low energies and at high energies probably by the inverse Compton (IC) component (Coppi 1999, Sikora \\& Madejski 2001, Krawczynski 2004).  From observations of blazars, it is known that they vary on the diverse time scales ranging from a few minutes to several years. Blazars variability can be broadly divided into three classes viz. intra-day variability (IDV) or intra-night variability or micro-variability, short term variability and long term variability. Variations in flux of a few tenth of a magnitude over the course of a day or less is often known as IDV (Wagner \\& Witzel 1995). Short and long term variabilities can have time scales from a few weeks to several months and several months to years, respectively.  First convincing optical IDV in blazar was reported by Miller et al. (1989), and since then variability of blazars on diverse time scales in radio to optical bands have been studied extensively and results have been reported in large number of papers (e.g. Carini 1990; Mead et al. 1990; Takalo et al. 1992, 1996; Heidt \\& Wagner 1996; Sillanp\\\"a\\\"a et al. 1996a, 1996b; Bai et al. 1998, 1999; Fan et al. 1998, 2002, 2007; Xie, et al. 2002; Ciprini et al. 2003, 2007; Gupta et al. 2004, 2008a, 2008b; Stalin et al. 2005 and references therein). In a recent paper, Gupta \\& Joshi (2005) have done statistical analysis of occurrence of optical IDV in different classes of AGNs. They divided their sample of 113 optical lights curves of blazars in three different time durations and found 64\\%(18/28), 63\\%(29/46), and 82\\%(32/39) blazars show IDV if observed for $\\leq$3h, 3h to $\\leq$6h, and $>$6h, respectively.  S5 0716$+$714 ($\\alpha_{2000.0} =$ 07:21:53.4, $\\delta_{2000.0} = +$71:20:36.4) is one of the brightest BL Lac object which has featureless optical continuum. The non detection of its host galaxy first sets a lower limit of redshift z $>$ 0.3 (Wagner et al. 1996), and then z $>$ 0.52 (Sbarufatti et al. 2005). Very recently, Nilsson et al. (2008) have claimed of its host galaxy detection which produces a ``standard candle'' value of $z = 0.31 \\pm 0.08$. Wagner \\& Witzel (1995) reported that the duty cycle of the source is one which implies that the source is almost always in the active state. The variability of S5 0716$+$714 has been studied in the complete EM spectrum on all time scales (e.g. Raiteri et al. 2003; Gupta et al. 2008a and references therein). Large and variable optical polarization in the source has been reported (Takalo et al. 1994; Fan et al. 1997; Impey et al. 2000). Since 1994, this source has been extensively monitored in optical bands. There are 5 major optical outbursts reported in the source: at the beginning of 1995, in late 1997, in the fall of 2001, in March 2004 and in the beginning of 2007 (Raiteri et al. 2003; Foschini et al. 2006; Gupta et al. 2008a). These 5 outbursts give possible period of long term variability of $\\sim$ 3.0$\\pm$0.3 years.  Compared to radio and optical bands, there are only a few attempts to search for NIR flux variability on diverse timescales in blazars (e.g. Mead et al. 1990; Takalo et al. 1992; Gupta et al. 2004; Hagen-Thorn et al. 2006 and references therein). Since blazars emit radiations in the complete EM spectrum, they are ideal candidates for multi-wavelength observations. But either due to unavailability of good quality NIR detectors or unavailability of low humidity observing sites at several 1-2 meter class NIR/optical telescopes around the world, there were no focused effort to search for NIR flux variability in LBLs in which SEDs synchrotron component peaks in NIR/optical bands. Now we have an excellent opportunity to carry out such observations from the Bohyunsan Optical Astronomy Observatory (BOAO), South Korea which has 1.8 meter telescope and is equipped with KASINICS (Korea Astronomy and Space Science Institute Near Infrared Camera System) (Moon et al. 2008). We have recently started our long term pilot project to search for flux variability on diverse time scales in LBLs. Our present and future planned observations will fill the gap between radio and optical bands and will give deep insight into the important and less studied NIR flux variability properties of LBLs. Simultaneous radio to gamma-ray observations will be useful to detect the synchrotron and IC component peaks of LBLs from the SEDs and will be useful to understand the emission mechanism of LBLs in the complete EM spectrum. With this motivation, we recently carried out J, H, Ks bands photometric observations of our first target, the BL Lac object S5 0716$+$714 which is a LBL, in 7 observing nights during April 2 $-$ April 9, 2007.  The paper is structured as follows: section 2 describes observations and data reduction methods, in section 3 we mention our results, the discussions and conclusions of the present work is reported in section 4. ",
        "conclusions": "From our multi-band NIR observations of the blazar S5 0716$+$714 in 7 observing nights in April 2007, genuine IDV in any of J, H and Ks is not detected. We noticed the existence of significant short term flux variability in the blazar from our observations. The total short term variation detected in our observations in J, H, Ks passbands are $\\sim$ 32.6\\%, 20.5\\% and 18.2\\%, respectively. Our data show significant variation in J$-$H color ($\\sim$ 11.9\\%) but H$-$Ks color variation was not detected. The difference in short term variations in H and Ks passbands is only 2.3\\% which caused no genuine H$-$Ks color variation. We have noticed variable spectral index (ranging from $-0.76$ to $-1.00$) with a mean value of $\\sim -0.88$ in our observations. The variable spectral index is mainly due to variation in J band flux. We also found correlated flux and spectral index (higher the flux, higher the spectral index). We observed the source in the post outburst state. Outburst of the source was reported by Gupta et al. (2008a) in their January $-$ February 2007 observations and they also noticed in their March 2007 observations, the source was becoming fainter than January $-$ February 2007.  S5 0716$+$714 was the target of three simultaneous multi-wavelength campaigns (Tagliaferri et al. 2003; Ostorero et al. 2006; Villata et al. 2008) and also several monitoring campaigns in single or two EM bands (e.g. radio-optical, optical and optical-X-ray) (Wagner et al. 1990; Quirrenbach et al. 1991; Sagar et al. 1999; Villata et al. 2000; Foschini et al. 2006; and references therein). It has shown radio and optical IDVs during all radio to optical campaigns (Heeschen et al. 1987; Wagner et al. 1990, 1996; Ghisellini et al. 1997; Sagar et al. 1999; Quirrenbach et al. 2000; Raiteri et al. 2003; Agudo et al. 2006; Gupta et al. 2008a, and references therein). It is the first IDV source in which simultaneous variations in radio and optical bands were detected which indicated a possible intrinsic origin of the observed IDV (Wagner et al. 1990; Quirrenbach et al. 1991). VLBI observations of the source over more than 20 years, show a very compact source at centimeter wavelengths with an evidence of a core-dominated jet extending several tens of milliarcseconds to the North (Eckart et al. 1986, 1987; Witzel et al. 1988; Polatidis et al. 1995; Jorstad et al. 2001). The X-ray observations have shown strong variations with short flares ($\\approx$ 1000s) detected with the {\\it ROentgen SATellite} (ROSAT) (Cappi et al. 1994). The source was also detected in hard X-rays up to 60 keV when observed after the outburst state of 2000 (Tagliaferri et al. 2003). It has been detected by {\\it Energetic Gamma-Ray Experiment Telescope} (EGRET) onboard the {\\it Compton Gamma-Ray Observatory} CGRO at GeV energies with steep $\\gamma-$ray spectrum (Hartman et al. 1999). But the soft $\\gamma-$ray part of its SED is poorly known and upper limit of the source detection in 3$-$10 MeV energy from the {\\it imaging COMPton TELescope} (COMPTEL) was reported by Collmar (2006). An exceptional energy sampling data of the blazar was obtained in the simultaneous multi-wavelength observing campaign in November 2003 (Ostorero et al. 2006). The source was very bright at radio frequencies and in a rather low optical state (R = 14.17 - 13.64). Significant short term variability and IDV were detected in the radio bands. The source was not detected by INTEGRAL in the observing campaign but the upper limit of the source emission in 3$-$200 keV was estimated. On September 2007, the source was detected in $\\gamma-$rays by the recently launched satellite {\\it Astro-rivelatore Gamma a Immagini LEggero} (AGILE) (Villata et al. 2008).  Several models have been developed to explain the IDV and short term variability in radio-loud AGNs viz. the shock-in-jet models, accretion disk based models (e.g. Wagner \\& Witzel 1995; Urry \\& Padovani 1995; Ulrich et al. 1997; and references therein). For blazars in the outburst state, IDV and short term variability are strongly supported by the jet based models of radio-loud AGNs. In general, blazars emission in the outburst state is nonthermal Doppler boosted emission from jets (Blandford \\& Rees 1978; Marscher \\& Gear 1985; Marscher et al. 1992; Hughes et al. 1992). On the other hand, IDV and short term variability of blazars in the low-state can be explained by the models based on some kind of instability in the accretion disk (e.g. Mangalam \\& Wiita 1993; Chakrabarti \\& Wiita 1993).  Here we rule out the possibility of emission from the accretion disk because it is expected to be relevant only in the source low-state (the source was in the post outburst period). The source was in the outburst state $\\sim$ 2 months before the present observations (Gupta et al. 2008a). In the low-state, jet emission is less dominant over the thermal emission from the accretion disk. According to the unified scheme of radio-loud AGNs, blazars are seen nearly face on, so, any fluctuations on the accretion disk should produce a detectable change in the emission characteristics.  The observed short term variability in the blazar S5 0716$+$714 is possibly explained by the jet based model known as turbulent jet model. According to this model, Marscher et al. (1992) suggested that the Reynolds number in the relativistic jet should be very high which will cause turbulent jet plasma. The shock will impinge upon regions of slightly different magnetic field strengths, densities and velocities, so the observed flux is expected to vary. The timescales of this proposed type of variations are shorter and their amplitudes larger at higher frequencies. In the observations reported here, we got variability amplitudes of $\\sim$ 32.6\\%, 20.5\\% and 18.2\\% in J, H, Ks bands, respectively, implying larger variations at higher frequencies. The relation between flux variation amplitude and frequency is confirmed from the SED and color variation.  Since the duty cycle of the source is 1, we may expect to detect IDV. However, we have not detected any genuine IDV in our observations. It might be due to less numbers of data points and short durations of the observing runs during each night, hence in the future we plan to observe the source for a longer time each night. This will allow us to collect more data points and possibly to have a higher signal-to-noise in order to have higher sensitivity to possible genuine IDVs."
    },
    "0809/0809.0668_arXiv.txt": {
        "abstract": "Mirror matter is a self-collisional dark matter candidate. If exact mirror parity is a conserved symmetry of the nature, there could exist a parallel hidden (mirror) sector of the Universe which has the same kind of particles and the same physical laws of our (visible) sector. The two sectors interact each other only via gravity, therefore mirror matter is naturally ``dark''. The most promising way to test this dark matter candidate is to look at its astrophysical signatures, as Big Bang nucleosynthesis, primordial structure formation and evolution, cosmic microwave background and large scale structure power spectra. ",
        "introduction": "The idea that there may exist a hidden mirror sector of particles and interactions with exactly the same properties  as our visible world was suggested long time ago by Lee and Yang~\\cite{mirror}, and the model with exact parity symmetry interchanging corresponding fields of two sectors was proposed many years later by Foot at al.~\\cite{mirror}. The two sectors communicate with each other only via gravity\\footnote{ There could be other interactions, as for example the kinetic mixing between O and M photons, but they are negligible for the present study.}. A discrete symmetry $G\\leftrightarrow G'$ interchanging corresponding fields of $G$ and $G'$, so called mirror parity, guarantees that two particle sectors are described by identical Lagrangians, with all coupling constants (gauge, Yukawa, Higgs) having the same pattern. As a consequence the two sectors should have the same microphysics. Once the visible matter is built up by ordinary baryons, then the mirror baryons would constitute dark matter in a natural way, since they interact with mirror photons, but not interact with ordinary photons.  The phenomenology of mirror matter was studied in several papers (for an extended list see the bibliografy of ref.~\\cite{okun50}), in particular the implications for Big Bang nucleosynthesis~\\cite{bcv,g_mir}, primordial structure formation and cosmic microwave background~\\cite{paolo,paolo1}, large scale structure of the Universe~\\cite{paolo2,ignavol-lss}, microlensing events (MACHOs)~\\cite{mirstar,mirMacho}.  If the mirror (M) sector exists, then the Universe along with the ordinary (O) particles should contain their mirror partners, but their densities are not the same in both sectors. In fact, the BBN bound on the effective number of extra light neutrinos implies that the M sector has a temperature lower than the O one, that can be naturally achieved in certain inflationary models~\\cite{dolgov-dnu}. Then, two sectors have different initial conditions, they do not come into thermal equilibrium at later epoch and they evolve independently, separately conserving their entropies, and maintaining approximately constant the ratio among their temperatures.  All the differences with respect to the ordinary world can be described in terms of only two free parameters in the model, \\begin{equation}\\label{mir-param} x \\equiv \\left( s' \\over s \\right)^{1/3} \\approx {T' \\over T} ~~~~~~~~~~ ; ~~~~~~~~~~ \\beta \\equiv \\Omega'_{\\rm b} / \\Omega_{\\rm b} ~~~~, \\end{equation} where $T$ ($T'$), $\\Omega_{b}$ ($\\Omega'_{b}$), and $s$ ($s'$) are respectively the ordinary (mirror) photon temperature, cosmological baryon density, and entropy density. The bounds on the mirror parameters are $ x < 0.7 $ and $ \\beta > 1 $, the first one coming from the BBN limit and the second one from the hypothesis that a relevant fraction of dark matter is made of mirror baryons.  As far as the mirror world is cooler than the ordinary one, $x < 1$, in the mirror world all key epochs (as are baryogenesis, nucleosynthesis, recombination, etc.) proceed in somewhat different conditions than in ordinary world. Namely, in the mirror world the relevant processes go out of equilibrium earlier than in ordinary world, which has many far going implications. ",
        "conclusions": "\\begin{figure}\\label{resume} \\includegraphics[height=.52\\textheight]{resume.eps} \\caption{Current status of the astrophysical research with mirror dark matter: solid lines mark what is already done, while dashed ones mark what is still to do.} \\end{figure}  Figure \\ref{resume} shows the current situation of the astrophysical research in presence of mirror dark matter. We have already investigated the early Universe (thermodynamics and Big Bang nucleosynthesis), and the process of structure formation in linear regime, that permit to obtain predictions, respectively, on the primordial elements abundances, and on the observed cosmic microwave background and large scale structure power spectra. In addition, we have studied the evolution of mirror dark stars, which, together with the mirror star formation, are necessary ingredients for the study and the numerical simulations of non linear structure formation, and of the formation and evolution of galaxies; furthermore, in future studies they will provide predictions on the observed abundances of MACHOs and on the gravitational waves background. Ultimately we will be able to obtain theoretical estimates to be compared with observations of gravitational lensing, galactic dark matter distribution, and strange astrophysical events still unexplained (as for example dark galaxies, bullet galaxy, ...).  Concluding, the astrophysical tests so far used show that mirror matter can be a viable candidate for dark matter, but we still need to complete the entire picture of the Mirror Universe.    \\begin{theacknowledgments} This work was supported by the Belgian Science Policy Office Inter University Attraction Pole VI/11 ``Fundamental Interactions''. \\end{theacknowledgments}"
    },
    "0809/0809.2492_arXiv.txt": {
        "abstract": "{Open clusters older than $\\sim4$\\,Gyr are rare in the Galaxy. Affected by a series of mass-decreasing processes, the stellar content of most open clusters dissolves into the field in a time-scale shorter than $\\sim1$\\,Gyr. In this sense, improving the statistics of old objects may provide constraints for a better understanding of the dynamical dissolution of open clusters.} {Our main purpose is to investigate the nature of the Globular cluster candidate FSR\\,1716, located at $\\ell=329.8^\\circ$ and $b=-1.6^\\circ$. We also derive parameters of the anti-centre open cluster Czernik\\,23 (FSR\\,834). Both objects have been detected as stellar overdensities in the Froebrich, Scholz \\& Raftery star cluster candidate catalogue.} {The analyses are based on near-infrared colour-magnitude diagrams and stellar radial density profiles. The intrinsic colour-magnitude diagram morphology is enhanced by a field-star decontamination algorithm applied to the 2MASS \\jj, \\hh, and \\ks\\ photometry.} {Isochrone fits indicate that FSR\\,1716 is more probably an old ($\\sim7$\\,Gyr) and absorbed ($\\aV=6.3\\pm0.2$) open cluster, located $\\approx0.6$\\,kpc inside the Solar circle in a contaminated central field. However, we cannot rule out the possibility of a low-mass, loose globular cluster. Czernik\\,23 is shown to be an almost absorption-free open cluster, $\\sim5$\\,Gyr old, located about 2.5\\,kpc towards the anti-centre. In both cases, Solar and sub-Solar ($[Fe/H]\\sim-0.5$) metallicity isochrones represent equally well the stellar sequences. Both star clusters have a low mass content ($\\la200\\,\\ms$) presently stored in stars. Their relatively small core and cluster radii are comparable to those of other open clusters of similar age. These structural parameters are probably consequence of the several Gyrs of mass loss due to stellar evolution, tidal interactions with the disk (and bulge in the case of FSR\\,1716), and possibly giant molecular clouds.} {Czernik\\,23, and especially FSR\\,1716, are rare examples of extreme dynamical survivors. The identification of both as such represents an increase of $\\approx10\\%$ to the known population of open clusters older than $\\sim4$\\,Gyr in the Galaxy.} ",
        "introduction": "\\label{Intro}  The Galaxy is an aggressive environment to star clusters in general, the open clusters (OCs) in particular. These stellar systems are continually harassed by a series of dynamical processes such as mass loss associated to stellar evolution, mass segregation and evaporation, tidal interactions with the Galactic disk and bulge, and collisions with giant molecular clouds. Combined over time, such processes tend to accelerate the dynamical evolution, which produces significant changes in the cluster structure and eroded mass functions. Eventually, most OCs end up completely dissolved in the Galactic stellar field or as poorly-populated remnants (\\citealt{PB07} and references therein).  Theoretical (e.g. \\citealt{Spitzer58}; \\citealt{LG06}), N-body (e.g. \\citealt{BM03}; \\citealt{GoBa06}; \\citealt{Khalisi07}), and observational (e.g. \\citealt{vdB57}; \\citealt{Oort58}; \\citealt{vHoerner58}; \\citealt{Piskunov07}) evidence indicate that the disruption-time scale near the Solar circle is shorter than $\\sim1$\\,Gyr. Around this region, the disruption time-scale depends on mass as $\\tdis\\sim M^{0.62}$ (\\citealt{LG06}), which for clusters with mass in the range $10^2 - 10^3\\ms$ corresponds to $\\rm75\\la\\tdis(Myr)\\la300$. In general, the effect of the relevant dynamical processes is stronger for the OCs more centrally located and the low-mass ones (see \\citealt{OldOCs} for a detailed discussion on disruption effects and time-scales). Indeed, OCs older than $\\sim1$\\,Gyr are preferentially found near the Solar circle and in the outer Galaxy (e.g. \\citealt{Friel95}; \\citealt{DiskProp}), where the frequency of potentially damaging dynamical interactions with giant molecular clouds and the disk is lower (e.g. \\citealt{Salaris04}; \\citealt{Upgren72}). Disruption efficiency increases critically towards the Galactic centre, to the point that the inner ($\\dgc\\la150$\\,pc) tidal fields can dissolve a massive star cluster in $\\sim50$\\,Myr (\\citealt{Portegies02}).  The above aspects considered, the natural expectation is that only a small fraction of the OCs can reach old ages, and that the successful ones should be preferentially found at large Galactocentric distances. In fact, of the $\\approx1000$ OCs with known age listed in the WEBDA\\footnote{\\em obswww.univie.ac.at/webda - \\citet{Merm03}} database, 180 are older than 1\\,Gyr, and only 18 ($\\approx2\\,\\%$) are older than 4\\,Gyr (see also \\citealt{OBB05a}; \\citealt{OBB05b}). Not surprisingly, most of the OCs older than 1\\,Gyr so far identified are located outside the Solar circle (see, e.g., the spatial distribution of OCs of different ages in Fig.~1 of \\citealt{OldOCs}).  \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1a.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1b.eps} \\end{minipage}\\hfill \\caption[]{Left panel: $4\\arcmin\\times4\\arcmin$ 2MASS \\ks\\ image of FSR\\,1716. Image provided by the 2MASS Image Service. The small circle indicates the central coordinates (cols.~2 and 3 of Table~\\ref{tab1}). Right panel: $7\\arcmin\\times7\\arcmin$ XDSS R image of Cz\\,23. Figure orientation: North to the top and East to the left.} \\label{fig1} \\end{figure*}  In this context, it is naturally expected that the discovery and derivation of astrophysical parameters of old OCs will better define the OC parameter space. Thus, a better understanding on the dynamical survival rate of star clusters in the tidal field of the Galaxy can be reached. Such parameters, in turn, can be used in studies of star formation and evolution processes, dynamics of N-body systems, and the geometry of the Galaxy, among others.  \\begin{table*} \\caption[]{General data on the clusters} \\label{tab1} \\tiny \\renewcommand{\\tabcolsep}{0.65mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{lcccccccccccccccc} \\hline\\hline &\\multicolumn{7}{c}{FSR2007}&&\\multicolumn{7}{c}{This paper}\\\\ \\cline{2-8}\\cline{10-16} Cluster&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$&Q&\\rch&\\rth&&Age&\\aV&\\ds&\\dgc&\\xgc&\\ygc&\\zgc\\\\ &(hms)&($^\\circ\\,\\arcmin\\,\\arcsec$)&($^\\circ$)&($^\\circ$)& &(\\arcmin)&(\\arcmin) &&(Gyr)&(mag)&(kpc)&(kpc)& (kpc)&(kpc)& (kpc)\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&&(9)&(10)&(11)&(12)&(13)&(14)&(15)\\\\ \\hline FSR\\,1716$^\\dagger$&16:10:33&$-$53:44:12&329.79&$-1.59$&2&1.2&5.9&&$7.0\\pm1.0$&$6.3\\pm0.2$ &$0.8\\pm0.1$&$6.6\\pm0.1$&$-6.6\\pm0.1$&$-0.37\\pm0.04$&$-0.02\\pm0.01$\\\\  FSR\\,1716$^\\ddagger$&16:10:33&$-$53:44:12&329.79&$-1.59$&2&1.2&5.9&&$12.0\\pm2.0$&$6.3\\pm0.4$ &$2.3\\pm0.3$&$5.4\\pm0.3$&$-5.3\\pm0.3$&$-1.13\\pm0.17$&$-0.06\\pm0.01$\\\\  Cz\\,23&05:50:07&$+$28:53:28&180.55&$+0.82$&1&0.8&4.2&&$5.0\\pm1.0$&$0.0\\pm0.1$ &$2.5\\pm0.1$&$9.7\\pm0.2$  &$-9.7\\pm0.1$&$-0.02\\pm0.01$&$+0.04\\pm0.01$\\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Coordinates (cols.~2 to 5), quality flag (col.~6), and core and tidal radii (cols.~7 and 8) measured in the \\hh\\ band are from \\citet{FSRcat}; Col.~10: reddening towards the cluster's central region (Sect.~\\ref{age}). Col.~11: distance from the Sun. Col.~12: cluster Galactocentric distance for $\\rs=7.2$\\,kpc (\\citealt{GCProp}). Cols.~13-15: coordinate components projected onto the Galactic plane. Cz\\,23 is the optical counterpart of FSR\\,834. Parameters of FSR\\,1716 are derived for the OC ($\\dagger$) or globular cluster ($\\ddagger$) interpretation (Sect.~\\ref{age}). \\end{list} \\end{table*}  In the present paper we study two stellar overdensities listed in the star cluster candidate catalogue of \\citet{FSRcat}, which turn out to be very old star clusters. They are FSR\\,1716 and FSR\\,834. The present work employs near-IR \\jj, \\hh, and \\ks\\ photometry obtained from the 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release (\\citealt{2mass1997}), available at {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} Point Source Catalogue (PSC). The spatial and photometric uniformity of 2MASS, which allow extraction of wide surrounding fields that provide high star-count statistics, are important to derive cluster parameters and probe the nature of stellar overdensities (e.g. \\citealt{ProbFSR}). For this purpose we have developed quantitative tools to statistically disentangle cluster evolutionary sequences from field stars in colour-magnitude diagrams (CMDs). Basically, we apply {\\em (i)} field-star decontamination to quantify the statistical significance of the CMD morphology, which is important to derive reddening, age, and distance from the Sun, and {\\em (ii)} colour-magnitude filters, which are essential for intrinsic stellar radial density profiles (RDPs), as well as luminosity and mass functions (MFs). In particular, field-star decontamination constrains more the age and distance, especially for low-latitude OCs (\\citealt{DiskProp}).  This paper is organised as follows. Sect.~\\ref{Target_OCs} contains basic properties and reviews literature data (when available) on both star cluster candidates. In Sect.~\\ref{2mass} we present the 2MASS photometry and build the stellar surface-density distribution in the direction of both objects. In Sect.~\\ref{Decont_CMDs} we build CMDs, discuss the field-star decontamination algorithm, and provide tests to the old age of both clusters. In Sect.~\\ref{age} we derive cluster fundamental parameters. Sect.~\\ref{struc} describes cluster structure by means of stellar RDPs. In Sect.~\\ref{MF} mass functions are built and cluster masses are estimated. In Sect.~\\ref{Discus} aspects related to the structure and dynamical state of the present clusters are considered. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  In this paper we study the nature of two stellar overdensities included in the catalogue of candidate star clusters of \\citet{FSRcat}, FSR\\,1716 and FSR\\,834. The former was suggested to be a globular cluster candidate by \\citet{FSRcat}, while the latter has the OC Cz\\,23 as optical counterpart. The analyses are based on field-star decontaminated 2MASS CMDs and stellar radial density profiles, with algorithms previously constructed by our group. Fundamental and structural parameters of the clusters are derived.  We present consistent evidence (e.g. CMD morphology, statistical tests, structural parameters, mass-function slope, and comparison with nearby OCs) that both objects are old star clusters.  FSR\\,1716 is significantly absorbed ($\\aV\\approx6.3$) and projected not far from the bulge. Its field-decontaminated CMD morphology is very similar to that of the $\\sim7$\\,Gyr well-known OC NGC\\,188. Indeed, its CMD can be well represented by isochrones with ages older than $\\sim6$\\,Gyr, both of Solar and sub-Solar ($[Fe/H]\\sim-0.5$) metallicity. We adopted the former as the metallicity of FSR\\,1716, because of its relatively central location. Alternatively, we cannot rule out the possibility that FSR\\,1716 is a low-mass, loose (Palomar-like) globular cluster. FSR\\,1716 is located inside the Solar circle, $\\approx0.6$\\,kpc (in the case of an OC) or $\\approx1.8$\\,kpc (GC).  The CMD morphology of Cz\\,23 is consistent with that of the $\\sim4$\\,Gyr old OC M\\,67. Similarly to FSR\\,1716, it can be well represented by Solar and sub-Solar metallicity isochrones, but with ages in the range $4-6$\\,Gyr. With the $\\approx5$\\,Gyr and Solar metallicity solution, we find that Cz\\,23 is projected nearly towards the anti-centre, located $\\approx2.5$\\,kpc outside the Solar circle.  The core and cluster radii of FSR\\,1716 and Cz\\,23 are small when compared to a set of open clusters in the Solar neighbourhood. However, such radii are comparable with those of other OCs of similar old age. Besides, the mass functions appear to be much flatter than Salpeter's IMF, especially FSR\\,1716, which seems to present an increasing depletion in the number of low-mass stars. As a consequence, the total masses presently stored in stars in both clusters are lower than $\\sim200\\,\\ms$. Such low values probably reflect the several Gyr-long period of mass loss due to stellar evolution, tidal interactions with the bulge (possibly in the case of FSR\\,1716), disk and giant molecular clouds. Actually, because of its low mass content and flat mass function, Cz\\,23 may be evolving into an open cluster remnant (e.g. \\citealt{PB07}).  Comprehensive catalogues of star cluster candidates, such as that of \\citet{FSRcat}, should be further explored with field-star decontamination algorithms and other tools, so that the nature of the candidates can be probed and the age derived. It is remarkable how the decontamination tool unveiled the intrinsic CMD sequences of FSR\\,1716, separating it from the crowded field population. Consequently, the characterisation of FSR\\,1716 and Cz\\,23 as OCs older than $\\sim4$\\,Gyr represents an important increase ($\\approx10\\%$) to the known population of such objects in the Galaxy. In particular, FSR\\,1716 is the most recent addition to the 8 open clusters older than 6\\,Gyr so far identified (WEBDA). In this sense, Cz\\,23 (FSR\\,834), and especially FSR\\,1716, can be considered as rare examples of extreme dynamical survivors in disk-regions where most open clusters are short-lived."
    },
    "0809/0809.2171_arXiv.txt": {
        "abstract": "{} {We have analyzed optical spectra of 25 X-ray sources identified as potential new members of the \\object{Taurus molecular cloud} (TMC), in order to confirm their membership in this star-forming region.} {Fifty-seven candidate members were previously selected among the X-ray sources in the XEST survey, having a 2MASS counterpart compatible with a pre-main sequence star based on color-magnitude and color-color diagrams. We obtained high-resolution optical spectra for 7 of these candidates with the SARG spectrograph at the TNG telescope, which were used to search for lithium absorption and to measure the H$\\alpha$ line and the radial and rotational velocities. Then, 18 low-resolution optical spectra obtained with the instrument DOLORES for other candidate members were used for spectral classification, for H$\\alpha$ measurements, and to assess membership together with IR color-color and color-magnitude diagrams and additional information from the X-ray data.} {We found that 3 sources show lithium absorption, with equivalent widths (EWs) of $\\sim 500$\\,m\\AA, broad spectral line profiles, indicating rotational velocities of $\\sim 20-40$\\,km\\,s$^{-1}$, radial velocities consistent with those for known members, and H$\\alpha$ emission. Two of them are classified as new weak-lined T~Tauri stars, while the EW ($\\sim -9$\\,\\AA\\ ) of the H$\\alpha$ line and its broad asymmetric profile clearly indicate that the third star (\\object{XEST-26-062}) is a classical T~Tauri star. Fourteen sources observed with DOLORES are M-type stars. Fifteen sources show H$\\alpha$ emission. Six of them have spectra that indicate surface gravity lower than in main sequence stars, and their de-reddened positions in IR color-magnitude diagrams are consistent with their derived spectral type and with pre-main sequence models at the distance of the TMC. The K-type star \\object{XEST-11-078} is confirmed as a new member on the basis of the strength of the H$\\alpha$ emission line. Overall, we confirm membership to the TMC for 10 out of 25 X-ray sources observed in the optical. Three sources remain uncertain.} {} ",
        "introduction": "\\label{intro}  The formation of stars and planets and the evolution of stellar properties during the early stages of their lives (internal structure, angular momentum, magnetic activity, and X-ray emission in low-mass stars) is one of the most intriguing problems in astrophysics. The fragmentation of large clouds and the subsequent gravitational collapse of molecular cores is a stochastic process that leads to a population of stellar and sub-stellar objects with a variety of masses. The details of these physical mechanisms are still not clear; turbulence and density fluctuations appear to be crucial \\citep{Klessen2001,Goodwin2004b}, but other factors probably play important roles, such as the presence of magnetic fields \\citep{Padoan1999} or the shock waves propagating from supernovae explosions that may determine the fragmentation of clouds \\citep{Truelove1997}.  Theoretical models of cloud fragmentation and core collapse should reproduce the \\emph{initial mass function} (IMF) of a star-forming region (SFR), that is the mass distribution of a stellar group just born. Significant differences between the IMFs derived for the \\object{Taurus-Auriga SFR} \\citep{Luhman2003Toro}, which hosts a distributed mode of star formation, and for much denser regions containing massive stars, such as \\object{Orion} and \\object{IC 348}  \\citep[][ respectively]{Muench2002,Luhman2003IC348}, have suggested that the star formation process may be affected by the physical conditions of the environment where stars form. However, previous studies of the TMC may not be complete at low masses, while the most recent surveys \\citep{Briceno2002,Luhman2003Toro,Luhman2004,Guieu2006,Luhman2006b} have allowed discovery of new low-mass stars and brown dwarfs, but were focused on limited portions of the region and are not spatially complete. In fact, studies of the TMC are also complicated by its large extension in the sky ($\\sim 100$ square degrees, also due to its distance of 140\\,pc), and require long observational campaigns with sufficiently deep exposures. More complete studies of the Taurus population are therefore needed to assess the IMF shape with greater confidence, especially at the very low-mass end.  We started a search for new members of the Taurus-Auriga SFR \\citep{ScelsiXEST2006}, based on X-ray data from the \\emph{{\\rm XMM}-Newton Extended Survey of the Taurus Molecular Cloud} \\citep[XEST, ][]{GuedelXEST2006} and on near-infrared data from the 2MASS point source catalog \\citep{Skrutskie2006}. This is a different approach with respect to the above mentioned searches for new Taurus members, where candidates were selected based on optical and IR data. Since intense X-ray emission \\citep[10 to $10^4$ times the solar level, see, e.g., ][]{Feigelson1999,Stelzer2001,Ozawa2005} is a ubiquitous characteristic of pre-main sequence solar-type stars and thanks to the relatively low interstellar absorption at these wavelengths, X-ray observations are particularly efficient at detecting the population of young objects, including very low-mass stars that are faint in the optical bands, and can therefore serve as a complement to optical/IR searches. Moreover, since non-accreting ``weak-lined'' T~Tauri stars (WTTSs) are generally more luminous in X rays \\citep[e.g. ][]{BriggsXEST2006} and less absorbed than classical T~Tauri stars (CTTSs), which are still surrounded by a thick circumstellar accretion disk, the former are expected to be selected more efficiently in X-ray surveys. This is particularly important, because it helps to reduce possible biases introduced by optical/IR surveys, which may favour the detection of CTTSs owing to their strong H${\\alpha}$ emission and IR excess, and hence allows us to properly address fundamental issues such as the estimate of disk lifetimes, with important consequences on our understanding of the evolution of the angular momentum during the earlier phases of stellar life and in the formation of planetary systems.  In this paper we present the results of the analysis of follow-up optical spectroscopy obtained at the Telescopio Nazionale Galileo (TNG) for 25 out of 57 candidate Taurus members selected in the previous work by \\citet{ScelsiXEST2006}. The paper is organized as follows. Section \\ref{sel} summarizes the main information about the XEST survey and the X-ray/IR based selection of potential new members; in Sect. \\ref{obs} we describe the follow-up TNG observations of the optically brightest candidates. In Sect. \\ref{opt_spec} we present the results of the optical spectroscopy, while in Sect. \\ref{IRdiagr} we determine the stellar properties. Finally, we discuss our results in Sect. \\ref{dis}. ",
        "conclusions": "\\label{dis}  In this work we employed optical spectroscopy, IR photometry, and X-ray data to characterize a sample of 25 sources, previously identified as candidate pre-main sequence stars. We confirm membership in the Taurus-Auriga star-forming region for 10 of them, while 12 sources are identified as main sequence or older stars in front of or behind the cloud. Three sources remain as uncertain cases. It is worth noting that we have confirmed membership for 5 candidates out of 9 with a high probability assigned by \\citet{ScelsiXEST2006} on the basis of the X-ray analysis; 3 more are the uncertain cases, and only one such source (\\object{XEST-27-084}) has been identified as an active M-type MS star. The percentage of new Taurus members with respect to the number of observed stars is at least 40\\%, up to $\\sim 50$\\% depending on the nature of the uncertain sources. Similar percentages of confirmed new TMC members have been obtained in other works previously mentioned (Sect. \\ref{intro}), where candidate selection was based on optical and infrared photometry. Hence, our work indicates that candidate selections based on X-ray emission and infrared photometry are as effective as selections based on optical and infrared photometry. We also note that higher percentages of confirmed members could be obtained if the selection is made considering only X-ray sources within the $M-L_{\\rm X}$ relation discussed below.  It is also interesting that 5 out of 10 newly discovered TMC members (i.e. the three lithium stars and 2 members observed at low-resolution) are variable in X-rays (see Table \\ref{tab:lista}). Short-term X-ray variability, often associated to flares, is a frequent characteristic of active stars \\citep{Guedel2003,Wolk2005,Caramazza2007}. The X-ray variability of pre-main sequence stars in the Taurus molecular cloud has recently been studied by \\citet{StelzerXEST2006}, who found that about half of a sample of 122 members detected in the XEST survey are variable. Our result is therefore consistent with the percentage of X-ray variable members in Taurus.  As explained in Sects. \\ref{sarg} and \\ref{IRdiagr_lrs_membri}, we found 8 new WTTS and 2 new CTTS, which is not surprising considering the initial X-ray-based candidate selection. The T~Tauri type has been deduced from the strength of the H$\\alpha$ emission and the shape of the line profile for the stars observed with SARG. The stars \\object{XEST-08-003}, \\object{XEST-08-033}, \\object{XEST-08-047}, \\object{XEST-16-045}, and \\object{XEST-17-059} also appear in the catalog of the \\emph{Taurus Spitzer Legacy Project} (Data Release 1); the photometry obtained with the IRAC and MIPS cameras onboard the \\emph{Spitzer Space Telescope}, together with the 2MASS measurements, gives SEDs typical of stars without significant IR excess for circumstellar material, and thus confirms their nature as weak-lined T~Tauri stars.  The new members have spectral types in the range $late$-K --- M so are low-mass stars, with masses estimated in the range $\\sim 0.2-0.6$\\,M$_{\\odot}$. These members fall in the same bins where the mass distribution of the $\\sim 150$ known Taurus members (within the same XMM fields) shows its bulk. However, the 25 sources observed at the TNG and analyzed here are the optically brightest stars of the initial sample selected by \\citet{ScelsiXEST2006}, and therefore other new members with masses of $\\sim 0.1$\\,M$_{\\odot}$, may be hidden among the 32 remaining candidates. Indeed, about 20 of them are very-low mass candidates ($M\\lesssim 0.2$\\,M$_{\\odot}$), and about ten of them are brown dwarf candidates\\footnote{The XEST has detected 8 out of 16 known brown dwarfs included in the survey.}, which will be studied in the continuation of our optical follow-up program.  \\begin{figure}[t!] \\begin{center} \\scalebox{0.5}{ \\includegraphics{sce08f7a.ps}} \\scalebox{0.5}{ \\includegraphics{sce08f7b.ps}} \\caption{\\emph{Upper panel}: Positions of the candidates in the $M-L_{\\rm X}$ diagram, with masses and X-ray luminosities estimated before the TNG observations, assuming they are all members. Circles and diamonds mark stars observed with SARG and DOLORES, respectively. A dot inside the symbol means the candidate had a higher probability of membership on the basis of X-ray data. Squares indicate new members confirmed from this work, and ``?'' mark uncertain cases. The dashed line is the best fit to the data for known TMC members and the dotted lines bracket the region where known members are found. \\emph{Lower panel}: $M-L_{\\rm X}$ diagram updated after the TNG observations. Circles mark new members from this work, with circle size indicating age, and crosses main sequence field stars. For the uncertain cases \\object{XEST-15-034} and \\object{XEST-08-014} both positions as a PMS or a field star are plotted. Light green lines are the saturation at 1\\,Myr, 10\\,Myr, and in the main sequence.} \\label{fig:Lx_M} \\end{center} \\end{figure} Figure \\ref{fig:Lx_M} shows the $M-L_{\\rm X}$ diagram for the sources studied in this work, for which mass estimates have been possible. In the upper panel, masses and X-ray luminosities were estimated by \\citet{ScelsiXEST2006} before the TNG observations and under the assumption that the candidates were indeed TMC members\\footnote{In this plot, the mass of \\object{XEST-11-078} had been estimated under the wrong assumption of spectral type~M and no IR excess.}. In this plot we report the relation $\\log\\,L_{\\rm X}=1.54\\,\\log\\,M+30.31$, best-fitting the data for known TMC members within the fields of the XEST survey \\citep{GuedelXEST2006}, and we also mark the new Taurus members from this study. All new TMC members are found within the relation derived from known members, while the five sources outside this relation (i.e. \\object{XEST-18-059}, \\object{XEST-27-022}, \\object{XEST-21-059}, \\object{XEST-04-060}, and \\object{XEST-11-035}) are confirmed as non-members. This result both confirms the validity of a correlation between mass and X-ray luminosity for the young stars of Taurus and suggests that the $M-L_{\\rm X}$ diagram may be used in future searches for new pre-main sequence stars as a further criterion for the identification of new candidate members.  The lower panel of the same figure updates the previous plot after the observations at the TNG and the analysis presented in this paper. In this plot, the new members are also distinguished by their age, as calculated in Sects. \\ref{IRdiagr_sarg} and \\ref{IRdiagr_lrs}. We did not include giant or subgiants field stars, since their masses and X-ray luminosities were not estimated. We marked the saturation curves $L_{\\rm X}=10^{-3}\\,L_{\\rm bol}$ relevant to different evolutionary stages, i.e. stars at 1\\,Myr, at 10\\,Myr, and on the main sequence \\citep[the relation between bolometric luminosity and mass was taken from the models by][]{Baraffe1998}. Interestingly, the slopes of the theoretical saturation curves for PMS stars are very similar to the slope of the $M-L_{\\rm X}$ relation found for Taurus members, suggesting that part of the observed spread around this best-fit relation is due to saturated stars at different ages. However, both the data in the plot and the narrower $L_{\\rm X}$ range spanned by the saturation curves in the $1-10$\\,Myr range suggest that other factors may play a role, such as flares, non-saturated stars (slow rotators), and the possibly lower $L_{\\rm X}$ of CTTS with respect to WTTS \\citep{Preibisch2005,BriggsXEST2006}.  Finally, it is also interesting to note that the saturation curve for main sequence stars steepens at $M\\sim 0.5-0.6$\\,M$_{\\odot}$, and explains why older stars may be found within the relation for TMC members. In fact, most of our confirmed field stars are close to, or in, the saturation regime."
    },
    "0809/0809.2940_arXiv.txt": {
        "abstract": "{In core-collapse supernovae, $\\nu$ and $\\overline\\nu$ are initially subject to significant self-interactions induced by weak neutral currents, which may induce strong-coupling effects on the flavor evolution (collective transitions). The interpretation of the effects is simplified when self-induced collective transitions are decoupled from ordinary matter oscillations, as for the matter density profile that we discuss. In this case, approximate analytical tools can be used (pendulum analogy, swap of energy spectra). For inverted $\\nu$ mass hierarchy, the sequence of effects involves: synchronization, bipolar oscillations, and  spectral split. Our simulations shows that the main features of these regimes are not altered when passing from simplified (angle-averaged) treatments to full, multi-angle numerical experiments.}   \\normalsize\\baselineskip=15pt  \\begin{figure}[b] \\center \\includegraphics[height=1.9in]{school.eps} \\caption{In a school of fish, individuals often show a collective behavior. Very dense neutrino gases, like those emerging from a core-collapse supernova, might show analogous features in flavor space.} \\label{school} \\end{figure} ",
        "introduction": " ",
        "conclusions": "We have studied supernova neutrino oscillations in a model where the collective flavor transitions (synchronization, bipolar oscillations, and spectral split) are well separated from later, ordinary MSW effects. We have performed numerical simulations in both single- and multi-angle cases, using continuous energy spectra with significant $\\nu$-$\\overline \\nu$ and $\\nu_e$-$\\nu_x$ asymmetry. The results of the single-angle simulation can be largely understood by means of an analogy with a classical gyroscopic pendulum. The main  observable effect appears to be the swap of final-state energy spectra, for inverted hierarchy, at a critical energy dictated by lepton number conservation. In the multi-angle simulation, details of self-interaction effects can change, but the spectral split remains a robust, observable feature. In this sense, averaging over neutrino trajectories does not alter the main effect of the self interactions. From the point of view of neutrino parameters, collective flavor oscillations in supernovae could be instrumental in identifying the inverse neutrino mass hierarchy, even for tiny $\\theta_{13}$.~\\cite{DigheMiri}  \\begin{figure}[b] \\center \\includegraphics[height=1.9in]{school2.eps} \\caption{A school of fish branching out in two different directions. Analogously, supernova neutrino polarization vectors might split up in flavor space, due to self-interaction effects.} \\label{school2} \\end{figure}"
    },
    "0809/0809.0497_arXiv.txt": {
        "abstract": "Caustics are a generic feature of the nonlinear growth of structure in the dark matter distribution. If the dark matter were absolutely cold, its mass density would diverge at caustics, and the integrated annihilation probability would also diverge for individual particles participating in them.  For realistic dark matter candidates, this behaviour is regularised by small but non-zero initial thermal velocities. We present a mathematical treatment of evolution from Hot, Warm or Cold Dark Matter initial conditions which can be directly implemented in cosmological N-body codes. It allows the identification of caustics and the estimation of their annihilation radiation in fully general simulations of structure formation. ",
        "introduction": "The idea that the dark matter might consist of a collisionless ``gas'' of weakly interacting, neutral particles was first published by \\cite{1973ApJ...180....7C} and \\cite{1976A&A....49..437S}, though earlier discussion can also be found in the textbook of \\cite{Zeldovich1971}. These authors proposed neutrinos as a promising dark matter candidate, and a claimed measurement of the electron neutrino mass at very nearly the expected value \\citep{1981Lyubimov} led to a flurry of interest in neutrino-dominated, so-called ``Hot Dark Matter'' universes.  This continued until detailed numerical simulations showed the predictions for low-redshift structure to be quite inconsistent with observation \\citep{1983ApJ...274L...1W}. Attention then shifted towards more exotic particle candidates for Warm or Cold Dark Matter \\citep{1982PhRvL..48..223P, 1982ApJ...263L...1P}.  If a cold collisionless gas evolves from near-uniform initial conditions under the influence of gravity, the nonlinear phases of growth generically involve caustics analogous to those formed when light propagates through a non-uniform medium. This connection was explored in some depth by Russian cosmologists interested in neutrino-dominated universes \\citep{1982GApFD..20..111A,1983SvPhU.139..153Z}. Caustics are also a very evident feature of the similarity solutions for cold spherical infall published by \\cite{1984ApJ...281....1F} and \\cite{1985ApJS...58...39B}. It was another 15 years, however, before \\cite{2001PhRvD..64f3515H} realised that dark matter annihilation could be very substantially enhanced in such caustics.  He showed that for absolutely cold dark matter the annihilation probability diverges logarithmically as a particle passes through a caustic, and that for realistic dark matter candidates this divergence is tamed by the small but finite initial thermal velocities of the particles. He argued that the annihilation radiation from dark halos might be dominated by emission from caustics.  Twenty years earlier \\cite{1982PAZh....8..259Z} had noted that thermal velocities limit the densities achievable in dark matter caustics, and the first rigorous calculation of annihilation rates in caustics was carried out by \\cite{2006MNRAS.366.1217M} for Bertschinger's (1985) similarity solution.  They found caustics to enhance the total annihilation flux substantially in the outer regions for plausible values of the initial dark matter velocity dispersion, but to be progressively less important at smaller radii.  It is unclear whether either of these results will apply in general, since the behaviour of the similarity solution is strongly influenced by its spherical symmetry (which reduces its phase-space dimensionality from 6 to 2) and by its lack of small-scale structure.  In the present paper we present a theoretical treatment of the growth of structure which shows how the geodesic deviation equation can be used to follow local phase-space structure in a Lagrangian treatment of nonlinear evolution. This formalism is well suited for implementation in N-body simulation codes, allowing the annihilation signal from caustics to be treated in full generality provided numerical artifacts from discretisation and integration error can be kept under control. We have presented results from a first implementation of this scheme in \\cite{2008MNRAS.385..236V}. A closely related but somewhat more complex scheme is described by \\cite{2005MNRAS.359..123A}. Here we complement this work by giving a fuller description of the mathematics behind the approach, in particular of the regularisation of caustic densities by the finite velocity dispersion of the dark matter. In the next section we describe our idealisation of the initial conditions for structure formation in WIMP-dominated cosmologies. Section 3 then presents useful general results for nonlinear evolution from these initial conditions. These are used in section 4 to describe the evolution of the local phase-space structure, in particular of its projection onto configuration space, following individual particle trajectories.  Caustic passages can be identified and the associated annihilation signal can be calculated explicitly. A final section discusses possible future uses of this approach. ",
        "conclusions": "In this paper we have developed the mathematical background to enable a relatively precise evaluation of the annihilation radiation from dark matter caustics in fully general simulations of the nonlinear growth of structure.  Our scheme allows the annihilation rate to be integrated along the trajectory of each simulation particle, including correctly the contributions from all the caustics in which it participates. Typically each particle experiences several such caustic passages in each orbit around the dark matter halo in which it resides \\citep[see][]{2008MNRAS.385..236V}. In order to include correctly the annihilation rate between particles which are members of {\\it different} streams, it is necessary to estimate a local coarse-grained density at the position of each particle, and to add in the contribution due to streams other than its own. This can be done, for example, using the SPH technique, since the smoothing this introduces does not bias the luminosity predicted from inter-stream annihilations. When a particle passes through a caustic occurring in a stream other than its own, the time-integrated annihilation probability is still correctly reproduced in the smoothed system.  By implementing these methods in high-resolution simulations of galaxy formation it should be possible to achieve a complete numerical description of the expected annihilation radiation, limited only by the ability to resolve the smallest collapsed clumps of dark matter. This latter limitation can be severe when attempting to predict the total annihilation radiation from a representative cosmological volume. For a standard supersymmetric neutralino, for example, the emission should be dominated by the smallest collapsed objects, with masses well below that of the Sun \\citep[e.g.][]{2003MNRAS.339..505T}. Recent work has shown, however, that this problem is much less severe when predicting the observability of annihilation radiation from the Solar System, which lies just 8 kpc from the centre of the Milky Way \\citep{2008arXiv0809.0894S}. According to these authors, less than 3 percent of the dark matter within 100 kpc of the Galactic Centre should be in small lumps; almost all the rest should be in extended streams of the kind discussed in this paper \\citep[see also][]{1999MNRAS.307..495H}. They argue that the highest signal-to-noise for detecting the annihilation signal will be that of the smooth dark matter distribution in the inner few kiloparsecs of the Galaxy; small subhalos will be significantly more difficult to detect. Application of the techniques we have presented here should allow a rigorous evaluation of an important and previously unresolved issue: whether the emission structure of this smooth component is significantly modified by caustic emission. In addition, they will allow an assessment of the expected morphology and observability of outer caustics around external galaxies, most notably the Andromeda nebula.  The authors thank St\\'ephane Colombi, Craig Hogan, Roya Mohayaee and Volker Springel for helpful discussions."
    },
    "0809/0809.1943_arXiv.txt": {
        "abstract": "We present an analysis of the spatial distribution of various stellar populations within the Large and Small  Magellanic Clouds.  We use optically selected stellar samples with mean ages between $\\sim9$ and $\\sim1000$~Myr, and existing stellar cluster catalogues to investigate how stellar structures form and evolve within the LMC/SMC.  We use two statistical techniques to study the evolution of structure within these galaxies, the $Q$-parameter and the two-point correlation function (TPCF).  In both galaxies we find the stars are born with a high degree of substructure (i.e. are highly fractal) and that the stellar distribution approaches that of the 'background' population on timescales similar to the crossing times of the galaxy ($\\sim80$~Myr \\& $\\sim150$~Myr for the SMC/LMC respectively).  By comparing our observations to simple models of structural evolution we find that 'popping star clusters' do not significantly influence structural evolution in these galaxies.  Instead we argue that general galactic dynamics are the main drivers, and that substructure will be erased in approximately the crossing time, regardless of spatial scale, from small clusters to whole galaxies.  This can explain why many young Galactic clusters have high degrees of substructure, while others are smooth and centrally concentrated.  We conclude with a general discussion on cluster 'infant mortality',  in an attempt to clarify the time/spatial scales involved. ",
        "introduction": "Most, if not all, stars are thought to be born in a {\\it clustered} or fractal distribution, which is usually interpreted as being due to the imprint of the gas hierarchy from which stars form (e.g. Elmegreen \\& Efremov~1996, Elmegreen et al.~2006, Bastian et al.~2007).  Older stellar distributions, however, appear to be much more smoothly distributed, begging the questions; 1) what is the main driver of this evolution? and 2) what is the timescale for the natal structure to be erased?  It has been noted that many nearby star forming clusters also have hierarchical structure seemingly dictated by the structure of the dense gas of natal molecular clouds (e.g. Lada \\& Lada 2003).  Using statistical techniques, Gutermuth et al. (2005) studied three clusters of varying degrees of embeddedness and demonstrated that the least embedded cluster was also the least dense and the least substructured of the three. That result hinted at the idea that the youngest clusters are substructured, but that dynamical interactions and ejection of the structured gas contributes to the evolution and eventual erasure of that substructure in approximately the cluster formation timescale of a few Myr (Palla \\& Stahler 2000). By having an accurate model of the spatial evolution of stellar structures we can approach a number of fundamental questions about star-formation, including what is the percentage of stars born in \"clusters\" and whether this depends on environmental conditions. Using automated algorithms on infrared Spitzer surveys of star-forming sites within the Galaxy (e.g. Allen et al. 2007), such constraints are now becoming possible.  In this contribution we present results of two recent studies on the evolution of structure in the LMC (Bastian et al.~2008) and SMC (Gieles et al.~2008).  \\begin{figure}[h] \\begin{center} \\includegraphics[width=3.4in]{cmdvi.ps} \\includegraphics[width=5.4in]{spatial-distribution.ps} \\caption{{\\bf Top:} A V-I vs. V CMD of stars in the LMC.  The selected boxes are shown, along with theoretical stellar evolutionary isochrones. {\\bf Bottom:} The spatial distribution of stars in boxes 1, 5, \\& 9 in the LMC.  The mean age of the stellar distributions is given at the top of each panel.}% \\label{fig:cmd} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0809/0809.4754.txt": {
        "abstract": "We present results of optical spectroscopic and photometric observation of the pre-main sequence stars associated with the cometary shaped dark cloud Lynds~1622, and $^{12}$CO and $^{13}$CO observations of the cloud. We determined the effective temperatures and luminosities of 14 pre-main sequence stars associated with the cloud from their positions in the Hertzsprung--Russell diagram, as well as constructed their spectral energy distributions using optical, 2MASS and \\textit{Spitzer} IRAC and MIPS data. We derived physical parameters of L\\,1622 from the molecular observations. Our results are not compatible with the assumption that L\\,1622 lies on the near side of the Orion--Eridanus loop, but suggest that L\\,1622 is as distant as Orion~B. At a distance of 400~pc the mass of the cloud, derived from our $^{12}$CO data, is 1100\\,M$_{\\sun}$, its star formation efficiency is $\\sim 1.8\\%$, and the average age of its low-mass pre-main sequence  star population is about 1~million years. ",
        "introduction": "\\label{Sect_1}  Studies of the large-scale properties of star forming regions are important for understanding the interstellar processes governing star formation. Census of young stellar objects (YSOs) born in a cloud, their mass, age and space distribution, together with the mass of the star forming cloud are key data for assessment of the efficiency and time scale of star formation.  The Orion~OB1 association and its associated molecular clouds represent a distinguished target for star formation studies, being the nearest giant star forming region where various stages of the star forming processes across the whole stellar mass spectrum, as well as interactions between the consecutive generations of stars can be studied in detail. The target of the present paper, Lynds~1622 and the neighbouring Lynds~1621 are small dark clouds at the low Galactic latitude edge of the giant molecular cloud Orion~B. These two clouds together are termed as {\\em Orion East\\/} by \\citet{HR72}, who discovered five H$\\alpha$ emission stars associated with L\\,1622. \\citet{Maddalena} found that the radial velocity  of this cloud is about $v_\\rmn{LSR} = +1$\\,km\\,s$^{-1}$, contrary to the characteristic radial velocity of +10\\,km\\,s$^{-1}$ of Orion~B. The asymmetric, cometary shape and the bright, ionized rim apparent in the SHASSA \\citep{Gaustad} image of L\\,1622 suggest its interaction with the hot stars of Ori~OB\\,1. \\citet{Casassus} present centimeter continuum image of L\\,1622, and attribute the radiation to spinning dust grains in a limb-brightened shell of the cloud, where the incident UV radiation from the Ori~OB1b association heats and charges the grains.  Low-mass star formation in L\\,1622 is indicated by several associated T~Tauri stars and candidates. The reflection nebula VDB~62, illuminated by the K2 type weak-line T~Tauri star HBC~515 (HD\\,288313, V1793~Ori) \\citep{Racine,HBC} can be found near the bright rim of the cloud. Spectroscopic and photometric variability of this star were studied in detail by \\citet{Mekkaden}, and it was found to be a triple system by \\citet{Reipurth08}. In addition to this WTTS, there are six known classical T~Tauri stars (CTTS) in L\\,1622: HBC~188 (LkH$\\alpha$~334), HBC~189 (LkH$\\alpha$~335), HBC~190 (LkH$\\alpha$~336), HBC~516 (LkH$\\alpha$~336c), HBC~191 (LkH$\\alpha$~337), and HBC~517 (L\\,1622--15) \\citep{HBC}. The LkH$\\alpha$~336 triple system was studied by \\citet{Correia}.  \\citet{OH83} detected 16 H$\\alpha$ emission stars (L\\,1622--1 to L\\,1622--16) in an objective prism survey over the surface of L\\,1622, including the above HBC stars except LkH$\\alpha$~336c and HD\\,288313. The pre-main sequence (PMS) nature of 11 objects has not yet been confirmed.  A further signpost of low-mass star formation is the Herbig--Haro object HH~122 \\citep{Reipurth89}, whose probable exciting source, HH\\,122~VLA~1 was detected by \\citet{Rodriguez}. \\citet{LM99} report four optically selected cores of the L\\,1621\\,/\\,L\\,1622 region: L\\,1621--1, L\\,1621--2, L\\,1622~A and L\\,1622~B. L\\,1622~B is associated with the optically invisible source \\textit{IRAS}~05522+0146, while \\citet{LMT01} list L\\,1622~A as an infall candidate. Recently \\citet{Reipurth_hb} identified 28 young stellar objects (YSOs) in the \\textit{Spitzer} IRAC images of L\\,1622, 10 of which they classified as protostars.  The velocity difference between Orion~B and L\\,1622 raises the question whether both clouds  are located in the same volume of space, or L\\,1622 is a foreground object, unrelated to the Orion~B molecular cloud.  \\citet{Maddalena} concluded that, in spite of the different velocities, Orion East is probably located at the same distance as Orion~B, regarding its cometary shape and bright rim, pointing toward Barnard's Loop. This conclusion was questioned by \\citet{Knude}, who studied the distance to interstellar extinction features by using spectral types of the Michigan Catalog and Tycho--2 photometry, and found an absorbing cloud located at some 160~pc from the Sun toward the line of sight of L\\,1622, suggesting that this cloud may be a foreground object. \\citet{Wilson05}, based on the parallaxes of three \\textit{Hipparcos\\/} stars, find that L\\,1622 is situated as close as 120\\,pc to the Sun. Nearby star forming clouds are of special interest because they allow us to study fine structural details and detect the lowest luminosity objects.  Motivated by these findings we started a spectroscopic and photometric study of the pre-main sequence stars and candidates associated with L\\,1622. Our aim is to assess the probable star forming history of the cloud, and see if the measured properties of these objects support the smaller or the larger distance.  In order to have a homogeneous data set on the young low mass stars born in L\\,1622 we obtained moderate resolution optical spectra and $VR_\\rmn{C}I_\\rmn{C}$ photometry of all pre-main sequence stars and candidates projected on L\\,1622. \\textit{Spitzer} archive data were used for studying the spectral energy distributions (SEDs) of the young stars. In order to find the physical properties of the cloud, compare the positions of the stars and the molecular gas, as well as estimate the efficiency of star formation  we analysed $^{12}$CO and $^{13}$CO observations of the cloud available in the data archive of the NANTEN radio telescope at Nagoya University. Our observations and results are described in Sect.~2, and discussed in Sect.~3. Section~4 gives a short summary of our findings. ",
        "conclusions": "\\label{Sect_concl}  We determined effective temperatures and luminosities of 14 pre-main sequence stars and candidates in the region of the dark cloud L\\,1622 and found that these quantities are incompatible with the assumption that L\\,1622 lies at the near side of the Orion--Eridanus Bubble, but suggest that this cloud is located at the distance of Orion~B. The derived ages of the stars at this distance are around 1 million~years. The shapes of the SEDs of the target stars support this age estimate. Star formation in L\\,1622 was probably triggered by the luminous stars of the association subgroup Ori~OB1b (Orion's Belt).  We present new $^{12}$CO and $^{13}$CO maps of L\\,1622. Accepting a distance of 400~pc we derived a mass of 1100\\,M$_{\\sun}$ for L\\,1622 from the $^{12}$CO observations. The total mass of the YSOs identified so far, assuming a mean mass of 0.5\\,M$_{\\sun}$ for the \\textit{Spitzer} sources not included in our spectroscopic survey, is $\\approx$20\\,M$_{\\sun}$, suggesting a star formation efficiency $SFE \\approx 1.8\\%$.  Several of our programme stars may be interesting targets for more detailed studies. The derived luminosity of the CTTS L\\,1622--6 suggests that this object is a binary or multiple system. The large inner hole in the accretion disc of this star, suggested by the models fitted to its SED, supports this hypothesis. The SEDs of L\\,1622--10A, L\\,1622--10B, and LkH$\\alpha$\\,336c suggest that the discs of these stars are seen nearly edge-on. L\\,1622--3  exhibits an extremely rich emission spectrum. This star might have been born in a small molecular clump  outside the main cloud. \\textit{IRAS}~05522+0146 is a Class~I source embedded in the dense core L\\,1622\\,B."
    },
    "0809/0809.1172_arXiv.txt": {
        "abstract": "{  We determine the metallicity distribution function (MDF) of the Galactic halo by means of a sample of 1638 metal-poor stars selected from the Hamburg/ESO objective-prism survey (HES). The sample was corrected for minor biases introduced by the strategy for spectroscopic follow-up observations of the metal-poor candidates, namely ``best and brightest stars first''. Comparison of the metallicities [Fe/H] of the stars determined from moderate-resolution (i.e., $R\\sim 2000$) follow-up spectra with results derived from abundance analyses based on high-resolution spectra (i.e., $R>20,000$) shows that the [Fe/H] estimates used for the determination of the halo MDF are accurate to within 0.3\\,dex, once highly C-rich stars are eliminated. We determined the selection function of the HES, which must be taken into account for a proper comparison between the HES MDF with MDFs of other stellar populations or those predicted by models of Galactic chemical evolution. The latter show a reasonable agreement with the overall shape of the HES MDF for $\\mbox{[Fe/H]} > -3.6$, but only a model of Salvadori et al. (2007) with a critical metallicity for low-mass star formation of $Z_{\\mathrm cr}=10^{-3.4}\\,Z_{\\odot}$ reproduces the sharp drop at $\\mbox{[Fe/H]} \\sim -3.6$ present in the HES MDF.  Although currently about ten stars at $\\mbox{[Fe/H]} < -3.6$ are known, the evidence for the existence of a tail of the halo MDF extending to $\\mbox{[Fe/H]} \\sim -5.5$ is weak from the sample considered in this paper, because it only includes two stars $\\mbox{[Fe/H]} < -3.6$.  Therefore, a comparison with theoretical models has to await larger statistically complete and unbiased samples.  A comparison of the MDF of Galactic globular clusters and of dSph satellites to the Galaxy shows qualitative agreement with the halo MDF, derived from the HES, once the selection function of the latter is included. However, statistical tests show that the differences between these are still highly significant. ",
        "introduction": "\\label{Sect:Intro}  One of the key observables for constraining models of the formation and chemical evolution of the Galaxy is the Metallicity Distribution Function (MDF) of the constituent stars of its various components (bulge, disk, halo). The MDF provides critical information on the enrichment history of those components with heavy elements. In the case of the halo, early enrichment may have been provided by the very first generations of massive stars, formed from material of primordial composition shortly after the Big Bang (i.e., Population III stars).  Models of Galactic chemical evolution need to be compared to an accurate (and precise) observed halo MDF to test their predictions, to constrain their various parameters (such as the effective yield, the star-formation rate and the IMF), and in order to obtain information on the properties of Population III stars that are responsible for the earliest enrichment. This is particularly important for the lowest metallicity tail of the MDF, which provides invaluable information on the earliest enrichment phases \\citep{Prantzos:2003}; for instance, it has been suggested that a minimum level of enrichment is required to form low-mass stars. This critical metallicity ranges between $10^{-4}\\mathrm{Z}_{\\odot}$ \\citep{Omukai:2000,Brommetal:2001,Bromm/Loeb:2003,Umeda/Nomoto:2003,Santoro/Shull:2006,Frebeletal:2007c} and $10^{-6}\\mathrm{Z}_{\\odot}$, the latter being applicable when dust grains are present \\citep{Schneideretal:2002,Schneideretal:2003,Schneideretal:2006,Omukaietal:2005,Tsuribe/Omukai:2006,Clarketal:2008}.   The precision of a derived halo MDF increases directly with the total number of observed metal-poor halo stars. Selection of such stars without the introduction of a kinematic bias (e.g., from among high proper motion stars) makes them of particular utility for examination of the relationships between the chemistry and kinematics of the halo. Early determinations of the halo MDF were based on small samples of globular clusters (\\citealt{Hartwick:1976}; $N=60$), or a mixture of halo subdwarfs and globular clusters (\\citealt{Bond:1981}; $N=90$ and $N=31$, respectively).  Problems with these samples arise not only from their small sizes, but also their inaccurate metallicities. Later studies employed significantly larger samples with spectroscopically-determined stellar abundances. For example, \\citet{Ryan/Norris:1991} used a sample of 372 kinematically-selected halo stars. \\citet{Ryan/Norris:1991} and \\citet{Carneyetal:1996} showed that the MDF peaks at a metallicity of $\\mathrm{[Fe/H]} = -1.6$ with wings from $\\mathrm{[Fe/H]} = -3.0$ to solar abundances.  The HK survey \\citep{BPSI,BPSII,TimTSS}, originated by Preston and Shectman, and greatly extended by Beers to include several hundred additional objective-prism plates, was, until the advent of the Hamburg/ESO Survey (HES; see below), the primary source of metal-poor candidates suitable for consideration of the halo MDF. With the assistance of numerous colleagues, medium-resolution spectroscopy of over 10,000 HK-survey stars was obtained, using 1.5--4\\,m class telescopes, over the past two decades. This led to the identification of thousands of stars with $\\mathrm{[Fe/H]} < -2.0$, as well as significant numbers of stars with $\\mathrm{[Fe/H]} < -3.0$.  \\begin{figure}[htbp] \\centering \\includegraphics[clip=true,bb=75 199 372 770,width=8.8cm]{10925f01.ps} \\caption{\\label{Fig:VmagBminV} Upper panel: Isochrones for an age of 12\\,Gyr and metallicities of $\\mbox{[Fe/H]}=-1$, $-2$, and $-3$ \\citep{Kimetal:2002}, and chosen colour cuts (see text for details); middle panel: $V$ magnitude distribution of the HES sample from which we construct the halo MDF; lower panel: $(B-V)_0$ distribution.} \\end{figure}  Another wide-angle spectroscopic survey is the HES. It was originally conceived as a survey for bright quasars \\citep{Reimers:1990,hespaperI,hespaperIII}; however, its data quality is sufficient to not only efficiently select quasars with redshifts of up to $z = 3.2$, but also various types of stellar objects, including metal-poor stars \\citep{Christliebetal:2008}. So far, several hundred new stars at $\\mathrm{[Fe/H]}<-3.0$ have been identified, including three stars that were confirmed by high-resolution spectroscopy to have $\\mathrm{[Fe/H]} < -4.0$: {\\FAC} ($\\mathrm{[Fe/H]} = -5.4$; \\citealt{Frebeletal:2005,Aokietal:2006,Frebeletal:2006a}); {\\CBB} ($\\mathrm{[Fe/H]} = -5.3$; \\citealt{HE0107_Nature,HE0107_ApJ,Besselletal:2004}); and HE~0557$-$4840 ($\\mathrm{[Fe/H]} = -4.8$; \\citealt{Norrisetal:2007}). It is perhaps of interest that the HK survey has not (to date) yielded any stars with $\\mathrm{[Fe/H]} < -4.0$ confirmed by high-resolution spectroscopy; this may be related to the fact that the HK survey reaches apparent magnitudes that are brighter than the HES, and as a result is dominated more than the HES by inner-halo stars.  The Sloan Digital Sky Survey (SDSS; \\citealt{Gunnetal:1998}, \\citealt{Yorketal:2000}), and in particular the Sloan Extension for Galactic Understanding and Exploration (SEGUE), has provided even larger samples of halo stars, as discussed by \\citet{Carolloetal:2007} and \\citet{Ivezicetal:2008}. The former emphasize the division of the halo into two structural components, an inner region with $R < 10$--$15$\\,kpc, and an outer region beyond that radius. These two components differ in stellar metallicities, stellar orbits, and spatial density profiles. As we discuss in Sect.~\\ref{Sect:Selection} below, the HES sample is dominated by inner-halo stars. We note that we hereafter refer to the inner halo as ``the halo'', unless indicated otherwise.  In spite of the very large sample of $\\sim 20,000$ stars used by Carollo et al., their coverage of the regime of very low metallicity is limited. According to their supplemental Fig.~4, they find only 3 stars with $\\mbox{[Fe/H]} < -3.0$ in their ``local sample'' of 10,123 stars. The main reason for this is that the stars of their sample were not selected to be metal-poor, but for the purpose of spectrophotometric and telluric calibration of the SDSS spectra.  Recent high-resolution spectroscopic follow-up of stars from the Carollo et al. sample (W. Aoki, priv. comm.) has indicated that the current version of the SEGUE Stellar Parameter Pipeline (SSPP; see \\citealt{Leeetal:2008a,Leeetal:2008b,AllendePrietoetal:2008}) is somewhat conservative in the assignment of stellar metallicity estimates, in the sense that stars assigned $\\mathrm{[Fe/H]} < -2.7$ by the SSPP are in reality more metal-deficient, on average, by on the order of 0.3\\,dex. A recent examination of the numbers of stars from the SDSS/SEGUE survey, taking into account this offset, suggests that up to several hundred stars with $\\mathrm{[Fe/H]} < -3.0$ are in fact present in the current SDSS sample of stars (including other categories of targets than just the calibration stars).  \\citet{Ivezicetal:2008} focus on the comparison between the inner halo and the disk. Since they rely on abundances determined from photometry, they cannot reliably determine metallicities of stars at $\\mbox{[Fe/H]} < -2$. Nevertheless, the metallicity map of some 2.5 million stars with photometric metallicies shown in Fig.~8 of Ivezic et al. indicates that there exist very large numbers of stars in SDSS consistent with $\\mathrm{[Fe/H]} < -2.0$. Follow-up spectroscopy is, at present, only available for a subset of them. Beers et al. (in preparation) discuss the MDF of the lowest metallicity stars found in SDSS/SEGUE. The total number of stars with $\\mathrm{[Fe/H]} < -2.0$, based on medium-resolution SDSS spectroscopy, is over 25,000 (i.e., five times the number discovered by the combination of the HK and HES).   This paper continues our series on the stellar content of the HES (\\citealt{HESStarsI}, Paper~I; \\citealt{HESStarsII}, Paper~II; \\citealt{HESFHBA}, Paper~III; \\citealt{Christliebetal:2008}, Paper~IV). We are mainly concerned with the low-metallicity tail of the halo MDF, which is constructed from a sample of 1638 metal-poor stars selected in the HES by quantitative criteria (Sect.~\\ref{Sect:Selection}). The follow-up observations and determination of the metallicities are described in Sect.~\\ref{Sect:FollowUp}. In Sect.~\\ref{Sect:MDFconstruction} we detail how the MDF was constructed. We discuss the shape of the halo MDF in Sect. \\ref{Sect:MDFshape}. Comparisons of the observed MDF with MDFs predicted by models of Galactic chemical evolution are presented in Sect.~\\ref{Sect:TheoryObservations}, and a comparison with the MDFs of the Galactic globular cluster system and dwarf spheroidal galaxies is presented in Sect.~\\ref{Sect:GCdSph}. The results are discussed in Sect.~\\ref{Sect:Conclusions}.   \\begin{table}[htbp] \\centering \\caption{Number of stars in each candidate class in the total sample of candidates, number of observed candidates, and number of accepted candidates after removal of emission line objects, ``peculiar'' objects (e.g., objects with continuous spectra) and all stars with a G-band index $\\mathrm{GP} > 6$\\,{\\AA}. In the last column, we list the scaling factors applied to the [Fe/H] histograms for each candidate class during the construction of the MDF (see Section~\\ref{Sect:MDFconstruction}).} \\label{Tab:CandidateStatistics} \\begin{tabular}{lrrrr}\\hline\\hline \\rule{0ex}{2.3ex} & \\multicolumn{3}{c}{Number of stars} &  \\\\\\cline{2-4} \\rb{Class}        & \\multicolumn{1}{c}{All} & \\multicolumn{1}{c}{Observed} \\rule{0ex}{2.3ex} & \\multicolumn{1}{c}{Accepted} & \\rb{Factor}\\\\\\hline \\texttt{mpca} &  201 &  123 &   105 & \\rule{0ex}{2.3ex} 1.63\\\\ \\texttt{unid} &  231 &  208 &   192 &                   1.11\\\\ \\texttt{mpcb} & 2006 & 1008 &   940 &                   1.99\\\\ \\texttt{mpcc} & 1275 &  432 &   401 &                   2.95\\\\\\hline Sum           & 3713 & 1771 &  1638 & \\rule{0ex}{2.3ex}\\\\\\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{Sect:Conclusions}  In Sect.~\\ref{Sect:TheoryObservations} we have shown that a reasonable agreement with the overall shape of the HES MDF can be obtained for $\\mbox{[Fe/H]} > -3.6$ by most models of Galactic chemical evolution, but only the model of Salvadori et al. with $Z_{\\mathrm cr}=10^{-3.4}\\,Z_{\\odot}$ reproduces the the sharp drop at $\\mbox{[Fe/H]} \\sim -3.6$ seen in the HES MDF. The lack of stars at $\\mbox{[Fe/H]} < -3.6$ is highly significant: The models typically predict that about ten such stars should be present in the HES sample, while only two are found. The significance of this discrepancy is reflected in the low probabilities for the MDFs predicted by the models and the HES MDF having the same parent distribution, as determined by KS-tests. It remains to be investigated whether the drop can be reproduced by modifying some of the assumptions of the models, or by adding further ingredients.  The HES sample discussed in this paper contains no objects with $\\mathrm{[Fe/H]} < -4.2$, but considering the abundance analyses of three additional stars in this metallicity range published in the recent literature, it is obvious that it exists. However, a thorough and quantitative comparison with theoretical MDFs has to await larger statistically complete and unbiased samples which include more stars with $\\mathrm{[Fe/H]} < -4.0$. Such samples will become available through new, deeper surveys for metal-poor stars that will commence in the near future; in particular, the Southern Sky Survey \\citep{Kelleretal:2007} and a survey to be conducted with the Chinese 4\\,m Large sky Area Multi-Object fiber Spectroscopic Telescope (LAMOST; \\citealt{Zhaoetal:2006}).  In the $\\Lambda$CDM picture, the Galactic halo was largely built out of disrupted satellite galaxies. If stars had already formed within them at the time of accretion, then the MDF of the Galactic halo and of the existing dSph galaxies should agree at the metal-poor end with regard to the presence of a weak tail of stars with $\\mbox{[Fe/H]} < -3.0$. It is thus encouraging for the $\\Lambda$CDM scenario that our analysis shows better agreement between the halo MDF and that of the dSph galaxies than claimed by \\citet{Helmietal:2006}. However, even if this were not the case, it would not necessarily be a strong contradiction to the $\\Lambda$CDM scenario. According to the semi-analytical models of \\citet{Salvadorietal:2008} and \\citet{Salvadori/Ferrara:2009}, the MDFs of dSph galaxies can differ quite significantly from each other, depending on their individual enrichment histories. Hence their MDFs can also be different from that of the Galactic halo. An important question remaining to be answered is how the elemental-abundance ratios of the dSph stars at $\\mbox{[Fe/H]} < -3.0$ compare with those of the Galactic halo stars.  Since the HES and the HK survey are in-situ surveys that predominantly sample the inner-halo population of the Galaxy (with $R < 15$\\,kpc), it is mandatory to consider the possibility that the (for now, poorly studied) outer-halo population of the Galaxy may indeed contain significant numbers of stars with $\\mathrm{[Fe/H]} < -3.6$, as might be indicated by the shift of the peak metallicity of the other-halo stars studied by \\citet{Carolloetal:2007} to $\\mathrm{[Fe/H]} = -2.2$, a factor of four lower than the peak metallicity of inner-halo stars. This possibility is being actively pursued by high-resolution spectroscopic follow-up of stars that are likely to be members of the outer-halo population, based on their kinematics, by a number of groups."
    },
    "0809/0809.3749.txt": {
        "abstract": "The growth of Jovian mass planets during migration in their protoplanetary disks is one of the most important problems that needs to be solved in light of observations of the small orbital radii of exosolar planets.  Studies of the migration of planets in standard gas disk models routinely show that the migration speeds are too high to form Jovian planets, and that such migrating planetary cores generally plunge into their central stars in less than a million years. %Recently, \\cite{Alibert05} found that %planetary migration has to be significantly slowed for the formation %of Jovian planets. Here, we propose a possible mechanism which %allows Jovian planet formation without slowing down migration %artificially. In previous work, we have shown that a poorly ionized, less viscous region in a protoplanetary disk called a {\\it dead zone} slows down the migration of fixed-mass planets. %we have pointed out the importance of dead zones in protostellar disks in significantly %slowing the migration of planets of fixed mass. In this paper, we extend our numerical calculations to include dead zone evolution along with the disk, as well as planet formation via accretion of rocky and gaseous materials. %including an opacity reduction effect in an envelope of the protoplanet. %as the planets migrate through the disk. %a wide range of important evolutionary effects on planetary growth and %migration including the evolution of disks with time as they accrete onto their central stars, %the consequent shrinkage of disks dead zones with time, the accretion of %rocky materials of the protoplanetary cores in a slower oligarchic phase, and the accretion %of gas from the disks as the planets migrate through the disks. Using our symplectic-integrator-gas dynamics code, we find that dead zones, even in evolving disks wherein planets grow by accretion as they migrate, still play a fundamental role in saving planetary systems. We demonstrate %by means of time-dependent simulations using our symplectic-integrator-gas dynamics code that Jovian planets form within $2.5$ Myr for disks that are ten times more massive than a minimum mass solar nebula with an opacity reduction and without slowing down migration artificially. %as in previous studies. Our simulations indicate that protoplanetary disks with an initial mass comparable to the minimum mass solar nebula (MMSN) only produce Neptunian mass planets. We also find that planet migration does not help core accretion as much in the oligarchic planetesimal accretion scenario as it was expected in the runaway planetesimal accretion scenario. Therefore we expect that an opacity reduction (or some other mechanisms) is needed to solve the formation timescale problem even for migrating protoplanets, as long as we consider the oligarchic growth. %rather than helping it in the oligarchic planetesimal accretion scenario. %They are saved from splashing into their stars by the dead zones, which %shrink with time from an intial radial extent of 13 AU to less than 2AU over a couple of million years. %Thus, with all of these physical effect included, we find %that dead zones, even in evolving disks wherein planets grow by accretion as they migrate, %still play a fundament role in saving planetary systems. We also point out a possible role of a dead zone in explaining long-lived, strongly accreting gas disks. ",
        "introduction": "The properties of nearly 300 recently discovered extrasolar planetary systems reveal that Jovian mass planets are often found at a scaled orbital radius of Mercury around their central stars \\citep[e.g.][]{Udry07}. Since none of the current theories of Jovian planet formation can explain {\\it in situ} formation of gas giants at these small distances from their stars, it is generally agreed that such planets were formed in the outer regions of protoplanetary disks and migrated through them to their current positions \\citep{Lin96}. Growth of giant planets takes place during this passage and both processes are terminated when most of the gas in the protoplanetary disk is either accreted onto the central star, or dissipated by photoevaporation \\citep[e.g.][]{Shu93,Hollenbach94,Alexander06b}. Observations of infrared to submm emission and gas accretion rates onto the central stars reveal that disk lifetimes are typically $1-10$ Myr \\citep[e.g.][]{Hartmann98,Muzerolle00,Andrews05,SiciliaAguilar06}.  The observed disk life-times raise two difficulties in planet formation theory. The first regards planet migration: the migration time scales of planets arising from the tidal interaction between a planet and the gas disk are shorter than the disk lifetimes. Therefore, unless they are stopped by some robust mechanism, planets plunge into their central stars within about a million years. %Why are there any planets at all? The second is %the accretion problem, which is specific to core-accretion model for Jovian planet formation: the formation time scales may be longer than, or comparable to, the disk life times. Therefore, unless the formation time scale is reduced somehow, the existence of giant planets cannot be explained by the core accretion scenario.  The core-accretion model posits that gas giants result from a two-stage process - the first being the formation of their rocky cores by the repetitive coagulation and agglomeration of the smaller bodies onto the larger ones (as in the terrestrial planet formation), and the second being the accretion of their massive gaseous envelope from the surrounding disk. %The accretion problem is that this second stage takes an %uncomfortably long time - comparable to the disk lifetime. % %Studies of the accretion problem go back to the pioneering work of %\\citet[][hereafter P96]{Pollack96} who performed numerical %simulations of the growth of a protoplanet core with an initial mass %of \\(0.6 \\ M_E\\) by calculating the planetesimal and gas accretion %rates in a self-consistent, interactive manner. They estimated that %the {\\it in situ} formation of Jupiter takes the uncomfortably long %time of  \\(8\\times 10^6\\).  A significant body of research has %reduced this somewhat (see below), but this is still uncomfortable.  In the first stage, the formation of the rocky core proceeds as bodies grow from a micron to a planetary size. One of the difficulties occurs when particles grow up to dm sizes. At this point, the collisional agglomeration may not be a preferred path to make the larger bodies \\citep{Langkowski08}. %(Blum \\& Wurm \\citep{Blum05}). Even if the sticking is still efficient and the bodies can keep on growing, the gas drag becomes non-negligible for such objects. Since the gas disk rotates at the sub-Keplerian speed, slightly slower than these growing bodies, the bodies feel the head wind, lose the angular momentum, and eventually migrate into the central star. Migration induced by gas drag becomes most efficient for meter-sized bodies, and becomes negligible again for the km-sized bodies since they are large enough not to be affected by the gas drag \\citep{Weidenschilling77}.  The gravitational instability in the planetesimal disk was suggested to avoid this meter-size barrier \\citep{Goldreich73}.  However, such a mechanism may be hindered by the Kelvin-Helmholtz instability \\citep{Cuzzi93,Weidenschilling95} unless the local solid-to-gas ratio is sufficiently high \\citep[e.g.][]{Sekiya98,Youdin02}. Recently, it was demonstrated that very rapid planetesimal growth up to Ceres-mass objects is possible via a streaming instability in such turbulent disks \\citep[e.g.][]{Johansen07,Youdin07}. Therefore, there is good justification to skip these early stages, and to assume that there are already km-sized planetesimals as well as planetary embryo(s) which are embedded in a gas disk.  In this paper, we will follow this approach as the previous studies did.  %More problematic is the need to reduce the gas accretion timescale %onto the rocky planetary core. %Yet another difficulty occurs before a protoplanet achieves a %crossover mass, above which gas accretion rate increases in a runaway fashion. In the second stage, protoplanetary cores keep on accreting planetesimals while they start developing gaseous envelopes. A cornerstone work done by \\cite{Pollack96} (P96 hereafter) identified three phases in giant planet formation. The first phase is a rapid core building phase, in which a protoplanetary core of $0.6 M_E$ grows up to \\(\\sim 10 M_E\\) within half a million years.  The second phase is a slow gas and planetesimal accretion phase which lasts until the crossover mass (for which the envelope mass is comparable to the core mass) is reached. The third phase is a rapid gas and planetesimal accretion phase, in which the planet quickly becomes a gas giant.  P96 showed that planet formation spends most time in the second phase that could last for several million years.  Since this is uncomfortably close to a typical disk life time, it was considered as one of the weakest points of core accretion scenario \\citep[e.g.][]{Boss97}. To shorten the second phase, protoplanets could either increase the gas accretion rate itself, or increase the planetesimal accretion rate and expand the gas feeding zone. % %One approach is to more carefully consider the mechanisms that could \\cite{Ikoma00,Hubickyj05} (hereafter INE00, and HBL05 respectively) took the former approach, and considered a reduced opacity in planetary envelopes. %\\citep[][hereafter INE00, and HBL05 respectively]{Ikoma00,Hubickyj05}. HBL05 used an updated version of the code by P96 with a more recent opacity table, and showed that smaller opacity, which mimics the dust settling and coagulation in the protoplanet's atmosphere, decreases the effective gas pressure of the protoplanetary envelope, and hence leads to the faster gas accretion. They showed that the {\\it in situ} formation timescale of Jupiter could be as short as 1 Myr if the opacity is \\(2\\%\\) of the interstellar value and the core mass is \\(10 M_E\\). % %Another factor that could in principle reduce the gas timescale is %that gas accretion occurs during migration \\citep[][hereafter %A05]{Alibert05}. On the other hand, \\cite{Alibert05} (hereafter A05) took the latter approach, and considered planet formation during migration.  In P96, planetesimal accretion slows when the planetesimals get depleted in the feeding zone, because %before the crossover mass for rapid gas accretion is %reached. Then a planet has to accrete gas and expand its feeding zone to further accrete planetesimals. Since the gas accretion in Phase 2 tends to be slower than the planetesimal accretion in Phase 1, it takes longer to accrete a similar amount of gas envelope to the core, and achieve a crossover mass. A05 overcame this problem by including disk evolution and planet migration in the model, so that the planetesimal feeding zone is constantly replenished. The gas accretion time shortens as the planetary mass increases, because it proceeds on the Kelvin-Helmholtz timescale, which is a steep function of planetary mass. They showed that the Jupiter formation timescale could be of the order of 1 Myr if the planet migrates. %Since the gas accretion proceeds on the Kelvin-Helmholtz timescale, which is a steep %function in planetary mass, the gas accretion time shortens as the planetary mass increases. However, as noted by \\cite{Ward97}, solving the accretion problem by migration in a standard disk model is a ``double-edged sword\", since it comes at the price of possibly losing the planets to their central star. Thus, A05 %did not address this problem and simply assumed that some unspecified mechanism would slow down the fast type I (pre-gap opening) migration by factors of \\(10-100\\) times compared to the migration estimated in a 3D disk by \\cite{Tanaka02}. %slower than the migration rate estimated for a %3D disk by \\cite{Tanaka02}.   %The disk models that were employed in the discussion above all %assume that a protostellar disk is turbulent throughout its entire %radial extent. Here, we propose a possible mechanism which can slow planet migration, and investigate planet formation in that context. In a standard disk with a smooth surface mass density and a standard viscosity expected from the MRI turbulence, both type I and II migration time scales are shorter than the disks' life times. %Disk accretion onto the central star is driven by viscosity in a %differentially rotating disk. The most popular source of such a viscosity is the magneto-rotational instability \\citep[MRI,][]{Balbus91}. However, it has been demonstrated by several groups that protoplanetary disks are not MRI active everywhere, but harbor extended regions known as the dead zones \\citep{Gammie96} where there is virtually no turbulence %%%%%%%%%%%%% Foot Note \\footnote{This arises in principal because operation of the  MRI requires good coupling between the gaseous disk and the magnetic field.  However, the ionization fraction in the inner regions of dense protoplanetary disks is typically so low that Ohmic diffusion prevents the growth of the MRI in such regions. The size of the dead zone can be computed once the source of disk ionization (e.g. stellar X-rays, cosmic rays) is known.  Assuming that the MRI is the dominant source of the disk's ``turbulent'' viscosity, the dead zone is nearly inviscid. The existence of a dead zone in a protoplanetary disk was first proposed by \\cite{Gammie96} who showed that there is a magnetically dead zone in the inner disk where cosmic rays cannot penetrate (assuming the cosmic ray stopping density is \\(98 \\ {\\rm g \\ cm^{-2}}\\)), and that such a region is sandwiched by the magnetically active surface layers where gas accretes toward the star efficiently.}. %%%%%%%%%%%%%% Foot Note END Many authors have studied the physical extent of the dead zones %has been studied by many authors by using different disk models and taking account of a variety of the ionization/recombination processes \\citep[e.g.][hereafter MP03, and MP06 respectively]{Glassgold97,Sano00,Fromang02,Semenov04,Matsumura03,Matsumura06}. These models suggest roughly the same size of the dead zones (from \\(<1\\) AU to \\(10-20\\) AU), indicating that a dead zone is a robust feature of a protoplanetary disk. Since this is a critical region of a disk both for planet formation and migration, it is of vital importance to investigate these problems in the context of protoplanetary disks with dead zones. % Recent numerical simulations of such a layered disk have shown that the expected viscosity parameter \\(\\alpha\\) is about \\(10^{-2}\\) in the active zone and \\(10^{-4}-10^{-5}\\) in the dead zone \\citep{Fleming03}. Therefore, if the MRI is the dominant source of the disk viscosity, a significant decrease in viscosity is expected within the dead zone, which most likely affects the rate of planet migration. Moreover, recent observations have revealed that disk's viscosity parameters may take a larger range than previously expected; \\(\\alpha=10^{-6}-10^{-1}\\) \\citep{Hueso05,Luhman07}, indicating that at least some observed disks may possess a low viscosity region like a dead zone.  %Of central importance for planet formation and migration is the fact %that dead zones can significantly slow the rate of migration of %planets. In an earlier paper, we have demonstrated that dead zones can significantly slow the rate of migraion of planets, and save planetary systems from plunging into the central stars \\citep[][hereafter MPT07]{Matsumura07}.   In particular, we found (1) type II migration (post-gap-opening migration) is slowed in the dead zone due to the low viscosity there, (2) even low-mass planets (\\(\\leq 10 M_E\\)), which are usually type I migrators (i.e. non gap-openers), may open a gap in the dead zone if the thermal condition is satisfied, and thus migrate slower there, and (3) type I migrators moving toward the dead zone can be stopped at the outer edge of the dead zone due to the jump in mass density there %at the outer edge of the dead zone, %increase inside it, which is a result of the slower advection speed inside the dead zone with respect to outside it. %gas accretion inside the dead zone %with respect to outside it.  %This is because one of the difficulties of planet formation is that the %migration timescale is much shorter compared to the formation %timescale, and thus growing planets are prone to being lost into the %central stars.  In this paper, we present a rather complete treatment of the formation and migration of Jovian planets within their evolving protoplanetary disks. %The mass evolution of the disk is a %consequence of viscous accretion which enables the transport of most %of the disk material into the central star in roughly 10 million %years. %(we end with a Jovian mass disk typically). We compute comprehensive time-dependent models for the growth of Jovian planets in the core accretion picture in viscously evolving gaseous disks.  Our calculations include updates on both stages of accretion (planestimals and gas), as well as the ability to follow planetary migration down to 0.1 AU. %for up to 10 million years. We approach this problem by means of time-dependent simulations that track both planetary accretion and migration through disks with evolving dead zones. %An important aspect of dead zone evolution that %should be taken into account is its There are two major effects of disk evolution in our models: (i) the gradual accretion of mass from the disk onto the star which limits the reservoir that is available to the growing Jovian planet; and (ii) the shrinkage of the size of the dead zone as a consequence of mass accretion onto the star (mainly) through the well coupled surface layers of the disk. This process reduces the size of the dead zone substantially with time --- perhaps leaving it only a few AU in extent after a million years. % %Our major findings are (1) when a planetary core starts its %migration and growth from outside the dead zone, the core is kept %outside the dead zone due to a surface mass density jump and grows %up to a Jupiter mass, and (2) when a planetary core starts from %inside the dead zone, the core is left outside the dead zone which %shrinks due to mass accretion through surface layers, and grows up %to a Jupiter mass. Jovian planets in both of these circumstances are %accreted within \\(7\\times 10^6\\) years.  We require models that are %initially an order of magnitude greater in column density than the %minimum mass solar nebula. % %, and (3) when a %planetary core starts well inside the dead zone, the core grows %inside the dead zone, opens a gap, and stops its gas accretion.  %the planet migration and formation by taking account of the dead zone. %--- poorly ionized region in a protostellar disk which %has no magneto-rotational instability (MRI) turbulence %\\citep{Gammie96}, and hence expected to be nearly inviscid. We first introduce our disk models and numerical methods in \\S 2. We also highlight a possible role of dead zones in disk accretion onto the central stars. Then we study planet formation in an inviscid disk, and compare our results with P96, and HBL05 (\\S 3). We also check the effect of the opacity, and compare the results with HBL05, which improved the work of P96 and further investigated the opacity effects. We then go on to compute planet formation and migration in an evolving disk, and compare the results with A05 in \\S 4. The culmination of our work is presented in \\S5 where we generalize our results and present planet formation in an evolving disk with a dead zone. Finally, we summarize our work in \\S 6. % %--- Section 2 ------------------------------------------------------------------------------ % ",
        "conclusions": "We have presented some new results on planet formation and migration in evolving disks with dead zones.  The most significant is that dead zones provide a natural way of saving planetary systems even as the planets migrate through disks whose properties change significantly over hundreds of thousands to millions of years. The dissipation of the disk does place interesting constraints on their masses - we find that only Neptunian mass planets can be formed in the MMSN-mass disk models.  Jovian planet formation requires more massive disks, which has also been suggested by other groups (e.g. P96 and HBL05), and was recently demonstrated by multiple planet formation simulations by \\cite{Thommes08}. The time scale that we find for Jovian planet formation in these more massive disks - about 2.5 Myr - is within observational limits of disk lifetimes. While not drastically reduced from previous estimates, our Jovian planet formation time incorporates many new aspects including migration in the presence of dead zones, effects of slower oligarchic growth that are more realistic than earlier accretion scenarios, and an effect of a reduced opacity.  Dead zones turn out to play a potentially important role as blockades against inward planetary motion.  Our planetary cores are generally kept outside of the dead zone - and therefore immersed in the region where they can ``feed'' more effectively on surrounding gas.  When the planets become massive enough to open a gap, the outwards directed torque associated with the density gradient of the outer edge of the dead zone weakens, and the planets migrate into the dead zone. Planets migrate slowly inside the dead zone due to the lower viscosity there compared to outside it.  The methods we used were straightforward to implement in our planetary/disk evolution code. In order to model dead zone evolution, we simply assumed that the dead zone is where the surface mass density is above the critical value (\\(\\Sigma>\\Sigma_{crit}=21\\), and \\(80 \\ {\\rm g \\ cm^{-2}}\\) for \\(\\Sigma_0=10^3\\) and \\(10^4 \\ {\\rm g \\ cm^{-2}}\\) respectively) that was determined from a stationary disk model in MP06.  Also, we have included gas accretion through the surface layers, which contributes to make a dead zone shrink very rapidly (see Fig. \\ref{fig3}). Combining these two conditions, we have found that the dead zone radius shrinks from \\(8.2\\) AU to \\(2\\) AU within 2 Myr for \\(\\Sigma_0=10^3 \\ {\\rm g \\ cm^{-2}}\\).  For the planet formation part of the code, we follow the approach by P96, but with a few differences.  First of all, we assume that the planetesimal accretion stage is mainly carried out by oligarchic growth, rather than by rapid runaway accretion used by P96. This is a reasonable assumption since the runaway growth most likely ceases long before the planetary mass reaches a few times \\(10^{-5} M_E\\) \\citep{Thommes03}, which is much smaller than our initial planetary mass (\\(0.6 M_E\\)), and switches to a slower oligarchic growth \\citep{Ida93,Kokubo98}. Depending on the size distribution of planetesimals, the rate of accretion could be significantly different. %Since we built our formation model based on the %results obtained by P96, this difference between planetesimal %accretion rates s that we choose a rather high surface mass density %for planetesimals (\\(\\Sigma_{solid}=1000(r/AU)^{-2}\\)). Secondly, we parameterize our gas accretion rate following INE00 and \\cite{DAngelo03}. %so that we can reproduce the results by P96 (see \\S 2.2 as well). Therefore, we are not calculating planetesimal and gas accretion rates in an interactive way as in P96, HBL05, or A05. However, our simulations show reasonable agreements with their results (see \\S 3). Also, since we take account of the subdisk accretion phase, the final stage of gas accretion slows down, rather than exponentially increases as in P96, or artificially cuts off as in HBL05. Thirdly, we include a 1D gas disk evolution and planet migration.  Our approach is similar to A05, but we don't include photoevaporation effects for this study, since this is likely to be negligible during planet formation. %dead zone helps a natural, rapid disk dispersal. %In other words, we cut off the planetesimal accretion artificially, and turn %on the gas accretion when the core mass reaches \\(10 M_E\\). %In P96, the core accretion slows down automatically due to the %depletion of planetesimals in the feeding zone. %Thirdly, the transition from the slow gas accretion to rapid gas %accretion is set by the condition: \\(r_{\\rm Bondi}\\sim h \\sim r_{\\rm %Hill}\\), while this transition naturally occurs in P96 as the core %mass becomes comparable to the gas envelope's mass.  This assumption %is reasonable, because beyond this point, all gas within the sphere %with a radius equal to the pressure scale height is gravitationally %dominated by the planet rather than the central star (\\(r_{\\rm %Hill}>h\\)), and the gas is bound to the planet rather than moving %freely (\\(r_{\\rm Bondi}>h\\), i.e. the gravitational energy is larger %than the thermal energy). Therefore, the mass accretion occurs %nonlinearly, and very quickly. The typical crossover mass that can %be estimated from the Hill condition is \\(\\sim 40 M_E\\). % %{\\bf We should think how the density jump changes when it becomes %Rayleigh unstable. Also corotation??} %  There are several aspects of planet formation that we have not included explicitly in this study. We list several below, but note that we do not expect their absence to strongly affect our results.  First of all, our simple torque prescription does not capture the nature of a turbulent disk properly. A number of numerical simulations have shown that turbulent fluctuations can cause torque fluctuations \\citep[e.g.][]{Laughlin04,Nelson04}, which leads to stochastic type I migration \\citep{Johnson06}. This results in even {\\it faster} migration for most planets, but also allows {\\it outward} migration for some of them \\citep{Johnson06}. The inclusion of such an effect may allow a planet outside the dead zone to ``jump the barrier'' provided by the edge of a dead zone. Even within a dead zone, turbulent fluctuations in upper and lower active layers may be able to affect planet migration significantly.  Also, we do not take account of the evolution of temperature profile of a gas disk. To better evaluate the disk evolution, we should include the radiative transfer to the code. This would affect not only the dead zone and disk evolution, but also planet migration. A recent work of \\cite{Paardekooper08} demonstrated that, in a non-isothermal disk, planet migration is preferentially outward at around 5 AU until the disk mass decreases significantly. Therefore, more precise treatment of a disk temperature may save type-I migrators effectively, even inside a dead zone.  Perhaps the most important simplification of our model is the 1D treatment of the disk, and the rather crude prescriptions for the dead zone, which led to a very sharp density transition at the outer dead zone radius (see Fig. \\ref{fig2}). This feature enabled a planet to stop its migration by balancing inner and outer torques.  However, such a sharp density gradient makes a disk locally Rayleigh unstable (i.e. an epicyclic frequency $\\kappa^2<0$), which may render the sharp transition from the dead zone to the outer active disk much smoother. Even when the disk is still locally stable, the Rossby wave instability can smooth the density gradient in a similar manner when pressure varies significantly over a few times the disk thickness \\citep[][and the references therein]{Li01}. This possibility should be checked carefully by numerical simulations.  We did not take account of the effect of the corotation torque either. This is a safe assumption most of the time, since the corotation torque depends on the gradient of the inverse of the specific vorticity ($\\partial (\\Sigma /B)/ \\partial r$) \\citep{Goldreich80,Masset01}, and therefore its effect becomes particularly important when the surface mass density changes sharply, e.g. at the edge of a dead zone. However, we do not expect this affects our results significantly, because inner and outer Lindblad torques balance with each other far away from the density jump, where the corotation torque is likely negligible, and planet migration stalls (also see Appendix B in MPT07). When Rayleigh, or Rossby wave instability makes the density gradient at the outer edge of a dead zone smoother, the corotation torque can become important. In such a case, a planet would not be stopped at the outer edge of a dead zone, but pulled into it instead.  The corotation torque is also found to be important for migration of sub-Jovian mass planets like Saturn \\citep{Masset03}. % affect planet migration for sub-Jovian mass planets like Saturn \\citep{Masset03}. In a standard disk, the time scale of the so-called type III migration for such planets is comparable to, or even shorter than, type I migration \\citep{Masset03}. Although this is most relevant to our RunH1 (Fig. \\ref{fig15}), we don't expect a significant change in our result. This is because type III migration is likely to switch to slower type II migration as soon as the planet enters a dead zone, where planets tend to open a wider gap at smaller mass due to its low viscosity.  Another simplifications is that we have only considered a protoplanetary disk with a single core, while the real disks are expected to have multiple cores.  Studies of this kind have been done by \\cite{Chambers06,Thommes08} for disks with no dead zone, and by \\cite{Morbidelli08} at the {\\it inner} edges of disks with dead zones. A recent paper by \\cite{Ida08ap}, took a similar approach to us, and studied a dead zone's effects on retaining icy grains, as well as protoplanetary cores. They reproduced the observed frequency and mass-period distribution of gas giants around solar-type stars with a moderate reduction in type I migration speed. %{\\bf Write a bit more...} %It would be interesting to study such an %interaction, and its effect on planet formation and migration.  Yet another simplification is that all of our planets in this study are on circular orbits. Planet-disk interaction has been proposed as a method to drive planetary eccentricity \\citep{Goldreich03}, and recent numerical simulations confirmed this for massive ($>M_J$) planets \\citep[e.g.][]{Masset04,DAngelo06}. %the overall effect may be dominated by eccentricity damping %Our results would be modified if we include the eccentricity effect. %This is also interesting in relation to the former point, because %the eccentricity could be enhanced either via planet-planet %interaction \\citep[e.g.][]{Rasio96}, or via planet-disk interaction %\\citep[e.g.][]{Goldreich03}. If a planetary eccentricity is enhanced significantly (\\(e\\gg h/r\\)), then planet migration could be slowed down \\citep[e.g.][]{Papaloizou00,Papaloizou02}, and mass growth rate could increase \\citep[e.g.][]{DAngelo06,Kley06}. % %Another point which has not been included explicitly in our study is %the variety of planetesimal sizes.  We have assumed a single size of %\\(10\\) km for planetesimals, and as a result, used a rather heavy %disk model (\\(\\Sigma_{solid}=1000(r/AU)^{-2}\\)). However, %\\cite{Rafikov04} suggested that the protoplanetary cores could be %assembled quickly by considering smaller planetesimals (\\(\\leq %0.1-1\\) km), because their random velocities excited by the cores %are damped by gas drag.  If about \\(\\sim 10 \\%\\) of the solid mass %is in these small planetesimals, his estimate shows that the %protoplanetary core of \\(10 M_E\\) can be assembled within roughly %\\(10^5-10^6\\) years at \\(1-10\\) AU for a minimum mass solar nebula %model. % %We will leave these investigations to future work.  We also employ simplified planetesimal disks.  First, we only consider single-size planetesimals ranging from 100 m to 100 km, while the real disks should have multiple-size planetesimals. Secondly, we don't include the planetesimals' effect on planet migration, which could potentially be important \\citep{Murray98}. Thirdly, we assume the dispersion-dominated random velocities for planetesimals with all sizes.  Since small planetesimals experience strong gas drag and achieve reduced random velocities, their mass accretion rate should be treated in the shear-dominated regime \\citep{Rafikov04,Chambers06}.  This may be important for 100 m-size planetesimals \\citep{Chambers06}, and the planetesimal accretion would be runaway, rather than oligarchic in such a case.  Finally, we don't include the effects of photoevaporation \\citep[e.g.][]{Shu93,Hollenbach94}, which is likely to be important during the last stage of disk evolution \\citep{Clarke01,Alexander06b}. However, photoevaporation becomes important only when disk mass becomes significantly low (\\(\\sim 1 M_J\\) or so), and therefore it is unlikely to affect our results.  We conclude by listing our major findings: \\\\  (1) Dead zones strongly evolve over the duration of the disk, starting with outer radii of order $10-15$ AU, and shrinking with time to of order an AU or less.  \\\\  (2) Protoplanets which are left outside the dead zone tend to migrate inward as they accrete planetesimals, and stop just outside the outer dead zone radius due to a steep surface mass density jump. %and grow there. Then such protoplanets migrate at the viscous time scale of the shrinking dead zone until they achieve gap-opening mass.  Once they open a gap, they enter the dead zone, and migrate inwards rather slowly due to the low viscosity there. \\\\ %the gap-opening mass is reached. \\\\  (3) The final mass of a planet is determined by disk mass.  In a minimum-mass solar nebula disk, it is difficult to get a planet more massive than Neptune.  We find Jovian planets form in disks that are ten times more massive than this. \\\\ %, which has implications for submillimeter surveys for disk masses. \\\\  (4) A Jupiter mass planet can form within \\(\\sim 2.5\\) Myr (see Fig. \\ref{fig15}) in a disk with a dead zone, by assuming standard type I migration. However, since planet migration does not help core accretion as much in oligarchic growth as in runaway growth scenario, we expect that the opacity reduction (or some other mechanism) is necessary to form a Jovian planet within a disk life time. \\\\  (5) Dead zones may help explain the existence of long-lived strong accretors. When a dead zone is present, the mass accretion rate is likely to take a nearly constant value until the dead zone disappears, rather than decreasing gradually as expected in a standard disk without a dead zone. Once this happens, the rest of the disk is dispersed on a few Myr time scale, unless photoevaporation and/or planet formation provide additional sources of disk dissipation. \\\\ % % %(2) Protoplanets which stay inside the dead zone tend to be left %outside the dead zone as it shrinks. These planets follow a similar %fate to the above case %(see Fig. \\ref{test924}). \\\\ % %(3) Protoplanets which start migrating far from the star do not %accrete enough planetesimals, and therefore do not form massive %planets (see Fig. \\ref{fig13}). \\\\ % %(4) We did not see the effect of decreasing opacity on planet %formation timescale. \\\\ % %(5) By mimicking the 3D effect of the torque, the migration slows %down as expected from \\cite[e.g.][]{Tanaka02,Menou04}.  Including %this effect, we see that... %"
    },
    "0809/0809.0338_arXiv.txt": {
        "abstract": "The discrete-dipole approximation (DDA) is a powerful method for calculating absorption and scattering by targets that have sizes smaller than or comparable to the wavelength of the incident radiation. The DDA can be extended to targets that are singly- or doubly-periodic. We generalize the scattering amplitude matrix and the $4\\times4$ Mueller matrix to describe scattering by singly- and doubly-periodic targets, and show how these matrices can be calculated using the DDA. The accuracy of DDA calculations using the open-source code DDSCAT is demonstrated by comparison to exact results for infinite cylinders and infinite slabs. A method for using the DDA solution to obtain fields within and near the target is presented, with results shown for infinite slabs. ",
        "introduction": "Electromagnetic scattering is used to study isolated particles, but increasingly to characterize extended targets ranging from nanostructure arrays in laboratories, to planetary and asteroidal regoliths. To model the absorption and scattering, Maxwell's equations must be solved for the target geometry.  For scattering by isolated particles with complex geometry, a number of different theoretical approaches have been used, including the discrete dipole approximation (DDA) \\cite{Purcell+Pennypacker_1973, Draine_1988, Draine+Flatau_1994, Draine_2000a}, also known as the coupled dipole approximation or coupled dipole method. The DDA can treat inhomogeneous targets and anisotropic materials, and has been extended to treat targets near substrates \\cite{Schmehl+Nebeker+Hirleman_1997,Paulus+Martin_2001}. Other techniques have also been employed, including the finite difference time domain (FDTD) method \\cite{Yang+Liou_2000,Taflove+Hagness_2005}.   For illumination by monochromatic plane waves, the DDA can be extended to targets that are spatially periodic and (formally) infinite in extent. This could apply, for example, to a periodic array of nanostructures in a laboratory setting, or it might be used to approximate a regolith by a periodic array of ``target unit cells'' with complex structure within each unit cell.  Generalization of the DDA (= coupled dipole method) to periodic structures was first presented by Markel \\cite{Markel_1993} for a 1-dimensional chain of dipoles, and more generally by Chaumet et al.\\ \\cite{Chaumet+Rahmani+Bryant_2003}, who calculated the electric field near a 2-dimensional array of parallelepipeds illuminated by a plane wave. Chaumet and Sentenac \\cite{Chaumet+Sentenac_2005} further extended the DDA to treat periodic structures with a finite number of defects.  From a computational standpoint, solving Maxwell's equations for periodic structures is only slightly more difficult than calculating the scattering properties of a single ``unit cell'' from the structure. Since it is now feasible to treat targets with $N\\gtsim 10^6$ dipoles (target volume $\\gtsim 200 \\lambda^3$, where $\\lambda$ is the wavelength), it becomes possible to treat extended objects with complex substructure.  The objective of the present paper is to present the theory of the DDA applied to scattering and absorption by structures that are periodic in one or two spatial dimensions. We also generalize the standard formalism for describing the far-field scattering properties of finite targets (the $2\\times2$ scattering amplitude matrix, and $4\\times4$ Mueller matrix) to describe scattering by periodic targets. We show how to calculate the Mueller matrix to describe scattering of arbitrarily-polarized radiation. The theoretical treatment developed here has been implemented in the open-source code DDSCAT 7 (see Appendix A).  The theory of the DDA for periodic targets is reviewed in \\S\\ref{sec:dda for periodic targets}, and the formalism for describing the far-field scattering properties of periodic targets is presented in \\S\\ref{sec:in the radiation zone}-\\ref{sec:mueller matrix}. Transmission and reflection coefficients for targets that are periodic in two dimensions are obtained in \\S\\ref{sec:T and R}.  The applicability and accuracy of the DDA method are discussed in \\S\\S\\ref{sec:cylinder},\\ref{sec:slab}, where we show scattering properties calculated using DDSCAT 7 for two geometries for which exact solutions are available for comparison: (1) an infinite cylinder and (2) an infinite slab of finite thickness. The numerical comparisons demonstrate that, for given $\\lambda$, the DDA converges to the exact solution as the interdipole spacing $d\\rightarrow0$. ",
        "conclusions": "The principal results of this study are as follows: \\begin{enumerate} \\item The DDA is generalized to treat targets that are periodic in one or two spatial dimensions.  Scattering and absorption of monochromatic plane waves can be calculated using algorithms that parallel those used for finite targets. \\item A general formalism is presented for description of far-field scattering by targets that are periodic in one or two dimensions using scattering amplitude matrices and Mueller matrices that are similar in form to those for finite targets. \\item The accuracy of the DDA for periodic targets is tested for two examples: infinite cylinders and infinite slabs.  The DDA, as implemented in DDSCAT 7, is accurate provided the validity criterion $|m|kd \\ltsim 0.5$ is satisfied. \\item We show how the DDA solution can be used to evaluate $\\bE$ and $\\bB$ within and near the target, with calculations for an infinite slab used to illustrate the accuracy of near-field calculations. \\end{enumerate}"
    },
    "0809/0809.0612_arXiv.txt": {
        "abstract": "A polarized gamma ray emission spread over a sufficiently wide energy band from a strongly magnetized astrophysical object like gamma ray bursts (GRBs) offers an opportunity to test the hypothesis of axion like particles (ALPs). Based on evidences of polarized gamma ray emission detected in several gamma ray bursts we estimated the level of ALPs induced dichroism, which could take place in the magnetized fireball environment of a GRB. This allows to estimate the sensitivity of polarization measurements of GRBs to the ALP-photon coupling. This sensitivity $\\gag\\le 2.2\\cdot 10^{-11}\\ {\\rm GeV^{-1}}$ calculated for the ALP mass $m_a=10^{-3}~{\\rm eV}$ and MeV energy spread of gamma ray emission is competitive with the sensitivity of CAST and becomes even stronger for lower ALPs masses. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0917.txt": {
        "abstract": "Highly supercritical accretion discs are probable sources of dense optically thick axisymmetric winds. We introduce a new approach based on diffusion approximation radiative transfer in a funnel geometry and obtain an analytical solution for the energy density distribution inside the wind assuming that all the mass, momentum and energy are injected well inside the spherization radius. This allows to derive the spectrum of emergent emission for various inclination angles. We show that self-irradiation effects play an important role altering the temperature of the outcoming radiation by about 20\\% and the apparent X-ray luminosity by a factor of $2\\div 3$. The model has been successfully applied to two ULXs. %The best-fit parameters of the model applied to two ULXs may be briefly %  characterized as moderate half-opening angles $\\theta_f \\sim 10\\div %  20^\\circ$ and % moderate ejection rates $\\mdot \\sim 10^2$ and small %  half-opening angles $\\theta_f \\sim 5\\div 10^\\circ$. %These values are %  however likely to be affected by mild comptonization effects making %  the actual spectrum harder. The basic properties of the high ionization HII-regions found around some ULXs are also easily reproduced in our assumptions. %However, the ejection rates derived from the %  observed spectra may be underestimated due to comptonisation effects %  that may be introduced as hardening factors similar to those used to %  describe the spectra of accretion discs in X-ray binaries. ",
        "introduction": "\\label{sec:intro}  Processes of gas accretion onto compact objects are studied since 60s when first X-ray sources were discovered \\citep{accretion_power}. Among several works describing the details of disc accretion in binary systems \\citet{ss73} was the most successfull. Their ``standard disc'' is still widely used to describe the thermal component of X-ray spectra of X-ray binaries. Standard disc model was worked out in several considerations such as high optical thickness, low geometrical thickness of the disc, etc. Among these assumptions one was that the power released in the accretion process does not exceed the Eddington luminosity. % and does not                 %% <- changed in response to comment #1 %therefore influence the dynamics of the  flow. This case is usually called subcritical accretion.  However, considering various phenomena such as growth of supermassive black holes and nova outbursts leads to the problem of supercritical accretion. It has been extensively investigated since 1980s when Abramowicz proposed it as a power source for active galactic nuclei \\citep{agn_tori}. Presently super-Eddington accretion is often applied to explain the observational properties of Ultraluminous X-ray Sources (ULXs,  see \\citet{roberts_review} for observational review) that are believed to be high-mass X-ray binaries with black holes accreting on thermal time scale.  % We do not discuss here existing alternatives to the supercritical % accretor (SA) hypothesis such as IMBHs \\citep{comil}, young supernovae % \\citep{fabbiano88} etc.  The first, outflow-dominated model of supercritical accretion, has been developed by Shakura and Sunyaev in the paper mentioned above. Authors assumed that all the inflowing gas above the critical accretion rate is being ejected from the disc in the form of a wind. Another version of supercritical regime is based on the relaxation of the locality condition in ``standard model'' which leads to  advective slim discs \\citep{agn_tori} or Polish doughnuts \\citep{abram78,kozl78,abram2004}. In reality, both processes % (as well as number of other, like the porosity of the accreting matter % \\citep{shaviv_porosity}, which effectively decrease the extinction % cross-section) work simultaneously.  A comprehensive model taking into account both for advection and ejection was recently developed by \\citet{poutanen}. Authors consider the structure of the disc in radial direction and estimate three characteristic temperatures relevant for the outcoming radiation (inner disc temperature, temperature at spherization radius and effective temperature of the wind photosphere) but do not study the processes of radiation transfer in the wind and do not calculate the outcoming spectra. %However, the authors concentrate on the spectrum produced by the accretion %disc only, neglecting the outflowing matter, which may be optically thick and %thus obscure disc emission.  Obviously, at high accretion rates the observational properties of the accretion flow are governed mostly by radiative transfer in the outflowing wind. %% In response to comment #3: The observational appearance of the pseudo-photosphere of the optically-thick wind of a supercritical accretion disc was considered by \\citet{nishi07}. The authors however assumed the optically-thick part of the wind fully adiabatic not taking into account radiative energy transfer that is expected to be important in the flow (see section \\ref{sec:diff}). %%  % Interplay between outflows and photon trapping is an unclear point in % the theory of SA. Currently, we have at least one example of a persistent supercritical accretor in our Galaxy -- the SS433 system, where most of accreting gas is being lost in a wind. %One should take into account that the local condition of %supercriticity means that the radiation field is strong enough to %develop a strong wind. The example of high-luminosity massive stars %possessing strong winds suggests that outflows must be very important %in the dynamics of supercritical accretion discs. Numerical simulations of this system partially support the outflow-dominated scenario. However, \\citet{okuda} states that his thorough 2D simulation fails to reproduce the outflow rate and jet collimation in SS433. In \\citet{ohsuga2005} a supercritical accretor accreting at $ 10^3$ Eddington rate appears a bright (about $10^{38}\\ergl$) hard X-ray source if viewed edge-on. \\citet{ohsuga2005} and \\citet{heinz} calculate the structure and emergent spectra considering only the inner parts of the flow ($R \\le 500 R_G$) not taking into account that the region considered is coated by accreting and outflowing matter both being optically thick. The outcoming spectra will strikingly differ from those calculated by \\citet{heinz}. %ejects about 10\\% %of the infalling material, which is much lower than observed for %SS433 \\citep{sklov81,blundell2001}, where the mass ejection rate is %roughly equal to the mass transfer rate. %This is the indication of some important effects in the dynamics of %super-critical accretion not being taken into account by %\\citet{ohsuga2005}. % See also discussion in section ~\\ref{sec:trep}  % For today SS433 remains the only proven example of a % persistent supercritical accretor (SA). %\\citet{katz86} supposed that a SS433-like object seen at low inclination angles %could exhibit properties similar to that of ULXs.  Optical and radio \\citep{blundell2001} observations allow to measure the mass ejection rate in the relatively slow ($1000\\div 2000\\,\\kms$) accretion disc wind seen in optical emission and absorption lines \\citep{ss2004} as well as the mass loss rate in the collimated mildly-relativistic jets launched along the disc axis. Infrared excess was used by \\citet{sklov81} and \\citet{vdh81} to estimate the mass ejection rate from SS433. No direct mass transfer rate measurements were ever made though it is usually supposed that it could not be much higher than the mass ejection rate $\\dot{M}_{ej} \\sim 10^{-4}\\, \\Msun\\, \\yr^{-1}$ \\citep{ss2004}.  The details of the processes  in the supercritical disc itself are unclear mainly due to the strong absorption in the wind. In fact even the photosphere of the wind is difficult to study due to high interstellar absorption (about $8\\magdot{\\,}$) making it impossible with the contemporary instrumentation to trace the Spectral Energy Distribution (SED) in the far ultraviolet (UV) spectral range were most of the radiation is expected to be emitted. The only and yet crude estimate of the blackbody temperature %% estimation? %by means of UV observations with {\\it HST} High-Speed Photometer %corresponds to the outer photosphere of the thick disc or the %pseudo-photosphere %% pseudo!! of SS433 was made by \\citet{dolan} and probably corresponds to the photosphere of the wind. Their measurements are consistent with a $(2\\div7) \\times 10^4\\,\\rm K$ blackbody source.  Both observations of SS433 \\citep{ss2004} and numerical simulations \\citep{okuda,ohsuga2005} support the conception that is traditionally called ``supercritical funnel''. It assumes that nearly all the matter accreted is being ejected in a form of a slow, roughly virial at spherization radius, dense wind. Due to centrifugal barrier two conical avoidance sectors with half-opening angles of $\\theta_f \\sim 30^{\\circ}$ \\citep{EGK} are formed along the accretion disc symmetry axis, filled with rarefied hot gas, which may be accelerated and collimated to form relativistic jets. In the inner parts of the wind this gas also can form a pseudo-photosphere (``funnel bottom'') at $R_{in} \\sim 10^9 \\rm cm$  \\citep{ss2004}. This radius is calculated in the consideration all the material ejected in the relativistic jets is uniformly distributed with respect to the polar angle. In case of any inhomogeneity of the jet material the inner radius becomes lower. %% in response to comment #6 The visible part of the wind may be therefore divided into three parts: funnel wall photosphere (or ``photocone''), the outer photosphere and the inner photosphere inside the funnel. %%  Another class of objects supposed to be supercritical accretors are extragalactic Ultraluminous X-ray sources (ULXs). \\citet{katz86} supposed that an object similar to SS433 seen at low inclination angles can appear a bright X-ray source with super-Eddington apparent luminosity. ULXs were discovered about that time by {\\it Einstein} (see \\citet{fabbiano88} and references therein). Though the nature of these sources remains unclear they are good candidates for the objects predicted by Katz \\citep{poutanen}.  The question about how representative is SS433 among the binary systems in the supercritical accretion regime in the observed Universe is difficult to answer. If one considers mass transfer in thermal timescale in a black hole high mass X-ray binary, the most relevant is the mass of the secondary that determines the timescale and hence the scaling for the mass accretion rate. It may be shown that in assumption the radius of the secondary scales with its mass as $R \\propto M^{1/2}$ and the black hole mass is close to $10~M_\\odot$ dimensionless thermal-timescale mass transfer rate depends on the secondary mass as:  $$ \\mdot \\simeq \\left(\\frac{M_2}{M_\\odot}\\right)^{5/2}. $$ %\\noindent  For the case of SS433 ($M_2 \\simeq 20\\,M_\\odot$) our estimate predicts $\\mdot \\sim 2000$, similar to the observed value \\footnote{Acutally, there are no direct measurements of the mass transfer rate for SS433. However, the observed stability of the orbital period and evolution-time considerations exclude significantly higher mass transfer rates \\citep{ss2004,gorss}.}. If the donor star is highly evolved, accretion rate is higher and less stable and the phase itself is shorter. %Masses lost during the SN explosions leading to BH formation are %  $M_1\\sim 1\\div 5M_\\odot$ \\citep{casares,bhtheor} exclude very %  low-mass companions with $M \\lesssim 1 M_\\odot$. Unfortunately, initial binary mass ratio distribution for massive stars is poorly known but there are indications that mass ratios close to 1 are much more probable \\citep{lucy,KoFr2007}. The thermal timescale mass transfer rates in high-mass X-ray binaries are therefore likely to be highly supercritical, $\\mdot \\sim 100\\div 10^4$. Lower yet supercritical rates may take place in case of wind accretion for example in WR+BH binaries \\citep{ic10_bauer}. Those are likely to form a certain sub-sample of evolved ULXs with moderately supercritical accretion rates. % relatively free from geometrical beaming effects and absorption in the wind. %As we will see below the X-ray luminosity depends on the mass %  accretion rate even weaker than the bolometric luminosity  %Note that the mass of the secondary is very unlikely lower than one %half of the mass of the progenitor of the primary. Otherwise the %system is disrupted after the first SN explosion. %Masses %of pre-SNe expected to form BHs after the explosion are %lost during the SN explosions leading to BH formation are %$M_1\\sim 1\\div 5M_\\odot$ \\citep{casares,bhtheor}, so the number of thermal-scale %accreting sources drops abruptly at $\\mdot \\sim 100$.  The paper is organized as follows. In the next section we construct a simple analytical model describing the internal structure of an optically thick wind flow that is likely to develop in the case of very high mass ejection (and accretion) rate \\mdot. Section \\ref{sec:irrad} is devoted to effects of self-irradiation that are likely to play very important role in our funnel solution. In section  \\ref{sec:sp} we consider the emergent SEDs for arbitrary inclination taking into account self-occultation effects. In section \\ref{sec:obs} we %present an \\xspec\\ model {\\tt sirf} and test it for two sets of publicly available X-ray data. We discuss the implications of the model and its early testing in section \\ref{sec:disc}. ",
        "conclusions": "Optically thick wind with irradiation is capable to explain the SEDs of ULXs in the standard X-ray band and even the high-energy curvature that is difficult to explain in the framework of unsaturated comptonisation cool-disc IMBH model \\citep{curva}. High-energy cut-off is predicted to appear at several \\keV. Higher observed values of $T_{in} \\sim 1\\div 2\\,\\keV$ may be explained by applying a hardening factor $\\sim 1\\div 3$ similar to those predicted for accretion discs in X-ray binaries.  Outcoming spectra observed at low inclinations are similar to the spectra of slim discs and resemble {\\it p-free} model spectra with the $p$ parameter close to $0.5$. %to vary in a slightly wider range ($p$ may be less than $0.5$). %In fat discs irradiation and occultation effects must be %very important as well. X-ray spectra of known ULXs are approximated equally good by our model, {\\it p-free} and standard disc + power law two component models. However the parameters are poorly constraint supporting the idea that the properties of the X-ray spectrum depend rather weakly on the accretion disc and wind parameters. Relatively high inner temperatures argue for the funnel interior to be transparent. In this assumption, $r_{in}$ parameter may be used to estimate the mass ejection rate that appears to be of the order $\\mdot \\sim 100\\div 300$ for the two sources analysed.  We stress the extreme importance of irradiation effects providing mild geometrical collimation of the observed X-ray radiation. In a simple assumption of local absorption and re-radiation of the absorbed energy we show that the temperature of the funnel wall surface is altered by about 20\\% in the inner parts of the funnel, and the outcoming apparent X-ray luminosity becomes about $2\\div 3$ times higher.  %Our model suggests the outflow contains all the accreted mass (though %the \\xspec\\ version easily accounts for any differences in the outflow %and accretion rates). We also suggest the wind velocity is close to the %virial velocity at the spherisation radius. % Due to that our results on the angular %dependence of the outcoming spectrum differ from the results by %\\citet{okuda} and \\citet{heinz} and predict much softer spectra and much lower %X-ray fluxes at high inclinations.  Photoionized nebulae are likely to be formed around supercritical accretors. Ionizing quanta production rates suggest that in most cases a supercritical accretor is capable to produce a photoionized HII-region with bright optical emission line luminosities $\\sim 10^{37}\\ergl$ but higher luminosities may appear as well.  %For high mass ejection rates $\\gtrsim 10^4$ greater part of the EUV quanta produced %by the photosphere of the wind are absorbed by the wind %matter. However, p %Photoionized nebulae are likely to be formed %near the symmetry axis of the funnel producing regions of high %ionization like those observed for IC342~X-1 \\citep{rgww} and HoIX~X-1 %\\citep{hoixmois}. Ionizing quanta production rates suggest that in %most cases a SA is capable to produce a photoionized HII-region with %optical emission line luminosities $\\sim 10^{37}\\ergl$ but higher %luminosities can appear as well   \\bigskip %"
    },
    "0809/0809.0754_arXiv.txt": {
        "abstract": "We have mapped the molecular gas content in the host galaxy of the strongly lensed high redshift quasar \\apm\\ ($z=3.911$) with the Very Large Array at 0.3$''$ resolution.  The \\aco\\ emission is clearly resolved in our maps. The \\aco\\ line luminosity derived from these maps is in good agreement with a previous single-dish measurement. In contrast to previous interferometer-based studies, we find that the full molecular gas reservoir is situated in two compact peaks separated by $\\lesssim$0.4$''$. Our observations reveal, for the first time, that the emission from cold molecular gas is virtually cospatial with the optical/near-infrared continuum emission of the central AGN in this source.  This striking similarity in morphology indicates that the molecular gas is situated in a compact region close to the AGN. Based on the high resolution CO maps, we present a revised model for the gravitational lensing in this system, which indicates that the molecular gas emission is magnified by only a factor of 4 (in contrast to previously suggested factors of 100). This model suggests that the CO is situated in a circumnuclear disk of $\\sim$550\\,pc radius that is possibly seen at an inclination of $\\lesssim$25$^\\circ$, i.e., relatively close to face-on.  From the CO luminosity, we derive a molecular gas mass of $M_{\\rm gas}$=1.3$\\times$10$^{11}$\\,M$_{\\odot}$ for this galaxy.  From the CO structure and linewidth, we derive a dynamical mass of $M_{\\rm dyn}$\\,sin$^2 i$=4.0$\\times$10$^{10}$\\,M$_\\odot$. Based on a revised mass estimate for the central black hole of $M_{\\rm BH}$=2.3$\\times$10$^{10}$\\,M$_\\odot$ and the results of our molecular line study, we find that the mass of the stellar bulge of \\apm\\ falls short of the local $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relationship of nearby galaxies by more than an order of magnitude, lending support to recent suggestions that this relation may evolve with cosmic time and/or change toward the high mass end. ",
        "introduction": "Several different populations of distant galaxies have been detected to date, out to a (spectroscopically confirmed) redshift of $z$=7.0 (see Ellis \\citeyear{ell07} for a review).  It is of fundamental importance to study such young galaxies in great detail in order to constrain their nuclear, stellar, and gaseous constituents and their physical properties and chemical abundances.  Understanding these different characteristics of high redshift galaxy populations and their progression through cosmic times is vital to develop a unified picture of galaxy formation and evolution. One key aspect in these high-$z$ galaxy studies are sensitive, high-resolution radio observations of the molecular gas phase, i.e., the raw material that fuels star formation.  Since its initial discovery at $z$$>$2 more than a decade ago (Brown \\& Vanden Bout \\citeyear{bv91}; Solomon et al.~\\citeyear{sol92}), observations of spectral line emission from interstellar molecular gas in distant galaxies have revolutionized our understanding of some of the most luminous objects that populate the early universe. This is due to the fact that the physical state of the molecular interstellar medium (ISM) plays a critical role in the evolution of a galaxy. The total amount of molecular gas in a galaxy determines for how long starburst activity can be maintained, while its temperature and density are a direct measure for the conditions under which star formation can occur (see Solomon \\& Vanden Bout \\citeyear{sv05} for a review).  Also, there is growing evidence that the observed relationships between black hole mass and galaxy bulge velocity dispersion ($M_{\\rm BH}$--$\\sigma_{\\rm bulge}$; Ferrarese \\& Merritt \\citeyear{fm2000}; Gebhardt et al.~\\citeyear{g2000}), black hole mass and host bulge concentration (i.e., Sersic index; Graham et al.~\\citeyear{gra01}; Graham \\& Driver \\citeyear{gd07}), and, ultimately, black hole mass and bulge mass ($M_{\\rm BH}$--$M_{\\rm bulge}$; Magorrian et al.~\\citeyear{mag98}; H\\\"aring \\& Rix \\citeyear{hr04}), may be a consequence of an active galactic nucleus (AGN) feedback mechanism acting on its surrounding material, at least in the most luminous systems (Silk \\& Rees \\citeyear{sr98}; Di Matteo et al.~\\citeyear{mat05}). Such an AGN-driven wind will also interact with the molecular gas, the material which will eventually form the stellar bulges and disks in galaxies at high redshift, but also feed the active nucleus itself. Observations of the dynamical structure and distribution of molecular gas in distant galaxies may even reveal the initial cause of event (e.g., mergers) for both star formation and AGN activity, an important test for recent cosmological models (e.g., Springel et al.~\\citeyear{spr05}; Narayanan et al.~\\citeyear{nar06}).  The flux magnification provided by gravitational lensing has been facilitated in studies of high redshift molecular gas emission since early on, leading to the detection of a number of quasar host galaxies (e.g., Barvainis et al.\\ \\citeyear{bar94}; Downes et al.\\ \\citeyear{dow99}) and submillimeter galaxies (SMGs; e.g., Frayer et al.\\ \\citeyear{fra98}; Greve et al.\\ \\citeyear{gre05}), but also substantially fainter populations such as Lyman-break galaxies (LBGs; Baker et al.\\ \\citeyear{bak04}; Coppin et al.\\ \\citeyear{cop07}). Also, some sources have large enough lens image separations to show structure even at only moderate angular resolution (Sheth et al.\\ \\citeyear{she04}; Kneib et al.\\ \\citeyear{kne05}).  High angular resolution observations of molecular gas in high-$z$ galaxies are, to date, exclusively obtained in the rotational transitions of carbon monoxide (CO). In the past few years, a number of SMGs (Genzel et al.~\\citeyear{gen03}; Downes \\& Solomon \\citeyear{ds03}; Tacconi et al.~\\citeyear{tac06}) and quasars (Alloin et al.~\\citeyear{all97}) at $z$$>$2 have been studied at a linear resolution of up to 5\\,kpc at the target redshifts (1\\,kpc=0.12$''$ at $z$=2). In some cases, gravitational lensing aids in zooming in further on these systems. However, the only telescope that currently allows to attain resolutions of sub-kpc to 1\\,kpc scale at even higher redshifts ($z$$\\gtrsim$4, where 1\\,kpc$\\gtrsim$0.14$''$) is the NRAO's Very Large Array (VLA)\\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.}. First sub-arcsecond resolution CO imaging with the VLA has been obtained toward the distant quasars BR\\,1202-0725 ($z$=4.69; Carilli et al.~\\citeyear{car02}), APM\\,08279+5255 ($z$=3.91; Lewis et al.~\\citeyear{lew02}), PSS\\,J2322+1944 ($z$=4.12; Carilli et al.~\\citeyear{car03}; Riechers et al.\\ \\citeyear{rie08a}), SDSS\\,J1148+5251 ($z$=6.42; Walter et al.~\\citeyear{wal04}), and BRI\\,1335-0417 ($z$=4.41; Riechers et al.\\ \\citeyear{rie08b}).  We here report on new, more sensitive high-resolution observations of \\aco\\ emission towards the strongly lensed quasar APM\\,08279+5255 ($z=3.911$). APM\\,08279+5255 was already\\footnote{Due to the source's enormous optical brightness, much older images may exist.}  imaged during the Palomar Sky Survey (PSS) in 1953, a decade before the discovery of the first quasar was reported in the literature (Schmidt \\citeyear{sch63}).  About 45\\,years later, it was `officially' discovered in an automatic plate measuring facility (APM) survey for distant cool carbon stars, and identified as a 15$^{\\rm th}$ magnitude, radio-quiet broad absorption line (BAL) quasar with an IRAS Faint Source Catalog (FSC) counterpart at a redshift of $z_{\\rm opt}$=3.87 (Irwin et al.~\\citeyear{irw98}).  APM\\,08279+5255 was found to have an unprecedented apparent bolometric luminosity of $L_{\\rm bol}$=7$\\times$10$^{15}$\\,L$_\\odot$, about 20\\% of which is emitted in the (far-)infrared (Lewis et al.~\\citeyear{lew98}). This extreme value was found early on to be due to strong gravitational lensing (Irwin et al.~\\citeyear{irw98}), producing three close (maximum separation $<$0.4$''$), almost collinear images in the optical/near-infrared (Ibata et al.~\\citeyear{iba99}; Egami et al.~\\citeyear{e2000}). Due to the lack of a detection of the lensing galaxy, and an indication for significant microlensing effects, the true nature of the gravitational lens configuration (and thus the magnification factor $\\mu_{L}$) in this system remains subject to debate (e.g., Egami et al.~\\citeyear{e2000}; Lewis et al.~\\citeyear{lew02}) despite the fact that APM\\,08279+5255 is one of the best-studied sources in the distant universe.  In addition to its enormous continuum brightness at basically every astronomically relevant wavelength, APM\\,08279+5255 is also one of the brightest CO sources at high redshift (Downes et al.~\\citeyear{dow99}). It also exhibits unusually bright HCN (Wagg \\etal\\ \\citeyear{wag05}) and HCO$^+$ (Garcia-Burillo \\etal\\ \\citeyear{gar06}) emission, and was recently also detected in \\ci\\ (Wagg \\etal\\ \\citeyear{wag06}) emission.  The properties of its molecular gas content have recently been modeled based on observations of the \\aco\\ up to \\kco\\ rotational ladder (`CO spectral line energy distribution', or SLED) with the NRAO Green Bank Telescope (GBT)\\footnote{The Green Bank Telescope is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.} and the IRAM 30\\,m telescope (Riechers et al.~\\citeyear{rie06}; Wei\\ss\\ et al.~\\citeyear{wei07}). We use a standard concordance cosmology throughout, with $H_0 = 71\\,$\\kms\\,Mpc$^{-1}$, $\\Omega_{\\rm M} =0.27$, and $\\Omega_{\\Lambda} = 0.73$ (Spergel \\etal\\ \\citeyear{spe03}, \\citeyear{spe07}). ",
        "conclusions": "The observed properties of \\apm\\ are undoubtedly peculiar. It thus is important to understand whether this peculiarity is due to a rare lens configuration, possibly leading to very high lensing magnification factors and significant differential lensing effects, or whether \\apm\\ is a source that is just moderately magnified by gravitational lensing and thus intrinsically extreme.  \\subsection{Multiwavelength Properties of \\apm}  \\apm\\ shows broad absorption lines in the optical/UV, and emission from high excitation lines like \\civ\\ and \\nv\\ close to the central nucleus (Irwin \\etal\\ \\citeyear{irw98}), kinematically blueshifted relative to the molecular gas and \\ci\\ emission by about 2500\\,\\kms\\ (Downes \\etal\\ \\citeyear {dow99}; Wagg \\etal\\ \\citeyear{wag06}). The quasar also shows relativistic X-ray broad absorption lines from even higher emission lines of iron from gas situated within the UV BAL region (Chartas \\etal\\ \\citeyear{cha02}; Hasinger \\etal\\ \\citeyear{has02}).  The BALs come from an outflow of highly ionized gas from the accretion disk, mostly driven by radiation pressure. This outflow distributes the accretion disk material over the central region of the quasar, and out into the host galaxy. The X-ray BAL region is likely a dust-free, high column density absorber responsible for shielding, as indicated by the detection of an iron K-shell absorption edge (Hasinger \\etal\\ \\citeyear{has02}). This shielding is responsible for the typical X-ray faintness of BAL QSOs.  The X-ray luminosity alone thus is not a good measure for the energy output of \\apm. The detection of strong iron lines also indicates that the X-ray BAL emission comes from significantly metal-enriched gas with possibly super-solar metallicities. While this strongly indicates that the gas has already been processed in a starburst cycle, it is biased toward the very central region (possibly $<$0.1\\,pc). It thus cannot be assumed ad hoc that the metallicity is super-solar over the whole scale of the galaxy. The continuum X-ray emission in \\apm\\ has a photon index that is in qualitative agreement with a radio-quiet quasar (RQQ). The optical/IR shape of the SED is also consistent with \\apm\\ being a RQQ. Even more so, it follows the radio-quiet locus of the radio--FIR correlation, which indicates that the copious amounts of dust in this source may be heated to a significant fraction by newly formed young stars, while the radio synchrotron emission may originate from supernova remnants (Irwin \\etal\\ \\citeyear{irw98}; Beelen \\etal\\ \\citeyear{bee06}). However, our analysis of the radio and dust properties indicates that most of the radio/FIR emission may well be AGN-related.  Within the unified theory of AGN galaxies, the smoothness of the IR spectrum, in particular the lack of a pronounced IR peak, indicates that the accretion disk in this radio-quiet quasar is seen close to face-on (Soifer \\etal\\ \\citeyear{soi04}). This agrees well with the enormous optical brightness and the flatness of the optical spectrum (Irwin \\etal\\ \\citeyear{irw98}; Egami \\etal\\ \\citeyear{e2000}), indicating a lack of obscuration and thus an angle of view which clearly lies within the ionization cone.  The hot dust in the parsec-scale accretion disk is heated by the central AGN, and extends outwards into a warm dusty torus, extending out to 100\\,pc scales.  The luminous AGN in \\apm\\ is the dominant power source out to large scales, which even may generate a large fraction of the huge (far-)infrared luminosity in this source by heating the dust.\\footnote{It is important to note that, due to the unique properties of \\apm, additional diagnostics are desirable to undoubtedly disentangle the heating sources of the gas and dust.}  The apparent FIR luminosity of \\apm\\ is $L_{\\rm FIR}$=(2.0 $\\pm$ 0.5) $\\times$ 10$^{14}\\,\\mu_{L}^{-1}$L$_\\odot$ (Beelen \\etal\\ \\citeyear{bee06}; Wei\\ss\\ \\etal\\ \\citeyear{wei07}).  Wei\\ss\\ \\etal\\ (\\citeyear{wei07}) have found that about 90\\% of the FIR luminosity are due to emission from warm 220\\,K dust which may be powered by the AGN, while only 10\\% are due to colder 65\\,K dust that is likely heated by star formation.  This would explain why a significant fraction of the dust in \\apm\\ on extended (100s pc) scales is warmer than in typical ULIRGs, or even other high-$z$, dust-rich QSOs. It would also explain why \\apm\\ is an outlier on the $L'_{\\rm CO}$--$L_{\\rm FIR}$ relation (Riechers \\etal\\ \\citeyear{rie06}), which is set by galaxies where the FIR luminosity is dominated by heating from star formation.  Subtracting the 90\\% of $L_{\\rm FIR}$ that may be powered by the AGN, \\apm\\ follows the $L'_{\\rm CO}$--$L_{\\rm FIR}$ relation quite well. However, note that even though the apparent FIR luminosity of \\apm\\ is extreme, this would indicate that only 3\\% of the total quasar power (as measured by $L_{\\rm bol}$) are re-emitted in the FIR wavelength regime (assuming no differential lensing effects play into the ratio). We find that up to 80--90\\% of the radio luminosity in \\apm\\ may be powered by the AGN, and only 10--20\\% by the starburst. Due to the fact that this ratio is similar to that found for the FIR luminosity, \\apm\\ still follows the radio--FIR correlation for star-forming galaxies.  \\apm\\ thus is a radio-quiet, optically luminous dust-rich quasar in which the central accretion disk and dust torus are likely seen close to face-on (note however the above concern about the inclination). Its high FIR luminosity can be explained by dust heating of the AGN, as can be the high temperature of the dust. It even explains the high CO excitation if the molecular gas is situated in a rather compact (and thus dense) circumnuclear ring seen at a low inclination toward the line of sight, just like the central region. The high dust and gas temperatures and relatively high densities also explain the exceptional HCN (and HCO$^+$) excitation in this system, as IR-pumping via the $\\nu_2$=1 vibrational bending mode becomes efficient at such high temperatures (Wei\\ss\\ \\etal\\ \\citeyear{wei07}).  Without shielding and self-shielding, the gas and dust would be even warmer than observed. Assuming that only the fraction of $L_{\\rm bol}$ re-radiated in the FIR reaches out to the molecular regions, the temperature at a radius of 125\\,pc would still be 200\\,K. Even at 550\\,pc, the temperature would not drop significantly below 100\\,K. The LVG-predicted temperatures for the cool, dense gas component ($T_{\\rm kin}=65$\\,K) thus require self-screening if the gas is indeed located in a 550\\,pc region around the AGN.\\footnote{Due to the fact that $T \\propto \\mu_L^{-1/4}$ and $\\mu_L \\propto r_{\\rm true}^2$, the $\\mu_L$=4 model requires more screening than the previous $\\mu_L$=100 models.} Note that all the above may be considered extreme (i.e., rarely observed). However, all of these effects can be explained physically without assuming extreme and/or differential lensing effects. The CO luminosity in \\apm\\ is relatively modest (compared to other high-$z$ sources) even without correcting for lensing. If $L_{\\rm FIR}$ is corrected assuming a lensing magnification factor of only 4, as suggested by our new model, it is not the highest observed value even without correcting for AGN heating. Both $L'_{\\rm CO}$ and $L_{\\rm FIR}$ are higher for the $z$=4.7 quasar BR\\,1202-0725, which is thought to be unlensed, but undergoing a massive, gas-rich merger event\\footnote{Both the CO and FIR emission in BR\\,1202-0725 are however distributed over at least two distinct components, which then each have comparable but lower luminosities than \\apm.} (Carilli \\etal\\ \\citeyear{car02}; Riechers \\etal\\ \\citeyear{rie06}).  \\subsection{The Black Hole in \\apm}  Correcting for $\\mu_{L}$=4 predicts $M_{\\rm BH}$=2.3 $\\times$ 10$^{10}$\\,M$_\\odot$ for \\apm. This appears to be extreme compared to more typical SMBH masses of a few times 10$^9$\\,M$_\\odot$ of high-$z$ quasars.  However, comparable or even higher black hole masses are found for BRI\\,1335-0417 ($z$=4.41) and BR\\,1202-0725 (however, the estimates for these sources are based on strongly absorbed \\civ\\ emission lines; Storrie-Lombardi \\etal\\ \\citeyear{sl96}; Shields \\etal\\ \\citeyear{shi06}). Moreover, Vestergaard (\\citeyear{ves04}) has found that for the range of black hole masses, an upper envelope of $M_{\\rm BH}$$\\sim$10$^{10}$\\,M$_\\odot$ is observed over a large range in redshift (4$\\lesssim$$z$$\\lesssim$6 for her sample). So, again, while \\apm\\ appears relatively extreme, even its AGN properties are far from being unique in a scenario assuming only $\\mu_{L}$=4. The extreme optical luminosity may then be explained by a luminous AGN episode of such a few times 10$^{10}$\\,M$_\\odot$ SMBH, which accretes at (super-)Eddington rates (within the uncertainties of the BH mass estimators).  Also, the formation of such a SMBH at the high mass end of the BH mass function by a redshift of 4 is consistent with `downsizing', which predicts that such high mass black holes form at the high density peaks at high redshift, and are built up rapidly by vigorous accretion (e.g., Di Matteo \\etal\\ \\citeyear{mat07}). Due to the fact that the SMBH in \\apm\\ is among the most massive black holes that are observed, it is likely to end up as a central cD galaxy of a massive cluster.  For such galaxies, recent cosmological simulations suggest that a large fraction of the mass of the black hole grows from mergers with other black holes rather than being accreted by a single progenitor black hole (Sijacki \\etal\\ \\citeyear{sij07}).  \\subsection{The Stellar Bulge of \\apm}  In addition to $M_{\\rm BH}$, we were also able to derive a limit on the mass of the stellar bulge ($M_{\\rm bulge}$) for \\apm. However, our dynamical mass estimate, which is used as the predictor for $M_{\\rm bulge}$, is only valid for the inner region of the galaxy where the bulk of the molecular gas is found. For the $\\mu_{L}$=4 model, this describes a region with a radius of 550\\,pc. For galaxies with $M_{\\rm BH}$$>$10$^9$\\,M$_\\odot$ used in local studies of the $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relation, it has been found that the host galaxies are typically giant elliptical galaxies with effective radii of 1.5--8\\,kpc (Tremaine \\etal\\ \\citeyear{tre02}; Faber \\etal\\ \\citeyear{fab97}). Assuming that the volume density of stellar luminosity (and thus mass) flattens with radius from $r^{-2}$ to $r^{-1}$ in the inner region of the elliptical host galaxy (Gebhardt \\etal\\ \\citeyear{geb96}), or even that it flattens into a `core', a significant fraction of the bulge mass is expected to be found inside the inner 550\\,pc. From the local $M_{\\rm BH}$--$M_{\\rm bulge}$ relationship, it is expected that $M_{\\rm bulge}$$\\sim$700\\,$M_{\\rm BH}$. For the inner region of \\apm, we have found $M_{\\rm bulge}$=3.4\\,$M_{\\rm BH}$. Although uncertain within a factor of a few, this value falls short by more than two orders of magnitude compared to the local estimate. Even if corrected for a 10\\,kpc radius host galaxy, a clear offset from the local $M_{\\rm BH}$--$M_{\\rm bulge}$ relationship remains\\footnote{Note that the local $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relationship is difficult to reconcile even with a $\\mu_{L}$=100 model, which leaves more room for $M_{\\rm bulge}$ within $M_{\\rm dyn}$ (i.e., $M_{\\rm bulge}$=50\\,$M_{\\rm BH}$), but also has to assume host galaxy size relations for less massive black holes.}. This is consistent with what is found for other quasars at even higher redshift studied with the same technique (Walter \\etal\\ \\citeyear{wal04}; Riechers et al.\\ \\citeyear{rie08a}, \\citeyear{rie08b}; see also discussions by Shields \\etal\\ \\citeyear{shi06}; Ho \\citeyear{ho07}). Based on different techniques, similar trends are found for various galaxy samples at low and intermediate redshift (e.g., Peng et al.\\ \\citeyear{pen06}; McLure et al.\\ \\citeyear{mcl06}; Woo \\etal\\ \\citeyear{woo06}; Salviander et al.\\ \\citeyear{sal07}). This may indicate a breakdown of the local $M_{\\rm BH}$--$M_{\\rm bulge}$ relationship toward high redshift, and that the mass correlation between SMBHs and their stellar bulges is not a universal property, but rather an endpoint of a long-lasting evolutionary process throughout cosmic times.  We thus conclude that the SMBH in \\apm\\ appears to already be largely in place, while the buildup of the stellar bulge is still in progress.   \\subsection{A Possible Evolutionary Scenario for \\apm}  If the assumptions made to derive the properties of \\apm\\ are correct, this result has several interesting implications. It appears as if SMBH growth and formation through gas dissipation and accretion (and possibly mergers) was the main characteristic during the early formation of this source (even relative to star formation).  Since the SMBH in \\apm\\ is already found at the high mass end of the black hole mass function, it is expected that most of the molecular gas present at $z$=3.9 is either blown out or formed into stars by $z$=0, while only a minor fraction is accreted onto the black hole.  Moreover, if \\apm\\ is to evolve into a galaxy that fulfills the $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relation by $z$=0, a significant fraction of the molecular gas has to actually end up in stars. In fact, even if the whole amount of molecular gas that is currently present in the host galaxy of \\apm\\ were to be converted into stars by $z$=0, it would not be sufficient to reach the local $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relation. If our picture of this galaxy is correct, this result implies that the buildup of the stellar bulge cannot be accomplished by only converting the observed molecular gas reservoir into stars, but relies to a significant part on the accumulation of stellar matter via other mechanisms, and/or the the accretion of additional star-forming material. Although further details are difficult to quantify at present, it thus is likely that a significant fraction of the spheroid will be produced via major and minor mergers (which are also required for the re-distribution of angular momentum). It has been suggested that a common, growth-limiting feedback mechanism acting on both star formation and black hole assembly is the physical explanation for the local $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relation.  Assuming that such a mechanism exists, quasar winds as a form of AGN feedback would be an obvious candidate for a system like \\apm\\ (Silk \\& Rees \\citeyear{sr98}). One difficulty with this assumption would then be that the SMBH in this source already accretes radiatively highly efficiently at (super-)Eddington rates (see above), where such feedback effects to regulate the further growth of the SMBH are expected to be very strong. If this effect were to also regulate star formation in the host galaxy, it would be expected to shut off both on a relatively short timescale by pushing the remaining gas outwards.  The build-up of the stellar bulge then can only proceed further through hierarchical merging and/or accretion of additional gas, which is expected to again feed both the SMBH and star formation. To end up anywhere near the local $M_{\\rm BH}$--$\\sigma_{\\rm bulge}$ relation through a self-regulated mechanism, \\apm\\ thus is required to form and/or accumulate stars more efficiently (relative to the black hole feeding/merger rate) than in the past. One possibility to build up stellar mass while avoiding significant SMBH growth may be through subsequent `{\\em dry}' mergers (e.g., Rines et al.\\ \\citeyear{rin07}), which may act in combination with other processes to assemble the massive host of \\apm\\ by $z$=0.   \\subsection{Our Current Picture of \\apm}  There are two main points that support the $\\mu_{L}$=4 lensing model proposed for \\apm.  First, the observed \\aco\\ properties give rather tight constraints on the allowed combinations of intrinsic CO disk size and magnification strength for the favoured range of low disk inclinations, assuming a relatively high CO area filling factor. The $\\mu_{L}$=4 model lies well within the range of these constraints. Second, the spatially resolved multiwavelength observations of \\apm\\ show similar morphologies of the lensed source from the X-ray to radio wavelength regime, probing sub-pc to hundreds of pc scales, while showing a different morphology out to kpc scales. The $\\mu_{L}$=4 model produces this quite naturally, predicting modest effects of differential lensing. The high lensing magnification models suggested so far do not reproduce the second characteristic. Also, for low inclinations, they predict a very compact CO disk of only about 100\\,pc radius. This is by a factor of a few less than in typical ULIRGs (Downes \\& Solomon \\citeyear{ds98}). The $\\mu_{L}$=4 model, on the other hand, produces a CO disk of 550\\,pc radius, which is in a more typical regime.  While some properties of \\apm, like the optical luminosity and expected strong tidal forces from the black hole, as well as the requirement of significant shielding of the molecular gas, may still favour a scenario in which the source is highly magnified by gravitational lensing, there thus are several arguments that favour models with low $\\mu_{L}$. It however is important to note that, if a high-$\\mu_{L}$ model was to be found that fits the observed morphological properties of \\apm\\ (see Krips et al.\\ \\citeyear{kri07} for a recent attempt), further studies would be necessary to exclude one scenario or the other.  Based on the existing observations, no definitive conclusion can be drawn -- the final decision can only be made by detecting the lens itself. The lensing model presented in this paper aids in developing a quite consistent picture of \\apm\\ based on the existing observations, describing the source as a dust and gas-rich galaxy with a very massive, active black hole. The deep gravitational potential of the central region of this source causes the molecular gas to be rather dense, while the strong, penetrating AGN radiation causes it (and the dust) to be rather warm.  The whole system may be seen relatively close to face-on (the inclination is however difficult to constrain).  It has a co-eval starburst which only contributes about 10\\% to the FIR dust heating, but still produces more than 500\\,M$_\\odot$\\,yr$^{-1}$ (Wei\\ss\\ \\etal\\ \\citeyear{wei07}, correcting for $\\mu_{L}$=4). The star formation in this galaxy takes place far off the black hole mass--bulge mass scaling relation of nearby galaxies, suggesting that the latter may be subject to significant evolution throughout cosmic times."
    },
    "0809/0809.5171_arXiv.txt": {
        "abstract": "{{}AGN feedback now appears as an attractive mechanism to resolve some of the outstanding problems with the ``standard'' cosmological models, in particular those related to massive galaxies. At low redshift, evidence is growing that gas cooling and star formation may be efficiently suppressed by mechanical energy input from radio sources. To directly constrain how this may influence the formation of massive galaxies near the peak in the redshift distribution of powerful quasars, z$\\sim 2$, we present an analysis of the emission-line kinematics of 3 powerful radio galaxies at z$\\sim 2-3$ (HzRGs) based on rest-frame optical integral-field spectroscopy obtained with SINFONI on the VLT. The host galaxies of powerful radio-loud AGN are among the most massive galaxies, and thus AGN feedback may have a particularly clear signature in these galaxies.  We find evidence for bipolar outflows in all HzRGs, with kinetic energies that are equivalent to 0.2\\% of the rest-mass of the supermassive black hole. Observed total velocity offsets in the outflows are $\\sim 800-1000$ km s$^{-1}$ between the blueshifted and redshifted line emission, and FWHMs $\\sim 1000$ km s$^{-1}$ suggest strong turbulence. Line ratios allow to measure electron temperatures, $\\sim 10^4$ K from [OIII]$\\lambda\\lambda\\lambda$4363,4959,5007 at z$\\sim 2$, electron densities ($\\sim 500$ cm$^{-3}$) and extinction (A$_V\\sim 1-4$ mag).  Ionized gas masses estimated from the H$\\alpha$ luminosity are of order $10^{10}$ M$_{\\odot}$, similar to the molecular gas content of HzRGs, underlining that these outflows may indicate a significant phase in the evolution of the host galaxy. The total energy release of $\\sim 10^{60}$ erg during a dynamical time of $\\sim 10^7$ yrs corresponds to about the binding energy of a massive galaxy, similar to the prescriptions adopted in galaxy evolution models. Geometry, timescales and energy injection rates of order 10\\% of the kinetic energy flux of the jet suggest that the outflows are most likely driven by the radio source. The global energy density release of $\\sim 10^{57}$ erg s$^{-1}$ Mpc$^{-3}$ may also influence the subsequent evolution of the HzRG by enhancing the entropy and pressure in the surrounding halo and facilitating ram-pressure stripping of gas in satellite galaxies that may contribute to the subsequent mass assembly of the HzRG through low-dissipation ``dry'' mergers.  } ",
        "introduction": "\\label{introduction} AGN feedback has now become a crucial process in most models of galaxy formation and evolution. Current models postulate that massive galaxies experienced an early epoch of vigorous star formation that was terminated by a nearly instantaneous, powerful blow-out phase triggered by the AGN \\citep[e.g.][]{silk98, scannapieco04, springel05, dimatteo05, croton06, bower06, hopkins06}. Such a mechanism would heat and remove most of the ambient gas of the galaxy, and thus reconcile the discrepancy between the hierarchical standard model of galaxy formation and observational constraints  Bolometric luminosities of powerful AGN at high redshift correspond to an energy output of $10^{60-61}$ erg during an AGN lifetime of $10^{7-8}$ yrs.  Although this is equal, or even exceeds the binding energy of a massive galaxy, this energy will only have an impact on the evolution of the host galaxy, if it is efficiently transferred to the ambient gas with a coupling efficiency of at least a few percent.  Observationally, evidence is growing that AGN feedback may be related mostly to radio-loud AGN. Perhaps the most spectacular examples at low redshift are the giant cavities in the X-ray halos of massive galaxy clusters, which appear to inject few $\\times 10^{60}$ erg into the surrounding gas within a few $\\times 10^7$ yrs, sufficient to suppress gas cooling \\citep{boehringer93,birzan04,rafferty06,mcnamara07}. \\citet{best05,best06} find more subtle, yet general evidence that radio sources may balance gas cooling in $\\sim 2000$ radio-loud early-type galaxies taken from the Sloan Digital Sky Survey and the NVSS and FIRST radio surveys. Interestingly, the fraction of radio loud galaxies seems to be a strong function of mass, suggesting that such radio-driven feedback will act predominantly in the most massive galaxies, in broad agreement with the predictions of galaxy evolution models. This is also supported by the observation that powerful radio galaxies are among the most massive galaxies at all redshifts \\citep{debreuck01,willott03,rocca04}. Since AGN feedback is expected to act predominantly in massive galaxies, we may expect that its influence will be most clearly expressed in particularly massive galaxies.  However, since massive early-type galaxies appear to have completed most of their growth at high redshift \\citep[e.g.,][and references therein]{rudnick03}, directly investigating the impact of jet-driven AGN feedback during the formation of massive galaxies requires direct observations of radio-loud galaxies at z$\\ge 1$.  Interestingly, longslit spectroscopy of powerful radio galaxies at z$\\ge 1$ (and in a few cases at lower redshifts) revealed strongly distorted emission line kinematics, with velocities often exceeding $\\sim 1000$ km s$^{-1}$ along the axis of the radio jet \\citep[e.g.,][]{tadhunter91,mccarthy96,evans98,baum00,villar99, inskip02,clark98,best97}.  Similarly, [OIII]$\\lambda$5007 narrow-band imaging of 4C19.71 at z$\\sim 3.6$ revealed a giant emission line region extending over $\\sim 74$ kpc \\citep{armus98}. This emission line region, corresponding to an ionized gas mass of $10^{8-9}$ M$_{\\odot}$ aligns roughly with the axis of the radio jet, which led \\citeauthor{armus98} to suspect that the jet and the nebulosity may be physically related. \\citet{heckman91a,heckman91b} also find this for radio-loud quasars. Moreover, Ly$\\alpha$ spectroscopy and line imaging revealed luminous emission line halos with sizes that appear related to the size of the radio jet, sometimes with kinematically more quiescent gas beyond \\citep[e.g.][]{villar03,villar06,villar07}.  However, the complexity of the velocity fields and relatively large spatial extent of the emission line regions made it difficult to robustly infer the global characteristics of these nebulae. Integral-field spectrographs now open a new avenue to study the underlying dynamical mechanisms with unprecedented robustness and across the full two-dimensional surface of the emission line nebulae. Over recent years, our understanding of galaxy evolution in the early universe has significantly advanced due to the discovery of large numbers of galaxies at similar redshifts \\citep[e.g.,][]{steidel96,smail02}, and during this time, modeling of galaxy and structure formation has improved. This allows us to reassess the emission line kinematics of powerful high-redshift radio galaxies (HzRGs). Near-infrared studies are particularly well suited to study the inner regions of HzRG extended emission line regions, allowing for relatively high spatial resolution of $\\sim 0.5$\\arcsec\\ even with seeing-limited observations, and opening a window to the rest-frame optical line emission where extinction is less severe than in the rest-frame UV.  \\citet{nesvadba06} find strong evidence for energetic outflows of ionized gas from rest-frame optical integral-field spectroscopy of the powerful radio galaxy MRC1138-262 at z$\\sim 2.2$. With a spatial extent of 30 kpc, relative velocities and line widths of FWHM$\\sim 1000$ km s$^{-1}$ across the source, corresponding to kinetic energies of few$\\times 10^{60}$ erg, this gas may be experiencing a feedback episode as dramatic as necessary to unbind a significant fraction of the ambient gas from the halo of a massive galaxy. The kinetic energy of the outflow corresponds to a few percent of the jet kinetic energy, and the dynamical timescale of a few $\\times 10^7$ yrs appears roughly similar to the age of the radio source. Moreover, the turbulent, luminous gas extends to the size of the radio jets. In contrast, very similar observations do not show evidence for such outflows in compact radio galaxies at similar redshifts \\citep{nesvadba07b}, which are likely the younger analogs to extended HzRGs. This agrees with expectations. If the entrainment rates are similar, then only a few percent of the ISM will already be affected by the radio source for ages less than $\\sim 10^6$ yrs.  Currently, the host galaxies of powerful, radio-loud AGN (HzRGs and quasars) represent the only galaxy population at z$\\ge 2$ with such extreme kinematics. To address the question whether MRC1138-262 is a ``one-off'', or whether AGN feedback is a common phenomenon in HzRGs, we continue the analysis of \\citet{nesvadba06,nesvadba07b}, adding another three HzRGs with near-infrared integral-field data at z$\\approx 2-3$. Comparing with the samples of \\citet[e.g.,][]{iwamuro03,debreuck00,humphrey08,baum00}, we find that their rest-frame UV-optical properties are within the typical range of powerful HzRGs. These galaxies have extended radio morphologies, and relatively similar radio power to MRC1138-262, ${\\cal P}_{1.4 GHz} \\sim 10^{26}$ W Hz$^{-1}$. Thus, they are among the most powerful radio galaxies observed at all redshifts. Given the observational difficulties in studying high-redshift galaxies, this seems the most promising way of identifying the fingerprints of powerful AGN-driven feedback and underlying physical processes directly from the gas kinematics of powerful AGN host galaxies at z$\\sim 2$.  However, these radio powers are not exceptional in the early universe. \\citet{willott01} studied the redshift evolution of the radio luminosity function and find a population of galaxies with radio powers ${\\cal P}_{151 MHz} \\sim 10^{26.5-29.5}$ W Hz$^{-1}$ at 151 MHz, which shows a strong peak near z$\\sim 2$ and rapidly declines towards lower redshifts. Finally, the radio emissivity is likely a strong function of environment and evolutionary stage of the synchrotron plasma \\citep[e.g.,][]{kaiser97a,kaiser97b}, causing changes in the radio luminosity of more than an order of magnitude. Hence, by selecting particularly powerful sources, we may bias our sample more in terms of the evolutionary stage or environment than the kinetic luminosity of the radio source.  Specifically, we observed MRC0316-257 at z$=3.13$, MRC0406-244 at z$=2.42$, and TXS0828+192 at z$=2.57$. Entrained ionized gas masses and kinetic energies seem very similar to MRC1138-262 in all three cases. This leads us to propose that AGN feedback may be a common phenomenon in powerful radio galaxies during the ``Quasar Era''. ",
        "conclusions": "We studied the kinematics and physical characteristics of the extended emission line regions in 3 powerful radio galaxies at z$\\sim 2$, using integral-field spectroscopy of the rest-frame optical emission lines. Our results are consistent with what is expected from giant, jet-driven AGN outflows, similar to what we found in an earlier study of the z$\\sim 2.2$ galaxy MRC1138-262 \\citep{nesvadba06}. Analysis of the emission and continuum line morphologies, velocity maps, and flow orientation relative to the orientation of the radio jet does not favor alternative interpretations, e.g., merging, rotation, or gas infall. Such outflows may be common  in powerful high-redshift radio galaxies (HzRGs) during the ``Quasar Era''. Overall, we find energetic outflows in all 4 z$\\sim 2$ powerful radio galaxies with extended jets for which we obtained rest-frame optical integral-field spectroscopy. Comparison with previous longslit spectroscopy \\citep{baum00} suggests similar properties in $> 10$ HzRGs at z$>1$.  The emission line regions extend over $\\sim 10$ kpc$\\times$30 kpc, and are elongated along the axis of the radio jets. Most of the complex morphologies of these galaxies, previously observed with broadband imaging, appears to be due to line emission. The line-free continuum emission extracted from the data cube is relatively faint and compact ($\\le 10$ kpc). Assuming simple case-B recombination, we find of order $10^{10}$ M$_{\\odot}$ of ionized gas, similar to the typical cold molecular gas content of HzRGs traced by CO line emission. From observed line ratios, we can constrain the physical properties of the gas. We detect [OIII]$\\lambda$4363 in a z$\\sim 2$ galaxy, allowing for a direct estimate of the electron temperature of $\\sim 10^4$ K. Electron densities are derived from the [SII]$\\lambda\\lambda$6716,6731 line ratio and are $\\sim 500$ cm$^{-3}$. From the measured H$\\alpha$/H$\\beta$ line ratio we estimate extinction of A(H$\\beta$)$\\sim 1-4$ mag.  The velocity fields are very similar in all 3 sources and are reminiscent of two bubbles expanding back-to-back from the AGN. Velocity offsets near the center of the host galaxy are large, 700$-$1000 km s$^{-1}$, but velocities within each bubble are relatively uniform. Line widths are of order FWHM$\\sim 1000$ km s$^{-1}$, indicating strong turbulence. Larger widths in the central regions are likely due to overlaps between the bubbles. Comparison with the orientation of the radio jets suggest that this gas is outflowing.  With kinetic energies of order $10^{60}$ erg required to power the observed emission line kinematics, these outflows may be sufficient to remove most of the interstellar medium of the host galaxies in a nearly explosive event. Observed kinetic energies correspond to $\\sim 0.2$\\% of the energy equivalent of the rest mass of the supermassive black holes, in good agreement with AGN feedback prescriptions used in galaxy evolution models  The following observations lead us to argue that the outflows are dominated by the radio jets: (1) Geometry: Good alignment with the radio jets and elongation along the jet axis. The gas is blueshifted on the side of the approaching radio jet. (2) Size: The sizes of the emission line regions do not reach beyond the radio lobes. \\citet{nesvadba07b} do not find extended outflows in galaxies with compact radio sources observed with SINFONI. Emission line gas outside the radio lobes observed in Ly$\\alpha$ is kinematically quiescent \\citep{villar03}. (3) Timescales: Dynamical timescales of the outflows are few $\\times 10^7$ yrs, similar to typical ages of radio jets. (4) Energies: Kinetic energies of the outflows correspond to about 10\\% of the jet kinetic energy.  Comparing with the z$\\sim 2$ radio luminosity function of \\citet{willott01}, we find that AGN winds like those observed may release about 9$\\times 10^{56}$ erg s$^{-1}$ Mpc$^{-3}$ in energy density from the massive host galaxies of powerful radio sources at z$\\sim 2$. This energy may have a long-term impact on the subsequent evolution of the AGN host galaxy as well as nearby galaxies and help maintaining the ``anti-hierarchical'' properties of massive galaxies in spite of possible subsequent merging with companion galaxies. Following \\citet{nath02,mccarthy08}, much of this energy may contribute to increasing the entropy and pressure in the environment of massive early type galaxies and clusters, and enhance the efficiency of dynamical interactions with galaxies in the environment of HzRGs, such as ram pressure stripping. Comparing with studies of ram-pressure stripping \\citep{boselli08} we find that the satellite galaxies may already be red, poor in gas and rich in metals when they are being accreted onto the central galaxy. In spite of subsequent growth by up to $\\sim 50\\%$ in stellar mass, we find that minor ``dry'' mergers with ram-pressure stripped satellites would not strongly influence the observed evolutionary properties of the central galaxy."
    },
    "0809/0809.1087_arXiv.txt": {
        "abstract": "The spectrophotometric calibration of surveys is a significant, but often neglected, issue when measuring the history of star formation by combining spectroscopic surveys conducted with different instruments. We describe techniques for photometric calibration of optical spectra obtained with the MMT's fiber-fed spectrograph, Hectospec.  The atmospheric dispersion compensation prisms built into the MMT's f/5 wide field corrector effectively eliminate errors due to differential refraction, and simplify the calibration procedure.  The procedures that we describe here are applicable to all 220,000+ spectra obtained to date with Hectospec because the instrument response is stable.  We estimate the internal error in the Hectospec measurements by comparing duplicate measurements of $\\sim$1500 galaxies.  For a sample of 400 galaxies in the Smithsonian Hectospec Lensing Survey (SHELS) with a median z=0.10, we compare line and continuum fluxes measured by Hectospec through a 1.5$^{\\prime\\prime}$ diameter optical fiber with those measured by the Sloan Digital Sky Survey (SDSS) through a 3$^{\\prime\\prime}$ diameter optical fiber. Agreement of the [OII] and H${\\alpha}$ SHELS and SDSS line fluxes, after scaling by the R band flux in the different apertures, suggests that the spatial variation in star formation rates over a 1.5 to 3 kpc radial scale is small.  The median ratio of the Hectospec and SDSS spectra, smoothed over 100 ${\\rm \\AA}$ scales, is remarkably constant to $\\sim$5\\% over the range of 3850 to 8000 ${\\rm \\AA}$. Offsets in the ratio of the median [OII] and H$\\alpha$ fluxes,  the equivalent width of H$\\delta$ and the continuum index d4000 are a few percent, small compared with other sources of scatter.  We also explore the  impact of atmospheric absorption.  Observing redwards of 6500 ${\\rm \\AA}$, it is impossible to remove the  effects of atmospheric absorption perfectly because the variation of absorption with wavelength is not resolved by a moderate dispersion spectrograph.  Thus  measurements of spectral line fluxes including H${\\alpha}$, and derived physical quantities including star formation rates, may have sizable systematic errors where the redshifted spectral features land on strong atmospheric absorption troughs. ",
        "introduction": "As spectroscopic and photometric surveys increase in quality and size, the physical parameters which govern  galaxy evolution become increasingly accessible. A range of measures of star formation rate densities from a variety of surveys provide an impressive outline of the star formation history of the universe \\citep[e.g.][]{Hopkins04, Ly07, Tresse07, Reddy08, Shioya08, Villar08}.  One of the many issues in reconstructing the star formation history is the relative calibration of surveys which use different measures (Hopkins 2004). After calibration, the scatter in the measured star formation rate density at fixed redshift remains large. Even for the same star formation indicator, inadequacy of the relative calibration of surveys with different instruments on different telescopes may be one source of scatter and systematic offset. Here we use the overlap of  two large spectroscopic samples from different telescopes with different fiber instruments to evaluate these issues for emission line fluxes, H$\\delta$ equivalent width, and the strength of the 4000 {\\rm \\AA} break. We demonstrate that with attention to technical detail, the median offsets in all of these measures are at the few percent level.  One of our overlapping samples is the SHELS survey {\\citep{Geller05} carried out with the Hectospec \\citep{Fabricant05} on the MMT; the other is derived from the Sloan Digital Sky Survey \\citep{DR6}. We describe the spectrophotometric calibration of Hectospec and estimate the internal errors in this calibration by examining 1467 unique pairs of measurements of a subset of the galaxies.  We estimate the external errors in the spectrophotometric calibration by comparisons with galaxies observed in common with the Sloan Digital Sky Survey \\citep[SDSS,][]{DR6}.  We extract the absolute calibration for our spectrophotometry from external photometry; we use the Deep Lensing Survey R-band photometry \\citep{Wittman06, Wittman02}.  This absolute calibration corrects for clouds, seeing, telescope tracking and guiding, as well as errors in astrometry and alignment, as long as these losses of light are independent of wavelength.  The spectrophotometric errors associated with differential refraction \\citep{Filippenko82} are largely eliminated because the MMT's wide-field corrector \\citep{Fabricant04} contains atmospheric dispersion compensation (ADC) prisms.  The atmospheric dispersion across Hectospec's wavelength range ($\\sim$3700 to 8500 {\\rm \\AA}), if not corrected by the ADC prisms,  would exceed the Hectospec's 1.5$^{\\prime\\prime}$ fiber diameter at a zenith angle of 45$^{\\circ}$, or an airmass of 1.4.  Section 2 describes the data and the procedures  we use to calibrate the Hectospec spectrophotometry. Because Hectospec's throughput is stable (Section 2.3), this calibration procedure is applicable to all Hectospec observations. In Section 3 we measure the internal errors in the Hectospec spectrophotometry by comparing repeated observations; in Section 4 we estimate the external errors in our spectrophotometry by comparing Hectospec and SDSS spectra. In Section 5 we discuss the impact of calibration issues on determination of the star formation history. We conclude in Section 6. ",
        "conclusions": "The most difficult spectrophotometric issues hampering astrophysical measurements such as the evolution of the star formation rate density affect all instruments, not just those with optical fibers.  These issues include: (1) calibration across different surveys and several instruments, (2) aperture effects on the derived star formation rate, and (3) the highly structured atmospheric absorption bands redwards of 6000 {\\rm \\AA}.  Because of Hectospec's stability and the MMT's f/5 corrector with ADC prisms that eliminate errors from differential refraction, a relatively simple procedure for flux calibration of Hectospec spectra, using infrequent observations of spectrophotometric flux standards for calibration of relative throughput and reference to R-band photometry for absolute throughput, works surprisingly well. Spectrophotometric calibration of instruments without ADC prisms requires observation of secondary spectrophotometric standards in each field. The success of the procedure we describe, when tested against well-calibrated SDSS spectra, demonstrates that short term extinction variations are not a serious issue.   Over its 4.5 year operational lifetime, the Hectospec has acquired over 220,000 spectra that can be calibrated using our procedure.  We recover the SDSS continuum shapes and line fluxes to $\\sim\\pm$10\\% for a subset of the SHELS galaxies also observed by the SDSS, allowing only a rescaling between the different aperture sizes based on R-band aperture magnitudes.  The median emission line fluxes, H$\\delta$ equivalent widths, and the amplitude of the continuum break, d4000, for the SDSS and SHELS spectra agree to within a few percent.  For high signal-to-noise SHELS spectra the typical errors in emission line fluxes are 18\\%.  We plan to use these measures to further explore spectroscopic investigation of the star formation history over the range of redshifts covered by the combined SDSS and SHELS surveys \\citep{Kewley09}."
    },
    "0809/0809.4177_arXiv.txt": {
        "abstract": "{ We have mapped the \\hcn emission from two spiral galaxies, NGC~2903 and NGC~3504 to study the gas properties in the bars. The \\hcn emission is detected in the center and along the bar of NGC~2903. The line ratio \\hcn/\\twelvecoonezero ranges from 0.07 to 0.12 with the lowest value in the center. The enhancement of \\hcn emission along the bar indicates a higher fraction of dense molecular gas in the bar than at the center. The mass of dense molecular gas in the center ($2.2\\times 10^7$ \\msun) is  about 6 times lower than that in the bar ($1.2\\times 10^8$ \\msun). The total star formation rate (SFR) is estimated to be 1.4 \\msunyr, where the SFR at the center is  1.9 times higher than that in the bar. The time scale of consumption of the dense molecular gas in the center is about $\\sim 3\\times 10^7$ yr which is much shorter than that in the bar of about 2 to 10$\\times 10^8$ yr.  The dynamical time scale of inflow of the gas from the bar to the center is shorter than the consumption time scale in the bar, which suggests that the star formation (SF) activity at the center is not deprived of fuel.  In the bar, the fraction of dense molecular gas mass relative to the total molecular gas mass is twice as high along the leading edge than along the central axis of the bar. The \\hcn emission has a large velocity dispersion in the bar, which can be attributed  partially to the streaming motions  indicative of shocks along the bar. In NGC~3504, the \\hcn emission is detected only at the center. The fraction of dense molecular gas mass in the center is about 15\\%. Comparison of the SFR with the predictions from numerical simulations suggest that NGC~2903 harbors a young type B bar with a strong inflow of gas toward the center whereas NGC~3504 has an older bar and has already passed the phase of inflow of gas toward the center. }{}{}{}{} ",
        "introduction": "The molecular interstellar medium in galaxies has been extensively studied through the rotational transitions of CO at millimeter wavelengths. These transitions are good tracers of the molecular gas mass  and represent the general distribution of molecular hydrogen (e.g. Young \\& Scoville 1982, Young \\& Devereux 1991). High-density gas tracer molecules, like HCN, add  relevant information to this view concerning the dense gas ($n_{H_2} > 10^4 cm^{-3}$). Observational evidence (Nguyen et al. 1992, Reynaud \\& Downes 1997, Kohno et al. 1999a) has suggested a close relationship between dense molecular gas and massive star formation in the centers of galaxies. In the center of the starburst galaxies, the HCN line emission is tightly correlated with the radio-continuum emission (see e.g. in NGC~1530, Reynaud \\& Downes 1997) and the total HCN luminosity correlates with the far-infrared (FIR) luminosity (Solomon et al. 1992, Gao \\& Solomon 2004), although contradictory results have been found, e.g. by Aalto et al. (1995). Since only very few galaxies have been mapped in HCN line emission at large scales, this correlation remains a source of speculation, especially  for milder star formation (SF),  typically 2 \\msunyr or less.  The SF in the centers of galaxies and in the spiral arms has been extensively studied.  Nevertheless the processes involved in the SF activity in the bar itself are  still not well  understood. Hydrodynamical and N-body simulations show inflow of gas along the leading edge after  a shock,  losing angular momentum (Athanassoula 1992). Sticky particle simulations of the gas in barred galaxies (Combes \\& Gerin 1985) produce the same configuration of enhanced molecular gas along the leading edge  due to crossing/crowding of orbits. The role of gas orbits in a barred potential on the star formation activity is still poorly understood.  Various observational studies using CO emission have found that molecular gas is located mainly along the leading edge of the bars (e.g Handa et al.  1990, Reynaud \\& Downes 1997, Downes et al. 1996, Sheth et al. 2002). However a contribution from large scale diffuse, unbound gas in the bar of the galaxy NGC~7479 has been suggested by H\\\"uttemeister et al. (2000). In the bar of NGC~1530, Reynaud \\& Downes (1998) have  detected large velocity gradients due to velocity jumps between the upstream regions and the shock along the leading edge. These shocks would inhibit the SF by destroying the giant molecular clouds because of a large shear.   The CO line is a tracer of molecular gas at low densities, whereas star formation occurs in very dense molecular clouds (Solomon et al. 1992, Paglione et al. 1995, Hatchell et al. 1998) which can be traced by the HCN transitions. To study the formation of dense molecular  gas  and its relation to  SF, the bars offer an unique dynamical system  to analyse the effects of a strong density wave on the molecular gas and SF. Moreover, numerical simulations and H$\\alpha$ observations (Martin \\& Friedli 1997, Verley et al. 2007) have shown a tight  correlation between the SF along the bars and the age of the bars, which suggests that  a similar correlation can be expected for the \\hcn emission in the bar because of the tight relationship between the SF and the dense molecular gas.  In this paper we present the result of the  HCN observations at the IRAM-30m telescope, along the bars of two spiral galaxies which have been chosen for their previous detection of HCN in the center. In  Sect. 2, the IRAM observations are described. The main results of the NGC~3504 and NGC~2903 observations are presented in Sect. 3 and 4 respectively. The SF activity is discussed  in Sect. 5 and a discussion about the molecular gas properties in the bar  is presented in Sect. 6. ",
        "conclusions": "  \\begin{itemize}  \\item \\hcn emission has been detected in the center of NGC~3504 and no emission is detected in the bar. The \\hcn to \\twelvecoonezero line ratio is 0.12. The mass of the dense molecular gas in the center is $2.1\\times10^8$ \\msun, which is about 15 \\% of the total mass of the molecular gas.   \\item In NGC~2903, \\hcn emission was detected in the center and  along the bar. The \\hcn to \\twelvecoonezero line ratio along the bar is found to vary from 0.07 to 0.12. This ratio is much lower in the center of the galaxy than in the bar on the northern side, suggesting that the fraction of dense gas in the bar is higher than that in the center. There is also an asymmetry in the \\hcn emission between the northern and the southern side of the bar.   \\item The total dense molecular gas mass in the bar of NGC~2903 is $\\sim 1.2\\times 10^8$ \\msun\\ which is about 6 times larger than the mass at the center of $\\sim 2.2\\times10^7$\\msun. The \\hcn emission is concentrated along the central axis and the leading edge of the bar. The fraction of dense molecular gas mass relative to the total molecular mass is higher along the leading edge  of the bar (11\\%) than on the bar axis ($\\sim$7\\%).  \\item The central velocity of \\hcn emission is systematically lower by about 5 to 20 \\kms in the leading edge of the bar than on the central axis of the bar. This drop is likely due to streaming motion along the bar. A large velocity dispersion of \\hcn emission ranging from 40 to 170 \\kms\\  is observed along the bar. Comparison with the  \\twelvecoonezero rotation curve suggests that this could be due to beam smearing of the rotational velocity over large scales, observed with a low spatial resolution.   \\item Intense SF activity seen in \\halfa\\  is  correlated with the dense molecular gas. The total star formation rate is estimated to be about 1.4 \\msunyr, with 0.67 \\msunyr\\ in the center and 0.36 \\msunyr\\ in the bar of NGC~2903. The consumption time of the dense molecular gas in the center of about 3$\\times 10^7$ yrs is much shorter than that along the bar, ranging from 2 to  10$\\times 10^8$ yrs. The dynamical time for the inflow of gas from the bar to the center is shorter than the consumption time scale suggesting that the inflow of the gas from the bar to the center will sustain SF activity in the center for a longer period.  \\item Comparison of the distribution of \\hcn in the bars of NGC~3504 and NGC~2903 with the results from the numerical simulations of Martin \\& Friedli (1997) suggest that the bars of these galaxies are at different stages of evolution. The bar in NGC~3504 is of type C with  age $\\ge$1 Gyr, with most of the gas already in its center. NGC~2903 harbors a young type B bar with an age between 200 and 600 Myr, with a strong inflow of gas toward the center.   \\end{itemize}"
    },
    "0809/0809.3082_arXiv.txt": {
        "abstract": "We present new \\spitzer, UKIRT and MMT observations of the blue compact dwarf galaxy (BCD) \\mrk, with an oxygen abundance of 12$+$log(O/H)\\,=\\,8.0. This galaxy has the peculiarity of possessing an extraordinarily dense nuclear star-forming region, with a central density of $\\sim$ 10$^6$ cm$^{-3}$. The nuclear region of \\mrk\\ is characterized by several unusual properties: a very red color \\jk\\ = 1.8, broad and narrow emission-line components, and ionizing radiation as hard as 54.9 eV, as implied by the presence of the \\oiv\\ 25.89\\,\\micron\\ line. The nucleus is located within an exponential disk with colors consistent with a single stellar population of age $\\ga$ 1 Gyr. The infrared morphology of \\mrk\\ changes with wavelength. IRAC 4.5\\,\\micron\\ images show extended stellar photospheric emission from the body of the galaxy, and an extremely red nuclear point source, indicative of hot dust; IRAC 8\\,\\micron\\ images show extended PAH emission from the surrounding ISM and a bright nucleus; MIPS 24 and 70 images consist of bright point sources associated with the warm nuclear dust; and 160 \\,\\micron\\ images map the cooler extended dust associated with older stellar populations. The IRS spectrum shows strong narrow Polycyclic Aromatic Hydrocarbon (PAH) emission, with narrow line widths and equivalent widths that are high for the metallicity of \\mrk. Gaseous nebular fine-structure lines are also seen. A CLOUDY model which accounts for both the optical and mid-infrared (MIR) lines requires that they originate in two distinct \\hii\\ regions: a very dense \\hii\\ region of radius $\\sim$ 580 pc with densities declining from $\\sim$ 10$^6$ at the center to a few hundreds cm$^{-3}$ at the outer radius, where most of the optical lines arise; and a \\hii\\ region with a density of $\\sim$ 300  cm$^{-3}$ that is hidden in the optical but seen in the MIR. We suggest that the infrared lines arise mainly in the optically obscured \\hii\\ region while they are strongly suppressed by collisional deexcitation in the optically visible one. The hard ionizing radiation needed to account for the \\oiv\\ 25.89\\,\\micron\\ line is most likely due to fast radiative shocks propagating in an interstellar medium. A hidden population of Wolf-Rayet stars of type WNE-w or a hidden AGN as sources of hard ionizing radiation are less likely possibilities. ",
        "introduction": "Introduction}  Among blue compact dwarf galaxies (BCDs), the dwarf emission-line galaxy Mrk 996 ($M_B$\\,=\\,$-16.9$) occupies a place apart because of the extreme electron density at the center of its star-forming region. {\\sl HST} $V$ and $I$ images show that the bulk of the star formation occurs in a compact roughly circular high surface brightness nuclear region  of radius $\\sim$ 340 pc, with evident dust patches to the north of it \\citep{thuan96}. The nucleus (n) is located within an elliptically (E) shaped low surface brightness (LSB) component, so that Mrk 996 belongs to the relatively rare class of nE BCDs \\citep{loose85}. The extended envelope shows a distinct asymmetry, being more extended to the northeast side than to the southwest side, perhaps the sign of a past merger. This asymmetry is also seen in the spatial distribution of the globular clusters around Mrk 996, these being seen mainly to the south of the galaxy.  Mrk 996 has a heliocentric radial velocity of 1622 km s$^{-1}$ which gives it a distance of 21.7 Mpc, adopting a Hubble constant of 75 km s$^{-1}$ Mpc$^{-2}$ and including a very small correction for the Virgocentric flow \\citep{thuanhi}. At the adopted distance, 1 arcsec corresponds to a linear size of 105 pc. \\citet{thuan96} found the extended LSB component to possess an exponential disk structure with a small scale length of 0.42 kpc. While Mrk 996 does not show an obvious spiral structure in the disk, there is a spiral-like pattern in the nuclear star-forming region, which is no larger than 160 pc in radius.  The UV and optical spectra of the nuclear star-forming region of Mrk 996 \\citep{izotov92,thuan96} show remarkable features, suggesting very unusual physical conditions. The He {\\sc i} line intensities are some 2 -- 4 times larger than those in normal BCDs. In the UV range, the N {\\sc iii}] $\\lambda$1750 and C {\\sc iii}] $\\lambda$1909 are particularly intense. Moreover, the line width depends on the degree of ionization of the ion. Thus low-ionization emission lines such as O$^+$, S$^+$ and N$^+$ have narrow widths, similar to those in other \\hii\\ regions, while high-ionization emission lines such as the helium lines, the O$^{++}$ and Ne$^{++}$ nebular lines and all auroral lines, show very broad line widths, $\\ga$ 500 km s$^{-1}$. Such correlations of line widths with the degree of excitation suggest different ionization zones with very distinct kinematical properties. \\citet{thuan96} found that the usual one-zone low-density ionization-bounded \\hii\\ region model cannot be applied to the nuclear star-forming region of Mrk 996 without leading to unrealistic helium and heavy-element abundances. Instead, they showed that a two-zone density-bounded \\hii\\ region model, including an inner compact region with a central density of $\\sim$ 10$^6$ cm$^{-3}$, some 4 orders of magnitude greater than the densities of normal \\hii\\ regions, together with an outer region with a lower density of $\\sim$ 450 cm$^{-3}$, comparable to those of other \\hii\\ regions, is needed to account for the observed line intensities. The large density gradient is probably caused by a mass outflow driven by the large population of Wolf-Rayet stars present in the galaxy. The gas outflow motions may account for the much broader line widths of the high-ionization lines originating in the dense inner region as compared to the low-ionization lines which originate in the less dense outer region. As for the high N {\\sc iii}] $\\lambda$1750, C {\\sc iii}] $\\lambda$1909 and He {\\sc i} line intensities, they can be understood by collisional excitation of these lines in the high density region. In the context of this model, the oxygen abundance of Mrk 996 is 12+log(O/H) = 8.0. Adopting 12+log(O/H) = 8.65 for the Sun \\citep{asplund05}, then Mrk 996 has a metallicity of 0.22 solar. Mrk 996 shows enhanced helium and nitrogen abundances, which can be accounted for by local pollution from Wolf-Rayet stars.  We present in this paper \\spitzer\\ \\citep{werner04} mid-infrared (MIR) observations of Mrk 996. The extraordinary UV and optical properties of this galaxy, along with the evident presence of dust patches in the star-forming region of Mrk 996 on {\\sl HST} optical images, make it a prime candidate for our Cycle 1 (PID 3139: P.I. Thuan) \\spitzer\\ observations. Our entire program consists of spectroscopic, photometric and imaging observations of 23 BCDs with metallicities ranging from 1/20 to 1/2 that of the Sun, and its main aim is to study star formation in metal-poor environments and to understand how star formation and dust properties change as a function of metallicity and other physical parameters. This paper is the second of our \\spitzer\\ series, the first paper being on Haro 3, the most metal-rich BCD in our sample \\citep{hunt06}.  We present in \\S2 our {\\sl Spitzer} IRAC, MIPS and IRS observations of Mrk 996 and their reduction. We discuss in \\S3 new complementary UKIRT near-infrared (NIR) imaging observations and optical MMT spectroscopic observations. In \\S4, we discuss our imaging results: the IR morphology of the disk of old stars, and the nature of the very red, bright and dense nuclear IR source in Mrk 996. We also discuss the extended PAH emission. In \\S5, we present our spectroscopic results: the PAH features and the IR fine-structure lines. In \\S6, we use the CLOUDY photoionization code \\citep{cloudy,ferland98} to model the observed optical and IR emission-line intensities. We show that it is necessary to postulate two \\hii\\ regions: one which is optically visible where the optical lines arise, and one which is optically hidden where the main part of the IR line emission arises.  We then discuss possible sources of hard ionizing radiation -- fast shocks, WNE-w stars or an AGN --  to account for the presence of the [O {\\sc iv}] $\\lambda$25.9$\\mu$m, line. We summarize our conclusions in \\S7. ",
        "conclusions": "}  We have acquired {\\sl Spitzer} MIR, UKIRT NIR and MMT optical observations of the blue compact dwarf galaxy \\mrk\\ to study its gas, dust and stellar content. This BCD, with a metallicity of about 1/5 that of the Sun, has the peculiarity of possessing an extremely dense nuclear star-forming region: its central density is $\\sim$ 10$^6$ cm$^{-3}$, some 4 orders of magnitude greater than the densities of normal \\hii\\ regions. We have obtained the following results:  \\noindent (1) The nucleus of \\mrk\\ is extremely red, with \\jk\\ = 1.8, and \\hk\\ = 1.0, probably due to very hot dust with a temperature between 600 and 1000 K. The optical spectrum of the BCD  shows the high-ionization lines to have both broad and narrow line components, and a trend of increasing line width with increasing ionization potential.  \\noindent (2) The $VIJHK$ colors of the underlying exponential disk are roughly consistent with the colors of a coeval stellar population with age $\\ga$ 1 Gyr.   \\noindent (3) Like most star-forming galaxies, \\mrk\\ is a composite entity in the IR. We see extended photospheric emission from evolved stars, compact hot dust continuum coming from the nuclear star-forming region at 4.5\\,\\micron, hot dust continuum and extended PAH emission coming mainly from the surrounding less dense ISM at 8\\,\\micron, compact small grain warm dust associated with the active star-forming nuclear region at 24\\,\\micron\\ and 70\\,\\micron, and cooler extended dust emission associated with older stellar populations at 160\\,\\micron.  \\noindent (4) The IRS spectrum (Fig. \\ref{fig:irs}) shows strong Polycyclic Aromatic Hydrocarbon (PAH) molecular emission, with features clearly detected at 5.7, 6.2, 7.7, 8.6, 11.2 and 12.7\\,\\micron. The PAHs in \\mrk\\ are predominantly neutral and small, similar to those found in normal spiral galaxies, suggesting that they reside in the general ISM and not in the star-forming region. The PAH emission features are relatively narrow and their equivalent widths are generally high for the metallicity of \\mrk, exceeding by more than one order of magnitude the values given for the mean EW(PAH)-metallicity relation derived by previous investigators.  \\noindent (5) Gaseous nebular line emission is seen. The IRS spectrum shows several fine-structure forbidden lines, including \\siv\\ $\\lambda$10.51, \\neii\\ $\\lambda$12.81, \\neiii\\ $\\lambda$15.55, \\siii\\ $\\lambda$18.71, 33.48 and \\oiv\\ $\\lambda$25.90\\,\\micron.  \\noindent (6) We have used CLOUDY to model the line-emitting region. To account for both the optical and MIR lines, two \\hii\\ regions are required: a) a very dense \\hii\\ region that is seen in the optical range (Model I) and b) an optically obscured \\hii\\ region ($A_V$ $\\ga$ 4 mag) with a constant number density of $\\sim$ 300 cm$^{-3}$, typical of other \\hii\\ regions (Model II). A two-zone model is required for the non-obscured \\hii\\ region (Model I): a) a very dense nuclear region where the broad optical line components arise; from a central value of $\\sim$ 10$^6$ cm$^{-3}$, the density decreases with distance as $r^{-2}$ until $r\\sim$ 100 pc; b) an outer zone for 100 pc $<r<$ 580 pc, with constant number density of $\\sim$ 400 cm$^{-3}$ . The density gradient is probably caused by large scale gas mass motions, powered by the stellar winds of Wolf-Rayet stars.  \\noindent (7) The UV and optical spectra show no evidence of the very hard ionizing radiation seen in AGN. The IRS spectrum does show however the presence of a faint \\oiv\\ line at 25.89\\,\\micron, indicating the presence of radiation as hard as 54.9 eV in \\mrk. This hard radiation is most likely due to fast radiative shocks in the ISM caused by outflowing stellar winds from WCE and WNL stars and/or by supernova remnants. A hidden AGN, or a population of hidden Wolf-Rayet stars of type WNE-w are are less likely sources of hard ionizing radiation."
    },
    "0809/0809.1939_arXiv.txt": {
        "abstract": "{Our knowledge of the properties of AGN, especially those of optical type-2 objects is very incomplete. Extragalactic source count distributions are dependent on the cosmological and statistical properties of AGN, and therefore provide a direct method of investigating the underlying source populations.} {We aim to constrain the extragalactic source count distributions over a broad range of X-ray fluxes and in various energy bands to test whether the predictions from X-ray background synthesis models agree with the observational constraints provided by our measurements. } {We have used 1129 XMM-{\\it Newton} observations at $|b|>20{^\\circ}$ covering a total sky area of 132.3 ${\\rm deg^2}$ to compile the largest complete samples of X-ray selected objects to date both in the 0.5-1 keV, 1-2 keV, 2-4.5 keV, 4.5-10 keV bands employed in standard XMM-{\\it Newton} data processing and in the 0.5-2 keV and 2-10 keV energy bands more usually considered in source count studies. Our survey includes in excess of 30,000 sources and spans fluxes from $\\sim$${\\rm 10^{-15}}$ to ${\\rm 10^{-12}\\,erg\\,cm^{-2}\\,s^{-1}}$ below 2 keV and from $\\sim$${\\rm 10^{-14}}$ to ${\\rm 10^{-12}\\,erg\\,cm^{-2}\\,s^{-1}}$ above 2 keV where the bulk of the CXRB energy density is produced.} {The very large sample size we obtained means our results are not limited by cosmic variance or low counting statistics. A break in the source count distributions was detected in all energy bands except the 4.5-10 keV band. We find that an analytical model comprising 2 power-law components cannot adequately describe the curvature seen in the source count distributions. The shape of the logN($>$S)-logS is strongly dependent on the energy band with a general steepening apparent as we move to higher energies. This is due to the fact that non-AGN populations, comprised mainly of stars and clusters of galaxies, contribute up to 30\\% of the source population at energies $<$2 keV and at fluxes $\\ge$${\\rm 10^{-13}\\,erg\\,cm^{-2}\\,s^{-1}}$, and these populations of objects have significantly flatter source count distributions than AGN. We find a substantial increase in the relative fraction of hard X-ray sources at higher energies, from $\\ge$55\\% below 2 keV to $\\ge$77\\% above 2 keV. However the majority of sources detected above 4.5 keV still have significant flux below 2 keV. Comparison with predictions from the synthesis models suggest that the models might be overpredicting the number of faint absorbed AGN, which would call for fine adjustment of some model parameters such as the obscured to unobscured AGN ratio and/or the distribution of column densities at intermediate obscuration. } {} ",
        "introduction": "The deepest X-ray surveys carried out to date by {\\it Chandra} (Chandra Deep Field North, CDF-N; Alexander et al.~\\cite{Alexander03} and Chandra Deep Field South, CDF-S; Giacconi et al.~\\cite{Giacconi02}, Lou et al.~\\cite{Lou08}) and XMM-{\\it Newton} (Hasinger et al.~\\cite{Hasinger01}) have resolved up to 90\\% of the Cosmic X-ray background (CXRB) at energies below $\\sim$5 keV into discrete sources reaching limiting fluxes of ${\\rm \\sim2\\times10^{-17}\\,erg\\,cm^{-2}\\,s^{-1}}$ in the 0.5-2 keV band and ${\\rm \\sim2\\times10^{-16}\\,erg\\,cm^{-2}\\,s^{-1}}$ in the 2-8 keV band (Bauer et al.~\\cite{Bauer04}).  However, above $\\sim$5 keV the fraction of CXRB resolved into sources is substantially lower (see e.g. Worsley et al.~\\cite{Worsley04}, Worsley et al.~\\cite{Worsley05}) although precise estimates are hampered by the remaining uncertainty in the absolute normalisation of the CXRB (see e.g. Cowie et al.~\\cite{Cowie02}). Additional uncertainties originate due to variations of the source counts between surveys arising from both the impact of the large scale structure of the Universe on the source distribution (Gilli et al.~\\cite{Gilli03}) and, more mundanely, on cross calibration uncertainties between different missions (Barcons et al.~\\cite{Barcons00}, De Luca \\& Molendi~\\cite{De Luca04}).  Follow-up campaigns targeted at the sources detected in deep-medium X-ray surveys have shown that at high Galactic latitudes Active Galactic Nuclei (AGN) dominate the X-ray sky. At bright X-ray fluxes (${\\rm\\gtrsim 10^{-14}\\,erg\\,cm^{-2}\\,s^{-1}}$), unabsorbed or mildly absorbed AGN, spectroscopically identified as type-1 AGN represent the bulk of the population (see e.g. Shanks et al.~\\cite{Shanks91}, Barcons et al.~\\cite{Barcons07}, Caccianiga et al.~\\cite{Caccianiga07}). At intermediate fluxes, absorbed AGN (optical type-2 AGN) at low redshifts (z$\\lesssim$1) become dominant, while at fluxes ${\\rm\\lesssim 10^{-16}\\,erg\\,cm^{-2}\\,s^{-1}}$ a population of `normal' galaxies starts to emerge (Barger et al.~\\cite{Barger03}, Hornschemeier et al.~\\cite{Hornschemeier03}, Bauer et al.~\\cite{Bauer04}).  \\begin{figure}[] \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{fig1.ps} \\caption{Sky distribution in Galactic coordinates of the selected observations. The high density of pointings in some areas of the sky correspond to planned surveys of relatively large sky areas (e.g. the XMM-{\\it Newton} Large Scale Survey, Pierre et al.~\\cite{Pierre04}).} \\label{skymap}% \\end{figure}  Although the nature of the sources that dominate the CXRB is reasonably clear, there are still large uncertainties in the cosmological and statistical properties of the objects, especially for type-2 AGN for which the redshift and column density distributions are rather poorly determined to date.  One of the most important open issues regarding the population of absorbed AGN is whether the relative fraction of obscured AGN varies with redshift or X-ray luminosity. Some results suggest that this fraction is independent of the X-ray luminosity and redshift (Dwelly \\& Page~\\cite{Dwelly06}), while others point to a decrease in the fraction of absorbed AGN with the X-ray luminosity (Ueda et al.~\\cite{Ueda03}, Barger et al.~\\cite{Barger05}, Hasinger et al.~\\cite{Hasinger05}, La Franca et al.~\\cite{La Franca05}, Akilas et al.~\\cite{Akylas06}, Della Ceca et al.~\\cite{Ceca08}) and/or an increase with redshift (La Franca et al.~\\cite{La Franca05}, Ballantine et al.~\\cite{Ballantine06}, Treister \\& Urry~\\cite{Treister06}). This issue has been recently investigated by Della Ceca et al.~(\\cite{Ceca08}) via the analysis of a complete spectroscopically identified ($\\sim$97\\% completeness) sample of bright X-ray sources (${\\rm >7\\times10^{-14}\\,erg\\,cm^{-2}\\,s^{-1}}$) selected in the 4.5-7.5 keV band. The sources in this study were bright enough to obtain their absorption properties from a detailed analysis of their X-ray spectra. They report a dependence of the fraction of obscured AGN on both the luminosity and redshift, the measured evolution being consistent with that proposed by Treister \\& Urry~(\\cite{Treister06}). A dependence of the fraction of absorbed AGN on the luminosity has also been suggested by observations in the optical and mid-infrared (Simpson et al.~\\cite{Simpson05}, Maiolino et al.~\\cite{Maiolino07}).  X-ray surveys are the best way to understand the properties (i.e. X-ray absorption distributions) and cosmological evolution (i.e. X-ray luminosity functions) of AGN and to test the predictions of the synthesis models of the CXRB (Treister \\& Urry~\\cite{Treister06}, Gilli et al.~\\cite{Gilli07}). However in order to provide strong observational constraints these analyses require complete samples with a high fraction of sources spectroscopically identified. This is a difficult and very time consuming task that can only be achieved for relatively small samples of objects. Source count distributions are dependent on the cosmological and statistical properties of AGN, and therefore provide a rather direct method of investigating the properties of the underlying source populations. Constraining the shape of the source counts is fundamental for cosmological studies of AGN, as it provides strong observational constraints for the synthesis models of the CXRB. The general shape of the source counts in the 0.5-2 keV and 2-10 keV bands is well determined from deep and medium depth X-ray surveys. The results of these analyses show that at fluxes $\\sim$${\\rm 10^{-15}-10^{-14}\\,erg\\,cm^{-2}\\,s^{-1}}$ the source count distributions can be reproduced well with broken power-law shapes (i.e two power-law, hereafter broken power-law, Baldi et al.~\\cite{Baldi02}, Cowie et al.~\\cite{Cowie02}, Cappelluti et al.~\\cite{Cappelluti07}, Carrera et al.~\\cite{Carrera07}, Brunner et al.~\\cite{Brunner08}). However, mostly due to poor statistics, a proper determination of the analytical form of the source count distributions is still unavailable, especially at high energies.  Deep pencil beam surveys are important to study the populations of sources at the faintest accessible fluxes and therefore they are best suited to constrain the faint-end slope of the source counts. However these observations only sample small sky areas and therefore suffer from significant cosmic variance. For example, the CDF-N and CDF-S source counts deviate by more than 3.9$\\sigma$ at the faintest flux levels (Bauer et al.~\\cite{Bauer04}). On the other hand, wide shallow surveys cover much larger areas of the sky and therefore are less affected by cosmic variance. However, they only sample sources at relatively bright fluxes which only contribute a small fraction of the CXRB emission. Surveys at intermediate fluxes, $\\sim$${\\rm 10^{-15}-10^{-13}\\,erg\\,cm^{-2}\\,s^{-1}}$, of the type reported here, fill the gap between deep pencil beam and shallow surveys and sample the fluxes at the break in the source count distributions, i.e. where the bulk of the CXRB energy density should be produced. These surveys are therefore appropriate to accurately determine the position of the break and bright-end slope of source count distributions.  In this paper we put strong constraints on the analytical shape of the extragalactic source count distributions in a number of different energy bands and over a wide range of fluxes. For this purpose we have compiled the largest complete samples of X-ray selected objects to date, ensuring that our results are not limited by low counting statistics or cosmic variance effects. Taking advantage of our large samples, we have investigated how the underlying population of X-ray sources changes as we move to higher energies. Finally we have used the new observational constraints provided by our analysis to check the predictions of current synthesis models of the CXRB.  \\begin{figure*}[!th] \\centering \\hbox{ \\includegraphics[angle=90,width=0.5\\textwidth]{fig2.ps} \\includegraphics[angle=90,width=0.5\\textwidth]{fig3.ps}} \\caption{Left: Distribution of Galactic hydrogen column density (in log units) along the line of sight taken from the 21 cm radio measurements of Dickey \\& Lockman~(\\cite{Dickey90}). Right: Distribution of the exposure times (after filtering).} \\label{nh_gal_dist}% \\end{figure*}   The paper is organised as follows: In \\textsection2 we describe the selection and processing of the XMM-{\\it Newton} observations (\\textsection2.1), the source detection procedure and criteria for selection of sources (\\textsection2.2 and \\textsection2.3), we discuss how we calculated the fluxes of the sources from their count rates (\\textsection2.4) and we explain how the sky coverage was calculated as a function of the X-ray flux (\\textsection2.5). In \\textsection3 we describe the different representations of source counts used in this work (\\textsection3.1), present source count distributions in different energy bands (\\textsection3.2 and \\textsection3.3), discuss the impact on our source count distributions of biases inherent in our source detection procedure (\\textsection3.4), describe the approach used to fit our distributions (\\textsection3.5) and discuss the fractional X-ray background contributed by our sources (\\textsection3.6). In \\textsection4 we summarise the X-ray spectral properties of our objects. The implications of our analysis for the cosmic X-ray background synthesis models are presented in \\textsection5. Finally, the summary and conclusions of our analysis are given in \\textsection6. Appendix A describes the empirical approach used to obtain the sky coverage as a function of the X-ray flux. In Appendix B we compare our source count distributions with those obtained using data taken directly from the {\\tt 2XMM} catalogue. ",
        "conclusions": "\\label{conclusions} We have used the largest samples of X-ray selected sources available to date to provide strong observational constraints on the X-ray source count distributions over a broad range of fluxes and at different X-ray energies. Our source lists were built from 1129 XMM-{\\it Newton} observations at high galactic latitudes, $|b|$$>20{^\\circ}$, covering a total sky area of 132.3 deg$^2$. We have focused our study on four 'narrow bands' and the two 'standard' energy bands, 0.5-2 keV and 2-10 keV, where we have in excess of 30,000 sources. Our data encompass roughly 3 decades of flux, from ${\\rm \\sim10^{-15}\\,erg\\,cm^{-2}\\,s^{-1}}$ to ${\\rm \\sim10^{-12}\\,erg\\,cm^{-2}\\,s^{-1}}$ at energies $\\lesssim$2 keV and more than 2 decades in flux at higher energies, $\\ge$2 keV, from ${\\rm \\sim10^{-14}\\,erg\\,cm^{-2}\\,s^{-1}}$ to ${\\rm \\sim10^{-12}\\,erg\\,cm^{-2}\\,s^{-1}}$. Our sources contribute more than $\\sim$60\\% of the CXRB intensity at energies below $\\sim$2 keV and $\\sim$40\\% above $\\sim$2 keV (although there is a marked decline in the fraction of CXRB resolved across the 2-10 keV bandpass). Thanks to the large size of the samples employed, our results are not limited by cosmic variance effects or low counting statistics. For the first time we have been able to investigate how the changing population of X-ray sources, as we move to different energies, modifies the shape of the measured distributions. The main results are summarised below: \\begin{enumerate} \\item A comparison with previous representative surveys at the fluxes of interest shows overall a good agreement. The largest discrepancies from the comparison were found at bright fluxes ${\\rm \\ge10^{-14}\\,erg\\,cm^{-2}\\,s^{-1}}$, where the results of the majority of the surveys are limited by low counting statistics. Although cross-calibration issues between missions might contribute to the scatter, especially at energies $\\ge$2 keV, we have seen that it cannot fully explain some of the observed discrepancies. We have also shown that different spectral assumptions made in converting count rates to fluxes can also introduce some additional scatter particularly above 2 keV, where the effective areas of the X-ray detectors vary strongly with the energy.  \\item Maximum likelihood fits to our distributions have been carried out using a broken power-law model. A break in the distributions is detected in all energy bands, except in the 4.5-10 keV band, where our survey does not go deep enough to detect the change in slope of the distribution. However, the results of the fits indicate that the measured curvature of our source count distributions cannot be well fitted with a simple broken power-law shape. A model with three power-law components provided a significantly better representation of the shape of our distributions across the set of energy bands.  \\item We find that our source count distributions become significantly steeper both at high and low fluxes as we move to higher energies, where we have shown that the contribution from objects with hard X-ray colours becomes more important. We explain this on the basis of a varying mix of the population of objects at different fluxes and energies and the different sampling of the spectra of AGN. Stars and clusters of galaxies become significantly less important as we move to higher energies, while sources with hard X-ray colours become substantially more important at high energies. We did not find any strong evidence that AGN with no soft flux dominate the source counts at the highest energies sampled by our survey.  \\item We have compared our distributions with the predictions from the CXRB synthesis models from Treister \\& Urry~(\\cite{Treister06}) and Gilli et al.~(\\cite{Gilli07}) which assume different absorption properties for the underlying population of AGN. Once we account for the contribution from clusters and stars at bright fluxes, the models seem to overpredict by 20-30\\% the source counts at energies below 2 keV. However our correction for the contribution from stars and clusters might contribute to the observed discrepancy. The two CXRB models predict different shapes for the source count distributions at faint fluxes, however both models overpredict our source counts (especially the Gilli et al.~\\cite{Gilli07} model) by 10-20\\% at energies below 4.5 keV and at faint fluxes. This result suggest that the synthesis models might overpredict the number of faint absorbed AGN. On the other hand the models seem to underpredict our 4.5-10 keV source counts at the faintest fluxes sampled by our analysis. This could be explained if, as suggested by deep X-ray surveys, the break in the source counts in this energy band is located at lower fluxes than those predicted by the models. The high statistical precision of source counts will allow fine tuning of some model parameters such as the obscured to unobscured AGN ratio and/or the details of the distribution of column densities at intermediate obscuration (${\\rm N_H=10^{22}-10^{23}\\,cm^{-2}}$). \\end{enumerate}"
    },
    "0809/0809.1108_arXiv.txt": {
        "abstract": "Of 104 AGN known to exhibit H$_2$O maser emission, X-ray data that enable estimation of column densities, or lower limits, are available for 42.  Contributing to this, we report analysis of new and archival X-ray data for 8 galaxies and collation of values for  three more.  Maser emission is indicative of large columns of cold gas, and in five of the eight new  cases, maser spectra point toward origins in accretion disks viewed close-to edge-on (a.k.a. ``disk-maser'' systems).  In these, we detect hard continuum and Fe K$\\alpha$ emission with equivalent widths on the order of 1\\,keV, which is consistent with Compton reflection, fluorescence by cold material, and obscuring columns $\\ga 10^{24}$\\,cm$^{-2}$.    Reviewing the full sample of 42, 95\\% exhibit N$_{\\rm H} >10^{23}$\\,cm$^{-2}$ and 60\\% exhibit N$_{\\rm H} >10^{24}$\\,cm$^{-2}$.  Half of these are now recognized to be disk masers (up from 13); in this sub-sample, which is likely to be more homogeneous vis-\\'a-vis the origin of maser emission, 76\\% exhibit N$_{\\rm H} >10^{24}$\\,cm$^{-2}$.  The probability of a common parent distribution of columns for disk-masers and other AGN masers is $\\la 3\\%$.   Because ground-based surveys of AGN to detect new disk masers are relatively unbiased with respect to X-ray brightness and comparatively inexpensive, they may also be efficient guides for the sensitive pointed X-ray observations required to identify Compton-thick objects outside of shallow surveys. ",
        "introduction": "Extragalactic H$_2$O maser emission at 22 GHz is known in 104 active galactic nuclei (AGN)(\\cite{zha06}, \\cite{kon07}, \\cite{braatz08}, Braatz et al. 2008, in preparation). In early analyses, the majority of the nuclei for which pointed X-ray observations exist  are Compton thick or nearly so, with columns (N$_{\\rm H}$) $\\ga 10^{24}$ cm$^{-2}$ \\citep{mad06, zha06, gree07}. \\citet{zha06} consider  all maser-host galaxies, where the emission is associated with either nuclear activity or star formation. In a sub-sample of 26 with apparent maser luminosity $>10$ L$_\\odot$ (a.k.a. ``megamasers''),   85\\% exhibited N$_{\\rm H}$ $>10^{23}$ cm$^{-2}$ and 50\\% exceeded $10^{24}$ cm$^{-2}$.  \\citet{mad06} focus on a probably more homogeneous sub-sample for which maser spectra suggest the emission specifically traces rotating, highly inclined disk structures close to the central engines.  (We refer to these as ``disk masers.'') Seven of 9 nuclei with disk masers were Compton thick.  It is plausible that AGN which host masers are characterized by high N$_{\\rm H}$  in general, because large reservoirs of molecular gas along the line of sight are required to produce substantial maser amplification \\citep{eli82}.  Among disk masers, prevalence of high N$_{\\rm H}$ may be anticipated because  the emission is believed to arise in inclined disk structures.   We have analyzed new, archival, and published X-ray data for maser AGN and examine the statistics of the updated sample, now about one third larger than that of \\citet{zha06}.  Masers show promise as signposts of Compton-thick AGN that are readily detected from the ground.  In $\\S2$, we describe new spectral analyses and estimation of N$_{\\rm H}$.  We discuss selection effects in the full sample and inferences about maser AGN in $\\S3$. ",
        "conclusions": "We have combined published estimates and limits on N$_{\\rm H}$ with results of our analyses to obtain a sample of 42 AGN (out of 104 known maser hosts).  The sample includes: (1) 31 AGN that exhibit megamaser and weaker emission \\citep{zha06}, (2) two AGN listed by  \\citet{zha06} but without known columns  -- NGC\\,591, Mrk\\,1, (3) five AGN treated here -- Tab.\\,2, and (4) four nuclei in which maser emission has been recently discovered -- NGC\\,17 \\citep{kon07}; NGC\\,3081, NGC\\,4253, and NGC\\,3783 (Braatz \\etal, in preparation). Columns are taken from Zhang \\etal -- 31 objects, except NGC\\,5256 \\citep[cf.][]{gua05}; this work -- 5 objects; \\citet{gua05} -- Mrk\\,1, NGC\\,17, and NGC\\,591; \\cite{ree04} -- NGC\\,3783; \\citet{mai98} -- NGC\\,3081; and \\cite{pounds03} --  NGC\\,4253. In order of precedence, we adopt results from SAX, XMM, Chandra and ASCA, except for NGC\\,3079, where we judged the XMM lower limit \\citep{cappi06} to be more consistent with SAX data  than larger limits based on SAX data alone \\citep{iyo01}.  Of the 42 AGN, 60\\% (25) are Compton thick (Tab.\\,\\ref{tabnh}; Fig.\\,\\ref{fig2}).  This fraction is consistent with that obtained for the smaller sample of \\citet{zha06}, 58\\% (18 of 31). Within the current sample of 42, a greater fraction of  recognized disk-maser galaxies are Compton thick, 76\\% (16 of 21).  The distribution of columns for disk masers is significantly skewed and non-Gaussian (Fig.\\,\\ref{fig2}).  Substitution (over time) of N$_{\\rm H}$ estimates for lower limits will exagerrate this.  In contrast, the distribution for AGN masers without evidence of origins in  edge-on disks includes a larger proportion of Compton-thin sources.  A Kolmogorov-Smirnov (KS) test indicates a $< 3.2\\%$ probability that  the two distributions are drawn from a single parent. To account for lower limits in our test, we constructed a Monte Carlo simulation of 1000 trials, sampled a uniformly distributed random variable of $10^{24}$ and $10^{26}$\\,cm$^{-2}$ for each galaxy with a limit, and compiled a distribution of KS statistics.   Boundaries of $10^{25.5}$ and  $10^{28}$ yielded probabilities of $\\la3\\%$ as well.  Distinctions between N$_{\\rm H}$ distributions for AGN that host acknowledged disk masers and those that do not may be still greater than estimated. Maser classifications depend on identification of spectroscopic markers: highly red and blue-shifted emission bracketing the systemic velocity.  Disks may be oriented close to edge-on, resulting in large columns toward the central engine, while these spectroscopic markers may nonetheless be absent (e.g., high-velocity masers may be beamed away from the observer due  to warping; cf. Miyoshi \\etal\\ 1995).  \\begin{deluxetable}{lccc} \\tablecolumns{4} \\tablewidth{3in} \\tablewidth{4.4in} \\tablecaption{Column Measurements Among Known Maser AGN} \\tablehead{ \\colhead{Sample} & \\colhead{$>10^{23}$ cm$^{-2}$} & \\colhead{$>10^{24}$ cm$^{-2}$} & \\colhead{No.} } \\startdata Disk masers\\tablenotemark{(a)}   & 100\\%  & 76\\%  & 21\\\\ All AGN masers                                & 95\\% & 60\\% & 42 \\\\ Megamasers\\tablenotemark{(b)}  & 87\\%  & 58\\%  & 31 \\\\ \\vspace{-0.08in} \\enddata \\tablenotetext{(a)}{In addition to five sources in Tab.\\,2, we recognize 16 other reported disk-maser AGN with accompanying estimates of N$_{\\rm H}$ or limits: NGC\\,591, NGC\\,1068, NGC\\,1386, NGC\\,2273, NGC\\,3079, NGC\\,3393, NGC\\,4051, NGC\\,4258, NGC\\,4388, NGC\\,4945, NGC\\,5728, IC\\,2560, Mrk\\,1, Mrk\\,1210, \\hbox{Circinus, and 3C\\,403.} } \\tablenotetext{(b)}{We include 26 megamasers with estimates of N$_{\\rm H}$ or limits as listed by \\citet{zha06}, four objects in Tab. 2 (NGC\\,1194, NGC\\,6926, Mrk\\,34, Mrk\\,1419), and Mrk\\,1. } \\label{tabnh} \\end{deluxetable}  \\citet{zha06} discussed an alternate sub-sample, selected based on apparent maser luminosity (i.e., megamasers).  Within the current sample of 42 AGN with masers, we count 58\\% (18 of 31) of recognized megamaser hosts are Compton thick (13 from \\citet{zha06}, plus NGC\\,1194, NGC\\,6926, Mrk\\,1, Mrk\\,1419, and Mrk\\,34), matching the incidence among maser AGN broadly (Tab.\\,\\ref{tabnh}).  We link the larger incidence of columns $>10^{24}$\\,cm$^{-2}$ among disk-masers to a simple model geometry (edge-on orientation) and contrast this with selection based on maser luminosity, which is nearly always estimated with the convenient but inaccurate assumption of isotropic emission, without knowledge of beam angles and orientation to the line of sight.  Columns $\\ga 10^{23}$ cm$^{-2}$ are more  common in maser AGN than in a sample of 22 hard X-ray (14-195\\,keV) selected, similarly nearby AGN observed with SWIFT  \\citep{win08}, which includes no maser hosts (95\\% vs 50\\%; Tab.\\,\\ref{tabnh}). However, for N$_{\\rm H}>10^{24}$\\,cm$^{-2}$, incidence rates are similar (60\\% vs 50\\%). For masers in all classes of AGN, we estimate a statistically similar frequency of Compton-thick columns to that reported by \\citet{gua05} and \\citet{ris99} among Seyfert\\,2 objects: $46\\%$ (22 of 48) and 65\\% (22 of 34), respectively.  However, this is somewhat below the rate for masers associated with  Seyfert\\,2 nuclei, 74\\% (20 of 27).  The difference is suggestive but comparison of column distributions of maser systems and local type-2 AGN is complicated by sample incompleteness in luminosity and distance, overlap between the X-ray and maser samples (e.g., 15/45 sources for Risaliti \\etal\\ 1999), and uncertainty in X-ray luminosities for obscured AGN (though aperture and extinction-corrected [OIII] luminosities could be used in place of X-ray estimates, e.g., Mel\\'endez \\etal\\ 2008).  Demonstrated correlation between incidence of maser emission and high N$_{\\rm H}$ may suggest a new strategy for future efforts to identify obscured AGN.  This is of particular interest if such objects are indeed substantial contributors to the Hard X-ray Background.  Heavily obscured  AGN typically have relatively low apparent luminosities and require long  integrations  (e.g., $> 100$ ksec  to obtain $>10^3$ counts for $cz=10^3-10^4$\\,km\\,s$^{-1}$).  In contrast, maser emission in AGN as distant as $cz\\sim10^4$\\,km\\,s$^{-1}$  can be detected  with integrations on the order of 1\\,ksec, using a well-instrumented 100m-dish antenna.  We characterize the relative efficiency of maser and pointed X-ray observations, $R$, by the ratio of Compton-thick AGN discoverable in a fixed time by each means, for a sample of optically identified type-2 objects.  We obtain $R\\sim 3$ from $[10^5 f_m M_{CT}] / [10^3 X_{CT}]$, where $f_m$ is the maser detection rate  ($\\sim 0.03$; Kondratko \\etal\\ 2006), $M_{CT}$ is the  Compton thick fraction of maser AGN ($\\sim 0.6$; Tab.\\,\\ref{tabnh}), and $X_{CT}$ is the fraction of type-2 AGN found to be Compton thick in X-ray studies ($\\sim 0.5$; e.g., Risaliti \\etal\\ 1999).  Deep, wide-field, hard X-ray surveys (e.g., NuSTAR) will extend synergy with maser studies, creating statistically significant X-ray and joint X-ray/radio samples.  Nonetheless, there will be practical constraints on survey fields ($5^\\circ\\times5^\\circ$ for NuSTAR), and all-sky radio observations with well-defined completeness criteria could usefully guide pointed follow-up to enrich X-ray samples.  Strengthened correlation between disk-maser emission and large obscuring columns begs the question, is maser gas in disks responsible for obscuration in type-2 objects \\citep{greenhill03,mad06}?  On the one hand, obscuration by cold disk gas at radii of 0.2-0.3 pc is inferred in NGC 4258 \\citep{fru05, herrnstein05}, and for Circinus along other lines of sight \\citep{greenhill03}.  On the other hand, time-variable N$_{\\rm H}$ data indicate obscuration at radii of a few$\\times0.01$ pc toward NGC\\,1365, NGC\\,4388, and NGC\\,4151 \\citep{elv04, ris07, puccetti07}.  Direct physical association of maser and obscuring gas volumes still appears to be case-specific.  Highly inclined structures that give rise to disk masers probably span a range of radii, at least extending inward with accretion flows.  Where these cross the line of sight to the central engine, and how extensively, likely depends on the detailed geometry of each system, requiring case study to establish."
    },
    "0809/0809.4207_arXiv.txt": {
        "abstract": "We present a new stellar evolution code and a set of results, demonstrating its capability at calculating full evolutionary tracks for a wide range of masses and metallicities. The code is fast and efficient, and is capable of following through all evolutionary phases, without interruption or human intervention. It is meant to be used also in the context of modeling the evolution of dense stellar systems, for performing live calculations for both normal star models and merger-products.  The code is based on a fully implicit, adaptive-grid numerical scheme that solves simultaneously for structure, mesh and chemical composition. Full details are given for the treatment of convection, equation of state, opacity, nuclear reactions and mass loss.  Results of evolutionary calculations are shown for a solar model that matches the characteristics of the present sun to an accuracy of better than 1\\%; a $1\\ \\Msun$ model for a wide range of metallicities; a series of models of stellar populations I and II, for the mass range $0.25$ to $64 \\Msun$, followed from pre-main-sequence to a cool white dwarf or core collapse. An initial final-mass relationship is derived and compared with previous studies. Finally, we briefly address the evolution of non-canonical configurations, merger-products of low-mass main-sequence parents. ",
        "introduction": "\\label{intro}  Simulating the evolution of a star requires the solution of a set of partial differential equations with boundary conditions at the center and surface, involving extensive input physics, such as equations of state, nuclear reactions, opacities, as well as recipes for treating convection, mass loss, or material mixing. This is accomplished by complex computer codes that are time consuming and depend on a large number of adjustable parameters, both physical and numerical, needed for dealing with evolutionary phases that are different in nature. From the formation to the death of a star, differences between evolutionary phases are so large, that studies are usually devoted to---and often codes are devised for---a specific part of a star's life, ignoring or simplifying, or suppressing others. So far, no code has been suited or applied to obtain complete, unabridged evolutionary tracks over the entire range of stellar masses and metallicities, although many have come close to accomplishing this task (e.g. \\citealt{1995MNRAS.274..964P}, \\citealt{1998MNRAS.298..525P}).  % For example, most (if not all) evolution codes crash at the helium core flash phase. Most of the stellar evolution codes do not solve simultaneously for the structure and the composition; this introduces serious errors in some critical phases whenever the mass grid $\\{m_i\\}$ changes, as it must eventually \\citep{2006MNRAS.370.1817S}. Our aim has been to develop a versatile and robust stellar evolution code that is free of such handicaps.  A further demand on the code is efficiency and speed. Furthermore, it should be capable not only of evolving any star through all phases without intervention, but also of dealing with peculiar objects. Such a code could be incorporated into an N-body code that deals with dense stellar systems, if not at present, then---given the rapid and continual advance in computing power---in the foreseeable future. The computation methods of N-body gravitating systems have undergone a revolutionary development owing to the work of \\citet{1963MNRAS.126..223A} (see review by \\citealt{1999PASP..111.1333A}) and gaining impetus in the past two decades (e.g. \\citealt{2003gmbp.book.....H}, \\citealt{2001MNRAS.323..630H}, \\citealt{2005MNRAS.363..293H}): not only have new algorithms been developed, capable of dealing with dense stellar systems (e.g. \\citealt{2001MNRAS.321..199P}, \\citealt{2004MNRAS.351..473P}), but also special hardware has been constructed under the GRAPE (GRAvity PipE) project \\citep{1997ApJ...480..432M}.  However, in order to render these sophisticated N-body calculations realistic, the effect of the structure and evolution of the constituent stars must be considered as well. This led, less than a decade ago, to the development of the MODEST (MOdelling DEnse STellar systems) project, whose aim is to combine N-body dynamics with the hydrodynamics of stellar collisions on the one hand, and with stellar evolution of the cluster population, on the other (see \\citet{2003NewA....8..337H}). So far, studies of stellar systems have resorted to short cuts based on sets of discrete pre-calculated evolutionary tracks: either interpolating between them, or using parametrized fit formulae. Clearly, this procedure is incapable of dealing with `non-canonical' stars, the outcome of collisions and mergers.  In this paper we thus present a new evolutionary code that we have developed with this aim in mind. The outline of the code and method of solution are presented in the next section, Section~\\ref{code}; the input physics is described in some detail in Section~\\ref{input-phys}, and results of representative calculations are discussed in Section~\\ref{res}. ",
        "conclusions": ""
    },
    "0809/0809.1436_arXiv.txt": {
        "abstract": "We present continued results from a wide-field, $\\sim$150 deg$^2$, optical photometric and spectroscopic survey of the northern part of the $\\sim$5 Myr-old Upper Scorpius OB Association. Photometry and spectral types were used to derive effective temperatures and luminosities and place newly identified association members onto a theoretical Hertzsprung-Russell diagram. From our survey, we have discovered 145 new low mass members of the association, and determined $\\sim$10\\% of these objects to be actively accreting material from a surrounding circumstellar disk. Based on comparison of the spatial distributions of low and high mass association members, we find no evidence for spatial segregation by mass within the northern portion of the association. Measured data are combined with pre-main sequence evolutionary models to derive a mass and age for each star.  Using Monte Carlo simulations we show that, taking into account known observational uncertainties, the observed age dispersion for the low mass population in USco is consistent with all stars forming in a single burst $\\sim$5 Myr ago, and place an upper limit of $\\pm$3 Myr on the age spread if the star formation rate has been constant in time. We derive the first spectroscopic mass function for USco that extends into the substellar regime, and compare these results to those for three other young clusters and associations. ",
        "introduction": "\\label{cha7:intro}  The review of \\citet{1987ARA&A..25...23S} laid the basic framework for how isolated stars are created. However, most stars do not form in isolation but rather in groups (i.e., clusters or associations; \\citealt{1993AJ....105.1927G}, \\citealt{1995AJ....109.1682L}, \\citealt{2000AJ....120.3139C}), and many details of the processes involved in stellar group formation remain unexplained.  In particular, do all members of an individual group form in a single burst, or is star formation a lengthy process? The distribution of stellar masses formed within a stellar group is known as the initial mass function (IMF).  Is the IMF universal, or does it vary with either star formation environment or time?  The phenomenon of rapid (1--2 Myr) clustered star-formation has been found in almost all nearby young associations (e.g., \\citealt{2003ARA&A..41...57L}). The large numbers of very young stars and apparent lack of more evolved (5--10 Myr-old) objects in star forming regions contrasts with ages of a few tens of megayears (e.g., Blitz \\& Shu 1980) inferred for molecular clouds. Either star-formation takes place for only a small fraction of the cloud lifetime, molecular clouds themselves live only a few megayears (e.g., Hartmann et al. 2001), or older group members have been missed in observational surveys which are often biased towards stars that are still actively accreting or which possess an optically thick disk. Observational surveys capable of detecting the entire age range present within a stellar group are needed if we are to begin to understand the time dependence of star formation within a molecular cloud.   The basic structure of the field star IMF in the mass range $\\sim$0.5 M$_\\odot$ to 10 M$_\\odot$ is well characterized by a power law with slope $dN/dM \\propto M^{-2.35}$, as derived originally by \\cite{1955ApJ...121..161S}. However, while the shape of the stellar mass function between $\\sim$0.5 M$_\\odot$ and 10 M$_\\odot$ is near universal, below $\\sim$0.5 M$_\\odot$ the shape of the mass function may be strongly dependent on environment within the parental molecular cloud (\\citealt{2003ApJ...593.1093L}, \\citealt{2004ApJ...610.1045S}). Young stellar clusters and associations are particularly valuable for examining the shape of the low mass IMF because contracting low-mass pre-main sequence stars and brown dwarfs are 2-3.5 orders of magnitude more luminous than their counterparts on the main sequence. Thus, despite the fact that they are farther away, young stars and brown dwarfs can be more readily detected than stars of equivalent mass in the field.   The Upper Scorpius OB Association (USco) is the closest (145 pc; \\citealt{1999AJ....117..354D}) young OB association to the Sun with 120 members more massive than $\\sim$1 M$_\\odot$, and an associated population of $\\sim$400 known low mass objects (e.g., \\citealt{1998A&A...333..619P}). Because of its poximity, and large population that spans the range from $<$0.02 M$_\\odot$ to $>$10 M$_\\odot$, USco is an ideal region in which to study the mass distribution of stars in OB associations. Furthermore, at an age of $\\sim$5 Myr \\citep{2002AJ....124..404P}, USco is young enough that we may be able to measure the ages of member stars to a precision of less than a million years, and old enough that an on-going episode of star formation lasting several millions of years could be detectable. A major difficulty faced by studies of the USco region is that the high mass members alone span $>$300 deg$^2$ on the sky. Obtaining a complete census of the association's low mass population is thus a formidable task as one must identify faint objects over a very large spatial region.  While there exist several techniques to identify young stars, many observational signatures (e.g., H$\\alpha$ emission or near-infrared excess) probe accretion processes and circumstellar material rather than characteristics intrinsic to the stars themselves.  Recently, deep, multicolor imaging surveys combined with spectroscopic follow-up have proven successful in identifying low mass stars and brown dwarfs without active accretion, as well as classical T-Tauri type objects. Young pre-main sequence (PMS) objects still undergoing contraction are cooler (i.e., redder in color) and/or more luminous (i.e., brighter in magnitude) than their main sequence counterparts.  In nearby regions, candidate PMS stars can be identified through photometric colors and magnitudes that are systematically different from those of the bulk field star population. Spectroscopic follow-up observations allow assessment of surface gravity diagnostics which can be used to distinguish bonafide young PMS stars from foreground field dwarfs and reddened background giants. Previous imaging and spectroscopic surveys in USco include studies by \\cite{2001AJ....121.1040P} and \\cite{2002AJ....124..404P}, whose work yielded 166 new low mass PMS objects. \\cite{2008MNRAS.383.1385L}, \\cite{2004AJ....127..449M}, and \\cite{2000AJ....120..479A} together identified $\\sim$70 members of USco with spectral type M6 or later. Thus far, over 400 low mass ($M<$0.6 M$_\\odot$) members have been identified in USco through X-rays, H$\\alpha$ emission, photometry and/or spectroscopy. However, most searches have been limited to small subregions (smaller than a few square degrees) or bright objects ($R$ $\\lesssim$ 16 mag). Given the derived mass function and number of observed high mass stars in USco, and assuming the high and low mass objects share the same spatial distribution, \\cite{2002AJ....124..404P} estimate the entire USco region should contain $>$1500 young, low mass objects with $M<$0.6 M$_\\odot$, most of which are yet to be discovered.  In Slesnick, Carpenter, \\& Hillenbrand (2006a; hereafter Paper~I), we introduced our wide-field photometric survey covering $\\sim$150 deg$^2$ of USco. In that work, we presented optical spectroscopic observations of 62 photometrically-selected candidate PMS stars, of which 43 were confirmed by us to be new USco members. In this companion paper we present spectra for an additional 178 candidates observed at either Palomar Observatory or Cerro-Tololo Inter-American Observatory (CTIO). In \\S\\ref{cha:6:sec:obs}, we give a brief overview of the photometric and spectroscopic surveys, and present new USco members discovered as a result of recent observations.  For the remainder of the paper beyond \\S\\ref{cha:6:sec:obs}, we discuss results for new members presented here together with those for association members identified in Paper~I. In \\S\\ref{cha:6:sec:discussion:sub:em}, we discuss H$\\alpha$ emission profiles and accretion signatures in USco. In \\S\\ref{cha:6:sec:discussion:sub:spat}, we compare the spatial distribution of the high and low mass populations. Finally, in \\S\\ref{discussion}, we construct an Hertzsprung-Russell (HR) diagram and present a discussion of the age and mass distributions derived for the members of USco presented in our work, as well as a comparison of distributions derived for other nearby star forming regions. ",
        "conclusions": "The most prominent emission line observed in the spectra of new members is H$\\alpha$ which, seen in the spectra of young stars and brown dwarfs, is created via one of several mechanisms. Weak, narrow H$\\alpha$ lines are presumed to originate from active chromospheres whereas strong, broad and/or asymmetric lines can be produced from high-velocity, infalling accretion or strong winds. \\cite{2003AJ....126.2997B} have proposed an empirical, spectral type-H$\\alpha$ equivalent width ($W$(H$_\\alpha$)) relation to describe the upper limit of non-accreting stars and brown dwarfs based on the chromospheric saturation limit observed in the Pleiades, $\\alpha$ Per, and IC 2391 open clusters.  These clusters are sufficiently old (50--125 Myr) that accretion should not be present, and any observed H$\\alpha$ emission is assumed to be produced entirely from chromospheric activity. Figure~\\ref{fig:cha6halpusco} plots measured H$\\alpha$ equivalent widths as a function of spectral type for the 145 members of USco presented here, shown with the \\cite{2003AJ....126.2997B} empirical accretor/nonaccretor division. Notably, every new member identified in our work shows H$\\alpha$ in emission. Many stars and brown dwarfs exhibit very strong H$\\alpha$ emission  (see also Table~1) at levels substantially above the accretor/nonaccretor division, and thus are possibly still undergoing active accretion.  We determined an empirical criterion for identifying objects with H$\\alpha$ excess emission based on our data. At an age of $\\sim$5 Myr, we assume the bulk of our sample is no longer accreting and compute median values of H$\\alpha$ emission as a function of spectral type using 1-sigma clipping to remove any bias from outliers. For most bins, we define a star to have an H$\\alpha$ excess if it exhibits emission at a level greater than 3$\\sigma$ above the median value for its spectral type, where $\\sigma$ is the dispersion about the median for stars at a given spectral type. Stars with such strong H$\\alpha$ emission are likely to be accreting and are assumed such for the remainder of this work. These sources are distinguished on Figure~\\ref{fig:cha6halpusco}.  For spectral type bins earlier than M4 and later than M7, we have identified fewer than 5 stars per spectral type. Thus, for these bins we do not have enough measurements to derive a statistically representative value for median H$\\alpha$ emission and we do not consider any stars in these bins to be in our sample of H$\\alpha$ excess, accreting sources. However, we note that two M8 stars (including the possible binary discussed in Paper~I \\S4.2) sit above the \\cite{2003AJ....126.2997B} accretor/nonaccretor dividing line.    Based on the above criterion, we find 15 objects to exhibit H$\\alpha$ emission with sufficient strength to be considered by us to be actively accreting.  Spectra for accreting stars and brown dwarfs are shown in Figure~\\ref{fig:cha6hap}. In addition to H$\\alpha$ emission, many of these stars (SCH J16014156-21113855, SCH J16222156-22173094, SCH J16150524-24593542, SCH J16033470-18293060, SCH J16075565-24432714, SCH J16284706-24281413, SCH J16103876-18292353, SCH J16110739-22285027, SCH J16060391-20564497) also exhibit He I (6678 \\AA) emission which is commonly seen in spectra of classical T-Tauri type objects. Two accreting sources lie very close (within $\\sim$1$^\\circ$) to the young ($\\lesssim$1 Myr) $\\rho$ Ophiuchi molecular cloud ($\\rho$Oph; $\\alpha$=16 25 35.118, $\\delta$= -23 26 49.84 J2000). However, because $\\rho$ Oph lies slightly in front of USco \\citep{2008ApJ...675L..29L}, if these stars were dynamically ejected $\\rho$ Oph members, we would expect to see them exhibit systematically higher luminosities than USco members of similar spectral type. Based on the derived HR diagram from these data (see \\S\\ref{cha:6:sec:discussion:sub:hr}) this phenomenon is not observed, and we include these two stars in our sample of accreting members of USco.  We find at an age of $\\sim$5 Myr (see \\S\\ref{cha:7:sec:usco}), 15/145 (or $\\sim$10$^{+3}_{-3}$\\%) of low mass association members (spectral type $\\leq$M7 and $\\geq$M4) are observed to be accreting based on the strength of H$\\alpha$ emission present in their spectra.  If we were to use the \\cite{2003AJ....126.2997B} accretion boundary instead of our own empirical classification, we would have determined 23/145 (16\\%) low mass stars and brown dwarfs in our sample are still accreting. In comparison, \\cite{2006A&A...446..485G} find $\\sim$65\\% (20/31) of 1 Myr-old low mass objects in the subclusters of Taurus to be actively accreting based on the strength of H$\\alpha$ emission observed in their spectra compared to the \\cite{2003AJ....126.2997B} accretion boundary. Thus, a significant fraction of very low mass stars and brown dwarfs must stop accreting between 1 and 5 Myr.   This conclusion is consistent with a median accretion lifetime of $\\sim$2--3 Myr for higher mass stars (\\citealt{2001ApJ...553L.153H}, \\citealt{2005astro.ph.11083H}).     Figure~\\ref{fig:cha6spath} shows the 2D spatial distribution of the 120 high mass members of USco identified in the Hipparcos survey \\citep{1999AJ....117..354D}.  This sample represents the complete population of known members more massive than $\\sim$1 M$_\\odot$. The density of high mass stars is roughly constant from 237$\\lesssim \\alpha \\lesssim$249$^\\circ$ and peaks at $\\delta \\sim$-24$^\\circ$. As can be seen, despite the large area of the Quest-2 survey, it still encompassed only the central $\\sim$13$^\\circ$ in RA and the northern $\\sim$12$^\\circ$ in DEC of the association.    In general, the low mass PMS stars presented here share a common spatial distribution with the high mass Hipparcos members.  Efforts to observe northwest of the Hipparcos stars largely yielded reddened field dwarfs rather than young association members. To correct for bias in the spatial area we observed spectroscopically, we first computed the percent of photometric candidates that we observed spectroscopically, in 1-degree bins, as a function of right ascension and declination. We found that we observed spectroscopically a maximum of $\\sim$25\\% of the photometric candidates with a given degree-wide spatial bin. We thus corrected every bin to a uniform 25\\% of candidates observed, and calculated the number of members we would have detected assuming we had observed 25\\% of the photometric candidates in each bin, and that the percentage of identified members relative to the number of stars observed spectroscopically at each spatial bin would remain unchanged. Figure~\\ref{fig:cha6spatp1} shows the resultant 1D spatial distributions for the low mass association members discussed here, together with those for the 56 Hipparcos stars that fall within our survey area. We conclude that the density of low mass association members found in the Quest-2 survey is roughly uniform in RA, and peaks at $\\delta \\sim$-25$^\\circ$ with stellar densities falling off beyond these values. We find no evidence for spatial segregation by mass in USco within the northern portion of the association.        In this section, we combine the spectral type and photometry of each new member to derive values for its luminosity and effective temperature, and place it onto an HR diagram. As described in \\S\\ref{obs:phot}, the final Quest-2 photometry is not on a standard magnitude system. Thus, because of the reliability and uniformity of the 2MASS survey, we chose to use $J$-band magnitudes and $(J-H)$ colors to derive luminosities. An empirical fit to BC$_J$ as a function of spectral type was determined from the observational data of \\cite{1996ApJS..104..117L} and \\cite{2002ApJ...564..452L} (spectral types M1-M6.5 and M6-L3, respectively). We derived intrinsic colors, extinction, and effective temperatures using the methods described in \\cite{2004ApJ...610.1045S}.  In Figure~\\ref{fig:cha6hrusco}, we present an HR diagram for the 145 low mass members of USco that we identified, shown with PMS model tracks and isochrones. The most commonly used PMS models for low mass stars and brown dwarfs are those derived by \\cite{1997MmSAI..68..807D} (hereafter DM97) and \\cite{1998A&A...337..403B}, which differ primarily in their atmospheric approximations and treatment of convection. Both models suggest similar mass ranges for our data of 0.02M$_\\odot$ $< M <$ 0.2M$_\\odot$, though predicted masses for individual objects can vary by up to 0.09M$_\\odot$ (60\\%). We have found that the slope of the DM97 isochrones provide a reasonable match to the derived HR diagram, whereas the \\cite{1998A&A...337..403B} models predict systematically younger stars at lower masses. A similar result was also noted by \\cite{2008ASPC..384..200H}.  These authors compared the slope of six different theoretical PMS isochrones as a function of binary fraction, and found that, assuming an intermediate binary fraction consistent with observations (e.g., \\citealt{2008ApJ...679..762K}), the DM97 models provide the best match to the observed HR diagram slope for stars in USco.  Thus, for the remainder of this work, we use the DM97 mass tracks and isochrones to derive mass and age for stars in our sample. All derived quantities are given in Table~2.         \\subsection{Age Distribution of the Low Mass Population in USco} \\label{cha:7:sec:usco}  Literal interpretation of the derived HR diagram (Figure~\\ref{fig:cha6hrusco}) reveals a population with median age of $\\sim$4.1 Myr, and a continuous spread of ages over $\\gtrsim$10 Myr. While this result $may$ be real, the continuous nature of the observed age distribution in USco could also be produced from uncertainties in observed parameters. Unlike the HR diagram, the observed surface gravity-sensitive spectral features do not show evidence for a large spread in age between association members.  Figure~\\ref{fig:cha7uscospec} shows spectra of two stars with spectral type M5 identified in the survey.  The top spectrum is that of the `youngest' M5 star observed spectroscopically at Palomar (SCH16054416-21550566), $\\sim$2.6 Myr-old based on its location on the HR diagram (see also Table~2); the bottom spectrum is of the `oldest' M5 star observed at Palomar (SCH16162599-21122315), $\\sim$14.4 Myr based on its location on the HR diagram. As discussed in SCH06 and \\cite{mythesis}, young stars less than a few megayears old exhibit systematically less Na I ($\\lambda$8190 \\AA) absorption than do intermediate-age stars ($\\sim$5--10 Myr) due to their lower surface gravity (see Figure~10 in Paper~I and Figure~6 in SCH06).  Thus, if the derived ages from the HR diagram are correct, the spectrum of the 14.4 Myr-old star should have noticeably stronger Na I absorption. For example, an M5 star at 1--5 Myr would have a Na-8190 index 5--10\\% larger than that observed for an M5 star at 10 Myr. The spectra presented in Figure~\\ref{fig:cha7uscospec} are nearly identical (and have near-identical measured Na-8190 indices of 0.89 and 0.90).  Based on analysis of spectral features, we classified these two stars as being roughly the same age. However, based on the differences in observed luminosity, interpreting the HR diagram literally, one would infer an age spread of $>$10 Myr between the two stars.   \\subsubsection{Comparison of the Observed Age Distribution to a Coeval Population} \\label{cha:7:sec:usco:coeval}  Due to the discrepancy between stellar ages derived from the HR diagram compared to those inferred from spectral surface gravity signatures, we sought to determine the statistical significance of the observed spread in ages for stars on the HR diagram. Effective temperatures and luminosities shown in Figure~\\ref{fig:cha6hrusco} are derived from observed $J$ magnitudes, $J-H$ colors, and spectral types. The age and mass of a star are inferred from effective temperature and luminosity using a set of theoretical isochrones and mass tracks.  Thus, uncertainties in measured photometry or spectral type are propagated into uncertainties in mass and age. Variations in distance and binarity can cause additional uncertainty in luminosity, and, hence, quantities derived from the HR diagram.  We explore first the possibility that the USco population could be coeval, and the apparent age spread in the HR diagram is a result of observational uncertainties combined with association depth and binarity effects. Using a similar method to that used in this work, \\cite{1999AJ....117.2381P} compute an age of $\\sim$5 Myr for the intermediate-mass members of USco. We used Monte Carlo techniques to generate a coeval population of 5 Myr-old stars and brown dwarfs at a distance of 145 pc. The input spectral type distribution was selected to mirror that of our observed distribution of stars, with 1000 stars simulated for every star observed. For each star in the simulated population, $J$- and $H$-band magnitudes were varied by adding random offsets drawn from a Gaussian distribution with a 1-sigma deviation of 0.025 mag, corresponding to the average uncertainty in the 2MASS photometry for observed stars.  To mimic the magnitude-limited data sample, we did not allow any star to be simulated below the photometric survey limits ($J$=16 and $H$=15.5).  Similarly, a random offset was added to the assumed spectral type, selected from a Gaussian distribution with 1-sigma errors of 0.5 spectral subtypes corresponding to the qualitative error of the optical spectral type determinations. Simulated spectral types were rounded to the nearest 0.25 subclasses to reflect the discreet nature of spectral type classification for new members discovered in our work. Using the new spectral type, for each simulated star, we re-derived the expected effective temperature, bolometric correction and intrinsic $J-H$ color using the methods described in \\cite{2004ApJ...610.1045S}.  The maximum distance spread (derived from secular parallax measurements) among members of the association with Hipparcos measurements is 50 pc (\\citealt{2002AJ....124..404P}; \\citealt{1999MNRAS.310..585D}).  In the simulation, we assumed a uniform spatial distribution over a box of this depth centered at 145 pc. A 33\\% binary fraction for stars across all simulated spectral types (M3--M8) was assumed, consistent with observational results of the binary frequency for low mass members of USco (\\citealt{2008ApJ...679..762K}, \\citealt{2005ApJ...633..452K}). All observed low mass binaries in USco have near equal masses ($m_{secondary}/m_{primary} \\gtrsim$0.6; \\citealt{2005ApJ...633..452K}), and thus, the assumption was made that all binaries were composed of two equal mass stars.  This assumption is somewhat liberal in that it will produce the largest possible dispersion in luminosity.   We compare quantitatively results of this simulation to the observed data in Figure~\\ref{fig:cha7ageusco}.  The red shaded histogram shows the age distribution derived from the Monte Carlo simulation, overplotted with a histogram of ages for the data (black hatched histogram). Both histograms have been normalized to unity at the peak for comparison, and all stars with log($T_{eff})<$3.4 (beyond which interpolation of the isochrones becomes unreliable) have been excluded from the data and the model results. The widths of the distributions are remarkably similar, given the simplistic nature of the Monte Carlo simulation.  For the simulated association, we find a mean apparent age of log($age$)=6.51$\\pm$0.40 whereas for the data (not including the 3 M8 stars) we find a mean age of log($age$)=6.53$\\pm$0.47. The data and the model differ significantly at the distribution tails, which is most likely caused either by an incorrect assumption of one of the model parameters, or by an inherent problem with the theoretical isochrones at very low masses \\citep{2004ApJ...604..741H}. To quantitatively assess how well the model reproduces the bulk of the data, excluding the distribution tails, we applied a $\\chi^2$ test of the central peaks of the distributions from log($age$)=6.0 to log($age$)=7.4.  This test yields a $\\sim$53\\% probability that the two distributions could have been drawn from the same population. Thus, given the uncertainties, the data are consistent with most stars in USco forming via a single burst $\\sim$5 Myr ago.  This result is somewhat surprising given the large extent of the association ($\\sim$35 pc across). Assuming a sound speed of $c_S$=0.2 km/s consistent with a T=10 K molecular cloud, and assuming the stars formed close to where they are observed today, the sound crossing time for the parental molecular cloud is $\\sim$85 Myr. Thus, one end of the cloud could not have `communicated' to the other in time to create a simultaneous burst of star formation.  This result does not, however, rule out the possibility that star formation happened simultaneously throughout the extent of the cloud because every part independently reached the threshold for star formation at the same time. Another scenario is that the stars formed much closer together and have since spread to their current positions.  However, this possibility can be ruled out from simple arguments. The velocity dispersion of the massive Hipparcos members is $\\sim$1.3 km/s \\citep{1999MNRAS.310..585D}.  In 5 Myr, the furthest an association member could travel at a speed of 1.3 km/s is $\\sim$6.5 pc, less than half the radius of the current association.  A similar discrepancy between a small observed age spread and the large spatial extent of the association has also been noted in the literature for USco's intermediate and high mass members \\citep{1999AJ....117.2381P}. To explain the disparity, \\cite{1999AJ....117.2381P} (and later \\citealt{2007IAUS..237..270P}) proposed a scenario in which star formation in USco was triggered by an external event in the form of a supernova explosion in the neighboring Upper Centaurus-Lupus association ($\\sim$70 pc away and $\\sim$17 Myr-old). They argue based on the structure and kinematics of large H I loops surrounding Sco Cen, that such an event is evidenced to have occurred $\\sim$12 Myr ago. If true, the explosion would have driven a shock wave that would have reached USco about $\\sim$5 Myr ago, consistent with the inferred age of USco's stellar population. Our results may imply a similar small age spread with an association age $\\sim$5 Myr and large spatial extent (see \\S\\ref{cha:6:sec:discussion:sub:spat}) for USco's lowest mass stars and substellar members. While our results do not prove the \\cite{1999AJ....117.2381P} hypothesis true, they do give support to the hypothesis' plausibility, and extend its validity to even the lowest mass association members.    \\subsubsection{Comparison of the Observed Age Distribution to a Uniform Distribution} \\label{cha:7:sec:usco:distrib}    From the simulation discussd in \\S\\ref{cha:7:sec:usco:coeval}, we have determined the observed age distribution in USco is consistent with all stars forming in a single burst 5 Myr ago.  We explore also the maximum age spread that can be inferred from our data assuming that the star formation rate has been constant in time. We repeated the original Monte Carlo simulation allowing age to vary in addition to spectral type, photometric error, distance and binarity.  We began with a 5 Myr population, and for each star added a random offset in age drawn from a uniform population between $\\pm$0.5 Myr, $\\pm$1 Myr, $\\pm$1.5 Myr, $\\pm$2 Myr, $\\pm$2.5 Myr, $\\pm$3 Myr, $\\pm$3.5 Myr, $\\pm$4 Myr, $\\pm$4.5 Myr, or $\\pm$5 Myr. This age offset had the effect of changing the starting $J$ and $H$ magnitudes. All other parameters were computed as described in \\S\\ref{cha:7:sec:usco:coeval}.  Figure~\\ref{fig:cha7ageuscomod} compares the observed age distribution derived from the HR diagram with the results from the Monte Carlo simulation. As expected, the peak of the simulated distribution decreases and more power is seen in the wings as a larger age spread is injected into the population. We have run a $\\chi^2$ test between central peaks of the data and the simulated model distributions between log($age$)=6 and log($age$)=7.4.  We find that the computed $\\chi^2$ probability remains high, between 50\\% and 60\\%, for simulations with an age spread of $\\leq \\pm$2 Myr, and then falls off rapidly with probabilities of $<$5\\% beyond $\\pm$3 Myr. Thus, we conclude that the observed low mass population of USco formed in $<$6 Myr, and most likely $<$4 Myr; this finding is much less than the $>$10 Myr age spread implied by literal interpretation of the HR diagram.        \\subsection{The Low-Mass IMF} \\label{cha:7:sec:usco:sub:imf}  In this section we use model tracks to derive the first spectroscopic mass function for stars and brown dwarfs in USco less massive than $\\sim$0.1 M$_\\odot$. The study by \\cite{2002AJ....124..404P} derived an IMF for the $stellar$ population in the association by combining the 166 spectroscopically confirmed low mass members ($M$$\\sim$0.8--0.1 M$_\\odot$) discovered via their 2dF survey with the 120 known Hipparcos members (constituting a complete sample of members more massive than $\\sim$1 M$_\\odot$) and 84 X-ray selected members $(M\\sim$0.8--2 M$_\\odot$; \\citealt{1999AJ....117.2381P}). The resultant mass function yielded a 3-segment power law function that was consistent with recent field star and other cluster IMF determinations above $\\sim$2 M$_\\odot$, but with a 2$\\sigma$ excess of stars above numbers seen in the field at lower masses.  The only previous survey that has attempted to derive an IMF for USco's substellar members is the photometric study by \\cite{2007MNRAS.374..372L}.  These authors observed $\\sim$6.5 deg$^2$ of USco at $ZYJHK$-bands as part of the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Cluster Survey.  From these data, the authors used the $J$-band luminosity function to derive a $photometric$ IMF in USco from 0.3--0.007 M$_\\odot$. They derived a mass function that was significantly flatter ($dN/dM \\propto M^{-0.3\\pm0.1}$) than the IMF derived by \\cite{2002AJ....124..404P} for similar mass ranges ($dN/dM \\propto M^{-0.9\\pm0.2}$ for 0.1$\\leq M$/M$_\\odot < $0.6). However, as noted in \\cite{2004ApJ...610.1045S}, photometrically derived IMFs suffer from degeneracies between mass, age, distance, extinction, and photometric excess present in the CMDs from which they are derived, and they cannot directly (only statistically) correct for field star contamination.    As was discussed in \\S\\ref{cha:7:sec:usco}, deriving quantities from spectroscopic data placed onto an HR diagram has its own set of uncertainties.  We have already shown that known uncertainties in observable magnitudes and spectral types can produce an apparent age spread of $>$10 Myr, and therefore must also determine if uncertainties in the data can produce a false spread in the observed mass distribution. To address this issue, we derived a theoretical IMF for a coeval population of stars at 5 Myr. We assumed a spectral type distribution consistent with that for our observed population, and assigned each star the mass that it would have at that spectral type if it were 5 Myr-old. Examination of the difference between the two mass distributions using a KS test reveals that the mass distributions derived from the data and the theoretical 5 Myr-old population are statistically consistent with each other (probability=12\\%). Because the mass tracks are close to vertical for low mass stars (see Figure~\\ref{fig:cha6hrusco}), the dominant source of uncertainty in derived masses will come from uncertainties in temperature, arising from uncertainties in spectral type ($\\pm$0.5 subtypes). Presuming we are equally likely to misclassify a star 0.5 subtypes too early as we are to misclassify it 0.5 substypes too late, the mass distribution will not change significantly. Thus, we conclude that the mass distribution derived from the HR diagram is robust to observable uncertainties. We note, however, that the mass distribution is strongly affected by systematics in the theoretical models, and results will vary dependent upon which set of mass tracks are used.       In Figure~\\ref{fig:cha7uscoimf}, we show the IMF derived using the DM97 models for our spectroscopic survey of low mass stars and brown dwarfs, covering the mass range $M\\sim$0.2--0.02 M$_\\odot$.  Because we have observed spectroscopically only $\\sim$15\\% of the photometric PMS star candidates, before constructing a mass function, it is necessary to determine that the spectroscopic sample presented here is representative of the USco low mass population as a whole.   We use a method similar to that employed by \\cite{2004ApJ...610.1045S} and determine the relative completeness of our sample by computing the number of stars observed spectroscopically compared to the number of photometric candidates present in uniformly-spaced bins perpendicular to the line of selection at the 1\\% data contour (see Figure~\\ref{fig:cha3cmdselect}). As can be seen, the number of observed stars falls off considerably for stars with magnitudes $r\\gtrsim$19.5 which could be observed only under the best seeing conditions. The spectroscopic survey completeness peaks at $\\sim$20\\%. We correct the IMF to uniform completeness at this level by determining the mass distribution of stars within each magnitude bin.  We then add stars to the relevant mass bins based on the fractional completeness of the magnitude bins they came from, relative to the maximum 20\\% completeness level.      In order to extend our sample to higher masses, we combined masses derived for our sample with those for stars from the survey by \\cite{2002AJ....124..404P}. We chose to use the \\cite{2002AJ....124..404P} sample because (1) it constitutes the largest sample of spectroscopically confirmed low mass stellar members, and (2) candidates were selected from an $I,R-I$ CMD in a manner similar to the methods employed to select our sample thus providing a complementary dataset. We specifically chose not to include the higher mass X-ray selected sources \\citep{1999AJ....117.2381P} due to the large difference in selection techniques. To provide the most analogous samples possible, we re-derived masses for the 160 stars in the \\cite{2002AJ....124..404P} survey with M spectral types by first converting published spectral types and 2MASS magnitudes to temperature and luminosity using the techniques described in \\S\\ref{cha:6:sec:discussion:sub:hr}, and then using the DM97 tracks and the same interpolation program used to generate masses for the sample of stars presented here. In combining samples from multiple sources, one must be careful to account for relative completeness. \\cite{2002AJ....124..404P} estimate spectroscopic completeness levels of 87\\% for stars with 0.2 $< M/M_\\odot <$0.8 and 67\\% for stars with mass $<$0.2 M$_\\odot$, whereas, our corrected survey is only 20\\% complete across all sampled mass bins. The number of stars in the \\cite{2002AJ....124..404P} survey were thus multiplied by either (20\\% complete/87\\% complete) for stars with 0.2 $< M/M_\\odot <$0.8, or (20\\% complete/67\\% complete) for stars with $M<$0.2 M$_\\odot$. The two surveys also cover very different areas.  \\cite{2002AJ....124..404P} looked at 9 deg$^2$ whereas the Quest-2 survey looked at $\\sim$150 deg$^2$.  However, as discussed in Paper~I, SCH06, and \\cite{mythesis}, the Quest-2 survey coverage within this area is not complete due to several failed CCDs, gaps between the CCDs, and incomplete $r-i$ color coverage. To account for differences in area, we scaled the IMF derived from the \\cite{2002AJ....124..404P} data to match, on average, the level of the IMF derived from stars in our survey in the -0.6$>$log(M/M$_\\odot$)$>$-1.0 mass bins. Due to the nature of magnitude limited surveys, the lowest mass bin (or bins, depending on bin-width) of any IMF is usually incomplete relative to the rest of the derived distribution; thus, for the IMF in the combined samples, we use the values derived from the Quest-2 data to populate the -1.0$>$log(M/M$_\\odot$)$>$-1.4 bins of the IMF.    Figure~\\ref{fig:cha7uscoimffull} shows IMFs for the low mass stellar and substellar population of the USco association. In total, the combined IMF contains $\\sim$377 stars with masses as high as $\\sim$0.6 M$_\\odot$ and as low as $<$0.02 M$_\\odot$. The derived mass function rises with a slope of $dN/dM \\propto M^{-1.13}$, peaks at $\\sim$0.13 M$_\\odot$, remains high with a secondary peak at M$\\sim$0.05 M$_\\odot$, and then gradually falls off through the substellar regime. Thus, we find the spectroscopic IMF for USco's low mass population to be very different from the photometric IMF derived by \\cite{2007MNRAS.374..372L} which had a flatter slope $dN/dM \\propto M^{-0.3}$ from 0.3$ < M/M_\\odot <$0.01, and peaked at $M\\sim$0.01 M$_\\odot$. The total mass inferred from the spectroscopic IMF over the mass range 0.6--0.02 M$_\\odot$ is $\\sim$48 M$_\\odot$. If we correct this number to account for the fact that the IMF shown here is only 20\\% complete over $\\sim$40\\% of the entire spatial extent of the Hipparcos stars, the total mass inferred for association members less massive than $\\sim$0.6 M$_\\odot$ is $\\sim$600 M$_\\odot$.    \\subsection{Comparisons of the Low Mass IMF between Star-Forming Regions} \\label{cha:7:sec:comp}   Diagnostic studies of stellar populations in different locations and at varying stages of evolution are needed to explore the possibility of a universal mass function for low mass objects. While one might expect that the IMF should vary with star formation environment, we do not yet have enough evidence to determine if such a variation exists. Aside from the work presented here, numerous studies have been carried out to characterize the low mass stellar and substellar mass functions of other young clusters and associations in a variety of environments. Because of the intrinsic faintness of low mass objects, most surveys are photometric. Authors then use a combination of theoretical models and statistical analysis to transform a group's color-magnitude diagram or luminosity function into a photometric IMF which may not accurately represent the underlying population. Spectroscopic surveys provide a more accurate assessment of the stellar membership and thus, of the underlying cluster population.  The substellar populations of several other young star-forming regions have been studied spectroscopically using techniques similar to those presented here, and we discuss results for USco in comparison to three other regions. \\cite{2003ApJ...593.1093L} studied the rich cluster IC 348 in Perseus. There have been several studies of the Taurus region, namely those by \\cite{2000ApJ...544.1044L}, \\cite{2002ApJ...580..317B}, \\cite{2003ApJ...590..348L}, and \\cite{2004ApJ...617.1216L}, which together constitute a complete sample within a selected region. \\cite{2004ApJ...610.1045S} used near-infrared imaging and spectroscopic observations to constrain the low mass IMF in the $\\sim$1 Myr-old Orion Nebula Cluster (ONC). The data for IC 348 \\citep{2003ApJ...593.1093L} and Taurus (\\citealt{2000ApJ...544.1044L}, \\citealt{2002ApJ...580..317B}, \\citealt{2003ApJ...590..348L}, and \\citealt{2004ApJ...617.1216L}) are published as complete across all mass ranges within the areas surveyed. \\cite{2004ApJ...610.1045S} employed similar techniques to those used here to correct the derived IMF to a uniform 40\\% completeness across all magnitude ranges. IC 348, Taurus, and the ONC are all young ($age\\lesssim$3 Myr), and, similar to USco, have not yet had time to lose members through dynamical ejection. Therefore, if the low mass IMF is universal, similar mass distributions should be observed in all four regions.  Our goal in this section is to assess the physical differences between stars formed in different star-forming environments. As described in \\S\\ref{cha:6:sec:discussion:sub:hr}, effective temperatures shown in Figure~\\ref{fig:cha6hrusco} are derived directly from observed spectral types using a theoretical or (in our case) empirical temperature scale. Observational evidence suggests that late-type young stars will have a similar mass across a large range of ages. For example, similar masses of $M\\sim$0.03--0.09 M$_\\odot$ have been measured for M6--M7 binary or triple compenents in both the ONC ($\\sim$1 Myr; \\cite{2006Natur.440..311S}) and AB Dor ($\\sim$75 Myr; \\cite{2007ApJ...665..736C}). Thus, spectral type distributions should provide a reasonable approximation to mass distributions for young low mass stars. We therefore begin this exercise by first comparing spectral type distributions found in different star forming regions, and thereby avoid introducing uncertainties that inherently arise when deriving stellar masses.  Spectral type distributions are shown in Figure~\\ref{fig:cha7imfalldm}. All four distributions are complete in a relative sense across spectral type bins for M-type stars. Distributions for members of the ONC and members of USco discovered with Quest-2 were corrected to a uniform completeness level in a manner similar to that described in \\S\\ref{cha:7:sec:usco:sub:imf}. We have included in the IC 348 and Taurus spectral-type distributions all M-type stars present in the IMF of the respective discovery paper. As can be seen, the distribution for USco is visually dissimilar to the other regions, even considering the small sample sizes and large error bars.  This difference could be due to real variation in the IMF, but could also be influenced by the difference in age of a few million years ($\\sim$5 Myr vs. $\\lesssim$3 Myr) between USco and the other regions. We performed $\\chi^2$tests between all four distributions to determine the probabilities that the distributions could have be drawn from the same population. The tests yielded a very small probability ($\\sim$10$^{-22}$) that the USco and ONC distributions could be drawn from the same population, and marginally-small probabilities at the 5--10\\% levels of USco arising from the same population as either Taurus or IC 348. However, we caution against taking results from the $\\chi^2$ test at face value.  The histograms for USco and the ONC have been scaled to account for incompleteness in magnitude and spatial area during the observations, and correctly and rigorously accounting for these scaling factors in the $\\chi^2$ test is a non-trivial and not straight-forward exercise.   The Taurus and IC348 populations, however, have not been scaled, and have an almost 100\\% chance of being drawn from the same population.  While this result argues for a common spectral-type distribution between these two regions amoung the lowest mass stars and brown dwarfs, substantial variation was found by \\cite{2003ApJ...593.1093L} for hotter, K-type stars.      We now compare directly mass distributions for the four selected star forming regions. As discussed above and in \\S\\ref{cha:6:sec:discussion:sub:hr}, the mass of a star is inferred from placement on an HR diagram using a set of theoretical mass tracks, and the exact method by which the mass is derived can vary significantly from study to study. Thus, direct comparison of the published mass data itself requires great caution given that different mass tracks, temperature scales, and interpolation methods were used for different studies. Quantitative assessment of the differences between theoretical PMS star models is beyond the scope of this work, as is analysis of differences between the numerous temperature scales employed by various authors who study young low mass stars. We have used published spectral types and 2MASS $J,H,K_S$ magnitudes to derive effective temperatures, luminosities, and hence masses in a manner consistent with that used in our work in the ONC and USco. We reiterate that masses for all stars presented in this section were derived in a consistent manner;  however, variations in the theoretical or empirical scales/tracks used could shift interpretation of the data.    IMFs derived from the DM97 tracks are shown in Figures~\\ref{fig:cha7imfalldmmass}. We find that the ONC, IC 348, and Taurus IMFs exhibit a fall-off of stars beyond $\\sim$0.1 M$_\\odot$. The mass distribution for USco, however, does not begin to turnover until $\\sim$0.05 M$_\\odot$. Thus, similar to the spectral type distribution, the IMF derived for USco shows an over-abundance of late-type, low mass stars, and suggests that the difference seen in USco's spectral type distribution did arise from a differnce in it's IMF, rather than from slight differences in age between USco and the other regions. Our work extends the claim by \\cite{2002AJ....124..404P} that USco may contain anomolously high numbers (compared to the field) of low mass stars to lower masses, and is suggestive that USco also contains relatively higher numbers of very low mass stars and brown dwarfs compared to other nearby star-forming regions.  USco differs from IC 348, Taurus and the ONC in that it contains $\\sim$50 very massive, O- and B-type stars, almost three times the number found in the ONC (none are found in either Taurus or IC 348).  Thus, the conclusion that USco may contain relatively more very low mass stars than other young regions may be related to its large number of high mass stars. For example, \\cite{2006ApJ...641..504A} have shown through numerical simulation that disks around lower mass stars are more susceptible to destruction (further out than $\\sim$10 AU) from dynamical interactions with surrounding stars of higher mass. It is possible that in a similar manner, low mass cores moving through a giant molecular cloud are more susceptible to stripping of their accretion envelope in the presence of higher mass cores. Our results suggest that, within a star forming region, either the presence of large numbers of very massive stars, or the environmental conditions that lead to numerous massive star formation, may play a large role in determining the low mass IMF. We note that this conclusion has been drawn from results for only four different environments.  Many more regions must be studied before a definitive explanation of the low mass IMF can be derived."
    },
    "0809/0809.4025_arXiv.txt": {
        "abstract": "We present the results of an extragalactic point source search using the five-year WMAP 41, 61 and 94 GHz (Q-, V- and W-band) temperature maps. This work is an extension of our designing and applying a CMB-free technique to extract point sources in the WMAP maps. Specifically, we have formed an internal linear combination (ILC) map of the three-band maps, with the weights chosen to remove the CMB anisotropy signal as well as to favor the selection of flat-spectrum sources. We have also constructed a filter to recover the true point source flux distribution on the sky. A total of 381 sources are found in our study at the $> 5\\sigma$ level outside the WMAP point source detection mask, among which 89 are ``new'' (i.e., not present in the WMAP catalogs). Source fluxes have been calculated and corrected for the Eddington bias. We have solidly identified 367 ($96.3\\%$) of our sources, the 1$\\sigma$ positional uncertainty of which is 2$^\\prime$. The 14 unidentified sources could be either extended radio structure or obscured by Galactic emission. We have also applied the same detection approach to simulated maps, which yielded 364$\\pm$21 detections on average. The recovered source distribution $N(>S)$ agrees well with the simulation input, which proves the reliability of this method. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) is designed to advance observational cosmology by making precise measurements of the cosmic microwave background (CMB) anisotropy (\\citealt{2003ApJ...583....1B}). The characterizing and removing of foreground emission play a vital role in the interpretation of CMB temperature maps. At small angular scales, extragalactic radio sources, especially those with flat spectra, constitute the most important foreground. Therefore, it is essential to have efficient source detection techniques developed and applied to the WMAP maps. Meanwhile, since WMAP provides the only all-sky millimeter survey, source extraction in its maps also opens a window to study extragalactic radio sources in the millimeter wavelength region.  With every WMAP data release, a catalog is provided of the brightest point sources in the WMAP maps. The detection procedure used by the WMAP science team features a filtering of the integration-time-weighted maps with $b_{\\ell}/(b_{\\ell}^{2}C_{\\ell}^{CMB}+C_{\\ell}^{noise})$ in harmonic space and a search for $> 5\\sigma$ peaks, where $b_{\\ell}$ is the transfer function of the WMAP beam response (\\citealt{2003ApJS..148...39P}, \\citealt{2007ApJS..170..263J}, \\citealt{2009ApJS..180..246H}), $C_{\\ell}^{CMB}$ is the CMB angular power spectrum and $C_{\\ell}^{noise}$ is the noise power. This procedure successfully generated 208, 323 and 390 point sources in the WMAP first year, three-year and five-year maps, respectively (\\citealt{2003ApJS..148...97B}, \\citealt{2007ApJS..170..288H}, \\citealt{2008arXiv0803.0577W}; hereafter, WMAP1, WMAP3 and WMAP5 catalogs). However, because of the limited angular resolution of WMAP, positive CMB excursions can be confused with point sources. In addition, the $C_{\\ell}^{CMB}$ term in the filter counts as a systematic noise that does not integrate down with observing time. This largely limits the increasing of the number of detections with increased sensitivity of the maps. To circumvent the CMB ``noise'', \\citet{2008ApJ...681..747C} introduced a CMB-free technique, which involves forming internal linear combination (ILC) of multi-frequency maps to suppress the CMB signal. The number of sources N found by applying this technique to the WMAP V- and W-band maps alone varied as  $t^{0.72}$ from one year to five years, in comparison to $N \\sim t^{0.39}$ between the WMAP1 and WMAP5 catalogs (\\citealt{2009ApJS..180..283W}).  In this paper, we extend this CMB-free technique and employ it to the WMAP five-year Q-, V- and W-band temperature maps. The main goal is to utilize the high signal-to-noise ratio Q-band data to find more flat-spectrum sources, which are known to dominate the 30 - 100 GHz frequency regime. ",
        "conclusions": "We present catalogs of point sources found in an ILC map constructed from the WMAP five-year Q-, V- and W-band maps. We have recovered 287 WMAP5 sources and five WMAP3 sources in our survey. Bypassing of CMB ``noise'' in the detection process brings us 89 new sources (23.4$\\%$ of the total). We have obtained flux density estimates for our sources in each of the three bands. Most of the new sources are found to have low or even negative flux estimates in at least one band, which is probably due to the coincidence of their locations with negative CMB fluctuations. This will generally lead to missing of these sources in single frequency band searches. We identify our sources by searching in NED for nearby low frequency radio sources and fitting the fluxes of our sources to their SEDs. Of all the sources we detected, we lack of solid identifications for 14. The flux estimates and number count distribution of our identified sources are comparable with those from the WMAP5 catalog, which proves our method to be complimentary to as well as consistent with the WMAP approach. Since this CMB-free technique responds rapidly as observing time increases (\\citealt{2009ApJS..180..283W}), we expect a fast increase of new detections in more years of WMAP maps."
    },
    "0809/0809.4355_arXiv.txt": {
        "abstract": "Hybrid inflation faces two well-known problems: the blue spectrum of the non-supersymmetric version of the model and the fine-tuning of the initial conditions of the fields leading to sufficient inflation to account for the standard cosmological problems. They are investigated by studying the exact two-fields dynamics instead of assuming slow-roll. When the field  values are restricted to be less than the reduced Planck mass, a non-negligible part of the initial condition space (around $15\\%$ depending on potential parameters) leads to successful inflation. Most of it is located outside the usual inflationary valley and organized in continuous patterns instead of being isolated as previously found. Their existence is explained and their properties are studied. This shows that no excessive fine-tuning is required for successful hybrid inflation. Moreover, by extending the initial condition space to planckian-like or super-planckian values, inflation becomes generically sufficiently long and can produce a red-tilted scalar power spectrum due to slow-roll violations. The robustness of these properties  is confirmed by conducting our analysis on three other models of hybrid-type inflation in various framework: ``smooth'' and ``shifted'' inflation in SUSY and SUGRA, and ``radion assisted'' gauge inflation. A high percentage of successful inflation for smooth hybrid inflation (up to $80\\%$) is observed. ",
        "introduction": "For almost a decade, the cosmic microwave background data have supported a cosmological concordance model, in which inflation~\\cite{Starobinsky:1980te,Guth:1980zm,Linde:1981mu,Linde:2005ht,Inflation+25}, an early phase of accelerated expansion, is the favored explanation for the origin of the primordial fluctuations\\footnote{Note however that some alternatives exist such as string gas cosmology~\\cite{Brandenberger:2006vv} or bouncing universe models \\cite{Peter:2006hx,Peter:2008qz}.}. For more than 25 years, many models of inflation have been proposed, from toy models to more realistic models based on various high energy physics frameworks \\cite{Lyth:1998xn,Linde:2005ht,Linde:2005dd,Kallosh:2007ig}. The incoming flow of cosmological data has however started to discriminate among the models. In particular, the last release of the WMAP 5-years data~\\cite{Komatsu:2008hk} favored a red tilted scalar power spectrum.  If some single-field models are still able to reproduce the current data, the presence of multiple scalar fields in all the high energy physics frameworks proposed today (Higgs fields in Grand Unified Theories, super-partners in Supersymmetry, moduli in string theory) makes it hard to imagine that the inflaton field is not coupled to any other scalars. The simplest (and yet motivated) known example of multi-field inflation is the hybrid one. The original hybrid model of inflation, proposed in~\\cite{Linde:1993cn,Copeland:1994vg}, had been introduced as an alternative way to end inflation and could be realized for sub-planckian field values unlike chaotic models. The key idea is to couple the inflaton field to a Higgs-type waterfall field which ends inflation by acquiring a non-vanishing Vacuum Expectation Value (vev). This model could be considered as realistic if employing a Higgs field and an extra singlet of some minimal extension of the Standard Model of particle physics. It also represents a toy model for many multi-field models of inflation in other frameworks. Indeed, hybrid(-type) models of inflation have been embedded in almost all high energy frameworks: in (extended) supersymmetry and supergravity~\\cite{Halyo:1996pp,Binetruy:1996xj,Dvali:1994ms, Kallosh:2003ux}, in Grand Unified Theories~\\cite{Jeannerot:1997is, Jeannerot:2003qv}, or various extra-dimensional theories~\\cite{Dvali:1998pa,Koyama:2003yc,Fukuyama:2008dv, Fairbairn:2003yx}.  When confronting the original hybrid inflation to the CMB data, it is however well known~\\cite{Lyth:1998xn} that the power spectrum tilt is blue which is now disfavored. This is only valid when slow-roll is assumed and when the vacuum energy density dominates the potential since in the other case, the potential becomes equivalent to a chaotic model. In this paper, we illustrate these properties and study the predictions for the spectral index of the model using the exact field dynamics. We find a new way to generate a red tilted spectrum due to non-trivial effects of the violation of slow-roll, give the two possible conditions on the parameters of the potential to generate a red spectrum of perturbations and discuss the field values that these conditions require.\\\\  Several fundamental questions about initial conditions for inflation are still open (see for example~\\cite{Goldwirth:1991rj,Vachaspati:1998dy,Kaloper:2002cs, Linde:1986fd,Tetradis:1997kp,Lazarides:1997vv,Mendes:2000sq}). In this paper, we will not address the important problem of spatial homogeneity of the fields~\\cite{Goldwirth:1991rj} and we will assume that the field values do not enter the self-reproducing inflationary regime~\\cite{Linde:1986fd}. Even when restricting to the classical approximation, the existence of a fine-tuning on the initial values of the fields was found, for hybrid inflation~\\cite{Tetradis:1997kp,Lazarides:1997vv,Mendes:2000sq}. (An opposite conclusion has been obtained~\\cite{Lazarides:1996rk} for the smooth hybrid inflation model. We will comment on this model at section~\\ref{sec:othermodels}). The space of initial conditions is described by regions in the plane $(\\phi_\\ui,\\psi_\\ui)$, where $\\phi_\\ui$ and $\\psi_\\ui$ denote the initial values of the inflaton and the waterfall field respectively. By fine-tuning of the initial conditions, one means that the regions leading to sufficient (60 e-folds or more) inflation have been found to be composed~\\cite{Tetradis:1997kp,Mendes:2000sq} of an extremely thin band around $\\psi_\\ui=0$ and a few apparently random points in the rest of the plane. Uncertainties remain on whether these points are of null measure~\\cite{Mendes:2000sq} or not~\\cite{Tetradis:1997kp}. The thin band is also considered as fine-tuned because $\\psi_\\ui$ has to be so close to 0 that any quantum fluctuations would shift its value outside the successful region~\\cite{Tetradis:1997kp}. This would be an important problem for hybrid-type inflation because it means that these models would not easily be the natural outcome of some pre-inflationary era (see however~\\cite{Calzetta:1992bp}).  Several papers have proposed some solutions to the fine-tuning problems. It has been proposed to replicate many times identically the inflaton sector~\\cite{Mendes:2000sq}, even though no motivations have been proposed for this replication. A similar idea had been employed to construct the N-flation model~\\cite{Dimopoulos:2005ac} but the replication in this context is not more natural~\\cite{Easther:2005zr}. It has also been proposed~\\cite{Mendes:2000sq,Underwood:2008dh} to embed hybrid inflation into a brane description. The induced modifications to the Friedmann-Lema\\^itre equations provide additional friction in the evolution of scalar fields. Thus slow-rolling is favored and more of the initial condition space gives rise to successful inflation. This friction can also be efficiently played by dissipative effects~\\cite{Ramos:2001zw}, when couplings between the inflaton and the waterfall field with a bath of other fields are assumed. Finally, it has been proposed~\\cite{Panagiotakopoulos:1997if} to solve this problem by accepting a short ($N\\sim 10$) phase of hybrid inflation and implementing a second one responsible for the generation of the primordial fluctuations, thus solving the horizon problem.  However, to our knowledge, little has been proposed to explain the properties of the (un-)successful space of initial conditions: discreteness, sub-dominance, size and limits. In this paper, we first show that super-planckian initial conditions always give rise to a sufficiently long phase of hybrid inflation and can produce a red-tilted power spectrum without the need of any fine-tuning. We provide a detailed analysis of the properties of the initial condition space, explain why parts of this space were thought to be discrete, and what are the field trajectories leading to these apparently isolated points. In particular, we show that they can be viewed as the ``anamorphosis'' (that is a deformed image) of the thin successful band. We also give the area of successful initial conditions  in the the plane $(\\phi_\\ui,\\psi_\\ui)$. Even when restricting the fields to sub-planckian values, we find around $15\\%$ of successful initial conditions and we discuss the effect of varying the different parameters of the potential. When going to super-planckian values, we confirm that this ratio tends to $100\\%$.  To prove the robustness of these results, we explore the space of initial conditions for three other hybrid-type models: the supersymmetric and supergravity ``smooth''~\\cite{Lazarides:1995vr,Lazarides:2007fh,Yamaguchi:2004tn} and ``shifted''~\\cite{Jeannerot:2000sv,Jeannerot:2002wt} models, as well as the ``radion assisted'' gauge inflation~\\cite{Fairbairn:2003yx}. The first two models are direct extensions of the F-term hybrid inflation \\cite{Dvali:1994ms} and are motivated by the fact that their inflationary valley is shifted away from $\\psi=0$, so that any harmful topological defects formed during the symmetry breaking induced by $\\psi$ would be diluted away. The last model is based on a hybrid-type potential even though constructed in 5D. Its main motivation resides in the fact, that by construction, the form of the potential is controlled (and thus protected) by gauge symmetries.  Before going any further, let us discuss the physical motivations of enlarging the space of initial conditions to super-planckian values. Even though this possibility has been moderately studied in \\cite{Lazarides:1997vv}, most previous works~\\cite{Tetradis:1997kp,Mendes:2000sq} restricted their analysis to initial values of the fields under the Planck mass. For non-supersymmetric four-dimensional theories, it was proposed in~\\cite{Linde:2005ht} that quantum gravity corrections are controlled as long as the energy density and the effective masses are sub-planckian. In this case, effective field theories can be an appropriate framework to describe fields of planckian-like amplitude. More recently, several models such as natural inflation~\\cite{Freese:1990rb}, or gauge inflation~\\cite{ArkaniHamed:2003wu,ArkaniHamed:2003mz,Kaplan:2003aj} also allow fields to be super-plankian. For example, in gauge inflation, the inflaton field is part of a gauge field and thus the form of the potential is protected by gauge symmetries, leaving non-renormalizable corrections highly constrained.  Let's turn to supersymmetric frameworks. Since uncontrolled non-renormalizable corrections to the superpotential and the K\\\"ahler potential appear, super-planckian fields are inevitably problematic for models constructed in the context of supersymmetry (SUSY) or supergravity (SUGRA). Global SUSY is only valid as long as all fields have an amplitude much smaller than\\footnote{Throughout this paper, we will denote the Planck mass by $\\mpl\\equiv G^{-1/2}\\simeq 1.2 \\times 10^{19}\\; \\mathrm{GeV}$ and the reduced Planck mass by $\\Mpl\\equiv (8\\pi G)^{-1/2}$.} $\\Mpl$. When closer to the reduced Planck mass (but still below), SUGRA corrections are important and supergravity is the correct framework to describe the model. Above $\\Mpl$, the non-renormalizable corrections become dominant: the non-renormalizability of  SUGRA prevents us from using it and a UV-complete theory is then necessary \\cite{Baumann:2009ni}. Finding and describing inflaton fields with Planckian displacements in string theories is, however, possible in certain sectors of string theories, though not always stable and their potential is not easily flat \\cite{Baumann:2009ni,Hsu:2003cy}. Several models of inflation constructed within string theories have been proposed (see for e.g. \\cite{HenryTye:2006uv,Kallosh:2007ig,Silverstein:2008sg,Baumann:2009ni} and refs therein) and some of them have a low-energy effective description that mimics hybrid inflation~\\cite{Dvali:1998pa,Linde:2005dd,Koyama:2003yc}. \\emph{These examples motivated us to assume that the effective inflationary potentials studied in this paper can also originate from frameworks in which it is safe to consider super-planckian fields.} As a conclusion, we have chosen to study them both in the sub-planckian and super-planckian field regimes. However, we will always restrict ourselves to sub-planckian energy densities.  The rest of the paper is organized as follows. In the section~\\ref{sec:originalhybrid}, we extensively study the original hybrid inflation model~\\cite{Linde:1993cn}. In particular, we discuss the validity of the one-field slow-roll approximation and show that violation of slow-roll conditions can strongly modify the dynamics. Using exact numerical methods, applied on the two-field potential, we provide a complete analysis of the space of initial conditions and revisit the above-mentioned fine-tuning problem. In Sec.~\\ref{sec:othermodels}, we test the robustness of our results on the three other models: SUSY/SUGRA smooth hybrid inflation, SUSY/SUGRA shifted hybrid inflation and radion-inflation. Our conclusions are drawn in Sec.~\\ref{sec:conclu} and some open questions are developed. ",
        "conclusions": "\\label{sec:conclu} Hybrid inflation is a class of models of inflation motivated by high energy physics. In these models, the inflaton field is assumed to be coupled to a Higgs-type auxiliary field that ends inflation by instability, when developing a non-vanishing expectation value. Two of its main well-known problems - the blue spectrum of the non-supersymmetric version of the model and the fine-tuning of the initial conditions of the fields - are re-analyzed.  First, we found that the original hybrid model can generate a red spectrum by two means.  As well-known, one way to have inflation takes place in the large field phase is to have the waterfall ending inflation in that phase. This requires a constraint on the critical value of the inflaton triggering the waterfall. We found a new criteria on the mass scale $\\mu$ so that a violation of the slow-roll conditions ensures the \\emph{non-existence of the small field phase of inflation}. In both cases, the spectral index generated is less than unity (see Fig.~\\ref{fig:ns}). However, we show that this requires in both cases a large initial value of the inflaton ($>\\mpl$), and therefore a realization of hybrid inflation in a regime away from the limit $\\phi \\ll \\mu$. This conclusion might reduce the appeal of this model.  When considering the full two-field potential, it was found~\\cite{Tetradis:1997kp,Lazarides:1997vv,Mendes:2000sq} that the original models suffer from a fine-tuning of the initial values of the fields to generate a sufficiently long inflationary phase. The space of successful inflation was thought to be composed of a extremely narrow band along the inflationary valley and some ``isolated scattered points'' which seemed randomly distributed and of null measure~\\cite{Mendes:2000sq}. It was therefore considered that these models suffered from some naturalness problem.  We have numerically integrated the exact equations of motion of both fields and studied in details the space of initial conditions. The study has been conducted for four different models of hybrid-type inflation in various frameworks: the original non-supersymmetric model (section~\\ref{sec:originalhybrid}), its extensions ``smooth'' and ``shifted'' hybrid inflation in global supersymmetry and supergravity and the ``radion assisted'' gauge inflation (section~\\ref{sec:othermodels}). As expected, we found that for sufficiently large initial values of the fields (planckian-like or super-planckian), enough e-folds of inflation are generated (see for e.g. Fig.~\\ref{fig:completegrid-hybrid}). This behavior is similar to the one of chaotic inflation~\\cite{Linde:1983gd}. We also studied the shape of the unsuccessful region and its dependence on the potential parameters: this property holds for any values of the parameters and all the models we considered, except when embedded in supergravity. Consequently, the first way to solve the fine-tuning problem of initial conditions for hybrid models is to formulate them in a framework for which it is safe to consider planckian-like or super-planckian initial values for the fields. This can be safe or problematic, depending on the framework used to build the model as detailed in the introduction.  Even if the considered model is formulated in a framework where the fields cannot safely be super-planckian, the unsuccessful region of initial conditions contains successful sub-regions. They correspond to special trajectories for which the velocity in field space becomes oriented along the inflationary valley after some oscillations at the bottom of the potential. Therefore the system climbs up the valley before slow-rolling back down, generating enough inflation. We find that these points form a complex structure, as represented in Fig.~\\ref{fig:anamorphosis}. They can be seen as the anamorphosis of the standard inflationary valley, and explain most of the successful initial conditions when restricting to sub-planckian fields. The relative area that these points occupy is typically of order of $15\\%$ for the original hybrid model. This value can go up to $25\\%$ for radion inflation and even above $70\\%$ for smooth inflation, even though these results depend on the values of the parameters of the potential (see Tab.~\\ref{tab:anamorph}). Moreover, even when supergravity corrections are included, these trajectories still exist, their proportion stays similar and they represent even more of the successful initial conditions. We would like to note that these percentages allow us to claim that the fine-tuning on hybrid inflation is less severe that found in the past. However, to translate this into an amount of fine-tuning for the model, it is necessary to compute a measure of the probability space. This is left for an extension of this work~\\cite{Clesse:2009rr}.\\\\  Several other questions remain open and are extensions of this work. We plan on investigating more deeply the statistical properties of the anamorphosis regions of the plane of initial conditions as well as the effects of initial velocities on this plane~\\cite{Clesse:2009r}. The study of the supersymmetric versions of hybrid inflation, the F-term~\\cite{Dvali:1994ms} and D-term~\\cite{Halyo:1996pp,Binetruy:1996xj} models are also left for a future study. These models could have a different dynamics from the models studied here since radiative corrections generate potentially important corrections to the tree level potential~\\cite{Clauwens:2007wc}.  Finally, our results illustrated that the successful realizations of hybrid inflation are not necessarily by fast-roll toward the inflationary valley as usually assumed but also radially (type-B trajectories), and in that sense chaotic like. Some aspects of the phenomenology of these purely two-field trajectories (power spectrum, generation of non-gaussianities, stochastic effects, reheating, impact on topological defect formation) are still unknown and should be studied."
    },
    "0809/0809.2090_arXiv.txt": {
        "abstract": "We report the results from our analysis of {\\it Suzaku} XIS (0.5--10~keV) and HXD/PIN (15--40~keV) observations of five well-known local ULIRGs: {\\em IRAS} $F$05189--2524, {\\em IRAS} $F$08572+3915, Mrk~273, PKS~1345+12, and Arp~220.  The XIS observations of $F$05189--2524 and Mrk~273 reveal strong iron lines consistent with Fe~K$\\alpha$ and changes in spectral shapes with respect to previous {\\it Chandra} and {\\it XMM-Newton} observations. Mrk~273 is also detected by the HXD/PIN at $\\sim$1.8-$\\sigma$.  For $F$05189--2524, modeling of the data from the different epochs suggests that the change in spectral shape is likely due to the central source switching off, leaving behind a residual reflection spectrum, or an increase in the absorbing column.  An increase in the covering fraction of the absorber can describe the spectral variations seen in Mrk~273, although a reduction in the intrinsic AGN luminosity cannot be formally ruled out.  The {\\it Suzaku} spectra of Mrk~273 are well fit by a $\\sim$94\\% covering fraction model with a column density of $\\sim$10$^{24}$~cm$^{-2}$. The absorption-corrected log[$L_{\\rm 2-10~keV}$/$L_{\\rm IR}$] ratio is consistent with those found in PG Quasars.  The change in the spectral shape on a time scale of a few years implies that the absorbing matter must be near the AGN ($\\sim$1~pc).  The 0.5--10~keV spectrum of PKS~1345+12 and Arp~220 seem unchanged from previous observations and their hard X-ray emission is not convincingly detected by the HXD/PIN.  The large column density derived from CO observations and the large equivalent width of an ionized Fe line in Arp 220 can be reconciled by an ionized reflection model.  $F$08572+3915 is undetected in both the XIS and HXD/PIN, but the analysis of unpublished {\\em Chandra} data provides a new measurement at low energies. ",
        "introduction": "\\label{sec:intro}  Ultraluminous Infrared Galaxies (ULIRGs) are galaxies with $L_{\\rm IR}$ = $L_{8-1000 \\mu {\\rm m}}$ $\\ge$ 10$^{12} L_{\\odot}$, equivalent to the minimum bolometric luminosity of QSOs.  The discovery of a large population of local ULIRGs is one of the most important legacies of the {\\em IRAS} satellite.  At luminosities above 10$^{12} L_{\\odot}$, the space density of ULIRGs in the local universe is greater than that of optically selected quasars with similar bolometric luminosity by a factor of $\\sim$1.5 (Sanders \\& Mirabel 1996). Thus ULIRGs represent the most common type of galaxy at these high luminosities. Systematic ground-based and space-based optical and near-infrared imaging studies have shown that local ULIRGs are almost always undergoing mergers (e.g., Sanders et al. 1988; Veilleux et al. 2002, 2006).  Sanders et al. (1988) suggested that these objects represent a dust-enshrouded phase that eventually evolves into optically-selected quasars. If this is true, ULIRGs take on a fundamental importance for the origin and evolution of quasars.  The primary energy source --- active galactic nucleus (AGN) versus starburst activity --- of ULIRGs is still under debate.  Optical and infrared observations show that most local ULIRGs are dominated by starbursts, but about 30\\% show evidence of AGNs \\citep{vei99a, vei99b}.  Furthermore, the ``warm'' infrared colors ($f_{25~\\mu {\\rm m}}$/$f_{60~\\mu {\\rm m}}$ $>$ 0.2) and quasar-like spectra of the more luminous objects imply that active nuclei are significant in these objects (e.g., Veilleux et al. 1995, 1999a,b, 2008; Genzel et al. 1998; Tran et al. 2001; Armus et al. 2007).  These results would be definitive if it were not for the possibility of severe obscuration.  Galaxy mergers cause massive in-falls of gaseous material towards the center of ULIRGs (e.g., Rupke, Veilleux, \\& Baker 2008), and column densities $\\ga$ 10$^{24}$ cm$^{-2}$ have been deduced from CO measurements (e.g., Downes \\& Solomon 1998; Evans et al. 2002). Observations at UV, optical, near-infrared, and even mid-infrared wavelengths may therefore not always probe the cores of these objects. High resolution ($\\ll$ 1''), high frequency radio observations can penetrate high columns and are an excellent probe of whether an AGN is present (e.g., Lonsdale, Smith, \\& Lonsdale 1993).  However, the bolometric luminosity in the radio band is insignificant, so radio data cannot test whether accretion onto a supermassive black hole (SMBH) is the dominant energetic process. We are left with hard X-rays.  Recent surveys with {\\it XMM-Newton} and {\\it Chandra} (e.g. Ptak et al. 2003; Franceschini et al. 2003; Teng et al. 2005) have found that $\\sim$40\\% of observed ULIRGs show signatures of AGNs.  The observed 2--10~keV luminosity of the surveyed sample is $\\sim$10$^{40}$--10$^{43}$~ergs~s$^{-1}$, with a majority of the sources having luminosities below 10$^{42}$~ergs~s$^{-1}$.  While the ratio log[$L_{\\rm 2-10~keV}$/$L_{\\rm IR}$] is small in nearby ULIRGs (from $-$4 to $-1$; e.g., Teng et al. 2005), this is not much smaller than that found in radio-quiet QSOs (from $-$3 to $-$1).  Moreover, absorption may be a factor even at these energies.  If the absorbing column exceeds $\\sim$10$^{24}$ cm$^{-2}$, the primary continuum emission is suppressed significantly by absorption and Compton down-scattering.  Thus observations at $\\ga$ 10 keV are best to detect Compton-thick AGNs. The sensitivity of {\\it Suzaku} at these high energies is well suited for the study of highly obscured sources like ULIRGs.  In this paper, we present {\\it Suzaku} XIS (0.5--10~keV) and HXD/PIN (15--40~keV) observations of five well-known local ULIRGs. In \\S 2, we discuss our sample. In \\S 3, we report the observations and describe the methods we used to reduce the data. In \\S 4, the results from our spectral analysis of the {\\em Suzaku} data are discussed. In \\S 5, we combine the {\\em Suzaku} data with earlier published and unpublished {\\em XMM-Newton} and {\\em Chandra} data to fine tune our spectral models.  The results of this study are summarized in \\S 6.  Throughout this paper, we adopt the cosmology of H$_0$=75~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega _{\\rm M}$=0.3, and $\\Omega _\\Lambda$=0.7. ",
        "conclusions": "\\label{sec:summary}  The results of our analysis of {\\it Suzaku} XIS (0.5--10~keV) and HXD/PIN (15--40~keV) observations of five of the brightest and best-known local ULIRGs ($F$05189--2524, $F$08572+3915, Mrk~273, PKS~1345+12, and Arp~220) have been presented and compared with earlier {\\em Chandra} and {\\em XMM}-Newton data. The results can be summarized as follows:  \\begin{enumerate}   \\item The XIS observations of $F$05189--2524 reveal a significant change in the observed 2--10~keV spectrum relative to previous {\\it Chandra} and {\\em XMM-Newton} observations.  The spectral variation in $F$05189--2524 suggests that the central source may have turned off, leaving behind a residual reflection component.  However, an increase in column density cannot be completely ruled out. The absorption-corrected 2--10~keV luminosity to infrared luminosity ratios of $F$05189--2524 during its ``high'' state is consistent with values observed in PG quasars.   \\item The XIS spectrum of Mrk~273 contains a strong Fe~K$\\alpha$ line and shows a change in the observed 2--10~keV spectrum relative to previous {\\it Chandra} and {\\em XMM-Newton} observations.  Mrk~273 is marginally (1.8-$\\sigma$) detected at high energies, with the HXD/PIN spectrum $\\sim$5.4\\% above the background.  A change in the covering fraction of the absorber best explains the spectral variations in Mrk~273, although a drop in the intrinsic AGN luminosity cannot be formally ruled out.  A column density of $\\sim$10$^{24}$~cm$^{-2}$ is derived from the {\\em Suzaku} data. The changes in spectral shape and covering fraction on a time scale of a few years suggest that the absorbing matter is $\\sim$1 pc from the central source.  The Mrk~273 spectrum is best modeled by a $\\sim$94\\% covering fraction model.  The absorption-corrected 2--10~keV luminosity to infrared luminosity ratios of Mrk~273 is consistent with values observed in PG quasars.  \\item $F$08572+3915 is undetected in the {\\it Suzaku} XIS and HXD/PIN observations.  The low X-ray count rate derived from unpublished {\\it Chandra} observations, combined with mid-infrared observations, suggests that this source is highly obscured.  \\item PKS~1345+12 is detected by the XIS.  No apparent flux or spectral variability is detected in its 0.5--10~keV spectrum relative to previous {\\it Chandra} observations.  The net 15--40~keV HXD/PIN spectrum is only 1.1$\\sigma$ above the total background, not strong enough for a meaningful full-band (XIS+HXD/PIN) spectral fitting.  Unlike other ULIRGs, the 0.5--2~keV spectrum of this source does not contain an obvious thermal MEKAL component.  Combining this result and the fact that the absorption-corrected 2--10~keV to infrared luminosity ratio of PKS 1345+12 is in agreement with those of PG quasars, the data suggest that the X-ray luminosity of this object is dominated by an AGN.  \\item Arp~220 is detected by the XIS, but not by the HXD/PIN.  Its 0.5--10~keV spectrum, including the iron complex, appears unchanged since previous {\\it Chandra} observations.  The X-ray continuum emission is in agreement with the possibility that a highly obscured AGN is present.  This interpretation of the data is consistent with previous CO observations.  The measurements of the iron emission line at 6.7~keV can be reconciled by the ionized reflection model.  \\item In all three cases where an optically thin thermal component is contributing to the soft X-ray emission detected by XIS ($F$05189--2524, Mrk~273, Arp~220), the temperature, $kT$ $\\sim$0.3--0.8~keV, is consistent with previous observations of ULIRGs.  \\end{enumerate}"
    },
    "0809/0809.2435_arXiv.txt": {
        "abstract": "We use the Cluster EFW experiment to measure the cross-shock electric field at ten low $\\beta$, quasi-perpendicular supercritical bow shock crossings on March 31, 2001.  The electric field data are Lorentz-tranformed to a Normal Incidence frame (NIF), in which the incoming solar wind velocity is aligned with the shock normal.  In a boundary normal coordinate system, the cross-shock (normal) electric field is integrated to obtain the cross shock potential.  Using this technique, we measure the cross-shock potential at each of the four Cluster satellites and using an electric field profile averaged between the four satellites.  Typical values are in the range 500-2500 volts.  The cross-shock potential measurements are compared with the ion kinetic energy change across the shock.  The cross-shock potential is measured to be from 23 to 236\\% of the ion energy change, with large variations between the four Cluster spacecraft at the same shock.  These results indicate that solar wind flow through the shock is likely to be variable in time and space and resulting structure of the shock is therefore nonstationary. ",
        "introduction": "At collisionless shocks, a large 'cross-shock' electric potential arises to oppose the incoming plasma flow.  This potential, and its corresponding electric field, have the sense to repel incoming ions (which comprise the bulk of the kinetic energy and momentum) and to reflect some ions, providing additional dissipation and also plays a role in the redistribution of the upstream flow and magnetic field to the downstream state; the physics of this energy partitioning is not yet fully understood.  In the Ohm's law sense, the cross-shock electric field arises from a combination of the Hall current, electron pressure gradients, and drag due to the small population of gyrating ions at the shock front.  There are few, reliable published measurements of the cross-shock potential, largely because it requires a DC (double-probe) electric field instrument and these have been deployed primarily on low-apogee magnetospheric missions. Formisano [1982] estimated a voltage drop of 140 and 240V at two shocks.  The double-probe instrument on ISEE was used to measure electric fields of up to 100 mV/m at the shock [{\\em Wygant et al.}, 1987] and a Normal Incidence Frame (NIF) electric potential of 420V [{\\em Scudder et al.}, 1986].  {\\em Eastwood et al.} [2007] used Cluster measurements to calculate a potential drop of 260V, which corresponded to the ion kinetic energy change across the shock.  Very large electric fields (600 mV/m), including parallel electric fields of 100 mV/m have been measured by the Polar spacecraft at a high Mach shock [{\\em Bale and Mozer}, 2007], however the single-spacecraft Polar mission does not allow for good estimates of shock frame transformations.  Furthermore, the bow shock is known to have large amplitude electric field structure from electron inertial scales [{\\em Walker et al.}, 2004] down to Debye scales [{\\em Bale et al.}, 1998]. Several authors have used a Liouville mapping of thermal electrons across the bow shock to estimate the deHoffman-Teller frame shock potential [viz. {\\em Schwartz et al.}, 1988], which can be related back to the NIF potential by a frame transformation.  {\\em Gedalin} [1997] calculated the expected shock potential (assuming pressure balance) and found that it should peak at an Alfven Mach number of around 2; this is where the magnetic compression begins to saturate, while the shock thickness continues to grow [{\\em Bale et al.}, 2003] with Mach numbers, leading to a small potential and more magnetic reflection at high Mach numbers.  The four Cluster spacecraft can be used to calculate a shock reference frame, which is important to transform correctly electric field data.  In this letter, we measure directly the cross-shock electric field and compute the NIF electric potential.  We find that the potential varies significantly between the different spacecraft at the same shock, indicating rapid temporal and/or spatial variations.  This seems to be consistent with expectations of shock 'reformation' [e.g. {\\em Krasnoselskikh et al.}, 2002; {\\em Hellinger et al.}, 2002 {\\em Matsukiyo and Scholer}, 2006] and indicates that plasma flow through the shock and particle heating and energization are likely to be bursty. ",
        "conclusions": "The last columns of Table 1 show $\\phi_A/E$ and $\\phi_A/\\Delta E$, the average cross-shock potential normalized to the upstream ion kinetic energy $E=1/2 m v^2_u$ and its change across the shock $\\Delta E = 1/2 m (v_u^2-v_d^2)$.  This is a measure of the ability of the shock to oppose the directed plasma flow; i.e. when the electric potential  $\\phi_A$ approaches the upstream ion energy $E$ the shock should turn back the {\\em entire} solar wind thermal ion population.  Figure 3 shows $\\phi/E$ with the black dots as $\\phi_A$ and the error bars representing the maximum and minimum $\\phi_i$ at each shock and plotted against Alfven Mach number; red dots show the values of  $\\phi_A/\\Delta E$ The electric potential can vary by nearly $100\\%$ in some cases and in three of the ten cases, $\\phi_A/E$ is greater than 1.   \\begin{figure}[h] \\centering \\includegraphics[width=80mm, height=75mm, scale=1,clip=true, draft=false]{figure3.pdf} \\caption{The cross-shock potential normalized to the upstream NIF ion kinetic energy $E = \\phi/(1/2 m v_u^2)$ plotted against Alfv\\'en Mach number.  The black dots are computed from the 'average' shock potentials $\\phi_A$ and the error bars show the range of potential variation between the 4 Cluster satellites.  Red dots show the ratio of $\\phi_A$ to the NIF ion kinetic energy change $\\Delta E = 1/2 m (v_u^2-v_d^2)$.  The dotted line is the analytical relationship from {\\em Gedalin} (1997).} \\end{figure}    {\\em Gedalin} [1997] estimated the cross-shock potential in the NIF analytically (assuming a monotonic shock profile and pressure balance) and found that the cross-shock potential peaks at small Alfven Mach numbers ($M_A \\approx$ 2), qualitatively in agreement with dHT potentials inferred from Liouville mapping of thermal electrons [{\\em Schwartz et al.,}, 1988]. Our NIF potentials show a similar trend (Figure 3), albeit with poor statistics; the dHT and NIF cross-shock ($\\hat{x}$) electric fields are related by a frame transformation $E_{dHT} = E_{NIF} + (E_y/B_x) B_y$.   Theoretical studies of shock reformation or nonstationarity predict that quasiperpendicular shock fronts should be unstable in certain regimes of Mach number and plasma $\\beta$ [e.g. {\\em Krasnoselskikh et al}, 2002; {\\em Hellinger et al.}, 2002; {\\em Matsukiyo and Scholer}, 2006]. In particular, the constraint that the incoming solar wind speed be larger than the whistler phase speed and/or the ion thermal speeds allows the development of a shock front instability that results in 'overturning' of the front with a characteristic timescale of the ion gyroperiod and a spatial scale of the ion gyroradius.   Recently {\\em Lobzin et al.} [2007], using Cluster observations, provided convincing evidence that high Mach number quasiperpendicular shocks are nonstationary, moreover, a quasi-periodic shock front reformation takes place which modulates the reflected ion population.     In the reformation scenario, the shock electric potential will oscillate to large values on timescales of the ion gyroperiod $\\tau_{ci}$. The time between shock crossings here (from spacecraft to spacecraft) ranges from 7 to 30 seconds, while the ion cyclotron period is $\\tau_{ci} \\approx$ 2 s; so it is plausible that these shocks are reforming rapidly and the four Cluster spacecraft each encounter the same shock at a different phase of the reformation cycle resulting in large variations in the measured potential.  This strongly varying electric potential should produce a modulated reflected ion flux, as observed by {\\em Lobzin et al.} [2007].  During this interval, the Cluster spacecraft were separated by 400-900 km, while $\\rho_i \\approx$ 200 km, so that the spacecraft-to-spacecraft variations may be spatial, rather than temporal.     This is the first multi-spacecraft study of the cross-shock electric potential (to our knowledge) and it shows that shock electric potentials vary largely on the timescale of the ion gyroperiod and/or spatial scales of ion gyroradii and that the Normal Incidence frame potential is often observed to be greater than the ion kinetic energy change across the shock.  These observations are consistent with the quasiperpendicular shock reformation scenario and suggest that the transmission and reflection at the shock front are a bursty phenomena.  This shock reformation may be due to an inherent instability (as described in references above) or due to solar wind driving, although there is no one-to-one signature of this in the upstream data [{\\em Maksimovic et al.}, 2003]."
    },
    "0809/0809.5072.txt": {
        "abstract": "We present a regularized maximum likelihood weak lensing reconstruction of the Deep Lens Survey F2 field (4 $\\rm{deg^{2}}$).  High signal-to-noise ratio peaks in our lensing significance map appear to be associated with possible projected filamentary structures.  The largest apparent structure extends for over a degree in the field and has contributions from known optical clusters at three redshifts $(z\\sim0.3,0.43,0.5)$.  Noise in weak lensing reconstructions is known to potentially cause ``false positives''; we use Monte Carlo techniques to estimate the contamination in our sample, and find that 10-25\\% of the peaks are expected to be false detections.  For significant lensing peaks we estimate the total signal-to-noise ratio of detection using a method that accounts for pixel-to-pixel correlations in our reconstruction.  We also report the detection of a candidate relative underdensity in the F2 field with a total signal-to-noise ratio of $\\sim5.5$. ",
        "introduction": "The mass function of galaxy clusters is a well known probe of cosmological parameters \\citep{haiman01}.  Since cluster mass is not a direct observable other proxies such as optical galaxy richness \\citep{koester07}, X-ray temperature \\citep{rosati02}, the Sunyaev-Zel'dovich decrement \\citep{carlstrom02}, or weak gravitational lensing \\citep{schneider96} are used to trace the mass distribution. To produce accurate constraints on cosmological parameters the mass-observable conversion must be properly calibrated in order to avoid potential biases.  With weak lensing, mass overdensities are selected from maps of the projected mass density and peaks of mass are identified as a function of their signal-to-noise ratio.  The observable in this case is the weak lensing `shear' which is independent of baryonic physics, and theoretically well understood.  This is a major strength that weak lensing has over other methods.  Known limitations of this method are projection effects due to large scale structure \\citep{hennawi05,white02} which can produce spurious detections, a low peak detection significance, complicated point-spread function (PSF) corrections, and the observational expense (since it requires deep imaging).  In spite of these limitations, forecasts for future imaging surveys predict that this method will still provide large samples of shear-selected galaxy clusters which will allow for constraints on cosmological parameters, including dark energy \\citep{marian06}. In addition to individual clusters, studies have also pointed out that the global statistics of dark matter peaks in weak lensing maps can be used to put constraints on cosmology \\citep{jain00}.  Within a given dataset cosmological information from peak statistics could also play a complementary role in breaking degeneracies in cosmological parameters from studies of cosmic shear \\citep{gavazzi07}.  The current generation of deep optical imaging surveys are creating the first ever maps of the dark matter distribution with weak lensing.  In \\citet{wittman06} we presented preliminary maps from the Deep Lens Survey (DLS), covering a total of area of $8$ $\\rm{deg}^{2}$ spread over five fields. Initial maps from the Canada-France-Hawaii Telescope Legacy survey covering four separate $1$ $\\rm{deg}^2$ fields were presented in \\citet{gavazzi07}. Recently, results from the Hubble Space Telescope COSMOS 2 $\\rm{deg}^2$ field were reported in \\citet{massey07}.  To reconstruct the projected mass density ($\\kappa$) in these surveys from weak lensing, all of these studies have implemented a `direct' reconstruction method based on or related to the technique of \\citet{kaiser93}.  This reconstruction method is computationally efficient but has a number of known limitations; for instance, it requires setting an arbitrary smoothing scale, and it is difficult to incorporate other lensing effects such as magnification \\citep{seitz98}.  Several different reconstruction techniques have been developed to overcome these issues, the most promising of which advocate reconstructing the deflection potential ($\\psi$) on a grid rather than directly determining $\\kappa$.  These techniques have primarliy been applied to pointed observations of individual clusters, for instance in \\citet{jee07} or \\citet{bradac06}. In this study we use a regularized maximum likelihood technique to reconstruct the projected mass distribution over a wide area, four square degree field from the DLS.  Our paper is organized as follows: In $\\S \\ref{sec:data}$ details of the imaging data, PSF correction, and the source galaxies used in our analysis are presented. In $\\S \\ref{sec:analysis}$ the maximum likelihood algorithm used in our reconstruction and its application to the DLS are discussed. In $\\S \\ref{sec:masspeaks}$ the structures observed in the lensing reconstruction are presented along with their likely optical counterparts.  In $\\S \\ref{sec:discussion}$ the expected false peak rate due to lensing shape noise is discussed as well as our method of determining the total signal-to-noise ratio of lensing peaks.  In $\\S \\ref{sec:summary}$ our results are summarized and we discuss future directions of our work. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{False Peak Rate} The noise in weak gravitational lensing reconstructions is influenced by several sources: the distribution in background galaxy positions, ellipticities, and redshifts, spatial variations in the sky background and PSF, and pixel-to-pixel correlations due to the reconstruction methods.  As a result, the noise distribution in the lensing maps is strictly non-Gaussian, and there is the potential for ``false positives''---peaks in the lensing map that do not correspond to any real objects.  Several previous studies have addressed the issue of the false peak rate in weak lensing mass reconstructions.  \\citet{hamana04} used N-body simulations and ray-tracing to construct catalogs of expected clusters to compare with them.  They reported that the completeness and contamination of the catalogs depended on the S/N threshold, and that even at relatively high thresholds ($\\sim 5\\sigma$), a 20-30\\% incompleteness and a similar contamination rate would be expected. \\citet{schirmer07} used a variation of the aperture mass statistic \\citep{schneider96} and estimates of the projected galaxy overdensity to estimate the noise contamination rate in a sample of ESO-WFI fields.  When they varied the S/N threshold, they found contamination rates of 30-80\\%, depending on the depth and seeing of the fields.  \\citet{hetterscheidt05} applied two variations of the aperture mass statistic to 50 deep Very Large Telescope fields; they set their thresholds so as to approximate a 20\\% contamination rate from noise peaks in their sample.  Even though these studies used very different reconstruction methods, data sets and selection criteria, they all concur in finding 20-40\\% contamination in the rough S/N range that we consider.  At lower S/N, the very large number of ``detections\" in both $\\kappa$ and randomized maps indicate a much higher contamination rate.  We also mention that \\citet{vanwaerbeke00} studied the noise in weak lensing reconstructions, however their analysis is not directly applicable to us since it does not apply to maximum likelihood reconstructions which use a regularization term.   To estimate the potential false peak rate in our high S/N sample due to lensing shape noise \\citep{kubo07} we performed two tests.  In our first test we created a mean map of the 100 Monte Carlo reconstructions described in $\\S \\ref{sec:peaks}$.  Each pixel in this map is the average value of this pixel over the individual 100 Monte Carlo map realizations.  An average signal-to-noise ratio map is created by dividing this map by the mean of the $\\kappa_{\\rm{rms}}$ map ($\\kappa_{\\rm{rms}}/10$) from $\\S \\ref{sec:peaks}$. Peaks are detected in this average signal-to-noise ratio map using exactly the same procedure used to analyze the original $\\kappa-\\rm{S/N}$ map.  The resulting peak histogram is shown in Figure \\ref{fig:falsepeaks} (Left).   We find that above our high signal-to-noise ratio cut ($\\nu>3.5$) three peaks are recovered.  In a second test we created individual signal-to-noise ratio maps from each of the 100 Monte Carlo realizations.  Here each individual Monte Carlo map is divided by the $\\kappa_{\\rm{rms}}$ map from $\\S \\ref{sec:peaks}$.  Peaks are detected in each of these individual signal-to-noise ratio maps using the same detection procedure described in $\\S \\ref{sec:peaks}$.  The resulting average histogram is shown in Figure \\ref{fig:falsepeaks} (Right).  Here we detect $\\sim 1$ false peaks above our high S/N cut. Both of the above tests indicate a potential false peak rate due to shape noise of 1-3 (or $\\sim 10-25\\%$) above our high signal-to-noise ratio cut $(\\nu>3.5)$.  This percentage compares reasonably well with an estimate based on the proximity of our peaks to overdensities of galaxies  in physical and redshift space.  Based on the analysis of the environment near the peaks, we see 2-3 (16-25\\%) of the structures appear not associated with or at $\\gtrsim 3\\arcmin$ from identified galaxy concentrations.  In the above tests we do not simulate the signal due to uncorrelated large scale structure along the line of sight which will increase the number of detected objects that are falsely interpreted as cluster candidate. The result of these tests therefore are only an estimate to the lower bound of the false peak rate.  We emphasize that although false peaks are expected above our high signal-to-noise ratio cut, the majority of our peaks appear to be associated with real clusters which also appear to trace the large-scale filamentary structure.  We note here that residuals in the PSF correction could also potentially cause false peaks.  In \\citet{kaiser93} type convergence maps this can be probed at some level by creating a B-mode map where source galaxies are rotated by 45 degrees relative to each grid point (e.g. \\citet{massey07}).  One limitation of our maximum likelihood algorithm however is that in its current implementation it cannot create a B-mode map.  To probe PSF systematics we have alternatively measured radial shear B-modes for our high significance peaks (Figure \\ref{fig:bmode}).  In the figure we show profiles for four of these peaks spread over the structures in the F2 field.  In each panel the profile is center on the peak coordinates given in Table \\ref{tab:clusters}.  The measured B-mode for each peak is consistent with zero, indicating that there is no residual contamination coming our PSF correction.  Profiles for the other high significance peaks in the field are similarly consistent with a zero B-mode signal.  \\subsection{Total S/N} \\label{sec:totsn} To be consistent with previous studies we have selected peaks from the $\\kappa-\\rm{S/N}$ map as described in $\\S \\ref{sec:peaks}$.  This provides a good method to globally select peaks from the map, however because our reconstruction technique does not track pixel-to-pixel correlations, neither the peak nor the integrated signal-to-noise ratio are good estimates of total signal-to-noise ratio of detection.  To calculate the total signal-to-noise ratio for peaks we use a different approach where the original $\\kappa$ map and the individual Monte Carlo maps are block averaged at different levels of block size. $\\kappa-\\rm{S/N}$ maps at each block size are then created in the same manner used to create the original $\\kappa-\\rm{S/N}$ map $(\\S \\ref{sec:peaks})$.  For small values of block size, due to the pixel-to-pixel correlations in the $\\kappa$ and $\\kappa_{\\rm{rms}}$ maps, both the signal and noise increase linearly, hence the peak signal-to-noise ratio ($\\nu$) will stay relatively constant.  When the block size becomes larger than the noise correlation length, noise will increase proportional to square root of the block size; therefore, the peak signal-to-noise ratio will grow as more $\\kappa$ pixels are used in the average.  At some level of block size the peak signal-to-noise ratio will be maximized when the pixel size matches the characteristic size of the cluster.  The peak signal-to-noise ratio at this block size can then be used to estimate the total signal-to-noise ratio of detection.  Past the maximum as the block size continues to increase we expect that the signal-to-noise ratio will decrease as more noise pixels are added.  An example of this technique is shown in Figure \\ref{fig:block}; here we show peak S/N vs. block size curves for three different peaks in our map.  The top curve is for the peak grouping 1 \\& 2, the middle curve is for peak 5, and the bottom curve is for the peak grouping of 8, 10, \\& 11.  In each of the curves the peak S/N reaches a maximum at certain level of block size, and the value of the peak S/N at that block size is used to approximate the total signal-to-noise ratio of detection.  This provides a near-optimal measure of the integrated signal-to-noise ratio for isolated clusters, subject to the constraint of finite pixel sizes.  One limitation of our total signal-to-noise ratio estimation is that for large block sizes some peaks can become blended with other peaks.  This behavior is exhibited in the top curve (Peaks 1 \\& 2) and the bottom curve (Peaks 8, 10, \\& 11).  In such cases we can only estimate the total signal-to-noise ratio for the combined system.  We list total S/N values using this method for the remaining peaks in Table \\ref{tab:clusters}.  Our method of estimating the total signal-to-noise ratio for lensing peaks is qualitatively similar to previously-used techniques  in this field. For example, the P-statistics introduced by Schirmer et al. (2007) use variations in the scale of the kernel to effectively identify peaks based on their signal at a smoothing scale matching the peak size, even though they do not take the additional step of identifying the S/N at this kernel with the total S/N of the peak.  Our method is also similar to the method used in \\citet{kaiser95}; however in that paper, the method was used to estimate the S/N ratio of faint galaxies rather than in estimating the significance of the lensing signal itself.   \\subsection{The Abell 781 Complex} \\label{sec:a781} As mentioned in $\\S \\ref{sec:clusters}$ there is a known X-ray cluster associated with the western portion of Abell 781 complex that does not appear in our reconstruction.  Our lensing map does hint at a peak at this position but this peak appears to occupy only 1-2 pixels in our map.  We recently reported a $1\\sigma-2\\sigma$ lensing detection of this cluster from an NFW fit in \\citet{sehgal08}.  The lensing significance of this cluster is low compared to what is expected from the X-ray mass of this cluster \\citep{sehgal08}.  The low lensing significance could possibly be due to a remaining systematic, or could potentially point toward an interesting physical effect which causes X-ray and lensing masses to be discrepant.  Observations of this region with a different telescope/imager are planned, and a detailed analysis will be presented in a future paper.   \\subsection{A Candidate Void?} \\label{sec:void} In $\\S \\ref{sec:peaks}$ we mentioned the detection of a candidate underdensity in the $\\kappa-{\\rm{S/N}}$ map with a significant negative signal-to-noise ratio.  Strictly speaking, this region is a relative underdensity since it is underdense relative to other structure in the map.  Our ability to detect a relative undersity is unaffected by the mass sheet degeneracy transformation \\citep{schneider95}.  The undersity is located at $\\rm{R.A.}=09^{\\rm{h}}20^{\\rm{m}}22.7^{\\rm{s}}$, $\\rm{decl.}=+29^{\\circ}54\\arcmin03.5\\arcsec$ and is detected with a peak significance $\\nu\\sim5$ and a total signal-to-noise ratio of $\\sim5.5$.  From our lensing reconstruction alone we cannot directly confirm that this underdensity is in fact a void.  A galaxy redshift survey, an alternative mapping of structure, is a well known method of directly tracing underdense regions or voids \\citep{geller89}.  We are conducting such a survey in the F2 field, the SHELS redshift survey \\citep{geller05}, which should provide additional insights into this underdense region.  Confirmation of an underdensity detected in a lensing reconstruction would demonstrate that lensing maps can trace mass underdensities, in addition to overdensities.  We note that an excess of negative peaks was reported in \\citet{miyazaki02}, however this excess has not been confirmed with a redshift survey.      In this work we have presented the application of a regularized maximum likelihood weak lensing reconstruction to the Deep Lens Survey F2 field.  Our lensing reconstruction reveals several projected structures in the F2 field which appear to lie along a network of potential filaments.  In the southwest section of our map we detect a structure which appears to lie at a redshift $z\\sim0.5$. Another structure, the Eastern-most structure in our reconstruction, contains the confirmed cluster Abell 781 and appears to extend southward for over a degree.  Known optical clusters at three redshifts $z\\sim0.3, 0.43, 0.5$ are potentially associated with significant lensing peaks along this structure. Near the center of our map we detect another significant peak at a redshift $z\\sim0.32$.  The cluster associated with this lensing peak lies at a roughly the same redshift as many of the peaks in the Eastern structure, however it is unclear if this structure is physically associated with any part of the Eastern structure.  These two structures also appear to surround a relative underdense region with $\\nu\\sim5$, and total signal-to-noise ratio $\\sim5.5$.  The noise in our reconstruction is estimated from 100 individual Monte Carlo realizations of the source galaxy catalog.  We find that generating Monte Carlo $\\kappa$ maps with this reconstruction technique is computationally expensive, but does allow us to create a statistically correct $\\kappa-\\rm{S/N}$ map.  Monte Carlo realizations also aid in estimating the number of expected false peaks due to lensing shape noise.  We performed two tests here using the Monte Carlo $\\kappa$ maps; in one test we created an average Monte Carlo significance map, and in another test we created individual Monte Carlo significance maps and averaged the results.  Both of these tests estimate $10-25\\%$ false positive detections due to shape noise in our high peak signal-to-noise ratio sample.  This is comparable to the number of peaks in our dataset which lack close counterparts in galaxy overdensities.   We have also presented a new method to estimate the total signal-to-noise ratio of detection for lensing peaks.  Because pixel-to-pixel correlations are not traced in our lensing reconstruction (or in any traditional direct reconstruction) the peak signal-to-noise ratio underestimates the total lensing significance.  Our method of estimating the total lensing significance uses block averaging of the $\\kappa$ map and the individual Monte Carlo maps to create $\\kappa-\\rm{S/N}$ maps with different levels of block size.  The peak signal-to-noise ratio will become maximized here when the peak matches the characteristic size of the cluster, and this peak signal-to-noise ratio can be used to estimate the total lensing significance.  We find that this method works quite well for isolated peaks, but for a grouping of peaks only the significance of the combined group can be measured.  In future work, photometric redshifts of source galaxies in the F2 field could allow us to make tomographic maps using this method and would also allow for the determination of the mass of significant lensing peaks.  Magnification information in the F2 field could also be incorporated into our reconstruction, potentially breaking the mass sheet degeneracy \\citep{broadhurst95}.  We are also pursing optical cluster finding in the DLS, however this is beyond the scope of our current study.   %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line.    %% Putting eqnarrays or equations inside the mathletters environment groups %% the enclosed equations by letter. For instance, the eqnarray below, instead %% of being numbered, say, (4) and (5), would be numbered (4a) and (4b). %% LaTeX the paper and look at the output to see the results.   %% This section contains more display math examples, including unnumbered %% equations (displaymath environment). The last paragraph includes some %% examples of in-line math featuring a couple of the AASTeX symbol macros.   %% The displaymath environment will produce the same sort of equation as %% the equation environment, except that the equation will not be numbered %% by LaTeX.    %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0809/0809.2603_arXiv.txt": {
        "abstract": "The connection between mass loss on the red giant branch (RGB) and horizontal branch (HB) morphology in globular clusters (GCs) has long been acknowledged but the mechanisms governing mass loss remains poorly understood from a theoretical perspective.  The present study uses synthetic HB models to demonstrate for the first time that $\\alpha$-enhancement and a simple relation between mass loss and metallicity can explain the entire range of HB morphology (characterized by the HB type index) observed in old, coeval GCs.  The mass loss-metallicity relation accounts naturally for the fact that the most metal poor GCs ([Fe/H] $< -2$) have redder HBs than is typical of GCs with $-2 <$ [Fe/H] $< -1.5$ without invoking younger ages. These results may prove useful in studying the contribution of HB stars to integrated light via stellar population synthesis. ",
        "introduction": "\\citet{z93} recommended that Galactic halo GCs be divided into two groups: one for old, roughly coeval GCs that follow the mean trend in [Fe/H] vs. HB type in the inner halo (the Old Halo; OH) and one for GCs that deviate from this trend by a significant amount (the Younger Halo; YH).  Due to the lack of very metal poor GCs in the inner halo, Zinn used more distant GCs to extend the line below [Fe/H] $< -2$, see his Figure 1.  Zinn's relation was empirically motivated and differs from later studies---such as those by \\citet[hereafter LDZ]{ldz2} and \\citet{re01} that are based on synthetic HB isochrones---in that it turns back to the red at the lowest metallicities ([Fe/H] $<$ -2). Trends derived from synthetic HB calculations with constant mass loss (LDZ) or mass loss that increases with [Fe/H] and age \\citep{re01} indicate that HB type (as defined by LDZ 1994) rises monotonically to +1 as [Fe/H] decreases and remains there for all lower values of [Fe/H]. \\citet{l91} showed that allowing mass loss to vary with [Fe/H], in a manner consistent with the observed trend in GCs, resulted in better agreement between his models and the pulsational properties of RR Lyrae stars in $\\omega$ Cen than was achieved by assuming constant mass loss. Aside from Lee's application to $\\omega$ Cen, the concept of searching for a mass loss relation that matches the observed trend in old, coeval GCs has not been explored.  The purpose of this Letter is to argue that it is possible to account for the fact that the most metal poor GCs have HB types less than +1 without requiring them to be significantly younger than more metal rich GCs if mass loss on the RGB varies as a function of [Fe/H] in a manner qualitatively similar to that proposed by \\citet{l91}. ",
        "conclusions": "\\citet{ca05} provides a thorough discussion of mass loss and HB morphology in GCs. Catelan considers a handful of empirically motivated formulae that describe the mass loss rate as a function of stellar parameters such as luminosity, radius, and surface gravity. Each predicts that mass loss increases smoothly with [Fe/H] \\citep[Figure 4]{ca05}. Some of the formulae are consistent with the observed difference in HB morphology between second parameter GC pairs, but not all, and no single formula is successful for each of the GC pairs considered \\citep[Figures 10-13]{ca05}. There is also a discrepancy between the formulae and observational efforts to measure mass lost by red giants in GCs by \\citet{or02} who found that mass loss is not strongly correlated with metallicity and only occurs near the tip of the RGB.  On the other hand, \\citet{or07} found that mass loss is evident at the level of HB and increases with luminosity up to the RGB tip in 47\\,Tuc.  The proposed mass loss-[Fe/H] relation should be straightforward to test observationally with mass loss estimates for a large enough GC sample.  These examples illustrate that mass loss remains a complex and poorly understood process.  It is by no means guaranteed that the information provided by standard stellar evolution models is sufficient to calculate mass loss as a function of age and metallicity.  The observed range of HB morphologies in OH GCs (which are assumed to be roughly coeval) of similar metallicity suggests that more than one factor must be involved.  It seems quite plausible that mass loss may be influenced at the stellar level (by, metallicity and rotation, for example) and at the cluster level (by cluster mass or concentration, for example).  Uncertainties in composition, coming from the lack of accurate abundance determinations in some lesser known GCs, or the possible presence of multiple populations with distinct compositions in a single GC, complicate the analysis and interpretation.  The situation is further complicated by the discovery of bi- and multi-modal HBs in some GCs, see \\citet{fe98} for examples and a review of the subject.  It has been suggested by, e.g. \\citet{da08} and references therein, that these (and perhaps all) GCs are comprised of chemically distinct populations that create the complex distribution of HB stars.  Complications arise even in GCs with uni-modal HBs. The analysis of M\\,3 by \\citet{va08} indicates that while the synthetic HB method developed by \\citet{ro73}, particularly the assumed Gaussian mass distribution, is `a reasonable first approximation, [it] fails to account for the detailed shape of M\\,3's HB mass distribution.'  \\citet{va08} also warn that there is a mismatch between observed and theoretical lifetime-luminosity relations on the HB.  In the absence of a firm theoretical foundation for mass loss in stars, it is still possible to further our understanding of this subject and to apply that knowledge.  One application of the results presented here is to the modeling of integrated light via stellar population synthesis.  Investigators such as \\citet{le00} and \\citet{le02} have explored the dependence of integrated light properties on synthetic HB morphology and compared them with observations of extragalactic GCs.  The results presented here may provide useful input for future studies along these lines.  To summarize, synthetic HB models that include the effects of $\\alpha$-enhancement and metallicity-dependent mass loss (via equations \\ref{dM} and \\ref{dMR}) can account for the entire distribution of HB morphologies observed in old, coeval GCs in the Galactic halo. Essentially the same result is obtained from two distinct sets of He-burning tracks, suggesting this is not a spurious or model-dependent result. As a corollary, it is recommended that NGC\\,5053, 5466, 6426, and 7078, recently associated with the YH, be considered members of the OH since their ages are consistent with other OH GCs and their HB morphologies can be explained without invoking substantially younger ages."
    },
    "0809/0809.4605_arXiv.txt": {
        "abstract": "{} {The Antares nebula is a peculiar emission nebula seen in numerous [\\ion{Fe}{ii}] lines and in radio free-free emission, probably associated with the \\ion{H}{ii} region caused by $\\alpha$ Sco B in the wind of $\\alpha$ Sco A. High-resolution spectra with spatial resolution were used to study the emission line spectrum, the physical nature of the nebula and to determine the mass-loss rate of the M supergiant $\\alpha$ Sco A.} {The Antares nebula was mapped with long-slit (10$\\arcsec$) and high-resolution ($R=80\\,000$) spectra using UVES at the VLT. The resulting 2-D images were used  to reconstruct a 3-D picture of the \\ion{H}{ii} region and its absolute location in space relative to $\\alpha$ Sco A.} {We found that the Antares nebula shows, in addition to numerous [\\ion{Fe}{ii}] lines, the Balmer line recombination spectrum H$_{\\alpha}$, H$_{\\beta}$ up to H$_{10}$, and [\\ion{N}{ii}] 6583/6548~{\\AA}, H$_{\\alpha}$ and [\\ion{N}{ii}]  with the same extent as seen in cm radio free-free emission. Combining velocity information from optical and GHRS/HST spectra with H$_{\\alpha}$ velocities, the \\ion{H}{ii} region is found to be located $\\sim 215$~AU behind the plane of the sky of  $\\alpha$ Sco A. From the H$_{\\alpha}$/[\\ion{N}{ii}] intensity ratio and the non-visibility of the [\\ion{O}{ii}] 3726/3729~{\\AA} lines we estimate a low mean electron temperature of $\\overline{T}_{\\rm e} = 4900$~K and an N abundance enhanced by a factor of $\\sim$ 3 due to the CNO cycle in $\\alpha$ Sco A. The shape and size of the \\ion{H}{ii} region yield a mean mass-loss rate of (1.05 $\\pm$ 0.3) $\\times$ 10$^{-6}$ M$_{\\odot}$\\,yr$^{-1}$. The [\\ion{Fe}{ii}] lines originate predominantly at the edges (rear and front) of the \\ion{H}{ii} region. UV continuum pumping as well as collisional excitation seem to be responsible for the observed iron lines. } {} ",
        "introduction": "The Antares nebula is a unique object in the sky. Antares \\mbox{($\\alpha$ Sco A, M 1.5 Iab)} and its blue visual companion $\\alpha$ Sco B (B 2.5 V) are known to be surrounded by a circumstellar envelope seen in absorption in both components which can be used to determine the mass-loss rate of the M supergiant (Deutsch \\cite{deu60}; Kudritzki \\& Reimers \\cite{kud78}; Hagen et al. \\cite{hag87}; Baade \\& Reimers \\cite{baa07}). The common envelope has also been known for a long time as an emission nebula with strong [\\ion{Fe}{ii}] lines (Wilson \\& Sanford \\cite{wil37}; Struve \\& Swings \\cite{str40}; Swings \\& Preston \\cite{swi78}). The most extensive study was that by Swings \\& Preston (\\cite{swi78}) based on high-resolution long-slit photographic Coud\\'e spectra taken with the Mount Wilson 100 inch and Palomar 200 inch telescopes. They found that the ``[\\ion{Fe}{ii}]-rich nebula'' is strong roughly $3\\farcs5$ around the B~star surrounded by a zone of weaker lines which may extend in the NW-SE direction up to $15\\arcsec$. It has been shown by Kudritzki \\& Reimers (\\cite{kud78}) that the [\\ion{Fe}{ii}] emission lines are probably associated with the \\ion{H}{ii} region formed by the B star within the stellar wind of the M supergiant.  The \\ion{H}{ii} region has  been detected and resolved at cm radio wavelengths with the VLA (Hjellming \\& Newell \\cite{hje83}). The maximum \\ion{H}{ii} region radio emission was found to be centered roughly $0\\farcs5$ from the B star on the line connecting the two stars.  Several questions remained unanswered. Why are the Balmer lines, expected in emission due to recombination within the \\ion{H}{ii} region, absent or weak compared to the [\\ion{Fe}{ii}] lines? Why are the ``classical'' \\ion{H}{ii} region emission lines absent or weak? Why are the [\\ion{Fe}{ii}] lines double west of B, but single east of B (between the two stars)? With the advent of UVES at the VLT it was obvious that progress in our understanding of the kinematics of the Antares nebula is now possible due to the high spectral resolution, high pointing accuracy, and stability of UVES. We used the spectrograph at a resolution of 80\\,000 to map the nebula in 2-D with 23 slit positions covering the whole nebula (Fig.~\\ref{slitpos}).  In section 2 we describe the observations and the data reduction which was difficult due to (unexpected) scattered light from the M supergiant even 10$\\arcsec$ away from the bright star. Section 3 gives a phenomenological description of the observations. In section 4 we present the analysis of the UVES spectra as well as a discussion of the obtained results. ",
        "conclusions": "\\subsection{The location of B relative to the plane of the sky} The velocity of the circumstellar \\ion{Ti}{ii} absorption lines in the line of sight had been used together with the assumption of a spherical wind with a constant expansion velocity to estimate a position of B $\\sim 600$ AU in front of the plane of the sky (Kudritzki \\& Reimers \\cite{kud78}). However, the observed episodic nature of mass loss of $\\alpha$ Sco A (4 absorption components) disagrees with these assumptions and the derived location must be considered as erroneous (Baade \\& Reimers \\cite{baa07}).  GHRS/HST spectra show that the \\ion{Al}{iii} absorption lines at +7.8 km\\,s$^{-1}$, which must be formed close to the B star, favor a location of B roughly 224 AU behind the M star (Baade \\& Reimers \\cite{baa07}). We can determine the B star position independently using the velocity of the H$_{\\alpha}$ line at the eastern boundary of the \\ion{H}{ii} region at the slit positions 1$\\farcs$4 and 1$\\farcs$9. The basic underlying assumption is that the M star wind velocity is not severely disturbed and that the width of the emission defines the front and rear side of the \\ion{H}{ii} region. The mean observed velocity in the line of sight is $\\sim 4.8$ km\\,s$^{-1}$, i.e.  7.8 km\\,s$^{-1}$ relative to the M supergiant. Using a constant M star  wind velocity of 20 km\\,s$^{-1}$ (Baade \\& Reimers \\cite{baa07}), we obtain a position angle of 23$^{\\circ}$ behind the M star with an uncertainty of about 5$^{\\circ}$. A measurement of the [\\ion{Fe}{ii}] lines of the \\ion{H}{i} region in front of the \\ion{H}{ii} region, where the lines are formed by UV pumping, gives a nearly identical result.   \\subsection{Geometry and binary orbit} Long-term observations of the visual binary have shown that the orbit is nearly in the line of sight (Hopmann \\cite{hop57}), and the binary separation has become smaller from 3$\\farcs$01 in 1930 (photographic plates, Hopmann \\cite{hop57}) to 2$\\farcs$86 in 1989 (interferometric measurements, McAlister et al. \\cite{mca90}), and to 2$\\farcs$74 in 2005 (lunar occultation, S\\^oma \\cite{som06}). The extrapolation yields a separation of 2$\\farcs$73 in February 2006. With the Hipparcos distance of 185~pc and assuming a circular orbit we obtain a semi-major axis of 549~AU. With 18 M$_{\\odot}$ and 7.2 M$_{\\odot}$ for $\\alpha$ Sco A and B, respectively (Kudritzki \\& Reimers \\cite{kud78}), Kepler's third law yields a binary period of 2562~yrs. The corresponding calculated orbital velocities are $v_{\\rm M} = -1.82\\,\\rm km\\,s^{-1}$ and $v_{\\rm B} = 4.56\\,\\rm km\\,s^{-1}$, with radial components $v_{\\rm M,r} = -1.68\\,\\rm km\\,s^{-1}$ and $v_{\\rm B,r} = 4.20\\,\\rm km\\,s^{-1}$. This is close to the observed values of -3 km\\,s$^{-1}$ and +3 km\\,s$^{-1}$ for the M and B star, respectively (Evans \\cite{eva67}; Kudritzki \\& Reimers \\cite{kud78}), if a small systemic velocity of -1.25 km\\,s$^{-1}$ is added. Unfortunately the observational data do not allow us to establish a reliable orbital solution. If we trust the 19th century measurements compiled by Hopmann (\\cite{hop57}) the orbit cannot be circular. However, for the present work the uncertainties of the binary orbit can be neglected.   \\subsection{The shape of the \\ion{H}{ii} region} The apparent shape of the H$_{\\alpha}$ emission nebula in the sky according to Fig.~\\ref{halp_emiss} is displayed in Fig.~\\ref{hii_region}, together with the radio map of Hjellming \\& Newell (\\cite{hje83}).  Details of the H$_{\\alpha}$ and [\\ion{Fe}{ii}] intensity projected onto the plane of the sky (2-D) have been presented in the sections 3.2.1 and 3.2.4. The emission nebula can be best described by a bow-shock shaped region with its front $\\sim 1 \\arcsec$ east of B which opens slowly to an extent of roughly 6$\\arcsec$ up to $\\sim 3 \\arcsec$ west of B. Beyond 3$\\arcsec$ the ``shock cone'' opens with filamentary structures to an extent of up to nearly 16$\\arcsec$ ($\\pm$ 8$\\arcsec$). In velocity space the ``double line'' character becomes more distinct, and typically the red line edge is sharply limited. We appear to look through a shock cone with a front and rear edge (cf. Fig.~\\ref{fe_emiss} on and slightly west of B).  With the assumption that the M star wind continues at constant velocity through the \\ion{H}{ii} region without being grossly disturbed, we can construct a 3-D image of the [\\ion{Fe}{ii}] emitting gas since each velocity in Fig.~\\ref{fe_emiss} corresponds in that case to a location relative to the plane of the sky (positive velocities behind, negative velocities in front of the plane of the sky). The resulting 2-D slices for each slit position are presented in Fig.~\\ref{fe_emiss_yz}: East of B the emission comes from behind the plane of the sky, as expected, since the B star and the surrounding \\ion{H}{ii} region are behind the plane of the sky. Between the position of B and $\\sim 3\\arcsec$ west of B, the emission comes from a region of cylindrical shape with a diameter of $\\sim 6 \\arcsec$. This extent is almost identical in the plane of the sky  and in the vertical direction. A gradual apparent shift of the [\\ion{Fe}{ii}] emitting gas is seen from 200~AU behind the plane at 1$\\farcs$4, slightly behind the plane of the sky (on average) to an extent of between 740~AU behind to 1480~AU in front of the plane of the sky. The shock cone apparently opens up beyond 3$\\arcsec$ west of B to larger distances in front of the plane of the sky.  \\subsection{Where are the ``normal'' forbidden \\ion{H}{ii} region lines?} The Antares nebula has been considered as ``peculiar'' from its detection (cf.\\ the extensive discussion in Struve \\& Zebergs \\cite{str62}) because its spectrum is dominated by numerous strong [\\ion{Fe}{ii}] lines while the classical \\ion{H}{ii} region lines [\\ion{O}{ii}] 3287~{\\AA}, [\\ion{O}{ii}] 3726/3729~{\\AA}, [\\ion{O}{iii}] 4957/5007~{\\AA}, and [\\ion{S}{ii}] are missing. Only the [\\ion{N}{ii}] lines are present. Since $\\alpha$ Sco B is with B 2.5 V of much later spectral class than the O stars normally exciting an \\ion{H}{ii} region, we expect a lower electron temperature. Swings \\& Preston (\\cite{swi78}) had estimated $T_{\\rm e}\\approx 4000$\\,K from the [\\ion{Fe}{ii}] excitation, while Kudritzki \\& Reimers (\\cite{kud78}) estimated $T_{\\rm e} = 4000$\\,K from the heating/cooling balance. The latter value , however, was based on normal \\ion{H}{ii} region cooling rates dominated by [\\ion{O}{iii}], [\\ion{O}{ii}] etc. and is not applicable according to the present work.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics*{9983fig7.eps}} \\caption{ Theoretical run of the electron temperature in different directions as a function of distance from the B star. The CLOUDY models are presented for 11 equidistant angles in the \\ion{H}{ii} region between 0{\\degr} and 180{\\degr}, where 0{\\degr} corresponds to the connecting line between the two stars. The density distribution used here and in Fig.~\\ref{slit_int} is $n \\sim r^{-2}$ according to constant wind expansion from the M star. } \\label{temperature} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics*{9983fig8.ps}} \\caption{ ``Look back'' contours (solid lines) of \\ion{H}{ii} region borders in a plane perpendicular to the plane of the sky caused by the combination of wind expansion and orbital motion of $\\alpha$ Sco B relative to $\\alpha$ Sco A shown for present and previous phases (-120 yrs to -360 yrs; 130 years being the wind travel time for the projected distance between A and B). This can be interpreted as a qualitative model for the comet like structure of the [\\ion{Fe}{ii}] emission regions. For clarification we overlay the projected brightness distribution of Fig.~\\ref{fe_emiss_yz}. } \\label{hii_hist} \\end{figure}   The key to an understanding of the missing [\\ion{O}{ii}], [\\ion{O}{iii}] lines with the simultaneous presence of the [\\ion{N}{ii}] lines is twofold: electron temperature and abundance. At first, the electron temperature of the \\ion{H}{ii} region must be so low that for normal or slightly reduced O/H ratios [\\ion{O}{ii}] 3726/3729~{\\AA} must be below the detection limit.  For [\\ion{O}{ii}] 3726/3729~{\\AA} we can give an upper limit: we have detected the Balmer line series up to H$_{10}$, H$_{11}$ (close to [\\ion{O}{ii}]) not being visible. Thus the H$_{10}$ intensity, as predicted by recombination theory, can be regarded as an upper limit to [\\ion{O}{ii}] which yields [\\ion{O}{ii}]/H$_{\\alpha} <$ 0.017 for 3726~{\\AA}. Models calculated for the Antares \\ion{H}{ii} region with the ionization code CLOUDY (version 07.02.00, last described by Ferland et al.\\ \\cite{fer98}) lead to $T_{\\rm e} < 5000$\\,K. At such a low $T_{\\rm e}$, [\\ion{N}{ii}]/H$_{\\alpha}$ would be $< 0.1$ (see also Fig. 3 in Mierkewicz et al. \\cite{mie06}) for solar abundances and thus below our detection limit, while the observed [\\ion{N}{ii}] 6583~{\\AA}/H$_{\\alpha}$ is between 0.3 and 0.4 in the central region of the \\ion{H}{ii} region. The solution to this puzzle is that the atmospheric (and wind) composition of the M supergiant must have been altered by the CNO cycle.  When a star ascends the red giant branch the first dredge-up may alter the CNO abundance significantly. The exact amount of the depletion of O and enhancement of N depends on the initial mass. As a consequence it is difficult to derive exact CNO abundances for $\\alpha$ Sco A on the basis of evolutionary tracks, since neither its current mass nor the mass it had on the main sequence are reliably known.  Recently Balachandran et al. (\\cite{bal06}) have published abundance measurements of a number of M type supergiants that suggest $\\rm [N/O] \\sim 0.7$ for $\\alpha$ Sco A.  While the CNO abundance ratios have not been measured in $\\alpha$ Sco A, Harris \\& Lambert (\\cite{har84}) have found that $^{17}$O and $^{13}$C are enriched so that there is also clear observational evidence for CNO cycle material in the envelope of $\\alpha$ Sco A. Our best model for the \\ion{H}{ii} region leads to $\\rm N/N_\\odot = 3.3 \\pm 1$ and $\\rm O/O_\\odot \\la 0.9 \\pm 0.1$, i.e $\\rm [N/O] = 0.56 \\pm 0.2$ which is consistent with the above-mentioned values for M supergiants. The final CLOUDY \\ion{H}{ii} region model used for Fig~\\ref{temperature} and Fig~\\ref{slit_int} has solar abundances except for CNO for which we used the above numbers.  The nondetection of [\\ion{S}{ii}] 6716/6731~{\\AA} can be understood with the same arguments. In classical, relatively cool \\ion{H}{ii} regions, [\\ion{S}{ii}]/[\\ion{N}{ii}] is typically $\\sim 0.3$ for solar abundances. With N being enhanced by a factor of more than 3, the expected [\\ion{S}{ii}]/[\\ion{N}{ii}] ratio is $< 0.1$, consistent with the nondetection in our spectra.  The ionization structure and temperature of the \\ion{H}{ii} region depend primarily on the number of incident Lyman continuum photons emitted by the B star. The most reliable estimate of the Lyman continuum luminosity $L_{\\rm Ly}$ comes from radio measurements of the \\ion{H}{ii} region by Hjellming \\& Newell (\\cite{hje83}). It is possible to adjust the effective temperature of the B star until the measured radio flux is reproduced. These calculations have to be performed iteratively, since the electron temperature $T_{\\rm e}$ of the \\ion{H}{ii} region is a priori unknown. Only the knowledge of $T_{\\rm e}$ allows us to quantify the relation between the Lyman continuum luminosity and the radio flux. A final model that satisfies all observational constraints yields a mean electron temperature of 4900~K. A detailed plot of the temperature distribution in the \\ion{H}{ii} region is shown in Fig.~\\ref{temperature}. It should be noted that \\ion{Fe}{ii} has a considerable impact on the \\ion{H}{ii} region and its electron temperature. As a consequence we apply the CLOUDY code with the more accurate \\ion{Fe}{ii} model using 371 levels as described by Verner et al. (\\cite{ver99}).  \\subsection{What excites the [\\ion{Fe}{ii}] lines?} The [\\ion{Fe}{ii}] lines have been observed in high-density regions such as the Orion Nebula (Bautista et al. \\cite{bau94}; Baldwin et al. \\cite{bal00}) and it has been shown by multilevel collisional-radiative \\ion{Fe}{ii} models that under the Orion Nebula conditions (at $n_{\\rm e} = 10^{4}\\,\\rm cm^{-3}$ and $T_{\\rm e} = 10^{4}\\,\\rm K$) the optical [\\ion{Fe}{ii}] lines (4814, 4277~{\\AA} etc.) are due to UV continuum pumping, while the infrared [\\ion{Fe}{ii}] lines are produced by collisional excitation (Baldwin et al. \\cite{bal96}, Verner et al. \\cite{ver00}). The strong [\\ion{Fe}{ii}] line at 4814~{\\AA} as presented in Fig.~\\ref{fe_emiss} is the result of a pumping route starting from level a$^4$F which is excited to level x$^4$F$^{\\rm o}$. The downward transition to b$^4$F followed by a transition back to the a$^4$F term finally produces the observed line. Details of this and other \\ion{Fe}{ii} pumping processes are discussed thoroughly by Verner et al. \\cite{ver00} (see especially their Fig.~7 showing all relevant routes of the continuum pumping).  In the case of the $\\alpha$ Sco wind we explicitly showed above that in the neutral, cool wind (\\ion{H}{i} region) the [\\ion{Fe}{ii}] lines are visible and are formed by UV pumping. So the question remains: Why is the contrast between \\ion{H}{ii} region and \\ion{H}{i} region (Figs.~\\ref{halp_emiss},~\\ref{fe_emiss}) so strong? At first, the [\\ion{Fe}{ii}] lines are strongest at the interface between the \\ion{H}{ii} and the \\ion{H}{i} region. Part of the explanation for this behavior is the fact that the \\ion{Fe}{iii} zone around the B star nearly fills the \\ion{H}{ii} region (in contrast to the smaller \\ion{Si}{iii} zone). This may lead to the triple velocity structure in [\\ion{Fe}{ii}] clearly visible at position $5\\farcs4$ and $5\\farcs9$ (Fig.~\\ref{fe_emiss}), where the central component comes from \\ion{Fe}{iii} recombination and the `edge' components from close to the ionization front.   \\begin{figure*}[h!tb] \\resizebox{\\hsize}{!}{\\includegraphics*{9983fig9.ps}} \\caption{Comparison of the velocity-integrated H$_\\alpha$ flux with CLOUDY simulations shown for a mass-loss rate of $1.05 \\times 10^{-6}$ M$_{\\odot}$\\,yr$^{-1}$.  Since we have no absolute fluxes, both the observed (solid line) and the theoretical fluxes (dashed line) are normalized.} \\label{slit_int} \\end{figure*}   Secondly, as has been shown by HST spectra, iron (and consequently \\ion{Fe}{ii}) is depleted onto dust by at least a factor of 10 in the M giant wind (Fig.~4 in Baade \\& Reimers \\cite{baa07}). The combination of density enhancement and grain destruction in the shock heated zone might be a further reason for the apparent [\\ion{Fe}{ii}] line enhancement at the \\ion{H}{ii}/\\ion{H}{i} interface. A further argument against pure collisional excitation of the [\\ion{Fe}{ii}] lines within the \\ion{H}{ii} region is our observation that the density enhanced shock zone is visible outwards to at least 7$\\arcsec$ west of B, far beyond the \\ion{H}{ii} region which extends to roughly 3$\\arcsec$ only.  The density enhancements made visible through [\\ion{Fe}{ii}] emission are apparently carried outwards by the M supergiant wind far beyond the \\ion{H}{ii} region and remain visible although \\ion{H}{ii} has recombined. Notice in particular the bright rear side and the fainter front side of the shock cone. The wind travel time $t_{\\rm w}$ for 550~AU is $\\sim 130$ yrs, to be compared with the recombination time $t_{\\rm rec}$ of roughly 18 yrs for a mass-loss rate of $10^{-6}$ M$_{\\odot}$\\,yr$^{-1}$ at the western border of the \\ion{H}{ii} region ($\\sim 3 \\arcsec$ or $\\sim 550$~AU west of B). Though $t_{\\rm rec}$ is short compared to the wind travel time, advection effects may explain the existence of weak H$_\\alpha$ emission outwards to $\\sim 6 \\arcsec$ west of B.  We have schematically sketched this behavior showing the combined effect of B star orbital motion and wind expansion on the movement of a ``former'' \\ion{H}{ii} region (as seen in Fe) relative to the M supergiant. The simple model (Fig.~\\ref{hii_hist}) reproduces the location of the front and the rear edge of former \\ion{H}{ii} region gas in velocity space. With increasing distance from the B star, the median velocity moves from $+$4 km\\,s$^{-1}$ 1$\\arcsec$ E of B over $+$3 km\\,s$^{-1}$ at B to 0 km\\,s$^{-1}$ W of B and $-$5 km\\,s$^{-1}$ at 7$\\arcsec$ W of B, consistent with the -- vertical to the plane of the sky -- cometary shaped \\ion{Fe}{ii} emission region as a result of the wind-carried density enhancements of the \\ion{H}{i}/\\ion{H}{ii} region interface.  From the above discussion it is obvious that a quantitative explanation of the [\\ion{Fe}{ii}] emission requires substantial additional effort: In a final step a non-stationary hydrodynamical model of the \\ion{H}{ii} region moving with the B star through the M star wind has to be computed. In a second step this model will be the basis for a 3D-NLTE radiative transfer model of \\ion{Fe}{ii} including the B star with the UV photons which, as we have shown, is at least partly responsible for the excitation of the [\\ion{Fe}{ii}] lines. Such a model is beyond the scope of the present paper.   \\subsection{Mass-loss rate} To quantify the mass-loss outflow we proceeded in two steps. Following Nussbaumer \\& Vogel (\\cite{nus87}) we calculated in a first step the shape of the \\ion{H}{ii} region as a function of the mass-loss rate assuming an electron temperature of 4000\\,K (Kudritzki \\& Reimers \\cite{kud78}) which is considered to be constant throughout the envelope. The resulting theoretical \\ion{H}{ii} region was projected onto the plane of the sky and compared to the observed H$_{\\alpha}$ emission along the slits. Thus we obtained a tentative mass-loss rate of $8 \\times 10^{-7}$ M$_{\\odot}$\\,yr$^{-1}$ that was used as the initial value for the subsequent step. In order to achieve a self-consistent model we used the ionization code CLOUDY (Ferland et al. \\cite{fer98}) and constructed a static \\ion{H}{ii} region considering all relevant heating and cooling processes. The resulting temperature distribution is shown in Fig.~\\ref{temperature}. Now a more elaborate estimate of the mass-loss rate is possible using the spatially dependent emissivity. The emergent intensity of the H$_{\\alpha}$ emission has been integrated for all slit positions as a function of the perpendicular distance from the central line. The best match with the observed intensity distribution yields mass-loss rates between $8.4 \\times 10^{-7}$ M$_{\\odot}$\\,yr$^{-1}$ and $1.26 \\times 10^{-6}$ M$_{\\odot}$\\,yr$^{-1}$. Fig.~\\ref{slit_int} shows the result for the mean value. There is a clear tendency that the required mass-loss rate increases for the outer slit positions. It is, however, questionable whether this finding reflects a true change in the mass outflow or is only an artifact of inhomogeneities or curved wind trajectories. It should be noted that the first slit position covers only part of the \\ion{H}{ii} region, resulting in a reduced emission.  We are aware that the static results of the CLOUDY code should be interpreted with care, but for optically thin media a correct implementation of the microphysics is more important than kinematic effects. Future work should include dynamical and time-dependent processes. Especially advection terms and hydrodynamical effects have to be considered carefully. Some of the observed features are clearly beyond the capabilities of a stationary model, as discussed above."
    },
    "0809/0809.1399_arXiv.txt": {
        "abstract": "Recovering images from optical interferometric observations is one of the major challenges in the field. Unlike the case of observations at radio wavelengths, in the optical the atmospheric turbulence changes the phases on a very short time scale, which results in corrupted phase measurements. In order to overcome these limitations, several groups developed image reconstruction techniques based only on squared visibility and closure phase information, which are unaffected by atmospheric turbulence. We present the results of two techniques used by our group, which employed coherently integrated data from the Navy Prototype Optical Interferometer. Based on these techniques we were able to recover complex visibilities for several sources and image them using standard radio imaging software. We describe these techniques, the corrections applied to the data, present the images of a few sources, and discuss the implications of these results. ",
        "introduction": "\\label{sec:intro}  %  Image reconstruction is one of the major challenges for optical interferometry. Unlike radio interferometry, where the atmosphere usully changes on long time scales, in optical interferometry the atmosphere fluctuates on times scales of a few ms, resulting in the corruption of the fringe visibility phases. Recently a lot of effort has been devoted by several groups in the development of imaging algorithms that use only squared visibilities ($V^2$) and  closure phases, which are uncorrupted quantities (see Ref.~\\citenum{Buscher94,Ireland06,Lawson04} for a few examples). However, these techniques have to combine multiple measurements, like in the case of closure phases, so they contain less information, resulting in higher noise.  Therefore, one of the challenges to optical interferometric imaging is to recover as much of the phase information as possible. Here we describe two techniques developed by our group using coherently integrated data from the Navy Prototype Optical Interferometer (NPOI). The first technique uses differential phases to recover the phase information in the H$\\alpha$ channel of Be stars, while the second one applies the phase self calibration method. These techniques allowed us to recover complex visibilities and image sources using radio interferometric reconstruction techniques. ",
        "conclusions": "In this paper we discussed the development of two techniques that used coherently averaged NPOI data to recover complex visibilities. This allowed us to image the sources using standard radio interferometry imaging methods. Using the first technique, differential phases, we were able to correct the instrumental and atmospheric contribution to the intrinsic H$\\alpha$ phases of $\\beta$ Lyrae, being able to image the line emission in this source and detect its movement relative to the continuum photocenter. Although this is a powerful technique, it requires a priori knowledge about the structure of the source at some wavelength, and has a limited number of applications. The second technique, self calibration, was used to correct the atmospheric effects to the observed phases, allowing us to image binary and triple systems. This technique is extensively used in radio interferometry, requires closure phases, but is applicable to a wider range of sources. Future plans include the application of this technique to other sources, like rapidly rotating stars, stars with spots and disks, among other sources.  We would like to point out that the techniques presented here make use of all the information collected by the interferometer and do not combine multiple measurements (e.g. closure phases), like in the case of imaging methods that use only $V^2$'s and closure phases. Consequently we should be able to obtain higher precision measurements and/or higher dynamic range images of astrophysical sources."
    },
    "0809/0809.3952_arXiv.txt": {
        "abstract": "We present a catalog of 1,172,157 quasar candidates selected from the photometric imaging data of the Sloan Digital Sky Survey (SDSS).  The objects are all point sources to a limiting magnitude of $i=21.3$ from 8417 deg$^2$ of imaging from SDSS Data Release 6 (DR6).  This sample extends our previous catalog by using the latest SDSS public release data and probing both UV-excess and high-redshift quasars.  While the addition of high-redshift candidates reduces the overall efficiency (quasars:quasar candidates) of the catalog to $\\sim80\\%$, it is expected to contain no fewer than 850,000 bona fide quasars --- $\\sim8$ times the number of our previous sample, and $\\sim10$ times the size of the largest spectroscopic quasar catalog.  Cross-matching between our photometric catalog and spectroscopic quasar catalogs from both the SDSS and 2dF Surveys, yields 88,879 {\\em spectroscopically} confirmed quasars.  For judicious selection of the most robust UV-excess sources ($\\sim500,000$ objects in all), the efficiency is nearly 97\\% --- more than sufficient for detailed statistical analyses.  The catalog's completeness to type 1 (broad-line) quasars is expected to be no worse than 70\\%, with most missing objects occurring at $z<0.7$ and $2.5<z<3.0$.  In addition to classification information, we provide photometric redshift estimates (typically good to $\\Delta z \\pm 0.3$ [$2\\sigma$]) and cross-matching with radio, X-ray, and proper motion catalogs.  Finally, we consider the catalog's utility for determining the optical luminosity function of quasars and are able to confirm the flattening of the bright-end slope of the quasar luminosity function at $z\\sim4$ as compared to $z\\sim2$. ",
        "introduction": "\\nocite{sch63}  The number of known quasars has grown exponentially since their discovery by Maarten Schmidt in 1963 (Fig.~\\ref{fig:fig1}).  There have been relatively frequent compilations of heterogeneous catalogs over the years and the 100, 1000, and 10000 quasar marks were reached in 1967, 1977, and 1998, respectively \\citep[see][and references therein]{hb93,vv06}.  Early quasar discoveries were often based on heterogeneous samples and/or previously existing photometric surveys, so the exact lineage of the growth of homogeneous samples is more difficult to trace.  However, the number of spectroscopically-confirmed, optically-selected quasars in a single homogeneous survey had certainly reached 100 by 1977 \\citep[e.g.,][]{mls77}. The 1000 quasar mark was broken during the Large Bright Quasar Survey (LBQS) in 1991 \\citep{mwa+91,hfc+95}.  The 2dF Quasar Redshift Survey (2QZ; \\citealt{bsc+00}) first cataloged 10,000 quasars by 2001 \\citep{csb+01}, soon followed by the Sloan Digital Sky Survey (SDSS; \\citealt{yaa+00}) Quasar Survey \\citep{shr+07}.  While the number of known quasars continues to grow at a rapid pace \\citep[e.g.,][]{shr+07}, the 100,000 object mark was broken years ahead of the extrapolated trend (see Fig.~\\ref{fig:fig1}) by this groups's {\\em photometric} sample in 2004 (\\citealt{rng+04}; hereafter Paper I). Quasar catalogs used for meaningful statistical analyses are almost always spectroscopic. This is in contrast to galaxies, for which a wealth of major statistical studies utilized purely photometric catalogs \\citep[e.g.,][]{mes+90}. Historically, this has been due to an inability to obtain $\\sim90\\%$ or greater star-quasar separation efficiency to match the typical star-galaxy separation readily obtainable from morphology. For instance, standard UV-excess (UVX) quasar selection \\citep[e.g.,][]{csb+01} is $\\sim50$\\% efficient and the SDSS's official quasar targeting efficiency is $\\sim80$\\% (at best) for bright ($i=19.1$) UVX sources \\citep{rfn+02}. The $\\sim$95\\% efficiency \\citep{rng+04,mbr+06} of our catalog thus heralded the era of statistically useful photometric star-quasar separation, opening up a new avenue for quasar studies.    Using the most recent SDSS data release \\citep{aaa+08}, this paper marks the next milestone by presenting a homogeneous photometric catalog of nearly one {\\em million} quasars.  Unfortunately, with our current approach, this trend will likely moderate in the near future, as this sample covers 8417 deg$^2$ to $i=21.3$ and there are only 41253 deg$^2$ in our sky.  On the other hand, large-scale synoptic surveys such as the Large Synoptic Survey Telescope (LSST; \\citealt{tys02}), the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS; \\citealt{kab+02}), and the Dark Energy Survey (DES; \\citealt{des05}) will, in the next decade, enable another order of magnitude gain by taking advantage of fainter photometric limits and quasar variability.  In the meantime, an alternative path allows us to anticipate an explosion in the number of obscured (so-called type 2) quasars \\citep{ant93}, which are expected to outnumber the type 1 quasars cataloged herein by up to a factor of 4-to-1 \\citep[e.g.,][]{lss+04,tuc+04,bh05,pwh+08,rzs+08}, and whose numbers will increase as the {\\em Spitzer Space Telescope} maps ever larger areas of sky during its warm mission.   The need for robust photometric classification is rapidly becoming apparent and will be an absolute necessity by the time LSST and Pan-STARRS are fully underway.  Even with multi-object spectrographs observing thousands of objects per square degree at a time, the small fields and relatively long exposure times mean that it will simply never be possible to obtain spectra of all of objects identified. In addition, new science goals nearly always demand increased sample size.  Indeed, this has been aptly demonstrated by previous work on the far smaller versions of this catalog. Much of the new science that used our catalogs detected subtle cosmological effects that were previously impossible without a large quasar catalog, but also highlighted the need for more extensive samples with which to study elusive aspects of cosmology and the quasar population.   For example, \\citet{mbr+06} explored quasar clustering using the Paper~I catalog --- the first such study of quasar evolution in a photometric catalog --- and found results consistent with spectroscopic surveys. This study was expanded in \\citet{mbn+07}, providing a luminosity baseline large enough to uniquely constrain topical models of quasar activity, but still with too few objects with which to constrain any luminosity dependence to quasar clustering. \\citet{hso+06} used the catalog to enhance their study of binary quasars, and detected the first definitive evidence for excess quasar clustering on small scales. In \\citet{mbr+07} we further examined small-scale quasar clustering, providing a homogeneous catalog of binary quasar candidates.  Myers et al.\\ (2008; in prep) present spectroscopic observations of pairs of photometric quasar candidates and are able to place only weak constraints on any redshift dependence to small-scale quasar clustering at $z < 2$, providing yet more impetus to produce a larger catalog over a wider redshift range. These papers on the clustering of our photometric quasars provided critical input to the clustering analysis done by \\citet{hlh+07}.  Cross-correlating with the cosmic microwave background, \\citet{gcn+06} and Giannantonio et al.\\ (2008, submitted) used the large number of photometric quasars to constrain dark energy using the Integrated Sachs-Wolfe (ISW) effect \\citep{sw67}, the first detection of the ISW effect using optically-selected quasars. These measurements represent one of the most robust measurements of dark energy at high redshift and are found to be consistent with predictions for flat $\\Lambda$CDM models (see Giannantonio et al. 2008). Finally, after many years of contradictory results in the field, \\citet{smr+05} used photometric quasars to categorically measure cosmic magnification bias, detecting the effect of gravitational lensing by foreground galaxies on quasar source counts at $\\sim8\\sigma$.  This paper is laid out as follows.  Section~2 briefly describes the data.  Section~3 reviews the Bayesian selection algorithm, discusses the changes from \\citet{rng+04}, and describes the construction of the training and test data sets.  The catalog itself (in Tables~1, 2, and 3) is presented in \\S~4.  Various catalog properties and diagnostics of the efficiency and completeness are described as is our prescription for limiting the catalog to particularly robust sub-samples.  We also discuss matching of the catalog to non-optical object catalogs and the determination of photometric redshifts. Finally, a rough analysis of the number counts and luminosity function are given in \\S~5. ",
        "conclusions": "Using a novel Bayesian algorithm we identify 1,172,157 quasars candidates from a sample of over 40 million SDSS point sources.  The overall efficiency of the catalog is $\\sim$80\\% and the catalog is expected to contain a minimum of 850,000 bona-fide quasars.  A UVX subsample, in excess of 500,000 objects has an expected efficiency of over 97\\%.  Additional information (redshift-dependent selection and radio, X-ray, and proper motion catalog matching) is provided in the catalog so that users can select sub-samples that are optimal for any particular follow-up investigation.  Photometric redshifts are estimated for the full sample and are expected to be accurate to $\\pm0.3$ roughly 80\\% of the time, with outliers being statistically well defined.  Cross-comparison with spectroscopically confirmed type 1 quasars in the COSMOS field suggests that the sample is at least 70\\% complete and may recover additional objects missed by X-ray and radio selection methods.  Careful analysis of the catalog could be used to create the deepest yet optical quasar luminosity function; simple arguments herein confirm the flattening of the QLF slope at $z\\sim4.25$ as compared with $z\\sim2$.  A final installment of this catalog will come after the seventh SDSS data release in the fall of 2008 and should bring the total number of quasars over the one million mark."
    },
    "0809/0809.3449_arXiv.txt": {
        "abstract": "Planet-planet scattering is the leading mechanism to explain the large eccentricities of the observed exoplanet population.  However, scattering has not been considered important to the production of pairs of planets in mean motion resonances (MMRs).  We present results from a large number of numerical simulations of dynamical instabilities in 3-planet systems.  We show that MMRs arise naturally in about five percent of cases.  The most common resonances we populate are the 2:1 and 3:1 MMRs, although a wide variety of MMRs can occur, including high-order MMRs (up to eleventh order).  MMRs are generated preferentially in systems with uneven mass distributions: the smallest planet is typically ejected after a series of close encounters, leaving the remaining, more massive planets in resonance.  The distribution of resonant planets is consistent with the phase-space density of resonant orbits, meaning that planets are randomly thrown into MMRs rather than being slowly pulled into them.  It may be possible to distinguish between MMRs created by scattering vs. convergent migration in a gaseous disk by considering planetary mass ratios: resonant pairs of planets beyond $\\sim$ 1 AU with more massive outer planets are likely to have formed by scattering.  In addition, scattering may be responsible for pairs of planets in high-order MMRs (3:1 and higher) that are not easily populated by migration.  The frequency of MMRs from scattering is comparable to the expected survival rate of MMRs in turbulent disks.  Thus, planet-planet scattering is likely to be a major contributor to the population of resonant planets. ",
        "introduction": "The current sample of exoplanets exhibits several interesting dynamical features (Butler \\etal 2006): here we focus on two of these.  First, there is a vast range of planetary eccentricities, from zero to $>0.9$, with a median of 0.2 (0.27 for planets past 0.1 AU that have not been affected by tides; Rasio \\etal 1996; Jackson \\etal 2008).  Second, mean motion resonances (MMRs) in multiple planet systems appear to be relatively common (e.g., Marcy \\etal 2001).  There are 31 currently-known multiple planet systems comprising 44 pairs of adjacent planets, of which ten (23\\%) show some evidence of resonances (Table 1).  However, the evidence for resonances is tentative for all but a few cases.  Dynamical instabilities in systems of two or more planets can explain the wide eccentricity distribution of exoplanets (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997).  Instabilities arise on timescales that are related to the planets' initial separation (Marzari \\& Weidenschilling 2002; Chatterjee \\etal 2008), and lead to close encounters between planets and subsequent ejections or mergers.  In the aftermath of close encounters, the surviving planets can statistically reproduce the observed eccentricity distribution of exoplanets (Adams \\& Laughlin 2003; Juric \\& Tremaine 2008; Chatterjee \\etal 2008).  MMRs are thought to arise primarily from convergent migration in gaseous protoplanetary disks (Snellgrove \\etal 2001; Lee \\& Peale 2002).  Indeed, models show that capture in the 2:1 and 3:2 MMRs is a particularly common occurrence (Thommes 2005; Pierens \\& Nelson 2008; Lee \\etal 2008).  However, MRI-derived turbulence can act to remove planets from resonance.  Adams \\etal (2008) estimate that only 1\\% of resonant systems should remain for a disk lifetime of 1 Myr.  In this paper we attempt to reconcile the planet-planet scattering scenario with the population of resonant exoplanets.  We numerically investigate dynamical instabilities in systems of three planets located at $\\sim$ 2-10 AU with a variety of mass distributions.  We find that MMRs are a common occurrence, arising in 5-10 percent of unstable systems.  Our simulations populate a range of MMRs, including the 2:1, 3:2, 3:1, 4:1 and extending up to much higher-order (Table 2).  MMRs are populated by scattering at random into stable regions; the density of resonant orbits (i.e., the fraction of phase space that undergoes resonant oscillations) is consistent with the scattered resonant systems.  We propose several ways to discriminate between scattering and convergent migration as the source of exoplanet MMRs. ",
        "conclusions": "Planet-planet scattering creates MMRs.  The typical path to a resonant system involves several close encounters between one smaller and two larger planets. After a relatively short time of instability, the smaller planet is destroyed, usually via ejection (78\\% of all cases) or collision (22\\%), leaving behind a pair of resonant planets.  Relatively weak instabilities are probably very common in planetary systems; one may even have occurred in our own Solar System (Thommes \\etal 1999).  Nonetheless, only a fraction of unstable systems produce resonant planets, typically 5-10\\%.  These systems have large libration amplitudes and occupy a range of low-and high-order MMRs (Table 2). Most of these resonances are indefinitely stable; we integrated the 170 resonant systems for an additional 1 Gyr and only 9 (5\\%) left the resonance, 4 cases leading to an additional system instability.  \\begin{deluxetable}{c|c|c|p{2cm}}[t] \\scriptsize \\tablewidth{0pt} \\tablecaption{Resonances from scattering simulations} \\renewcommand{\\arraystretch}{.6} \\tablehead{ \\\\ \\colhead{Set} & \\colhead{Nsims ---} & \\colhead{N(\\%) in} & \\colhead{MMRs}\\\\ \\colhead{ } & \\colhead{unstable(\\%)} & \\colhead{MMRs} & \\colhead{(\\%)}} \\startdata Mixed1 & 965--569 (59\\%) & 27 (4.7\\%) & 2:1 (1.6\\%), 3:1 (1.6\\%), 4:1 (0.7\\%), 5:1, 6:1, 7:2, 9:2\\\\ Mixed2 & 982--744 (76\\%) & 52 (7\\%) & 3:2 (0.8\\%), 2:1 (2.4\\%), 3:1 (1.1\\%), 5:3, 4:1, 5:2, 5:1, 7:3, 6:1, 7:2, 8:3, 9:4, 10:3, 12:5, 11:2, 14:5\\\\ Mequal:3$M_J$ & 368--241 (65\\%) & 1 (0.4\\%) & 7:1\\\\ % Mequal:$M_J$ & 452--232 (51\\%) & 4 (1.7\\%) & 2:1, 3:1, 4:1 (0.9\\%)\\\\ Mequal:$M_{Sat}$ & 390--362 (93\\%) & 14 (3.9\\%) & 2:1 (1.9\\%), 3:1 (0.6\\%), 5:2, 5:1 (0.6\\%), 6:1, 7:1\\\\ Mequal:30$\\mearth$ & 367--365 (99\\%) & 10 (2.7\\%) & 2:1 (1.9\\%), 3:1, 11:6\\\\ Mgrad:JSN & 250--206 (82\\%) & 13 (6.3\\%) & 2:1 (1\\%), 3:1 (1.9\\%), 4:1, 5:2, 5:1, 7:3, 8:3, 13:2\\\\ Mgrad:NSJ & 245--221 (90\\%) & 30 (14.6\\%) & 2:1 (9.2\\%), 3:1 (2.3\\%), 5:2, 7:3, 7:2, 11:5\\\\ Mgrad:3JJS & 250--150 (60\\%) & 4 (2.7\\%) & 3:1, 4:1 (1.3\\%), 11:3\\\\ Mgrad:SJ3J & 245--219 (89\\%) & 16 (6.5\\%) & 3:1 (5.7\\%), 4:1 \\enddata \\end{deluxetable}  It may be possible to tell apart resonant exoplanets created via scattering from those created via convergent migration.  In fact, only one resonant system appears to be inconsistent with a scattering origin due to its very low-amplitude libration (GJ 876; Marcy \\etal 2001).  If two planets are trapped in the 2:1 or 3:2 MMR and the inner planet is the more massive, then migration can be stopped or even reversed (Masset \\& Snellgrove 2001; Crida \\& Morbidelli 2007).  However, if the outer planet is the more massive then inward migration continues.  In contrast, scattering produces planets in a variety of MMRs (including the 3:2 and 2:1) with a wide range in mass ratios and a preference for the outer planet to be more massive.  Thus, scattering is likely to be responsible for systems past $\\sim$ 1 AU with 2:1 or 3:2 resonant planets and a more massive outer planet.  The HD 128311 and $\\mu$ Arae systems are good candidates for creation via scattering (Table 1; see also S\\'andor \\& Kley 2006).  Several extra-solar systems show tentative evidence for high-order MMRs -- 4:1, 5:1 and even 7:1 (Table 1; Johnson \\etal 2007; Correia \\etal 2005; Gregory 2007; Short \\etal 2008).  No study to date has shown that migration could capture planets in MMRs of higher order than 2:1, although we encourage expanded studies of this process.\\footnote{Highly-damped bodies can undergo resonant shepherding by the 6:1 or even 8:1 MMRs (Raymond \\etal 2006; Mandell \\etal 2007).  However, as bodies grow the damping decreases, and shepherded planets do not survive in resonance.}  Scattering produces a wide range of high-order resonances (Table 2).  Thus, if the current candidate high-order MMR systems are confirmed (Table 1), then scattering is likely to be the responsible mechanism.  Turbulence in gaseous disks may destroy MMRs, leaving perhaps only $\\sim$ 1\\% of planet pairs in resonance (Adams \\etal 2008).  This effect is stronger for higher-order MMRs.  However, the timescale for MMR destruction is sensitive to the strength of MRI turbulence which is very uncertain. In particular, if the MRI is fully or partially suppressed by the low ionization fraction in the inner protoplanetary disk (Gammie 1996) then the survival prospects for resonant planets would be improved.  Nonetheless, if only a few percent of systems remain in resonance, then scattering and migration may provide a comparable number of MMRs.  Our simulations do not account for any dissipation.  However, instabilities may occur while some gas remains in the disk (Moeckel \\etal 2008).  In that case, MMRs could still result from scattering (Lee \\etal 2008) and perhaps have smaller libration amplitudes.  In conclusion, we have identified a new mechanism for the creation of exoplanet systems in MMRs.  Unfortunately the current data do not allow a conclusive determination of a resonance, let alone precise descriptions of the resonant argument oscillation.  Nonetheless, our scattering model has important distinctions from the convergent migration model: high-order MMRs, large-amplitude resonant libration, and low-order MMRs with $M_{inner}/M_{outer} < 1$.  As the orbital properties of exoplanets are better determined, it should be possible to distinguish between these scenarios.  \\vskip .2in We thank Google for access to their machines.  We are grateful to Greg Laughlin, Dimitri Veras, and an anonymous referee for helpful input. S.N.R. was supported by the NASA Postdoctoral Program administered by Oak Ridge Associated Universities through a contract with NASA.  R.B. acknowledges support from NASA's PG\\&G grant NNG05GH65G and NASA Terrestrial Planet Finder Foundation Science grant 811073.02.07.01.15."
    },
    "0809/0809.2954_arXiv.txt": {
        "abstract": "Accurate knowledge of the telescope's point spread function (PSF) is essential for the weak gravitational lensing measurements that hold great promise for cosmological constraints.  For space telescopes, the PSF may vary with time due to thermal drifts in the telescope structure, and/or due to jitter in the spacecraft pointing (ground-based telescopes have additional sources of variation). We describe and simulate a procedure for using the images of the stars in each exposure to determine the misalignment and jitter parameters, and reconstruct the PSF at any point in that exposure's field of view. The simulation uses the design of the SNAP{\\footnote{http://snap.lbl.gov}} telescope. Stellar-image data in a typical exposure determines secondary-mirror positions as precisely as $20\\,{\\rm nm}$. The PSF ellipticities and size, which are the quantities of interest for weak lensing are determined to $4.0 \\times 10^{-4}$ and  $2.2 \\times 10^{-4}$ accuracies respectively in each exposure, sufficient to meet weak-lensing requirements.  We show that, for the case of a space telescope, the PSF estimation errors scale inversely with the square root of the total number of photons collected from all the usable stars in the exposure. ",
        "introduction": "\\label{sec:intro}  The accelerated expansion of the universe is one of the most puzzling astrophysical discovery of the century. The proposed explanations are dark energy, modified gravity and feedback from density fluctuations. To explore the mystery, a few large astronomical surveys are underway (DES{\\footnote{http://www.darkenergysurvey.org}}, ACT{\\footnote{http://www.physics.princeton.edu/act}}, SPT{\\footnote{http://pole.uchicago.edu}}) or in planning stages (SNAP, DESTINY{\\footnote{http://destiny.asu.edu}}, LSST{\\footnote{http://www.lsst.org}}, EUCLID{\\footnote{The merger of DUNE (http://www.dune-mission.net) and SPACE \\citep{SPACE}}}). The sensitivity of these surveys to the expansion of the universe comes from both cosmological distances and growth of density perturbations as a function of the cosmic time or redshift. The probes utilized are Type-Ia supernova, weak gravitational lensing (WL), baryon acoustic oscillations, galaxy cluster counting (selected by optical or using Sunyave-Zeldovic effect), and the integrated Sachs-Wolfe effect (ISW). Among all these probes, weak lensing is potentially the most rewarding one if systematics are well under control. The dominating systematics for weak lensing measurements include galaxy shape measurement errors \\citep{Huterer06,STEP1,STEP2, Stabenau07, Amara07}, photometric redshift errors \\citep{Huterer06, Maetal06, MaBernstein08}, uncertainties of the matter power spectrum \\citep{Huterer05, Bernstein08}, galaxy intrinsic alignment \\citep{King02, King03, Heymans03, King05, Mandelbaum06,Heymans06,Hirata07, Bridle07, Lee08, Joachimi08}, and higher order effects such as reduced shear \\citep{White05, Dodelson06, Shapiro08}, Born approximation \\citep{Cooray02, Shapiro06}, and source clustering \\citep{Bernardeau98, Hamana01, Schneider02}.  This paper is concerned with reducing the systematic errors in galaxy shape measurements.  For future weak lensing surveys, the tolerable RMS multiplicative calibration error on WL shear is about $10^{-3}$ \\citep{Huterer06,Amara07}. Additive errors in galaxy shear should also be held to $<10^{-3.5}$. Mis-estimation of the PSF will propagate into systematic errors in the shear.  The size of the PSF must therefore be determined to better than 1 part in $10^3$ to avoid unacceptable multiplicative shear error; likewise the PSF ellipticity must be known to $10^{-3}$ or better to avoid unacceptable additive shear systematic \\citep{PaulinHenriksson}.  Unfortunately the PSF of real telescopes changes with time, as well as with color and field position. Effects that could change the PSF include thermal expansion of mirrors and support structures; pointing jitter due to structural vibrations and tracking errors; for ground-based telescopes there are additionally gravity loading, atmospheric distortions, and wind loading. Careful engineering of the telescope, mount, and housing can minimize these effects, but there are limits to what can be achieved with finite resources. Even for a space telescope it is prohibitively expensive to guarantee PSF stability to $<1$ part in $10^3$. On the other hand there will be stars in each image to diagnose the PSF behavior during each exposure. Recall that the WL analysis requires part-per-thousand {\\em knowledge} of the PSF at each location and each exposure, not necessarily that the PSF be {\\em constant} to this level. In this work, we study how well the PSF can be constrained using stellar images, using the proposed space-based SNAP telescope as a case study.   The outline of the paper is as follows. We present an overview of the task we are trying to accomplish in \\S \\ref{sec:task}. In \\S \\ref{sec:PSF}, we provide details of the modeling of the SNAP PSF. We describe the fitting procedure for the misalignment and jitter parameters in \\S \\ref{sec:procedure}. The results of the fit are presented in \\S \\ref{sec:result} and we conclude in \\S \\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  Our simulation of ``morphometry'' for the SNAP telescope demonstrates that the $\\approx2000$ well-measured stars in a typical exposure contain sufficient information to reduce the errors in the modeled PSF ellipticity and size to $4.0 \\times10^{-4}$ and $2.2 \\times10^{-4}$ respectively, giving significant margin to meet the $\\approx10^{-3}$ level needed to reduce weak-lensing systematic errors below statistical errors of future surveys. For the SNAP telescope design and focal plane, we find $e^{\\rm RMS} \\approx 1.8 / \\sqrt{N_\\gamma}$ and $\\sigma_\\star^{\\rm RMS} \\approx 1.1 / \\sqrt{N_\\gamma}$.  PSF estimation error in morphometry will be only part of the shape-measurement error budget, so this margin is important. Other potential source of errors in PSF estimation include charge transfer inefficiency (CTI), data compression artifacts, and chromatic PSF dependence that causes galaxies' PSFs to differ from stellar PSFs. Shape measurement errors can also arise in the deconvolution process even if the PSF is known precisely \\citep{STEP1, STEP2}. To have a successful weak lensing mission, the sum of these errors must meet the weak lensing requirement.  We have also assumed that the only time-variable aspects of the PSF are the secondary mirror alignment, and small pointing jitter.  The SNAP spacecraft is engineered to take advantage of the extremely stable space environment so that these are the only relevant degrees of freedom. A ground based observatory would suffer in addition the effects of wind, gravity loading, and seeing, which are complicated to model, potentially involving a large number of degrees of freedom.  We have seen in the SNAP case that adding 6 jitter degrees of freedom to the PSF model makes the PSF model errors twice as large as when we fit for only 5 misalignment parameters, so it seems likely that the PSF modeling performance will be degraded by larger number of parameters in the system."
    },
    "0809/0809.2486_arXiv.txt": {
        "abstract": " ",
        "introduction": "Our view of the formation of the Milky Way has changed dramatically since the discoveries of complex substructures in the Galaxy. The Milky Way has four spatially and kinematically distinct components; thin disk, thick disk, halo, and the bulge. However, recent discoveries of irregular structures like the Sagitarrius dwarf tidal stream and the Monoceros stream show that the formation of galaxies is not a steady process resulting in a smooth distribution of stars. Instead, galaxies are constantly shaped by the infalling smaller galaxies (see Juric et al. 2008 and the references there in).  The Sloan Digital Sky Survey ($SDSS$) and the Sloan Extension for Galactic Understanding and Exploration ($SEGUE$) are providing photometry and spectroscopy for a very large number of stars that can be used to study the stellar populations in the Galaxy. Such a study involving $\\approx20,000$ stars with the SDSS spectra demonstrated that the Galactic halo is likely to have two components; an inner and an outer component (Carollo et al. 2007). These two components seem to have different spatial density profiles, stellar orbits, and metallicities. \\cite{Carolloetal2007} suggest that the outer halo formed through dissipationless chaotic merging of smaller subsystems within a pre-existing dark matter halo, supporting the complex galaxy formation scenarios. \\cite{Ivezicetal2008} use the SDSS photometry data to study the metallicity distribution of disk and halo stars. However, their dataset is not deep enough to test the dichotomy of the halo.  The Canada$-$France$-$Hawaii Telescope Legacy Survey ($CFHTLS$) provides a valuable new source to study the metallicity distribution of the Galactic stellar populations. The third data release of the Deep survey now provides $ugriz$ photometry down to a limiting magnitude of 26. $CFHTLS$ lack spectroscopic observations, which would provide the most accurate metallicities. However, imaging data can also be used to deduce stellar metallicities through the traditional UV excess method (Wallerstein 1962; Sandage 1969). The metal absorption lines mostly affect the UV part of the spectrum, therefore metal$-$rich and metal$-$poor stars with the same effective temperature (or the same $B-V$ color) can be differentiated based on their UV colors. In addition, imaging data enables us to obtain a complete flux-limited sample of stars and push the studies of stellar metallicities to fainter magnitudes than the spectroscopic surveys. Here we use the $CFHTLS$ data for the D4 Field to perform such an analysis. Section~2 describes the data and the selection of point sources. The derivation of metallicites from the photometric data and its application to the $CFHTLS$ $ugriz$ system is discussed in Section~3, whereas the results from this analysis are discussed in Section~4.  \\section {The Data}  The CFHTLS D4 field is located at RA = $22^h 15^m 31^s$, DEC = $-17^{o} 43^{'} 56^{''}$ (J2000), ${\\it l}=39^{o}$, and ${\\it b}=-53^{o}$. The TERAPIX T003 public data release for the D4 field covers an area of 0.86 deg$^2$ and provides photometry with limiting AB magnitudes of $u=$ 26.5, $g=$ 26.3, $r=$ 26.4, $i=$ 26.0, and $z=$ 25.0, respectively. The high galactic latitude of this field is similar to the fields targeted by the $SDSS$.  Without follow-up spectroscopy, we have to rely on morphological classification to select point sources. \\cite{Schultheisetal2006} and \\cite{Limbozetal2008} showed that the star$-$galaxy separation is possible through the use of the half$-$light radius (hereafter, HLR) measurements from the $i$-band images. HLR corresponds to the radius that encloses 50\\% of the total flux from the source. Figure 1 shows the HLR versus $i$-band magnitudes for the sources in the D4 field. Objects with $i < 17$ mag are saturated, and therefore have unreliable HLR measurements. Using spectroscopic data from the VLT$-$VIRMOS Deep Survey ($VVDS$; Le Fevre et al. 2005), \\cite{Limbozetal2008} demonstrated that the spectroscopically confirmed stars have HLR($i)<2.82$ and that the star$-$galaxy separation is reliable down to $i=21$ mag. Since the $CFHTLS$ goes five magnitudes deeper than this, we now look into pushing the point source classification limit to fainter magnitudes.  \\begin{figure}[t!] \\resizebox{\\hsize}{!} {\\includegraphics[]{f01.eps}} \\caption{HLR versus $i$-band for the D4 field. Note that objects with $i < 17.00$ are saturated. Star$-$galaxy separation seems to be robust down to $i$ = 22 mag. The solid horizontal line marks the boundary for selection of point sources ($HLR < 2.82$).} \\label{figure1} \\end{figure}  Figure~2 presents $(u-g)$ versus $(g-r)$ color$-$color diagrams for four different magnitude bins. The top left panel shows the color$-$color diagram for point sources (black points) and resolved sources (red points) for $i=20-21$ mag, where the point source classification seems reliable. This panel reveals a tight stellar sequence, and confirms the expectations that stars can be identified reliably down to 21 mag. The other panels in this figure show the color$-$color diagrams for $i=21-22, 22-23$ and $23-24$ mag. These three panels show that the stellar sequence is still clearly visible for $i=21-22$, but the galaxy contamination starts to become a problem for $i>22$ mag. We count the number of sources that are classified as point-like and that lie to the left of the stellar sequence, and compare it with the number distribution of resolved sources with HLR($i)> 2.82$ to estimate contamination of the stellar sequence by galaxies. We estimate that the galaxy contamination is less than 4\\% in the range $(g-r)=0.1-0.7$ mag for $i=21-22$ mag. An analysis of the number distribution of resolved and un-resolved sources in the other panels show that the galaxy contamination ranges from 7\\% to 23\\% for $i=22-23$ mag, and even higher for $i>23$ mag.  Therefore, we restric our sample to the sources with $i < 22$ mag  and $HLR < 2.82$. Using the above criteria, we classify 7348 sources as stars.   \\begin{figure}[b!] \\centering \\includegraphics[scale=0.4, angle=0]{f02.eps} \\caption{$(u-g)$ versus $(g-r)$ color$-$color diagram for different $i$-filter magnitude limits. Black points shows stars ($HLR < 2.82$), whereas red points represent galaxies ($HLR \\ge 2.82$).} \\end{figure}  We use the dust maps of \\cite{Schlegeletal1998} to de-redden the photometry. E($B-V$) for the D4 field ranges from 0.023 to 0.031 mag, with a mean of 0.027 mag. The small range of reddening correction for the D4 field shows that this correction does not have a signifincant effect on our analysis. Perhaps the most important correction required for the photometry is the correction from the $MegaCam$ system to the standard $ugriz$ filter system. Studying the differences between the $ugriz$ system and the $MegaCam$ system, \\cite{Clemetal2008} find that the $MegaCam$ $griz$ filters are consistent with the standard system, though the $u$ band data need to be corrected using a third order polynomial. The $MegaCam$ $u-$band filter was selected to take advantage of the better UV transparency of the $CFHT$ and it is slightly redder than the $SDSS$ $u-$ band filter. \\cite{Clemetal2008} compare theoretical isochrones in the standard system and the corrected $MegaCam$ system and they find that the isochrones overlap each other within 0.01 mag for both dwarfs and giants with colors $u-g$ = 0.0 to +3.0 mag and $g-i=-0.5$ to +3.5 mag, and with metallicities [Fe/H] = $-2.3$ to 0.0. We use the transformation given by \\cite{Clemetal2008} (see their Figure 5) to convert the $MegaCam$ $u-$band photometry into the standard $u-$band system. In addition, a correction is required to transform the photometry from the primed system to the $SDSS$ 2.5 m telescope's unprimed $ugriz$ system. We use the transformations given by \\cite{Tuckeretal2006} to transform the photometry to the unprimed system.  The apparent $g$ magnitude distribution for our sample of stars is shown in Figure 3a. The number of stars increases with increasing magnitude until $g=$ 22.1 mag, which we use as the completeness limit. For a $\\approx$ G5 main$-$sequence star with $(g-r)$ = 0.4 mag (see Table 1 in Ivezic et al. 2008), this limit corresponds to a Galactocentric distance of 21 kpc and a Galactic plane distance of $\\approx$17 kpc. Figure~3b shows the $(g-r)$ color distribution for our sample. There are two distinct peaks located at $\\approx$ 0.20 mag and 1.15 mag, corresponding to F stars and late type M stars, respectively. We select our sample of F and G type stars by restricting the sample to the stars with $0.20 < (g-r) < 0.60$ (shown by vertical dashed lines). This results in a sample of 1315 F and G stars. Majority of these stars lie above the synthetic $(u-g)$ vs. $(g-r)$ relation of \\cite{Pickles1998}, indicating that they are metal$-$poor disk and halo stars.  \\section {The Metallicity Distribution}  Using $SDSS$ spectroscopy and photometry, \\cite{Ivezicetal2008} have developed a metallicity estimation method based on the $(u-g)$ and $(g-r)$ colors. Their method reproduces the metallicities obtained from low resolution $SDSS$ spectroscopy with a root-mean-square scatter of 0.2 dex. They also derive a photometric parallax relation using $ugriz$ photometry of globular clusters. The agreement between the $CFHT/MegaCaM$ and the $SDSS$ $ugriz$ filter sets, after correcting the $MegaCam$ $u-$band photometry (Clem et al. 2008), enables us to use their photometric metallicity and parallax methods on our dataset. Photometric metal abundances of the F and G dwarfs with $0.2<(g-r)\\leq0.4$, are determined via the following equation of \\cite{Ivezicetal2008}.  \\begin{eqnarray} [Fe/H]=- 4.37-8.56(u-g)+15.5(g-r)\\nonumber\\\\ -39.0(u-g)(g-r)+23.5(u-g)^2+20.5(g-r)^2\\nonumber\\\\ +12.1(u-g)^2(g-r)+7.33(u-g)(g-r)^2\\nonumber\\\\ -10.1(u-g)^3-21.4(g-r)^3 \\end{eqnarray}  For G dwarfs with $(g-r) > 0.4$, we replaced the $(u-g)$ color in the above equation with $(u-g)-2(g-r)+0.80$. We estimate the absolute magnitudes of our sample of stars using the following relations from \\cite{Ivezicetal2008}.  \\begin{eqnarray} M_{r}(g-i,[Fe/H])=M_{r}(g-i)+\\Delta M_{r}([Fe/H]) \\end{eqnarray} \\begin{eqnarray} \\Delta M_{r}([Fe/H]) = 4.50-1.11[Fe/H]-0.18[Fe/H]^2 \\end{eqnarray} \\begin{eqnarray} M_{r}(g-i)=-5.06+14.32(g-i)-12.97(g-i)^2\\nonumber\\\\ +6.127(g-i)^3-1.267(g-i)^4+0.0967(g-i)^5 \\end{eqnarray}  \\begin{figure} [t!] \\centering \\resizebox{\\hsize}{!} {\\includegraphics[]{f03.eps}} \\caption{(a) $g$ magnitude distribution. Drop in $g= 22.10$ is taken as the completeness limit. (b) $(g-r)$ distribution. The presence of the two main peaks at $(g-r)\\sim0.20$ and $\\sim1.15$ are seen. Vertical dashed lines show the color index range $0.20 < (g-r) < 0.60$, which corresponds to F and G type stars.} \\end{figure}  \\begin{figure} [h!] \\centering \\resizebox{\\hsize}{!} {\\includegraphics[]{f04.eps}} \\caption{(a) $(u-g)$ color and (b) [Fe/H] distribution for our sample of 1315 F and G type stars. Solid lines represent two Gaussian distributions fitted to the data.} \\label{figure4} \\end{figure}  \\begin{figure}[h] \\includegraphics[scale=0.45, angle=0]{f05.eps} \\caption{The Metallicity distribution as a function of the distance from the plane of the Galaxy.} \\label{figexample} \\end{figure}  Figure 4 displays the $u-g$ color distribution and the derived metallicities for our sample of F and G stars. Both of these distributions are bimodal, and they are fairly well fitted with two gaussians. \\cite{Ivezicetal2008} also found a bimodal $u-g$ color distribution for the $SDSS$ sample with two peaks at $u-g\\approx 0.9$ and 1.1 mag. The two peaks observed in the color distribution of our sample are located at $u-g= 0.86 \\pm 0.13$ and $1.10 \\pm 0.41$ mag. The agreement between the colors of our sample of F and G stars and that of the $SDSS$ F and G stars shows that the transformations between the $MegaCam$ photometry and the standard $ugriz$ system are reliable. The two peaks in the $u-g$ color distribution correspond to metal$-$poor halo stars and more metal$-$rich and closer disk stars, respectively. The width of 0.13 mag is significantly larger than the median $u-g$ error of 0.02 mag at $g=22$ mag, providing a measure of the intrinsic variations in the $u-g$ color and the metallicity distribution.  The observed peaks in the $u-g$ color distribution translates into two peaks at [Fe/H] = $-0.78 \\pm 0.39$ and $-1.62 \\pm 0.82$ in the metallicity distribution. Figure 5 shows this distribution as a function of increasing distance from the Galactic plane. The mean metal abundances for the fitted gaussians are also listed in Table 1. The mean metal abundance shifts as a function of galactic plane distance. The thick disk dominates at $z = 1-4$ kpc, and a small contribution is still visible up to $8-12$ kpc, after which only the halo stars with [Fe/H] $\\approx -1.7$ are visible. ",
        "conclusions": "We have performed a morphological selection of point sources and analyzed the metallicity distribution of distant F and G type stars in the $CFHTLS$ D4 field. We have used $MegaCam$ photometry transformed to the standard $ugriz$ filter system, and used the photometric parallax and metallicity relations from the $SDSS$ to obtain metallicities and distances for our sample of 1315 F and G type stars.  Limiting our study to $z$ = 7 kpc, the distance limit used by \\cite{Ivezicetal2008}, we find that our estimates for the mean metallicity of the thick disk ([Fe/H] = $-0.77$) and halo ($-$1.42 dex) are consistent with the results from the $SDSS$ Data Release 6 photometry ($-$0.80 dex and $-$1.46 dex). Since the $CFHTLS$ data is deeper than the $SDSS$, we are also able to check for trends in the thick disk and halo metallicity distribution beyond 7 kpc. Unlike the Ivezic et al. study, we do not find a trend in metallicity for the thick disk between 1 and 4 kpc. However, the significance of the trend observed in the $SDSS$ dataset comes from the stars within 1 kpc of the Galactic plane. Unfortunately, those stars are saturated in the $CFHTLS$ observations, and the metallicity trend for the thick disk stars further away than 1 kpc from the plane is consistent with the $SDSS$ results.  We observe a decline in metallicity of the halo with increasing distance from the Galactic plane. We suggest that this decline may be due to an increasing contribution from the very metal poor ([Fe/H]$<-2$) halo stars, but we cannot rule out systematic problems in our analysis. Proper motion data from the $CFHTLS$ observations will be available at the end of the survey period. These data will be useful to study the kinematic properties of our sample of stars and constrain the contribution of very metal poor stars."
    },
    "0809/0809.5090_arXiv.txt": {
        "abstract": "We analyze three-band imaging data of the giant elliptical galaxy \\esog\\ from the \\textit{Hubble Space Telescope} Advanced Camera for Surveys (ACS).  This is the nearest known strongly lensing galaxy, and it resides in the center of the poor cluster Abell S0740 at redshift $z{\\,=\\,}0.034$.  Based on magnitude, color, and size selection criteria, we identify a sample of 15 ultra-compact dwarf (UCD) galaxy candidates within the ACS field.  This is comparable to the numbers of UCDs found within similar regions in more nearby clusters (Virgo, Fornax, Hydra).  We estimate circular half-light radii $\\rc$ from 2-D \\sersic\\ and King model fits and apply an upper cutoff of 100~pc for our UCD selection.  The selected galaxies have typical \\sersic\\ indices $n{\\,\\approx\\,}1.5$, while larger sources with $\\rc{\\;>\\,}100$~pc are more nearly exponential, perhaps indicating that the latter are dominated by background disk galaxies. Many of the UCD candidates are surrounded by a faint ``fuzz'' of halo light, which may be the remnants of stripped material, and there is some evidence for intrinsic flattening of the UCDs themselves.  An apparent separation in size between the most compact UCDs with $\\rc<17$~pc and larger ones with $\\rc>40$~pc may hint at different formation mechanisms. We do not find any M32 analogues in this field. The colors of the UCD candidates span the range from blue to red globular clusters, although the brightest ones are predominantly red. The UCD candidates follow the flattened, elliptical distribution of the globular clusters, which in turn follow the galaxy halo light, suggesting a common evolution for these three components. Planned follow-up spectroscopy can determine which candidates are truly members of Abell S0740 and how similar they are in distribution to the globulars. ",
        "introduction": "A new class of stellar system has emerged in recent years.  Due to the size of these objects, being larger than average globular clusters (GCs) and smaller than dwarf galaxies, they have been dubbed ultra-compact dwarf galaxies, or UCDs (Phillips \\etal\\ 2001). They are typically a few $\\times 10^7$ \\msun\\ in mass, with effective radii in the range 10-100~pc. First discovered in the Fornax Cluster (Hilker et al.\\ 1999; Drinkwater et al.\\ 2000), UCDs have now been found in significant numbers in the Virgo, Centaurus, and Hydra clusters (\\hasegan\\ \\etal\\ 2005; Mieske et al 2007; Wehner \\& Harris 2007), all systems within $\\sim\\,$50~Mpc of the Local Group. They are apparently very rare outside of galaxy clusters (Evstigneeva \\etal~2007b).  As they have absorption line spectra and appear to be transitional between GCs and early-type dwarfs (cf. \\hasegan\\ \\etal\\ 2005), there are two basic ideas for the nature of UCDs: they are related to globulars, or to dwarf galaxies.  More specifically, UCDs may be the largest members of the rich GC populations found inside galaxy clusters (Mieske \\etal\\ 2002), possibly growing to such large size through dissipational merging early in their lifetimes (Fellhauer \\& Kroupa 2002).  Or, they may be the small, tidally stripped remains of nucleated dwarf galaxies on orbits that carried them too close to the center of the cluster potential (Bekki \\etal\\ 2001; Drinkwater \\etal\\ 2003).  This latter explanation has come to be known as galaxy ``threshing,'' but the idea has been around for many years. Bassino \\etal\\ (1994) numerically simulated the evolution of nucleated dwarfs in Virgo and showed that stripped nuclei could constitute a large fraction of M87's very rich GC system, while larger UCD-like remnants would occur farther out.  Recent \\textit{Hubble Space Telescope} (\\hst) imaging has revealed nuclei in a much higher percentage of Virgo early-type dwarfs than previously thought (\\cote\\ \\etal\\ 2006).  Thus, stripping of nucleated dwarfs may account for both UCDs and many of the GCs in the center of cluster potentials.   \\begin{figure*} \\epsscale{0.92} \\plotone{f1.eps} \\caption{\\textit{Hubble Space Telescope} ACS/WFC image of ESO\\,325-G004, showing about $3\\farcm0{\\,\\times\\,}3\\farcm3$ of the field at the observed orientation.  This color composite was constructed by the Hubble Heritage Team (STScI/AURA) from our imaging in the F475W $(g)$, F625W $(r)$, and F814W $(I)$ bandpasses.} \\label{fig:pic} \\end{figure*}   Some evidence based on the color-magnitude sequence of UCDs suggests that they may be an extension of the red GC component to brighter magnitudes (Wehner \\& Harris 2007).  The UCDs in the Virgo and Fornax clusters also have spectroscopic metallicities and $\\alpha$-element enhancements consistent with their being the high-mass mass extreme of the red GC population (Evstigneeva \\etal\\ 2007a; Mieske et al.\\ 2006), and less consistent with simple versions of the threshing model. Estigneeva \\etal\\ (2008) attempted to distinguish between the two formation scenarios on the basis of the structural properties of UCDs in the nearby Virgo and Fornax clusters measured using the \\hst\\ High Resolution Channel. Even with such high resolution measurements, the data were consistent with either explanation, although more detailed predictions of the size evolution of the nuclei during threshing are needed to test this scenario.  The relatively low velocity dispersions of cluster UCD populations are expected in either model (e.g., Bekki 2007).  However, detailed comparison between the spatial distributions of large samples of UCDs and their possible nucleated dwarf progenitors in clusters may help uncover their evolutionary histories (e.g., Goerdt \\etal\\ 2008; Thomas \\etal\\ 2008).  Given the difficulty in distinguishing between the formation scenarios, further UCD surveys can provide valuable information on the properties of this new type of stellar system.  A larger sample of groups and clusters is especially useful for constraining environmental effects on the formation of UCDs.  Here, we present a search with \\hst\\ for UCD candidates near \\esog, the central giant elliptical in Abell S0740.  This is one of the systems in the supplementary list of  poor clusters tabulated by Abell \\etal\\ (1989) that did not meet the lowest richness criteria of the original Abell (1958) catalogue. The cluster velocity dispersion is only $\\sim\\,$300 \\kms\\ (see plot in Smith \\etal\\ 2005), similar to that of Fornax, where UCDs were first discovered. The absolute $V$ magnitude of \\esog\\ is $M_V=-23.2$, making it about 60\\%~more luminous than M87, or 2.5 times the luminosity of NGC\\,1399 in Fornax.  At $z=0.034$, \\esog\\ is also the closest known gravitational lens and has both dynamical and lensing mass estimates (Smith \\etal\\ 2005). This makes it an interesting target for UCD searches, since it is a very massive, dominant elliptical in a poor cluster or rich group environment.  The following section describes our data in detail.  In \\S\\,\\ref{sec:selection}, we present our photometric and size measurements and discuss the selection of UCD candidates.  The properties of the UCD candidates are discussed in \\S\\,\\ref{sec:props} and compared with those of GCs and other objects in the field.  The final section summarizes the results. Throughout this paper, we use the WMAP 3-year cosmology results (Spergel et al. 2007) and assume a distance modulus for \\esog\\ of $(m{-}M) = 35.78$ mag, or a luminosity distance of 143 Mpc, and an angular scale of 0.65 kpc arcsec$^{-1}$.  This translates to an image scale of about 33 pc pix$^{-1}$ for our observations with Advanced Camera for Surveys Wide Field Channel (ACS/WFC). ",
        "conclusions": "We have presented an analysis of three-band ACS/WFC imaging to search for possible UCDs near the lensing galaxy \\esog\\ in Abell S0740.  This is an interesting target for a UCD search because it is a massive central elliptical in a poor cluster environment with a velocity dispersion similar to that of Fornax.  We selected objects based on their having magnitudes brighter than 99\\% of the expected GCs population, color appropriate for an early-type or population II system at this redshift, ellipticity less than 0.5, and circular half-light radii in the 10-100~pc range. The radii were measured using both the Galfit and Ishape programs. We found 15 good UCD candidates meeting the selection criteria, comparable to the expectations from previous searches.  In addition, we presented a sample of larger compact galaxies with radii in the range 100-300~pc, if they are located within the cluster.  These objects appear to be a mix of irregular background galaxies and larger versions of the cluster UCDs.  We did not find any counterparts of M32 in this field. The mean \\sersic\\ index for the UCD candidates is around 1.5, which is marginally higher than the value $\\sim\\,$1 found for the larger compact galaxies.  This may indicate that the latter objects are dominated by background disk-like galaxies, while the former group is mainly comprised of UCDs in the cluster.  Most of the UCD candidates and larger compact galaxies have visible surrounding halo light, consistent with galaxy threshing models. There is also evidence that most UCDs are intrinsically flattened, as none of the 15 UCD candidates has a fitted ellipticity $\\epsilon<0.16$.   The magnitude-size and color-magnitude diagrams show general continuity in the distributions of these parameters from GCs to the UCDs candidates. For our limited sample of UCD candidates, we find $\\rc\\sim L^{0.5}$.  This is an intriguing proportionality, as it implies a roughly constant surface brightness for UCDs of different sizes.  The better determined relation from Evstigneeva \\etal\\ (2008) is somewhat steeper, but consistent within the errors. There may be a bifurcation in the UCDs between those with sizes similar to GCs and larger ones with $\\rc>40$~pc, suggesting different origins for these two groups.  However, because of the small number of objects, the significance of the observed gap in \\rc\\ is only 83\\%. Therefore, although suggestive, it remains inconclusive. The colors of UCD candidates are weighted towards the red compared to the expected \\gIcolor\\ GC color distribution.  Several of the bright compact galaxies with sizes in the 100-300~pc range have colors near the expected peaks of the GC color distribution.  It would be useful to know if these objects are also in the cluster, and what may be their relation to the UCDs.    The majority of UCD candidates align along the major axis of \\esog, similar to the spatial distribution of the bright GCs.  Because of the small numbers involved, this result is not highly significant, but follow-up spectroscopy can provide a confirmed sample of UCDs; it will be interesting to see if these are mainly along the galaxy's major axis. These findings may appear to support a scenario in which the UCDs are the high-luminosity extension of the GC system.  However, as discussed in the Introduction, the true situation is probably more complex, and many red GCs may actually have their origin as stripped nucleated dwarfs, clouding the distinction between the main UCD formation scenarios. It would be useful to discover how the number of UCDs in complete surveys of many different clusters scales with the GC population of the central galaxy. We are currently completing a more detailed analysis of the GC population in this cluster and other similar fields from the same \\hst\\ program.  We also plan to obtain spectroscopy for all our UCD candidates to see what fraction of them are indeed associated with \\esog.  The additional information from these studies should provide further insight into the origin of UCDs and their connection to the GC populations."
    },
    "0809/0809.5059_arXiv.txt": {
        "abstract": "We present observations and analysis of nine dwarf irregular galaxies (dIs)  in the M81 Group taken with the Advanced Camera for Surveys aboard the Hubble Space Telescope.  The nine galaxy sample (the Garland, M81 Dwarf~A, DDO 53, \\hoix, \\hoi, DDO 165, NGC 2366, \\hoii, and IC 2574) spans 6 magnitudes in luminosity, a factor of 1000 in current star formation rate, and 0.5 dex in metallicity allowing us to study star formation (SF) under a variety of different conditions. Here we use color-magnitude diagrams of resolved stellar populations to study the star formation histories (SFHs) of these galaxies. Based on luminosity we divide the sample into faint and bright galaxies, with a dividing line of M$_{B}$ $=$ $-$15, and then analyze the similarities and differences in the SFHs, birthrate parameters, fraction of stars formed per time interval, and spatial distribution of stellar components.  Comparing these parameters as a function of luminosity, we find only minor differences in SF characteristics among the M81 Group dIs despite a wide range of physical properties.  We extend our comparison to select dIs in the Local Group (LG), with similar quality photometry,  and again find only minor differences in SF parameters.  The lack of a clear trend in SF parameters over a wide range of diverse environments suggests that SF in low mass systems may be dominated by stochastic processes.   The fraction of stars formed per time interval for an average M81 Group and LG dI is consistent with a constant SFH.  However, individual galaxies can show significant departures from a constant SFH. Thus, we find this result underlines the importance of stochastic SF in dIs. In addition to a statistical comparison, we compare possible formation scenarios of the less luminous and candidate tidal dwarfs in the M81 Group.  The SFHs and the lack of an overdensity of associated red stars suggest that the Garland and \\hoix\\ are not dIs and are potentially tidal dwarf galaxies.  Interestingly, a noteworthy difference between the LG and the M81 Group is the lack of tidal dwarf candidates in the LG. ",
        "introduction": "\\subsection{The M81 Group Dwarf Galaxy Project}  The ability to translate resolved stellar populations into star formation histories (SFHs) has become an increasingly powerful method for understanding how star formation (SF) shapes galaxy evolution \\citep[e.g.,][]{gza05}. The power of the Hubble Space Telescope (HST) is truly being realized as we are able to resolve the stellar components of increasingly distant galaxies. With the continuing expansion of the sample of galaxies with measured SFHs from resolved stars \\citep[e.g., the ACS Nearby Galaxy Survey Treasury (ANGST),][]{dal08}, we are able to study how a wide variety of galaxies evolve both individually and in their respective environments.  Of particular interest are star forming galaxies that sample the faint end of the galactic luminosity function.  Characteristics of SF in this regime are not as well studied as brighter, more massive galaxies, due to completeness limits of large galaxy surveys (e.g., SDSS). For example, \\citet{lee07} and \\citet{ken08} empirically find an increased dispersion in the \\halpha\\ equivalent width (EW) measurements, the normalized SFR over the past 5 Myr, of star forming galaxies fainter than $M_{B}$ $\\sim$ $-$15 compared to brighter galaxies.  The cause behind the increased dispersion is not clear, and we could simply be observing the galaxies at different points in their SF duty cycle.  Alternatively, different physical mechanisms could be responsible for the observed dispersion. There are several candidate theories that have distinct signatures: stochastic effects \\citep[cf.,][]{mue76, ger78}, internal feedback \\citep[stellar winds, supernovae, e.g.,][]{pel04, sti07}, external interactions \\citep[mergers, tidal influence, e.g.,][]{tom72, gne99}, or a combination thereof.  Direct observational evidence of recent SF episodes can shed light into the role of each process.  Specifically, dwarf irregular galaxies (dIs) are excellent environments for probing the mechanisms which regulate SF.  By number, dIs dominate the faint star forming galaxy count and span a wide range in mass, luminosity, metallicity, and SFR \\citep{mat98}.  They are solid body rotators \\citep{ski88}, so that features resulting from past SF episodes are preserved and not destroyed as they are in larger galaxies with differential rotation \\citep{ski96}.  Thus, SFHs based on resolved stellar populations of dIs directly probe the physical mechanisms that govern SF in low mass galaxies.  Additionally, dIs are useful for understanding galaxy interactions within a group environment.  Because SFHs trace the evolution of each galaxy over the history of the universe, we can measure the fraction of stellar mass formed at different epochs, which is directly comparable to models of galaxy interactions within groups \\citep[e.g.,][]{ric05, orb08}.  At more recent times, SFHs of dIs can be used to quantify patterns and duty cycles of SF over the past $\\sim$ 1 Gyr and compare them to theoretical predictions and amongst different groups of galaxies.  Comparing nearby galaxy groups, an interesting feature of the M81 Group is the presence of suspected tidal dwarfs \\citep[e.g.,][]{mak02, kar04}, which have no counterparts in the Local Group.  A three body interaction between larger galaxies in the M81 Group (M81, M82, and NGC 3077) occurred $\\sim$ 300 Myr ago, and there have been numerous observations of \\hi\\ tidal debris leftover from this event \\citep[e.g.,][]{van79,yun94}.  Young stellar systems have been observed in these tidal remnants, and we may be witnessing the formation of new dwarf galaxies (i.e., tidal dwarfs) in the nearby universe \\citep[e.g.,][]{wal99b,mak02}.  In this paper, we introduce the observations, color-magnitude diagrams (CMDs), and SFHs of nine dIs in the M81 Group based on HST/ACS imaging.  While our focus is on the recent temporal SFHs, we can also place loose constraints on the ancient SFHs and use red and blue star density maps to trace locations of past SF episodes.  Using the recent SFHs, we are able to quantify and compare the strength and duration of SF episodes amongst the sample and to similarly observed galaxies in the LG.  We use this paper to establish analysis techniques that we anticipate applying to a larger sample in the future.  Here, we present a limited comparison of the SFHs to observations in other wavelength regimes (e.g., \\halpha , ultraviolet (UV), infrared (IR), \\hi), and in our future work we will compare the spatially resolved recent SFHs to \\hi\\ observations and calibrate different SFR indicators to our CMD based SFHs.    \\subsection{The M81 Group Dwarf Galaxy Sample}  The M81 Group presents an excellent environment for the study of resolved stellar populations and SFHs.  The M81 Group is a prominent group of galaxies close enough to the Local Group \\citep[with distances typically $\\sim$ 3.8 Mpc,][]{kar02} that the HST/ACS easily resolves the bright, young stellar populations.  Further, the kinematic state of the M81 Group is very dynamic.  \\hi\\ observations and related models of the M81 Group \\citep{yun94, yun99} reveal a recent ($\\sim$ 300 Myr) interaction between the major member galaxies and the presence of tidal streams and debris.  A three dimensional map as well as detailed structural and kinematical information of the M81 Group are presented in \\cite{kar02}.  The sample of nine M81 Group dIs was selected from galaxies which had excellent ancillary observations, starting with galaxies observed as part of the Spitzer Infrared Nearby Galaxies Survey \\citep[SINGS;][]{ken03}. There are eight M81 Group dIs in both the SINGS and The \\hi\\ Nearby Galaxy Survey \\citep[THINGS;][]{wal08}. This list was augmented by adding NGC 2366 (which had similar quality ancillary observations, although not part of the SINGS survey), the Garland (for which HST/ACS observations already existed, GO-9381, PI: Walter), and UGC 4483, which is a dwarf starburst galaxy \\citep[cf.,][]{sea94, vss98}.  Two galaxies were subsequently dropped from the sample. UGC 4483 was observed on December 31, 2006, but the observation failed; this was shortly followed by the failure of the ACS, so we have no new observations for it and it was dropped from the sample.  M81 Dwarf~B was dropped from the sample after  we found it to be significantly more distant than the other galaxies in agreement with other recent analysis \\citep[e.g., 5.3 Mpc,][]{kar02}.  The final sample of nine M81 Group dIs spans a wide range of properties including 6 magnitudes in luminosity, a factor of 1000 in current star formation rate (SFR), and 0.5 dex in metallicity (see Table \\ref{tab1}).  In Figure \\ref{ha_ew} we have highlighted the dIs from this paper (M81 Group dIs in blue and LG dIs in black) in the context of the Local Volume star-forming galaxy sample \\citep{lee07, ken08}.   The galaxies here are shown in the M$_{B}$$-$\\halpha\\ equivalent width (EW) plane, which traces mass and current normalized SFR, respectively.  \\citet{lee07} find an empirical transition at M$_{B}$ $\\sim$ $-$15 based on the dispersions in \\halpha\\ EW. Our sample of dIs straddles both sides of this division, allowing us to probe differences in SF characteristics in both luminosity regimes over a wide range of current SFRs.  Within our sample, we have two candidate tidal dwarf galaxies, the Garland and \\hoix.  We will adopt that label throughout this paper, recognizing that the status of these features as true galaxies is controversial \\citep[e.g.,][]{vand98, mak02, kar02} and they may simply be stellar systems forming in the outskirts of their larger companions (see Figure \\ref{hi_map}).  We chose to exclude the candidate tidal dwarfs from our statistical analysis because of we find them to be fundamentally different than dIs in the sample. ",
        "conclusions": "From HST/ACS imaging we use resolved stellar populations to study the temporal and spatial properties of nine diverse dIs in the M81 Group.  Photometry from the HST/ACS images was used to create CMDs and trace the temporal evolution of each galaxy via measured SFHs.  From the CMD, we isolated red (old) and blue (young) stars and created stellar density maps to better understand the stellar spatial components.  Notably, we analyze two candidate tidal dwarf galaxies, the Garland and \\hoix, and find that they are clearly different from any of the other galaxies in the sample, both in their SFHs and spatial stellar distributions, despite having similar luminosities.  They are particularly interesting because they are two of the nearest candidate tidal dwarf galaxies in the universe, as we do not have similar galaxies in the Local Group.  In addition to analyzing the galaxies on an individual basis, we divided the sample into two groups, based on an empirical separation in luminosity of star forming galaxies in the Local Volume \\citep{lee07, ken08}.  Based on this division, we compute the fraction of stars formed and birthrate parameters, and compare these values between the two regimes, finding very little difference.  We extend our comparison of dIs to the LG, selecting LG galaxies with similar quality photometry for comparison.   Again, comparing the M81 Group dIs and LG dIs, we find very little difference in their SF characteristics (i.e., fraction of stars formed and birthrate parameters).  Analyzing the composite population of the M81 and LG dIs, we find that the average values of the fraction of stars formed in both luminosity regimes are consistent with that of a constant SFH.   Importantly, we note that individual galaxies are not well modeled by a constant SFH.  We interpret this result to mean that the fraction of stars formed in each galaxy is essentially random (or stochastic), but that with a larger sample, the mean values converge on those consistent with a constant SFH.  We searched for trends in the birthrate parameters as a function of luminosity, but found none, supporting the idea that stochastic SF is important in dIs.  With our future work, we will explore the same type of analysis with a larger sample, as well as comparing spatially resolved SFHs to multi wavelength observations to give us further insight into the impact of SF on the evolution of low mass galaxies."
    },
    "0809/0809.2023_arXiv.txt": {
        "abstract": "{} {We aim to characterize the properties of the prestellar and protostellar condensations to  understand the star formation processes at work in a young HII region} {We have obtained maps of the 1.25mm thermal dust emission and the molecular gas emission  over a region of $20\\arcmin \\times 10 \\arcmin$ around the Trifid Nebula (M20), with the IRAM 30m and the CSO telescopes as well as in the mid-infrared wavelength with ISO and SPITZER. Our survey is sensitive to features down to $\\rm N(\\htwo) \\sim 10^{22}\\cmmd$ in column density.} {The cloud material is distributed in fragmented dense gas filaments ($\\rm n(\\htwo)$ of a few $10^3\\cmmt$) with sizes ranging from  1 to $10\\pc$. A massive filament, WF, with properties typical of Infra Red Dark Clouds, connects M20 to the W28 supernova remnant. We find that these filaments pre-exist the formation of the Trifid and were originally self-gravitating. The fragments produced are very massive (typically 100\\msol\\ or more) and are the progenitors of the cometary globules observed at the border of the \\Hp\\ region. We could identify 33 cores, 16 of which are currently forming stars. Most of the starless cores have typical \\htwo\\ densities of a few $10^4\\cmmt$. They are usually gravitationally unbound and have low masses of a few $\\msol$.  The densest starless cores (several $10^5\\cmmt$) are located in condensation TC0, currently hit by the ionization front, and  may be the site for the next generation of stars. The physical gas and dust properties of the cometary globules  have been studied in detail and have been found very similar. They all are forming stars. Several intermediate-mass protostars have been detected in the cometary globules and in the deeply embedded cores. Evidence of clustering has been found in the shocked massive cores TC3-TC4-TC5.} {M20 is a good example of massive-star forming region in a turbulent, filamentary molecular cloud. Photoionization appears to play a minor role in the formation of the cores. The observed fragmentation is well explained by MHD-driven instabilities and is usually not related to M20. We propose that the nearby supernova remnant W28 could have triggered the formation of protostellar clusters in nearby dense cores of the Trifid.} ",
        "introduction": "Many observations suggest that H{\\small II} regions could play an important role in spreading star formation throughout the Galaxy. As reviewed in Elmegreen (2002), most of the embedded young stellar clusters in the solar neighbourhood are adjacent to H{\\small II} regions excited by more evolved stars~: Orion, Perseus OB2, the Rosette nebula, W3-5, M17, to name but a few. Yamaguchi et al. (1999) estimate that 10\\%-30\\% of stars in mass in the Galaxy are formed under the influence of adjacent H{\\small II} regions. In particular, the condensations encountered at the border of the HII regions, the so-called bright-rimmed Globules and cometary globules (CGs), have long been identified as stellar nurseries or ``incubators'' (Sugitani et al. 1991,1994; Lefloch et al. 1997; Dobashi et al. 2001). These condensations are local clumps which emerge from the expanding molecular layer surrounding the nebula. Several theoretical and observational studies  have shown the important role played by the Ly-continuum radiation of the nebula in the evolution of these condensations, which is now well understood (Bertoldi, 1989; Lefloch \\& Lazareff 1994, 1995). Systematic studies show that these ``stellar factories'' tend to form more massive and luminous (proto)stars, with luminosities $\\geq 10^3\\lsol$ (Sugitani et al. 1991, Dobashi 2001). It is estimated  that more than 15\\% of the massive stars form in CGs. The question of whether the mechanical and radiative luminosities of OB stars can trigger the formation of subsequent generations of stars, in particular of massive stars, in these condensations is still an open question.  The ``collect and collapse'' model, first proposed by Elmegreen \\& Lada (1977), is particularly interesting as it enables the formation of massive condensations and subsequently of massive objects out of an initially uniform medium. Systematic studies have shown that this mechanism was at work in some isolated Galactic \\Hp\\ regions and most likely responsible for the massive YSOs observed at the border of the nebulae (Deharveng et al. 2003,2006; Zavagno et al. 2007). In the case of a turbulent environment, the molecular gas is no longer homogeneous but distributed in clumpy and filamentary structures; it is not clear however what are the most efficient processes that account for the formation of the subsequent generations of stars. Several groups have investigated the impact of \\Hp\\ regions on the second-generation star formation, and proposed two main mechanisms of induced star formation, either direct triggering from \\Hp\\ regions and other pressure mechanisms associated with hot stars, or gravity-driven streaming of gas along the filaments to build dense cores (Nakamura, Hanawa \\& Nakano, 1993; Tomisaka, 1995; Fiege \\& Pudritz 2000; Fukuda \\& Hanawa, 2000).   We have started a comprehensive study of a young \\Hp\\ region, the Trifid Nebula (M20), small enough to allow a complete census of the young stellar and protostellar population, to address this question. The Trifid Nebula appears as a small dusty nebula in a young evolutionary age, at an heliocentric distance of $1.68\\kpc$ (Lynds et al. 1985). Mapping of the thermal dust and molecular gas emission of the nebula at millimeter wavelengths revealed the presence of several massive protostellar cores (dubbed TC1 to TC4), some of which are exposed to the ionizing radiation of the exciting star (Cernicharo et al. 1998, CL98). ISOCAM observations confirmed the extreme youth of the protostellar sources detected in the cores TC3 and TC4 (see also Lefloch \\& Cernicharo 2000, hereafter LC00). So far, only the cometary globule TC2 has been analyzed in detail  from its line and continuum emission at mid-infrared, millimeter and centimeter wavelengths (Lefloch et al. 2002, herafter L02); its protostellar content could not be characterized from the ISOCAM data, however.  The last phases of protostellar evolution in the nebula have been identified thanks to ISOCAM and SPITZER. Deep images obtained by SPITZER in the range $3.5-24\\mu m$ with IRAC and MIPS have revealed the population of low-mass objects which formed together with the exciting star of the nebula (Rho et al. 2006). In addition to 32 protostellar candidates, about 120 young stellar objects (Class II) have been detected, which are surrounded by warm dust, presumably distributed in protoplanetary disks. Analyzing the mid-infrared and centimeter emission of the central stellar cluster members, we found that their disks (components B, C, D) undergo copious photoionization (Lefloch et al. 2001, hereafter L01; Yusef-Zadeh et al. 2005).  In this article, we report on a large-scale survey of the cold dust emission around the Trifid Nebula, complemented with molecular line emission data. Our systematic mapping has allowed us to determine the evolutionary status of the SPITZER Class~0/I candidates; it reveals the presence of stellar incubators in the dense molecular material surrounding the \\Hp\\ region, and the progenitors of the protostellar cores. This allows us to draw a full evolutionary sequence for the incubators in the Trifid Nebula, from the early stage when they are embedded in low-density material of the parental cloud, as illustrated  by the massive pre-Orion cores TC3-TC4, to the late stage when the parental cocoon is fully dissipated, leaving the protoplanetary disks exposed to the ionizing radiation of the O star, as observed by L01 and Yusef-Zadeh et al (2005).  The article is organized as follows~: after presenting the observational data (Sect.~2), we summarize the main results of the continuum survey (Sect.~3). We then present in Sects.~4-5-6 detailed studies of typical cores which have formed in the Western filament(TC0), the Eastern lane (TC1) and in the South of the nebula (TC2 and TC9). We discuss in Sect.~7 the implications on the star formation history of the Trifid. The main conclusions are summarized in Sect.~8. ",
        "conclusions": "Our mapping of the Trifid Nebula at millimeter, submillimeter  and mid-infrared wavelengths has allowed us to characterize the gas and dust distribution. Thirty three cores have been identified, which cover all the stages of protostellar evolution, from the early prestellar to the late protoplanetary phases. Sixteen cores are found to display evidence of star formation.  The cores are found to lie at various stages of photoionization, from  the deeply embedded molecular phase to the cometary phase and the late stages of photoevaporation. Taking into account also the previously studied objects like TC2 (LC02), all cometary globules display very similar properties. They have typical sizes of 0.1 to $0.2\\pc$, \\htwo\\ densities ranging from  $\\approx 10^5\\cmmt$ to $10^6\\cmmt$, and masses of a few \\msol\\ up to a few tens $\\msol$. They appear to form at most one intermediate-mass object. Hence, we conclude that photoionization does not increase the star formation efficiency in the individual cores. We find evidence of clustering only in the shocked, dense condensations, such as TC3 and TC4 (previously studied in LC00) and TC5.  Comparison of the line and continuum emission properties shows hardly any evidence for molecular depletion inside the condensations. Only moderate molecular depletion has been reported towards the densest sources (TC3; LC00). The mid-infrared $5-17\\mu m$ flux detected with ISOCAM towards  TC1 and TC9 is consistent with the emission of a hot central source, with dimensions typical of protoplanetary disks, which is absorbed by the parental envelope of cold dust and gas. ISOCAM spectroscopy has revealed the presence of large amounts of $\\rm CO_2$ ices in cometary globule TC1, thereby unveiling the potentially rich composition of icy grain mantles in protostellar environments in the Trifid. As the ice composition reflects the various processes that take place on grains and the chemical evolution of the protostellar envelope, mid-infrared observations should allow us to investigate how  the chemical composition of the circumstellar envelopes varies along with protostellar evolution in the \\Hp\\ region.  The cloud material is distributed in filamentary structures of relatively high density (a few $10^3 \\cmmt$) and gas column density $\\rm N(H)$ of a few $10^{22}\\cmmd$. Their properties compare well with those of IRDCs and the Orion filaments reported by Maddalena et al. (1986). The length of these structures ranges from about $1\\pc$ to $15\\pc$ or more. Since some of them extend far beyond the ionization front of the \\Hp\\ region, these filamentary structures cannot have formed in a \"collect and collapse\" manner. They were pre-existing the birth of M20 and reflect the initial conditions of the parent molecular cloud.  Among the four filaments studied in greater detail by us, two of them are confined by self-gravity (WF and TC4f), whereas the other two are presently not. Taking into account the photoevaporation of the surface layers of the latter, they were probably self-gravitating too before the onset of photoionization. All these filaments are fragmented in condensations with typical masses of a few hundreds of $\\msol$. The fragmentation of the filaments can be accounted for by MHD-driven instabilities, although the exact importance of the magnetic field in the region remains to be quantified. Comparison with numerical simulations shows that the impact of M20 on the filaments is too recent to have triggered their fragmentation, with the exception of the TC4f filament. It may well be, however, that this fragmentation occurred at an earlier stage or was accelerated under the influence of high energy shock interactions with young supernova remnants. It is not clear if the mass distribution of the fragments in the filament WF results from the propagation of instabilities inside the filament, hence triggering the formation of the subsequent generations of cores. Conversely, the small-scale fragmentation reported in the filament TC4f provides a convincing example of a filament compressed and destabilized in the expansion of the \\Hp\\ region, well accounted for by the modelling of Elmegreen (1989). The mass of the individual fragments formed in the TC4f filament is rather low ($\\sim 10\\msol$), whereas the mass of the objects in WF is typically ten times higher.  The star formation going on in the dense cores TC3 and TC5 is characterized by very short dynamical timescales (about $10^4\\yr$), comparable to the age of the nearby W28 supernova remnant. TC5 is a good candidate to investigate the possibility of SNR shocks as a trigger of star formation. Therefore, the environmental conditions appear to have played a very important role in determining  the star forming conditions in the Trifid Nebula. The kinematics of the gas will be investigated in a forthcoming paper, in order to search for signatures of these environmental conditions and of possible interactions with the highly energetic phenomena which have recently taken place in the neighbourhood."
    },
    "0809/0809.2812_arXiv.txt": {
        "abstract": "{} {In our ongoing search for close and faint companions around T Tauri stars in the Chamaeleon star-forming region, we here present observations of a new common proper motion companion to the young T-Tauri star and Chamaeleon member CT Cha and discuss its properties in comparison to other young, low-mass objects and to synthetic model spectra from different origins.} {Common proper motion of the companion and CT Cha was confirmed by direct Ks-band imaging data taken with the VLT Adaptive Optics (AO) instrument NACO in February 2006 and March 2007, together with a Hipparcos binary for astrometric calibration. An additional J-band image was taken in March 2007 to obtain color information for a first classification of the companion. Moreover, AO integral field spectroscopy with SINFONI in J, and H+K bands was obtained to deduce physical parameters of the companion, such as temperature and extinction. Relative flux calibration of the bands was achieved using photometry from the NACO imaging data.} {We found a very faint (Ks\\,=\\,14.9 mag, Ks$_{0}$\\,=\\,14.4 mag) object, just $\\sim$\\,2.67\\arcsec~northwest of CT Cha corresponding to a projected separation of $\\sim$\\,440\\,AU at 165\\,$\\pm$\\,30\\,pc. We show that CT Cha A and this faint object form a common proper motion pair and that the companion is by $\\geq$\\,4\\,$\\sigma$ significance not a stationary background object. The near-infrared spectroscopy yields a temperature of 2600\\,$\\pm$\\,250\\,K for the companion and an optical extinction of $A_{\\rm V}$=\\,5.2\\,$\\pm$\\,0.8\\,mag, when compared to spectra calculated from Drift-Phoenix model atmospheres. We demonstrate the validity of the model fits by comparison to several other well-known young sub-stellar objects.} {We conclude that the CT Cha companion is a very low-mass member of Chamaeleon and very likely a physical companion to CT Cha, as the probability for a by chance alignment is $\\leq$\\,0.01. Due to a prominent Pa-$\\beta$ emission in the J-band, accretion is probably still ongoing onto the CT Cha companion. From temperature and luminosity ($\\log(L_{bol}/L_{\\odot})$=\\,--2.68\\,$\\pm$\\,0.21), we derive a radius of R=\\,2.20$_{-0.60}^{+0.81}$\\,R$_{\\mathrm{Jup}}$. We find a consistent mass of M=\\,17\\,$\\pm$\\,6\\,M$_{\\mathrm{Jup}}$ for the CT Cha companion from both its luminosity and temperature when placed on evolutionary tracks. Hence, the CT Cha companion is most likely a wide brown dwarf companion or possibly even a planetary mass object.} ",
        "introduction": "%   \\object{CT Cha} (aka HM\\,9), introduced in the 65th Name-List of Variable stars by \\citet{1981IBVS.1921....1K}, was originally found by \\citet{1973ApJ...180..115H} as an emission-line star in Chamaeleon, exhibiting variations in its H$\\alpha$\\,line from plate to plate and showing partial veiling \\citep{1980AJ.....85..444R}.  While the star was first classified as a T Tauri star by \\citet{1987MNRAS.224..497W}, it was later found to be a classical T Tauri star by \\citet{1990ApJS...74..575W} and \\citet{1992ApJ...385..217G} from IRAS data. \\citet{2000ApJ...534..838N} found evidence of a silicate feature disk ($L_{sil}=10^{-2}\\,L_{\\odot}$, $L_{sil}/L_{\\ast}=0.014$), using ISO data.  The variations in the H$\\alpha$\\,line were later interpreted as accretion signatures when \\citet{1998ApJ...495..385H} measured a mass accretion rate of log{\\,\\.{M}}$=-8.28\\,M_{\\odot}/a$. Additional variations in infrared \\citep{1979MNRAS.187..305G} and optical photometry can possibly be explained by surface features on CT Cha at a rotation period of 9.86 days, as found by \\citet{1998A&AS..128..561B}.  All additional properties of the K7 \\citep{1988A&AS...76..347G} star CT Cha, such as its age ranging from 0.9 Myr \\citep{2000ApJ...534..838N} to 3 Myr \\citep{1993ApJ...416..623F}, as well as its equivalent width of the lithium absorption line of W$_{\\lambda}$(Li)\\,=\\,0.40\\,$\\pm$\\,0.05\\,\\AA\\, \\citep{2007A&A...467.1147G}, its radial velocity of 15.1\\,$\\pm$\\,0.1 km/s \\citep{2006A&A...448..655J} and proper motion (Table\\,\\ref{table:2}), are consistent with a very young member of the Cha I star-forming region, having an age of 2\\,$\\pm$\\,2\\,Myr. ",
        "conclusions": "We derived a luminosity of $\\log(L_{bol}/L_{\\odot})$=\\,-2.68\\,$\\pm$\\,0.21 for the CT Cha companion from the extinction corrected values Ks$_{0}$= 14.37\\,$\\pm$\\,0.32\\,mag (2006), J$_{0}$= 15.25\\,$\\pm$\\,0.39\\,mag, and Ks$_{0}$= 14.31\\,$\\pm$\\,0.32\\,mag (2007) \\cite[from our A$_{V}$ and extinction law by][]{1985ApJ...288..618R} using a bolometric correction of B.C.$_{K}$\\,=\\,3.15\\,$\\pm$\\,0.15\\,mag from \\citet[][for spectral type M8--L0]{2004AJ....127.3516G} and adopting a distance of 165\\,$\\pm$\\,30\\,pc. From the luminosity and temperature, we calculated the radius to be R=\\,0.23$_{-0.06}^{+0.08}$\\,R$_{\\odot}$ or 2.20$_{-0.60}^{+0.81}$\\,R$_{\\mathrm{Jup}}$.  We considered several possibilities for deriving a final mass. Since our S/N is too low in the very blue region of the J-band, we did not use the gravity-sensitive sodium index from \\citet{2007ApJ...657..511A}. We likewise cannot use the depth of the alkali lines, because they are still in the process of being included in the {\\sc Drift-Phoenix} atmospheric models, which are the only models we found able to reproduce the J, H, and K continua of the CT Cha companion without any additional assumptions. As a result, we use evolutionary models to get a mass estimation of the object, stressing, however, that such models are uncertain up to at least $\\sim$\\,10 Myrs due to unknown initial conditions \\citep{2005astro.ph..9798C}. For an age of 2\\,$\\pm$\\,2\\,Myr, we find from luminosity and temperature masses of 8 -- 23 M$_{Jup}$, 13.5 -- 60 M$_{Jup}$ \\citep{2003A&A...402..701B} and 9 -- 23 M$_{Jup}$, 12.5 -- 50 M$_{Jup}$ \\citep{2000ApJ...542..464C} and $\\sim$\\,9 -- 22 M$_{Jup}$, $\\sim$\\,11 -- 60 M$_{Jup}$ \\citep{1997ApJ...491..856B}, respectively. The best-fitting models are from \\citet{1997ApJ...491..856B} and give for an age of 2\\,Myr a temperature of 2540\\,K, a surface gravity of $\\log{g}$= 3.9, and a luminosity of $\\log(L_{bol}/L_{\\odot})$= -2.64 a mass of 17 M$_{Jup}$. The possible mass intervals from luminosity and temperature of all evolutionary models overlap, so we conclude from the intersection of these intervals that our companion candidate has a mass of 17\\,$\\pm$\\,6\\,M$_{Jup}$.  To obtain a limit from direct observational evidence, we can compare the CT Cha companion with the eclipsing double-lined spectroscopic binary brown dwarf 2MASS J05352184\u2013-0546085 found by \\citet{2006Natur.440..311S} with known masses from first principles and of a comparable age to CT Cha. From this comparison we conclude that, due to much lower luminosity, lower temperature, and radius, the CT Cha companion must be lower in mass than all the components in 2MASS J05352184\u2013-0546085, i.e.\\,$\\leq$\\,35\\,M$_{\\mathrm{Jup}}$ \\citep{2007ApJ...664.1154S}.   From the comparison to objects in the Upper Scorpius association \\citep{2008MNRAS.383.1385L}, we found an optical spectral type of the CT Cha companion of M8 -- L0. We note that our object is very likely a member of the Cha I star-forming region, because it is co-moving with the primary, a member of Cha I, and because the equivalent width of the two combined potassium (K I) lines around 1.25\\,$\\mu$m of 9\\,$\\pm$\\,5\\,\\AA \\ is smaller, just as for USco J160714-232101, pointing to a lower age than $\\sim$\\,5\\,Myr since this object is a member of the Upper Scorpius association \\citep{2008MNRAS.383.1385L}. From our fits to the {\\sc Drift-Phoenix} models, we found a likely metallicity of [Me/H]=0.0. Although this value is not very well constrained, it is consistent with the median value $\\langle[Fe/H]\\rangle$ of -0.11 found for some members of Chamaeleon \\citep{2008A&A...480..889S}.  At a projected distance of $\\sim$\\,440 AU, the system is probably still below the long-term stability limit of $\\sim$\\,700 AU for a K7 primary star of 0.7\\,M$_{\\odot}$, following the argumentation of \\citet{1987ApJ...312..367W} and \\citet{2003ApJ...587..407C} \\citep[see also][for a discussion]{2005AN....326..701M}.  Brown dwarf companions of slightly higher mass have also been found as companions to much older stars and at a variety of projected separations. Examples are GJ 229 B \\citep{1995Natur.378..463N,1995Sci...270.1478O} at a projected separation of $\\sim$\\,45\\,AU and the binary brown dwarf companion Eps Indi B \\citep{2003A&A...398L..29S} at a projected separation of $\\sim$\\,1460\\,AU from their primaries. The currently most similar co-moving system to CT Cha A and its companion at a higher age is, however, HD 3651 B \\citep{2006MNRAS.373L..31M} and its host star found to have a projected separation of $\\sim$\\,480\\,AU.  The prominent Pa-$\\beta$ emission line in the J-band is most likely due to ongoing accretion onto the CT Cha companion, produced in the shock-heated magnetospheric accretion flow \\citep{1998AJ....116..455M,2004A&A...417..247W,2005MNRAS.358..671K} and thus another sign of its youth. While this magnetospheric accretion model predicts a correlation between emission line strengths of hydrogen lines and accretion rate \\citep{1998AJ....116..455M}, as used to determine accretion rates from the H-$\\alpha$ emission line, \\citet{2001A&A...365...90F} find no such agreement for the Pa-$\\beta$ emission line of approximately half of the 49 studied classical T-Tauri stars \\citep[see][for a discussion]{2004A&A...417..247W}. Interestingly, we did not find any emission in the Br-$\\gamma$ line around 2.166 \\AA . While a possible emission of Br-$\\gamma$ might be present below the detection limit, the same behavior, having only Pa-$\\beta$ emission, was also found for the companion of GQ Lup \\citep{2007A&A...463..309S}.   Very interesting is the high extinction value of $A_{V}$=\\,5.2\\,$\\pm$\\,0.8\\,mag of the CT Cha companion compared to the extinction value of CT Cha A, which we determined with the J-, H-, and K-band 2MASS photometry partly given in Sect.~2.3 and in comparison to the colors of main-sequence stars given in \\citet{1995ApJS..101..117K}. We find a best fit for $A_{V}\\sim$\\,1.3\\,mag and spectral type K7 consistent with the spectral type K7 by \\citet{1988A&AS...76..347G}. This is in good agreement with $A_{V}\\sim$\\,1.4\\,mag found by \\citet{1998A&A...338..977C} from I\\,-\\,J color excess, but inconsistent with the $A_{V}\\sim$\\,0.1\\,mag found by the same authors using star counts.  This behavior is very similar for CHXR 73 A and its companion CHXR 73 B, which has % similar properties as the CT Cha companion and thus both are probably similar in nature. We found an extinction of CHXR 73 B of $A_{V}$=\\,12.6\\,$\\pm$\\,2.1\\,mag from fitting the spectrum of CHXR 73 B to {\\sc Drift-Phoenix} models, afterwards matching in the highest signal-to-noise areas the spectra of the CT Cha companion. We find the same best temperature 2600\\,K fitting both objects, while the luminosities are consistent within their 1\\,$\\sigma$ error bars. Even the difference of primary to secondary optical extinction is similar, being 3.1 -- 4.7\\,mag for CT Cha and 1.7 -- 5.9\\,mag for CHXR 73, derived from comparison of our CHXR 73 B value to the value of CHXR 73 A $A_{J}$=\\,1.8\\,mag \\citep{2004ApJ...602..816L} and using the extinction of KPNO-Tau 4 $A_{V}$=\\,2.45\\,mag \\citep{2007A&A...465..855G} (see Sect.~3.5). However, we can still not exclude from our spectral comparison that the deviant slopes in the blue part of the H-band and the red part of the K-band (see Fig.~\\ref{spectrum_CHXR73}), as well as the deep potassium lines at $\\sim$\\,1.25\\,$\\mu$m of CHXR 73 B, are indeed real and connected to a higher age and surface gravity of CHXR 73 B.  This difference in the extinction of the objects can have any of the following three causes. (1) The extinction in the young Chamaeleon I star-forming region is very clumpy and changes drastically on short spatial scales. (2) Although a close to pole-on  disk around a sub-stellar object would make no significant contribution to the spectrum in the near-infrared, an object bearing a close to edge-on disk could be reddened in this spectral region due to e.g. reflection and absorption by the disk. (3) The companions of CHXR 73 and CT Cha are both a projection effect and are actually situated further to the back of the Cha I cloud, still co-moving as they are members of the star-forming region, sharing the same space motion in a moving group, but not being physically bound to their primary objects. This can hardly be reconciled, however, with the probabilities of finding a by chance alignment estimated to be $\\sim$\\,0.001 for CHXR 73 B by \\citet{2006ApJ...649..894L} and $\\sim$\\,0.01 for the CT Cha companion by us (see Sect.~2.2).  We conclude from our mass estimate of 17\\,$\\pm$\\,6\\,M$_{Jup}$ that the CT Cha companion is most likely a brown dwarf, we can, however, not exclude the possibility that the CT Cha companion is a planetary mass object. There is yet no consensus or definition for \\textit{planets} around other stars, or for the upper mass limit of planets. In general we can follow the discussion in \\citet{2006ApJ...649..894L} saying that CHXR 73 B was probably not formed in a planetary fashion as its separation from the primary is very high, as is the case for the CT Cha companion, and cloud fragmentation was shown to be able to produce even isolated objects of similar masses. However we would like to stress that this argumentation assumes that the companions actually formed at the current separation to their primaries. We cannot yet rule out that the objects were scattered outward, as supposed for the sub-stellar companion of GQ Lup by \\citet{2006ApJ...637L.137B} and \\citet{2006A&A...451..351D} and supported by planet-planet interaction models as e.g. from \\citet{2008ApJ...678..498N}, where planets can be scattered to highly excentric orbits and large semi-major axis by interaction.  New evolutionary models by \\citet{2007ApJ...655..541M} show that objects created by core accretion have lower luminosities at very young ages than expected before. As CT Cha A and its companion have an estimated age of 2\\,$\\pm$\\,2\\,Myr, it is thus unlikely that the companion was formed by core accretion."
    },
    "0809/0809.3871_arXiv.txt": {
        "abstract": "The magnetorotational instability (MRI) is believed to be an efficient way to transport angular momentum in accretion discs. It has also been suggested as a way to amplify magnetic fields in discs, the instability acting as a nonlinear dynamo. Recent numerical work has shown that a large-scale magnetic field, which is predominantly azimuthal, axisymmetric and has zero net flux, can be sustained by motions driven by the MRI of this same field.  Following this idea, we present an analytical calculation of the MRI in the presence of an azimuthal field with a non-trivial vertical structure. In the limit of small vertical wavelengths, we show that magnetorotational shearing waves have the form of vertically localized wavepackets that follow the classical MRI dispersion relation to a first approximation. We determine analytically the spatiotemporal evolution of these wavepackets and calculate the associated mean electromotive force (EMF), which results from the correlation of the velocity and magnetic field perturbations. The vertical structure of the azimuthal field results in a radial EMF that tends to reduce the magnetic energy, acting like a turbulent resistivity by mixing the non-uniform azimuthal field.  Meanwhile, the azimuthal EMF generates a radial field that, in combination with the Keplerian shear, tends to amplify the azimuthal field and can therefore assist in the dynamo process.  This effect, however, is reversed for sufficiently strong azimuthal fields, naturally leading to a saturation of the dynamo and possibly to a cyclic behaviour of the magnetic field. We compare these findings with numerical solutions of the linearized equations in various approximations, and show them to be compatible with recent nonlinear simulations of an MRI dynamo. ",
        "introduction": "The magnetorotational instability (MRI) is believed to be responsible for turbulent motion and angular momentum transport in accretion discs, at least those that are sufficiently ionized \\citep{BH91a,BH98}.  In its simplest version it appears as a linear instability of a rotating shear flow in the presence of a uniform magnetic field \\citep{V59}.  The nonlinear outcome, at sufficiently large Reynolds and magnetic Reynolds numbers, is magnetohydrodynamic (MHD) turbulence \\citep{HGB95}.  More importantly, the MRI is also believed to act as a dynamo in accretion discs, meaning that the inductive effect of the motions driven by the instability is able to sustain the magnetic field against resistive decay and so to perpetuate the turbulent motion even in the absence of an externally imposed magnetic field \\citep{HGB96}.  The dynamo is fully nonlinear because the Lorentz force is essential to the operation of the MRI.  The MRI dynamo can be studied in the well known shearing-box model \\citep{HGB95}, which is a local representation of a differentially rotating disc, with boundary conditions that conserve the box-averaged magnetic field. When the initial conditions include a magnetic field with zero volume average, the MRI can be initiated but it has the opportunity to annihilate the magnetic field and thereby suppress the turbulence.  In order for the magnetic field and the turbulence to be sustained indefinitely, dynamo action must occur.  Recent numerical simulations of the MRI with zero box-averaged magnetic field have shown that the outcome depends in an important way on the numerical method and resolution unless explicit diffusion coefficients (viscosity and resistivity) are included and the relevant dissipative scales are resolved in the calculation \\citep{FP07,FPLH07}.  These results cast doubt on the operation and efficiency of the nonlinear dynamo in accretion discs, especially in the regime of small magnetic Prandtl number (Pm).  It is possible that the shearing box is too restrictive a model to describe the behaviour of accretion discs, as it conserves the box-averaged field.  Different vertical boundary conditions that allow horizontal flux to escape may assist in the operation of the MRI dynamo \\citep{BNST95}.  However, in the light of the results described below, we think this model is a good way to isolate the dynamo process from other effects, and can be seen as a minimal setup in which to study this process.  Steady solutions of the MHD equations representing a self-sustaining MRI dynamo process have been obtained by numerical continuation methods \\citep{ROC07} in a rotating shear flow between conducting walls.  The magnetic field has zero net flux and is predominantly toroidal (azimuthal).  In a condition of marginal stability, the MRI generates steady nonaxisymmetric motion.  This induces a poloidal (radial and vertical) magnetic field which, through the action of Keplerian shear, sustains the toroidal field.  These solutions reveal the workings of the dynamo but so far they have been found only for large Pm, perhaps because the method is restricted to steady states.  We have recently identified and analysed a cyclical behaviour of the large-scale magnetic field in an MRI simulation in a shearing box with zero net flux \\citep{LO08a}.  The processes involved in the cycle are related to those occurring in the wall-bounded steady solutions of \\citet{ROC07} and are probably relevant to the operation of the nonlinear dynamo in accretion discs.  However, there is an important role for shearing waves, which are the typical nonaxisymmetric solutions obtained in the shearing box without any walls \\citep{GL65}. The nonaxisymmetric instabilities are transient in nature and difficult to analyse, yet they appear to be responsible for sustaining the magnetic field and limiting its growth.  A common feature of these calculations is that the magnetic field is dominated by a large-scale, axisymmetric, azimuthal component.  We can think of this as a `mean' field if we consider a process of azimuthal averaging.  At a simple level the dynamo operates because the MRI generates radial field from azimuthal, while shear creates azimuthal field from radial.  However, the first process requires detailed understanding.  In linear theory, the MRI of an axisymmetric azimuthal field gives rise to a nonaxisymmetric radial magnetic field perturbation that averages to zero because of its wavelike form. Since we are looking for an axisymmetric effect, we are interested in the average second-order effect resulting from the correlations of the velocity and magnetic field perturbations. This is described by the mean electromotive force (EMF) of the MRI solutions.  We interpreted the results of the simulations in \\citet{LO08a} by carrying out linear numerical calculations of the evolution of shearing waves in the presence of a large-scale azimuthal field with a sinusoidal dependence on the vertical coordinate, $B_x\\propto\\cos kz$. The waves undergo transient growth as a result of the MRI.  We analysed the mean EMF associated with the waves and its feedback on the large-scale field.  The radial EMF was found to have a resistive effect, reducing the energy of the azimuthal field.  The azimuthal EMF $\\mathcal{E}_x$ was found to generate a radial field which, in combination with the Keplerian shear, also affects $B_x$.  When $B_x$ is relatively weak, $\\mathcal{E}_x$ results in an amplification of the field; when it is stronger, the sign of $\\mathcal{E}_x$ is reversed and $B_x$ is reduced by the process.  This leads to a cyclical behaviour in which the azimuthal field strength oscillates irregularly about a characteristic value.  The purpose of this paper is to make an analytical and more general investigation into the behaviour of shearing waves in the presence of a vertically non-uniform azimuthal magnetic field.  We aim to explain the systematic behaviour of the EMF, which is an essential part of the dynamo process; it is not our objective here to give a complete explanation of the dynamo. In order to make analytical progress we examine an asymptotic regime in which a separation occurs between the vertical length-scales of the large-scale field and of the shearing wave.  We develop a general theory of vertically localized waves in Section~2 and derive expressions for the associated EMF.  In Section~3 we compare these with the results of numerical calculations of the linearized equations.  We summarize our findings and draw conclusions in Section~4. ",
        "conclusions": "In this paper, we have investigated the magnetorotational instability in the presence of an azimuthal magnetic field with a non-trivial vertical structure. We have shown first that, in the limit of perturbations of small vertical wavelength, the instability takes the form of wavepackets that are vertically localized near the local maxima of the MRI growth rate according to the standard dispersion relation.  We have determined analytically the spatiotemporal evolution of these wavepackets.  Computing the electromotive force that arises from the correlation of the velocity and magnetic field perturbations, we have investigated what would be the feedback of the MRI wavepackets on the large-scale field. We have found that the dominant process is a dissipative effect on the mean azimuthal field $B_x(z)$, which can be understood to a first approximation as a turbulent resistivity resulting from the mixing of the non-uniform azimuthal field. We have demonstrated that the MRI shearing waves can also have a feedback on the mean radial field $B_y(z)$, leading to the possibility of dynamo action when combined with the Keplerian shear. We have compared our analytical results with linear numerical calculations, showing that our analysis captures most of the physics involved in these MRI shearing waves. Finally, we have found that these results are consistent with the closure model presented by \\cite{LO08a} for nonlinear dynamo cycles in accretion discs.  We emphasize that the findings presented here are only linear results that have many limitations. Hence, when some EMFs are referred to as providing a `turbulent resistivity', one should understand that some energy is transferred from the large-scale field to the perturbation, but the total energy is conserved. When the system becomes nonlinear, these perturbations couple with themselves and the energy will eventually cascade towards small scales to be dissipated by molecular processes. Therefore, the mechanism by which the energy is actually dissipated is intrinsically nonlinear, and is not covered by our work. Moreover, the amplitude of the quasilinear effects cannot be directly computed from our analysis, first because this amplitude will depend on the initial excitation of the waves, which is related to the strength of the underlying turbulent motions and is therefore unknown, and secondly because turbulent motions will modify the linear shearing waves as they grow and eventually destroy them. Since this destruction process will depend on the intensity of the underlying turbulence, the saturation level of the shearing waves is unpredictable, and may even be time-dependant. Therefore, our analysis can help us to understand the main effects due to the presence of the MRI in a turbulent flow, but it cannot predict the saturation values, which must be computed from nonlinear simulations. As an example, we cannot directly conclude from (\\ref{IntegralJ3}) that the dynamo effect should become stronger as we increase $k_x$, because in a turbulent flow the nonlinear coupling will also become more efficient as we increase $k$.  Despite these limitations, one can still deduce several important features from this linear analysis, the most important one being probably the dynamo feedback carried by $\\mathcal{E}_x$. This feedback allows us to close the loop given by the azimuthally averaged equations (\\ref{BxEvol})--(\\ref{ByEvol}), following a kind of $\\alpha\\Omega$ dynamo concept, although we do not consider the azimuthal EMF to arise from an $\\alpha$~effect. Note however that since this feedback creates a large-scale radial field $B_y(z)$, this large-scale field should also be included in the linear analysis. We have chosen to neglect $B_y$ compared to $B_x$, assuming that the dynamo feedback is smaller than the effect of the shear. This assumption is compatible with MRI turbulence in zero-net-flux simulations in which $\\sqrt{\\langle B_x^2 \\rangle }\\gg \\sqrt{\\langle B_y^2 \\rangle}$ \\citep{SHGB96}. This feedback is therefore a natural explanation to the persistence of a magnetic field in zero-net-flux MRI simulations, but does not constitute on its own a self-sustaining mechanism, since one still needs an initial perturbation to excite non-axisymmetric waves. One may also note that our results are found in the ideal MHD limit, and therefore do not provide any information about the dependence of the feedback on the magnetic Prandtl number. To investigate this effect, we have tried introducing a small amount of resistivity and viscosity ($<10^{-3}$) in our numerical calculations. It appears that the main results are not drastically modified by these effects, and no strong magnetic Prandtl number dependence is detectable from our linear analysis. This result may imply that the dynamo process itself does not depend strongly on the magnetic Prandtl number, although the way turbulence is maintained in the system does. This conclusion is also compatible with simulations with a non-zero mean vertical flux \\citep{LL07}, in which no dynamo process is required to sustain the magnetic field but a strong magnetic Prandtl number dependence is still found.  It has already be pointed out by \\cite{BNST95} that MRI turbulence may act as an accretion disc dynamo. Using shearing box simulations with specific vertical boundary conditions, they have shown that large-scale azimuthal magnetic fields may be produced by MRI turbulence over long time scales. This idea has also been explored by \\cite{ROC07} with a stationary nonlinear solution involving the MRI in a rotating plane Couette flow. More recently, we have found a dynamo cycle in MRI simulation \\citep{LO08a} which motivated this linear analysis. Naturally, it would be tempting to say that all these numerical results are, at least partially, explained by our results. However, to confirm this conjecture, one has to reproduce all these numerical results, and check how the EMF profiles are correlated to the large-scale fields. If our result proves to be more general, it would be the first step toward a general closure model for global disc simulations, involving a self-consistent mechanism to generate magnetic fields."
    },
    "0809/0809.3089_arXiv.txt": {
        "abstract": "We present a polarimetric map of a $20\\arcmin \\times 20\\arcmin$ area toward the Galactic center. The polarization of point sources has been measured in the $J$, $H$, and $K_S$ bands using the near-infrared polarimetric camera SIRPOL on the 1.4 m telescope IRSF. One percent or better accuracy of polarization degree is achieved for sources with $J<14.5$, $H<13.5$, and $K_S<12.0$. Comparing the Stokes parameters between high extinction stars and relatively low extinction ones, we have obtained a polarization originating from magnetically aligned dust grains at the central region of our Galaxy of at most 1$-$2 kpc. The distribution of the position angles shows a peak at $\\sim$20\\degr, nearly parallel to the Galactic plane, suggesting a toroidal magnetic configuration. The derived direction of the magnetic field is in good agreement with that obtained from far-infrared/submillimeter observations, which detect polarized thermal emission from dust in the molecular clouds at the Galactic center. Our results show that by subtracting foreground components, near-infrared polarimetry allows investigation of the magnetic field structure {\\it at } the Galactic center. ",
        "introduction": "\\label{sec:Intro}   The discovery of polarized radio emission extending perpendicular to the Galactic plane, such as nonthermal radio filaments \\citep{Yusef84} and polarized plumes \\citep{Tsuboi86}, has provided early evidence for a substantial poloidal component of the magnetic field in the central region of our Galaxy. High resolution observations at radio wavelengths have revealed more filaments of similar orientations \\citep[e.g.,][]{LaRosa04,Yusef04}, where the magnetic fields were found to be predominantly aligned along the filaments. The accumulation of these observational results has led to the hypothesis that most of the volume of the Galactic center (GC) is permeated by a {\\it poloidal} magnetic field.   Evidence for the existence of a {\\it toroidal} magnetic structure comes from far-infrared (FIR) and submillimeter (sub-mm) observations, which detect polarized thermal emission from magnetically aligned dust grains. The rotation axis of the dust grain aligns with the magnetic field, and thus the polarization of the thermal emission indicates the magnetic field direction (the measured direction of the E-vector is orthogonal to the magnetic field). The magnetic fields inferred from these observations \\citep[e.g.,][]{Werner88,Morris92,Hildebrand93} run parallel to the Galactic plane. Recent linear polarization observations of the 450-$\\mu$m continuum suggest a large-scale toroidal magnetic field extending over a region of $170 \\times 30 $ pc \\citep{Novak03}.  \\citet{Chuss03} found an interesting dependence of the magnetic field direction on the sub-mm flux. In low density regions, the field aligns generally perpendicular to the Galactic plane, while in high density regions, the field has a toroidal geometry. One explanation for this is that the global magnetic field in the early Galaxy was initially in a poloidal configuration, however the gravitational energy density in dense molecular clouds was strong enough to distort the poloidal field into a toroidal one, though it was insufficient in the low density regions.   A model which can connect the poloidal and toroidal magnetic fields was proposed by \\citet{Uchida85}, and was extended to a more realistic case by \\citet{Shibata86}. This magnetohydrodynamic model was developed to explain the GC lobes found by \\citet{Sofue84}. Since magnetic flux is generally frozen into matter, differential rotation and infall can shear an initially poloidal field into a toroidal one. Consequently, a toroidal field is developed close to the Galactic plane, while the field is vertical at high Galactic latitude.   The \\citet{Uchida85} model predicts that the toroidal component will generally be more dominant closer to the Galactic plane. A map of the direction of the magnetic field could therefore be a simple indicator of how well this model works in the GC. However, dust emissivity is high in dense, warm clouds, and thus observations of dust emission in FIR/sub-mm wavelengths are strongly limited to such regions, which show patchy distributions in the GC region.   In this paper, we present near-infrared (NIR) polarimetry of point sources toward the GC covering a much larger region of the sky than previous observations of this region. We demonstrate that NIR polarization can provide information on the magnetic field structure not only in the Galactic disk, but also in the central region of our Galaxy. ",
        "conclusions": "\\label{sec:Disc}  \\subsection{Previous Infrared Polarimetry toward the GC} \\label{subsec:PrevNIR}   NIR polarimetry for diffuse emission and point sources toward the GC has been conducted since the 1970's \\citep{Maihara73}. From the polarization angle aligned nearly along the Galactic plane, and the wavelength dependence of polarization well fitted by a power law \\citep{Nagata94}, it has been interpreted that the polarization is of interstellar origin dominated by dust that lies in the Galactic disk \\citep[e.g.,][]{Hough78,Kob80,Lebofsky82}.  One of the deepest NIR polarimetric observations toward the GC was carried out by \\citet{Eckart95} for a small field of $\\sim 13\\arcsec \\times 13\\arcsec$. They established significant polarization for 160 sources fainter than 13 mag in the $K$ band. The mean flux-weighted polarization is 4 \\% at 25\\degr, nearly parallel to the Galactic plane. A similar result, $4.1 \\pm 0.6 \\%$ at $30 \\pm 10\\degr$, was also obtained by \\citet{Ott99}. \\citet{Eckart95} concluded that most of the polarizations are caused via absorption by aligned dust grains in the Galactic plane.   A change of the magnetic field configuration along the line of sight toward the GC has been pointed out by \\citet{Kob83}. The diagram of position angles in the $K$ band versus $H-K$ for 15 discrete sources \\citep[Fig. 5,][]{Kob83} showed that the position angles are smaller and less ordered for sources with $H-K<1.0$, while those with $H-K>1.0$ are confined to a relatively narrow range around $20\\degr$. They concluded that this may indicate a change of the magnetic field direction at a distance corresponding to $H-K=1.0$ (5 kpc or more according to their calculation).  We can also identify a change of the position angle in the $K_S$ bands at $H-K_S \\sim 1.0$ (see Fig. \\ref{fig:DistHKPK}). As shown in the color-color diagram in Fig. \\ref{fig:Col2}, most of the stars with $H-K_S \\ga 1.0$ have the color of giants (i.e., their unreddened positions are on the locus of giants), and the strong peak in the $H-K_S$ histogram (the top panel in Fig. \\ref{fig:Col2}) suggests that they are distributed in the Galactic bulge. Hence, this change of the position angle indicates a transition of the magnetic field configuration along the line of sight, between the Galactic disk and bulge.   \\subsection{``Red'' and ``Blue'' Stars and Their Separation in the Line of Sight}  To discriminate between foreground disk stars and those in the Galactic bulge, the criterion $H-K_S=1.0$ is applied in our analysis because most stars with $H-K_S>1.0$ are attributed as giants in the Galactic bulge. The histogram of $H-K_S$ and the color-color diagram of point sources with photometric errors of less than 0.1 mag in the three bands are shown in Fig. \\ref{fig:Col2}. Nearby stars can be found around the giants' and dwarfs' loci (thick and thin curves) without reddening, while most of the redder stars, which are located at $H-K_S \\ga 1.0$ and $J-H \\ga 2.5$, have a color expected for reddened giants. Toward the GC, the number of stars in the bulge is larger than in the exponential disk by a factor of $\\sim 50$ \\citep{Wainscoat92}, and thus most of the stars at $H-K_S \\ga 1.0$ and $J-H \\ga 2.5$ are giants in the Galactic bulge.   The bottom panel of Fig. \\ref{fig:CMDRedBlue} shows the location of the red and blue stars in the $K_S$ vs. $H - K_S$ color-magnitude diagram. The theoretical Padova isochrones \\citep{Girar02} for solar metallicity with an age of 10 Gyr are also plotted in the color-magnitude diagram. The isochrones are put at the distance of the GC \\citep[7.5 kpc;][]{Nishi06b} and shifted along the reddening vector \\citep[$A_{K_S}/E_{H-K_S} = 1.44$;][]{Nishi06a} with $K_S$ extinctions of 0, 1, 2, 3, 4, and 5 mag. The distances to the red and blue stars are difficult to estimate accurately, since we do not know the distribution of the interstellar dust along the line of sight; however, most of the giants in the Galactic bulge can be detected for at least $K_S < 15$, and for these stars, we can find a clear peak in the $H-K_S$ histograms. Fig. \\ref{fig:CMDAll} clearly shows that our observations are essentially equally sensitive to stars on the near ($H-K_S \\la 2.0$) and far ($H-K_S \\ga 2.0$) sides of the Galactic bulge for $H \\leq 15.0$. The clear peaks found in the $H-K_S$ histograms are thus due to the spatial distribution of stars, not due to a distance effect. This indicates that ${(H-K_S)}_{\\mathrm{peak}}$ roughly corresponds to a real peak of the spatial distribution of giants along the line of sight. The color-magnitude diagram of Fig. \\ref{fig:CMDRedBlue}  also tells us that a large number of blue stars are distributed around the distance where $1.0 \\la A_{K_S} \\la 2.0$, and most of the red stars are further than the distance which suffers from the extinction of $A_{K_S} = 2.0$.    From the distribution of $H-K_S$ color, we estimate the depth of the region where we have mapped the magnetic configuration shown in Fig. \\ref{fig:PABGChuss}. For each sub-field, we calculated $H-K_S$ color differences between the peak in the $H-K_S$ histogram and the mean colors of blue and red stars, i.e., $(H-K_S)_{\\mathrm{peak}} - \\langle (H-K_S)_{\\mathrm{blue}} \\rangle$ and $\\langle (H-K_S)_{\\mathrm{red}} \\rangle -(H-K_S)_{\\mathrm{peak}}$. These color differences and corresponding extinction $A_{{K_S}_{\\mathrm{blue}}} = 1.44 \\times [(H-K_S)_{\\mathrm{peak}} - \\langle (H-K_S)_{\\mathrm{blue}} \\rangle]$ and $A_{{K_S}_{\\mathrm{red}}} = 1.44 \\times [\\langle (H-K_S)_{\\mathrm{red}} \\rangle - (H-K_S)_{\\mathrm{peak}}]$, where 1.44 comes from $A_{K_S}/E_{H-K_S} = 1.44$ \\citep{Nishi06a}, show the amount of dust extinction in the region where we have obtained polarimetric information. First, the amount of extinction toward the peaks in the $H-K_S$ histograms can be calculated to be $A_{{K_S}_{\\mathrm{peak}}} = 1.44 \\times [(H-K_S)_{\\mathrm{peak}}-(H-K_S)_0]$, where the mean intrinsic color $(H-K_S)_0$ of red giants is assumed to be $\\sim 0.2$ (see Fig. \\ref{fig:CMDRedBlue}). This corresponds to the amount of extinction up to the GC in each sub-field. Next, the extinctions $A_{{K_S}_{\\mathrm{blue}}}$ and $A_{{K_S}_{\\mathrm{red}}}$ are converted to the actual distances from the GC, using a simple model. To make conservative estimates, we use the model by \\citet{Davies97}, who derived a rather extended distribution of dust. \\citet{Davies97} showed that the extinction to the GC associated with cool diffuse interstellar dust can be calculated using \\begin{eqnarray} A = C \\times \\int_0^{R_0} e^{-r/\\alpha_{\\mathrm{d}}} dr, \\label{eq:Extinction} \\end{eqnarray} where $\\alpha_{\\mathrm{d}} \\approx 5.3$ kpc is a scale-length of the dust distribution in the radial direction, $C$ is a constant depending on the wavelength and dust density, and $R_0$ is the distance between the GC and the Sun \\citep[$R_0 = 7.5$ kpc;][]{Nishi06b}. The distances from the GC $x_{\\mathrm{blue}}$ and $x_{\\mathrm{red}}$ corresponding  to $A_{{K_S}_{\\mathrm{blue}}}$ and $A_{{K_S}_{\\mathrm{red}}}$, respectively, can be estimated with the equations \\[ \\Bigl( \\int_0^{x_{\\mathrm{blue,red}}} e^{-r/\\alpha_{\\mathrm{d}}} dr \\Bigr) \\Bigg/ \\Bigl( \\int_0^{R_0} e^{-r/\\alpha_{\\mathrm{d}}} dr \\Bigr) = A_{{K_S}_{\\mathrm{blue,red}}} \\Big/ A_{{K_S}_{\\mathrm{peak}}}.\\] We obtained the average and standard deviation of $x_{\\mathrm{blue}}$ as 0.5 and 0.1 kpc, and those of $x_{\\mathrm{red}}$ as 1.0 and 0.5 kpc, respectively. We found a long-side tail in the $x_{\\mathrm{red}}$ histogram, which enlarges the average and the standard deviation. This estimation suggests that the polarization shown in Fig. \\ref{fig:PABGChuss} occurs between the average distances of $(R_0 - 0.5)$ kpc and $(R_0 + 1.0)$ kpc from the Sun, arising probably from the central 1$-$2 kpc region of our Galaxy.   In reality, the central part of the Galaxy harbors a strong concentration of gas and dust called the ``Central Molecular Zone'' \\citep[$R \\la 200$ pc;][]{Morris96}, which approximately corresponds to the concentration of stars called the ``Nuclear Bulge'' \\citep[$R \\la 300$ pc;][]{Mezger96,Serabyn96}. According to \\citet{Launhardt02}, stars belonging to the Nuclear Bulge dominate in the central part of the Galaxy (see their Fig. 14). Therefore, a large portion of the stars we have detected is located in the Nuclear Bulge, and the polarization at the GC originates mostly within the central few hundred pc.    \\subsection{Magnetic Field Configuration at the GC}   As shown in Fig. \\ref{fig:PABGChuss}, we present the magnetic field configuration at the central region of our Galaxy by discriminating between the polarizations due to the disk and GC origin. The peak of the histogram of the position angle is $\\sim 20\\degr$ (see Fig. \\ref{fig:HistPAFGBG}), almost parallel to the Galactic plane. This coincidence, and the deficiency of $\\theta_{\\mathrm {R-B}}$ around $-60\\degr$, the angle perpendicular to the Galactic plane, show the basically toroidal geometry of the magnetic field. We cannot find a clear systematic dependence of position angle on Galactic latitude (from $b \\approx -0\\fdg27$ to $+0\\fdg18$), suggesting that there is no systematic transition of the magnetic field direction in this region.   The direction of dust grain alignment at the GC has been investigated from polarized dust emission in the mid- and far-infrared and sub-mm wavelengths. The magnetic field implied by the emission is generally parallel to the Galactic plane in the circumnuclear disk \\citep[e.g.,][]{Werner88,Morris92,Hildebrand93}. As seen in Fig. \\ref{fig:PABGChuss}, the magnetic field configuration we obtained at the GC shows a good agreement globally with those obtained by \\citet{Dotson00}, \\citet{Novak00}, and \\citet{Chuss03}, which are the highest angular resolution polarimetry data sets in the FIR/sub-mm wavelengths. The local features also show an excellent agreement: an X-shaped feature extends from $(\\Delta \\alpha, \\Delta \\delta) = (-3\\arcmin, +10\\arcmin)$ down through ($+5\\arcmin, 0\\arcmin$) \\citep[described in \\S 3.1.4,][]{Chuss03} is also confirmed in our map. The configuration at M$-$0.13$-$0.08 around ($-1\\arcmin, -5\\arcmin$) is also reproduced. The polarized FIR/sub-mm emission comes from molecular clouds, which are known to be located in the GC. Therefore we conclude that the position angles derived from our NIR polarimetry represent the direction of the aligned dust grains {\\it in} the GC.   \\citet{Chuss03} has suggested that the magnetic field aligns generally perpendicular to the Galactic plane in low density regions, while the field has a toroidal configuration in high density regions. They explain this correlation using the idea that in underdense regions, the magnetic field energy density can support itself against gravitational forces, preserving a primordial poloidal magnetic field. In overdense regions like molecular clouds, on the other hand, the gravitational forces are strong enough to shear the magnetic field into a direction parallel to the Galactic plane. In this context, lower density regions should have a poloidal configuration. However, in our vector map, the field shows a predominantly toroidal direction even at locations where the FIR/sub-mm emission is too weak for polarimetry [at the southeastern corner in Fig. \\ref{fig:PABGChuss}, see also Fig. 1 in \\citet{Chuss03} and Fig. 2 in \\citet{Novak03}.] The weak FIR/sub-mm emission suggests a paucity of dense clouds along the lines of sight, and the polarization in this direction can be considered to be interstellar in origin. Hence the ``interstellar'' magnetic field at the GC is dominated by a toroidal configuration, and the primordial poloidal magnetic field, if it existed, is not preserved today in this region.    Several radio filaments exist in our observed field, and three of them are prominent: the GC Radio Arc \\citep{Yusef84}, and the Northern and Southern Threads \\citep[also known as G0.08+0.15 and G359.96+0.09;][]{Morris85}, which are shown in Fig. \\ref{fig:PABGChuss}. Polarization studies have confirmed that the emission from the filaments is strongly linearly polarized, and that the internal magnetic field orientations are parallel to the long axes of the filaments \\citep{Tsuboi86,Lang99}. The simplest interpretation of these observations, combined with the discovery of more filaments \\citep[e.g.,][]{LaRosa04,Yusef04}, is that poloidal magnetic fields are pervasive throughout the central few hundred pc. Although the polarizations originating from the filaments cannot be detected in our observation, those toward the surrounding fields can be detected. The filaments have a width of less than $\\sim 10\\arcsec$ \\citep{Lang99}, and probably have a depth similar to their width. On the other hand, the size of the grids of our analysis is $2\\arcmin \\times 2\\arcmin$, and the polarization shown in Fig. \\ref{fig:PABGChuss} is the average of $\\sim$1$-$2 kpc from the GC along the line of sight. Hence most of the stars we detected do not show a polarization originating from the filaments. However, we can detect the average polarization near the line of sight toward the filaments. As shown in Fig. \\ref{fig:PABGChuss}, most of the position angles of the grids including the filaments align nearly perpendicular to them rather than parallel. Therefore, Fig. \\ref{fig:PABGChuss} suggests that the average magnetic field has a toroidal configuration even around the sight-lines toward the filaments. We note again that the polarization is the average along the line of sight and does not originate from the area close to the filaments.    \\subsection{NIR Polarization as a New Tool for Mapping the GC Magnetic Field}   We have shown that the polarization of starlight can be a probe of the magnetic field near the GC. \\citet{Morris98} enumerated five different ways in which the magnetic field near the GC has been studied: morphology, polarization angle, Faraday rotation of the radio continuum, Zeeman effect, and polarized dust emission in FIR/sub-mm wavelengths. The polarization of starlight has not been employed for such investigations.  The optical polarization of starlight can trace the structure of the local magnetic field, but cannot detect stars near the GC due to large extinction. NIR polarimetry has been carried out toward the GC prior to our observations (see \\S \\ref{subsec:PrevNIR}), but division of polarization into several components (e.g., originating from the Galactic disk and the central region) has not been done previously. The wide field-of-view of the NIR polarimeter SIRPOL, and the statistical treatment of tens of thousands of stars enable us to study the magnetic field near the GC; that is, the NIR polarimetry of starlight is a new way to investigate the magnetic field.   NIR polarimetry has the advantage of providing information about the magnetic field at locations where FIR/sub-mm emission is weak. NIR polarization of starlight is attributed to extinction along the line of sight by aligned dust grains, while FIR/sub-mm polarization is due to emission from the aligned dust. Hence, NIR polarimetry can investigate the magnetic field in regions where FIR/sub-mm polarimetry is absent, if background stars exist. This advantage is clearly shown in Fig. \\ref{fig:PABGChuss}. We have detected polarization at positions where blue bars are not shown. The distribution of the position angles including such low emission regions shows a globally toroidal magnetic configuration at the GC."
    },
    "0809/0809.0333_arXiv.txt": {
        "abstract": "We consider radiation emitted by the jitter mechanism in a Blandford-McKee self-similar blastwave. We assume the magnetic field configuration throughout the whole blastwave meets the condition for the emission of jitter radiation and we compute the ensuing images, light curves and spectra. The calculations are performed for both a uniform and a wind environment. We compare our jitter results to synchrotron results. We show that jitter radiation produces slightly different spectra than synchrotron, in particular between the self-absorption and the peak frequency, where the jitter spectrum is flat, while the synchrotron spectrum grows as $\\nu^{1/3}$. The spectral difference is reflected in the early decay slope of the light curves. We conclude that jitter and synchrotron afterglows can be distinguished from each other with good quality observations. However, it is unlikely that the difference can explain the peculiar behavior of several recent observations, such as flat X-ray slopes and uncorrelated optical and X-ray behavior. ",
        "introduction": "The synchrotron external shock model (Meszaros \\& Rees 1997) has been used with great success to model GRB afterglows since their discovery (Wijers, Rees \\& Meszaros 1997; Piran 1999; Panaitescu \\& Kumar 2001; Zhang \\& Meszaros 2004). However, its physical foundations have been put under debate, principally because a coherent equipartition magnetic field is assumed to be generated at the shock front. In addition, the fraction of internal energy that is given to the magnetic field is assumed independent of the location in the blastwave (the distance from the shock) and of the properties of the shock (e.g. its Lorentz factor).  The two latter assumptions have been tested phenomenologically on afterglow data. Rossi \\& Rees (2003) considered a magnetic field that is allowed to decay behind the shock front. They concluded that afterglow observations require a magnetic field that has an approximately uniform intensity throughout the blastwave (or at least within a distance $R/\\Gamma^2$ from the shock front). Yost et al. (2003) studied the possibility that the fraction of the internal energy stored in the magnetic field component grows with the decrease of the Lorentz factor of the fireball. They concluded that this effect could explain a sample of radio afterglows that have flat late time radio decay.  Even though these works present interesting constraints on the behavior of the magnetic field, they do not address the mechanism that generates it, its structure, and whether the radiation mechanism is really synchrotron or not. Analytic (Medvedev \\& Loeb 1999) and numerical work (Silva et al. 2003; Nishikawa et al. 2003; Frederiksen et al. 2004; Medvedev et al. 2005; Spitkovsky 2008; Chang, Spitkovsky \\& Arons 2008) have started to unveil recently the mechanism by which a magnetic field component is generated behind a relativistic collisionless shock front. All the studies find that the magnetic field is generated through some kind of two-stream instability, such as the Weibel instability (Weibel 1959). Magnetic fields are generated behind the shock with an energy of the order of 30 per cent of the total available internal energy. They are observed to subsequently decay and, possibly, to stabilize at a level of a fraction of a per cent of the internal energy. Whether they decay even further deep into the shocked material is a matter of debate. The structure of the field is observed to be highly tangled and planar behind the shock front and to grow more coherent deep into the blastwave.  The correlation length of the magnetic field right behind the shock is so short that the synchrotron approximations do not hold and a new radiative regime sets in. This form of radiation, called jitter radiation (Medvedev \\& Loeb, 1999; Medvedev 2000, 2006 or, sometimes, diffusive synchrotron radiation, Fleishman 2006ab), has been studied analytically, semi-analytically and numerically (Medvedev \\& Loeb 1999; Medvedev 2000; Frederiksen et al. 2004; Medvedev 2006; Fleishman 2006a; Medvedev et al. 2007; Workman et al. 2008). In this work we apply the theory of jitter radiation to compute the emission from a Blandford \\& McKee (1976) blastwave, taking into account its structure, the special relativistic effects, and the light propagation time. The results include spectra, lightcurves and images, analogously to what is done by Granot, Piran \\& Sari (1999a) for the synchrotron case.  This paper is organized as follows: In \\S~2 we describe the numerical methods applied to integrate the radiation spectrum over the BM solution and over surfaces of equal arrival time; in \\S~3 we describe our results and in \\S~4 we summarize and discuss them. ",
        "conclusions": "We have calculated the spectra and light curves produced by jitter radiation from a relativistic fireball expanding into either an ISM (constant density) or wind ($\\rho \\propto r^{-2}$) external medium including self absorption and electron cooling, and then compared these results to results for synchrotron radiation in identical cases. We improved upon our previous computations (Medvedev et al. 2007) by performing the integration over the blastwave and taking into account the light travel time of the photons.  The primary difference between jitter and synchrotron radiation is that jitter radiation produces a flat spectral slope below the peak frequency, rather than the $\\omega^{1/3}$ slope seen in synchrotron radiation.  Jitter and synchrotron radiation are not distinguishable in the optically thick regime or above the peak frequency, only in between.  The flat spectral slope of jitter radiation has the effect of decreasing the slope of the light curve by a factor of $t^{-1/2}$ compared to synchrotron radiation in the relevant frequency and temporal intervals.  This means that the slope of the light curve, while the observed frequency is below the peak frequency, decreases from $t^{1/2}$ for synchrotron to $t^{0}$ for jitter in the ISM case and decreases from $t^{0}$ to $t^{-1/2}$ for the wind case.  This indicates that the jitter-ISM and synchrotron-wind cases would not be distinguishable from a single light curve as they both have a $t^{0}$ slope.  In the last four years, {\\it Swift} observations, combined with ground-based optical and radio telescopes have put the standard synchrotron external shock model under scrutiny. The new observations have shown temporally flat X-ray light curves (Nousek et al. 2006), chromatic breaks (e.g., Huang et al. 2007), and uncorrelated multi-band behavior (e.g., Blustin et al. 2006; Dai et al. 2007). The light curves and spectra of jitter and synchrotron radiation are different. However, the differences appear to be too subtle to explain the atypical afterglow observations summarized above. Even though jitter radiation can produce flat lightcurves, it does so in all wavebands and not only in the X-rays, and cannot explain the steep decay observed, for example, in GRB 070110 (Troja et al. 2007). Chromatic breaks can be produced by jitter radiation as well as by synchrotron, the only difference being the amount of change in the pre- and post-break decays. Pending a more quantitative comparison of the jitter light curves and spectra with the data, it seems that jitter is a viable radiation mechanism to explain the behavior of GRB afterglows. It shares, however, all the problems of synchrotron in the explanation of some of the most recent data. Some possible explanation and a review of the difficulties can be found, for example, in Ghisellini et al. (2007); Liang, Zhang \\& Zhang (2007); Uhm \\& Beloborodov (2007); Panaitescu (2008).   \\begin{figure*} \\psfig{file=figure1.eps} \\caption{{Schematic diagram of the region from which photons reach a distant observer simultaneously (Granot, Piran \\& Sari, 1999a).  The line A-B represents a line of sight along which the radiative transfer equation is solved to find $I_{\\omega}(y)$. } \\label{fig:1}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure2.eps} \\caption{{Comparison of afterglow spectra at $t=500$ seconds.  Spectra plotted are for radiation from a shock expanding into a wind or ISM-type medium produced by jitter and synchrotron radiation.  Both jitter spectra are flat between the peak frequency and optically thick cutoff, while the synchrotron spectra increase as $\\omega^{1/3}$ before reaching a broad peak.  The wind cases are optically thick at lower frequencies. } \\label{fig:2}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure3.eps} \\caption{{Comparison of surface brightness at $1$~eV at $t=1000$ seconds for the same 4 models of Fig.~\\ref{fig:2}.  In all cases, the emission is concentrated towards the outer edge.  Normalized transverse distance is defined as $y/y_{max}$ with $y$ defined as in Fig.~\\ref{fig:1}.} \\label{fig:3}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure4.eps} \\caption{{Simulated light curves at $h\\nu=10^{-3}$~eV ($\\lambda=1.24$~mm) for the 4 models of Fig.~\\ref{fig:2}.  All 4 models are optically thick at this frequency.  The ISM models decrease in luminosity as $t^{1/2}$, while the wind models increase as $t^{1}$.} \\label{fig:5}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure5.eps} \\caption{{Simulated light curve at $h\\nu=1$~eV ($\\lambda=1240$~nm) for the 4 models of Fig.~\\ref{fig:2}.  The synchrotron-ISM model luminosity increases at early times as $t^{1/2}$.  The synchrotron-wind and jitter-ISM models are both initially flat before transitioning to a power-law decay as the peak frequency decreases below $h\\nu=1$~eV.  The jitter-wind model initially decreases as $t^{-1/2}$ before transitioning to a steeper power-law decay.  } \\label{fig:6}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure6.eps} \\caption{{Comparison of light curves for jitter-wind model at different energies.  The transition time between the luminosity scaling as $t^{-1/2}$ and $t^{-(3p-1)/4}$ is dependent on frequency and gives the appearance of a chromatic break in a power-law decay light curve.  The time of this break ranges from about $t=10$s to $t=10^5$s for the frequencies shown in the figure.} \\label{fig:7}} \\end{figure*}  \\begin{figure*} \\psfig{file=figure8.eps} \\caption{{Comparison of light curves for jitter-wind model in the X-ray ($h\\nu=1$~keV), visible ($\\lambda=575$~nm) and radio ($\\nu=24.2$~GHz) bands, observed from $t=60$s to $t=10^5$s.  The x-ray light curve is a single powerlaw decay, the visible is a broken powerlaw and the radio, in the optically thick regime, is a powerlaw increase.} \\label{fig:9}} \\end{figure*}"
    },
    "0809/0809.2793_arXiv.txt": {
        "abstract": "It has been observed that Cepheids in the Magellanic Clouds have lower masses for the same luminosity than those in the Milky Way.  The model, from \\cite{Neilson2008}, of pulsation--driven mass loss for Cepheids is applied to theoretical models of Cepheids with metallicity consistent with the Milky Way and Large and Small Magellanic Clouds.   The mass--loss model is analyzed using the metallicity correction of the Period--Luminosity relation to compare the ratio of mass loss of Cepheids with lower metallicity to that of Cepheids with solar metallicity.  It is determined that mass loss may be larger for the lower metallicity Cepheids,  counterintuitive to radiative driving estimates.  Also the mass--loss rates of theoretical Cepheid models are found to be up to $5\\times 10^{-9}$ for Galactic Cepheids, $5\\times 10^{-8}$ for Large Magellanic Cloud Cepheids, and  $2\\times 10^{-7}M_\\odot /yr$ for Small Magellanic Cloud Cepheids.  It is argued that mass loss increases as metallicity decreases for Cepheids with periods less than $20$ days and that mass loss decreases for longer periods.  Assuming dust forms in the wind of a Cepheid at some distance, the infrared excess of the models is computed, finding the infrared brightness is approximately a magnitude larger due to mass loss. The infrared magnitudes are compared to recently published Period--Luminosity relations as a test of our predictions. ",
        "introduction": "The study of classical Cepheid variables has been a cornerstone of stellar and Galactic astrophysics since the discovery of the Period--Luminosity relation \\citep{Leavitt1912}.  Observations of light curves and period provide information for testing stellar structure and evolution while the Cepheid Period--Luminosity relation makes them ideal standard candles.  It has been by understanding the physical structure and mechanism of pulsation that Cepheids have become powerful tools, but there are still phenomena being discovered.  One phenomenon was the recent discovery of infrared excesses from circumstellar material surrounding Galactic Cepheids \\citep{Kervella2006,Merand2006,Merand2007}.  It has been suggested that the cause of the infrared excess is dust forming in a wind from the Cepheids \\citep{Kervella2006}. There is additional evidence for mass loss, albeit circumstantial and wide ranging.    Analysis of IRAS observations \\citep{McAlary1986, Deasy1988} infers mass--loss rates up to $10^{-6} M_\\odot/yr$ for some Cepheids, although radio observations have placed upper limits of $10^{-9}$--$10^{-7}M_\\odot/yr$ \\citep{Welch1988}. Also \\cite{Rodrigues1992} determined an upper limit of $\\dot{M} \\le 7\\times 10^{-10}M_\\odot/yr$ for SU Mus and \\cite{Bohm-Vitense1994} found $\\dot{M} \\approx 5\\times 10^{-5} M_\\odot/yr$ for $l$ Car.  Furthermore the discovery of circumstellar material surrounding $l$ Car, Polaris, $\\delta$ Cep, Y Oph \\citep{Kervella2006, Merand2006, Merand2007} as well as RS Pup and SU Cas being associated with nebulae \\citep{Havlen1972, Turner1984} provide more evidence for mass loss.  Theoretical evidence for mass loss based on hydrodynamics is presented by \\cite{Willson1986} and \\cite{Brunish1989}. In \\citet[hereafter Paper I]{Neilson2008} we developed an analytic method to calculate mass--loss rates for Cepheids in which the mass loss is driven by the combination of pulsation and radiation. The method was applied to observed Galactic Cepheids where the mass--loss rates are predicted to range from $10^{-10}$--$10^{-7}M_\\odot/yr$. This implies that over the lifetime of Cepheid evolution, typically $10$ to $30 Myr$, a significant amount of mass can be lost.  Mass loss is interesting in Cepheids not only because it produces infrared excess and may affect the period of pulsation, but it has also been proposed as a solution to the Cepheid mass discrepancy, which is the difference between mass estimates of Cepheids by stellar evolution models and by pulsation models.  Historically, the mass determinations differed by as much as  $40\\%$ \\citep{Cox1980}.   The introduction of the OPAL opacities \\citep{Rogers1992} brought a partial resolution to the mass discrepancy problem by reducing it to order $10\\%$ \\citep*{Moskalik1992}, however, the discrepancy is still an issue. \\cite{Sebo1995} found a mass discrepancy of about $20\\%$ for Cepheids with periods ranging from  $2.7$ to $30$ days in the Large Magallanic Cloud (LMC) cluster NGC 1850; likewise \\cite*{Beaulieu2001} found similar results for Cepheids from the OGLE database for the LMC as well as a larger mass discrepancy for Small Magallanic Cloud (SMC) Cepheids.  \\cite{Brocato2004} found masses of Cepheids in the LMC cluster NGC 1866 to be smaller than predicted by evolution by almost $30\\%$ for short period Cepheids. The only dynamical mass estimates for Cepheids are from observations of Galactic Cepheids in binary systems; these are all found to have masses less than that predicted by evolutionary calculations \\citep{Bohm-Vitense1998, Evans2006}.  The current status of the mass discrepancy for Galactic Cepheids was presented by \\cite{Caputo2005}, who give the discrepancy as a function of period.  For short periods, the mass discrepancy is approximately $10\\%$ -- $20\\%$, and this decreases with increasing period.  \\cite{Keller2008} argues against this conclusion based on an analysis of the behavior of the mass--luminosity relation of Cepheids at large masses, and he proposes the mass discrepancy is constant with respect to mass, about $20\\%$.  The results differ because of the stellar evolution models used; \\cite{Caputo2005} use stellar evolution models that do not include a prescription for mass loss on the main sequence or on the red giant branch while the \\cite{Keller2008} results use stellar evolution models that do include mass loss.  Mass loss affects the evolution of the massive stars $M > 10M_\\odot$ and causes a non--linear mass--luminosity relation for Cepheids,  the various mass--luminosity relations are shown in Figure 2 of \\cite{Keller2008}.  \\cite{Keller2002} explore the mass discrepancy of LMC Bump Cepheids using non--linear pulsation models and find a mass discrepancy of about $20\\%$.  \\cite{Keller2006} determined that the mass discrepancy for SMC Cepheids is about $20\\%$ for Cepheids with mass $5$--$7M_\\odot$, while the discrepancy for LMC Cepheids is about $17\\%$, determined by matching MACHO and OGLE observations with non--linear pulsation models with varying mass and metallicity.  This observation has been verified for the LMC  \\citep{Testa2007}, and it also agrees qualitatively with the analysis of \\cite{Cordier2002}.  While mass loss is one possible cause of the mass discrepancy, a second suggested source is convective core overshooting in the main sequence progenitors of Cepheids \\citep{Huang1983}.  There is no direct evidence of this in Cepheids other than the mass discrepancy, coupled with the argument that mass loss is negligible over the lifetime of stars \\citep{deJager1988, Girardi2000}.  However there is evidence for convective core overshooting from stellar isochrones fit to stellar cluster observation \\citep[for example][]{Mucciarelli2007,Chiosi2007,Rosvick1998}. \\cite*{Schroder1997} used observations of $\\zeta$ Aurigae systems  and \\cite*{Ribas2000} used observations of detached binaries to constrain convective core overshooting.  Representing convective core overshooting as $\\Lambda_{\\rm{CCO}} = \\alpha H_{\\rm{P}}$, where $H_{\\rm{P}}$ is the pressure scale height, the free parameter $\\alpha$ produces a match to the observations for $\\alpha \\approx 0.1$--$0.3$.  \\cite{Ribas2000} find larger convective core overshooting for the binaries V380 Cyg, and HV 2274.  This evidence is circumstantial because the parameter describing convective core overshoot is obtained by matching models that  might ignore other physical processes.  The purpose of this article is to investigate the effect of metallicity on mass loss in Cepheids using the analytic model from Paper I.  In particular, we wish to explore the potential for Cepheid mass loss in the Large and Small Magellanic Clouds.  Understanding Cepheid mass loss in the Clouds will provide insight into the source of the mass discrepancy. The goals of this work are to differentiate how mass loss and convective core overshoot would contribute to the mass discrepancy, as well as to predict observational consequences of mass loss, such as infrared excess.  We start this exploration by deriving the metallicity dependence of the analytical mass--loss formulation from Paper I  in Section 2.  In Section 3, the analytical model is applied to Cepheids in the LMC and SMC to predict the potential behavior at lower metallicity.  Mass--loss rates are calculated in Section 4 for theoretical models of Galactic, LMC and SMC Cepheids and the resulting infrared luminosities  are predicted in Section 5. The dependence of the mass discrepancy on pulsation--driven mass loss is explored in Section 6. ",
        "conclusions": "In this article, we have investigated the hypothesis that mass loss is important for the lower metallicity Cepheids in the LMC and SMC. This has been done by rewriting the analytic method for computing pulsation--driven mass--loss rates as a ratio of mass--loss rates for two Cepheids that are alike in effective temperature and luminosity but with different metallicities, and allowing the other parameters to vary, such as the radius and the period of pulsation.  This leads to a parameter space that can be tested to understand the potential behavior of Cepheids at lower metallicity.  Over the majority of the parameter space the relative mass--loss of the two Cepheids is less than unity, reflecting the explicit dependence on metallicity.  However, there exists two regimes where the relative mass--loss may be larger for lower metallicity.  The first regime is where the relative mass is less than unity and the relative radius is larger than unity, and thus the ratio of the pulsation periods may be greater than unity.  The other regime is where the pulsation period of the lower metallicity Cepheid is less than that of the Cepheid with solar metallicity and the ratios of the mass and radius are approximately unity.  There is an additional parameter constraining this result.  The mass loss is only larger for the lower metallicity Cepheid if the mass loss is significant for the similar solar--metallicity Cepheid. In other words, if mass loss is significantly amplified by pulsation for the solar metallicity Cepheid then the mass loss may be further amplified by pulsation for the lower metallicity Cepheid.  We tested if the parameter space that caused the pulsation amplification of mass loss is consistent with with properties of LMC and SMC Cepheids.  This was done using the metallicity correction of the Period--Luminosity relation to determine periods of LMC and SMC Cepheids relative to Galactic Cepheids with the same luminosity.  Also we used the Period--Radius relation to determine how the radius changes for the lower metallicity Cepheids.  The metallicity correction suggests that lower metallicity Cepheids have longer pulsations periods than Galactic Cepheids and the longer pulsation periods imply that the radii of lower metallicity Cepheids are larger than Galactic Cepheids for the same luminosity.  These behaviors are consistent with mass--loss rates being further amplified implying mass loss is important for LMC and SMC Cepheids.  We have calculated mass--loss rates for theoretical models of Large, and Small Magellanic Cloud, and Galactic Cepheids using the method presented in Paper I. It should be noted that the theoretical models use stellar evolution calculations that do not include mass loss in earlier stages of evolution.  This affects the Cepheid mass--luminosity relation.  It is not clear how the pulsation--driven mass--loss rates of massive ($M>10M_\\odot$) Cepheids would change if we use the non--linear mass--luminosity relation of \\cite{Keller2008}.  The non--linear mass--luminosity relation suggests that at large mass the luminosity is less for a given mass than what is predicted by \\cite{Caputo2005}.  This would suggest that pulsation--driven mass--loss rate would be less, but it is unclear what the pulsation period would be in this case.  The pulsation period of a Cepheid following the non--linear mass--luminosity relation may be smaller than the period of a Cepheid following the linear mass--luminosity relation.  If the period is smaller then the shocks that help drive mass--loss are more efficient and the mass--loss rate may increase.  It would be very interesting to test this possibility.  In the calculation of the mass--loss rates, it was found that the peak mass--loss rates in the second crossing of the instability strip in the SMC is almost a factor of $10$ larger than in the LMC, which in turn, is almost $10$ times larger than that in the Milky Way. Furthermore it was found that the period of about $10$--$20$ days is the point where the dependence of mass loss on metallicity changes sign. This result may be interpreted as the mass loss being amplified significantly for the shorter period Cepheids and radiative driving being the dominant driving mechanism in longer period Cepheids.  The larger mass--loss rates for lower metallicity Cepheids is consistent with the conclusion that mass loss may be amplified more by pulsation at lower metallicity.  It is believed that dust forms in the wind of a Cepheid, generating an infrared excess.  From the predicted mass--loss rates of theoretical models of Cepheids, we compute Color Magnitude Diagrams as an observational prediction.  The result, however, depends on both the dust--to--gas ratio and the mass--loss rate.  The color excess computed here is thus the maximum value because the dust--to--gas ratio used in this work is the maximum value.  The dust--to--gas ratio may be a few times to an order of magnitude less, implying the color excess may be about a factor of 2--3 less.  The predicted $3.6$ $\\mu m$ luminosity computed from the sum of the stellar and circumstellar shell luminosities was found to be  consistent with the observed  infrared Period--Luminosity relation at $3.6$ $\\mu m$ \\citep{Ngeow2008, Freedman2008} as well as possibly agreeing with the set of Cepheids from \\cite{Ngeow2008} that were considered outliers with much larger luminosities due to blending effects.  The ability of mass loss to explain the mass discrepancy has also been explored.  It is shown that pulsation--driven mass loss is a possible explanation, but there still is not enough information to say this with certainty.   The most difficult aspect to explain is the $17\\%$ mass discrepancy at large mass where the predicted mass--loss rates are smaller than the necessary values of $10^{-6}M_\\odot  /yr$. However, the $17\\%$ mass discrepancy is an upper limit and the mass discrepancy at large masses range from $17\\%$ down to zero.  At smaller masses the pulsation--driven mass loss is consistent with a mass discrepancy of $17\\%$.  For the LMC and SMC the mass--loss rates increase,  implying mass loss can also explain the metallicity dependence of the mass discrepancy.  The results of this work suggest that it is possible for mass--loss rates of lower metallicity Cepheids in the  LMC and SMC to be greater than the mass--loss rates of similar Galactic Cepheids. This result is counterintuitive, one might expect mass loss to scale as the metallicity.  However, the pulsation period, radius and other parameters that describe a Cepheid and its mass loss also have dependencies on the metallicity, and the combination of these parameters and the dependencies cause the mass--loss rates to be potentially more significant.  Furthermore, the amount of mass loss predicted for Magellanic Cloud Cepheids is consistent with the observed behavior of the mass discrepancy of Cepheids as a function of metallicity from \\cite{Keller2006} as well as the dependence on mass \\cite{Caputo2005}.  Therefore mass loss cannot be discounted as a possible contribution to the mass discrepancy.  The amount of mass loss calculated using the pulsation mass--loss model is significant but there is not enough information to argue mass loss alone resolves the mass discrepancy.    This needs to be tested by with a large number of nonlinear pulsation models of Cepheids on both the first and second crossing to account for the effect of mass loss of the evolution of the period."
    },
    "0809/0809.3383_arXiv.txt": {
        "abstract": "{% Observational data and theoretical calculations show that significant small-scale substructures are present in dark molecular clouds.  These inhomogeneities can provide precious hints on the physical conditions inside the clouds, but can also severely bias extinction measurements.  We present \\textsc{Nicest}, a novel method to account and correct for inhomogeneities in molecular cloud extinction studies.  The method, tested against numerical simulations, removes almost completely the biases introduced by sub-pixel structures and by the contamination of foreground stars. We applied \\textsc{Nicest} to 2MASS data of the Pipe molecular complex.  The map thereby obtained shows significantly higher (up to $0.41 \\mbox{ mag}$ in $A_K$) extinction peaks than the standard \\textsc{Nicer} (Lombardi et al.\\ 2001) map.  This first application confirms that substructures in nearby molecular clouds, if not accounted for, can significantly bias extinction measurements in regions with $A_K > 1 \\mbox{ mag}$; the effect, moreover, is expected to increase in more distant molecular cloud, because of the poorer physical resolution achievable. ",
        "introduction": "\\label{sec:introduction}  Although the details of star and planet formation are very little known, there is a large consensus that these objects are created inside dark molecular clouds from the contraction and fragmentation of dense, cold cores.  As a result, in the last decades molecular clouds have been studied in detail using many different techniques, from optical number counts \\citep{1923AN....219..109W, 1937dss..book.....B}, to radio observations of carbon monoxide (CO) molecules \\citep{1970ApJ...161L..43W,1982ApJ...262..590F}, to near-infrared (NIR) extinction measurements \\citep{1994ApJ...429..694L}.  Molecular hydrogen and helium represent approximately $99\\%$ of the mass of a cloud, but the absence of a dipole moment in these molecules makes them virtually undetectable at the low temperatures ($T \\sim 10 \\mbox{ K}$) present in these objects.  Hence, the (projected) mass distribution of dark clouds is usually inferred from the distribution of relatively rare tracers (such as CO or dust grains) under the assumption that these are uniformly distributed in the cloud.  As pointed out by \\citet{1994ApJ...429..694L}, the reddening of background stars in NIR bands is a simple and reliable method to study the dust distribution and thus the hydrogen column density.  Compared to optical bands, NIR bands are less affected by extinction and are less sensitive to the physical properties of the dust grains \\citep{1990ARA&A..28...37M}, and thus their color excess can be directly translated into a hydrogen column density.  In a series of papers we reconsidered and improved the NIR color excess (\\textsc{Nice}) technique, by generalizing it to use three or more bands (\\textsc{Nicer}, see \\citealp{2001A&A...377.1023L}) and by taking a maximum-likelihood approach \\citep{2005A&A...438..169L}.  Although the \\textsc{Nice} and \\textsc{Nicer} techniques have been very successful in studying molecular clouds \\citep[see, e.g.,][]{2001Natur.409..159A}, they are plagued by two complications: small-scale inhomogeneities and contamination by foreground stars.  NIR studies typically assume a uniform extinction over (small) patches of the sky; however, in presence of substructures such as steep gradients, unresolved filaments, or turbulence-induced inhomogeneities \\citep{1994ApJ...429..694L, 1981MNRAS.194..809L, 2004ApJ...615L..45H}, the background stars used to estimate the cloud extinction would not be uniformly distributed in the patch, but would be preferentially observed in low density regions.  Similarly, in presence of foreground stars, their \\textit{relative\\/} fraction increases with the dust column density.  Unfortunately, both effects will bias \\textit{all\\/} color excess estimators toward low column densities; moreover, the bias will be especially important in the dense regions of molecular clouds, which is where the star formation takes place.  In this paper we describe in details the repercussions of small scale inhomogeneities and foreground stars in extinction measurements, and propose a simple method to obtain estimates of the column density in presence of both effects that, under simple working hypotheses, is asymptotically unbiased.  Our method does not rely on any (usually poorly verified) assumption regarding the small scale structure of the cloud; moreover, it can be applied to a general class of NIR extinction measurements (marked spatial point processes).  The paper is organized as follows.  In Sect.~\\ref{sec:cloud-substructure} we introduce a general smoothing technique virtually used in all cases to create smooth, continuous maps from the discrete pencil-beam extinction measurements carried out for each background star; we also show that in two typical applications (moving average and nearest neighbours smoothing) a bias is expected.  In Sect.~\\ref{sec:over-weighting-high} we propose a new method that can correct this bias without making any assumption on the underlying form of the cloud substructure.  This new method has been tested with numerical simulations, as described in Sect.~\\ref{sec:simulations}, and with a real-case application, presented in Sect.~\\ref{sec:sample-application}.  We discuss and comment the results obtained in Sect.~\\ref{sec:discussion}, and finally we briefly present our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   The main results obtained in this paper can be summarized in the following points: \\begin{itemize} \\item We discussed the effects of small scale inhomogeneities in the NIR extinction maps based on color excess methods, and showed that large inhomogeneities can significantly bias standard extinction maps toward low column densities. \\item We proposed a new estimator for $\\hat A$, \\textsc{Nicest}, and we showed that it is (i) equivalent to the usual estimator in low-density regions, (ii) \\textit{unbiased\\/} and \\textit{robust\\/} for any value of $A$, i.e.\\ insensitive to the actual form and amount of substructure present in the cloud. \\item We showed that the new estimator is also suitable in presence contamination by foreground stars. \\item We tested \\textsc{Nicest} against numerical simulations, and showed that it effectively reduces by a large factor the biases due to substructures and to foreground stars.  We also tested an alternative approach, the nearest neighbor, and showed that the results obtained from this interpolation are still severely biased. \\item We applied \\textsc{Nicest} to the Pipe nebula and showed a few preliminary properties of the resulting extinction map.  We also noted a direct connection between the bias of the \\textsc{Nicer} and the $\\Delta^2$ map defined by \\citep{Lombardi08a}. \\end{itemize}"
    },
    "0809/0809.1665_arXiv.txt": {
        "abstract": "We present an analysis of a new \\xmm\\ observation of the Seyfert 1 Galaxy NGC~985. The \\epic\\ spectra present strong residuals to a single power-law model, indicating the presence of ionized absorbing gas and a soft excess. A broad-band fit to the \\epic\\ and \\rgs\\ spectra shows that the continuum can be well fit with a power-law ($\\Gamma\\approx 1.57$) and a blackbody component ($kT\\approx0.09$ keV). The \\rgs\\ can be modeled either with two or three absorption components. In the two absorber model the low-ionization one, with $\\log$U$\\approx.05$ and $\\log$N$_H\\approx21.08$ accounts for the presence of the Fe M-shell unresolved transition array (Fe VII-XIII), and the high ionization component ($\\log$U$\\approx1.31$ and $\\log$N$_H\\approx21.99$) is required by the presence of several Fe L-shell transitions. The data suggest the presence of a third ionized component with higher ionization, so that the Fe L-shell absorption features are produced by two different components (one producing absorption by Fe XVII-XX, and the other absorption by Fe~XX-XXII). However, the presence of the third absorbing component cannot be detected by means of an isolated absorption line in a significant way, so we consider this detection only as tentative. Interestingly, all ionization components have similar kinematics (with outflow velocities $\\sim 280$ km s$^{-1}$).  In addition, whether two or three absorbers are considered, the components appear to be in pressure balance within $1\\sigma$. These results give further support to the idea that warm absorbers in AGN consist of a two or three-phase medium. We note that, while in the model with only two absorbers one of them (the high ionization component) lies on an unstable branch of the thermal equilibrium curve, in the model with three absorbers all of the components lie on stable branches of the curve. This gives further plausibility to a multi-phase absorber. ",
        "introduction": "}  Highly ionized (or `warm', WA) absorbers are observed in about half of X-ray spectra of broad line active galactic nuclei (AGN), both Seyfert 1s (e.g.  Halpern 1984; Reynolds et al. 1997, George et al. 1998) and quasars (Piconcelli et al. 2005). The absorption lines of these systems are blueshifted with respect to the optical emission lines (with velocities of the order of a few hundreds to $\\sim 2000$ km s$^{-1}$), implying outflowing winds. Their frequency, combined with evidence for transverse flows (Elvis 2000; Crenshaw et al. 2003; Arav 2004) suggest that WAs are actually ubiquitous in AGN, but become directly visible in absorption only when our line of sight crosses the outflowing  material.  Ionized absorption is also observed in the UV band.  However, the exact relation between the X-ray absorbers and Narrow Absorption Line (NAL) UV systems is still uncertain. These absorbers must be related, as exactly the same AGNs show both (Mathur et al. 1995), and in many cases, with similar outflow velocities and similar ionization state.  Several studies have shown that the absorbers can be described in a simple way. It has been shown that X-ray ionized absorbers can be well modeled by   including only two or three absorbing components (Krongold et al. 2003, 2005a;  Netzer et al. 2003; Blustin et al. 2003;  Steenbrugge et al. 2005). It has also been found that these components  appear to be in pressure balance with each other (Krongold et al. 2003,2005a,2007; Netzer et al. 2003). This has been interpreted as evidence of a multi-phase medium for the structure of the absorber. Alternatively, it has also been suggested that we could be looking at a single radially stratified medium with constant pressure (Rozanska et al. 2006). Other studies suggest that the warm absorber could consist of a continuous radial flow of ionization structures (Behar et al. 2003; Ogle et al. 2004; Steenbrugge et al. 2005), though variability studies seem to rule out this possibility (Krongold et al. 2005b, 2007). An intermediate case consisting of a double peaked continuous distribution of material has also been reported for Mkn 279 (Constantini et al. 2007), although Fields et al. (2007) reports that these data is consistent with a multiphase medium, once metalicity effects are taken into account. The exact nature of AGN winds still remains a mystery.   The nature of the wind directly impacts our understanding of both the structure and physics of the active nucleus itself, and the effect of AGNs on  their larger scale galactic and extragalactic environment. One of the key questions is where do these absorbers originate. Several studies have set lower limits on the location of the ionized outflows, placing them at parsecs from the central ionizing source (Behar et al. 2003, Netzer et al. 2003), and suggesting an origin in the putative obscuring torus (Krolik \\& Kriss et al. 2001, Blustin et al. 2005). However, other studies have found WA much closer in, at sub-pc scales (Pounds et al. 2003; Kaastra et al. 2004, Krongold et al. 2007), leading to the suggestion that warm absorbers originate from the accretion disk (Krongold et al. 2007). The mass loss rate inferred for these outflows is still uncertain by orders of magnitude and is strongly debated, as this rate depends critically on the distance of the outflow to the central source.  \\subsection{The Ionized Absorber in NGC 985}   NGC 985 (Mrk 1048) is a Seyfert 1 galaxy located at redshift 0.04274$\\pm 0.00005$ (12814$\\pm 15$ km s$^{-1}$; Arribas et al. 1999, based on stellar absorption features, see Krongold et al. 2005 for a detailed description of NGC 985).    An ionized absorber has been observed in the X-ray spectra of this object; first through low resolution data (Brandt et al. 1994 in a ROSAT-PSPC spectrum of this source; Nicastro et al. 1998, 1999, with an ASCA observation). High resolution observations carried out with the HETGS on board the {\\em Chandra} X-ray telescope (and performed $\\sim$1 year before the \\xmm\\ observations discussed here) have confirmed the presence of the absorber (Krongold et al. 2005 hereafter K05). In this analysis, it was found that two absorbing components were required to fit the spectrum of NGC 985. Furthermore, the gas pressure of the two components was consistent, suggesting that the absorber could be formed by a two-phase wind. The UV band also shows the presence of absorption lines produced by ionized gas (Arav 2002). Five components are observed in outflow, with velocities ranging from 780 to 243 km s$^{-1}$. A sixth component is observed inflowing at -140  km s$^{-1}$.   In this paper we present the analysis of the ionized absorber in NGC985 carried out on XMM-Newton X-ray spectra. In \\S  \\ref{par:data} we describe the reduction of the data. In \\S \\ref{par:analysis}, we present the data analysis, including a detailed model of the ionized absorber in NGC 985.  In \\S \\ref{sec:disc} we discuss our results, and compare them with those obtained with the {\\em Chandra} data by K05. Finally, in \\S \\ref{sec:conc} we present our conclusions. ",
        "conclusions": "}   The \\xmm\\ data of NGC 985 shows the presence of an ionized (warm) absorber in the X-ray spectra  of this source. The data can be modeled with two or maybe three absorbing components. The presence of a third absorber is statistically significant in the models, however we could not identify  a discrete significant absorption feature required by this component (and not modeled with only two absorbers). Therefore, we consider the detection of this component as tentative. The absorbing components (whether two or three different absorbers are present) are in pressure balance (within $1\\sigma$) and have outflow velocities indistinguishable from each other.  Thus, our results give further support to the idea that ionized absorbers in AGN form a multi-phase medium, where a hot gas component is confining the other(s). However,  in the model with only two absorbers, one is located in an unstable branch of this curve. On the other hand, in the three-phase model all the components lie in stable regions of the thermal equilibrium curve. Thus, only the model with three absorbing components allows a multi-phase absorber to survive for long periods of time. In this case, the multi-phase medium would not dissolve even if the crossing time is much larger than the timescale in which the medium adapts to mechanical perturbations."
    },
    "0809/0809.1948_arXiv.txt": {
        "abstract": "{Methanol masers are associated with young high-mass stars and are an important tool for investigating the process of massive star formation.} {The recently discovered methanol maser ring in G23.657$-$00.127 provides an excellent ``laboratory'' for a detailed study of the nature and physical origin of methanol maser emission, as well as parallax and proper motion measurements.} {Multi-epoch observations of the 12.2\\,GHz methanol maser line from the ring were conducted using the Very Long Baseline Array. Interferometric observations with milliarcsecond resolution enabled us to track single maser spots in great detail over a period of 2 years.} {We have determined the trigonometric parallax of G23.657$-$00.127 to be 0.313$\\pm$0.039\\,mas, giving a distance of $3.19^{+0.46}_{-0.35}$\\,kpc.  The proper motion of the source indicates that it is moving with the same circular velocity as the LSR, but it shows a large peculiar motion of $\\approx35$ km\\,s$^{-1}$ toward the Galactic center.} {} ",
        "introduction": "Two spectral lines of methanol masers, 6.7\\,GHz and 12.2\\,GHz, are strongly associated with young high--mass stars or protostars and enable one to investigate the stellar environment on scales of 100--1000\\,AU in great detail (Menten \\cite{menten}; Moscadelli et al. \\cite{moscadelli}; Minier et al. \\cite{minier00}). Interferometric studies of methanol masers have been carried out for more than ten years in order to constrain {\\it how} and {\\it where} massive stars are being born (Norris et al. \\cite{norris}; Phillips et al. \\cite{phillips}; Walsh et al. \\cite{walsh98}; Dodson et al. \\cite{dodson}). We started interferometric studies of the 6.7\\,GHz masers discovered in the Toru\\'n unbiased survey towards the Galactic plane (Szymczak et al. \\cite{szymczak02}).  One source, G23.657$-$00.127, was imaged during the European VLBI Network (EVN) session in 2004 when eight antennas were working for the first time at 5\\,cm wavelength. These observations revealed a new class of spherically symmetric methanol maser sources (Bartkiewicz et al. \\cite{bartkiewicz}). The ``ring'' of masers in G23.657$-$00.127 has a mean radius of 127\\,mas. Its circular geometry makes it a unique laboratory for detailed study of a methanol maser in an isolated massive star forming region. However, the morphology and velocity signature of the maser spots at 6.7\\,GHz did not allow us to determine unambiguously the origin of the ring--like structure. Therefore we started VLBI observations of the 12.2\\,GHz methanol maser line following its detection with the Toru\\'n 32\\,m antenna.  In this paper we report multi--epoch observations of G23.657$-$00.127 at 12.2\\,GHz using the NRAO\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.} Very Long Baseline Array (VLBA). We focus on the measurement of the source parallax and proper motion, providing a reliable distance. Maser variability and internal proper motion studies for both the 6.7 and 12.2\\,GHz lines will be presented in a separate paper. ",
        "conclusions": "Using VLBA observations of the 12.2 GHz methanol masers in G23.657$-$00.127, we have determined an accurate distance of 3.19$^{+0.46}_{-0.35}$\\,kpc to the maser ring. At this distance, the radius of the ring corresponds to a linear size of 405 AU. We have also measured the proper motion of the source, which we use to determine its 3-dimensional motion in the Galaxy. The source moves with a similar rotation speed as the LSR, but shows a high peculiar motion of 30--40 km\\,s$^{-1}$\\, toward the Galactic center. This large peculiar motion might be induced by interactions with the Galactic bar and is responsible for the large error in the kinematic distance."
    },
    "0809/0809.5100_arXiv.txt": {
        "abstract": "We present spectroscopy of 15 star-forming BzK galaxies (sBzKs) with $\\KAB \\lesssim 23$ in the Subaru Deep Field, for which $\\Ha$ and some other emission lines are detected in 0.9 to 2.3 $\\micron$ spectra with a resolution of $R$=500. Using $\\Ha$ luminosities, we obtain star formation rates (SFRs), and then specific SFRs (SSFRs) dividing SFRs by stellar masses, which are derived from SED fitting to $BVRi'z'K$ photometry. It is found that sBzKs with higher stellar masses have larger SFRs. A negative correlation is seen between stellar mass and SSFR, which is consistent with the previous results for $z \\sim 2$ galaxies. This implies that a larger growth of stellar mass occurs in less massive galaxies. In addition, gas-phase oxygen abundances, 12+log(O/H), are derived from the ratio of [\\ion{N}{2}]($\\lambda6584$) to $\\Ha$ using the N2 index method. We have found a correlation between stellar mass and oxygen abundance in the sense that more massive sBzKs tend to be more metal rich, which is qualitatively consistent with the relation for UV-selected $z \\sim 2$ galaxies. However, the metallicity of the sBzKs is $\\sim 0.2$ dex higher than that of UV-selected galaxies with similar stellar masses, which is significant considering the small uncertainties. The sBzKs in our sample have redder $R-K$ colors than the UV-selected galaxies. This galaxy color-dependence in the oxygen abundance may be caused by older or dustier galaxies having higher metallicities at $z \\sim 2$. ",
        "introduction": "Recent studies suggest that the era of $z \\sim 2$ is a turning point in galaxy formation and evolution. The cosmic star formation rate (SFR) begins to drop at $z \\sim 1-2$ from a flat plateau at higher redshifts \\citep{dic03,fon03,ly07}. Significant evolution of the Hubble sequence occurred at $z \\sim 1-2$ \\citep{kaj01}. It is also found that the number density of QSOs has a peak at $z \\sim 2$ \\citep{ric06}. These facts suggest that dramatic changes in the galaxy population occurred at $z \\sim 2$, which are important to study in further detail.  Galaxy properties are determined from spectral characteristics such as emission lines, absorption lines and strong continuum breaks. At $z \\sim 2$, the strongest of these features fall in the near-ultraviolet (NUV) and near-infrared (NIR) wavebands. NUV and NIR observations are thus essential to reveal detailed properties of $z \\sim 2$ galaxies, but these observations encounter many difficulties, including poor sky transparency except for some atmospheric windows, and strong OH airglow emission lines and thermal emission from the atmosphere. This is why the era of $z \\sim 2$ had been called a {\\lq}redshift desert{\\rq} until recently. However, recent advances in wide-field imaging in the NUV and NIR, along with multi-object spectrographs, are beginning to fill in this {\\lq}desert{\\rq} (e.g., Subaru Deep Field (SDF) survey, GOODS-North and South, UKIDSS, GMASS, and MUSYC \\citep{mccarthy99,red06b,erb06c,hay07,dad07,gra07,lan07,hal08,kri08,ly08}).  A photometric selection technique using $B-z$ and $z-K$ colors has been proposed to choose an unbiased sample of $z \\sim 2$ galaxies (BzK galaxies: \\citet{dad04}). Using these colors to find galaxies with a Balmer break between $z$ and $K$, the selected galaxies are then classified into star forming galaxies (sBzKs) and passively evolving galaxies (pBzKs). Thanks to wide and deep imaging surveys in the NIR as well as the optical, large new BzK samples are enabling us to measure their statistical properties, such as clustering strength \\citep{kon06,hay07,bla08}. The multiwavelength data also show that sBzKs are very actively forming stars\\citep{dan06,dad07}. On the other hand, due to lack of spectroscopic observations, there are relatively few accurately known redshifts, which are needed to derive accurate stellar masses and star formation rates from SED fitting to the multiwavelength data. Also, their physical properties such as metallicity are still not well-known.  More spectroscopy has recently become available for $z \\sim 2$ galaxies selected by their UV colors. \\citet{erb06a,erb06b} presented a mass-metallicity ($M$-$Z$) relation and SFRs for UV-selected galaxies at $z \\sim 2$ using 114 NIR spectra. They found a $M$-$Z$ correlation at $z \\sim 2$ as there is in the local universe, with the more massive galaxies being more metal rich. Although the correlation has a similar slope to the local one, it is shifted to lower metallicity by $\\sim 0.3$ dex. They have also found that SFRs derived from $\\Ha$ luminosities are in good agreement with those from de-reddened rest-UV luminosities, with an average dust-corrected SFR of $\\sim 30 \\Msun {\\rm yr^{-1}}$.  However, the rest-UV color selection misses about half of $z \\sim 2$ galaxies, which have a large amount of dust extinction \\citep{kon06}. The spectroscopic properties of this dusty population have yet not been clearly revealed. $K$-limited samples are better suited than UV-limited ones to investigate the relations between stellar mass and other fundamental properties at $z \\sim 2$, because they approximately correspond to samples limited by stellar mass, without excluding galaxies with large dust extinction. \\citet{kri08} carried out a NIR spectroscopic survey for $K$-bright galaxies at $z \\sim 2.3$. They suggest that studies with broadband photometric data alone may overestimate the number of massive galaxies at $z \\sim 2$, and underestimate the evolution of the stellar mass density, again indicating the importance of spectroscopy.  In this paper, we present NIR spectroscopy for 44 BzK galaxies in the SDF, and their spectroscopic properties. The structure of this paper is as follows. The photometric and spectroscopic data are described in \\S \\ref{sec;data}. In \\S \\ref{sec;analysis}, we examine the properties of the detected emission lines, including spectroscopic redshifts and fluxes. We then investigate active galactic nuclei (AGN) contamination in our spectroscopic sample in \\S \\ref{sec;agn}. In \\S \\ref{sec;discussions}, physical properties are derived from the emission lines, and then discussed. A summary is given in \\S \\ref{sec;conclusion}. Throughout this paper, magnitudes are in the AB system, and we adopt cosmological parameters of $h=0.7$, $\\Omega_{m0}=0.3$ and $\\Omega_{\\Lambda 0}=0.7$. ",
        "conclusions": "Emission line information enables us to estimate SFR and metallicity. In the local universe, more massive galaxies tend to have lower SFRs and higher metallicities \\citep[e.g.,][]{gal05}. It is important to examine whether a similar relation exists among these properties for high redshift galaxies.  In this section, we first describe the properties obtained from multiwavelength photometric data by SED fitting method. Then, spectroscopic properties are discussed, along with the photometric properties.   \\subsection{Properties from SED fitting} \\label{sec;sedfit} Properties such as stellar mass and dust extinction, $E(B-V)$, are obtained from fitting synthetic model spectra of continuous bursts \\citep{bru03} to the $BVRi'z'K$ photometry, assuming a standard \\citet{sal55} initial mass function (IMF), solar abundance, and the empirical dust extinction of \\citet{cal00}. These are shown in Table \\ref{table;photo}. Stellar mass is one of the most important parameters of galaxy properties and the most robust one obtained from SED fitting. Although sBzKs are well-fitted by the continuous burst model, fits using the $\\tau$ model, where the SFR declines exponentially with time, would make the stellar masses slightly smaller than those with the continuous burst model. Note that SEDs of the simple stellar population (SSP) model are irrelevant in this work, since sBzKs are still star-forming at $z \\sim 2$. It is also noted that if a \\citet{cha03} IMF is used, stellar masses (and SFRs) are smaller by a factor of 1.8 than those with a Salpeter IMF. Thus, stellar masses are determined within a factor of $\\sim 2.5$. Fitted $E(B-V)$s are mainly used to calibrate the dust extinction of emission line fluxes in section \\ref{sec;sfr}. As discussed in section \\ref{sec;agn}, almost all SEDs of the sBzKs are dominated by stellar emission with a clear 1.6$\\mu m$ bump. Therefore, any AGN contribution to the SED is small enough that parameters obtained from the SED fitting reflect properties of the stellar populations.   \\subsection{Star Formation Rates} \\label{sec;sfr} \\subsubsection{SFR derived from H$\\alpha$ and UV continuum} \\label{sec;sfrha} SFRs for sBzKs are derived from the luminosity of $\\Ha$ using the equation given in \\citet{ken98}; \\begin{equation} {\\rm SFR}_{\\Ha}[\\Msun {\\rm yr^{-1}}] = 7.9 \\times 10^{-42} L_{\\Ha} [{\\rm ergs\\ s^{-1}}], \\end{equation} assuming solar abundance, Salpeter IMF with mass limits of 0.1 and 100 $\\Msun$, and a constant SFR. $\\Ha$ luminosities are corrected for the dust extinction by equation (3) in \\citet{cal00}, $E(B-V) = 0.44 E_{\\rm gas}(B-V)$. The amount of nebular extinction is then estimated from $E_{\\rm gas}(B-V)$ using the extinction law given in \\citet{cal00}. The median amount of extinction of $\\Ha$ is 2.83 magnitude.  The amount of extinction at $\\Ha$ can be directly estimated from the apparent $\\Ha/\\Hb$ ratio. However, we do not apply this method for our sample, because only three galaxies have detection of both lines and even for these three the amount of extinction is not strongly constrained due to a large uncertainty in the $\\Hb$ flux.  The rest-UV luminosity is also used as a tracer of SFR. We derive SFR from the continuum at 1500\\AA, $L_{1500}$, which is predicted by the best-fit model to the $BVRi'z'K$ photometry. $L_{1500}$ is corrected for the dust extinction with Calzetti extinction law and converted into a SFR using the equation given in \\citet{dad04}; \\begin{equation} {\\rm SFR}_{\\rm UV}[\\Msun {\\rm yr^{-1}}] = \\frac{L_{1500}[{\\rm ergs\\ s^{-1}\\ Hz^{-1}}]}{8.85} \\times 10^{-27}. \\end{equation}  \\subsubsection{Comparison between ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$} \\label{sec;sfruv} Figure \\ref{fig;sfr_ha_uv} compares the SFR from the rest-UV luminosity with that from $\\Ha$. A correlation is seen between the two SFRs, but ${\\rm SFR}_{\\Ha}$ tends to be larger than ${\\rm SFR}_{\\rm UV}$ by a factor of three. The median ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$ are 182 and 61 $\\Msun {\\rm yr^{-1}}$, respectively. Also, this discrepancy between SFR measured by $\\Ha$ and from shorter wavelength observations is quite similar to what \\citet{hicks02} found in their sample of IR-selected galaxies.  On the other hand, \\citet{erb06b} report that ${\\rm SFR}_{\\Ha}$ for UV-selected $z\\sim 2$ galaxies is in good agreement with ${\\rm SFR}_{\\rm UV}$, which is inconsistent with our result. However, their observed $\\Ha$ luminosities are similar to ours. That is, the median $\\Ha$ luminosities of our sample and that of \\citet{erb06b} are 2.0 and 2.2 times $10^{42}$ ergs s$^{-1}$, respectively. This disagreement on resulting SFRs with \\citet{erb06b} may be due to different corrections for dust extinction at $\\Ha$. They used the color excess for the stellar continuum to de-redden $\\Ha$, that is, $E_{\\rm gas}(B-V)=E(B-V)$, which implies that their amount of dust correction for $\\Ha$ luminosity is systematically lower than ours by 1.6 magnitude on average. This leads to a difference in derived SFRs of a factor of 4.4. Using the same method as \\citet{erb06b} for de-reddening $\\Ha$, our correlation between ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$ becomes consistent with that of \\citet{erb06b}. In this case, the median ${\\rm SFR}_{\\Ha}$ is 50 $\\Msun {\\rm yr^{-1}}$.  \\citet{dad04} report that most color excesses of 24 sBzKs in the K20/GOODS area, which are derived from $U$ to $K$ SEDs, fall within a range of $0.2 \\lesssim E(B-V) \\lesssim 0.5$, which is consistent with our result (Table \\ref{table;photo}). The median $E(B-V)$ in our sample is 0.38. This implies that our estimation of $E(B-V)$ by SED fitting is reliable.  This difference between ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$ implies two possibilities. One is that a Salpeter IMF is unsuitable, if the equation of \\citet{cal00}, $E(B-V) = 0.44 E_{\\rm gas}(B-V)$, is assumed to be valid for star-forming galaxies at $z \\sim 2$. A more top-heavy IMF than Salpeter would account for the difference, since nebular emission results from ionizing photons shortward of the Lyman limit. The other is that the equation of \\citet{cal00} cannot be applied to $z \\sim 2$ galaxies. In any case, in what follows we use ${\\rm SFR}_{\\Ha}$ for sBzKs in the comparison with their stellar masses. Even if ${\\rm SFR}_{\\rm UV}$ were used instead, the discussion in the next section would remain unchanged.  \\begin{figure}[tb] \\begin{center} \\plotone{f8.eps} \\caption{Comparison between ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$. ${\\rm SFR}_{\\Ha}$s are derived from $\\Ha$ luminosities, and then ${\\rm SFR}_{\\rm UV}$s are from the luminosities at 1500 \\AA \\  estimated from the best-fitted SEDs. Both SFRs are corrected for dust extinction. The dotted line shows that ${\\rm SFR}_{\\rm UV}$ is equal to ${\\rm SFR}_{\\Ha}$. } \\label{fig;sfr_ha_uv} \\end{center} \\end{figure}   \\subsubsection{Relation to stellar mass} \\label{sec;RtoMstar} We examine the relation between ${\\rm SFR}_{\\Ha}$ and stellar mass. Figure \\ref{fig;mstar_sfr} suggests that sBzKs with higher stellar masses have larger SFRs, which is qualitatively consistent with \\citet{dad07}. However, the slope of our relation is shallower than theirs. One cause for this disagreement in the slope may be a difference between ${\\rm SFR}_{\\Ha}$ and ${\\rm SFR}_{\\rm UV}$. Figure \\ref{fig;sfr_ha_uv} suggests the difference between these two estimates of SFR is larger at lower SFRs. Since less massive sBzKs have lower SFRs, this difference results in shallower slope of the correlation between stellar mass and ${\\rm SFR}_{\\Ha}$. Other cause may be a difference in the method used to derive stellar masses. \\citet{dad07} have used the empirical relation between the stellar mass and $K$ given in \\citet{dad04}, which is calibrated based on $z-K$ color. We compare the stellar masses derived from a SED-fitting with those from the relation of \\citet{dad04}, finding that the difference between the two estimates of stellar mass is larger at smaller stellar mass. The stellar masses from the relation are larger by a factor of 2 -- 3 in $\\sim 10^{10} \\Msun$. Because this relation is derived using bright $K<22$ BzKs, the stellar masses of faint BzKs may be overestimated compared with those from SED fitting, resulting in the shallower correlation between stellar mass and ${\\rm SFR}_{\\Ha}$.  It is also possible that a flux-limited bias, discussed in \\S\\S 3.2 and 3.3, produces this disagreement. If galaxies with $\\left<E(B-V)\\right>=0.38$ and $F(\\Ha)=4.2\\times10^{-17} {\\rm ergs\\ s^{-1}\\ cm^{-2}}$, which is the lowest $\\Ha$ flux we have detected, are at $\\left<z\\right>\\sim1.7$, the SFR would be $\\sim87\\Msun {\\rm yr^{-1}}$. This suggests that it is likely that we cannot detect any emission lines for galaxies with SFR $\\lesssim$ 87 $\\Msun {\\rm yr^{-1}}$. Moreover, it is found that ${\\rm SFR}_{\\Ha}$s is three times larger than ${\\rm SFR}_{\\rm UV}$s on average. Combining this with the relation between stellar mass and ${\\rm SFR}_{\\rm UV}$ given in Fig. 17 of \\citet{dad07}, we find that the limiting SFR of 87 $\\Msun {\\rm yr^{-1}}$ corresponds to a stellar mass of $\\sim 1\\times10^{10}\\Msun$. These estimates indicate that it is possible that no galaxies with $\\lesssim 1\\times10^{10}\\Msun$ and $\\lesssim 87 \\Msun {\\rm yr^{-1}}$ are detected due to flux-limited bias. However, even if Figure \\ref{fig;mstar_sfr} is limited to a region with SFRs higher than the estimated detection limit, a mild correlation between stellar mass and SFR can be seen. If there are less massive galaxies with lower SFRs, as found in \\citet{dad07}, the correlation would be steeper.   \\citet{erb06b} have also reported a similar correlation between SFR and stellar mass for UV-selected $z \\sim 2$ galaxies, but their correlation is shifted to much lower SFR than ours. This is mainly due to differences in the assumed IMF as well as dust correction, as discussed in section \\ref{sec;sfruv}. In \\citet{erb06b}, a Chabrier IMF is used to estimate stellar masses and SFRs. SFRs derived with a Chabrier IMF are smaller than those with a Salpeter IMF by a factor of 1.8. These differences make our SFR 7.9 times larger than that of \\citet{erb06b}. After correcting these systematic differences, the comparison of the results between sBzKs and UV-selected galaxies indicates that the former tend to have slightly higher SFRs than UV-selected galaxies with the similar stellar masses. This reflects the intrinsic difference between the the dust amount in the two populations--only sBzKs with higher SFRs may be easily detected, due to the larger dust extinction than in UV-selected galaxies.  The positive correlation between SFR and stellar mass may be attributed to a larger amount of gas in more massive galaxies, as suggested by \\citet{dad07}.  \\begin{figure}[tb] \\begin{center} \\plotone{f9.eps} \\caption{The relation between ${\\rm SFR}_{\\Ha}$ and stellar mass. Stellar masses are derived from SED fitting to $BVRi'z'K$ photometry. The dotted line shows the limiting SFR of $87 \\Msun {\\rm yr^{-1}}$ estimated from the lowest detected flux on the assumption that the galaxy is at $<z>=1.7$ and suffer a typical color excess of 0.38. } \\label{fig;mstar_sfr} \\end{center} \\end{figure}  \\subsubsection{Specific SFR} \\label{sec;ssfr} We derived specific SFRs (SSFRs) from ${\\rm SFR}_{\\Ha}$s and stellar Masses.  SSFR is an indicator of the time scale of star formation activity. Figure \\ref{fig;mstar_ssfr} shows SSFR as a function of stellar mass. A negative relation is seen between stellar mass and SSFR, which is consistent with previous results for $z \\sim 2$ galaxies \\citep{red06a,erb06b}. This negative relation implies that a larger growth of stellar mass occurs in less massive galaxies. In comparison with \\citet{erb06b}, it should be noted that we apply a different IMF and dust correction as described in the previous section. In Figure \\ref{fig;mstar_ssfr}, the results of \\citet{erb06b} are converted into those with the same manner as ours.  We must consider the effect of flux-limited bias on this correlation, as discussed in the previous section. However, it seems that there is a genuine negative correlation even if the range of validity is limited to stellar masses larger than $1\\times10^{10}\\Msun$. The steep relation between stellar mass and SSFR results from the shallow correlation between stellar mass and SFR, which shows that the increase in SFR is not as rapid as that in stellar mass. If we miss galaxies with low SFRs of $\\lesssim87\\Msun {\\rm yr^{-1}}$ as discussed in section \\ref{sec;RtoMstar}, that is, the correlation between stellar mass and SFR is steeper, the relation with SSFR would be shallower than that in Figure \\ref{fig;mstar_ssfr}.  \\begin{figure}[tb] \\begin{center} \\plotone{f10.eps} \\caption{The relation between the specific ${\\rm SFR}_{\\Ha}$ and the stellar mass. Filled circles are our results, open triangles show those of \\citet{red06a} for sBzKs in their UV-selected sample, and crosses are those of \\citet{erb06b} for UV-selected sample. All stellar masses and SSFRs are converted into those derived with Salpeter IMF and the same dust correction as we apply. The dotted line shows the track of galaxies with a constant SFR of 87$\\Msun {\\rm yr^{-1}}$. } \\label{fig;mstar_ssfr} \\end{center} \\end{figure}   \\subsection{Equivalent Width} \\label{sec;ew} The rest-frame equivalent widths of $\\Ha$ are derived from Gaussian fitting. Both $\\Ha$ and continuum fluxes have been corrected for dust extinction. As discussed in \\citet{erb06b}, one can infer the star formation history of galaxies from the equivalent width. While the flux of $\\Ha$ reflects the current star forming activity, the continuum flux emitted from stars reflects the past star formation. That is to say, the equivalent width represents how actively a galaxy is forming stars compared with the past. Thus, the equivalent width is a similar indicator to SSFR. Figure \\ref{fig;ew_sfr} shows the equivalent width as a function of the SSFR, indicating a mild correlation. This supports the result of the previous section. Less massive galaxies are more actively forming stars, which results in a larger growth of stellar mass.   \\begin{figure}[tb] \\begin{center} \\plotone{f11.eps} \\caption{The relation between specific star formation rate and $\\Ha$ equivalent width in rest frame. Rest-frame equivalent widths of $\\Ha$ are derived from line and continuum fluxes after correction for dust extinction. } \\label{fig;ew_sfr} \\end{center} \\end{figure}   \\subsection{Metallicity} \\label{sec;metallicity} Average metallicity is one of the most important properties of a galaxy. Since gas and stellar metallicities are related to the past star formation, they are crucial for understanding the star formation history and galaxy evolution. It is well-known that metallicity is correlated with stellar mass in the local universe \\citep[e.g.,][]{tre04}. \\citet{erb06a} have reported a similar correlation to the local one for UV-selected galaxies at $z \\sim 2$.  However, their mass-metallicity ($M$-$Z$) relation at $z \\sim 2$ is shifted to lower metallicity by $\\sim 0.3$ dex. This $M$-$Z$ relation at $z \\sim 2$ obtained by \\citet{erb06a} is widely used as representative of $z \\sim 2$ star forming galaxies in comparison with both observational and theoretical studies \\citep{liu08,mai08,bro07,der07,fin08}. However, because the UV selection misses star forming galaxies with a large dust extinction, we cannot be confident that the $M$-$Z$ relation of \\citet{erb06a} applies to all star-forming galaxies at $z \\sim 2$. Here, we determine a $M$-$Z$ relation for sBzKs. Our sBzK sample, which approximately corresponds to a stellar mass-limited sample, should be more suitable for deriving the $M$-$Z$ relation at $z \\sim 2$.  It is well known that electron temperature reflects gas metallicity \\citep{kew02,kob04,erb06a}. Because auroral lines, which are used to derive electron temperature, are weak, and observed in only galaxies with low metallicities, it is very difficult even in local galaxies to estimate the metallicity of galaxies from them. We can only use the ratio of strong emission lines which are emitted from different ionization levels to derive metallicity for high-$z$ galaxies. At present, there are various metallicity diagnostics, such as the $R_{23}$ and N2 methods, to estimate gas-phase oxygen abundance, 12+log(O/H). One problem with these diagnostics is a disagreement between the resulting oxygen abundances inferred from the different diagnostics. Even if the same diagnostic is used, some different calibrations give a large difference in the resulting metallicities \\citep{kew02,kob04}. The same diagnostics and calibration should be used to properly compare with other results.  We derive 12+log(O/H) using the N2 index method proposed by \\citet{sto94}, where N2 is defined as the flux ratio of [\\ion{N}{2}]($\\lambda6584$)/$\\Ha$; \\begin{equation} {\\rm N2} \\equiv \\log\\frac{[{\\rm NII}](\\lambda6584)}{\\Ha}. \\end{equation} The gas-phase oxygen abundances are derived with the following relation given in \\citet{pet04}; \\begin{equation} \\begin{array}{l} 12+\\log({\\rm O/H}) = 8.90 + 0.57\\times{\\rm N2}. \\end{array} \\end{equation} The reasons we use this relation are that it requires only two lines of $\\Ha$ and [\\ion{N}{2}], and the result can be compared directly with that of \\citet{erb06a}, where the same relation was used to derive the gas abundance. Also, it should be noted that the lines of $\\Ha$ and [\\ion{N}{2}] are close enough that the difference in dust extinction is negligible. The $1\\sigma$ dispersion in the abundance derived from the relation is 0.18 \\citep{pet04}.  Figure \\ref{fig;mstar_metal} shows the resulting relation between stellar mass and metallicity. The broken line shows the abundance for galaxies with the flux ratio of [\\ion{N}{2}]/$\\Ha$=0.63, where N2 saturates \\citep{kew02}, suggesting that the linear relation between the oxygen abundance and N2 index is acceptable up to the abundance of $\\sim8.79$. As discussed in section \\ref{sec;agn}, the three sBzKs with a large abundance, ID21193, ID22424, and ID31340, may harbor AGNs, because they have flux ratios of [\\ion{N}{2}]/$\\Ha$ larger than 0.63. However, their SEDs suggest that their continuum is dominated by stellar emission, and that the AGN contribution is small. Even if the three AGN candidates are excluded from our sample, our conclusions in the following discussion are unchanged.  We derive upper limits of metallicity for sBzKs without a secure detection of [\\ion{N}{2}], assuming that the peak flux is 3 times sky noise, and that the line width is the same as that of $\\Ha$. These galaxies are shown by crosses in the figure. Also, we estimate mean metallicity of the five sBzKs without [\\ion{N}{2}] detection. We derive it from the stacked spectrum using the equations (3) and (4), finding that the mean abundance is 8.55. This is shown by the double circle in the figure, whose stellar mass is the average of those of the five sBzKs.   \\begin{figure*}[htb] \\begin{center} \\plotone{f12.eps} \\caption{Stellar mass vs. metallicity relation. Circles and crosses show our results. Filled circles show those of sBzKs with the detection of both $\\Ha$ and [\\ion{N}{2}]. Crosses show upper limits of sBzKs without the detection of [\\ion{N}{2}], which are derived using the $3\\sigma$ upper limit on [\\ion{N}{2}] flux. The double circle shows the mean abundance of the five sBzKs without a detection of [\\ion{N}{2}] against the mean stellar mass. Open squares show UV-selected $z \\sim 2$ galaxies, and open triangles are those of the local $\\sim 53,000$ SDSS galaxies \\citep{erb06a}. All abundances are estimated using the same N2 index method. The dotted line is the solar oxygen abundance of 12+log(O/H)=8.66 \\citep{asp04}. The broken line shows the abundance for galaxies with the flux ratio of [\\ion{N}{2}]/$\\Ha$=0.63. Stellar masses of the \\citet{erb06a} are converted into those derived with Salpeter IMF. } \\label{fig;mstar_metal} \\end{center} \\end{figure*}   Figure \\ref{fig;mstar_metal} shows that more massive sBzKs tend to be more metal-rich, a trend qualitatively consistent with that of \\citet{erb06a}. However, the metallicities of sBzKs with [\\ion{N}{2}] detection are higher by $\\sim 0.2$ dex than those of UV-selected galaxies with the same stellar mass at the similar redshift. Considering the uncertainties, a significant shift exists at the 2.7$\\sigma$ level between the two $M$-$Z$ relations. If three AGN candidates are excluded from the sample, the confidence level in the shift slightly decreases, to the 1.8$\\sigma$ level.  We again emphasize that the N2 flux ratios are converted into the oxygen abundances using the same relation as \\citet{erb06a}. We make sure that the N2 indices are larger than those of \\citet{erb06a} for galaxies with the given stellar masses, suggesting that the cause of the difference in $M$-$Z$ relation is not a systematic difference in our abundance estimate methodology.  \\citet{pet04} have found a relation between gas oxygen abundance and O3N2 index, which is defined as ${\\rm O3N2}$ $\\equiv$ $\\log$[([\\ion{O}{3}]($\\lambda5007$)/$\\Hb$) / ([\\ion{N}{2}]($\\lambda6584$)/$\\Ha$)]. We check the abundance with the O3N2 diagnostics for three sBzKs with the necessary lines observed. The abundances from the O3N2 ratios are lower than those from the N2 ratios by $\\sim 0.13$ dex. This is consistent with \\citet{erb06a}, who found that the abundances from O3N2 are 0.17 dex lower on average than those from N2. In addition, it is not likely that the systematic difference in estimates of stellar masses accounts for the difference in the $M$-$Z$ relation. The stellar masses of \\citet{erb06a} are converted into those derived with Salpeter IMF. As discussed in section \\ref{sec;sedfit}, the systematic difference of stellar masses should be less than a factor of 2. This indicates that there is no way that it is the cause of the difference in $M$-$Z$ relation.  We consider the possibility that our $M$-$Z$ relation suffers from a selection bias that only sBzKs with high metallicities are plotted, since metallicities are properly measured only for sBzKs with both $\\Ha$ and [\\ion{N}{2}] lines detected. The stacking analysis for sBzKs without [\\ion{N}{2}] detection suggests that there are also sBzKs with lower metallicities (Figure \\ref{fig;mstar_metal}), supporting the idea that selection bias cannot be completely ruled out. This may imply that the mean metallicities of sBzKs with given stellar masses are slightly lower than those plotted in Figure \\ref{fig;mstar_metal}. However, the mean metallicity of only sBzKs without [\\ion{N}{2}] detection is similar to that of UV-selected galaxies, implying that sBzKs would still be more metal rich on average than UV-selected galaxies. This fact suggests that it is likely that the UV selection misses star forming galaxies at $z\\sim2$ with comparatively higher metallicities, and that metallicities of star forming galaxies at $z \\sim 2$ have a large dispersion.   Therefore, we conclude that the oxygen abundance of sBzKs is greater than that of UV-selected $z \\sim 2$ galaxies. As described in section \\ref{sec;sfr}, $K$-selected star forming galaxies at $z\\sim2$ have slightly higher SFRs but similar to UV-selected galaxies, if the same correction for dust extinction is applied, while their metallicities are significantly different. This fact may imply that the two star forming galaxy populations at $z \\sim 2$ pass through different star forming histories before this epoch.    This difference in the $M$-$Z$ relation between sBzKs and UV-selected galaxies may be related to the difference in color of the two populations. Figure \\ref{fig;k_rk} shows $R-K$ color as a function of $K$ for our sBzK sample and the UV-selected galaxy sample of \\citet{erb06a}. This figure shows that the sBzK spectroscopic sample galaxies are redder than the UV-selected galaxies. Indeed, the sBzK ID31569, which is as blue as UV-selected galaxies, is located on the $M$-$Z$ relation of \\citet{erb06a}. The difference in $R-K$ color may imply that galaxies with different colors have different ages and/or dust contents, since the 4000\\AA \\ break falls between the $R$ and $K$-bands. This fact indicates that star forming galaxies at $z \\sim 2$ with redder colors have larger metal abundances, and that the $M$-$Z$ relation of \\citet{erb06a} may not represent the entire population of $z \\sim 2$ galaxies.   Another possibility is that the difference in redshifts of galaxies between the two samples accounts for the difference in the $M$-$Z$ relation. While the median redshift of galaxies in our sample is 1.7, the mean redshift is $\\left<z\\right>=2.2$ in the \\citet{erb06a} sample. Because the galaxies in our sample are on average at a lower redshift, the fact that our sBzKs have higher metallicities than UV-selected galaxies of \\citet{erb06a} may plausibly be attributed to cosmic evolution. However, this interpretation would require a surprisingly rapid evolution of the $M$-$Z$ relation over 800 Myr.  We plot results of star forming galaxies at $z \\sim 0$ from SDSS data in Figure \\ref{fig;mstar_metal} \\citep{tre04,erb06a}. It should be noted that we cannot directly discuss redshift evolution of the $M$-$Z$ relation, since these SDSS galaxies are not necessarily descendants of sBzKs. Our $M$-$Z$ relation seems to be similar to that of SDSS galaxies \\citep{tre04}, but we may not necessarily claim that the metallicities of sBzKs are the same as those of the local SDSS galaxies. There is the possibility that the true metallicities of the SDSS galaxies are higher than those derived by the N2 method, if the saturation of the N2 index at high metallicity is taken into account. Thus, the oxygen abundances of sBzKs can be as high as those of the SDSS galaxies. Then, our result that the metallicities of sBzKs are close to solar abundance is consistent with the assumption of the SEDs with solar abundance in SED fitting.  It is also found that there are $M$-$Z$ correlations at other redshifts. \\citet{sav05} have reported the relation for NIR-selected galaxies at $z\\sim0.7$, finding that the metallicities of galaxies with a given stellar mass have a dispersion of $\\sim0.2$ dex around the mean abundances. The correlation with large scatter is similar to what we found for galaxies at $z\\sim2$. \\citet{mai08} have found the relation shifted to much lower metallicity for 9 Lyman Break Galaxies at $z\\sim3.5$, suggesting that the metallicity of galaxies decreases with increasing redshift. However, since metallicity measurements for high-$z$ galaxies are very limited, especially at $z>2$, the $M$-$Z$ relations of high-$z$ galaxies obtained so far may have significant statistical and systematic uncertainties. Therefore, in future, it is essential to derive more reliable $M$-$Z$ relations at high redshifts from a large, unbiased sample.  \\begin{figure}[tb] \\begin{center} \\plotone{f13.eps} \\caption{The color magnitude diagram of $R-K$ vs. $K$ for our sBzK spectroscopic sample (filled circles), the whole sBzK sample in the SDF (dots) and the NIRSPEC sample of \\citet{erb06c} (open triangles). } \\label{fig;k_rk} \\end{center} \\end{figure}"
    },
    "0809/0809.0987_arXiv.txt": {
        "abstract": "The atmosphere of the Sun is characterized by a complex interplay of competing physical processes: convection, radiation, conduction, and magnetic fields. The most obvious imprint of the solar convection and its overshooting in the low atmosphere is the granulation pattern. Beside this dominating scale there is a more or less smooth distribution of spatial scales, both towards smaller and larger scales, making the Sun essentially a multi-scale object.  Convection and overshooting give the photosphere its face but also act as drivers for the layers above, namely the chromosphere and corona.  The magnetic field configuration effectively couples the atmospheric layers on a multitude of spatial scales, for instance in the form of loops that are anchored in the convection zone and continue through the atmosphere up into the chromosphere and corona. The magnetic field is also an important structuring agent for the small, granulation-size scales, although (hydrodynamic) shock waves also play an important role --- especially in the internetwork atmosphere where mostly weak fields prevail.  Based on recent results from observations and numerical simulations, we attempt to present a comprehensive picture of the atmosphere of the quiet Sun as a highly intermittent and dynamic system. ",
        "introduction": "Observations of the solar atmosphere reveal a wealth of different phenomena, which occur over an extended range of different temporal and spatial scales. This is not surprising, considering the fact that already basic parameters such as gas density and temperature span many orders of magnitude, from the convection zone below the photosphere to the corona.  At a first look, it may thus appear rather hopeless to construct an overall picture that can account for all the phenomena.  At a closer look, however, many connections between apparently independent phenomena can be found, ultimately implying a multitude of couplings through the atmosphere.  In addition, there seems to be a hierarchical arrangement of approximately selfsimilar convective motions, with the granulation pattern embedded in increasingly larger meso- and supergranulation patterns.  The key to a comprehensive picture of the solar atmosphere thus lies in relaxing too strict and oversimplified concepts, even when they are didactically nicer than the reality.  The solar atmosphere should not be seen as a static stack of layers but rather as intermittent domains that are dynamically coupled together.  One example is magnetic flux structures (or ``flux tubes'') fanning out with a wine-glass geometry.  Such regular building blocks put certain constraints on the implied atmospheric structure, which can make it difficult to fit in other observational findings.  Accepting that magnetic field structures are far less regular offers room for a more generally valid comprehensive picture. This trend became more and more obvious during the recent years, both from the observational and theoretical side \\citep[see, e.g.,][and many more]{2007ASPC..368...49C, 2006ASPC..354..331G, 2007ASPC..369..193H, 2006ASPC..354..259J, 2007ASPC..368...27R, 2007msfa.conf..321S}.  The advantages of a relaxed picture can be seen from the example of the quiet Sun chromosphere above internetwork regions, which in itself is a complex and intriguing phenomenon \\citep[see, e.g.,][ and many more]{2006ASPC..354..259J,2006ASPC..354..276R, 2004Natur.430..536D, 1999ApJ...517.1013L}.  Despite tremendous progress, there are still many open questions concerning its structure, dynamics and energy balance.  Recent observations now prove -- beyond any doubt -- the chromosphere to be a highly dynamic and intermittent layer.  The internetwork chromosphere is the product of a dynamic interplay of shock waves and magnetic fields.  This picture, which was already suggested by many earlier investigations, offers a key to resolve some apparent contradictions that lead to much confusion in the past.  A prominent example concerns the observation of carbon monoxide \\citep[see, e.g.,][]{ayres02}, which now can be explained as an integral part of a dynamic and intermittent atmosphere \\citep{2007A&A...462L..31W,2006ASPC..354..301W,2005A&A...438.1043W}.  And still the chromosphere cannot be investigated without also taking into account the layers above and below.  The shock waves, which are so essential at least for the lowest, weak-field parts of the chromosphere, are generated in the layers below, while  significant amounts of mass and energy are exchanged between the chromosphere and the corona above. Obviously, the whole atmosphere must be seen as an integral phenomenon.  In the following sections, we report on a selection of results from observations and numerical simulations, which will help us put together an updated, revised view of the structure of the quiet Sun atmosphere. ",
        "conclusions": "The solar atmosphere is a very dynamic and inhomogeneous multi-scale system. Its individual components are coupled; some of them even show a kind of hierarchical self-similarity. Examples are the observational imprints of sub-surface convection, with a continuous spectrum of scales from below granulation scales to above supergranulation scales, and magnetic fields, which also exhibit similar features over a large range of spatial scales.  Despite great progress on the theoretical and observational sides, which go hand in hand, we are still missing an ultimate, comprehensive picture of the quiet Sun atmosphere. But at least we can now see what is needed for a corresponding numerical simulation. First, the computational domain should be large enough to encompass a few supergranulation cells while the spatial resolution must still be high enough to capture important processes that occur on scales smaller than granulation. The vertical couplings make it necessary to consider an extensive height range. The corona and chromosphere can only be treated realistically when including the important driving motions in the layers below, i.e. in the photosphere and (at least) the upper part of the convection zone. While many simplifying assumptions can be made for the lower parts of such a model, the layers above the (middle) photosphere require a numerically complicated and thus computationally expensive non-equilibrium modeling approach, e.g., a realistic treatment of hydrogen ionization etc.. The production of such a comprehensive model -- and analogous models for, e.g., active regions -- is thus very involved and can be regarded as one of the current challenges in (computational) solar physics.  On the observational side, we must continue to push forward the instrumental possibilities towards higher resolution in the spatial, temporal, and spectral domains -- all at the same time. In addition, the (further) development and exploitation of advanced diagnostics is needed to derive a seamless tomography of the atmosphere as an integral phenomenon.  Until we succeed to reach these ambitious goals, we are left with a number of open questions. Of particular interest for the quiet Sun are, amongst others:  \\begin{itemize} \\item How does the weak internetwork field connect with the stronger network field? What does the magnetic field look like just below the ``canopy''? And can we talk about a ``canopy'' even in a wider sense after all? \\item How do the propagating fluctospheric shock waves interact with the stronger (``canopy'') field? Is it mostly a ``passive'' refraction / reflection at the ``boundaries'' of flux concentrations \\citep[e.g.][]{2002ApJ...564..508R, 2007AN....328..323S} % or also ``active'' distortion/displacement/compression of magnetic field? How and where does mode conversion take exactly place under realistic conditions? \\item The question of the coupling between the atmospheric layers is closely connected to the heating mechanism question or, better said, to the question of the atmospheric energy balance, not only for the Sun but also for other stars. Amongst other things, the ongoing controversy concerning the heating mechanism of the quiet Sun chromosphere \\citep[e.g.,][]{2005Natur.435..919F} has important implications for stellar activity in general. \\end{itemize}   \\textbf{\\ \\\\ \\noindent{}Acknowledgments:} We would like to thank O.~Steiner, K.~Schrijver, B.~de~Pontieu, M.~Carlsson, V.~Hansteen, R.~Rutten, and H.~Peter for helpful discussions. SWB was supported with a Marie Curie Intra-European Fellowship of the European Commission (6th Framework Programme, FP6-2005-Mobility-5, Proposal No. 042049). {\\AA}N acknowledges support from the Danish Natural Science Research Council and from the Danish Center for Scientific Computing."
    },
    "0809/0809.5046_arXiv.txt": {
        "abstract": "A supermassive black hole ejected from the center of a galaxy by gravitational wave recoil carries a retinue of bound stars -- a ``hypercompact stellar system'' (\\rgc). The numbers and properties of \\rgcs\\ contain information about  the merger histories of galaxies, the late evolution of binary black holes, and the distribution of gravitational-wave kicks. We relate  the structural properties (size, mass, density profile) of \\rgcs\\ to the properties of their host galaxies and to the size of the kick, in two regimes: collisional ($\\mh\\lap 10^7\\msun$), i.e. short nuclear relaxation times; and collisionless ($\\mh\\gap 10^7\\msun$), i.e. long nuclear relaxtion times. \\rgcs\\ are expected to be similar in size and luminosity to globular clusters but in extreme cases (large galaxies, kicks just above escape velocity) their stellar mass can approach that of ultra-compact dwarf galaxies. However they differ from all other classes of compact stellar system in having very high internal velocities. We show that the kick velocity is encoded in the velocity dispersion of the bound stars. Given a large enough sample of \\rgcs, the distribution of gravitational-wave kicks can therefore be empirically determined. We combine a hierarchical merger algorithm with stellar population models to compute  the rate of production of \\rgcs\\ over time and the probability of observing \\rgcs\\ in the local universe as a function of their apparent magnitude, color, size and velocity dispersion, under two different assumptions about the star formation history prior to the kick. We predict that $\\sim 10^2$ \\rgcs\\ should be detectable within 2 Mpc of the center of the Virgo cluster and that many of these should be bright enough that their kick velocities (i.e. velocity dispersions) could be measured with reasonable exposure times. We discuss other strategies for detecting \\rgcs\\ and speculate on some exotic manifestations. ",
        "introduction": "A natural place to search for supermassive black holes (SMBHs) is at the centers of galaxies, where they presumably are born and spend most of their lives. But it has become increasingly clear that a SMBH can be violently separated from its birthplace as a result of linear momentum imparted by gravitational waves during strong-field interactions with other SMBHs \\citep{Peres:62,Bekenstein:73, Redmount:89}. The largest net recoils are produced from configurations that bring the two holes close enough together to coalesce. Kick velocities following coalescence can be as high as $\\sim 200$ \\kms\\ in the case of nonspinning holes \\citep{Gonz:Nonspin:07,Sopuerta:07}; $\\sim 4000$ \\kms\\ for maximally spinning, equal mass BHs on initially circular orbits \\citep{Campanelli:07,Gonz:Spin:07,Herrmann:07,Pollney:07, Tichy:07,Brugmann:08,Dain:08,Baker:08}; and even higher, $\\sim 10,000$ \\kms, for black holes that approach on nearly-unbound orbits \\citep{Healy:08}. Since escape velocities from the centers of even the largest galaxies are $\\lap 2000$ \\kms\\ \\citep{Merritt:04}, it follows that the kicks can in principle remove SMBHs completely from their host galaxies. While such extreme events may be relatively rare \\citep[e.g.][]{SB:07,Schnittman:07}, recoils large enough to displace SMBHs at least temporarily from galaxy cores -- to distances of several hundred to a few thousand parsecs -- may be much more common \\citep{Merritt:04,MQ:04,GM:08,KM:08b}.  Komossa et al. (2008) reported the detection of a recoil candidate. This quasar exhibits a kinematically offset broad-line region with a velocity of 2650 km s$^{-1}$ , and very narrow, restframe, high-excitation emission lines which lack the usual ionization stratification  -- two key signatures of kicks. In addition to spectroscopic signatures \\citep{Naked,Bonning:07}, recoiling SMBHs could be detected by their soft X-ray, UV and IR flaring \\citep{Shields:08,Lippai:08,SK:08} resulting from shocks in the accretion disk surrounding the coalesced SMBH. Detection of recoiling SMBHs  in this way is contingent on the presence of gas. But only a small fraction of {\\it nuclear} SMBHs exhibit signatures associated with gas accretion, and a SMBH that has been displaced from the center of its galaxy will only shine as a quasar until its bound gas has been used up \\citep{Loeb:07}. The prospect that the SMBH will encounter and capture significant amounts of gas on its way out are small \\citep{Kapoor:76}.  A SMBH ejected from the center of a galaxy will always carry with it a retinue of bound stars. The stars can reveal themselves via tidal disruption flares or via accretion of gas from stellar winds onto the SMBH \\citep[][hereafter Paper I]{KM:08a}. The cluster of stars is itself directly observable, and that is what we discuss in the current work. The linear extent of such a cluster is fixed by the magnitude $\\vk$ of the kick velocity and by the mass of the SMBH: \\begin{subequations} \\begin{eqnarray} \\rk &\\equiv& G\\mh/\\vk^2 \\\\ &\\approx& 0.043\\ {\\rm pc} \\left({\\mh\\over 10^7\\msun}\\right) \\left({\\vk\\over 10^3\\ {\\rm km\\ s}^{-1}}\\right)^{-2}\\, . \\end{eqnarray} \\label{eq:rk} \\end{subequations} Reasonable assumptions about the density of stars around the binary SMBH prior to the kick (Paper I) then imply a total luminosity of the bound population comparable to that of a globular star cluster.  In this paper we discuss the properties of these ``hyper-compact stellar systems'' (\\rgcs) and their relation to host galaxy properties. Our emphasis is on the prospects for detecting such objects in the nearby universe at optical wavelengths, and so we focus on the properties that would distinguish \\rgcs~from other stellar systems of comparable size or luminosity. As noted in Paper I, a key signature is their high internal velocity dispersion: because the gravitational force that binds the cluster comes predominantly from the SMBH, of mass $10^6\\lap M_{\\rm BH}/\\msun\\lap 10^9$, stellar velocities will be much higher than in ordinary stellar systems of comparable luminosity. Other signatures include the small sizes of \\rgcs\\ (unfortunately, too small to be resolved except for the most nearby objects); their high space velocities (due to the kick); and their broad-band colors, which should resemble more closely the colors of galactic nuclei rather than the colors of uniformly old and metal-poor systems like globular clusters.  As we discuss in more detail below (\\S2), a remarkable property of \\rgcs\\ is that they encode, via their internal kinematics, the velocity of the kick that removed them from their host galaxy. A measurement of the velocity dispersion of the stars in a \\rgc\\ is tantamount to a measurement of the amplitude of the kick -- independent of how long ago the kick occurred; the black hole mass; and the space velocity of the \\rgc\\ at the moment of observation. This property of \\rgcs\\ opens the door to an empirical determination of the distribution of gravitational-wave kicks.  The outline of the paper is as follows. \\S 2 derives the relations between the structural parameters of \\rgcs -- mass, radius, and internal velocity dispersion -- given assumed values for the slope and density normalization of the stellar population around the SMBH just before the kick. In \\S3, models for the evolution of binary SMBHs are reviewed and their implications for the pre-kick distribution of stars are described. These results, combined with the relations derived in \\S2, allow us to relate the structural parameters of \\rgcs~to the global properties of the galaxies from which they were ejected. \\S4 discusses the effect of post-kick dynamical evolution of the \\rgcs~on their observable properties. Stellar evolutionary models are used to predict the luminosities and colors of \\rgcs~and their post-kick evolution in \\S5, and in \\S6, the evolutionary models are combined with models of hierarchical merging to estimate the number of \\rgcs~to be expected per unit volume in the local universe as a function of their observable properties. \\S7 discusses search strategies for \\rgcs\\ and various other observable signatures that might be uniquely associated with them. In \\S8 we briefly discuss the inverse problem of reconstructing the distribution of recoil velocities from an observed sample of \\rgcs. \\S9 sums up and suggests topics for further investigation. ",
        "conclusions": "1. Supermassive black holes (SMBHs) kicked out from the centers of galaxies by gravitational wave recoil are accompanied by a cluster of bound stars with mass $\\sim 10^{-2}$ times the black hole mass or less, and  radius $\\sim 10^1$ pc or less -- a ``hyper-compact stellar system'' (\\rgc).  2. \\rgcs\\ have density profiles that can uniquely be calculated given the kick velocity and given the stellar distribution prior to the kick. The density at large distances from the SMBH falls off as $\\sim r^{-4}$.  3. Internal (rms) velocities of \\rgcs\\ are very high, $\\sim 10^2-10^3$ \\kms, and comparable to their kick velocities. Their overall velocity distributions are extremely non-Gaussian.  4. HCSSs could be distinguished photometrically from foreground red giants, based on their bluer (lower) values of $B-V$ for a given $B-K$. They also should appear redder than low-metallicity GCs with comparable ages $\\gtrsim 1$ Gyr.  5. With a simplified cosmological merger model, we are able to estimate expected number counts and luminosity distributions of HCSSs in the local universe. Detection of perhaps $10^2$ HCSSs should be possible in the Virgo cluster alone, although only a few may be bright enough to allow high S/N spectroscopy and provide solid confirmation of their extreme velocity dispersions.  6. Some \\rgcs\\ may already exist in survey data of compact stellar sysetms in the Fornax, Virgo and Coma galaxy clusters.  7. Because the kick velocity of a \\rgc\\ is related in a simple way to its measured velocity dispersion, the distribution of gravitational wave kicks can be empirically determined from a sufficiently large sample of \\rgcs.  Paper I \\citep{KM:08a} first derived the basic properties of \\rgcs\\, including their compactness and high internal velocity dispersions, and presented a possible route to detection via off-nuclear tidal disruption flares. The current paper derives the intrinsic properties of \\rgcs\\ in a more complete and general way and relates those properties to the properties of the host galaxy. Together, these two papers complement the growing number of recent papers that discuss gas-related signatures of gravitational wave recoil. While this paper was being submitted, we learned of a related work by O'Leary \\& Loeb (2009) which argues that many thousands of low-mass \\rgcs\\ should be present in the Milky Way halo.   \\bigskip\\bigskip"
    },
    "0809/0809.2982_arXiv.txt": {
        "abstract": "Using the gradient expansion method of Rigopoulos, Shellard and van Tent which treats cosmological perturbations as gradients on top of a homogeneous and isotropic FRW background, we study the production of nongaussianities in the roulette model of inflation. % Investigating a number of trajectories within this two-field model of inflation, we find that while the superhorizon influence of the isocurvature modes on the curvature bispectrum produces nonzero contribution to $\\fNL$, the effect is negligible next to the standard inflationary prediction $|\\fNL| \\sim n_s - 1$. % This is the case in both the squeezed and equilateral configurations of the bispectrum, although the former is slightly larger in the trajectories under consideration. ",
        "introduction": "Recent advances in vacuum stabilization in string theory, most notably the KKLT \\cite{KKLT} and the large volume  \\cite{largevol} varieties, have put stringy realizations of cosmic inflation on a much firmer footing.  Only when all moduli have been stabilized does it make sense to study which of them might be candidates for inflation. String theory provides an abundance of moduli fields, for example the size and shape of the extra dimensions, and thus many potential opportunities for inflation.  These include single-field models, such as ref.\\ \\cite{kahler}, and multiple-field models  such as the racetrack model \\cite{racetrack}, the N-flation scenario \\cite{nflation}, or the swiss-cheese scenario \\cite{swiss}.  Multiple field inflation has the advantage of  more easily providing some of the nonstandard observational features that one would like for helping to discriminate between theories. The extra  degrees of freedom allow the production of isocurvature modes (perturbations transverse to the classical trajectory) and it has become apparent \\cite{largenongauss} that these may give rise to large nongaussianities in the cosmic microwave background (CMB) temperature fluctuations.  One such model is the roulette scenario \\cite{roulette}, in which the Calabi-Yau manifold (the compactified extra dimensions of string theory) relaxes from an initial excited state towards a  minimum of its potential. The large volume compactification that is used ensures that this minimum exists for large ranges of the microscopic parameters. In addition, unlike KKLT, it does not require  tuning the constant term in the superpotential to very small values, and it gives a natural expansion parameter, the inverse volume $1/\\cyvol$, providing a controlled $\\alpha'$ expansion. In this particular model, the last four-cycle and corresponding axionic partner to relax act as slow rolling scalar fields which drive the final stage of inflation.  Specializing to a specific set of microscopic parameters, we explored the various trajectories of inflation available within the context of this two-field model. The prediction of observable parameter values were in accordance with known results from CMB and large scale structure survey data. The influence of isocurvature perturbations, which seed inhomogeneieties between the species of fields driving inflation, was furthermore found to be quite important.  In addition, recent claims of detection of nongaussianity in the WMAP CMB data, specifically a nonvanishing nonlinearity parameter $f_{NL}$ \\cite{yadav} have sparked a large interest in deviations from gaussianity in the spectrum of primordial fluctuations as an additional observable that must be predicted by a successful theory of the early universe. Given that precision measurements of $f_{NL}$ from experiments such as Planck will soon be available, it is all the more important that the mechanisms governing the production and evolution of primordial nongaussianities be well understood.  We will first provide an overview of the roulette model, followed by a discussion of primordial nongaussianities from inflation. In Section \\ref{sec:results} we will give the main results of our paper, which are the numerical caluclations of nongaussianities in the curvature perturbation. The important aspects of our calculations, which follow \\cite{quantitative}, are presented in Appendix A. ",
        "conclusions": "We have studied a model of \\kahler moduli inflation built from a realistic construction of Type IIB string theory, using the formalism developed in \\cite{nonlin, largenongauss,quantitative}.  We confirm previous results for the power spectrum and the superhorizon influence of isocurvature modes for typical inflationary trajectories. The main new undertaking is a search for nongaussian perturbations from highly curved trajectories.  Such deviations from gaussianity can originate from the superhorizon evolution of the second-order curvature perturbation and its interaction with isocurvature perturbations.  When the full spectrum is considered, roulette inflation predicts a smooth power spectrum with a slight red-tilt, in excellent agreement with estimates based on the latest WMAP data. In addition, as shown by \\cite{lalak}, the superhorizon influence of isocurvature modes can come to dominate the scalar curvature power spectrum via the relation \\eqref{gradcons}. This is the case for inflationary trajectories with large curvature, which can be pictured as a ``projection'' of the isocurvature perturbation modes onto the adiabatic direction. This contribution, which can account for over 90\\% of the curvature power spectrum, must therefore be considered in the context of these multiple-field inflationary models.  Concerning the issue of nongaussianity, we did not find examples in which $f_{NL}$, produced by the mechanism considered in section \\ref{sec:results}, could be large enough to  be observable by future missions; we obtained results which did not exceed the level of $f_{NL}$ expected in conventional single-field inflation models.  It was hoped that trajectories with  very sharp curving in field space would have yielded larger $f_{NL}$ values, but two considerations could make this difficult. First, if the trajectory relaxes to a straight, effectively single-field form before the end of inflation, $\\fNL$ damps to zero,  as illustrated in figure 1 of ref. \\cite{quantitative}. We found trajectories with moderate curving until the end of inflation, which prevented the total erasure of $f_{NL}$.   Multiple-field (``multi-brid'') hybrid inflation models such as discussed in refs.\\ \\cite{multibrid,Naruko:2008sq} use this effect to produce sizable nongaussianities. Second, large deviation from Gaussianity from this method is correlated with large isocurvature modes being projected onto the curvature direction. If one is able to find models with large $\\fNL$, it must be verified that the power-spectrum is not overly amplified and distorted by this effect.  This phenomenon is also seen to play an important role in nongaussianity generated by the waterfall fields in hybrid inflation models \\cite{hybrid}.  However, it is possible that a more complete search of the model's parameter space would reveal examples with larger values of $\\fNL$. In that case, running bispectra ({\\it i.e.}, which depend on the magnitude of the average scale of $k_i$, as opposed to the shape-dependence), such as we find in the roulette model, could provide an interesting discriminator between models.  Recent developments in the detection of nongaussianity via large-scale structure, which would probe $f_{NL}$ at smaller scales,  promise to give additional observational handles on such a dependence \\cite{LoVerde, LSS}.   It may be interesting to examine other such moduli inflation scenarios that arise once the assumptions of a strict hierarchy of scale are relaxed (for instance, the model of ref.\\ \\cite{vijay} in which the second dynamical field is the inverse overall volume, rather than the axionic partner). A more general study of two-field inflationary models with such an exponentially flat potential would furthermore reveal how generic the above-mentioned behavior is."
    },
    "0809/0809.0677_arXiv.txt": {
        "abstract": "Mirror matter is a promising self-collisional dark matter candidate. Here we study the evolution of thermodynamical quantities in the early Universe for temperatures below $\\sim 100$ ~MeV in presence of a hidden mirror sector with unbroken parity symmetry and with gravitational interactions only. This range of temperatures is interesting for primordial nucleosynthesis analyses, therefore we focus on the temporal evolution of number of degrees of freedom in both sectors. Numerically solving the equations, we obtain the interesting prediction that the effective number of extra-neutrino families raises for decreasing temperatures before and after Big Bang nucleosynthesis; this could help solving the discrepancy in this number computed at nucleosynthesis and cosmic microwave background formation epochs. ",
        "introduction": "In 50's Lee and Yang~\\cite{Lee:1956qn} suggested that the existence of a hidden sector of particles and interactions with exactly the same properties as our visible world (except for right-handed, instead of left-handed, weak interactions) could restore the total symmetry of physical laws. Many years later Foot at al.~\\cite{Foot:1991bp} proposed a model with exact parity (mirror) symmetry between the two sectors, that are described by the same Lagrangian and coupling constants, and consequently have the same microphysics. Since the only interaction linking the two worlds is gravity\\footnote{ There could be other interactions, as for example the kinetic mixing between ordinary and mirror photons~\\cite{Foot:2000vy}, but they are very low, and are neglected in the present study. }, mirror baryons constitute a dark matter candidate in a natural way, as they interact with mirror photons, but not with ordinary ones.  The phenomenology of mirror matter was studied in several papers (for an extended list see the bibliography of~Ref.~\\cite{Okun:2006eb}); in particular the implications for Big Bang nucleosynthesis~\\cite{Berezhiani:2000gw,HPLW2008}, primordial structure formation and cosmic microwave background~\\cite{Ciarcelluti:2003wm,Berezhiani:2003wj,Ciarcelluti:2004ij,Ciarcelluti:2004ik,Ciarcelluti:2004ip,HPLW2008}, large scale structure of the Universe~\\cite{Ciarcelluti:2003wm,Berezhiani:2003wj,Ciarcelluti:2004ij,Ciarcelluti:2004ip,HPLW2008,Ignatiev:2003js}, microlensing events (MACHOs)~\\cite{Blinnikov:1996fm,Foot:1999hm,Mohapatra:1999ih,Berezhiani:2005vv}, interpretation of DAMA experiment~\\cite{Foot:2008nw,Bernabei:2008yi} are related to this study.  If the mirror (M) sector exists, then the Universe along with the ordinary (O) particles should contain their mirror partners, but their densities are not the same in both sectors. Indeed, the BBN bound on the effective number of extra neutrino families implies that the M sector has a temperature lower than the O one, as naturally obtained in some inflationary models~\\cite{Berezhiani:1995am}. Then, two sectors have different initial conditions, they do not come into thermal equilibrium at later epoch and they evolve independently, separately conserving their entropies and maintaining approximately constant the ratio among their temperatures.  All the differences with respect to the O world can be described in terms of only two free parameters, defined as: \\begin{equation} \\label{mir-param} x \\equiv \\left( s' \\over s \\right)^{1/3} \\;; ~~ \\beta \\equiv \\Omega'_{\\rm b} / \\Omega_{\\rm b} \\;, \\end{equation} where $T$ ($T'$), $\\Omega_{b}$ ($\\Omega'_{b}$), and $s$ ($s'$) are respectively the O (M) photon temperature, cosmological baryon density, and entropy density.\\footnote{From now on all particles and parameters of the mirror sector will be marked by $'$ to distinguish them from the ones related to the observable or ordinary world. } The bounds on the mirror parameters are $ x < 0.7 $ and $ \\beta > 1 $, the first one coming from the previously mentioned BBN limit and the second one from the hypothesis that a relevant fraction of dark matter is made of M baryons. In general, during most of the history of the Universe, we can approximate \\begin{equation} \\label{xmir} x \\equiv \\left( s' \\over s \\right)^{1/3} = \\left[q'(T') \\over q(T) \\right]^{1/3}{T' \\over T} \\approx {T' \\over T} \\;, \\end{equation} where now $q(T)$ and $q'(T)$ are respectively the O and M entropic degrees of freedom, and the M photon temperature $T'(T)$ is a function of the O one. This approximation is only valid if the temperatures of two sectors are not too different, or alternatively if we are away enough from crucial epochs, like $e^+$-$e^-$ annihilation~\\cite{Berezhiani:2000gw}. But we are just investigating the range of temperatures interested by this phenomenon, and in fact one of the aims of this paper is to exactly compute the trends of these thermodynamical quantities when the previous approximation is no longer valid.  The present study is required for different reasons. First of all, we need a detailed study of thermodynamical evolution of the early Universe in order to obtain a reliable description of the thermal cosmic history in presence of M dark matter. Secondly, an accurate study of the trend of number of degrees of freedom is necessary for any future simulation of nucleosynthesis of primordial elements in both sectors, that may be compared with observations. In particular, it becomes even more important in view of the recent proposed interpretation of the DAMA/LIBRA experiment in terms of interactions with M heavy elements~\\cite{Foot:2008nw,Bernabei:2008yi}. Thirdly, claims for variations of the effective number of extra neutrino families computed at BBN ($\\sim 1$ MeV) and cosmic microwave background (CMB) formation ($\\lsim 1$ eV) epochs require investigations for possible variations of these numbers as consequences of physics beyond the Standard Model (see the recent Ref.~\\cite{Mangano:2006ur} and references therein for previous works).  The plan of the paper is as follows. In next section we introduce the thermodynamical equilibrium in the early mirror Universe and the equations that govern its evolution. In section \\ref{sec_Num_calc} we present the numerical solutions of these equations and discuss the results. Finally, our main conclusions are summarized in section \\ref{sec_conclusions}. ",
        "conclusions": "\\label{sec_conclusions}  In this paper we have investigated the consequences of the existence of a mirror sector of particles and interactions with exact parity symmetry on the thermodynamics of the early Universe for temperatures below $\\sim 100$ MeV. In particular we studied in detail the evolution of the entropic and energetic degrees of freedom in both sectors.  To perform these calculations we first found the system of equations governing the thermodynamical equilibria in the two sectors, and then we numerically solved them in order to obtain the mirror temperature $T'$ corresponding to any ordinary temperature $T$; once these temperatures were known, we were able to work out the entropic and energetic densities, and thus the numbers of degrees of freedom and the effective numbers of neutrino families in both sectors. We found that the evolution of the ordinary sector is influenced by the mirror particles, that provide extra degrees of freedom, and, due to the electron-positron annihilation process happening in the mirror sector, changes with time for temperatures around the BBN epoch. This makes crucial the present investigation in order to study in details the BBN process in presence of mirror dark matter. Furthermore, the special feature that models containing the mirror sector change the number of equivalent neutrinos during the Universe evolution, if considered together with preexisting models of cosmic microwave background (CMB) and large scale structure power spectra with mirror dark matter, furnishes a possible interpretation of the discrepancy in the effective number of neutrino families indirectly ``observed\" at BBN and CMB epochs.  In the mirror sector we obtained very high values of degrees of freedom, as predicted, and again they are variable for temperatures below some MeV. These data provide the necessary input for a detailed study of BBN in the mirror sector. It is required in order to obtain the exact amounts of primordial heavy mirror elements, that are necessary for a correct interpretation of non-linear processes of structure formation and of the new DAMA/LIBRA results."
    },
    "0809/0809.2502_arXiv.txt": {
        "abstract": " ",
        "introduction": "The accelerating cosmic expansion first inferred from the observations of distant type Ia supernovae \\cite{Riess:1998cb} has strongly confirmed by some other independent observations, such as the cosmic microwave background radiation (CMBR) \\cite{Spergel:2006hy} and Sloan Digital Sky Survey (SDSS) \\cite{:2007wu}. An exotic form of negative pressure matter called dark energy is used to explain this acceleration. The simplest candidate of dark energy is the cosmological constant $\\Lambda$, whose energy density remains constant with time $\\rho_\\Lambda = \\Lambda / {8\\pi G}$ and whose equation of motion is also fixed, $w_{\\Lambda} = P_\\Lambda / \\rho_\\Lambda = -1$ ($P_\\Lambda$ is the pressure) during the evolution of the universe. The cosmological model that consists of a mixture of the cosmological constant and cold dark matter is called LCDM model, which provides an excellent explanation for the acceleration of the universe phenomenon and other existing observational data. However, as is well know, this model faces two difficulties, namely, the 'fine-tuning' problem and the 'cosmic coincidence' problem. The former also states: Why the cosmological constant observed today is so much smaller than the Plank scale, while the latter states: Since the energy densities of dark energy and dark matter scale so differently during the expansion of the universe, why they are at the same order today? To alleviate or even solve these two problems, many dynamic dark energy models were proposed such as the quintessence model rely on a scalar field minimally interacting with Einstein gravity. Here 'dynamic' means the equation of state of the dark energy is no longer a constant but slightly evolves with time. Despite considerable works on understanding the dark energy have been done, the nature of dark energy and its cosmological origin are still enigmatic at present.\\\\  On the other hand, the problem of discriminating different dark energy models is now emergent. In order to solve this problem, a sensitive and robust diagnostic for dark energy is a must. The statefinder parameter pair $\\{r, s\\}$ introduced by Sahni et al.\\cite{Sahni:2002fz} and Alam et al.\\cite{Alam:2003sc} is proven to be useful tools for this purpose. The statefinder probes the expansion dynamics of the universe through high derivatives of the scale factor $\\ddot a$ \\& $\\dddot a$ and is a natural next step beyond the Hubble parameter $H\\equiv \\dot a/a$ and the deceleration parameter q which depends upon $\\ddot a$. The statefinder pair $\\{ r, s\\}$ is defined as \\begin{equation}\\label{statefinder pair defined} r\\equiv\\frac{\\dddot a}{aH^3}, \\quad s\\equiv\\frac{r-1}{3(q-1/2)} . \\end{equation} The statefinder pair is a 'geometrical' diagnostic in the sense that it is constructed from a space-time metric directly, and it is more universal than 'physical' variables which depends upon properties of physical fields describing dark energy, because physical variables are, of course, model-dependent. Usually one can plot the trajectories in the $r-s$ plane corresponding to different dark energy models to see the qualitatively different behaviors of them. The spatially flat LCDM scenario corresponds to a fixed point $\\{ r, s \\} = \\{1,0\\} $ in this diagram. Departure of a given dark energy model from this fixed point provides a good way of establishing the \"distance\" of this model from LCDM. As demonstrated in refs.\\cite{Sahni:2002fz, Alam:2003sc, Zimdahl:2003wg} the statefinder can successfully differentiate between a wide variety of dark energy models including the cosmological constant, quintessence, the Chaplygin gas, braneworld models and interacting dark energy models.\\\\  One can plot the current locations of the parameters $r$ and $s$ corresponding to different models in statefinder parameter diagrams by theoretical calculating in these models. And on the other hand it can also be extracted from experiment data such as the SNAP( SuperNovae Acceleration Probe ) data, with which combined the statefinder parameters can serve as a versatile and powerful diagnostic of dark energy in the future. In this letter, we apply the statefinder diagnostic to the Ricci dark energy model(RDE). It is found that the evolution behavior of the statefinder parameters in this model is much like that in quiessence models, but in a very special case the statefinder diagnostic fails. In Section II, we will briefly review RDE model, and apply the diagnostic to it in Section III. In the last section we will give some conclusions. In the case of that the statefinder diagnostic fails, we apply a new diagnostic called the Om diagnostic proposed recently to RDE model in Appendix A. ",
        "conclusions": "In this letter, we have apply the statefinder diagnostic to the Ricci dark energy model, and plot the trajectories in the $r-s$ and $r-q$ planes in the case of $\\alpha=0.46$ and $f_0=0.65$. Here we have used the values of $\\alpha$ and $f_0$ that founded in ref.\\cite{Gao:2007ep}, but other values are also possible. Different values of $\\alpha$ will determine different evolutions of the statefinder parameters, so the determining of $\\alpha$ from more precise data provided by future experiments will be needed. In a very special case that $\\alpha = 0.5$, the statefinder pair $\\{r, s\\}$ fails to discriminate LCDM model and RDE model, because they give the same fixed point $r=1, s=0$ in the $r-s$ diagram. The difference of these two models is the evolution of the equation of state $w$, which is a constant that equals $-1$ in the former and a time-dependent variable in the latter. Recently, a new diagnostics called the $Om$ diagnostic of dark energy is proposed in ref.\\cite{Sahni:2008xx}. In Appendix A, we will apply this diagnostic to RDE model in the case of $\\alpha=0.5$ in order to differentiate LCDM model and RDE model. The result indicates that this new diagnostic really works well for this purpose."
    },
    "0809/0809.1737_arXiv.txt": {
        "abstract": "{We report on a multiwavelength observation of the blazar 3C~454.3 (which we dubbed {\\it crazy diamond}) carried out on November 2007 by means of the astrophysical satellites \\agile{}, \\igr{}, \\swi{}, the \\webt{} Consortium, and the optical-NIR telescope \\rem{}.} Thanks to the wide field of view (FoV) of the \\agile{} satellite and its prompt alert dissemination to other observatories, we obtained a long (three weeks), almost continuous \\gray coverage of the blazar 3C~454.3 across fourteen decades of energy. {This broad-band monitoring allows us to study in great detail light curves, correlations, time-lags and spectral energy distributions (SEDs) during different physical states.} {Gamma-ray data were collected  during an \\agile{} pointing towards the Cygnus Region. Target of Opportunity (ToO) observations were performed to follow-up the \\gray observations in the soft and hard X-ray energy bands. Optical data were acquired continuously by means of a pre-planned \\webt{} campaign and through a \\rem{} ToO repointing.} {\\source{} is detected at a $\\sim 19-\\sigma$ level during the 3-week observing period, with an average flux above 100~MeV of $ F_{\\rm E>100MeV} = (170 \\pm 13) \\times 10^{-8}$\\,\\phcmsec. The \\gray spectrum can be fit with a single power-law with photon index $\\Gamma_{\\rm GRID} = 1.73 \\pm 0.16$ between 100~MeV and 1~GeV. We detect significant day-by-day variability of the \\gray emission during our observations, and we can exclude that the fluxes are constant at the 99.6\\% ($\\sim 2.9 \\sigma$) level. The source was detected typically around 40 degrees off-axis, and it was substantially off--axis in the field of view of the AGILE hard X-ray imager. However, a 5-day long ToO observation by \\igr{} detected \\source{} at an average flux of about $F_{\\rm 20-200\\,keV} = 1.49 \\times 10^{-3}$\\,\\phcmsec with an average photon index of $\\Gamma_{\\rm IBIS} = 1.75 \\pm 0.24$ between 20--200~keV. \\swi{} also detected \\source{} with a flux in the 0.3--10~keV energy band in the range $(1.23-1.40) \\times 10^{-2}$\\,\\phcmsec{} and a photon index in the range $\\Gamma_{\\rm XRT} = 1.56-1.73$. In the optical band, both WEBT and REM show an extremely variable behavior in the $R$ band. A correlation analysis based on the entire data set is consistent with no time-lags between the \\gray and the optical flux variations.} {Our simultaneous multifrequency observations strongly indicate that the dominant emission mechanism between 30~MeV and 30~GeV is dominated by inverse Compton scattering of relativistic electrons in the jet on the external photons from the broad line region.} ",
        "introduction": "\\label{3c454:intro} Among active galactic nuclei (AGNs), blazars show intense and variable \\gray emission above 100~MeV \\citep{Hartman1999:3eg}. Variability timescale can be as short as few days, or last few weeks. They emit across several decades of energy, from the radio to the TeV energy band. Blazar spectral energy distributions (SEDs) are typically double humped with a first peak occurring in the IR/Optical band in the so-called {\\it red blazars} (including Flat Spectrum Radio Quasars, FSRQs, and Low-energy peaked BL Lacs, LBLs) and at UV/X-rays in the so-called {\\it blue blazars} (including High-energy peaked BL Lacs, HBLs) and it is commonly interpreted as synchrotron radiation from high-energy electrons in a relativistic jet. The second SED component, which peaks at MeV--GeV energies in red blazars and at TeV energies in blue blazars, is commonly interpreted as Inverse Compton (IC) scattering of soft seed photons by relativistic electrons. A recent review of the blazar emission mechanisms and energetics is given in \\cite{Celotti2008:blazar:jet}. Alternatively, the blazar SED can be explained in the framework of the hadronic models, where the relativitic protons in the jet are the primary accelerated particles, emitting \\gray radiation by means of photo-pair and photo-pion production (see \\citealt{Mucke2001:hadronic,Mucke2003:hadronic} for a recent review on hadronic models).  Multiwavelength studies of variable \\gray blazars have been carried out since the beginning of the 1990s, thanks to the \\cgro{} Observatory. Nevertheless, only a few objects were detected on a timescale of about two weeks in the \\gray energy band, and simultaneously monitored at different energies, obtaining a multifrequency coverage. Among the FSRQs detected at energies above 100~MeV by the \\egret{} telescope on-board the Compton Gamma-Ray Observatory \\citep{Hartman1999:3eg}, \\source{} (PKS~2251$+$158; $z=0.859$) is certainly one of the most active at high energy. In the \\egret{} era, it was detected in 1992 during an intense \\gray flaring episode \\citep{Hartman1992:3C454iauc, Hartman1993:3C454_EGRET} when its flux $F_{\\rm E>100MeV}$ was observed to vary within the range $(0.4-1.4) \\times 10^{-6}$\\,photons\\,cm$^{-2}$\\,s$^{-1}$. In 1995, a 2-week campaign detected a \\gray flux $< 1/5$ of its historical maximum \\citep{Aller1997:3C454_EGRET}.  In 2005, \\source{} underwent a  major flaring activity in almost all energy bands (see \\citealt{Giommi2006:3C454_Swift}). In the optical, it reached $R=12.0$\\,mag \\citep{vil06} and it was detected by \\igr{} at a flux\\footnote{Assuming a Crab-like spectrum.} level of $\\sim 3 \\times 10^{-2}$\\,photons\\,cm$^{-2}$\\,s$^{-1}$ in the 3--200~keV energy band \\citep{Pian2006:3C454_Integral}. Since the detection of the exceptional 2005 outburst, several monitoring campaigns were carried out to follow the source multifrequency behavior \\citep{vil06,vil07,rai07,rai08a,rai08b}. During the last of these campaigns, 3C~454.3 underwent a new optical brightening in mid July 2007, which triggered observations at all frequencies.  In July 2007, \\citet[][hereafter V08]{Vercellone2008:3C454_ApJ} reported the highest \\gray flare from \\source{}. During the period 2007 July 24--30 the average \\gray flux was $F_{\\rm E>100MeV} = (280 \\pm 40) \\times 10^{-8}$\\,\\phcmsec, about a factor of two higher than in 1995. No emission was detected by Super-AGILE in the energy range 20--60~keV, with a 3-$\\sigma$ upper limit of $2.3 \\times 10^{-3}$\\,\\phcmsec.  By means of \\agile{} preliminary flux estimate \\citep{Vercellone2007:atel1160}, \\cite{Ghisellini2007:3C454:SED} compared the \\source{} SEDs obtained during three multiwavelength campaigns (2000, 2005 and 2007). The 2007 data show that the \\gray high state occurred during a weaker optical/X-ray flux compared with the 2005 flare.  In this paper (Paper I) we present the results of a multiwavelength campaign on \\source{} during a long-lasting \\gray activity period between 2007 November 10 and  December 1. Preliminary \\gray results were distributed in \\citet{Chen2007:ATel3C454}, while radio-to-optical and UV data were published in \\citet{rai08a}. A companion paper (Paper II, Donnarumma et al., in preparation) will describe the \\agile{} multiwavelength campaign during December 2007. In Sect.~\\ref{3c454:ml} we present the multiwavelength campaign on \\source{}; in Sect.~\\ref{3c454:grid} through~\\ref{3c454:optical} we present the \\agile/GRID, \\igr/IBIS, \\swi/XRT, \\webt{} and \\rem{} data analysis, respectively; in Sec.~\\ref{3c454:sim:data} we present the simultaneous multiwavelength light-curves and SEDs. In Sect.~\\ref{3c454:disc} and~\\ref{3c454:conc} we discuss our results and draw our conclusions. Throughout this paper the quoted uncertainties are given at the 1--$\\sigma$ level, unless otherwise stated. ",
        "conclusions": "\\label{3c454:conc} The \\agile{} mission is particularly suited to monitor a large number of potential \\gray sources. The \\agile{} pointings during the month of November 2007,  despite being centered approximately in the Cygnus region of the Galactic plane, revealed the very prominent \\gray activity of the blazar \\source{}. The \\agile{} detection of this blazar prompted a series of important multiwavelength observations. The electromagnetic emission of 3C 454.3 could be determined with unprecedented coverage over 14 orders of magnitude in energy during a period that included a substantial fraction of the months of November and December, 2007. Results of the AGILE data and related multifrequency campaign carried out in December, 2007, will be presented in a forthcoming paper.  We reported in this paper the main results of our \\agile{} and associated multifrequency campaigns during the month of November, 2007. Our results can be summarized as follows: \\begin{enumerate} \\item The \\gray emission from \\source{} dominated the whole extragalactic sky as monitored by AGILE during its first year of scientific operations. \\item Our \\gray data show remarkable variability on a daily timescale for a Flat Spectrum Radio Quasar. \\item Emission in the optical range appears to be correlated with that at \\gray energies above 100 MeV. \\item Variability in the soft and hard X-ray range is less sensitive to the short-timescale variations of the optical flux. \\item The average \\gray spectrum during the whole campaign is substantially harder than that reported in previous observations. \\item We determined the SEDs for episodes of relatively high \\gray emission. \\item  Our results support the idea that the dominant emission mechanism in \\gray energy band is the inverse Compton scattering of external photons from the BLR clouds scattering off the relativistic electrons in the jet. \\end{enumerate}"
    },
    "0809/0809.1447_arXiv.txt": {
        "abstract": "{ This study investigates some of the consequences of representing the sky by a rectangular grid of pixels on the dynamic range of images derived from radio interferometric measurements. In particular, the effects of image pixelization coupled to the CLEAN deconvolution representation of the sky as a set of discrete delta functions can limit the dynamic range obtained when representing bright emission not confined to pixels on the grid. {\\reffont  Sky curvature effects on non-coplanar arrays will limit the dynamic range even if strong sources are centered on a pixel in a ``fly's eye'' representation when such pixel is not located at the corresponding facet's tangent point.  Uncertainties in the response function of the individual antennas as well as in the calibration of actual data due to ionospheric, atmospheric or other effects will limit the dynamic range even when using grid-less subtraction (i.e. in the visibility domain) of strong sources located within the field of view of the observation. } A technique to reduce these effects is described and examples from an implementation in the Obit package are given. {\\reffonta Application of this technique leads to significantly superior results without a significant increase in the computing time. } \\keywords {Techniques: image processing, Techniques: interferometric} } ",
        "introduction": "With the new generation of high sensitivity interferometers to come on-line in the next few years (EVLA, ALMA, LOFAR) wide--field imaging will be necessary to achieve the sensitivity possible with these instruments. The problem is especially acute at low frequencies ($<10$ GHz) where every field of view will contain several relatively bright sources at any time. The sensitivity of instruments such as LOFAR or the EVLA at lower frequencies may be compromised much of the time by artifacts due to the bright sources in the field if these artifacts are not reduced to an acceptable level. This paper describes artifacts arising from using pixelated images to describe the sky as well as a technique for reducing them. All data manipulations discussed in this report used the Obit package (\\cite{OBIT}, http://www.cv.nrao.edu/$\\sim$bcotton/Obit.html).   \\section {Effects of Pixelated Images} It is generally convenient to represent the sky seen by an imaging interferometer as a set of pixel values on a rectangular grid. This is a good match to the widely used CLEAN deconvolution technique which represents the sky as a set of delta functions located at the centers of cells on such a grid. A commonly used measure of the quality of an image is its ``dynamic range,'' generally defined to be the ratio of the brightest pixel in an image to the RMS pixel-to-pixel fluctuation in areas devoid of emission. {\\reffonta Application of this criterion is generally straightforward as the response of the primary beam leads to mostly empty regions in the images sufficiently far from the pointing center. This convention is adopted in the following.}  \\begin{figure*}[!t] \\centerline{ \\includegraphics[height=3.5in,angle=-90]{9104fg1a.eps} \\includegraphics[height=3.5in,angle=-90]{9104fg1b.eps} } \\caption{ {\\reffonta Contour plots of example CLEAN images derived from noiseless point source model data restored using delta functions for the components to demonstrate the actual distribution of locations of CLEAN components. The plots show the same region and have the same contour levels, factors of powers of 2 x 0.1\\% of the model flux density; negative contours are dashed. \\hfil\\break {\\bf Left:} The point source was located approximately midway between cells in both dimensions. CLEANing used 1000 components. \\hfil\\break {\\bf Right:} The point source was located exactly on the central pixel. Only the central pixel contains emission; CLEANing used 200 components. } } \\label{Example} \\end{figure*}   One limitation of the pixelization technique is that emission not confined to points on the grid cannot be represented exactly and the CLEAN technique will approximate such structure by a potentially infinite series of alternating positive and negative delta functions. The problem is particularly severe when the image contains very bright, unresolved emission as is common in the radio sky. A combination of the limited support of the actual representation and the effects of the finite precision of digital computers will limit the dynamic range obtained by introducing artifacts in the derived image. {\\reffonta This effect has long been recognized as a problem \\citep{Briggs1992, Perley1999B}. \\citet{Briggs1992} describe the result as due to the discontinuity in the visibility function of an off--center point source at the boundary of the sampled region of the u--v plane which requires an infinite number of components in the image plane to model it accurately. \\citet{Briggs1992} estimate that this effect will limit dynamic range to $\\approx 1000$.  This effect is easily understood for an {\\reffont unresolved} source located between pixels. In order to model a point source between pixels, the deconvolution must add emission in the adjacent pixels. To counteract the resultant broadening of the source, negative emission must then be added around the source. An example of this effect is clearly shown in Figure \\ref{Example} which compares the results of CLEANing a model point source both exactly centered on a pixel and offset between pixels. For the offset source, the off--source RMS is 2.3 $\\times 10^{-4}$ of the model source flux density (dynamic range = 4300) while for the centered source, the off--source RMS is 4.7 $\\times 10^{-10}$ of the model source flux density (dynamic range = 2.1$\\times 10^{9}$). In the latter case, the result is limited by the precision of 32--bit digital arithmetic. }  In the past, uncentered point sources have not been a particularly serious problem as the bright source was usually the object under investigation and centering it on a pixel was straightforward. This is not the case for surveys or future observations where the source(s) of interest may be faint sources in the presence of multiple, much stronger sources whose locations on the imaging grid are not easily controlled. {\\reffonta Processing of the NRAO VLA Sky Survey (NVSS) \\citep{NVSS} used partial pixel shifts to align the imaging grid to a single exceptionally bright ($>0.5$ Jy) source whenever such a source was present in the field. }  A second limitation of the pixelization technique is that the rectangular grid is flat whereas the sky is not, see \\cite{Cornwell1992}. In the case of an array confined to a plane during synthesis, such as the E-W linear Westerbork Synthesis Radio Telescope, a projection of the sky is possible which avoids this problem; but, in the general case, this is not possible. This effect is referred to in the following as the co-planarity problem. If the field of interest is small enough, the curvature of the sky can be negligible. Several techniques have been developed to solve the more general problem. Some of these are: \\begin{itemize} \\item Full 3-D imaging\\\\ Interferometric measurements are made in visibility space described by the coordinate set (u,v,w). These measurements can be convolved onto a 3D grid and Fourier transformed into a 3D image. The celestial sphere is a spherical surface in this 3D image. A 3D deconvolution followed by projection onto a plane is possible but in practice this is sufficiently expensive in computing resources that it is not used. See \\cite{Perley1999A} for details. \\item Fly's Eye\\\\ The curved surface of the celestial sphere can be approximated by a mosaic of facets, each tangent to the celestial sphere and of sufficiently limited extent that the error introduced is {\\reffont negligible}. However, the errors increase quadratically with the distance of a cell to the corresponding tangent point and this can still limit the dynamic range. See ``the Polyhedron Method'' in \\cite{Cornwell1992} for details. \\item W projection\\\\ It is possible to correct for the diffractive effects on the wavefront as it propagates from the antenna closer to the source to the farther on each baseline. This correction is made to the convolution of the visibility data onto the grid prior to Fourier transformation. This allows a single, flat 2D grid to represent the curved sky. However, the derived ``dirty'' image is no longer strictly a convolution. See \\cite{Cornwell2005} for details. \\end{itemize} In the following, only adaptations of the Fly's Eye technique are considered.  \\section {Wide--field Imaging with Fly's Eye and Catalog of Sources} The technique used for wide--field imaging in the following tests is as follows. A circular field of view to be completely imaged is specified by its radius from the pointing center. The data are examined and the cell spacing (if not specified) is picked on the basis of the longest baseline in the data {(\\reffonta one quarter of the smallest fringe spacing)} and the size of a facet ``undistorted''  by co-planarity effects determined from the maximum extent of any baseline in the selected data in the direction of the pointing center. {\\reffonta The radius of an undistorted region is adapted from \\citet{Thompson99} as $0.33\\ \\sqrt{1/maxW}$ radians where $maxW$ is the maximum value of the ``w'' in the data set in wavelengths. } A ``Fly's Eye'' tessellation of {\\reffonta circular } regions in a hexagonal pattern is then defined which fully covers the field of view and a set of square images enclosing these circular regions defines a mosaic of facets.  In general, the imaged field of view does not enclose all of the sky to which the array elements have significant gain, so that sources are visible outside of this fully--imaged field of view. To include such sources, facets are added to the mosaic centered on the positions of outlying sources obtained from a catalog (currently a stripped--down version of the NVSS) which are deemed to have an apparent brightness above a given {\\reffonta user-specified} threshold based on an assumed spectral index and a model of the antenna gain pattern. {\\reffonta  These additional facets do not need to be contiguous with those covering the field of view wanted. } At high frequencies, a catalog could be generated from the WMAP catalog of point sources \\citep{WMAP}. {\\reffonta Accurate positions and flux densities are not required, only a list of positions around which a bright source {\\bf might} appear, as the decision to auto--center is based on the results of the initial CLEAN.}  The variant of CLEAN used in the following discussion is the ``visibility based'' or ``Cotton-Schwab'' CLEAN \\citep{Schwab1983, Cotton-SW89}, in which the Fourier transform of the estimate of the sky is iteratively subtracted from the visibility data and the residual image re-derived. This allows multiple independent ``facets'' to be imaged on the sky. All facets in the mosaic to be deconvolved are formed and a quality measure based on both peak (residual) brightness and extended emission ({\\reffonta 0.95 $\\times$ peak absolute value residual in the clean window + 0.05 $\\times$ the average residual - as used in the AIPS package\\footnote{http://www.aips.nrao.edu}). } is used to determine which facet is to be CLEANed next. Components are selected from this facet and then subtracted from the visibility data. The next highest quality measure facet is re-imaged and if it still maintains its status, it is CLEANed, otherwise, the next highest facet is re-imaged and tested, etc. {\\reffonta The CLEAN is stopped when one of two conditions is satisfied, 1) the total number of CLEAN components reaches a user specified limit or 2) the maximum absolute value residual in the CLEAN window of all facets is less than a user specified minimum, usually of order of the anticipated RMS in the image. }  After CLEAN is finished, components are (optionally) convolved with a Gaussian approximation to the instrumental PSF and restored to the facet from which they were subtracted as well as to any overlapping facet containing their positions. After restoration of the subtracted components, all facets are projected (``flattened'') onto a grid covering the specified field of view.  {\\reffont Because strong ``flanking sources'' can be observed with dedicated pointings prior to reducing the data at hand and can thus be available in a catalog, it would seem possible (in principle) to subtract them from the measured visibilities prior to processing of the target field.  However, such visibility-based subtraction is hindered by foreground effects on the calibration as well as by imprecise knowledge of the response of the primary beam of the antennas which modify the apparent position and flux density of the cataloged sources.  Efforts to fit these effects in the uv-plane have had limited success.  The technique has been shown to work reasonably well for only a few sources using simulated data and appears to require vast computing resources to handle even a few sources \\citep{Voronkov04}. }  \\section {AutoCenter Technique} In order to minimize pixelization related artifacts using the Fly's Eye technique, it is desirable to locate each strong, point--like source {\\reffonta exactly} on the pixel that is at the tangent point of the facet that contains it. {\\reffonta This is accomplished by first identifying these sources in an initial imaging step and then adding facets for strong sources not already very close to a facet center. } It is also necessary to restrict CLEAN from assigning any components to any overlapping regions of other facets.  An initial CLEAN is used to determine which objects in the field, if any,  have sufficiently bright emission to warrant being centered on a special facet. The CLEAN needs to be deep enough that an accurate measure of the centroid of the source can be made to subsequently center it on a pixel. {\\reffont Currently,} the implementation in the Obit package CLEANs to a factor of 0.1 of {\\reffonta the user--specified auto--center brightness threshold above which artifacts are expected to be above the noise level.} The initial CLEAN level needed depends on the uv coverage, and dynamic range but generally a factor of 100 to 1000 below the peak is adequate. Also, since the residuals are not needed, it is not necessary to derive a full set of residual images at the end of the initial CLEAN. The decision {\\reffonta that AutoCentering is needed } is based on the highest pixel value in the CLEAN-able portions of the initial dirty images.  Peak brightnesses are determined from the sum of the CLEAN components within a given radius (cells within 62.5\\% of the FWHM of the PSF in the current implementation) of each component derived from each facet. If this sum exceeds the {\\reffonta auto--center brightness} threshold level, then the centroid of the peak is determined from the CLEAN components by a moment analysis. A {\\reffont small} facet, currently 96$\\times$96 pixels centered at the derived centroid, is added to the working imaging mosaic. It is possible that the same source may appear in several overlapping facets in the initial {\\reffont fly's eye tiling} so it is necessary to ensure that a given strong source {\\reffont is added only once}, a coincidence of better than one-half of the FWHM of the synthesized beam (typically two cells) is considered to correspond to the same source. If a source is already within 0.5 pixel of the tangent pixel of the enclosing facet, a new facet is not created, but the facet is re-centered (if not already within 0.01 cell).  In order to ensure that CLEAN assigns no components to the re-centered source appearing in overlapping facets of the initial mosaic, each of the facets is examined to see if it contains the position of the source to be re-centered. In any facet in which the re-centered source appears in the CLEAN-able region, a round ``unbox,'' currently of radius 33 pixels, is added to that facet centered on the position of the corresponding source. This size is slightly smaller than the size of the initial cleanable region in the new facet added, a radius of 38 pixels, to ensure that no pixel will be excluded a priori from a component search {\\reffont given the non-coplanarity of the original and the new facets}. These choices are somewhat arbitrary but seem to work well. An ``unbox'' is like a normal CLEAN window except that any enclosed pixel will not be considered as a location for CLEAN components even if located inside of another regular CLEAN box,  i.e. the CLEAN process ignores any pixels inside of an unbox. Pixels inside of unboxes are also excluded from statistical estimates such as maximum, minimum and RMS. If any re--centering operations are required, the initial CLEAN is repeated. After this re--centering, all prior CLEAN components are discarded {\\reffont before beginning a CLEAN}.  Because the accuracy of the determination of the centroid of a source to be re-centered is adversely affected by the limited CLEAN and the very effects this technique is trying to correct, some iteration may be in order. Images in which high dynamic range is desired generally benefit as well from one or more iterations of self--calibration. At the beginning of each {\\reffonta CLEAN}, the centroids of {\\reffonta auto--centered sources from the previous CLEAN are} checked to see if they are sufficiently close to the tangent pixel. If the first (reduced) moment of CLEAN components within 1.5 pixels of the central pixel is offset by more than 0.01 pixel, then the image is re-centered and all components from the previous CLEAN are discarded. This allows an iterative refinement of the centroid of the peak and improves significantly the results over a single estimate.  \\section {Examples} The following sections give two sets of examples processed using this technique. The first involves simulated data and the second, actual VLA observations of 3C84.  \\begin{table}[!t] \\centering \\caption{Offset Source Dynamic Ranges} \\vskip 0.1in \\begin{tabular}{|c|c|c|}   \\hline \\hline Offset(pixels) &  DR$^a$  &  DR\\_corr$^b$  \\\\ \\hline 0     &  $>1.0\\times 10^{23}$   & $>1.0\\times 10^{23}$  \\\\ 0.01  &  $260\\times 10^{3}$     &   $9.2\\times 10^{6}$  \\\\ 0.02  &  $130\\times 10^{3}$     &   $10.1\\times 10^{6}$ \\\\ 0.03  &  $94\\times 10^{3}$    &   $10.3\\times 10^{6}$ \\\\ 0.05  &  $60\\times 10^{3}$    &   $8.9\\times 10^{6}$  \\\\ 0.1   &  $29\\times 10^{3}$    &   $8.9\\times 10^{6}$  \\\\ 0.141 (0.1,0.1)$^c$ &  $17\\times 10^{3}$    &   $4.9\\times 10^{6}$  \\\\ 0.2   &  $17\\times 10^{3}$    &   $11.5\\times 10^{6}$ \\\\ 0.283 (0.2,0.2)$^c$ &  $9.7\\times 10^{3}$     &   $4.4\\times 10^{6}$  \\\\ 0.3   &  $12\\times 10^{3}$    &   $10.2\\times 10^{6}$ \\\\ 0.5   &  $10\\times 10^{3}$    &   $9.6\\times 10^{6}$  \\\\ 0.707 (0.5,0.5)$^c$ &  $6.2\\times 10^{3}$    &   $7.1\\times 10^{6}$  \\\\ \\hline 10    &  $65\\times 10^{3}$    & $69.6\\times 10^{6}$   \\\\ 20    &  $16\\times 10^{3}$    & $79.1\\times 10^{6}$   \\\\ 30    &  $7.3\\times 10^{3}$    & $66.3\\times 10^{6}$   \\\\ 50    &  $2.7\\times 10^{3}$    & $8.1\\times 10^{6}$    \\\\ 100   &  $0.74\\times 10^{3}$    & $0.4\\times 10^{6}$    \\\\ \\hline \\end{tabular} \\hfill\\break Notes:\\hfill\\break $^a$ Dynamic range without autoCenter\\hfill\\break $^b$  Dynamic range with autoCenter\\hfill\\break $^c$  Offsets along both coordinates\\hfill\\break \\label{pixel_offset} \\end{table}  \\begin{figure}[!t] \\centering \\includegraphics[height=3.0in,angle=-90]{9104fig2.eps} \\caption{ Dynamic Range (DR) obtained on simulated, noiseless data as a function of fractional pixel offset. A combination of offsets on one and both axes are included. The ``+'' symbols represent CLEANing without using autoCenter and the ``$\\ast$'' represent the results of using autoCenter. } \\label{frac_pixel} \\end{figure} \\begin{figure}[!t] \\centering \\includegraphics[height=3.0in,angle=-90]{9104fig3.eps} \\caption{ Dynamic Range (DR) obtained on simulated, noiseless data as a function of large, whole pixel offset offsets to test non co-planarity effects. The ``+'' symbols represent CLEANing without using autoCenter and the ``$\\ast$'' represent the results of using autoCenter. } \\label{whole_pixel} \\end{figure}  \\subsection {Simulated Data} Simulated data have the advantage that their properties are known which makes it simple to separate artifacts from source structure. In these tests, the model used was a single 1 Jy point source with no noise or other corruptions added. In order to test the effects of fractional pixel offsets, the source was located at the phase center of the data {\\reffont and the data imaged with} a series of fractional pixel shifts added to the center of the image. The artificial data set used the same uv--plane sampling as a 74 MHz VLA data set which consisted of 12 frequency channels of 122 kHz bandwidth.  The tests were performed in the Obit package.  The CLEAN used the visibility--based technique and proceeded for 1000 iterations with a loop gain of 0.1 and a CLEAN window of radius 10 cells centered on the peak. The model visibility computation used the  ``DFT'' method which is more accurate than the ``GRID'' method as implemented in Obit (and AIPS). Several iterations of imaging/re-centering were done in order to refine the estimate of the centroid. The cell spacing used was 0.25 of the FWHM of the fitted Gaussian restoring beam so the image was reasonably over--sampled. The components removed by the CLEAN procedure were not restored and the RMS of the pixel values in the final residual image was used to derive the dynamic range. Each of the tests was then repeated turning on the autoCenter mode and the corresponding dynamic range {\\reffont determined in turn}. A combination of offsets on a single axis and on both axes are included. The results are shown in Table \\ref{pixel_offset} and Figure \\ref{frac_pixel}.  A second set of tests explored the effects of non co--planarity by inserting a series of large, whole--pixel offsets to the position of the point source model and using a procedure like the one described above for small pixel offsets. The results are shown in Table \\ref{pixel_offset} and Figure \\ref{whole_pixel}. All offsets except the one of 100 pixels are within the ``undistorted field of view'' of a facet.  While these tests are not exhaustive, it is clear that fractional pixel offsets of bright point--like sources can limit the dynamic range to  $\\sim10^4$ and non co--planar effects can limit the dynamic range to  $\\sim10^3$ even with a moderately conservative limit on facet size. Applying the autoCenter technique improved the dynamic range by typically a factor of 50 to 1000 for the fractional pixel offset tests and typically by a factor of 1000 in the large pixel offset tests.  Figure \\ref{frac_pixel} shows a fair amount of scatter in the corrected dynamic range achieved. We believe this to be the result of residual errors in the centroiding as offsets on two axes produced lower dynamic range than comparable offsets on a single axis. All corrected dynamic range {\\reffont values} were substantially better than that obtained using a 0.01 pixel offset without correction. At very high dynamic range the accuracy of the centroiding needs to be exceedingly precise. This is possible in this test as there are no systematic errors and the model is exactly a point source. In the real sky, resolution may be a problem even for a source whose size is a very small fraction of the PSF; position errors of less than 0.0025 of the PSF seem to be capable of limiting dynamic range so resolution on similar scales might be a problem. In Figure \\ref{whole_pixel} the efficacy of the correction seems to diminish with increasing pixel offset. This may be as much a problem with the simulated data as with the imaging; the larger position shifts needed to model a source with a substantial offset from the pointing center will result in loss of numerical precision.  Non co--planar effects coupling to the pixelization can limit significantly the dynamic range achieved. A source observed 20\\% of the way to the edge of its imaging facet suffers comparable dynamic range loss to a central source observed 0.2 cells from the closest pixel. This indicates that imaging using the Fly's Eye technique needs to be applied with bright sources at the center of a facet (tangent point) and not merely on a pixel. A coarser grid spacing will likely lead to larger errors than presented here.  \\subsection {Actual Data} \\begin{figure*}[!t] \\centering \\includegraphics[height=3.5in]{9104fg4a.eps} \\centerline{ \\includegraphics[height=3.5in]{9104fg4b.eps} \\includegraphics[height=3.5in]{9104fg4c.eps} } \\caption{ 3C84 with contour interval of powers of $\\sqrt{2}$. The same contouring relative to the peak in the image is used for all plots. \\hfil\\break {\\bf Top:} 3C84 observed at the pointing center. The most negative contour (dashed) is at -8 mJy/beam \\hfil\\break {\\bf Bottom Left:} Portion of wide-field image with 3C84 at the half power of the antenna pattern. The most negative contour (dashed) is at -11 mJy/beam \\hfil\\break {\\bf Bottom Right:} Like Left except using the autoCenter technique. The most negative contour (dashed) is at -4 mJy/beam } \\label{3C84} \\end{figure*}  A test using real data and wide-field imaging was made using the VLA at 1.4 GHz and observations of 3C84 (peak = 24.6 Jy). 3C84 was observed both at the pointing center and at the half power point of the beam. The observations were made in spectral mode with 15 channels of 390 kHz bandwidth. The data were bandpass calibrated in addition to the amplitude and phase calibration and the edge channels were excluded from processing. Wide--field 3671$\\times$3671 pixel images were made to cover the primary beam of the antennas using the Fly's Eye technique and a 37 facet mosaic. In all cases, Obit task Imager {\\reffont determined and} applied amplitude and phase self--calibration to optimize the dynamic range. The data set with 3C84 at the half power of the antenna power pattern was imaged both with and without the autoCenter technique. In these images, 3C84 was half--way to the edge of its facet (co--planarity limit) and 0.4 of a cell from the nearest grid cell in Right Ascension and 0.3 of a cell in Declination. Sampling was 4 pixels per beam. Contour plots of a portion of the images around 3C84 are shown in Figure \\ref{3C84}.  \\begin{table*}[t] \\caption{3C84 Dynamic Range} \\begin{center} \\begin{tabular}{|l|r|r|r|r|r|} \\hline\\hline Image &  Peak       & Near RMS & Far RMS & Near DR & Far DR \\\\ &  Jy         & Jy       & Jy      &         &  \\\\ \\hline center$^1$    & 24.6 & 0.0020 & 0.00122 & 12289 & 20145 \\\\ half$^2$      & 12.0 & 0.0024 & 0.00121 &  5029 &  9933 \\\\ half/auto$^3$ & 11.6 & 0.0010 & 0.00055 & 11745 & 21163 \\\\ \\hline \\end{tabular} \\end{center} \\hfill\\break Notes:\\hfill\\break $^1$ 3C84 observed at pointing center,\\hfill\\break $^2$ 3C84 observed at half power,\\hfill\\break $^3$  3C84 observed at half power using autoCenter. \\label{3C84DR} \\end{table*}  It is immediately obvious from Figure \\ref{3C84} that the autoCenter technique helps improve the dynamic range of the image with 3C84 well away from the pointing center. This is explored further in Table \\ref{3C84DR} which gives the relevant image statistics. The RMS was determined in 601$\\times$601 pixel windows either centered on 3C84 (``Near RMS'') or far from 3C84 (``Far RMS''). The near RMS values were determined using a histogram analysis.  As can be seen from Table \\ref{3C84DR}, using the autoCenter technique doubles the effective dynamic range of the image, even in regions far from the obvious artifacts caused by 3C84 and does even better near the source. The autoCenter image has comparable dynamic range to the observation with 3C84 on axis. Especially in the neighborhood of 3C84, it is clear that other systematics {\\reffont are limiting the dynamic range. Indeed, given the extended emission surrounding 3C84 the images are limited by uv-coverage as the test observations lacked the necessary short spacings.  We find that this limitation is more severe than other systematic effects such as bandpass mismatches, pointing errors and beam squint.  We have re-imaged the data after discarding baselines shorter than 12~k$\\lambda$ and we have obtained a dynamic range that is $\\sim 20$\\% higher (again using the autoCenter technique) although the extended flux that surrounds the core of 3C84 is, of course, lost to the baseline restriction.}  {\\reffonta The difference in execution times for processing with and without the autoCenter technique depends on the details, i.e., number of self--calibrations, structure in the field, etc.; but is seldom significant. The cost of the extra, shallow CLEAN to locate sources to be re-centered can be partially or totally compensated by less time spent modeling artifacts in subsequent CLEANs. In the test presented in this section, processing without autoCentering actually took 1\\% longer than with autoCentering. } ",
        "conclusions": "We have demonstrated that the dynamic range obtained from imaging interferometric observations can be adversely affected by pixelization of the images. Image pixelization effects can limit dynamic range to about $10^4$ even for point sources and non co--planar effects can limit the dynamic range to about $10^3$ even with a moderately conservative limit on facet size. The higher dynamic range needed for the EVLA and LOFAR, where wide fields of view with numerous bright sources will be common, need improved techniques to circumvent these limits.  The autoCenter technique presented here has shown factors of 50 to 1000 improvement in images made from simulated data with no noise or systematic errors. An improvement greater than a factor of 2 was achieved in images made from real observations of the bright source 3C84, even far from the obvious artifacts due to the source. {\\reffonta Due to the lower level of artifacts, processing using this technique on the 3C84 test presented above used marginally less computer time than without. }  Fractional pixel corrections are more difficult to implement in a single image w-projection method. A simple shift can center a single source, also varying the pixel spacing could center two sources and a rotation could add a third but centering a larger number of sources will not be possible. The introduction of separate, simultaneous ``w--projection'' grids to accommodate such strong sources would seem to reduce the benefits of the ``w--projection'' technique\\footnote{Alternatively, a finite number of tangent facets could be used to supplement the ``w--projection'' grid and ``unboxes'' used in a manner similar to the implementation discussed here.  Again, this would seem to reduce the benefits of the ``w--projection'' technique.}. {\\reffont Even grid-less subtraction of strong sources will require some shifting of their cataloged position due to foreground and (variable) instrumental effects.  One such procedure discussed in the literature \\citep{Voronkov04} seems to be much more computationally expensive than the procedure discussed in this paper, even when dealing with a small number of sources.}  The technique described here can be applied to an arbitrary number of point--like sources. The tests presented here suggest that high dynamic range imaging of bright extended sources needs a better set of basis functions than the delta functions on grid cells that are used by CLEAN."
    },
    "0809/0809.1671_arXiv.txt": {
        "abstract": "Remote observing of exoplanetary atmospheres is now possible, offering us access to circulation regimes unlike any of the familiar Solar System cases.  Atmospheric circulation models are being developed to study these new regimes but model validations and intercomparisons are needed to establish their consistency and accuracy. To this end, we present a simple Earth-like validation of the pseudo-spectral solver of meteorological equations called IGCM (Intermediate General Circulation Model), based on Newtonian relaxation to a prescribed latitudinal profile of equilibrium temperatures. We then describe a straightforward and idealized model extension to the atmospheric flow on a hot Jupiter with the same IGCM solver. This shallow, three-dimensional hot Jupiter model is based on Newtonian relaxation to a permanent day-night pattern of equilibrium temperatures and the absence of surface drag. The baroclinic regime\\footnote{Contours of pressure and density are misaligned in a baroclinic flow, while they are aligned in a barotropic flow. Strongly baroclinic flows are susceptible to baroclinic instabilities which, in essence, are slanted versions of convection \\citep[see the review by][]{Showman2007}.} of the Earth's lower atmosphere is contrasted with the more barotropic regime of the simulated hot Jupiter flow. For plausible conditions at the 0.1-1 bar pressure level on HD 209458b, the simulated flow is characterized by unsteadiness, subsonic wind speeds, a zonally-perturbed superrotating equatorial jet and large scale polar vortices. Violation of the Rayleigh-Kuo inflexion point criterion on the flanks of the accelerating equatorial jet indicates that barotropic (horizontal shear) instabilities may be important dynamical features of the simulated flow. Similarities and differences with previously published simulated hot Jupiter flows are briefly noted. ",
        "introduction": "Since the first discoveries of extrasolar giant planets around nearby sun-like stars with Doppler velocimetry (Mayor \\& Queloz 1995; Marcy \\& Butler 1996), extrasolar planet research has experienced a spectacular series of observational breakthroughs. In recent years, progress has been particularly rapid with a subset of close-in extrasolar planets found to transit the disk of their host star \\citep[e.g.,][]{Charbo2008,Deming2008}.   The first transit observation of a hot Jupiter (Charbonneau et al.\\ 2000; Henry et al.\\ 2000) has been followed by many such measurements, including a growing number of new close-in planet discoveries based on transit searches \\citep[e.g.,][]{Torres2008}. For the brightest nearby systems with transiting planets, this has also enabled the detection of the planetary thermal flux occulted at secondary eclipse (e.g., Deming et al. 2005, 2006; Charbonneau et al.  2005; Harrington et al. 2007), the detection of day/night temperature variations through IR phase curve monitoring (e.g. Harrington et al. 2006; Knutson et al. 2007, 2008b; Cowan et al. 2007), IR spectral measurements (Grillmair et al. 2007; Richardson et al. 2007; Knutson et al. 2008a) as well as a variety of additional constraints based on transit spectroscopic studies (Tinetti et al. 2007; Ehrenreich et al. 2007; Swain et al. 2008; Barman 2007; Redfield et al. 2008; Pont et al. 2008).  Overall, these results have contributed to a shift in focus from the detection of extrasolar planets to the characterization of their physical atmospheric properties \\citep[e.g.,][]{Seager2005,Burrows06,Marley2007,Fortney2007,Barman2008,Burrows2008}.   The need to interpret these astronomical data reliably, so as to infer the physical conditions present in distant exoplanetary atmospheres, has also fueled a growing atmospheric modeling effort.  While plane-parallel radiative models have been the tools of choice to interpret these data until now, they are likely insufficient in the case of hot Jupiter/Neptune atmospheres.  Indeed, due to their short orbital separations, all these exoplanets are expected be tidally-locked to their parent star (or, in some cases, pseudo-synchronized). Tidally-locked planets are subject to an unusual situation of permanent, asymmetric day/night radiative forcing, leading to heat redistribution by atmospheric motions.  Current data on hot Jupiters already indicate that existing one-dimensional radiative transfer models fail to capture the multi-dimensional nature of this atmospheric regime (e.g., Seager et al. 2005; Knutson et al. 2007; 2008b; Fortney et al. 2006) and, when included, the effects of atmospheric heat redistribution are usually accounted for only indirectly \\citep[e.g.,][]{Iro05,Seager2005,Barman2008,Burrows2008,Fortney2008}. To adequately address this circulation regime and interpret the growing data set, multi-dimensional, coupled radiation-hydrodynamics models of these atmospheres are required \\citep[see the review by][]{Showman2007}.   Hot Jupiters have high atmospheric temperatures, slow rotation rates and an unusually permanent pattern of asymmetric day/night radiative forcing. Recognizing that this constitutes an interesting problem in atmospheric dynamics, several groups have explored plausible circulation regimes on these planets using different modeling approaches and assumptions \\citep{Showman2002,Cho2003,Menou2003,C&S05,Cooper2006,Langton2007,Cho2006,Dobbs-Dixon2008,Langton2008,Showmanetal08}. These investigations have exhibited significant diversity in flow results, for which there currently is no simple unifying explanation \\citep[see the discussion in][]{Showman2007}.   Atmospheres in motion are non-linear, presumably turbulent flows. In addition, radiatively active species in an atmosphere are advected by the flow and in doing so non-linearly couple the flow to its effective source of radiative forcing. Simulated atmospheric flows are thus generally expected to be sensitive to various numerical and physical details of any specific implementation.  In the exoplanetary context, with rather limited direct information on the atmospheres studied, this means that simple parameterized models isolating key features of the simulated flow are important in helping us understand the general behavior of remote atmospheric flows. With this in mind, we present here idealized atmospheric models which emphasize dynamical aspects under linear forcing conditions. These simplified models do not address any detailed aspect of the radiative or chemical structure of the atmospheres under consideration.   In general, different numerical implementations of a specific atmospheric problem will not lead to identical results. For this reason, model validations and inter-comparisons are standard practices in atmospheric science. This is particularly true of the complex hydrodynamic solvers known as dynamical cores \\citep[e.g.,][]{HeldSuar94}.  The development of dynamical cores to solve the equations satisfied by an hydrostatic atmosphere in motion has been a major enterprise. Dynamical cores currently used in meteorological and climate models are the results of years of refinements to guarantee stability, efficiency, and accurate conservation of mass, momentum and energy \\citep[e.g.,][]{HoskSimm75}. In the exoplanetary context, different modeling groups have used various hydrodynamic solvers. Some are new and little tested while others are old, well-tested but only validated under conditions appropriate for Solar System planetary atmospheres. As a result, intercomparisons and validations of dynamical cores on idealized atmospheric problems tailored for exoplanets will be important in the future to assess, both qualitatively and quantitatively, how reliable our interpretations of exoplanetary data can be.   As a first step in this direction, we present here a simple Earth-like validation of the IGCM dynamical core (in \\S3), followed by a direct extension of this model to the atmospheric flow on a hot Jupiter (in \\S4). ",
        "conclusions": "In this work, we have presented a simple Earth-like general circulation model based on the IGCM dynamical core. We used this model to contrast the baroclinic circulation regime of the Earth's lower atmosphere with the more barotropic circulation regime that emerges from a straightforward extension of the model to the atmospheric flow on a hot Jupiter.  The distinction between these two major (barotropic and baroclinic) regimes of atmospheric circulation is an important one in ``geophysical'' fluid dynamics. For instance, both regimes are relevant to the Earth's atmosphere and critical to our understanding of its general circulation, with a baroclinic (lower-level) troposphere and a barotropic (higher-level) stratosphere.    Various factors contribute to the degree of baroclinicity/barotropicity of an atmospheric flow. The more stably-stratified an atmosphere is (i.e. the more ``radiative'' it is, as opposed to convective, to use the language of stellar physics), the larger its external Rossby deformation radius is, the weaker baroclinic instability growth is and thus the more barotropic the circulation regime will be (e.g. Pedlosky 1987; Cho et al. 2008).  The strong external irradiation experienced by hot Jupiter atmospheres creates rather strongly-stratified temperature profiles in their photospheric regions (e.g. Seager \\& Sasselov 1998; Sudarsky et al. 2000; Iro et al. 2005; Barman et al. 2005; Fortney \\& Marley 2007). This relative vertical stability, together with slow (synchronized) rotation and high atmospheric temperatures, leads to large external Rossby deformation radii \\citep{Showman2002,Cho2003,Menou2003,Cho2006} and favors a barotropic circulation regime (with vertically aligned horizontal motions in the various atmospheric layers). The lack of surface drag on the atmospheric flow, for otherwise identical forcing conditions, also favors a barotropic regime as horizontal shear tends to inhibit the development of baroclinic instabilities \\citep[e.g.,][]{JamesGray86,James1987,Robinson1997} For all these reasons, a single-layer, vertically-integrated barotropic treatment of horizontal motions in hot Jupiter atmospheres may be justified \\citep{Cho2003,Cho2006,Menou2003,Salby1989}.     Since the distinction between barotropic and baroclinic regimes depends on the degree of atmospheric vertical stratification, which is known for the Earth but {\\it a priori} unknown for remote exoplanets, the results from our shallow hot Jupiter model should only be interpreted as suggestive that this regime is relevant to hot Jupiter atmospheric flows. A more systematic exploration of circulation regimes on hot Jupiters will be needed to address this issue more thoroughly.  The shallow hot Jupiter model presented here is rather specific and idealized in a number of important ways. For instance, adopted values for the profile of relaxation temperatures and the radiative relaxation time are rather arbitrary.  The presence of deeper atmospheric layers and their interaction with the modeled layers has also been ignored in this shallow model.    Nevertheless, the model may capture important dynamical features of the atmospheric flow on hot Jupiters. In particular, the simulated flow has a number of similarities with comparable results reported in the literature, together with noticeable differences. As we have already emphasized, our simulated flow is characterized by a broad, zonally-perturbed superrotating equatorial wind, large scale polar vortices, unsteadiness and subsonic wind speeds. The emergence of a super-rotating (eastward equatorial) wind in this hemispherically-forced flow is consistent with the results of \\citet{Showman2002,C&S05,Showmanetal08,Dobbs-Dixon2008,Langton2008}. The broad width of this equatorial wind and the presence of counter-jets at midlatitudes is also consistent with the specific results of \\citet{Showmanetal08} and \\citet{Cho2003,Cho2006}. As discussed in detail by \\citet{Cho2006}, however, the typically opposite (westward) direction of the equatorial wind that emerges in equivalent-barotropic simulations may point to some limitation of that approach. On the other hand, the strong zonal disturbances of the equatorial wind in our model is very reminiscent of a similar pattern of large-amplitude Rossby waves discussed by Cho et al. (2003; 2008; see also Langton \\& Laughlin 2007). This is qualitatively different from the zonally-symmetric equatorial flow reported by \\citet{Showman2002,C&S05,Showmanetal08}. As Fig.~\\ref{fig:seven} illustrates, these zonal disturbances of the equatorial wind are closely related to the flow unsteadiness observed in our model \\citep[see also][]{Langton2008}.   While our shallow hot Jupiter flow exhibits large-scale polar vortices, their dynamical nature is quite distinct from that of the circumpolar vortices discussed by \\citet{Cho2003,Cho2006}.  The polar vortices shown in Fig.~\\ref{fig:seven}, for instance, are anti-cyclonic, not geostrophically-balanced and consistently located on the planetary night-side, somewhat eastward of the anti-stellar point. This is in contrast with the geostrophically-balanced, cyclonic polar vortices that exhibit systematic longitudinal translations in the equivalent-barotropic flows described by \\citet{Cho2003,Cho2006}. We interpret this difference as being due to the forcing-dominated nature of polar vortices in our hot Jupiter model, which starts at rest, rather than a dynamical origin like in the flows of \\citet{Cho2003,Cho2006}, where polar vortices emerge from turbulent initial conditions subject to an energy cascade to large scales under the constraint of potential vorticity conservation (see \\citealt{Cho2003,Cho2006} for a discussion; see also \\citealt{Langton2007}). We note that, even though our simulated flow is clearly unsteady (see Fig.~\\ref{fig:seven}), a possible consequence of the different nature of polar vortices in the present shallow hot Jupiter model could be a reduced level of disk-integrated variability, from less dynamically active polar vortices, by comparison to what equivalent-barotropic results have indicated so far \\citep{Cho2003,Menou2003,Cho2006,Rauscher2007,Rauscher2008}.  An important difference between our results and several comparable studies reported in the literature is the consistently subsonic value of wind speeds in our shallow hot Jupiter model.  In this respect, our results stand out by comparison with those of \\citet{Showman2002,C&S05,Showmanetal08,Dobbs-Dixon2008,Langton2008}. It is presently unclear what is the origin of this fundamental discrepancy. We simply note here that one element of answer might be the emergence of barotropic (horizontal shear) instabilities in our shallow model, which appear to result from the acceleration of the equatorial wind and its flank counterjets and could possibly limit the asymptotic wind speeds in our model. More work, including model inter-comparisons, is needed to clarify this point.    The magnitude of wind speeds on hot Jupiters is an open problem. Unlike the atmospheres of terrestrial planets, giant planet atmospheres lack the large-scale sink of energy and momentum that is associated with friction on a solid surface.\\footnote{On the Earth and in our simple Earth-like model, for instance, the atmosphere reaches a global state of momentum balance with the bulk planetary rotation through surface drag, with positive and negative contributions depending on the easterly or westerly nature of surface winds \\citep[e.g.,][]{Holton1992,JamesGray86}. } As a result, the only sources and sinks of energy and momentum in a hot Jupiter flow as simulated here are the Newtonian relaxation, which represents large-scale sources and sinks of radiation, and the hyperdissipation, which operates on small scales. The absence of a clearly identified large-scale sink of energy (ground friction) makes giant planet atmospheres possibly more difficult to understand than their terrestrial counterparts.  Indeed, dissipation in the flow, which ultimately determines asymptotic wind speeds, is then more strongly dependent on the flow itself, via its large-scale coupling to radiation and its effective turbulent dissipation on small scales. \\citet{Goodman09} has recently suggested that internal friction between atmospheric layers could play an important role in the hot Jupiter context.   The large variety of flow behaviors found so far in distinct hot Jupiter studies suggests that, in addition to model validations such as our simple Earth-like model, inter-comparisons of hydrodynamic solvers on identical, well-defined atmospheric circulation problems may be necessary to build a solid understanding of these new circulation regimes. The shallow hot Jupiter model presented here could be used as a first step in this direction."
    },
    "0809/0809.3397_arXiv.txt": {
        "abstract": "{ Recent VLT SINFONI observations of the close environments ($\\sim 30$~pc) of nearby AGNs have shown that thick gas tori and starbursts with ages between $10$ and $150$~Myr are frequently found. By applying these observations to a previously established analytical model of clumpy accretion disks, we suggest an evolutionary sequence for starburst and AGN phases. Whereas the observed properties of the gas tell us about the current state of the torus, the starburst characteristics provide information on the history of the torus. In the suggested evolution, a torus passes through 3 different phases predetermined by an external mass accretion rate. Started by an initial, short, and massive gas infall, a turbulent and stellar wind-driven $Q \\sim 1$ disk is formed in which the starburst proceeds. Once the supernovae explode the intercloud medium is removed, leaving a massive, geometrically thick, collisional disk with a decreasing, but still high-mass accretion rate. When the mass accretion rate has significantly decreased, the collisional torus becomes thin and transparent as the circumnuclear disk in the Galactic center of the Milky Way. Variations on this scenario are possible either when there is a second short and massive gas infall, in which case the torus may switch back into the starburst mode, or when there is no initial short massive gas infall. All observed tori up to now have been collisional and thick. The observations show that this phase can last more than $100$~Myr. During this phase the decrease in the mass accretion rate within the torus is slow (a factor of 4 within $150$~Myr). The collisional tori also form stars, but with an efficiency of about $10$\\,\\% when compared to a turbulent disk. ",
        "introduction": "} In the unification scheme for active galactic nuclei (AGN) the central massive black hole is surrounded by a geometrical thick gas and dust torus (see, e.g., Antonucci 1993). If the observer's line-of-sight crosses the torus material, the AGN is entirely obscured from near-IR to soft X-rays and only visible at X-ray energies if the gas column density is not too high (Seyfert 2 galaxies). On the other hand, if the torus is oriented face-on with respect to the observer, the central engine is visible (Seyfert 1 galaxies). The spectral energy distributions (SEDs) of most quasars and AGN in Seyfert galaxies have a pronounced secondary peak in the mid-infrared (mid-IR) (e.g. Sanders et al. 1989; Elvis et al. 1994), which is interpreted as thermal emission by hot dust in the torus. The dust is heated by the primary optical/ultraviolet (UV) continuum radiation, and the torus extends from the dust sublimation radius outwards (Barvainis 1987).  The geometrical thickness of the torus in the gravitational potential of the galactic nucleus implies a vertical velocity dispersion of about $50$-$100$~km\\,s$^{-1}$. If one assumes that the disk is continuous, i.e. thermally supported, this corresponds to a temperature of $\\sim 10^5$~K. Since this is beyond the dust sublimation temperature ($\\sim 10^3$~K), thick tori have to be clumpy or must be supported by additional forces other than thermal pressure. Krolik \\& Begelman (1988) proposed a clumpy torus model where the clumps have supersonic velocities. Vollmer et al. (2004) and Beckert \\& Duschl (2004) elaborated this model in which orbital motion can be randomized if magnetic fields permit the cloud collisions to be sufficiently elastic. Vollmer et al. (2004) found that the circumnuclear disk (CND) in the Galactic center (G\\\"{u}sten et al. 1987) and obscuring tori share the same gas physics, where the mass of clouds is  in the range 20 - 50 $M_\\odot$ and their density close to the limit of disruption by tidal shear. A change in matter supply and the dissipation of kinetic energy can turn a torus into a CND-like structure and vice versa. Any massive torus will naturally lead to sufficiently high mass accretion rates to feed a luminous AGN.  If and how efficient these clumpy tori form stars is an open question. The large majority of observational studies probed the nuclear star formation on scales of a few hundred parsecs (see, e.g., Sarzi et al. 2007, Asari et al. 2007, Gonz\\'ales Delgado \\& Cid Fernandes 2005, Cid Fernandes et al. 2004). These studies resulted in a general view that about $30$\\,\\% - $50$\\,\\% of the sample AGNs are associated with recent (ages less than a few 100~Myr) star formation on these scales. Thanks to the high spatial resolution of the near-infrared adaptive optics integral field spectrograph SINFONI, it has only recently become possible to study the environments of AGN on the 10~pc scale. Davies et al. (2007) analyzed star formation in the nuclei of nine Seyfert galaxies at spatial resolutions down to $0.085''$. They found recent, but no longer active, starbursts in the central regions which occurred $10$ - $300$~Myr ago. Moreover, Hicks et al. (2008) were able to measure the rotation and dispersion velocity of the molecular gas in these galaxies using the $2.12$~$\\mu$m H$_2$ (1-0)S(1) line. Surprisingly, all molecular gas tori have high velocity dispersions and are therefore geometrically thick.  In this article we compare % the observations of Davies et al. (2007) and Hicks et al. (2008) with the expectations of analytical models of clumpy accretion disks developed in Vollmer \\& Beckert (2002, 2003) and Vollmer et al. (2004). In a first step, we test % whether these models are able to describe observations. In a second step, these models allow us to investigate the scenario hypothesized by Davies et al. (2007) where sporadic, short-lived starbursts are due to short massive accretion events in the central region, followed by more quiescent phases until there is another episode of accretion. This scenario is corroborated by a closer look at the nucleus of NGC~3227 where Davies et al. (2006) found signs of a past starburst ($\\sim 40$~Myr ago) and a presently quiescent gas torus with a Toomre parameter $Q>1$. ",
        "conclusions": ""
    },
    "0809/0809.4262_arXiv.txt": {
        "abstract": "Gravitational wave emission by coalescing black holes (BHs) kicks the remnant BH with a typical velocity of hundreds of $\\kms$. This velocity is sufficiently large to remove the remnant BH from a low-mass galaxy but is below the escape velocity from the Milky Way (MW) galaxy. If central BHs were common in the galactic building blocks that merged to make the MW, then numerous BHs that were kicked out of low-mass galaxies should be freely floating in the MW halo today.  We use a large statistical sample of possible merger tree histories for the MW to estimate the expected number of recoiled BH remnants present in the MW halo today. We find that hundreds of BHs should remain bound to the MW halo after leaving their parent low-mass galaxies. Each BH carries a compact cluster of old stars that populated the core of its original host galaxy. Using the time-dependent Fokker-Planck equation, we find that present-day clusters are $\\lesssim 1\\,$pc in size, and their central bright regions should be unresolved in most existing sky surveys. These compact systems are distinguishable from globular clusters by their internal (Keplerian) velocity dispersion greater than one hundred km~s$^{-1}$ and their high mass-to-light ratio owing to the central BH.  An observational discovery of this relic population of star clusters in the MW halo, would constrain the formation history of the MW and the dynamics of BH mergers in the early Universe.  A similar population should exist around other galaxies and may potentially be detectable in M31 and M33. ",
        "introduction": "During the final coalescence of two black holes (BHs), gravitational waves (GWs) are emitted unisotropically and carry away linear momentum, thus kicking the merger BH remnant in the opposite direction \\citep{1962PhRv..128.2471P,1973ApJ...183..657B,1983MNRAS.203.1049F}. The resulting kick velocity of typically hundreds of $\\kms$ depends on the mass ratio of the BHs as well as the spin and orientation of the binary before coalescence \\citep{2006ApJ...653L..93B,2007PhRvL..98w1102C,2007ApJ...659L...5C,2007PhRvD..76f1502T}. Such kicks can alter both the the population of nuclear BHs in galactic bulges \\citep{2004ApJ...606L..17M,2006MNRAS.368.1381L,2007ApJ...663L...5V,2007ApJ...667L.133S,2008arXiv0805.1420B} as well as the core of stars in the bulge itself \\citep{2008ApJ...678..780G}. The discovery of a remnant population of recoiled BHs in the present-day Universe can provide a new window into the merger statistics and spin distribution of the BHs \\citep{2006MNRAS.368.1381L,2007ApJ...663L...5V}, as well as test general relativity in the strong regime.  Recoiling BHs may be detected through their unique flaring \\citep{2008ApJ...676L...5L,2008ApJ...682..758S,2008arXiv0802.3556S} or observed as spatially and kinematically offset quasars \\citep{2004ApJ...606L..17M}.  If the BH binary was surrounded by an accretion disk, the ejected remnant BH would carry the disk with it and shine as a quasar \\citep{Loeb}.  A search for spectral shifts in the broad lines of quasars relative to the narrow emission lines of their parent galaxies resulted mainly in upper limits \\citep{2007ApJ...666L..13B} and possibly one suggested candidate \\citep{2008ApJ...678L..81K} for an ejected quasar. Unfortunately, these quasars lose their source of gas for accretion, and have short lifetimes \\citep{Loeb,2008arXiv0805.1420B,2008ApJ...687L..57V}.  Past studies of the observational consequences of GW recoil have generally focused on massive BHs with a mass $\\msmbh \\gtrsim 10^7\\,\\msun$.  Such BHs reside in the most massive galaxies, and require the greatest and rarest kick velocities in order to escape their parent galaxies \\citep{2007ApJ...662L..63S}.  But since the kick velocity from BH mergers is independent of the total mass of the system, the dynamical consequences are most prominent for less massive BHs which tend to reside in low-mass galaxies with shallower potential wells and lower escape velocities \\citep{2004ApJ...606L..17M,2005MNRAS.358..913V,2006ApJ...653L..93B,2006MNRAS.368.1381L,2007ApJ...663L...5V,2008arXiv0807.4702T}. In fact, most BHs with $\\msmbh \\lesssim 10^5\\,\\msun$ will likely be kicked out of their parent galaxies as a result of major mergers \\citep{2006ApJ...653L..93B}. During the hierarchical build-up of the Milky-Way (MW) galaxy, mergers of low-mass galaxies were common and should have resulted in a population of freely floating BHs \\citep{2005MNRAS.358..913V}. Since the gravitational potential of an overdense region does not evolve dramatically during the growth of cosmological structure \\citep{2006astro.ph..3360L} the region that eventually collapsed to make the MW was able to trap those BHs with the most common kick velocities ($\\lesssim 500~{\\rm km~s^{-1}}$) even before the MW had formed.  Each of the ejected BHs carries with it a star cluster that used to populate the core of its parent galaxy \\citep{2008ApJ...678..780G, 2008ApJ...678L..81K}.  In this paper we examine the observational signatures of the ejected star clusters which are expected to be floating in the MW halo.  The discovery of a relic population of the star clusters attached to recoiled BHs would provide a new window for cosmology, allowing one to constrain the merger history of the MW as well as the formation history of the first population of massive BHs.  {\\new Previous studies have looked at populating the MW halo with recoiling BHs as analysed again here \\citep{2004ApJ...606L..17M,2005MNRAS.358..913V,2006MNRAS.368.1381L}, through three body encounters \\citep{2005MNRAS.358..913V}, as well as a variety of other processes.  Other sources of wandering BHs are the remnants of Population III stars \\citep{2003MNRAS.340..647I,2005PhRvL..95a1301Z} as well as the the direct collapse of baryons into a BH \\citep{2005PhRvD..72j3517B,2006MNRAS.368.1340M}.  In both of these cases, however, the BH either is presently located in the center of a galaxy or is ``naked'' without any stellar companions.  Such BHs are likely to only be located by their interaction with the matter that surrounds them \\citep{2004MNRAS.354..629I,2004MNRAS.354..443I,2005MNRAS.358..913V,2005ApJ...628..873M,2005PhRvD..72j3517B}. Additionally, dwarf galaxies may be tidally stripped of their outer stars, leaving behind a BH surrounded by a massive cluster of stars similar to some globular clusters \\citep{1988IAUS..126..603Z,1994ApJ...431..634B}.  In contrast, the new source of BHs presented in this paper uniquely have a population of bound stars that are only a fraction of the mass of the central BH.  }  In \\S~\\ref{sec:hier} we consider the formation history of the MW to estimate the number and distribution of ejected BHs that may be found in its halo today. We then follow the long term evolution of the star clusters that are attached to them in \\S~\\ref{sec:relax}.  In \\S~\\ref{sec:observe} we discuss the observational signatures of these systems.  Finally, \\S~\\ref{sec:disc} summarizes our main results and their implications. ",
        "conclusions": "\\label{sec:disc}  Based on a large statistical ensemble of merger tree histories for the MW, we have found that hundreds of GW recoiled BHs should reside within the MW halo today. The BHs have a mass $\\msmbh \\gtrsim 10^3\\,\\msun$, and should be surrounded by a cluster of stars that were tightly bound to the BH when it was ejected.  The most massive BH weighs $\\sim 1-6 \\times 10^5\\,\\msun$, and is surrounded by a compact star cluster $\\lesssim 1\\,$pc in size with a stellar velocity dispersion of hundreds of km~s$^{-1}$.  High-resolution adaptive optics imaging along with spectroscopy of the star clusters can constrain the distance and mass of the BH in the center of the cluster based on orbits of individual stars for the closest clusters \\citep{2005ApJ...620..744G, 2005ApJ...628..246E}.   The number of recoiled BHs in the MW is most sensitive to the fraction of low mass galaxies that harbored BHs in their centers, as well as to the merger history of such galaxies during the formation of the MW. We emphasize that most of these galactic ``building blocks'' were made at high redshifts and were disrupted during the assembly process of the MW; hence, they likely had different properties than the remaining dwarf satellites of the MW today.  The future discovery of the population of ejected star clusters will provide a unique probe of the early history of the Milky Way, as well as the distribution and evolution of low mass BHs. It would also open a new window to exploring the low-mass end of the population of nuclear BHs in high-redshifts galaxies.  \\citet{2008arXiv0809.5046M}, whose work we learned of after submission of this manuscript, arrived at a similar conclusion independently, with a slightly different application to the nearby Virgo cluster."
    },
    "0809/0809.0261_arXiv.txt": {
        "abstract": "Several aspects of mathematical astrobiology are discussed. It is argued that around the time of the origin of life the handedness of biomolecules must have established itself through an instability. Possible pathways of producing a certain handedness include mechanisms involving either autocatalysis or, alternatively, epimerization as governing effects. Concepts for establishing hereditary information are discussed in terms of the theory of hypercycles. Instabilities toward parasites and possible remedies by invoking spatial extent are reviewed. Finally, some effects of early life are discussed that contributed to modifying and regulating atmosphere and climate of the Earth, and that could have contributed to the highly oxidized state of its crust. ",
        "introduction": "Since the early days of nonlinear dynamics and non-equilibrium thermodynamics it has been clear that one of the ultimate applications of this theory might be to facilitate an understanding of the transition from non-living to living matter. The main reason is obviously that living systems are very far from equilibrium---as indicated by the high degree of order and hence the low entropy of living systems relative to their exterior.  Already in 1952 Turing \\cite{Turing} proposed the idea that chemical reaction-diffusion systems might provide a tool for studying biochemical pattern formation which have increased our understanding of the laws of nature far from equilibrium, where life occurs. This idea was followed up in the late 1960s by Prigogine \\cite{Prigogine,Prigogine2} who suggested that dissipative structures have great importance in establishing a physical description of living matter. A general theory of autocatalytic molecular evolution was developed in 1971 by Eigen \\cite{Eigen71} who argued that in a single micro-environment only a single handedness can result from a single event. In particular the famous chicken and egg problem that occurs in biology at different levels was identified as a Hopf bifurcation. A Hopf bifurcation describes the spontaneous emergence of an oscillating solution once some stability threshold has been crossed. The mathematics of this is familiar to any physicist, but it requires that the equations describing the relevant physics are known. In biology it is not even clear that the various phenomena can be described by equations. A first detailed attempt in this direction was indeed that of Eigen. However, the equations governing the emergence of life are only phenomenological ones. Nevertheless, these approaches are invaluable in that they help giving the origin of life question a mathematical basis.  One of the earliest anticipated forms of life that is still similar to present life is the RNA world \\cite{Gilbert}, whereby simple RNA molecules with functional behavior self-reproduces using genetic information encoded either in the same or in other participating RNA molecules. Obviously, there are tremendous difficulties given that RNA is too complicated a molecule to be synthesized abiologically. A significantly simpler molecule is peptide nucleic acid or PNA \\cite{Nilsen93}, where the backbone consists of peptide instead of sugar phosphate. Nevertheless, the difficulty of producing RNA remains.  There is no firm idea where on Earth such molecular replication may have taken place. Frequently discussed scenarios include hydrothermal vent systems \\cite{Russell06}, but also beach scenarios are discussed that are subject to tides leading to cyclic changes in concentration \\cite{Lathe} as well as to repeated wetting and drying \\cite{Bywater}.  An early experiment that contributed significantly to the research into origins of life was the Urey--Miller experiment \\cite{Miller53} that demonstrated the spontaneous production of amino acids in a reducing atmosphere consisting of H$_2$O vapor, CH$_4$, NH$_3$, and H$_2$ with an energy supply in the form of sparks. More recent experiments also allow for the presence of CO$_2$, which now seems unavoidable on the early Earth given that it is continuously replenished through outgassing by volcanoes.  In the following we review some aspects where there has been considerable cross-fertilization between astrobiology and nonlinear dynamics. We begin by discussing a phenomenon that is believed to have taken place around the time of the origin of life, namely the establishment of a definitive handedness of biomolecules that is inherent to DNA and RNA ({\\sc D}-form) and to amino acids ({\\sc L}-form). Next we discuss constraints on the evolution of hereditary information, and finally review some models that characterize the alteration of the terrestrial environment by early life. In addition to any of these physical effects there are random fluctuations leading inevitably to local imbalances between the concentrations of molecules of {\\sc D}- and {\\sc L}-form. In the following we discuss mechanisms that can lead to an exponential amplification of the enantiomeric excess. For a recent review of these ideas see Ref.~\\cite{Plasson07}. ",
        "conclusions": "Astrobiology has developed into a rapidly growing research field involving expertise from a number of neighboring disciplines. Nonlinear dynamics and nonequilibrium thermodynamics find applications in all these subfields. Here, we have elaborated on a few such aspects. Closest to the onset of life is perhaps the emergence of homochirality of biomolecules. Given that RNA has been proven to form longer polymers only in a homochiral environment, one would expect that homochirality must be a prerequisite to the emergence of life at the level of a replicating RNA world. On the other hand, the very mechanism causing the polymerization to terminate, namely the enantiomeric cross-inhibition, can also be the mechanism responsible for causing 100\\% homochirality by destroying RNA molecules whose chirality is already in the minority. This would however requite the possibility of auto-catalysis, which can be avoided in another scenario where a closed peptide system is kept away from equilibrium by continuous activation of amino acids.  Chemically speaking, the stabilization of a definite chirality is in some models similar to the subsequent establishment of hereditary information in that catalysis plays a crucial role. Furthermore, in both cases the possibility of chemistry in an extended system is crucial. On the one hand, spatial extent gives rise to the possibility of coexistence of life forms of opposite handedness on the early Earth. On the other hand, spatial extent can be critical in allowing the system to find unpopulated locations faster than being overwhelmed by the effects of parasites that tap the same resources that are required for the maintenance and development of hereditary information.  Finally, life is invariably coupled to come kind of metabolism that is ultimately powered by solar energy. This clearly affects the environment by reducing carbon and oxidizing  the crust of the Earth and, over the last two billion years also the atmosphere. How much these alterations of the environment are due to biological processes is less obvious. However, it is clear that biological factors greatly speed up weathering on the Earth. The extent of biologically induced alterations of the continental crust, for example, may therefore best be tested using quantitatively accurate model calculations. The outcome may ultimately hinge on energetic considerations and on the efficiency of photosynthesis as a solar energy collector.  With the scope of being able to explore in the near future not only the planets and other celestial bodies in the solar system in much more detail, but also planets of other planetary systems, the research in astrobiology quickly develops into a field that will be driven more and more by new data, making this field less susceptible to speculation. It is therefore important to be prepared for upcoming discoveries in this field. Finally, it should be emphasized that astrobiology is efficient in communicating science to the general public, which may provide additional boost to the field."
    },
    "0809/0809.0782_arXiv.txt": {
        "abstract": "{We present a detailed determination of the astrophysical parameters of the chromospherically active binary star EI\\,Eridani. Our new radial velocities allow to improve the set of orbital elements and reveal long-term variations of the barycentric velocity. A possible third-body orbit with a period of $\\approx 19$\\,years is presented. Absolute parameters are determined in combination with the {\\em Hipparcos} parallax. EI\\,Eri's inclination angle of the rotational axis is confined to $56\\fdg0 \\pm 4\\fdg5$, its luminosity class {\\sc iv} is confirmed by its radius of $2.37 \\pm 0.12$\\,\\Rsun. A comparison to theoretical stellar evolutionary tracks suggests a mass of $1.09 \\pm 0.05$\\,\\Msun\\ and an age of $\\approx6.15$~Gyr. The present investigation is the basis of our long-term Doppler imaging study of its stellar surface. } ",
        "introduction": "EI\\,Eridani = HD 26337 (G5\\,{\\sc iv}, $P\\nm{rot}\\ = 1.947$ days, $V$ = 7.1 mag) is a rapidly rotating ($v\\sin i = 51$ \\kms) active, non-eclipsing, single-lined spectroscopic binary star and as such a typical RS CVn star. RS CVn stars are close but detached binary systems that rotate synchronously due to tidal forces. This anomalously fast rotation is believed to be responsible for the high activity level found on these stars and constitute valuable laboratories to study the evolved Sun in action.  \\cahk emission -- the classical fingerprint of magnetic activity -- on EI\\,Eri was first noted by \\citet{bidelman73} and confirmed later by \\citet{fekel80} who classified it as moderate in strength, class C on the qualitative emission scale by \\citet{hearn79}. Line-profile data of \\cahk and \\halpha\\ emission lines were obtained by \\citet{kgs:fekel90} and \\citet{ff:montes94}. \\citet{fekel:hall82} classified EI~Eri as a single-lined RS~CVn star when they detected its light variability, with an amplitude being almost 0\\fm2 in $V$ and a photometric and orbital period of around two days. Later, \\citet{fekel:moffett86} determined the orbital period to be 1\\fd9472 while \\citet{hall:osborn87} detected a photometric period of $1.945 \\pm 0.005$ days from $UBV$ photometry. The \\halpha\\ line is in absorption, as in most RS\\,CVn stars, but highly filled in with chromospheric emission \\citep[see][]{smith:bopp82} and quite variable in strength \\citep{fekel:moffett86,zboril:oliveira05}. These variations are assumed to be caused by rotational modulation of the chromospheric \\halpha\\ emission, presumably due to evolving plage regions on the stellar surface moving in and out of sight. \\citet{fekel:quigley87}, \\citet{pallavicini:randich92} and \\citet{randich:gratton93} reported the presence of moderate amounts of lithium. \\citet{hall:osborn87} already noted season-to-season changes in the photometric period ($\\sim$1\\%), in the light-curve amplitude (0\\fm07 - 0\\fm20) and in the mean brightness ($\\sim$10\\%), likely indicating latitude and/or longitude changes of the location of starspots. \\citet{kgs:hall89} report seasonal changes of the photometric period of 0.043\\,days within three seasons. \\citet{rodono:cutispoto92} suggested a possible cycle period of about 10 years. \\citet{kgs:bartus97} summarized the first 16 years of photometric data and found a possible $11 \\pm 1$ years cycle of the mean brightness while \\citep{olah:kgs02} refined this to be 12.2 years. Recent observations in other spectral regions were presented by \\citet{osten:brown02}, \\citet{garcia:paredes03} and \\citet{cardini05}.  With its large rotational velocity and an intermediate inclination, EI\\,Eri is an ideal candidate for {\\em Doppler imaging\\/}.  This technique \\citep[see][]{rice02} allows the reconstruction of the surface spot distribution of rapidly rotating stars by using the relation between wavelength position across an absorption line and spatial position across the stellar disk. However, the rotational period of 1.947 days makes it difficult to obtain spectral coverage for a complete rotation period, and ideally, when using a single observing site, three weeks of observations are needed to provide good phase coverage for a Doppler image of EI\\,Eridani.  Doppler images for the 1984-87 period were presented by \\citet{hatzes:vogt92}. \\citet{kgs90} and \\citet{kgs:rice91} published Doppler images from 1987-88. All images show a large asymmetric spot at the pole or at high latitudes with several appendages and occasionally equatorial spots. The polar spot is long lived but exhibits significant changes of its size and shape, while the equatorial spots change within weeks. Furthermore, \\citet{donati:semel97} detected a clear Zeeman signature of the magnetic field of the primary star.  In this paper, we present spectroscopic and photometric observations obtained at KPNO and NSO in the years 1988 -- 1998 and with the MUSICOS~98 campaign, and perform an investigation of the absolute dimensions of EI\\,Eri in order to prepare for a Doppler imaging study. In two forthcoming paper, we will present our Doppler imaging results and address the question of spot lifetimes, long-term cyclic behavior \\citep[][ hereafter ``paper {\\sc ii}'']{paper2} and differential rotation \\citep[][ hereafter ``paper {\\sc iii}'']{paper3}. ",
        "conclusions": "The following results were obtained: \\begin{itemize} \\item Improved orbital parameters were determined with HJD = 2,448,054.7109 + 1.9472324 $\\times\\ E$, $\\gamma$ = 21.6\\,\\kms, $K_1$ = 26.8\\,\\kms. The barycenter velocity $\\gamma$ exhibits variations of about 8\\,\\kms\\ that are possibly cyclic. We interpret these variations due to a third stellar component and suggest a possible orbital solutions for the tertiary with a period of around 6700 days (19 years) and an amplitude of 4 -- 6\\,\\kms. An instrumental origin cannot be definitely ruled out. However, a thorough investigation of the orbit solutions of the individual observing runs and their radial velocity standard stars confirmed its reality in our data. The proposed period is also supported by the smaller {\\em O--C} radial velocity residuals. The third component would thereby surround the known binary system in a wide orbit. Corresponding values for the tertiary mass and the orbital parameters of the binary--tertiary system are quoted (Tab~\\ref{tab:tertiary}). However, its plausibility cannot be asserted unless a recurrence is observed or higher-quality RV measurements are available. \\item From the \\emph{Hipparcos} parallax, we obtain a luminosity of $4.60 \\pm 0.35$\\,\\Lsun\\ and a radius of $2.37 \\pm 0.12$\\,\\Rsun, independent from the inclination. \\item With our new parameters and by using evolutionary tracks from \\citet{pietrinferni:cassisi04}, we locate the position of the primary in the HR diagram and determine its mass to $1.09 \\pm 0.05$\\,\\Msun\\ and an age of $6.15 \\pm 0.50$\\,Gyr. \\item The inclination is found to be $i = 56.0 \\pm 4.5\\degr$. \\item Using the observed mass function of the inner binary and assuming $i\\nm{orb} \\equiv i\\nm{*}$, we derive the mass of the unseen secondary: $M_2 = 0.23 \\pm 0.04$\\,\\Msun\\ corresponding to an absolute magnitude of \\MV $\\approx 12 \\pm 1$ \\citep{baraffe:chabrier98}. \\end{itemize}"
    },
    "0809/0809.0057_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} In this note, we consider the observational constraints on some cosmological models by using the 307 Union type Ia supernovae (SNIa), the 32 calibrated Gamma-Ray Bursts (GRBs) at $z>1.4$, the updated shift parameter $R$ from WMAP 5-year data (WMAP5), and the distance parameter $A$ of the measurement of the baryon acoustic oscillation (BAO) peak in the distribution of SDSS luminous red galaxies with the updated scalar spectral index $n_s$ from WMAP5. The tighter constraints obtained here update the ones obtained previously in the literature. ",
        "introduction": "\\label{sec1} Recently, some observational data have been updated or have become available. In~\\cite{r1,r2}, the Wilkinson Microwave Anisotropy Probe (WMAP) collaboration released their 5-year observational data (WMAP5). The data of Cosmic Microwave Background (CMB) anisotropy have been significantly improved. Also, in~\\cite{r3,r4}, the Supernova Cosmology Project (SCP) collaboration released their new dataset of type~Ia supernovae (SNIa), which was called the Union compilation. The Union compilation contains 414 SNIa and reduces to 307 SNIa after selection cuts. This 307 SNIa Union compilation is the currently largest SNIa dataset.  On the other hand, Gamma-Ray Bursts (GRBs) were proposed as a complementary probe to SNIa recently~\\cite{r5,r6,r7,r8}. GRBs have been advocated as standard candles since several empirical GRB luminosity relations were proposed as distance indicators. However, there is the so-called circularity problem in the direct use of GRBs to probe cosmology~\\cite{r5}. Recently, a new idea to calibrate GRBs in a completely {\\em cosmology independent} manner has been proposed~\\cite{r9,r10}, and the circularity problem can be solved. The main idea is that of the cosmic distance ladder. Similar to the case of calibrating SNIa as the secondary standard candles by using Cepheid variables, which are primary standard candles, we can also calibrate GRBs as standard candles with a large amount of SNIa. Following the calibration method proposed in~\\cite{r10}, the distance moduli $\\mu$ of 32 calibrated GRBs at redshift $z>1.4$ are derived in~\\cite{r11}. Now, one can use them to constrain cosmological models {\\em without} circularity problem. See~\\cite{r10,r11} for details. As argued in~\\cite{r8,r32}, the observations at $z>1.7$ are fairly important to distinguish cosmological models and break the degeneracies between the parameters. In this note, we try to combine GRBs with the conventional datasets to constrain cosmological models. Although the number of GRBs is small and the systematic and statistical errors are very large so that their contribution to the constraints would be not so significant, this is still a beneficial exploration.  Here, we consider the observational constraints on some cosmological models by using the 307 Union SNIa compiled in~\\cite{r3}, the 32 calibrated GRBs at $z>1.4$ compiled in Table~I of~\\cite{r11}, the updated shift parameter $R$ from WMAP5~\\cite{r1}, and the distance parameter $A$ of the measurement of the baryon acoustic oscillation (BAO) peak in the distribution of SDSS luminous red galaxies~\\cite{r12} with the updated scalar spectral index $n_s$ from WMAP5~\\cite{r1}.  We perform a $\\chi^2$ analysis to obtain the constraints on the parameters of cosmological models. The data points of the 307 Union SNIa compiled in~\\cite{r3} and the 32 calibrated GRBs at $z>1.4$ compiled in Table~I of~\\cite{r11} are given in terms of the distance modulus $\\mu_{obs}(z_i)$. On the other hand, the theoretical distance modulus is defined as \\be{eq1} \\mu_{th}(z_i)\\equiv 5\\log_{10}D_L(z_i)+\\mu_0, \\ee where $\\mu_0\\equiv 42.38-5\\log_{10}h$ and $h$ is the Hubble constant $H_0$ in units of $100~{\\rm km/s/Mpc}$, whereas \\be{eq2} D_L(z)=(1+z)\\int_0^z \\frac{d\\tilde{z}}{E(\\tilde{z};{\\bf p})}, \\ee in which $E\\equiv H/H_0$ and $H$ is the Hubble parameter; ${\\bf p}$ denotes the model parameters. The $\\chi^2$ from the 307 Union SNIa and the 32 calibrated GRBs at $z>1.4$ are given by \\be{eq3} \\chi^2_{\\mu}({\\bf p})=\\sum\\limits_{i} \\frac{\\left[\\mu_{obs}(z_i)-\\mu_{th}(z_i)\\right]^2}{\\sigma^2(z_i)}, \\ee where $\\sigma$ is the corresponding $1\\sigma$ error. The parameter $\\mu_0$ is a nuisance parameter but it is independent of the data points. One can perform an uniform marginalization over $\\mu_0$. However, there is an alternative way. Following~\\cite{r13,r14}, the minimization with respect to $\\mu_0$ can be made by expanding the $\\chi^2_{\\mu}$ of Eq.~(\\ref{eq3}) with respect to $\\mu_0$ as \\be{eq4} \\chi^2_{\\mu}({\\bf p})=\\tilde{A}-2\\mu_0\\tilde{B}+\\mu_0^2\\tilde{C}, \\ee where $$\\tilde{A}({\\bf p})=\\sum\\limits_{i}\\frac{\\left[\\mu_{obs}(z_i) -\\mu_{th}(z_i;\\mu_0=0,{\\bf p})\\right]^2}{\\sigma_{\\mu_{obs}}^2(z_i)}\\,,$$ $$\\tilde{B}({\\bf p})=\\sum\\limits_{i}\\frac{\\mu_{obs}(z_i) -\\mu_{th}(z_i;\\mu_0=0,{\\bf p})}{\\sigma_{\\mu_{obs}}^2(z_i)}\\,, ~~~~~~~~~~ \\tilde{C}=\\sum\\limits_{i}\\frac{1}{\\sigma_{\\mu_{obs}}^2(z_i)}\\,.$$ Eq.~(\\ref{eq4}) has a minimum for $\\mu_0=\\tilde{B}/\\tilde{C}$ at \\be{eq5} \\tilde{\\chi}^2_{\\mu}({\\bf p})= \\tilde{A}({\\bf p})-\\frac{\\tilde{B}({\\bf p})^2}{\\tilde{C}}. \\ee Since $\\chi^2_{\\mu,\\,min}=\\tilde{\\chi}^2_{\\mu,\\,min}$ obviously, we can instead minimize $\\tilde{\\chi}^2_{\\mu}$ which is independent of $\\mu_0$. Note that the above summations are over the 307 Union SNIa compiled in~\\cite{r3} and the 32 calibrated GRBs at $z>1.4$ compiled in Table~I of~\\cite{r11}. On the other hand, the shift parameter $R$ is defined by~\\cite{r15,r16} \\be{eq6} R\\equiv\\Omega_{m0}^{1/2}\\int_0^{z_\\ast} \\frac{d\\tilde{z}}{E(\\tilde{z})}, \\ee where the redshift of recombination $z_\\ast=1090$ which has been updated in~\\cite{r1}, and $\\Omega_{m0}$ is the present fractional energy density of pressureless matter. The shift parameter $R$ relates the angular diameter distance to the last scattering surface, the comoving size of the sound horizon at $z_\\ast$ and the angular scale of the first acoustic peak in the CMB power spectrum of the temperature fluctuations~\\cite{r15,r16}. The value of $R$ has been updated to $1.710\\pm 0.019$ from WMAP5~\\cite{r1}. The distance parameter $A$ is given by \\be{eq7} A\\equiv\\Omega_{m0}^{1/2}E(z_b)^{-1/3}\\left[\\frac{1}{z_b} \\int_0^{z_b}\\frac{d\\tilde{z}}{E(\\tilde{z})}\\right]^{2/3}, \\ee where $z_b=0.35$. In~\\cite{r17}, the value of $A$ has been determined to be $0.469\\,(n_s/0.98)^{-0.35}\\pm 0.017$. Here the scalar spectral index $n_s$ is taken to be $0.960$ which has been updated from WMAP5~\\cite{r1}. So, the total $\\chi^2$ is given by \\be{eq8} \\chi^2=\\tilde{\\chi}^2_{\\mu}+\\chi^2_{CMB}+\\chi^2_{BAO}, \\ee where $\\tilde{\\chi}^2_{\\mu}$ is given in Eq.~(\\ref{eq5}), $\\chi^2_{CMB}=(R-R_{obs})^2/\\sigma_R^2$ and $\\chi^2_{BAO}=(A-A_{obs})^2/\\sigma_A^2$. The best-fit model parameters are determined by minimizing the total $\\chi^2$. As in~\\cite{r18}, the $68\\%$ confidence level is determined by $\\Delta\\chi^2\\equiv\\chi^2-\\chi^2_{min}\\leq 1.0$, $2.3$ and $3.53$ for $n_p=1$, $2$ and $3$, respectively, where $n_p$ is the number of free model parameters. Similarly, the $95\\%$ confidence level is determined by $\\Delta\\chi^2\\equiv\\chi^2-\\chi^2_{min}\\leq 4.0$, $6.17$ and $8.02$ for $n_p=1$, $2$ and $3$, respectively.  In sections~\\ref{sec2}, \\ref{sec3} and~\\ref{sec4}, we consider the joint constraints on single-parameter, two-parameter and three-parameter cosmological models respectively, by using the 307 Union SNIa compiled in~\\cite{r3}, the 32 calibrated GRBs at $z>1.4$ compiled in Table~I of~\\cite{r11}, the updated shift parameter $R$ from WMAP5~\\cite{r1}, and the distance parameter $A$ of the measurement of BAO peak in the distribution of SDSS luminous red galaxies~\\cite{r12} with the updated scalar spectral index $n_s$ from WMAP5~\\cite{r1}. Note that we also present the constraints without GRBs for comparison. A brief summary is given in section~\\ref{sec5}. ",
        "conclusions": "\\label{sec5} In this note, we consider the observational constraints on some cosmological models by using the 307 Union SNIa compiled in~\\cite{r3}, the 32 calibrated GRBs at $z>1.4$ compiled in Table~I of~\\cite{r11}, the updated shift parameter $R$ from WMAP5~\\cite{r1}, and the distance parameter $A$ of the measurement of the baryon acoustic oscillation (BAO) peak in the distribution of SDSS luminous red galaxies~\\cite{r12} with the updated scalar spectral index $n_s$ from WMAP5~\\cite{r1}. The tighter constraints obtained here update the ones obtained previously in the literature (e.g.~\\cite{r13,r14,r18,r23,r28,r29,r30,r31}).  It is worth noting that GRBs are potential tools which might be powerful to probe the cosmic expansion history up to $z>6$ or even higher redshift. Of course, due to the large scatter and the lack of a large amount of well observed GRBs, there is a long way to go in using GRBs extensively and reliably to probe cosmology. The cosmology independent calibration method of GRBs proposed in~\\cite{r10} is an important step towards this end. The works of this note and~\\cite{r11} are beneficial explorations. Along with the accumulation of well observed GRBs with much smaller errors, we believe that a bright future of GRB cosmology is awaiting us. Combining the calibrated GRBs with other probes such as SNIa, CMB, large-scale structure and weak lensing, we can learn more about the mysterious dark energy.    \\begin{center} \\begin{figure}[tbp] \\centering \\includegraphics[width=0.5\\textwidth]{cplanew.eps} \\caption{\\label{fig4} The $68\\%$ and $95\\%$ confidence level contours in the $w_0-w_a$ plane for the CPL model. The results for the cases with and without GRBs are indicated by the black solid lines and the red dashed lines, respectively. The best-fit parameters for the cases with and without GRBs are also indicated by a black solid point and a red box, respectively.} \\end{figure} \\end{center}   \\vspace{-10mm}  %"
    },
    "0809/0809.0327_arXiv.txt": {
        "abstract": "We present new radial velocity measurements of 586 stars in a one-degree field centered on the open cluster  M46, and the planetary nebula NGC 2438 located within a nuclear radius of the cluster. The data are based on medium-resolution optical and near-infrared spectra taken with the AAOmega spectrograph on the Anglo-Australian Telescope. We find a velocity difference of about 30 km~s$^{-1}$ between the cluster and the nebula, thus removing all ambiguities about the cluster membership of the planetary nebula caused by contradicting results in the literature. The line-of-sight velocity dispersion of the cluster is 3.9$\\pm$0.3 km~s$^{-1}$, likely to be affected by a significant population of binary stars. ",
        "introduction": "Any physical associations discovered between planetary nebulae (PNe),  the short-lived but spectacular late evolutionary stage of small and intermediate mass stars (between 1--8 M$_\\odot$), and star clusters would be a valuable discovery that provides a means of establishing  accurate astrophysical parameters for the nebulae through fixing distances and progenitor ages from cluster isochrones. Accurate distances are particularly useful, from which one can infer PNe physical properties such as the absolute magnitude of the central stars, accurate physical dimensions and fluxes. Also, they would provide excellent calibrators for the surface brightness--radius relation (Frew \\& Parker 2006, Frew 2008). Whereas PNe have been found in 4 globular clusters of the Milky Way (M15, M22, Pal 6 and NGC~6441; Jacoby et al. 1997), none has been reported in the literature as an unambiguous member of a much younger open cluster (OC). The interest in the latter case is not only due to being able to determine independent distances to individual nebulae, but also because in a young open cluster the progenitor of a now-visible PN will be a reasonably constrained higher mass star than those in globular clusters. This fact offers the opportunity to calibrate the initial-to-final mass relation of stars on a broad range of masses, usually done by modelling white dwarf populations in  open clusters (e.g. Weidemann 2000, Dobbie et al. 2006).  Recently, Majaess, Turner \\& Lane (2007; MTL07) and Bonatto, Bica \\& Santos Jr (2008; BBS08) have performed detailed investigations of possible physical associations between PNe and OCs. MTL07 considered the cluster membership for 13 PNe that are located in close proximity to open clusters lying in their lines of sight and listed another 16 PNe/open cluster coincidences, which might contain physically associated pairs. However, they noted that we have yet to establish a single association between a PN and an open cluster based on a correlation between their full set of physical parameters, including the three key parameters of radial velocity, reddening, and distance that need to be in good agreement if an association is to be viable. BBS08 used near-infrared colour-magnitude diagrams and stellar radial density profiles to study PN/open cluster association for four pairs. They concluded that the best, but still only probable, cases are those of NGC~2438/M46 and PK~167$-$01/New Cluster 1. However, Parker et al. (in prep.) have uncovered another compelling case in an old open cluster that may prove the best candidate yet for a true OC-PN association.  NGC~2438 is a well-known annular PN located about 8 arcminutes from the core of the bright open cluster M46 (=NGC~2437; BBS08). Despite its brightness, the cluster was relatively unstudied until recently (e.g. Cuffey 1941; Stetson 1981). Recently published cluster parameters are relatively well-determined, e.g. $E(B-V)$=0.10-0.15, D =1.5-1.7 kpc and an age of 220-250 Myr (Sharma et al. 2006, MTL07, BBS08). The estimated turnoff mass is about 3.5 M$_\\odot$ (BBS08). In addition to the possible association with NGC~2438, M46 is also thought to host the well-studied post-AGB candidate OH~231.8+4.2 (Jura \\& Morris 1985).  Early studies of the radial velocity of NGC~2438 and M46 (Cuffey 1941; O'Dell 1963) indicated a difference of $\\Delta v_{\\rm r}\\approx$30 km~s$^{-1}$ between the PN and cluster stars, which suggested that the pair constitues a spatial coincidence only. Three red giants in the cluster have systemic velocities (Mermilliod et al. 1989, 2007) identical to that of cluster dwarf members obtained by Cuffey (1941). However, Pauls \\& Kohoutek (1996) rekindled interest in the possibility of the PN/open cluster association when they found similar velocities for both, although based on a small number of stars. Both MTL07 and BBS08 pointed out the importance of measuring sufficient stellar radial velocities for the cluster and the PN to establish if the proximity is real or only chance superposition.  Previous distance estimates to the cluster and PN  gave similar results though the distance estimates to the OC are far more accurate than the crude statistical estimates for the PN distance provided by Zhang (1995). A distance can be estimated based on the H$\\alpha$ surface brightness -- radius (SB-$r$) relation observed for PNe (e.g. Frew \\& Parker 2006; Frew 2008).  Using an updated version of the relation first presented in Pierce et al. (2004), a distance of 1.4$\\pm$0.4 kpc is estimated.  Using instead the relation applicable to bipolar and bipolar core PNe (Frew, Parker \\& Russeil 2006), D = 1.9 $\\pm$ 0.5 kpc.  These values are in broad agreement with other recent SB-r determinations in the radio domain, e.g. Van de Steene \\& Zijlstra (1995), D = 1.7 kpc, Zhang (1995), D = 2.1 kpc, Phillips (2004), D = 1.2 kpc, and Stanghellini, Shaw \\& Villaver (2008), D=1.2 kpc, and are all compatible with the cluster distances (MTL07, BBS08).  Because of the recently revived interest in possible OC-PN associations, we obtained new multi-object spectroscopic observations of M46 and NGC~2438, where a putative association has still not been satisfactorily resolved due to ambiguities and confusion in the literature. Here we present a statistically significant radial velocity sample of 586 stars within 0.5 degrees of the cluster centre. The data leave no doubt that the PN is not a member of the cluster. ",
        "conclusions": "\\subsection{Cluster membership of the planetary nebula}  The Spitzer image of NGC~2438 extracted from the archive is very reminiscent of that of  M57 with the multiple outer shells of molecular hydrogen seen to extend well ouside the bounds of the well known optical image. Interestingly here, a deep optical H$\\alpha$  image from the SuperCOSMOS H-alpha Survey (Parker et al. 2005), Fig.\\ \\ref{pnobs} clearly show faint optical emission too that matches the inner shell and part of the faint outer shell to  the west seen in the mid-infrared (see also fig.\\ 7 in Corradi et al. 2003). The Spitzer data also reveal evidence of an interesting possible interaction of the molecular material with the ISM at the south-eastern edge of the PN, pointing more or less away from the center of the open cluster.  We present the measured PN emission line velocities in Table\\ \\ref{pnvr}. The numbers show a very good agreement between the northern and the southern edges, suggesting a mean velocity of 78$\\pm$2 km~s$^{-1}$. The only outlier is the P16~8502 line, however, that feature was significantly affected by the residual skylines at $\\lambda\\lambda8504.6-8505.1$ \\AA\\ and their removal during the data reduction. The mean velocity value is also in perfect agreement with the average mid-point of  the double-peaked [O III] lines in the centre (77.5$\\pm$1 km~s$^{-1}$). We therefore adopt $v_{\\rm PN}$=78$\\pm$2 km~s$^{-1}$ as the radial velocity of NGC~2438, which is in excellent agreement with other published velocities in the literature (77 km~s$^{-1}$, Campbell \\& Moore 1918;  74$\\pm$4 km~s$^{-1}$, Meatheringham, Wood \\& Faulkner 1988; 74$\\pm$5 km~s$^{-1}$, Corradi et al. 2000), as well as an unpublished determination of  73$\\pm$6 km~s$^{-1}$ from a SAAO 1.9-m long-slit spectrum (Frew 2008). On the other hand, from the relative velocity difference between the two peaks in the [O III]  profiles we measured an expansion velocity of 21.0$\\pm$0.2 km~s$^{-1}$, virtually identical, for instance, to what Corradi et al. (2000) reported from high-resolution H$\\alpha$ and [O III] spectra (21 km~s$^{-1}$). In summary, our PN measurements draw a picture that agrees exceptionally well with other results in the literature, confirming the reliability of the determined velocities and the quoted uncertainties.  \\begin{table} \\begin{center} \\caption{\\label{pnvr} Measured PN velocities. See the text for a discussion of the uncertainties.} \\begin{tabular}{lrrr} \\hline Line   & PN central & PN north & PN south\\\\ & (km~s$^{-1}$) & (km~s$^{-1}$) & (km~s$^{-1}$)\\\\ \\hline H$\\beta$ 4861 &         &   76.1  &  80.1  \\\\ $[$O III$]$ 4959 & 58.7, 100.2    &  78.6        & 79.8    \\\\ $[$O III$]$ 5007 &  55.6, 98.2   &  77.6   & 77.6   \\\\ P16 8502       &       &  68.1        &  58     \\\\ P15 8545       &       &  75.5      &   76.5      \\\\ P14 8598       &       &  72.1      &   84.4   \\\\ P13 8665        &      &  74.3        &  79.5   \\\\ P12 8750       &       &  73.1      &   76.5   \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure} \\begin{center} \\leavevmode \\includegraphics[width=8cm]{fig5.eps} \\end{center} \\caption{The histogram of stellar radial velocities for M46. The arrow shows the mean center-of-mass velocity of three red giant binaries published by Mermilliod et al. (1989, 2007), while the smooth lines represent the Gaussian fits of the field and the cluster. About half of the observed stars belong to the cluster.} \\label{m46hist} \\end{figure}  Fig.\\ \\ref{m46hist} compares the mean PN velocity to the histogram of the newly derived cluster star velocities that represent the most extensive and accurate such data for M46 obtained to date. We confirm the early results that NGC~2438 has a relative velocity of about 30 km~s$^{-1}$ with respect to the cluster (O'Dell 1963), hence the {nebula is not a bound member of the cluster} despite being located approximately at the same $\\sim$1.5-1.7 kpc distance (MTL07, BBS08). In Fig.\\ \\ref{m46hist} we also put an arrow at the mean center-of-mass velocity (48.5 km~s$^{-1}$) of three red giant binaries  measured by Mermilliod et al. (1989, 2007). The excellent agreement between the maximum of the histogram (49 km~s$^{-1}$ for the highest value, 48 km~s$^{-1}$ for the centroid of the fitted Gaussian) and the very accurate CORAVEL data confirms both the cluster membership of those systems and the quoted accuracy of our single-epoch velocity measurements. We also note that while this paper was under review, a new cluster radial velocity was published by Frinchaboy \\& Majewski (2008), who measured +46.9$\\pm$1.0 km~s$^{-1}$ from 19 member stars. The nice agreement gives further support to our results.  We examined the possibility that the progenitor of NGC~2438 was a runaway star that has been ejected from M46, since it is the only option for maintaining possible physical association of the cluster and the PN. Runaway stars escape from star clusters with velocities of typically 30 km~s$^{-1}$ (Blaauw 1961; Hoogerwerf, de Bruijne \\& de Zeeuw 2001) and one cannot ad hoc exclude the possibility that something similar has happened in M46. However, we feel this to be quite unlikely. If we assume a binary-supernova scenario as a possible explanation, then the SN must have happened 170-200 Myr ago (i.e. the cluster age minus the evolutionary age of a progenitor star of 8 M$_\\odot$, the lower limit for a core-collapse SN). A star travelling at 30 km~s$^{-1}$ for this time traverses 5-6 kpc in the absence of other forces, so that it should have left the apparent vicinity of the cluster a long time ago, unless its radial velocity is almost perfectly aligned along the line-of-sight. The other possibility could be a very recent ejection, most likely within the last couple of million years, as part of a binary-binary dynamical ejection event. However, the dynamical ejection scenario (DES) requires a high-density environment in which close encounters of binaries can happen more frequently: young open clusters or OB  associations forming massive stars (Hoogerwerf, de Bruijne \\& de Zeeuw 2001). M46 is neither young nor has an exceptional high density compared to other intermediate-age open clusters. Ultimately, future proper motion measurements could be used to distinguish between the option of chance projection and the option of runaway star.  In summary, while past membership of the progenitor cannot be completely ruled out, we can safely conclude against present membership of the PN in the cluster.  \\subsection{Evidence of a significant binary population of the cluster}  We also noted the surprisingly large velocity dispersion of the cluster, indicated by the broad cluster peak of the histogram. While some of it might be explained by the degraded velocimetric accuracy for the hotter stars (cf. Sect. 3.1), the full range between 35 km~s$^{-1}$ and 60 km~s$^{-1}$ seems to be too large as purely due to measurement errors. We quantified the velocity dispersion by fitting two Gaussians to the histogram, one broad and one narrow to represent the galactic background and the cluster members, respectively. The resulting line-of-sight velocity dispersion of the cluster is $\\sigma_{\\rm los}=$3.9$\\pm$0.3 km~s$^{-1}$ (removing 78 stars with $T_{\\rm eff}>$10000 K changes the result to $\\sigma_{\\rm los}=$3.8$\\pm$0.2 km~s$^{-1}$). This is a very high value for a relatively old open cluster and suggests the presence of many binaries. The velocity dispersion of the broad Gaussian is 18.4 km~s$^{-1}$, which is consistent with field stars predominantly in the thin disk  (Veltz et al. 2008).  As a simple exercise, we estimated the dynamical mass of the cluster using the  equation $M_{\\rm dyn}=\\eta R_{\\rm hl} \\sigma_{\\rm los}^2 /G,$ where $R_{\\rm hl}$ is the half-light radius, $G$ is the gravitational constant, and $\\eta$ is a dimensionless constant (Spitzer 1987). With this we also make the assumption that the cluster is in virial equilibrium, which is most likely the case, given that clusters older than $\\sim$50 Myr are expected to be in virial equilibrium (e.g. Goodwin \\& Bastian 2006). A rough estimate for the half-light radius of M46 can be taken as the core radius  of 3.3 pc (Sharma et al. 2006). Similarly, about half of the likely members in our spectroscopic sample lie within 3.4 pc to the cluster centre. Assuming $R_{\\rm hl}$=3.3 pc and the canonical $\\eta=9.75$, the result is $M_{\\rm dyn}=$1.1$\\times$10$^5$ M$_{\\odot}$.  On the other hand, adopting $V_{\\rm tot}$=6.1 mag from SIMBAD and $V-M_{\\rm V}\\approx$11.2 mag, the resulting  absolute brightness $M_V=-5.1$ mag corresponds to $L_{\\rm V}/L_\\odot\\approx$8800, leading to $L_{\\rm V}/M_{\\rm dyn}\\lesssim$0.1, which is way too low for any stellar cluster (see, e.g., fig. 5 in Goodwin \\& Bastian 2006). Therefore, the dynamical mass must be grossly overestimated, which is a well-known effect when a significant binary population exists in a cluster (for a recent study see, for instance, Kouwenhoven \\& de Grijs 2008). Other evidence that points in this direction was mentioned by Sharma et al. (2006), who noted that the broad main-sequence in the colour-magnitude diagram could be due to the presence of binary stars. In that case orbital motion can introduce extra velocity scatter, which can lead to significantly overestimated  dynamical masses.  One caveat is the neglected effect of the radial velocity errors on the velocity  dispersion. If the real errors are much larger than the quoted formal errors from the cross-correlation (cf. Sect. 3.1), then the velocity dispersion could be dominated by measurement errors, hence getting only an upper limit to $\\sigma_{\\rm los}$ and $M_{\\rm dyn}$ and a lower limit on $L_{\\rm V}/M_{\\rm dyn}$. Our experiences with the same instrument and data on other, mostly globular, star clusters (Kiss et al. 2007; Sz\\'ekely et al. 2007) showed that high S/N AAOmega spectra can reproduce published radial velocities with the quoted $\\pm$1--2 km~s$^{-1}$ accuracy for stars cooler than $\\sim$6000 K. For M46, we have neither repeated observations nor independent measurements with other instruments to check the precision and the accuracy of the velocities, but with the good S/N for most of the data, we are confident that the dispersion is not dominated by the measurement errors."
    },
    "0809/0809.2608_arXiv.txt": {
        "abstract": "We present a study of 15 long-duration $\\gamma$-ray burst (GRB) host galaxies at $z>2$.  The GRBs are selected with available early-time afterglow spectra in order to compare interstellar medium (ISM) absorption-line properties with stellar properties of the host galaxies.  In addition to five previously studied hosts, we consider new detections for the host galaxies of GRB\\,050820 and GRB\\,060206 and place 2-$\\sigma$ upper limits to the luminosities of the remaining unidentified hosts.  We examine the nature of the host galaxy population and find that (1) the UV luminosity distribution of GRB host galaxies is consistent with expectations from a UV luminosity weighted random galaxy population with a median luminosity of $\\langle L(UV)\\rangle=0.1\\,L_*$; (2) there exists a moderate correlation between UV luminosity and Si\\,II $\\lambda\\,1526$ absorption width, which together with the observed large line widths of $W(1526)> 1.5$ \\AA\\ for a large fraction of the objects suggests a galactic outflow driven velocity field in the host galaxies; (3) there is tentative evidence for a trend of declining ISM metallicity with decreasing galaxy luminosity in the star-forming galaxy population at $z=2-4$; (4) the interstellar UV radiation field is found $\\approx 35-350\\times$ higher in GRB hosts than the Galactic mean value; and (5) additional galaxies are found at $\\aapl 2''$ from the GRB host in all fields with known presence of strong Mg\\,II absorbers, but no additional faint galaxies are found at $\\aapl 2''$ in fields without strong Mg\\,II absorbers.  Our study confirms that the GRB host galaxies (with known optical afterglows) are representative of unobscured star-forming galaxies at $z>2$, and demonstrates that high spatial resolution images are necessary for an accurate identification of GRB host galaxies in the presence of strong intervening absorbers. ",
        "introduction": "Gamma-ray bursts (GRBs) are among the most energetic events in the universe.  In particular, long-duration GRBs are believed to originate in the catastrophic death of massive stars (e.g.\\ Woosley 1993; Paczy\\'nski 1998; Bloom \\etal\\ 2002; Stanek \\etal\\ 2003; Hjorth \\etal\\ 2003; see Woosley \\& Bloom 2006 for a recent review).  Since massive stars evolve rapidly, long-duration GRBs should probe instantaneous star formation out to the highest redshifts (e.g.\\ Wijers \\etal\\ 1998) with the afterglows serving as signposts to starburst galaxies in the distant universe.  Many bursts are followed by optical afterglows that can briefly exceed the absolute brightness of any known quasar by orders of magnitude (e.g.\\ Akerlof \\etal\\ 1999; Kann \\etal\\ 2007; Bloom \\etal\\ 2008) and serve as bright background sources for probing intervening gas along the line of sight.  Early-time, high-resolution spectroscopy of GRB afterglows have revealed numerous absorption features produced by ground-state and excited-state ions in the interstellar medium (ISM) of the host galaxies (e.g.\\ Vreeswijk \\etal\\ 2004; Prochaska, Chen, \\& Bloom 2006; Vreeswijk \\etal\\ 2007; D'Elia \\etal\\ 2008). Detailed studies based on comparisons of absorption-line strengths have yielded accurate constraints on the host ISM properties, including gas density, temperature, chemical composition, and kinematics of the GRB host environment (e.g.\\ Fynbo \\etal\\ 2006; Savaglio 2006; Prochaska \\etal\\ 2007a).  Specifically, roughly 50\\% of known GRBs at $z>2$ are found in ISM of neutral gas column density $N(\\hI)>10^{21}$ \\cmjj\\ (Jakobsson \\etal\\ 2006; Chen \\etal\\ 2007a). In addition, the host ISM generally exhibit moderate chemical enrichment with a median metallicity of $> 1/10$ solar, although with a substantial scatter over the range from $1/100$ to $\\sim 1/2$ solar values (Fynbo \\etal\\ 2006a; Savaglio 2006; Prochaska \\etal\\ 2007a).  Comparisons of different ionic abundances also show that there exists a large differential depletion with $[{\\rm Zn}/{\\rm Fe}]> +0.6$ dex and $\\alpha/{\\rm Fe} \\aapg 0.4$ dex, confirming the presence of a large amount of gas mass and a chemical enrichment history dominated by massive stars (Savaglio 2006; Prochaska \\etal\\ 2007a).  Finally, there is a lack of molecular gas despite the presence of a large $N(\\hI)$ (Fynbo \\etal\\ 2006; Tumlinson \\etal\\ 2007).  Interpretations of these absorption-line data are not straightforward.  In particular, the observed low-metal content in the GRB host ISM, as opposed to optically selected luminous starburst galaxies (e.g.\\ Shapley \\etal\\ 2004), can be explained if the progenitor stars originate in the outskirts of a luminous galaxy or if the host galaxies are underluminous and have on average lower metallicities.  A local-galaxy analogue of this is seen with the so-called luminosity--metallicity relation (e,g.\\ Tremonti \\etal\\ 2004 for nearby galaxies).  Recently, Fynbo \\etal\\ (2008) considered both scenarios and showed that the metallicity distribution of GRB hosts is consistent with the expectation that these host galaxies represent a weighted star-forming galaxy population according to the on-going star formation rate (SFR).  In addition, the large atomic gas column density in contrast to the lack of molecular gas may be due to either an enhanced UV radiation field in the star-forming regions near the GRB progenitor or a relatively low dust and metal content indicated in the absorption-line data (e.g.\\ Tumlinson \\etal\\ 2007; Whalen \\etal\\ 2008).  Recent detections of the 2175-\\AA\\ dust absorption feature in GRBs\\,070802 (Kr\\\"uhler \\etal\\ 2008; Eliasdottir \\etal\\ 2008) and 080805 (Jakobsson \\etal\\ 2008) offer an important test for this scenario.  Regardless, these issues have direct impact on our understanding of both the GRB progenitors and star formation physics. Supplemental imaging and spectroscopic observations of the host galaxies are necessary for accurate interpretations of the absorption-line measurements.  The transient nature of optical afterglows allows deep imaging and spectroscopic studies of galaxies close to the lines of sight, including the hosts, when the afterglows disappear (c.f.\\ M{\\o}ller \\etal\\ 2002a; Chen \\& Lanzetta 2003 for searches of damped \\lya\\ absorbing galaxies along quasar sightlines).  At $z>2$, where accurate ISM abundance measurements are available based on afterglow absorption-line spectroscopy, only four host galaxies have been unambiguously identified (see Savaglio \\etal\\ 2008 for a compilation). Comparison studies between known ISM properties from afterglow absorption spectroscopy and the observed morphology and luminosity of the host galaxies have been published individually for GRB\\,000926 (Castro \\etal\\ 2003), GRB\\,011211 (Vreeswijk \\etal\\ 2006), GRB\\,021004 (Fynbo \\etal\\ 2005), and GRB\\,030323 (Vreeswijk \\etal\\ 2004).  The four GRB host galaxies together show an order-of-magnitude scatter in their rest-frame UV luminosity and ISM metallicity.  At $z<2$, where the majority of common ISM absorption features occur at ultraviolet wavelengths and spectroscopic observations become challenging on the ground, $>40$ GRB host galaxies have been identified in late-time imaging follow-up.  These known host galaxies provide important insights for understanding the nature of GRB progenitors at intermediate redshifts.  First, Chary \\etal\\ (2002) compared the estimated SFR and total stellar mass for 12 GRB host galaxies and found that the host galaxies have on average higher SFR per unit stellar mass than local starburst galaxies.  Le Floc'h \\etal\\ (2003) examined the optical and near-infrared $R-K$ colors and rest-frame $B$-band luminosity function of 15 GRB host galaxies at $\\langle z\\rangle\\approx 1$. They found that these host galaxies have on average bluer colors and fainter luminosity ($\\langle L_{B}\\rangle \\approx 0.1\\,L_*$) than random star-forming galaxies at the same redshift range.  Additional mid-infrared imaging observations by these authors showed that this is not due to dust extinction (Le Floc'h \\etal\\ 2006).  Similarly, Christensen \\etal\\ (2004) studied the optical and near-infrared spectral energy distributions (SEDs) of ten GRB host galaxies at $\\langle z\\rangle=0.85$. They found based on a comparison with $\\sim 1000$ galaxies identified at a similar redshift range in the Hubble Deep Fields that GRB host galaxies have on average younger stellar age and shorter characteristic star-forming time scale.  In addition, Fruchter \\etal\\ (2006) compared HST images of $>40$ GRB host galaxies at $\\langle z\\rangle\\approx 1$ with the host galaxies of core-collapse supernovae (SNe).  They found that the majority of GRB host galaxies exhibit irregular morphology, unlike the host galaxies of core-collapse SNe. Recently, independent studies by Castro Cer\\'on \\etal\\ (2008) and Savaglio \\etal\\ (2008) show that $\\langle z\\rangle\\aapl 1$ GRB host galaxies contain on average lower stellar mass than field star-forming galaxies.  Together, these results show that GRB host galaxies at $\\langle z\\rangle\\approx 1$ represent a relatively young dwarf population that have experienced recent on-going star-forming episodes.  While GRB host galaxies based on the intermediate-redshift sample appear to be underluminous and low mass systems, it is not clear whether the long-duration GRBs originate preferentially in relatively metal dificient star-forming regions (c.f.\\ Wolf \\& Podsiadlowski 2007; Modjaz \\etal\\ 2008). A low-metallicity environment is favored by popular progenitor models so that the progenitor star can preserve high spin and a massive stellar core to produce a GRB (e.g.\\ Hirschi \\etal\\ 2005; Yoon \\& Langer 2005; Woosley \\& Heger 2006). Chemical abundance measurements for $z\\aapl 2$ galaxies have been based primarily on emission line observations of associated H\\,II regions. A subset of the host galaxies at $z<1$ have been followed up spectroscopically for measuring emission-line fluxes.  A mean metallicity of roughly $1/4$ solar is found but with a large scatter (e.g.\\ Sollerman \\etal\\ 2005; Modjaz \\etal\\ 2008; Savaglio \\etal\\ 2008).  The accuracy of emission-line-based abundance estimates depends sensitively on the accuracy of the calibrations between different line diagnostics (e.g.\\ Kewly \\& Dopita 2002; Skillman \\etal\\ 2003; Kennicutt \\etal\\ 2003).  Whether or not there exists a maximum metallicity for forming a GRB remains an open question.  We have carried out an optical and near-infrared imaging survey of fields around 15 GRBs at $z>2$.  The GRBs are selected to have early-time afterglow spectra in order to compare ISM absorption-line properties with stellar properties.  The goal is to identify the host galaxies and constrain their rest-frame UV and optical luminosities. The primary objectives are (1) to quantify the luminosity distribution of the GRB host galaxy populations and investigate whether or not the GRB host galaxies trace the typical star-forming galaxies at high redshift and (2) to examine whether there exists a correlation between the ISM metal content and host luminosity.  The starburst nature of GRB hosts makes this galaxy sample a unique laboratory for studying star formation physics and stellar feedback at high redshift.  We adopt a $\\Lambda$CDM cosmology, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda = 0.7$, with a dimensionless Hubble constant $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ throughout the paper.  \\setcounter{footnote}{0} ",
        "conclusions": "We have carried out an optical and near-infrared imaging survey of the fields around 15 GRBs at $z>2$.  The GRBs are selected to have early-time afterglow spectra in order to compare ISM absorption-line properties with stellar properties.  The redshifts of the GRBs span a range from $z=2.04$ to $z=4.05$, and the neutral hydrogen column densities of the GRB host ISM span a range from $\\log\\,N(\\hI)=16.9$ to $\\log\\,N(\\hI)=22.6$ (Figure 1).  In addition to the five previously known GRB host galaxies, we report new detections for the host galaxies of GRB\\,050820 and GRB\\,060206. The seven identified GRB host galaxies have rest-frame UV absolute magnitudes spanning from $M_{AB}(U)-5\\,\\log\\,h=-15.2$ to $M_{AB}(UV)-5\\,\\log\\,h=-19.8$ mag.  For sources with a known dust-to-gas ratio from afterglow absorption-line analysis, we correct the rest-frame UV magnitude for dust extinction assuming an SMC extinction law.  The inferred SFR spans a range from 0.6 to 3.8 $h^{-2}$ M$_\\odot$ yr$^{-1}$.  We are able to constrain the underlying stellar populations for three host galaxies, GRB\\,000926, GRB\\,021004, and GRB\\,050820, based on comparisons of the observed optical and near-infrared broad-band colors and a suite of stellar population synthetic models .  We estimate total stellar masses of between $M_*=(0.9-2.1)\\times 10^9\\ h^{-2}$ M$_\\odot$ for these hosts.  We also place 2-$\\sigma$ upper limits for the rest-frame luminosities of the remaining eight GRB host galaxies based on the depths in available optical and near-infrared images.  Finally, high spatial resolution images from the HST allow us to deblend the GRB host galaxies from foreground absorbers and to unveil a range of rest-frame UV morphology between compact (e.g.\\ GRB\\,060206) and extended (e.g.\\ GRB\\,050820) emission features of the hosts.  A summary of known absorption-line properties and stellar properties of the 15 GRB host galaxies is presented in Table 3.  Combining early-time, high-resolution afterglow spectra and the results of late-time imaging survey of the GRB fields allows us to address a number of issues regarding both the nature of GRB progenitor environment and star-forming physics in distant starburst galaxies.  \\subsection{The Luminosity Distribution of GRB Host Galaxies}  First, we examine the luminosity distribution of GRB host galaxies above $z=2$ based on the survey result of our sample.  The goal is to characterize the nature of GRB host galaxies based on comparisons of their luminosity distribution and the luminosity distribution of a randomly selected sample from the field galaxy population.  We incorporate the non-detections in our follow-up imaging survey by evaluating the cumulative maximum fraction ${\\cal F}_{\\rm max}$ of GRB host galaxies that are fainter than a given UV absolute magnitude $M_{\\rm max}(UV)$ (solid histogram in Figure 19).  For the host of GRB\\,030429, we have a constraint only for the rest-frame $B$-band magnitude.  We infer its corresponding UV magnitude based on the mean color of $\\langle {\\rm UV}-B\\rangle=1.22$ (with an r.m.s. scatter of 0.3 mag) observed for luminous starburst galaxies at $z=2-3$ in Shapley \\etal\\ (2005).  For the host of GRB\\,060607, we have estimated limiting magnitudes both in the rest-frame UV and $B$ bands.  We adopt the more sensitive limit based on the conversion of ${\\rm UV}-B=1.22$\\footnote{We note that while there is no apparent trend between $UV-B$ and $M_{UV}$ in the luminous starburst sample of Shapley \\etal\\ (2005), it is possible that that fainter dwarf starburst galaxies may be bluer.  Given that only two fields (GRB\\,030429 and GRB\\,060607) are affected by this conversion, we find that the results presented in this section are not sensitive to the adopted $UV-B$ color.}.  The empirical observations of the sample of 15 GRB fields confirm the previous understanding for $z\\sim 1$ GRB hosts (e.g.\\ Le Floc'h \\etal\\ 2003): as much as 70\\% of long-duration GRBs may originate in galaxies fainter than $0.1\\,L_*$.  At $z=2-4$, the field galaxy population is now well characterized by a Schechter luminosity function $\\phi(M)$ at rest-frame UV wavelengths with $M_{AB*}-5\\,\\log\\,h=-20.2\\pm 0.1$, $\\phi_*=(4\\pm 0.6)\\times 10^{-3}\\ h^{3}\\,{\\rm Mpc}^{-3}$, and a faint-end slope $\\alpha=-1.7\\pm 0.1$ (Bouwens \\etal\\ 2007; Reddy \\etal\\ 2008).  If the GRB host galaxies are representative of the field galaxy population, then we expect that the fraction of host galaxies found in a luminosity interval is proportional to the space density of galaxies in the luminosity range.  The expected cumulative maximum fraction of the host galaxies versus UV magnitude can be estimated following \\begin{equation} {\\cal F}_{\\rm max}[M>M_{\\rm max}]\\propto \\left [ \\sum_{i=1}^{n} \\phi(M_i)\\,dM + \\sum_{j=1}^{m} \\int_{M_{min}}^{M_{\\rm max}} \\phi(M) dM \\right ], \\end{equation} where the first term extends over $n$ known host galaxies that have $M_{AB}(UV)>M_{\\rm max}$, the second term extends over $m$ unidentified host galaxies that may be as luminous as $M_{\\rm max}$, and $\\phi(M)$ is the galaxy luminosity function.  The luminosity distribution is expected to resemble the galaxy luminosity function with a dominant fraction attributed to faint dwarf galaxies (e.g.\\ Jakobsson \\etal\\ 2005; Fynbo \\etal\\ 2008).  We test this hypothesis by calculating ${\\cal F}_{\\rm max}$ for a sample of 15 random galaxies that share the known luminosities of the seven identified GRB host galaxies and the empirical 2-$\\sigma$ upper limits for the remaining eight GRB host galaxies.  We experiment with different $M_{\\rm min}$ (with corresponding minimum luminosity $L_{\\rm min}$) in Equation (1).  The results for $L_{\\rm min}=0.01\\,L_*$ and for $L_{\\rm min}=0.005\\,L_*$ are shown as the dashed curves in Figure 19.  Smaller $L_{\\rm min}$ models predict a larger contributions from fainter galaxies.  The model expections for different adopted $L_{\\rm min}$ clearly deviate from the observed distribution.  The hypothesis that the host galaxies of long-duration GRBs trace random field galaxies is rejected at $> 99$\\% confidence level based on a Kolmogorov--Smirnov (KS) test.  Next, under the assumption that GRBs trace instantaneous star formation, galaxies with higher on-going SFR are expected to have higher probability to host a GRB event.  Applying the rest-frame UV luminosity as a measure of the on-going SFR, we modify Equation (1) to include a UV luminosity weighting, $\\int L\\,\\phi(L)\\,d\\,(L)$, for calculating the expected ${\\cal F}_{\\rm max}$, \\begin{eqnarray} {\\cal F}_{\\rm max}[M>M_{\\rm max}] &\\propto& \\sum_{i=1}^{n} 10^{-0.4\\,(M_i-M_*)}\\,\\phi(M_i)\\,dM \\nonumber \\\\ &+& \\sum_{j=1}^{m} \\int_{M_{min}}^{M_{\\rm max}} 10^{-0.4\\,(M-M_*)}\\,\\phi(M) dM . \\end{eqnarray} The results are shown as solid and dotted curves in Figure 19.  The observations are best described under this hypothesis for $L_{\\rm min}=0.001\\,L_*$ (solid curve) based on a KS test, but we cannot rule out models with $L_{\\rm min}=0.01\\,L_*$ (lower dotted curve) or $L_{\\rm min}=0.0001\\,L_*$ (upper dotted curve).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.3,angle=270]{magUV_mod.eps} \\end{center} \\caption[]{Cumulative distribution of the maximum fraction of GRB host galaxies that are fainter than a given UV absolute magnitude (solid histogram), based on the seven identified host galaxies and eight fields with upper limits in our imaging sample.  The curves show model expectations based on different hypothesis for the origin of GRB host galaxies.  We consider two scenarios: (1) GRB host galaxies are representative of the general galaxy population and (2) GRB host galaxies originate preferencially in galaxies of higher star formation rate.  Adopting the best-fit UV luminosity function of star-forming galaxies at $z=3-4$ from Bouwens \\etal\\ (2007) and Reddy \\etal\\ (2008), we show the expectations for scenario (1) in dashed curves for two different minimum luminosity cutoffs, $L_{\\rm min}=0.01\\,L_*$ (bottom) and $L_{\\rm min}=0.005\\,L_*$ (top).  The expectation for scenario (2) with $L_{\\rm min}=0.001\\,L_*$ is shown in the solid curve, and the dotted curves are for $L_{\\rm min}=0.01\\,L_*$ (lower) and $L_{\\rm min}=0.0001\\,L_*$ (upper).  The comparisons provide empirical evidence supporting the expectation that GRB host galaxies form an SFR weighted field galaxy sample.} \\end{figure}  In summary, the sample of seven known GRB host galaxies and eight upper limits for unidentified hosts allows us to determine the cumulative maximum fraction of the $z>2$ host galaxy population as a function of rest-frame UV magnitude.  Adopting rest-frame UV luminosity as a measure of on-going SFR, we find that the empirical sample is best described by a SFR-weighted sample of the field galaxy population. Models that do not include SFR weighting can be ruled out at $>99$\\% confidence level.  Based on the best-fit SFR-weighted model, we estimate a median luminosity for the GRB host galaxies at $\\approx 0.1\\,L_*$.  A similar analysis has been presented in Jakobsson \\etal\\ (2005), who derived constraints for the luminosity function of GRB host galaxies under the assumption that the observed brightness distribution of the hosts follows a luminosity weighted field galaxy population. Our study differs from the approach of Jakobsson \\etal\\ in that our analysis is based on a uniform set of photometric measurements from our own imaging survey.  Then we adopt the known UV luminosity function of the field galaxy population (with a faint-end slope of $\\alpha=-1.7$) and examine different hypotheses for generating the observed GRB host galaxy sample.  We confirm that the host galaxy population is representative of a UV luminosity weighted sample.  The difference between the host luminosity function of Jakobsson \\etal\\ and the best-fit luminosity of star-forming galaxies at $z>2$ may be due to uncertainties in published photometric data of the host galaxies in their sample and uncertainties in the line-of-sight absorption properties.  We note that although the selection criterion of our GRB sample is based on available afterglow spectra that presumably includes only GRBs with relatively bright optical afterglows, the broad range in the isotropic energy release of the GRBs (see Table 1 in Chen \\etal\\ 2007a) indicates that the GRBs in our sample are not an overly biased portion of the long-duration GRB population.  That is, we have selected a representative subsample of the unobscured GRB host galaxy population.  However, the presence of 20\\% dark bursts that do not have optical afterglows found (e.g.\\ Tanvir \\& Jakobsson 2007) suggests that some GRBs originate in heavily obscured star-forming regions that are not included in our sample and are likely missed in the UV selected field galaxy sample as well.  To constrain the fraction of dust obscured star-forming galaxies requires an independent study of the fields around dark bursts.   \\subsection{The Luminosity--Metallicity Relation in GRB Host Galaxies}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4,angle=0]{gdla_w1526.eps} \\end{center} \\caption[]{Absolute UV magnitude of GRB host galaxies versus rest-frame absorption equivalent width of the Si\\,II $\\lambda\\,1526$ transition in the hosts as observed in afterglow spectra.  For this study, we consider only host DLAs because the lyman limit absorbers do not include neutral ISM in the host galaxies.  The Si\\,II $\\lambda\\,1526$ transition is typically saturated at rest-frame absorption equivalent width $W(1526)>0.3$ \\AA.  The line width therefore provides a measure of the velocity field of cool clouds along the line of sight through a halo (e.g.\\ Prochaska \\etal\\ 2008). We find a moderate trend with stronger $W(1526)$ appearing in more luminous hosts at nearly 95\\% confidence level.} \\end{figure}  We combine known absorption-line properties with estimated galaxy UV luminosity to investigate the gas kinematics and chemical enrichment in the ISM of GRB host galaxies.  The goals of this study are (1) to examine the physical processes that determine the observed absorption-properties in the host ISM, (2) to investigate whether there is a metallicity cutoff in GRB host galaxies as favored by various theoretical models (e.g.\\ Hirschi \\etal\\ 2005; Yoon \\& Langer 2005; Woosley \\& Heger 2006), and (3) to probe the luminosity--metallicity relation below the magnitude limit of most previous studies at $z>2$ (e.g.\\ Erb \\etal\\ 2006; Maiolino \\etal\\ 2008).  We first compare rest-frame absorption equivalent widths of the Si\\,II $\\lambda\\,1526$ transition $W(1526)$ observed in the host with the absolute UV magnitude.  For this study, we consider only host DLAs because the lyman limit absorbers do not include neutral ISM in the host galaxies.  The Si\\,II $\\lambda\\,1526$ transition is typically saturated at $W(1526)>0.3$ \\AA.  The line width provides a measure of the velocity field of cool clouds along the line of sight through a galactic halo, rather than the total gas column density (e.g.\\ Prochaska \\etal\\ 2008).  We adopt the $W(1526)$ measurements published in Prochaska \\etal\\ (2008).  Including additional measurement for GRB\\,060206 in public afterglow spectra from Subaru Science Data Archive (Aoki \\etal\\ 2006), we have assembled eight GRB hosts galaxies for this study.  Figure 20 shows that there exists a relatively significant trend with stronger Si\\,II transitions appearing in more luminous hosts. Including upper limits, we find based on a generalized Kendall test that the probability of a positive correlation between $M_{UV}$ and $W(1526)$ is nearly 95\\%.  While the positive correlation is similar to the mass-metallicity relation seen in field galaxies (e.g.\\ Erb \\etal\\ 2006; Maiolino \\etal\\ 2008) , we note that half of the host galaxies have $W(1526)\\aapg 1.5$ \\AA, corresponding to a velocity width of $\\sim 300$ \\kms.  Attributing the large line width to gravitational motion of clouds within the host dark matter halos would require a halo mass $> 10^{11}\\ h^{-1}\\,{\\rm M_\\odot}$, comparable to the mass scale found for halos that host luminous starburst galaxies at $z\\sim 2$ (e.g.\\ Conroy \\etal\\ 2008) but at odds with expectations for starburst galaxies of low stellar mass and low luminosity.  But because $M_{UV}$ serves as a measure of the on-going SFR, the observed $M_{UV}$ versus $W(1526)$ correlation implies that the velocity field observed in GRB host galaxies is driven by galactic outflows.   Next, we examine whether there exists a correlation between host luminosity and ISM metallicity, in comparison to what is known for luminous starburst galaxies published by Erb \\etal\\ (2006).  We note two apprarent caveats in our study.  First, chemical abundances of the GRB host galaxies in our sample are determined for the cold neutral medium using absorption-line techniques, whereas the metallicities of field galaxies are determined from the integrated emission-line fluxes of their H\\,II regions.  Although extensive studies have yet to be undertaken to examine possible systematic differences between absorption- and emission-line abundance measurements, available evidence based on limited studies of nearby starburst galaxies have yielded consistent metallicity measurements using either absorption-line or emission-line techniques (see Russell \\& Dopita 1992 and Welty \\etal\\ 1997, 1999 for the Large and Small Magellanic Clouds; Lecavelier des Estangs \\etal\\ 2004 for I\\,Zw 18; and Schulte-Ladbeck \\etal\\ 2005 and Bowen \\etal\\ 2005 for SBS\\,1543$+$593).  In the few cases where emission and absorption line measurements have been compared in galaxies at $z > 2$, the two methods have been found to give consistent answers to within a factor of $\\sim 2$ (e.g.\\ Pettini 2006).  Second, absorption-line measurements are for the ISM along the afterglow line of sight, whereas emission-line observations are integrated measurements averaged over the entire galaxies.  A large metallicity gradient is commonly seen in nearby galaxies with a slope varying between $d\\,Z/d\\,r \\approx -0.02$ dex / kpc to $-0.07$ dex / kpc (e.g.\\ Zaritsky \\etal\\ 1994; van Zee \\etal\\ 1998; Kennicutt \\etal\\ 2003).  Therefore, a metallicity measurement for individual sightlines may not be representative of the global mean value of the entire host galaxy.  We note, however, that $z>2$ galaxies are relatively compact with typical half light radii of $\\approx 2$ kpc (Bouwens \\etal\\ 2004; Law \\etal\\ 2007).  Previous observations (e.g.\\ Bloom \\etal\\ 2002; Fruchter \\etal\\ 2006) and current findings for GRB\\,050820 and GRB\\,060206 show that the GRBs occur within $\\aapl 2$ kpc radius of their host galaxies.  This is also consistent with the large $N(\\hI)$ observed in the afterglow spectra with the exception of four non-DLAs. Spatial variation of the observed metallicity across distant star-forming galaxies is therefore not expected to exceed 0.2 dex.  We proceed with an analysis that compares the ISM metal content of GRB galaxies with those of known starburst galaxies at $z\\approx 2$.  Figure 21 shows the the luminosity--metallicity relation reproduced from Erb \\etal\\ (2006) for starburst galaxies at $\\langle z\\rangle=2.3$ (open squares).  The oxygen abundances are measured using the $N2$ index, which is a measure of the flux ratio between the observed [N\\,II] $\\lambda\\,6584$ to H$\\alpha$ lines\\footnote{See Erb \\etal\\ 2006 for a discussion of possible systematic uncertainties associated with the $N2$ index.  While it is known that the $N2$ index saturates at solar metallicity, the saturation does not affect our comparison because GRB host galaxies appear to have mostly sub-solar abundances (e.g.\\ Prochaska \\etal\\ 2007a).}.  The rest-frame $B$-band magnitudes published in Erb \\etal\\ have been converted to UV magnitude according to $UV=B+1.22$ to facilitate comparisons with the GRB host galaxy population and with predictions from numerical simulations (crosses in Figure 21, reproduced from Fynbo \\etal\\ 2008).  The $UV-B$ color conversion is estimated based on the mean $R-K$ colors for these galaxies in Shapley \\etal\\ (2005).  We include in Figure 21 measurements for nine GRB host galaxies for which measurments (solid points) or constraints (open circles) on the ISM metallicity are available.  Only known host DLAs are considered in the comparison here.  Left arrows represent the fields (GRB\\,050401, GRB\\,050922c and GRB\\,050730) for which no hosts have been found and 2-$\\sigma$ upper limits to the intrinsic luminosity are shown.  The lower limits represent those GRB fields for which only moderate resolution afterglow spectra are available and the observed absorption-line strengths represent a lower limit to the intrinsic values (e.g.\\ Prochaska 2006).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4,angle=0]{magz_hiz.eps} \\end{center} \\caption[]{The luminosity--metallicity relation in distant galaxies.  Solid points with errorbars are GRB host galaxies with accurate absorption-line metallicity measurements, while open circles are GRB host galaxies with constraint on their ISM metallicities due to either insufficient spectral resolution in the afterglow spectra or unknown ionization fraction of the ISM.  Open squares are emission-line metallicity measurements for luminous starburst galaxies at $z\\approx 2$ from Erb \\etal\\ (2006).  The oxygen abundance was evaluated using the observed [N\\,II] to H$\\alpha$ line ratio.  We have converted the rest-frame $B$-band magnitude to UV magnitude according to $UV=B+1.22$, which is estimated based on the mean $R-K$ colors for these galaxies in Shapley \\etal\\ (2005).  Crosses represent the predicted luminosity--metallicity relation for starburst galaxies at $z\\sim 3$ (Fynbo \\etal\\ 2008) based on numerical simulations by Sommer-Larsen \\& Fynbo (2008).  } \\end{figure}  Three interesting features are seen in Figure 21.  First, the majority of the GRB host galaxies are fainter than the faintest starburst galaxies studied in magnitude-limited surveys (open squares) and offer a promising probe to extend the studies of ISM metal enrichment to fainter luminosity limits than existing faint galaxy surveys (Djorgovski \\etal\\ 2004).  Second, the metallicities found in GRB host galaxies span two orders of magnitude from $-2.2$ dex below solar to $\\approx -0.2$ below solar, while the rest-frame UV luminosity of the galaxies extends from $\\approx 1/2\\, L_*$ to fainter than $0.05\\,L_*$. There is no apparent upper metallicity cutoff in the sample of GRB host galaxies at $z>2$.  Finally, we include predictions from numerical simulations performed by Sommer-Larsen \\& Fynbo (2008) that include SNe feedback.  The crosses are reproduced from Fynbo \\etal\\ (2008) for comparison with observations.  Despite the presence of lower limits in the metallicities of three GRB hosts, there is a moderate trend that suggests a steeper luminosity--metallicity relation than what is seen in simulations.  We note that this is unlikely due to a selection bias against metal enriched faint galaxies based on the study presented in \\S\\ 5.1, which shows that the GRB host galaxies are consistent with an SFR selected field galaxy sample.  Different theoretical studies have been carried out to understand the key astrophysical processes that determine the luminosity--metallicity relation found in galaxies from different epochs (e.g.\\ Dekel \\& Woo 2003; Tassis \\etal\\ 2008; Brooks \\etal\\ 2007; Finlator \\& Dav\\'e 2008).  The steep decline in ISM metallicity toward fainter magnitudes may be interpreted as a signature of SNe driven outflows that removes metals more effectively in lower-mass halos than in more massive ones (e.g.\\ Dekel \\& Woo 2003), or as a signature of inefficient star formation in low mass halos due to low gas density (e.g.\\ Tassis \\etal\\ 2008; Robertson \\& Kravtsov 2008).  The model expectations presented in Figure 20 includes SNe feedback but no ISM radiation field to account for subsequent destruction of molecules in star-forming regions.  The observed steeper luminosity--metallicity relation in faint galaxies and the absence of molecular gas in GRB host ISM (Tumlinson \\etal\\ 2007) suggests that an enhanced ISM radiation field from young stars may be important to effectively reduce subsequent star formation and ISM chemical enrichment in dwarf galaxies.  \\subsection{Empirical Constraints for the ISM Radiation Field in GRB Host Galaxies}  To understand the lack of molecular gas in GRB host ISM, Tumlinson \\etal\\ (2007) performed a comparison between observations and predictions from a grid of models that cover a parameter space spanned by four unknown properties.  These include gas density, metallicity, clouds size and interstellar radiation field.  The authors showed that the lack of molecular gas in the host ISM can be understood by a combination of low metallicity (low grain production rate) and high interstellar UV radiation field in the host ISM.  Recently, Whalen \\etal\\ (2008) carried out a set of numerical simulations to examine whether the absence of molecular gas is due to the afterglow radiation field or an enhanced interstellar UV radiation field from a pre-existing H\\,II region where the progenitor star resides.  Taking into account known metallicities from afterglow absorption analysis, Whalen \\etal\\ concluded that, similar to what is found in the Magellanic Clouds, ISM radiation fields of intensities up to 100 times the Galactic mean are necessary to explain the absence of H$_2$ in the host ISM.  With the resolved galaxy morphologies seen in the HST images of GRB\\,000926, GRB\\,030323, GRB\\,050820A, and GRB\\,060206, we can measure directly the mean radiation field in the ISM of the host galaxies at near UV wavelengths ($\\approx 1500-2000$ \\AA).  The mean UV radiation intensity is determined by averaging the observed UV flux over the extent of each host galaxy.  For comparison with models, we convert from the observed mean radiation intensity at near UV wavelengths to the far UV Lyman-Warner band ($11.8-13.6$ eV) based on the Galactic UV radiation spectrum of Gondhalekar \\etal\\ (1980).  The estimated far UV radiation fields $I_0$ of individual host galaxies are presented in Table 3.  Adopting a mean Galactic radiation field at far UV wavelengths of $I_0=2.5\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ (Gondhalekar \\etal\\ 1980), we find that, despite a moderate SFR, the ISM radiation field in the GRB host galaxies spans a range over $35 - 350 \\times$ the Galactic mean value due to a relatively compact size.  The observations provide empirical support for the previous theoretical understanding (e.g.\\ Whalen \\etal\\ 2008) that strong UV radiation fields in low-metallicity ISM environment may be a dominant factor for suppressing the formation of molecules and therefore subsequent star formation in dwarf galaxies.  \\subsection{Properties of Foreground Mg\\,II Absorbing Galaxies}  A puzzling observation of the absorption properties along GRB lines of sight is the apparent overdensity of strong (rest-frame absorption equivalent width $W(2796)>1$ \\AA) Mg\\,II absorbers (Prochter \\etal\\ 2006).  Afterglow spectra exhibit on average $\\approx 4\\times$ more strong Mg\\,II absorbers than random QSO spectra (Prochter \\etal\\ 2006; Sudilovsky \\etal\\ 2007), although such over abundance is not seen in C\\,IV absorbers at somewhat higher redshift (Sudilovsky \\etal\\ 2007; Tejos \\etal\\ 2007).  Various scenarios have been considered to explain this discrepancy, including dust extinction due to the presence of these absorbers that biases observations of QSO sightlines, the GRBs being gravitationally lensed by the foreground absorbers, and the absorbers being intrinsic to the GRBs (Prochter \\etal\\ 2006; Porciani \\etal\\ 2007).  However, none of these scenarios alone is found sufficient to explain the observed overabundance along afterglow sightlines.  A different scenario has been proposed by Frank \\etal\\ (2007), who consider different beam sizes between QSOs and GRB afterglows as a possible explanation, but two important consequences are associated with this scenario.  First, a partial covering of Mg\\,II gas is expected along QSO sightlines.  Second, a skewed frequency distribution of Mg\\,II absorbers is also expected along GRB sightlines.  None is confirmed in empirical data (e.g.\\ Pontzen \\etal\\ 2007).  We refer the readers to Porciani \\etal\\ (2007) for a detailed discussion.  Late-time HST images of the fields around known GRBs offer a detailed view of the galaxy environment around known foreground Mg\\,II absorbers along the sightlines toward the GRBs (e.g.\\ Pollack \\etal\\ 2008).  Six of the 15 GRBs sightlines in our sample have known strong Mg\\,II absorbers of $W(2796) > 1$ \\AA\\ in the foreground (e.g.\\ Prochter \\etal\\ 2006), three of which have deep HST images available. Including the field around GRB\\,060418 (Pollack \\etal\\ 2008), we find that in all four cases the HST images have uncovered at least one galaxy at angular separation $\\Delta\\,\\theta \\aapl 1''$ from the afterglow sightlines.  In contrast, such galaxies at small $\\Delta\\,\\theta$ are clearly absent in fields with no known strong Mg\\,II absorbers (Figure 22).  A summary of the known Mg\\,II absorbers and possible candidate absorbing galaxies is presented in Table 4.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.35,angle=270]{gaden.ps} \\end{center} \\caption[]{Cumulative number of galaxies versus angular distance to the afterglow line of sight for GRB fields with (top) and without (bottom) intervening strong Mg\\,II absorbers found in afterglow spectra.  All seven fields have high spatial resolution HST images available for identifying galaxies brighter than $AB(R)=28$.  In all four fields with known intervening strong Mg\\,II galaxies, we find additional objects at $\\Delta\\,\\theta<2\\arcsec$ from the GRB afterglow position.  In contrast, none is seen at this small angular separation in fields without known Mg\\,II absorbers.} \\end{figure}  Although only one of the candidate galaxies in Table 4 has been spectroscopically confirmed (the Mg\\,II absorber at $z=0.84$ toward GRB\\,030429), the finding of additional faint galaxies along the GRB sightlines strongly disfavors the absorbers being intrinsic to the GRBs.  In addition, these candidate galaxies exhibit a broad range of optical-infrared colors (c.f.\\ the blue colors of galaxies toward GRB\\,050820 shown in Table 2 and Figure 9 and those found toward GRB\\,060418 in Pollack \\etal\\ 2008).  Dust extinction does not appear to be a uniform factor across different GRB sightlines that results in the higher incidence of strong Mg\\,II absorbers. Comparisons of the absorbing galaxy properties found along afterglow and QSO sightlines (O'Meara \\etal\\ 2006; Becker \\etal\\ in preparation) should provide further insights for understanding the differential incidences of Mg\\,II absorbers.  The morphology of additional faint emission near the absorbing galaxy toward GRB\\,030429 (Figure 6) is suggestive of a lensed event.  Additional optical imaging and spectroscopic data are necessary to test the lensing hypothesis for this source.  A conclusive answer to the observed overdensity of foreground Mg\\,II absorbers along GRB sightlines requires a larger sample of imaging and spectroscopic data of the candidate absorbing galaxies.  Based on the current finding, however, we caution that the presence of intervening galaxies at small angular distances to the GRBs introduces non-negligible contamination for identifying GRB host galaxies based on imaging data alone.  High spatial resolution images are crucial for resolving the host galaxies from foreground absorbers.  While the properties of the Mg\\,II absorbers are beyond the scope of this paper, we emphasize that in the absence of spectroscopic observations it will be necessary to take into account line-of-sight absorption-line properties in afterglow spectra when evaluating the uncertainty of an imaging identification of GRB host galaxies."
    },
    "0809/0809.5028_arXiv.txt": {
        "abstract": "We study the role of fundamental constants in an updated recombination scenario. We focus on the time variation of the fine structure constant $\\alpha$, and the electron mass $m_e$ in the early Universe, and put bounds on these quantities by using data from CMB including WMAP 5-yr release and the 2dFGRS power spectrum. We analyze how the constraints are modified when changing the recombination scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1117_arXiv.txt": {
        "abstract": "We give an overview of our recent integral-field-unit spectroscopy of luminous extended emission-line regions (EELRs) around low-redshift quasars, including new observations of 5 fields. Previous work has shown that the most luminous EELRs are found almost exclusively around steep-spectrum radio-loud quasars, with apparently disordered global velocity fields, and little, if any, morphological correlation with either the host-galaxy or the radio structure. Our new observations confirm and expand these results. The EELRs often show some clouds with velocities exceeding 500 \\kms, ranging up to 1100 \\kms, but the velocity dispersions, with few exceptions, are in the 30--100 \\kms\\ range. Emission-line ratios show that the EELRs are clearly photoionized by the quasars. Masses of the EELRs range up to $>10^{10}$ \\msun.  Essentially all of the EELRs show relatively low metallicities, and they are associated with quasars that, in contrast to most, show similarly low metallicities in their broad-line regions. The two objects in our sample that do not have classical double-lobed radio morphologies (3C\\,48, with a compact-steep-spectrum source; Mrk\\,1014, radio-quiet, but with a weak compact-steep-spectrum source) are the only ones that appear to have recent star formation. While some of the less-luminous EELRs may have other origins, the most likely explanation for the ones in our sample is that they are examples of gas swept out of the host galaxy by a large-solid-angle blast wave accompanying the production of the radio jets. The triggering of the quasar activity is almost certainly the result of the merger of a gas-rich galaxy with a massive, gas-poor galaxy hosting the supermassive black hole. ",
        "introduction": "\\citet{Bor84} and \\citet{Bor85} carried out an important spectroscopic survey to obtain off-nuclear spectra of 9 radio-loud quasars and 3 QSOs\\footnote{Here we consider only their high-luminosity sample to be consistent with the sample of \\citet{Sto87} which we will discuss in the following paragraphs.}. In addition to finding more quasars with extended emission, an important discovery of this program was the strong correlation between luminous optical extended emission and radio morphology (or radio spectral index). \\citeauthor{Bor85} found that all of the six quasars with strong extended emission lines are steep-spectrum radio sources, and the other six objects\\footnote{Contrary to these original papers, 3C\\,48, a compact steep-spectrum radio source with a powerful radio jet extending 0\\farcs7 to the north, is in fact surrounded by a luminous EELR, as was later shown by \\citet{Sto87}.} where such extended emission is weak or absent are either radio-quiet or flat-spectrum radio-loud sources. There is a strong correlation between radio spectral index and radio morphology, since extended radio structures such as radio lobes are dominated by optically thin synchrotron radiation, which produces a steep power-law spectrum, while unresolved radio cores generally shows a flat spectrum due to synchrotron self-absorption (or thermal bremsstrahlung radiation from a disk wind, see \\citealt{Blu07}). This extended emission$-$spectral index correlation can thus be translated into an extended emission$-$radio morphology correlation---luminous extended emission was found exclusively in quasars having extended radio structures, most of which show Fanaroff-Riley type II (FR II) twin-jet morphologies \\citep{Fan74}.  This extended emission$-$radio morphology correlation was confirmed by the \\othree\\ narrow-band imaging survey of \\citet{Sto87}, but it evolved into a more complicated picture as a result of their $4\\times$ larger sample and their seamless spatial sampling. The authors found that (1) not all steep-spectrum ($\\alpha_{\\nu} > 0.5$) lobe-dominated quasars are surrounded by luminous EELRs---only 10 out of the 26 ($38\\pm12$\\%) have EELRs with \\leothree\\ $ \\geq L_{\\rm cut} = 5\\times10^{41}$ \\ergs (hereafter EELR quasars), where \\leothree\\ is the total \\othree\\ $\\lambda5007$ luminosity within an annulus of inner radius 11.2 kpc and outer radius 44.9 kpc centered on the nucleus and $L_{\\rm cut}$ is determined by the highest \\leothree\\ upper limits of the objects where extended emission is not detected (see Fig.~\\ref{introfig:corr}$a$); and (2) none of the seven flat-spectrum ($\\alpha_{\\nu} < 0.5$) core-dominated quasars and only one of the 14 ($7\\pm7$\\%) optically selected QSOs show an EELR with \\leothree\\ $> L_{\\rm cut}$. To illustrate this correlation more clearly, we have created a new figure (Fig.~\\ref{introfig:corr}$a$) based on the original data of \\citet{Sto87}.  The sample observed by \\citet{Sto87} included almost 2/3 of all optically luminous QSOs known at the time with $z\\le0.45$ and $\\delta > -25\\arcdeg$. It is still nearly complete for radio-loud quasars within these limits but clearly incomplete for radio-quiet QSOs, which typically comprise about 90\\%\\ of the total population. Nevertheless, the radio-quiet objects in the sample are likely typical of the population of luminous, low-redshift, UV-selected QSOs.  \\begin{figure*}[!t] \\epsscale{0.55} \\plotone{f1a.eps} \\plotone{f1b.eps} \\vskip 0.1in \\plotone{f1c.eps} \\plotone{f1d.eps} \\caption[Correlations between extended \\othree\\ emission and quasar properties]{ ({\\it a}) Correlation between extended \\othree\\ luminosity (\\leothree) and radio spectral index ($\\alpha_{\\nu}$), which correlates strongly with radio morphology. ({\\it b}) Correlation between \\leothree\\ and nuclear \\othree\\ luminosity (\\lnothree). ({\\it c}) \\leothree\\ versus quasar absolute {\\it B} magnitude. ({\\it d}) \\lnothree\\ versus $\\alpha_{\\nu}$. Open and filled circles represent objects with solid detections of extended \\othree\\ emission. Open triangles show objects where extended emission was undetected or ambiguous, they therefore represent upper limits of \\leothree. Upper limits of \\lnothree\\ are drawn by arrows. {\\it Solid lines} separate radio-quiet quasars (QSOs), flat-spectrum core-dominated quasars, and steep-spectrum lobe-dominated quasars. {\\it Dashed lines} mark $L_{\\rm cut} = 5\\times10^{41}$ \\ergs, which is our minimum \\othree\\ luminosity for any extended emission to be called an EELR. So the green filled circles highlight the {\\it EELR quasars}, while the open circles and triangles together are considered as {\\it non-EELR quasars} (even though the open circles indicate detection of weak extended emission). The spectral indices of radio quiet objects (\\ie\\ QSOs) have been arbitrarily set to be $-0.1$ for illustration purposes. {\\it Dotted lines} indicate \\lnothree\\ = $6.5\\times10^{42}$ \\ergs, below which apparently no EELRs were found. Original data came from \\citet{Sto87} and \\citet{Ver06} and all except $\\alpha_{\\nu}$ have been converted to our concordance cosmology. } \\label{introfig:corr} \\end{figure*}  \\citet{Bor85} also noticed that the presence of strong extended emission was frequently accompanied by broader and bumpier nuclear Balmer lines, weaker nuclear \\fetwo\\ emission, and stronger nuclear narrow lines. The first two correlations are often regarded as by-products of the extended emission$-$radio morphology correlation, as there exist well-established correlations between radio morphology and nuclear broad-line profile and \\fetwo\\ emission \\citep{Mil79,Ste81}. The last correlation seems to be an intrinsic one, and it was confirmed by \\citet{Sto87}, who found that strong extended \\othree\\ emission is present only for objects with strong nuclear narrow-line emission. Specifically, all of the 11 EELR quasars in \\citet{Sto87} have nuclear \\othree\\ $\\lambda5007$ luminosity (\\lnothree) of at least $6.5\\times10^{42}$ \\ergs, which is clearly seen in Fig.~\\ref{introfig:corr}$b$.  The nuclear \\othree\\ $\\lambda5007$ luminosity has frequently been used as a proxy for the bolometric luminosity or accretion rate of AGN \\citep[\\eg][]{Hec04}, since it correlates with numerous multiwavelength broad-band luminosities \\citep[\\eg][]{Mul94}. However, since (1) the scatter of the \\lnothree$-$continuum luminosity correlation is quite large \\citep[see \\eg][their Fig.~14]{Zak03}, and (2) the \\citet{Sto87} quasar sample spans a limited range in \\lnothree\\ ($3\\times10^8$ to $10^{10}$ \\lsun), the existence of the apparent minimum \\lnothree\\ seems not to reflect a requirement of a minimum quasar accretion rate, instead, it may merely be a requirement on the availability of gaseous material in the nuclear region---in order to form an EELR, a quasar may have to have a fair amount of gas in the nuclear NLR (given a density of $10^3$ \\cc\\ and \\othree\\ $\\lambda5007$/H$\\beta$ = 10,  \\lnothree\\ = $6.5\\times10^{42}$ \\ergs\\ corresponds to an ionized mass of $4.4\\times10^6$ \\msun). In fact, if we replace \\lnothree\\ with the absolute {\\it B} magnitude of the quasar ($M_B$; compiled from \\citealt{Ver06}) or continuum luminosity at a certain rest-frame wavelength, then this correlation immediately disappears (Fig.~\\ref{introfig:corr}$c$), confirming that quasar accretion rate is not a dominant factor in the formation of an EELR.  Combining the two correlations together, one would expect to achieve a high success rate in identifying EELR objects. In Figure~\\ref{introfig:corr}$d$, nuclear \\othree\\ luminosity is plotted against radio spectral index for the same 47 objects for which \\citet{Sto87} obtained narrow-band \\othree\\ images. If one pre-selects only quasars with both $\\alpha_{\\nu} > 0.5$ and $L_{\\rm N\\,[O\\,III]} > 6.5\\times10^{42}$ \\ergs, then roughly half (10/19) of the sample would turn out to be EELR quasars, as opposed to $<23\\%$ (11/47)\\footnote{When radio-quiet quasars are included in this kind of calculation, the derived percentage of EELR quasars should be treated as upper limits since \\citet{Sto87}'s survey included a disproportionally small number of radio-quiet quasars compared to the radio-loud ones. Note that there are a total of 1876 quasars known at $z \\leq 0.45$ and $\\delta > -25^{\\circ}$, among which less than 10\\% are radio-loud \\citep{Ver06}.}, 38\\% (10/26), or $<46\\%$ (11/24) if one performs the pre-selection on the basis of no {\\it a priori} assumption, $\\alpha_{\\nu}$, or $L_{\\rm N\\,[O\\,III]}$, respectively.  Although the identification rate of EELRs can be improved significantly through this kind of exercise, one inevitably encounters this outstanding puzzle---why do only half of the otherwise identical quasars end up having luminous EELRs while the other half do not? The EELR quasars and the non-EELR quasars not only have similar radio morphology and nuclear narrow-line luminosity (given that they have been preselected with both criteria), but they also look remarkably similar in terms of the redshift range, the broad-band quasar luminosity, the luminosity of the host galaxy and the central black hole mass. It is natural to suspect that there is at least one concealed parameter at work, the difference in which results in the two distinct groups. This parameter has recently been identified. Quite unexpectedly, it turned out to be the gas metallicity of quasar broad-line regions (BLRs): quasars showing luminous EELRs have significantly lower metallicity BLRs ($Z \\lesssim 0.6$ \\zsun) than other quasars ($Z >$ \\zsun) \\citep{Fu07a}.  \\subsection{The Origin of the Extended Gas}\\label{introsec:orig}  The suggestion that the extended ionized gas is organized into a rotating envelope or disk \\citep{Wam75} was soon ruled out by the a counter example of 4C\\,37.43, where blue-shifted clouds were detected on both sides of the nucleus \\citep{Sto76}; instead, \\citeauthor{Sto76} proposed that the ionized gas was ejected from the quasar itself, as strong outflows had been observed in a number of high-redshift quasars as blue-shifted absorption lines. But this scenario did not draw much attention. As 3C\\,249.1 and Ton\\,202 happened to be the first targets to be imaged in their redshifted \\othree\\ line, and both showed structures similar to the galactic tidal tails in the simulations of \\citet{Too72}, \\citet{Sto83} argued that the extended gas is tidal debris from a merger of two gas-rich galaxies and the merger is also responsible for triggering the quasar activity. However, a problem with this scenario was noticed soon afterward.  The density of the ionized gas was known to range from tens to a few hundreds \\cc, as indicated by both the density-sensitive \\otwo\\ $\\lambda3727$ doublet profile \\citep{Sto76} and ionization-state-sensitive \\otwo\\ $\\lambda3727$/\\othree\\ $\\lambda5007$ ratio in combination with the projected distance to the nucleus \\citep{Fab87,Cra88}. Such a high density implies both a small size of the ionized cloud and a high pressure, which is unlikely to be balanced by that of the interstellar medium of a normal galaxy. An ionized cloud of this density would dissipate within a sound-crossing time of $\\sim10^4$ years unless it were confined gravitationally or by the pressure of a hot surrounding medium. As there are not enough stars at the locations of the gaseous ``tidal tails\" to gravitationally confine the gas, \\citet{Fab87} suggested that the EELRs are embedded within a hot high-pressure medium. In this scenario, the warm ionized gas comprising the EELRs condenses from a hot halo sometimes seen even in poor clusters. The most direct proof for such a cooling flow would be a detection of diffuse X-ray$-$emitting gas. The X-ray luminosity can be used to estimate the density of the hot medium at the locations of the EELRs ($\\sim10$ kpc), assuming a typical halo density profile and a typical temperature for the X-ray$-$emitting gas (\\eg\\ $kT$ = 1 keV, or $T$ = $1.2\\times10^7$ K). Unfortunately, since quasars themselves are luminous X-ray sources and known EELR quasars typically have a redshift of $\\sim$0.3, such an observation demands (1) a large-aperture X-ray telescope, (2) an arcsec-scale angular-resolution to confine most of the quasar emission within $\\sim$10 kpc, and (3) an accurate knowledge of the instrument point-spread-function to remove the quasar spill-over; therefore it had to wait until {\\it Chandra} X-ray observatory was launched near the turn of the millenium.  Contrary to the expectation, deep {\\it Chandra} X-ray images of three EELR quasars failed to detect the hot X-ray halo from which the warm gas is supposed to condense, and the upper limits on the density of the hot gas were too low to allow sufficient cooling \\citep{Sto06b}. In fact, the pressure of the hot gas, as implied by the density upper limits, is comparable to that of the ISM in our Galaxy ($n T \\approx 10^4$ K \\cc). To reiterate, without a confinement of a high-pressure surrounding medium, the warm ionized gas in the EELRs would dissipate quickly if its density is significantly greater than 1 \\cc.  A somewhat related possibility that avoids these problems is that of a cold accretion flow along Mpc-scale filamentary structures \\citep{Ker05}.  We postpone commenting on this option until our discussion in \\S~\\ref{sumsec:toy}.  Following the approach pioneered by \\citet{Vie92} and \\citet{Bin96}, \\citet{Rob00} found that a photoionization model of a mixed medium with both optically thin and thick components not only overcomes the outstanding difficulties in fitting certain line ratios  suffered by previous single-phase models in the extended ionized gas in the radio galaxy 3C\\,321, but it also results in an ionizing power more consistent with far-infrared (FIR) and radio observations of the hidden central engine where there is less extinction. \\citet{Sto02} showed that the EELR of quasar 4C\\,37.43 is also better modeled as a two-phase medium, consisting of a matter-bounded diffuse component with a unity filling factor ($n \\sim 2$ \\cc, $T \\sim15,000$ K), in which are embedded small, dense clouds ($n \\sim 500$ \\cc, $T \\sim 10,000$ K). Without assuming that most of the mass resides in a high-density medium, which raises the confinement crisis, this two-phase model can explain the high densities inferred from both the \\otwo/\\othree\\ ratio and the \\otwo\\ doublet ratio, as the \\otwo\\ emission is almost entirely produced by the dense clouds even though they account for less than 1\\% of the total mass.  \\citet{Sto02} also obtained a global \\othree\\ velocity map of the 4C\\,37.43 EELR from spectroscopy with an image slicer. The velocity field appears largely disordered, consistent with the results from earlier integral field spectroscopy of similar objects \\citep{Dur94,Cra97,Cra00}. Some faint clouds show velocities as high as 700 \\kms, which are unlikely to be solely gravitational in origin. As the embedded dense clouds would have to be constantly resupplied because of their short lifetimes, \\citeauthor{Sto02} suggested that shocks were the most likely mechanism for regenerating the small dense clouds. It was also known that two of the EELR quasars have FIR colors similar to those of ultraluminous infrared galaxies (ULIRGs), which are virtually all starburst galaxies triggered by major mergers \\citep{San96}. The host galaxies of both quasars had been confirmed to have starburst or recent poststarburst stellar populations (3C\\,48 and Mrk\\,1014, \\citealt{Can00a,Can00b}). Combining all these pieces, \\citet{Sto02} suggested that the extended emission-line gas has been expelled during a merger and is being shocked by a galactic superwind driven by a starburst triggered by the merger.  So three decades after EELRs were first recognized, the origin of the ionized gas and the physical forces that controlled its distribution remained uncertain: (1) the cooling flow model seemed discredited following the {\\it XMM-Newton} high-resolution spectroscopy \\citep{Pet03} and the {\\it Chandra} imaging program \\citep{Sto06b}, but a new possibility had arisen---that the gas could have originated in a cold accretion flow funneled along intergalactic filaments \\citep{Ker05}, (2) the tidal debris model had to incorporate a new physical mechanism (\\eg\\ shocks from a starburst superwind) to explain the existence of small dense clouds embedded in the otherwise low-density photoionized medium, and (3) the old quasar ejection model \\citep{Wam75,Sto76} had been rejuvenated because of the accumulating circumstantial evidence of shocks propagating in the immediate environments of EELR quasars. This evidence includes the two-phase photoionization model of the brightest condensation in the EELR of 4C\\,37.43 \\citep{Sto02} and the discrete X-ray sources detected at distances within 40 kpc from two EELR quasars \\citep[3C\\,249.1 and 4C\\,37.43,][]{Sto06b}. There was thus some evidence for an outflow, but it was unclear what was driving it---a starburst, or the quasar itself, or both?  \\begin{figure*}[!t] \\epsscale{0.285} \\plotone{f2a.eps} \\plotone{f2b.eps} \\plotone{f2c.eps} \\plotone{f2d.eps} \\vskip 0.1in \\plotone{f2e.eps} \\plotone{f2f.eps} \\plotone{f2g.eps} \\plotone{f2h.eps} \\caption[Finding charts of GMOS observations]{ Narrow-band \\othree\\ $\\lambda5007$ images of 3C\\,48, Mrk\\,1014, 3C\\,249.1, Ton\\,616, Ton\\,202, 4C\\,37.43, PKS\\,2251+11, and 3C\\,79 (panels $a$ to $h$). Overlaid are the fields of the GMOS IFUR observations ({\\it solid boxes}) and those of the GMOS IFU2 observations ({\\it long-dashed rectangles}). \\hst\\ linear-ramp-filter images are used if available. The radio jet direction is described by the solid lines for all of the objects except for 3C\\,48 and Mrk\\,1014, where the radio sources are less than 2\\arcsec\\ across. All images have been scaled to show the same physical size (150 kpc on a side) at the redshift of the object which is labeled under the scale bar. The images of 3C\\,48 and Ton\\,616 have been allowed to wrap around to show high surface brightness peaks. In the image of 4C\\,37.43, the long-dashed rectangle shows the mosaicked field of four IFU2 pointings, the overlapping regions of which are shown by the tick marks on the inner edge.} \\label{introfig:obs} \\end{figure*}  The ultimate goal of this and our previous papers \\citep{Fu06,Sto07,Fu07a,Fu07b,Fu08} has been to address these fundamental questions. The recent availability of integral-field units (IFUs) on large telescopes has made a dramatic improvement in our ability to do so. Most of the previous investigations on quasar EELRs have relied on either narrow-band imaging of the entire structure or slit spectroscopy of only a small fraction of the whole extended emission. IFUs allow spectroscopy in a two-dimensional area, combining the advantages of the two traditional observational methods. The earliest integral field spectrograph used in studying quasar EELRs was TIGER \\citep{Bac95}, which was a lenslet-based instrument mounted on the 3.6-m Canada France Hawaii Telescope (CFHT). It offered a $7\\arcsec\\times7\\arcsec$ field of view (FOV) with a spatial sampling of 0\\farcs4. \\citet{Dur94} observed three EELR quasars in the \\citet{Sto87} sample with TIGER (all of which are included in our sample). Due to the limited spectral coverage of the instrument, their spectra included only the range from H$\\beta$ $\\lambda4861$ to \\othree\\ $\\lambda5007$. The main result of this study was that the velocity fields of EELRs were rather chaotic. TIGER was soon superseded by ARGUS \\citep{Van95}, which was also hosted by CFHT but had a ``lenslet+fiber\" design. ARGUS covered a $\\sim8\\farcs5\\times12\\farcs4$ hexagonal FOV with a sampling of 0\\farcs4. Thanks to its optical-fiber design, ARGUS offered a wider wavelength coverage than did TIGER, allowing \\otwo\\ $\\lambda3727$ and \\othree\\ $\\lambda5007$ to be observed simultaneously. \\citet{Cra97,Cra00} performed ARGUS integral field spectroscopy of seven radio-loud quasars that were known to possess spatially extended emission-line gas, among which three are in the \\citet{Sto87} sample and two are in common with our sample. Besides confirming the chaotic velocity fields, \\citeauthor{Cra97} found that all of the extended nebulae had very similar \\otwo/\\othree\\ ratios independent of location. These pioneering studies using IFUs on CFHT were nevertheless restricted in signal-to-noise ratio (S/N) and in resolution because of the size of the telescope. ",
        "conclusions": "\\label{chap:sum}  Here we bring together our results from IFU spectroscopy of the EELRs around 7 quasars and 1 FR II radio galaxy, as have been presented in this and previous papers  \\citep{Fu06,Fu07b,Fu08,Sto07}. We summarize these results for several key properties and interpret them in terms of a tentative ``toy model'' for the formation of quasar EELRs.  \\subsection{Physical Properties of Extended Emission-Line Regions}  \\subsubsection{Kinematics}  The velocity structure of the ionized gas is locally ordered but globally disordered in all of the 7 quasars, consistent with the results of previous work \\citep{Dur94,Cra00,Sto02}. We are not aware of any simple dynamical models that can fit the data. The only velocity map that resembles a rotating disk is that of the radio galaxy 3C\\,79; however, even in this case, this is true only when the fainter clouds are ignored.  The radial velocity maps have shown that most of the extended nebular emission is confined within 300 \\kms\\ of the quasar's systemic velocity, which is defined by the nuclear NLR. Nevertheless, extremely high velocities are observed in fainter or presumably less massive clouds in 5 of the 7 quasars---3C\\,249.1 ($+550$ \\kms), 4C\\,37.43 ($-600$ \\kms), Mrk\\,1014 ($-1100$ \\kms), 3C\\,48 ($-500$ \\kms), and Ton\\,616 ($+440$ \\kms). We believe such high velocities are unlikely to be solely gravitational in origin.  The velocity dispersion maps look quite quiescent with typical values between 30 and 100 \\kms. Three of the 7 EELRs show clouds with velocity dispersions significantly greater than 100 \\kms: in 3C\\,249.1 there is a cloud to the SE of the nucleus with $\\sigma \\sim 310$ \\kms; in Mrk\\,1014 there are two clouds to the E and W of the nucleus with $\\sigma \\sim 280$ \\kms; in 3C\\,48 the high velocity cloud 0\\farcs25 north of the nucleus shows a dispersion of $\\sim$390 \\kms. In all of the three cases, there seems to be a clear interaction between a radio jet and the cloud: the 3C\\,249.1 cloud is aligned with the radio jet to the SE, the Mrk\\,1014 clouds coincide with the two radio knots on either side of the core, and the high-velocity cloud in 3C\\,48 is near the base of the 0\\farcs7 radio jet. However, no such kinematic signatures were seen in PKS\\,2251+11 where a positional correspondence between a radio hot spot and one of the brightest EELR cloud has long been noticed. Thus this is most likely a chance alignment; this conclusion is also supported by the lack of evidence of shocks from line intensity ratios.  \\subsubsection{Ionization Mechanism}  There is now little doubt that EELRs are photoionized by the central quasars. We reached this conclusion because photoionization by a power-law continuum source is the only model that fits the data.  Shock ionization and the ``shock + precursor\" model cannot predict the correct line ratios; this result is especially clear in the \\othree/H$\\beta$ vs. \\otwo/H$\\beta$ diagram \\citep[][their Fig.~5$a$]{Fu07b}. The high \\othree/H$\\beta$ ratios of the EELRs can only be reproduced by the ``shock + precursor\" model when the shock speed exceeds 450 \\kms, which is inconsistent with the low velocity dispersions measured in the clouds (typically 30 $< \\sigma <$ 100 \\kms). Further, the model predicts an \\otwo/H$\\beta$ ratio $\\sim2\\times$ higher than the observed values.  Photoionization by star-forming regions has also been ruled out as a dominant ionization mechanism. For instance, the EELRs lie firmly above the maximum star-forming curve in the BPT diagrams of \\othree/H$\\beta$ vs. \\stwo/H$\\alpha$ and \\othree/H$\\beta$ vs. \\ntwo/H$\\alpha$ (\\eg\\ Figs.~\\ref{fig:3c48o2} and \\ref{fig:3c48n2}). The cloud 3C\\,48-$a$ represents the only case where in addition to the quasar young stars may also have contributed significantly to the ionization (\\S~\\ref{sec:3c48metal}).  \\subsubsection{Metallicity}  EELRs consist of clouds with low metallicities ($Z \\lesssim$ 2/3 \\zsun, on the \\citealt{And89} scale) as implied by their low \\ntwo/H$\\alpha$ line ratios. These clouds are significantly more metal-poor than the gas in similar environments, \\eg\\ the NLRs of type-2 Seyfert galaxies and radio galaxies at similar redshifts. Power-law photoionization models with low metallicities can correctly predict fluxes of optical lines other than \\ntwo\\ as well (see Fig.~\\ref{fig:3c48o2_grid}). Four of the seven quasars in our sample were observed by the \\hst\\ FOS in the UV (3C\\,249.1, Ton\\,202, 4C\\,37.43, and PKS\\,2251+11), both the \\nfive\\ $\\lambda1240$/\\cfour\\ $\\lambda1549$ and \\nfive/\\hetwo\\ $\\lambda1640$ ratios of their BLRs indicate a metallicity less than 0.5 \\zsun \\citep{Fu07a}, remarkably consistent with the metallicity derived from the extended gas in a completely different density regime.  The only radio-quiet quasar, Mrk\\,1014, is also the only one where extended metal-rich gas is detected. It seems that the gas in this quasar has not been well mixed, as a bright EELR cloud to the east shows a low metallicity that is similar to that of the other EELRs in the sample. We suspect the clouds to the west have been enriched by recent star-forming activities. Evidence for a recent starburst in the host galaxy of Mrk\\,1014 was reported by \\citet{Can00b}.  \\subsection{Star Formation in the Host Galaxies of EELR Quasars}  The host galaxies of a couple of our EELR quasars have been studied in great detail by \\citet{Can00a,Can00b}. 3C\\,48 and Mrk\\,1014 were included in their sample because they fall into an intermediate region between quasars and ULIRGs in a FIR color-color diagram \\citep{Lip94}, suggesting significant star-forming components in their host galaxies. Deep imaging of their host galaxies has shown clear evidence that these are cases of major mergers in their final stages. Deep spectroscopy revealed young stellar populations, confirming that FIR colors are indicative of star formation. In 3C\\,48, \\citeauthor{Can00a} found vigorous ongoing star formation and post-starburst populations with ages up to 114 Myr. In Mrk\\,1014, only post-starburst populations with ages between 180 and 290 Myr were seen. Notice that 3C\\,48 and Mrk\\,1014 are the only two objects in our sample that do not have classical FR II double-lobe morphologies.  Similar studies do not exist for the rest of the quasars in our sample. Nevertheless, an additional two EELR quasars have been detected at 60 $\\mu$m by {\\it IRAS} and one by ISO---3C\\,249.1, 4C\\,37.43, and 3C\\,323.1. Although solid detections are only available at 25 $\\mu$m and 60 $\\mu$m, their large $\\alpha(60,25)$ spectral indices place them firmly above the transition region defined by \\citet{Can01}. The FIR colors\\footnote{When calculating these spectral indices, photometric data from both {\\it IRAS} and ISO have been considered. At 25 $\\mu$m and 60 $\\mu$m the priority list is {\\it IRAS} solid detection $>$ ISO solid detection $>$ {\\it IRAS} upper limits, and at 100 $\\mu$m only {\\it IRAS} upper limits were used.} are $\\alpha(60,25) = -0.35, -0.23, -0.05$ and $\\alpha(100,60) > -0.63, -2.10, -2.24$ for 3C\\,249.1, 4C\\,37.43, and 3C\\,323.1, respectively. If ULIRG-unlike FIR colors are indicative of no recent star formation, then we would expect that the host galaxies of these EELR quasars will be devoid of young stellar populations less than a few Myr old. An indirect proof of this conjecture comes from our study of the host galaxy of 3C\\,79 \\citep{Fu08}.  Given that 3C\\,79 is a type-2 EELR quasar with FR II morphology, the unique geometry of the radio galaxy offers a much clearer view of the host galaxy of an EELR quasar. We found that the host galaxy consists of an intermediate-age (1.3 Gyr) stellar population (4\\% by mass) superimposed on a 10 Gyr old population, \\ie\\ there has been no recent star formation.  In summary, the host galaxies of EELR quasars can have a large range of star formation histories. For  Mrk\\,1014 and 3C\\,48, stars have been vigorously forming in the past few hundred Myr; but for 3C\\,79, and, apparently, for at least many of the FR II quasars in our sample, significant star formation has occurred only in the remote past ($\\gtrsim 1$ Gyr ago).  \\subsection{A Toy Model of Quasar Extended Emission-Line Regions}\\label{sumsec:toy}  A scenario for producing quasar EELRs must be able to account for all of the following compelling results: (1) the correlation between extended nebular emission and radio spectral index (\\S~\\ref{introsec:corr}), (2) the correlation between extended nebular emission and quasar BLR metallicity \\citep{Fu07a}, (3) the lack of detailed morphological correspondence between extended nebula and the host galaxy or the radio structure, (4) the chaotic velocity distribution but relatively calm velocity dispersions of the extended nebulae, (5) the frequent detection of extremely high velocity clouds ($V >$ 400 \\kms), and (6) the large masses of the EELRs ($M \\sim 10^9$--$10^{10}$ \\msun; \\citealt{Fu06,Fu07b,Fu08,Sto07}).  As mentioned in \\S~\\ref{introsec:orig}, there were three remaining at least somewhat plausible origins for the extended gas, and we now examine each one of them with the current set of observational results: \\begin{enumerate} \\item Cold accretion of intergalactic gas through Mpc-scale filaments. Although this model can easily explain the low metallicity gas comprising the EELRs, it is difficult for it to explain the quasar activity and the similarly low metallicity of the BLRs. Cold accretion by itself is unlikely to activate the central black hole as there is no likely mechanism to provide the drastic reduction of the angular momentum of the accreted gas that would be required to force a significant amount of it to within the accretion radius of the black hole.  \\item Tidal debris from a galactic merger. It is difficult for this model to explain the clouds with velocities higher than reasonable for gravitational kinematics ($\\gtrsim 400$ \\kms), which have been detected in 5 of the 7 quasars, often at a considerable distance from the quasar and in locations not in the path of a radio jet (3C\\,249.1, Ton\\,616, 4C\\,37.43). In addition, it is difficult to see how this tidal picture would give a situation where {\\it no} luminous EELRs have been found associated with flat-spectrum core-dominated quasars, and those associated with radio-quiet QSOs are relatively rare and of generally lower luminosity than those found around steep-spectrum radio sources. Nevertheless, most of the quasars with luminous EELRs do show independent evidence for recent strong interactions or mergers \\citep{Sto87}, and this evidence has to be taken into account in any reasonable physical model. But it is significant that we never see any indication of an underlying stellar component with a distribution similar to that of a luminous EELR.  \\item Remnants of galactic superwinds driven by either the quasar itself or a coeval starburst. The starburst superwind picture has been found problematic both in 3C\\,249.1 on the ground of the derived momentum injection rate \\citep{Fu06} and in 3C\\,79 based on the lack of young stellar populations \\citep{Fu08}. In addition, starburst-driven superwinds reflect mass ejection from type-II supernovae, and thus they are expected to be metal enriched (it takes only $\\sim$0.1 Gyr to reach a gas metallicity of 1 \\zsun\\ from the primordial abundances; \\citealt{Fri98}). This contradicts the dominance of metal-poor gas in the environments of EELR quasars. It would be especially difficult to explain the metal-poor gas in the quasar BLRs, which are within 0.1 pc of the central black holes. Finally, a starburst-driven superwind cannot explain the fact that luminous EELRs are found virtually exclusively around steep-spectrum radio-loud quasars.  On the other hand, EELRs as remnants of superwinds driven by quasars themselves do not suffer from these difficulties. \\end{enumerate}  We clearly need to combine aspects of the two last points into our over-all picture.  Both strong interactions or mergers, on the one hand, and the presence of a steep-spectrum radio source, on the other, each in general appear to be necessary, but not sufficient, ingredients for the production of a luminous EELR. There is considerable evidence that the merging companion supplies the gas \\citep{Fu07a}, but that the distribution of the ionized gas is controlled by a massive outflow \\citep{Sto02,Sto07,Fu06,Fu07b,Fu08}. We suggest the following overall picture for the formation and the fate of the EELRs around quasars with large FR II radio jets: A large, late-type galaxy with a mass of a few $\\times10^{10}$  M$_{\\odot}$ of low-metallicity gas has recently merged with the gas-poor quasar host galaxy, which has a $\\sim10^9$ M$_{\\odot}$ black hole at its center. Low metallicity gas on the outskirts of the merging companion is driven to its center and is completely mixed with any higher metallicity gas there well before the final coalescence \\citep[\\eg][]{Kew06b}. The merger triggers the current episode of quasar activity, including the production of FR II radio jets. The initiation of the jets also produces a wide-solid-angle blast wave that sweeps most of the gas from the encounter out of the galaxy. This gas is immediately photoionized by UV radiation from the quasar when the blast wave clears the dust cocoon enshrouding the central engine. Turbulent shocks produce high-density ($\\sim400$ cm$^{-3}$) filaments or sheets in the otherwise low-density ($\\sim1$ cm$^{-3}$) ionized medium which is in hydrodynamical equilibrium with a hot ($T \\sim 10^7$ K) diffuse medium, maintaining a low ionization state. The quasar will eventually shut itself off due to starvation, and some high velocity clouds will escape into the inter-galactic medium, but most of the warm gas will stay and recombine to neutral hydrogen (HI) $\\sim10^4$ yr after the quasar shuts down. The chaotic kinematics will be washed away as the gas settles into a rotating disk in a crossing time of $\\sim$0.5 Gyr. But before it settles down, it might look much like the HI gas in the nearby early-type galaxy NGC\\,1023 \\citep{Mor06}. The dynamically settled gas could probably account for a significant fraction of the dispersed HI in giant early-type galaxies.  Note that, in 3C\\,48, we may be witnessing an early phase in one of these episodes (\\S~\\ref{sec:3c48}; \\citealt{Sto07}). The high-velocity gas currently concentrated near the base of the radio jet constitutes a wide-solid-angle outflow with velocities ranging up to at least 800 \\kms\\ and a total mass of ionized gas of at least $\\sim10^9$ \\msun, and quite likely much more.  This picture naturally explains the EELR---radio morphology correlation, since the superwind has been assumed to be associated with the formation of the jets. But the reason for the EELR---low gas metallicity correlation is much less clear. We can speculate that it may have something to do with the lower radiative coupling such gas would have to the quasar radiation field, allowing more efficient accretion. Specifically, a higher metallicity will lower the accretion rate of material towards the center, because both the higher opacity of the gas and larger amount of dust will couple the gas more efficiently to the radiation field of the quasar. Such a lowered accretion rate may delay the spin up of the black hole, consequently delaying the formation of the radio jet, assuming that the jet production is determined by the black hole spin \\citep{Bla90}. Hence, most of the gas may have time to form stars before the jet is launched, leaving an insufficient amount of gas to form an EELR. It is interesting to note that related considerations have been invoked recently to explain the link between low metallicity and long duration $\\gamma$-ray bursts \\citep{Fru06}.  Whether this particular explanation for the low metallicities in the BLRs of EELR quasars is correct or not, the correlation itself clearly suggests {\\it some} physical mechanism by which the injection of a large amount of low-metallicity gas into a system with a super-massive black hole can result in a large-scale, quasar-powered superwind. This sort of situation is likely to have been relatively common in the early universe, when mergers were more frequent and gas-rich galaxies more numerous. This form of feedback might therefore have been of some importance in restricting the growth of the most massive galaxies at early epochs."
    },
    "0809/0809.4504_arXiv.txt": {
        "abstract": "Dielectronic recombination (DR) of singly charged ions is a reaction pathway that is commonly neglected in chemical models of molecular clouds. In this study we include state-of-the-art DR data for He$^+$, C$^+$, N$^+$, O$^+$, Na$^+$, and Mg$^+$ in chemical models used to simulate dense molecular clouds, protostars, and diffuse molecular clouds.  We also update the radiative recombination (RR) rate coefficients for H$^+$, He$^+$, C$^+$, N$^+$, O$^+$,   Na$^+$, and Mg$^+$ to the current state-of-the-art values.   The new RR data has little effect on the models. However, the inclusion of DR results in significant differences  in gas-grain models of dense, cold molecular clouds for the evolution of a number of surface and gas-phase species.   We find differences of a factor of 2 in the abundance for 74 of the 655 species at times of $10^4$--$10^6$~years in this model when we include DR. Of these 74 species, 16 have at least a factor of 10 difference in abundance.  We find the largest differences for species formed on the surface of dust grains.  These differences are due primarily to the addition of C$^+$ DR, which increases the neutral C abundance, thereby enhancing the accretion of C onto dust.  These results may be important for the warm-up phase of molecular clouds when surface species are desorbed into the gas phase.  We also note that no reliable state-of-the-art RR or DR data exist for Si$^+$, P$^+$, S$^+$, Cl$^+$, and Fe$^+$. Modern calculations for these ions are needed to better constrain molecular cloud models. ",
        "introduction": "\\label{sec:intro}  Investigating the physics and chemistry of molecular clouds (also known as quiescent cores) is crucial if one is to understand the processes that ultimately lead to star formation. Such studies are also important in the field of astrobiology since these clouds are the birthplace of the first organic molecules.  Chemical models used to describe the evolution of molecular clouds typically include hundreds of species.   At molecular cloud temperatures ($\\lesssim 100$~K) these atoms and molecules are either neutral or singly ionized. The abundances of these species are, in turn, governed by thousands of reactions that are highly non-linear. These derived abundances are sensitive to the accuracy of the rate coefficients involved. The implication of uncertainties in the rate coefficients used in chemical models has been investigated by \\citet{Roue96a}, \\citet{Vasy04a,Vasy08a}, \\citet{Wake05a,Wake06a}, and others. However, these studies cannot account for reaction processes that are not included in the models in the first place.  To the best of our knowledge, the only  recombination process of free electrons with atomic ions included in any molecular cloud simulation is radiative recombination (RR). Recent calculations by \\citet{Badn06a}\\footnote{http://amdpp.phys.strath.ac.uk/tamoc/RR/} have improved the accuracy of the rate coefficients of these reactions for a number of ions. The alternative pathway of dielectronic recombination (DR) has previously not been included in molecular cloud models. This is largely due to the bulk of published DR calculations and experiments being valid only for plasmas of higher temperature. Recently, however, theoretical calculations by \\citet{Badn03a}\\footnote{http://amdpp.phys.strath.ac.uk/tamoc/DR/} have probed DR at the low temperature regimes of molecular clouds. At these temperatures, the DR rate coefficient can be a factor of 5 or more greater than the RR rate coefficient for certain systems.  In the present paper, we both update the RR data and include DR for the relevant singly charged atomic ions.  We do this for several molecular cloud models under a variety of initial conditions.  We use  Nahoon \\citep{Wake04a} to simulate dense clouds and protostars, the OSU gas-grain code \\citep{osu,Garr06a} to simulate dense clouds,  and the Meudon Photodissociation Region (PDR) code \\citep{pdr,lepe06a} to simulate both diffuse and dense PDRs. We compare the species abundances with those calculated with and without DR included  in order to show the effects on the models. The remainder of the paper is organized as follows: In Sec.~\\ref{sec:recom} we review the recent developments in our understanding RR and DR. Section~\\ref{sec:models} outlines the chemical models we use to simulate dense clouds, diffuse clouds, and protostars. In Sec.~\\ref{sec:results} we present the results of including DR in these chemical models. Concluding remarks are given in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  This work has investigated the importance of new RR and DR rate coefficients in molecular cloud models. We have updated the RR rate coefficients of H$^+$, He$^+$, C$^+$, N$^+$, O$^+$, Na$^+$, and Mg$^+$ to the current state-of-the-art in 3 different molecular cloud chemical models. We have run these models under different conditions and compared the abundances of the species present with and without the new RR rate coefficients. We find that the new RR data have no significant effects on the model results (i.e., less than a factor of 2 effect). We have also included DR for He$^+$, C$^+$, N$^+$, O$^+$, Na$^+$, and Mg$^+$ in the 3 chemical models, finding some sizable differences (greater than a factor of 2) in the evolution of certain species.  We find that DR has the greatest effect  for  dense cloud models, particularly when surface chemistry is included. Of the models we ran here, only the OSU gas-grain code includes the extensive surface chemistry that allows species to be accreted onto grains and subsequently react with other surface species.  Nahoon does not include this surface chemistry and the Meudon code only allows the formation of H$_2$ on surfaces. DR has its greatest influence for species that form on grain surfaces.   These species may be particularly important in the warm-up phase of molecular clouds when they become desorbed into the gas phase. We find DR of C$^+$  to be primarily responsible for the abundance changes in the model.  This is because charged species are assumed not to accrete onto grains and adding DR increases the gas-phase neutral C abundance available for first accreting onto grains and then participating in the surface chemistry.  We conclude by noting that there are some ions for which no state-of-the-art RR or DR data exist, specifically Si$^+$, P$^+$, S$^+$, Cl$^+$, and Fe$^+$. Given the fine structure present in the ground terms of Si$^+$, Cl$^+$, P$^+$, and Fe$^+$ one would expect a significant DR component to the total electron-ion recombination rate coefficient at low temperatures for these ions.  These 4 elements also have first ionization potentials below that of H (13.6~eV). Given that they would thus be ionized in molecular clouds, the effect of increasing their recombination rates may be important.  Modern calculations of both RR and DR rate coefficients for these ions are needed for generating reliable molecular cloud chemical models."
    },
    "0809/0809.1321_arXiv.txt": {
        "abstract": "{Supersonic turbulence is ubiquitous in the interstellar medium and plays an important role in contemporary star formation.} {To perform a high-resolution numerical simulation of supersonic isothermal turbulence driven by compressive large-scale forcing and to analyse various statistical properties.} {The compressible Euler equations with an external stochastic force field dominated by rotation-free modes are solved with the piecewise parabolic method. Both a static grid and adaptive mesh refinement is used with an effective resolution $N=768^{3}$.} {After a transient phase dominated by shocks, turbulence evolves into a steady state with root mean square Mach number $\\approx 2.2\\ldots 2.5$, in which cloud-like structures of over-dense gas are surrounded by highly rarefied gas. The index of the turbulence energy spectrum function $\\beta\\approx 2.0$ in the shock-dominated phase. As the flow approaches statistical equilibrium, the spectrum flattens, with $\\beta\\approx 1.9$. For the scaling exponent of the root mean square velocity fluctuation, we obtain $\\gamma\\approx 0.43$ from the velocity structure functions of second order. These results are well within the range of observed scaling properties for the velocity dispersion in molecular clouds. Calculating structure functions of order $p=1,\\ldots,5$, we find for \\emph{all} scaling exponents significant deviations from the Kolmogorov-Burgers model proposed by Boldyrev. Our results are very well described by a general log-Poisson model with a higher degree of intermittency, which implies an influence of the forcing on the scaling properties. The spectral index of the quadratic logarithmic density fluctuation $\\beta_{\\delta}\\approx 1.8$. Contrary to previous numerical results for isothermal turbulence, we obtain a skewed probability density function of the mass density fluctuations that is not consistent with log-normal statistics and entails a substantially higher fraction of mass in the density peaks than implied by the Padoan-Nordlund relation between the variance of the density fluctuations and the Mach number.} {Even putting aside further complexity due to magnetic fields, gravity or thermal processes, we question the notion that Larson-type relations are a consequence of universal supersonic turbulence scaling. For a genuine understanding, it seems necessary to account for the production mechanism of turbulence in the ISM.} ",
        "introduction": "Today it is agreed that supersonic turbulence occurs over a large range of length scales in the interstellar medium (ISM) and, particularly, within molecular clouds \\citep{ElmeScal04}. This is inferred from different observations of molecular line broadening which indicate, firstly, a velocity dispersion comparable to or greater than the speed of sound and, secondly, a power-law dependence of the velocity dispersion, $\\Delta v\\propto \\ell^{\\,\\alpha}$, in the range of length scales $0.1\\,\\mathrm{pc}\\lesssim\\ell\\lesssim10\\,\\mathrm{pc}$ of the associated structures \\citep{Larson81,FalPu92,MieBal94,BruntHey02b}. This has been interpreted as a manifestation of the universality of interstellar turbulence \\citep{HeyBrunt04} and bears important consequences on the theory of turbulence-regulated star formation, where Larson-type relations for the velocity dispersion are explained as a direct consequence of the self-similarity of the turbulence cascade, with possible modifications of the scaling exponent due to magnetic fields, self-gravity or thermal processes \\citep{PadoanNordlund02,LowKless04,BallKless07,KeeOst07}.  The question of the origin of supersonic interstellar turbulence remains largely open \\citep{ElmeScal04,BallKless07}. For the numerical simulation of interstellar turbulence, there is also the difficulty that energy injection into the interstellar medium and processes leading to star formation may occur on vastly disparate scales, which cannot be encompassed by a single numerical simulation. This is why two major approaches to the problem have emerged: Either the generation of turbulence by instabilities at relatively large scales is computed or statistically isotropic turbulence is simulated in a periodic box as a simple model for the small-scale regime. Typical examples for the first approach are the simulations of colliding flows by  \\citet{HeitSly06}, \\citet{VazGom07} and \\citet{HenneAud07}. Recent simulations of driven turbulence in a periodic box were presented by \\citet{PadNord07} and \\citet{KritNor07}. While colliding flow simulations follow the idea of controlling the flow evolution in the computational domain via initial and boundary conditions only, the meaning of turbulence simulations with a driving force are implicitly based on the notion of an \\emph{upper} numerical cutoff scale. This means that one imagines turbulence being produced on scales larger than the size of the computational box and the driving force is intended to mimic the interactions with the larger scales in a heuristic fashion. Indeed, decomposing the hydrodynamic equations into a large-scale part and a small-scale part entails coupling via an energy flux that, in the most simplified manner, can be modelled as a random volume force acting upon the small-scale part \\cite{Frisch}. Such a volume force can produce shear or compression by exciting solenoidal (divergence-free) or dilatational (rotation-free) modes of the velocity field.  In this article, we follow a hybrid approach. A statistically isotropic stochastic force field that excites mostly dilatational modes is applied in our simulations. Thereby, converging supersonic flows are produced randomly at large scales. From the collisions of these flows in combination with a small solenoidal component of the stochastic force, turbulence is generated and the flow evolves towards a steady state, in which energy injection due to the forcing is balanced by numerical dissipation. We refer to this scenario as compressively driven turbulence. Apart from the dominance of the rotation-free component, our method differs in another aspect form the forcing applied, for instance, by \\citet{KritNor07} and in many earlier simulations. Whereas they imprint a certain large-scale structure by a steady random force that is derived from some initial velocity field and merely replenishes the energy of the flow, we start with isothermal gas of uniform density at rest and let turbulence develop dynamically. The spatiotemporal correlation properties of the stochastic force field guarantee that large-scale structures evolve on a time scale that is in accordance with their characteristic velocity and size. Moreover, a statistically isotropic state is approached regardless of the initial condition, because temporal correlations decay exponentially. There is a common point of view though that the mode of energy injection in driven turbulence simulations is not important for statistical properties of the resulting flow \\citep{Boldyrev+02}. The statistical analysis of our simulation data, however, cast doubt on his claim. Our results indicate that the dynamics of compressible turbulence at length scales significantly smaller than the energy-containing scale is affected by the forcing. For this reason, the notion of fully developed turbulence--i.~e., turbulence that is not affected by imposed constraints such as boundary conditions, forcing or viscosity--is possibly meaningless in the highly compressible regime. Of course, we cannot rule out that universal statistics might be asymptotically approached toward length scales smaller than the resolution of our simulation. We note, however, that universality of the small-scale dynamics has been questioned even in the case of incompressible turbulence \\citep{AlexaMin05,CasGual07}.  The isothermal approximation is expected to apply to the interiors of molecular clouds, where the temperature of the cold phase of the interstellar gas is close to its minimum \\citep{BallKless07}. In molecular clouds, there are prominent magnetic fields. Although they have indisputably a large impact on fragmentation properties \\citep{HenneTeys08}, it appears that scaling exponents are not greatly altered by magnetic fields, at least in the super-Alfv\\'{e}nic case \\citep{Boldyrev+02,PadJim04,PadNord07}. Consequently, our analysis of the velocity statistics of compressively driven turbulence has potential relevance to molecular cloud turbulence.  At larger scales in the ISM, various heating and cooling processes bring about thermal instabilities. These are included in the colliding flow simulations by \\citet{HeitSly06}, \\citet{VazGom07} and \\citet{HenneAud07}. We will consider driven turbulence in thermally bistable gas in a follow-up paper. In the future, we also plan to include self-gravity and magnetic fields. By successively including more complex physics, we expect to achieve a thorough understanding of the effects stemming from certain physical processes. The numerical treatment of processes such as cooling or self-gravity poses the problem of covering a wide range of length scales, at least in some regions of the flow. Using grid-based codes, a possible solution is adaptive mesh refinement (AMR). The simulation of turbulence with AMR was pioneered by \\citet{KritNor06}. Recently, an AMR simulations of self-gravitating turbulence have been presented \\citep{OffKrum07}. As a preliminary study, we also performed an AMR simulation, with the rate of compression (i.~e., the rate of change of the negative divergence) and the vorticity modulus as control variables for refinement. Velocity and density statistics are well reproduced in comparison to the simulation with a static grid, but, using only one level of refinement, we have not achieved a significant gain in terms of performance in the case of isothermal supersonic turbulence. The same difficulty was encountered by \\citep{KritNor07} for more than twice the resolution we used.  The article is organised as follows. Section~\\ref{sc:numerics} explains numerical techniques, in particular, the stochastic forcing and the method of refinement employed in the AMR simulation. In Section~\\ref{sc:results}, we begin our presentation of simulation results with global statistics and interpret some properties via three-dimensional renderings of the mass density and the vorticity in the stage of turbulence production. In the following, we analyse the probability density functions of the mass density and the vorticity. Next we turn our attention to the spectra of the turbulence energy and the density fluctuations. Finally, we derive scaling exponents from velocity structure functions. The results are related to each other and compared to observational results, theoretical predictions as well as results from other numerical simulations in Section~\\ref{sc:discussion}. In the conclusion, we discuss possible implications on molecular cloud turbulence and star formation and give an outlook to future work. ",
        "conclusions": "We performed simulations of supersonic turbulence in isothermal gas driven by a mostly compressive stochastic force field. The analysis of statistical properties of the velocity and mass density revealed several remarkable results:  \\begin{enumerate} \\item The turbulence energy spectrum function is close to the spectrum of Burgers turbulence in the initial phase of the flow evolution and subsequently approaches a flatter spectrum with an index $\\beta\\approx 1.9$. Within the error bars, a comparable index $\\beta_{\\delta}\\approx 1.8$ is found for the spectrum function of $\\delta^{2}$, where $\\delta=\\ln(\\rho/\\rho_{0})$ is the logarithmic density fluctuation. \\item In the statistically stationary state, the scaling exponent of the RMS velocity fluctuation $\\gamma=\\zeta_{2}/2\\approx 0.43$, which is comparable to typical exponents inferred from measurements of the velocity dispersion in molecular clouds. \\item The relative scaling exponents $Z_{p}=\\zeta_{p}/\\zeta_{3}$ up to order $p=5$ deviate from the Kolmogorov-Burgers model for supersonic turbulence. Varying only the codimension $C$ of dissipative structures, the best fit to our data implies $C\\approx 0.71$, whereas $C=1$ in the Kolmogorov-Burgers model. Fitting the general log-Poisson hierarchical structure model, we find $C\\approx 1.23$ and the scaling index $\\Delta\\approx 1.0$. This can be interpreted as shocks being the \\emph{most singular} dissipative structures. In the Kolmogorov-Burgers model, $\\Delta=2/3$ is assumed. \\item Compared to other simulations, our analysis point toward non-universal scaling properties, i.~e., the velocity statistics cannot be accommodated within a one-parameter family of models. There appears to be an additional degree of freedom related to the large-scale forcing. \\item The time-averaged probability density function of $\\delta$ is not consistent with a log-normal distribution of the mass density. We find a skewed distribution that is broader than the log-normal distribution predicted by \\citet{PadoanNordlund02} on the basis of the RMS Mach number. Large deviations of the mass fraction in over-dense gas compared to log-normal distributions are implied. \\end{enumerate}  The assumption of a log-normal mass density distribution has been used for analytical predictions of the rate of turbulence-regulated star formation \\citep{PadoanNordlund02,KrumKee05,HenneChab08}. With the distribution of the mass density found in our simulation, the resulting star formation rate would differ substantially form the log-normal case because the mass converted into stars by gravitational collapse mainly depends on the high-density tail of the probability distribution function. However, this result is not conclusive yet, because a wider range of parameters has to be studied, self-gravity and magnetic fields have to be included, and, most importantly, the dependence of the mass density distribution and clump mass spectra on the numerical resolution has to be clarified. In part, this is the subject of ongoing work, where turbulence simulations with purely solenoidal and compressive forcing are systematically compared for numerical resolutions up to $1024^{3}$ \\citep{SchmFeder08,FederKless08}. Moreover, this complementary study is essential to corroborate the non-universal velocity scaling properties discussed in this article.  The observational results of \\citet{BruntHey02b} point towards a scaling exponent $\\alpha$ that is \\emph{significantly greater} than $0.5$ in many instances. This is hard to explain on the basis of supersonic turbulence in isothermal gas. Stiffer scaling exponents might be caused, for instance, by self-gravity \\citep{BigDia89}. On the other hand, \\citet{HilyFal08} found scaling properties that are indicative of nearly incompressible, shear-dominated turbulence in some molecular clouds. Possibly, molecular clouds are a rather inhomogeneous class of objects. From that point of view, the situation is reminiscent of the observation of type Ia supernovae, which were recognised to be much more diverse than it had appeared in the beginning. To understand how the mass density distribution and the two-point statistics of the turbulent velocity at various scales in the interstellar medium comes about, it is vital to include additional physics such as self-gravity, radiative cooling and magnetic fields also in numerical simulations of turbulence with stochastic forcing. To this end, the application of adaptive methods will be indispensable. \\citet{OffKrum07} succeeded in performing simulations of self-gravitating isothermal turbulence with AMR. Radiative post-processing of their numerical data revealed discrepancies with observed core line widths both for driven and for decaying turbulence. This suggests that not only gravity but also thermal processes are important for the dynamics of molecular cloud turbulence. With the AMR simulation presented in this article, we have attempted a first step toward a methodology that is based on the dynamics of turbulent flows. \\citet{IapiAdam08} have already presented a multi-level AMR application using these techniques. The development is far from being complete yet, but it appears to be promising for a self-consistent treatment of turbulence in the complex scenarios encountered in star-forming clouds."
    },
    "0809/0809.3262_arXiv.txt": {
        "abstract": "A search for recoiling supermassive black hole candidates recently yielded the best candidate thus far, SDSS~J092712.65+294344.0 reported by Komossa et al. Here we propose the alternative hypothesis that this object is a supermassive black hole binary.  From the velocity shift imprinted in the emission-line spectrum we infer an orbital period of $\\sim 190$ years for a binary mass ratio of 0.1, a secondary black hole mass of $10^8~{\\rm M}_\\odot$, and assuming inclination and orbital phase angles of $45^\\circ$.  In this model the origin of the blueshifted narrow emission lines is naturally explained in the context of an accretion flow within the inner rim of the circumbinary disk. We attribute the blueshifted broad emission lines to gas associated with a disk around the accreting secondary black hole. We show that, within the uncertainties, this binary system can be long lived and thus, is not observed in a special moment in time. The orbital motion of the binary can potentially be observed with the $VLBA$ if at least the secondary black hole is a radio emitter. In addition, for the parameters quoted above, the orbital motion will result in a $\\sim 100~{\\rm km\\,s^{-1}}$ velocity shift of the emission lines on a time scale of about a year, providing a direct observational test for the binary hypothesis. ",
        "introduction": "The ``mass loss'' and gravitational rocket effect in the aftermath of the coalescence of a supermassive black hole binary (SBHB) can have profound effects on the properties of the nucleus of the host galaxy and give rise to unique observational signatures \\citep{rr89, milos05, loeb07, sb08, lippai08, sk08, kl08, moh08, devecchi08,bl08,gm08}.  Since the emission of gravitational waves in the last stages of the merger is not symmetric in general, the product of the merger can receive a significant recoil velocity, up to a few $\\times \\,100~{\\rm km\\,s^{-1}}$ for low black hole spins or spin axes aligned with the orbital axis \\citep{herrmann07a, herrmann07b, baker06, baker07, gonzalez07b, koppitz07, rezzolla08}. In the special case of maximally spinning, equal mass supermassive black holes (SBHs) with their spin vectors in the orbital plane and directed opposite to each other, the recoil speed can be up to $\\sim$ 4000~ ${\\rm km\\,s^{-1}}$ \\citep{campanelli07b} . The fraction of coalescences expected to produce a remnant recoiling at $V>$ 1000~${\\rm km\\,s^{-1}}$ is about $10\\, \\%$, assuming black hole spins of $cJ/GM^2 = 0.9$, mass ratio range of 0.1 $<$ q $<$ 1, and arbitrary spin orientations \\citep{sb07,campanelli07a,baker08}.  Since the escape speed from most galaxies is less than 2000~${\\rm km\\,s^{-1}}$ \\citep{merritt04}, if high velocity recoils are common, there should be many ``empty nest\" galaxies, without a central SBH. This is in contrast with the observation that almost all galaxies with bulges have a central SBH \\citep[for summary see][]{ff05}. However, if following a galactic merger the two SBHs find themselves in a gas rich environment, gas accretion torques will act to align the spin axes with the orbital axis and thus, reduce the post-merger kick to a value well below the galactic escape speed. This effect is expected to increase the chance of retention of recoiling SBHs by their host galaxies \\citep{bogdanovic07}. Gas-poor mergers, on the other hand can lead to large recoil speeds that launch the final black hole out of the potential well of the host bulge or host galaxy.  \\citet{bonning07} searched the Sloan Digital Sky Survey (SDSS) archive for merger products with large velocity offsets from their host galaxy. Among nearly 2600 objects, they found no convincing evidence for recoiling SBHs in quasars, placing an upper limit of 0.2~$\\%$ on the probability of kicks with projected speeds greater than 800~${\\rm km\\,s^{-1}}$ and 0.04~$\\%$ on the probability of kicks with projected speeds greater than 2500~${\\rm km\\,s^{-1}}$. More recently, Komossa et al. (2008), hereafter \\citet{komossa08}, reported the discovery of one such candidate, a SBH receding from its host galaxy at a high projected speed of 2650 ${\\rm km\\,s^{-1}}$. This detection has very important implications for the cosmological evolution of SBHs. However, it is also remarkable given the likelihood of high velocity kicks and a narrow observational window (in cosmological terms) associated with this class of objects \\citep{bl08}.  Combined with several observational curiosities, they call the recoil idea into question (see discussion in \\S\\,3). Here we propose that the peculiar spectroscopic properties of SDSS~J092712.65+294344.0 \\citep{am07} can be explained in the context of a SBHB model. In \\S\\,\\ref{S_model} we describe the proposed model, in \\S\\,\\ref{S_conclude} we discuss its physical basis, revisit the recoiling black hole model, and close by considering observational tests. ",
        "conclusions": "\\subsection{The Physical Picture}\\label{S_picture}  The evolution of a binary orbit is determined by the relative efficiency of processes that can transport orbital angular momentum, such as stellar and gas dynamical processes and in later stages, the emission of gravitational radiation \\citep{bbr}. In the binary model considered here stellar processes are expected to be inefficient and to operate on time scales comparable to the Hubble time \\citep{berczik06,sesana07,perets07}. A scenario in which large amounts of cold gas are present within the sub-parsec binary orbit is not favored in the context of current understanding of these systems \\citep{armitage02, milos05, mm06}. Feedback from the AGN and binary torques are expected to efficiently heat and disperse the gas. As a consequence, the gaseous dynamical friction does not affect the orbital evolution of the binary in this phase.  The gas outside of the SBHB orbit, in the circumbinary disk, can exert torques on the binary and \\citet{cuadra08} find that only binaries with mass $\\lesssim 10^{7} \\Msun$ can coalesce within a Hubble time due to this effect \\citep[see also][]{hayasaki08}, while more massive SBHBs, like the one assumed here, evolve on longer time scales. If stellar and gas dynamical mechanisms for angular momentum transport are inefficient, the evolution of the sub-parsec binary orbit will be determined by the emission of gravitational radiation and the assumption of the circular orbit is justified. Note however that if high, near-Eddington mass accretion rates onto the binary are plausible, interactions with the circumbinary disk may drive the evolution of the orbital separation and eccentricity on time scale $\\sim {\\rm few}\\times t_{\\rm visc}\\approx 10^8\\,{\\rm yr}$, shorter than $t_{\\rm gw}$ (where $t_{\\rm visc}$ and $t_{\\rm gw}$ are, respectively, the viscous time scale of the disk at the inner edge and the gravitational wave decay time). Realistically, $t_{\\rm visc}$ will be longer because the structure of the disk and consequently, accretion rate will be affected by the presence of the binary. If so, the time scale given by $t_{\\rm visc}$ can be considered a conservative lower limit on the life span of the binary and the assumption of circularity (or more precisely, moderate eccentricity) is still a plausible one. The time scale for orbital decay of the binary due to the emission of gravitational waves is \\begin{equation} t_{\\rm gw} \\approx 1.4\\times10^{10}\\;{M_{2,8}\\over (1+q)^5} \\left({0.1\\over q}\\right)^2 \\left({\\sin i\\over \\sin 45^\\circ}\\;{\\sin \\phi\\over \\sin 45^\\circ}\\right)^8\\;{\\rm yr}\\, . \\label{eq3} \\end{equation} Since $t_{\\rm gw}$ is a sensitive function of $\\sin i$ and $\\sin \\phi$, it may be shorter than the Hubble time if we are observing the binary at low inclination or close to conjunction, or alternatively if the binary orbit has eccentricity e $>$ 0.1.  \\begin{figure}[t] \\epsscale{1.0} \\plotone{f1.ps} \\caption{ Illustration of the innermost region of a circumbinary disk after the binary has cleared a low density ``hole'' in the center (top view, not drawn to scale). In the context of the model proposed here, the r-NELs are associated with the circumbinary disk, b-BELs with the disk surrounding the less massive secondary SBH, and b-NELs with the accretion streams flowing from the inner edge of the circumbinary disk towards the disk of the secondary. Accretion occurs preferentially on the secondary SBH, rendering the primary quiescent or much fainter in comparison. \\label{fig_sketch}} \\end{figure}  We thus assume that the binary described here is surrounded by a circumbinary disk, as illustrated in Figure~\\ref{fig_sketch}. As the binary orbit decays, the inner rim of the disk follows it inward until the time scale for orbital decay by gravitational radiation becomes shorter than the viscous time scale. In the model presented here, $t_{\\rm gw} > t_{\\rm visc}$, implying that the binary has not detached from the circumbinary disk and can still draw matter from it. Moreover, detailed calculations have shown that for small mass ratios (q $\\lesssim {\\rm 1/few}$) accretion occurs preferentially onto the lower mass object which, as a consequence, will be more luminous and easier to detect \\citep{al96,gould00}. For example, \\citet{hayasaki07} find a significant inversion of accretion rates, $\\dot{M}_2/\\dot{M}_1 = 3.25$ for a $q = 0.5$ binary. This effect is consistent with the picture proposed here that the single set of broad emission lines observed in J0927 is associated with the secondary SBH. On the other hand, a lower limit of $q > 0.01$ can be placed on the binary mass ratio, given our choice $M_2 = 10^8 \\Msun$ and the expected upper limit to black hole masses, $M_{\\rm max} \\sim 10^{10} \\Msun$ \\citep{nt08}.  In addition to simulations of SBHBs, the accretion flow between a circumbinary disk and the binary has been modeled in simulations of disks surrounding stellar, T~Tauri binaries \\citep{gunther02, gunther04}. The dynamics of these two types of systems are, in fact, quite similar. In particular, a common result of the above simulations is that the accretion on individual binary members is mediated by one or more accretion streams, implying that these may arise as a general property of circumbinary accretion flows. The detailed structure and spectroscopic signatures of such flows are unknown, nevertheless, we propose the following simple picture for the system considered here. Because the gas in accretion streams flowing towards the secondary SBH has a nonzero angular momentum, it is expected to form a small accretion disk surrounding the secondary SBH prior to plunging into it. Given the uniform and monotonic orbital evolution of the binary, this may be a long lived and stable phenomenon, where the rate of accretion onto $M_2$ (i.e., the surface density of the disk surrounding the black hole) will be determined by the flow rate of matter in the gas streams. \\citet{mm06} find the flow rate in such a binary system to be of order of $10\\%$ of the accretion rate onto a single black hole with a mass equal to that of the binary. Some of the gas in the flow is accreted by the black holes, while the rest develops eccentric orbits, may collide with the inner rim of the circumbinary disk and leave the binary system in form of the high velocity outflows \\citep{armitage02,mm06}. While the details of this process remain to be modeled, we suggest in the context of our proposed scenario that a negligibly small fraction of the gas will be accreted by the primary black hole, given the angular momentum ``barrier'' that the gas experiences. Consequently, an AGN associated with $M_1$ may either be faint or have short lived accretion phases. This implies that for the binary mass ratio of $q = 0.1$, the accretion rate onto the secondary is less than or equal to its Eddington limit ($\\dot M_2\\lesssim {\\dot M_{\\rm Edd,\\,2}}$). Thus, the necessary conditions to establish a long term accretion process on $M_2$ exist, though the realistic physical picture may be more complex due to radiative feedback from the AGN.  The accretion flow within the Roche lobe of the secondary would resemble the accretion flow onto a single supermassive black hole in an AGN, i.e., its optical spectroscopic signature will be that of an AGN, exhibiting a spectrum with a blue, featureless continuum and broad, permitted emission lines. This picture is consistent with the properties of the observed b-BELs --- if the FWHM of these lines are interpreted to roughly represent the size of the emission region around the secondary, it follows that the size of this region is $\\sim 0.1a$ (assuming $i = 45^\\circ$). Using the \\citet{eggleton83} approximation we estimate the effective Roche lobe radius of the secondary to be $R_{L2} / a = 0.21$ for $q = 0.1$. Thus, the BLR is bound to the secondary SBH, consistent with the expectation that any gas beyond the Roche lobe of the secondary should be tidally truncated by the primary.  Binaries with moderate eccentricities and mass ratios that are not extreme are expected to truncate their circumbinary disks at an inner radius of about twice the binary semi-major axis, and hence, the streams from the circumbinary disk to the secondary will flow over a region of size comparable to that of the binary orbit. The stream properties will be intermediate between the physical properties of the secondary's BLR and the NLR of the host galaxy; \\citet{dotti08}, for example, estimate an {\\it average} density in such a circumbinary region in the range $2-8\\times 10^6\\,{\\rm cm^{-3}}$.  Of course the density of the gas streams themselves will be higher but they will likely be surrounded by lower-density envelopes resulting from expansion or ablation of the gas due to illumination and heating by the accreting black hole. The gas in the streams will be photoionized by the AGN continuum and the resulting emission-line spectrum may have the following properties: $(i)$ It will consist of permitted lines, as well as forbidden lines from the lower-density parts of the flow, such as the ionized skin of the streams.  Due to proximity in velocity space, the lines from the stream should be shifted to approximately the velocity of the broad emission lines from the vicinity of the secondary.  $(ii)$ The profiles of the lines from this stream will be narrower than the broad, permitted lines and asymmetric, since they trace the {\\it emissivity weighted} distribution of the spatially confined stream of gas\\footnote{The asymmetry may arise since in general case the stream geometry and velocity will not be symmetric with respect to the observer.  While the exact profile shapes of these lines will depend on the photoionization effects (not considered here), these are not expected to give rise to symmetric line profiles, in general.}.  $(iii)$ The line shifts, and perhaps also the asymmetries will be variable over the orbital cycle of the binary.  We propose that the b-NELs, with their variety of shifts and widths, originate in this part of the flow. Circumstantial evidence in support of the proposed picture is also provided by the relative intensities of the b-NEL and r-NEL lines. By comparing the relative line intensities reported by \\citet{komossa08} with the photoionization models of \\citet{nagao01,nagao02} we find that the observed [\\ion{Ne}{3}]/[\\ion{O}{2}] ratios suggest a higher density (by 1--2 orders of magnitude) in the b-NLR compared to the r-NLR. Moreover, the [\\ion{O}{3}]/H$\\beta$ ratios suggest a higher ionization parameter in the former region by a factor of several, which is also reflected in the [\\ion{Ne}{3}]/[\\ion{Ne}{5}] ratio. These differences in density and ionization are consistent with our hypothesis that the r-NLR is associated with gas in the nucleus of the host galaxy while the b-NLR is associated with denser gas that is part of the accretion flow onto the secondary black hole. The above conclusions are subject to some uncertainty because the results of photoionization models are sensitive to a variety of input parameters, such as the spectral energy distribution of the ionizing continuum.  \\subsection{The Recoiling Black Hole Scenario Revisited}\\label{S_revisit}  Let us consider the characteristic time scales relevant for the recoiling SBH model with the parameters adopted by \\citet{komossa08}. The time during which a recoiling SBH of mass $\\sim 6 \\times 10^8\\,\\Msun$ appears as an AGN can in principle be close to $10^9$~yr, if the mass of the accretion disk carried along by the SBH is comparable to its own mass \\citep{loeb07}. This cannot be the case in J0927, if its SBH charges away from the host galaxy at nearly the maximum speed predicted by numerical relativity. To accommodate a high velocity kick, the mass of the disk should be small enough not to slow down the SBH, implying that $t_{\\rm AGN} \\ll 10^9$~yr. Indeed, \\citet{bl08} calculate that the mass of the gas disk carried by the recoiling SBH in J0927 is $\\sim 2 \\%$ of the SBH mass, which would be accreted in only $\\sim 10^7$~yr, if the accretion rate is 0.1 of the Eddington limit.  If the FWHM of the [\\ion{Ne}{3}] line is indicative of the expansion of the accretion disk due to outward transport of angular momentum, as suggested in the context of the recoiling SBH model, then the final disk radius is $\\sim$7 times larger than that immediately after the recoil. This factor is based on a comparison of the FWHM of the [\\ion{Ne}{3}] line \\citep[reported by][]{komossa08} and the line of sight recoil velocity\\footnote{Assuming that the recoil velocity and the FWHM of the [\\ion{Ne}{3}] line indicate the truncation radius of the disk before and after the expansion. Larger factors are expected if the radius of marginal self-gravity is considered instead to obtain the former value.}.  A disk expansion should be accompanied by a drop in surface density and an accretion rate reduction by a factor of $> 7^2$ (assuming the same disk scale height before and after the expansion). Even if the initial luminosity of this system was close to the Eddington limit, it is now a factor of 100 lower, implying a minimum SBH mass of $\\sim5\\times10^9\\,\\Msun$, given the estimated bolometric luminosity. It also follows from conservation of angular momentum that in process of disk expansion a sizable portion of the disk mass will be accreted, $M_{\\rm acc} \\sim M_{\\rm disk} / \\sqrt{7}$. Given the SBH recoil velocity and the estimated accretion time scale of $\\sim 10^7$~yr \\citep{bl08}, this implies that, by the time this fraction of disk mass is accreted, the receding AGN should be at least few tens of kiloparsecs away from the host galaxy. This poses a challenge for the recoiling SBH model given the requirement for the AGN to power the observed emission lines in both its own BLR and NLR and the NLR bound to the host galaxy. More specifically, the ionization parameter of the r-NLR should be considerably lower than what is observed (see our comparison with photoionization models in \\S\\,\\ref{S_picture}). We note that origin of the widths and shifts of the [NeIII] and other b-NELs in an outflow or the swept up ISM gas cannot be excluded. In such a scenario the geometry and velocity distribution of the line-emitting gas is likely to be very complex.  A detection of an AGN recoiling at nearly the maximum speed predicted by numerical relativity requires a combination of coincidences: (a) its recoil velocity vector must lie close to our line of sight, (b) the parameters of the pre-coalescence binary should have been such as to give a recoil speed close to the maximum, and (c) we should be observing the object in a relatively narrow time window after the recoil. This would also imply a larger population of systems at lower recoil speeds, which were not found by \\citet{bonning07}.  Indeed, \\citet{dotti08} calculate that SBHBs with parameters similar to those considered here can be $\\sim$100 times intrinsically more common than the recoiling SBHs. A related question, then, is why have there been no more discoveries of SBHBs in the SDSS. The observational biases in the case of subparsec SBHBs are not understood due to uncertainties in the structure and properties of their nuclear regions. The Doppler shift signature may not be observable in all subparsec binaries. The majority of them may either be quiescent or they may exhibit a different signature that we still do not recognize, thus, making the expected (a posteriori) discovery rate for this class of objects difficult to evaluate.   \\subsection{Possible Observational Tests}\\label{S_tests}  The observational property of J0927 that presents the biggest challenge to any model are its two sets of narrow emission-lines. It seems that by coincidence the [\\ion{O}{3}]~$\\lambda$4959 line at $z=0.713$ overlaps with the [\\ion{O}{3}]~$\\lambda$5007 line at $z=0.698$. Moreover, the low signal-to-noise ratio (S/N) at the blue end of the original SDSS spectrum makes the line profiles difficult to characterize. New optical spectra with a higher S/N and spectral resolution are needed to characterize the line profiles and make better measurements of some line strengths (especially [\\ion{Ne}{5}]). Complementary near-IR spectra in the J-band could show the profile of the H$\\alpha$ line and additional narrow, forbidden lines of important diagnostic value.  Spectroscopic monitoring on time scales of years to decades could provide the most direct test of the binary hypothesis.  According to the SBHB model, the emission lines associated with the BLR of the secondary SBH (b-BELs) will shift at a rate of $$ {du_2\\over dt} \\approx 88\\; {(1+q)^2\\over M_{2,8}} \\left({q\\over 0.1}\\right) \\left({\\sin i\\over \\sin 45^\\circ}\\right)^{-3} \\hbox to 4 em{\\hss} $$ \\begin{equation} \\hbox to 4 em{\\hss} \\times \\left({\\sin \\phi\\over \\sin 45^\\circ}\\right)^{-4} \\left({\\cos \\phi\\over \\cos 45^ \\circ}\\right)\\;{\\rm km~s^{-1}~yr^{-1}}\\, . \\label{eq_shift} \\end{equation} Thus, at $\\phi=45^{\\circ}$ the b-NELs would shift by $\\sim 100~{\\rm km~s}^{-1}$ in about a year\\footnote{However, note that there is a range of non-extreme parameter values for which the normalization factor in equation~\\ref{eq_shift} takes a lower value and consequently, the monitoring period may need to be longer.}. Spectroscopic and imaging observations at high angular resolution with the {\\it Hubble Space Telescope} the ({\\it HST}) are necessary in order to measure the redshift of the host galaxy, determine whether its morphology shows signs of a recent merger, and provide a check against the possibility of projection of two otherwise unrelated AGNs along the line of sight. Also, in case of the recoiling SBH, {\\it HST} may resolve an off center AGN, depending on the orientation of the recoil velocity vector with respect to the line of sight. In the case of a binary scenario, if the two black holes (or at least the secondary) are radio emitters, their orbital motion may be detected with the $VLBA$. At the redshift of J0927, the orbital separation of the binary with parameters scaled as in eq~(\\ref{eq1}), translates to an angular separation on the sky of about $\\sim 20\\,\\mu{\\rm as}$. This is comparable to the astrometric precision achieved with the VLBA and for a higher mass binary may approach the expected spatial resolution of the {\\it Square Kilometre Array}.  The confirmation of either model would be a major step towards our understanding of subparsec binaries or recoiling SBHs. A binary holds great promise for revealing the physics of gas in the nuclear region and the accretion signatures in the pre-coalescence phase of its evolution. Discovery of a recoiling SBH, on the other hand, has important implications for understanding the demographics of SBHs and their cosmological spin evolution. Both scenarios have direct ramifications for the rate of SBHB coalescences and the exploration of the binary parameter space, of large importance for gravitational wave observatories such as the {\\it Laser Interferometer Space Antenna}."
    },
    "0809/0809.2465_arXiv.txt": {
        "abstract": "{ {The Hubble Legacy Archive (HLA) aims to create calibrated science data from the Hubble Space Telescope archive and make them accessible via user-friendly and Virtual Observatory (VO) compatible interfaces.  It is a collaboration between the Space Telescope Science Institute (STScI), the Canadian Astronomy Data Centre (CADC) and the Space Telescope - European Coordinating Facility (ST-ECF). Data produced by the Hubble Space Telescope (HST) instruments with slitless spectroscopy modes are among the most difficult to extract and exploit. } {As part of the HLA project, the ST-ECF aims to provide calibrated spectra for objects observed with these HST slitless modes. } {In this paper, we present the HLA NICMOS G141 grism spectra.  We describe in detail the calibration, data reduction and spectrum extraction methods used to produce the extracted spectra.  The quality of the extracted spectra and associated direct images is  demonstrated through comparison with near-IR imaging catalogues and existing near-IR spectroscopy.} { The output data products and their associated metadata are publicly available (\\url{http://hla.stecf.org/}) through a web form, as well as  a VO-compatible interface that enables flexible querying of the archive of the \\nspectra\\ NICMOS G141 spectra. \\new{In total, spectra of \\ntargets\\  unique targets are included.}} ",
        "introduction": "Three of the current Hubble Space Telescope instruments include built-in slitless spectroscopic modes: the Space Telescope Imaging Spectrograph (STIS), the Near Infrared Camera and Multi-Object Spectrometer (NICMOS), and the Advanced Camera for Surveys (ACS).  The main advantage of slitless spectroscopy is that spectra can in principle be obtained from all objects within the field of view of an instrument. The main disadvantages are that the size of the object limits the achievable spectral resolution, that spectra might overlap, and that the background is relatively high because no slit mask prevents the full sky background from illuminating every pixel of the detector. In addition, extracting one-dimensional spectra from such data is a complex process and is often achieved using highly interactive procedures.  A fraction of the slitless data in the HST archive has been collected to obtain spectra of specific objects, but spectra of objects other than the primary targets have in most cases not been extracted or analysed.  The goal of the ST-ECF HLA project is to extract spectra of all objects that have been observed with HST slitless spectroscopy modes and to serve these spectra through an archive with associated descriptions of the spectra, such as how much one spectrum is contaminated by those of nearby objects.  Each of the HST slitless spectrograph modes provides a different capability to HST.  STIS and the ACS prisms cover the UV part of the spectrum that is not accessible from the ground. The NICMOS grisms cover IR wavelengths that can partially be observed from the ground, but where the background is much lower in space. The ACS optical channels covering the optical wavelength range also benefit from a combination of lower background and higher spatial resolution from space. \\new{The STIS grating provides the highest  resolving power among the HST instruments with slitless spectroscopy modes. }  The data analysis for each of the spectrographs presents a significant challenge, and specialised, mostly interactive tools are available for individual instruments such as NICMOS \\citep{nicmoslook}. For ACS a set of non-interactive software tools, called aXe, has been developed which, on the basis of a catalogue of the objects on the associated direct image, extracts one and two-dimensional spectra \\citep{axe1}. This software package was, for example, extensively used for extracting ACS Wide Field Camera slitless spectra in the Hubble Ultra Deep Field (HUDF) \\citep{axe2}.  For the HLA project, we have produced a pipeline that is designed to extract spectra from large numbers of datasets in an unsupervised manner and can be tailored to particular instruments.  We chose to start the slitless spectroscopy HLA project with the NICMOS G141 grism dataset. After a brief description of the available data in the archive (Sect.~\\ref{sec:data}), the basic reduction steps applied to the NICMOS images are described in Sect.~\\ref{sec:images}, followed by a detailed discussion of the spectrum extraction procedure (Sec~\\ref{sec:extract}). In Sec~\\ref{sec:cal}, we show a new set of NICMOS spectrum calibration data.  Our extraction pipeline, PHLAG, is described in (Sect.~\\ref{sec:phlag}).  The extensive set of metadata that allows the archive of spectra to be served through Virtual Observatory (VO) interfaces is summarised in Sect.~\\ref{sec:metadata}.  The calibrated spectra were subject to quality assessment in terms of their astrometric and spectrophotometric properties and shown by internal and external comparisons to be well-calibrated (see Sect.~\\ref{sec:qc}).  \\begin{figure}[t] \\includegraphics[scale=0.48]{grismexample} \\caption{Example of one matching pair of a F160W image (left, image n4k6j1a4q) and the corresponding G141 grism image (right, image n4k6j1zyq)  from HLA dataset N4K6IZZCQ }\\label{fig:grismexample} \\end{figure} ",
        "conclusions": "The HLA NICMOS grism project provides a database of low resolution H-band IR spectra. Most of the spectra have never been previously extracted from the HST NICMOS grism data. The database is useful for work on cool stars and emission line galaxies at $z$ between 1.1 and 1.9, which are readily detected with the NICMOS grism. The calibration of the spectra is based on a new analysis of available calibration data and extraction of the spectra for isolated point source should be close to optimum.  Confused spectra and spectra of extended objects are identified in the data release and are also of high quality.  The absolute and relative flux calibration of the spectra is better than 10\\%, and the wavelength calibration better than 5 nm.  The data release is accompanied by a wide range of auxiliary data and is available through several interfaces. We anticipate a revised data release based on a new version of the STScI CALNICA pipeline in early 2009."
    },
    "0809/0809.0526_arXiv.txt": {
        "abstract": "The supernova (SN) neutronization phase produces mainly electron ($\\nu_e$) neutrinos, the oscillations of which must take place within a few mean-free-paths of their resonance surface located nearby their neutrinosphere. The state-of-the-art on the SN dynamics suggests that a significant part of these $\\nu_e$ can convert into right-handed neutrinos in virtue of the interaction of the electrons and the protons flowing with the SN outgoing plasma, whenever the Dirac neutrino magnetic moment be of strength $\\mu_\\nu < 10^{-11} \\, \\mu_{\\rm B}$, with $\\mu_{\\rm B}$ being the Bohr magneton. In the supernova envelope, part of these neutrinos can flip back to the left-handed flavors due to the interaction of the neutrino magnetic moment with the magnetic field in the SN expanding plasma (Kuznetsov \\& Mikheev 2007; Kuznetsov, Mikheev \\& Okrugin 2008), a region where the field strength is currently accepted to be $B \\gtrsim 10^{13}$ ~G. This type of $\\nu$ oscillations were shown to generate powerful gravitational wave (GW) bursts (Mosquera Cuesta 2000, Mosquera Cuesta 2002, Mosquera Cuesta \\& Fiuza 2004, Loveridge 2004). If such double spin-flip mechanism does run into action inside the SN core, then the release of both the oscillation-produced $\\nu_\\mu$s, $\\nu_\\tau$s and the GW pulse generated by the coherent $\\nu$ spin-flips provides a unique emission offset $\\Delta T^{emission}_{\\rm GW} \\leftrightarrow \\nu = 0$ for measuring the $\\nu$ travel time to Earth. As massive $\\nu$s get noticeably delayed on its journey to Earth with respect to the Einstein GW they generated during the reconversion transient, then the accurate measurement of this time-of-flight delay by SNEWS + LIGO, VIRGO, BBO, DECIGO, etc., might readily assess the absolute $\\nu$ mass spectrum. ",
        "introduction": "The determination of the absolute values of neutrino masses is certainly one of the most difficult problems from the experimental point of view \\citep{Bilenky03}. One of the main difficulties of the issue of determining the $\\nu$ masses from solar or atmospheric $\\nu$ experiments concerns the ability of $\\nu$ detectors to be sensitive to the species mass-square difference instead of so doing to the $\\nu$ mass itself. In this paper we introduce a model-independent novel nonpareil method to achieve this goal. We argue that a highly accurated and largely improved assessment of the $\\nu$ mass-scale can be directly achieved by measurements of the delay in {\\it time-of-flight} between the $\\nu$s themselves and the GW burst generated by the asymmetric flux of neutrinos undergoing coherent (Pantaleone 1992) helicity (spin-flip) transitions during either the neutronization phase, or the relaxation (diffusion) phase in the core of a type II SN explosion. Because special relativistic effects do preclude massive particles of traveling at the speed of light, while massless do not (the { {\\it graviton}} in this case), the measurement of this $\\nu$ time lag leads to a direct accounting of its mass. We posit from the start that two bursts of GW can be generated during the PNS {\\it neutronization phase} through spin-flip oscillations: a) one signal from the early conversion of active $\\nu$s into right-handed partners, at density $\\rho \\sim $ few $10^{12}$ g~cm$^{-3}$, via the interaction of the Dirac neutrino magnetic moment (of strength $\\mu_\\nu < (0.7 - 1.5) \\times 10^{-12} \\, \\mu_{\\rm B}$, with $\\mu_{\\rm B}$ being the Bohr magneton) with the electrons and the protons in the SN outflowing plasma. Specifically, the neutrino chirality flip is caused by the scattering via the intermediate photon (plasmon) off the plasma electromagnetic current presented by electrons: $\\nu_L e^- \\longrightarrow \\nu_R e^-$, protons: $\\nu_L p^+ \\longrightarrow \\nu_R p^+$, etc. b) a second signal in virtue of the reconversion process of these sterile $\\nu$s back into actives some time later, at lower density, via the interaction of the neutrino magnetic moment with the magnetic field in the SN envelope. The GW characteristic amplitude, which depends directly on the luminosity and the mass square-difference of the $\\nu$ species partaking in the coherent transition \\citep{pantaleone92}, and the GW frequency of each of the bursts are computed. Finally, the time-of-flight delay $\\nu \\leftrightarrow {\\textrm GW}$ that can be measured upon the arrival of both signals to Earth observatories is then estimated, and the prospective of obtaining the $\\nu$ mass spectrum from such measurements is discussed. ",
        "conclusions": "In most SN models \\citep{burrows96,mezzacappa,beacom01,vanputten02} the neutronization burst is a well characterized process of {\\it intrinsic} duration $\\Delta T \\simeq 10$~ms, with its maximum occurring within $3.5 \\pm 0.5$~ms after core collapse \\citep{mayle87,walker1987A,vanputten02,burrows96}. This timescale relates to the detectors approximate sensitivity to $\\nu$ masses beyond the mass limit \\begin{equation} m_\\nu > 6.7 \\times 10^{-2}~{\\rm eV} \\left[{\\frac{2.2 ~\\rm Mpc}{D}} \\,{\\frac{\\Delta T}{\\rm 10~ms}}\\right]^\\frac{1}{2} \\left(\\frac{|\\vec{p}|}{10~\\rm MeV}\\right). \\label{mass-sensitivity} \\end{equation} A threshold in agreement with the current bounds on $\\nu$ masses \\citep{fukuda98}.  \\begin{table*} \\caption{\\label{tbl-2}{Time delay between GW and ($|\\vec{p}| = 10$~MeV) $\\nu$ bursts from a SN neutronization, as a function of $\\nu$ mass and distance. } } \\begin{tabular}{ccccccc} \\\\ \\hline\\hline \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\nu$ flavor } & \\multicolumn{1}{c}{{$\\nu$ Mass } } & \\multicolumn{1}{c}{ GC } & \\multicolumn{1}{c}{ LMC } & \\multicolumn{1}{c} { M31 } & \\multicolumn{1}{c} { Source } \\\\ \\vspace{3pt} {   } & {   } & { [eV] } & { [10 kpc] } & { [55 kpc]} & { [2.2 Mpc] } & { [11 Mpc] } \\\\ \\hline \\hline {   }  & { $\\nu_1$ }  & { $10^{-3}$  } & { $5.15\\times 10^{-9} $ } & { $2.83\\times 10^{-8} $ } & { $1.13 \\times 10^{-6}$ } & $5.66\\times 10^{-6}$ \\\\ { {$\\Delta T^{arrival}_{\\rm GW \\leftrightarrow \\nu }$ } } & { $\\nu_2$ } & { $1.0$  } &  { $5.15\\times 10^{-3} $ } & { $2.83\\times 10^{-2}$ } & { 1.13 } & 5.66 \\\\ { [s]  } & { $\\nu_3$ }  & { $2.5$ } &  { $0.32 $  } & { $1.7$ } & { 68.8 } & 344.0  \\\\ \\hline \\hline \\end{tabular} \\end{table*}  Nearby SNe will somehow be seen. Apart from GW and $\\nu$s, $\\gamma$-rays, x-rays, visible, infra-red, or radio signals will be detected. Therefore, their position on the sky and distance ($D$) may be determined quite accurately, including; if far from the Milky Way, their host galaxy \\citep{Ando05}. Besides, the Universal Time of arrival of the GW burst to three or more gravitational radiation interferometric observatories or resonant detectors will be precisely established \\citep{schutz86,arnaud02}. The uncertainty in the GW timing depends on the signal-to-noise ratio ($SNR$) as $\\Delta T ({\\textrm GW}|_{ D = 10 \\textrm kpc}) \\sim 1.45 \\tau/SNR \\sim 0.15$~ms, with $\\tau \\sim 1$~ms the {\\it rms} width of the main GW peak \\citep{arnaud02}. Meanwhile, the type of $\\nu$ and its energy and Universal Time of arrival to $\\nu$ telescopes of the SNEWS network will be highly accurately measured \\citep{Antonioli99,beacom99}. The $\\nu$ timing uncertainty is $\\Delta T_\\nu^{\\rm max} = \\sigma_{flash}(N_\\nu)^{-1/2}$, with $\\sigma_{flash} \\sim (2.3 \\pm 0.3)$~ms, and $N_\\nu$ the event statistics (proportional to $D$). This leads to the SN distance-dependent uncertainty in the $\\nu$ mass:  $\\delta m^2_\\nu \\propto \\Delta T_\\nu^{\\rm max}/D \\sim 0.5-0.6$~eV$^2$ \\citep{arnaud02}, which implies a $m_\\nu \\sim 7 \\times 10^{-1}$~eV, which is consistent with our previous estimate (\\ref{mass-sensitivity}). Hence, those $\\nu$s and their spin-flip conversion signals must be detected.  Therefore, the left-hand-side of Eq.(\\ref{delay}), i.e., the time-of-flight delay $\\Delta T_{\\rm GW \\leftrightarrow \\nu}$, will be measured with a very high accuracy. With these quantities a very precise and stringent assessment of the absolute $\\nu$ mass-eigenstate spectrum will be readily set out by means not explored earlier in astroparticle physics: An inedit technique involving not only particle but also GW astronomy. For instance, at a 10 kpc distance, e.g., to the galactic center (GC in Fig. \\ref{GW-sensitivity}), the resulting time delay should approximate: $ \\Delta T_{\\rm GW \\longleftrightarrow \\nu} = 5.2 \\times 10^{-3} \\;$s, for a flavor of mass $m_{\\nu} \\leq 1~{\\rm eV}$ and $|\\vec{p}| \\sim 10$~MeV. A SN event from the GC or Large Magellanic Cloud (LMC) would provide enough statistics in SNO, SK, etc. $\\sim 5000-8000$ events, so as to allow for the definition of the $\\nu$ mass eigenstates \\citep{beacom01}. Farther out $\\nu$ events are less promising in this perspective, but we stress that one $\\nu$ event collected by the planned Megaton $\\nu$ detector, from a large distance source, may prove suffice, see further arguments in \\citep{Ando05}."
    },
    "0809/0809.4243_arXiv.txt": {
        "abstract": "We study the acoustic properties of the solar chromosphere in the high-frequency regime using a time sequence of velocity measurements in the chromospheric \\caii~854.2 nm line taken with the Interferometric Bidimensional Spectrometer (IBIS). We concentrate on quiet-Sun behavior, apply Fourier analysis, and characterize the observations in terms of the probability density functions (PDFs) of velocity increments. We confirm the presence of significant oscillatory fluctuation power above the cutoff frequency and find that it obeys a power-law distribution with frequency up to our 25~mHz Nyquist limit. The chromospheric PDFs are non-Gaussian and asymmetric and they differ among network, fibril, and internetwork regions. This suggests that the chromospheric high-frequency power is not simply the result of short-period waves propagating upward from the photosphere but rather is the signature of turbulence generated within the chromosphere from shock oscillations near the cutoff frequency. The presence of this pervasive and broad spectrum of motions in the chromosphere is likely to have implications for the excitation of coronal loop oscillations. ",
        "introduction": "\\label{sec:intro}  The role played by high-frequency acoustic oscillations in the heating of the solar chromosphere and corona has been addressed in numerous studies. Because acoustic oscillations above the cutoff frequency freely propagate upward, they have long been considered candidates for transporting to the outer solar atmosphere some of the abundant mechanical energy generated below the photosphere \\citep[see review by][]{1996SSRv...75..453N}.  The primary test for this theory has been to measure the high-frequency oscillations in the photosphere and to estimate whether they contain sufficient acoustic flux to balance the chromospheric and coronal losses. For example, \\cite{2002A&A...395L..51W} and \\cite{ 2007SoPh..242....9A} have looked for power in velocities and intensities measured in photospheric lines. Other authors  \\citep{2005A&A...430.1119D, 2004ApJ...617L..89D} have used {\\em Transition Region and Coronal Explorer} ({\\em TRACE}) UV continuum images to measure power spectra of the intensity fluctuations in the upper photosphere. \\citet{2006ApJ...646..579F} also used {\\em TRACE} but deduced the upward acoustic flux through comparison with hydrodynamic simulations, finding that the energy carried by high-frequency waves is insufficient to balance the chromospheric losses. This analysis was recently repeated by \\citet{2007PASJ...59S.663C} using the Solar Optical Telescope aboard {\\em Hinode}.  Chromospheric oscillations have generally been measured using the \\caii~H\\,and\\,K and infrared triplet lines. Imaging measurements using broad filters in H\\,and\\,K are mostly dominated by the photospheric contribution within the filter bandpass and do not provide a clear indication of chromospheric behavior. Traditional slit spectroscopy gave evidence of significant oscillatory power above the acoustic cutoff frequency \\citep[e.g.][]{1963AnAp...26..368E,1974ApJ...189..359L, 1978A&A....70..345C, 1981A&A....97..310M, 1990A&A...228..506D}. The breakthrough simulation by \\citet{1997ApJ...481..500C} of the \\caii\\ H spectrogram sequence of \\citet{1993ApJ...414..345L} % demonstrated that photospheric oscillations produce chromospheric shocks through upward propagation, but the role they play in chromospheric heating and their energetics are still under debate \\citep[e.g.][]{2007ApJ...671.2154K}.  In this Letter we analyze a time series of chromospheric and photospheric velocity images, obtained at high spatial resolution with the Interferometric Bidimensional Spectrometer \\citep[IBIS;~][]{2006SoPh..236..415C, paper2}. We investigate the scaling behavior of high-frequency velocity fluctuations in chromospheric regions with different magnetic topologies, to study the basic properties of chromospheric turbulence. In Sec.~\\ref{sec:obs} we describe the observations and discuss the velocity power spectra, in Sect.~\\ref{sec:turbulence} we present the analysis of statistical properties of velocity increments, and in in Sec~\\ref{sec:conclusions} our conclusions are summarized. ",
        "conclusions": "\\label{sec:conclusions}  We performed Fourier analysis of spatially and temporally resolved measurements in the \\caii~854.2 nm line. We find that high-frequency velocity fluctuations in the lower chromosphere follow a clear power-law behavior, suggesting turbulence in the chromospheric plasma. We studied the PDFs of the velocity increments as a function of the time lag in order to characterize this turbulence for each of the three defined chromospheric regions (i.e. network, fibril, and internetwork), which showed stronger intermittency in the network areas.  A high-frequency tail is also seen in the photospheric velocity up to a frequency of almost 15 mHz. This is limited by the higher noise level relative to the observed velocities in the photosphere compared to the chromosphere due to the strong density decrease with height. Our observed photospheric power spectra show a smooth monotonic decrease with frequency, without an increase or structure above the cutoff frequency as predicted, for example, by \\citet{1997A&A...324..717T}  Although it is possible that the high-frequency velocities seen in the chromosphere are simply due to the upward propagation of short-period acoustic waves, we find that the chromospheric power spectra and the PDFs are significantly different from the photospheric power spectrum and PDF, indicating that the velocities in the chromosphere at these frequencies are not merely the result of upward-propagating waves. As \\citet{2002ESASP.505..293C} point out, only a small percentage of the power present in the photosphere at high-frequencies would be expected to survive up to higher layers because of radiative damping.  We suggest instead that the observed chromospheric turbulence is generated by the acoustic shocks that are present at this height because of the steep vertical density gradients and that the non linear shock processes produce the cascade of energy to higher frequencies. This provides a mechanism for depositing energy into the chromosphere that is more evenly distributed both temporally and spatially than the shock occurrences themselves. Since shocks are observed to be abundant (with different characteristics) in both the network and internetwork regions of our data \\citep{2008arXiv0807.4966V}, similar turbulent cascades are indeed to be expected in both regions. The resemblance of the network and fibril power spectra may result from the direct connection of the magnetic field from the network into the fibrils.  The presence of non linear turbulent dissipation in the chromosphere is an important consideration in the modeling of this layer of the solar atmosphere. The presence of non linearities and the generation of high-frequency fluctuations there suggest that traditional oscillation analysis, based on the phase and coherence between the photospheric and chromospheric fluctuations, may not properly capture the complexities of the dynamic behavior. Extrapolations of the photospheric energy spectrum to higher layers \\citep{2007ApJS..171..520C, 2007ApJ...662..669V} should take into account the modification of the energy spectrum in the chromosphere \\citep{1977SoPh...52..163W}; this modification is due to processes occurring close to the acoustic cutoff frequency.  Finally, the presence of significant mechanical energy in the chromosphere with a broad spectrum of time scales may play an important role in coronal heating \\citep{1994ApJ...435..482P}. The observed high-frequency plasma motions are conveniently generated near the surface where the plasma $\\beta$ is close to 1. This allows for efficient conversion of these purely acoustic motions into a variety of wave modes, through interaction with the pervasive magnetic field, and provides many additional routes for the continued outward propagation of this acoustic energy. If coronal loops are driven at frequencies comparable to the Alfv\\'en crossing time (typically $10\\!-\\!100$ sec), the amplitude of Alfv\\'enic perturbations may increase resonantly and may efficiently dissipate the wave energy in an impulsive fashion \\citep{1984ApJ...277..392H, 1997ApJ...490..442M, 2004PhRvL..92s4501N}. Acoustic shocks are presumably ubiquitous over the solar surface and may therefore provide an additional source of energy for heating the coronal plasma."
    },
    "0809/0809.3655_arXiv.txt": {
        "abstract": "{The atmospheres of (exo) planets and moons, as well as reflection nebulae, contain in general independently scattering particles in random orientation and are often supposed to be plane-parallel. Relations are presented for the (bidirectional) reflection function and several related functions of such a medium in case the directions of incidence and reflection both tend to horizontal directions. The results are quite general. The medium may be semi-infinite or finite, with or without a reflecting surface underneath, and vertically homogeneous or inhomogeneous. Some approximative formulae for the reflection function of a plane-parallel medium with independently scattering particles in random orientation, including Lambert's law, may be very inaccurate if the directions of incidence and reflection are both nearly horizontal. \\\\  \\noindent keywords: planets and satellites - scattering - radiative transfer - reflection nebulae} ",
        "introduction": "A well--known subject in astrophysics concerns multiple scattering of (electromagnetic) radiation in an extended medium containing small, independently scattering particles \\citep[see e.g.][]{Ambarzumian1943,1950ratr.book.....C, 1975QB603.A85S6213,1980vandeHulst,1974SSRv...16..527H, 2004Hovenier,2006Mishchenko}. Examples of such media are provided by the atmospheres of (exo)planets and satellites, as well as reflection nebulae and protoplanetary disks. The medium is often supposed to be locally plane--parallel, so that one can use the theory developed for a plane--parallel atmosphere, i.e.\\ a horizontally homogeneous atmosphere of infinite horizontal extent. The radiation (which we will also call light) coming from a distant source, like the Sun or a star, may illuminate the top of the atmosphere and then be scattered by the particles inside before leaving the atmosphere at the top in all upward directions. This is called reflected radiation and, neglecting polarization, the angular distribution of its (specific) intensity can be expressed by means of the so--called reflection function.  The reflection function is an important fundamental property of an atmosphere, normalized so that it is identically equal to one for a perfectly white surface following Lambert's law. Once the reflection function of an atmosphere has been obtained, one readily finds the angular distribution of the intensity of the reflected radiation for any angular distribution of incident radiation at the top. In the literature this function has a variety of names, such as reflection coefficient \\citep[][]{1975QB603.A85S6213,1997Yan}, bidirectional reflection function \\citep[][]{1999JQSRT..63..409M} and reflectance factor \\citep[][]{1993tres.book.....H}.  Numerous theoretical and numerical studies of reflection functions have been reported in books and papers. Yet, very little attention was given to the limiting case when the directions of incident and reflected radiation both tend to horizontal directions. This case not only provides more insight into the angular distribution of the reflected radiation, but it is also important for exact and approximate computations, in particular when discontinuities are involved. This was shown for the intensity of the reflected radiation, first when polarization is neglected \\citep[][]{2006JQSRT.101....1H} and later when it is taken into account \\citep[][]{2007JQSRT.107...83H}. As far as the reflection function is concerned, an interesting statement was made by \\citet{1980vandeHulst}, namely that the reflection function becomes infinitely large for horizontal directions of both incidence and reflection. But he restricted himself to the azimuth--independent terms in a Fourier series expansion and he did not give any evidence or clarification regarding this statement.  The principal aim of this paper is to present a comprehensive treatment of the behavior of the reflection function and related functions when the directions of incidence and reflection both tend to horizontal directions. The organization of this paper is as follows. Some basic concepts and definitions are discussed in Sect.~\\ref{sect2}. In Sect.~\\ref{sect3}, the limiting process of directions of incidence and reflection tending to horizontal directions, while keeping the azimuth difference of the directions fixed, is considered for, respectively, the reflected intensity, the reflection function and orders of scattering of both. Simple examples are given in Sect.~\\ref{sect4} for vertically homogeneous as well as inhomogeneous media. Section~\\ref{sect5} is devoted to functions that are related to the reflection function. The azimuth dependence of the incident and reflected light is treated in Sect.~\\ref{sect6}. Approximations are discussed in Sect.~\\ref{sect7} and some concluding remarks are presented in Sect.~\\ref{sect8}. ",
        "conclusions": "\\label{sect8}  Numerous complicated equations occur in the theory of multiple light scattering in homogeneous and inhomogeneous plane--parallel atmospheres \\citep[see e.g.][]{1950ratr.book.....C,1975QB603.A85S6213, 1980vandeHulst,1997Yan,2004Hovenier,2006Mishchenko}. Since there is always a possibility that printed equations contain errors and their derivations are not always given, it is useful to have simple checks available like letting $\\mu$ and $\\mu_0$ approach zero and comparing the results with expressions in this paper.  On performing model computations it is usually very helpful to know and understand what happens with the reflection function or a related function in limiting cases \\citep[][]{1983Icar...55..187I}. This kind of knowledge is provided in this paper even for a complicated model of an inhomogeneous atmosphere with an arbitrarily reflecting surface underneath. In particular one should be prudent in the proximity of discontinuities like those presented in this paper.  We have shown that a discontinuity exists for the reflection function which may hamper interpolation and extrapolation. This problem may be by-passed by multiplying the reflection function by a simple function of $\\mu$ and $\\mu_0$, like $\\mu+\\mu_0$ (cf.\\ Eq.~\\ref{eq9}), before performing the interpolation or extrapolation \\citep[][]{2000JQSRT..64..173K}. This is illustrated in Fig.~\\ref{fig2} by the curve marked $R/R_1$ since in the case considered this curve equals $8 \\mu R(\\mu,\\mu)$.  We have also shown that great care should be exercised with using approximative formulae for the reflection function, since they may lead to large errors for nearly horizontal directions. This holds, for instance, for the \"rapid guess formula\" and the very popular Lambert reflection law, which is often used, e.g.\\ for cloudy atmospheres of planets \\citep[see e.g.][]{2006Kok}. Although approximative formulae exist that do not give large errors for nearly horizontal directions, more accurate results are possibly obtained by computing a few orders of scattering for such directions."
    },
    "0809/0809.0595_arXiv.txt": {
        "abstract": "The nearby long GRB 060614 was not accompanied by a supernova, challenging the collapsar model for long-duration GRBs and the traditional classification scheme for GRBs. However, Cobb et al. have argued that the association of GRB 060614 and its host galaxy could be chance coincidence. In this work we calculate the probability for a GRB to be randomly coincident with a galaxy on the sky, using a galaxy luminosity function obtained from current galaxy surveys. We find that, with a magnitude limit that current telescopes can reach and an evolving galaxy luminosity function obtained from VVDS, the probability for chance coincidence of a GRB with a galaxy of redshift $<1.5$ is about several percent. These results agree with previous estimates based on observed galaxies. For the case of GRB 060614, the probability for it to be coincident with a $z<0.125$ galaxy by angular separation $<0.5^{\\prime \\prime}$ is $\\approx 0.02\\%$, indicating that the association of GRB 060614 and its host galaxy is secure. If the telescope magnitude limit is significantly improved in future, the probability for GRB-galaxy association will be considerably large, making it very problematic to identify a GRB host based only on the superposition of a GRB and a galaxy on the sky. ",
        "introduction": "\\label{sect:intro}  Observation of host galaxies of gamma-ray bursts (GRBs) is very important for understanding the nature of GRBs. Current observations reveal that long duration GRBs occur in star-forming galaxies, consistent with the general belief that long GRBs are produced by the death of massive stars \\citep[and references therein]{con05,fru06,tan07,wai07}. The discovery of the connection between long GRBs and core-collapse Type Ibc supernovae \\citep[and references therein]{gal98,li06,woo06b} supports the collapsar model of long GRBs (MacFadyen \\& Woosley 1999; MacFadyen, Woosley \\& Heger 2001).  In contrast, short-duration GRBs are found in both early- and late-type galaxies, similar to the situation of Type Ia supernovae. The rate of star formation in the host galaxies of short GRBs is often lower than that in the hosts of long GRBs \\citep[and references therein]{ber06}. So far no supernovae have been found to be associated with short GRBs.  The difference in the observed host properties for short and long GRBs supports the idea that short and long GRBs have different progenitors. Long GRBs are believed to arise from the death of massive stars (the collapsar model)---most likely the Wolf-Rayet stars since all observed supernovae associated with GRBs are Type Ic, while short GRBs are more likely produced by the merger of compact stars---neutron star-neutron star merger and black hole-neutron star merger (Li \\& Paczy\\'nski 1998; O'Shaughnessy, Kalogera \\& Belczynski 2007).  However, the above scenario is challenged by the observation of GRB 060614. This is a long burst with a duration $\\sim 100$~s, with a host galaxy at redshift $z=0.125$ \\citep{del06,fyn06,geh06}. For a long GRB that has such a low redshift it is expected that a supernova associated with it should be observed. However, despite extensive observation on its host, no supernova has been found down to limits fainter than any known Type Ic SN and hundreds of times fainter than the archetypal SN 1998bw that accompanied GRB 980425. This challenges the ordinary GRB classification scheme based on GRB durations and the general belief that long GRBs are produced by the core-collapse of massive stars \\citep{zha06,wat07}.  In fact, except its duration, GRB 060614 is much like a short GRB in many aspects. Besides the fact that it has no associated supernova, GRB 060614 has a vanishing spectral lag that is typical of short GRBs \\citep{man07}. Its lightcurve has a very hard and short-duration initial peak, followed by an extended soft emission. \\citet{zha07} have shown that if this burst had an eight time smaller total energy, it would have been detected by BATSE as a marginal short-duration GRB, and would have properties in the {\\em Swift} BAT and XRT bands similar to GRB 050724.  GRB 060505 has a duration $\\sim 4$ s and a host galaxy at $z=0.089$. No supernova has been detected at the location of this burst also \\citep{fyn06,ofe06}, so GRB 060505 has also been considered as a long GRB without a supernova. It has been argued that GRB 060505 is indeed a short burst \\citep[see, however, Th\\\"one et al. 2008; McBreen et al. 2008]{ofe06,lev07}. Models for ``long GRBs without supernovae'' have been proposed (King, Olsson \\& Davies 2007).  On the other hand, it has been suggested that GRB 060614 and its host galaxy was just a coincidence rather than a physical association \\citep{cob06b}. By counting the number of galaxies observed by SMARTS in a field centered on the burst, \\citet{cob06b} showed that the probability for a chance superposition of GRB 060614 and a galaxy along the line of sight is $\\sim 1\\%$. This probability is high enough to cause that several cases of chance superposition may have happened for {\\em Swift} GRBs. This conclusion is enforced by a more detailed study by \\citet{cob08}.  The results of Cobb et al. have raised an important question in identifying GRB host galaxies based only on the superposition of a GRB and a galaxy on the sky \\citep{leva07}. For a telescope with very high sensitivity, it would observe many galaxies on the sky, then the probability for a GRB to be aligned with a galaxy could be high. Then, unavoidably, some GRB hosts discovered in this approach might be superficial, i.e. they are not physically related to the GRBs. Cobb et al. obtained their results by using galaxy survey data. It is interesting to verify these results with an independent, more theoretical approach.  In this paper, we calculate the probability for a GRB to be coincident with a galaxy on the sky using galaxy luminosity functions and compare the results with that of Cobb et al. obtained with different ways. Then, we use our results to assess at what a level we can trust the GRB host galaxies that have been found so far. The UVOT on {\\em Swift} can resolve a source to sub-pixels ($\\sim 0.2^{\\prime\\prime}$). Hence, in our calculations we regard a GRB as a point source.  The approach that we adopt has the benefit of extending beyond the limit of current surveys and to broader types of problems. For example, with slight modification it can be applied to the calculation of the probability of Ly$\\alpha$ forests in the spectra of quasars and GRBs which has important applications in probing the high-$z$ Universe \\citep{loe02}. ",
        "conclusions": "We have calculated the probability for a GRB to be coincident with a galaxy on the sky, using the luminosity function and the radius-luminosity relation derived from the SDSS and VVDS surveys.  Since there is not a reliable luminosity function available to higher redshifts, the probability is calculated only up to a redshift $z\\sim 1.5$ (Fig. \\ref{probvvds}). The results are in agreement with that of \\citet{cob06b} and \\citet{cob08} which were obtained with different approaches. The total probability at $z=1.5$ is a few percent.  We have also calculated the probability of chance coincidence with a criterion that a GRB is considered to be associated with a galaxy if the distance from the GRB to the galaxy center is smaller than $1^{\\prime \\prime}$ (Fig.~\\ref{probvvds}, dashed line). This probability is larger than that calculated with $R_{90}$ for $z>0.7$ (Fig.~\\ref{probvvds}, solid line), caused by the fact that for a fixed angular separation the corresponding linear separation increases with $z$ and $R_{90}$ decreases with $z$.  Although the chance probability is small, it warns us that identifying a GRB host based only on the superposition of a GRB with a galaxy on the sky is dangerous. So far about 350 GRBs have been detected by {\\em Swift}, our results imply that several chance coincidence of a GRB with a galaxy might have already happened. As a result, some GRB hosts that have been found might be superficial. However, for the case of GRB 060614, calculation of the chance superposition of it and a $z<0.125$ galaxy with separation $<0.5^{\\prime\\prime}$ leads to a probability $P = 0.02\\%$, consistent with the result of \\citet{gal06}. This small probability indicates that the association of GRB 060614 and its host is secure.  Obviously, a secure identification of a GRB's host would be obtained by (1) the superposition of the GRB with a galaxy; and (2) the afterglow of the GRB and the host candidate give rise to the same measured redshift.  We have also calculated the probability directly from the data of SDSS, following the approach of Cobb et al. The results are presented in Appendixes A, which agree with our analytical results."
    },
    "0809/0809.2073_arXiv.txt": {
        "abstract": "We present nine epochs of {\\it Hubble Space Telescope} optical imaging of the bipolar outflow from the pre-main sequence binary XZ Tauri.  Our data monitors the system from 1995 -- 2005 and includes emission line images of the flow.  The northern lobe appears to be a succession of bubbles, the outermost of which expanded ballistically from 1995 -- 1999 but in 2000 began to deform and decelerate along its forward edge.  It reached an extent of 6\\asec from the binary in 2005.  A larger and fainter southern counterbubble was detected for the first time in deep ACS images from 2004.  Traces of shocked emission are seen as far as 20\\asec south of the binary.  The bubble emission nebulosity has a low excitation overall, as traced by the [S~II]/H$\\alpha$ line ratio, requiring a nearly comoving surrounding medium that has been accelerated by previous ejections or stellar winds.  Within the broad bubbles there are compact emission knots whose alignments and proper motions indicate that collimated jets are ejected from each binary component.  The jet from the southern component, XZ Tau A, is aligned with the outflow axis of the bubbles and has tangential knot velocities of 70 -- 200 km s$^{-1}$.  Knots in the northern flow are seen to slow and brighten as they approach the forward edge of the outermost bubble.  The knots in the jet from the other star, XZ Tau B, have lower velocities of $\\sim$ 100 km s$^{-1}$.  To explain the observations of the outer bubble, we propose that the XZ Tau A stellar jet underwent a large velocity pulse circa 1980.  This ejection quickly overtook older, slower-moving ejecta very near the star, producing a $\\sim$ 70 km s$^{-1}$ shock in a hot (T$\\sim$ 80,000 K), compact ``fireball''.  The initial thermal pressure of this gas parcel drove the expansion of a spherical bubble.  Subsequent cooling caused the bubble to transition to ballistic expansion, followed by slowing of its forward edge by mass-loading from the pre-shock medium.  Repeated pulses may explain the multiple bubbles seen in the data. Collimated jets continue to flow through the bubble's interior, and with the fading of the original fireball they are becoming the primary energizing mechanism for the emission line structures.  Future evolution of the flow should see the outer bubble structures fade from view, and the emergence of a more typical Herbig-Haro jet/bowshock morphology.  We present a preliminary numerical model of a pulsed jet to illustrate this scenario. ",
        "introduction": "XZ Tauri (HBC 50, Haro 6-15) is a pre-main sequence binary system located in the L1551 molecular cloud at a distance of 140 pc.  Both components of the 0\\farcs 3 binary have strong emission lines that are characteristic of classical T Tauri stars (Hartigan \\& Kenyon 2003; White \\& Ghez 2001).  These studies have determined spectral types of M3-M3.5 for the southern component and M1.5-M2 for the northern one.  The spectrum of the S component in our data (hereafter XZ Tau A) indicates that it is a typical T Tauri star, while the N component (XZ Tau B) has large amount of spectral veiling and numerous emission lines, especially Ca~II (Hartigan \\& Kenyon 2003).  The relative brightnesses of the two components fluctuates at optical wavelengths, with the S component 1--2 mags brighter in 1995 and 1997 {\\it Hubble Space Telescope} ({\\it HST}) images (Krist et al.  1997; White \\& Ghez 2001) and the N component brighter in 1996 and 2000 (White \\& Ghez 2001; Hartigan \\& Kenyon 2003).  The XZ Tau system is an X-ray source with strong variability and a hard spectrum (Carkner et al. 1996; K\\\"onig et al. 2001; Giardino et al. 2006).  It is also a 3.6 cm radio continuum source (Rodriguez et al. 1994).  XZ Tau was identified by Mundt et al. (1990) as the source of a bipolar, collimated outflow (HH 152) along a position angle of $\\sim15^{\\circ}$ with radial velocities of up to $\\sim80$ km s$^{-1}$, and the blueshifted lobe to the N.  Welch et al. (2000) found an expanding shell of $^{13}$CO emission $\\sim0.1$ pc across and centered on XZ Tau, which may be related to this outflow.  Krist et al. (1997, 1999) resolved the inner parts of the outflow using PSF-subtracted {\\it HST} images, finding an unusual bubble of emission nebulosity within 5\\asec of the star.  Marked changes were seen in this bubble over 1995 -- 1998: the onset of limb brightening, the dissipation of a bright knot within the bubble, and large outward proper motions ($\\sim$ 150 km s$^{-1}$) indicative of a dynamical age of only 20 years.  However, they were unable to determine if the bubble energetics were dominated by the energy of an initial blast or the ongoing input of energy from a continuous jet.  To measure flow accelerations and monitor the growth and decay of emission knots, multiple epochs and high spatial resolution are needed.  {\\it HST} is a particularly valuable tool for the study of Herbig-Haro flows such as XZ Tau's. For sources at 140 pc distance, its $\\sim$0\\farcs 05 resolution corresponds to linear scales of $10^{12}$ m, small enough to resolve the cooling zones of radiative shocks. This resolution is also sufficient to measure proper motions of $\\sim$100 km s$^{-1}$ on temporal baselines of only one year.  We report here the results of an imaging campaign spanning 1995 -- 2005 that monitored the evolution of the XZ Tau outflow.  The new data include emission line images that reveal the excitation structure of the XZ Tau bubble, and deeper, wider-field images taken with the {\\it HST} Advanced Camera for Surveys (ACS).  Some of our data has also been discussed by Coffey et al. (2004). ",
        "conclusions": "The nine epochs of HST imaging reveal a complex combination of outflows from both stars and their interaction with the surrounding medium.  Each star has a collimated, bipolar jet.  Large bubbles are aligned along the XZ Tau A jet axis.  Knots in the XZ Tau A jet move through these bubbles and are seen catching up with the apex of the outer bubble, a few arcseconds N of the binary.  The S side of the flow shows a faint counterpart to the outermost N bubble, and two interior, compact, bow-shaped knots that suddenly appeared in 2004.  Evidence for a more extended flow to the N is seen in ground-based images, and to the S extending to distances of 20\\asec where the H$\\alpha$-A structure is found oriented perpendicular to the jet.  In comparison, the visible portion of the XZ Tau B jet consists of only knots N4 and S4 that show proper motions about half those of the fastest XZ Tau A knots.  It is clear that XZ Tau A is the dominant outflow source in the system, which is suprising given that XZ Tau B is much more photometrically variable and has a more prominent emission line spectrum.   \\subsection{Origin of the Bubbles}  The XZ Tau bubbles are unlike the large bowshocks seen in other Herbig-Haro flows and numerically modeled by various workers.  A prototypical bowshock represents the interaction region where a narrow, fast-moving jet impacts material along the flow axis.  The impacted gas can be either ambient interstellar material, or slower-moving ejecta output previously by the source. The brightest emission and highest excitation is found at the bowshock apex, where normal incidence of the flow along the jet axis creates the largest shock velocity difference.  Lower shock velocities at oblique angles along the sides of the bowshock produce fainter emission and weaker excitation along the trailing bowshock wings.  Examples include the HH 1 and HH 34 systems (Hester et al. 1998; Reipurth et al. 2002), and the H$\\alpha$ emission structures at northern limit of the HL Tau jet in Figure 2.  In comparison, the XZ Tau bubbles lack the characteristic peak brightness and excitation at their apex; the outer bubble becomes limb-brightened, fades, decelerates, and deforms along its leading edge over a period of less than 10 years; and the flow has a large opening angle $\\sim$30\\deg near the source.  Taken together, these features suggest that the outer bubble is largely ``coasting'' on the initial momentum with which it formed, with little ongoing energy deposition from the flow of a continuous jet.  We now sketch out a scenario for the origin and evolution of the N outer bubble in the XZ Tau outflow.  It cannot be the result of a wide-angle, ballistic ejection from the star given that it appears limb-brightened in 1998-1999 along more than just its forward edge.  Instead, the spherical symmetry of the bubble in the early epochs suggests radial expansion of an initially hot, compact gas parcel carried along in the general outward flow from XZ Tau A.  The observed transverse proper motions of the bubble (70 km s$^{-1}$; see Table 6) suggest a characteristic initial energy for the bubble.  We postulate that the XZ Tau A stellar jet underwent a large velocity pulse circa 1980.  This ejection quickly overtook older, slower-moving ejecta very near the star, producing a $\\sim$ 70 km s$^{-1}$ shock in a hot (T$\\sim$ 80,000 K), compact ``fireball''.  The initial pressure of this gas parcel created a expanding hot bubble within the flow, but rapid cooling terminated this early inflationary stage and left a ballistically expanding cloud carried along at the mean longitudinal velocity of the flow. Surrounding material swept up in the radial expansion of the bubble led to the formation of a shell, which took on the appearance of a limb-brightened emission bubble when the shell became dense enough for enhanced cooling and/or enough time had elapsed for the post-shock H$\\alpha$/[S~II] cooling zone to form.  The forward edge of the bubble has preferentially slowed and deformed as it encounters higher downstream densities found in the wakes of previous ejections along the flow axis.  Subsequent jet velocity pulses have led to the formation of additional shells trailing the bright outer bubble.  The southern bubble, seen only in 2004, is sufficiently similar to the N one in appearance to suggest a common origin in a symmetric velocity pulse in bipolar jets. The S bubble is slightly larger and centered further from the binary than the N one, and is much fainter.  It lacks the complex internal knots seen within the N bubble.  Because the S bubble is detected at only one epoch, it is unknown if it transitioned from an internally-brightened to limb-brightened nebula as the N bubble did.  It is clear that the N bubble has encountered material dense enough to decelerate and deform its forward edge, creating the cusps.  The faintness and relative simplicity of the S bubble suggest that it is expanding into a more rarified and uniform medium than the northern one.  Coffey et al. (2004) studied the 1995 -- 2001 epochs from our dataset and suggested that XZ Tau B was an eruptive, EXor-type star based on its spectral lines and photometric variability and thus was the likely source of the bubbles.  Without using PSF-subtracted images or tracing knot motions through later epochs, they did not recognize that both stars have jets and that the bubbles are aligned with the XZ Tau A jet.  They modeled the formation of the outer bubble as the result of a 250 km s$^{-1}$ stellar wind with an opening angle of 22\\deg, with the collimated jet having no significant effect.  This scenario, however, does not suitably explain the appearance of the bubble in 1995, the multiple bubbles and their curvatures, or the finite apparent length of the northern outflow.  The interaction of slow and fast moving knots from XZ Tau A seem to better explain these characteristics.  \\subsection{The Jets}  The extended emission bubbles of XZ Tauri enclose multiple compact emission knots which attest to the continued flow of bipolar jets after the bubble ejection events.  The position and proper motions of knots N4 and S4 point to XZ Tau B as their source and a flow PA of 36\\deg.  Both steadily brightened by a factor of $\\sim$5 between 2000 and 2004.  All of the other observed knots are associated with the jet from XZ Tau A along PA 15\\deg, and show more transient brightness variation.  On timescales of just a year or two, knots N1, N2, N3, and S3 completely faded from view.  This requires relatively high gas densities within the knots so that rapid cooling and recombination can take place.  Conversely, S5 and S6 appeared suddenly in 2004 as the brightest and fastest-moving knots in the entire outflow, but were entirely absent two years previously.  The jet's encounter with a sharp density discontinuity in the pre-shock medium at this epoch (likely a structure in the wake of prior ejections) would account for this.  S5 appears the most bowshock-like of all the jet knots, showing both extended curved wings and a bright spot at its apex.  \\subsection{A Trial Model for the XZ Tau A Outflow}  Numerical simulations of Herbig-Haro jets tend to focus on the structure of the classical jet/bowshock working surface at large propagation distances from the outflow source.  For comparision to the XZ Tau bubbles, simulations of a very young pulsed jet are needed in the region very close to the outflow source.  We have begun to explore such models using a third-order WENO (weighted essentially non-oscillatory) method (Shu 1995) for supersonic astrophysical flow simulations (Ha et al. 2005).  WENO schemes are high-order finite difference methods designed for nonlinear hyperbolic conservation laws with piecewise smooth solutions containing sharp discontinuities like shock waves and contacts.  Locally smooth stencils are chosen via a nonlinear adaptive algorithm to avoid crossing discontinuities whenever possible in the interpolation procedure.  The weighted ENO (WENO) schemes use a convex combination of all candidate stencils, rather than just one as in the original ENO method.  \\begin{deluxetable}{lccc} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Outflow Simulation Parameters} \\tablehead{ \\colhead{Parameter} & \\colhead{Units} & \\colhead{Jet} & \\colhead{Ambient} } \\startdata Gas density & H atoms cm$^{-3}$ & 500 & 50 \\\\ Gas velocity & km s$^{-1}$ & 200 & 0 \\\\ Gas temperature & K & 1000 & 10000 \\\\ Sound speed & km s$^{-1}$ & 3.8 & 12 \\\\ \\enddata \\end{deluxetable}  We simulated an approximation to the northern jet outflow and a pair of expanding bubbles, including the effects of radiative cooling using the cooling function from Figure 8 of Schmutzler \\& Tscharnuter (1993).  Two-dimensional simulations were performed on a $750 \\Delta x \\times 750 \\Delta y$ grid spanning $10^{11}$ km on each side.  The jet has a width of $10^{10}$ km, and inflows at Mach 15 with respect to the sound speed in the light ambient gas and Mach 55 with respect to its own internal soundspeed.  Other flow parameters of the simulation are summarized in Table 8, and atomic gas ($\\gamma = 5/3$) is assumed.  The leading edge of the resulting outer bubble propagates initially at 150 km s$^{-1}$, decaying to 90 km s$^{-1}$ at 24 years, with an average velocity over this period of 120 km s$^{-1}$.  The jet was pulsed: a first pulse was turned on for 0.4 yr, and then the jet inflow was turned off for 8.6 yr while the first pulse propagated.  Next a second pulse was turned on for 0.4 yr creating a second bubble, and then the jet inflow was turned off for 5.6 yr.  Figure 12a depicts the flow after the two initial pulses.  A third pulse was turned on for 0.4 yr, and the jet inflow was turned off for 5.6 yr; finally a fourth pulse was turned on for 0.4 yr, and then the jet inflow was turned off for 2.6 yr. Figure 12b depicts the flow after the fourth pulse.  The third and fourth pulses create internal shocks {\\em within}\\/ the bubbles that formed from the first and second bubbles.  The third pulse has almost merged with the second pulse and the third pulse has just propagated into the first bubble at 24 yr.  These simulations produce limb-brightened bubbles and internal knots like those seen on the N side of the outflow from XZ Tau A, supporting the idea that a jet pulsed by strong velocity variations is the basic underlying mechanism of the flow.  The diameter of the simulated bubbles agrees with the Hubble images: roughly $6 \\times 10^{10}$ km ($\\sim$ 3\\asec).  A wide-angle wind is not necessary to match the morphology of the imaged flows.  Several model details remain unoptimized, however: the initially static ambient medium makes the shock in the outer bubble stronger and less spherically symmetric than seen in XZ Tau; the jet width in the simulations is oversized compared to the knots observed in the HST images; and we do not yet have specific predictions for the optical emission line fluxes."
    },
    "0809/0809.2559_arXiv.txt": {
        "abstract": "We present new high resolution spectroscopy from which we derive abundances and radial velocities for stars in the field of the open cluster Tombaugh 2, which has been suggested to be one of a group of clusters previously identified with the Galactic Anticenter Stellar Structure (also known as the Monoceros stream). Using VLT/FLAMES with the UVES and GIRAFFE spectrographs, we find a radial velocity (RV) of $\\langle V_{r}\\rangle = 121 \\pm 0.4$ km s$^{-1}$ using eighteen Tombaugh 2 cluster stars; this is in agreement with previous studies, but at higher precision. We also make the first measurement of Tombaugh 2's velocity dispersion, which is $\\sigma_{int} = 1.8 \\pm 0.3$ km s$^{-1}$. Our abundance analysis of RV-selected members finds that Tombaugh 2 is more metal-rich than previous studies have found; moreover, unlike the previous work, our larger sample also reveals that stars with the velocity of the cluster show a relatively large spread in chemical properties (e.g., $\\Delta$[Fe/H] $> 0.2$). This is the first time a possible abundance spread has been observed in an {\\it open} cluster, though this is one of several possible explanations for our observations. While there is an apparent trend of [$\\alpha$/Fe] with [Fe/H], the distribution of abundances of these ``RV cluster members'' also may hint at a possible division into two primary groups with different mean chemical characteristics --- namely ($\\langle$[Fe/H]$\\rangle$, $\\langle$[Ti/Fe]$\\rangle$) $\\sim$ ($-$0.06, $+$0.02) and ($-$0.28, $+$0.36). Isochrone fitting to the colour-magnitude distribution of apparent Tombaugh 2 members yields an age of 2.0 Gyr, $E(B-V) = 0.3$, and $(m-M)_0 = 14.5$ or $d = 7.9$ kpc for both populations --- parameters that are within the range of previous findings. Based on position and kinematics Tombaugh 2 is a likely member of the GASS/Monoceros stream, which makes Tombaugh 2 the second star cluster within the originally proposed GASS/Monoceros family after NGC2808 to show some evidence for internal population dispersions.  However, we explore other possible explanations for the observed spread in abundances and two possible sub-populations, with the most likely explanation being that the metal-poor ([Fe/H] = $-$0.28), more centrally-concentrated population being the true Tombaugh 2 clusters stars and the metal-rich ([Fe/H] = $-$0.06) population being an overlapping, and kinematically associated, but ``cold''  ($\\sigma_V < 2$ km s$^{-1}$) stellar stream at $R_{gc} \\ge 15$ kpc. ",
        "introduction": "Tombaugh 2 (To2; at Galactic coordinates [$l,b$] = [232.8,$-$6.9]$^{\\circ}$) is an old, distant open cluster in the outer Galactic disc. We originally targeted this cluster as part of an ongoing program \\citep[e.g.,][]{carraro07} to explore in depth a set of star clusters proposed by \\citet{pmf04} --- on the basis of their position and kinematics --- to be possibly associated with the Monoceros stream (Mon; Newberg et al. 2002, Ibata et al. 2003, Yanny et al. 2003), also known as the Galactic anticenter stellar structure (GASS; Rocha-Pinto et al 2003, Crane et al. 2003). GASS/Mon was discovered as an overdensity of stars near the Galactic plane that seems to wrap around the outer parts of the Galactic disc and is frequently explained as tidal debris from the disruption of a dwarf galaxy on a low inclination orbit (e.g., Crane et al. 2003, Yanny et al. 2003,  Penarrubia et al. 2005). To2 has also been associated with the reputed dwarf galaxy in Canis Major (CMa; Bellazzini et al. 2004), which is suggested to be a progenitor of the Mon/GASS structure (Martin et al. 2004a). However, Rocha-Pinto et al. (2006) argue that a more likely progenitor of Mon/GASS may lie somewhere in the region of the sky covered by the former constellation Argo Navis, and that the overdensity reported near CMa is correlated to a window in the dust extinction there; in this case, To2 would be farther displaced from the putative Mon/GASS progenitor. In any of these circumstances To2 is interesting as a potential member of a star cluster system accreted from another galaxy by the Milky Way.  Subsequently, our spectroscopic analysis of this cluster reveals To 2 to be interesting in its own right as the first known open cluster exhibiting possible evidence for multiple populations or an internal population spread, evidence we present and explore here No well studied open cluster has shown a spread in abundances of major elements \\citep[e.g.,][study of M67]{Randich06}. % The possible association of To2 to Mon is a characteristic shared with another unusual star cluster showing multiple populations, the globular cluster NGC 2808 (Crane et al. 2003, Martin et al. 2004a).\\footnote{It should be noted that Casetti-Dinescu et al. (2007) have excluded NGC 2808 from being a part of the Canis Major/Mon structures based on the model of Pe\\~narrubia et al. (2005).  However if Monoceros/GASS and Canis Major are not related or if NGC 2808 is like $\\omega$ Cen (as suggested by Casetti-Dinescu et al. 2007) and has its own tidal stream, the derived NGC 2808 orbit remains consistent with the tilted spatial configuration found in outer disc star clusters (Frinchaboy et al. 2004).} But whereas a few other globular clusters with multiple populations are now known (see \\S6.2), to date no multi-population open clusters have been reported.  Nevertheless, having two unusual, multi-population clusters potentially associated with the GASS/Mon/CMA/Argo structures is intriguing and may point to a mechanism for their creation ... if we can first establish that association {\\it and} determine what GASS/Mon/CMA/Argo is.  The nature and/or reality of the proposed Mon ``tidal debris stream'' and the CMa overdensity have been called into question and are currently a matter of great debate. Ibata et al. (2003) originally proposed that Mon could be related to warps in the outer disc, while Momany et al. (2004, 2006) have argued that much of the observed stellar overdensity associated with Mon --- and particularly all of that associated with CMa (at $l \\sim 240^{\\circ}$) --- is due to the warping and flaring of the Galactic disc, and that no ``extra-Galactic'' component is needed to account for the apparent overdensities in the third quadrant.  This conclusion has been contested by Martin et al. (2004b) on the basis of radial velocities (but cf. Momany et al. 2006), while the discovery of  ``blue plume'' stars in this part of the sky has been used to argue further for the presence of a dwarf galaxy nucleus in CMa (Bellazzini et al. 2004, Martinez-Delgado et al. 2005, Dinescu et al. 2005, Butler et al. 2007). However, these young stars have also been more prosaically attributed to the presence of spiral arm structure (Carraro et al. 2005, Moitinho et al. 2006).  And Grillmair (2006) has quite clearly identified tidal streams across the Galactic anticenter region with relatively low inclination (35$^{\\circ}$) and at a similar distance and main sequence turnoff colour as Mon, though he concludes on the basis of the best fit orbits that these streams have nothing to do with Mon or CMa.  What is clear from all of the debate, and a commonly expressed sentiment by all sides, is that further study, particularly detailed spectroscopic analysis, is needed before the complicated nature of the outer disc, its warp, flare, spiral arm structure and possibly co-located tidal debris and halo substructure, can be confidently disentangled. Thus, a number of studies \\citep[e.g.,][]{yong04,yong2,pmf05} are being conducted to explore further the kinematical and chemical properties of Monoceros and the outer Galactic disc.  Our own venture in this regard starts by focusing on high resolution spectroscopic investigation of old, open star clusters of the outer disc, including those hypothesized to be parts of the debated overdensities.  In \\citet{carraro07} we present high resolution spectroscopy of the first five outer disc open clusters observed in our survey (Ruprecht~4, Ruprecht~7, Berkeley~25, Berkeley~73 and Berkeley~75) and derive detailed abundance analyses and kinematics for these systems. Here, we discuss separately the case of the open cluster To2, which has surfaced as a potentially unusual system among the clusters of the outer disc. In this paper we obtain new estimates of the radial velocities (\\S 3) and abundances ([Fe/H] and [Ti{\\hspace{0.1cm}\\scriptsize\\rm I}/Fe]; \\S 4) of stars in the To2 field.  We discuss the unusual abundance results for To2 in \\S5, and address the possible origins of the observed mixed chemistry, its implications for To2 and the outer disc, and the issue of the possible connection of To2 to the Monoceros/GASS structure in \\S6.  \\begin{figure*} \\includegraphics[width=170mm]{figHIST.RV.eps} \\caption{Radial velocity distribution of all stars observed with the VLT. The cluster is clearly visible as the peak near $\\sim 121$ km s$^{-1}$.\\label{rvhist} } \\end{figure*} ",
        "conclusions": "While this study was intended to verify or refute the connection of To2 with the Monoceros stream/Galactic anticenter stellar structure (GASS), our work has produced more questions than answers.  The observed spread in To2 metallicity and $\\alpha$-abundances was not expected and cannot be simply explained. To date, no other open cluster is known to have a metallicity spread, and it is difficult to understand how a low-mass open cluster might have been able to retain gas after an initial starburst to self-enrich. Next, we review our findings and possible scenarios that could explain the data   \\subsection{Are both populations part of Tombaugh 2?}  First, we ask the question of whether both the MP and MR groups among the To2 RV members are likely to be truly part of the cluster, and then investigate the likelihood that either population is simply contamination of the RV member sample.   \\subsubsection{Is the metal poor population part of Tombaugh 2?}  To2 has been proposed to be part of the putative GASS cluster system \\citep{pmf04,martin04a}, and the MP population is similar, though slightly more metal rich than, other studied Monoceros candidate clusters such as Saurer 1 \\citep{fp02,carraro03,carraro04}, Berkeley 29 \\citep{carraro04,yong04}, and Palomar 1 \\citep{rosen98,ata07}.  These clusters all have ages of approximately 4--6 Gyr,  metallicities between $-$0.7 and $-$0.4 dex, and are all $\\alpha$--enhanced.  Our To2 MP group is younger though similarly moderately metal poor ([Fe/H] $\\sim$ $-$0.3) and also $\\alpha$--enhanced.  This is at least circumstantial evidence that the MP group might be the more likely member population. This subgroup also has a metallicity roughly consistent with previous metallicity determinations for To2 by \\citet{brown96} and \\citet{friel02} --- though it is not impossible that these other researchers could have also been hit with a spate of bad luck in terms of misleading target selection.  On the other hand, it must also be admitted that the MP group both exhibits more scatter in terms of metallicity properties than the MR population and has a smaller number of members --- though this number increases when we account for the additional ``MP'' members found in the \\citet{brown96} and \\citet{friel02} samples that are not part of our sample. Perhaps the strongest support for the supposition that the MP group is a ``real'' part of the cluster comes from the fact that the MP group is strongly centrally concentrated (Figs. 4 and 5) and that it has a marginally smaller velocity dispersion than the MR group, as shown in Table \\ref{compare}.    \\subsubsection{Is the metal rich population part of Tombaugh 2?}  The MR population is indeed less centrally concentrated that the MP population, and in fact has a distribution function not unlike our spectroscopic sampling function, as shown in Figure \\ref{XYplot} and \\ref{INTGplot}.  Given the small number statistics overall, it may not be impossible for us to have been successful at achieving a uniform ``success'' fraction (i.e. in identifying RV members) over the radial range sampled, even by chance. Indeed, the overall distribution of RV members very closely tracks the sampling function. Though Figure \\ref{XYplot} clearly shows a tighter MP distribution than MR population, it is not inconceivable for there to exist a metallicity gradient in the system to yield a differential radial distribution --- on the other hand, the one needed (i.e. more metal poor towards the centre) is quite unusual and not seen elsewhere.   On the other hand, the MR group has more members and is more tightly clumped in [Fe/H].  Given the tight RV distribution, chemical properties, and sheer numbers it seems difficult to believe that the MR population is not some coherent Simple Stellar Population (SSP), whether part of the cluster or a ``contaminant.''  We show below that this population is quite unlike that of the disc field star population (which is not an SSP).  If the true To2 stars are those in the MR group (and excluding the MP group) then To2 is less consistent with being part of the putative GASS cluster system, as it would be inconsistent with the observed metallicity distribution of Monoceros (Crane et al 2003; Sbordone et al 2005; de Jong et al. 2007). In this case, To2 would be the most distant metal-rich Milky Way disk open cluster found to date. It seems unlikely that a 2 Gyr old cluster with near solar [Fe/H] could have formed ``normally'' in the outer disc, given the observed, more depleted chemistry of the surrounding disc stars \\citep[e.g.,][]{yong2,yong3}; this would suggest that the cluster must have formed elsewhere (e.g., in a dwarf galaxy) and been deposited in the outer disc.  This is not out of line with current CDM models for the formation of Milky Way-like galaxies, which suggest that discs form from the outside via continued merging (e.g., Abadi et al 2003). Stars as metal rich as our MR group are indeed found in paradigm accretion events, like the Sagittarius dwarf (e.g., Smecker-Hane \\& McWilliam 2002, Monaco et al. 2005, Chou et al. 2007).   \\subsubsection{Can either population simply be contamination?}  If we believe that only one of the two metallicity populations is truly part of To2, then the other population must be a contaminant of some kind.  Here we address how likely it is that either population may be part of the Galactic disc.  We have already seen (\\S \\ref{model}) that the predicted contamination rate of normal disc stars into our RV membership range based on the Besancon model (see \\S \\ref{model}) is already smaller (by factor of 10$^2$) than either our MP or MR samples --- thus, it seems very unlikely that either population can be a statistical, ``fluke'' superposition of a large number of disc stars with just the right RV. From Poisson statistics, the chances that we would get 7 (e.g., MP) or 11 (e.g., MR) stars as contaminants when the models predict 0.02 stars is $1.1 \\times 10^{-5}$ \\% and $2.3 \\times 10^{-11}$ \\%, respectively. And this ignores the fact that other studies have identified additional stars in the To2 that we would categorize as ``MP.''  Based on the radial distribution of members, the MP population seems more likely to be truly part of the cluster.  This could then imply that the MR group is a contaminant population.  However, stars this metal rich (nearly solar metallicity) at $R_{gc} \\ge 15$ kpc are not consistent with the measured metallicities of red giants in the outer disc \\citep{yong2} as well as outer disc Cepheid stars \\citep{yong3}, which both suggest that the median disc metallicity at this Galactocentric distance should be [Fe/H] $\\sim$ -0.4. These studies of outer disc tracers find no stars as metal rich as our MR group, even for the younger Cepheid populations.   This then might suggest that the MP group could be a contaminating field population, since the chemistry of this group is more in line with that of the \\citet{yong2} and \\citet{yong3} stars in the outer disc.  In turn, this implies that the tight radial concentration of the MP stars is a statistical anomaly. However Monte Carlo simulations show that the probability for a random disc population to achieve the observed level of concentration, taking into account our sampling function, is 0.1\\%. It also implies that previous spectroscopic studies had similarly ``bad luck'' in picking targets to represent To2.   \\subsection{Does Tombaugh 2 really have a metallicity spread?}   The cumulative evidence for To2 --- the nearly identical velocities and velocity dispersions of the stars when divided into MP and MR groups, the positions of the stars in the CMD, the unlikeliness that these stars are field contamination, etc. --- suggests that all of the RV ``members'' in our To2 spectroscopic sample could truly be members of the system. If all RV ``members'' are indeed members of To2 then we are finding evidence for a {\\it significant} metallicity spread within To2.  That we also seem to find a similar bifurcation of [Ti{\\hspace{0.1cm}\\scriptsize\\rm I}/Fe] with [Fe/H], rather than a scatter-diagram in Figure 3 is further evidence of a possible coherent chemical enrichment history. It is also possible, based on an apparent ``break'' in both [Fe/H] {\\it and} [Ti{\\hspace{0.1cm}\\scriptsize\\rm I}/Fe], that the To2 chemical and age distribution suggests the possibility that there may be two populations rather than a ``trend.''  If these populations of varying chemistry do belong to To2, what formation scenarios could create such a unique open cluster?     \\subsubsection{Is To2 actually two ``very close'' clusters?}  What is the possibility that two clusters with the same RV lie at nearly the same distance and have nearly the same metallicity and could also be aligned along the same line of sight? Though such a configuration could produce the observed results in the To2 field, and even account for the apparent ``bimodal'' behavior of the chemical distribution, it seems extremely unlikely given the small number of clusters (less than 20) known this far from the Galactic centre. It is the case that a superposition of clusters would increase the ``discoverability'' of this distant ``double cluster'', but the individual clusters would have to be very close together along the line of sight given that a significant broadening of the main sequence is not seen in the available photometry. However new higher quality photometry would be needed to fully test this case \\citep[as was done with the {\\it Hubble Space Telescope} in the case of NGC 2808;][]{n2808}.  There are a few cases of overlapping clusters, with the most-well studied overlapping pair being NGC 1750 and NGC 1758, with ages of 200 $\\pm$ 50 and 400 $\\pm$ 100 Myr respectively \\citep{gjt98b}. For NGC 1750 and NGC 1758 proper motions were required to be able to separate the two clusters separated by $\\sim 130$ pc along the line of sight \\citep{gjt98b}, but with little noticeable differences seen in the overlapping CMD sequences \\citep{gjt98a}. These clusters are the oldest pair of confirmed overlapping clusters, and, given the age and distance differences, \\citet{gjt98b} conclude that the clusters do not constitute a binary system. However for To2, the fact that mean velocity derived from high-resolution RVs is the same, and different than the Galactic disc trend, makes the possibility of it being part of a ``double cluster'' remote.     \\subsubsection{Does To2 represent the merging of two clusters?}  Another explanation for our results in the To2 field might be that two clusters with nearly the same age and different metallicities could have have collided and merged. Given that open clusters rotate with the Galactic disc, the small relative velocity difference between two clusters with similar orbits makes prolonged near encounters between close clusters possible. Additionally, the relatively large velocity dispersion of To2, $\\sigma_{int} = 1.8$ km s$^{-1}$ (e.g., M67 has $\\sigma_{int} = 0.8$--0.96 km s$^{-1}$; Girard et al. 1989; Frinchaboy \\& Majewski 2007) might suggest that the cluster could have been dynamically heated at some point, possibly by a merger of clusters. A simple calculation based on merging of two equal-mass eliptical galaxies \\citep{bt}, shows that for a two-body encounter with two equally massed clusters (assuming each is $10^4$ M$_{\\sun}$) finds that mergers are possible if two clusters pass within less than 2 pc given a differential velocity of 5 km s$^{-1}$. However as pointed out in the argument above, given the small number of distant clusters the chances of such an occurrence are very small; nevertheless, the merging of two clusters cannot be ruled out as a possible explanation for the observed chemical peculiarities.   \\subsubsection{Has To2 had multiple episodes of star formation?}  The cumulative evidence for To2 --- the nearly identical velocities and velocity dispersions of the stars when divided into MP and MR groups, the positions of the stars in the CMD, the unlikeliness that these stars are field contamination, etc. --- suggests that all of the RV ``members'' in our To2 spectroscopic sample could truly be members of the system. If all RV ``members'' are indeed members of To2 then we are finding evidence for a {\\it significant} metallicity spread within To2.  That we also seem to find a similar bifurcation of [Ti{\\hspace{0.1cm}\\scriptsize\\rm I}/Fe] with [Fe/H], rather than a scatter-diagram in Figure 3 is further evidence of a possible coherent chemical enrichment history. It is also possible, based on an apparent ``break'' in both [Fe/H] {\\it and} [Ti{\\hspace{0.1cm}\\scriptsize\\rm I}/Fe], that the To2 chemical and age distribution suggests the possibility that there may be two populations rather than a ``trend.''  No other open cluster has been found to have an internal metallicity spread. On the other hand, a small number of {\\it globular} clusters have been found with metallicity spreads (e.g., $\\omega$ Cen; Norris, Freeman, \\& Mighell 1996; Majewski et al. 2000; Carraro \\& Lia 2000; Frinchaboy et al. 2002; Villanova et al. 2007), multiple internal populations \\citep[NGC 1851 and NGC 2808; Milone et al. 2007,][]{n2808}, or to represent one of a series of populations in a more complex structure (M54 in the Sagittarius dwarf spheroidal galaxy; Sgr dSph; Siegel et al. 2007). Interestingly, such unusual clusters typically have been associated with the possibility of tidal accretion events.  Could To2 be the remains of something that was once larger --- a globular cluster or dwarf galaxy? This seems like the only way that To2 could have had two episodes of star formation, since the mass of a typical open cluster is generally much too small to retain and self-enrich gas after an initial star burst. The mass of a long-lived open cluster is typically a few $10^3$--$10^4$ M$_{\\sun}$ \\citep[e.g., M67 has a current mass $\\sim 2 \\times 10^3 \\; {\\rm M}_{\\sun}$;][]{hurley05}. A mass of at least $10^7$ M$_{\\sun}$ would be needed for To2 to have retained enough gas to have a second epoch of star formation (e.g., the current mass of $\\omega$ Cen is $2.8 \\times 10^6$; D'Souza \\& Rix 2005). Thus, this would require To 2 either to have had a significant dark matter content, which is ruled out by the velocity dispersion (and in any case would have made To 2 unique among star clusters, open or globular), or to have been initially much larger and then stripped down to its current size. Fortunately for this scenario, the disc is a dynamically brutal environment for clusters and dwarf galaxies, which are expected to undergo vigorous stripping.  Indeed, this is the prevailing scenario to explain the presently visible multiple populations in the $\\omega$ Cen ``globular cluster'' system.  The two apparent sub-populations in To2 do seem to show a chemical trend evoking self-enrichment, with the metal-poor, $\\alpha$-enhanced population formed first, followed by formation of a more metal-rich, non-$\\alpha$-enhanced population. The above ``evolution'' is what is expected from closed-box chemical evolution models. Unfortunately, two other observational facts contradict the self enrichment scenario: (1) The derived isochrone ages for the two populations are the same, within the errors of the fitting, which are $\\pm$ 0.5 Gyr. (2) It is typical in self-enrichment scenarios for the more metal rich population, formed in the bottom of the potential well more recently, to be more centrally concentrated.  Thus, if To2 was the tidal remnant of a larger stellar system one would expect that the ``younger'' metal-rich population would be more centrally concentrated, as is found in the case of $\\omega$ Cen (Pancino et al. 2000) and dSph galaxies (e.g., Sculptor; Tolstoy et al 2004).  However given the available data, we find that the metal-poor stars are more centrally concentrated, in marked contrast to large systems undergoing tidal stripping (Tolstoy et al 2004).  \\subsubsection{A ``hybrid'' solution?}  We have seen how a self-enrichment scenario seems to be contradicted by the relative radial and age distributions of the MP and MR populations.  And we have seen that it is unlikely that To2 represents either dynamically or apparently merged/overlapped systems. Yet the strongly matched velocities and velocity dispersions of the two systems suggests a dynamical connection of the two populations, and there is at least some metallicity and velocity mismatch of these populations with the surrounding disc stars.  There remains one possible scenario that could account for most of the observed properties in the To2 field, and which has a known prototype: The Sgr dSph stream and its cluster system.  Any small spectroscopic survey of a Sgr star cluster, particularly one of those Sgr clusters near the core of the dSph (e.g., Terzan 7, Terzan 8 or Arp 2), would produce a number of the same phenomena we observe in our small spectroscopic survey of the To2 field:  A centrally concentrated, older core of more metal-poor stars that are part of the star cluster, seen against a more broadly distributed, typically younger and more metal rich population of stars, generally unbound, tidally stripped from the parent dSph galaxy, and all with the same radial velocity.  This could be an analogous situation to what we observe in the field of To2, with the primary difference being that we are explaining the properties of an open, not globular, cluster.  Thus, we hypothesize this as the leading explanation for what we are finding in the To2 field:  We propose that the MP population represents the true cluster, as most strongly supported by the radial concentration of these stars, and that the MR population represents stars that are part of a distinct parent system, probably a tidally disrupted dwarf galaxy, that are spread more uniformly across our survey field (and beyond).  Such a scenario is consistent with prevailing models of disc formation, which indeed suggest that discs continually grow by accreting ``subhalos'' around the edge (e.g., Abadi et al 2003, Brook et al. 2004). It also accounts for how we can have two populations with identical RVs, but differing metallicities, and with the more metal-rich population being more extended.  That To2 has been already hypothesized to be a part of the GASS/Monoceros ``tidal stream'' \\citep{pmf04,martin04a} provides additional support for this proposed scenario, although the velocity dispersion of GASS/Mon is much larger than our MR population (Crane et al. 2003, Martin et al 2006). Nevertheless, it is possible that other tidal streams could exist and have contributed star clusters to the outer disc.    \\subsection{Final remarks}   While we have discussed a number of possible explanations for our unusual findings in our survey of the To2 field, the only hypothesis that seems to explain most of the results in a natural way is that To2 is represented by our MP population and that the MR population represents stars from a parent dwarf galaxy that is contributing To2 to our Galactic disc.  All other hypotheses --- self-enrichment, a merger of two clusters, a superposition of two clusters --- have serious problems or are very improbable.  It is also yet unclear whether To2 may be part of the  proposed GASS/Mon system. Surely it is worth mentioning that if To2 actually has two different populations, then it is truly an unusual object. The one clear finding from this work is that further work is needed on the outer Galactic cluster system, in particular To2.  A much larger spectroscopic sample that includes both high precision velocities and detailed abundances would be especially useful. Moreover, the properties of outer disc stars are needed to better constrain the existence, characteristics and origin of the GASS/Monoceros system and other tidally accreted systems in the outer Galactic disc. While RVs and abundances may not provide an explicit separation of an accreted satellite from the outer disc, it is nonetheless essential that the kinematical and more importantly abundance trends in the outer disc are known.  Without clear knowledge of the abundance patterns with radius, especially at large distances and in the second and third Galactic quadrants, one cannot explicitly use abundances to distinguish whether the GASS/Monoceros stream is distinct from the outer Galactic disc.  The present paper is one among many efforts % that aim to correct the deficiency of abundance information known for the outer Milky Way disc."
    },
    "0809/0809.0901_arXiv.txt": {
        "abstract": "$N$-body simulations have unveiled several apparently universal properties of dark matter halos, including a cusped density profile, a power-law pseudo phase-space density $\\rho/\\sigma_r^3$, and a linear $\\beta-\\gamma$ relation between the density slope and the velocity anisotropy. We present a family of self-consistent phase-space distribution functions $F(E,L)$, based on the Dehnen-McLaughlin Jeans models, that incorporate these universal properties very accurately. These distribution functions, derived using a quadratic programming technique, are analytical, positive and smooth over the entire phase space and are able to generate four-parameter velocity anisotropy profiles $\\beta(r)$ with arbitrary asymptotic values $\\beta_0$ and $\\beta_\\infty$. We discuss the orbital structure of six radially anisotropic systems in detail and argue that, apart from its use for generating initial conditions for $N$-body studies, our dynamical modeling provides a valuable complementary approach to understand the processes involved in the formation of dark matter halos. ",
        "introduction": "\\label{introduction.sec} \\defcitealias{2007A&A...471..419B}{Paper~I}  In theoretical astrophysics the steady increase in computational power has sparked a proportional interest and progress in the study of large-scale structure formation. In particular, as $N$-body simulations of cold dark matter halos have become more detailed, several ''universal'' properties have emerged. We highlight three important characteristics.  Firstly, numerous cosmological studies \\citep[e.g.][]{1991ApJ...378..496D, 1994ApJ...434..402C, 1996ApJ...462..563N, 1997ApJ...477L...9F, 1997ApJ...490..493N, 1997ApJ...485L..13C, 1998ApJ...499L...5M, 1999MNRAS.310.1147M, 2000ApJ...529L..69J} revealed similar density profiles over several orders of magnitude in halo mass, with a central cusp and an a $\\rho(r) \\propto r^{-3}$ falloff at large radii. These generalized NFW models can be described by a general 3-parameter family, referred to as the Zhao models (or $\\alpha\\beta\\gamma$-models) \\citep{1990ApJ...356..359H,1996MNRAS.278..488Z}.  They are defined by the density \\begin{equation}\\label{zhao} \\rho(r) = \\frac{2^{(\\gamma_\\infty-\\gamma_0)/\\eta}\\,\\rho_\\txs}{\\left(r/r_\\txs\\right)^{\\gamma_0} \\left(1 + \\left(r/r_\\txs\\right)^\\eta\\right)^{(\\gamma_\\infty-\\gamma_0)/\\eta}}, \\end{equation} or in terms of the logarithmic slope, \\begin{equation} \\gamma(r) = -\\frac{\\txd\\ln\\rho}{\\txd\\ln r}(r) = \\frac{\\gamma_0 + \\gamma_\\infty \\left(r/r_\\txs\\right)^\\eta}{1 + \\left(r/r_\\txs\\right)^\\eta}. \\end{equation} Although in recent years alternative profiles based on the S\\'ersic law have produced equally good results \\citep{2004MNRAS.349.1039N, 2005ApJ...624L..85M}, the generalized NFW models remain very popular and successful to represent dark matter halos.  A second relation was found by \\citet{2001ApJ...563..483T}. These authors identified that the quantity $\\pseudo(r)=\\rho/\\sigma^3(r)$, which has become known as the pseudo phase-space density, behaves as a power law over 2-3 orders of magnitude in radius inside the virial radius, \\begin{equation} \\pseudo(r) \\propto r^{-\\alpha}. \\end{equation} Other studies \\citep[e.g.][]{2004MNRAS.351..237R,2004MNRAS.352.1109A} have confirmed the scale-free nature of $\\pseudo(r)$, and their results indicate that its slope lies in the range $\\alpha=1.90\\pm0.05$. This property is remarkable since the density $\\rho(r)$ nor the velocity dispersion $\\sigma(r)$ separately show a power-law behavior.  Finally, the velocity anisotropy profiles $\\beta(r)$ of dark matter systems also evolve toward a similar shape, steepening gradually from isotropic in the center to radially anisotropic in the outer regions. \\citet{2006NewA...11..333H} suggest a nearly linear relation between the logarithmic density slope $\\gamma(r)$ and the velocity anisotropy profile, based on various types of equilibrated simulations. They proposed the $\\beta-\\gamma$ relation \\begin{equation} \\beta(\\gamma) \\simeq 1 - 1.15(1 + \\gamma/6). \\end{equation}  Several theoretical studies have been made to investigate whether solutions of the Jeans equation exist that encompass the observed properties of dark matter halos. In particular, \\citet{2005MNRAS.363.1057D} investigated the anisotropic Jeans equation constrained by a slightly different form of the pseudo phase-space density, namely $\\pseudo_r(r)=\\rho/\\sigma_r^3$ with $\\sigma_r(r)$ the radial velocity dispersion.  They found a special solution, namely an analytical self-consistent potential-density pair of the form (\\ref{zhao}) with an anisotropy profile \\begin{equation} \\beta(r) = \\frac{\\beta_0+\\beta_\\infty(r/r_\\txa)^{2\\delta}}{1+(r/r_\\txa)^{2\\delta}}, \\label{betagen} \\end{equation} that also has an exactly linear $\\beta-\\gamma$ relation (in other words, $r_\\txa=r_\\txs$ and $2\\delta = \\eta$).  As these Jeans models satisfy the three universal relations mentioned above, we can consider them as representative models for realistic dark matter halos.  The goal of this paper is to take the analytical study of dark matter systems a step further, i.e.\\ to look for full dynamical models that encompass the universal properties found in $N$-body simulations. In concreto we will look for phase-space distribution functions (DFs) $F(\\vec{r},\\vec{v})$ that self-consistently generate the required density, potential, and anisotropy profiles encountered in dark matter halos. Such dynamical models would provide a very useful complementary approach to gain insight into the structure of dark matter halos.  We shall focus on spherical models, for which a dynamical description simplifies substantially as in this case the DF can be expressed as a function $F(\\calE,L)$ of the binding energy and the angular momentum. But even if we limit ourselves to the spherical case, it is not straightforward to obtain such dynamical models.  A simple approach is to solve the Jeans equation and approximate the velocity dispersion profiles by a multivariate Gaussian \\citep{1993ApJS...86..389H}. \\citet{2004ApJ...601...37K} demonstrated however that these systems are far from equilibrium, thus losing their initial dynamical structure as they evolve to non-Gaussian velocity distributions in subsequent $N$-body simulations. We therefore need the tools to derive full self-consistent equilibrium models, governed by DFs with a sufficiently general velocity anisotropy profile.  Simple analytical models are found for only a few special cases, such as the Plummer, Hernquist, isochrone, and $\\gamma$-models \\citep[e.g.][]{1987MNRAS.224...13D, 1990ApJ...356..359H, 2002A&A...393..485B, 2006AJ....131..782A, 2007MNRAS.375..773B}, none of which are able to describe realistic dark matter halos.  For general potential-density pairs, \\citet{1979PAZh....5...77O} and \\citet{1985AJ.....90.1027M} provided an algorithm to construct dynamical models with a specific velocity anisotropy profile of the form (\\ref{betagen}), with $\\beta_0=0$, $\\beta_\\infty=1$ and $\\delta=1$, leaving $r_\\txa$ as a free parameter. While the Osipkov-Merritt method has been widely adopted \\citep[e.g.][]{2000ApJS..131...39W, 2001MNRAS.321..155L, 2004ApJ...601...37K}, comparison with simulations shows that such anisotropy profiles are too steep to describe dark matter \\citep{2005MNRAS.363..705M}. Halos are not completely radial at infinity (i.e.\\ $\\beta_\\infty<1$) and the linear $\\beta-\\gamma$ relation can only be realized with a lower transition rate $\\delta<1$. On a dynamical note, the associated DFs are of the form $F(Q)$ with $Q=\\calE - L^2/2r_\\txa^2$, creating an unphysical cut-off boundary for orbits with $Q<0$. Hence, the Osipkov-Merritt framework is too limited to generate realistic dark matter systems, and a more extensive method is needed.  Recently, \\citet{2008MNRAS.tmp..719W} presented an interesting approach to generate dynamical models for potential-density pairs with a more general anisotropy profile.  They proposed to express the DF as a separable function of the form $f_\\calE(\\calE)\\,f_L(L)$ where $f_L(L)$ is a double power-law function with three parameters $\\beta_0$, $\\beta_\\infty$, and $L_0$. Once their values have been determined, the function $f_\\calE(\\calE)$ is derived from the observed density profile by a numerical inversion. This technique yields a three-parameter anisotropy profile that resembles eq.~(\\ref{betagen}), where $L_0$ has a similar role as $r_\\txa$, and a fixed transition rate $0.5<\\delta<1$. In this manner, the authors were able to construct models with an NFW density with velocity dispersion profiles that agreed with their dark matter simulations.  In this paper, we present a technique that enables us to obtain dynamical models with exactly the four-parameter anisotropy profile~(\\ref{betagen}).  We demonstrated in a previous paper \\citep[][hereafter \\citetalias{2007A&A...471..419B}]{2007A&A...471..419B} how this can be achieved, by postulating a separable parameterized form of the so-called augmented density, which provides an equivalent description of a dynamical system. In certain special cases the transformation from the augmented density to the DF is analytically tractable. We were thus able to derive a family of DFs for the generalized Plummer models (also called $\\alpha$-models or Veltmann models, \\citet{1979AZh....56..976V}) with a linear $\\beta-\\gamma$ relation.  Since the generalized Plummer potential-density pairs form a special class of the Zhao models~(\\ref{zhao}), we will now demonstrate that this result is a first step toward more representative dark matter profiles.  In particular, we seek a family of DFs for the Dehnen-McLaughlin systems, because of their unique property to satisfy a universal density, a power law $\\pseudo_r(r)$, and a linear $\\beta-\\gamma$ relation. We use a quadratic programming technique \\citep{1989ApJ...343..113D} to build DFs as a linear combination of base functions of the form derived in \\citetalias{2007A&A...471..419B}.  In this manner, we demonstrate that it is indeed possible to generate full dynamical models for the Dehnen-McLaughlin Jeans models, thus providing DFs that encompass the observed properties of dark matter halos.  Our paper is organized as follows. In Section~\\ref{preliminaries.sec} we describe the notion of a dynamical model and we briefly summarize the main aspects of the Dehnen-McLaughlin halos.  Next we outline our modeling technique in Section~\\ref{modeling.sec}: we explain the quadratic programming algorithm and we recapitulate the functions that we derived in \\citetalias{2007A&A...471..419B}. With these components, we can build a library to apply the QP-method to the Dehnen-McLaughlin halos. In Section~\\ref{results.sec} we present our results for a set of systems with different velocity anisotropy profiles, and we discuss the moments, phase-space DFs, and energy and angular momentum distributions for these models. Finally, we formulate our conclusions in Section~\\ref{conclusions.sec}. ",
        "conclusions": "\\label{conclusions.sec}  In this paper we presented a set of anisotropic distribution functions $F(\\calE,L)$ for the Dehnen-McLaughlin Jeans models. We constructed these DFs as a linear combination of base functions, which we fitted to the halo density by means of a quadratic programming algorithm. The base functions were specifically designed for this task, derived from a separable augmented density that generates exactly the general four-parameter velocity anisotropy profiles~(\\ref{betagen}). We demonstrated that the resulting fits are very accurate, from the center to large radii, far beyond the range of $N$-body simulations.  This method has several advantages. The DFs can be written as a sum of a double series, without numerical integrations or inversions. Consequently, the DFs and all subsequent moments can be computed with high accuracy. Moreover, the advanced form of the base functions makes the QP-fitting very fast, requiring only a small number of library components.  In this manner, we have constructed a family of dynamical models that incorporate the observed features in $N$-body simulations. We summarize their properties: \\begin{enumerate} \\item{The models a priori have the universal properties encountered in $N$-body studies of dark matter halos, which formed the building bricks of the Jeans models by \\citet{2005MNRAS.363.1057D}: they generate a universal density profile, a power-law pseudo phase-space density $\\pseudo_r(r)$, four-parameter velocity anisotropy profiles with arbitrary values of $\\beta_0$ and $\\beta_\\infty$, and a linear $\\beta-\\gamma$ relation. In particular, we analyzed six models that are isotropic in the center and radially anisotropic at large radii.} \\item{The DFs are physical, i.e.\\ they are non-negative over the entire phase space. This means that the Dehnen-McLaughlin Jeans models can actually be realized by self-consistent dynamical systems. In addition, the DFs are continuous and smooth functions and, contrary to the popular Osipkov-Merritt models, they fill the entire accessible phase space.} \\item{The energy distributions are monotonously decreasing functions, and like the radial kurtosis profiles, these functions are nearly independent of $\\beta_\\infty$. This suggests that $\\rho(r)$, $\\pseudo_r(r)$, $\\kappa_r(r)$, and $N(\\calE)$ have a universal form, caused by the same physical processes.} \\end{enumerate} In addition, we have also been able to generate dynamical models with a non-linear $\\beta-\\gamma$ relation, i.e.\\ $\\delta \\neq \\eta/2$ and $r_\\txa \\neq r_\\txs$, albeit with higher $\\chi_N^2$ values.  From these successful results for the Dehnen-McLaughlin halos, we expect equally adequate fits for other Zhao models, such as the NFW halos. Other profiles, such as the S\\'ersic-type densities, might require a modified family of base functions.  Such a systematic study could unravel more hitherto hidden properties of dark matter halos, and the various connections between these characteristics.  Our set of dynamical models is not only useful for a purely theoretical analysis. The DFs can also serve to generate initial conditions with Monte Carlo simulators \\citep[e.g.][]{2004ApJ...601...37K,2007MNRAS.375.1157B} to investigate the physical processes within these equilibrium models by means of controlled numerical $N$-body simulations. Finally, our algorithm and base functions can be used in other dynamical studies, using different data moments than the density. This could lead to a significant improvement in the dynamical modeling of observed stellar systems, such as clusters of galaxies.  \\appendix"
    },
    "0809/0809.2303_arXiv.txt": {
        "abstract": "The physical processes involved in the advective-acoustic instability are investigated with 2D numerical simulations. Simple toy models, developed in a companion paper, are used to describe the coupling between acoustic and entropy/vorticity waves, produced either by a stationary shock or by the deceleration of the flow. Using two Eulerian codes based on different second order upwind schemes, we confirm the results of the perturbative analysis. The numerical convergence with respect to the computation mesh size is studied with 1D simulations. We demonstrate that the numerical accuracy of the quantities that depend on the physics of the shock is limited to a linear convergence. We argue that this property is likely to be true for most current numerical schemes dealing with SASI in the core-collapse problem, and could be solved by the use of advanced techniques for the numerical treatment of the shock. We propose a strategy to choose the mesh size for an accurate treatment of the advective-acoustic coupling in future numerical simulations. ",
        "introduction": "Most of our knowledge about the possible consequences of SASI on the core-collapse problem has been built, over the last 5 years, on the results of multidimensional numerical simulations \\citep[e.g.][]{Blondin+03,Scheck+04,Burrows+06,BM07,MJ07,Iwakami+08}. Whether or not SASI can contribute to overcome the explosion threshold, to kick the neutron star and alter its spin is still debated. In addition to the fundamental uncertainties associated with the equation of state of dense matter or the numerical treatment  of neutrino transport, some difficulties are simply related to multidimensional hydrodynamics \\citep{Blondin+03, Ohnishi+06, BM06, BM07, Iwakami+08}. This latter difficulty is partly due to the complexity of the mechanism underlying SASI, which is at best unfamiliar, and possibly also affected by the different numerical techniques used by different groups. The present study aims at improving our understanding of the instability mechanism at work by studying the advective-acoustic instability in the highly simplified set up introduced in the first paper of this series \\citep[][~hereafter paper I]{F08}. We note that a debate exists about the nature of this mechanism, as witnessed by \\citet[][~hereafter BM06]{BM06},~\\citet[][~hereafter FGS07]{Foglizzo+07}, \\citet{Laming07}, \\citet{YF08} and \\citet{Laming08}. Thus we believe that a better understanding of the advective-acoustic instability in simple examples can help recognise it in more complex situations. The separation of the advective-acoustic cycle into two separate problems is necessary in order to identify, between advected and acoustic perturbations, the consequences of each one on the other, as seen on Fig.~7 of \\citet{Blondin+03} or Figs.~11-12 of \\citet{Scheck+08}. In paper I, the following questions were answered through a perturbative analysis: \\par (i) what are the amplitudes of the entropy and vorticity waves generated by a shock perturbed by an acoustic wave propagating against the flow, towards the shock? \\par (ii) what is the amplitude of the acoustic wave generated by the deceleration of an entropy/vorticity wave through a localised gravitational potential?  The first purpose of our study is thus to check the results of the perturbative analysis presented in paper I through numerical experiments, thus providing concrete examples of the coupling processes involved.  The second purpose of this study is to gain confidence in the results of more elaborate numerical simulations by assessing their accuracy using our simple set up. The 2D numerical simulations of BM06 showed some globally good agreement with the perturbative analysis of FGSJ07. The typical error on the growth rate and the oscillation frequency of SASI, around $30\\%$, was not small though. Could this be a concern for the many other simulations which use a coarser mesh size? We wish to evaluate quantitatively, using our simple toy model, to what extent the advective-acoustic instability can be affected by numerical resolution.  The paper is organised as follows. In Sect.~2, the set up of the simulations is described and the numerical codes are presented. Sect.~3 illustrates qualitatively the two coupling processes involved in the advective-acoustic instability using 2D simulations. A quantitative analysis of these simulations is also performed which validates both the perturbative analysis and the numerical technique. In Sect.~4, we evaluate the rate of numerical convergence with respect to the mesh size, using a series of 1D numerical simulations. While the accuracy of the acoustic feedback produced by the flow gradients is quadratic with respect to the mesh size, the accuracy of the entropy wave produced by the shock depends on the mesh size only linearly. The linear phase of the full problem is simulated in Sect.~5, where the oscillation frequency and growth rate are compared to the results of the perturbative analysis. The consequences of these numerical difficulties for the simulations of core-collapse supernovae are discussed in Sect.~6. ",
        "conclusions": "\\label{Sec5}  \\begin{itemize} \\item A toy model has been used to illustrate through numerical experiments the coupling processes described in mathematical terms in paper I. Despite the high degree of simplification of our toy model, in particular the adiabatic hypothesis and the very local character of the deceleration region, these simulations can help us build our intuition about the physics of the advective-acoustic instability and better recognise it when present in numerical simulations. \\item The results of the perturbative approach have been confirmed quantitatively by our numerical simulations. \\item We have studied the effect of the mesh size on the accuracy of the numerical calculation. This will prove useful in the future to improve the reliability of the hydrodynamical part of simulations involving SASI in the core-collapse problem. We have proposed a conservative estimate of the desired mesh size close to the neutron star, which guarantees that the dominant acoustic feedback from advected perturbations is correctly taken into account. \\item The difficulties associated with the numerical treatment of the shock have direct consequences on the accuracy with which the flow resulting from SASI is calculated: without a special numerical effort, the convergence of the computation of the growth time and oscillation frequency of SASI is reduced to first order even if the numerical scheme converges with a higher order away from the shock. Among the published simulations of SASI, only the 2D simulations with the finest grid seem to be able to estimate the entropy and vorticity production at the shock with a $<10\\%$ accuracy. The importance of an accurate treatment of SASI in the core-collapse problem may make it worth implementing advanced techniques for the numerical treatment of the shock in future simulations, such as the level set method for example \\citep{SS03}.  \\end{itemize}"
    },
    "0809/0809.2629_arXiv.txt": {
        "abstract": "{ A quasar wind model is proposed to describe the spatial and velocity structure of the broad line region. This model requires detailed photoionization and magnetohydrodynamic simulation, as the broad line region it too small for direct spatial resolution. The emission lines are Doppler broadened, since the gas is moving at high velocity. The high velocity is attained by the gas from a combination of radiative and magnetic driving forces. Once this model is complete, the model predictions will be tested against recent microlensing data in conjunction with diverse existing observations. ",
        "introduction": "Broad emission lines are a prominent spectral features of quasars. The physical state and geometry of this region is poorly understood since the broad line region is not directly resolvable.  The broadening of these emission lines is so extreme, that speeds of up to 0.1$c$ are observed. Temperature and number density limits can be inferred from observations of the relative strengths of broad emission lines. The temperature inferred from observations is too low for the broadening process to be thermal, therefore the gas producing these emission lines is moving at high speed relative to the quasar rest frame.  This model proposes that the gas is accelerated outward by photon and magnetohydrodynamic forces. % Direct spatial resolution of this region is not possible. In order to determine the geometry, origin and nature of this gas, simulation is required. ",
        "conclusions": ""
    },
    "0809/0809.3200.txt": {
        "abstract": "\\begin{center} {\\bf ABSTRACT} \\end{center} Many string constructions have a classical no-scale structure, resulting in a one-parameter model (OPM) for the supersymmetry breaking soft terms. As a highly constrained subset of mSUGRA, the OPM has the potential to be predictive. Conversely, if the observed superpartner spectrum at LHC is a subset of the OPM parameter space, then this may provide a clue to the underlying theory at high energies.  We investigate the allowed supersymmetry parameter space for a generic one-parameter  model taking into account the most recent experimental constraints.  We find that in the strict moduli scenario, there are no regions of the parameter space which may satisfy all constraints. However, for the dilaton scenario, there are small regions of the parameter space where all constraints may be satisfied and for which the observed dark matter density may be generated.  We also survey the possible signatures which may be observable at the Large Hadron Collider (LHC).  Finally, we compare collider signatures of OPM to those from a model with non-universal soft terms, in particular those of an intersecting $D6$-brane model. We find that it may be possible to distinguish between these diverse scenarios at LHC. ",
        "introduction": "With the dawn of the Large Hadron Collider (LHC) era, the prospects for the discovery of new physics may finally be arriving. In particular, whatever physics is responsible for stabilizing the electroweak scale should be discovered.  Signals of the favored mechanism, broken supersymmetry, may be observed as well as the Higgs states required to break the electroweak symmetry. However, at present there is no theory which may uniquely predict the masses of the superpartners should they be observed at LHC   In principle, it should be possible to derive all known physics in a top-down approach directly from a more fundamental theory such as string theory, as well as potentially predicting new and unexpected phenomena. Conversely, following a bottom-up approach, one may ask if it is possible to deduce the origin of new physics given such a signal at LHC. For example, in the case of low-energy supersymmetry, it may be possible from the experimental data to deduce the structure of the fundamental theory at high energy scales which determines the soft-supersymmetry breaking terms and ultimately leads to radiative electroweak symmetry breaking (REWSB)~\\cite{Ellis:1982wr, Ellis:1983bp, AlvarezGaume:1983gj}.  No-scale supergravity (nSUGRA)~\\cite{DVNS} is such a framework where it is possible to naturally explain REWSB and correlate it with the gravitino mass, or more generally, the effective SUSY breaking scale. In the simplest no-scale models, the gravitino mass $m_{3/2}$ remains undetermined at the classical level, and is instead fixed by radiative corrections to be near the electroweak scale~\\cite{DVNS}. Thus, in this framework, we find that the scale of supersymmetry breaking is correlated with the electroweak scale~\\cite{DVNS}. Another striking feature of nSUGRA is that the cosmological constant vanishes at tree-level. Although it is presently known that the cosmological constant is in fact non-zero, its very small value is still consistent with the no-scale framework with small corrections. Furthermore, it is well-known that the K\\a\"ahler moduli of Type I, IIB orientifold, and heterotic string compactifications have a classical no-scale structure~\\cite{JLDN, VKJL, ABIM}.  It has been shown that in Type IIB orientifold compactifications, this type of supersymmetry breaking corresponds to turning on RR and NS fluxes, which are generically present in order to cancel tadpoles as well as to stabilize closed-string moduli.  Indeed, this is the case in the so-called large volume models~\\cite{BBCQ}~\\cite{JCQS}. This combined with the generic appearance of the no-scale structure across many string compactifications leads to the idea that supersymmetry breaking is moduli dominated.   In string models, supersymmetry breaking is typically performed in a hidden sector as well as through the universal moduli and dilaton fields.  For a given string compactification, the precise nature of the supersymmetry breaking is determined by model-dependent calculations.  However, at present there are no specific string compactifications which completely satisfy all theoretical criteria which are desired in such a model.  Thus, a model-independent approach is perhaps wiser at the present time.  The most studied model of supersymmetry breaking is minimal supergravity (mSUGRA), which arises from adopting the simplest ansatz for the K\\a\"ahler metric, treating all chiral superfields symmetrically.  In this framework,~~$\\mathcal{N}=1$ supergravity is broken in a hidden sector which is communicated to the observable sector through gravitational interactions.  Such models are characterized by the following parameters: a universal scalar mass $m_0$, a universal gaugino mass $m_{1/2}$, the Higgsino mixing $\\mu$-parameter, the Higgs bilinear $B$-parameter, a universal trilinear coupling $A_0$, and tan~$\\beta$.  One then determines the $B$ and $|\\mu|$ parameters by the minimization of the Higgs potential triggering REWSB, with the sign of $\\mu$ remaining undetermined.  Thus, we are left with only four parameters.  Although, this is one of the most generic frameworks that can be adopted, and many string compactifications typically yield expressions for the soft terms which are even more constrained due to the no-scale structure which emerges naturally in these theories assuming that supersymmetry breaking is dominated by the K\\a\"ahler moduli and/or dilaton.  In particular, in such nSUGRA models, we will generically have $m_{0} = m_{0}(m_{1/2})$ and $A = A(m_{1/2})$. This reduces the number of free parameters compared to mSUGRA down to two, $m_{1/2}$ and tan$\\beta$.  In fact, adopting a strict no-scale framework, one can also fix the $B$-parameter as $B=B(m_{1/2})$, and thus we are led to a {\\it one-parameter} model where all of the soft terms may be fixed in terms of $m_{1/2}$.   If we assume that the supersymmetry breaking is triggered by some of the moduli fields in a given string compactification, namely the dilaton $S$ and the three K\\a\"ahler moduli $T$ which obtain VEVs $\\left\\langle F_S \\right\\rangle$ and $\\left\\langle F_T \\right\\rangle$ respectively, a simple expression for the scalar masses may be adopted: \\begin{equation} \\tilde{m}^2_i = m^2_{3/2} (1 + n_i \\mbox{cos}^2 \\theta), \\end{equation} with $\\mbox{tan}\\theta = \\left\\langle F_S \\right\\rangle / \\left\\langle F_T\\right\\rangle$ and where $m_{3/2}$ is the gravitino mass and $n_i$ are the modular weights of the respective matter fields.  In order to obtain universal scalar masses, which are highly suggested by the required absence of FCNC~\\cite{Ellis:1981ts}, there are two possible cases which may be considered:  (i)  setting $\\theta = \\pi/2$ so that $\\left\\langle F_S \\right\\rangle >> \\left\\langle F_T\\right\\rangle$; or (ii) setting all $n_i$ to be the same ($n_i=-1$) and $\\theta=0$ so that $\\left\\langle F_T \\right\\rangle >> \\left\\langle F_S\\right\\rangle$ so that all scalar masses vanish at the unification scale.  The first of these two cases is referred to as the special {\\it dilaton} scenario, \\begin{equation} \\label{eqn:dilaton} m_{0} = \\frac{1}{\\sqrt{3}}m_{1/2}, \\ \\ \\ \\ \\ A = -m_{1/2}, \\ \\ \\ \\ \\ B = \\frac{2}{\\sqrt{3}}m_{1/2}. \\end{equation} while the second is referred to as the strict {\\it moduli} scenario, \\begin{equation} \\label{eqn:moduli} m_{0} = 0, \\ \\ \\ \\ \\ A = 0, \\ \\ \\ \\ \\ B = 0. \\end{equation}  For many string compactifications, especially those within the free-fermionic class of models in particular those with a flipped $SU(5)$ gauge group~\\cite{AEHN}, the soft-terms will have such a form.  Interestingly, soft terms for heterotic M-theory compactifications with moduli dominated supersymmetry breaking take the form~\\cite{TJLI} \\begin{eqnarray} &m_{1/2} = \\frac{x}{1+x}m_{3/2} \\\\ \\nonumber &m_{0} = \\frac{x}{3+x}m_{3/2} \\\\ \\nonumber &A = - \\frac{3x}{3+x}m_{3/2} \\end{eqnarray}  \\noindent while for dilaton dominated supersymmetry breaking they take the form \\begin{eqnarray} &m_{1/2} = \\frac{\\sqrt{3}m_{3/2}}{1+x} \\\\ \\nonumber &m_{0}^{2} = m_{3/2}^{2} - \\frac{3m_{3/2}^{2}}{(3+x)^{2}}x(6+x) \\\\ \\nonumber &A = - \\frac{\\sqrt{3}m_{3/2}}{3+x}(3-2x) \\end{eqnarray}  \\noindent which reduce to the above moduli and dilaton scenarios in the limit $x \\rightarrow 0$, where \\begin{equation} x \\propto \\frac{(T+\\overline{T})}{S+\\overline{S}} \\end{equation}  In addition, the so-called large-volume models have been studied extensively~\\cite{BBCQ}~\\cite{JCQS} in recent years and the generic soft terms for this framework have been calculated in~\\cite{CAQS}. These models involve Type IIB compactifications where the moduli are stabilized by fluxes and quantum corrections to the K\\a\"ahler potential generate an exponentially large volume. This exponentially large volume may lower the string scale to an intermediate level which can be in the range $m_{s} \\sim 10^{3-15}$ GeV. In such models, the soft terms can take the form \\begin{eqnarray} \\centering &m_0 = \\frac{1}{\\sqrt{3}}M \\nonumber \\\\ &A_0 = -M \\nonumber \\\\ &B = - \\frac{4}{3}M \\end{eqnarray} where $M$ is a universal gaugino mass, which are essentially identical to the special dilaton scenario given above. However, this framework is different from our analysis in that the string scale may be lower than what is usually taken for the grand unification scale $\\approx 2.1 \\times 10^{16}$~GeV.  Thus, in this scenario the observed unification of the gauge couplings when extrapolated to high energies is merely coincidental.  Of course, this then also affects the running of the soft-masses resulting in different superpartner spectra than what would otherwise be obtained.  More generally, no-scale moduli-dominant scenarios of SUSY breaking are favored by F-theory~\\cite{Aparicio:2008wh}, as are models with a flipped SU(5) gauge group~\\cite{Beasley:2008kw}.  In this work, we identify the regions of the supersymmetry parameter space for a generic one-parameter model which are allowed by current experimental constraints and survey the signatures which may be observable at LHC. We find that in the strict moduli scenario, there are no regions of the parameter space which may satisfy all constraints. However, for the dilaton scenario, there are small regions of the parameter space where all constraints may be satisfied and for which the observed dark matter density may be generated.  The model is thus a highly constrained subset of mSUGRA, which allows the model to potentially be predictive. Conversely, if the superpartner spectrum actually observed at LHC lies within the OPM parameter space, then this may provide a strong clue to the underlying theory at high energy scales.   Finally, we simulate the different LHC signatures for this model and compare to those for an intersecting $D6$-brane model which possesses many desirable phenomenological characteristics~\\cite{RIBM}. We find certain signatures may indicate there are distinguishing phenomenological characteristics between these different types of constructions. ",
        "conclusions": "We have updated and surveyed the allowed parameter space of the one-parameter model. Our motivation for studying this model stems from the commonality of the universal soft supersymmetry breaking ansatz across multiple types of string compactifications. These include weak coupled and heterotic M-theory vacua, as well as Type IIB flux vacua, in particular the so-called large-volume compactification models.  By performing a comprehensive scan of the entire parameter space and filtering the results according to experimental constraints, the allowed parameter space was obtained. In the strict moduli dominant case, we found that there is no parameter space which can satisfy all of the constraints, whereas there is a small parameter space allowed for the dilaton scenario.  We identified the probable squark and gluino interactions and presented cascade decay modes that will produce specific favorable collider signatures. The dominant component of these favorable signatures are tau and hadronic jets. In future work, we plan to study the observable signatures at LHC for this model in somewhat more depth.  We compared the collider signatures of the one-parameter model to a model with non-universal soft terms, in particular an intersecting $D6$-brane model with interesting phenomenological properties. We found that for one particular intersecting $D$6-brane pattern of mass hierarchies, there are possible distinguishing characteristics between these two classes of models. Although there may also be a lot of overlap in the observable signatures of these two models, there are regions of the parameter space of each class which may give strikingly different observable signatures which may be used to distinguish them. Thus, it may be possible for LHC to say something about the structure of the underlying theory at high-energies, e.g. universality vs. non-universality.   We plan to investigate this possibility more deeply in an upcoming paper."
    },
    "0809/0809.3133_arXiv.txt": {
        "abstract": "{} {Spectroscopic and photometric properties of low and high-z supernovae Ia (SNe Ia) have been analyzed in order to achieve a better  understanding of their diversity and to identify possible SN Ia sub-types.} {We use wavelet transformed spectra in which one can easily measure spectral features. We investigate the \\ion{Si}{II} 4000 equivalent width ($EW_w\\lbrace\\ion{Si}{II}\\rbrace$). The ability and, especially, the ease in extending the method to SNe at high-$z$ is demonstrated.} {We applied the method to 110 SNe~Ia and found  correlations between $EW_w\\lbrace\\ion{Si}{II}\\rbrace$  and parameters related to the light-curve shape for 88 supernovae with available photometry. No evidence for evolution of $EW_w\\lbrace\\ion{Si}{II}\\rbrace$ with redshift is seen. Three sub-classes of SNe~Ia were confirmed using an independent cluster analysis with only light-curve shape, colour, and $EW_w\\lbrace\\ion{Si}{II}\\rbrace$.} {SNe from high-$z$ samples seem to follow a similar grouping to nearby objects. The $EW_w\\lbrace\\ion{Si}{II}\\rbrace$ value measured on a single spectrum may point towards  SN~Ia sub-classification, avoiding the need for expansion velocity gradient calculations.} ",
        "introduction": "The peak luminosities of type Ia supernovae (SNe Ia) are one of the best distance indicators at high redshifts. Consequently, due to the relatively small dispersion in their light-curves, they are used for cosmological parameter estimation (e.g., ~\\citealt{riess98,perlmutter99}). However, some SNe seem to depart from standard behaviour and deserve further attention. The most representative examples are the sub-luminous 1991bg-like objects, the unusual SNe 2002ic and 2005gj\\footnote{Some authors argue that SN 2002ic may be a type Ic supernova instead of Ia (e.g. \\citealt{benetti06,wang04}), and the SN 2005gj classification as type Ia/IIn is still doubtful (\\citealt{prieto07,trundle08}).},  and the even more outstanding SN 2003fg (\\citealt{howell06}). The existence of a single family of SNe~Ia is thus still debatable (see \\citealt{filippenko97} for a review on SN (in)homogeneity). Hence separating sub-types of SNe~Ia or attempting a continuous parametrization of spectra and/or light-curves including the full diversity of SNe~Ia observed up to now could reduce the scatter in the Hubble diagram and improve their use in cosmology.   In order to quantify the spectral differences between SNe Ia,~ \\citet{nugent95} proposed a ratio between the depths of absorption features of \\ion{Si}{II} at 5972$\\textrm{\\AA}$ and 6355$\\textrm{\\AA}$. This ratio was also found to correlate with the absolute magnitude of SNe~Ia and the light-curve shape parameter. We further investigate this idea exploring a consistent method of measuring SN features on transformed spectra using wavelets as will be explained below. One of the major concerns regarding the use of SNe~Ia in cosmology is a possible systematic difference between the low-$z$ and high-$z$ samples (\\citealt{blondin06,garavini07a,bronder08}). Reasons for such evolution are generally thought to be related to changes in metallicity or composition of the progenitor or circumstellar medium, progenitor mass or delay times. It is already known that progenitor age is a relevant issue for variability in SN peak luminosity (\\citealt{prieto08,gallagher05}) and that metallicity might evolve with redshift (see recent work of ~\\citealt{ellis08} and reference therein). We shall demonstrate that the equivalent width of the \\ion{Si}{II} 4000 feature, $EW_w\\lbrace\\ion{Si}{II}\\rbrace$ defined below, coupled with light-curve parameters, is an indicator of possible sub-classes within SNe~Ia, consistent for both low and high-$z$ SNe. ",
        "conclusions": "The use of SNe~Ia for cosmology relies on empirical calibration techniques on the light-curves and K-correction. It would be reassuring to have a usable indicator measured directly from SN spectra that can be used as an independent calibrator of the luminosity of an SN Ia event. Applying a straightforward transformation using wavelets, we were able to estimate, in a consistent manner on a fair number of SNe~Ia, the equivalent width of the \\ion{Si}{II} 4000 feature which previously showed potential use for cosmology. At low redshift, we were able to automatically distinguish three classes of SNe~Ia previously found by other authors using much more information, such as the expansion velocity gradient.   The same wavelet-based  approach was applied to high redshift data, revealing an analogous clustering tendency to that found for low redshift supernovae. Yet it is not clear whether the three sub-classes, found in the low-$z$ and verified  in the high-$z$ sample, are completely distinct or come from one continuous family. If so,  this  could possibly be the remnant of an extra  parameter or a more complex modelling of the supernova data. The implementation of our  method using recently available much larger spectral samples and  the use of different SN Ia sub-classes in cosmological analysis is now in preparation.   Although the method presented here is promising and we are planning to apply it to other spectral features,  we remark that line ratios, pseudo-equivalent widths and their current derivatives are not optimal to extract information on lower signal-to-noise spectra. Parameters calculated from spectral features, such as equivalent widths, still consider only local information and each feature independently. New indicators based on combining local information, such as wavelet coefficients from the whole spectra, can provide a better characterisation of the supernovae."
    },
    "0809/0809.0866.txt": {
        "abstract": "Spin parameters of stellar-mass black holes in X-ray binaries are currently being estimated by fitting the X-ray continuum spectra of their accretion disk emission.  For this method, it is necessary to know the inclination of the X-ray-producing inner region of the disk.  Since the inner disk is expected to be oriented perpendicular to the spin axis of the hole, the usual practice is to assume that the black hole spin is aligned with the orbital angular momentum vector of the binary, and to estimate the inclination of the latter from ellipsoidal modulations in the light curve of the secondary star.  We show that the inclination of the disk can be inferred directly if we have both spectral and polarization information on the disk radiation.  The predicted degree of polarization varies from 0\\% to 5\\% as the disk inclination changes from face-on to edge-on.  With current X-ray polarimetric techniques the polarization degree of a typical bright X-ray binary could be measured to an accuracy of 0.1\\% by observing the source for about 10 days.  Such a measurement would constrain the disk inclination to within a degree or two and would significantly improve the reliability of black hole spin estimates.  In addition, it would provide new information on the tilt between the black hole spin axis and the orbital rotation axis of the binary, which would constrain any velocity kicks experienced by stellar-mass black holes during their formation. ",
        "introduction": "X-ray polarimetry is an uncharted frontier with great promise for the exploration of the behavior of accreting black holes and for the measurement of black hole spin.  To date, only small X-ray polarimeters have been flown and the only result of consequence has been the measurement of the polarization of the Crab Nebula (Novick et al.\\ 1972; Weisskopf et al.\\ 1978).  Anticipating the launch of larger and more sensitive X-ray polarimeters in the future, this paper discusses possible uses of polarization observations to study thermal emission from accretion disks around stellar-mass black holes in X-ray binaries.  Polarization measurements at radio and optical wavelengths attest to the power of polarimetry.  In radio astronomy, the measurement of polarization is trivial.  However, its consequences have been of central importance since Alf\\'ven \\& Herlofson (1950) first suggested the synchrotron mechanism as the source of celestial radio emission.  For example, studies of this strongly polarized radiation has allowed the mapping of magnetic fields that are threaded throughout stellar coronae, supernova shocks, relativistic jets, and the interstellar- and intergalactic media.  At optical wavelengths, the measurement of polarization is more difficult, but the payoffs have also been large. For example, polarimetric observations of NGC~1068 provided a spectacular confirmation of the unified model of AGN (Antonucci \\& Miller 1985).  Polarization measurements are most challenging at X-ray wavelengths; however, sure instrumental approaches are known (\\S6), and the discovery space and potential rewards are large.  X-ray polarimetric information -- direction and degree -- increases the parameter space used to investigate compact objects from the current two -- spectra and time variability -- to four independent parameters that models must satisfy (Rees 1975; Lightman \\& Shapiro 1975; M\\'esz\\'aros et al.\\ 1988).  X-ray polarimetry can provide qualitatively new information on source geometry and magnetic fields on spatial scales comparable to a black hole event horizon.  In this paper, we follow the pioneering work of \\citet{sta77}, \\citet{con77}, and \\cite{con80}, and we narrow our focus to consider solely the problem of polarized thermal emission generated by a thin accretion disk around a stellar-mass black hole.  The theory developed here of course applies as well to thermal accretion disks around supermassive black holes.  We do not consider coronal effects that can generate an X-ray reflection spectrum \\citep[e.g.,][]{ros99,dov04} and Comptonized components of emission \\citep[e.g.,][]{gie99}.  The polarization model presented here extends our workhorse accretion disk model {\\sc kerrbb} (Li et al.\\ 2005), which we have used to estimate the spins of four stellar-mass black holes via fits to their thermal continuum spectra (Shafee et al.\\ 2006; McClintock et al.\\ 2006; Liu et al.\\ 2008).  The extension of {\\sc kerrbb} presented herein includes a simple electron scattering atmosphere based on an $\\alpha$-viscosity prescription of the stress (Shakura \\& Sunyaev 1973) and a treatment of the polarized radiation field via the Stokes parameters.  Radiative transfer effects will be considered in a future paper.  In \\S\\ref{motivation}, we discuss our motivations for this work, focusing in particular on the application of polarization measurements for estimating the spins of accreting black holes.  In \\S\\ref{polar}, we briefly review the polarization of disk emission induced by electron scattering in an accretion disk atmosphere, assuming a semi-infinite plane-parallel medium. In \\S\\ref{polar_remote}, we present the mathematical scheme for calculating the polarization as observed by a remote observer, after the polarized disk emission propagates through the curved spacetime of a Kerr black hole. In \\S\\ref{application}, we apply our polarization calculations to the measurement of black hole spin, and in \\S\\ref{observation}, we briefly discuss the prospects for observing the predicted polarization signal from accreting stellar-mass black holes. Finally, in \\S\\ref{conclusion}, we summarize the results.  Some mathematical details related to the thin disk model we employ and the propogation of the polarization vector in Kerr spacetime are presented in Appendices \\ref{disk_model} and \\ref{p_prop}. ",
        "conclusions": "\\label{conclusion}  In order to secure the measurement of black hole spin using the continuum-fitting method, it is essential to obtain an independent determination of the inclination of the inner accretion disk $i_{\\rm disk}$.  We have shown that this appears to be entirely feasible via X-ray polarimetry.  Our models predict polarizations of up to $\\sim 5$\\% that vary monotonically with inclination, while sensitivities of $\\sim 0.1$\\% are achievable with quite modest SMEX-class polarimeters.  Such instruments are expected to be capable of routinely making measurements of $i_{\\rm disk}$ over a wide range of inclinations with a precision of $\\sim 1^{\\circ}$.  Meanwhile, with full attention it will be possible to achieve measurements of similar quality for the orbital inclination angle $i_{\\rm orb}$, given recent advances in optical/NIR instrumentation (e.g., Vernet et al.\\ 2007) and adaptive optics (e.g., van Dam et al.\\ 2006).  Once reliable values of these two inclination angles have been obtained for a sample of black hole binaries, the question will simply be: is $i_{\\rm disk} \\approx i_{\\rm orb}$, or is it not?  There are of course significant hurdles that must be cleared.  For example, the model of the disk atmosphere presented here includes only an approximate treatment of the effects of absorption.  In the manner of Davis \\& Hubeny (2006), one must model the non-LTE effects, Compton scattering, and the opacities due to ions of the abundant elements. Magnetohydrodynamic models of thin disks in general relativity, which include radiation and polarization, are the ultimate theoretical goal, and progress is now being made on this front \\citep{bec08,nob08,rey08,sha08b}. Then, there are observational complications to consider.  For example, even in the thermal dominant state there exists some remnant scattered coronal emission, which must be modeled.  Some other examples of possible nettlesome sources of polarization include self-irradiation and reflection in the disk, scattering by interstellar grains, and the effects of a global component of magnetic field in the disk.  Clearly, the greatest hurdle is getting an X-ray polarimeter into space.  Thirty years has passed since the tiny {\\it OSO-8} crystal instrument successfully measured the polarization of the Crab Nebula at 2.6 keV to be $19.2 \\pm 1.0$\\% at a position angle of $156\\fdg4 \\pm 1\\fdg4$ (Weisskopf et al.\\ 1978).  It is surely now time to open the polarimetric channel for serious exploration and discovery.  In this paper we have shown how a polarimeter can secure the measurement of black hole spin.  This is just one exciting application of X-ray polarimetry and there are numerous others.  In the study of black holes, polarimetry furthermore promises to define the geometry of the emitting elements and constrain source models in decisive ways (\\S1).  This is true, for instance, for modeling the relativistic iron line (Miller 2007), or relativistic jets (Mirabel \\& Rodr\\'iguez 1999), or the mysterious power-law component that extends unbroken to MeV energies (Grove et al.\\ 1998), or the global oscillations that sometimes modulate up to one quarter of the total accretion power (McClintock \\& Remillard 2006).  In modeling these and other black-hole phenomena, a central question is the geometry of the Comptonizing coronal source, which is vaguely and variously described as a sphere or a slab or a lamp post.  Polarimetry will provide the best, and often only, clue to the actual geometry of the corona and other key structures (M\\'esz\\'aros et al.\\ 1988; Blandford et al.\\ 2002)."
    },
    "0809/0809.1191_arXiv.txt": {
        "abstract": "Just as a rotating magnetised neutron star has material pulled away from its surface to populate a magnetosphere, a similar process can occur as a result of neutron-star pulsations rather than rotation. This is of interest in connection with the overall study of neutron star oscillation modes but with a particular focus on the situation for magnetars. Following a previous Newtonian analysis of the production of a force-free magnetosphere in this way \\citet{timokhin_00}, we present here a corresponding general-relativistic analysis. We give a derivation of the general relativistic Maxwell equations for small-amplitude arbitrary oscillations of a non-rotating neutron star with a generic magnetic field and show that these can be solved analytically under the assumption of low current density in the magnetosphere. We apply our formalism to toroidal oscillations of a neutron star with a dipole magnetic field and find that the low current density approximation is valid for at least half of the oscillation modes, similarly to the Newtonian case. Using an improved formula for the determination of the last closed field line, we calculate the energy losses resulting from toroidal stellar oscillations for all of the modes for which the size of the polar cap is small. We find that general relativistic effects lead to shrinking of the size of the polar cap and an increase in the energy density of the outflowing plasma. These effects act in opposite directions but the net result is that the energy loss from the neutron star is significantly smaller than suggested by the Newtonian treatment. ",
        "introduction": "\\label{introd} \\noindent   Study of the internal structure of neutron stars (NSs) is of fundamental importance for subatomic physics since these objects provide a laboratory for studying the properties of high-density matter under very extreme conditions. In particular, there is the intriguing possibility of using NS oscillation modes as a probe for constraining models of the equation of state of matter at supranuclear densities. It was suggested long ago that if a NS is oscillating, then traces of this might be revealed in the radiation which it emits \\citep{pacini_74, tsygan_75, boriakoff_76, bisnovatyi_95, ding_97, duncan_98}. Recently, a lot of interest has been focussed on oscillations of magnetized NSs because of the discovery of gamma-ray flare activity in Soft Gamma-Ray Repeaters (SGRs) which are thought to be the very highly magnetised NSs known as magnetars \\citep[for recent review on the SGRs see][]{woods_06, watts_07}. The giant flares in these objects are thought to be powered by global reconfigurations of the magnetic field and it has been suggested that the giant flares might trigger starquakes and excite global seismic pulsations of the magnetar crust \\citep{thompson_95, thompson_01, schwartz_05, duncan_98}. Indeed, analyses of the observations of giant flares have revealed that the decaying part of the spectrum exhibits a number of quasi-periodic oscillations (QPOs) with frequencies in the range from a few tens of Hz up to a few hundred Hz \\citep{israel_05, strohmayer_06, watts_06} and there has been a considerable amount of theoretical effort attempting to identify these with crustal oscillation modes \\citep{glampedakis_06, samuelsson_07, levin_07, sotani_07_a, sotani_07_b}. While there is substantial evidence that the observed SGR QPOs are caused by neutron star pulsations, there is a great deal of uncertainty about how stellar surface motion gets translated into the observed features of the X-ray radiation \\citep{strohmayer_08, strohmayer_06, timokhin_07b}. To make progress with this, it is necessary to develop a better understanding of the processes occurring in the magnetospheres of oscillating neutron stars.  Standard pulsars typically have magnetic fields of around $10^{12}$ G while magnetars may have fields of up to $10^{14}-10^{15}$ G near to the surface. Rotation of a magnetized star generates an electric field: \\begin{equation} \\label{e_rot} E^\\mathrm{rot} \\sim \\frac{\\Omega R}{c}B \\ , \\end{equation} where  $ B $ is the magnetic field strength, $ c $ is the speed of light and $ \\Omega $ is the angular velocity of the star with radius $ R $. Depending on the rotation velocity and the magnetic field strength, the electric field may be as strong as $ 10^{10} $ V cm$^{-1}$ and it has a longitudinal component (parallel to $B$) which can be able to pull charged particles away from the stellar surface, if the work function is  sufficiently small, and accelerate them up to ultra-relativistic velocities. This result led \\citet{goldreich_69} to suggest that a rotating NS with a sufficiently strong magnetic field should be surrounded by a magnetosphere filled with charge-separated plasma which screens the accelerating electric field and thus hinders further outflow of charged particles from the stellar surface. Even if the binding energy of the charged particles is sufficiently high to prevent them being pulled out by the electric field, the NS should nevertheless be surrounded by charged particles produced by plasma generation processes \\citep{sturrock_71,ruderman_75}, which again screen the longitudinal component of the electric field. These considerations led to the development of a model for pulsar magnetospheres which is frequently called the ``standard model'' \\citep[an in depth discussion and review of this can be found in, e.g.,][]{michel_91,beskin_93, beskin_05}.  Timokhin, Bisnovatyi-Kogan $\\&$ Spruit (2000) (referred to as TBS from here on) showed that an oscillating magnetized NS should also have a magnetosphere filled with charge-separated plasma, even if it is not rotating, since the vacuum electric field induced by the oscillations would have a large radial component which can be of the same order as rotationally-induced electric fields. One can show this quantitatively by means of the following simple arguments. To order of magnitude, the radial component of the vacuum electric field generated by the stellar oscillations is given by \\begin{equation} E^\\mathrm{osc}\\sim \\frac{\\omega \\xi}{c}B \\ , \\end{equation} where $\\omega$ is the oscillation frequency and $\\xi$ is the displacement amplitude. Using this together with Eq.(\\ref{e_rot}), it follows immediately that the electric field produced by oscillations will be stronger than the rotationally induced one for sufficiently slowly-rotating neutron stars, having \\begin{equation} \\Omega \\lesssim \\frac{\\omega \\xi }{ R } \\ . \\end{equation} For stellar oscillations with $\\xi / R \\sim 0.001 $ and $\\omega \\sim 1$ kHz, the threshold is $\\Omega \\sim 1 $ Hz. Within this context, TBS developed a formalism extending the basic aspects of the standard pulsar model to the situation for a non-rotating magnetized NS undergoing arbitrary oscillations. This formalism was based on the assumption of low current densities in the magnetosphere, signifying that the influence of currents outside the NS on electromagnetic processes occurring in the magnetosphere is negligibly small compared to that of currents in the stellar interior. This assumption leads to a great simplification of the Maxwell equations, which then can be solved analytically. As an application of the formalism, TBS considered toroidal oscillations of a NS with a dipole magnetic field, and obtained analytic expressions for the electromagnetic field and charge density in the magnetosphere. (Toroidal oscillations are thought to be particularly relevant for magnetar QPO phenomena.) They found that the low current density approximation (LCDA) is valid for at least half of all toroidal oscillation modes and analyzed the energy losses due to plasma outflow caused by these modes for cases where the size of the polar cap (the region on the stellar surface that is crossed by open magnetic field lines) is small, finding that the energy losses are strongly affected by the magnetospheric plasma. For oscillation amplitudes larger than a certain critical value, they found that energy losses due to plasma outflow were larger than those due to the emission of the electromagnetic waves (assuming in that case that the star was surrounded by vacuum). Recently, \\citet{timokhin_07} considered spheroidal oscillations of a NS with a dipole magnetic field, using the TBS formalism, and found that the LCDA again holds for at least half of these modes. Discussion in \\citet{timokhin_07} also provided some useful insights into the role of rotation for the magnetospheric structure of oscillating NSs.  The TBS model was a very important contribution and, to the best of our knowledge, remains the only model for the magnetosphere of oscillating NSs available in the literature. However, it should be pointed out that it does not include several ingredients that a fully consistent and realistic model ought to include. Most importantly, it does not treat the magnetospheric currents in a fully consistent way: although it gives a consistent solution for around half of the oscillation modes, the remaining solutions turn out to be unphysical and, as TBS pointed out, this is a symptom of the LCDA failing there. Also, rotation and the effects of general relativity can be very relevant; in particular, several authors have stressed that using a Newtonian approach may not give very good results for the structure of NS magnetospheres \\citep[see, e.g.,][]{beskin_90, muslimov_92, mofiz_00, morozova_08}. However, a more realistic model would naturally be more complicated than the TBS one whose relative simplicity can be seen as a positive advantage when using it as the basis for further applications.  The aim of the present paper is to give a general relativistic reworking of the TBS model so as to investigate the effects of the changes with respect to the Newtonian treatment. We derive the general relativistic Maxwell equations for arbitrary small-amplitude oscillations of a non-rotating spherical NS with a generic magnetic field configuration and show that they can be solved analytically within the LCDA as in Newtonian theory. We then apply this solution to the case of toroidal oscillations of a NS with a dipole magnetic field and find that the LCDA is again valid for at least half of all toroidal oscillation modes, as in Newtonian theory. Using an improved formula for the determination of the last closed field line, we calculate the energy losses resulting from these oscillations for \\textit{all} of the modes for which the size of the polar cap is small and discuss the influence of GR effects on the energy losses.  The paper is organized as follows. In Section 2 we introduce some definitions and derive the quasi-stationary Maxwell equations in Schwarzschild spacetime as well as the boundary conditions for the electromagnetic fields at the stellar surface. In Section 3 we sketch our method for analytically solving the Maxwell equations for arbitrary NS oscillations with a generic magnetic field configuration. In Section 4 we apply our formalism to the case of purely toroidal oscillations of a NS with a dipole magnetic field and also discuss the validity of the LCDA and the role of GR effects. In Section 5 we calculate the energy losses due to plasma outflow caused by the toroidal oscillations. Some detailed technical calculations related to the discussion in the main part of the paper are presented in Appendices A-C.  We use units for which $c = 1$, a space-like signature $(-,+,+,+)$ and a spherical coordinate system $(t,r,\\theta ,\\phi)$. Greek indices are taken to run from 0 to 3 while Latin indices run from 1 to 3 and we adopt the standard convention for summation over repeated indices. We indicate four-vectors with bold symbols ({\\it e.g.} ${\\boldsymbol u}$) and three-vectors with an arrow ({\\it e.g.} ${\\vec u}$). ",
        "conclusions": "In this paper, we have described our general relativistic model for the force-free magnetosphere of an oscillating, non-rotating neutron star. Our approach is based on the previous Newtonian model developed by TBS and focuses on toroidal modes which are thought to be particularly relevant for magnetar QPO phenomena. We have taken the spacetime geometry to be spherically symmetric and have neglected any modifications of it caused by the electromagnetic fields, the stellar oscillations and the magnetospheric plasma. Within this context, we have derived the general relativistic Maxwell equations for arbitrary small-amplitude oscillations of a neutron star with a generic magnetic field configuration and have shown that, as in the Newtonian case, they can be solved analytically for the force-free configuration of the electromagnetic fields (i.e. $E_{||} \\ll B $ ) under the low current density approximation (LCDA). We have applied our formalism to small-amplitude toroidal oscillations of a neutron star with a dipole magnetic field and have found that the LCDA is valid for at least half of these modes in GR, as in the Newtonian calculations of TBS. We have also discussed the contribution of the GR effects to our solution, finding that they lead to an increase in the absolute values of the electromagnetic fields and the space charge (SC) density near to the stellar surface.  We have calculated the energy losses due to plasma outflow resulting from these oscillations, focusing on cases where the size of the polar cap is small so that one can expand the Maxwell equations as Taylor series in powers of $ \\theta $, retaining only the two lowest order terms. This approach leads to a great simplification and allowed us to perform a thorough analysis of the solution. We have found that in GR, the polar cap is smaller than in Newtonian theory and have shown that this is due to the change in the geometry of the dipole magnetic field and the gravitational redshift of the energy of the outflowing plasma. Also, we found that the oscillation modes which have small $ \\theta_0 $, all have $ m' < 3 $ as in the Newtonian case.  The total energy loss resulting from the stellar oscillations causing plasma outflow through the polar cap region, is determined through an integral over the whole polar cap area, and so it depends on both the kinetic energy density of the outflowing plasma and the surface area of the polar cap. Although GR effects lead to some increase in the energy density of the outflowing plasma (due to the increase in the surface magnetic field strength for a given magnetic moment), the area of the polar cap is smaller in GR and we have found that the increase in the energy density of the outflowing plasma cannot compensate for the shrinking in size of the polar cap. Therefore the total energy losses for the toroidal oscillation modes are significantly smaller in GR than in Newtonian theory.  In conclusion, we point out that while our calculations represent an advance with respect to previous ones, they still do not include a number of very important aspects which would be necessary for a realistic description of these phenomena. Most importantly, as noted above, they do not take account of electric currents flowing in the magnetosphere. Inclusion of these in a consistent way would require solving a version of the non-linear ``pulsar equation'' \\citep{michel_73} for oscillating neutron stars, possibly by adopting a numerical approach similar to that of \\citet{contopoulos_99} \\citep[see also][]{gruzinov_05, timokhin_06} or performing {\\it time-dependent} simulations of the magnetosphere \\citep{spitkovsky_06, komissarov_06, mckinney_06}. This will be the subject of future investigations."
    },
    "0809/0809.2471_arXiv.txt": {
        "abstract": "Non-commutative black holes are characterised by a minimum mass which would result in a remnant after the Hawking evaporation ends. We numerically study the decay of neutral non-commutative black holes for up to ten spatial dimensions and typical parameters that would make their production possible at the LHC. Neglecting possible accretion mechanism, we find that decay-times are extremely short. ",
        "introduction": "\\setcounter{equation}{0} One of the most striking predictions of General Relativity is the existence of black holes, that is event horizons which surround and hide the point-like or otherwise singular classical sources. On general quantum-mechanical arguments, one however expects that singular sources are smeared out and there are indeed many theoretical reasons to believe the fabric of space-time can be effectively described by non-commutative geometry at short scales (see, for example Section~2 of Ref.~\\cite{nicolini} and Ref.~\\cite{szabo}). Such modifications to General Relativity must necessarily have an impact on the structure of black holes. \\par The non-commutative analogue of the Schwarzschild black hole metric was found in Ref.~\\cite{nicolini06} (see also Ref.~\\cite{nicolini05}) and subsequently generalised to include the electric charge~\\cite{nicolini07} and extra-spatial dimensions~\\cite{rizzo06,nicolini08}. A common feature for all such metrics is that there exists a minimum mass for the black hole which solely depends on the non-commutative length $\\ell$~\\footnote{This length is related to the non-commuative parameter $\\theta$ of Ref.~\\cite{nicolini} by $\\ell=2\\,\\sqrt{\\theta}$.} and the fundamental scale of gravity $M_g=\\ell_g^{-1}$ (for a comprehensive and updated review, see Ref.~\\cite{nicolini}). The former is usually identified with the Planck length $\\lp$ and the latter with the Planck mass $\\mpl$ in four space-time dimensions, but $\\ell$ could actually be much larger and $M_g$ much lower, possibly around $1\\,$TeV, if extra-dimensions exists~\\cite{add,rs}. This has opened up the possibility of having micro-black holes (mBH) with a mass of a few TeV's~\\cite{banks} that might be produced and detected at the Large Hadron Collider (LHC)~\\cite{dimopoulos,ch02,mbh}. \\par The standard picture of mBH at the LHC is that they should be produced with a mass of around $10\\,$TeV and very rapidly decay completely~\\cite{dimopoulos} via the Hawking effect~\\cite{hawking}. Generally small (but possibly significant) corrections to this scenario arise from the use of the microcanonical picture~\\cite{mfd}, which essentially enforces energy conservation throughout the evaporation process and suppresses the emission for small mBH mass~\\cite{ch02,c07}. If the non-commutative description of mBH given in Ref.~\\cite{nicolini} is correct, one however expects that a (stable) remnant will be left at the end of the Hawking evaporation~\\cite{hosse}, which might exit the accelerator and detectors. In this context, and for phenomenological purposes, it is therefore important to estimate the typical decay-times for such objects. For example, the issue of mBH life-times has recently been debated in Refs.~\\cite{mangano1}. \\par Since it is widely accepted that mBH discharge very quickly, via several processes, including the Schwinger mechanism at least for $3\\le d\\le 5$~\\cite{nicolini08}, we shall here consider only the neutral metrics of Refs.~\\cite{nicolini,rizzo06} for $d=5,\\ldots,10$ spatial dimensions (noting that higher values of $d$ seem to be favoured~\\cite{d>4}). In order to have mBH produced at the LHC, we need a minimum mBH mass of order $1\\,$TeV, which in turn implies that $\\ell_g\\sim \\ell$. Further, we shall see that the typical temperature of such mBH's always remains smaller than their mass, and microcanonical corrections can thus be neglected. The evolution of the mBH mass is then determined by solving the relevant equation numerically for initial mBH mass of order $10\\,$TeV and an upper bound for the decay-time obtained. \\par We shall use units with $c=\\hbar=1$. ",
        "conclusions": "\\setcounter{equation}{0} We have estimated the decay-times of non-commutative mBH's that might be produced at the LHC and found that they should evaporate nearly instantaneously (see the upper bound in Eq.~\\eqref{td}), much the same as is expected according to the ``canonical'' scenario of Ref.~\\cite{dimopoulos}. We can consequently conclude that non-commutative mBHs would evaporate within the detectors, since they can propagate at most a few nanometers away from their point of production during the evaporation. \\par The present results are similar to previous estimates of decay-times for usual (``commutative'') mBH's in the ADD brane-world~\\cite{add}, described according to the microcanonical picture~\\cite{mfd}, for which life-times were obtained on the order of $10^{-17}\\,$sec or shorter~\\cite{ch02}. Remarkably, our conclusions are also comparable with the recently estimated life-times of mBH's derived from Generalised Uncertainty Principles in Ref.~\\cite{scardigli}. Actually, this is not surprising, since the dependence of the temperature on the mBH mass is roughly similar in all of these cases~\\cite{c07}. Of course, the mBH's of Ref.~\\cite{ch02} evaporate completely (thus reaching zero temperature for vanishing mass) and their life-time is actually dominated by the latest stages (when $m\\lesssim 5\\,$TeV), whereas the end-point of the process here, as well as in Ref.~\\cite{scardigli}, is a (presumably) stable remnant of mass $m_{\\rm min}\\simeq 1\\,$TeV and zero Hawking temperature. \\par Let us conclude by mentioning that no accretion mechanism has been included in our analysis of the evaporation, which is fully consistent given the very short decay-times. What happens afterwards and the final fate of the remnant are still open questions. \\subsection*"
    },
    "0809/0809.0278_arXiv.txt": {
        "abstract": "We use 2D numerical simulations and eikonal approximation, to study properties of MHD waves traveling below the solar surface through the magnetic structure of sunspots. We consider a series of magnetostatic models of sunspots of different magnetic field strengths, from 10 Mm below the photosphere to the low chromosphere. The purpose of these studies is to quantify the effect of the magnetic field on local helioseismology measurements by modeling waves excited by sub-photospheric sources. Time-distance propagation diagrams and wave travel times are calculated for models of various field strength and compared to the non-magnetic case. The results clearly indicate that the observed time-distance helioseismology signals in sunspot regions correspond to fast MHD waves. The slow MHD waves form a distinctly different pattern in the time-distance diagram, which has not been detected in observations. The numerical results are in good agreement with the solution in the short-wavelength (eikonal) approximation, providing its validation. The frequency dependence of the travel times is in a good qualitative agreement with observations. ",
        "introduction": "Local helioseismology (time-distance helioseismology, acoustic holography and ring-diagram analysis) provides valuable information about the physical properties of solar sub-surface layers \\citep[\\eg ][]{Duvall+etal1993, Kosovichev1999, Kosovichev2002, Kosovichev+etal2000, Zhao+Kosovichev2003, Braun+Lindsey2000, Hill1988, Habler+etal2000}. Time-distance helioseismology \\citep{Duvall+etal1993, Kosovichev+Duvall1997} makes use of wave travel times measured for wave packets traveling between various points on the solar surface through the interior. By inversion of these measurements, variations of the wave speed and velocities of mass flows can be recovered below the visible solar surface. Inversion results for the travel times have been obtained for quiet Sun regions as well as for magnetic active regions including sunspots. In most of these cases, the variations of the travel times were assumed to be caused by mass flows and magneto-sonic speed fluctuations along the wave paths. It is known that sunspots possess strong magnetic fields with a complicated structure in the visible layers of the Sun where the Doppler measurements used by local helioseismology are taken \\citep{Solanki2003}. Consequently, such magnetic field can cause important effects on helioseismic waves, beyond the perturbation theories employed for helioseismic data analysis \\citep{Kosovichev+Duvall1997, Birch+Kosovichev2000, Gizon+Birch2002}. Magnetic field can change the acoustic excitation rate and spectral properties of solar oscillations \\citep[see the recent work by][]{Jacoutot+etal2008};  produce new wave modes; change wave propagation speeds; change the wave propagation paths and reflection characteristics at the near-surface layers.      The influence of magnetic field on the interpretation of local helioseismology measurements has not been fully explored. Theoretical efforts have been made by \\citet{Crouch+Cally2003}, \\citet{Cally2005, Cally2006}, \\citet{Schunker+Cally2006}, \\citet{Cally+Goossens2008}, \\citet{Schunker+etal2008}  to include mode conversion and model waves in magnetized structures by means of analytical and ray-path theories. These studies confirm the potential importance of magnetic effects. A more complete understanding of the problem can be reached by direct forward modeling of helioseismic data, by solving numerically the MHD equations in realistic magnetic configurations \\citep[\\eg, ][]{Gizon+etal2006, Khomenko+Collados2006, Hanasoge2008, Cameron+etal2007, Parchevsky+Kosovichev2008}. In these papers, the MHD equations, governing the problem, are solved for magnetic field configurations resembling a sunspot or a part of a sunspot. The results have demonstrated how magnetic fields change the speed, amplitude and spectral properties of the waves, and also produce wave transformation and refraction in the sunspot atmosphere.     Observational helioseismology over active regions has demonstrated repeatedly  that waves travel faster through sunspots in comparison with quiet Sun measurements \\citep[\\eg ][]{Duvall+etal1993, Braun1997}. This fact is still the subject of present investigations and its explanation is still far from clear. Although, given that sunspots possess an intense field at photospheric levels, direct or indirect the influence of the magnetic field is generally accepted \\citep[\\eg ][]{Hindman+etal1997}. At sub-photospheric levels the influence of the magnetic field becomes less and less important with depth as pressure forces dominate over magnetic forces. This is why the phenomena related to magnetic field are referred to as ``surface effects''. Explanations based on wave perturbations due to scattering phase shifts \\citep{Braun1997, Lindsey+Braun2005a, Lindsey+Braun2005b} taking place within a few Mm below the photosphere or local wave speed variations below sunspots \\citep{Kosovichev+etal2000} have been suggested to explain the different travel times measured in sunspots and in the quiet photosphere.  In this paper, we study wave properties and variations of travel times  with the help of numerical simulations and Eikonal solutions for wave propagation in sunspot-like magnetic structures in stratified surface and sub-surface layers. We  perform 2D numerical simulations of the waves produced by a single source for a series of model sunspots \\citep{Khomenko+Collados2008}. We study various types of MHD waves and their propagation properties numerically. In addition, for the fast magneto-acoustic mode we obtain a solution in the eikonal approximation, which is often important for correct interpretation of helioseismic results  \\citep{Gough2007}. Our numerical results clearly indicate that the fast mode is the primary source of helioseismic signals. We study the frequency dependence of the wave travel times and compare the results of the eikonal approach with our numerical simulations and observational data. Our results provide insight into the properties of helioseismic waves in magnetic regions and are important for interpretation of time-distance helioseismology measurements. ",
        "conclusions": "In this paper, we present an investigation of the influence of magnetic field on helioseismic wave propagation in active regions below sunspots. In order to better understand the physics of the waves in complex magnetic field configurations, we perform a single source experiment and study the propagation properties of the different wave modes excited by such a source, as well as the wave travel times. We have carried out our calculations in a series of sunspot models with different field strength and for two source positions located in magnetically and acoustically dominated regions. The results of our 2D simulations are compared to an asymptotic solution using the eikonal approximation. The results of both methods, numerical simulations and eikonal approximation, are found to be in a good qualitative agreement. Our approach has allowed us to investigate the effects of the magnetic field strength in sunspots on the travel-time variations obtained there from observations using the time-distance helioseismology techniques. In addition, the frequency dependence of such measurements has been studied.  The following conclusions summarize our study: \\begin{itemize}   \\item The wave source located immediately below the surface in a non-magnetic standard solar model excites a mixture of $p$, $f$ and $g$ mode waves. When the magnetic field is present, the same source excites a mixture of fast and slow MHD waves. The fast modes represent an analog of the acoustic $p$ modes, modified by the magnetic field.  \\item In addition to slow MHD waves excited directly by the source, there is another wave type with low vertical wavelength, propagating with speed close to that of the fast waves horizontally across the sunspot. We call them surface magneto-gravity waves. Several arguments listed in Sect. 4 indicate that these waves are different from either fast or slow waves. In our simulations these waves propagate below the visible surface. This makes their detection complicated in spectral observations.  \\item Slow MHD wave modes form an additional group of ridges in the time-distance diagram with much lower propagation speed. Their properties depend on the magnetic field strength of the model, as well as on the height in the atmosphere where the velocity is measured. The slow mode ridges on the time-distance diagrams are well isolated and do not affect the travel time measurements in magnetic regions.  \\item  The helioseismic waves below sunspots are speed up by the magnetic field. The travel times of these waves are shorter by, about 20--40 seconds compared to the waves in the quiet Sun. Changing the photospheric magnetic field strength from 0.9 to 2.4 kG produces a variation of the travel times by about 10--15 seconds.   \\item Magnetic field produces a strong frequency dependence of the wave travel times. High frequency waves travel faster than  low frequency waves. This happens because the propagation path of high frequency fast waves is modified strongly by the rapid increase of the Alfv\\'en speed in the top layers, since these waves penetrate higher up in the atmosphere. Low frequency waves are reflected below the photosphere due to the acoustic cut-off frequency and do not reach the heights where $v_A > c_S$ and the magnetic fields produces strong effects. The travel time differences change from $-20$ to $-35$ seconds between 3 and 5 mHz, in agreement with observations  \\citep[\\eg ][]{Rajaguru2008}.  \\end{itemize}  The analysis presented here is limited by the 2D approximation. As a next step in our work we plan to relax this assumption and to analyze 3D simulations, with a wave excitation by either single or multiple sources. This will allow us to build a more realistic picture and to complete our understanding of the interaction of helioseismic waves with magnetic fields of sunspots."
    },
    "0809/0809.2242.txt": {
        "abstract": "We study a scenario that large non-Gaussianity arises from the baryon asymmetry of the Universe. There are baryogenesis scenarios containing a light scalar field, which may result in baryonic isocurvature perturbations with some amount of non-Gaussianity.  As an explicit example we consider the Affleck-Dine mechanism and show that a flat direction of the supersymmeteric standard model can generate large non-Gaussianity in the curvature perturbations, satisfying the observational constraints on the baryonic isocurvature perturbations. The sign of a non-linearity parameter, $f_{\\rm NL}$, is negative, if the Affleck-Dine mechanism accounts for the observed baryon asymmetry; otherwise it can be either positive or negative. ",
        "introduction": "\\label{sec:1} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   The WMAP results~\\cite{Komatsu:2008hk} provided strong support for the inflation; the observed primordial fluctuations are consistent with nearly scale-invariant, adiabatic and Gaussian density perturbations, as predicted by a simple class of inflation models. Those predictions are derived based on a simple but crude assumption that it is only the inflaton that acquires sizable quantum fluctuations during inflation. Its apparent success, however, does not necessarily mean that such a non-trivial condition is commonly met in the landscape of the inflation theory.  It is perhaps natural to expect that there are many scalar fields in Nature. If some of them are light during inflation, they acquire quantum fluctuations, which may leave slight deviation from the above properties in e.g. the cosmic microwave background (CMB) anisotropy, such as the isocurvature perturbations and/or sizable non-Gaussianity. Interestingly, it was recently reported that large non-Gaussianity was detected by the analysis on the WMAP 3yr data \\cite{Yadav:2007yy}. The latest WMAP 5yr data seem to have the same tendency, although the vanishing non-Gaussianity is allowed within 95\\% C.L. \\cite{Komatsu:2008hk}.  Those hints on non-Gaussianity may be originated from such additional light scalars.  The non-Gaussian features in the observed CMB can arise from either adiabatic or isocurvature density perturbations. While the former was extensively studied in e.g. the curvaton \\cite{Lyth:2001nq,Lyth:2002my}/ungaussiton \\cite{Suyama:2008nt} scenarios, much less attention was paid to the latter case since this possibility was noted in Refs.~\\cite{Linde:1996gt,Boubekeur:2005fj}.  Recently Sekiguchi, Suyama and the current authors systematically studied how the non-Gaussian isocurvature perturbations exhibit themselves in the CMB temperature fluctuations \\cite{Kawasaki:2008sn}. In particular, it turned out that the non-Gaussianity in the isocurvature perturbations is enhanced at large scales, which may be confirmed (or refuted) by the current and future observations.  In Ref.~\\cite{Kawasaki:2008sn} we also studied the QCD axion as an example, which generates isocurvature perturbations in the cold dark matter (CDM).  In this paper we focus a scenario that baryonic isocurvature perturbations possess non-Gaussian properties.    One of the promising candidates for providing a theory beyond the standard model (SM) is supersymmetry (SUSY), and the supersymmetric SM (SSM) contains many flat directions. Those flat directions can play important roles in cosmology; they generate the baryon asymmetry of the Universe via the Affleck-Dine (AD) mechanism \\cite{Affleck:1984fy,Dine:1995kz}, and may account for dark matter by deforming into $Q$-balls \\cite{Coleman:1985ki,Kusenko:1997zq,Dvali:1997qv,Enqvist:1997si}.  If a flat direction remains light during inflation, it acquires quantum fluctuations, leading to the baryonic isocurvature fluctuations \\cite{Linde:1985gh,Enqvist:1998pf,Kasuya:2008xp}.  Moreover, as shown in \\cite{Kasuya:2008xp}, the phase direction of the flat direction generically remains flat in most inflation models in supergravity. Thus, we expect that the baryonic isocurvature perturbations as well as the associated non-Gaussianity are generically present in the AD mechanism.    In this paper we study the non-Gaussian property of the baryonic isocurvature perturbations produced in the AD mechanism.  We find that the resultant non-Gaussianity has distinctive features, compared to those produced in the curvaton/ungaussiton mechanism.  In terms of a non-linerity parameter, $f_{\\rm NL}$, to be defined later, the AD mechanism predicts a negative value of $\\fnl$ if the mechanism accounts for the observed baryon asymmetry of the Universe; otherwise, $f_{\\rm NL}$ becomes positive or negative, depending on the sign of baryon asymmetry created by the AD mechanism. We would like to emphasize that the AD mechanism can generate a significant amount of non-Gaussianity while accounting for the {\\it total} baryon asymmetry of the Universe.    This paper is organized as follows.  In Sec.~\\ref{sec:NGiso} we briefly summarize the calculation of non-Gaussianity from isocurvature perturbations.  Then we discuss non-Gausianity generated by the AD mechanism in Sec.~\\ref{sec:NGAD}.  Sec.~\\ref{sec:conclusion} is devoted to discussion and conclusions.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   We have focused on the AD mechanism in this paper, but there are other baryogenesis scenarios that may induce large non-Gaussianity in the baryon asymmetry. For instance, it is known that the spontaneous baryogenesis~\\cite{Cohen:1988kt} can lead to the baryonic isocurvature perturbations, since in the original model a non-vanishing chemical potential is due to a slow-rolling scalar field. In order to estimate non-Gaussianity, however, one has to specify how the chemical potential for the baryon number arises. In the case of the spontaneous baryogenesis using a flat direction of the SSM~\\cite{Chiba:2003vp}, the resultant non-Gaussianity has almost the same features as that in the case of the AD mechanism.  Another example is non-thermal leptogenesis using the right-handed sneutrino condensate, $\\tilde{N}$~\\cite{Murayama:1993em}. Suppose that $\\tilde{N}$ is light and fluctuating around the origin during inflation. Such a set-up may occur without fine-tuning, because the origin is the symmetry-enhanced point, and this is exactly what is considered in the ungaussiton scenario \\cite{Suyama:2008nt}.  The $\\tilde{N}$ can generate non-Gaussianity in both the adiabatic and baryonic isocurvature perturbations.  If the mass of the right-handed sneutrino is heavy, the baryonic isocurvature perturbations will become more important than the adiabatic one. In this scenario, the non-Gaussianity becomes large only if the non-thermal leptogenesis accounts for a tiny fraction of the baryon asymmetry.  It is also possible to consider non-Gaussianity in other type of isocurvature perturbations.  For instance, a large lepton asymmetry may have isocurvature perturbations with some amount of non-Gaussianity. If they are generated by the AD field with a lepton number \\cite{Kawasaki:2002hq}, non-Gaussianity arises from the fluctuations of the phase component, as in the case of the AD mechanism.  However, such isocurvature perturbations will affect the CMB temperature fluctuations in a different way, and so does the associated non-Gaussianity. We leave this issue for future work.  So far we have considered up to the three-point correlation functions. It is straightforward to extend our analysis to the correlation functions of higher order. In particular, when the linear perturbation is negligible,\\footnote{ To generate large non-Gaussianity while satisfying the constraint on the amplitude of the isocurvature perturbation, the linear term in Eq.~(\\ref{eq:S-expansion}) must be suppressed (see also Fig.~\\ref{fig:contour}.) } we  have a consistency relation between $\\fnl$ and $\\tau_{\\rm NL}$ as in the case of the ungaussiton~\\cite{Suyama:2008nt}, the latter of which is defined as a non-linearity parameter of the four-point correlation function.    In summary, in this paper we have studied a scenario that non-Gaussian baryonic isocurvature fluctuation is produced from the AD mechanism in SUSY. We have found that the AD mechanism can create large non-Gaussianity without conflicting with the current isocurvature constraint. We have seen that for $|f_{\\rm NL}^{\\rm (iso)}|\\gtrsim 1$ the Hubble parameter during inflation must be larger than about $10^{12}$~GeV, and that $|f_{\\rm NL}^{\\rm (iso)}|$ is bounded as $|f_{\\rm NL}^{\\rm (iso)}|< 60$ to satisfy the isocurvature constraint. Interestingly, as opposed to many known mechanisms for generating sizable non-Gaussianity~\\footnote{ % It is pointed out that negative $f_{\\rm NL}$ can be obtained in the curvaton scenario if the curvaton potential deviates from quadratic one \\cite{Enqvist:2008gk}. % }, the AD field can generate large non-Gaussianity even if it is the main component for generating the baryon asymmetry. The non-linerity parameter $f_{\\rm NL}^{\\rm (iso)}$ is negative if the AD mechanism is responsible for the total baryon asymmetry of the Universe, while it can be either positive or negative otherwise.  In our scenario the non-Gaussianity is necessarily accompanied with some amount of the isocurvature perturbations, therefore, if there are indeed large non-Gaussianity arising from the isocurvature perturbations, the future (or even on-going) observations will detect the isocurvature perturbations. It would be very interesting if the non-Gaussianity tells us about the origin of the baryon asymmetry, which is difficult to be probed, otherwise.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.1763.txt": {
        "abstract": "We show that a coupling between chameleon-like scalar fields and photons induces  linear and circular polarization in the light from astrophysical sources. In this context chameleon-like scalar fields includes those of  the  Olive-Pospelov (OP) model, which describes a varying fine structure constant. We determine  the form of this polarization numerically and give   analytic expressions  in two useful limits.  By comparing the predicted signal with current observations we are able to improve the constraints on the chameleon-photon coupling and the coupling in the OP model by over two orders of magnitude. It is argued that, if observed, the distinctive form of the chameleon induced circular polarization would represent a smoking gun for the presence of a chameleon. We also report a tentative statistical detection of a chameleon-like scalar field from observations of starlight polarization in our galaxy. \\pacs{04.50.Kd, 97.10.Ld, 14.80.Mz, 98.80.Cq} ",
        "introduction": "\\label{sec:intro} Extensions of the Standard Model of particle physics, such as string theory, introduce many new scalar fields which are not seen in the Standard Model. Such scalar fields are commonly invoked to explain the observed acceleration of the universe,  as inflation \\cite{Inflation} or dark energy \\cite{Peebles} fields, or to cause variations in fundamental constants \\cite{Uzan}.   If new scalar fields do indeed exist in the Universe, it is important to understand the properties of the theoretical models that describe them, e.g. the interactions of the scalar fields with themselves and with matter, which may give rise to additional observable effects that could be tested and constrained by experiments.  In this article we consider the effect of scalar field theories with a self-interaction  potential, $V(\\phi)$, and couplings to matter and light on observations of the polarization of light from astrophysical sources.  These scalar field theories are described by  the action: \\begin{eqnarray} S&=&\\int d^4x \\sqrt{-g}\\left(\\frac{1}{2\\kappa_4^2}R- \\frac{1}{2}g^{\\mu\\nu}\\partial_\\mu\\phi \\partial_\\nu \\phi -V(\\phi) \\right. \\label{action} \\\\ && \\left.-\\frac{B_{F}(\\phi/M_{0})}{4} F^2\\right)+ \\sum S_{\\rm m}^{(i)}  B_{i}(\\phi/M_{0}) g_{\\mu\\nu},\\psi_m^{(i)}) \\nonumber \\end{eqnarray} where the $S_m^{(i)}$  are  the matter actions for the matter fields $\\psi^{(i)}_m$, and the functions $B_{i}(\\phi/M_{0})$ and $B_F(\\phi/M_{0})$  determine the couplings of the scalar field, $\\phi$, to  the $i^{\\rm th}$ matter species, $\\psi_{i}$, and to the photon field respectively.  A scalar field with such couplings to matter fields might be expected to give rise to fifth force effects or violations of the weak equivalence principle.  In this  article we are specifically interested in a  scalar field, $\\phi$, which is light in  relatively low density regions such as galaxies, galaxy clusters and  the inter-galactic medium.  More precisely, in these regions we require the mass,  $m_{\\phi}$, of small perturbations  about the background  value of the scalar field, $\\phi_{\\rm b}$, to satisfy $m_{\\phi} \\lesssim  10^{-11}\\eV/c^2$.  Hence, any force mediated by $\\phi$ would have a  range $\\lambda_{\\phi} = 1 /   m_{\\phi} \\gtrsim 20\\km$.  Additionally  we  require that coupling between  photons and $\\phi$ in these regions is relatively strong: $$g_{\\phi\\gamma\\gamma} = \\frac{1}{M}  = \\left.{\\rm d}\\ln B_{F}/{\\rm d} \\phi\\right\\vert_{\\phi = \\phi_b}  \\gtrsim 10^{-11}\\,{\\rm GeV}^{-1}.$$ Therefore even if the coupling to matter is much weaker than the coupling to photons, since roughly $10^{-4}$ of the mass of nuclei  is due to electromagnetic interactions the  $\\phi$-mediated force between individual nuclei in these  backgrounds will be at least $10^{7}$ times the strength of  gravity on scales smaller than $\\lambda_{\\phi}$.  One might, at first glance,  conclude that a scalar field theory with these  properties is already strongly ruled out by laboratory constraints,  e.g. \\cite{EotWash04, EotWash06, Irvine, Colorado, Stanford1,  Stanford2},  on the strength of fifth forces.  Specifically,  measurements of the displacement of a micro-machined silicon cantilever using a fibre interferometer, reported in \\cite{Stanford2}, require that  a Yukawa-type fifth force with strength $10^{7}$ times that  of gravity has a range $\\lambda_{\\phi} < 5 \\mumetre$ with 95\\% confidence.   This, however, does not rule out the models we wish to consider  as neither the strength nor the  range of the $\\phi$-mediated force are necessarily the same in the relatively  high density environment of the laboratory as they are in the  low density background of space.  In recent years, two classes of  models have arisen that allow a scalar field that  is strongly interacting in low density environments  and yet is currently undetected in laboratory tests: the chameleon model,  \\cite{chamKA, chamstrong}, and the Olive-Pospelov (OP) model  \\cite{OP}.  Both models are described in detail in the following Section.  The mechanism by which these models avoid laboratory tests can be understood by extremizing  Eq. (\\ref{action}) with respect to $\\phi$ to give the  following field equation: \\begin{equation} \\square \\phi = V_{{\\rm eff},\\phi}(\\phi, T_{\\rm m}, F^2/4) \\label{phiField} \\end{equation} where \\begin{equation} V_{\\rm eff}(\\phi; T_{i}, F^2) = V(\\phi) + \\frac{B_{F}(\\phi)}{4} F^2 - \\frac{\\ln B_{m}(\\phi)}{2}  T_{\\rm m}, \\label{Veff} \\end{equation} and $T_{\\rm m} = g_{\\mu \\nu}T_{\\rm m}^{\\mu \\nu}$ is the trace of the energy momentum tensor  for matter: $$ T_{\\rm m}^{\\mu \\nu} = \\frac{2}{\\sqrt{-g}} \\frac{\\delta S_{\\rm m}}{\\delta g_{\\mu \\nu}}. $$ For non-relativistic matter $T_{\\rm m} \\approx -\\rho_{\\rm m}$, where $\\rho_{\\rm m}$ is the energy  density of the matter.  Both the chameleon and OP models play  the scalar field potential, $V(\\phi)$, off against the  matter couplings, $B_{F}$ and $B_{\\rm m}$, to make the vacuum expectation value (VEV), and  hence the properties of the field, depend strongly on the local density  of matter.  For convenience we shall refer to such a  scalar field  $\\phi$ as the chameleon or chameleon field, however our analysis  applies equally well to  both the  chameleon and OP models.  In this analysis we posit a universal coupling to the different matter  species i.e. $B_{i}(\\phi/M_{0}) = B_{m}(\\phi/M_{0})$.   Although it is not required by either model, we make this assumption because it simplifies the analysis whilst having little effect our  conclusions.  The best constraints on $M_0$ come from the requirement that corrections to particle  physics are small, which limits $M_0 \\gtrsim 10^{4}\\,{\\rm GeV}$ \\cite{chamPVLAS}.  A coupling between matter fields and the chameleon potentially causes violations of the weak equivalence principle (WEP) and other fifth force effects such as an effective alteration to Newton's inverse square law.  A coupling between photons and chameleons introduces additional observable phenomena for the chameleon field. If such a coupling has super-gravitational strength, it can result in a non-negligible  conversion of photons to chameleons and vice versa. The detectable effects associated with this conversion are similar to those predicted  for axion-like-particles (ALPs) which interact with light  \\cite{chamPVLAS}.  Mixing requires the interaction between two photons and one scalar particle, and so the effects of the mixing are most likely to be seen when a photon, or a scalar particle is passing through an external electromagnetic field.  The chameleon-photon coupling induces  both birefringence and dichroism \\cite{chamPVLAS,chamPVLASlong} in a coherent photon beam passing through an external magnetic field.  These effects could be detected by laboratory searches, such as the  polarization experiments PVLAS, Q\\&A, and BMV \\cite{Zavattini:2005tm,Zavattini:2007ee,Chen:2006cd,Robilliard:2007bq}, that are sensitive to new hypothetical particles with a small mass and coupling to photons.  Such experiments can constrain the coupling $M$ in the chameleon model, indeed  the otherwise anomalous detection of birefringence with a $5.5\\,{\\rm T}$ magnetic  field by PVLAS \\cite{Zavattini:2007ee}, could, at  least in principle, be explained by the presence of a chameleon  field \\cite{chamPVLASlong}. For the most widely studied class of potentials,  the PVLAS data was found to rule out $M \\lesssim 2 \\times 10^{6}\\,\\GeV$ \\cite{chamPVLASlong}. In the OP model the mass of the scalar field  in the laboratory is too large to produce a  detectable effect in these experiments.  If chameleons exist and  couple to photons,  then they could, as suggested in  \\cite{chamJar, RingJar}, be trapped  and slowly converted back into photons resulting in a long lived  chameleonic afterglow.  A number of experiments, most notably  GammeV \\cite{GammeV}, are searching or aiming to search for this  afterglow effect.  The GammeV chameleon search recently  announced its first results which, for models with $m_{\\phi}  < 10^{-3}\\eV$ in the interior of the experiment, ruled out $2.4 \\times 10^{5}\\GeV  < M < 3.9 \\times 10^{6}\\GeV$ \\cite{GammeVResults}. Ultimately  GammeV may be sensitive to $M \\lesssim 10^{8}\\,{\\rm GeV}$  \\cite{chamJar}, and an optimal sensitivity for afterglow searches of  $M < 10^{10}\\GeV$ is feasible within the constraints of currently  available technology \\cite{chamJar}.  Indeed for any of the effects  associated with the coupling to photons to be large enough  to be detected in the laboratory, either now or in the foreseeable  future, one must have $M \\lesssim 10^{10}\\,{\\rm  GeV}$. Such laboratory constraints do not, however, apply to the OP models.  These laboratory experiments need to be performed in a very good approximation to a vacuum otherwise the chameleon becomes too heavy to have a noticeable effect.  A complimentary approach to testing chameleon-photon couplings is to look for the effects of the coupling  in observations of astronomical objects.  The densities of interstellar space are typically very low and so the effects of the chameleon may be significant. Light from all astronomical objects  travels a significant distance through  magnetic fields in galaxies, galaxy clusters and possibly  in the intergalactic medium, before reaching the earth.   However astronomical magnetic fields are typically made up of large numbers of randomly oriented magnetic domains - a very different scenario to the well controlled constant magnetic fields of laboratory experiments.  In contrast to laboratory tests, such astrophysical effects should be see in the OP as well as the chameleon models.  The coupling between photons and chameleons means that photon number is not conserved, however, as the flux of photons emitted by astronomical objects is difficult to determine, measurements of flux cannot be used to bound the parameters of the chameleon model.  In the  following sections we show how the coupling between photons and chameleons  generates polarization in the light from astronomical objects.  Therefore measurements of polarization can be used to constrain the parameters of the chameleon model because the intrinsic polarization of astronomical objects is often very well constrained.   Astronomers are interested in measuring polarization because it can provide information both about the source of the radiation and about any magnetic fields present between the source and the earth.  Very precise astronomical polarization measurements are therefore available, and can be used to constrain the chameleon model.  This article is organized as follows: in \\S \\ref{sec:model} we introduce and provide further details of the two classes of model to which our analysis applies i.e the chameleon and OP models. The coupling of a  chameleon-like-particle to photons, is essentially the same as that  which is assumed for a scalar axion-like particle (ALP). There is  a great deal of literature concerning constraints (both local and astrophysical)  on ALPs. However, the density dependent mass of a chameleon field,  allows chameleon theories to  evade the tightest of these constraints.  In \\S \\ref{sec:ALPs}, we review previous constraints on ALPs,  and consider to what extent they do, or do not, apply to  chameleon-like models.  In \\S \\ref{sec:Optics}, we consider how the existence  of a chameleon-like field alters the polarization of light from astrophysical  objects  as it passes through an astrophysical magnetic field, and  derive the form of the induced polarization.  In \\S \\ref{sec:Magnetic} we discuss the observed and predicted  properties of the different types of large scale astrophysical magnetic fields.  In \\S \\ref{sec:obs}, we apply the results of the previous sections,  and use astrophysical polarization observations to  constrain the chameleon to photon coupling. We find that such  measurements place the tightest constraints yet on this coupling. Applying the analysis to starlight polarisation in our galaxy we find a tentative statistical detection of a chameleon-like scalar field. Finally we summarize our results in \\S \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  Theories of physics beyond the standard model typically predict the existence of new scalar fields.  If these scalar fields do exist it is important to understand both their self interactions and their interactions with the other fields present in the model in order to test and constrain the theory.  In this article we have studied the results of a coupling between scalar fields and photons on observations of astrophysical objects.  Specifically we have studied the scalar fields of the chameleon and Olive-Pospelov models, which are strongly interacting in low density environments yet currently undetected in the laboratory.  For simplicity we refer to both types of scalar field as chameleons.  If the chameleon field couples to photons then in the presence of a background magnetic field  the chameleon mixes with the component of the photon polarized orthogonally to the direction of the magnetic field.  We have studied the effect of this mixing on light beams passing through a large number of randomly oriented homogeneous magnetic domains, in order to predict the effects of chameleon-photon mixing on observations of light from astrophysical objects.  Typically both linear and circular polarization are induced in the light beam by mixing in such an environment.  We found analytic solutions to the equations describing the mixing in two important limits.  In the weak mixing limit the polarization fractions induced by chameleon photon mixing are highly wavelength dependent. If the light is not polarized at the source the averaged values of the total and circular polarization scale as $N\\Pphi$, that is as the product of the number of domains traversed and the probability of mixing in any one domain.  This limit is generally appropriate when one is considering the chameleon induced polarizations at wavelengths longer than roughly 1000\\AA.  In the maximal mixing limit, which applies when the chameleon-photon coupling is strong and the wavelength is sufficiently short, little or no circular polarization is produced by the mixing, but the production of linear polarization is at its strongest.  The distribution of the total polarization fraction after mixing in a large number of domains is independent of the parameters of the chameleon model, and instead depends only on the initial polarization of the light.  The average value of the total polarization fraction is always greater than $(\\pi/2)-1 \\approx 0.57$ in the maximal mixing limit.  Numerical simulations confirm the analytic analysis.  In particular they clearly demonstrate the existence of two wavelength scales $\\lambda_{\\rm osc} = \\lambda_{\\rm crit}/N = 4\\pi^2/\\vert m_{\\rm eff}^2 \\vert L_{\\rm path}$, and $\\lambda_{\\rm max} = \\lambda_{\\rm crit}{\\rm max} \\left(1, \\frac{BL}{\\pi \\sqrt{N} M}\\right)$ which determine the shape of the polarization signal.  Here $B$ and $L$ are the strength and domain size of the magnetic field. $g_{\\phi \\gamma \\gamma} = M^{-1}$ is the coupling between two photons and the scalar field. The linear polarization is greatest for $\\lambda \\lesssim \\lambda_{\\rm max}$, and the circular polarization is peaked for $\\lambda_{\\rm osc}\\lesssim\\lambda\\lesssim  \\lambda_{\\rm max}$.  Both polarization fractions are highly frequency dependent for $\\lambda \\gtrsim \\lambda_{\\rm osc}$.  This highly oscillatory behaviour means that observations of polarization at these wavelengths must be performed with a sufficiently good spectral resolution if any chameleon induced signal is to be resolved.  We have considered a wide variety of astrophysical observations and have used these to constrain the parameters of the chameleon model.  From observations of starlight polarization in our galaxy we show that at the 99\\% confidence level $M>1.1\\times10^9\\mbox{ GeV}$, which is an improvement of over two orders of magnitude on the previous best constraints on $M$.  The equivalent constraint on the Olive-Pospelov model is given in (\\ref{OPconstraint}).  Both constraints could, however, be evaded if the potential of the scalar field, $V(\\phi)$, is chosen so that the field is sufficiently heavy in regions with the density of our galaxy: $m_{\\phi} \\gg 10^{-11}\\eV$.  Constraints from objects outside the galaxy are limited by our lack of knowledge about a possible intergalactic magnetic field, $B_{\\rm IGM}$. If, however, $B_{\\rm IGM} \\approx 10^{-9}\\,{\\rm G}$ and is coherent over roughly ${\\rm Mpc}$ scales, then the lower bounds on $M$ and the OP model coupling scale are raised by roughly two orders of magnitude.  The circular polarization signal predicted from chameleon-photon mixing in a large number of randomly oriented magnetic domains was shown to be very distinctive.  Its frequency dependence is unlikely to have been caused by any other physical process, particularly as astrophysical objects do not normally produce significant amounts of circular polarization. To date, no observations of astrophysical circular polarization have yet been made with sufficient accuracy to allow us to search for a chameleon signal. Nonetheless, we have shown how future observations of circular polarization in the wavelength range $O(1)-O(1000)$\\AA\\ would be a smoking gun for chameleon-photon coupling.  We have also reported a seemingly strong statistical preference in observations of starlight polarization in our galaxy for the presence of a chameleon-like field. Precisely, at the 99\\% confidence level, we find \\begin{equation} \\left(\\frac{|B|L}{2M}\\right)_{\\rm rand}=(6.27\\pm 1.91)\\times 10^{-2} \\end{equation} where $B$ and $L$ are the strength and domain size of the random component of the galactic magnetic field.  Formally, the central value deviates from zero (the value for a theory without a chameleon) by more than $10\\sigma$.  It must be stressed, however, that this is only a preliminary analysis and we have only performed it for three out of a possible 121 objects. Before a detection could be claimed with any confidence, a full study of the possible backgrounds and systematics for these observations that could bias one towards larger values of $1/M$ would have to be undertaken.  Based on initial numerical simulations of data, it does, however, seem unlikely that polarization due to interstellar dust would produce such a strong signal.  In summary: astrophysical polarization measurements currently provide the strongest constraints on any coupling between photons and the scalar field, for many  chameleon and chameleon-like theories such as the Olive-Pospelov model, improving on previous constraints by more than two orders of magnitude.  Furthermore, future measurements of linear and, in particular, circular polarization at short wavelengths (i.e. $\\lesssim 2000$\\AA)  could provide one of the best tools in the continuing search for such scalar fields.  \\vspace{0.5cm}  \\noindent{\\bf Acknowledgements:} CB, ACD and DJS are supported by STFC. We are grateful to J. D. Barrow, Ph. Brax, C. van de Bruck, D. F. Mota, C. Spyrou and W. Sommerville for helpful conversations.  \\appendix"
    },
    "0809/0809.0939_arXiv.txt": {
        "abstract": "Bayesian Inference is a powerful approach to data analysis that is based almost entirely on probability theory. In this approach, probabilities model {\\it uncertainty} rather than randomness or variability. This thesis is composed of a series of papers that have been published in various astronomical journals during the years 2005-2008. The unifying thread running through the papers is the use of Bayesian Inference to solve underdetermined inverse problems in astrophysics. Firstly, a methodology is developed to solve a question in gravitational lens inversion - using the observed images of gravitational lens systems to reconstruct the undistorted source profile and the mass profile of the lensing galaxy. A similar technique is also applied to the task of inferring the number and frequency of modes of oscillation of a star from the time series observations that are used in the field of asteroseismology. For these complex problems, many of the required calculations cannot be done analytically, and so Markov Chain Monte Carlo algorithms have been used. Finally, probabilistic reasoning is applied to a controversial question in astrobiology: does the fact that life formed quite soon after the Earth constitute evidence that the formation of life is quite probable, given the right macroscopic conditions? ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1424_arXiv.txt": {
        "abstract": "We present Gemini near-infrared adaptive optics imaging and spectroscopy of a planetary mass candidate companion to \\planethost, a roughly solar-mass member of the 5~Myr-old Upper Scorpius association. The object, separated by 2.22\\arcsec\\ or 330~AU at $\\sim$150~pc, has infrared colors and spectra suggesting a temperature of $1800_{-100}^{+200}$~K, and spectral type of L4$_{-2}^{+1}$. The $H$- and $K$-band spectra provide clear evidence of low surface gravity, and thus youth. Based on the widely used DUSTY models, we infer a mass of $8_{-2}^{+4}~M_{\\rm Jupiter}$. If gravitationally bound, this would be the lowest mass companion imaged around a normal star thus far, and its existence at such a large separation would pose a serious challenge to theories of star and planet formation. ",
        "introduction": "Since 1995, over 300 planets have been identified around stars other than the Sun,\\footnote{See {\\it The Extrasolar Planets Encyclopaedia} at \\url{http://exoplanet.eu/}} revealing a remarkable diversity in the properties of planets as well as the architecture of planetary systems. Yet, the three hitherto successful planet detection techniques -- radial velocity monitoring, transit searches and microlensing surveys -- are generally limited to planets in orbits much smaller than the full radial extent of our solar system. Thus, the discoveries to date, while dramatic, can only provide an incomplete picture of the overall planet population. Direct imaging can help complete the census by uncovering planets at wide separations. Furthermore, it enables extensive follow-up studies to characterize the planets. However, due to the small angular separation and the high brightness contrast between a planet and its star, direct imaging has not succeeded thus far \\citep[e.g.][and references therein]{lafreniere07}. At present, our best hope lies in targeting nearby young stars for newborn planets.  During their formation, giant planets and brown dwarfs (BDs), like stars, generate heat from gravitational contraction. But, unlike stars, which eventually reach core temperatures sufficient for hydrogen fusion, these objects are left without a means of producing energy, and thus rapidly cool down and become dimmer with time. Therefore younger planets and BDs, as companions to stars, are much easier to detect directly than their older counterparts. Accordingly, the lowest mass companions imaged so far all orbit stars younger than $\\sim$30~Myr; several of them (e.g. CHXR~73~B, \\citealp{luhman06}; AB~Pic~B, \\citealp{chauvin05b}; DH~Tau~B, \\citealp{itoh05}; GQ~Lub~B, \\citealp{neuhauser05}) have masses slightly above the deuterium burning limit of 13~$M_{\\rm Jupiter}$, which is often used as a boundary to differentiate planets from the more massive BDs\\footnote{{\\it Working Group on Extrasolar Planets of the International Astronomical Union}, see \\url{http://www.dtm.ciw.edu/boss/definition.html}}, though some researchers prefer a distinction based on the formation mechanism, whereby planets form within a circumstellar disk and BDs form through cloud fragmentation. One lower mass companion ($\\sim$8~$M_{\\rm Jupiter}$) orbits the young 25~$M_{\\rm Jupiter}$ BD 2MASSW~J1207334-393254, rather than a star \\citep{chauvin04, mohanty07}; other similar systems might exist \\citep[e.g.][]{bejar08}. All of these low-mass sub-stellar companions are located at large separations from their primaries, in sharp contrast with solar system planets and extra-solar planets detected by indirect methods, possibly reflecting a fundamental difference in their formation mechanism.  Here we report the direct imaging discovery of a 6--12~$M_{\\rm Jupiter}$ candidate companion to a young solar-mass star, \\planethost, in the nearby ($145\\pm20$~pc), 5~Myr-old Upper Scorpius association \\citep{preibisch02}. ",
        "conclusions": "\\label{sect:discussion}  Although our photometry and spectroscopy establish that the candidate companion has low gravity and a mass below 0.012~$M_\\odot$, they do not prove that it is physically bound to the primary star rather than a free-floating planet in the association. Based on integration of their best-fit mass function, \\citet{preibisch02} estimate that the entire Upper Scorpius population comprises 2525 stars more massive than $>0.1$~$M_\\odot$, distributed over an area of $\\sim$150~deg$^2$. Assuming very conservatively that there are as many free-floating planets in the association as there are stars $>0.1$~$M_\\odot$, the probability that one would fall within 2.5\\arcsec\\ from any of the 85 stars we have observed would be only 0.002. Thus this scenario is unlikely. Verification of common proper-motion over the next few years will nevertheless be important, although it would not readily confirm physical association of the two objects given the small internal velocity dispersion of the association. The latter would require detection of orbital motion, which could take several years given the small orbital motion expected ($\\sim$2~mas~yr$^{-1}$).  Previous direct imaging surveys for planets around nearby solar-type stars have put upper limits of $\\sim$6\\% on the fraction of stars with planets more massive than 5~$M_{\\rm Jupiter}$ at separations over $\\sim$50~AU \\citep[][and references therein]{lafreniere07}. Our single candidate in a sample of $\\sim$85 stars is consistent with these results, and confirms that planets on wide separations are rare also at ages of a few Myr.  Compared to known star-planet systems, the inferred mass for the candidate companion of \\planethost\\ is near the high end of the range. What stands out even more is the large separation of $>$300~AU. The simplest explanations would be that there is no physical association or that the companion has recently been ejected. Both appear unlikely, however, and thus it seems worthwhile to speculate about possible formation scenarios. In situ growth via core accretion \\citep{pollack96} appears unlikely: the formation timescale would greatly exceed the age of the system, even if a disk could survive long enough. On the other hand, at smaller separations this mechanism could build a $\\sim$5~$M_{\\rm Jupiter}$ planet in 1--5~Myr \\citep{pollack96,alibert05}, which might then migrate outwards through interactions with a disk or other massive planets (either can be quite rapid; \\citealt{peplinski08,veras04}). A problem with this alternative, however, is that a planet formed through core accretion may not be able to reach the observed temperature \\citep{marley07}. In situ formation via gravitational instability \\citep{boss97} is another possibility, but it would require a rather massive ($\\gtrsim$0.2-0.35 M$_\\star$) disk. Typical disks around T Tauri stars contain only 0.01--0.1 M$_\\star$ of material \\citep{scholz06}, but more massive disks may exist during earlier proto-stellar phases. Still, the thermal state of the disk may prevent fragmentation even at radii as large as required here \\citep{matzner05}. Lastly, this system may have formed like a binary star, through the fragmentation of a pre-stellar core.  Current star formation simulations, however, find it difficult to make binaries with extreme mass ratios and very low mass components \\citep{bate03}.  Future observations could look for additional closer-in companions, evidence of a large debris disk, and whether the object we found is in a highly eccentric or nearly circular orbit. If bound, the very existence of the \\primary\\ system poses a challenge to theories of planet and star formation, and may well suggest that there is more than one mechanism in nature for producing planetary mass companions around normal stars."
    },
    "0809/0809.0172_arXiv.txt": {
        "abstract": "Precision Doppler velocity measurements from the Anglo-Australian Telescope reveal a planet with a 9.4$\\pm$0.4 year period orbiting the M1.5 dwarf GJ\\,832.  Within measurement uncertainty the orbit is circular, and the minimum mass (\\msini) of the planet is 0.64$\\pm$0.06\\,\\mjup.  GJ\\,832 appears to be depleted in metals by at least 50\\% relative to the Sun, as are a significant fraction of the M dwarfs known to host exoplanets. GJ\\,832 adds another Jupiter-mass planet to the known census of M dwarf exoplanets, which currently includes a significant number of Neptune-mass planets. GJ\\,832 is an excellent candidate for astrometric orbit determination with $\\alpha \\sin i = 0.95$ mas. GJ\\,832b has the second largest angular distance from its star among radial velocity detected exoplanets (0.69 arc sec) making it a potentially interesting target for future direct detection. ",
        "introduction": "\\label{intro}  Most of the known exoplanets orbit late-F, G, or early-K dwarfs, with masses ranging from 0.7 to 1.2\\,\\msun.  There are nearly 2,000 such stars within 50\\,pc brighter than V=8.  In contrast there are just a handful of M dwarfs brighter than V=8.  Although M dwarfs make up 70\\% of nearby stars, their faintness in the optical makes them  difficult targets for which to obtain precision Doppler velocities. As a result they make up as little as 5\\% of the current planet search targets, and only 11 exoplanets have been found to date, orbiting a total of 7 M dwarfs. Just over half of these M dwarf planets are less massive than Neptune, leading to speculation that M dwarfs typically host either fewer planets than G dwarfs \\citep{johnson07}, or lower mass planets as a result of their smaller proto-planetary disks \\citep{laughlin04,il05,kk08}.  The primary parameters in planet formation theory are the mass of the central star, the mass of the protoplanetary disk, and the metallicity of the system.  While it is now well established for late-F, G, and K dwarfs that metal-rich stars are enhanced in planets relative to metal-poor stars \\citep{g1,g2,g4,g5,g3,sb,s1,s04,g6,re,fv,bond}, it has been harder to establish the importance of stellar mass on planet formation since most of the stars under survey lie in the relatively narrow mass range encompassed by late-F, G and K dwarfs.  The searches which {\\em are} being done for planets orbiting M dwarfs, will ultimately provide the data needed to see if the metallicity-planet relation extends down to the M dwarf regime, and whether the mass distribution of exoplanets formed around M dwarfs is similar to, or different than, that for more massive host stars. M dwarf Doppler surveys, therefore,  have the power to address some of the most important questions in exoplanetary science, as they extend the mass range of potential exoplanet host stars down to 0.3\\,\\msun.  We report here a new extrasolar planet in a long period orbit with eccentricity consistent with zero, discovered by the Anglo-Australian Planet Search (AAPS).  The AAPS program is described in Section 2.  The characteristics of the host star and our Doppler measurements are presented in Section 3. A discussion follows. ",
        "conclusions": "GJ\\,832 at a distance of 4.93pc is one of the nearest known exoplanetary systems. The combination of the small distance and relatively long period gives a large angular distance from the star of 0.69 arc seconds for an edge-on circular orbit. This is exceeded only by $\\epsilon$ Eri among radial velocity detected exoplanets, and only six other systems exceed 0.2 arc sec. GJ\\,832b is therefore a potentially interesting target for direct detection, although the high constrast with the star \\citep[likely to be $<10^{-8}$;][]{Burrows04} still makes this an extremely challenging observation.  GJ\\,832 is an excellent candidate for astrometric orbit determination. The astrometric orbit semimajor axis is $\\alpha \\sin i = 0.95$ mas, which is comparable to that of $\\epsilon$ Eri for which an astrometric orbit was determined by \\cite{benedict06} and larger  than that of GJ 876 which also has an astrometric orbit determination \\citep{benedict02}. The astrometric orbit would enable the inclination to be determined, removing the current $\\sin i$ uncertainty on the mass.  Seven M dwarfs (including GJ\\,832) are currently known to host as many as 11 exoplanets, and these are listed in Table \\ref{mdwarfs} (see table notes for references). As noted earlier, determining the metallicities of M dwarfs is notoriously difficult -- published metallicity estimates are available for several of the known exoplanet host M dwarfs, and these are listed in the Table. In addition, we have also derived for all seven M-dwarfs a photometric metallicity estimate, using  the technique of \\citet{Bonfils05a}, which has the advantage of being uniform over all these M dwarfs. On average the \\citet{Schiavon97} metallicities appear to be systematically 0.3-0.4\\,dex lower than those derived from the Bonfils et al. calibration. The \\citet{bean06} metallicities are similarly on the metal-poor side of the Bonfils et al. results, though not by as much ($\\approx$ 0.2\\,dex). In general, the metallicity trends are similar across all three calibrations, and it is clear there is a metallicity spread across the observed M dwarf exoplanet hosts.  Based on these metallicity estimates, it would appear that four of the current M dwarf exoplanet host stars are somewhat metal-poor, two have about solar metallicity, and one is slightly metal rich. Given the well known correlation between stellar metallicity and observed exoplanet frequency for F, G, and K dwarf host stars, this metallicity distribution for M dwarf host stars is quite unexpected. Whhile the numbers of systems are small there is no obvious difference in metallicity between the stars hosting Jupiter-mass planets and those hosting Neptune-mass planets.  The correlation between high stellar metallicity and planets for late-F, G, and early-K dwarfs points toward the core accretion model for planet formation. But there does not appear to be strong evidence to date that M dwarf planet formation is strongly correlated with high metallicity. This is puzzling, particularly in view of the fact that M dwarfs probably have lower mass protoplanetary disks, and therefore would need even higher metallicity than a F, G or K dwarf to provide enough solid material (silicates and ices) to build a planetary core. Obviously it must be kept in mind that measuring metallicities for M dwarfs is problematic, and that even the \\citet{Bonfils05a} calibration (though empirically based and moderately robust) is only good to $\\pm$0.2\\,dex. Nonetheless it is interesting to consider possible means by which M dwarf exoplanets could be formed in such a manner as to {\\em not} display the strong metallicity correlation seen in FGK dwarfs. One initially attractive explaneation is that since  M dwarfs are essentially immortal on a Hubble time scale, the vast majority of nearby M dwarfs could be old metal poor stars. Unfortunately such an explanation would appear unlikely. The study of M dwarf kinematics has an extensive and venerable history \\citep[e.g.][]{wielen77, wu95, rhg95} which has contributed to the creation of extensive and sophisticated models of the stellar populations present in the Solar Neighbourhood \\citep[e.g. the Besancon models of][]{robin2003}. More recently, the availability of huge numbers of M dwarf spectra from the SDSS survey have enabled sophisticated tests of the kinematics of the Besancon models by \\citet{bochanski07}, and  have substantially born out the Besancon model predictions for M dwarfs. Those models indicate that dominant solar neighbourhood M dwarfs will be thin disc members with ages almost uniformly spread between 0.1 and 10\\,Gyr. Thick disc M dwarfs (which would indeed be expected to have systematically lower metallicities) will be present at much lower densities \\citep[around a factor of one twentieth or less;][]{robin2003}, and the probability that they would make up four of the seven M dwarf exoplanet hosts would seem to be negligibly small.  An alternative explanation could be that M dwarf planets might form primarily via the disk instability mechanism \\citep[see e.g.][and references therein]{boss08}, rather than via core accretion, which would make their formation probability more or less independent of metallicity.   Six of the eleven exoplanets known to orbit M dwarfs have minimum masses less than 0.1\\,\\mjup.  In contrast, only nine planets with \\msini $<$ 0.1\\,\\mjup\\ have been found among the 216 Doppler velocity planets with B-V $<$ 1.2 in the ``Catalog of Nearby Exoplanets'' \\citep{butler06b}. With a minimum mass of 0.64$\\pm$0.06\\,\\mjup, GJ\\,832b is the fifth jovian mass planet found orbiting an M dwarf.  The most massive M dwarf planet yet found is 1.93\\,\\mjup.  Since massive planets are by far the easiest ones to find, planets of more than 2\\,\\mjup\\ orbiting within 3\\,AU of M dwarfs must be rare, occuring less than around once per 300 M dwarfs."
    },
    "0809/0809.2800_arXiv.txt": {
        "abstract": "The distributions of galaxy properties vary with environment, and are often multimodal, suggesting that the galaxy population may be a combination of multiple components.  The behaviour of these components versus environment holds details about the processes of galaxy development.  To release this information we apply a novel, nonparametric statistical technique, identifying four components present in the distribution of galaxy H$\\alpha$ emission-line equivalent-widths. We interpret these components as passive, star-forming, and two varieties of active galactic nuclei. Independent of this interpretation, the properties of each component are remarkably constant as a function of environment.  Only their relative proportions display substantial variation.  The galaxy population thus appears to comprise distinct components which are individually independent of environment, with galaxies rapidly transitioning between components as they move into denser environments. ",
        "introduction": "It has long been recognised that galaxies may be divided into at least two distinct sub-populations. Originally this division was based on visual appearance.  Most galaxies can be morphologically classified as either elliptical or spiral.  Finer classification is possible, discretizing an apparently continuous variation in galaxy appearance.  However the dichotomy between elliptical and spiral morphology is more pronounced than the variations within each class.  Subsequently, it has been discovered that several other, more quantitative, galaxy properties are distributed unevenly or in a multi-modal manner.  The colour distribution of SDSS galaxies is strongly bimodal \\citep{2001AJ....122.1861S}. Galaxies in the ``red'' and ``blue'' modes can be roughly identified as those with elliptical and spiral morphology, respectively \\citep{2002AJ....124..646H,2006MNRAS.368..414D}.  Whereas morphology reflects the dynamical state of galaxies, colour is related to their star-formation history, particularly over the last $\\la 10^9$ years.  The colour bimodality thus implies a division of the galaxy population into blue galaxies, which have recently formed stars, and red galaxies, which have not.  Such a bimodality in the star-formation properties of galaxies has also been observed using more direct measures of current star-formation, such as emission-line strength \\citep{2004MNRAS.348.1355B}.  The position of the red and blue galaxy sequences in the colour--luminosity or colour--stellar mass planes display only a weak dependence on environment.  However, the relative proportions of galaxies in the two sequences vary strongly. In regions with a higher local galaxy density the fraction of galaxies on the red sequence is higher \\citep{2004ApJ...615L.101B,2004AIPC..743..106B,2006MNRAS.373..469B}.  It remains a matter of debate whether colour is more closely related to environment than morphology.  Some claim that trends in morphology versus environment can be mostly explained via a morphology--colour relation which is almost independent of environment \\citep{2006MNRAS.366....2W,2006astro.ph..8353B,2006astro.ph.10171B,2007MNRAS.376L...1W}. However, other studies oppose this view \\citep{2007ApJ...658..898P}, and it has been clearly shown that the colour and morphology bimodalities behave differently with respect to environment and stellar mass \\citep{2008arXiv0805.2612B}.  There are growing indications that, in a fraction of the galaxy population, star-formation must be terminated rapidly \\citep{2004MNRAS.348.1355B,2006MNRAS.373..469B}.  The emission lines in galaxy spectra provide a way of measuring the level of current star formation on a timescale of $\\la 10^7$ years.  They therefore trace rapid star formation variations more sensitively than colour.  Another important property of emission lines is that they are produced by active galactic nuclei (AGN), in addition to star formation.  AGN are present in many galaxies, and are thought to be produced by accretion of material onto the super-massive black holes which appear to reside at the centre of most, if not all, galaxies \\citep{1998Natur.395A..14R}.  Recently a variety of studies have suggested that AGN strongly influence star-formation in their host galaxies, and thus play an important role in defining the galaxy population \\citep{2003MNRAS.346.1055K,2005MNRAS.364.1337S,2006MNRAS.365...11C,2006MNRAS.370..645B}. The potential presence of an AGN contribution complicates the traditional usage of emission-lines as an indicator of star formation rate (SFR).  However, it also presents an opportunity to study these two interdependent processes, star-formation and AGN, through the distribution of a single quantity.  Galaxies with contrasting properties are found to be distributed differently in space.  Elliptical galaxies cluster together more strongly than spirals \\citep{2000ApJ...545....6B,2001ApJ...554..857G}. Similarly, red galaxies are preferentially found in denser environments than blue galaxies \\citep{2005ApJ...630....1Z}.  We have a well developed theory for how structure forms in the cosmos, at least in terms of the underlying cold dark matter which dominates the mass density \\citep{2005Natur.435..629S}.  Baryonic matter is expected to be similarly distributed, in broad terms.  This theory thus explains the range of galaxy environments observed.  However, the properties of galaxies as a function of environment is a much more complicated issue, depending on the detailed physics of galaxy formation and evolution.  By studying trends in the galaxy population with environment we can learn about these physical processes.  There has been a logical progression in studies of galaxy properties as a function of environment.  Early work was based on dividing galaxies into simple classes and looking at variations in the fractions of galaxies of each class in bins of environment \\citep{1985ApJ...288..481D}.  As galaxy samples grew, this moved on to examining trends in the mean properties of galaxies as a smooth function of local galaxy density \\citep{2002MNRAS.334..673L,2003ApJ...584..210G}. A significant development was fitting to the data functions that describe the distribution of galaxies in two classes \\citep{2004ApJ...615L.101B,2004AIPC..743..106B,2006MNRAS.373..469B}. Most of the approaches employed so far have relied upon enforcing a predefined view of how to divide or classify the galaxy population in increasingly complex ways.  However, our understanding of the physical processes at work is highly uncertain and does not provide a sufficient basis to make this decision.  Our only guide is the data itself.  A natural next step is thus to turn to nonparametric methods, where the components of the population are deduced consistently from the data itself.  Recently, several studies have performed multivariate statistical analyses on datasets containing a wide variety of galaxy properties, in order to identify components of the galaxy population, and determine which properties are most important for identifying to which component a galaxy belongs \\citep{2005MNRAS.363.1257E,2006MNRAS.373.1389C}.  Such studies are highly informative, but become complicated when one wishes to determine the behaviour of the identified components versus another variable.  In this paper we are primarily concerned with variation in the components of the galaxy population as a function of environment. The statistical method we present below may be straightforwardly applied to multivariate datasets.  However, for simplicity, in the present work we consider the environmental dependence of just one galaxy property.  Nevertheless, even with this elementary approach, we are able to learn much about the galaxy population.  \\begin{figure*} \\includegraphics[clip=True,trim = 3cm 22.3cm 9.5cm 3.3cm,width=1.0\\textwidth]{fig1.ps} \\caption{\\label{fig:w13} The distribution of transformed H$\\alpha$ equivalent width ($\\W$) for (left) low and (right) high density environments.  The histogram displays the data, with Poisson uncertainties indicated by the grey shading.  The red, purple, green and blue lines show the components derived by applying the NMR technique.  The brown line gives the sum of these components, which is clearly a good representation of the data.} \\end{figure*} ",
        "conclusions": ""
    },
    "0809/0809.3940_arXiv.txt": {
        "abstract": "s{We report a stochastic mechanism of particle acceleration from first principles in an environment having properties like those of Radio Lobes in AGNs. We show that energies $\\sim 10^{20}$ eV are reached in $\\sim 10^6$ years for protons. Our results reopen the question regarding the nature of the high-energy cutoff in the observed spectrum: whether it is due solely to propagation effects, or whether it is also affected by the maximum energy permitted by the acceleration process itself.} ",
        "introduction": "The search for the origin of the UHECRs still represents one the major challenges of theoreti\\-cal astrophysics. Theoretical models may be divided into two classes: the so-called ``bottom-up''scenarios~\\cite{ta04}, in which a specific process of acceleration in a particular astrophysical object leads to UHEs; and the so-called ``top-down'' prescription~\\cite{bs00}, in which UHE particles are produced through the decay of super-heavy dark matter particles or by collision among cosmic strings or by topological defects. The recent measurement by Auger, demonstrating a low fraction of high-energy photons in the CR distribution, rule out the top-down models, in which the UHECRs represent the decay products of high-mass particles created in the early Universe~\\cite{semikoz07}. The top-down models based on topological defects, however, are still compatible with the current data and could be constrained by future experiments.  Recently a steepening in the UHECR spectrum has been reported by both the HiRes~\\cite{H07} and Auger~\\cite{auger2007} collaborations. This result may be a strong confirmation of the predicted Greisen-Zatsepin-Kuzmin (GZK) cutoff due to photomeson interactions between the UHECRs and low-energy photons in the cosmic microwave background (CMB) radiation~\\cite{greisen66,zatsepin66}.  The most telling indicator for the possible origin of these UHECRs is the discovery by Auger of their clustering towards nearby ($\\sim 75$ Mpc) AGNs along the supergalactic plane. However, the question remains open regarding the mechanism of acceleration to such high energies and on the origin of the observed cutoff in the spectrum, i.e., if it is due solely to the GZK effect, or whether it also points to an intrinsic limit to the acceleration efficiency.  UHECRs generation scenarios include the so-called first-order Fermi acceleration in GRBs, Pulsar Wind Bubbles, and also relativistic second order Fermi acceleration~\\cite{f49,g08}. We report here a treatment of particle acceleration in the lobes of radio-bright AGNs from first principles~\\cite{fm08}, considering the acceleration of charged particles via random scatterings (a second-order process) with fluctuations in a turbulent magnetic field. ",
        "conclusions": "In view of the very good match between our theoretical simulation and the Auger observations, it is worth emphasizing that this calculation was carried out without the use of several unknown factors often required in approaches involving a hybrid Boltzmann equation to obtain the phase-space particle distribution. In addition, we point out that the acceleration mechanism we have invoked here is sustained over 10 orders of magnitude in particle energy, and the UHECRs therefore emerge naturally---without the introduction of any additional exotic physics---from the physical conditions thought to be prevalent within AGN giant radio lobes.  As the Auger observatory gathers more data and improves the statistics, our UHECR source identification will continue to get better. Eventually, we should be able to tell how significant the GZK effect really is, and whether the cutoff in the CR distribution is indeed due to propagation effects, or whether it is primarily the result of limitations in the acceleration itself. Given the fact that energies as high as $\\sim 10^{20}$ eV may be reached within typical radio lobes, it is possible that both of these factors must be considered in future refinements of this work."
    },
    "0809/0809.3171_arXiv.txt": {
        "abstract": "While iron emission lines are well studied in black hole systems, both in X-ray binaries and Active Galactic Nuclei, there has been less of a focus on these lines in neutron star low-mass X-ray binaries (LMXBs).  However, recent observations with {\\it Suzaku} and {\\it XMM-Newton} have revealed broad asymmetric iron line profiles in 4 neutron star LMXBs, confirming an inner disk origin for these lines in neutron star systems.  Here, we present a search for iron lines in 6 neutron star LMXBs.  For each object we have simultaneous {\\it Chandra} and {\\it RXTE} observations at 2 separate epochs, allowing for both a high resolution spectrum, as well as broadband spectral coverage.  Out of the six objects in the survey, we only find significant iron lines in two of the objects, GX~17+2 and GX~349+2. However, we cannot rule out that there are weak, broad lines present in the other sources.  The equivalent width of the line in GX~17+2 is consistent between the 2 epochs, while in GX~349+2 the line equivalent width increases by a factor of $\\sim$3 between epochs as the source flux decreases by a factor of 1.3.  This suggests that the disk is highly ionized, and the line is dominated by recombination emission.  We find that there appears to be no specific locations in the long-term hardness-intensity diagrams where iron emission lines are formed, though more sources and further observations are required. ",
        "introduction": "Accretion disks around compact objects in AGN and X-ray binaries exist in strong gravity around these objects, and thus can provide a means to study gravity in the strong-field regime.  In many of these objects, an Fe K-shell fluorescence emission line is observed in the X-ray spectra \\citep[e.g.][]{white86,nandra97,reynolds97,asai00,miller04}.  These lines are formed in the inner accretion disk and are observed to be broadened due to strong Doppler shifts and gravitational redshifts in that region \\citep[see][for a review of iron lines]{miller07}.  As these lines are shaped by the motions of gas in the disk, they are an important diagnostic of the inner accretion disk radius in these objects.  Observations of black holes (in both AGN and X-ray binaries) have revealed some of these lines to be highly asymmetric \\citep[e.g.][]{tanaka95,fabian02,miller04} due to the extreme motions of gas in the inner disk around black holes.  While iron lines in black holes are now well-studied, iron lines in neutron star low-mass X-ray binaries (LMXBs) have not been studied in the same detail.  Nevertheless, iron lines in neutron star LMXBs are well known, and these lines are significantly weaker than in black holes \\citep[e.g.][]{white85,white86,hirano87,asai00,disalvo05}. Thus, the line shapes could not be studied in detail, and could be modeled as just a Gaussian.  Only with recent sensitive {\\it XMM-Newton} and {\\it Suzaku} observations have broad, asymmetric iron lines been observed in neutron stars, and these observations highlight the potential to use iron lines in neutron stars as a probe of the inner disk radius; the inner disk radius is an upper limit on the neutron star radius \\citep{bhattacharyya07,cackett08}.  The use of these lines as a disk diagnostic also allows for a test of the standard neutron star accretion picture described by color-color diagrams and kilo-Hertz quasi-periodic oscillations (kHz QPOs).  These kHz QPOs are seen to vary in frequency as the neutron star changes state and moves around the color-color diagram \\citep[see][for a comprehensive review]{vanderklis06}.  The kHz QPOs have frequencies that may be associated with the orbital frequency of the inner disk and there are several proposed mechanisms for their production \\citep[e.g.][]{miller98,stella99,lamb01,abramowicz03}.  By combining color-color diagrams, QPOs and iron lines, it may be possible to break the degeneracies in the models.  For instance, if they are due to orbital motions in the disk, then a change kHz QPO frequency may correspond to a change in the inner accretion disk radius.  However, in disk resonance models this would not be expected.  If the inner disk radius can be measured independently using iron lines, there is the potential to test these models.  In fact, recent observations of iron lines in 3 systems (Ser~X$-$1, GX~349+2 and 4U~1820$-$30) show a good agreement between the measured inner disk radius from the iron lines and the inferred inner disk radius from kHz QPOs \\citep[when the highest observed kHz QPO frequency is used;][]{cackett08}. It is therefore important to identify as many neutron stars with iron lines as possible, to determine whether such a method is possible.  The location on hardness-intensity (or alternatively color-color) diagrams, where iron lines are present, may also be key, and in particular how this corresponds to where kHz QPOs are seen.  Moreover, iron lines in neutron star LMXBs can provide a test of iron line models for black holes.  In several objects extremely broad lines are observed from black holes.  It has been proposed that the extreme width of the lines requires that the black holes are spinning \\citep[e.g.][]{iwasawa96,wilms01,fabian02,miller_xtej1650_02,miller04}, and thus allowing the disk to extend closer in than the inner-most stable orbit for a non-spinning black hole.  Of course, in neutron stars the solid surface of the star prevents the accretion disk from extending any further in than that.  Any iron lines in neutron star LMXBs should therefore be significantly less broad than those in spinning black holes.  Results from recent sensitive observations do find that the neutron star lines are narrower than in the proposed spinning black holes, supporting that hypothesis \\citep{bhattacharyya07,cackett08}.  As part of the {\\it Chandra} HETGS Z/Atoll Spectroscopic Survey \\citep[CHAZSS, see][for related work]{cackett1820} we have obtained observations of 6 neutron star LMXBs.  Each source in this survey was observed at 2 separate epochs using the {\\it Chandra} High Energy Transmission Grating Spectrometer (HETGS).  Simultaneously, we observed these sources with the {\\it Rossi X-ray Timing Explorer (RXTE)}.  In this paper we present the search for Fe K$\\alpha$ emission lines in these sources and we examine how the presence, or lack of iron lines corresponds to position in hardness-intensity diagrams.  In Section \\ref{sec:data} we describe the observations and data reduction.  Section \\ref{sec:results} details our analysis and results, and in Section \\ref{sec:disc} we discuss the implications of the iron lines that we find. ",
        "conclusions": "\\label{sec:disc} We present {\\it Chandra} and {\\it RXTE} observations of six neutron star low-mass X-ray binaries at two separate epochs (4U~1636$-$53, 4U~1735$-$44, 4U~1820$-$30, GX~17+2, GX~340+0 and GX~349+2).  We detect iron emission lines in the {\\it Chandra} HEG spectra of both GX~17+2 and GX~349+2 at both epochs.  No iron lines are detected in the {\\it Chandra} spectra of the other four sources, although we do not have the sensitivity to rule out that weak, broad lines are present.  Comparing the position of the sources in hardness-intensity diagrams (using the {\\it RXTE} data) we find that from these observations there appears to be no region of the diagram where the lines are preferentially observed. However, we note that we only sample a small number of states here, and more sensitive observations are needed to confirm this. Such a finding is interesting though since we expect that the hardness-intensity diagrams are related to the mass accretion rate and the accretion disk-corona interaction. For example, changes from bright, soft (disk dominated) states to fainter, hard (corona dominated) states may indicate changes in mass accretion rate and the prominence of the disk or the corona.  One might expect to see changes in the iron lines as a function of the prominence of the disk or corona.  The 3 Z sources observed are seen across most locations in the color-color and color-intensity diagrams.  GX~17+2 is seen at both the normal branch (NB)/horizontal branch (HB) transition as well as the NB/flaring branch (FB) transition. GX~340+0 is seen mainly on the NB, with part of the first observation on the FB. GX~349+2 is observed on the FB in both observations.  Thus, iron lines are detected in the {\\it Chandra} spectra on both the FB, and at the NB/HB and NB/FB transitions.  Previous {\\it BeppoSAX} observations of GX~17+2 and GX~349+2 \\citep{disalvo00,disalvo01} support our finding that iron lines are seen in no special location on the hardness-intensity (or color-color) diagram.  During the {\\it BeppoSAX} observations of GX~17+2 the source was in the horizontal and normal branches.  When the data was split based on state, an iron line with an equivalent width in the approximate range $30-40$ eV was observed in all spectra \\citep{disalvo00}.  For the {\\it BeppoSAX} GX~349+2 observations the source was observed on the flaring branch, and both non-flaring and flaring spectra showed an iron line \\citep{disalvo01}.  We do not see any iron lines in the {\\it Chandra} spectra in any of the atoll sources in our sample.  The majority of the observations are spread out over the softest part of the color-color diagram - the banana branch, although we do observe 4U~1636$-$53 in the hard (island) state.  This indicates that we do not observe strong lines in either spectrally soft or hard states in the atolls.  However, we cannot rule out that weak, broad lines are present in the spectra as has been seen with more sensitive observations of 4U~1820$-$30 \\citep{cackett08}. In fact, while we don't detect iron lines during the banana branch here, observations of other neutron star LMXBs, such as the atoll 4U~1705$-$44, do show iron lines during the banana branch \\citep{disalvo05,piraino07}, supporting that iron lines can be seen in a wide variety of spectral states.  One interesting proposition is that of using sensitive observations of iron lines to help constrain models for kHz QPOs.  Additionally one can combine the measurement of both the disk velocity (from iron lines) and orbital frequency (from kHz QPOs, assuming that is their origin) to determine the neutron star mass \\citep{cackett08}.  For this to be possible we would need to see both iron lines and kHz QPOs simultaneously.  Here we have detected both an iron line and twin kHz QPOs in the case of GX~17+2. However, the observation of GX~17+2 is not sensitive enough to detect any broad red wing of the iron line, and thus we cannot get a reliable measure of the inner accretion disk radius by modeling the iron line.  Nevertheless, this is promising for further studies using more sensitive iron line observations, such as can be obtained with {\\it Suzaka}.  There are several possibilities for the source of ionizing flux that gives rise to the iron emission lines.  It may be irradiation by a corona, by the base of a compact jet, or by the boundary layer.  These observations allow us to investigate these possibilities.  For instance, the hardest states in the hardness-intensity diagram are usually interpreted as being the state where the corona is most dominant, thus we may expect to have stronger iron lines in harder spectral states.  Or, if the ionization is too high in the hardest states, the iron lines may get stronger as the source gets softer, but would be expected to weaken or even disappear in the disk dominated states.  In addition, the irradiation may be due to the boundary layer, or the base of a compact jet.  If the irradiation is due to the base of a compact jet, then we would expect to see iron lines when the jet is present, and not when the jet disappears.  We find that the iron lines detected in GX~17+2 and GX~349+2 are both strongest when the sources are closest to the NB/FB transition (though the changes in equivalent width are not highly significant). This is similar to what was found by \\citet{disalvo00, disalvo01} for these sources. Comparing the 10-20 keV flux with the line equivalent widths, we find that in both GX~17+2 and GX~349+2 the equivalent width increases with decreasing 10-20 keV flux.  This can be explained if the gas in the disk is very highly ionized already and thus the emission is mostly due to recombination.  In such a scenario a drop in hard flux should lower the ionization leading to more recombinations and a stronger line.  The fact that we see the line energies consistent with the He-like 6.7 keV iron line is in agreement with this.  The absence of strong lines in sources with short periods (such as 4U~1820$-$30) is also consistent with such a picture, as the disks should be highly ionized in these sources.  X-ray and radio observations of neutron star LMXBs have associated jet formation and location on the color-color/hardness-intensity diagrams \\citep[e.g.,][]{penninx88,hjellming90a,hjellming90b,migliari06,migliari07}.  From a detailed study of GX~17+2, \\citet{migliari07} find that radio emission is not detected in the FB and turns on as the source transitions into the NB.  The radio flux density increases through the NB and is strongest during the HB.  We have concurrent radio observations of the sources in this study with the VLA (Cackett et al., in preparation), and thus can look for any correlations between radio flux density and the iron lines.  For GX~349+2 we do not detect the source at 5 GHz in either observation (consistent with its position on the FB), yet for GX~17+2 we have flux densities of $1.80\\pm0.05$ mJy and $0.9\\pm0.07$ mJy at 5 GHz for observations concurrent with observations 1 and 2, respectively.  Thus, in GX~17+2 we see higher radio flux densities at the NB/HB state than in the NB/FB state \\citep[consistent with the picture of][]{migliari07}.  Note though that while there is a significant change in radio flux between observations the equivalent width of the iron line in GX~17+2 does not change significantly.  In addition, given that we do not detect any radio emission from GX~349+2 yet we do see an iron line this would tend not to favor a compact radio jet as the source of ionizing flux for the line emission.  However, it is not clear whether a disk corona or the boundary layer is the more likely source of irradiation."
    },
    "0809/0809.1668_arXiv.txt": {
        "abstract": "Observing Delta Scuti stars is most important as their multi-frequency spectrum of radial pulsations provide strong constraints on the physics of the stars interior; so any new detection and observation of these stars is a valuable contribution to asteroseismology. While performing uvby-beta photoelectric photometry of some RR Lyrae stars acquired in 2005 at the Observatorio Astron\\'omico Nacional, M\\'exico, we also observed several standard stars, HD115520 among them. After the reduction this star showed indications of variability. In view of this, a new observing run was carried out in 2006 during which we were able to demonstrate its variability and its nature as a Delta Scuti star. New observations in 2007 permitted us to determine its periodic content with more accuracy. This, along with the uvby-beta photoelectric photometry allowed us to deduce its physical characteristics and pulsational modes. ",
        "introduction": "In [1] confirmed  the membership of HD 115520 to the $\\delta$ Scuti class. It had been considered as a standard star in a 2005 observing run. From the relatively large scatter shown, [2] considered it to be variable candidate. With this in mind, new data were acquired in two new nights in 2006 which established it as a $\\delta$ Scuti star. In the present paper we present new observations which were performed in 2007 with the same instrumentation over a period of four nights and which have served to determine its periodic content. The frequencies found explain the behavior of both seasons separated by more than one year. ",
        "conclusions": ""
    },
    "0809/0809.3165_arXiv.txt": {
        "abstract": "{We have performed 2.5D and 3D simulations of conical jets driven by the rotation of an ordered, large-scale magnetic field in a stratified atmosphere.  The simulations cover about three orders of magnitude in distance to capture the centrifugal acceleration as well as the evolution past the \\Alfven surface.  We find that the jets develop kink instabilities, the characteristics of which depend on the velocity profile imposed at the base of the flow. The instabilities are especially pronounced with a rigid rotation profile, which induces a shearless magnetic field. The jet's expansion appears to be limiting the  growth of \\Alfven mode instabilities. } ",
        "introduction": "Strong magnetic fields on large scales may play an essential, active role in the formation and evolution of jet-like outflows.  The general idea is that a poloidal magnetic field, embedded in a plasma and anchored e.g. in an accretion disk or a black hole, is forced into rotation at the anchor point, a toroidal field develops and the plasma is accelerated by what can be interpreted as a centrifugal force in a corotating frame \\citep{1982Blandford,1996Spruit,2000Koenigl}.  However, this magnetocentrifugal acceleration is only effective up until the \\Alfven surface, defined as the surface where the flow velocity equals the \\Alfven velocity.  Beyond this point, the magnetic field will be strongly wound-up. Such a field configuration is potentially unstable with respect to certain MHD instabilities.  MHD jets are susceptible to a variety of instabilities.  Kelvin-Helmholtz (KH) instabilities are fed by the relative kinetic energy between the jet and the ambient medium.  They can distort the jet surface only (ordinary modes) or the whole beam \\citep[e.g.][]{1991Birkinshaw}, provoking shocks, mixing with ambient material and possibly a disruption of the jet \\citep{1995Bodo,1998Bodo}.  The presence of strong magnetic fields, poloidal or toroidal, is expected to hamper the growth of KH instabilities \\citep{1992Appl,1999Keppens}.  The free energy associated with the toroidal magnetic field is responsible for another class of instabilities, which is traditionally known as current-driven (CD) and of notorious importance in controlled fusion devices \\citep[for an introduction, see, e.g.][]{1987Freidberg,1978Bateman}.  The relevance for magnetized astrophysical jets has been pointed out by \\citet{1993Eichler,1997Spruit,1998Begelman} and others.  Among CD instabilities, $m=1$ kink instabilities are the most effective.  An ideal kink mode is characterized by helical displacements of the cylindrical cross sections of a plasma column.  It is expected to grow on a dynamical time scale with respect to an \\Alfven wave crossing the unstable column. The susceptibility is strongly dependent on the magnetic pitch, a measure for the degree of wind-up.  Kink instabilities might destroy the ordered, symmetric state of a jet, leading to its disruption or, through magnetic reconnection, the associated dissipation of magnetic fields and steepening of the magnetic pressure gradient, to its acceleration \\citep{2002Drenkhahn,2006Giannios}.  Different kinds of instability can mix and interact. For example, \\citet{2002Baty} show how CD instabilities can stabilize KH vortices at the jet surface. For this work, we used conditions under which CD kink instabilities are expected to dominate (low plasma-$\\beta$, small magnetic pitch).  For a self-consistent study of kink instabilities, numerical simulations need to be carried out in 3D.  \\citet{1996Lucek} did so using a simple model in which a toroidal magnetic field configuration was allowed to expand into a uniform atmosphere. This generated a jet which was subject to kink instabilities. \\citet{2004Nakamura} performed 3D simulations of MHD jets in variously stratified atmospheres, finding that they can develop kink-like distortions in the trans-\\Alfvenic region.  Laboratory experiments of MHD jets have been performed by \\citet{2005Hsu}, confirming that the magnetic pitch plays a crucial role for the formation of kink instabilities.   \\subsection{Effects of jet expansion}  Jets from protostars, and especially AGN and microquasars, expand in width $d$ by orders of magnitude after passing through their \\Alfven radius. In an expanding flow there is no clean separation between time dependence due to instability and that due to the expansion itself, making the question of stability less well defined. Analytical studies thus tend to focus on instabilities in a cylindrical geometry, with constant diameter (such as in the ``magnetic tower'' picture of \\citealt{2003Lynden}). Expansion has strong consequences on the behavior of instabilities, however, compared with jets modeled as cylinders of constant width.  First, there is the tendency for the toroidal (azimuthal, around the jet axis) component of the magnetic field to dominate in an expanding jet. From the induction equation, the poloidal and toroidal  components of the field vary as $B_\\mathrm{p}\\sim 1/d^2$ and $B_\\varphi\\sim 1/d$ respectively (for constant jet velocity). Expansion thus causes a continual increase of the ratio $B_\\varphi/B_\\mathrm{p}$. Even when dissipation were to decrease the toroidal field at some point, the ratio increases again on further expansion. Free energy available in the toroidal field thus remains the dominant form of magnetic energy, and one may expect the question of stability and dissipation to remain relevant on all length scales. It also follows that the nonlinear development of instabilities in an expending jet is expected to be very different from the constant-diameter case.  Secondly, expansion has a strong effect even on the conditions for occurrence of instability. It has a stabilizing effect, since magnetic instabilities become ineffective when their signal speed (the \\Alfven speed) drops below the lateral expansion speed of the jet. This is discussed further below.   \\subsection{Rationale of the calculations}  The aim of the calculations reported here is to study how kink instability operates under these conditions of expansion of the jet over several orders of magnitude in width.  The degree of instability to be expected in a jet driven by a rotating magnetic field is intimately tied to the way it is collimated.  If, instead of being cylindrical, the jet has a non-vanishing opening angle $\\vartheta$, the expected incidence of instability depends on the details of the dependence of $\\vartheta$ on distance $r$.  An opening angle increasing with distance reduces instability, while for an asymptotically vanishing opening angle instability must always set in at some distance, if it was not present already from the start (see discussion in Sect.~\\ref{sec:growthexp}).  The setup in the simulations presented here produces jets in the intermediate case of an (approximately)  constant opening angle: a ``conical'' outflow. It turns out that in this case the presence of instabilities and their amplitude depends on secondary conditions such as the rotation profile imposed at the base of the flow, hence it is a good test case for the incidence of instabilities.  Since the observed jets travel over such large distances, even marginal forms of instability can become effective. An important goal of the present calculations is therefore to cover a large range in distance, about 3 orders of magnitude. This is achieved by the use of a grid adapted to the approximately conical shape of the jet.     \\subsection{Expected instability growth in expanding jets} \\label{sec:growthexp}  In the following, we estimate how the sideways expansion affects the growth of instabilities in broadening jets. In spherical coordinates $(r,\\vartheta,\\varphi)$, we assume that the jet radius (distance to the jet's central axis) is given by \\begin{equation} R = R' \\left( \\frac{r}{r'} \\right)^\\alpha \\quad \\text{with} \\quad R' = r' \\sin\\vartheta' \\label{} \\end{equation} where the prime stands for evaluation at a reference distance $r'$, for which we take a distance somewhat beyond the \\Alfven radius. The flow has then approximately reached its asymptotic speed $v_r \\approx \\const$, and the magnetic field has become predominantly azimuthal. In the absence of magnetic dissipation due to instabilities, the field strength then varies as $B_\\varphi \\propto R^{-1}$ (magnetic flux conservation) and the density as $\\rho \\propto R^{-2}$ (mass conservation). Since the growth rate $\\varGamma_\\mathrm{g}$ is expected to scale with the \\Alfven crossing rate $\\vAphi/R$, we introduce a dimensionless instability rate $\\varkappa$ of order unity: \\begin{equation} \\varGamma_\\mathrm{g} = \\varkappa \\frac{\\vAphi}{R} = \\varkappa \\frac{\\vAphi}{R'} \\left( \\frac{r}{r'} \\right)^{-\\alpha} \\label{} \\end{equation} where $\\vAphi=B_\\varphi/\\sqrt{4\\pi \\rho}$ is the azimuthal \\Alfven velocity. The expansion rate is estimated by \\begin{equation} \\varGamma_\\mathrm{e} = \\frac{\\de \\ln R}{\\de t} = \\frac{1}{R} \\frac{\\de r}{\\de t} \\frac{\\de R}{\\de r} = \\frac{\\alpha v_r}{r} . \\label{} \\end{equation} We find \\begin{equation} \\frac{\\varGamma_\\mathrm{g}}{\\varGamma_\\mathrm{e}} = \\frac{\\varkappa}{\\upsilon \\alpha \\sin \\vartheta'} \\left( \\frac{r}{r'} \\right)^{1-\\alpha} \\label{} \\end{equation} with $\\upsilon \\coloneqq v_r/\\vAphi \\approx \\const$ according to the ballistic approximations mentioned above.  Consequently, the instability growth rate dominates at some distance $r$ for a collimating jet ($\\alpha<1$). Decollimation ($\\alpha>1$), on the other hand, thwarts the growth of instabilities. A conical jet ($\\alpha = 1$) constitutes a limiting case where all depends on the combination of parameters $\\varkappa/(\\upsilon\\, \\sin \\vartheta')$, which is of order unity.  A numerical simulation is necessary to find out whether the instability or expansion prevails.  The paper is organized as follows. In Sect.~\\ref{sec:model}, we introduce the magnetocentrifugal jet model and account for the assumptions made in our simulations.  A detailed description of the numerical setup, the coordinate system and the scale-free units employed in the analysis is given in Sect.~\\ref{sec:methods}.  In Sect.~\\ref{sec:simulatedcases} we give the parameters of the simulated cases and in Sect.~\\ref{sec:results} we present the results.  There, we start by making predictions on the characteristics of instabilities by examining the relevant properties of our simulated jets. We proceed with an analysis of the instabilities that actually appeared and complete with looking for effects on the jets' dynamics. We finish with a discussion and conclusions in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have simulated magnetocentrifugally driven, conical jets over a range in distance of 1000 times the initial jet radius, in both 3D and axisymmetric 2.5D.  The calculations extend to a factor of about 5--10 beyond the \\Alfven surface.  The 3D jets developed non-axisymmetric instabilities of the kink kind.  The violence of the instabilities depends on the rotation profile applied at the base.  With a rigid rotation ($\\proptoo R$) profile, the perturbations grow to much larger amplitudes than with a Keplerian ($\\proptoo R^{-1/2}$) profile. We suspect that the reason for the differing behavior lies in the magnetic shear, defined as the variation of the magnetic pitch with distance to the axis. In the rigid rotation case, there is virtually no shear as opposed to the Keplerian case, for which the pitch increases with distance from the axis, see Fig.~\\ref{fig:pitch}.  A shear-free configuration is expected to be unstable to non-resonant modes, whereas a configuration with increasing pitch is expected to be unstable to modes with a resonant surface inside the jet \\citep{2000Appl,2000Lery}.  This fits well with what we observe in the simulations, viz. that the kink is confined inside the jet in the Keplerian case. Heuristically speaking, the differing behavior could be attributed to the fact that the outer (high $\\vartheta$) layers of the jet, which are more stable (higher magnetic pitch), damp internally arising modes in the Keplerian case.  In both cases, the longitudinal wavelength of the instabilities is $\\simm5$ times larger than the value of the magnetic pitch near the axis. The relation is qualitatively consistent with the findings of \\citet{2000Appl} for a cylindrical jet. The growth time of the instabilities is on the order of the \\Alfven crossing time.  The exact relation is difficult to determine, because the crossing time as well as the location of the resonant surface can only be estimated.  As the azimuthal magnetic field strength and with it the azimuthal \\Alfven speed decrease past the \\Alfven surface, opposing parts of the jets become causally disconnected from each other. Thus, the jet expands too fast for \\Alfven mode instabilities to grow. The effect is amplified if the jet is diverging.  Recollimation, on the other hand, should boost the growth of instabilities.  As found in other studies, the conversion of magnetic enthalpy (Poynting flux) to kinetic energy is fairly efficient, on the order 70\\%. Dissipation of magnetic fields by internal instabilities is expected to contribute additional acceleration of the flow \\citep{2002Drenkhahn}.  The calculations do not show a clear signature of this process.  It seems that either the observed region is too small, and/or the numerical dissipation of magnetic fields is too weak. Also, from a macroscopic point of view, the instabilities were not violent enough to bring together fields with an antiparallel component, as is necessary for magnetic reconnection to occur.  Moreover, most of the magnetic enthalpy was already converted in the magnetocentrifugal acceleration process. Therefore, even if there was magnetic dissipation, the effect would not be dramatic. Nevertheless, we found that the magnetic field gets significantly distorted by the instabilities.  This should facilitate magnetic field dissipation further downstream but it may be necessary to extend the calculations to larger distances to see the effect.  It is tempting to compare the instability-related structures in the simulations with structures in observed jets. The 3D jet structure in Fig.~\\ref{fig:rhot}, for example, is reminiscent of the semi-regular patterns seen in \\Halpha images taken of outflows from young stellar objects (YSO) like HH 34 \\citep{2002Reipurth}.  There are, however, other possible interpretations of the observed structure. The wiggles in YSO jets could also be the result of a precessing or orbitally moving source \\citep[and references therein]{2002Masciadri}. The symmetric nature of structures often seen in jet and counterjet \\citep[e.g. HH 212][]{2001Wiseman} suggests a modulation of the outflow speed or mass flux originating at the source of the outflow rather than an instability developing further away. The irregularities caused by the instabilities studied here are possibly more important for internal magnetic energy release inside the jet than for major observable structures like the knots and wiggles in YSO jets, though they are likely to contribute to these at some level as well."
    },
    "0809/0809.4329_arXiv.txt": {
        "abstract": "{\\em AstraLux} is the Lucky Imaging camera for the Calar Alto 2.2-m telescope and the 3.5-m NTT at La Silla. It allows nearly diffraction limited imaging in the SDSS {\\em i'} and {\\em z'} bands of objects as faint as {\\em i'}=15.5\\,{\\em mag} with minimum technical effort.  One of the ongoing {\\em AstraLux} observing programs is a binarity survey among late-type stars with spectral types K7 to M8, covering more than 1000 targets on the northern and southern hemisphere. The survey is designed to refine binarity statistics and especially the dependency of binarity fraction on spectral type. The choice of the SDSS {\\em i'} and {\\em z'} filters allows to obtain spectral type and mass estimates for resolved binaries.  With an observing efficiency of typically 6 targets per hour we expect to complete the survey in mid-2009. Selected targets will be followed up astrometrically and photometrically, contributing to the calibration of  the mass-luminosity relation at the red end of the main sequence and at visible wavelengths. ",
        "introduction": "Observations of binary stars can not only constrain stellar formation theories, which have to reproduce the observed binary frequency, but are the only way to reliably determine stellar masses by orbit fitting techniques. While binary statistics are well established for solar-like stars, this is not entirely true for the low-mass and very low-mass end of the main sequence. Here large binarity surveys can still provide substantial improvements of our knowledge of the binary fraction and its dependency on spectral type, giving indirect evidence of a possible discontinuity in the low-mass initial mass function \\cite{Thies:2007}.  Apart from that, dynamical mass estimates of low-mass stars are the key ingredient to the calibration of mass-spectral-type and mass-luminosity relations at the red end of the main sequence. While stellar evolution models reproduce near-infrared magnitudes and colours of low-mass stars quite satisfyingly nowadays, this is not the case for visible wavelengths $<$1\\,$\\mu$m.  This is the motivation for the large {\\em AstraLux} M~dwarf survey, designed to find and characterise late-spectral-type binary and multiple systems at visible wavelengths. With more than 1000 M~dwarfs targeted by the initial survey programme, this is by far the largest of its kind. The Lucky Imaging cameras {\\em AstraLux} at the Calar Alto 2.2-m telescope and {\\em AstraLux Sur} at ESO's NTT, La Silla, provide the necessary observing efficiency to conduct such a survey in relatively short time. ",
        "conclusions": ""
    },
    "0809/0809.1373.txt": {
        "abstract": "We deal with the attempts to measure  the Lense-Thirring effect with the Satellite Laser Ranging (SLR) technique applied to the existing LAGEOS and LAGEOS II terrestrial satellites and to the recently approved LARES spacecraft. According to general relativity, a central spinning body of mass $M$ and angular momentum $\\bds S$ like the Earth generates a gravitomagnetic field which induces small secular precessions of the orbit of a test particle geodesically moving around it. Extracting this signature from the data is a demanding task because of many classical orbital perturbations having the same pattern as the gravitomagnetic one, like those due to the centrifugal oblateness of the Earth which represents a major source of systematic bias. The first issue addressed here is: are the so far published evaluations of the systematic uncertainty induced by  the bad knowledge of the even zonal harmonic coefficients $J_{\\ell}$ of the multipolar expansion of the Earth's geopotential  reliable and realistic? Our answer is negative. Indeed, if the differences $\\Delta J_{\\ell}$ among the even zonals estimated in different Earth's gravity field global solutions from the dedicated GRACE mission are assumed for the uncertainties $\\delta J_{\\ell}$  instead of using their covariance sigmas $\\sigma_{J_\\ell}$, it turns out that the systematic uncertainty $\\delta\\mu$ in the Lense-Thirring test with the nodes $\\Omega$ of LAGEOS and LAGEOS II may be up to 3 to 4 times larger than in the evaluations so far published ($5-10\\%$) based on the use of the sigmas of one model at a time separately.  The second issue consists of the possibility of using a different approach in extracting the relativistic signature of interest from the LAGEOS-type data. The third issue is the possibility of reaching a realistic total accuracy of $1\\%$ with LAGEOS, LAGEOS II and LARES, which should be launched in November 2009 with a VEGA rocket.   While LAGEOS and LAGEOS II fly at altitudes of about 6000 km, LARES will be likely placed at an altitude of 1450 km. Thus, it will be sensitive to much more even zonals than LAGEOS and LAGEOS II. Their corrupting impact has been evaluated with the standard Kaula's approach up to degree $\\ell=60$ by using $\\Delta J_{\\ell}$ and $\\sigma_{J_{\\ell}}$; it turns out that it may be as large as some tens percent. The different orbit of LARES may also have some consequences on the non-gravitational orbital perturbations affecting it which might further degrade the obtainable accuracy. %The cm-level accuracy in reconstructing in the Root-Mean-Square (RMS) sense the cross-track component of the orbit of the LAGEOS-type satellites, %affected by the gravitomagnetic force, yields a $\\approx 7\\%$ error over each orbital arc 7 days long. ",
        "introduction": "In the weak-field and slow motion approximation,  the Einstein field equations of general relativity get linearized resembling to the Maxwellian equations of electromagntism. As a consequence, a gravitomagnetic field $\\bds B_{\\rm g}$, induced by the off-diagonal components $g_{0i}, i=1,2,3$ of the space-time metric tensor related to the mass-energy currents of the source of the gravitational field, arises \\citep{MashNOVA}. The gravitomagnetic field affects orbiting test particles, precessing gyroscopes, moving clocks and atoms and propagating electromagnetic waves \\citep{Rug,Scia04}. Perhaps, the most famous gravitomagnetic effects are the precession  of the axis of a gyroscope \\citep{Pugh,Schi} and the Lense-Thirring\\footnote{According to an interesting historical analysis recently performed by \\citet{Pfi07}, it would be more correct to speak about an Einstein-Thirring-Lense effect.} precessions \\citep{LT} of the orbit of a test particle, both occurring in the field of a central slowly rotating mass like, e.g., our planet.   Direct, undisputable measurements of such  fundamental predictions of general relativity are not yet available.   The measurement of the gyroscope precession in the Earth's gravitational field has been the goal of the dedicated space-based\\footnote{See on the WEB http://einstein.stanford.edu/} GP-B mission \\citep{Eve, GPB} launched in 2004 and carrying onboard four superconducting gyroscopes; its data analysis is still ongoing. The target accuracy was originally $1\\%$, but it is still unclear if the GP-B team will succeed in reaching such a goal because of some unmodelled effects affecting the gyroscopes: 1) a time variation in the polhode motion of the gyroscopes and 2) very large classical misalignment torques on the gyroscopes.  In this Chapter we  will focus on the attempts to measure of the Lense-Thirring effect in the gravitational field of the Earth; for Mars and the Sun see \\citep{LTmars,LTkrogh,LTreplytokrogh} and \\citep{LTsun}, respectively. Far from a localized rotating body  with angular momentum ${\\bds S}$  the gravitomagnetic field can be written as \\eqi \\bds B_{\\rm g} = -\\rp{G}{c r^3}\\left[\\bds S -3\\left(\\bds S\\bds\\cdot{\\hat{r}}\\right){\\hat{r}}\\right],\\eqf   where $G$ is the Newtonian gravitational constant and $c$ is the speed of light in vacuum. It acts on a test particle orbiting with a velocity $\\bds v$ with the non-central acceleration \\citep{Sof} \\eqi \\bds A_{\\rm LT} = -\\rp{2}{c}{\\bds v}\\bds\\times\\bds B_{\\rm g}\\eqf which induces secular precessions of the longitude of the ascending node $\\Omega$ \\begin{equation}\\dot\\Omega_{\\rm LT} = \\rp{2 G S}{c^2 a^3 (1-e^2)^{3/2}},\\lb{let}\\end{equation} and the argument of pericentre $\\omega$ \\begin{equation}\\dot\\omega_{\\rm LT} = -\\rp{6 G S\\cos i}{c^2 a^3 (1-e^2)^{3/2}},\\lb{LT_o}\\end{equation} of the orbit of  a test particle. In \\rfr{let} and \\rfr{LT_o} $a$ and $e$ are the semimajor axis and the eccentricity, respectively, of the test particle's orbit and $i$ is its inclination to the central body's equator. The semimajor axis $a$ fixes the size of the ellipse, while its shape is determined by the eccentricity $0\\leq e<1$; an orbit with $e=0$ is a circle. The angles $\\Omega$ and $\\omega$ establish the orientation of the orbit in the inertial space and in the orbital plane, respectively. $\\Omega$, $\\omega$ and $i$ can be viewed as the three Euler angles which determine the orientation of a rigid body with respect to an inertial frame. In Figure \\ref{plo} we illustrate the geometry of a Keplerian orbit. % % % % % \\begin{figure} \\caption{Keplerian orbit. The longitude of the ascending node $\\Omega$ is counted from a reference X direction  in the equator of the central body, assumed as reference plane $\\{{\\rm X,Y}\\}$,  to the line of the nodes which is the intersection of the orbital plane with the equatorial plane of the central body. It has mass $M$ and proper angular momentum $\\bds S$. The argument of pericentre $\\omega$ is an angle in the orbital plane counted from the line of the nodes to the location of the pericentre, here marked with $\\Pi$. The time-dependent position of the moving test particle of mass $m$ is given by the true anomaly $f$, counted anticlockwise from the pericentre's position. The inclination between the orbital and the equatorial planes is $i$. Thus, $\\Omega,\\omega, i$ can be viewed as the three (constant) Euler angles fixing the configuration of a rigid body, i.e. the orbit which in the unperturbed Keplerian case does  change neither its shape nor its size, in the inertial $\\{{\\rm X,Y,Z}\\}$ space. Courtesy by H.I.M. Lichtenegger, IWF, Graz.} \\label{plo} \\includegraphics{orbit_spin.eps} \\end{figure} % % % % %    In this Chapter we will critically discuss the following  topics % \\begin{itemize} \\item Section \\ref{grav}. The realistic evaluation  of the total accuracy in the test performed in recent years with the existing Earth's artificial satellites LAGEOS and LAGEOS II \\citep{Ciu04,Ciu06,Ries08}. LAGEOS was put into orbit in 1976, followed by its twin LAGEOS II in 1992; they are passive, spherical spacecraft entirely covered by retroreflectors which allow for their accurate tracking through laser pulses sent from Earth-based ground stations according to the Satellite Laser Ranging (SLR) technique \\citep{slr}. They orbit at altitudes of about 6000 km ($a_{\\rm LAGEOS} =12270$ km, $a_{\\rm LAGEOS\\ II}=12163$ km) in nearly circular paths ($e_{\\rm LAGEOS}=0.0045$, $e_{\\rm LAGEOS\\ II}=0.014$) inclined by 110 deg and 52.65 deg, respectively, to the Earth's equator. The Lense-Thirring effect for their nodes amounts to about 30 milliarcseconds per year (mas yr$^{-1}$) which correspond to about 1.7 m yr$^{-1}$ in the cross-track direction\\footnote{A perturbing acceleration like $\\bds A_{\\rm LT}$ is customarily projected onto the radial $\\bds{\\hat r}$, transverse $\\bds{\\hat \\tau}$ and cross-track $\\bds{\\hat \\nu}$ directions of an orthogonal  frame comoving with the satellite \\citep{Sof}; it turns out that the Lense-Thirring node precession affects the cross-track component of the orbit according to $\\Delta\\nu_{\\rm LT} \\approx a\\sin i\\Delta\\Omega_{\\rm LT}$ (eq. (A65), p. 6233 in \\citep{Cri}).} at the LAGEOS altitudes.  The idea of measuring the Lense-Thirring node rate with the just launched LAGEOS satellite, along with the other SLR targets orbiting at that time, was put forth by \\citet{Cug78}. Tests have started to be effectively performed later by using the LAGEOS and LAGEOS II satellites \\citep{tanti}, according to a strategy  by \\citet{Ciu96} involving the use of a suitable linear combination of the nodes $\\Omega$ of both satellites and the perigee $\\omega$ of LAGEOS II. This was done to reduce the impact of the most relevant source of systematic bias, i.e. the mismodelling in the even ($\\ell=2,4,6\\ldots$) zonal ($m=0$) harmonic coefficients $J_{\\ell}$ of the multipolar expansion of the Newtonian part of the terrestrial gravitational potential due to the diurnal rotation (they induce secular precessions on the node and perigee of a terrestrial satellite much larger than the gravitomagnetic ones. The $J_{\\ell}$ coefficients cannot be theoretically computed but must be estimated by suitably processing long data sets from the dedicated satellites like CHAMP and GRACE; see Section \\ref{grav}): the three-elements combination used allowed for removing the uncertainties in $J_2$ and $J_4$. In \\citep{Ciu98} a $\\approx 20\\%$ test was reported by using the\\footnote{Contrary to the subsequent models based on the dedicated satellites CHAMP (http://www-app2.gfz-potsdam.de/pb1/op/champ/index$\\_$CHAMP.html) and GRACE (http://www-app2.gfz-potsdam.de/pb1/op/grace/index$\\_$GRACE.html), EGM96 relies upon multidecadal tracking of SLR data of a constellation of geodetic satellites including LAGEOS and LAGEOS II as well; thus the possibility of a sort of $a-priori$ `imprinting' of the Lense-Thirring effect itself, not solved-for in EGM96, cannot be neglected.} EGM96 Earth gravity model \\citep{Lem98}; subsequent detailed analyses showed that such an evaluation of the total error budget was overly optimistic in view of the likely unreliable computation of the total bias due to the even zonals \\citep{Ior03,Ries03a,Ries03b}.  An analogous, huge underestimation turned out to hold also for the effect of the non-gravitational perturbations \\citep{Mil87} like the direct solar radiation pressure, the Earth's albedo, various subtle thermal effects depending on  the physical properties of the satellites' surfaces and their rotational state \\citep{Inv94,Ves99,Luc01,Luc02,Luc03,Luc04,Lucetal04,Ries03a}, which the perigees  of LAGEOS-like satellites are particularly sensitive to. As a result, the realistic total error budget in the test reported in \\citep{Ciu98} might be as large as $60-90\\%$ or (by considering EGM96 only) even more.  The observable used by \\citet{Ciu04} with the GRACE-only EIGEN-GRACE02S model \\citep{eigengrace02s} and by \\citet{Ries08} with other more recent Earth gravity models was the following linear combination\\footnote{See also \\citep{Pav02,Ries03a,Ries03b}.} of the nodes of LAGEOS and LAGEOS II, explicitly computed by \\citet{IorMor} following the approach put forth by \\citet{Ciu96} \\begin{equation} f=\\dot\\Omega^{\\rm LAGEOS}+ c_1\\dot\\Omega^{\\rm LAGEOS\\ II }, \\lb{combi}\\end{equation} where \\begin{equation} c_1\\equiv-\\rp{\\dot\\Omega^{\\rm LAGEOS}_{.2}}{\\dot\\Omega^{\\rm LAGEOS\\ II }_{.2}}=-\\rp{\\cos i_{\\rm LAGEOS}}{\\cos i_{\\rm LAGEOS\\ II}}\\left(\\rp{1-e^2_{\\rm LAGEOS\\ II}}{1-e^2_{\\rm LAGEOS}}\\right)^2\\left(\\rp{a_{\\rm LAGEOS\\ II}}{a_{\\rm LAGEOS}}\\right)^{7/2}.\\lb{coff}\\end{equation} The coefficients $\\dot\\Omega_{.\\ell}$ of the aliasing classical node precessions \\citep{Kau} $\\dot\\Omega_{\\rm class}=\\sum_{\\ell}\\dot\\Omega_{.\\ell}J_{\\ell}$ induced by the even zonals  have been analytically worked out up to $\\ell=20$ in, e.g. \\citep{Ior03}; they yield $c_1=0.544$. The Lense-Thirring signature of \\rfr{combi} amounts to 47.8 mas yr$^{-1}$. The combination  \\rfr{combi} allows, by construction, to remove the aliasing effects due to the static and time-varying parts of the first even zonal $J_2$. The nominal (i.e. computed with the estimated values of $J_{\\ell}$, $\\ell=4,6...$) bias due to the remaining higher degree even zonals would amount to  about $10^5$ mas yr$^{-1}$; the need of a careful and reliable modeling of such an important source of systematic bias is, thus, quite apparent. Conversely, the nodes of the LAGEOS-type spacecraft are directly affected by the non-gravitational accelerations at a $\\approx 1\\%$ level of the Lense-Thirring effect \\citep{Luc01,Luc02,Luc03,Luc04,Lucetal04}. For a comprehensive, up-to-date overview of the numerous and subtle issues concerning the measurement of the Lense-Thirring effect see, e.g., \\citep{IorNOVA}. % % % \\item Section \\ref{approach}. Another approach which could be followed in extracting the Lense-Thirring effect from the data of the LAGEOS-type satellites. % % \\item Section \\ref{larez}. The possibility that the LARES mission, recently approved by the Italian Space Agency (ASI), will be able to measure the Lense-Thirring node precession with an accuracy of the order of $1\\%$.  In \\citep{vpe76a,vpe76b}  it was proposed to measure the Lense-Thirring precession of the nodes $\\Omega$ of a pair of counter-orbiting spacecraft to be launched in terrestrial polar orbits and endowed with drag-free apparatus. A somewhat equivalent, cheaper version of such an idea was put forth in 1986 by \\citet{Ciu86} who proposed to launch a passive, geodetic satellite in an orbit identical to that of LAGEOS  apart from the orbital planes which should have been displaced by 180 deg apart. The measurable quantity was, in the case of the proposal by \\citet{Ciu86}, the sum of the nodes of LAGEOS and of the new spacecraft, later named LAGEOS III, LARES, WEBER-SAT, in order to cancel to a high level of accuracy the corrupting effect of the multipoles of the Newtonian part of the terrestrial gravitational potential which represent the major source of systematic error (see Section \\ref{grav}). Although extensively studied by various groups \\citep{CSR,LARES}, such an idea was not implemented for many years. In \\citep{Ioretal02} it was proposed to include also the data from LAGEOS II by using a different observable.  Such an approach was proven in \\citep{IorNA} to be potentially useful in making the constraints on the orbital configuration of the new SLR satellite less stringent than it was originally required in view of the recent improvements in our knowledge of the classical part of the terrestrial gravitational potential due to the dedicated CHAMP and, especially, GRACE  missions. % For recent updates of the LARES/WEBER-SAT %mission, including recently added additional goals in fundamental physics and related criticisms, see %\\citep{LucPao01,Ior02,Ciu04b,Ciu06b,IorJCAP,Ior07c,Ior07d}.  Since reaching high altitudes and minimizing the unavoidable orbital injection errors is expensive, it was explored the possibility of discarding LAGEOS and LAGEOS II using a low-altitude, nearly polar orbit for LARES \\citep{LucPao01,Ciu06b}, but in \\citep{Ior02,Ior07c} it was proven that such alternative approaches are not feasible. It was also suggested that LARES would be able to probe alternative theories of gravity \\citep{Ciu04b}, but also in this case it turned out to be impossible \\citep{IorJCAP,Ior07d}.  The stalemate came to an end when ASI recently made the following official announcement  (http://www.asi.it/SiteEN/MotorSearchFullText.aspx?keyw=LARES): ``On February 8, the ASI board approved funding for the LARES mission, that will be launched with VEGA\u0092s maiden flight before the end of 2008. LARES is a passive satellite with laser mirrors, and will be used to measure the Lense-Thirring effect.''  The italian version of the announcement yields some more information specifying that LARES, designed in collaboration with National Institute of Nuclear Physics (INFN), is currently under construction by Carlo Gavazzi Space SpA; its Principal Investigator (PI) is I. Ciufolini and its scientific goal is to measure at a $1\\%$ level the  Lense-Thirring effect in the gravitational field of the Earth. Concerning the orbital configuration of LARES, %the ASI website says about VEGA that %(http://www.asi.it/SiteEN/ContentSite.aspx?Area=Accesso+allo+spazio): %``[...] VEGA can place a 15.000 kg satellite on a low polar orbit, 700 km from the Earth. By lowering the orbit inclination it can launch heavier payloads, whereas diminishing the payload mass it can achieve higher orbits. [...]'' In one of the latest communication to INFN, Rome, 30 January 2008, \\citep{INFN} writes that LARES will be launched with a semimajor axis of approximately 7600 km and an inclination between 60 and 80 deg. More precise information can be retrieved in Section 5.1, pag 9 of the document Educational Payload on the Vega Maiden Flight Call For CubeSat Proposals, European Space Agency, Issue 1 11 February 2008, downloadable at http://esamultimedia.esa.int/docs/LEX-EC/CubeSat$\\%$20CFP$\\%$20issue$\\%$201.pdf. It is  written there that LARES will be launched into a circular orbit with altitude $h=1200$ km, corresponding to a semimajor axis $a_{\\rm LARES}=7578$ km, and inclination $i=71$ deg to the Earth's equator. Latest information\\footnote{See on the WEB http://www.esa.int/esapub/bulletin/bulletin135/bul135f$\\_$bianchi.pdf.} point towards a launch at the end of 2009 with a VEGA rocket in a circular orbit inclined by 71 deg to the Earth's equator at an altitude of\\footnote{I thank Dr. D. Barbagallo (ESRIN) for having kindly provided me with the latest details of the orbital configuration of LARES.} 1450 km corresponding to a semimajor axis of $a_{\\rm LR}=7828$ km. More or less the same has been reported by \\citet{INFN2} to the INFN in Villa Mondragone, 3 October 2008. \\end{itemize} ",
        "conclusions": "The so far published evaluations of the total systematic uncertainty induced by the even zonal harmonics of the geopotential in the Lense-Thirring test with the combined nodes of the SLR LAGEOS and LAGEOS II satellites  are  likely optimistic. Indeed, they are all based on the use of the covariance sigmas, more or less reliably calibrated, of the covariance matrices of various Earth gravity model solutions used one at a time separately in such a way that the model X yields an error of $x\\%$, the model Y yields an error $y\\%$, etc. Instead,  comparing the estimated values of the even zonals for different pairs of models   allows for a more conservative evaluation of the real uncertainties in our knowledge of the static part of the geopotential. As a consequence, the uncertainty in the Lense-Thirring signal is about $3-4$ times larger than the figures so far claimed ($5-10\\%$), amounting to various tens percent ($37\\%$ for the pair EIGEN-GRACE02S and ITG-GRACE03s, about $25-30\\%$ for the other most recent GRACE-based solutions).  Concerning the extraction of the Lense-Thirring signal from the data of the LAGEOS-type satellites, different approaches with respect to the one followed so far should be implemented in order to do something really new. For instance, the gravitomagnetic force should be explicitly included in the dynamical force models of the LAGEOS satellites and an ad-hoc parameter should be estimated in the least-square sense in addition to those determined so far without modelling the Lense-Thirring effect. Moreover, also the variation of the values of all the other estimated parameters with and without the gravitomagnetic force modelled should be inspected along with their mutual correlations.  Applying the strategy of the difference of the estimated even zonals to the ongoing LARES mission shows that reaching a $1\\%$ measurement of the Lense-Thirring effect with LAGEOS, LAGEOS II and LARES maybe difficult. Indeed, since LARES will orbit at a lower altitude with respect to the LAGEOS satellites, more even zonal harmonics are to be taken into account. Assessing realistically their impact is neither easy nor unambiguous. Straightforward calculations up to degree $\\ell = 60$ with the standard Kaula's approach yield errors as large as some tens percent; the same holds if the sigmas of the covariance matrices of several global Earth's gravity models are used. Such an important point certainly deserves great attention. Another issue which may potentially degrade the expected accuracy is the impact of some non-gravitational perturbations which would have a non-negligible effect on LARES  because of its lower orbit. In particular, the secular decrease of the inclination of the new spacecraft due to the neutral and charged atmospheric drag induces an indirect bias in the node precessions by the even zonals which, in the case of LARES, should be of the order of $\\approx 3-9\\%$ yr$^{-1}$. %The gravitomagnetic force affects the cross-track component of the orbit of a satellite; the accuracy in reconstructing it, in the RMS sense, for %the LAGEOS-type satellites, assumed to be of the order of 1 cm over a typical orbital arc 7 days long, would yield a $\\approx 7\\%$ error in the %Lense-Thirring effect per arc."
    },
    "0809/0809.3215_arXiv.txt": {
        "abstract": "We discuss some cosmological implications of extensions of the Standard Model with hidden sector scalars coupled to the Higgs boson. We put special emphasis on the conformal case, in which the electroweak symmetry is broken radiatively with a Higgs mass above the experimental limit. Our refined analysis of the electroweak phase transition in this kind of models strengthens the prediction of a strongly first-order phase transition as required by electroweak baryogenesis. We further study gravitational wave production and the possibility of low-scale inflation as well as a viable dark matter candidate. ",
        "introduction": "The Standard Model of particle physics (SM) is nowadays considered to be an effective theory valid only up to a certain physical cutoff scale. Even though there exist a large variety of extensions of the SM, models with a hidden sector have recently attracted some attention. We will consider models with additional scalar fields that might transform non-trivially under a hidden gauge group but which are singlets under the SM gauge group. The only renormalizable interaction of such scalars with the SM occurs via the Higgs sector, which in this case serves as a portal to the hidden sector~\\cite{hidden_sec}.  In this paper we are concerned with some of the possible cosmological implications of hidden sector extensions of the SM. This is a continuation of the study of the electroweak breaking and phase transition presented in Ref.~\\cite{Espinosa:2007qk} and we will provide some technical details that were omitted there. In addition we will present an analysis of other cosmological implications, namely gravitational wave production and dark matter abundance. We also comment on the possibility of low-scale inflation and present a calculation of the bubble wall velocity in case of a first-order electroweak phase transition. As in Ref.~\\cite{Espinosa:2007qk} we pay special attention to the classically conformal case which, for a strong coupling between the hidden sector scalars and the Higgs field, can be consistent with the mass bounds on the Higgs particle.  The paper is organized as follows. In Section~\\ref{sec_model}, the model is presented, both at zero and finite temperature.  In Section~\\ref{sec_pheno} the cosmological implications of the model mentioned above are discussed and we conclude in Section~\\ref{sec_concl}. ",
        "conclusions": "}  We have studied several cosmological implications of Standard Model extensions with hidden-sector scalars. In particular, we strengthen the results of \\cite{Espinosa:2007qk} finding that in models with a moderate number of hidden-sector scalars, $N_S \\approx 12$, the electroweak phase transition is generically of first-order as long as the Higgs mass is not much larger than the electroweak scale and the coupling to the hidden sector is substantial, $\\zeta \\gtrsim 0.9$. An interesting feature of the model is that this property persists even if the theory is classically conformal invariant and the electroweak scale is induced by dimensional transmutation. This was already emphasized in Ref.~\\cite{Espinosa:2007qk}. We find that the phase transition is in a large portion of the parameter space strong enough to suppress the sphaleron process after the phase transition, $\\phi/T \\gtrsim 1.0$ as required by electroweak baryogenesis. Besides, we find that sizable production of gravitational radiation requires a tuning of the parameters at the percent level.  Besides a strong first-order phase transition, viable electroweak baryogenesis requires sizable CP violation. Electroweak baryogenesis in non-SUSY models typically utilizes a Higgs VEV that has a changing complex phase during the phase transition. One useful ingredient hence seems to be to complexify the present scalars and to allow for scalar VEVs, but still this would not induce a change in the complex phase of the Higgs VEV such that the introduction of a second Higgs doublet seems unavoidable. Alternatively, one can introduce an additional source of CP violation in the quark sector (see e.g. Ref.~\\cite{Bodeker:2004ws}) but undoubtedly CP violation arising from the hidden sector would be much more appealing in our model.  Concerning dark matter, we find that the abundance required by the concordance model can be provided by hidden-sector scalars in two different regimes. In the first, the hidden-sector scalars have moderate couplings but large masses $M_S \\gtrsim 1$ TeV. In the second, the hidden-sector scalars are rather light, $M_S \\lesssim M_W$. In this case, the scalars cannot annihilate into W-bosons, which greatly enhances the dark matter abundance. Notice that this scenario is compatible with scalars that obtain their mass solely by electroweak symmetry breaking. Nevertheless, neither type of scalar can contribute significantly to the strength of the phase transition, such that a viable dark matter candidate cannot substantially improve the prospects of electroweak baryogenesis compared to the SM. Hence, a simultaneous solution of the dark matter and baryogenesis problems of the Standard Model close to electroweak scales either requires a large number of scalars (in which case we found $N_S \\gtrsim 50$), or several types of scalars in the hidden sector with non-uniform masses and/or couplings to the Higgs sector."
    },
    "0809/0809.3745_arXiv.txt": {
        "abstract": "{The spectra of the planetary nebulae \\object{NGC\\,3242 and NGC\\,6369} are reanalysed using spectral measurements made in the mid-infrared with the {\\em Spitzer Space Telescope} and the {\\em Infrared Space Observatory (ISO)}. The aim is to determine the chemical composition of these objects.  We also make use of {\\em International Ultraviolet Explorer (IUE)} and ground based spectra. These elliptical PNe are interesting because they are well-studied, nearby, bright objects and therefore allow a reasonably complete comparison of this type of nebulae. Abundances determined from the mid-infrared lines, which are insensitive to electron temperature, are used as the basis for the determination of the composition, which are found to differ somewhat from earlier results. The abundances found, especially the low value of helium and oxygen, indicate that the central star was originally of rather low mass. The abundance of phosphorus has been determined for the first time in NGC\\,3242. The electron temperature in both of these nebulae is roughly constant unlike NGC\\,6302 and NGC\\,2392 where a strong temperature gradient is found.  The temperature of the central star is discussed for both nebulae. Finally a comparison of the element abundances in these nebulae with the solar abundance is made. The low abundance of Fe and P is noted and it is suggested that these elements are an important constituent of the nebular dust. } ",
        "introduction": "\\object{NGC\\,3242} (PK 261.0+32.0) is a bright planetary nebula with a low radial velocity and is located at a rather high galactic latitude. A photograph of the nebula is shown in Fig.\\,1. The nebula has a bright inner ellipsoidal shell of about 28\\arcsec x 20\\arcsec. This is surrounded by an almost spherical halo of 46\\arcsec x 40\\arcsec with considerably lower emission. Both the inner and outer regions contain little structure.  The nebula is located about 32 degrees above the galactic plane and has little or no extinction. Because of its brightness it is clear that it is a nearby nebula. An expansion distance is known for this nebula to be about 0.5 kpc (Terzian, \\cite{terz}; Mellema, \\cite{mell}).  The nebula has a rather bright central star (V=12.43) which has been studied by several authors. Pauldrach et al. (\\cite{paul}) have compared the stellar spectrum with a model atmosphere and conclude that the star has an effective temperature T$_{eff}$=75\\,000 K which is in between the hydrogen Zanstra temperature of about 59\\,000 K, and somewhat less than the ionized helium Zanstra temperature which is close to 91\\,000 K (Gathier \\& Pottasch, \\cite{gath}).  Tinkler \\& Lamers (\\cite{tinkler}) prefer to assign a higher effective temperature of T$_{eff}$=94\\,000 K. to the star based on the high ionized helium Zanstra temperature. We shall discuss the stellar temperature later in the paper.  NGC\\,6369 (PK 002.4+05.8) is located in the direction of the galactic bulge but it is undoubtably a nearby PN for two reasons. First, it is a rather large nebula, with a diameter of about 32\\arcsec. Second, it is the third brightest PN in the sky, at least at radio frequencies. It cannot however be very close because the extinction is quite high and its radial velocity is also high (-101 km/s). The nebula can be described as a bright ring with an outer diameter of 32.4\\arcsec and an inner diameter of about 13\\arcsec. There is faint emission both to the east and the west of the main structure. The nebula is shown in Fig.\\,2. The central star is clearly visible. It is classified as spectral type WC4 and has a magnitude V=15.91 (Gathier \\& Pottasch \\cite{gath}). Its hydrogen Zanstra temperature is about 70\\,000 K and ionized helium Zanstra temperature is 106\\,000 K (Monteiro et al. \\cite{mont}). Judging from these temperatures it would appear that the central star of NGC6369 is somewhat hotter than that of NGC3242. Yet as we shall see, the HeII lines are considerably stronger in NGC3242. This will be discussed further in Sect.6 on the basis of the excitation of other elements found in the nebulae.  The purpose of this paper is to study the element abundances in these nebula with the help of mid-infrared spectra, in the hope that the chemical abundances will shed some light on the evolution of these nebulae. In recent years abundances in NGC\\,3242 have been studied by Tsamis et al. (\\cite{tsamis}), by Barker (\\cite{barker}), by Henry et al.  (\\cite{henry}), by Aller \\& Czyzak (\\cite{aller}) and by Krabbe \\& Copetti (\\cite{krab}). For NGC\\,6369 abundances have been studied by Aller \\& Keyes (\\cite{ allerkey}), Pena et al. (\\cite{pena}) and Monteiro et al. (\\cite{mont}). All of these groups use optical nebular spectra (taken by themselves) and in the case of NGC\\,3242 ultraviolet {\\em IUE} spectra were also used. For NGC\\,6369 ultraviolet {\\em IUE} spectra were taken but because of the very high extinction in the direction of this nebula they are underexposed and unusable.  Barker (\\cite{barker}) has measurements of NGC\\,3242 taken at five different positions in the nebula. He uses low dispersion {\\em IUE} measurements taken with the small aperture (3\\arcsec~diameter) which are very noisy.  Henry et al. (\\cite{henry}) use low dispersion {\\em IUE} measurements taken with the large aperture (10\\arcsec x 23\\arcsec).  We have measured the spectrum of both NGC\\,3242 and NGC\\,6369 in the mid-infrared with the IRS spectrograph of the {\\em Spitzer Space Telescope} (Werner et al. \\cite{werner}) and with the ISO spectrograph. The use of the mid-infrared spectrum permits a more accurate determination of the abundances.  The reasons for this have been discussed in earlier studies (e.g. see Pottasch \\& Beintema \\cite{pott1}; Pottasch et al.  \\cite{pott2}, \\cite{pott4}; Bernard Salas et al. \\cite{bernard}), and can be summarized as follows.  1) The intensity of the infrared lines is not very sensitive to the electron temperature nor to possible extinction effects. 2) Use of the infrared line intensities enable a more accurate determination of the electron temperature for use with the visual and ultraviolet lines. 3) The number of observed ionization stages is doubled.  This paper is structured as follows. First the Spitzer and {\\em ISO} spectra of both nebulae are presented and discussed (in Sect.\\,2). Then the intrinsic H$\\beta$ flux is determined using both the measurements of the infrared hydrogen lines and the radio continuum flux density (Sect.\\,3). The visible spectrum of the nebula is presented in Sect. 4 together with a new reduction of the ultraviolet ({\\em IUE}) spectrum of NGC\\,3242 This is followed by a discussion of the nebular electron temperature and density and the chemical composition of both nebulae (Sect.\\,4). A comparison of the resulting abundances with those in the literature is given in Sect.\\,5. In Sect.\\,6 a possible evolution of these nebulae is compared to other PNe. In Sect.\\,7 the central star is discussed especially in relation to the nebular spectrum.  In Sect.\\,8 a general discussion and concluding remarks are given. ",
        "conclusions": "The nebular abundances of nine elements have been determined for the PN NGC\\,3242; abundances of the same elements, except for carbon, have been found for NGC\\,6369. Both nebulae are classified as having an elliptical shape. The abundances found for both nebulae are very similar. For helium, oxygen and neon they are essentially the same as in the Sun.  This is consistent with the expectation that these PNe are rather local. This is because the gradient in the stellar abundances as a function of distance from the galactic center (e.g. see Pottasch \\& Bernard-Salas \\cite{pbs}) would lead to the expectation that a different abundance prevails when the position of these PNe is non-local.   The state of evolution of these nebulae is considered in conjunction with five other local PNe with elliptical shape. A comparison is made between the observed abundances and those expected using the stellar evolution models calculated by Karakas (\\cite{karakas}). This comparison leads to the conclusion that NGC\\,3242 must be a descendent of a low mass star, probably between 1M\\smallsun~and 1.5M\\smallsun. NGC\\,6369 probably has a slightly higher initial mass, but this conclusion is less certain because the abundance of carbon is unknown.  The stellar temperature and luminosity is considered with the help of the Zanstra method, the Energy Balance method and the observed nebular excitation.  For NGC\\,3242 we find T$_s$=80,000 K and L=730L\\smallsun~and for NGC\\,6369 T$_s$=70,000 K and L=2400L\\smallsun. The temperatures are in quite good agreement with what is given in the literature. The luminosities, which are distance dependent, are considerably lower than given by the theories of Kudritzki et al. (\\cite{kud}) or Pauldrach et al. (\\cite{paul}). This has already been discussed by Napiwotzki (\\cite{napi}) in a general way and by Surendiranath \\& Pottasch (\\cite{nath}) for the case of NGC\\,6826. This is evidence that all of the nebulae listed in Table 15 have central stars of quite low luminosity, probably between 500L\\smallsun~and 2000L\\smallsun.  The abundances of the other elements, S, Cl, P and Fe are all lower than solar. For the first two of these elements the difference with respect to the sun is a factor of two to three. In this case we cannot distinguish between the possibility that this is caused by a poor determination of the solar abundance or by the possibility that important amounts of these elements have condensed in the form of dust. For the remaining elements, P and Fe, the underabundance with respect to the Sun is larger; it is likely that both of these elements have formed molecules which have condensed in the form of dust. This is probably true for all the PNe listed in Table 15."
    },
    "0809/0809.1159.txt": {
        "abstract": "We present detailed observations of a $z\\sim1.99$ cluster of submillimeter galaxies (SMGs), discovered as the strongest redshift spike in our entire survey of $\\sim$100 SMGs across 800 square arcmin. % survey over 8 fields on the sky -- It is the largest blank-field SMG concentration currently known %from all surveys and has $<0.01$\\% chance of being drawn from the underlying selection function for SMGs. We have compared UV observations of galaxies at this redshift, where we find a much less dramatic overdensity, having an 11\\% chance of being drawn from its selection function. We use this $z\\sim1.99$ overdensity to compare the biasing of UV- and submm-selected galaxies, and test whether SMGs could reside in less overdense environments, with their apparent clustering signal being dominated by highly active merger periods in modest mass structures. % We discuss the probable mechanisms for the apparently different bias we see at the two wavelengths. %, and argue for a modest mass %galaxy cluster seen at a particularly formative period in its evolution. This impressively active formation phase in a low mass cluster is not something seen in simulations, although we propose a toy model using {\\it merger bias} which could account for the bias seen in the SMGs. % While enhanced buildup of stellar mass appears characteristic of other high$-z$ galaxy clusters, neither the UV- nor submm-galaxies in this structure exhibit larger stellar masses than their field galaxy counterparts (although the excess of SMGs in the structure represents a larger volume-averaged stellar mass than the field). Our findings have strong implications for future surveys for high-z galaxies at long wavelengths such as SCUBA2 and {\\it Herschel}. %\\footnote{SCUBA2, Herschel, CCAT, ALMA}. We suggest that since these surveys will select galaxies during their episodes of peak starbursts, they could probe a much wider range of environments than just the progenitors of rich clusters, revealing more completely the key events and stages in galaxy formation and assembly.  %surveys for high$-z$ clustered galaxies at long wavelengths (SCUBA2, Herschel), where one is biased to selecting the most luminous episode in a given galaxy regardless of its timescale or environment. % ",
        "introduction": "\\label{txt:intro}  Submillimeter galaxies (SMGs) have provided an efficient probe of ultraluminous ($>10^{12}$~L$_\\odot$) star formation activity in the distant Universe (e.g, Smail, Ivison \\& Blain 1997; Blain et al.\\ 2002; Borys et al.\\ 2003; Coppin et al.\\ 2008; Knudsen et al.\\ 2008), with the bright submm detection (rest frame far-infrared -- FIR) providing an unambiguous signal of large dust masses heated predominantly by young stars rather than AGN (e.g., Almaini et al.\\ 1999; Fabian et al.\\ 2000; Alexander et al.\\ 2005; Menendez-Delmetre et al.\\ 2007; Pope et al.\\ 2008). % Piecing together the detailed properties of SMGs has been difficult because of the inherent dust obscuration to the regions of the largest bolometric output. High spatial resolution radio observations using the MERLIN interferometer (Chapman et al.\\ 2004a; Biggs \\& Ivison 2008) provided early indications that the dust and gas in SMGs was more compact than typical galaxies, yet still resolved on scales three times larger than in local ULIRGs. Once verified directly through high resolution CO gas studies of SMGs (Tacconi et al.\\ 2006, 2008) this was used to argue that SMGs are scaled up ULIRGs. This radio and CO mapping of SMGs % high resolution radio (Chapman et al.\\ 2004) has suggested that the locus of far-IR emission is often offset from UV/optical emission, consistent with findings (Chapman et al.\\ 2005 -- C05) that even dust-corrected UV estimates  of star formation rates in SMGs under-predict the true star formation rates (SFRs) by factors $>100\\times$ on average. %it has only been through CO linewidths for a sizable sample of SMGs (Frayer et al.\\ 1998, Neri et al.\\ 2003, Greve et al.\\ 2005) have suggested large dynamical masses, $>10^{11} M_\\odot$, though not nearly as large as the halos in which they reside based on their clustering properties (Blain et al.\\ 2004). %ZZZ - Need a reference here to some mass-r_0 link since it's the intro? % The large dynamical masses are associated with large ($1-2\\times10^{11}$~M$_\\odot$) stellar masses %as measured from the rest-frame $K$-band (Borys et al.\\ 2005; Hainline 2008; Swinbank et al.\\ 2008). Finally, the gas masses and SFRs define an average gas exhaustion timescale $<$50~Myrs, suggesting a duty cycle of $\\sim$10 for an underlying population of galaxies hosting the submm-luminous events. These studies bring together a picture of SMGs as typically representing massive, gas rich galaxies, which ought to be building elliptical galaxies that are found preferentially  in  clusters at the present day (Swinbank et al.\\ 2006, 2008).  The strongly clustered populations of red $K$-selected galaxies at $z=2-4$ (e.g., Daddi et al.\\ 2003) and of luminous high-$z$ radio galaxies (HzRG -- Brand et al.\\ 2005) suggest they are formed in the highest density perturbations at early epochs, implying that they are the progenitors of local massive early-type galaxies and $z\\sim1$ extremely red objects (EROs) which themselves have similarly strong clustering (e.g., Brown et al.\\ 2005). HzRGs are believed to host massive black holes; thus may represent some of the most massive galaxies at these epochs (e.g., de Breuck et al.\\ 2000) and could trace some of the highest density environments. The identi\u00decation of modest excesses of companion galaxies around HzRGs in the optical/near-infrared (e.g., Rottgering et al.\\ 1996), X-ray (Pentericci et al.\\ 2002; Smail et al.\\  2003a), and submm (Ivison et al. 2000; Stevens et al. 2003), supports  their association with overdense structures.  Studies to date have suggested that  the most FIR-luminous, distant galaxies  are also clustered very strongly (Blain et al.\\ 2004; Borys et al.\\ 2004; Scott et al.\\ 2002, 2006; van Kampen et al.\\ 2005; Farrah et al.\\ 2006; Blake et al.\\ 2006; Gilli et al.\\ 2007). At face value this would indicate that they occupy the largest dark-matter halos ($\\sim$10$^{13}$~M$_\\odot$). However selection of galaxies by the most FIR-luminous specimens may be prone to finding the brief periods when systems of lower matter overdensity are rapidly evolving (e.g.\\ Scannapieco \\& Thacker 2003; Furlanetto \\& Kamionkowski 2006). Indeed burst timescales as short as 10~Myr have been estimated for some SMGs (e.g., Smail et al.\\ 2003b). Blain et al.\\ (2004) calculated that, complex biases aside, SMGs have a clustering length which is consistent with a form of evolution ensuring their properties subsequently match the clustering length typical of evolved red galaxies at $z\\sim1$ and finally of rich clusters of galaxies at the present epoch.  % Figure projection of z=2 structure \\begin{figure} \\centering \\includegraphics[width=7.9cm,angle=0]{f1.ps} \\caption{\\footnotesize Positions of all UV-selected galaxies an SMGs/SFRGs in the survey region are shown as small symbols, while those lying in the $z=1.99$ structure are large symbols. Offset projected Mpc are shown on the axes for reference at $z=2$. Galaxies in the $z=1.99$ structure is generally widely separated on the sky, with $>$Mpc separations between most SMGs/SFRGs except for two close ($\\sim100$~kpc) SMG/SFRG pairs in two cases. } \\label{fig1} \\end{figure}  % % Figure N(z) % \\begin{figure} \\centering \\includegraphics[width=7.5cm,angle=0]{f2a.ps} \\includegraphics[width=7.5cm,angle=0]{f2b.ps} \\caption{\\footnotesize {\\bf (top panel)} Redshift histograms in GOODS-N (bins of $\\delta z= 0.02$). from the surveys of Reddy et al.\\ (2006) in UV-selected BX/BM galaxies, {\\it Team Keck Redshift Survey}  (Wirth et al.\\ 2004), and Hawaii redshift survey (Cowie et al.\\ 2004), the SMG surveys of Chapman et al.\\ (2003, 2005), and the SFRG surveys of Chapman et al.\\ (2004) and Casey et al.\\ (2008). The z=1.99 structure can be seen as the highest peak of galaxies in the $z\\sim2$ region. The selection functions for SMGs and BX/BMs are depicted normalized to the number of sources in each sample. The SFRGs are included in the SMG selection function (see text). There are a few overlapping BX/BM, TKRS and SMGs/SFRGs which are effectively double-counted in the plot. %{\\bf (bottom panel)} A zoom in around the $z=1.99$ structure, where overlapping galaxies in the various samples have been removed, and the histogram now represents a cumulative histogram of all populations. % (notably the two OFRGs at z=2.09), but none Note that the BX and SMG/OFRG galaxy samples in the  $z=1.99$ structure are completely non-overlapping, and the structure could arguably be more significant than other BX galaxy peaks in the combined BX/SMG galaxy sample. {\\bf (lower panel)} A zoomed in view of the structure with smaller redshift bins ($\\delta z= 0.005$). The spike galaxies listed in Table~1 are shown by the vertical bars.} \\label{fig2} \\end{figure}  %ZZZ - why pick 100 arcmin? I think 100 arcmin %is about right to get one of these guys - since 100 arcmin is ~50 Mpc.  However, as pointed out by Blain et al.\\ (2004),  this is inconsistent with the expected density of rich clusters. It  is very unlikely (0.05\\% chance) that any field of order 10\\arcmin$\\times$10\\arcmin\\  includes the progenitor of a moderately rich cluster (volume density $\\rho= 10^{-7}$Mpc$^{-3}$), even in the long pencil beam from $z=$1-3 of a radio-SMG survey. This apparent  paradox can only be resolved by %this is inconsistent with the progenitor of a relatively rich cluster of galaxies intersecting any field out to a redshift $\\sim3$ covering a solid angle of order 100 arcmin$^2$ on a side, an apparent paradox that can only be resolved by reducing the frequency of such associations by a factor of ten ($\\rho=10^{-7}$Mpc$^{-3}$ observed for moderately rich clusters compared with  $10^{-6}$Mpc$^{-3}$ implied by SMGs). Blain et al.\\ concluded that either the dark matter halo masses are less extreme than inferred from a picture of simple 1-1 biasing at a given mass, or some other biasing effect must be at work. % that SMGs appear to cluster as strongly Adelberger et al.\\ (2005a) presented an alternative explanation in the inherent bias of the cluster calculation methodology, however statistical analyses of the SMG clustering through other methods is consistent with the Blain et al.\\ (2004) calculation (e.g., Blake et al.\\ 2006). The alternative explanation, that SMGs often reside in less overdense environments (and that ULIRG clustering could be dominated by the coordination of highly active periods across modest mass environments), is tested here, using UV observations of the richest overdensity found  in Blain et al.\\ (2004) and Chapman et al.\\ (2005). %our own spectroscopic surveys.>> REFS.  In \\S2 we describe the sample properties, with \\S2.1 assessing the various selection functions, and a developing a statistically robust measure of the galaxy overdensity. In \\S2.2 we analyze the properties of the galaxies lying in the overdense structure, including their star formation rates at a number of wavelengths, their luminosity function, and their angular positions on the sky. In \\S3 we discuss the theoretical expectations of the measured redshift space contrasts, propose an explanation for the differential bias seen at different wavelengths (\\S3.1), and discuss the implications for future surveys at submm wavelengths (\\S3.2). We assume a $\\Lambda$-CDM cosmology (Spergel et al.\\ 2003) with $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, so that 1\\,arcsec corresponds to 8.4\\,kpc physical size at $z=2.0$. ",
        "conclusions": "% %How do the 24um inferred SFRs for SMGs+SFRGs differ from the UV-selected galaxies (which are completely orthogonal to the SMGs+SFRGs in this structure). % %Discussion of the SA22 $z=3.09$ result, and possible larger extent of the cluster. Could we have somehow just found a low density edge of a super-cluster? %Are we finding a limb of a filamentary structure where cold flows stimulate more SMGs etc.? % We have described an apparent cluster of SMGs and RGs at $<z>=1.99$ in the GOODS-N field, the strongest known association of SMGs (Blain et al.\\ 2004; Chapman et al.\\ 2005). The implied matter over-density from the redshift space contrast $\\delta^z_g\\simeq10$ is is $\\sim10^{15}$~M$_\\odot$. Searching the same volume in UV-selected galaxies, we find a much reduced overdensity ( $\\delta^z_g\\simeq2.5$). % and the implied dark matter halo mass scale is $<$3$\\times$ less. The low SFRs of UV-galaxies and much longer star formation timescales on average compared with the SMGs and SFRGs suggest they should be more dependable for defining the mass scale of this structure. % We conclude that we are observing a modest mass overdensity (which likely will not even virialize by $z=0$), which is undergoing a particularly active star formation phase. %(to be explained in some heuristic manner, since it isn't an obvious byproduct of any simulations)  We have characterized the bolometric luminosity function of the $z\\sim1.99$ spike galaxies compared to that estimated for a large sample of UV-selected star forming galaxies (Reddy et al.\\ 2008), finding that it is almost entirely in the SMG population where an excess is seen. The clustered environment allows for a useful comparison of SFR and stellar mass indicators without K-correction and model galaxy uncertainties. Our new radio measurements for the spike galaxies detect significant numbers of the lower luminosity UV-selected galaxies in addition to all the SMGs and SFRGs, and radio-derived SFRs are generally consistent with those derived from {\\it Spitzer} 24$\\mu$m fluxes.  %However, timescale of the luminosity must be key. While the most overdense forming environments at $z>1.5$ have been observed to have large numbers of ultra-luminous galaxies associated, %(e.g., the SSA22 $z=3.09$ proto-cluster has 8 known SMGs -- Geach et al.\\ 2006), our results here suggest that if one searches for galaxy clusters using the highest luminosity systems, one is prone to finding the brief periods when systems of lower matter overdensity are rapidly evolving. SMGs don't obviously trace the most massive structures in the Universe, as simulations and constraints to date have suggested they should. The success of our toy model with merger bias explaining the basic elements of our observations  suggests that simulations may not currently be accurately reflecting the processes occurring in the formative phases of galaxy clusters. Further work to embed these aspects of merger bias in simulations will therefore be of great interest in understanding galaxy cluster assembly. Investigation of similar systems in upcoming surveys at long wavelengths will be of great interest to explore the full range of environments associated with strong overdensities of ultra-luminous galaxies."
    },
    "0809/0809.4486_arXiv.txt": {
        "abstract": "Using the local helioseismic technique of ring diagram  we analyze the frequencies of high--degree {\\it f}- and {\\it p}-modes derived from both velocity and continuum intensity data observed by MDI. Fitting the spectra with asymmetric peak profiles, we find that the asymmetry associated with velocity line profiles is  negative for all frequency ranges agreeing with previous observations while the asymmetry of the intensity profiles shows a complex and frequency dependent behavior. We also observe systematic frequency differences between intensity and velocity spectra at the high end of the frequency range,  mostly above 4~mHz. We infer that this difference arises from the fitting of the intensity rather than the velocity spectra. We also show that the frequency differences between intensity and velocity  do not vary significantly from the disk center to the limb when the spectra are fitted with the asymmetric profile and conclude that only a part of the background is correlated with the intensity oscillations. ",
        "introduction": "Different helioseismic instruments, both from ground and space, observe the Sun in different observables. Due to the different techniques used by these instruments, it is possible to measure the solar oscillations simultaneously either as variations in the photospheric velocity or as intensity fluctuations. It is therefore important to confirm that the oscillation mode parameters measured from both the intensity and velocity agree with each other to a high degree of precision. However, the initial measurement of low degree {\\it p}-mode frequencies from Doppler velocity ($V$) and continuum intensity (I$_{\\rm c}$) observations from Michelson Doppler Imager (MDI) instrument on board {\\it Solar and Heliospheric Observatory} (SOHO) showed systematic differences. A comparison of 108-day power spectra between $V$ and \\ic showed a weighted mean difference of $-0.1~\\mu$Hz for $\\ell=0$, and $-0.16~\\mu$Hz for $\\ell=1$ modes  \\citep{toutain97}. Since the apparent frequency shift between an oscillation observed in velocity and intensity cannot be a property of the mode, it must arise from systematic errors while calculating the frequencies from the observed power spectra. Hence it was argued that the source of the systematic difference could be due to the opposite asymmetry effect seen between the velocity and intensity power spectra \\citep{duvall93}. \\cite{app98} also presented a similar evidence using VIRGO and SOI/MDI data. Around the same time \\cite{harvey98} reported that the intermediate degree modes observed in $V$ and total spectral intensity also show different central frequencies and observed that the apparent differences could be as large as 50~$\\mu$Hz close to the acoustic cut-off frequency. However, the analysis of \\cite{nigam98b}, using an asymmetric line profile-fitting formula, illustrated that the frequency difference between $V$ and \\ic in the intermediate degree range is much smaller compared to that obtained by fitting a symmetric Lorentzian profile. Using the same asymmetric line profile-fitting formula, \\cite{toutain98} re-analyzed the data from MDI and showed that the  frequency differences between $V$ and \\ic are considerably reduced. \\cite{gav99} have also analyzed data from different instruments and  have argued that the reported  frequency shift is merely an artifact of the reduction technique.  Renewed interest in the topic began when local helioseismic techniques were developed  to study the properties of high-degree modes in localized regions.  \\cite{bab01} analyzed azimuthally averaged (2-d) power spectra and inferred  that the eigenfrequencies obtained using the asymmetric peak profiles agree reasonably well with each other while the use of symmetric profiles gives significant differences between frequencies computed using continuum intensity and velocity or continuum intensity and line-depth spectra. In order to gain further information for high-degree and high-frequency modes, \\cite{jain03} analyzed the high-resolution GONG+ data.  These authors also compared the azimuthally averaged power spectra of a region centered on the equator and reported that the frequency dependence of the frequency shift between $V$ and I is negligible below the acoustic cutoff frequency around 5.3 mHz and substantial above the cutoff. These results supported the finding of \\cite{harvey98}. However, the conclusion is based on the visual comparison of the peak frequency of the power spectra and  may not necessarily be a true measure of the shift due to the reversal of the asymmetry and different background power between $V$ and \\ic spectra.  It is now well established that line asymmetry of the solar power spectra alters the peak frequencies that are obtained under the assumption that the lines are symmetric (e.g. \\cite{ab99, ba00}. However, the cause of the opposite asymmetry between the velocity and intensity spectra still remains inconclusive.  The current understanding is that the reversal in the sign of asymmetry between the $V$ and \\ic spectra is due to the influence of the solar background noise that is correlated with the source \\citep{rv97, nigamall98, severino01} and the level depends on the characteristic granulation. On the other hand, the numerical simulation \\citep{geo03} indicates that the reversal is produced by the radiative transfer effects. Since the physics of the correlated noise is not yet fully understood and the spatial leakage signature for $V$ and I is  different due to their center-to-limb variations,  our objective is to examine the frequency dependence of the observed asymmetry and differences in eigenfrequencies between velocity and intensity observations as a function of the radial distance from the disk center to the limb. A preliminary investigation of a similar nature using azimuthally averaged power spectra is reported in \\cite{sct06}. However  the present  analysis differs from all earlier ones since here we use the three-dimensional (3-d) power spectrum, which is associated with flow fields, while the azimuthally averaged spectrum has no flow fields associated with it.  The rest of the paper is organized as follows: We briefly describe the data and analysis technique in Section 2, while the results are described in Section~3. Finally, we summarize the main conclusions in Section~4. ",
        "conclusions": "Using the local helioseismic technique of ring diagram analysis we have studied the high-degree {\\it f}- and {\\it p}-mode frequency differences between velocity and continuum intensity data obtained from the MDI instrument during a quiet period. Since the spectra are known to be asymmetric, we  fitted the three-dimensional spectra with an asymmetric model profile based on the known form of the asymmetry observed in velocity and intensity. The asymmetry obtained from fitting the velocity spectra agrees with the previous results while that of intensity shows a frequency dependent behavior. We further notice significant differences between the asymmetry obtained from fitting the three-dimensional and azimuthally averaged power spectra.  Fitting the three-dimensional disk center power spectra with a symmetric Lorentzian profiles leads to frequency differences on the order of 10~$\\mu$Hz or less up to $\\nu \\le$ 4~mHz  between intensity and velocity. This difference is found to be increasing systematically from the disk center to the limb. The use of asymmetric profiles leads to frequency differences that are smaller than the differences resulting from the symmetric fits. However, systematic differences still remain at the high end of the frequency range, mostly above 4~mHz. We demonstrate that this difference arises from the fitting of the intensity rather than the velocity spectra. We speculate that the form of asymmetry used in the model for the intensity spectra is over-simplified and does not adequately represent the real profile at all frequency ranges.  We also conclude that the frequency differences between velocity and intensity do not vary significantly from the disk center to the limb. However, variations between intensity spectra are higher than seen in velocity. This supports the idea that only a part of the background is correlated with the intensity oscillations. In this context several authors (e.g.  \\cite{severino01, bhc04, toutain06}) have advocated the importance of fitting the intensity-velocity cross spectrum along with a multicomponent model of the coherent background. This provides an impetus for future work using the high-resolution observations from {\\it Solar Dynamics Observatory}."
    },
    "0809/0809.1430_arXiv.txt": {
        "abstract": "We study the formation of galaxies in a large volume ($50\\hmpc$, $2 \\times 288^3$ particles) cosmological simulation, evolved using the entropy and energy conserving smoothed particle hydrodynamics (SPH) code \\newgad. Most of the baryonic mass in galaxies of all masses is originally acquired through filamentary ``cold mode'' accretion of gas that was never shock heated to its halo virial temperature, confirming the key feature of our earlier results obtained with a different SPH code \\citep{keres05}. Atmospheres of hot, virialized gas develop in halos above $2-3\\times 10^{11} \\msun$, a transition mass that is nearly constant from $z=3$ to $z=0$. Cold accretion persists in halos above the transition mass, especially at $z \\geq 2$.  It dominates the growth of galaxies in low mass halos at all times, and it is the main driver of the cosmic star formation history. Our results suggest that the cooling of shock-heated, virialized gas, which has been the focus of many analytic models of galaxy growth spanning more than three decades, might be a relatively minor element of galaxy formation. At high redshifts, satellite galaxies have gas accretion rates similar to central galaxies of the same baryonic mass, but at $z<1$ the accretion rates of low mass satellites are well below those of comparable central galaxies. Relative to our earlier simulations, the \\newgad simulations predict much lower rates of ``hot mode'' accretion from the virialized gas component. Hot accretion rates compete with cold accretion rates near the transition mass, but only at $z \\leq 1$.  Hot accretion is inefficient in halos above $\\sim 5\\times 10^{12}\\msun$, with typical rates lower than $1\\msunyr$ at $z \\leq 1$, even though our simulation does not include AGN heating or other forms of ``preventive'' feedback.  Instead, the accretion rates are low because the inner density profiles of hot gas in these halos are shallow, with long associated cooling times.  The cooling recipes typically used in semi-analytic models can overestimate the accretion rates in these halos by orders of magnitude, so these models may overemphasize the role of preventive feedback in producing observed galaxy masses and colors.  A fraction of the massive halos develop cuspy profiles and significant cooling rates between $z=1$ and $z=0$, a redshift trend similar to the observed trend in the frequency of cooling flow clusters. ",
        "introduction": "The wealth of high quality observational galaxy data at different epochs, from the earliest times to the present, has rapidly grown in the last several years. This provides a great opportunity to calibrate and improve our theoretical understanding of galaxy formation and evolution. One of the remaining challenges in theoretical extragalactic astronomy is the origin of the bimodality in galactic properties: massive galaxies that are typically red, elliptical systems with very little star formation, and lower mass galaxies that are typically blue, disky, star forming systems (e.g. \\citealt{kauffmann03a, baldry04}). Usually one invokes various feedback processes to shut down star formation in massive galaxies, which enables them to populate the red sequence \\citep[see review in][]{hopkins08b}. In addition, without feedback in the low mass galaxies, a large fraction of baryons would cool and form stars, even at very early times \\citep{white78}, contrary to the small fraction of baryons in observed galaxies \\citep[e.g.][]{bell03}.  To understand where, when, and to what degree we need feedback during galaxy formation and evolution, we first have to understand how galaxies are supplied with baryons, and in particular with the gas that provides the fuel for star formation. In our previous work \\citep[K05 hereafter]{keres05}, we showed that it is the smooth accretion of intergalactic gas that dominates the global galactic gas supply, not accretion by merging, and that it proceeds via two stages. First, gas is accreted through filamentary streams, where it remains relatively cold before it reaches the galaxy \\cite[K05]{katz03}. This process, which we call cold mode accretion, is very efficient, and because the gas does not need to cool it falls in on approximately a free-fall time. This means that the baryonic growth closely tracks the growth of the dark matter halo, albeit with a slight time delay. Cold mode accretion dominates the global growth of galaxies at high redshifts and the growth of lower mass objects at late times. As the dark matter halo grows larger, a larger fraction of the infalling material shock heats to temperatures close to the virial temperature. In the denser, central regions a fraction of this hot gas is able to cool. This process, virial shock heating followed by radiative cooling, is the one highlighted in classic papers on galaxy formation theory \\citep[e.g.][]{rees77, silk77, white78, white91}, and to differentiate it from the previous mode we call it hot mode accretion. In massive halos, this hot mode accretion not only supplies massive galaxies with fresh gas but could also be identified as cluster cooling flows \\citep[e.g.][]{fabian94}, although observations suggest that cooling in most clusters does not actually reach down to galactic temperatures, perhaps because of heating by AGN or some other feedback mechanism \\citep[e.g.][]{mcnamara07} . The dominant accretion mode depends on halo mass, and the transition between these two regimes occurs at a mass of about $M_{\\rm halo}=2-3\\times 10^{11}\\msun$. At higher redshifts, cold filaments are able to survive within halos above this transition mass, i.e. halos dominated by hot halo gas. Overall, the bulk of the baryonic mass in galaxies is accreted through the cold accretion mode.  Cold accretion owes its existence to the short cooling times present in low mass halos near the virial radius, which prevents the development of a stable  accretion shock.  The detailed criteria required to prevent the formation of a virial shock in low mass idealized systems have been derived by \\citet{binney77} and \\citet{birnboim03}. However, in cosmological simulations the situation is more complex than in these simple models. For example, the spherical symmetry assumed in the case of \\cite{birnboim03} will not be valid owing to the presence of dense, cold filaments that dramatically enhance the density contrasts in the gas and shorten the cooling times. In halos with masses near the transition mass from cold to hot mode, this allows cold filamentary flows to still supply the central galaxy with gas even though some of the infalling gas shock heats to near the virial temperature.  There are two main classes of galaxy formation feedback processes: those that prevent gas from entering a galaxy in the first place, ``preventive'' feedback, and those that expel a fraction of the gas that does manage to enter the galaxy, ``ejective'' feedback. AGN ``radio mode'' heating is an example of preventive feedback \\citep[e.g.][]{best05,croton06} and winds driven by supernovae are an example of ejective feedback \\citep[e.g.][]{dekel86,springel03a}. The effectiveness of these two feedback types will vary depending on the dominant accretion mode in the halo hosting the galaxy. Because cold mode halos have very little halo gas outside of the cold dense filaments, it is not clear if feedback processes can drastically affect the accretion of gas. Ejective feedback, such as galactic winds driven by supernovae, could still lower the masses of galaxies by expelling already accreted material. Such winds, however, might be stopped by the quasi-spherical, hot halos that surround hot mode galaxies. Conversely, preventive feedback, like the AGN radio mode, is likely only effective in hot mode halos, where it can prevent the hot, dilute gas that is in quasi-static equilibrium from cooling. Similar constraints apply for the ``quasar wind mode'' of AGN feedback \\citep{dimatteo05}, which could possibly be a mixture of these two feedback modes. At first, a relatively small fraction of galaxy mass is ejected in the quasar wind, but later the energy released from the black hole accretion could in principle provide ``preventive'' feedback if the quasar was active within a halo that has a hot quasi-spherical atmosphere. In summary: ejective feedback is likely to be most effective in cold mode halos, and preventive feedback is likely to be effective only in hot mode halos. An exception to this rule is ``pre-heating'' of intergalactic gas by photoionization \\citep{efstathiou92, quinn96, thoul96}, or gravitational shocks \\citep{mo05}, which is a form of preventive feedback that will mostly affect low mass, cold mode halos.  Current cosmological simulations have insufficient resolution (often by large margins) to predict these feedback processes from first principles, though they may incorporate simplified recipes designed to achieve these effects. It is, therefore, important to gain a general understanding of when and where these two classes of feedback processes are needed to explain the observed properties of galaxies. In this paper, we examine the growth of galaxies by gas accretion and the dependence of this process on galaxy mass, halo mass, and redshift. In a companion paper \\citep[Paper II, hereafter]{keres09}, we compare the observable properties of our simulated galaxies, such as stellar masses and specific star formation rates, to current data.  In K05 we showed that the mass of the transition from cold to hot mode accretion  does not depend much on the resolution of the cosmological simulation.  However, we showed that increasing the mass and spatial resolution substantially reduced the amount of hot mode accretion.  This suggested that the simulations from K05 had too much hot mode accretion, perhaps the result of numerical problems inherent in the Smoothed Particle Hydrodynamics (SPH) formulation that we used (see \\citeauthor{springel02} \\citeyear{springel02}, SH02 hereafter, for a brief review). How much gas is able to cool from the hot atmosphere remained, therefore, an unanswered question in K05. In this paper, we use the \\newgad SPH code \\citep{springel05a} because it does not suffer from numerical overcooling in massive halos, as we demonstrate in Kere\\v{s} et al. (2009b, in preparation) and in \\citet{mythesis}, and because it parallelizes more efficiently to large numbers of processors. The SPH algorithm implemented in \\newgad, which simultaneously conserves entropy and energy, is a reliable method for simulations of cosmologically relevant volumes where galaxies are often only resolved with a small number of particles.  In our previous simulations, this low resolution allowed a mixing of the cold and hot gas phases, which led to an artificial increase in the cooling rates and an overestimate of the hot mode accretion \\citep{pearce01, croft01}. In our new \\newgad simulations, we show that a large fraction of the massive, hot mode halos cannot cool their gas, owing to nearly uniform density cores that develop in these halos.  However, as discussed below (\\S\\ref{sec:acc_rates}), \\newgad simulations may still suffer numerical problems that distort gas accretion rates, albeit at lower magnitude; fully understanding these effects will require tests beyond those in this paper. Relative to K05, we also increase the simulation volume by about an order of magnitude, allowing much better statistics for high mass halos and galaxies. We use a similar but slightly better mass resolution.  Here we present theoretical predictions of galaxy masses and accretion rates in cosmological simulations without strong feedback. The energy from supernovae pressurizes the interstellar medium, but it does not drive large scale outflows. We do not include any other feedback mechanisms that could prevent the accretion of gas or that could expel gas from galaxies in the mass range of interest in this paper. The only exception is that we include an extragalactic UV background, which affects the intergalactic medium and reduces gas accretion in the lowest mass halos.  In the Paper II we confirm earlier findings that the amount of material that cools onto galactic component is too high in simulations such as ours, which causes galaxies at any mass, but especially at the low and high mass end to be more massive than their observed counterparts. There we argue that an efficient high redshift feedback in low mass galaxies (perhaps with the addition of low redshift AGN feedback) is the key for getting the correct galaxy masses.  Many of the results in this paper and its companions (Paper II, and Keres et al. 2009b (in preparation)) first appeared as part of D. Kere\\v{s}'s PhD thesis at University of Massachusetts, Amherst \\citep{mythesis}.   In \\S\\ref{sec:simulations} we describe the basic parameters and setup of our new simulations. In \\S\\ref{sec:accretion} we describe the smooth accretion of gas, both globally and for individual objects. In \\S\\ref{sec:halos} we describe the properties of non-cooling and cooling halos and conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Our principal results are based on a cosmological SPH simulation of a $50\\hmpc$ (comoving) box in a $\\Lambda$CDM universe, evolved with $288^3$ dark matter and $288^3$ gas particles using the code \\newgad.  The combination of volume and particle number allows us to resolve (with $> 64$ particles) galaxies of baryonic mass $\\mgal \\ga 6 \\times 10^9\\msun$ while following the formation of several cluster-size halos with $\\mhalo \\ga 10^{14}\\msun$. Relative to the main simulation analyzed by K05, our present simulation has an eleven times larger volume, and it uses the entropy-conserving SPH formulation of \\cite{springel02}, which should provide a more accurate treatment of multi-phase gas than the PTreeSPH code employed by K05.  We incorporate photoionization by the UV background and thermal feedback on a 2-phase galaxy interstellar medium, but we do not explicitly add galactic winds or track metal enrichment, so this is a ``minimal feedback'' simulation.  Confirming the key result of K05, we find that galaxies with baryonic masses below $2-3 \\times 10^{10}\\msun$ or halo masses below $2-3 \\times 10^{11}\\msun$ gain most of their mass through ``cold'' accretion of gas that has never been close to the halo virial temperature (see also \\citealt{kay00,fardal01,katz03}). In halos below $\\sim 2.5 \\times 10^{11}\\msun$, the majority of intergalactic gas is colder than 250,000$\\,$K, while in higher mass halos the majority of gas is close to the virial temperature; the transition mass stays nearly constant from $z=3$ to $z=0$. Similar behavior is found in spherically symmetric calculations \\citep{birnboim03,dekel06}, and the likely cause of the transition is that the short cooling times in low mass halos prevent the formation of a stable virial shock (\\citealt{birnboim03}; see also \\citealt{binney77,white91}).  However, as shown by K05 and illustrated more comprehensively here, halos near and above the transition mass have hot gas halos penetrated by cold gas filaments, which may continue to feed the central galaxy.  Similar results are found in adaptive mesh refinement simulations (\\citealt{ocvirk08,dekel09}; see also \\cite{dekel06} figure 6). At high redshift these cold filaments penetrate close to the halo center, while in high mass halos at $z<1$ they are usually truncated or disrupted before reaching the center. The difference likely reflects the higher physical densities and shorter associated cooling times of the high redshift filaments that prevent the expansion of a stable shock \\citep{birnboim03} in these high density regions.  The key difference between our present results and those of K05 is that {\\it hot} accretion rates are much lower, a consequence of changing to the more accurate, entropy-conserving SPH formulation.  Because of the low hot accretion rates, the growth of high redshift galaxies ($z \\geq 2$) is dominated by cold accretion at all galaxy and halo masses. At lower redshifts, massive halos experience a combination of hot accretion and ``cold drizzle,'' accreting lumps of cold gas below our 64-particle galaxy resolution threshold. Our resolution tests here and those of \\cite{naab07} suggest that some or even most of this ``cold drizzle'' may be a numerical artifact.  Eliminating it would significantly affect the predicted star formation rates and colors of massive galaxies, but it would not have much impact on their predicted stellar masses (Paper II). At {\\it all} galaxy masses, the majority of galaxy baryonic mass in our simulation was originally acquired by cold accretion, either directly onto the main progenitor or onto a satellite that merged with it.  Galaxy accretion rates decline systematically from $z=4$ to $z=2$ to $z=1$ to $z=0$.  For $10^{10} \\msun < \\mgal < 10^{11}\\msun$, there is a trend of increasing accretion rate with increasing mass at all redshifts, though it becomes flatter at low $z$.  At higher masses, accretion rates have large scatter and are sensitive to cold drizzle. With our adopted star formation prescription, empirically motivated by the Kennicutt-Schmidt law \\citep{kennicutt98,schmidt59}, star formation closely tracks gas accretion with a short delay; K05 showed this on a global basis, and here we show it on a galaxy-by-galaxy basis.  At each redshift, more massive galaxies tend to have higher star formation rates but (moderately) lower specific star formation rates.  Satellite galaxies orbiting within the virial radius of larger halos experience continuing accretion and star formation. The differences in the accretion rates of central and satellite galaxies of the same mass are small at high redshift, while at low redshift a significant fraction of satellites have much lower accretion rates.  More detailed analysis of the satellite accretion and merger rates (based on the K05 simulation) appears in \\citet{simha08}.  The low hot accretion rates in our simulated massive halos arise because these halos typically have cores in their hot gas density profiles (or, more precisely, because they break from an approximately $r^{-2}$ profile to a much shallower profile in their inner regions).  By $z=0$, a modest fraction of the $M>10^{14} M_\\odot$ halos develop cuspy density profiles with cooling cores, evolution that is reminiscent of Vikhlinin et al.'s (\\citeyear{vikhlinin07}) finding that cooling flow clusters are fairly common at $z=0$ but much rarer at $z>0.5$.  We speculate that the cores are produced by shocks during the chaotic, rapid assembly phase of halos, and that the more quiescent evolution during the slow accretion phase allows cooling to establish a cusped profile. The accretion rates of central galaxies in our massive halos are predicted with reasonable accuracy by instantaneous cooling calculations that use the actual simulated gas density profiles, but cooling calculations that assume isothermal gas in equilibrium in an NFW halo or singular isothermal halo may overestimate the cooling rates by orders of magnitude.  Our simulation predicts low hot accretion rates even though it does not include AGN heating or other forms of preventive feedback. While such feedback may be needed to explain the absence of cold gas in cooling flow clusters \\citep{kaastra01,peterson03}, semi-analytic models may overestimate its global importance by assuming incorrect halo gas profiles (and thus excessive cooling).  The comparison to observations in Paper II shows that this simulation, despite its low hot accretion rates, overpredicts the galaxy baryonic mass function.  The amount of baryons that is locked in galaxies at $z=0$ and the amount of baryons locked in stars at $z=0$ is about a factor of 2-3 higher than observed. This is caused by simulated galaxies being too massive at all masses, but the exact difference is a function of galaxy mass with the lowest and the highest mass galaxies being the most problematic. Matching observations therefore requires either the suppression of cold accretion (which seems unlikely on physical grounds) or substantial amounts of ejective feedback that returns gas to the IGM before it forms stars. Such ejective feedback is also needed to explain the observed enrichment of the IGM (e.g., \\citealt{oppenheimer06,oppenheimer08}). Ejected, metal-enriched gas will affect the cooling of the IGM in massive halos, and it will reduce the rate at which mergers add stars and gas to massive galaxies.  We will investigate the impact of winds and metal-enrichment on galaxy growth in future work with simulations that include these effects.  However, our present results suggest that the cooling of shock-heated, virialized gas, which has been the focus of many analytic models of galaxy growth spanning more than three decades, might be a relatively minor element of galaxy formation.    We thank V. Springel for making the \\newgad public and K. Finlator and B. Oppenheimer for help during the modification of the public version of the code. D.K acknowledges the support from the ITC Fellowship at the Institute for Theory and Computation at the Harvard College Observatory. We also acknowledge support from NSF grant AST-0205969 and from NASA grants NAGS-13308 and NNG04GK68G."
    },
    "0809/0809.2928_arXiv.txt": {
        "abstract": "An approach for measuring linear X-ray polarization over a broad-band using conventional spectroscopic optics is described. A set of multilayer-coated flats reflect the dispersed X-rays to the instrument detectors. The intensity variation as a function of energy and position angle is measured to determine three Stokes parameters: I, Q, and U. By laterally grading the multilayer optics and matching the dispersion of the gratings, one may take advantage of high multilayer reflectivities and achieve modulation factors over 80\\% over the entire 0.2 to 0.8 keV band. A sample design is shown that could be used with a small orbiting mission. ",
        "introduction": "\\label{sec:intro}  %  Although polarization is a fundamental property of electromagnetic radiation, there have been no astronomical measurements of polarization in the X-ray band in over 30 yr.  There has been no lack of workable concepts, including simple designs such as the Polarimeter for Low Energy X-ray Astrophysical Sources (PLEXAS) that was proposed by Marshall et al.\\ (2003) to use multilayer-coated mirrors tuned to 0.25 keV as Bragg reflectors.\\cite{plexas} An excellent review of the history and prospects for astronomical polarimetry in the 0.1-10 keV band is presented by Weisskopf et al. (2006)\\cite{weisskopf06}. Weisskopf et al.\\ rightly argue that the PLEXAS design had a narrow bandpass, reducing its attractiveness for astrophysical observations because one expects polarization to be energy dependent, so a wide bandpass is desired.  Marshall (2007\\cite{2007SPIE.6688E..31M}; hereafter, Paper I) described a method to overcome this limitation by using transmission gratings to disperse in the incoming X-rays.  Then, a multilayer-coated reflector would be used to modulate the signal in a way depending on the polarization. Two approaches were suggested.  The first was a simple modification of a dispersive spectrometer, such as a grating spectrometer proposed for Constellation-X (Con-X).  The second approach was intended to provide a means for obtaining polarimetric data by inserting optics forward of the focal plane of a long focal length telescope such as XEUS.  Similar to one of the approaches suggested in Paper I, this paper is limited to the case where the polarization is determined by measuring spectral intensity along different dispersed spectra. In this case, we consider a configuration that would be appropriate for a small orbiting observatory. ",
        "conclusions": ""
    },
    "0809/0809.2269_arXiv.txt": {
        "abstract": "{The next generation of ground-based Imaging Air Cherenkov Telescopes will play an important role in indirect dark matter searches. In this article, we consider two particularly promising candidate sources for dark matter annihilation signals, the nearby dwarf galaxies Draco and Willman~1, and study the prospects of detecting such a signal for the soon-operating MAGIC~II telescope system as well as for the planned installation of CTA, taking special care of describing the experimental features that affect the detectional prospects. For the first time in such studies, we fully take into account the effect of internal bremsstrahlung, which has recently been shown to considerably enhance, in some cases, the gamma-ray flux in the high energies domain where Atmospheric Cherenkov Telescopes operate, thus leading to significantly harder annihilation spectra than traditionally considered. While the detection of the spectral features introduced by internal bremsstrahlung would constitute a smoking gun signature for dark matter annihilation, we find that for most models the overall flux still remains at a level that will be challenging to detect, unless one adopts somewhat favorable descriptions of the smooth dark matter distribution in the dwarfs.} ",
        "introduction": "Even though basically nothing is known about its underlying nature, compelling evidence has accumulated in recent years that the building block of the observed structures in the universe is a new form of non-baryonic, cold dark matter (DM), with the fraction of DM to critical density being $\\Omega_{\\mbox{\\tiny{DM}}}\\sim 0.233 \\pm 0.013$ according to the most recent estimates \\cite{Komatsu:2008hk}. A class of particularly well-motivated DM candidates arises in many extensions of the standard model of particle physics in the form of weakly interacting massive particles (WIMPs) that, thermally produced in the early universe, automatically acquire a relic density matching the observed DM abundance today (for reviews, see \\cite{Jungman:1995df,Bergstrom:2000pn,Bertone:2004pz,Taoso:2007qk}). In the following, we will restrict our analysis to the supersymmetric neutralino, which is maybe the best-motivated and the most widely studied DM candidate.  \\par The search for DM is a vast field where different experimental techniques can provide complementary information. At the Large Hadron Collider (LHC), the detection of new supersymmetric particles through missing energy signatures, would provide valuable information about the composition of the neutralino and its role as a DM candidate (see, e.g. \\cite{Battaglia:2004mp,Baltz:2006fm}). The neutralino could also be observed through nuclear recoils in direct detection experiments. A new generation of cryogenic detectors is aiming to improve the already impressive recent limits reported by Angle et al.~\\cite{Angle:2007uj} and Ahmed et al.~\\cite{Ahmed:2008eu}, thereby probing significant regions of the underlying parameter space. Finally, DM can be searched for through \\emph{indirect detection}, i.e.~by looking for primary and secondary products (mainly gamma rays, neutrinos, positrons and anti-protons) of DM annihilation. Our attention will be focused on the detection of high-energy gamma rays.   In this field, Imaging Air Cherenkov Telescopes (IACTs) play a leading role, exploring the possibility of tracing Cherenkov photons from electromagnetic showers produced by gamma rays impinging on the Earth's atmosphere. Ground-based gamma ray astronomy is a very active field that has matured rapidly in recent years. In the near future, new installations with improved sensitivity will investigate those regions in the sky where DM signatures can be expected. In particular MAGIC~II, already in the commissioning phase, will perform observations in the northern hemisphere, while a new generation of  IACTs is currently in the design phase, the Cherenkov Telescope Array (CTA), where an array of several tens of telescopes will have improved angular and energy resolution, as well as a very low energy threshold and a greatly increased sensitivity.  \\par Dwarf spheroidal galaxies (dSphs) are the faintest known astrophysical objects believed to be dominated by DM \\cite{Gilmore:2007fy}. So far, almost all detected dSphs  are located in the Local Group, a fact that is most likely related to the low luminosity of these objects, ranging from $330$~L$_\\odot$ to $3\\times10^7$~L$_\\odot$ \\cite{Mateo:1998wg,DaCosta:1999,Simon:2007dq,Geha:2008zr}. Typically, dSphs are satellite galaxies hosted by the halo of a larger and more massive galaxy that, in some cases, strongly influences their evolution: tidal stripping,  e.g., can deprive a dSph from stars in the outer  region (see, e.g., \\cite{Munoz:2005be,McConnachie:2006nb}). While dSphs share many characteristics with stellar globular clusters --- a similar number of stars (from tens to a few thousand) and a  similar range in luminosity --- they usually have a larger size, a fact that implies the presence of a DM halo of the order of  $10^6-10^8$~M$_{\\odot}$. The resulting mass--to--light ratios are typically larger than 10~M$_{\\odot}/$L$_\\odot$ and can even reach values up to $10^3$~M$_{\\odot}/$L$_\\odot$, suggesting that these objects are particularly interesting targets for indirect DM searches (for previous studies, see e.g., \\cite{Bergstrom:2005qk,Strigari:2007ma,Colafrancesco:2006he,SanchezConde:2007te,Strigari:2006rd}).  \\par Let us mention in passing that the number of detected dSphs is around one order of magnitude below what is expected from $N$-body simulations of structure formation in $\\Lambda$CDM cosmologies, which has become known as the \\emph{missing satellites problem} \\cite{Klypin:1999uc,Moore:1999nt}. While many interpretations have been suggested, including the effect of decaying DM (see, e.g., \\cite{Cembranos:2005us,Kaplinghat:2005sy,Borzumati:2008zz} and references therein), the most natural explanation would be that the missing dSphs have so far simply escaped detection. In fact, it has been proposed that dSphs below a certain mass cannot efficiently accrete baryons, which would explain the intrinsic faintness of these objects \\cite{Strigari:2007ma}. The claimed discrepancy has recently also been mitigated considerably by the discovery of a bunch of  new ultra-faint galaxies in the SDSS data \\cite{Simon:2007dq}.  \\par In this paper, we provide an update on the possible future  of the indirect detection of DM with IACTs. In particular, we discuss the prospects for the detection of DM in dSphs with MAGIC II  and make tentative predictions in view of the expected characteristics of CTA. We choose to focus on two dSphs, Draco and the recently discovered ultra--faint Willman~1. While the choice of Draco is supported by the large set of available data, Willman~1 will be discussed as a very promising target for future detection.  \\par A key ingredient in estimating the gamma ray flux for IACTs  is the computation of the energy spectrum for photons produced by neutralino annihilations. In Ref.~\\cite{Bringmann:2007nk}, it was recently shown that internal bremsstrahlung (IB) provides an important, hitherto neglected, contribution to the spectrum. While the effect of IB is largely model dependent, it appears in general as a pronounced ``bump'' at energies close to the kinematic cutoff at the  neutralino mass. The importance of this effect is two--fold: first, the flux at high energies, where IACTs are most sensitive, is significantly increased; secondly, the introduction of spectral features allows an easier discrimination of a DM source from potential astrophysical sources located in the vicinity, whose spectrum is usually a featureless power law. This is the first time that IB is fully included in studying prospects of DM detection for IACTs.  \\par The paper is organised as follows: in Section~\\ref{sec:sources},  we will briefly review some of the main properties of Draco and Willman~1 and discuss how their DM profiles are constrained by observations. The following Section~\\ref{sec:iact} is dedicated to a short  description of the IACT technique, with particular attention to the features that are relevant for DM observations. In Section~\\ref{sec:flux}, we introduce a set of five benchmark neutralinos, representatives for the relevant regions of the parameter space, and estimate the corresponding flux based on the technical specifications for MAGIC~II and CTA. Results are discussed in Section~\\ref{sec:results}, and in Section~\\ref{sec:conclusions} we present our conclusions. ",
        "conclusions": "\\label{sec:conclusions} In this article, we have computed the prospects of detecting gamma rays from annihilating DM in two nearby dwarf galaxies, Draco and Willman~1, for the upcoming ground--based Imaging Atmospheric Cherenkov Telescopes MAGIC~II and the CTA telescope array (the latter still being in the early design phase). We have focused our analysis on a set of five benchmark models, representatives for the parameter space of neutralino DM in the mSUGRA framework, and paid special attention to describing those telescope features that are most relevant in this context. For the first time in this kind of analysis, we have fully taken into account the contributions from radiative corrections that were recently reported by Bringmann et al.~\\cite{Bringmann:2007nk}. As it turned out, in fact, taking realistic DM spectra has an important impact on the analysis and, although common practice, \\emph{it can be a rather bad approximation to simply assume a featureless DM spectrum like from $b\\bar b$ fragmentation and/or to only focus on the total flux above a given energy threshold $E_0$} in these kind of studies. The basic underlying reason for this is that \\emph{realistic DM annihilation spectra show a harder energy dependence than the sensitivity of IACTs}. Once detected, clear spectral features would, of course, have the additional advantage of providing a rather fool-proof way of discriminating DM spectra against astrophysical background sources -- which is even more important in view of the still rather large astrophysical uncertainties involved.  Although these effects do provide a considerable enhancement of the detectional prospects, the expected flux from dSphs remains at a level that, for conservative scenarios, will be challenging to detect with the next generation of IACTs. This, rather than the angular resolution of these instruments, is the reason why the potential of IACTs to discriminate between different DM profiles in  dSphs is limited even in the case of the detection of an annihilation signal; the  eventual disentanglement between cored and cuspy profiles is probably more promising to perform at other wavelengths.  On the other hand, if one adopts the most optimistic astrophysical configurations that are compatible with current observational data of Willman~1, i.e. a favourable DM profile and an $\\mathcal{O}(10-100)$ flux enhancement due to the existence of substructures, \\emph{all} of our benchmark models approach the reach of at least the CTA which, for the models studied here, is a factor of $6-8$ more sensitive to the annihilation signal than MAGIC~II (this is, of course, independent of the source). The most promising case of our analysis turns out to be a neutralino in the funnel region, characterised by no sizeable IB contributions to its spectrum but a rather large annihilation rate; the second best case is a neutralino from the coannihilation region, making up for its small annihilation rate with enormously large radiative corrections.  Having demonstrated that the prospects of indirect DM detection through gamma rays do depend on the details of the annihilation spectrum, and thus the underlying particle nature, it would be interesting to perform similar analyses also for other targets of potential DM annihilation. Another further direction of extending the present analysis would be to perform a full scan over the parameter space of viable models. Finally, we have stressed that the very concept of sensitivity of an IACT depends on the spectrum that is observed; in the context of DM searches, this is particularly important as DM annihilation spectra can significantly deviate from the usually assumed Crab-like spectrum. While we have provided a first estimate of how to proceed in such a case, it would be warranting to perform a dedicated analysis, using the full power of state-of-the-art MC tools, in order to accurately determine the importance of this effect.  Finally, we would like to mention that even in the case of negative detection, IACTs could in principle put interesting upper limits on the flux which in turn would translate into constraints on the combined space of astrophysical and particle physics parameters. Though much smaller than for other sources like, e.g., the galactic center, the main uncertainty in this case lies in the overall scale of the flux as determined by the details of the DM distribution. This, unfortunately, will therefore greatly obstacle any stringent constraint from null searches on the particle physics nature of DM for quite some time ahead.  To conclude, nearby dwarf galaxies -- and in particular Willman~1 -- are very interesting and promising targets for DM searches with the next generation of IACTs. An excellent performance of these experiments, in particular in terms of the sensitivity at energies slightly below the DM particle mass, will be paramount in such searches. In fact, given the low level of fluxes involved, a factor of 2 in sensitivity might decide whether a signal will be seen or not. Complementary to such demanding requirements on the experiments, the above discussion should also have made clear that  it will be very important to collect more astrophysical data and to improve the theoretical understanding of how DM is distributed in order to reduce the still unpleasantly large astrophysical uncertainties involved."
    },
    "0809/0809.4552_arXiv.txt": {
        "abstract": "{ This paper is the last in a series investigating low-redshift galaxy systems identified by the matched-filter technique in a moderately deep $I-$band survey. In this paper we present new redshifts for 747 galaxies in 23 ESO Imaging Survey (EIS) cluster fields. We use the ``gap''-technique to search for significant overdensities in redshift space for identifying groups/clusters of galaxies corresponding to the original EIS matched-filter cluster candidates. In this way we spectroscopically confirm systems in 10 of the 23 cluster candidate fields with a matched-filter estimated redshift $z_{MF}=0.3-0.4$ and with spectroscopic redshifts in the range from $z=0.158$ to $z=0.534$, with the observations favouring the confirmation of systems at the lower redshift end. After careful analysis of the redshift distribution, one system was split into two very close clumps in redshift space. We find that the systems identified in the present paper span a broad range of velocity dispersion and richness. The measured one-dimensional velocity dispersion range from $175\\mathrm{km/s}$ to $497\\mathrm{km/s}$, consistent with the values obtained in previous papers using much larger samples for systems over the same redshift range. Both undersampling and contamination by substructures contribute to the uncertainty of these measurements.  The richness range corresponds to clusters with an estimated total luminosity in the range $12L^*<L<65L^*$, but these estimates are very uncertain as are their relation to the velocity dispersion (mass) of the systems. From the analysis of the colours of the galaxy populations we find that $\\sim$60\\% of the spectroscopically confirmed systems have a ``significant'' red sequence.  We find that the colour of the red sequence galaxies matches passive stellar evolution predictions. With this paper we complete our spectroscopic survey of the fields of 58 EIS cluster candidates with estimated redshifts $z\\leq0.4$. We have measured a total of 1954 galaxy redshifts in the range $z=0.0065$ to $z=0.6706$. Of the 58 systems we confirm 42 ($\\sim$75\\%) with redshifts between $z=0.095$ and $z=0.534$.  } ",
        "introduction": "Clusters of galaxies have long been recognised as important targets for cosmology and astrophysics, both to constrain cosmological parameters and for evolutionary studies of large-scale structure and of cluster members. However, to carry out these studies large samples of clusters of galaxies with well-defined selection criteria, representative of the entire population of clusters and spanning a wide redshift range are required.  Over the past decades a number of systematic efforts to assemble catalogues of galaxy clusters \\cite[e.g. ][]{gunn86,postman96,scodeggio99, gladders01, gonzalez01, bahcall03, olsen07, vanbreukelen2006, koester2007} have been undertaken based on large optical and infrared imaging surveys in one or more passbands, using different search techniques. This has greatly expanded the available samples both at low and high redshifts, as compared to those detected in X-rays \\citep[e.g. ][]{rosati98a,bohringer00}. Unfortunately, systematic studies determining the yield of real bound systems from these image based searches has not progressed as fast, mainly because spectroscopic data covering large areas of the sky are difficult to assemble especially at intermediate and high redshifts. Furthermore, each detection technique is based on some different characteristic of the cluster such as its luminosity function, galaxy population or gas content. Therefore, it is important to understand the impact that these underlying assumptions may have on the detection and on the estimates of redshift and richness.  In order to be able to use cluster samples in statistical studies, especially to constrain cosmological parameters, it is important not only to detect the systems but also be able to establish how reliable estimates for the redshift and richness are, how well they relate to the underlying mass of the system and the selection function of the sample. To determine that requires both extensive spectroscopic follow-up and the use of realistic numerical simulations from which mock galaxy catalogues can be built.  Here we continue the effort of analysing the ESO Imaging Survey (EIS) Cluster Candidate Catalogue constructed using a matched-filter algorithm \\citep{olsen99a,olsen99b,scodeggio99}. A first estimate of the selection function for this catalogue based on simplistic simulations was given in \\cite{olsen00}.  In this paper we concentrate on the results of a comprehensive spectroscopic follow-up of these cluster candidates, covering a broad range of redshifts. While firm conclusions cannot be drawn due to the small size of the sample, the issues that may affect the use of galaxy clusters for statistical studies can be explored and will serve as a guide for further investigation based on on-going numerical simulations.  This paper is part of a series reporting the results obtained from a spectroscopic follow-up of selected cluster candidate fields drawn from the original EIS cluster candidate sample assembled by \\cite{olsen99a,olsen99b,scodeggio99}. This sample was constructed from the EIS $I-$band survey over 17~square degrees of the sky using the matched-filter technique. Out of 302 identified cluster candidates we have so far studied more than 50 candidate cluster fields chosen to contain candidates at different estimated redshift bins. So far we have probed fields with candidates at low \\citep[$z_{MF}\\sim0.2$, ][hereafter Paper~I, II and III]{hansen02,olsen03,olsen05}, intermediate \\citep[$0.5\\leq z_{MF}\\leq0.7$, ][]{ramella00} and high redshifts \\citep[$z_{MF}\\geq0.6$, ][]{benoist02, olsen05b} using different telescopes and spectrographs. In this paper we extend the previous work at the low-redshift end by carrying out spectroscopic observations of 23 fields around cluster candidates selected in the estimated redshift bin $0.3\\leq z_{MF}\\leq0.4$.   The paper is structured as follows: Sect.~\\ref{sec:obs} gives an overview of the observations and data reduction. Sect.~\\ref{sec:results} reviews the identification of systems in redshift space as well as the procedure adopted for associating the redshift groups to the EIS detections. Sect.~\\ref{sec:properties} describes the photometric properties of the spectroscopically confirmed systems including an analysis of the colour properties of the galaxy populations.  Finally, in Sect.~\\ref{sec:conclusions} a brief discussion of the conclusions of the results obtained in the present paper is presented, followed by a summary of the results of the entire series in Sect.~\\ref{sec:summary}. ",
        "conclusions": "  \\begin{itemize}  \\item The mean redshift of the confirmed sample is 0.24 higher than 0.18 measured in Paper~III. This indicates that we have successfully identified systems at higher redshift, even though at a lower rate, which as indicated throughout the paper was mostly due to a bright limiting magnitude of the spectroscopic data.  \\item Consistent with our findings in Paper~III, $\\sim$60\\% of the identified overdensities in redshift space with photometric data in two passbands show evidence for a red sequence among their spectroscopic members with a colour consistent with their measured redshift. This lends credence to the interpretation that these overdensities are indeed bound systems with a well-defined population of early-type galaxies. On the other hand, the remaining cases may not be totally discarded as they may still be spiral-rich bound systems. A final conclusion will require a considerable improvement in sampling, a hard task considering the angular extent of these systems.  \\item The distribution of velocity dispersions is consistent with that obtained by the much larger SDSS survey.  \\item The range of richness of our confirmed systems is in good agreement with those of \\citet{bahcall03} analysing a subset of the SDSS data using a similar implementation of the matched-filter algorithm. As previous authors we find that richness as estimated by the matched filter is not an adequate proxy for the mass of the system.  \\end{itemize}     In this series of four papers we have measured 1954 galaxy redshifts in the range 0.0065--0.6706 over 58 fields around EIS candidate clusters identified using the matched-filter algorithm and with estimated redshifts below $z_{MF}=0.4$. For a total of 42 cases we were able to associate the candidate with density enhancements in redshift space with mean redshifts between 0.095 and 0.534.  This represents a yield of $\\sim$75\\% for the matched-filter technique, which is consistent with the original estimates by \\cite{olsen99a} based on simulations and also comparable to the estimates reported by \\citet{kim02}. The method tends to overestimate the system redshift reaching an offset of 0.1 at $z_{MF}\\lesssim0.4$. The number of galaxies with concordant redshifts range from four to 35. The one-dimensional velocity dispersions of the identified density enhancements, which we have referred to as systems, vary from very low values up to $\\sim$1400~km/s, with the equivalent number of $L^*$ galaxies varying from 11 to 65. Due to the undersampling of the systems by our spectroscopic observations, these values, in general, have large errors and cannot by themselves distinguish between chance enhancements, small groups and rich clusters. Based on the results of this series we find that the identified clusters consist of substructures in more than 35\\% of the cases, impacting the matched-filter richness estimates.  For 34 of the confirmed groups and clusters colour information is available and was used to search for a red sequence. For 20 systems, corresponding to $\\sim$60\\%, we were able to detect one using the spectroscopically confirmed members.  While having a red sequence is not a necessary condition for a bound system, its existence definitely lends further credence to the detection. On the other hand, the matched-filter technique may detect systems with a large fraction of late-type galaxies or where the early-type galaxies are affected by recent starbursts that may occur during merging of groups or clusters. The confirmation of this idea will require more extensive spectroscopic surveys than the one carried out in the present program.  This series complements the work carried out by \\cite{ramella00,benoist02,olsen05b} who carried out spectroscopic observations of EIS candidate clusters at intermediate and high redshifts, obtaining confirmation yields, redshift offsets relative to the matched-filter estimates and velocity dispersions similar to those reported here.  Combining the results of this series with those of the previous papers, we have obtained roughly 2500 redshifts for a total of 74 cluster candidate fields out of 302 in the original candidate list, spanning the redshift range from 0.2 to 1.3, being perhaps one of the most extensive spectroscopic follow-up works of its kind. Taken altogether, these results show that the matched-filter technique leads to the detection of real density enhancements in redshift space. It should be used in conjunction with other cluster search techniques for the optimal identification of clusters of galaxies from optical and near-infrared imaging data. This is extremely important considering that large surveys such as UKIDSS, Dark Energy Survey and VISTA are either on-going or are envisioned to start in the next few years providing unprecedented samples of candidate clusters. Understanding the differences and possible biases of different detection techniques is particularly important for determining the completeness of the observed samples and be able to use, for instance, number counts of clusters of galaxies as a tool to constrain cosmological parameters as proposed by several new dark energy projects.  On-going and future surveys will provide considerable more information enabling the application of other detection algorithms based on for instance colours or photometric redshifts. However, regardless of the detection algorithm optical cluster surveys will always have to rely on extensive spectroscopic follow-up both for confirmation, membership assignment, subclustering and for fully characterising the systems.  Some of these issues are addressed in the present work."
    },
    "0809/0809.1694_arXiv.txt": {
        "abstract": "A small number of quasars exhibit interstellar scintillation on time-scales less than an hour; their scintillation patterns are all known to be anisotropic. Here we consider a totally anisotropic model in which the scintillation pattern is effectively one-dimensional.  For the persistent rapid scintillators J1819$+$3845 and PKS1257$-$326 we show that this model offers a good description of the two-station time-delay measurements and the annual cycle in the scintillation time-scale. Generalising the model to finite anisotropy yields a better match to the data but the improvement is not significant and the two additional parameters which are required to describe this model are not justified by the existing data. The extreme anisotropy we infer for the scintillation patterns must be attributed to the scattering medium rather than a highly elongated source. For J1819$+$3845 the totally anisotropic model predicts that the particular radio flux variations  seen between mid July and late August should repeat between late August and mid November, and then again between mid November and late December as the Earth twice changes its direction of motion across the scintillation pattern.  If this effect can be observed then the minor-axis velocity component of the screen and the orientation of that axis can both be precisely determined. In reality the axis ratio is finite, albeit large, and spatial decorrelation of the flux pattern along the major axis may be observable via differences in the pairwise fluxes within this overlap region; in this case we can also constrain both the major-axis velocity component of the screen and the magnitude of the anisotropy. ",
        "introduction": "There are three quasars which are known to exhibit large-amplitude radio flux variations on time-scales of less than an hour: PKS0405$-$385 (Kedziora-Chudczer et al 1997), whose variations are intermittent, and the persistent variables PKS1257$-$326 (Bignall et al 2003) and J1819$+$3845 (Dennett-Thorpe \\&\\ de~Bruyn 2000).  The short time-scale and the strong frequency dependence of the observed variations indicate that interstellar scintillation is the cause (Kedziora-Chudczer et al 1997). This conclusion is compellingly reinforced by the observed annual cycle in variability time-scale which is effected by the rotation of the Earth's orbital velocity vector (Dennett-Thorpe \\&\\ de~Bruyn 2003; Bignall et al 2003). Similarly the two-station time-delay measurements (Dennett-Thorpe \\&\\ de~Bruyn 2002; Bignall et al 2006) explicitly demonstrate the spatial modulation of flux which is inherent in scintillation. The annual cycle in the characteristic time-scale of the scintillations provides strong constraints on the velocity of the scattering material; in turn this constrains the distance to the scattering screen if one identifies the scintillation time-scale with the time taken to cross a Fresnel zone. Better constraints on the screen distance are obtained by constructing models and attempting to match them to all aspects of the data for a given source; the distances deduced in this way are surprisingly small -- of order 10pc from Earth (Dennett-Thorpe \\&\\ de~Bruyn 2003; Bignall et al 2003) -- so very high brightness temperatures are not required for the sources.  The same annual cycle in the scintillation time-scale also provides evidence for strong anisotropy in the scintillation pattern (Dennett-Thorpe \\&\\ de~Bruyn 2003; Bignall et al 2003). The scintillation time-scale is affected by both anisotropy in the scattering material and in the source structure. Quasars often exhibit elongated structure (``jets'') on milliarcsecond and larger scales (e.g. Walker, Benson \\&\\ Unwin 1987) so it is likely that they do have elongated structure on the sub-milliarcsecond angular scales relevant to interstellar scintillation. However, the spectral purity of the observed light-curves -- i.e. their quasi-sinusoidal nature -- argues that the scattering screens must also have anisotropic structure (Rickett, Kedziora-Chudczer \\&\\ Jauncey 2002; see also Bignall et al 2003), and that the major-to-minor axis ratio  is large ($> 4$). Unfortunately the spectral purity of the light-curves does not offer a sensitive test of the level of anisotropy once the axis ratio becomes large.  Because they are so nearby, yet they produce large amplitude variations, we know that the screens reponsible for intra-hour variability in quasars must be regions of very strong scattering -- i.e. they must scatter radio-waves through large angles -- and if the local interstellar medium is representative then similar, more distant screens could make a substantial contribution to the total scattering seen on other lines-of-sight. For compact radio quasars this suggests that smaller amplitude variability on longer time-scales should be relatively common and this expectation is qualitatively consistent with the results of flux monitoring of large numbers of compact extragalactic radio sources (Lovell et al 2003).  A further indication that similar screens are widely distributed in the Galaxy is the observation of  parabolic arcs in the ``secondary spectra'' (power spectra of the dynamic spectra) of radio pulsars (Stinebring et al 2001). Modelling of this phenomenon indicates that localised, strongly-scattering screens are responsible, that the scattering is anisotropic and that the screens are  hundreds of parsec distant from the Sun (Stinebring et al 2001; Cordes et al 2006; Walker et al 2004; Putney \\&\\ Stinebring 2006). It therefore seems that the three known intra-hour variable quasars can offer insights into the structure of the broader interstellar medium and it is important that we understand as much as we can about the properties of these nearby scattering screens.  As we have just noted, it is already clear that the major-to-minor axis ratio of the scintillation patterns is large for the three intra-hour variable quasars. Here we seek to determine whether existing data admit the possibility of the patterns being so anisotropic that the behaviour is effectively one-dimensional and, if so, whether that is the preferred model. In \\S2 we use published data on  the annual cycle in scintillation time-scale and the two-station time-delay experiments for both J1819$+$3845 and PKS1257$-$326 to test our models. We find that infinite anisotropy in the scintillation patterns is  consistent with and slightly preferred by existing data, and in \\S3 we interpret this result as one-dimensional scattering rather than the influence of anisotropy in the source structure. In \\S3 we further note that extreme flux-pattern anisotropy leads us to expect the J1819$+$3845 light-curves seen at one time of year to repeat at two other times, as the observer moves back-and-forth across the same region of the scintillation pattern. If the major-to-minor axis ratio of the pattern is $R\\la10^5$ then spatial  decorrelation along the major-axis direction may be measurable and flux monitoring of J1819$+$3845 can then be used to determine the length-scales and velocity components of the screen along both major-  and minor-axis directions. ",
        "conclusions": "At present there is no evidence for finite anisotropy in the scattering screens responsible for the intra-hour variations of PKS1257$-$326  and J1819$+$3845 and only a one-dimensional model is appropriate for each source.  In the case of J1819$+$3845 a strictly one-dimensional model predicts that the particular  scintillations seen between late August and mid November each year should also be seen before and after this period as the Earth's velocity across the flux pattern changes sign at these times and the observer moves back and forth across the same part of the pattern.  If repeating flux patterns can be identified in the data they will permit precise determination of the major-axis orientation and the minor-axis velocity component. To the extent that decorrelation of the pair-wise fluxes can be quantified we can measure the major-to-minor axis ratio, $R$, and the major-axis velocity component of the screen. Decorrelation is expected to be observable only if $R\\la 10^5$."
    },
    "0809/0809.3002_arXiv.txt": {
        "abstract": "The Large Area Multi-Object Fiber Spectroscopic Telescope (LAMOST) is a dedicated spectroscopic survey telescope being built in China, with an effective aperture of 4 meters and equiped with 4000 fibers. Using the LAMOST telescope, one could make redshift survey of the large scale structure (LSS). The baryon acoustic oscillation (BAO) features in the LSS power spectrum provide standard rulers for measuring dark energy and other cosmological parameters. In this paper we investigate the meaurement precision achievable for a few possible surveys: (1) a magnitude limited survey of all galaxies, (2) a survey of color selected red luminous galaxies (LRG), and (3) a magnitude limited, high density survey of $z<2$ quasars. For each survey, we use the halo model to estimate the bias of the sample, and calculate the effective volume. We then use the Fisher matrix method to forecast the error on the dark energy equation of state and other cosmological parameters for different survey parameters. In a few cases we also use the Markov Chain Monte Carlo (MCMC) method to make the same forecast as a comparison. The fiber time required for each of these surveys is also estimated. These results would be useful in designing the surveys for LAMOST. ",
        "introduction": "Baryon acoustic oscillations (BAO) before recombination left wiggling features in the matter power spectrum. At large scale where the evolution of the density perturbation is still linear, such features is preserved in the galaxy power spectrum \\citep{EH98,M99,EH97}. This provides us with a well calibrated standard ruler, which enables precise measurement of cosmological parameters, especially the dark energy parameters \\citep{EW04,E05a,SE03,SE07,K06,L03,GB07,BG03,B06}. Like cosmic microwave background (CMB), the physics involved is relatively clean and well understood, hence it is arguably less affected by unknown systematic errors, which might undermine empirical-rule based methods such as the type Ia supernovae (SNIa) and the cluster abundance measurements.  Recently, such a BAO feature have been observed in the 2dF galaxy redshift survey (2dFGRS) \\citep{2dF05} and the Sloan Digital Sky Survey (SDSS) survey data\\citep{E05b,H05a,H05b,T06,P07b,O07, CG08a,CG08b}. A number of more powerful BAO surveys are being planned. Some examples are WiggleZ \\citep{G07}, SDSS-3 (BOSS) \\footnote{{\\it http://cosmology.lbl.gov/BOSS/, http://www.sdss3.org/ }}, HETDEX \\footnote{{\\it http://www.as.utexas.edu/hetdex/}}, WFMOS \\citep{BNE05}, ADEPT \\footnote{{\\it http://www7.nationalacademies.org/ssb/be\\_nov\\_2006\\_bennett.pdf}}, Space/Euclid \\citep{C08} in the optical/IR and FAST \\footnote{{\\it http://www.bao.ac.cn/LT/}}, HSHS\\citep{PBP06}, SKA \\citep{AR05} \\footnote{{\\it http://www.skatelescope.org/}} in the radio.  In this paper, we make forecasts on cosmological constraints from BAO measurement in future surveys with the {\\it Large Sky Area Multi-Object Spectroscopic Telescope}, or LAMOST \\footnote{{\\it http://www.lamost.org/}}\\citep{C98,R08}, a telescope being built in China. The LAMOST is a 4 meter Schmidt telescope with a field of view of 20 $\\deg^2$. Equipped with 4000 optical fibers that are individually positioned by computer during observation, it is able to take spectra of 4000 targets simultanously. More details of the LAMOST telescope are described in Appendix A. The LAMOST is ideally suited for conducting large scale redshift surveys.  We expect that once LAMOST becomes operational, it will carry out several different survey projects. The design of such surveys should maximize the scientific output for a given amount of observing fiber time. Some obvious choices include a magnitude limited general galaxy redshift survey, and a magnitude limited quasar survey. For the purpose of BAO measurements, it is desirable to have additional surveys for targets at relatively higher redshift. Often found to be the bright central galaxies of galaxy groups and clusters, the luminous red galaxies (LRG) \\citep{E01,BD03,Z05} can be detected at higher redshifts than a typical galaxy for a given apparent magnitude. Quasars can be observed at even higher redshifts. In the SDSS surveys, quasars were very sparse, making it difficult to use in BAO measurement \\citep{S05,S07}. However, it is expected that the density of observable quasars with the LAMOST is much higher, since a lower limiting flux can be reached. Therefore, we also consider quasar samples in our study, with quasar themselves as the tracer of the large-scale structure (LSS). The absorption lines (Ly$\\alpha$ forest) in the quasar spectra also provide useful information \\citep{W03,M06,ME07}, which will be considered in our future work. We describe these surveys in more detail (including the selection of the sample, the required observing fiber hours, and the estimate of bias of the sample) in Appendix B.   We will make forecasts of the measurement errors on the dark energy equation of state and other cosmological parameters with these surveys, and also estimate the resources required for these surveys. We will first use the Fisher matrix method \\citep{SE03,SE07}. Then for a few cases we also use a Markov Chain Monte Carlo (MCMC) simulation to make the forecast. The advantages of the Fisher matrix method are that it can easily be used to explore large parameter space and that the relationship between the input and output is very clear. The MCMC method does not assume the likelihood to be Gaussian and thus can probe the full shape of the likelihood surface.    Previous studies have shown that it is feasible with the LAMOST to obtain better precisions in cosmological parameters than ongoing surveys like the SDSS \\citep{FCY00,SSF06,LXFZ08}. The study by \\cite{FCY00} were conducted before the importance of BAO in dark energy measurement was widely recognized. \\cite{SSF06} considered the Alcock-Paczynski test \\citep{AP79} in real space \\citep{MS03}, with a method different from the one we use.  In all of these studies, a single survey of galaxies up to $B=20.5$ were considered, assuming an total survey area of $15000 \\deg^2$, a total number of galaxies $10^7$, and a galaxy sample bias of 1.  In this paper, we try to make a {\\it realistic} assessment of LAMOST surveys. We consider different samples that could be collected by the LAMOST telescope within a reasonable amount of observation time. The surveyed sky area is assumed to be about $8000 \\deg^2$, for which the SDSS photometric survey catalogue is currently available for target selection. We use the published SDSS luminosity function to estimate the number of targets. In order to accurately estimate the measurement error for the given sample, it is necessary to know the bias of the sample. We use the halo model to estimate the bias of the samples.  It should be noted that in the methods employed here \\citep{SE03, SE07}, the treatment of the non-linear effect and redshift distortion is very crude. Furthermore, the assumption of scale independent clustering bias may be overly idealistic. Recently, \\cite{CS07} studied the BAO in both the power spectrum and correlation function by using the renormalized perturbation theory(RPT, \\cite{CS06a,CS06b}). They showed that mode-coupling typically leads to percent-level shifts in the acoustic peak of correlation function. \\cite{SSS08} have also shown with both numerical simulation and analytical calculation that the peak is moved by these effects. \\cite{A08} proposed a new method to extract the unbiased estimation of sound horizon scale based on a large N-body simulation and semi-analytical galaxy formation model. \\cite{S08} also presented high signal-to-noise ratio measurements of acoustic scale from large volume simulations, and obtained a robust measurements of acoustic scale with scatter close to that predicted by the methods utilized in this paper. There have also been works on scale dependent halo and galaxy bias \\citep{SSS07}. In analysing the real data to extract cosmological information, a more refined treatment is required to account for all these effects.   We review the methods for error forecasting in \\S~2. \\S~2.1 is devoted to the Fisher matrix method, which was developed by \\cite{SE03,SE07} for BAO measurements; \\S~2.2 is devoted to the MCMC method; in \\S~2.3 we discuss how to estimate the bias of the sample for the planned surveys. After a brief description of the data that we assume LAMOST could collect, we present our forecast on the BAO measurement precision in \\S~3, and conclude in \\S~4. In the appendices, we describe the characteristics of the LAMOST telescope (Appendix A), and the details of our survey design (Appendix B), including the generic main survey (Appendix B1), the LRG survey (Appendix B2), and the QSO survey (Appendix B3). We estimate the density and bias of each sample, and also discuss the fiber time required for completing each survey.   Throughout the paper, we adopt the flat $\\Lambda$CDM, with the WMAP three year \\citep{SB07} best fitted parameter values (~$\\Om_bh^2=0.0223, ~ \\Om_mh^2=0.128, ~ \\Om_k=0,~ h=0.73,~ w_0=-1,~ w_a=0,~ n_s=0.958,~ A_S=2.3 \\times 10^{-9},~ \\tau=0.092$~) as our fiducial model. ",
        "conclusions": "In this paper we forecast the measurement precision of dark energy equation of state from the BAO measurement expected from LAMOST surveys. We investigate three types of sample targets: magnitude limited surveys of all types of galaxies (main survey), a survey of color-selected luminous red galaxies (LRG), and surveys of quasars (QSO). For the main survey and the QSO survey, we consider a few different magnitude limits, all deeper than the SDSS spectroscopic survey. For the LRGs, we have used the MegaZ sample as our reference sample. The required fiber hours to complete each survey are also estimated (Table\\ref{tab:summary}).   For each sample, we predict the redshift distribution, and use the halo model to estimate their clustering bias. We then calculate the effective volume of each survey as a function of wavenumber $k$, the statistical errors in the measured matter power spectrum, and the resulting errors in distance scale measurements with the BAO used as standard rulers. The constraints on the dark energy equation of state parameters are derived using two methods, the Fisher matrix method and the MCMC method.   Our results yield useful insights on how to design the surveys. We find that the MAIN1 sample, which is one magnitude deeper than the SDSS main sample has an effective volume about three and half times larger than the SDSS main sample. Similarly, with the LAMOST LRGs, an effective volume of about four times of the SDSS LRG can be obtained. This effective volume is almost as large as the much more time-consuming MAIN2 sample (two magnitudes deeper than the SDSS main sample). The QSO samples can also yield high effective volume at large scales, but it declines rapidly at small scales. The improvement on the measurement of dark energy equation of state depends not only on the large scale effective volume, but also on the effective volume at relatively small scales. Our investigation shows that although the effective volume of the LAMOST LRG sample is comparable to the MAIN2 at large scales, the latter's figure of merit for dark energy measurement is greater. The QSO sample, on the other hand, generally has lower figure of merit. However, the number of targets and required observing fiber hours of the QSO samples are also much smaller than the galaxy samples, so it should not be regarded as a competing project of the galaxy surveys. The efficient design for using LAMOST survey to constrain dark energy parameters is to have a MAIN1 survey, an LRG survey supplemented by a QSO survey.   We also studied the impact of spectral resolution. The figure of merit improves by only about 1\\%-3\\% for the galaxy samples when one switches the spectral resolution from 1000 to 2000. For quasars the figure of merit improvement is about 15\\%. From the point of view of dark energy constraints, it is unnecessary to achieve a high spectral resolution, although high resolution spectra are useful for other scientific objectives.  Finally, we also used Monte Carlo simulations to make the same forecast to test the accracy of the Fisher matrix result. The figure of merits of the two estimate roughly agree, although the contours obtained from the Monte Carlo method are not exactly ellipses, but rather asymmetric.   Overall, the LAMOST surveys could achieve an figure of merit (as defined by \\cite{W08}) of 5-10 from each individual galaxy samples. For a constant $w$ dark energy model, the expected precision on $w$ is about 6.9\\%. For time-dependent $w$ models, the LAMOST survey would also help to significantly improve the constraints on the equation-of-state parameters, although the expected errors on $w_0$ and $w_a$ are not tight enough to pin down the nature of dark energy."
    },
    "0809/0809.1602.txt": {
        "abstract": "\\noindent{\\it  As several large scale interferometers are beginning to take data at sensitivities where astrophysical sources are predicted, the direct detection of gravitational waves may well be imminent. This would  open the gravitational-wave window to our Universe, and should lead to a much improved understanding of the most violent processes imaginable; the formation of black holes and neutron stars following core collapse supernovae and the merger of compact objects at the end of binary inspiral.   } ",
        "introduction": "Gravitational waves are ripples of spacetime generated as masses are accelerated. It is one of the central predictions of Einstein's general theory of relativity but despite decades of effort these ripples in spacetime have still not been observed directly. Yet we have strong indirect evidence for their existence from the excellent agreement between the observed inspiral rate of the binary pulsar PSR1913+16 and the theoretical prediction (better than 1\\% in the phase evolution). This provides confidence in the theory and suggests that ``gravitational-wave astronomy'' should be viewed as a serious proposition. This new window onto  universe will complement our view of the  cosmos  and will  help  us unveil the fabric of spacetime around  black-holes, observe directly the formation of black holes or  the merging  of  binary systems consisting  of  black holes  or neutron  stars,  search for rapidly spinning neutron  stars, dig deep into the very early moments of the origin of the universe, and look at the very center of the galaxies where  supermassive black holes weighing millions of  solar masses  are  hidden.   Secondly, detecting gravitational  waves is important for our understanding of the fundamental   laws   of  physics;   the proof   that gravitational waves exist will verify a fundamental 90-year-old prediction of general relativity. Also, by comparing the arrival times of light and gravitational waves from,  e.g., supernovae, Einstein's prediction   that    light    and gravitational waves  travel at the  same  speed  could  be checked.  Finally, we  could  verify that  they  have  the polarization predicted by general relativity.    These expectations follow from the comparison between gravitational and electromagnetic waves. That is : a) While electromagnetic waves are radiated by individual particles, gravitational waves are due to non-spherical bulk motion of matter. In essence, this means that the information carried by electromagnetic waves is stochastic in nature, while the gravitational waves provide insights into coherent mass currents. b) The electromagnetic waves will have been scattered many times. In contrast, gravitational waves interact weakly with matter and arrive at the Earth in pristine condition. This means that gravitational waves can be used to probe regions of space that are opaque to electromagnetic waves. Unfortunately, this weak interaction with matter also makes the detection of gravitational waves an extremely hard task. c) Standard astronomy is based on deep imaging of small fields of view, while gravitational-wave detectors cover virtually the entire sky. d) The wavelength of the electromagnetic radiation is  smaller than the size of the emitter, while  the wavelength of a gravitational wave is usually larger than the size of the source. This means that we cannot use gravitational-wave data to create an image of the source. In fact, gravitational-wave observations are more like audio than visual.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0809/0809.0237_arXiv.txt": {
        "abstract": "\\noindent The orbital dynamics of most planetary satellites is governed by the quadrupole moment from the equatorial bulge of the host planet and the tidal field from the Sun. On the Laplace surface, the long-term orbital evolution driven by the combined effects of these forces is zero, so that orbits have a fixed orientation and shape. The ``classical'' Laplace surface is defined for circular orbits, and coincides with the planet's equator at small planetocentric distances and with its orbital plane at large distances. A dissipative circumplanetary disk should settle to this surface, and hence satellites formed from such a disk are likely to orbit in or near the classical Laplace surface.  This paper studies the properties of Laplace surfaces. Our principal results are: (i) if the planetary obliquity exceeds $68.875^\\circ$ there is a range of semimajor axes in which the classical Laplace surface is unstable; (ii) at some obliquities and planetocentric distances there is a distinct Laplace surface consisting of nested eccentric orbits, which bifurcates from the classical Laplace surface at the point where instability sets in; (ii) there is also a ``polar'' Laplace surface perpendicular to the line of nodes of the planetary equator on the planetary orbit; (iv) for circular orbits, the polar Laplace surface is stable at small planetocentric distances and unstable at large distances; (v) at the onset of instability this polar Laplace surface bifurcates into two polar Laplace surfaces composed of nested eccentric orbits. ",
        "introduction": "\\noindent In his study of Jupiter's satellites, \\cite{lap05} recognized that the combined effects of the solar tide and the planet's oblateness induced a ``proper'' inclination in satellite orbits with respect to Jupiter's equator. He remarked that this proper inclination increases with the distance to the planet, and defines an orbital plane for circular orbits that lies between the orbital plane of the planet's motion round the Sun and its equator plane. This plane is called the Laplace plane.  More generally, the Laplace plane is usually defined as the plane normal to the axis about which the pole of a satellite's orbit precesses\\footnote{Unfortunately, the term is sometimes also applied to the invariable plane, the plane perpendicular to the total angular momentum of an $N$-body system.}. The ``Laplace surface'' is the locus of all orbits that do not precess (i.e., the secular motion of the node and apse vanishes).  In the most common situation, we consider circular satellite orbits around an oblate planet with non-zero obliquity, traveling around the Sun. The Laplace surface is then determined by the competition between the interior quadrupole potential from the equatorial bulge and the external quadrupole potential from the Sun. Close to the planet, the ``classical'' Laplace surface nearly coincides with the planetary equator, while at large distances it nearly coincides with the planetary orbital plane. The transition between these two orientations occurs near the ``Laplace radius,'' defined below in equation (\\ref{eq:laprad}).  The Laplace surface is important because it traces the shape expected for a thin gas disk or dissipative particulate ring surrounding the planet. Thus it is not surprising that many planetary satellites orbit close to the Laplace surface; those that do probably formed from a circumplanetary gas disk while those that do not were presumably either captured from heliocentric orbits or experienced unusual events in their past history.  The purpose of this paper is to study the properties of the Laplace surface, including the stability of is generating orbits and its generalization to eccentric orbits. Although the Laplace surface has been known and studied for over two centuries, we believe that many of the results we present are new. ",
        "conclusions": "\\noindent We have examined the orbital dynamics of planetary satellites under the combined influence of the quadrupole moment from the planet's equatorial bulge and the tidal field from the Sun.  The Laplace equilibria are orbits in which the secular evolution due to these forces is zero. They represent the orbits in which planetary satellites formed from a circumplanetary gas disk should be found.  Laplace equilibria exist for both circular and eccentric orbits. At a given semimajor axis, the orbit normals to the circular Laplace equilibria lie along three orthogonal directions, two of them in the principal plane defined by the planet's spin axis and the normal to its orbit around the Sun (the ``circular coplanar'' equilibria). One of the two circular coplanar equilibria is always unstable. The warped surface swept out by the other circular coplanar equilibrium orbits as the semimajor axis varies is called the classical Laplace surface (eq.\\ \\ref{eq:sola}). The classical Laplace surface coincides with the equator of the planet at small distances and with the orbital plane of the planet at large distances. Orbits in the classical Laplace surface are stable if the planetary obliquity $\\phi_\\odot<68.875^\\circ$, while in the range $68.875^\\circ < \\phi_\\odot < 90^\\circ$ there is a range of semimajor axes near the Laplace radius $r_L$ (eq.\\ \\ref{eq:laprad}) in which the surface is unstable. The third Laplace equilibrium for circular orbits (the ``orthogonal'' or ``polar'' equilibrium) corresponds to orbits that cross over the planet's pole, and these are stable if and only if the semimajor axis is less than $2^{-1/5}r_L$ (eq.\\ \\ref{eq:al}).  At some obliquities and semimajor axes eccentric Laplace equilibria also exist. In the ``coplanar-coplanar'' equilibria both the eccentricity and angular-momentum vectors lie in the principal plane (Figures \\ref{fig:stablecc1} and \\ref{fig:coco}). The ``orthogonal-coplanar'' equilibria are stable, eccentric polar orbits that bifurcate from the circular orthogonal equilibria at the semimajor axis $2^{-1/5}r_L$ where these become unstable. Other eccentric equilibria exist but are unstable.  The use of the secular equations of motion (\\ref{eq:milank}) should be legitimate so long as the precession times for the eccentricity and angular-momentum vector are much longer than the orbital period of the satellite. A sufficient condition for this is that the satellite semimajor axis is much larger than the planetary radius and much smaller than the planet's Hill radius (cf.\\ Table \\ref{tab:outer}). A possible limitation of the secular equations is that they neglect the evection resonance, where the apsidal precession rate equals the mean motion of the planet around the Sun. For the outer planets, the evection resonance typically occurs at $\\sim 0.2r_L$. \\cite{tw98} have stressed the role of the evection resonance in the history of the lunar orbit.  Several unanswered questions remain: (i) What is the structure of the four-dimensional phase space of the equations of motion (\\ref{eq:milank}), and what fraction of the orbits are chaotic? These issues could be explored through surfaces of section, although the exploration would be laborious since there is a separate surface of section for each value of the Hamiltonian, semimajor axis, and planetary obliquity. (ii) What is the nature of the evolution of satellites in the classical Laplace surface that migrate into an unstable region (Fig.\\ \\ref{fig:orbit})? (iii) What are the shape and properties of dissipative disks in the range of semimajor axes where the classical Laplace surface is unstable? Does the dissipation stabilize the classical Laplace surface? Does the disk follow the eccentric coplanar-coplanar Laplace surface? Or is there no steady-state disk? Hints that eccentric structures can occur in dissipative disks include the eccentric structures seen in the planetary ring systems of Uranus and Saturn, debris disks around young stars, and galactic nuclei such as M31. (iv) What is the analog of the Laplace surface in accretion disks around black holes, where Lense--Thirring precession rather than a quadrupole potential is the dominant non-Keplerian perturbation from the central body?"
    },
    "0809/0809.5138_arXiv.txt": {
        "abstract": "{ In some galaxies, the central velocity dispersion, $\\sigma$, is depressed with respect to the surroundings. This sigma-drop phenomenon may have different physical origins, bearing information about the internal dynamics of the host galaxy. In this article, we stress the importance also of observational artifacts due to the $\\sigma$-metallicity degeneracy: when a spectrum of a population is compared with a template of miss-matched metallicity, the velocity dispersion may be wrongly estimated. A sigma-drop may appear in place of a metallicity peak.  The discussion is illustrated using VLT/FORS spectra of diffuse elliptical galaxies. Some of the sigma-drop galaxies reported in the literature may be analysis artifacts.  } ",
        "introduction": "In a fraction of galaxies, mostly spirals but also ellipticals, the velocity dispersion is depressed in the center \\citep[e. g.][]{comeron08}. However, simple density profile models, like the de Vaucouleurs profile, predict a continued increase inwards. In the very center, approaching the central massive black hole, a further steep increase is even expected \\citep{DR90}.  \\begin{figure} \\includegraphics[width=8cm]{fig1.eps} \\caption{ Monte Carlo simulations and $\\chi^2$ map. The bottom panel shows a $\\chi^2$ map projecting the parameter space on the $\\sigma$ vs. metallicity plane. For each point of a grid, the other free parameters, cz, age and multiplicative polynomial are optimize and the map shows the variations of the resulting  $\\chi^2$ values. The iso-$\\chi^2$ contours are elongated along a direction mixing the two parameters. The top panel shows a Monte-Carlo experiment, simulating the effect of a Gaussian noise comparable to the FORS observations (S/N=70). Each of the 3000 points represent the analysis of a different realization of the noise. The cloud is elongated in a diagonal direction, showing the $\\sigma$ - metallicity degeneracy. } \\label{fig:1} \\end{figure}    These observed $\\sigma$-drops may have various physical explanations. They may indicate the absence of a central mass concentration \\citep{DR90} or more generally a flat density core. The latter is certainly expected in diffuse elliptical galaxies (dE) with shallow density profiles, flatter than exponential. Such examples may be found in \\citet{SP02}. Still, in the list of $\\sigma$-drop galaxies, the spiral galaxies are the majority and since the stellar kinematics is less studied in spirals it is quite probable that most of the $\\sigma$-drops occur in spirals. Some studies suggest up to 50 \\% of $\\sigma$-drops among the spirals \\citep{comeron08}.  In spiral galaxies, a $\\sigma$-drop may result from a young stellar population formed from a rotating gas disk funneled toward the center by the dynamical action of a bar \\citep{wosniak03}.  Simulations show that this feature may persist 1 Gyr until the newborn population mixes with the underlying population and its luminosity fades. A sustained gas supply and star formation may produce a long-lived cold stellar disk.  In this article we investigate another phenomenon affecting the detection of $\\sigma$-drop galaxies: the $\\sigma$-metallicity degeneracy. ",
        "conclusions": "We have shown that some $\\sigma$-drops may be artifacts of the analysis method. They may result from a peak in the metallicity, which is actually often expected in the centers of the galaxies.  This naturally does not mean that all the $\\sigma$-drops are of this nature, but that individual observations must be examined with caution. One has to ask the questions: (i) Is the analysis made with a fixed-metallicity template, or with an additive polynomial? (ii) Is the magnitude of the effect comparable with the sigma-metal\\-li\\-ci\\-ty degeneracy?  In the literature, we found two samples of $\\sigma$-drop galaxies:~\\citet{emsellem08} and \\citet{comeron08}.  Half of the 20 $\\sigma$-drop galaxies were detected in various studies using fixed-metallicity templates and the magnitude of the drops are generally consistent with them resulting from the sigma-metallicity degeneracy. Some of these studies are reducing the mismatch with an additive term, but it is difficult to determine their reliability without testing the actual analysis program.  Only recent studies, like \\citet{dumas07}, \\citet{emsellem01}, and some other based on SAURON data use an optimal template fit to each single spectrum. Unfortunately they use as a reference the \\citet{jones98} stellar library which has a limited coverage in metallicity and may therefore not match perfectly the observations.  We conclude that the fraction of $\\sigma$-drop is certainly\\break overestimated, and that the $\\sigma$-metallicity degeneracy should not be neglected.  Fitting in the same time the kinematics and the chemical characteristics of the population is definitely the best solution to correctly reconstruct stellar population properties."
    },
    "0809/0809.0001_arXiv.txt": {
        "abstract": "We present an 8.5-hour simultaneous radio, X-ray, UV, and optical observation of the L dwarf binary 2MASSW J0746425+200032. We detect strong radio emission, dominated by short-duration periodic pulses at 4.86 GHz with $P=124.32\\pm 0.11$ min.  The stability of the pulse profiles and arrival times demonstrates that they are due to the rotational modulation of a $B\\approx 1.7$ kG magnetic field.  A quiescent non-variable component is also detected, likely due to emission from a uniform large-scale field.  The H$\\alpha$ emission exhibits identical periodicity, but unlike the radio pulses it varies sinusoidally and is offset by exactly 1/4 of a phase.  The sinusoidal variations require chromospheric emission from a large-scale field structure, with the radio pulses likely emanating from the magnetic poles.  While both light curves can be explained by a rotating mis-aligned magnetic field, the 1/4 phase lag rules out a symmetric dipole topology since it would result in a phase lag of 1/2 (poloidal field) or zero (toroidal field).  We therefore conclude that either (i) the field is dominated by a quadrupole configuration, which can naturally explain the 1/4 phase lag; or (ii) the H$\\alpha$ and/or radio emission regions are not trivially aligned with the field. Regardless of the field topology, we use the measured period along with the known rotation velocity ($v{\\rm sin}i\\approx 27$ km s$^{-1}$), and the binary orbital inclination ($i\\approx 142^\\circ$), to derive a radius for the primary star of $0.078\\pm 0.010$ R$_\\odot$. This is the first measurement of the radius of an L dwarf, and along with a mass of $0.085\\pm 0.010$ M$_\\odot$ it provides a constraint on the mass-radius relation below 0.1 M$_\\odot$.  We find that the radius is about $30\\%$ smaller than expected from theoretical models, even for an age of a few Gyr.  The origin of this discrepancy is either a breakdown of the models at the bottom of the main sequence, or a significant mis-alignment between the rotational and orbital axes. ",
        "introduction": "\\label{sec:intro}  Radio observations conducted over the past several years have uncovered a substantial fraction of magnetically active low mass stars and brown dwarfs.  Both quiescent and flaring emission are present, with luminosities that remain unchanged down to at least spectral type $\\sim {\\rm L3}$, in contrast to the declining activity seen in X-rays and H$\\alpha$ \\citep{bbb+01,ber02,brr+05,bp05,ber06,ohb+06,had+06,adh+07,pol+07,hbl+07,aob+07,adh+08,bgg+08,bbg+08,had+08}. Depending on the nature of the radio emission mechanism, the inferred magnetic field strengths are $\\sim 0.1-3$ kG with order unity filling factors \\citep{ber06,had+08}.  Long term monitoring of several objects in the spectral type range M8--L3.5 has further shown that the fields are generally stable for at least several years (\\citealt{brr+05,ber06,bgg+08}; but see also \\citealt{adh+07}), providing an important constraint on the lifetime of magnetic dynamos in fully convective objects.  The strength, scale, and stability of the inferred fields is in good agreement with recent results from phase-resolved spectropolarimetry (spectral type $\\sim {\\rm M4}$; \\citealt{dfc+06,mdp+08}) and Zeeman broadening of FeH lines (spectral type $<{\\rm M9}$; \\citealt{rb06,rb07}), as well as with the most recent numerical dynamo simulations (e.g., \\citealt{bro08}).  Equally important, three ultracool dwarfs to date have been observed to produce periodic radio emission, with periods of 184 min (2MASS J00361617+1821104; \\citealt{brr+05}), 118 min (\\tvlm\\footnotemark\\footnotetext{This object also exhibits periodic H$\\alpha$ emission, with the same period as observed in the radio \\citep{bgg+08}.}; \\citealt{hbl+07}), and 170 min (LSR\\,1835+3259; \\citealt{had+08}).  In the case of 2M\\,0036+18 the periodic emission is sinusoidal, while in the latter two objects it is in the form of short duration pulses (duty cycle of a few percent).  In all three cases the observed periods are in good agreement with the known rotation velocities ($v{\\rm sin}i$), indicating that the radio periodicity traces the stellar rotation.  As a result, in addition to allowing a measurement of the magnetic field properties, radio observations provide a unique opportunity to measure the radii of late-M and L dwarfs through the combination of rotation period and velocity.  However, since only $v{\\rm sin}i$ values are known for these objects, there is an inherent degeneracy between the radius and the inclination of the rotation axis, $i$. Using a typical range of $R\\approx 0.09-0.11$ R$_\\odot$ for late-M and L dwarfs, the allowed range of inclinations span a relatively wide range of $\\sim 60-90^\\circ$ \\citep{brr+05,had+08}.  To break this degeneracy, and thus measure the radius directly, it is desirable to observe objects for which the inclination can be estimated.  This is the case for binary systems with well determined orbital parameters if we make the reasonable assumption that the rotation and orbital axes are aligned \\citep{hal94}.  Along with an estimate of the mass, we can thus place constraints on the mass-radius relation for ultracool dwarfs.  Here we present simultaneous radio, X-ray, optical, and UV observations of one such system, the L dwarf binary 2MASSW J0746425+200032 (hereafter, \\2m0746).  These observations are part of a long-term project to study the field properties of ultracool dwarfs \\citep{brr+05,bgg+08,bbg+08}.  While no X-ray or UV emission are detected, the radio and H$\\alpha$ emission exhibit clear periodicity, which along with the known orbital inclination and $v{\\rm sin}i$ allow us to explore for the first time the radius and magnetic field topology of an L dwarf. ",
        "conclusions": "\\label{sec:conc}  We presented simultaneous radio, X-ray, optical, and UV observations of the L dwarf binary \\2m0746.  Both the radio and H$\\alpha$ light curves are dominated by periodic emission, with $P\\approx 124$ min. The agreement between the two periods indicates that both arise from the same binary member, which we assume to be the primary star.  We note that high angular resolution Very Long Baseline Array observations may pinpoint the origin of the radio emission.  The observed period is well-matched to the rotation velocity of \\2m0746, and the assumption that the rotation axis is aligned with the orbital inclination allows us to infer a radius of $R=0.078\\pm 0.010$ R$_\\odot$.  Combined with the estimated mass of $0.085\\pm 0.010$ M$_\\odot$ \\citep{bdk+04}, we find that the radius is about 30\\% smaller than predicted from theoretical evolutionary models.  It is unclear from this single object whether this is the result of a breakdown in the models, or an indication that the rotation and orbital axes are severely misaligned.  Both possibilities have important implications for our understanding of ultracool dwarfs and their formation mechanism.  The combination of periodic radio and H$\\alpha$ emission further allows us to explore the magnetic field topology.  The radio periodicity is in the form of short duration pulses with a duty cycle of $0.8\\%$ and a uniform sense of circular polarization, while the H$\\alpha$ emission is sinusoidal and leads the radio emission by 1/4 of a phase.  These emission properties can be explained in the context of a rotating mis-aligned magnetic field, but the observed phase lag rules out a symmetric dipole field.  Instead, we conclude that either the field is dominated by a quadrupole configuration, or the emission regions are not trivially aligned with a dipole field.  In both scenarios the quiescent and non-variable radio emission most likely arises from the largest scales of the field, which are least affected by rotation and inclination effects.  We note that the availability of simultaneous radio and optical data is crucial to our understanding of the field geometry.  If only one band was available to us (as in all previous cases of detected radio pulses), we would have concluded that the dipole topology is the most likely scenario.  Instead, the field configuration of \\2m0746\\ appears to be somewhat more complex than those inferred for mid-M dwarfs from phase-resolved optical spectropolarimetry \\citep{dfc+06,mdp+08}, although the vast difference in techniques may result in sensitivity to different field structures.  Regardless of the exact configuration, our observations support the general conclusions of recent numerical dynamo simulations \\citep{bro08}, which predict large-scale fields at low Rossby numbers.  Indeed, we estimate that the Rossby number for \\2m0746\\ is very low, $\\sim 4\\times 10^{-4}$.  Future observations of \\2m0746\\ will allow us to determine the stability timescale of the magnetic field configuration, a crucial constraint on dynamo models.  Previous observations of the L3.5 dwarf 2M\\,0036+18 \\citep{brr+05,had+08} indicated stability on a timescale of at least $\\sim 3$ yr.  In addition, continued surveys for radio emission from low mass stars and brown dwarfs, particularly in binary systems, are warranted.  As demonstrated here, these observations, particularly in conjunction with optical spectroscopy, can provide important constraints not only on the mass-radius relation below $\\sim 0.1$ M$_\\odot$, but also on the magnetic field topology of fully convective objects."
    },
    "0809/0809.0692_arXiv.txt": {
        "abstract": "Naively, one would expect longer papers to have larger impact (i.e., to be cited more). I tested this expectation by selecting all ($\\sim 30,000$) refereed papers from A\\&A, AJ, ApJ and MNRAS published between 2000 and 2004. These particular years were chosen so papers analyzed would not be too ``fresh'', but at the same time length of each article could be obtained via ADS. I find that indeed longer papers published in these four major astronomy journals are on average cited more, with a median number of citations increasing from 6 for articles 2--3 pages long to about 50 for articles $\\sim 50$ pages long.  I do however observe a significant ``Letters effect'', i.e. ApJ and A\\&A articles 4 pages long are cited more than articles 5--10 pages long. Also, the very few longest ($>80$ pages) papers are actually cited less than somewhat shorter papers.  For individual journals, median citations per paper increase from 11 for $\\sim 9,300$ A\\&A papers to 14 for $\\sim 5,300$ MNRAS papers, 16 for $\\sim 2,550$ AJ papers, and 20 for $\\sim 12,850$ ApJ papers (including ApJ Letters and Supplement). I conclude with some semi-humorous career advice, directed especially at first-year graduate students. ",
        "introduction": "There have been a number of publications analyzing patterns in citations in the astronomical literature, such as ``Productivity and impact of astronomical facilities: Three years of publications and citation rates'' (Trimble \\& Ceja 2008), ``The Importance of Being First: Position Dependent Citation Rates on arXiv:astro-ph'' (Dietrich 2007), ``Demographic and Citation Trends in Astrophysical Journal Papers and Preprint'' (Schwarz \\& Kennicutt 2004) or ``Patterns in Citations of Papers by American Astronomers'' (Trimble 1993), just to mention a few.  Scanning abstracts of these publications, some career advice emerges:  \\begin{enumerate}  \\item{} ``There are hot topics (cosmology, exoplanets) and not so hot topics (binary stars, planetary nebulae)'' and ``service to the community, in the form of catalogues and mission descriptions, is rewarded, at least in citation numbers, if not always in other ways.'' (Trimble \\& Ceja 2008)  \\item{} ``We find that e-prints appearing at or near the top of the {\\tt astro-ph} mailings receive significantly more citations than those further down the list. This difference is significant at the 7 sigma level and on average amounts to two times more citations for papers at the top than those further down the listing.'' (Dietrich 2007)  \\item{} ``Papers posted on {\\tt astro-ph}  are cited more than twice as often as those that are not posted on {\\tt astro-ph}.'' (Schwarz \\& Kennicutt 2004)  \\item{} ``It still pays to be mature, prizewinning theorist, working on cosmology or high-energy astrophysics at a prestigious institution.'' (Trimble 1993)  \\end{enumerate}  Some of that advice might be easier to implement than other, and sometimes a finely timed, Wednesday afternoon 04:00:03 p.m. EST submission to {\\tt astro-ph}  produces less than desired effect (e.g., Stanek 2003, private communication).  However, to my knowledge, there has been no discussion of how long a paper should be to maximize its impact. Is it better to write several shorter papers or one longer paper? And serving as a referee, one of the standard questions is ``Can this paper be made shorter?'' But maybe it should be made longer? We have all heard a famous quote, often attributed to Mark Twain: ``If I Had More Time I Would Write a Shorter Letter,'' but actually from Blaise Pascal: ``Je n'ai fait celle-ci plus longue que parce que n'ai pas eu le loisir de la faire plus courte.'' Is not at all clear this is a valid advice when writing astronomical papers.  To alleviate that glaring omission, I have decided to investigate the citation vs.~paper length correlation for astronomical papers. In Section 2 I describe the data, namely citation and page length statistics for about 30,000 refereed astronomical papers published between 2000 and 2004 in ApJ, A\\&A, AJ and MNRAS. In Section 3 I discuss the results, and in Section 4 I conclude with some career advice, especially useful for first-year graduate students in astronomy.  \\addtocounter{footnote}{1} ",
        "conclusions": ""
    },
    "0809/0809.5248_arXiv.txt": {
        "abstract": "We construct a secular theory of a coplanar system of $N$-planets not involved in strong mean motion resonances, and which are far from collision zones. Besides the point-to-point Newtonian mutual interactions, we  consider the general relativity corrections to the gravitational potential of the star and the innermost planet, and also a modification of this potential by the quadrupole moment and tidal distortion of the star. We focus on hierarchical planetary systems.  After averaging the model Hamiltonian with a simple algorithm making use of very basic properties of the Keplerian motion, we obtain analytical secular theory of the high order in the semi-major axes ratio. A great precision of the analytic approximation is demonstrated with the numerical integrations of the equations of motion.  A survey regarding model parameters (the masses, semi-major axes, spin rate of the star) reveals a rich and non-trivial dynamics of the secular system. Our study is focused on its equilibria. Such solutions predicted by the classic secular theory,  which correspond to aligned (mode~I) or anti-aligned (mode~II) apsides, may be strongly affected by the gravitational corrections. The so called true secular resonance, which is a new feature of the classic two-planet problem discovered by Michtchenko \\& Malhotra (2004), may appear in other, different regions of the phase space of the generalized model. \\corr{ We found bifurcations of mode~II which emerge new, yet unknown in the literature, secularly unstable equilibria and a complex structure of the phase space. These equilibria may imply secularly unstable orbital configurations even for initially moderate eccentricities. } The point mass gravity corrections can affect the long term-stability in the secular time scale,  which may directly depend on the age of the host star through its spin rate. We also analyze the secular dynamics of the $\\upsilon$~Andromede system in the realm of the generalized model. Also in this case of the three-planet system, new secular equilibria may appear. ",
        "introduction": "The currently known sample of extrasolar planets\\footnote{See www://exoplanet.eu, http://exoplanets.org, http://exoplanets.eu} comprises of many multiple-planet configurations. Their architectures are diverse and, usually, very different from the Solar system configuration. Some of them consist of so called hot-Jupiters or hot-Neptunes, with  semi-major axes $\\sim 0.05$~au and the orbital periods of a few days.  The multiple-planet systems may also contain companions in relatively distant orbits. Such systems are non-resonant (or hierarchical).  The discovery of multi-planet systems and their unexpected orbital diversity initiated interest in their secular evolution. Generally, the analysis of secular interactions, including tidal effects, general relativity (GR), and quadrupole moment (QM) corrections to the Newtonian gravitation (NG) of point masses, follows the theory of compact multiple stellar systems \\citep{Mardling2002,Nagasawa2005}.  The secular evolution  of planetary systems has been already intensively studied in this context by many authors. For instance, \\cite{Adams2006a} consider the GR corrections in the sample of detected extrasolar systems. They have shown that the GR interactions can be particularly important for the long-term dynamics of  short-period planets. Also \\cite{Adams2006b,Mardling2007} study the evolution of such objects taking into account the tidal circularization of their orbits, and secular excitation of the eccentricity of the innermost planet by more distant companions. These secular effects imply constraints on the orbital elements of the innermost orbits, including yet undetected Earth-like planets, which are expected to be found in such systems. The  quadrupole moment  and relativistic effects may detectably affect the light-curves of stars hosting transiting planets in a few years time-scale  \\citep{MiraldaEscude2002,Agol2005,Winn2005}.  In  general, the interplay of apparently subtle perturbations with the  point mass Newtonian gravity depends on many physical and orbital parameters, and it may lead to non-trivial and interesting dynamical phenomena. \\corr{Still, the problem is of particular importance for studies of the long-term stability of the Solar system \\citep[see very recent works by][]{Benitez2008,Laskar2008}.}  Because a fully general study of such effects would be a very difficult task, we introduce some simplification to the planetary model. We skip tidal effects caused by dissipative interactions of extended star and planetary bodies, in particular  the tidal friction (TF) damping the eccentricity and the tidal distortion (TD) that modifies planetary figures,  and could influence their gravitational interaction with the parent star. The tidal effects are typically much smaller than the leading Newtonian, relativistic and quadrupole moment contributions \\citep{Mardling2007}. Also their time-scale is usually much longer than of the most  significant conservative effects.    In this work, we focus on the secular dynamics of planetary systems over an intermediate time scale relevant to the  conservative perturbations.  Our goal is to obtain a qualitative picture of the secular system that may be useful as the first order approximation to more general model of the long-term planetary dynamics.  Moreover, we show in this paper that even if we skip the non-conservative  tidal effects then the GR and/or QM {\\em corrections} to the point mass Newtonian gravity can be much more significant for the secular dynamics than the mutual NG interactions alone. These effects may change qualitatively  the view of the dynamics of planetary systems  as predicted by the Laplace-Lagrange theory \\citep{Murray2000} or its recent versions \\citep[e.g.,][]{Lee2003,Michtchenko2004,Michtchenko2006,Libert2005,Veras2007,Migaszewski2008a}.  The model without dissipative effects may be investigated  with the help of conservative Hamiltonian theory. Assuming that planetary orbits are well separated and the system is far from mean motion resonances and collision zones, we can apply the averaging proposition \\citep{Arnold1993} to derive the long term evolution of their orbital elements. This approach can be classified among secular theories having a long history in the context of the Solar system \\citep{Brouwer1961,Murray2000}.  The present work is a step towards a generalization of the secular model of coplanar 2-planet system by \\cite{Michtchenko2004} and its analytical version applied to $N$-planet system in \\citep{Migaszewski2008a}.  To explore the phase space globally without restrictions on eccentricities, we simplify the equations of motion through the averaging of perturbations to the Keplerian motion, with the help of semi-numerical method proposed by \\cite{Michtchenko2004}. To obtain analytical results, we also use  a very simple averaging algorithm that can be applied to perturbations dependent on the mutual distance of interacting bodies \\citep{Migaszewski2008a}. These works demonstrate that the secular evolution can be precisely described in wide ranges of the orbital parameters, including eccentricities up to 0.8--0.9, for well separated (hierarchical)  orbits with small ratio of the semi-major axes, $\\alpha\\sim 0.1$.  Here, the secular NG theory is very helpful to derive the more general model including the GR and QM  effects. Basically, without the averaging, the only possibility of investigating the long-term secular dynamics relies on numerical solutions of the equations of motion \\citep{Mardling2002}. However, due to extremely different time scales which are related as days (the orbital periods of inner planets) to $10^3$--$10^6$~years of the secular orbital evolution, the numerical integrations are of very limited use when we want to investigate large volumes of initial conditions rather than isolated orbits. In such a case,  the analytical theory can be very helpful to get a deep insight into the secular dynamics. Moreover, when necessary, particular solutions can be studied in detail with the help of the direct numerical integrations.  The plan of this paper is the following. In Sect.~2 we formulate the generalized model of a coplanar, $N$-planet system, accounting for the relativistic and quadrupole moment corrections to the NG-perturbed motion of the innermost planet. To make the paper self-consistent, we average out the perturbations to the Keplerian motion  with the help of the averaging algorithm described in \\citep{Migaszewski2008a}. In Sect.~3 we apply the analytical and numerical tools to study the influence of the GR and QM corrections on the secular evolution. In particular, we recall the concept of the so called representative  plane of initial conditions \\citep{Michtchenko2004}. We focus on the  search for stationary solutions in the averaged and reduced systems (periodic orbits in the full systems), and we investigate their stability and bifurcations with the help of phase diagrams.  In this section we also derive interesting conclusions on the stability of the unaveraged systems. In Sect.~4, we study the phase space of two-planet systems in wide ranges of parameters governing their orbital configurations (masses, semi-major axes ratios, eccentricities) and physical parameters (flattening of the star). To illustrate the application of the secular theory to multi-planet configurations, we consider the three-planet $\\upsilon$~Andr system, and we investigate the QM (stellar rotation) influence on its secular orbital evolution. ",
        "conclusions": "In this work, we consider a generalized secular theory of coplanar, $N$-planet system. Extending the model analyzed in the recent works devoted to the secular planetary dynamics  with mutual Newtonian point-to-point interactions \\citep[e.g.,][]{Michtchenko2004,Libert2005,Gallardo2005,Migaszewski2008a},  we consider the influence of the general relativity and quadrupole moment of the parent star on the secular dynamics of the innermost planet and stability of the whole planetary system. In general, these corrections to the classic model still do not cover all physics governing the dynamics of such systems. In some cases (for instance, of the short-period hot-Jupiters), the tidal, dissipative torques acting between the inner planet and the star may be significant for the orbital evolution. However,  our main  goal is rather to extend the classic model with the perturbations that are conservative and may be well modeled in the realm of the Hamiltonian mechanics than to build a complete, general secular theory. Still, this approach is useful to a wide class of systems, when the tidal interactions may be regarded as secondary effects, or are acting during much longer characteristic time-scale than the GR and QM perturbations. In reward, for paying the price of less general model, we may investigate the secular dynamics in a global manner.  Our analytic model follows assumptions required by the averaging theorem. Technically, the averaging has been done with  the help of a very simple method. This algorithm relies on  appropriate change of integration variables. It does not incorporate any classic Fourier expansion of the perturbing function.  We obtain a very precise analytic model of the coplanar, $N$-planet system in terms of the semi-major axes ratio. It can be regarded as a generalization of the recent analytic secular theories  of the classic model investigated in many recent papers \\citep[e.g.,][]{Ford2000,Lee2003,Michtchenko2004,Libert2005}. On the other hand,  our work also covers, to some extent, the global dynamics  of the generalized model studied in \\cite{Mardling2002,Nagasawa2005} with the help of the Gauss/Lagrange planetary equations of motion. We stress, however, that our investigations are devoted to  more narrow class of systems (regarding the conservative perturbations).  A general conclusion which can be derived on the basis of the generalized theory is quite unexpected. Even in a case when the orbital parameters cannot be regarded as {\\em extreme}, the corrections to the classic Hamiltonian stemming from the general relativity and the quadrupole moment of the star, may affect the secular dynamics dramatically. Not only  the structure of the phase space of the secular model changes, and new branches  of stationary solutions appear. These solutions may bifurcate within small relative ranges of the parameters (for instance, when the spin of the parent star is changing). We show that there is no simple and general recipe to predict the behavior of the secular system, when the perturbations are ``switched on''. The secular dynamics of  the generalized model becomes extremely complex and rich. For some combinations of the system parameters, the notion of the GR and QM  effects as {\\em corrections}  to the point-mass  NG interactions does not seem proper anymore.  In some cases, these effects may be {\\em more important} for the secular dynamics than the mutual, point mass Newtonian interactions between the planets.  We also show that these effects may be significant for the dynamical stability of planetary systems both in the short-term and in the secular time scales. For instance, the QM generated perturbations may directly depend on the star age and its physical parameters ($k_L$, $R_0$). In turn, these effects may strongly influence the structure of the phase space and can imply short-term, strong chaotic orbital evolution during a few secular periods.  The direct tests of the analytic theory are very encouraging.  The results justify its great accuracy. The precision of the analytic calculations is  very important for studying the global dynamics of hierarchical systems. The alternative numerical approach would require huge CPU time because the hierarchical planetary systems evolve during very different time scales. Then the CPU requirements of the direct numerical integrations are  by orders of magnitude larger than those ones needed by the analytic formulae."
    },
    "0809/0809.4975_arXiv.txt": {
        "abstract": "The Cryogenic Rare Event Search with Superconducting Thermometers Phase II (CRESST-II) at the L.N.G.S in Italy is searching for Dark Matter using low-temperature calorimeters. These detectors allow to discriminate different particles by simultaneous measurement of phonons and scintillation light. The sensors used consist of superconducting tungsten thin-film thermometers, which measure the thermal effect of the phonons created in an attached absorber crystal. It has been observed that the scintillation  of the \\cawo absorber degrades during the process of depositing the tungsten film. In order to prevent this, a new technique for producing the detectors was investigated. This technique might also be valuable by expanding the range of scintillator materials suitable for producing a Dark Matter detector. ",
        "introduction": "\\subsection{Introduction} Since the observations of Zwicky in the 1930's the nature of Dark Matter remains unknown. Experiments like WMAP~\\cite{Spe06} could only deliver more evidence on its existence but could not clarify its nature. The weakly interacting massive particle (WIMP) is a well-motivated candidate for Dark Matter~\\cite{jungman}.\\par \\subsection{CRESST-II} The CRESST-II experiment aims for direct detection of WIMPs  scattering off nuclei in a scintillating absorber crystal~\\cite{Ang05}. The target material is chosen in order to optimize the cross section for spin-independent coherent WIMP-nucleus scattering, whose cross section scales quadratically with the atomic mass $A$.\\par The energy transfered by a scattering is typically only in the order of \\unit[10]{keV}. This, together with an extremely low interaction rate of less than 10 events per kilogram of absorber material and year of exposure requires very sensitive detectors, which allow the rejection of background events~\\cite{jungman}.\\par In order to reject the background radiation, CRESST-II uses scintillating crystals as target material. A schematic view of the detector can be seen in Figure \\ref{fig:schemaDetector}. The energy deposited in a scintillating absorber crystal by a particle interaction excites both phonons and scintillation light. The relative amount of light depends on the nature of the detected particle. As WIMPS, in comparison to highly ionizing particles, cause a relatively low light output, the optimization of detection of the scintillation light is of utmost importance.\\par \\begin{figure} \\centering \\includegraphics[width=6cm]{kiefer_fig_schemaDetector.pdf} \\caption{Schematic view of the detector.}\\label{fig:schemaDetector} \\end{figure} The scintillation light is absorbed by a silicon coated sapphire crystal, exciting phonons in the crystal lattice. Each of the two crystals of a detector module, \\cawo and sapphire, is connected to a superconducting phase transition thermometer for measurement of the phonon signals.\\par Such a thermometer consists of a thin film of superconducting material (tungsten) which is deposited on the crystal. Phonons are absorbed in the film and increase the temperature of its electron system. The film is thermally stabilized in its phase transition from the superconducting to normal conducting state, therefore its resistance is extremely sensitive to variations in temperature. Phonons that are absorbed in the film increase the temperature of its electron system and thus its resistance. The change in resistance is read out by a sensitive SQUID circuit~\\cite{Ang05,Antonio,Lang:2008fa}.\\par The upcoming EURECA experiment is going to use particle detection and discrimination techniques developed and studied in the CRESST-II experiment~\\cite{EureKraus}. ",
        "conclusions": "Detectors with glued thermometers have proven to operate in a stable manner and to deliver pulses which can be used to measure a spectrum. They bring many advantages in terms of fabrication of the detectors without significantly degrading the energy spectrum.\\par Nevertheless there is a loss of signal quality in the phonon channel. In the case of CRESST, this loss can be accepted, as the increase in scintillation light is expected to improve the discrimination capabilities.\\par Gluing thermometers onto scintillator crystals may provide a possibility to use scintillator materials which up to now have not been suitable for experiments like CRESST because of difficulties in depositing a thermometer. This might enlarge the choice of scintillating absorber materials for CRESST and the upcoming EURECA experiment."
    },
    "0809/0809.3849_arXiv.txt": {
        "abstract": "The polarization of light across individual spectral lines contains information about the circumstellar environment on very small spatial scales. We have obtained a large number of high precision, high resolution spectropolarimetric observations of Herbig Ae/Be, Classical Be and other emission-line stars collected on 117 nights of observations with the HiVIS spectropolarimeter at a resolution of R=13000 on the 3.67m AEOS telescope. We also have many observations from the ESPaDOnS spectropolarimeter at a resolution of R=68000 on the 3.6m CFH telescope. In roughly $\\sim$2/3 of the so-called ``windy'' or ``disky'' Herbig Ae/Be stars, the detected H$_\\alpha$ linear polarization varies from our typical detection threshold near 0.1\\% to over 2\\%. In all but one HAe/Be star the detected polarization effect is not coincident with the H$_\\alpha$ emission peak but is detected in and around the obvious absorptive part of the line profile. The qu-loops are dominated by the polarization in this absorptive region. In several stars the polarization varies in time mostly in the absorptive component and is not necessarily tied to corresponding variations in intensity. This is a new result not seen at lower resolution. In the Be and emission-line stars, 10 out of a sample of 30 show a typical broad depolarization effect but 4 of these 10 show weaker effects only visible at high resolution. Another 5 of 30 show smaller amplitude, more complex signatures. Six stars of alternate classification showed large amplitude (1-3\\%) absorptive polarization effects. These detections are largely inconsistent with the traditional disk-scattering and depolarization models. ",
        "introduction": "Understanding how stars interact with their environment is a general problem important for many astrophysical systems. Many types of stars are surrounded by shells, envelopes or disks of material. Such circumstellar material often participates in accretion, polar outflows, winds and disk-star interactions. In young stars, these processes directly effect star and planet formation and evolution. Even for the closest young stars we study here, the stellar radius is an important spatial scale that is smaller than 0.1 milliarcseconds and it will not be imaged directly even by the next generation of telescopes. Other methods like interferometry can yield useful constraints, but the indirect techniques of spectroscopy and spectropolarimetry have much to offer. In particular, spectropolarimetric measurements can put unique constraints on the geometry of the system as well as thermodynamic and radiative environments of circumstellar material.  High-resolution linear spectropolarimetry measures the change in linear polarization across a spectral line. Typical spectropolarimetric signals are very small, often a few tenths of a percent change in polarization across a spectral line. Measuring these signals requires very high signal to noise observations and careful control of systematics to measure signals at this level. In order to address these issues, we built a dedicated spectropolarimeter for the High-resolution Visible and Infrared Spectrograph (HiVIS) on the 3.7m AEOS telescope and performed a telescope polarization calibration and cross-instrument comparison (Harrington et al. 2006, Harrington \\& Kuhn 2008). In addition, new and archival observations with the ESPaDOnS spectropolarimeter will be presented which will complement our HiVIS observations.  Many models show spectropolarimetry is a potentially useful probe of circumstellar environments at small spatial scales. Circumstellar disks, rotationally distorted winds, magnetic fields, asymmetric radiation fields (optical pumping), and in general, any scattering asymmetry can produce a change in linear polarization across a spectral line (cf. McLean 1979, Wood et al. 1993, Harries 2000, Ignace et al. 2004, Vink et al. 2005a, Kuhn et al. 2007). These spectropolarimeric signatures can directly constrain the physical properties of the circumstellar environment once sufficiently detailed models are developed. There have been several studies of young stars to date using medium resolution spectropolarimetry (Vink et al. 2002, 2005b, Mottram et al. 2007). These studies focused on young stars and are based on only a few nights of observations. Though spectropolarimetric signatures have been detected in many classes of star, a many-epoch, high-resolution comparison of different stellar classes has yet to be done.  We collected a large number of stellar observations on over 100 nights from 2004 to 2008, monitoring 29 Herbig Ae/Be stars and 30 Be and emission-line stars. These stellar types have very different circumstellar environments with various combinations of winds, disks, accretion, and jets (cf. Porter \\& Rivinius 2003 and Waters \\& Waelkens 1998). However, there is still much debate over the details of these processes. Thomson scattering theory is the current framework used to describe linear spectropolarimetric signatures. A ``depolarization\" scattering theory of spectropolarimetric line profiles was originally developed for Be stars (McLean 1979). Another scattering model we will call the ``disk-scattering\" model has been advanced for other star types. However, neither model significantly matches most of our Herbig Ae/Be observations, and many observations of other star types.  The observing campaign and basic data processing will be presented in section 2 along with the spectroscopy of the Herbig Ae/Be stars. A quick overview of current spectropolarimetric theories will be presented in section 3. The HAe/Be spectropolarimetry will be detailed in section 4. All the HiVIS and ESPaDOnS data will be presented with individual examples for each star. Section 5 will present the spectroscopy and spectropolarimetry of the Be and Emission-line stars for comparison. A comparison with other stellar types will also be made. These observations will then be compared and contrasted in the discussion of section 6. ",
        "conclusions": "This paper presents a large amount of HiVIS observations taken on over 100 nights on 29 Herbig Ae/Be stars, 30 Be and emission-line targets as well as several alternatively classified objects to compile a linear spectropolarimetric survey far larger in scope than any survey performed to date.  Archival data as well as our own observations with the ESPaDOnS spectropolarimeter were used to supplement the HiVIS survey and to prove the robustness of the combined dataset. The ESPaDOnS observations represent an order-of-magnitude gain in resolution from any survey done to date. The observing campaign from 2004 to 2008 was outlined in section 2. The spectroscopic observations of the Herbig Ae/Be stars showed, as expected, a wealth of line morphologies. The H$_\\alpha$ line of these stars is typically strong and quite variable. The emission lines can be a few to thirty times continuum with many different absorptive components on top of the emission. Spectroscopic measurements for these stars over many time-scales were presented and discussed. A few short-term variability studies were reported illustrating the dynamic circumstellar environment in these stars.  Current theories of spectropolarimetric line profiles, based on disk or envelope scattering effects were discussed in section 3. The current disk-scattering theory predicts symmetric, broad, double-peaked spectropolarimetric signatures in the orbiting thin-disk case. The corresponding qu-loops show position-angle changes that are equally broad in wavelength. In the case of outward radial motion, the disk-scattering theory predicts a shift of the spectropolarimetric profiles toward the red side of the spectral line. The depolarization theory is essentially a dilution of continuum polarization by unpolarized emission. This effect produces broad signatures that are roughly inversely proportional to the amount of emission.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.35\\linewidth, angle=90]{f34a.eps} \\includegraphics[width=0.35\\linewidth, angle=90]{f34b.eps} \\\\ \\includegraphics[width=0.35\\linewidth, angle=90]{f34c.eps} \\includegraphics[width=0.35\\linewidth, angle=90]{f34d.eps} \\\\ \\includegraphics[width=0.35\\linewidth, angle=90]{f34e.eps} \\includegraphics[width=0.35\\linewidth, angle=90]{f34f.eps} \\caption{A side-by-side comparison of the spectropolarimetric morphology of individual stars. From left to right: {\\bf a)} The ESPaDOnS archive data for MWC 158 on February 9th, 2006. This is a ``disky\" HAeBe star with a strong spectropolarimetric effect in the central absorption. {\\bf b)} The ESPaDOnS archive data for MWC 480 on February 7th, 8th, and August 13th 2006 This is a ``windy\" HAeBe star with a strong spectropolarimetric effect in the blue-shifted absorption. {\\bf c)} The HiVIS data for MWC 143 - a Be type star showing a broad spectrpolarimetric signature that shows no effect across the central absorption. {\\bf d)} The HiVIS data for $\\alpha$ Col - a Be star with a smaller, more complex spectropolarimetric signature. {\\bf e)} The ESPaDOnS archive data for 51 Oph on March 20th 2008. This is a HAe star with a more complex spectropolarimetric effect that is not only in the central absorption. {\\bf f)} The ESPaDOnS archive data for 3 Pup on February 7th \\& 8th, 2006. This is an A-type emission-line star with a very strong, complex spectropolarimetric signature. The morphology is completely different between these stars. The polarization-in-absorption of the HAe/Be stars is seen in a) and b). There are many Be/emission-line stars with broad spectropolarimetric signatures like c) but a significant number with smaller more complex detections similar to d). Within each class there are exceptions to these general trends, as in e) 51 Oph. These morphologies are also worth exploring in other types of stars, as was shown by 3 Pup in f).} \\label{fig:swp-sidebyside} \\end{center} \\end{figure*}   In section 4, the spectropolarimetric survey of the Herbig Ae/Be stars was presented. The survey showed that roughly two thirds of the HAeBe stars showed polarization effects in the absorptive components of the H$_\\alpha$ line. Both the windy-type and disky-type stars showed this polarization-in-absorption effect in the blue-shifted or line-center absorptive components respectively. The magnitude of the effect was typically 0.1\\% to 2\\% with several stars showing complex morphologies and complex geometries in qu-space. Only one star, HD 179218, showed a spectropolarimetric effect across the entire H$_\\alpha$ line. These observations are difficult to explain using current scattering theories. The depolarization effect causes a broad, monotonic change in polarization that is roughly inversely proportional to emission and is a linear extension in qu-space. The disk-scattering effect shows broad, symmetric morphology for orbiting material and a shift toward red wavelengths for out-flowing material. Both of these effects do not fit the observed morphology of polarization-in-absorption on the blue-shifted side of the windy-stars or the polarization-in-central-absorption-only for the disky-stars. The amplitude of the observed effects is also a consideration. Scattering theory has no natural amplitude, but most of the detections fall in the 0.3\\% to 1.5\\% range.  A  survey of 30 Be and emission-line stars was performed and presented in section 5. This was also complemented by studies of 6 other-type stars such as post-AGB and RV-Tau types, as well as an in-depth look at 3-Pup. These Be and emission-line stars had a very different type of spectropolarimetric morphology. In 10 of 30 stars there was a broad, monotonic spectropolarimetric signature present across the entire H$_\\alpha$ line. This type of signature is typical of the depolarization effect. However, in these 10 broad detections, four showed more complex morphologies in addition to the simple depolarization tied to absorptive components of the emission line. In another 5 of 30 stars there were antisymmetric or other more complex spectropolarimetric signatures that are inconsistent with depolarization theory. This smaller survey of Be stars is in good agreement with the large collection of previous work done to date on these stars at lower resolution. The broad spectropolarimetric morphology is reproduced, but the higher resolution, higher signal-to-noise and larger sample of this survey shows that there are still some additional uncertainties about the exact form of the spectropolarimetric effects in some of these stars. One third of the detections (5/15) do not fit the simple depolarization framework. Four of 10 broad signatures show significant deviations of the line from the simple depolarization morphology in absorptive components that would only have been seen at high resolution with high signal-to-noise.  The inability of scattering theories to fit the ``polarization-in-absorption\" seen in the Herbig Ae/Be stars is good evidence for intrinsic absorptive polarization effects. This survey has already inspired optical pumping calculations (Kuhn et al. 2007) that demonstrate such a mechanism.  Detailed models based on any paradigm assuming any realistic circumstellar environments do not yet predict spectropolarimetric line profiles. The best example of this is the ``polarization-in-absorption\" seen primarily in the Herbig Ae/Be stars or the more complicated profiles observed in Be and other emission-line stars. Unusual linear spectropolarimetric profiles were seen in both HAeBe and Be classes as well as in other A-type stars such as 3 Pup and $\\epsilon$ Aur. An intrinsic absorptive polarization mechanism seems essential  for explaining many of these observations as the Herbig Ae/Be detections are clearly inconsistent with disk-scattering theories. However, as this large survey shows, there is a broad range of linear spectropolarimetric ``signatures''  in these stellar systems waiting to be deciphered."
    },
    "0809/0809.4691_arXiv.txt": {
        "abstract": "We report on an {\\it XMM-Newton} observation of the $z=1.055$ quasar and Giga-hertz Peaked Spectrum (GPS) source 3C 287.  Our 62.3~ksec observation provides an exceptional X-ray view of a prominent member of this important subclass of active galactic nuclei (AGN).  The X-ray spectra of 3C 287 are consistent with a simple absorbed power-law with a spectral index of $\\Gamma=1.72\\pm 0.02$.  Our fits imply a bolometric luminosity of $L = 5.8\\pm 0.2 \\times 10^{45}~{\\rm erg}~ {\\rm s}^{-1}$ over the 0.3--10.0 keV band; this gives a mass lower limit of $M_{BH min} \\geq 4.6 \\times 10^{7}~{\\rm M_{\\odot}}$ assuming X-rays contribute 10\\% of the bolometric luminosity and radiation at the Eddington limit.  Iron emission lines are common in the X-ray spectra of many AGN, but the observed spectra appear to rule-out strong emission lines in 3C~287. The simple power-law spectrum and absence of strong emission lines may support a picture where our line of sight intersects a relativistic jet.  Milliarcsecond radio imaging of 3C 287 appears to support this interpretation.  We discuss our results in the context of different AGN sub-classes and the possibility that GPS sources harbor newly-formed black hole jets. \\smallskip ",
        "introduction": "Connections between accretion disks and relativistic jets in black hole systems are anticipated theoretically, but observational constraints on this coupling remain elusive.  Some progress has been made by tracking X-ray and radio flux variations in stellar-mass black holes, for instance, where transient outbursts evolve over weeks and months (e.g. Gallo, Fender, \\& Pooley 2003).  However, observations of stellar-mass black holes average over vastly more dynamical timescales and much larger regions than observations of the supermassive black holes powering AGN.  In a sense, every observation is an average, and every image is coarse.  Giga-hertz Peaked Spectrum (GPS) sources and Compact Steep-Spectrum (CSS) sources are well-known AGN subclasses.  Beyond the flux-based criteria used to separate these classes, surveys reveal an important size difference: GPS sources are generally less than 1 kpc in projected linear size, while CSS sources can be 20 times larger (see, e.g., Edwards \\& Tingay 2004). These small size scales imply relatively young jets from new activity.  Compact Symmetric Objects (CSOs) -- a subclass of GPS and CSS sources -- represent an even more interesting regime in which to explore disk--jet connections.  Proper motion studies of some CSO sources imply jets that have only become active within the last $10^{3}$ years (Conway 2002).  3C 287 is a well-known radio--loud quasar at $z=1.055$ that falls into the GPS category, and perhaps also into the CSO category (Kellerman et al.\\ 1998).  Our preliminary analysis of a prior {\\it Chandra} observation suggested the presence of a broad, relativistic iron disk line.  Such lines are excellent probes of the inner disk (e.g. Miller 2007), making 3C 287 an even more exciting target in which to explore accretion flows and to study disk--jet connections.  GPS quasars do not necessarily have the same X-ray properties as GPS galaxies.  GPS galaxies show large X-ray absoption columns while GPS quasars show significantly less absorption (Siemiginowska et al. 2008).  X-ray absorption in GPS galaxies may result from an orientation coincident with an obscuring \"torus\", possibly blocking the central engine from direct observation (Guainazzi et al. 2006).  Although a modest survey of GPS sources has been performed with Chandra (Siemiginowska 2003, 2008), deep X-ray observations of these sources are rare.  In the following sections, we report the results of fits to the spectra of 3C 287 as measured using the {\\it XMM-Newton} EPIC/pn, MOS-1, MOS-2, and OM cameras.  Fits to these spectra suggest power-law behavior typical of CSS sources and the absence of strong Fe K emission lines.  We discuss these results as well as their implications on emission mechanisms and our line of sight to 3C 287. ",
        "conclusions": "In the sections above, we have detailed our analysis of a deep exposure of 3C 287 made with {\\it XMM-Newton}.  In the 0.3--10.0~keV observed frame (approximately 0.6--20~keV in the source frame), the spectrum is well represented by a simple absorbed power-law model.  In AGN that are viewed at a modest inclination, it is common to observe either an X-ray warm absorber (likely a disk wind; see, e.g., Reynolds 1997) or a soft excess (perhaps also tied to the disk; e.g. Crummy et al.\\ 2006), and an Fe K emission line due to hard X-ray illumination of the disk, the torus, or both.  Fits made with a broken power-law were made in order to test for evidence of either a soft excess or warm absorption at low energy.  However, these fits do not provide a significantly improved description of the data.  Similarly, fits with Gaussian and disk line functions place limits on the strength of both narrow and broad emission lines.  It is also worth noting that our results place a limit of approximately $5 \\times 10^{19}$~atoms~${\\rm cm}^{-2}$ on any neutral absorbing gas along our line of sight to 3C 287.  This limit is about a factor of two lower than the measurement reported by Siemiginowska et al.\\ (2003).  Although models where the source is confined by dense gas (Snellen et al.\\ 2000, Alexander 2000) have not been completely ruled out, new work argues against confinement (Guainazzi et al.\\ 2004, Morganti 2007).  Our X-ray results are consistent with emission related to a jet that dominates over the typical X-ray emission associated  \\centerline{~\\psfig{file=f4.eps,width=3.3in}~} \\figcaption[h]{\\footnotesize A broad-band spectral energy distribution of 3C 287 is shown above.  X-ray and optical points are from our {\\it XMM-Newton} observations.  Infrared points are from 2MASS observations in J and H.  The single radio point is from a VLA observation at 5~GHz.} \\medskip  \\noindent with the accretion process.  The emission we observe could be accounted for through a combination of synchrotron and synchrotron self-Comptonization in the jet, hot spots, and compact lobes.  A line of sight coincident with the jet is suggested by deep milliarcsecond radio imaging of 3C 287.  VLBA imaging at 15~GHz reveals an extremely compact core in 3C 287, with a minimal extension of less than 4 milliarcseconds, or less than about 20~pc in the source frame (Kellermann et al.\\ 1998).  This core is among the smallest and least extended in the survey of AGN conducted by Kellerman et al.\\ (1998), and is strongly suggestive of a jet that lies primarily along our line of sight.  In many AGN, the presence of a strong, narrow iron fluorescence line is tied to X-ray irradiation of a parsec-scale torus (e.g. Antonucci 1993, Urry \\& Padovani 1995, Yaqoob \\& Padmanabhan 2004, Nandra 2006).  Thus, although the radio core is remarkably small by the standards of jet sources, it is still large enough to block our line of sight to the innermost region (or even the torus) in 3C 287. A possible difficulty with this inpretation is that emission consistent with the big blue bump is observed. This may indicate that the jet does not block all of the disk emission along our line of sight.  The question arises, then, why extreme blazar-like variability is not observed in 3C 287 (and similar sources).  A partial answer may lie in the likely nature of CSS and GPS sources. Proper motion studies of some sources in this class suggest that these AGN have only been powering jets for approximatly $10^{3}$ years (Conway 2002), however, a young radio source may have higher efficiency of converting the jet power into radiation.  In such young sources the X-ray emission may originate in compact lobes and hot spots (Stawarz et al 2007) if the line of sight is not aligned directly with a jet.  In this model there is no expectation of the X-ray variability typically observed in blazars. Note that there is no change in X-ray luminosity between the {\\it Chandra} and {\\it XMM-Newton} observations (Siemiginowska et al 2008).  The broad-band spectral energy distribution  shown in Figure 4 may further support this interpretation.  The source is more radio-loud than typical radio-loud quasars (Elvis et al 1994, Richards, 2006, Siemiginowska et al 2008), while its X-ray emission is relatively weak with $\\alpha_{ox}=1.7$. This suggests a weak contribution from the standard emission component (for example a parsec scale jet, Blazejowski et al. 2004) observed in radio-loud quasars.  We conclude that 3C 287 may be a ``proto-blazar'' candidate. Dedicated multi-wavelength monitoring of sources like 3C 287 over a period of years can help to better reveal the nature of this source.  \\vspace{0.25in}  We thank the anonymous referee for insightful comments that improved the paper.  JMM gratefully acknowledges funding from NASA, through the {\\it XMM-Newton} guest observerations program.  Partial support for this work was provided by the National Aeronautics and Space Administration through Chandra Award Number GO2-3148A and GO5-6113X and under the contract NAS8-39073."
    },
    "0809/0809.2624_arXiv.txt": {
        "abstract": "We present interferometric observations in the $^{12}$CO\\,(2--1) line and at 1.3\\,mm dust continuum of the low-mass protostellar binary system in the cometary globule CG\\,30, using the Submillimeter Array. The dust continuum images resolve two compact sources (CG\\,30N and CG\\,30S), with a linear separation of $\\sim$\\,8700\\,AU and total gas masses of $\\sim$\\,1.4 and $\\sim$\\,0.6\\,$M_\\odot$, respectively. With the CO images, we discover two high-velocity bipolar molecular outflows, driven by the two sources. The two outflows are nearly perpendicular to each other, showing a quadrupolar morphology. The northern bipolar outflow extends along the southeast (redshifted, with a velocity up to $\\sim$\\,23\\,km\\,s$^{-1}$) and northwest (blueshifted, velocity up to $\\sim$\\,30\\,km\\,s$^{-1}$) directions, while the southern pair has an orientation from southwest (blueshifted, velocity up to $\\sim$\\,13\\,km\\,s$^{-1}$) to northeast (redshifted, velocity up to $\\sim$\\,41\\,km\\,s$^{-1}$). The outflow mass of the northern pair, driven by the higher mass source CG\\,30N, is $\\sim$\\,9 times larger than that of the southern pair. The discovery of the quadrupolar molecular outflow in the CG\\,30 protobinary system, as well as the presence of other quadrupolar outflows associated with binary systems, demonstrate that the disks in (wide) binary systems are not necessarily co-aligned after fragmentation. ",
        "introduction": "Binarity/multiplicity is frequent among protostellar systems (see e.g., Looney et al. 2000). To probe these protostellar binary (protobinary) systems, high angular resolution observations of gas and optically thin dust emission at (sub-)\\,millimeter (mm) wavelengths are needed, because protostellar systems are deeply embedded in molecular cloud cores and are surrounded by infalling envelopes and accretion disks. However, these observations were long limited by the low angular resolution of single-dish telescopes and have only become possible with the recent availability of large (sub-)\\,mm interferometers. During the past few years, interferometric mm observations have enabled us to discover several protobinary systems (Launhardt 2004), and there is an increasing number of interferometric studies of binarity in the protostellar stage (e.g., Chen et al. 2007). Nevertheless, due to the large effort and observing time required, the number of known and well-studied protobinary systems is still very small.  CG\\,30 (also known as BHR\\,12, Bourke et al. 1995; or DC\\,253.3$-$1.6, Hartley et al. 1986) is a bright-rimmed cometary globule located in the Gum Nebula region at a distance of $\\sim$\\,400\\,pc (Reipurth 1983). With high resolution submm (SCUBA; Henning et al. 2001) and mm (ATCA; Chen et al. 2008, hereafter Paper\\,I) dust continuum observations, we have identified in this globule a wide protobinary system with a projected separation of $\\sim$\\,21\\farcs7 ($\\sim$\\,8700\\,AU) and bolometric luminosities of 13.6\\,$\\pm$\\,0.8 and 4.3\\,$\\pm$\\,0.5\\,$L_\\odot$ for the northern and southern protostars, respectively. The two protostars appear to drive their own bipolar protostellar jets, which are roughly perpendicular to each other, as seen in the H$_{2}$\\,1--0\\,$S$(1) (Hodapp \\& Ladd 1995) and the $Spitzer$ IRAC 4.5\\,$\\mu$m images (Paper\\,I). However, information on the molecular outflows in this protobinary system has been limited, although single-dish observations in the CO line (Nielsen et al. 1998) and CS line (Garay et al. 2002) had shown a large-scale (unresolved) bipolar outflow in the east-west direction. In this Letter, we report Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA; Ho et al. 2004) observations in the $^{12}$CO\\,(2--1) line towards CG\\,30 with angular resolution sufficient to allow us, for the first time, to map the distribution of outflowing molecular gas from each member of the protobinary system. ",
        "conclusions": "\\subsection{Outflow Physical Parameters}  Thanks to the versatility of the SMA correlator, the $^{13}$CO\\,(2--1) line was observed simultaneously with the $^{12}$CO\\,(2--1) line towards CG\\,30, which allows us to estimate the outflow mass. For source CG\\,30N, the $^{13}$CO\\,(2--1) blueshifted emission is detected at velocities from $\\sim$\\,$-$2 to $\\sim$\\,3\\,km\\,s$^{-1}$, while redshifted emission is from $\\sim$\\,8 to $\\sim$\\,12\\,km\\,s$^{-1}$; For source CG\\,30S, however, no $^{13}$CO\\,(2--1) outflow gas is found, although weak emission is detected around systemic velocity. The outflow masses are then calculated following Cabrit \\& Bertout (1990): for the southern bipolar outflow, as well as high-velocity component of the northern pair where $^{13}$CO\\,(2--1) is not detected, the $^{12}$CO\\,(2--1) emission is assumed to be optically thin, and the mass is calculated at each channel where outflow emission is detected, and then summed; for the low-velocity component of the northern pair, a correction for line opacity is performed at velocities where both the $^{12}$CO/$^{13}$CO\\,(2--1) emissions are detected. In the calculations, we assume LTE conditions, an excitation temperature of 20\\,K, an abundance ratio of [H$_2$]/[$^{13}$CO] equal to 8.9\\,$\\times$\\,10$^{5}$ (Cabrit \\& Bertout 1992), and an isotopic ratio of [$^{12}$CO]/[$^{13}$CO] = 71 (Wilson \\& Rood 1994). No correction for inclination effects was applied in the calculations.  The derived outflow masses, as well as other outflow properties (i.e., momentum $P$, energy $E$, mass-loss rate $\\dot{M}$$_{\\rm out}$, force $F_{\\rm m}$, and mechanical luminosity $L_{\\rm m}$), are listed in Table~2. A simple comparison between the two bipolar outflows shows that the northern pair, driven by the higher (circumstellar envelope) mass source CG\\,30N, is much stronger than the southern pair ($\\sim$\\,9 times in mass and $\\sim$\\,3 times in energy). The outflow parameters derived in CG\\,30 are relatively small in comparison with the general range of such parameters for sources of similar mass (see Wu et al. 2004). Nevertheless, we want to mention that the parameters in Table~2 only refer to the compact outflows detected in the SMA maps. The nature of the ``extended outflow\" components is unclear yet and requires confirmation by combining the SMA data with single-dish data to recover extended missing flux. Therefore, these parameters represent lower limits. Furthermore, as suggested by the two larger-scale jets, mosaicked CO observations are probably needed to fully map the two outflows.  \\subsection{Binarity v.s. Quadrupolar Outflows}  The two bipolar outflows found in CG\\,30 resemble a combined quadrupolar morphology. A similar outflow morphology was also found in several other (low-mass) sources, like e.g., IRAS\\,16293--2422 (Mizuno et al. 1990), HH\\,288 (Gueth et al. 2001), and L\\,723 (Lee et al. 2002 and references therein). It is important to note that in all these sources protobinary systems have been discovered by high-resolution interferometric observations (IRAS\\,16293, Mundy et al. 1992; HH\\,288, Gueth et al. 2001; L\\,723, Launhardt 2004). This strongly supports the hypothesis that many quadrupolar outflows\\footnote{Quadrupolar outflows are sources in which four outflow lobes are observed and seem to be driven by the same protostellar core (seen in current mm single-dish telescopes).} are actually two independent outflows driven by a binary system (Anglada et al. 1991), although several other scenarios have been proposed (see Arce et al. 2007). Furthermore, assuming that binary systems are formed through the fragmentation of collapsing protostellar cores (see Goodwin et al. 2007), these quadrupolar outflows demonstrate that the disks in (wide) binary systems are not necessarily co-aligned after fragmentation.  On the other hand, if binarity among protostars occurs as frequently as among pre-main sequence stars, one would expect to find many quadrupolar/multipolar protostellar outflows. Such kind of outflows are however rare in our observations. One possible reason is that protostellar disks (jets) are preferentially aligned during fragmentation, and they may then drive one common large-scale outflow. An example is the source HH\\,30 (see Guilloteau et al. 2008 and references therein). At the same time, the outflow strength of protostars is proportional to the circumstellar envelope mass (Bontemps et al. 1996). Since binary protostars appear to have very unequal circumstellar masses (Launhardt 2004), their relative outflow strengths are expected to be also very unequal. Therefore, another possible reason is that the weaker outflow may often not be detected with current single-dish telescopes that do not have sufficient angular resolution and sensitivity. The discovery of an equal-mass protobinary system in L\\,723 (Launhardt 2004) and IRAS\\,16293 (Mundy et al. 1992) and of a weak secondary outflow from the lower-mass component in the unequal-mass protobinary systems CG\\,30 (this work), CB\\,230 (Launhardt 2004), and BHR\\,71 (Bourke 2001, Parise et al. 2006) supports this explanation. However, high-resolution outflow data so far exist only for very few systems, and it is unclear yet why multipolar protostellar outflows are so rare in our observations."
    },
    "0809/0809.5142_arXiv.txt": {
        "abstract": "The impact of particle production during inflation on the primordial curvature perturbation spectrum is investigated both analytically and numerically. We obtain an oscillatory behavior on small scales, while on large scales the spectrum is unaffected. The amplitude of the oscillations is proportional to the number of coupled fields, their mass, and the square of the coupling constant. The oscillations are due a discontinuity in the second time derivative of the inflaton, arising from a temporary violation of the slow-roll conditions. A similar effect on the power spectrum should be produced also in other inflationary models where the slow-roll conditions are temporarily violated. ",
        "introduction": "\\label{Introduction}  Inflation is considered one of the most promising candidates to explain the statistical features of the observable Universe revealed by the detection of the cosmic microwave background (CMB)~\\cite{Spergel:2003cb,Spergel:2006hy} and  large scale structure surveys such as the Sloan Digital Sky Survey (SDSS)~ \\cite{sdss}.%  In the simplest models, inflation is an exponential expansion period of the Universe, driven by scalar field, called inflaton, slowly rolling down its potential. One of its main predictions is a primordial curvature perturbation spectrum of the form $P_R(k)=k^{n_s-1}$, where $n_s$ is the so-called scalar spectral index.  Various extensions of the simplest models have been proposed, and in this paper we will consider the coupling of the inflaton to another massive scalar field, which has been previously investigated by different groups, leading to apparently conflicting results~\\cite{Chung:1999ve,Elgaroy:2003hp}. The first group was in fact obtaining just a local peak in the power spectrum for scales which leave the horizon around the time of particle production, while the second obtained a small scale oscillatory behavior leading to a step between large and small scales. We study the problem both analytically and numerically, obtaining an intermediate result which confirms the oscillation of the power spectrum on small scales, but without any step. We also derive an analytical approximation for the curvature power spectrum, which clearly shows the dependency of the amplitude and period of the oscillations on the mass and number of the coupled fields and the coupling constant.  We mention that the presence of a temporary non-slow-roll stage during inflation and its effect on the scalar and tensor perturbation spectrum as well as the resulting CMB anisotropy have been investigated in \\cite{Boyanovsky:2006pm,Destri:2008fj}. Our analysis may be regarded as a special case where the perturbation spectrum can be studied analytically, and confirms other general studies \\cite{Starobinsky:1992ts,Gong:2005jr,Joy:2007na} of the effects of singularities in the inflaton potential. See also a recent paper by Joy et al.~\\cite{Joy:2008qd} for comparizon with the WMAP data.  The paper is organized as follows. In Section~\\ref{model}, we briefly describe the main features and motivations of the model we study. In Section~\\ref{calculation}, we describe our analysis and present the numerical results for the spectrum of the curvature perturbation. In Section~\\ref{approximation}, we adopt an analytical approximation and derive the spectrum in both large and small $k$ limits. In Section~\\ref{conclusion}, we summarize the results obtained and provide some ideas about possible extensions. ",
        "conclusions": "\\label{conclusion}  We have studied the impact of the coupling of the inflaton to a scalar field on the primordial curvature perturbation spectrum. We found the presence of an oscillatory behavior on small scales, for the modes which leave the horizon after the time of particle production. We have also presented a very good analytical approximation for the evolution of small scale modes. This method can be applied to the general case of a sudden change in the inflaton potential, leading to a temporary violation of the slow roll conditions.  The amplitude of the oscillations is proportional to the number of coupled fields $N$ and the square root of their mass $m_{\\varphi}^{1/2}$, and to the square of the coupling constant $g^2$. The period of the oscillations is $k_0\\pi$, where $k_0$ is the wavelength that crosses the horizon right at the time of the particle production.  On large scales, $k<k_0$, the power spectrum is virtually unaffected by the particle production, contrary to what was claimed in a previous numerical investigation~\\cite{Elgaroy:2003hp}. This is because the violation of the slow-roll conditions can affect only on those modes close to $k=k_0$ on superhorizon scales, and the curvature perturbation is conserved on sufficiently large superhorizon scales, $k\\ll k_0$, no matter what occurs there.  Our results are quite general in the sense that in models where slow-roll conditions are temporarily violated, the spectrum will have oscillations on scales smaller than the mode which leaves the horizon at time of transition, while it will remain unchanged on the larger scales. The presence of such features in the observed CMB spectrum could help to determine the magnitude and the lapse of periods during which the slow-roll conditions are violated, although it may be difficult in practice to distinguish such a feature of a primordial origin from a similar feature due to intermediate astrophysical processes happening before and/or after recombination."
    },
    "0809/0809.0883_arXiv.txt": {
        "abstract": "We have cross-correlated the SDSS DR3 \\citet{Schnei05} quasar catalog with the XMM-Newton archive.  Color and redshift selections ($g - r$ $\\geq$ 0.5 and 0.9 $<$ z $<$ 2.1) result in a sample of 17 red, moderate redshift quasars.  The redshift selection minimizes possible contamination due to host galaxy emission and Ly$\\alpha$ forest absorption. Both optical and X-ray information are required to distinguish between the two likely remaining causes of the red colors: 1) dust-reddening and 2) an intrinsically red continuum.  We find that 7 of 17 quasars can be classified as probable `intrinsically red' objects.  These 7 quasars have unusually broad MgII emission lines ($<$FWHM$>$=10,500 km s$^{-1}$), moderately flat, but unabsorbed X-ray spectra ($<\\Gamma>$=1.66$\\pm$0.08), and low accretion rates ($\\dot{M}/\\dot{M_{Edd}} \\sim$ 0.01). We suggest low accretion rates as a a possible physical explanation for quasars with intrinsically red optical continua.  We find that 8 of 17 quasars can be classified as dust-reddened.  Three of these have upper-limits on the absorption column from X-ray spectral fits of N$_H$ = 3-13 x 10$^{22}$ cm$^2$, while the other five quasars must be absorbed by at least N$_H$ = 10$^{23}$ cm$^2$ in order to be consistent with a comparably selected $\\alpha_{ox}-l_{uv}$ distribution.  Two objects in the sample are unclassified. ",
        "introduction": "Until recently, quasars have typically been selected as point-like objects with colors bluer than stars (e.g. U $ - $ B $<$ -0.44, Schmidt \\& Green 1983)\\nocite{SG83}. This limited the possible spectral energy diagrams (SEDs) that known quasars could have. A composite SED shows that quasars' blue colors are due to the \u2018Big Blue Bump\u2019 (BBB) continuum feature that dominates the optical and UV spectrum \\citep{MS82, Elvis94}. However, surveys using different selection mechanisms discovered that quasars can have red colors as well, indicating a diminished or missing BBB.  While in many cases, this reddening is due to dust obscuration, some cases may be due to changes in the accretion disk thought to power the optical/UV continuum.  We identify such a population in this paper.  Red quasars were first discovered in the \\citet{Web95} sample of radio-loud quasars.  \\citet{Web95} suggested dust-reddening as the cause of the red $B_J-K$ colors and estimated that red colors could cause as much as 80\\% of the quasar population to go undetected in existing optical surveys.  Some 200 members of a new population of red, radio-quiet quasars were discovered in the near-infrared (1-2 $\\mu$m) with the Two Micron All Sky Survey (2MASS) by \\citet{Cutri02}.  The space density of the red 2MASS quasars was found to be at least equal to that of optical and UV-selected quasars.  \\cite{Cutri02} also suggested dust-reddening to explain the red colors of the 2MASS red quasars because the near-IR color distribution is consistent with a reddening of $A_V$ = 1-5 magnitudes.  In addition, $\\sim$10\\% of the predominantly radio-quiet sample is highly polarized ($P > 3\\%$) \\citep{Smith02}.  Supporting this conclusion, Chandra and XMM-Newton observations of red 2MASS AGN find spectra that are either unusually hard or absorbed by N$_H$\u00ad = 10$^{21}$-10$^{23}$ cm$^{-2}$ \\citep{Wilkes02, Wilkes05, Urr05}.  Unlike UV-excess based surveys, the Sloan Digital Sky Survey (SDSS) allows detection of red quasars in the optical band because the algorithm for selecting candidates for spectroscopy uses a four-dimensional multicolor selection criterion \\citep{Rich03} similar to that of the 2dF QSO Redshift Survey \\citep{Croom04}. This allows the SDSS to select all objects lying outside of the stellar locus, including atypically red quasars.  \\citet{Rich03} explored the color distribution of SDSS quasars and found that while a population of intrinsically red quasars exists, the majority of the red tail of the color distribution is explained by dust reddening.  In a later study \\citep{Hopkins04}, the curvatures of the SDSS spectra were found to be best-fit by a dust extinction curve similar to that of the Small Magellanic Cloud (SMC) \\citep{Prevot84}.  However, only mildly dust-reddened quasars can be detected (E(B-V) $<$ 0.5) in the SDSS before dust extinction causes them to drop below the detection threshold of the survey \\citep{Rich03}.  While dust-reddening is a common explanation for red colors in quasars, other possible explanations include: (1) host galaxy contamination, (2) absorption by the Ly$\\alpha$ forest, (3) optical synchrotron emission superimposed on a normal, blue spectrum to create a red spectrum \\citep{SR96, Francis00, Whit01}, and (4) an intrinsically red continuum. The redshift selection in this paper minimizes contamination via options (1) and (2).  X-ray observations used in conjunction with optical information can constrain the amount of intrinsic absorption in a source, thereby allowing an investigation of dust-reddened versus intrinsically red optical continua. Previous studies have found intrinsically red quasar candidates in addition to dust-reddened quasars.  For example, \\citet{Hall06} selected a sample of 12 red quasars from the SDSS, but Chandra X-ray observations show evidence for NH absorption (N$_H\u00ad \\sim~10^{22}$ cm$^{-2}$) for only 4 of their 12 quasars.  The other seven quasars show no evidence for even moderate intrinsic absorption and three of these show evidence for an intrinsically red optical continuum.  In a similar vein, \\citet{Risaliti03} observed 16 optically-selected, X-ray weak quasars with Chandra.  While the sample was not selected to be red, the average color of the sample is redder than that of the parent Hamburg Quasar Survey ($\\Delta(B-R) \\sim$ 1 vs. $\\Delta(B-R) \\sim$ 0.5).  The X-ray weak quasars have a flat average photon index, $<\\Gamma> = 1.5$ (where $\\Gamma$ = -$\\alpha$ + 1 for F$_{\\nu} \\propto \\nu^{\\alpha}$) and 12 of the 16 quasars have steeper optical-to-X-ray indices than normal. (The optical-to-X-ray continuum is characterized by $\\alpha_{ox}$ = log(F$_{2keV}$/F$_{uv}$) / log($\\nu_{2keV}$/$\\nu_{uv})$, where F$_{2keV}$ and F$_{uv}$ are the intrinsic flux densities at 2 keV and 2500 $\\mbox{\\AA}$, respectively \\citep{Tan79}.)  While X-ray absorption remains a possibility, \\citet{Risaliti03} argue that these quasars are intrinsically underluminous in the X-rays.  The evidence, however, is not conclusive due to the low signal-to-noise for many of the Chandra spectra.  Larger surveys including high quality optical and X-ray spectra are now possible with the advent of the SDSS and the XMM-Newton archive of X-ray observations, both of which cover large areas of the sky.  SDSS data can constrain dust-reddening in the optical, while XMM-Newton observations constrain absorption in the X-rays.  In this paper we investigate the optical and X-ray properties of 17 red SDSS quasars.  \\S2 introduces the data sources and sample selection and \\S3 covers the optical and X-ray reduction and data analysis.  In \\S4, we classify the quasars as dust-reddened, intrinsically red or unclassified, and in \\S5 we discuss low accretion rates as a possible explanation for the steep optical/UV spectra of intrinsically red quasars.  We assume throughout the paper that H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have cross-correlated the SDSS DR3 Quasar Catalog \\citep{Schnei05} with the XMM-Newton archive, selecting the reddest ($g - r$ $\\geq$ 0.5), moderate-redshift (0.9 $<$ z $<$ 2.1) quasars.  We obtain a sample of 17 quasars, 16 of which are detected in the X-rays.  Using both optical and X-ray data to constrain dust-reddening and absorption allowed us to distinguish between obscured and intrinsically red quasars, although two cases remain ambiguous.  Eight quasars are dust-reddened in the optical and, while the X-ray data prefer an unabsorbed power-law, the upper-limits are high enough that X-ray absorption at the level expected for an SMC dust-to-gas ratio is allowed.  For the three quasars with high enough X-ray S/N to fit a power-law + intrinsic absorption model, we obtain upper-limits of 3 - 13 x 10$^{22}$ cm$^{-2}$. For five sources, the X-ray S/N is too low to fit spectral models.  However we can obtain a lower-limit to N$_H$ by comparing the $\\alpha_{ox}-l_{uv}$ distribution of the red quasar sample to that of \\citet{Strat05}, a recent study that uses a similar selection method, aside from the red color cut.  If the two samples come from the same parent distribution, an absorbing column of at least 10$^{23}$ cm$^{-2}$ is required for the five red quasars.  Seven quasars display no evidence of X-ray absorption, and dust-reddening is contraindicated by the continuum shape as determined by the optical photometry.  These quasars seem to form a group with intrinsically red power-laws in the optical/UV.  It seems that these intrinsically red power-laws may be caused by lower accretion rates: low accretion rates are derived from M$_{BH}$ and L$_{bol}$ estimates, although the M$_{BH}$ values may be biased high by the red SED's.  Moreover the unusually broad MgII emission lines place these quasars on the $\\Gamma$- FWHM(MgII) relation, also suggesting lower accretion rates.  The intrinsically red quasars are a substantial population with extreme characteristics that are well-suited to X-ray follow-up.  Since 7 intrinsically red quasar candidates have been found in the 1\\% of the SDSS DR3 quasar catalog that overlaps with archival XMM-Newton observations, as many as 700 intrinsically red candidates may exist in the complete catalog.  Since the DR6 has more than doubled the number of SDSS quasars, we can expect 1600 intrinsically red quasar candidates, 16 of which have been observed serendipitously with the XMM-Newton.  The intrinsically red quasar candidates in this paper also merit further observations in order to confirm the lack of dust-reddening and pinpoint the physical cause of the red optical power-laws. For example, NIR (e.g. JHK) spectra would include H$\\alpha$ and H$\\beta$, and so could give an independent measure of dust-reddening via the Balmer decrement.  NIR measurements will also constrain a low-temperature disk model, as well as L$_{bol}$ and so $\\dot{M}$.  With optical and X-ray spectra already available, the intrinsically red quasar candidates described in this paper are well-suited to more complex SED modeling."
    },
    "0809/0809.0897_arXiv.txt": {
        "abstract": "We present a three-dimensional, fully parallelized, efficient implementation of ionizing UV radiation for smoothed particle hydrodynamics (SPH) including self-gravity. Our method is based on the SPH/tree code VINE. We therefore call it iVINE (for Ionization + VINE). This approach allows detailed high-resolution studies of the effects of ionizing radiation from e.g. young massive stars on their turbulent parental molecular clouds. In this paper we describe the concept and the numerical implementation of the radiative transfer for a plane-parallel geometry and we discuss several test cases demonstrating the efficiency and accuracy of the new method. As a first application, we study the radiatively driven implosion of marginally stable molecular clouds at various distances of a strong UV source and show that they are driven into gravitational collapse. The resulting cores are very compact and dense exactly as it is observed in clustered environments. Our simulations indicate that the time of triggered collapse depends on the distance of the core from the UV source. Clouds closer to the source collapse several $10^5$ years earlier than more distant clouds. This effect can explain the observed age spread in OB associations where stars closer to the source are found to be younger. We discuss possible uncertainties in the observational derivation of shock front velocities due to early stripping of proto-stellar envelopes by ionizing radiation. ",
        "introduction": "\\label{intro} As hydrodynamical simulations become more and more advanced one of the key issues is the successful implementation of additional physics like the effects of radiation. Prominent applications are for example the reionization of the early universe (for a comparison of methods see \\citealt{2006MNRAS.371.1057I} and references therein).  In our present day universe ionizing radiation still plays a vital role. UV-radiation from massive, young stars ionizes their surrounding. The hot, ionized gas then expands into the cold, neutral gas and thus drives shock fronts into the parental molecular clouds.  Up to now it is not fully understood if this violent feedback enhances or hinders star formation. \\citet{1977ApJ...214..725E} proposed that the shock front builds up dense regions by sweeping up the cold gas, which then eventually collapse due to gravitational instability and form stars. This is called the 'collect and collapse model' (see also the review by \\citealt{1998ASPC..148..150E}). Another situation arises when preexisting, dense structures (e.g. molecular cloud cores) that are gravitationally marginally stable get compressed by the approaching front and start collapsing. This is commonly called radiation driven implosion (see e.g. \\citealt{1982ApJ...260..183S}).  Observations provide widespread evidence for these scenarios (\\citealt{1989ApJ...342L..87S}, \\citealt{1996AJ....111.2349H}). More recent observations indicate triggered star formation on the edges of HII-regions e.g. in the Orion clouds \\citep{2002A&A...393..251S}, the Carina nebula \\citep{2000ApJ...532L.145S}, M16 \\citep{2002ApJ...568L.127F}, M17 \\citep{2002ApJ...577..245J}, 30 Dor \\citep{2002AJ....124.1601W} and the SMC \\citep{2007arXiv0710.1352G}. \\citet{2005A&A...433..565D} report triggered star formation in samples of more distant HII regions. Besides these quite complex large scale regions there have been numerous observations of bright rimmed cometary globules. These are small isolated clouds with a clear head to tail structure with the dense heads pointing towards an ionizing source \\citep{1991ApJS...77...59S}. Their morphology enables a direct comparison to simulations. In particular, the properties of individual young stellar objects (YSOs) surrounding OB-associations can be determined precisely. YSOs are observed in the mass range from T Tauri ($0.1-3 \\textrm{M}_{\\odot}$) up to Herbig Ae Be ($2-8\\textrm{ M}_{\\odot}$) stars (see e.g. \\citealt{2007ApJ...657..884L}, \\citealt{2007arXiv0711.1515S}). The velocity of the shock front triggering the star formation is calculated from the age difference of the stars and their relative distance. These estimates are in the range of a few km/s (e.g. \\citealt{2004A&A...414.1017T}, \\citealt{2007ApJ...654..316G}).  Numerous simulations on the topic of cloud evaporation and sequential star formation have been performed. \\cite{1982A&A...108...25Y} (and references therein) published a series of two-dimensional simulations on the gas dynamics of HII-regions, especially on champagne flows, where a stream of hot gas breaks through the border of cold, confining gas. Subsequently, \\citet{1995ApJ...451..675E} presented two-dimensional, grid-based simulations showing that the expansion of an HII-region into the surrounding cloud can trigger star formation. \\citet{2003MNRAS.338..545K} demonstrated with a three-dimensional SPH code, that a marginally stable molecular cloud core can be triggered into collapse when exposed to strong UV radiation. With a more detailed description of radiation implemented into an SPH code \\cite{2006MNRAS.369..143M} could reproduce the observed features of the Eagle Nebula, including the photodissociation regions and the temperature profile. Using a three-dimensional grid-based scheme, \\citet{2006ApJ...647..397M} simulated the HII-region excavated by a point source of UV-radiation.  They find remarkably similar morphologies and physical properties when comparing their models to observations. Furthermore simulations with an SPH-code by \\citet{2005MNRAS.358..291D} and a grid code by \\citet{2007ApJ...668..980M} showed that a turbulent interstellar medium surrounding an O-star allows the ionizing radiation to efficiently expel most of the nearby gas. Only the denser regions resist and continue to collapse.  However none of the authors so far described ionizing radiation as an efficient trigger for star formation. There is only weak evidence by \\citet{2007MNRAS.377..535D}, that the external irradiation of a collapsing cloud by a point source can indeed increase the star formation efficiency from 3\\% to 4\\% when compared to a control run without radiation. For a review of feedback processes we refer the reader to \\citet{2007arXiv0711.4047M}. For completeness we would like to refer to recently published implementations for ionizing radiation into an SPH code by \\citet{2008MNRAS.389..651P}, where the photons of a source are followed along cones, and \\citet{2008MNRAS.386.1931A}, where the radiation is followed via a Monte Carlo ray-tracing scheme  All these studies demonstrate that there is a strong connection between the UV-radiative feedback from massive stars and the observed morphologies of the ambient molecular cloud gas. Yet, a quantitative discussion of the interaction between UV-radiation and turbulent molecular clouds is still missing. To advance our understanding, we introduce iVINE, the fully parallel implementation of UV-radiation in the parallel tree-SPH-code VINE (\\citealt{2008arXiv0802.4245W}, \\citealt{2008arXiv0802.4253N}). This efficient tool permits high resolution simulations of molecular clouds in the vicinity of strong UV sources such as an O-star or association.  The paper is structured as follows. The physical model and its implementation are described in Section \\ref{NM}, followed by a detailed comparison of the scheme with analytical results (Sec.~\\ref{NT}). We apply the new method to the radiatively driven implosion of a marginally stable molecular cloud core and compare three simulations with different initial UV fluxes to observations (Sec~\\ref{simulations}). In Section~\\ref{summary} we summarize and discuss the results. ",
        "conclusions": "\\label{summary} We present iVINE, a new implementation of UV-radiation into the tree-SPH code VINE. It uses a plane-parallel geometry which renders the code most suitable to perform high resolution studies of the small scale effects of e.g. ionization and turbulence in the surrounding of young massive stars.  It is efficiently parallelized and very fast, as only 2\\%-8\\% of the total computational time are used for the calculation of ionization. The comparison with analytic solutions shows that iVINE treats time-dependent ionization as well as the resulting heating effects precisely and convergently.  We base our numerical implementation of ionizing radiation on several assumptions. First, we use a simplified prescription for the radiative transfer by e.g assuming a monochromatic flux. Second, we neglect UV absorption by dust, which would lower the total UV flux. Third, we do not include a full treatment of recombination zones. In our simulations the ionized gas which gets shaded is assumed to recombine immediately. In addition, the gas in the shaded regions does not get heated by irradiation from the hot gas surrounding it. These effects require a precise time-dependent treatment of heating and cooling processes by ionization and recombination as well as a treatment for the scattering of photons. An implementation of this effects is planned in a future version of the code.  As an application we investigate radiation driven implosion of a marginally stable Bonnor-Ebert sphere. We show that these spheres are indeed driven into gravitational collapse. The resulting cores are in the observed mass range. They are more compact and a factor of $\\approx 10$ more dense than it would be expected in a more quiescent environment. This fact fits very well with the observations of star formation in a clustered environment.  By comparing simulations with three different UV-fluxes we show that there is a clear dependence of the final mass and the age of the collapsed core on the position of the preexisting density enhancement relative to the Str{\\\"o}mgren radius. Our findings that the onset of star formation is delayed by $0.08-0.4$Myr, depending on the position, are in good agreement with observations of the age spread in bright rimmed clouds. The velocity of the triggering shock is an order of magnitude higher than the observational estimates. This discrepancy has been noted before. We suggest that this can be attributed to the ionizing radiation stripping the envelope from a Class 0/I star. Thereby it might be classified as an Class II/III star, leading to an higher age-spread between the observed objects. Correcting for this effect would increase the estimated velocity of the shock front and thus lead simulations and observations towards agreement."
    },
    "0809/0809.2540_arXiv.txt": {
        "abstract": "We briefly review  the standard derivation of the spectra of cosmological perturbations in the simplest models of inflation. We then consider models with several scalar fields, described by  Lagrangians with an arbitrary dependence on the  kinetic terms. We  illustrate our general formalism with the case of multi-field DBI inflation. ",
        "introduction": "Inflation has now become a standard paradigm to describe the physics of the very early universe, and as cosmological data keep improving, one can hope to get more and more clues about the very early universe. So far, the simplest models of inflation are compatible with the data but it is instructive to study more refined models for at least two reasons. First, because models inspired by high energy physics are usually more complicated than the simplest phenomenological inflationary models. Second, because these generalized models will  give us an idea of how much the future data will be able to pin down some specific region in parameter space.  In the first part of this contribution, we present a brief summary of the standard approach to compute the cosmological perturbations in the simplest inflationary models. This summary is mainly based on the pedagogical introduction \\cite{cargese} where the reader will find more details.  In the second part, we show how the standard results are modified when several scalar fields play a role during inflation. We present  recent results for very general multi-field inflationary models, allowing for non-standard kinetic terms. This generalization is  motivated by efforts  to connect  string theory and inflation and we focus our attention on multi-field DBI (Dirac-Born-Infeld) inflation. ",
        "conclusions": "\\label{sec:conclusion} In this contribution, we have presented  a general analysis of cosmological perturbations in multi-field inflationary models, allowing for non standard kinetic terms. This approach generalizes much more restrictive situations  considered previously in the literature and is motivated by recent constructions of inflationary models in the context of string theory, where multiple fields and non standard kinetic terms are very common, a typical example being DBI inflation.  As we have tried to show, multi-field inflation is a very rich playground, where entropy modes can play a significant role. The most important consequence of entropy modes is the possibility to  modify the curvature perturbation, on large scales, in contrast with single field inflation. This means that the adiabatic fluctuations, which we observe today in the CMB, could come originally from entropy perturbations produced during multi-field inflation.  \\vspace{0.3cm} {\\bf Acknowledgment:} I would like to thank the organizers for inviting me to a very friendly and stimulating conference.       \\vspace{1cm}"
    },
    "0809/0809.2295_arXiv.txt": {
        "abstract": "The projected distribution of stars in the Small Magellanic Cloud (SMC) from the Magellanic Clouds Photometric Survey is analysed. Stars of different ages are selected via criteria based on $V$ magnitude and $V-I$ colour, and the degree of  `grouping' as a function of age is studied.  We quantify the degree of structure using the two-point correlation function and a method based on the Minimum Spanning Tree and  find that the overall structure of the SMC is evolving from a high degree of sub-structure at young ages ($\\sim$10\\,Myr) to a smooth radial density profile. This transition is gradual and at $\\sim$75\\,Myr the distribution is statistically indistinguishable from the background SMC distribution. This time-scale corresponds to approximately the dynamical crossing time of stars in the SMC. The spatial positions of the star clusters in the SMC show a similar evolution of spatial distribution with age. Our analysis suggests that stars form with a high degree of (fractal) sub-structure, probably imprinted by the turbulent nature of the gas from which they form, which is erased by random motions in the galactic potential on a time-scale of a galactic crossing time. ",
        "introduction": "\\label{sec:intro} The\\blfootnote{{\\bf$^\\star$}E-mail: mgieles@eso.org} majority of stars form in clustered environments (e.g. \\citealt*{2000prpl.conf..151C} and \\citealt{2003ARA&A..41...57L}, henceforth LL03), but only a small fraction ($\\lesssim10$ per cent) of stars ends up in bound star clusters which survive the embedded phase.  It was noted much earlier that almost all stars in the Galactic disc form in (unbound) OB associations  (e.g. \\citealt{1957PASP...69...59R, 1979ApJS...41..513M}) and that residual gas expulsion is the most probable explanation (e.g. \\citealt*{1980ApJ...235..986H, 1984ApJ...285..141L}). LL03 introduced the term `infant mortality' for the rapid dissolution of embedded clusters at young (few Myrs) ages.  More recently, the infant mortality scenario has also been used to explain the strong drop in the age distribution, \\dndt, around $\\sim$10-20\\,Myr of clusters in the Antennae galaxies and M51 (\\citealt*{2005ApJ...631L.133F} and  \\citealt{2005A&A...431..905B}, respectively).  \\citet{2007ApJ...658L..87P} introduced a new approach for studying infant mortality, by comparing the spatial distribution of stars of different stellar types in the nearby spiral galaxy NGC~1313. They found that O-type stars are more strongly grouped than B-type stars, which is what is expected if all stars form in clusters and the majority disperses on time-scales comparable to the life-time of O-stars (few Myrs).  The term infant mortality is now used in widely varying contexts. First, the time-scales involved range from $~3\\,$Myr (LL03) up to a Gyr \\citep{2005ApJ...631L.133F,2006ApJ...650L.111C}. Second, the term `star cluster' is used for irregular, dynamically unmixed, groups of embedded stars with total masses of $10-100\\,\\msun$ (e.g.~LL03), for compact  (few pc) and massive (few times $10^5-10^6\\,\\msun$) clusters in interacting galaxies (e.g.~\\citealt{1996AJ....112.1839S}) and for  large scale ($\\gtrsim100\\,$pc) stellar associations \\citep{2007ApJ...658L..87P}.  The Small Magellanic Cloud (SMC) cluster system has been heavily under debate recently. \\citet{2005AJ....129.2701R} presented a new catalogue of SMC star clusters and showed the existence of a declining \\dndt. This was  interpreted as infant mortality working on a 3 Gyr time-scale by \\citet{2006ApJ...650L.111C}.  However, \\citet*{2007ApJ...668..268G} showed that the decline  is in fact due to detection incompleteness at higher ages,  later confirmed by an independent analysis \\citep{2008MNRAS.383.1000D}.   This study aims at tying the above issues together to give a more coherent picture of the evolution of structures in the SMC. Inspired by the work of \\citet{2007ApJ...658L..87P}, the catalogue of five million stars detected in the Magellanic Clouds Photometric Survey (MCPS) \\citep{2002AJ....123..855Z} is used to select stars of different ages. In section~\\ref{sec:data} the data and the selection criteria of different regions are described. In section~\\ref{sec:clustering} two techniques are employed in order to search for stellar groups and quantify their evolution with age. The implications and conclusion are presented in  section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have studied the evolution of sub-structure in the distributions of stars of different ages in the Small Magellanic Cloud (SMC). Using two different methods, namely the two-point correlation function (TPCF), and the statistical $Q$-parameter, we find that the stellar distribution evolves to the   reference distribution, that is, a smooth power-law density profile,  on a time-scale of $\\sim$75\\,Myr.  Our results qualitatively agree with what  \\citet{2007ApJ...658L..87P} found for stars in NGC~1313. However, we have also quantitatively determined the time-scale in which `stellar groupings' disappear ($\\sim$75\\,Myr). In the scenario of \\citet{2007ApJ...658L..87P} all stars form in clusters/groups, which dissolve due to gas expulsion. The super-virial velocities of stars after gas removal are of the order of a few $\\kms$. Combined with the radius of the SMC galaxy ($\\rsmc\\simeq2\\,$kpc), the time it would take for stars that travel from a dissolving cluster to cover one galactic radius is $\\sim$1\\,Gyr,  which is much longer than the 75\\,Myr we find. \\citet{2008MNRAS.386..826E} measured a velocity dispersion, $\\sigma$, of $30\\,\\kms$ for (massive) field stars in the SMC, which combined with the radius of the SMC results in a dynamical crossing time, $\\tcr=\\rsmc/\\sigma\\simeq75\\,$Myr. This is the same as the time-scale for structure dispersal that we found in our study, making random motions in the galactic potential a plausible explanation for the erasure of structure.  \\citet{2008A&A...482..165G} showed that only a small fraction ($\\sim$3 per cent) of stars that are formed in the SMC end up in  clusters that survive past the embedded phase, i.e. those that can be identified as genuine star clusters in optical studies. Whether this is due to a 97 per cent infant mortality rate of embedded clusters or due to a large fraction of stars forming in the field, can not be judged from the current data.  We do know that already at very young ages ($\\sim$10\\,Myr) the majority of stars are not in dense clusters and these `field stars'  still show the  fractal imprints of the natal gas, which is being erased dynamically  by random motions in the galaxy potential.  In an accompanying paper \\citep{bastianetal08} we further develop and explain the techniques and conclusions reported here, and compare these results to those found for the Large Magellanic Cloud."
    },
    "0809/0809.0403_arXiv.txt": {
        "abstract": "We present the results of a 10.5\\,yr, volume limited (28\\,Mpc) search for supernova (SN) progenitor stars. In doing so we compile all SNe discovered within this volume (132, of which 27\\% are type Ia) and determine the relative rates of each sub-type from literature studies. The core-collapse SNe break down into 59\\% II-P and 29\\% Ib/c, with the remainder being IIb (5\\%), IIn(4\\%) and II-L (3\\%). There have been 20 II-P SNe with high quality optical or near-IR pre-explosion images that allow a meaningful search for the progenitor stars. In five cases they are clearly red supergiants, one case is unconstrained, two fall on compact coeval star clusters and the other twelve have no progenitor detected. We review and update all the available data for the host galaxies and SN environments (distance, metallicity and extinction) and determine masses and upper mass estimates for these 20 progenitor stars using the STARS stellar evolutionary code and a single consistent homogeneous method. A maximum likelihood calculation suggests that the minimum stellar mass for a type II-P to form is $m_{min}=8.5^{+1}_{-1.5}$\\msol\\ and the maximum mass for II-P progenitors is $m_{max}=16.5\\pm1.5$\\msol, assuming a Salpeter initial mass function holds for the progenitor population (in the range $\\Gamma = -1.35^{+0.3}_{-0.7}$). The minimum mass is consistent with current estimates for the upper limit to white dwarf progenitor masses, but the maximum mass does not appear consistent with massive star populations in Local Group galaxies. Red supergiants in the Local Group have masses up to 25\\msol\\ and the minimum mass to produce a Wolf-Rayet star in single star evolution (between solar and LMC metallicity) is similarly 25-30\\msol. The reason we have not detected any high mass red supergiant progenitors above 17\\msol\\ is unclear, but we estimate that it is statistically significant at 2.4$\\sigma$ confidence. Two simple reasons for this could be that we have systematically underestimated the progenitor masses due to dust extinction or that stars between 17-25\\msol\\ produce other kinds of SNe which are not II-P. We discuss these possibilities and find that neither provides a satisfactory solution. We term this discrepancy the ``red supergiant problem'' and speculate that these stars could have core masses high enough to form black holes and SNe which are too faint to have been detected. We compare the $^{56}$Ni masses ejected in the SNe to the progenitor mass estimates and find that low luminosity SNe with low $^{56}$Ni production are most likely to arise from explosions of low mass progenitors near the mass threshold that can produce a core-collapse. ",
        "introduction": "Stars which are born with masses above a critical threshold mass of around 8\\msol\\ have long been thought to produce supernovae when their cores collapse at a point when nuclear burning no longer provides support against gravity. Supernovae were first suggested to be a new class of astrophysical phenomena by \\cite{1934PNAS...20..254B} and since then detailed study has allowed them to be split into physical types : the thermonuclear explosions and core-collapse supernovae (CCSNe). The CCSNe form when their cores evolve to iron white dwarfs and detailed stellar evolutionary models predict a minimum mass for this to occur of between 7-12\\msol \\citep{2003ApJ...591..288H,2004MNRAS.353...87E,2007A&A...476..893S,2008ApJ...675..614P} . The CCSNe form a diverse group in terms of their spectral and photometric properties and the classification scheme has arisen primarily based on the appearance of their optical spectra, but also supplemented with their photometric behaviour \\citep{1997ARA&A..35..309F}. There are the H-rich type II SNe, which are sub-classified into II-P (plateau lightcurves), II-L (linear decline lightcurves), IIn (narrow emission lines), and some peculiar events, generically labelled II-pec. The H-deficient SNe are split into Ib and Ic depending on whether He is visible and a hybrid class of IIb (type II events which metamorphose into Ib SNe) has also been uncovered. The evolutionary stage of the progenitor star (i.e. its position in the HR diagram and chemical composition of its atmosphere) very likely dictates what type of CCSN is produced.  Determining what types of stars produce which types of SNe and what mass range can support a CCSN is a major goal in modern studies of these explosions. The first attempts relied primarily on linking the results of stellar evolutionary models to the spectral (and photometric) evolution of the SNe. \\citep[e.g.][]{1980ApJ...237..541A,1976ApJ...207..872C} More recently \\cite{2003ApJ...582..905H} and \\cite{2003MNRAS.346...97N} have studied the lightcurves and ejecta velocities of II-P SNe to estimate the mass of the material ejected. These studies tend to favour quite large masses for progenitor stars, with Hamuy suggesting ranges of 10-50\\msol\\ and Nadyozhin 10-30\\msol. However when SN~1987A exploded a new opportunity arose. The SN was in a galaxy close enough that its progenitor star could be easily identified. A blue supergiant star of around 15-20\\msol\\ was identified and it is clear that this object no longer exists \\citep{1987Natur.328..318G,1989A&A...219..229W,1992PASP..104..717P} The fact that it was a compact blue supergiant was a key factor in enabling the community to understand the event as a whole. A progenitor was also detected for the next closest explosion (SN~1993J in M81) and the binary nature of the progenitor helped understand the physical reason behind the II-b type given to the event \\citep{1994AJ....107..662A,1993Natur.364..509P}.   Studies of the environments of CCSNe after discovery have been ongoing for many years with authors looking for correlations between the ages of starformation regions and the type of SNe that occur. For example \\cite{1992AJ....103.1788V} and \\cite{1996AJ....111.2017V} suggested that there was no clear difference in the spatial distributions of type Ib/c and type II SNe compared to giant \\Hii\\ regions in their host galaxies. They concluded that they hence arose from parent populations of similar mass. However recently both \\cite{2006A&A...453...57J} and \\cite{2007arXiv0712.0430K} suggest that the stripped type Ic SNe are more likely to follow regions of either high surface brightness or high H$\\alpha$ emission than type II SNe. A larger sample of nearby core-collapse events and H$\\alpha$ correlations has been compiled by \\citep{2008MNRAS.390.1527A} indicating that there is a progressive trend for Ib/c SNe to be associated with H$\\alpha$ emission regions, in the sense that Ic show the closest association with galactic H$\\alpha$ emission, then comes the Ib and then type II.  While these efforts are very valuable to discern differences in progenitor channel, it is difficult to assign definitive mass ranges to the progenitor systems from spatial correlations alone.   A much more direct way to determine the type of star that exploded is to search directly for progenitors in images of the host galaxies taken before explosion. The ease of access to large telescope data archives makes this search feasible for nearby events. There are now a number of groups around the world that are competitively searching for progenitor stars in such archive images. Particularly with the  {\\em Hubble Space Telescope} (HST) it has become possible to resolve massive stellar populations out to at least  20Mpc. In this case the information on each progenitor is much more detailed and quantitative than can be achieved with the unresolved environment studies, but there are fewer events which allow such a study.  Early work with HST concentrated on looking at the environments of SNe \\citep{1996AJ....111.2047B,1999AJ....118.2331V}. But it has now become possible to search directly for progenitor stars and carry out concerted campaigns on the nearest events. Three of the first type II SNe to have excellent HST and ground-based pre-explosion images were the II-P SNe 1999gi, 1999em and 2001du \\citep{2001ApJ...556L..29S,2002ApJ...565.1089S,2003MNRAS.343..735S,2003PASP..115..448V}. In all there was no detection of a progenitor, but meaningful limits were derived. \\cite{2003ApJ...594..247L} and \\cite{2002AJ....124.2490L} subsequently showed the importance of having reliable distances to the host galaxies of the SN progenitors in order to constrain the upper mass limits. Efforts to find progenitors continued \\citep{2003MNRAS.343..735S,2003PASP..115....1V} until the first confirmation of a red supergiant progenitor of a type II-P explosion in prediscovery HST and Gemini-North images \\citep{2004Sci...303..499S}. Unambiguous detections in HST images require the SN to be located on the pre-explosion images with accuracies of around 10 milliarcseconds, which requires follow-up images of the SN to be taken at HST or corrected adaptive optics ground-based resolution \\citep{2005MNRAS.360..288M,2005ApJ...630L..29G,2008MNRAS.391L...5C} The next clear and unambiguous detection of a progenitor star was also a red supergiant in NGC5194 \\citep[SN2005cs ; ][]{2005MNRAS.364L..33M,2006ApJ...641.1060L}. And the recent discovery of SN2008bk  in NGC7793 (3.9\\,Mpc) has produced detections in $IJHK$ of a red progenitor star \\citep{mattila08bk}. The near-IR spectral energy (SED) distribution of the SN2008bk progenitor is the best sampled SED yet of any red supergiant progenitor and matches a late M4I spectral type with moderate extinction of $A_{V} \\simeq 1$ and an initial  mass around 8-9\\msol. There have been claims of the detections of others such as 2004A \\citep{2006MNRAS.369.1303H}, 2004et \\citep{2005PASP..117..121L}, 2006my and 2006ov \\citep{2007ApJ...661.1013L} and we review these in this paper. An additional method that has much to offer this field is locating SNe directly coincident with compact, coeval star clusters. If the age of the cluster can be determined then a main-sequence turn-off age, and turn-off mass can lead directly to an estimate of the progenitor star mass \\cite[e.g. 2004dj in NCG2403 as studied by][]{2004ApJ...615L.113M}. The detection and characterisation of progenitor stars has the potential to directly link the type of star to the SN explosion characteristics, (e.g. the amount of $^{56}$Ni synthesised in explosive O and Si burning and the total energy of the explosion) and also to set quantitative limits on progenitor mass ranges. The latter is of great interest to compare to the highest mass white dwarf progenitors and to the stars that form neutron stars and black holes after core-collapse.  Work in this field has progressed substantially in the last 8 years and compilations of progenitor properties have been made by \\cite{2003MNRAS.343..735S,2007ApJ...656..372G,2007ApJ...661.1013L,2008arXiv0802.0456K}. However these compilations are somewhat ad-hoc, incomplete and potentially biased as they do not define the selection criterion for inclusion rigorously. In addition the results are based on different methods for estimating the progenitor masses and upper limits (in terms of measurement and the theoretical models employed). The goal of this series of two papers is to define the selection criteria for inclusion (a volume and time limited survey) and to determine the physical parameters of the progenitor stars (luminosities and masses, or limits thereon) in a homogeneous and consistent way. Only then is it possible to reliably estimate the population parameters. This paper specifically deals with the type II SNe and all but two of the final sample of twenty have been confirmed as type II-P. A companion paper will discuss the stripped SNe of types IIb, Ib and Ic which are drawn from the same volume and time limited survey (Crockett et al. 2008, in prep).  In addition to the supernovae with pre-explosion images we discuss in the paper there are three SNe with progenitor detections that fall outside either our time or volume definition. Those are SN~1987A, SN~1993J and SN~2005gl. Discussions of the first two are well documented in the literature \\cite[e.g.][]{1989A&A...219..229W,1994AJ....107..662A,2002PASP..114.1322V,2004Natur.427..129M}; 1987A will be discussed later in this paper and the implications of 1993J will be discussed in the 2nd paper in this series. SN 2005gl is a type IIn in NGC266 at approximately 66\\,Mpc, hence although a progenitor is detected by \\cite{2007ApJ...656..372G} it does not fall within our distance limit. \\cite{2007ApJ...656..372G} suggest that this was a very massive star with $M_{V} \\simeq -10.3$ and very likely a luminous blue variable. Although with only a detection in a single filter the blue colour is not confirmed, and there is not enough data to determine if the source was indeed variable. This is evidence that very massive stars do explode as bright SNe and is a point we will return to in the discussion.  This paper starts with defining the sample from which the targets with high quality pre-explosion images are drawn. A consistent and homogeneous analysis method is then defined and justified. We then review previous detections (and add some new data) to build the data required for analysis and then statistically analyse the results. We follow this with an extensive discussion. ",
        "conclusions": "This paper presents a consistent and homogeneous analysis of the constraints on progenitor stars of II-P SNe, within a volume and time limited search. The work builds and enhances the discoveries and limits presented in the literature. There are now enough data that a statistical study of progenitor properties is feasible. This represents the culmination of a a 10.5 year search for progenitor stars and the conclusions are as follows:  \\begin{itemize} \\item We have compiled all SNe discovered within a strict volume and time limit and reviewed the SN types reported in the literature. This gives an very good estimate of the relative rate of CCNSe in the Local Universe, at metallicities between LMC and solar. \\item Of the 55 type II-P SNe found, 20 have HST pre-explosion (or excellent ground-based) images available to search for progenitor stars. We summarise the data presented in the literature to date on these SNe and carry out a consistent and homogeneous analysis of all events. Three groups of events are discussed - those with probable single star progenitors detected, those with upper limits to their luminosities and masses, and those falling on compact, unresolved coeval clusters \\item The masses and mass limits of each are determined using our STARS evolutionary models and a statistical analyses presented of the final masses. A maximum likelihood analysis suggests that the minimum mass for a type II-P SN is $m_{min}=8.5^{+1}_{-1.5}$\\msol. This is consistent with current estimates of the maximum mass that will produce a white dwarf in young clusters. \\item We have not detected any progenitors above 16\\msol. Assuming a Salpeter IMF, the most likely maximum mass for a II-P SN is $m_{max}=16.5\\pm1.5$\\msol.This is not particularly sensitive to the IMF slope within the typical variations known in the Local Universes of $\\pm0.7$. \\item We suggest that there is a discrepancy between this maximum mass and our knowledge of massive star evolution. Red supergiants between 17 - 30\\msol\\ are not detected as progenitors but are predicted by theory to exist and this is supported by stellar population studies. We term this discrepancy the ``red supergiant problem''. It is unlikely to be due to IMF variations and possible explanations include the possibility that we have systematically under estimated the stellar luminosities and masses due to foreground extinction or that the gap is filled with the other flavours of type II SNe (e.g. II-L, IIn and IIb). However neither of these solutions are supported by current data. We suggest that these objects may be forming black-holes with faint, or non-existent SN explosions. \\item Although low-luminosity SNe are already known (e.g. 1999br, 2005cs) we suggest that these events are not the missing SNe - rather our analysis supports the interpretation that they are stars at the minimum mass limit for SNe. \\item We review the information on extinctions toward these SNe and their progenitors and compare it with that toward red supergiant populations in the LMC. We suggest that high extinction toward the SNe progenitors is unlikely to be the cause of the lack of detections of massive supergiants. \\item The search for progenitor stars should continue for every nearby SN that has deep, high resolution pre-explosion images. In particular the missing high mass red supergiants should be sought in both optical, near infra-red and mid infra-red images. If the limit for black hole formation is as low as 17\\msol\\ then it bodes well for surveys for disappearing stars as suggested by \\cite{2008arXiv0802.0456K}.  \\end{itemize}"
    },
    "0809/0809.5226_arXiv.txt": {
        "abstract": "We consider a multidimensional cosmological model with nonlinear quadratic $R^2$ and quartic $R^4$ actions. As a matter source, we include a monopole form field, D-dimensional bare cosmological constant and tensions of branes located in fixed points. In the spirit of the Universal Extra Dimensions models, the Standard Model fields are not localized on branes but can move in the bulk. We define conditions which ensure the stable compactification of the internal space in zero minimum of the effective potentials. Such effective potentials may have rather complicated form with a number of local minima, maxima and saddle points. Then, we investigate inflation in these models. It is shown that $R^2$ and $R^4$ models can have up to 10 and 22 e-foldings, respectively. These values are not sufficient to solve the homogeneity and isotropy problem but big enough to explain the recent CMB data. Additionally, $R^4$ model can provide conditions for eternal topological inflation. However, the main drawback of the given inflationary models consists in a value of spectral index $n_s$ which is less than observable now $n_s\\approx 1$. For example, in the case of $R^4$ model we find $n_s \\approx 0.61$. ",
        "introduction": "\\setcounter{equation}{0}   Recently, very elegant idea of inflation has achieved spectacular success in explaining the acoustic peak structure seen in CMB (see e.g. \\cite{Bennett}). It is very difficult to explain correctly the observable large-scale structure formation without taking into account the stage of early inflation. There is a big number of different models of inflation. However, the most of them are poorly related to fundamental physics. In these models, the stage of inflation occurs due to a special form of scalar field potential. Here, the origin of scalar field and the form of its potential is usually remained out of the scope of these papers. However, it is well known that scalar fields has naturel origin from higher-dimensional theories. They are geometrical moduli (radions, gravexcitons) which are related to the shape of internal spaces (to scale factors of the internal spaces). After dimensional reduction to four dimensions, scalar field potential is completely defined by the topology and matter content of original higher-dimensional model \\cite{RZ,GZ1}. Thus, it is of undoubted interest to realize inflation in these models (see e.g. \\cite{KSS,LW,CKLQ} in string theory and \\cite{GZ2} in multidimensional cosmological models and references therein).    On the other hand, scalar fields with corresponding potentials originate naturally from nonlinear gravitational models where Lagrangian is a function of scalar curvature $f(R)$. It is well known that such models are equivalent to linear ones plus scalar field. This scalar field corresponds to an additional degree of freedom of nonlinear models. For motivation of these theories, see review \\cite{SF}. For example, among others higher-order gravity theories, $f(R)$ theories are free of ghosts and of Ostrogradski instability \\cite{Woodard}. These models attract great attention because can provide the late-time acceleration of our Universe due to a special form of scalar field potentials (see e.g. \\cite{SF,NO,GZBR} and references therein), which is interesting alternative to the cosmological constant. These models can also provide the stage of early inflation both in four-dimensional (see the pioneering paper by Starobinsky \\cite{Star} and numerous references in \\cite{SF,NO}) and multidimensional \\cite{GZBR,GMZ1,BR} cases.  In our paper we combine both of these approaches. Namely, we consider multidimensional models with nonlinear action. To start with, we investigate the most simple linear multidimensional model. We show that such model can provide power-law inflation. Unfortunately, it takes place for the branch where internal space is decompactified \\footnote{Each time when we consider multidimensional cosmological models we should remember about the problem of the internal space stabilization/compactification. If such stabilization is absent we confront with the variation of four-dimensional fundamental constants. General method of the internal space stabilization was described in \\cite{GZ1} and was applied after that to numerous models. In the present paper, we follow also this method.}. Then, to get inflation of the external space with subsequent stabilization of the internal spaces, we turn to multidimensional nonlinear models with quadratic and quartic nonlinearities. First, we obtain the conditions of the internal space compactification (stabilization). Second, for corresponding effective potentials, we investigate the possibility of the external space inflation. We show that in the quadratic and quartic models we can achieve 10 and 22 e-folds, respectively. These numbers are sufficient to explain the present day CMB date, but not enough to solve the horizon and flatness problems. However, 22 e-foldings is rather big number to encourage the following investigation of the nonlinear multidimensional models to find theories where this number will approach to 50-60 e-folds. Even more, this number (50-60) can be reduced in models with long matter dominated stage followed by inflation with subsequent decay into radiation. Precisely this scenario takes place for our models where we find that e-folds can be reduced by 6 if the mass of decaying scalar field $m\\sim 1$TeV. So, we believe that the number of e-folds is not a big problem for proposed models. The main problem consists in spectral index $n_s \\approx 0.61$ (for the quartic model) which is less than observable $n_s \\approx 1$. A possible solution of this problem may consist in more general form of the nonlinearity $f(R)$. For example, it was observed in \\cite{Ellis} that simultaneous consideration quadratic and quartic nonlinearities can flatten the effective potential. We postpone the investigation of this problem for our following paper.  To conclude, we want to indicate two interesting features of models under consideration. First, the quartic model can provide the topological inflation. Here, due to quantum fluctuation of scalar fields, inflating domain wall has fractal structure (inflating domain wall will contain a number of new inflating domain walls and each such domain walls will contain again a new inflating walls etc \\cite{Linde}). So, we arrive at the eternal inflation. Second, obtained solution has the property of the self-similarity transformation (see Appendix B). It means that in the case of zero minimum of the effective potential and fixed positions for extrema in $(\\varphi,\\phi)-$plane, the change of the hight of extrema results in rescaling of the dynamical characteristics of the model (graphics of the number of e-folds, scalar fields, the Hubble parameter and the acceleration parameter versus synchronous time) along the time axis. The decrease (increase) of hight in $c$ times ($c$ is a constant) leads  to the stretch (shrink) of these figures along the time axis in $\\sqrt{c}$ times.  The paper is structured as follows. In Sec. II we consider the linear (on scalar curvature) model. The nonlinear quadratic $R^2$ and quartic $R^4$ models are investigated in Sec. III and Sec. IY, respectively. Here, we find the range of parameters where the internal space is stabilized and investigate a possibility for the external space inflation. A brief discussion of the obtained results is presented in the concluding Sec. Y.  The Friedmann equations for multi-component scalar field models are reduced to the system of dimensionless first order ODEs in Appendix A. In Appendix B, we show that the dynamical characteristics (e.g. the Hubble parameter, the acceleration parameter) of considered nonlinear models satisfy the self-similarity condition. ",
        "conclusions": ""
    },
    "0809/0809.4016_arXiv.txt": {
        "abstract": "\\noindent We have shown previously that many of the properties of persistent accretion-powered millisecond pulsars can be understood if their X-ray emitting areas are near their spin axes and move as the accretion rate and structure of the inner disk vary. Here we show that this ``nearly aligned moving spot model'' may also explain the intermittent accretion-powered pulsations that have been detected in three weakly magnetic accreting neutron stars. We show that movement of the emitting area from very close to the spin axis to $\\sim\\,$10$\\arcdeg$ away can increase the fractional rms amplitude from $\\la\\,$0.5\\%, which is usually undetectable with current instruments, to a few percent, which is easily detectable. The second harmonic of the spin frequency usually would not be detected, in agreement with observations. The model produces intermittently detectable oscillations for a range of emitting area sizes and beaming patterns, stellar masses and radii, and viewing directions. Intermittent oscillations are more likely in stars that are more compact. In addition to explaining the sudden appearance of accretion-powered millisecond oscillations in some neutron stars with millisecond spin periods, the model explains why accretion-powered millisecond oscillations are relatively rare and predicts that the persistent accretion-powered millisecond oscillations of other stars may become undetectable for brief intervals. It suggests why millisecond oscillations are frequently detected during the X-ray bursts of some neutron stars but not others and suggests mechanisms that could explain the occasional temporal association of intermittent accretion-powered oscillations with thermonuclear X-ray bursts. ",
        "introduction": "\\label{intro}  Observations made using the \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}) have led to the discovery of nine neutron stars in low-mass X-ray binary systems (\\mbox{LMXBs}) that produce persistent accretion-powered X-ray oscillations with periods equal to their millisecond spin periods. These accretion-powered millisecond X-ray pulsars (\\mbox{APMXPs}) are thought to be weakly magnetic stars that have been spun up by accreting angular momentum (see \\citealt{lamb08a} and references therein). Recently, three other neutron stars in \\mbox{LMXBs} have been found to produce accretion-powered millisecond X-ray oscillations intermittently. These \\textit{intermittent} \\mbox{APMXPs} are the focus of this \\textit {Letter}.  An intermittent accretion-powered X-ray oscillation was first detected in HETE~J1900.1$-$2455 \\citep{gall07}. The 377-Hz oscillation disappeared about two months after it was discovered, even though the star was still bright. Oscillations at $\\sim\\,$376~Hz have recently been discovered in X-ray bursts from this star (\\citealt{watt09}). A 442-Hz accretion-powered X-ray oscillation was serendipitously detected in a $\\sim\\,$500~s interval during a longer \\textit{RXTE} observation of SAX~J1748.9$-$2021 \\citep{gavr07}. An independent analysis of all the currently available \\textit{RXTE} data on this X-ray star detected a 442-Hz oscillation in 11 intervals during its 2001 and 2005 outbursts \\citep{alta08, patr09}. The oscillation appeared and disappeared on a timescale of several hundred seconds. A search of the 1.3~Ms of \\textit{RXTE} data available on Aql~X-1 discovered an accretion-powered X-ray oscillation at its 550-Hz spin frequency for 150~s during the peak of its 1998 outburst \\citep{case08}. The spin frequency of Aql~X-1 was previously known from its X-ray burst oscillations \\citep{zhan98}. Due to the poor temporal coverage of most \\mbox{LMXBs} and the sparseness of the intermittent oscillations discovered so far, the total number of intermittent \\mbox{APMXPs} in the Galaxy is unknown.  The accretion-powered oscillations of the known intermittent \\mbox{APMXPs} share several common characteristics. All have nearly sinusoidal waveforms and fractional rms amplitudes $\\la\\,$3\\%. The second harmonic (first overtone) of the spin frequency was rarely detected in HETE~J1900.1$-$2455; when it was, its amplitude was $\\la\\,$0.4\\% (\\citealt{gall07}). No second harmonic has been detected in the other two intermittent \\mbox{APMXPs}, with upper limits \\mbox{$\\sim\\,$0.4\\%--0.9\\%} (\\citealt{patr09, case08}). The amplitudes of the accretion-powered oscillations produced by HETE~J1900.1$-$2455 and SAX~J1748.9$-$2021 seem to increase shortly before or after their X-ray bursts (\\citealt{gall07, gavr07, alta08}).   \\begin{figure*}[t!] \\hspace{0pt} \\includegraphics[scale=0.33]{fig1a.eps} \\hspace{6pt} \\includegraphics[scale=0.33]{fig1b.eps} \\vspace{-6pt} \\caption{Fractional rms amplitudes of the full bolometric waveform (left) and its second harmonic component (right) as a function of the observer's inclination, for the spot inclinations shown. These waveforms are for two stable, isotropically emitting, 25$\\arcdeg$-radius antipodal spots on the surface of a $1.4M_\\odot$ star with a radius of $5\\,GM/c^2$, spinning at 400~Hz.} \\label{fig:fractional-amplitudes} \\end{figure*}   The nearly aligned moving spot model was originally proposed to explain the  properties of the persistent \\mbox{APMXPs} (see \\citealt{lamb06, lamb07b, lamb08b}; \\citealt{lamb09}, hereafter P1). In this model, the X-ray emitting areas are close to the spin axis but move around as the accretion rate and structure of the inner disk vary. As explained in P1, the model also has implications for the nuclear-powered oscillations observed in most of these pulsars during their thermonuclear X-ray bursts. Here we show that the model can also explain the properties of the intermittent \\mbox{APMXPs}. The reason is that small changes in the latitude of the emitting area can cause oscillations to appear and disappear if the emitting area is very close to the star's spin axis (see also \\citealt{lamb08b, boutloukos08}). In contrast to P1, which presented results for a wide range of spot inclinations, this paper focuses on spot inclinations $\\la\\,$15$\\arcdeg$. ",
        "conclusions": "\\label{discussion}  As explained in P1, the magnetic poles of accreting neutron stars are expected to move close to the spin axis as the stars spin up to millisecond periods. The resulting magnetic field will channel accreting gas toward the axis, causing the X-ray emitting area to be close to the spin axis. In this \\textit {Letter} we have shown that if the emitting area is centered within a few degrees of the spin axis, the observed flux modulation at the spin frequency will often be so small for a wide range of viewing angles that it will be undetectable using current instruments. Rapid fluctuations in the position of the emitting area will contribute to making the oscillation at the spin frequency undetectable. These effects may explain why accretion-powered oscillations have have been detected in so few accreting neutron stars with millisecond spin periods (see also P1).  We have shown further that if the emitting area wanders toward the spin equator by $\\sim\\,$10$\\arcdeg$, the oscillation can change from being undetectable to being easily detectable, with amplitudes comparable to those observed in the intermittent \\mbox{APMXPs}. In most cases, the amplitudes of the second and higher harmonics would be detectable only when the amplitude of the fundamental exceeds several percent.  Our computations show that this behavior occurs for a variety of spot geometries and radiation beaming patterns, and a wide range of spot radii, stellar masses and radii, and observer inclinations. Infrequent episodes of detectable accretion-powered oscillations at the spin frequency are more likely for stars that are more compact, because such oscillations will generally have lower amplitudes in these stars and are therefore less likely to be detectable.  In addition to providing a possible explanation for the sudden appearance of accretion-powered oscillations at the spin frequency, our model predicts that otherwise persistent oscillations may occasionally disappear if the center of the emitting area moves too close to the spin axis. Such \\textit {oscillation dropouts}, which could occur on a variety of timescales, have not yet been reported, but may be present in the existing data on known \\mbox{APMXPs} or found in data on \\mbox{APMXPs} discovered in the future. We encourage searches for such dropouts.  The basic physical picture of accreting millisecond X-ray pulsars proposed here and in P1 may also explain the observed pattern of millisecond oscillation detections in the X-ray bursts of different types of neutron stars.  Assume for the sake of argument that the magnetic poles of most weak-field accreting neutron stars in LMXBs are so closely aligned with their spin axes and the accretion flow so symmetric around the rotation pole that accretion-powered X-ray oscillations are undetectable (see P1). If the surface magnetic field is strong enough to create a symmetric pattern of nuclear burning around the spin axis (again see P1), X-ray burst oscillations will also be undetectable.  In P1 we conjectured that the magnetic fields of the persistent APMXPs SAX~J1808.4$-$3658 and XTE~J1814$-$338 are slightly asymmetric and strong enough to confine nuclear fuel, causing their nuclear burst oscillations to be locked or nearly locked to their accretion-powered oscillations, whereas the magnetic fields of other persistent APMXPs are too weak to force this.  The model of the intermittent APMXPs proposed here also suggests why accretion-powered oscillations and X-ray bursts may be associated in time, as appears to be the case in SAX~J1748.9$-$2021 \\citep{alta08, patr09}. The appearance of accretion-powered oscillations signals a shift in the impact point of the accreting matter. On the one hand, such a shift could be due to a change in the accretion flow through the inner disk. This is likely to change the thermal structure and fuel loading in the outer layers of the neutron star, which might trigger a nuclear burst. On the other hand, a nuclear burst is expected to temporarily change the accretion flow in the inner disk (see, e.g., \\citealt{mill96} and references therein), changing the impact point of the accreting matter and thereby possibly causing accretion-powered oscillations to become detectable."
    },
    "0809/0809.3894_arXiv.txt": {
        "abstract": "A search for a dark matter (DM) annihilation signal into $\\gamma$-rays toward the direction of the Canis Major (CMa) overdensity is presented. The nature of CMa is still controversial and one scenario represents it as a dwarf galaxy, making it an interesting candidate for DM annihilation searches. A total of 9.6 hours of high quality data were collected with the H.E.S.S. array of Imaging Atmospheric Cherenkov Telescopes (IACTs) and no evidence for a very high energy $\\gamma$-ray signal is found. Upper limits on the CMa dwarf galaxy mass of the order of 10$^{9}$ M$_{\\odot}$ are derived at the 95$\\%$ C.L. assuming neutralino masses in the range 500 GeV - 10 TeV and relatively large annihilation cross-sections. Constraints on the velocity-weighted annihilation cross section $\\langle\\sigma v\\rangle$, are calculated for specific WIMP scenarios, using a NFW model for the DM halo profile and taking advantage of numerical simulations of hierarchical structure formation. 95$\\%$ C.L. exclusion limits of the order of 5 $\\times$ 10$^{-24}$ cm$^{3}$ s$^{-1}$ are reached in the 500 GeV - 10 TeV DM particle mass interval, assuming a total halo mass of 3 $\\times$ 10$^{8}$ M$_{\\odot}$. ",
        "introduction": "It is widely believed that around one third of the energy content of the Universe is made of cold matter, of which only a small fraction is luminous and is of baryonic origin. In the framework of the Cold Dark Matter (CDM) scenarios, most of the matter is composed of non-baryonic Weakly Interacting Massive Particles (WIMPs). The standard model of particle physics does not provide a natural and suitable candidate for the Dark Matter (DM) particle and many theories beyond the standard model have been proposed to explain its origin and properties (see \\citep{Bertone} for a recent review). In various models, the self-annihilation of WIMPs gives a $\\gamma$-ray continuum emission resulting from the hadronization of primary annihilation products. The spectral features of such a DM annihilation radiation might help to distinguish it from ordinary astrophysical sources \\citep{Bergstrom1,Bergstrom2}. The indirect detection of DM through very high energy (VHE) $\\gamma$-rays is one of the best ways to probe the astrophysical nature of the DM and the H.E.S.S. array of Cherenkov telescopes, dedicated for detection of VHE $\\gamma$-rays in the 100 GeV-10 TeV energy regime, is an interesting instrument for this purpose.\\\\ Many astrophysical objects, ranging from DM clumps to galaxy clusters are expected to lead to DM particle annihilation signals detectable with sufficiently sensitive instruments. Regions of high concentration of DM are good candidates to search for such annihilations and the Galactic Centre (GC) was first considered. H.E.S.S. observations of the GC region \\citep{SgrA*} revealed a source of VHE $\\gamma$-ray emission (HESS J1745-290) but ruled out the bulk of the signal as of DM origin \\citep{Rolland}. Prospects for indirect detection in the elliptical galaxy M87 at the center of the Virgo cluster were also investigated. The time variability of the VHE $\\gamma$-ray signal observed by H.E.S.S. \\citep{M87} gave clear evidence that the signal was not of a sole DM origin. There are also other candidates with high DM density in relative proximity that might lead to detectable DM annihilation signals. Satellite dwarf galaxies of the Milky Way (MW) such as Sagittarius, Draco or Canis Major are popular targets, owing to their relatively low astrophysical background \\citep{EvansFerrer}. Indeed, dwarf spheroidal galaxies usually consist of stellar populations with no hot or warm gas and no cosmic rays, and are among the most extreme DM-dominated environments. A null result concerning the search for DM toward the Sagittarius dwarf spheroidal galaxy (Sgr dSph) direction was published by the H.E.S.S. collaboration \\citep{Moulin}. Constraints on the parameter spaces of two popular WIMPs models, namely the R-parity conserving Minimal Supersymmetric extension of the Standard Model (MSSM) and Kaluza-Klein (KK) scenarios with K-parity conservation, were derived. A null result was also established by the MAGIC collaboration \\citep{MAGIC} and the WHIPPLE collaboration \\citep{WHIPPLE}, when searching for a DM annihilation signal toward the Draco dwarf galaxy. Upper limits on the velocity-weighted annihilation cross-section of DM particle were derived in the framework of minimal SUper GRAvity (mSUGRA) models and MSSM models, respectively.\\\\ The present paper reports the search for a DM annihilation signal towards the direction of the CMa overdensity with the H.E.S.S. array of Cherenkov telescopes. The paper is organized as follows: in Sec 2 the controversial nature of the CMa overdensity is briefly discussed; in Sec 3 the analysis of the data is presented, while in Sec 4 the predictions for DM annihilation into $\\gamma$-rays in the CMa overdensity are discussed. Constraints on the WIMP velocity-weighted annihilation rate, as well as on the CMa total mass, are given. ",
        "conclusions": "The CMa overdensity is the subject of many debates over whether it is a dwarf galaxy or the warp and flare of the Galactic outer disk. Considering the first scenario, its relative proximity makes it potentially the best region for searches of a DM annihilation signal. However, the lack of observational data prevents the precise modelling of its density profile. Assuming a NFW profile and a mass content of 3 $\\times$ 10$^{8}$ M$_{\\odot}$ within its tidal radius, typical of dwarf galaxies, H.E.S.S. is close to exclude a few pMMSM scenarios with higgsino-like neutralinos, but does not reach the necessary sensitivity to test models compatible with the WMPA+SDSS constraint on the CDM relic density. In the case of DM made of B$^{(1)}$ particle from KK models with extra dimensions, no constraints are obtained."
    },
    "0809/0809.2156_arXiv.txt": {
        "abstract": "We inspect the coupled dependence of physical parameters of the Sloan Digital Sky Survey galaxies on the small-scale (distance to and morphology of the nearest neighbor galaxy) and the large-scale (background density smoothed over 20 nearby galaxies) environments. The impacts of interaction on galaxy properties are detected at least out to the neighbor separation corresponding to the virial radius of galaxies, which is typically between 200 and 400 $h^{-1}$ kpc for the galaxies in our sample. To detect these long-range interaction effects it is crucial to divide galaxy interactions into four cases dividing the morphology of target and neighbor galaxies into early and late types. We show that there are two characteristic neighbor-separation scales where the galaxy interactions cause abrupt changes in the properties of galaxies. The first scale is the virial radius of the nearest neighbor galaxy $r_{\\rm vir,nei}$. Many physical parameters start to deviate from those of extremely isolated galaxies at the projected neighbor separation $r_p$ of about $r_{\\rm vir,nei}$. The second scale is at $r_p \\approx 0.05 r_{\\rm vir,nei} = 10 -20 h^{-1}$ kpc, and is the scale at which the galaxies in pairs start to merge. We find that late-type neighbors enhance the star formation activity of galaxies while early-type neighbors reduce it, and that these effects occur within $r_{\\rm vir,nei}$. The hot halo gas and cold disk gas must be participating in the interactions at separations less than the virial radius of the galaxy plus dark halo system. Our results also show that the role of the large-scale density in determining galaxy properties is minimal once luminosity and morphology are fixed. We propose that the weak residual dependence of galaxy properties on the large-scale density is due to the dependence of the halo gas property on the large-scale density. ",
        "introduction": "According to the currently popular $\\Lambda$CDM model of structure formation, galaxies should form hierarchically. Numerical simulations demonstrate that galaxy-sized objects form through numerous interactions and mergers (e.g. Toomre \\& Toomre 1972; Hernquist 1992, 1993; Naab \\& Burkert 2003; Robertson et al. 2006; Maller et al. 2006). Therefore, the outcomes of galaxy-galaxy interactions and mergers are of key importance in understanding the hierarchical picture of galaxy formation. However, the impact of interactions and mergers on galaxy properties is still not known well even though there have been many studies.  One of the earliest works on this problem is Larson \\& Tinsley (1978) who showed that galaxies selected from the Atlas of Peculiar Galaxies (Arp 1966) have enhanced star formation (SF) rate compared to typical galaxies. Since that time, many studies have shown observational evidence of significant changes in galaxy properties due to interactions and/or mergers, for instance, in SF activity (Kennicutt \\& Keel 1984; Kennicutt et al. 1987; Bushouse, Werner \\& Lamb 1988; Barton et al. 2000, 2003, 2007; Lambas et al. 2003; Sanchez \\& Gonzalez-Serrano 2003; Nikolic, Cullen \\& Alexander 2004; Hernandez-Toledo et al. 2005; Geller et al. 2006; Li et al. 2008), galaxy structure parameters (Nikolic et al. 2004; Patton et al. 2005; Hernandez-Toledo et al. 2005; Coziol \\& Plauchu-Frayn 2007; Kacprzak et al. 2007), luminosity ratio (Woods et al. 2006; Woods \\& Geller 2007), and stellar mass ratio (Ellison et al. 2008). Recent  $N$-body simulations have also shown that the general features of the observations can be modeled (e.g. Barnes \\& Hernquist 1996; Mihos \\& Hernquist 1996; Tissera et al. 2002; Perez et al. 2006a,b; Di Matteo et al. 2007).  However, previous studies have not always been in agreement with one another. Yee \\& Ellingson (1995) and Patton et al. (1997) found no significant difference between the mean properties of isolated and paired galaxies. Bergvall et al. (2003) found no clear difference between isolated galaxies and interacting/merging systems in their global SF rate in Optical/near-IR bands. Allam et al. (2004) identified a set of merging galaxies in the SDSS data and found only a weak positive correlation in color for the merging pairs. Hernandez-Toledo et al. (2005) analyzed the $BVRI$ images of 42 elliptical/lenticular galaxies, and claimed that the structural effects of interactions on E/S0s are minor, in contrast to disk galaxies involved in interactions. de Propris et al. (2005) analyzed galaxies in the Millennium Galaxy Catalog and found merging galaxies are only marginally bluer than noninteracting galaxies, showing an excess of both early and late types but a deficiency of intermediate type spirals. Smith et al. (2007) did not find any enhancement in Spitzer mid-infrared color depending on pair separation by using pre-merger interacting galaxy pairs selected from the Arp Atlas.  In addition, in terms of the degree and scale of the effects of interactions on galaxy properties, the results also have not always agreed. Lambas et al. (2003) studied the galaxy pairs in the Two Degree Field Galaxy Redshift Survey (2dFGRS) data and reported that SF in galaxy pairs is significantly enhanced over that of isolated galaxies only when the projected separation $r_p<25 h^{-1}$ kpc and radial velocity difference $\\Delta v<100$ km s$^{-1}$. Nikolic et al. (2004) reported using the SDSS data that the mean SF rate is significantly enhanced for galaxy pairs with $r_p < 30$ kpc. But for late types, the enhancement extended out to 300 kpc. They also noted that the SF rate slightly decreased with increasing $\\Delta v$, and the light concentration was lowest at $r_p=75$kpc and then increased rapidly inward. Alonso et al. (2006) measured SF rate of the galaxies in the 2dFGRS and SDSS data and found that the SF rate was strongly enhanced when $r_p <100 h^{-1}$ kpc and $\\Delta V<350$ km s$^{-1}$, which was more effective in low and intermediate density environments. These discrepancies as listed above are expected mainly because the effects of interactions/mergers between different types of galaxies on their final products are different. Woods \\& Geller (2007) found that for blue-blue major pair sample in the SDSS, there exists clear correlation between specific SFR and pair separation, and for red-red pair, there is none (cf. Tran et al. 2005; van Dokkum 2005; Bell et al. 2006). Li et al. (2008) found that for the most strongly star-forming systems, tidal interactions are the dominant trigger of enhanced star formation and the enhancement is a strong function of separation less than 100 kpc.  In addition to such small-scale environment, the large-scale environment has been known to be one of the determining factors of galaxy properties. It is now well-known that galaxy properties correlate with the large-scale background density at low redshift (Hogg et al. 2003; Goto et al. 2003; Balogh et al. 2004a,b; Tanaka et al. 2004; Kauffmann et al. 2004; Blanton et al. 2005a; Croton et al. 2005; Weinmann et al. 2006; Park et al. 2007). Deep redshift surveys have extended these studies to high redshift (Cucciati et al. 2006; Elbaz et al. 2007; Cooper et al. 2007, 2008; Poggianti et al. 2008). A number of papers (Kauffmann et al. 2004; Blanton et al. 2005a; Quintero et al. 2006; Ball, Loveday, \\& Brunner 2008) claimed from an analysis of the SDSS data that structural properties of galaxies are less closely related to the (large-scale) environment than are their masses and SF related parameters such as color and SF rate. However, the structural parameters such as concentration and Sersic index are not true measure of morphology (Bamford et al. 2008) and indeed, van der Wel (2008) pointed that morphology and structure are intrinsically different galaxy properties and structure mainly depends on galaxy mass whereas morphology mainly depends on environment.  Several papers tried to use sophisticated automated morphology classifications for the study of the relationship between morphology and environment (Goto et al. 2003; Park \\& Choi 2005; Allen et al. 2006; Ball et al. 2008). Many studies reported that the SF rate of galaxies is a strongly decreasing function of the large-scale density and that there is a critical density for SF activity (Gomez et al. 2003; Tanaka et al. 2004).  Park et al. (2007), however, found that this trend was mostly because morphology and luminosity are strong functions of the large-scale density. They made an extensive study of the environmental dependence of various physical properties of galaxies on large- and small-scale densities, and concluded that morphology and luminosity are major fundamental parameters. Once morphology and luminosity are fixed, other galaxy properties are almost independent of the large-scale density. The large-scale density dependence of these properties reported by many previous and current studies merely reflects the correlations of the properties with morphology and luminosity rather than independent correlations with environment. Park et al. (2007) also found that galaxy morphology changes sensitively across the nearest neighbor distance of a few hundred kpc. This work has been extended by Park, Gott, \\& Choi (2008) who studied galaxy morphology as a function of the nearest neighbor separation at fixed large-scale background density and found that this characteristic scale corresponds to the virial radius of the neighbor galaxy. Park et al. (2008)'s findings and claims can be summarized as follows.  1. The effects of galaxy interactions reach farther than the distance previously thought. They reach at least out to the galaxy virial radius, namely a few hundred kpc for bright galaxies.  2. The result of galaxy interactions can be very different depending on the morphological type of the nearest neighbor galaxy when the pair separation is less than the virial radius. Without separating the neighbor galaxies into different morphological types, one will find the effects of galaxy interactions are diverse but negligible on average. The dependence on neighbor's morphology disappears at separations farther than the virial radius.  3. The fact that, at fixed large-scale density, the morphology of a galaxy is more likely to be a late type as it approaches a late-type neighbor, suggests that galaxies can transform their morphology from early types to late types through close encounters with cold gas-rich neighbors.  4. In most places of the universe, except for the regions within massive clusters of galaxies, the well-known morphology-density relation originated largely due to the effects of galaxy-galaxy interactions. Galaxy morphology appears to depend on the large-scale density mainly because the mean separation between galaxies is statistically correlated with the large-scale density.  5. A series of close interactions and mergers transform galaxy morphology and luminosity classes in such a way that galaxies on average evolve to become bright early types. The transformation speed depends on the large-scale density.  The fourth result is supplemented by the recent work of Park \\& Hwang (2008) who found, in the special case of galaxies located within the virial radius of massive clusters, galaxy properties depend on both the distance to the nearest neighbor and  clustercentric radius.  The purpose of this paper is to extend Park et al. (2008)'s work by looking at the dependence of various other properties of galaxies, as well as morphology and luminosity, on small and large scale environmental factors. We use a volume-limited sample of the SDSS galaxies whose morphology is accurately classified. The environment is specified by the small-scale (distance to the nearest neighbor galaxy and morphology of the nearest neighbor galaxy) and large-scale (the background density) factors. It is hoped that the effects of galaxy interactions can be understood in fuller detail in this three-dimensional environmental parameter space. ",
        "conclusions": "We have inspected the dependence of eight physical parameters of galaxies on the small-scale (the nearest neighbor distance, the nearest neighbor's morphology) and the large-scale (background density smoothed over 20 nearest galaxies with $M_r < -19.0$) environments. We have also studied the kinematic properties of the galaxies in pairs. We found the impact of interaction on galaxy properties are detectable at least out to the pair separation corresponding to the virial radius of (the neighbor) galaxies in our sample, which is mostly between 200 and 400 $h^{-1}$ kpc. It was crucial to divide galaxy interactions into four cases depending on morphology of target and neighbor galaxies in order to detect these long-range interaction effects.  Our major results are as follows.  1. There are two characteristic pair-separation scales where the breaks in the dependence of galaxy properties on $r_p$ are observed. The first scale is the virial radius of the nearest neighbor galaxy $r_{\\rm vir,nei}$. All parameters studied, except for $R_{\\rm Pet}$, start to deviate from those of extremely isolated galaxies at $r_{p} \\sim r_{\\rm vir,nei}$ in the case of late-type galaxies, in particular. The second scale is the merger scale which is about $0.05 r_{\\rm vir,nei}$. This corresponds to $10 - 20h^{-1}$ kpc for the galaxies in our sample.  2. The SF activity of galaxies is enhanced when the nearest neighbor is a late type, but reduced when the neighbor is an early type. These effects occur within the virial radius of the neighbor galaxy. These are strong evidence for hydrodynamic interactions within the virial radius of the galaxy plus dark halo system during encounters.  3. The dependence of galaxy properties on $\\rho_{20}$ is strong only for luminosity (Fig. 5) and morphology (Fig. 3). All other parameters show weak or negligible dependence on $\\rho_{20}$ once luminosity and morphology are fixed. We have inspected the small subtle dependence on $\\rho_{20}$ whenever detectable. For example, the $u-r$ color of late-type galaxies has a weak residual dependence on $\\rho_{20}$ at fixed $r_p$. We suggest that galaxies in higher density environment maintain hotter and denser halo gas due to some internal heating mechanisms and external confining material, which can be the reason for the large-scale density dependence of morphology and SF activity parameters.  4. At fixed large-scale background density, galaxies with larger pair separations have higher luminosity. Such dependence exists mainly at the nearest neighbor distance larger than the virial radius of the neighbor. This is interpreted as evidence for the on-going process of luminosity transformation through mergers.     In a forthcoming paper (Park \\& Hwang 2008) we will examine the dependence of the SDSS galaxies associated with the Abell clusters on the nearest neighbor separation and the clustercentric radius. The latter is the large-scale environmental parameter replacing $\\rho_{20}$ here. This work will extend the present work to the extreme situation where the large-scale background density itself exceeds the virialized density. We are also studying the effects of the nearest neighbor on galaxy properties using higher redshift samples like the GOODS and DEEP2 samples (Hwang \\& Park 2008) to understand galaxy evolution due to galaxy-galaxy interactions in the high redshift universe."
    },
    "0809/0809.0479_arXiv.txt": {
        "abstract": " ",
        "introduction": "On 4 July 2005 UT, comet 9P/Tempel 1 crashed into the Deep Impact spacecraft. The spacecraft was located in the path of the comet in order for the impact to create a large crater, releasing material from the inside of the nucleus, to enable study of interior ices. In order to understand the changes which were observed, it was necessary to have observations of the comet prior to the impact.  In this paper, we detail low-resolution spectroscopic observations obtained of comet 9P/Tempel 1 in its 1983, 1989, 1994 and 2005 apparitions using spectrographs on the 2.7m Harlan J. Smith telescope of McDonald Observatory. The observations of 2005 include some made on the night of impact plus the two nights following the impact.  Comet 9P/Tempel 1 is a Jupiter family comet with a 5.5 year period.  Its orbit is such that it has alternating good and poor viewing apparitions. The apparitions 11 years apart have very similar viewing geometry. Over the span of our observations, the perihelion date changed slightly and the perihelion distance moved outwards from 1.49 to 1.51\\,{\\sc au} (see Table~\\ref{orbit}).  All of our data were obtained pre-perihelion with the exception of the data from 5 and 6 July 2005. \\begin{table}[h!] \\caption{Orbital Parameters}\\label{orbit} \\vspace*{10pt} \\centering \\begin{tabular}{ccccc} \\hline Year & Perihelion & Perihelion & $e$ & $i$ \\\\ & Date & Distance ({\\sc {\\sc au}}) & & (degrees) \\\\ \\hline 1983 & Jul. 9.797 & 1.491117 & 0.520898 & 10.5571 \\\\ 1989 & Jan. 4.443 & 1.496725 & 0.519684 & 10.5462 \\\\ 1994 & Jul. 3.314 & 1.494151 & 0.520255 & 10.5518 \\\\ 2000 & Jan. 2.617 & 1.500047 & 0.518953 & 10.5413 \\\\ 2005 & Jul. 5.348 & 1.506434 & 0.517490 & 10.5303 \\\\ \\hline \\multicolumn{5}{c}{Data from NASA/JPL's Horizons web interface} \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "The goal of the Deep Impact spacecraft mission was to study the interior of a comet by excavating material from deep within the nucleus. Groussin {\\it et al.} (2007) analyzed IR spectra obtained from the flyby spacecraft to determine that the thermal inertia was low, ``most probably $<50$ W K$^{-1}$ m$^{-2}$ s$^{1/2}$\". This low thermal inertia implies that the sublimation of volatiles generally occurs close to the surface.  The exact depth of the crater produced by the impact is unknown, but probably included ``deep\" materials (Schultz {\\it et al.} 2007). \\nocite{groussinetal2007, schultzetal2007}  This unique event was monitored by observatories world-wide as a complement to the data which were obtained by the spacecraft (Meech {\\it et al.} 2005). \\nocite{medeepimpact} In order to understand the event, it was necessary to place the impact monitoring observations into the context of the usual behavior of this comet.  Our observations, reported in this paper, allow not only for an understanding of the state of the comet at the time of the impact, but also allow for an understanding of the natural changes we see in this comet.  Comet Tempel 1 rotates relatively slowly, with a rotation period of 1.701$\\pm$0.014 days (A'Hearn {\\it et al.} 2005). \\nocite{ahdeepimpact} Thus, our observations on a single night and a single slit orientation were never more than 5\\% of the rotation period.  Thus, even if there had been a strong active region, it would not change position much relative to our slit during a single night.  Major jets show up as moving ``bumps\" in the gas distribution as the material flows outward (Cochran and Trout 1994). \\nocite{cotr94} Observations on subsequent nights might have viewed different regions, but we do not see any differences from night-to-night in our data.  In addition to the smoothly varying rotational modulation, the comet would undergo sporadic outbursts (A'Hearn {\\it et al.} 2005; Meech {\\it et al.} 2005; Farnham {\\it et al.} 2007). \\nocite{fadeepimpact}  However, according to Table~1 of Farnham {\\it et al.}, there were no outbursts on our dates of observations.  Thus, the outbursts would have little affect on our observations.  Therefore, we can use our observations to comment on the properties of the gas production of comet Tempel 1 during the season leading up to the impact and also over the preceding two decades. Overall, comet Tempel 1 is a moderate producer of gas. IR observations with the Deep Impact spacecraft found that water ice ``is restricted to three discrete and relatively small areas\" (Sunshine {\\it et al.} 2007) (this is not altogether consistent with the conclusion of Ferrin (2007) that Tempel 1 is a young comet). \\nocite{sunshine2007,fe2007} Our values for the production rates of the various optically observed emissions are in good agreement with Schleicher (2007), with small differences due to differences in fluorescence efficiencies, scale lengths and outflow velocities used in the modeling. Lara {\\it et al.} (2006) also observed comet Tempel 1 during 2005. Their derived production rates are lower than ours and Schleicher's by a factor of a few for C$_{2}$, C$_{3}$, and CN, the species they observed.  We cannot explain this difference. \\nocite{laetal06}  Within the error bars, the comet produced the same {\\it relative} amounts of all of the species that are in our bandpass for the three well observed apparitions, even while the overall amount of gas was decreasing. We see the same abundance ratios in the spectra from 4 July 2005, the night of impact. The impact was an inherently non-steady state event while the Haser model assumes steady state production.  Therefore, the production rates for July 2005 must be used cautiously.  On 4 July, the observations were obtained with the first 40 minutes after the impact so ejecta from the impact would only affect our inner few pixels (at 0.5 km sec$^{-1}$ ouitflow, gas from the crater will flow outwards 1200km in 40 minutes, or less than 1 pixel).  The outer part of the slit would still have the ambient signal and this is what we would be measuring predominantly.  By 24 hours later, the ejected material would have traveled across our slit.  The increase in production resulted from freshly exposed surface (Jackson {\\it et al.} 2008, in preparation).  Inspection of the gas distribution shows that within our slit, the gas distribution has the same general shape on 4 through 6 July as on other nights (Fig~\\ref{distrib}). Rauer {\\it et al.} (2006) found a ''bump\" of material in observations from 4 and 5 July.  However, along the Sun-comet line (their Fig. 4), their profiles look the same for the two nights. Within the scatter, ours do too, as seen in Fig. 4.  We did not have a similarly oriented slit angle on 6 July. \\nocite{raueretal2006} Thus, while a Haser model production rate does not account for the impulsive nature of the event, it allows for comparison with previous data. The freshly released material has the same relative composition as the ambient material. As long as the impact excavated material deep enough that it reached fresh, unprocessed material, the implication of our findings is that the comet does not differentially lose one kind of ice versus another as it passes through the inner Solar System repeatedly.  This also means that prior studies of cometary composition via observations of the normal coma emissions are probing the original composition of these primitive bodies. \\begin{figure} \\vspace{4in} \\special{psfile=\"distrib-r.ps\" vscale=55 hscale=55 voffset=-30 hoffset=20} \\caption[fig3]{The column densities from 4 and 5 July 2005 UT are compared, with the data from 4 July scaled upwards to match the level of the data from 5 July.  Note that the distribution of the gas with cometocentric distance is not the same on the east (negative direction) and west sides of the comet.  The slit position angle was 20 degrees off of the extended heliocentric radius vector so the slit is essentially along the Sun-comet line.  Within the scatter of the 4 July data, there are no difference in the profiles for 4 and 5 July on either side of the optocenter. }\\label{distrib} \\end{figure}   This rich data set, which spans more than two decades and four apparitions allows us to look at secular changes in the comet. We find that the total gas abundances declined from 1983 to 2005, with decreases of factors of 1.3--2.7, depending on the species. This is in good agreement with Schleicher (2007).  Except in cases where a comet was totally disrupted (e.g. 1999 S4 (LINEAR)) or its orbit changed significantly (e.g. 81P/Wild 2), such large changes of gas production have not been seen before.  The simplest explanation is that some region which was active on the surface of the comet in 1983 was not well illuminated in 2005.  However, A'Hearn {\\it et al.} (2005) found that the obliquity of the comet was only 11$^\\circ$ and the viewing geometry was very similar in 1983, 1994 and 2005.  Thus, the only way the illumination of a particular active region could change would be if that active region was very near a pole.  This cannot be confirmed with the flyby data since the encounter period was very short.  Such a polar jet has been seen in comet Borrelly (Soderblom {\\it et al.} 2002). \\nocite{soderetal:ds1}  Other mechanisms, such as the ``sealing\" of an active region from the fall-back of dust after perihelion, cannot be ruled out.  We found that the gas production was not symmetric around perihelion but peaked almost two months prior to perihelion. Because we observed only at monthly intervals, we cannot comment on whether different species peaked at different times, something that was noted by Schleicher (2007). Such asymmetries are common in cometary activity and generally indicate that small regions of the surface are active and see changing illumination in the course of their orbit. Indeed, as noted above, the Deep Impact spacecraft observations did not observe distinct icy regions as sources of the activity.  Overall, we find that comet Tempel 1 has abundance patterns which are similar to the vast majority of comets and can be classified as a ``normal\" comet.  These abundance patterns do not change with depth, as evidenced from the post-impact spectra. Therefore, as this comet continues to evolve, it is unlikely to change its abundances patterns."
    },
    "0809/0809.4636_arXiv.txt": {
        "abstract": "{} {Systematic surveys to search for exoplanets have been mostly dedicated to solar-type stars sofar. We developed in 2004 a method to extend such searches to earlier A--F type dwarfs and started spectroscopic surveys to search for planets and quantify the detection limit achievable when taking into account the stars properties (Spectral Type, \\vsini) and their actual levels of intrinsic variations. We give here the first results of our southern survey with {\\small HARPS}. } {We observed 185 A--F (\\bv~in the range [$-$0.1; 0.6]) stars with {\\small HARPS} and analysed them with our dedicated software. We use several criteria to probe different origins for the radial-velocity variations -- stellar activity (spots, pulsations) or companions: bisector shape, radial-velocity variations amplitudes and timescales. } {1) 64\\,$\\%$ of the 170 stars with enough data points are found to be variable. 20 are found to be binaries or candidate binaries (with stars or brown dwarfs). More than 80\\,$\\%$ or the latest type stars (once binaries are removed) are intrinsically variable at a 2\\,\\ms precision level. Stars with earlier spectral type (\\bv~$\\leq$ 0.2) are either variable or associated to levels of uncertainties comparable to the RV rms observed on variable stars of same \\bv. 2) We have detected one long-period planetary system (presented in another paper) around an F6IV--V star. 3) We have quantified the jitter due to stellar activity and we show that taking into account this jitter in addition to the stellar parameters (spectral type, \\vsini), it is still possible to detect planets with {\\small HARPS} with periods of 3 days (resp. 10 days and 100 days) on 91\\,$\\%$ (resp. 83\\,$\\%$, 61\\,$\\%$) of them. We show that even the earliest spectral type stars are accessible to this type of search! , provided they have a low projected rotational velocity and low levels of activity. 4) Taking into account the present data, we compute the actually achieved detection limits for 107 targets and discuss the limits as a function of \\bv. Given the data at hand, our survey is sensitive to short-period (few days) planets and to longer ones (100 days) at a lower extent (latest type stars). We derive first constrains on the presence of planets around A--F stars for these ranges of periods. } {} ",
        "introduction": "Since the discovery of the first exoplanet around a solar-like star in 1995, more than 250 planets have been found by radial-velocity (RV) surveys (Jean Schneider, http://exoplanet.eu). These surveys have generally focused on late-type stars (later than F8). However, knowing about the presence of planets or brown dwarfs (hereafter BDs) around more massive objects is mandatory if one wishes to investigate the impact of the mass of the central stars on the planetary formation and evolution processes.  There are theoretical indications that the mass of the planets increases with the mass of the parent star, at least for low mass stars (\\cite{ida05}) and that the frequency of giant planets increases linearly with the parent star mass for stars between 0.4 and 3\\,\\Msun~(\\cite{kennedy08}), with \\eg 6\\,$\\%$ frequency of giant planets around 1\\,\\Msun~and 10\\,$\\%$ frequency around 1.5\\,\\Msun. More numerous and massive planets are consistent with what we could expect from a disk surface density increasing with stellar mass. On the other hand the shorter lifetimes of the systems as well as the lack of solid material close to the star could reduce the number of planets. Clearly, several parameters probably impact the occurence and properties of planets around massive stars, and they have not yet been fully explored.  The data to test the models are still quite limited as the largest and earliest, now long-lasting surveys had focused on solar type, main-sequence (MS) stars. In recent years, some efforts have been made nevertheless to search for planets around stars with various masses: less massive, M-type stars on the one hand, and more massive stars on the other hand. The search for planets around M stars seems sofar to confirm the previously mentionned expectations from theoretical works (see \\eg \\cite{bonfils05}; \\cite{butler06}). As far as massive stars are concerned, the available observations are still very limited. Massive MS stars have been removed from early surveys as it was generally thought that their spectra (few lines, usually broadened by stellar rotation) would not allow planet detection and indeed, the classical RV measurements technics (based on the cross correlation of the actual spectra with a binary spectral mask corresponding to a star with an appropriate spectral type and \\vsini~= 0\\,\\kms) fail to measure the RV of these stars. This lead some groups to study instead ``retired'' early-type, either low-mass ($\\leq$ 1.6\\,\\Msun) giants, intermediate-mass (1.6--2\\,\\Msun) subgiants or clump giants (1.7--3.9\\,\\Msun) (see \\eg {\\bf \\cite{hatzes05}}; {\\bf \\cite{niedzielski07}}; {\\bf \\cite{johnson06}}; {\\bf \\cite{johnson07}}; \\cite{lovis07}; \\cite{sato08}). These stars indeed have cooled out and also rotate more slowly due to coupling of stellar winds and magnetic fields; they therefore exhibit more numerous, narrower lines, which is adequate for classical RV measurements technics, and their level of activity (jitter) is relatively low (10--20 m\\,s$^{\\rm -1}$; \\cite{hekker06}) for giants and 10\\,\\ms~for subgiants (\\cite{johnson07} and ref. therein; \\cite{sato08}). The data available today are still limited compared to those available on solar-type stars and less than 20 planets have been found sofar in total around these evolved stars. So far, the planets found around K subgiants stars with M $\\geq$ 1.5\\,\\Msun~are located at distances larger than {\\bf 0.8}\\,AU (\\cite{johnson07}); this led these authors to conclude that close-in planets are rare, in agreement with some theoretical predictions on disks depletion timescales (\\cite{burkert07}). However, the impact of the post MS evolution of the stars on closer-in planets has not been explored yet for these stars. Concerning giant stars, all planets found so far have relatively long periods, the closest ones beeing reported at less than 0.7\\,AU from 2 giants, in addition to a previously reported planet at 0.7\\,AU by \\cite{sato03}. Numerical simulations by the same authors suggest that planets with orbits inside 0.5--1\\,AU around 2--3\\,\\Msun~stars could be engulfed by the central stars at the tip of RGB due to tidal torque from the central stars. According to them, if one assumes then that most of the clump giants are post RGB stars, there is then a risk that closer planets, if present before had disappeared since the star evolved. In summary, eventhough there are some hints that hot Jupiters are not present around retired stars, it is recognized that data are still needed to definitely confirm this point. Also, as stellar evolutionary processes may have impacted the presence of planets close to the stars, it is acknowledged that data are needed on A--F MS stars (see eg \\cite{burkert07}; see also \\cite{li08}). We note finally that short-period planets have indeed been found around F5--F6 MS stars through transit observations (see http://exoplanet.eu).  We developed a few years ago a dedicated software to extract the RV data around early type MS stars. The method consists in correlating, in the Fourier space, each spectrum and a reference spectrum built by summing-up all the available spectra for this star. We had shown earlier that with this approach, and taking into account the stars \\bv~and projected rotational velocities, it is possible to find planets around A--F type stars (\\cite{galland05a}). However, the price to pay is that more measurements are needed to find planets around A and early F dwarfs than for late F and G--K dwarfs, because of the relatively higher uncertainties in the RV measurements due to higher \\vsini~and higher effective temperature, and also because of the possible presence of pulsations or spots in the case of late F stars, the impact of which has not been quantified so far. (Note that spots or pulsations become also a limiting factor in the case of later-type stars, as well as pulsations, if one looks for low mass planets.) We started then systematic searches for low-mass companions to A--F type stars, with {\\small HARPS} in  the southern hemisphere and with {\\small ELODIE} and then {\\small SOPHIE} at OHP in the northern hemisphere. We found so far a 9.1\\,\\Mjup~(minimum mass) planet orbiting ($a$ = 1.1\\,AU) an F6V star, with \\vsini~= 12\\,\\kms~(\\cite{galland05b}). Very promisingly, we also detected a 21\\,\\Mjup~brown dwarf orbiting ($a$ = 0.2\\,AU) a pulsating A9V star with \\vsini~= 50\\,\\kms~(\\cite{galland06}); noticeably, in that case, we could disentangle stellar pulsations from the presence of a low-mass companion.  In parallel, we developed detailed simulations of stellar activity (spots) in order to estimate more quantitatively than what was available so far (\\cite{saar97}; \\cite{hatzes02}) the impact of such stellar activity on RV data and other observables (bisectors, bisectors velocity-span, photometry). We have showed that if the star \\vsini~is smaller than the spectrograph spectral resolution, depending on their location with respect to line of sight, depending also on their size, spots with realistic sizes can produce RV variations and bisector velocity-span variations quite similar to those of low mass planets. Hence, low amplitude (level of typically 20\\,\\ms~or less) planetary-like RV and bisector velocity span variations cannot alone definitely prove the presence of planets around low \\vsini~G--K stars (\\cite{desort07}) and additional criteria are mandatory to rule out spots: photometry, activity evaluation down to levels relevent to explain the amplitude of RV variations, precise knowledge of the star rotaional period, etc. The situation in that respect is much more favorable in the case of earlier-type stars as they rotate statistically faster and hence the bisector criteria can apply.  The present paper is devoted to our southern hemisphere survey. The sample, observations, measurements and diagnostics are provided in Sect.~\\ref{observ}. The results concerning the stellar variability as well as the quantitative impact on planet detectability around the early-type stars are presented in Sect.~\\ref{results}. Finally we give and discuss in Sect.~\\ref{detectionlim} the detection limits obtained in the present survey. ",
        "conclusions": "Based on the observation of a large number of A--F type stars (170), we have been able to measure their jitters, and derive for the first time estimations of the detection limits that can be expected on average on those stars with \\bv~in that range [$-$0.1; 0.6] (once previously known $\\delta$\\,Scuti and $\\gamma$\\,Doradus stars are removed) for 3 periods: 3, 10 and 100-day periods. We have shown that at the precision provided by {\\small HARPS}, most of the stars are variable in RV, and the impact of the RV jitter, due to either spots or pulsations is generally non negligeable on planet detectability. However, assuming that planets are detectable if the amplitude of the induced RV variation is larger than 3$\\times {\\rm rms}$ (this threshold defines the achievable detection limits, which depends on the star and the spectrograph used), we have shown that even when taking into account the jitter, giant planets can still be found around these stars in most cases. This is not only true for the stars with \\bv~$\\geq$ 0.3 for which we can find either short or long period planets, with masses as small as 0.02\\,\\Mjup~(case of short period) for the latest type stars but also for dwarfs with \\bv~$\\leq$ 0.3: for such stars short-period planets can still be found around those with relatively low projected rotational-velocity and low level of activity, with masses down to 0.5\\,\\Mjup~(best case). This survey has allowed to identify for the first time the stars that are best suited for further searches for planets around massive dwarfs.  We have also shown that given the data available, the present survey is sensitive to short-period planets (hot Jupiters) and only partially sensitive to larger periods (up to 100\\,days). We found in particular one 2-planet system with periods larger than 100\\,days around one late-type star. Whenever possible (107 stars), we computed for each star the detection limits actually achieved and showed that when enough data are available, the achieved detection limit is set by the $3 \\times {\\rm rms}$ threshold. We indeed reached such limits for early-type as well as for late-type stars. We finally derive first estimates of the presence of short-period planets around these A--F stars. We showed for instance that less than 5\\,$\\%$ of the latest-type stars (\\bv~$\\geq$ 0.4) host P=3 days-period planets with masses 1\\,\\Mjup~or more. Such statistics are not constraining enough to allow interesting comparisons with later type stars or with model predictions; but as soon as more data become available, the statistics can be straightforwardy improved.  Finally, we note that to compute these detection limits, we did not try to average out the spectra over timescales associated to the frequencies of intrinsic stellar variations. This approach would of course allow to decrease significantly the detectable masses. As it would require lots of telescope time, it should be probably kept for stars with the highest scientific interest."
    },
    "0809/0809.1013_arXiv.txt": {
        "abstract": "The Casimir densities are investigated for a massive spinor field in de Sitter spacetime with an arbitrary number of toroidally compactified spatial dimensions. The vacuum expectation value of the energy-momentum tensor is presented in the form of the sum of corresponding quantity in the uncompactified de Sitter spacetime and the part induced by the non-trivial topology. The latter is finite and the renormalization is needed for the first part only. The asymptotic behavior of the topological term is investigated in the early and late stages of the cosmological expansion. When the comoving lengths of the compactified dimensions are much smaller than the de Sitter curvature radius, to the leading order the topological part coincides with the corresponding quantity for a massless fermionic field and is conformally related to the corresponding flat spacetime result with the same topology. In this limit the topological term dominates the uncompactified de Sitter part and the back-reaction effects should be taken into account. In the opposite limit, for a massive field the asymptotic behavior of the topological part is damping oscillatory.\\\\  \\bigskip  \\noindent PACS numbers: 04.62.+v, 04.20.Gz, 04.50.-h, 11.10.Kk ",
        "introduction": "Many of high energy theories of fundamental physics are formulated in higher dimensional spacetimes and it is commonly assumed that the extra dimensions are compactified. In particular, the idea of compactified dimensions has been extensively used in supergravity and superstring theories. From an inflationary point of view universes with compact spatial dimensions, under certain conditions, should be considered a rule rather than an exception \\cite{Lind04}. These models may play an important role by providing proper initial conditions for inflation (for physical motivations of considering compact universes see also \\cite{Star98}). As it was argued in \\cite{McIn04}, there are many reasons to expect that in string theory the most natural topology for the universe is that of a flat compact three-manifold. The quantum creation of the universe having toroidal spatial topology is discussed in \\cite{Zeld84} and in references \\cite{Gonc85} within the framework of various supergravity theories. In addition to the theoretical work, there has been a large activity to search for signatures of non-trivial topology by identifying ghost images of galaxies, clusters or quasars. Recent progress in observations of the cosmic microwave background provides an alternative way to observe the topology of the universe \\cite{Levi02}. If the scale of periodicity is close to the particle horizon scale then the changed appearance of the microwave background sky pattern offers a sensitive probe of the topology.  The compactification of spatial dimensions leads to a number of interesting quantum effects which include instabilities in interacting field theories \\cite{Ford80a}, topological mass generation \\cite{Ford79,Toms80a,Toms80b} and symmetry breaking \\cite{Toms80b,Odin88}. In the case of non-trivial topology the boundary conditions imposed on fields give rise to the modification of the spectrum for vacuum fluctuations and, as a result, to the Casimir-type contributions in the vacuum expectation values of physical observables (for the topological Casimir effect and its role in cosmology see \\cite{Most97,Bord01}). A characteristic feature of models with compactified dimensions is the presence of moduli fields which parametrize the size and the shape of the extra dimensions and the Casimir effect has been used for the generation of the effective potential for these moduli. The Casimir energy can also serve as a model for dark energy needed for the explanation of the present accelerated expansion of the universe (see \\cite% {Milt03,Eliz06,Gree07} and references therein).  Quantum field theory in de Sitter (dS) spacetime has been extensively studied during the past two decades. Much of early interest was motivated by the questions related to the quantization of fields on curved backgrounds. dS spacetime has a high degree of symmetry and numerous physical problems are exactly solvable on this background. The importance of this theoretical work increased by the appearance of \\ the inflationary cosmology scenario \\cite{Lind90}. In most inflationary models an approximately dS spacetime is employed to solve a number of problems in standard cosmology. During an inflationary epoch quantum fluctuations in the inflaton field introduce inhomogeneities and may affect the transition toward the true vacuum. These fluctuations play a central role in the generation of cosmic structures from inflation. More recently astronomical observations of high redshift supernovae, galaxy clusters and cosmic microwave background \\cite{Ries07} indicate that at the present epoch the Universe is accelerating and can be well approximated by a world with a positive cosmological constant. If the Universe would accelerate indefinitely, the standard cosmology would lead to an asymptotic dS universe. Hence, the investigation of physical effects in dS spacetime is important for understanding both the early Universe and its future.  One-loop quantum effects for various spin fields on the background of dS spacetime have been discussed by several authors (see, for instance, \\cite% {Cher68}-\\cite{Birr82} and references therein). The effects of the toroidal compactification of spatial dimensions in dS spacetime on the properties of quantum vacuum for a scalar field with general curvature coupling parameter are investigated in \\cite{Saha07,Bell08,Saha08b} (the quantum effects in braneworld models with dS spaces and in higher-dimensional brane models with compact internal spaces were discussed in \\cite{dSbrane,Flac03}). In the present talk, based on \\cite{Saha08,Beze08} we present the results on the Casimir effect for a fermionic field on background of dS spacetime with spatial topology $\\mathrm{R}^{p}\\times (\\mathrm{S}^{1})^{q}$.  The paper is organized as follows. In the next section the plane wave eigenspinors in $(D+1)$-dimensional dS spacetime with an arbitrary number of toroidally compactified dimensions are constructed. In section \\ref{sec:EMT} these eigenspinors are used for the evaluation of the vacuum expectation value of the energy-momentum tensor in both cases of the fields with periodicity and antiperiodicity conditions along compactified dimensions. The behavior of these quantities is investigated in asymptotic regions of the parameters. The special case of topology $\\mathrm{R}^{D-1}\\times \\mathrm{% S}^{1}$ with the corresponding numerical results is discussed in section \\ref% {sec:Special}. The main results are summarized in section \\ref{sec:Conc}. ",
        "conclusions": "\\label{sec:Conc}  We have investigated the Casimir densities for a massive spinor field in $(D+1)$-dimensional dS spacetime with an arbitrary number of toroidally compactified spatial dimensions assuming that the field is prepared in the Bunch-Davies vacuum state. For the evaluation of the vacuum expectation value of the energy-momentum tensor the mode-summation procedure is employed. In this approach we need the corresponding eigenspinors satisfying appropriate boundary conditions along the compactified dimensions. These eigenspinors are constructed in section \\ref% {sec:EigFunc} for fields with both periodicity and antiperiodicity conditions. By using the eigenspinors and applying to the mode-sum the Abel-Plana formula, the VEV for the spatial topology $\\mathrm{R}^{p}\\times (% \\mathrm{S}^{1})^{q}$ is presented in the form of the sum of the corresponding quantity in the topology $\\mathrm{R}^{p+1}\\times (\\mathrm{S}% ^{1})^{q-1}$ and of the part which is induced by the compactness of $(p+1)$% th dimension. For the field with periodicity conditions the topological part is given by formula (\\ref{TopTll}). The corresponding formula for the field with antiperiodicity conditions is obtained from that for the field with periodicity conditions inserting the factor $(-1)^{n}$ in the summation over $n$ and replacing the definition for $k_{\\mathbf{n}_{q-1}}^{2}$ by (\\ref% {knq-1Tw}).  The topological part is finite and the renormalization procedure for the VEV\\ of the energy-momentum tensor is reduced to that for the uncompactified dS spacetime. This part is time-dependent and breaks the dS symmetry. The corresponding vacuum stresses along the uncompactified dimensions coincide with the energy density and, hence, in the uncompactified subspace the equation of state for the topological part of the energy-momentum tensor is of the cosmological constant type. For a massless fermionic field the problem under consideration is conformally related to the corresponding problem in the Minkowski spacetime with spatial topology $\\mathrm{R}% ^{p}\\times (\\mathrm{S}^{1})^{q}$ and we have the standard relation $\\langle T_{k}^{l}\\rangle _{\\mathrm{c}}=a^{-(D+1)}(\\eta )\\langle T_{k}^{l}\\rangle _{\\mathrm{c}}^{\\mathrm{(M)}}$ between the topological terms.  For a massive fermionic field, in the limit when the comoving length of a compactified dimension is much smaller than the dS curvature radius, the topological part in the VEV of the energy-momentum tensor coincides with the corresponding quantity for a massless field and is conformally related to the VEV in toroidally compactified Minkowski spacetime. In particular, the topological part in the vacuum energy density is positive for an untwisted fermionic field. This limit corresponds to the early stages of the cosmological evolution and the topological part dominates over the uncompactified dS part. At these stages the back-reaction effects of the topological term are important and these effects can essentially change the dynamics of the model. In the opposite limit, when the comoving lengths of the compactified dimensions are large with respect to the dS curvature radius, in the case of a massive field the asymptotic behavior of the topological part is damping oscillatory and to the leading order is given by formula (\\ref{TllSmall}). These formula describes the behavior of the topological part in the late stages of the cosmological expansion. As the uncompactified dS part is time-independent, we have similar oscillations in the total VEV as well. Note that this type of oscillatory behavior is absent for a massless fermionic field."
    },
    "0809/0809.1680_arXiv.txt": {
        "abstract": "{ A systematic study of the nuclear emission of a sample of 97 spirals in isolated galaxy pairs with mixed morphology (E+S) shows that: 1) AGN activity is found in 40\\% of the spiral galaxies in these pairs, 2) Only one out of the 39 AGN found has type 1 (Broad line Component) activity, and 3) AGN tend to have closer companions than star forming galaxies. These results are at odds with a simple Unified Model for Seyferts, where only obscuration/orientation effects are of relevance, and neatly support an evolutionary scenario where interactions trigger nuclear activity, and obscuration/orientation effects may be complementary in a certain evolutionary phase. ",
        "introduction": "\\label{sec:intro}  One of the outstanding problems in the understanding of Active Galactic Nuclei (AGN) phenomenon is the role of the cicumgalactic environemnt in the triggering of the central engine. In an attempt to elusidate this question, in the past 10 years several efforts have focused in the study of the environment of AGNs, from a few kiloparsecs around the galactic nucleus to some hundreds of kiloparsecs around the host galaxy. Most of the investigations have dealt with samples of Seyfert galaxies, because these are the closest clearly non-thermal dominated active nuclei. LINERs are easy to observe, however, the nature of the dominating emission mechanism is not well established yet (Krongold et al 2003; Gonz{\\'a}lez-Mart{\\'{\\i}}n et al. 2006).  It has been suggested that Seyfert 2 galaxies are in interaction with the same frequency than star-forming galaxies (Krongold et al. 2002), while Seyfert 1 galaxies are in interaction less frequently, comparably as often as non-active galaxies (Dultzin-Hacyan et al. 1999, Krongold et al. 2001). The most recent studies confirm these findings considering physical companions, i. e. not only from companions selected frrom statitiscal considerations but from actual measurements of radial velocities for the neighboring galaxies (Koulouridis et al., 2006a 2006b).   In all the above analyses the environment of well defined samples of active vs. non-active galaxies were compared. In the present paper we adopt a complementary approach. We study the incidence of nuclear activity in a well defined sample of interacting galaxies. We focus on the sample of isolated mixed-morphology galaxy pairs studied by Hern\\'andez-Toledo et al. (1999, 2001), obtained from the Catalog of Isolated Pairs in the Northern Hemisphere (KPG; Karachentsev 1972).  These pairs are a unique laboratory to study the effect of tidal forces in triggering nuclear activity because they are relatively simple systems where a gas rich galaxy interacts with a gas poor one. In such systems a clean interpretation of the origin and evolution of the gaseous component is possible.  In this paper we present high resolution, long-slit, spectroscopy for the spiral component of 97 mixed morphology pairs from the {\\it Catalog of Isolated Pairs} (CPG; Karachentsev 1972). For details on the properties of the mixed pair sample see Hern\\'andez-Toledo et al. (1999, 2001), and Gonzalez et al. (2008, herafter G08). ",
        "conclusions": "\\label{sec:discussion}  The most striking result in our analysis is that only  1 out of the 39 AGN  can be classified as type 1. Even if in the present analysis we cannot separate LINERs from Seyferts, this result is very much at odds with the expectation of the so-called Unified Model (UM). A frequency of 2.6\\% for the presence of type 1 activity is too low to be explained with an obscuration\\/orientation effect alone. Ho et al. (1997) found that $\\sim$20\\% of their sample of nearby low-luminosity AGN (including LINERs and Seyferts) observed with HST presented broad lines (the expectation of the UM is that $\\sim$60\\% of Seyferts should be type 1).  Our results also show a clear connection between interaction and nuclear activity, in particular of type 2 activity. This fact follows naturally from the evolutionary model developed by Krongold et al. (2002, 2003). In this model, the onset of both nuclear activity and enhanced nuclear star formation is due to the tidal forces that induce the infall of large amounts of gas into the nucleus. These tidal effects are produced by an interaction with a nearby (similar mass) companion.  In the first stages of the interaction, a Starburst should dominate the emission. As more material falls into the innermost regions, the onset of non-thermal activity begins. At this stage, the non-thermal power is low, and there is full obscuration of the nuclear region (including the broad line clouds), therefore only a type 2 AGN is observed. During this stage, the model predicts a high probability for the detection of a nearby companion. As even more material falls, the accretion rate is increased, and partial obscuration is expected (as the initial spherical distribution of obscuring dust should flatten to form a torus), so both types of AGN can be observed. In the final stage, the dust is sweeped away by both the winds of massive evolved stars and the AGN itself. In this stage a naked type 1 AGN neatly emerges. A timescale of 10$^9$ yr is required to erase all evidence of possible past interactions (either through the completion of a merging event between two galaxies of similar mass, or through the dissolution of an unbounded interacting companion). This naturally explains why type 2 AGN are often found in interacting systems, while type 1 AGN are rarely found with companions. We notice that this evolutionary scenario does not exclude the effect of a possible obscuring/inclination combined effect, such as the one postulated in the UM. Our scenario includes the UM, but only at a given evolutionary stage.  The above evolutionary scenario explains in a straight forward way the results found in this work for the spiral component of the mixed morphology pairs. This implies that interactions play a key role in triggering nuclear activity."
    },
    "0809/0809.4546_arXiv.txt": {
        "abstract": "Today, the generation of magnetic fields in solar-type stars and its relation to activity and rotation can coherently be explained, although it is certainly not understood in its entirety.  Rotation facilitates the generation of magnetic flux that couples to the stellar wind, slowing down the star.  There are still many open questions, particularly at early phases (young age), and at very low mass.  It is vexing that rotational braking becomes inefficient at the threshold to fully convective interiors, although no threshold in magnetic activity is seen, and the generation of large scale magnetic fields is still possible for fully convective stars.  This article briefly outlines our current understanding of the rotation-magnetic field relation. ",
        "introduction": "The rotational evolution of stars is the result of the complex interaction of several fundamental processes. First, the molecular cloud contracts conserving initial angular momentum spinning up the central object. Angular momentum can be stored in a disc, which may brake the rotation of the central object. After the disc is dissipated, the star can contract reaching the highest rotation rate after several ten million years of lifetime. Solar-type stars, i.e., stars with convective envelopes, start generating magnetic fields that couple to the stellar wind. The interaction between charged particles in the wind and the magnetic field generates a torque braking the star's rotation. In the case of the Sun, braking has led to a rotation rate of about 1 revolution every month.  According to the rotation-activity relation, rapidly rotating stars produce strong magnetic fields generating a strong magnetic torque that brakes the star. This leads to slower rotation, which in turn weakens the magnetic field production and braking is weakening, too. At young ages in open clusters, we observe rapidly rotating, very active stars, while the (single) field stars generally are slowly rotating and only weakly active. This means that in principle rotation and activity can tell about the age of a star \\citep[e.g.,][]{Barnes03, Barnes07}.  The connection between rotation, stellar wind, magnetic fields, and magnetic activity is reviewed in this splinter session summary. First, we give an overview about the current picture on rotation in both clusters and the field, i.e., in young and old stars. Next, we discuss results from direct and indirect magnetic field measurements and their connection to stellar wind theory. In the last part, we give a summary on the theoretical work on magnetic field generation through stellar dynamos. Low-mass stars and in particular the regime where stars become completely convective currently present a rather puzzling picture of the connection between magnetism, activity, and rotation. Thus, low-mass stars are in the focus of our summary. ",
        "conclusions": "Our current picture of magnetic field generation, rotation, and stellar activity may be summarized as follows:  \\begin{enumerate} \\item Rotation rates are available for a wide range of masses and ages.  Measurements of \\emph{projected} rotation velocities extend far into the brown dwarf regime, but direct measurements of rotational periods are lacking at very low masses. \\item We observe a sharp break in rotation around the threshold where stars become fully convective.  This probably indicates a breakdown of wind braking. \\item Magnetic field measurements as well as activity tracers like X-rays, H$\\alpha$, and radio emission show now obvious break at the convection boundary.  However, around spectral type M7 normalized activity strongly weakens and the relation between radio and X-ray emission breaks down. \\item Apparently, very low mass stars can have strong large-scale magnetic fields yet only little wind braking. This remains an unresolved problem. \\item A key for understanding spindown is a theoretical understanding of wind braking.  However, it is still a challenge for magnetic stellar wind theory to reliably calculate the wind torque for a range of stellar parameters.  Furthermore, the wind torque is affected by the mass loss rate, so it is very important that we get measurements of mass loss rates and continue to improve mass loss theory. \\item Efforts to theoretically understand magnetic field generation evolved from the solar dynamo to the larger class of stellar dynamos, in particular to fully convective ones in absence of a tachocline.  First models successfully reproduce magnetic field generation, but it certainly is still a long way to understanding magnetic dynamos in very cool stars. \\end{enumerate}     \\begin{theacknowledgments} We thank the organizers of CS15 for giving us the opportunity to hold this session. \\end{theacknowledgments}"
    },
    "0809/0809.1155_arXiv.txt": {
        "abstract": "{} {Estimating Alfv\\'en speeds is of interest in modelling the solar corona, studying the coronal heating problem and understanding the initiation and propagation of coronal mass ejections (CMEs). } {We assume here that the corona is in a magnetohydrostatic equilibrium and that, because of the low plasma $\\beta$, one may decouple the magnetic forces from pressure and gravity. The magnetic field is then described by a force-free field for which we perform a statistical study of the magnetic field strength with height for four different active regions. The plasma along each field line is assumed to be in a hydrostatic equilibrium. As a first approximation, the coronal plasma is assumed to be isothermal with a constant or varying gravity with height. We study a bipolar magnetic field with a ring distribution of currents, and apply this method to four active regions associated with different eruptive events. } {By studying the global properties of the magnetic field strength above active regions, we conclude that (i) most of the magnetic flux is localized within 50 Mm of the photosphere, (ii) most of the energy is stored in the corona below 150 Mm, (iii) most of the magnetic field strength decays with height for a nonlinear force-free field slower than for a potential field. The Alfv\\'en speed values in an isothermal atmosphere can vary by two orders of magnitude (up to 100000 km/s). The global properties of the Alfv\\'en speed are sensitive to the nature of the magnetic configuration. For an active region with highly twisted flux tubes, the Alfv\\'en speed is significantly increased at the typical height of the twisted flux bundles; in flaring regions, the average Alfv\\'en speeds are above 5000 \\kms~and depart strongly from potential field values.  } {We discuss the implications of this model for the reconnection rate and inflow speed, the coronal plasma $\\beta$ and the Alfv\\'en transit time.} ",
        "introduction": "\\label{sec:intro}  In \\citeyear{alf42}, \\citeauthor{alf42} discovered that magnetohydrodynamic (MHD) waves are of three kinds: two magnetoacoustic waves (slow and fast modes) which are compressible and can be subject to damping, and the so-called Alfv\\'en wave which propagates along the magnetic field and is incompressible. The Alfv\\'en wave propagates at the Alfv\\'en speed given by: \\begin{equation} v_A = \\frac{B}{\\sqrt{\\mu_0 \\rho}} \\end{equation} where $B$ is the magnetic field strength and $\\rho$ is the density of the plasma. Since Alfv\\'en's work on MHD waves, the Alfv\\'en speed has been shown to be an important quantity in understanding many physical processes in the solar atmosphere such as (i) the coronal heating problem, (ii) the (non-) equilibrium state of the corona, (iii) the initiation and propagation of coronal mass ejections (CMEs).  The coronal heating problem aims basically to identify and understand the mechanisms responsible for sustaining the high temperature of the corona (about 1 MK) compared with the photosphere (about 4600 K). It is thought that the heating mechanism for the corona is of magnetic origin: the magnetic energy stored in the corona is released and converted into heat either by magnetic reconnection or by dissipation of waves \\cite[see review by][and reference therein]{kli06}. The Alfv\\'en speed appears as a tuning parameter in both heating processes. For instance, in a reconnection model such as the Sweet-Parker model \\citep{swe58, par63}, the outflow speed is of the order the Alfv\\'en speed of the inflow region. Therefore, the measurement of the local Alfv\\'en speed in the corona is an estimate of the outflow speed in the case of a possible reconnection process. Observationally, \\citet{nar06} and \\citet{nag06} have analysed numerous EUV and soft X-ray flares to deduce an inflow speed of a few tens of \\kms~and a reconnection rate between 10$^{-3}$ and 10$^{-1}$ of the Alfv\\'en speed. These results are consistent with the estimates made by \\cite{der96}. In terms of MHD waves, the mode conversion between slow and fast magnetoacoustic waves is the efficient when the sound speed is equal to the Alfv\\'en speed for a disturbance propagating along the magnetic field. Nevertheless under coronal conditions, the heating by dissipation of Alfv\\'en wave is more likely than by magnetoacoustic waves. The dissipation of shear Alfv\\'en waves generated in the low corona can be considered as a plausible heating mechanism especially when the dynamics of Alfv\\'en waves is dominated by phase-mixing \\citep{hey83, hoo97, nak97} or resonant absorption \\citep{gro88, poe89, poe90}. The damping of shear Alfv\\'en waves by phase-mixing is associated with the existence of a gradient in the Alfv\\'en speed and inhomogeneities in the background stratification. Resonant absorption is an other viable mechanism for coronal heating when the driving velocity of the magnetic structure is equal to the local Alfv\\'en speed. Recently \\cite{mcl04, mcl06} showed that fast magnetoacoustic waves can be dissipated near a null point in a low-$\\beta$ plasma and that Alfv\\'en waves are more likely dissipated near separatrices. Both fast and Alfv\\'en waves contribute to Joule heating in the corona localized near topological elements. Furthermore, \\cite{lon07a} showed how fast magnetoacoustic waves maybe launched by transient reconnection.  A coronal magnetic configuration can be considered to be in equilibrium if the time scale of the evolution $\\tau_{eq}$ is longer than both the Alfv\\'en time scale $\\tau_{A}$ (or Alfv\\'en transit time) and the reconnection time scale $\\tau_{rec}$ \\citep[e.g.,][]{ant87, pri99}. For a typical length of 100 Mm and an Alfv\\'en speed of 2000 km$\\cdot$s$^{-1}$, $\\tau_{A}$ is of 1 min. The reconnection time $\\tau_{rec}$ is of few seconds. The transit time $\\tau_{eq}$ is often assumed to be of the order of 10-20 min in the corona for a typical length of 100 Mm and an upper limit of the plasma velocity of 100 km$\\cdot$s$^{-1}$ corresponding to a sub-sonic coronal flow. In the limit of $\\tau_{eq} \\gg \\tau_{A} \\gg \\tau_{rec}$, the magnetic configuration can relax to a minimum energy state, i.e. a linear force-free field according to \\cite{wol58}. Then the evolution of the magnetic field can be described as a series of linear force-free equilibria \\citep{hey84}. If $\\tau_{eq} > \\tau_{A} \\gg \\tau_{rec}$, the magnetic configuration can relax to a nonlinear force-free equilibrium. \\cite{reg06} showed that the evolution of an active region can be well described by a series of nonlinear force-free equilibria, the time span between two equilibria being of 15 min.  In the different classes of CMEs, two main ingredients are often considered: the existence of a twisted flux tube or sheared arcade that can store mass and magnetic energy, and the existence of a current sheet above or below the twisted flux tube or sheared arcade in order to change the connectivity of the field lines by magnetic reconnection. Different models exist. The classical CSHKP model \\citep{car64, stu68, hir74, kop76} assumes the formation of a current sheet below a rising twisted flux tube in the pre-flare phase leading to the eruption of the flux rope and the formation of post-flare loops. According to \\cite{car82} and contrary to \\cite{kop76}, the evolution of post-flare loops is better described by  slow-shock waves (ohmic heating). In the breakout model \\citep{ant99, mac04}, the increase of shear in the underlying field lines of a quadrupolar configuration leads to field lines opening at the top of the magnetic system. This model leads to the formation of a flux rope in 3D \\citep{lyn05}.  The development of sheared arcades and/or the instability of a flux rope are then a prerequisite to a CME. The coronal sheared arcades are formed by shearing motions of magnetic polarities in the photosphere. The formation of unstable flux ropes and/or the destabilisation of an existing flux rope in the corona are often considered: kink instability \\citep[e.g.,][]{hoo81}, flux emergence \\citep[e.g.,][]{ama00, che00}, cancellation of flux \\citep[e.g.,][]{par94, pri94}, tether-cutting \\citep[e.g.,][]{moo80}. As noticed above, the reconnection process is an important step in the CME: the speed and the acceleration of a CME is related to the local Alfv\\'en speed and to the reconnection rate. Therefore, the nature of CMEs can be inferred from the determination of the local Alfv\\'en speed: to know where the reconnection is more likely to happen and to estimate the speed of the CME.  To our knowledge, few attempts have been made to characterise the Alfv\\'en speed in a coronal magnetic configuration. \\cite{der96} performed a study of the Alfv\\'en speed in active regions, quiet-Sun regions and coronal holes based on magnetic field estimates and/or potential extrapolations, the density and loop length being derived from coronal observations. \\cite{war05} derived the Alfv\\'en speed associated with the propagation of wave disturbances by modelling the magnetic field of an active region by a magnetic dipole superimposed on that of the quiet Sun as in \\cite{man03} and constraining the density by observations. Following \\cite{lin02}, the local Alfv\\'en speed in an isothermal atmosphere with a constant gravity is assumed to decrease with height until about 1.5 solar radii ($\\sim$ 350 Mm above the surface) and then to increase. When a solar wind component is added to the modelled atmosphere \\citep{sit99}, \\cite{lin02} showed that the Alfv\\'en speed decreases with height but is still consistent with Alfv\\'en speed values in the low corona ($<$ 350 Mm). Nevertheless, the variation with height and the dynamical range of the local Alfv\\'en speed are not well-known in the low corona.  As the physical processes mentioned above rely on the nature of magnetic configurations in the corona, the first part of this paper is dedicated to a determination of the magnetic field strength and its variation from the photosphere to the corona. The coronal magnetic field is not easily measurable. Several attempts have been made to measure the magnetic field in prominences via the Hanle effect \\citep{ler84}, or above active regions from gyro-emission frequencies \\citep{whi97}. Recently the possibility of measuring the magnetic field based on infrared wavelengths has been investigated as well as the inference of the coronal magnetic field strength from the properties of type II bursts \\citep{cho07}. \\citet{pol75} derived the magnetic field strength of several active region field lines using a potential field extrapolation. Based on this method the magnetic field strength decreases with height faster than $1/r$ from the photosphere and depends on the total magnetic flux and the spatial distribution of polarities on the photosphere. The significant results in measuring the coronal magnetic field have been summarized by \\citet{kou04}. Using a nonlinear force-free field method to extrapolate the observed photospheric magnetic field into the corona, we study the global properties of the magnetic field in the corona for different active regions. In \\cite{reg07a, reg07}, the authors described the geometry of field lines as well as the variations of energy density with height. In this paper, we focus on the distributions of the magnetic field strength with height revealing different behaviours between different active regions.  The second part of this paper is to derive Alfv\\'en speeds in the corona. The assumption of our model is that the corona satisfies a magnetohydrostatic equilibrium given by: \\begin{equation} - \\vec \\nabla p + \\rho \\vec g + \\vec j \\wedge \\vec B = \\vec 0, \\label{eq:mhs} \\end{equation} describing the balance between plasma pressure $p$, gravity and magnetic forces. Several magnetohydrostatic models already exist and are based on the extrapolation of the magnetic field into the corona satisfying Eqn.(\\ref{eq:mhs}) with or without gravity \\citep{low85, low86, low91, low92, low93, neu95, low95, neu97}. Recently \\cite{wie06a} have developed an optimization scheme to address this problem. If we assume that the magnetic forces dominate the corona, so that there is a force-free equilibrium for the coronal magnetic field and a hydrostatic equilibrium for the distribution of the plasma along each field line: \\begin{equation} (\\vec \\nabla \\wedge \\vec B) \\wedge \\vec B = o(\\beta) \\label{eq:ffo} \\end{equation} and \\begin{equation} \\vec B \\cdot \\left(- \\vec \\nabla p + \\rho \\vec g\\right) = 0. \\label{eq:hs} \\end{equation} These equations imply that there is no feedback of the coronal plasma on the magnetic field. However the magnetic field strongly influences the plasma by determining the shapes of coronal loops. Under coronal conditions, we assume that the plasma $\\beta$ is much less than 1: Eqn~(\\ref{eq:ffo}) is a force-free equilibrium and Eqn.~(\\ref{eq:hs}) is the hydrostatic equilibrium. Solving these equations, we investigate the values and distribution of the Alfv\\'en speed in the low corona assuming that the corona is  isothermal. We only consider the global properties of the Alfv\\'en speed and not the Alfv\\'en speed profiles along individual field lines.  In Section~\\ref{sec:obs_mod}, we first describe the observed active regions, the magnetic field model as well as the isothermal atmosphere model required to compute the Alfv\\'en speeds. In Section~\\ref{sec:model_param}, we then study the field strength and Alfv\\'en speed in a twisted bipolar field. We apply this model to four solar active regions in Section~\\ref{sec:ar}. We discuss the implications of these results on the physics of active regions and CMEs in Section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  \\begin{figure*} \\centering \\includegraphics[width=.49\\linewidth]{0362_f13a.eps} \\includegraphics[width=.49\\linewidth]{0362_f13b.eps} \\includegraphics[width=.49\\linewidth]{0362_f13c.eps} \\includegraphics[width=.49\\linewidth]{0362_f13d.eps} \\caption{Plasma $\\beta$ for the nonlinear force-free field (solid line) and the potential field (dashed line) as a function of height for four active regions. Red (resp. blue) curves are for the varying (resp. constant) gravity model.} \\label{fig:beta} \\end{figure*} In this paper, we address two questions related to the global properties of the magnetic field and the plasma in the corona: \\begin{itemize} \\item[1-]{How does the magnetic field strength and Alfv\\'en speed vary with height in the solar corona?} \\item[2-]{What can we learn from the estimate of the coronal Alfv\\'en speed for some key solar physics problems (coronal heating, magnetic reconnection, wave propagation, initiation and development of CMEs)? } \\end{itemize}  In order to tackle the above questions, we have developed a method to determine the plasma properties in a magnetic field above active regions. We assume that an active region can be modelled in two steps: a nonlinear force-free equilibrium for the magnetic field, and a hydrostatic equilibrium between pressure and gravity forces describing the plasma properties. To solve the latter equation, the corona is assumed here to be isothermal with either a constant gravity or a gravitational field varying with height.  In Sections~\\ref{sec:bdip} and~\\ref{sec:bstr}, we investigated the evolution of the average field strength with height. Interestingly, the log-log plots showed that the decay of the magnetic field with height does not follow a power law, even for a potential field. This fact suggests that the decay of the field strength is strongly influenced by the geometry of the field (e.g., toroidal geometry for the bipolar field), by the distribution of electric currents and their typical height of influence, and by the complexity of the magnetic distribution on the photosphere for active regions. This is a well-known property of the magnetic field. Based on the data reported by \\cite{pol75}, \\cite{van78}  plotted the measured magnetic field strength as a function of height for three different active regions and assuming current-free magnetic configurations. They obtained similar curves to those shown in this paper in which the log-log plot can be fitted by a parabola. In addition, they fitted the distribution of field strength by two power laws of index -1 and -3 suggested by their theoretical work: -1 for a line current along the inversion line of the magnetic field and -3 for a dipolar potential field decay. \\cite{van78}  concluded that (i) the electric currents strongly influence the magnetic configuration when the power law index is less than -1, (ii) above a certain height the magnetic field decays rapidly (probably with a power law index of -3). More importantly, the authors concluded that, if the magnetic field below the $z^{-1}$ decay is decaying with a power law index less than -1, then the configuration is stable. This condition can be compared with the conditions required to trigger a kink instability \\citep[e.g.,][]{hoo81} or a torus instability \\citep{tor06}. Recently,  \\cite{liu08} investigated the evolution of the magnetic field strength with height and compared with the nature of the eruptive phenomenon. He derived the magnetic field from a large-scale potential field extrapolation based on the PFSS model (Potential Field Source Surface, e.g., \\citeauthor{sch69} \\citeyear{sch69}). For heights between 42 Mm and 140 Mm, the power law index is about -1.5 whatever the onset mechanism. Combined with our study, this result suggests that the behaviour of the magnetic field in the  low atmosphere plays the main role in triggering flares and CMEs, and also that electric currents flowing along field lines as in the \\nlff~model affect strongly the power law index and therefore the stability of the magnetic configuration. \\begin{figure*} \\centering \\includegraphics[width=.49\\linewidth]{0362_f14a.eps} \\includegraphics[width=.49\\linewidth]{0362_f14b.eps} \\includegraphics[width=.49\\linewidth]{0362_f14c.eps} \\includegraphics[width=.49\\linewidth]{0362_f14d.eps} \\caption{Average Alfv\\'en transit time for the potential (dashed line) and nonlinear force-free (solid line) fields using as typical scale the pressure scale-height $H$ (bottom curves) or the length of a semi-circular loop with a diameter of the scale of the computational box (top curves). Red (resp. blue) curves are for the varying (resp. constant) gravity model.} \\label{fig:tau} \\end{figure*}  \\subsection{Reconnection rate}  As mentioned in Section~\\ref{sec:intro}, the Alfv\\'en speed is an important parameter in understanding the magnetic reconnection processes. The reconnection is characterised by the reconnection rate (or external Alfv\\'en Mach number) as follows: \\begin{equation} M_e = \\frac{v_i}{v_A}, \\end{equation} where $v_i$ is the inflow speed. This definition does not depend on the nature of the diffusion region where the magnetic reconnection occurs. To distinguish between different regimes of reconnection, the magnetic Reynolds number is introduced as follows: \\begin{equation} R_{me} = \\frac{\\eta}{L_e v_A}, \\end{equation} where $\\eta$ is the magnetic diffusivity, and $L_e$ is a typical length of the diffusion region \\citep[e.g.,][]{pri00}. If $R_{me}^{-1} < M_e < R_{me}^{-1/2}$, we have a slow reconnection regime of the Sweet-Parker type. If $M_e > R_{me}^{-1/2}$, the reconnection regime is fast such as in the \\cite{pet64} model. For a typical astrophysical plasma, the reconnection rate in the fast regime is in the range 0.01 to 0.1. The square of the reconnection rate is also the ratio of kinetic energy to magnetic energy of the inflow region assuming that the diffusion region is small enough to have the same field strength in the inflow and outflow regions. For a given inflow speed, the higher is the magnetic energy stored in the region, the smaller is the reconnection rate.  If we assume that the reconnection rate is between 10$^{-3}$ and 10$^{-1}$ \\citep[see e.g.,][]{nar06}, we can infer the inflow speed required to trigger the reconnection. The results for the four studied active regions are summarized in Fig.~\\ref{fig:inflow}, where we plot the inflow speed as a function of height for different reconnection rates ($M_e= 0.1, 0.01, 0.001$) for both potential (dashed curves) and \\nlff~(solid curves) fields. The variation with height of the inflow  speed is consistent with the variation of the Alfv\\'en speed as described in Fig.~\\ref{fig:va_iso}. From these curves, we can deduce a threshold on the reconnection rate in order to obtain inflow speeds in agreement with observations \\citep[e.g.,][]{nar06, nag06}. For instance, if a reconnection process occurs in AR 8151, the reconnection rate should be of the order or below 0.01 which gives typical inflow speeds below 50 \\kms.   We also note that the minimum of the inflow speed or Alfv\\'en speed (see Section~\\ref{sec:model_param}) is the location in the low corona where the reconnection process is more likely to happen. The conclusion is valid in a statistical sense because we only consider average values of the inflow speed and because, in order for it to occur, the reconnection process needs a particular magnetic topology and a diffusion region which are beyond the scope of this article.  \\subsection{Coronal plasma $\\beta$}  An important quantity in plasma physics is the plasma $\\beta$, the ratio of the gas pressure to the magnetic pressure. When $\\beta \\ll 1$, the plasma is dominated by the magnetic field. Nevertheless, the plasma $\\beta$ in the solar atmosphere varies greatly with height \\citep{gar01}. In addition, we note that the plasma $\\beta$ varies from one structure to another and even coronal structures, such as filaments, can have $\\beta \\geq 1$. As mentioned in Section~\\ref{sec:model_param}, the plasma $\\beta$ can be expressed as follows: \\begin{equation} \\beta = \\frac{2~c_s^2}{\\gamma v_A^2} \\end{equation} for an isothermal atmosphere ($\\gamma = 1$). In Fig.~\\ref{fig:beta}, we plot the variation of the average plasma $\\beta$ with height for the four active regions described in Section~\\ref{sec:obs_ars}. For both potential and \\nlff~models, the plasma $\\beta$ is less than 1. This result is consistent with the typical nature of the coronal active region magnetic field. The plasma $\\beta$ is larger for the potential field because the average Alfv\\'en speed decreases with height more rapidly than for the \\nlff~field. The maximum of the plasma $\\beta$ is located at the same height as the minimum of Alfv\\'en speed (see Section~\\ref{sec:model_va}) and also  satisfies Eqn.~(\\ref{eq:min_va}). The plasma $\\beta$ tends to zero for a constant gravitational field, whilst it increases towards 1 for the varying gravity model. The latter is then in agreement with plasma $\\beta$ measurements reported by \\cite{gar01}.  \\subsection{Average Alfv\\'en transit time}  The Alfv\\'en transit time is an important quantity for the propagation of Alfv\\'en waves and for the relaxation of a magnetic configuration. In this study of global properties of Alfv\\'en speeds, we derive the average Alfv\\'en transit time associated with two different lengths: \\begin{equation} \\tau_{A}^{(H)} = \\frac{H}{v_A} \\end{equation} and \\begin{equation} \\tau_{A}^{(\\pi L)} = \\frac{\\pi L}{v_A}, \\end{equation} where $v_A$ is the Alfv\\'en speed, $H$ is the pressure scale-height and $L$ is the length of the computational box. We suppose $\\tau_{A}^{(H)}$ (resp. $\\tau_{A}^{(\\pi L)}$) is a minimum (resp. maximum) of the Alfv\\'en transit time. The larger the Alfv\\'en transit time is, the more stable the equilibrium is. In Fig.~\\ref{fig:tau}, we plot the different Alfv\\'en transit times for the four active regions. At the base of the corona (5 Mm), the maximum Alfv\\'en time is 750s for AR 8151, 3600s for AR 8210, 70s for AR 9077, 300s for AR 10486. These results mean that the most stable active  region is AR 8210 associated with confined C-class flares, and the least stable one is AR 9077 exhibiting post-flare loops. It is worth noticing that the average Alfv\\'en transit time is an interesting quantity only when discussing the global or statistical properties of  magnetic configurations. A more appropriate quantity is the Alfv\\'en transit time along a single field line which is beyond the scope of this article. Of importance for MHD modelling, the Alfv\\'en transit time along field lines is not a constant with height which makes it difficult to estimate the scaling of the time evolution with this quantity.  Despite the simple assumptions of this model, we have derived interesting properties of the Alfv\\'en speed in the solar corona. In a forthcoming paper, we will focus on the properties of individual field lines."
    },
    "0809/0809.4894_arXiv.txt": {
        "abstract": "We present the results of three~body simulations focused on understanding the fates of intermediate mass black holes (IBH) that drift within the central 0.5 pc of the Galaxy. In particular, we modeled the interactions between pairs of $4000 {\\rm M}_{\\odot}$ black holes as they orbit a central black hole of mass $4 \\times 10^6 {\\rm M}_{\\odot}$. The simulations performed assume a Schwarzschild geometry and account for Chandrasekhar dynamical friction as well as acceleration resulting from energy lost due to gravitational radiation. We found the branching ratio for one of the orbiting IBHs to merge with the CBH was 0.95 and is independent of the inner IBH's initial eccentricity as well as the rate of sinking.  This, coupled with an infall rate of $\\sim 10^7$ yrs for an IBH to drift into the Galactic center, results in an IBH-CBH merger every $\\lesssim 11$ Myrs.  Lastly we found that the IBH-IBH-CBH triple body system ``resets'' itself, in the sense that a system with an inner IBH with an initially circular orbit generally left behind an IBH with a large eccentricity, whereas a system in which the inner IBH had a high eccentricity ($e_0 \\sim 0.9$) usually left a remnant with low eccentricity. Branching ratios for different outcomes are also similar in the two cases. ",
        "introduction": "\\label{sec:Intro}  Starting with the discovery of quasars \\citep{Schmidt63, Greenstein64, Greenstein64b}, we have gradually come to appreciate the existence of massive, dark objects in the centers of almost all galaxies \\citep{Bell69, Richstone98}, although their origins are still unclear. The nature of such a central dark object is best constrained in the center of our own Galaxy. Over the last decade, the infra-red monitoring of the Galactic center has led to severe dynamical constraints on the nature of the Milky Way's central object \\citep{Eckart97, Ghez98, Schodel02, Ghez05}, leaving little doubt that it is a black hole with mass $3.7 \\pm 0.2 \\times 10^6 {\\rm M}_{\\odot}$.  An equally interesting result to emerge from these studies is that the stellar population within the central parsec of the galaxy contains an unexpected component of young, massive stars \\citep{Sanders92, Morris93, Genzel97, Ghez03, Paumard06}. The origin of these stars is a question of interest because the strong tidal fields close to the central black hole (CBH) are large enough to shear apart a normal molecular cloud at the locations where the stars are seen \\citep{Morris93}. Two principal classes of alternative models have been investigated -- one which postulates that the stars form from the gravitational instability of a quasar-like accretion disk around the CBH during an earlier period of high nuclear activity \\citep{Levin03} and another which proposes that the stars migrate inwards as part of a stellar cluster which sinks towards the CBH by dynamical friction and is eventually broken up by the strong tidal field \\citep{Gerhard01}. A recent addition to this latter scenario is the possibility that such a cluster contains an intermediate mass black hole \\citep{Hansen03}, which serves to bind the cluster tighter and slow the internal dynamical relaxation, allowing the cluster to survive long enough to deposit the stars in their observed locations.  Each of these scenarios have their proponents, and the arguments for and against are presented elsewhere (e.g. \\citep{Paumard06, Berukoff06}). The goal of this paper is to examine the implications the cluster infall scenario for the merger history of the central black hole. In particular, once the intermediate mass black hole (IBH) is stripped of its coterie of cluster stars, it continues to sink towards the center due to dynamical friction. We wish to examine the final fate of such an IBH and how that fate is reached. We will explore in \\S\\ref{sec:Merge} reasons to believe that the rate of infall of such IBH into the central parsec may actually be larger than the rate at which an individual IBH would merge with the CBH. This then opens the possibility of multiple IBH in orbit, so that in \\S\\ref{sec:Orbs} we examine the gravitational interactions of such CBH-IBH-IBH triples and explore the probabilities of various outcomes. The results of this study have a variety of interesting applications, not the least of which is a prediction for the kinds of signals we might expect to observe with future space based gravitational wave observatories. We discuss this further in \\S\\ref{sec:Discuss}. ",
        "conclusions": "\\label{sec:Discuss}  The model problem discussed here was motivated by the desire to understand the consequences of a steady rate of infall of IBH into the Galactic center, a potential consequence of theories for the origins of the young stars observed there. Perhaps surprisingly, we find that the probability of a merger between the CBH and one of the IBH is greater than the probability of an ejection (figure \\ref{fig:probs}). Coupling the observed branching ratio of 0.65 for an IBH-CBH merger, along with the infall rate $\\sim 10^7$ yrs (consistent with the ages of  the young stars found in the central parsec), results in an IBH-CBH merger occurring approximately once every 11 Myrs. We have also found similar results for both enhanced and reduced dynamical friction ($t_{fric} \\sim 10^6$ yrs and $t_{fric} \\sim 10^8$ yrs) where the differences are associated with statistical error and the total time for infall and IBH-IBH interactions to occur. Depending on the rate of IBH infall, such mergers may have contributed significantly to the mass growth of the CBH \\citep{Ebisuzaki01}.  The eccentricity of a surviving IBH is determined by several factors. Dynamical friction will circularise the orbit (e.g. \\citet{Berukoff06}), although this depends on the background density profile. Indeed, eccentricity can be pumped if there is an evacuated core, so that dynamical friction operates over only part of the orbit, as we have seen in this calculation and has been verified in numerical simulations such as \\citet{LockBaum}. In our case, these effects are superceded by the strong interactions between the pair of IBH, both in the form of resonant interactions and close passages and scattering. However, we find that a system that starts with an inner IBH on an eccentric orbit is likely to leave a survivor on a low eccentricity orbit, and vice versa. Taken together, this implies that the outcome of a series of IBH infalls can be considered as a chain of three-body encounters such as the ones described here, unless the rate of infall is considerably faster than one every $10^7$~years.  In conclusion, although the stellar scattering responsible for dynamical friction cannot cause an IBH-CBH merger in a Hubble time, the dynamical interactions between two or more IBH result in ejection or merger on timescales $\\sim 10^7$~years, so that the number of IBH inhabiting the central parsec today should be small in number, even if the process of infall is a regular feature of Galactic center life."
    },
    "0809/0809.1438_arXiv.txt": {
        "abstract": "We study the formation and dynamical evolution of clusters with multiple stellar generations. Observational studies have found that some globular clusters host a population of second generation (SG) stars which show chemical anomalies and must have formed from gas containing matter processed in the envelopes of first generation (FG) cluster stars. We study the SG formation process by means of 1D hydrodynamical simulations, starting from a FG already in place and assuming that the SG is formed by the gas ejected by the Asymptotic Giant Branch (AGB) stars.  This gas collects in a cooling flow into the cluster core, where it forms SG stars. The SG subsystem emerging from this process is initially strongly concentrated in the cluster innermost regions and its structural properties are largely independent of the FG initial properties. We also present the results of a model in which pristine gas contributes to the SG formation. In this model a very helium-rich SG population and one with a moderate helium enrichment form; the resulting SG bimodal helium distribution resembles that observed for SG stars in NGC~2808.  By means of $N$-body simulations, we then study the two-population cluster dynamical evolution and mass loss.  In our simulations, a large fraction of FG stars are lost early in the cluster evolution due to the expansion and stripping of the cluster outer layers resulting from early mass loss associated with FG SN ejecta.  The SG population, initially concentrated in the innermost cluster regions, is largely unscathed by this early mass loss, and this early evolution leads to values of the number ratio of SG to FG stars consistent with observations.  We also demonstrate possible evolutionary routes leading to the loss of most of the FG population, leaving an SG-dominated cluster.  As the cluster evolves and the two populations mix, the local ratio of SG to FG stars, initially a decreasing function of radius, tends to a constant value in the inner parts of the cluster.  Until mixing is complete, the radial profile of this number ratio is characterized by a flat inner part and a declining portion in the outer cluster regions. ",
        "introduction": "\\label{sec:introduction} The common paradigm that Globular Clusters (GCs) are examples of ``Simple Stellar Populations'' (SSP), assemblies of stars born all at the same time and sharing identical chemical composition, had until a few years ago only one noticeable exception: $\\omega$~Cen, the most massive cluster, whose stars show a large spread in metallicity. Today, the situation is much more complex, and it is now clear that practically none of the GCs so far examined in detail is an SSP. Even when all the stars in a single GC share the same Fe abundance, star-to-star differences in the abundances of lighter elements are found. These differences concern light-element abundances: in a relevant fraction of cluster stars, they look like the result of nuclear processing through proton capture reactions at high temperature (CN, ON, NeNa and MgAl cycles). In particular, O-Na and Mg-Al anticorrelations are found for stars at different stages of their evolution, from the red giant branch to the main sequence (MS) \\citep[e.g.][]{gratton-ar2004}.  Halo field stars with the same metallicity do not show such huge variations \\citep{kra94}. This leads us to attribute the anomalous composition to the fact that these stars formed in a dense GC environment, and that nuclear processing in a first stellar generation (hereafter FG), has been followed by a subsequent episode of star formation, creating a second stellar generation (hereafter SG) from matter including this processed gas.  An additional ingredient to the complexity of the new scenario has been added by recognizing that there are at least two clusters, \\ocen\\ and NGC~2808, in which the MS morphology shows that a non negligible fraction of the stars ($\\sim$25\\% in \\ocen, $\\sim$15\\% in NGC~2808) populate a ``blue MS'' which can only be interpreted as a very helium-rich MS \\citep{bedin2004, norris2004, piotto2005, dantona2005}. The extreme helium population must also be very {\\it homogeneous}, both in helium and metal content, since in both clusters the blue MS is well detached from the rest of the main sequence. \\cite{piotto2007} have shown that the MS of NGC~2808 is actually made up of three different sequences: a normal MS, a population with somewhat enriched helium (Y$\\sim$ 0.30) and a very helium-rich one (Y$\\sim$0.35--0.40). In \\ocen\\ the ``normal'' MS is actually made up of sequences with different metal content, as also shown by the complexity of the subgiant branch(es) \\citep[e.g.][]{villanova2007}. In addition to \\ocen and NGC~2808, the two massive clusters NGC~6441 and NGC~6388 apparently also harbor a very high helium population, including $\\sim 10$\\% and $\\sim 20$\\% of stars respectively.  In this case the evidence comes  from the high-\\teff\\  extension of the Horizontal Branch (HB) \\citep{caloi2007, busso2007,yoon2008}, a very peculiar occurrence at the high metallicity of these clusters \\citep{rich1997}, and it is not clear whether the very high helium population is as homogeneous as it is in the previous cases.  Observationally, it appears that these extreme Y sequences are present only in the most massive clusters \\citep{piotto2007}. A moderately helium rich population ($0.26 < \\rm{Y} < 0.30$), on the other hand, is present in most GCs, largely independent of the total cluster mass. The observational spectroscopic evidence resides mainly in the existence of similar Na--O anticorrelation ---also among the unevolved stars--- in all the clusters so far examined \\citep[see, e.g., the summary by][]{carretta2006} although the most extreme Na--O anomalies are present only in clusters that also show extreme high-\\teff\\ extensions of the HB \\citep[see][]{carretta2007}.  These results, based generally on spectroscopy of a few stars, have been recently confirmed by the analysis of larger samples of 30--100 stars per cluster (Carretta et al., 2008) The indications coming from the analysis of the CNO anomalies are similar \\citep[for a summary, see e.g.][]{cohenetal2005}.  Photometrically, this non-extreme population can not be easily isolated in the MS photometry, mainly because the moderately helium rich isochrones merge with the standard isochrones in the turnoff region \\citep{dantona2002,salaris2006}, but also because, in general, these SG stars do not have an unique value of Y \\citep[e.g.][]{lee2005,caloi2007}.  Nevertheless, these stars show up quite prominently in the HB morphology, which amplifies the mass differences (due to the initial helium content) among stars of the same age \\citep{dantona2002,daca04}.\\footnote{The age difference between FG and SG is in any case so small with respect to the total GC age that it is irrelevant for the turnoff evolution.}  It is important to notice that, according to the detailed HB modelling by  \\cite{caloi2007} \\citep[see also][]{dantonacaloi2008},  also in NGC~6441 and NGC~6388 the total SG population  includes $\\sim 60 \\%$ of the total number of stars, and is therefore much larger than just the very high helium fraction population.  The interpretation of the HB morphologies in terms of helium enrichment, coupled with the spectroscopic information indicating that the helium-rich samples probably coincide with the sample of stars showing the Na--O anomalies \\citep[e.g.][]{carretta2006}, lead us to conclude that, in the clusters so far examined, the percentage of stars of the SG is generally $\\sim$40--60\\% \\citep{dantonacaloi2008}.  For the clusters having mainly blue HBs, there are even indications that they may contain only---or mainly---SG stars \\citep{dantonacaloi2008}. In an analysis of the relative magnitude locations of the turnoff, red giant ``bump'' and HB, \\cite{caloi-dantona2005} have provocatively suggested this possibility for M~13, the cluster showing the broadest range of chemical anomalies \\citep{sneden2004}.  In addition, some clusters in which the tightness of the evolutionary sequences in the HR diagram \\citep[e.g. the relatively small cluster NGC 6397,][]{king1998,richer2006} or the short extent of the HB \\citep[e.g. the massive GC M 53, ][]{dantonacaloi2008} indicate that they indeed represent SSPs, might be dominated by a {\\it homogeneous} SG. For example, in NGC~6397 only three out of fourteen scarcely evolved stars examined by \\cite{carretta2005} are normal in their nitrogen content. In the others, [N/Fe]$\\sim 1 -1.7$, showing CN and ON processing in stars that should not show it. The blue HB of M~53 can easily be described as uniparametric \\citep{dantonacaloi2008}, but its integrated spectrum shows a strong nitrogen enhancement with respect to field halo stars \\citep{liburstein2003}.  In conclusion, some ``normal'' blue--HB GCs (with typical current masses of a few $\\times 10^5\\msun$) should be considered self-enriched, and any model for the formation and evolution of multiple population clusters must be able to explain the observed large SG population and the cluster-to-cluster variations in the SG/FG number ratio.  Both massive Asymptotic Giant Branch (AGB) stars \\citep{cottrell1981,ventura2001} or ``fast rotating'' massive stars (FRMS) \\citep{mm2006,decressin} have been proposed as possible progenitors of the SG. The most critical point, common to both scenarios, is the following: for a standard Initial Mass Function (IMF) of the FG, such as a Salpeter or Kroupa IMF, the gas ejected by the massive stars during the phase of H-burning, or the gas contained in the envelopes of massive AGBs \\citep{daca04, bekki-norris2006}, is too scarce to form a large SG population.  Two solutions to this conundrum have been suggested: either the system had a standard IMF but it was initially much more massive---{\\it at least a factor 10 more massive}, under very conservative assumptions (see Sect. 3.1)---than today's GC \\citep{daca04, bekki-norris2006, prch06, dantona2007, decressin2007b}, or the system evolves at constant mass and the IMF is highly anomalous \\citep[][see also \\cite{dosills07} for a dynamical study]{daca04,decressin2007b} and contains a significant number of FG polluters. As we will discuss in Sect. \\ref{subsec:imf}, a large initial mass is necessary also in this second scenario if a cluster must have today an approximately equal number of long-lived SG and FG stars (where we define long-lived stars as those having initial masses in the range 0.1 $M_{\\odot}$--0.8$M_{\\odot}$).  We reconsider here the problem of GC formation to explain not only the most peculiar very helium-rich population (15-20\\% in three GCs, 25\\% in \\ocen), but the whole SG population (about 40--60\\% of the current GC stellar content, and possibly even more in the ``blue HB'' clusters).  In this work, we follow the idea that the intermediate--mass stars of the FG, starting at 8M$_{\\odot}$ and extending at most down to $\\sim$4--5M$_{\\odot}$, lose by stellar winds and the ejection of planetary nebulae matter processed by the hot CNO cycle via Hot Bottom Burning (Na- and Al-rich, O- and Mg-poor) during their AGB evolution, and that an SG forms either from the pure AGB ejecta or from ejecta mixed with pristine gas. While previous work examined the reliability of this model from a chemical point of view, here we focus mainly on its hydrodynamical and dynamical consistency.  In Sect. \\ref{sec:sgprog} we first briefly summarize the state of the art of chemical modelling of the SG, we discuss the reasons why we focus our attention on the scenario in which the AGB stars are the polluters and emphasize the possible role of super-AGB stars.  In Sect. \\ref{sec:models} and \\ref{sec:res}, we address this  problem by means of 1D hydro simulations.  We show that in the massive AGB epoch, the slow winds collect in a cooling flow driving AGB ejecta to the center of the cluster, where they form SG stars. Simulations exploring the dependence of our results on the cluster mass, IMF, initial concentration as well as the role of SNe Ia and other extra energy sources are also presented in this section.  In Sect. \\ref{sec:pristine} we explore a model in which the pristine matter is partially re-accreted after the SN II explosion stage. The stellar population forming after the accretion shows a lower helium abundance relative to the stars formed earlier from the pure ejecta of the super-AGBs. This may explain the presence of the very high-helium population observed in a few, very massive GCs.  In Sect. \\ref{sec:dyn0} we present the results of a number of $N$-body simulations exploring the subsequent cluster dynamical evolution.  We show that the early cluster expansion driven by mass loss due to SN ejecta leads to a strong preferential loss of FG stars.  Our simulations show that dynamical evolution can lead to values of the ratio of the number of SG to FG stars consistent with those suggested by observational studies. We also show possible evolutionary routes leading to the loss of most of the FG stars, leaving an SG-dominated, cluster. ",
        "conclusions": "In this paper we have explored a scenario for the formation and the dynamical evolution of multiple stellar generations in globular clusters.  A successful model for the origin and the evolution of clusters with multiple stellar generations needs to satisfy a number of observational constraints concerning the current chemical properties, relative abundance and structural properties of the different stellar populations.  After reviewing the main observational results on multiple population clusters, we have pointed out that the key issues to address concern the general origin of SG stars (among which those with very high helium content observed in the most massive clusters are only a fraction) and the identification of the evolutionary processes leading to clusters with a similar number of SG and FG stars (or even to clusters dominated by SG stars).  In order to address these different constraints and be able to model all the relevant aspects of multiple population clusters, we have combined the input from stellar evolution models to constrain the origin  of the gas from which SG stars might form, hydrodynamical simulations to model the process of SG star formation and predict its initial structural properties and, finally, $N$-body simulations to explore the subsequent dynamical evolution of the two populations.  The main steps and results of this investigation are as follows.  \\subsection*{1.~Chemical Abundances} We have reviewed all the issues concerning the super-AGB and massive AGB stars as the source of the polluted gas from which SG stars could form.  The hypothesis that AGB stars might provide the gas for the SG formation has recently received some criticism in the literature based on possible inconsistencies between models and the observed chemical abundances in SG stars.  However, in our review, we have pointed out that the results of chemical processing during the Hot Bottom Burning phases are highly affected by the uncertainties in the basic input physics of these models, and that there is an ample range of plausible input parameters which can provide yields consistent with observations. In addition, recent models show that the super-AGB evolution can produce the helium content required for the very helium-rich population mentioned above. During the super-AGB epoch, between the end of SN II explosions and the beginning of the SN Ia explosions, there are no mechanisms preventing the formation of a cooling flow and the formations of the very helium-rich SG population from super-AGB winds.  \\subsection*{2.~Hydrodynamical Simulations} We have carried out a number of hydrodynamical simulations to explore the dynamics of gas from super-AGB and AGB stellar winds and the process of formation of the SG population in the assumption that the SNe II belonging to the FG population clear the GC of its pristine gas.  \\begin{enumerate}  \\item Our simulations show that this gas collects in a cooling flow into the innermost regions of the cluster where it forms SG stars. The cluster emerging from the hydrodynamical simulations is one with a SG strongly concentrated in the inner core of a more extended FG population. Initially the FG stars are the dominant stellar population with a total mass that must be about ten times larger than the total mass of the SG stars. The initial mass of the FG stars needs to be large enough to provide enough stellar mass return and form a substantial amount of SG stars.  \\item We have explored the dependence of our results on the initial mass of the cluster, the stellar IMF and the initial cluster size and concentration.  SG stars are always concentrated in the cluster innermost regions and the structure of the newly formed SG subsystem is largely independent of the cluster size and concentration.  \\item Starting with a initial FG mass ten times smaller than the value adopted in our standard model, leads to a SG total mass also ten times smaller but does not alter the shape of the SG density profile or its concentration.  For simulations with a flatter IMF the larger amount of gas supplied by AGB winds implies that a smaller star formation efficiency is needed to form the required mass of SG stars.  If the requirement that the current mass of long-lived SG and FG stars are similar is maintained, a large FG initial mass is needed also for a flatter IMF. In this case the number of long-lived FG stars is smaller, and a large value of $\\mi$ is still needed to have a current mass of long-lived FG stars consistent with observations.  \\item The effects of SNe Ia and additional energy sources on the SG formation process have been investigated. Our simulations show that, for a ``normal'' (e.g. Kroupa) IMF of the FG stars, the stellar explosions quickly evacuate most of the gas returned by the AGB stars, leading to a substantial drop in the gas content of the GC and to a suspension of the SG star formation.  Given the uncertainties on the SN Ia rate, the interruption of the SF may occur in the time interval $40<t<100$ Myr. Flat IMFs produce an increment in the amount of returned gas, leading to an enhancement of the radiative losses and, possibly, to the inability of the SNe Ia to get rid of the gas. In this case SG stars continue to form and their chemical abundances should show the signature of the pollution by gas from SN ejecta.  This however does not seem to be confirmed observationally and flat IMFs are thus disfavoured.  \\item We have further considered the possible presence of a diffuse energy source due to X-ray binaries and/or planetary nebulae. We show that for the standard model the cooling flow and the associated SG star formation can be halted by energy sources with a total power larger than $10^{38}$ erg s$^{-1}$.  \\item We have also explored the possibility that the FG SNe II do not vent away all the pristine gas, and that part of it remains gravitationally bound to the cluster and falls back onto it after an initial ousting.  The SG stars forming in the cluster central regions before this pristine gas reaches the cluster core originate from pure ejecta of the super-AGB stars which is very helium rich. These stars may form the very helium-rich population present in the most massive GCs. Once the pristine gas has fallen back in the cluster, SG stars form in an ambient medium in which the helium abundance is ``diluted,'' and give rise to the moderately helium rich population.  \\end{enumerate}  \\subsection*{3.~Stellar Dynamical Simulations} The subsequent dynamical evolution of the cluster has been studied by means of $N$-body simulations.  \\begin{enumerate}  \\item One of the critical issues for viability of the scenario presented here is the identification of a mechanism leading to the escape of a large fraction of the initial FG population.  For clusters as massive as required by our hydrodynamical models, the mass lost in one Hubble time due to two-body relaxation is negligible.  Following observational and theoretical studies suggesting that initial mass segregation may be imprinted in clusters by the star formation process or produced dynamically early in the cluster evolution, we have studied the evolution of initially segregated clusters.  In our simulations, a large fraction of FG cluster stars are lost early in the cluster evolution due to the expansion and stripping of the cluster outer layers resulting from early mass loss associated with FG SN ejecta. Initial FG mass segregation plays a key role since the heating and the expansion due to the loss of SN ejecta are augmented by the preferential removal of the mass from the inner regions of the cluster.  \\item The population of escapers is dominated by FG stars. The SG population, initially concentrated in the innermost cluster regions, is largely unscathed by the early mass loss.  We find that the early cluster evolution and mass loss can lead to a significant loss of FG stars and to current values of $\\popms$, consistent with observations ($0.5<\\popms<1.5$).  We have also demonstrated possible evolutionary routes leading to the loss of most of the FG population and leaving a SG-dominated cluster. Clusters initially filling their Roche lobe with higher degrees of initial mass segregation and/or losing impulsively a larger fraction of SN ejecta undergo more expansion and may lose most of their FG stars.  \\item After the initial phase of expansion and mass loss, cluster evolution is driven by two-body relaxation which results in the spatial mixing of the two populations.  Following the hydrodynamical simulations, in the $N$-body initial conditions SG stars are concentrated in the inner regions of the cluster.  The radial profile of the number ratio of SG to FG stars, $f_{MS}(r)$ is initially a decreasing function of radius.  We find that, as the cluster evolves and the two populations mix, $f_{MS}(r)$ flattens in the inner parts of the cluster. The central flat region of the $f_{MS}(r)$ radial profile expands on a relaxation timescale as mixing proceeds.  In general, unless mixing is complete, we predict that $\\popms(r)$ is characterized by a flat inner part and a declining portion in the outer regions of the cluster. \\end{enumerate}  \\par\\smallskip\\noindent A more extended survey of hydrodynamical and $N$-body simulations to further exlore the dynamics and the implications of the scenario presented here as well as for a more detailed comparison with observations is currently in progress and will be presented in future papers."
    },
    "0809/0809.3926_arXiv.txt": {
        "abstract": "{The high sensitivity of {\\it JWST} will open a new window on the end of the cosmological dark ages.  Small stellar clusters, with a stellar mass of several $\\times 10^6 {\\rm M_\\odot}$, and low-mass black holes (BHs), with a mass of several $\\times 10^5 {\\rm M_\\odot}$ should be directly detectable out to redshift $z=10$, and individual supernovae (SNe) and gamma ray burst GRB afterglows are bright enough to be visible beyond this redshift.  Dense primordial gas, in the process of collapsing from large scales to form protogalaxies, may also be possible to image through diffuse recombination line emission, possibly even before stars or BHs are formed.  In this article, I discuss the key physical processes that are expected to have determined the sizes of the first star-clusters and black holes, and the prospect of studying these objects by direct detections with {\\it JWST} and with other instruments.  The direct light emitted by the very first stellar clusters and intermediate--mass black holes at $z>10$ will likely fall below {\\it JWST}'s detection threshold.  However, {\\it JWST} could reveal a decline at the faint--end of the high--redshift luminosity function, and thereby shed light on radiative and other feedback effects that operate at these early epochs. {\\it JWST} will also have the sensitivity to detect individual SNe from beyond $z=10$. In a dedicated survey lasting for several weeks, thousands of SNe could be detected at $z>6$, with a redshift distribution extending to the formation of the very first stars at $z\\gsim 15$.  Using these SNe as tracers may be the only method to map out the earliest stages of the cosmic star--formation history.  Finally, we point out that studying the earliest objects at high redshift will also offer a new window on the primordial power spectrum, on $\\sim 100$ times smaller scales than probed by current large--scale structure data.} ",
        "introduction": "\\label{sec:intro}  The formation of the first astrophysical objects and the subsequent epoch of reionization are at the frontiers of research in Astronomy. Recent years have seen significant progress both in our theoretical understanding and in observational probes of this transition epoch in the early universe.  Through a combination of methods, including measurements of the cosmic microwave background (CMB) anisotropies, culminating in the precise determination of the temperature and polarization power spectra from the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} experiment \\cite{bennett03,spergel03,spergel07,page07,dunkley08,komatsu08}, Hubble diagrams of distant Supernovae \\cite{riess,perlmutter}, and various probes of large scale structure (see, e.g., references in \\cite{bahcall99,tegmark04,slosar07}) the key cosmological parameters have been determined to high accuracy.  Because of the emergence of a concordance ($\\Lambda$CDM) cosmology, we can securely predict the collapse redshifts of the first non--linear dark matter condensations: $2-3\\sigma$ peaks of the primordial density field on mass scales of $10^{5-6}~{\\rm M_\\odot}$, corresponding to the cosmological Jeans mass, collapse at redshifts $z=15-20$.  Of course, in addition to the cosmological parameters describing the average background universe ($\\Lambda$CDM), one also needs a description of the seed fluctuations, to make such predictions. The nearly scale--invariant nature of the initial fluctuation power spectrum has also been empirically confirmed by the above--mentioned large--scale structure data.  However, the predictions for the earliest halos rely on extrapolating the primordial power spectrum to a wavenumber of $k\\sim 10~h{\\rm Mpc^{-1}}$, a scale that is 2--3 orders of magnitude smaller than the smallest scale directly probed by current data (e.g. \\cite{viel07} and references therein).  It is possible that the small--scale power differs substantially from the extrapolated value. For example, in the case of warm dark matter (WDM), the power can be reduced by many orders of magnitude on the relevant scales, and, to zeroth order, the first generation of halos predicted in $\\Lambda$CDM would not exist \\cite{bho01}.  Recent simulations indicate that the formation of the first halos would indeed be delayed \\cite{osheawdm,yoshidawdm}, and the details of the collapse dynamics modified, possibly ultimately affecting the properties of the first stars \\cite{gaotheunswdm}.  Within the $\\Lambda$CDM paradigm, however, robust predictions can be made, using three--dimensional simulations \\cite{yoshida04,springel05}, rather than semi--analytical methods \\cite{ps74,smt01}. Such predictions are now limited mainly by the $\\approx 5\\%$ uncertainty in the normalization of the primordial power spectrum, $\\sigma_{8h^{-1}}$ (e.g. \\cite{slosar07}), which translates essentially to a $\\approx 5\\%$ uncertainty in the collapse redshift of the first structures.  On the observational side, the recent measurement of the optical depth to electron scattering ($\\tau_e=0.09 \\pm 0.03$) by {\\it WMAP} suggests that the first sources of light significantly ionized the intergalactic medium (IGM) at redshift $z\\sim 11\\pm 3$ \\cite{page07,spergel07}.  The Sloan Digital Sky Survey (SDSS) has uncovered a handful of bright quasars at redshifts as high as $z=6.41$ (see \\cite{becker01,fan01,fan01,fan03,fan04,fan06a} or the recent review by \\cite{fan06}).  The luminosity of these sources rivals those of the most powerful quasars at the peak of their activity at $z\\simeq 2.5$.  A sizable population of $z\\sim 6$ galaxies have also been identified in the past several years (see, e.g., \\cite{ellis08} for a recent comprehensive review), many of them in broad--band photometry in deep, wide-area {\\it Hubble Space Telescope} fields \\cite{dickinson04,malhotra05,bouwens06,stanway04,stanway07}. Narrow--band searches have also been successful in discovering numerous $z>5$ Lyman--$\\alpha$ emitting galaxies.  Starting with the initial success of \\cite{weymann98}, 10m class telescopes have now identified about two dozen such objects out to $z\\approx 6.7$ (see \\cite{mr04} for a compilation of datasets; and the more recent Subaru survey results by, e.g., \\cite{taniguchi05,kashikawa06}). In summary, the above observational data leave little doubt that the first galaxies and quasars were indeed well in place by redshift $z\\simeq 6$, but not in significant numbers prior to redshift $z\\simeq 15$.  These developments have been complemented by progress in theoretical work focusing on the cooling and collapse of gas in primordial low--mass halos (see reviews by \\cite{zhcarn,h2review} and references therein). A broad picture has emerged, identifying several processes that are important in the formation of the first stars and black holes (BHs) out of the baryonic gas in the earliest halos.  The key ingredient in this picture is the abundance of ${\\rm H_2}$ molecules that form via gas--phase reactions in the early universe (this beautifully simple fact was recognized as early as in 1967 by \\cite{saslaw}). The first objects form out of gas that cools via ${\\rm H_2}$ molecules, and condenses at the centers of virialized dark matter ``minihalos'' with virial temperatures of $T_{\\rm vir}\\sim 200$K \\cite{htl96,tetal97}. Detailed numerical simulations have shown convergence toward a gas temperature $T\\sim 300~{\\rm K}\\,$ and density $n \\sim 10^{4} \\, {\\rm cm}^{-3}$, dictated by the thermodynamic properties of ${\\rm H_2}$ \\cite{ABN02,BCL02,yoshida03,yoshida08}, which allows the collapse of a single clump of mass $10^{2}-10^{3} \\,{\\rm M_\\odot}$ at the center of the halo. The soft UV and X--ray radiation emitted by the stars and black holes formed in the first handful of such clumps provide prompt and significant global feedback on the chemical and thermal state of the IGM.  A relatively feeble early background radiation field can already have strong effects on gas cooling and ${\\rm H_2}$ chemistry, and can produce significant global impact, affecting the formation of the bulk of the first astrophysical structures, and therefore the reionization history of the IGM \\cite{hrl97,har00,mach01,mach03,oh03,rgs02a,km05,FL05,MBH06,AS07,jetal07}. The observational data summarized above suggests that such feedback processes indeed shaped the reionization history, and hints that star--formation in early minihalos was suppressed \\cite{hb06}. ",
        "conclusions": "\\label{sec:physics}  In the context of $\\Lambda$CDM cosmologies, it is natural to identify the first dark matter halos as the hosts of the first sources of light.  A simplified picture for the emergence of the first sources of light, and consequent reionization is as follows: the gas in dark matter halos cools and turns into ionizing sources (stars and/or black holes), which produce UV (and possibly X--ray) radiation, and drive expanding ionized regions into the IGM.  The volume filling factor of ionized regions then, at least initially, roughly tracks the global formation rate of dark halos.  Eventually, the ionized regions percolate (when the filling factor $F_{\\rm HII}$ reaches unity), and the remaining neutral hydrogen in the IGM is rapidly cleared away as the ionizing background builds up.  However, a soft UV radiation background at photon energies of $<13.6$eV, as well as possibly a soft X--ray background at $\\gsim 1$keV, will build up before reionization, since the early IGM is optically thin at these energies.  We now briefly describe the various effects that should determine the evolution of the first sources and of reionization at high redshifts. For in--depth discussion, we refer the reader to extended reviews (e.g. \\cite{BL01}), and to reviews focusing on the roles of ${\\rm H_2}$ molecules in reionization \\cite{h2review}, on the effect of reionization on CMB anisotropies \\cite{HK99}, and on progress in the last three years in studying reionization \\cite{zhcarn,ferrara07}.  In the discussions that follow, it will be useful to distinguish three different types of dark matter halos, which can be roughly divided into three different ranges of virial temperatures, as follows:  \\begin{tabbing} \\hspace{0.2cm} \\= $300\\,{\\rm K}$   $\\lsim$ \\= $T_{\\rm vir}$ \\=  $\\lsim  10^4\\, {\\rm K}$ \\hspace{1cm} \\= (Type II -- susceptible to ${\\rm H_2}$--feedback)\\\\ \\> $10^4\\,{\\rm K}$  $\\lsim$ \\> $T_{\\rm vir}$ \\>  $\\lsim 2\\times 10^5\\,{\\rm K}$        \\> (Type Ia -- susceptible to photo--heating feedback)\\\\ \\>                          \\> $T_{\\rm vir}$ \\>  $\\gsim 2\\times 10^5\\,{\\rm K}$        \\> (Type Ib -- gas can fall in, even in the face of photo--heating) \\end{tabbing}  We will hereafter refer to these three different types of halos as Type II, Type Ia, and Type Ib halos.  The motivation in distinguishing Type II and Type I halos is based on ${\\rm H_2}$ molecular vs atomic H cooling, whereas types ``a'' and ``b'' reflect the ability of halos to allow infall and cooling of photoionized gas.  While the actual values of the critical temperatures, and the sharpness of the transition from one category to the next is uncertain (see discussion below), as we argue below, star formation in each type of halo is likely governed by different physics, and each type of halo likely plays a different role in the reionization history.  In short, Type II halos can host ionizing sources only in the neutral regions of the IGM, and only if ${\\rm H_2}$ molecules are present; Type Ia halos can form new ionizing sources only in the neutral IGM regions, but can do so by atomic cooling, irrespective of the ${\\rm H_2}$ abundance, and Type Ib halos can form ionizing sources regardless of the ${\\rm H_2}$ abundance, and whether they are in the ionized or neutral phase of the IGM.  It is useful to next provide a summary of the various important physical effects that govern star--formation and feedback in these halos.  \\begin{table*}[t]\\small \\caption{\\label{table:collapse} Collapse redshifts and masses of 1, 2 and 3$\\sigma$ dark matter halos with virial temperatures of $T_{\\rm vir}=100$ and $10^4$K, reflecting the minimum gas temperatures required for cooling by ${\\rm H_2}$ and neutral atomic H, respectively.}  \\begin{center} \\begin{tabular}{lrrr} $T_{\\rm vir}$(K) & $\\nu$ & $z_{\\rm coll}$ & $M_{\\rm halo}({\\rm M_\\odot})$ \\\\ \\hline 100    & 1   &  7    & $2\\times10^5$ \\\\ & 2   &  16   & $5\\times10^4$ \\\\ & 3   &  26   & $3\\times10^4$ \\\\ $10^4$ & 1   &  3.5  & $4\\times10^8$ \\\\ & 2   &  9    & $1\\times10^8$ \\\\ & 3   &  15   & $6\\times10^7$ \\\\ \\hline \\multicolumn{4}{l}{} \\\\ \\end{tabular}\\\\[12pt] \\end{center} \\end{table*}  {\\em ${\\rm H_2}$ Molecule Formation.} In $\\Lambda$CDM cosmologies, structure formation is bottom--up: the earliest nonlinear dark matter halos form at low masses.  Gas contracts together with the dark matter only in dark halos above the cosmological Jeans mass, $M_{\\rm J}\\approx 10^4$. However, this gas can only cool and contract to high densities in somewhat more massive halos, with $M\\gsim M_{\\rm H2} \\equiv 10^5\\msun[(1+z)/11]^{-3/2}$ (i.e. Type II halos), and only provided that there is a sufficient abundance of ${\\rm H_2}$ molecules, with a relative number fraction at least $n_{\\rm H2}/n_{\\rm H}\\sim 10^{-3}$ \\cite{htl96,tetal97}.  Because the typical collapse redshift of halos is a strong function of their size, the abundance of ${\\rm H_2}$ molecules is potentially the most important parameter in determining the onset of reionization.  For example, 2$\\sigma$ halos with virial temperatures of $100$K appear at $z=16$, while 2$\\sigma$ halos with virial temperatures of $10^4$K (Type Ia halos, in which gas cooling is enabled by atomic hydrogen lines) appear only at $z=9$. Table~\\ref{table:collapse} summarizes the collapse redshifts and masses of dark halos at these two virial temperatures in the concordance cosmology.  As a result, the presence or absence of ${\\rm H_2}$ in Type II halos makes a factor of $\\sim$2 difference in the redshift for the onset of structure formation, and thus potentially effects the reionization redshift and the electron scattering opacity $\\tau$ by a similar factor. In the absence of any feedback processes, the gas collecting in Type II halos is expected to be able to form the requisite amount of ${\\rm H_2}$ \\cite{htl96,tetal97}.  However, both internal and external feedback processes can alter the typical ${\\rm H_2}$ abundance (see discussion below).  {\\em The Nature of the Light Sources in the First Halos.} A significant uncertainty is the nature of the ionizing sources turning on inside halos collapsing at the highest redshifts. Three dimensional simulations using adaptive mesh refinement (AMR; \\cite{ABN02}) and smooth particle hydrodynamics (SPH; \\cite{BCL02,yoshida08}) techniques have followed the contraction of gas in Type II halos at high redshifts to high densities. These works have shown convergence toward a temperature/density regime of ${\\rm T \\sim 200~{\\rm K}\\,}$, ${\\rm n \\sim 10^{4} \\, {\\rm cm}^{-3}}$, dictated by the critical density at which the excited states of ${\\rm H_2}$ reach equilibrium population levels.  These results have suggested that the first Type II halos can only form unusually massive stars, with masses of at least $\\sim 100 \\, {\\rm M_\\odot}$, and only from a small fraction ($\\lsim 0.01$) of the available gas.  Such massive stars are effective producers of ionizing radiation, making reionization easier to achieve.  In addition, it is expected that the first massive stars, forming out of metal--free gas, have unusually hard spectra \\cite{TS2000,BKL01,schaerer02}, which is important for possibly ionizing helium, in addition to hydrogen.  An alternative possibility is that a similar fraction of the gas in the first Type II halos forms massive black holes \\cite{hl97,madauetal04,rog05}.  In fact, when massive, non--rotating metal--free stars end their life, they are expected to leave behind stellar--mass seed BHs, unless their mass is in the range of 140-260 ${\\rm M_\\odot}$ \\cite{Heger03}. These early black holes can then accrete gas (possibly only after a delay -- if the progenitor star clears the host halo of gas, the BH will have to await until the host halo merges with another, gas--rich halo), acting as ``miniquasars''. The miniquasars will produce a hard spectrum extending to the soft X--rays, which could be important in catalyzing ${\\rm H_2}$ formation globally (see discussion below).  {\\em Efficiencies and the Transition from Metal Free to ``Normal'' Stars.} Another fundamental question of interest is the efficiency at which the first sources inject ionizing photons into the IGM.  This can be parameterized by the product $\\epsilon_* \\equiv N_\\gamma f_* f_{\\rm esc}$, where $f_* \\equiv M_*/(\\Omega_{\\rm b}M_{\\rm halo}/\\Omega_m)$ is the fraction of baryons in the halo that turns into stars; $N_\\gamma$ is the mean number of ionizing photons produced by an atom cycled through stars, averaged over the initial mass function (IMF) of the stars; and $f_{\\rm esc}$ is the fraction of these ionizing photons that escapes into the IGM.  It is difficult to estimate these quantities at high--redshifts from first principles, but the discussion below can serve as a useful guide.  Although the majority of the baryonic mass in the local universe has been turned into stars \\cite{fukugita98}, the global star formation efficiency at high redshifts was likely lower.  To explain a universal carbon enrichment of the IGM to a level of $10^{-2}-10^{-3}~{\\rm Z_\\odot}$, the required efficiency, averaged over all halos at $z\\gsim 4$, is $f_*=2-20\\%$ \\cite{hl97}. However, the numerical simulations mentioned in \\S~2.2 above suggest that the fraction of gas turned into massive stars in Type II halos is $f_*\\lsim 1\\%$.  The escape fraction of ionizing radiation in local starburst galaxies is of order $\\sim 10\\%$.  The higher characteristic densities at higher redshifts could decrease this value \\cite{dsf00,wl00}, although there are empirical indications that the escape fraction in $z\\sim3$ galaxies may instead be higher \\cite{spa01}, at least in some galaxies \\cite{newfesc}. Radiation can also more readily ionize the local gas, and escape from the small Type II halos which have relatively low total hydrogen column densities ($\\lsim 10^{17}~{\\rm cm^{-2}}$), effectively with $f_{\\rm esc}=1$ in the smallest halos \\cite{wan04}.  The ionizing photon yield per proton for a normal Salpeter IMF is $N_\\gamma\\approx 4000$. However, if the IMF consists exclusively of massive $M\\gsim 200\\,{\\rm M_\\odot}$ metal--free stars, then $N_\\gamma$ can be up to a factor of $\\sim 20$ higher \\cite{BKL01,schaerer02}. The transition from metal--free to a ``normal'' stellar population is thought to occur at a critical metallicity of $Z_{\\rm cr}\\sim 5\\times 10^{-4}{\\rm Z_\\odot}$, above which cooling and fragmentation becomes efficient and stops the IMF from being biased toward massive stars \\cite{BKL01}.  It is natural to associate this transition with that of the assembly of halos with virial temperatures of $>10^4$K (Type Ia halos). Type II halos are fragile, and likely blow away their gas and ``shut themselves off'' after a single episode of (metal--free) star--formation. They are therefore unlikely to allow continued formation of stars with metallicities above $Z_{\\rm crit}$. Subsequent star--formation will then occur only when the deeper potential wells of Type Ia halos are assembled and cool their gas via atomic hydrogen lines. The material that collects in these halos will then have already gone through a Type II halo phase and contain traces of metals.  As we argue in \\S~\\ref{sec:feedback} below, there exists an alternative, equally plausible scenario.  Most Type II halos may not have formed ${\\rm any}$ stars, due to global ${\\rm H_2}$ photodissociation by an early cosmic soft--UV background.  In this case, the first generation of metal--free stars must appear in Type Ia halos. Halos above this threshold can eject most of their self--produced metals into the IGM, but, in difference from Type II halos, can retain most of their gas \\cite{maclowferrara99}, and can have significant episodes of metal--free star formation.  These halos will also start the process of reionization by driving expanding ionization fronts into the IGM. The metals that are ejected from Type Ia halos will reside in these photoionized regions of the IGM. As discussed in \\S~\\ref{sec:feedback} below, photoionization heating in these regions may suppress gas infall and cooling, causing a pause in the formation of new structures, until larger dark matter halos, with virial temperatures of $T_{\\rm vir}\\gsim 2\\times 10^5$K (Type Ib halos) are assembled.  The material that collects in Type Ib halos will then have already gone through a previous phase of metal--enrichment by Type Ia halos, and it is unlikely that Type Ib halos can form significant numbers of metal--free stars."
    },
    "0809/0809.1865_arXiv.txt": {
        "abstract": "{The presence of titanium oxide (TiO) and vanadium oxide (VO) gas phase species is searched for in the atmosphere of the hot Jupiter \\hd.} {We compared a model for the planet's transmitted spectrum to multi-wavelength eclipse-depth measurements (from 3\\,000 to 10\\,000~\\AA), obtained by Sing et al. (2008a) using archived HST-STIS time series spectra. We make use of these observations to search for spectral signatures from extra absorbers in the planet atmosphere between 6\\,000 and 8\\,000~\\AA.} {Along with sodium depletion and Rayleigh scattering recently published for this exoplanet atmosphere, an extra absorber of uncertain origin, redward of the sodium lines, resides in the atmosphere of the planet. Furthermore, this planet has a stratosphere experiencing a thermal inversion caused by the capture of optical stellar flux by absorbers that resides at altitude. Recent models have predicted that the presence of TiO and VO in the atmosphere of \\hd\\ may be responsible for this temperature inversion. Although no specific TiO and VO spectral band head signatures have been identified unambiguously in the observed spectrum, we suggest here that the opacities of those molecules are possible candidates to explain the remaining continuous broad band absorption observed between 6\\,200 and 8\\,000~\\AA. To match reasonably well the data, the abundances of TiO and VO molecules are evaluated from ten to one thousand times below solar. This upper limit result is in agreement with expected variations with altitude due to depletion effects such as condensation.}  {} ",
        "introduction": "\\label{sec:intro}  The discovery of transiting extrasolar giant planets (EGP) has opened the window to the direct detections and characterization of their atmospheres. Because of the wavelength-dependent opacities of absorbing species, measurement of relative changes in eclipse depth as a function of wavelength during primary transit has the potential to reveal the presence (or absence) of specific chemical species (\\cite{Seager2000a}; \\cite{Hubbard01}; \\cite{Brown01}). In the case of \\hd's atmosphere, the transmission spectroscopy method led  to the detection of sodium (Charbonneau et~al. 2002). The small atmospheric sodium signature has made its detection from ground-based telescopes difficult. Only recently, Redfield et al. (2008) and Snellen et al. (2008) reported the detection of Na\\,{\\sc i} in the atmosphere of HD189733b and HD209458b, respectively. Absorptions of several percents for H\\,{\\sc i} Lyman-$\\alpha$, O\\,{\\sc i} and C\\,{\\sc ii} have been measured in the hydrodynamically escaping upper atmosphere (\\cite{Madjar03}, 2004, 2008). More recently, Rayleigh scattering by H$_2$ molecules has been identified (Lecavelier et~al. 2008b) and a temperature-pressure (TP) profile with inversion was derived (Sing et~al. 2008b). This temperature inversion leads to high temperature at both low and high pressure. Note that this temperature bifurcation was very well predicted by atmospheric models of strongly irradiated planets (\\cite{Hubeny03}). In the lower part of the atmosphere ($\\sim$30 mbar), the temperature is found to be in the range $1\\,900 -2\\,400$~K (Lecavelier et~al. 2008b), which corresponds to the M/L/T brown dwarf regime, as expected for a hot Jupiter such as \\hd\\ (\\cite{Kirkpatrick05}).   In the cool atmosphere of sub-stellar objects like brown dwarfs and close-in EGPs, the strongest absorption features are expected to be those of alkali metals like sodium and potassium (\\cite{Seager2000a}; \\cite{Sudarsky00}; \\cite{Burrows00}). This is due to the condensation and the rain out of elements in the atmosphere which clears the atmosphere of most of the metals and keeps the less refractory alkali metals (\\cite{Burrows99}; \\cite{Lodders99}; \\cite{Burrows02}; \\cite{Sudarsky03}). Thus, in contrary to the spectrum of HD189733b (Pont et al. 2007; Lecavelier des Etangs et al. 2008a), the spectrum of \\hd\\ should be dominated by absorption from sodium Na\\,{\\sc i} and potassium K\\,{\\sc i}.   However, depending on the effective temperature, a large number of diatomic and polyatomic molecules are predicted to be present according to various models of brown dwarf and hot Jupiter (\\cite{Burrows99}; \\cite{Lodders99}; \\cite{Allard01}; \\cite{Lodders02}; \\cite{Hubeny03}; \\cite{Burrows06}). Among those molecules and at a temperature above 1\\,800 K, titanium oxide (TiO) and vanadium oxide (VO) in gas phase equilibrium are most probably present with a high abundance in strongly irradiated planet atmospheres (Seager et al. 1998; \\cite{Hubeny03}; \\cite{Fortney07}). Furthermore, the low albedo measurements of \\hd\\ (\\cite{Rowe06}) rule out most of the absorbents, but TiO and VO. Recently, the interpretation of visible observations obtained during the transit of this planet (Sing et al. 2008a), as well as the interpretation of near infrared observations taken during its secondary eclipse (Knutson et al. 2008) require that HD 209458b experience an inversion in the atmospheric \\TP\\ and a have stratosphere (Sing et al. 2008b; Burrows et al. 2007c; Burrows et al. 2008). Such an inversion could be due to the absorption of the incident visible light by TiO and VO molecules (Fortney et al. 2007). Ongoing models (Burrows et al. 2008) claim that this inversion must be caused by the capture of the incident optical stellar flux by a stratospheric absorber that could potentially be TiO and VO. Similarly other studies (\\cite{Hubeny03}; \\cite{Fortney07}) have more particularly investigated the effects of TiO and VO which absorb much of the incoming stellar radiation high in the atmospheres of the hottest CEGPs.  In this work, we use {\\it STIS} spectra obtained during planetary transit at two spectral resolutions (low and medium) to search for the direct evidence of the presence of TiO and VO in the planet's atmosphere. Both datasets are combined to extend the measurements over the entire optical regime to quantify other possible absorbers appearing in the transmission spectrum. Rayleigh scattering and sodium absorption have been proposed to explain the spectral features observed between 3\\,000~\\AA\\ and 6\\,000~\\AA\\ (Lecavelier des Etangs et al. 2008b; Sing et al. 2008a; 2008b) yet, additional absorptions remain unattributed from 6\\,100~\\AA\\ up to 8\\,000~\\AA. These absorptions cannot be due to optically thick high altitude clouds, which would otherwise mask the detected signature of Rayleigh scattering (Lecavelier des Etangs et al. 2008b). Since TiO and VO molecules are expected to be present in high abundance for the range of temperatures and pressures for strongly irradiated hot Jupiter, their opacities could largely contribute to the observed absorption in this wavelength domain (Seager et al.1998; \\cite{Burrows99}; \\cite{Sharp07}).   After a brief description of the observation and interpretation status (Sect.~\\ref{sec:obs}), we present, in Sect.~\\ref{sec:model}, the model together with the method used to calculate the atomic and molecular opacities. Finally, we estimate the contribution of  TiO and VO to the observed transmission spectrum for several abundance scenarios (Sect. ~\\ref{sec:analysis}) and discuss the consequences on the atmospheric chemistry and aeronomy of HD209458b (Sect.~\\ref{sec:discuss}). ",
        "conclusions": "\\label{sec:discuss}  We have highlighted that the \\hd\\ atmosphere is optically thick at low pressures and this requires new absorbers between 6\\,000-8\\,000~\\AA. The observation of the Na\\,{\\sc i} in the atmosphere exclude the presence of clouds at pressure under $30$ mbar in the observed limb of the planet. We have shown that Li\\,{\\sc i} and K\\,{\\sc i} absorption cannot reproduce the broad band absorptions observed. Nonetheless, we derive upper limits for their abundances at $2\\times 10^{-3}$ solar for K\\,{\\sc i} and $2\\times 10^{-1}$ for Li\\,{\\sc i}.   The Barman (2007) analysis of the Knuston et al. (2007) results and Sing et al. (2008b) analysis both agree that remaining absorbents are required to explain the continuous broad band absorption observed between 6\\,200 and 8\\,000~\\AA, since the radii derived in both reductions are significantly above the model with no TiO nor VO molecules (see Fig.~\\ref{fig_atoms}). As a consequence, the presence of remaining absorbents, possibly TiO and VO, is mandatory to explain the observations. This result does not depend on the \\TP\\ adopted. The abundances of TiO and VO are only poorly constrained by any \\TP. We decided to use the \\TP\\ derived by Sing et al. (2008) as an example.  In that case, the lower part of the \\TP, i.e below the condensation curves, is the most critical part for the determination of the TiO and VO abundances. The abundances of TiO and VO molecules are mainly constrained by the level of the observed AD curve between 6\\,500 and 8\\,000~\\AA\\ and by the pressure of the lower point. Indeed, for this lower point, the TiO and VO lines are mostly pressure broadened. Thus, the abundances of these species are mainly sensitive to the pressure and not to the temperature in the lower part of this profile. The \\TP\\ above the condensation curves constrains mostly the spectral features, especially the 6\\,250~\\AA\\ one, as seen for the case two level of TiO and VO. However, we do not detect any TiO/VO narrow spectrale features, thus any \\TP\\ above the condensation curve can be adopted. With the \\TP\\ used in this study, we have shown that no clear solutions emerge wether condensation is taken into account, or if two levels of TiO and VO molecules are considered (see Table~\\ref{table:Chi2}). Thus, other profiles could led to the same conclusion for the presence of TiO and VO and/or of potential additional absorbers.   Although no typical TiO and VO spectral signatures have been identified unambiguously in the observed spectrum, we suggest that the opacities of those molecules are the best candidates to explain the remaining continuous broad band absorption observed between 6\\,200 and 8\\,000~\\AA. Theoretical absorption with models including TiO/VO with abundances below solar were evaluated to match the data reasonably well. Using the \\TP\\ from \\cite{Sing08b}, we derived upper limits for the TiO and VO abundances. The model without TiO or with TiO but no VO molecules give the worst $\\chi_{2}$ solutions. We found that the abundance of TiO should be around $10^{-4}$ to $10^{-3}$ solar, and  the abundance  VO around $10^{-3}$ to $10^{-2}$ solar. The fits become marginally worse on the blue side of the sodium wing when adding TiO and VO (see Fig.~\\ref{fig_molecules}). This is because the \\TP\\ used in our work was obtained by fitting the observation on the blue side of the sodium line with a model which contains only Rayleigh scattering and sodium opactity (Sing et al. 2008b) which helped limit the number of free parameters in the fit, such that a convergent solution could be found easilly. This model fits precisely the observed data for this part of the spectrum (see the dotted line in Figs.~\\ref{fig_atoms} and ~\\ref{fig_molecules}). The addition of the TiO and VO opacities in this model, contributes to slightly decrease the quality of the fit on the blue side of the sodium line (see continous line in Fig.~\\ref{fig_molecules}). This is, however, of marginal consequence on the TiO and VO evaluations in our study since these are mostly controlled by the main part of the observations in the redder range of the spectrum. Considering the quality of the observations this does not affect the TiO and VO evaluations. Hence, the use of the \\TP\\ with only Rayleigh and sodium calculated by Sing et al.2008 is sufficient to derive reasonably good estimates of the TiO and VO abundances.   Two layers of TiO/VO in the atmosphere are better suited to explain the 6\\,250~\\AA\\ spectral feature, however the $\\chi^2$ is not significantly improved (Table~\\ref{table:Chi2}). Note that, the 6\\,250~\\AA\\ spectral feature is seen in the two independent spectra obtained at both low resolution, as mentioned here, and at medium resolution as observed by \\cite{Sing08a} in the G750M STIS spectra.   The tholins, polyacetylenes, or various non-equilibrium compounds at high-altitude  could also be responsible for the  temperature inversion. However, the study of these molecules requires a full non-equilibrium chemical model which is not the purpose of that work. TiO and VO molecules are very good high-altitude absorber candidates (\\cite{Hubeny03}; \\cite{Fortney07}). Fortney et al.\\ (2008) highlighted the importance of gaseous TiO and VO opacity in their model of highly irradiated close-in giant planets. They define two classes of irradiated atmospheres. Those which are warm enough to have a strong opacity due to TiO and VO gases (``pM Class'' planets), and those that are cooler (``pL Class'' planets) dominated by Na\\,{\\sc i} and K\\,{\\sc i}. Our possible detection of TiO/VO and Na\\,{\\sc i} in the hot atmosphere of \\hd\\ confirm that this planet is located in the transition region between the two classes defined by these authors. As observed here, due to the presence of TiO/VO, \\hd's atmosphere absorbs incident energy between 6\\,500~\\AA\\ and 8\\,000~\\AA. Consequently, this planet has a hot stratosphere (around 2\\,500~K) with a temperature inversion (\\cite{Burrows07c}). As another consequence, because of the thermal emission of the energy trapped by the TiO/VO absorption, the planet appears very bright during the mid infrared secondary eclipse (\\cite{Knutson08}).  The first titanium condensates appearing at high temperature and high pressure are Ti$_3$O$_5$ and CaTiO$_3$ (\\cite{Burrows99}; \\cite{Lodders02}). The condensation of vanadium starts at lower temperature (1\\,600 K) than for titanium (1\\,800 K). The vanadium condenses into solid VO and then into V$_2$O$_3$. Thus, depletion of TiO should start at lower altitude than for VO. Below 1\\,600K, TiO, VO and major refractory elements are absent, leaving monoatomic Na\\,{\\sc i} and K\\,{\\sc i} to dominate the spectrum. Nevertheless, sodium and potassium chlorides become increasingly abundant with decreasing temperature, especially for KCl which is the dominant K-bearing compound. Condensed potassium depletes the atmosphere of atomic K\\,{\\sc i}, as seen with Na\\,{\\sc i}, leading to reduced signatures. This could explain why Na\\,{\\sc i} is seen in abundance in the low resolution data, but not K\\,{\\sc i}. In the lower warmer atmosphere, where wide atomic K line wings would be observable even at low resolution, TiO and VO likely make up the surrounding continuum further masking the signature.  An alternative \\TP\\ with Na ionization has been also proposed (Sing et al. 2008b). A model with Na ionization leaves a wide range of temperatures possible in the middle atmosphere. Within the framework of equilibrium chemistry, this profile presents a lower temperature gradient \\textbf{(Fortney et al. 2003)}. The minimum temperature is below the TiO/VO condensation curves, and above the sodium condensation curve. We did not enter into a detailed study of TiO/VO abundances with these \\TP; however, we note that these profiles would have the effect to increase the abundance of TiO and VO molecules in the middle atmosphere. See Sing et al. (2008b) for more details about ionization.    If present in the upper part of the atmosphere, TiO/VO molecules are not fully depleted by condensation. Thermochemical equilibrium calculations with rainout (\\cite{Burrows99}, \\cite{Lodders02}, \\cite{Hubeny03}) have shown that TiO and VO can exist at high-$T$/low-$P$ points such as in the upper part of our \\TP. The presence of TiO and VO at high-$T$/high-$P$ points leads to a situation in which there is two levels of TiO/VO. This has been already proposed by Hubeny et~al. (2003) to explain the presence of a temperature inversion in strongly irradiated planets atmosphere's. Since the \\TP\\ used here crosses twice the Ti and V condensation curves, two levels containing TiO/VO molecules in gas phase could be separated by a middle level that is free of those molecules, as proposed by Hubeny et al. 2003. Nonetheless, a cold trap region is usually expected to deplete the upper low-$P$ region. Flushing out the upper atmosphere of TiO/VO would rule out 6\\,250~\\AA\\ feature as being due to TiO. However, for now, it has not been proven that cold-trap effect would deplete those molecules, especially for the hydrodynamically escaping type of upper atmosphere (\\cite{Madjar03}, 2004, 2008; Lecavelier 2004; Lecavelier et al. 2007; \\cite{Ehrenreich08}).  New observations with higher signal-to-noise ratio and better resolution, together with improved chemical models, are required to address TiO and VO spectral absorption features and detect specific band heads.         \\newpage \\begin{figure*} \\includegraphics[width=12cm]{fig_Atoms.ps}   \\caption{The low resolution STIS measurements of the planetary transit absorption depth (AD) corrected from limb-darkenning effects and binned by 60 pixels (histogram). The observed 1$\\sigma$ errorbar is plotted above the spectrum. The radii derived by Knutson et al.(2007) with the corresponding error bars and spectrale bins are plotted (squares). The dashed line corresponds to the best fit model assuming Rayleigh scattering and sodium absorption with the physical T-P profile plotted in Fig.~\\ref{fig_TP}. The difference between the dashed curve and the observed AD spectrum for wavelength over 6\\,000~\\AA\\ indicates that remaning absorbents should be presents in this spectral domain. Overplotted in continuous line, is the same model, with atomic lithium and potassium with solar abundances.This model cannot reproduce the observations above 6\\,200~\\AA.} \\label{fig_atoms} \\end{figure*} \\newpage    \\newpage \\begin{figure*} \\centering \\includegraphics[angle=90,width=10cm]{TPprofile_escape.ps} \\caption{The atmospheric Temperature-Pressure profile (\\TP) with error bars derived by \\cite{Sing08b} used in this study. The dashed lines correspond to the condensation curves for titanium, vanadium and sodium. The three points of this T-P profile are derived from the fit of the observed absorption depth curve. The hot point at a pressure of 0.05~Bar is imposed by the Rayleigh scattering (See Lecavelier et al. 2008).} \\label{fig_TP}% \\end{figure*}   \\newpage \\begin{figure*} \\centering \\includegraphics[angle=90,width=10cm]{sigma_TiO_VO.ps}  \\caption{The $\\log_{10}$ of the monochromatic cross-section $\\sigma$ (cm$^2$) as a function of wavelength for the vibration-rotation transitions of TiO and VO. The contribution due to different isotopes is included. TiO has a strong absorption feature shortward of 7\\,500~\\AA, and has a strong peak near of 6\\,200~\\AA\\ as observed in Fig.~\\ref{fig_atoms}.} \\label{fig_sigma}% \\end{figure*}   \\newpage \\begin{figure*} \\includegraphics[width=17cm]{fig_TiO_VO_Na_Rayleigh_VariousAbund_style4.eps} \\newline \\caption{Same as Fig.~\\ref{fig_atoms} with fits obtain with various models (continuous thick line). The observed 1$\\sigma$ error bars are plotted around 7\\,800~\\AA\\ (see also Fig.10 of Sing et al. 2008). The dashed line shows the fit obtained using only the model with Rayleigh scattering and sodium absorption with the T-P profile (Fig.~\\ref{fig_TP}). A remaining  broad band absorption appears in between 6\\,200 to 8\\,000~\\AA\\ which could be attributed to TiO and VO opacities. a: Model with a constant mixing ratio of TiO. A constant solar abundance of TiO is excluded (upper dashed-dotted line). The best fit gives an abundance of $(8\\pm0.6)\\times 10^{-4}$ solar (red solid line). b: Model with a constant mixing ratio of TiO and VO. Solar abundance for VO is also excluded. The red solid line is the best fit which includes TiO and VO. c: TiO, VO, and condensation. The altitudes of condensation for titanium and for vanadium are both plotted (dash dotted lines) together with the best fit (red solid line). d: Best fit for two distinct levels containing TiO and VO (See Sect.~\\ref{sec:TiO_VO_cond2}). No clear solution emerge in between the three last models (red solid line).} \\label{fig_molecules} \\end{figure*} \\newpage"
    },
    "0809/0809.3784_arXiv.txt": {
        "abstract": "{We examine the X-ray luminosity scaling relations of 31 nearby galaxy clusters from the Representative {\\it XMM-Newton\\/} Cluster Structure Survey (\\rexcess). The objects are selected only in X-ray luminosity, optimally sampling the cluster luminosity function. Temperatures range from 2 to 9 keV, and there is no bias toward any particular morphological type. To reduce measurement scatter we extract pertinent values in an aperture corresponding to $\\Rv$, estimated using the tight correlation between $Y_X$ (the product of gas mass and temperature) and total mass. The data exhibit power law relations between bolometric X-ray luminosity and temperature, $Y_X$ and total mass, all with slopes that are significantly steeper than self-similar expectations. We examine the possible causes for the steepening, finding that structural variations have little effect and that the primary driver appears to be a systematic variation of the gas content with mass. Scatter about the relations is dominated in all cases by the presence of cool cores. The natural logarithmic scatter about the raw X-ray luminosity-temperature relation is about 70 per cent, and about the X-ray luminosity-$Y_X$ relation it is 40 per cent. Systems with more morphological substructure show similar scatter about scaling relations than clusters with less substructure, due to the preponderance of cool core systems in the regular cluster subsample. Cool core and morphologically disturbed systems occupy distinct regions in the residual space with respect to the best fitting mean relation, the former lying systematically at the high luminosity side, the latter lying systematically at the low luminosity side. Simple exclusion of the central regions serves to reduce the scatter about the scaling relations by more than a factor of two. The scatter reduces by a similar amount with the use of the central gas density as a third parameter. Using $Y_X$ as a total mass proxy, we derive a Malmquist bias-corrected local luminosity-mass relation and compare with other recent determinations. Our results indicate that luminosity can be a reliable mass proxy with controllable scatter, which has important implications for upcoming all-sky cluster surveys, such as those to be undertaken with {\\it Planck} and {\\it eROSITA}, and ultimately for the use of the cluster population for cosmological purposes. } ",
        "introduction": "The X-ray luminosity is an observationally attractive quantity because of the relative ease with which it can be measured, and thus it is a key parameter for cosmological applications of the galaxy cluster population. For a fully virialised cluster formed through pure gravitational collapse, the X-ray luminosity $L$ is determined solely by the mass and distribution of gas in the intracluster medium (ICM), and the X-ray temperature $T$ is determined by the depth of the potential well in which the ICM rests. Correlations between these two basic quantities were found in the very early days of X-ray observations of clusters, even while the thermal nature of the emission was still under debate \\citep{mitchell77,mushotzky78,ht79}. Initial results from these works suggested that the slope of the luminosity temperature relation was steeper than expected from gravitational collapse alone. The launch of the {\\it EXOSAT} and {\\it Einstein} observatories enabled the first systematic studies of large samples of clusters \\citep{es91,david93}, and the subsequent launch of {\\it ROSAT, ASCA} and {\\it Ginga} allowed further investigation with increasingly better quality data.  The density squared ($n_e^2$) dependence of the X-ray emission means that luminosity measurements are very sensitive to the exact physics of the gas near the cluster core. Mechanisms such as rapid radiative cooling or merging can change the thermodynamic state of this core gas, introducing scatter into the various luminosity scaling relations. Since our knowledge of the absolute extent of the scatter limits the constraints that can be put on cosmological models with the cluster population, the magnitude of the scatter, its source(s), and how to correct for it constitute some of the most important open issues in the study of clusters \\citep[see e.g.][]{lh05}. {\\citet{fabian94} were the first to note that the offset of a cluster from the mean relation was connected to the presence of a cool core, motivating examination of methods to correct for this effect. \\citet{mark98} derived quantities corrected for the presence of cooling cores by exclusion of the emission from the central region and the introduction of a second spectral component in the temperature estimation; \\citet{ae99} determined the luminosity temperature relation using clusters specifically chosen to have weak or non-existent cool cores. More recent efforts have aimed at reducing scatter by using the peak surface brightness as a third parameter \\citep{ohara06}. At the same time it had long been suspected that merging events also contributed to scatter about the mean relation. This effect has been  investigated with increasingly sophisticated numerical simulations, which have shown that while major mergers can indeed boost both luminosity and temperature, the boosting appears to be short lived and the net movement in the luminosity temperature plane is approximately parallel to the mean relation \\citep{rs01,rt02, hart07}.  In the present paper we re-investigate the luminosity scaling relations with \\rexcess\\ \\citep{boehringer07}, a sample of 33 local ($z < 0.2$) clusters drawn from the REFLEX catalogue \\citep{reflex}, all of which have been observed with a single satellite, \\xmm, with the aim of minimising effects due to instrumental cross calibration uncertainties. The unique sample selection strategy, in which clusters have been selected by luminosity only, in such a way as to have close to homogeneous coverage of luminosity space, delivers an optimal sampling of the luminosity function of the cluster population with no bias towards any particular morphological type. Moreover, distances were optimised so that the angular scale of the objects is such that $\\sim R_{500}$ falls well within the \\xmm\\ field of view, allowing detailed local background modelling to be undertaken, increasing the precision of measurements at large radii as compared to more nearby clusters which often fill the field of view. Since the basic selection criterion is X-ray luminosity, \\rexcess\\ should be representative of any local, unbiased high-quality X-ray survey, of the type applicable to testing of cosmological models.  In the following, we first use the representative nature of the \\rexcess\\ sample to investigate the raw luminosity scaling relations, finding that the slopes are steeper than expected if the gas is heated purely by gravitational processes and that cooling cores are the dominant contributor to scatter about them. Dividing the data into subsamples, we investigate the effect of cool cores and morphological disturbance on a cluster's position with respect to the mean relation, finding that the former lie systematically at the high luminosity side, and the latter lie systematically at the low luminosity side. We then investigate two different methods to minimise scatter: simple exclusion of the central region, and use of the central gas density as a third parameter in scaling law fitting, finding that both methods result in a significant reduction in the dispersion about the best fitting relations. Lastly, we examine the physical causes of the steep slope of the luminosity scaling relations, concluding that variations of gas content with total mass are most likely the dominant reason why these are steeper than expected. Our Appendix details the \\rexcess\\ survey volume calculations and a first attempt at correcting for Malmquist bias in the luminosity-mass relation.  \\begin{table*} \\begin{center} \\caption{{\\footnotesize Cluster properties. The luminosity is the bolometric [0.01-100] keV luminosity.  All quantities are calculated assuming $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$, and $h_0 = 0.7$.}}\\label {tab:tab1} \\centering \\begin{tabular}{l l l r l r l r r r c c} \\hline \\hline  \\multicolumn{1}{c}{Cluster} & \\multicolumn{1}{c}{$z$} & \\multicolumn{1}{c}{$T_1$} & \\multicolumn{1}{c}{$L_1$} & \\multicolumn{1}{c}{$T_2$} & \\multicolumn{1}{c}{$L_2$} & \\multicolumn{1}{c}{$T_3$} & \\multicolumn{1}{c}{$Y_X$} & \\multicolumn{1}{c}{$\\Rv$} & \\multicolumn{1}{c}{$R_{\\rm det}$} & \\multicolumn{1}{c}{CC} & \\multicolumn{1}{c}{Disturbed} \\\\  \\multicolumn{1}{c }{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c }{(5)} & \\multicolumn{1}{c}{ (6)} & \\multicolumn{1}{c }{(7)} & \\multicolumn{1}{c}{(8)} & \\multicolumn{1}{c}{(9)} & \\multicolumn{1}{c}{(10) } & \\multicolumn{1}{c}{(11) } & \\multicolumn{1}{c}{(12) } \\\\  \\hline \\\\  RXC\\,J0003.8+0203 & 0.0924 & $3.85_{-0.09}^{+0.09}$ & $ 1.88_{-0.01}^{+0.01}$ & $3.64_{-0.09}^{+0.09}$ & $ 1.16_{-0.01}^{+0.01}$ & $3.87_{-0.10}^{+0.10}$ & $ 7.69_{-0.26}^{+0.26}$ &  876.7 & 0.84 & \\ldots & \\ldots \\\\ RXC\\,J0006.0-3443 & 0.1147 & $5.03_{-0.19}^{+0.19}$ & $ 4.13_{-0.05}^{+0.05}$ & $4.60_{-0.16}^{+0.21}$ & $ 3.18_{-0.05}^{+0.05}$ & $5.18_{-0.20}^{+0.20}$ & $22.74_{-1.21}^{+1.22}$ & 1059.3 & 0.93 & \\ldots & \\checkmark \\\\ RXC\\,J0020.7-2542 & 0.1410 & $5.69_{-0.11}^{+0.11}$ & $ 6.52_{-0.04}^{+0.04}$ & $5.24_{-0.15}^{+0.15}$ & $ 4.07_{-0.04}^{+0.04}$ & $5.55_{-0.13}^{+0.13}$ & $22.41_{-0.63}^{+0.65}$ & 1045.3 & 1.07 & \\ldots & \\ldots\\\\ RXC\\,J0049.4-2931 & 0.1084 & $3.09_{-0.10}^{+0.10}$ & $ 1.78_{-0.02}^{+0.02}$ & $2.79_{-0.11}^{+0.11}$ & $ 1.00_{-0.02}^{+0.02}$ & $3.03_{-0.12}^{+0.12}$ & $ 5.09_{-0.24}^{+0.25}$ &  807.8 & 0.93 & \\ldots & \\ldots \\\\ RXC\\,J0145.0-5300 & 0.1168 & $5.53_{-0.13}^{+0.13}$ & $ 5.00_{-0.03}^{+0.03}$ & $5.51_{-0.16}^{+0.16}$ & $ 3.88_{-0.03}^{+0.03}$ & $5.63_{-0.14}^{+0.14}$ & $26.61_{-0.87}^{+0.89}$ & 1089.3 & 1.23 & \\ldots & \\checkmark \\\\ RXC\\,J0211.4-4017 & 0.1008 & $2.07_{-0.00}^{+0.07}$ & $ 0.81_{-0.01}^{+0.01}$ & $2.02_{-0.06}^{+0.06}$ & $ 0.48_{-0.01}^{+0.01}$ & $2.07_{-0.05}^{+0.05}$ & $ 2.03_{-0.06}^{+0.06}$ &  685.0 & 1.33 & \\ldots & \\ldots \\\\ RXC\\,J0225.1-2928 & 0.0604 & $2.47_{-0.06}^{+0.15}$ & $ 0.51_{-0.01}^{+0.01}$ & $2.61_{-0.16}^{+0.16}$ & $ 0.31_{-0.01}^{+0.01}$ & $2.67_{-0.13}^{+0.13}$ & $ 2.00_{-0.12}^{+0.12}$ &  693.9 & 0.91 & \\ldots & \\checkmark \\\\ RXC\\,J0345.7-4112 & 0.0603 & $2.19_{-0.04}^{+0.04}$ & $ 0.77_{-0.01}^{+0.01}$ & $2.15_{-0.08}^{+0.08}$ & $ 0.37_{-0.01}^{+0.01}$ & $2.30_{-0.06}^{+0.09}$ & $ 1.91_{-0.06}^{+0.09}$ &  688.4 & 0.89 & \\checkmark & \\ldots \\\\ RXC\\,J0547.6-3152 & 0.1483 & $6.02_{-0.11}^{+0.11}$ & $ 8.97_{-0.04}^{+0.04}$ & $5.68_{-0.11}^{+0.11}$ & $ 5.76_{-0.04}^{+0.04}$ & $6.06_{-0.14}^{+0.14}$ & $35.54_{-0.99}^{+1.02}$ & 1133.7 & 1.32 & \\ldots & \\ldots \\\\ RXC\\,J0605.8-3518 & 0.1392 & $4.56_{-0.05}^{+0.05}$ & $ 9.54_{-0.04}^{+0.04}$ & $4.81_{-0.12}^{+0.12}$ & $ 4.26_{-0.04}^{+0.04}$ & $4.91_{-0.11}^{+0.11}$ & $22.39_{-0.63}^{+0.66}$ & 1045.9 & 1.17 & \\checkmark & \\ldots \\\\ RXC\\,J0616.8-4748 & 0.1164 & $4.22_{-0.10}^{+0.10}$ & $ 2.38_{-0.02}^{+0.02}$ & $4.16_{-0.12}^{+0.12}$ & $ 1.88_{-0.02}^{+0.02}$ & $4.17_{-0.11}^{+0.11}$ & $11.81_{-0.41}^{+0.39}$ &  939.2 & 1.12 & \\ldots & \\checkmark \\\\ RXC\\,J0645.4-5413 & 0.1644 & $6.95_{-0.13}^{+0.13}$ & $18.88_{-0.10}^{+0.10}$ & $6.97_{-0.19}^{+0.19}$ & $11.39_{-0.09}^{+0.09}$ & $7.27_{-0.18}^{+0.18}$ & $71.61_{-2.33}^{+2.35}$ & 1280.0 & 1.28 & \\ldots & \\ldots \\\\ RXC\\,J0821.8+0112 & 0.0822 & $2.68_{-0.09}^{+0.09}$ & $ 0.77_{-0.01}^{+0.01}$ & $2.44_{-0.12}^{+0.12}$ & $ 0.54_{-0.01}^{+0.01}$ & $2.84_{-0.10}^{+0.10}$ & $ 3.34_{-0.15}^{+0.15}$ &  755.9 & 0.93 & \\ldots & \\ldots \\\\ RXC\\,J0958.3-1103 & 0.1669 & $5.34_{-0.21}^{+0.21}$ & $11.56_{-0.15}^{+0.15}$ & $5.85_{-0.40}^{+0.45}$ & $ 5.25_{-0.16}^{+0.16}$ & $6.30_{-0.44}^{+0.50}$ & $28.04_{-2.30}^{+2.67}$ & 1077.4 & 0.78 & \\checkmark & \\ldots \\\\ RXC\\,J1044.5-0704 & 0.1342 & $3.41_{-0.03}^{+0.03}$ & $ 7.42_{-0.02}^{+0.02}$ & $3.52_{-0.05}^{+0.05}$ & $ 3.00_{-0.02}^{+0.02}$ & $3.57_{-0.05}^{+0.05}$ & $11.77_{-0.19}^{+0.19}$ &  931.9 & 1.09 & \\checkmark & \\ldots \\\\ RXC\\,J1141.4-1216 & 0.1195 & $3.31_{-0.03}^{+0.03}$ & $ 3.75_{-0.01}^{+0.01}$ & $3.40_{-0.06}^{+0.06}$ & $ 1.70_{-0.01}^{+0.01}$ & $3.54_{-0.05}^{+0.05}$ & $ 8.60_{-0.15}^{+0.16}$ &  885.2 & 1.25 & \\checkmark & \\ldots \\\\ RXC\\,J1236.7-3354 & 0.0796 & $2.70_{-0.05}^{+0.05}$ & $ 1.03_{-0.01}^{+0.01}$ & $2.57_{-0.03}^{+0.11}$ & $ 0.61_{-0.01}^{+0.01}$ & $2.73_{-0.01}^{+0.09}$ & $ 3.27_{-0.02}^{+0.15}$ &  753.5 & 0.99 & \\ldots & \\ldots \\\\ RXC\\,J1302.8-0230 & 0.0847 & $2.97_{-0.07}^{+0.06}$ & $ 1.38_{-0.01}^{+0.01}$ & $2.92_{-0.07}^{+0.09}$ & $ 0.83_{-0.01}^{+0.01}$ & $3.44_{-0.07}^{+0.07}$ & $ 6.07_{-0.18}^{+0.19}$ &  842.1 & 1.22 & \\checkmark & \\checkmark \\\\ RXC\\,J1311.4-0120 & 0.1832 & $8.91_{-0.08}^{+0.08}$ & $36.06_{-0.08}^{+0.08}$ & $8.24_{-0.13}^{+0.13}$ & $15.13_{-0.07}^{+0.07}$ & $8.44_{-0.12}^{+0.12}$ & $88.18_{-1.50}^{+1.51}$ & 1319.2 & 1.31 & \\checkmark & \\ldots \\\\ RXC\\,J1516.3+0005 & 0.1181 & $4.51_{-0.06}^{+0.06}$ & $ 4.12_{-0.02}^{+0.02}$ & $4.18_{-0.08}^{+0.08}$ & $ 2.77_{-0.02}^{+0.02}$ & $4.48_{-0.07}^{+0.07}$ & $15.81_{-0.31}^{+0.30}$ &  989.9 & 1.29 & \\ldots & \\ldots \\\\ RXC\\,J1516.5-0056 & 0.1198 & $3.55_{-0.07}^{+0.07}$ & $ 2.31_{-0.02}^{+0.02}$ & $3.40_{-0.08}^{+0.08}$ & $ 1.77_{-0.02}^{+0.02}$ & $3.74_{-0.09}^{+0.10}$ & $11.08_{-0.36}^{+0.41}$ &  927.0 & 1.37 & \\ldots & \\checkmark \\\\ RXC\\,J2014.8-2430 & 0.1538 & $4.78_{-0.05}^{+0.05}$ & $21.06_{-0.07}^{+0.07}$ & $5.63_{-0.11}^{+0.11}$ & $ 7.52_{-0.07}^{+0.07}$ & $5.73_{-0.10}^{+0.10}$ & $39.89_{-0.82}^{+0.78}$ & 1155.3 & 1.09 & \\checkmark & \\ldots \\\\ RXC\\,J2023.0-2056 & 0.0564 & $2.71_{-0.09}^{+0.09}$ & $ 0.61_{-0.01}^{+0.01}$ & $2.46_{-0.12}^{+0.12}$ & $ 0.40_{-0.01}^{+0.01}$ & $2.72_{-0.09}^{+0.09}$ & $ 2.81_{-0.12}^{+0.13}$ &  739.5 & 0.86 & \\ldots & \\checkmark \\\\ RXC\\,J2048.1-1750 & 0.1475 & $4.65_{-0.07}^{+0.13}$ & $ 5.13_{-0.03}^{+0.03}$ & $4.59_{-0.08}^{+0.08}$ & $ 4.40_{-0.03}^{+0.03}$ & $5.01_{-0.11}^{+0.11}$ & $26.91_{-0.80}^{+0.81}$ & 1078.0 & 1.48 & \\ldots & \\checkmark \\\\ RXC\\,J2129.8-5048 & 0.0796 & $3.81_{-0.15}^{+0.15}$ & $ 1.46_{-0.02}^{+0.02}$ & $3.64_{-0.12}^{+0.16}$ & $ 1.19_{-0.02}^{+0.02}$ & $3.88_{-0.14}^{+0.14}$ & $ 8.67_{-0.41}^{+0.40}$ &  900.6 & 0.93 & \\ldots & \\checkmark \\\\ RXC\\,J2149.1-3041 & 0.1184 & $3.26_{-0.04}^{+0.04}$ & $ 3.56_{-0.02}^{+0.02}$ & $3.40_{-0.08}^{+0.08}$ & $ 1.58_{-0.02}^{+0.02}$ & $3.50_{-0.07}^{+0.07}$ & $ 8.65_{-0.32}^{+0.32}$ &  886.6 & 1.26 & \\checkmark & \\ldots\\\\ RXC\\,J2157.4-0747 & 0.0579 & $2.46_{-0.08}^{+0.08}$ & $ 0.45_{-0.01}^{+0.01}$ & $2.30_{-0.06}^{+0.10}$ & $ 0.37_{-0.01}^{+0.01}$ & $2.76_{-0.07}^{+0.07}$ & $ 3.07_{-0.11}^{+0.11}$ &  751.5 & 0.97 & \\ldots & \\checkmark \\\\ RXC\\,J2217.7-3543 & 0.1486 & $4.86_{-0.09}^{+0.09}$ & $ 6.12_{-0.03}^{+0.03}$ & $4.45_{-0.09}^{+0.09}$ & $ 3.70_{-0.03}^{+0.03}$ & $4.65_{-0.08}^{+0.10}$ & $20.32_{-0.47}^{+0.54}$ & 1022.6 & 1.33 & \\ldots & \\ldots \\\\ RXC\\,J2218.6-3853 & 0.1411 & $5.84_{-0.11}^{+0.11}$ & $ 9.43_{-0.06}^{+0.06}$ & $5.88_{-0.15}^{+0.20}$ & $ 5.60_{-0.06}^{+0.06}$ & $6.16_{-0.19}^{+0.19}$ & $34.36_{-1.33}^{+1.30}$ & 1130.1 & 1.04 & \\ldots & \\checkmark \\\\ RXC\\,J2234.5-3744 & 0.1510 & $7.78_{-0.15}^{+0.15}$ & $19.15_{-0.11}^{+0.11}$ & $6.95_{-0.14}^{+0.14}$ & $12.36_{-0.10}^{+0.10}$ & $7.30_{-0.12}^{+0.12}$ & $70.43_{-1.54}^{+1.51}$ & 1283.2 & 1.15 & \\ldots & \\ldots \\\\ RXC\\,J2319.6-7313 & 0.0984 & $2.22_{-0.03}^{+0.03}$ & $ 2.00_{-0.02}^{+0.02}$ & $2.48_{-0.08}^{+0.08}$ & $ 0.97_{-0.01}^{+0.01}$ & $2.52_{-0.07}^{+0.07}$ & $ 4.37_{-0.16}^{+0.16}$ &  788.7 & 1.11 &\\checkmark & \\checkmark \\\\  \\\\ \\hline \\end{tabular} \\end{center}  Columns: (1) Cluster name; (2) $z$: cluster redshift; (3) $T_1$: spectroscopic temperature of the $R < \\Rv$ region in keV; (4) $L_1$: luminosity in the $R < \\Rv$ region in units of $10^{44}$ erg s$^{-1}$; (5) $T_2$: spectroscopic temperature in the $[0.15-1]\\,\\Rv$ region in keV; (6) $L_2$: luminosity in the $[0.15-1]\\, \\Rv$ region in units of $10^{44}$ erg s$^{-1}$; (7) $T_3$: spectroscopic temperature in the $[0.15-0.75]\\,\\Rv$ region in keV; (8) $Y_X$ in units of $10^{13}\\, M_\\odot\\ {\\rm keV}$; (9) $\\Rv$ in kpc; (10) ratio of the detection radius of the surface brightness profile at $3\\sigma$ significance, $R_{\\rm det}$, to $\\Rv$; (11) systems classified as cool cores on the basis of central density vs. cooling time (see Sect.~\\ref{sec:subsamples}); (12) systems classified as disturbed on the basis of the centre shift parameter $\\langle w \\rangle$ (see Sect.~\\ref{sec:subsamples}). \\end{table*}  We adopt a $\\Lambda$CDM cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$, and all uncertainties are quoted at the 68 per cent confidence level. All logarithmic quantities are given to base $e$, and the quantity $L$ refers to the bolometric [0.01-100 keV] X-ray luminosity. ",
        "conclusions": "We have presented a detailed study of the luminosity scaling relations of \\rexcess, a galaxy cluster sample selected by X-ray luminosity in such a way as to optimally sample the cluster X-ray luminosity function. \\rexcess\\ contains objects of different dynamical states with a range of core X-ray properties, allowing us to investigate the effect of the presence of these systems on the scaling relations. The homogeneous nature of the sample data, which have all been observed with the same satellite to approximately the same depth, combined with an analysis approach based on extraction of relevant quantities within scaled apertures, has been designed to minimise measurement scatter. We found the following results:  \\begin{itemize}  \\item The slope of the luminosity-temperature, luminosity-$Y_X$ and luminosity-mass relations are all steeper  at greater than 99 per cent confidence, than expected for self-similar gravitational collapse scenarios.  \\item  The dependence of the radially averaged structure factor $\\hat{Q}(T)_{| \\rm bin}$ on temperature, where both quantities have been estimated for the first time within $\\Rv$, is not significant either when measured from all emission, or from core-excluded emission. This suggests that, contrary to previous results, structural variation cannot be a significant contributor to the steepening of the relations unless there is a very strong temperature/mass dependence of gas clumpiness. Furthermore, the scatter in $\\hat{Q}(T)_{| \\rm bin}$ measured in the $[0.15-1]\\,\\Rv$ aperture is $\\sim 15$ per cent, illustrating the remarkable structural similarity of the present sample in the outer regions.  \\item There is strong evidence for a decrease in gas mass content in poorer systems relative to higher mass systems, or in other words, a dependence of the gas mass fraction on total mass. This effect is clearly seen in the $M_{\\rm gas}-T$ relation of the present sample \\citep{croston08} and has been seen in many previous samples. Using the total mass determined from published hydrostatic estimates of a combined sample of clusters and groups spanning $10^{13}$--$10^{15}\\, M_\\odot$, we find that the gas mass fraction depends on total mass such that $f_{\\rm gas} \\propto M^{0.21\\pm 0.03}$. The trend in the \\rexcess\\ data, when total masses are estimated using the $M$--$T$ scaling relation, is similar.  This dependence is the dominant cause of the steepening of the X-ray luminosity scaling relations.  \\item For the whole sample, the scatter of the X-ray luminosity-temperature relation derived from all emission interior to $\\Rv$ is up to 75 per cent depending on the exact fitting method; the scatter is less than 40 per cent for the X-ray luminosity-$Y_X$ relation. The scatter is not strongly compatible with a Gaussian distribution in either case, the distribution being characterised by a tail caused by the presence of cool core systems.  \\item Cooling core systems, the one-third of the sample with the highest central density, describe the high luminosity envelope and contribute the majority of the variance to the relations. A fit only to the cool core systems suggests that they are offset from the relation by a simple normalisation factor. However, there is nearly two times more logarithmic scatter about the X-ray luminosity-temperature relation for cool core systems compared to that for non cool core systems. This suggests that there are large differences in the core structure even for cool core systems as a class.  \\item Systems exhibiting morphological substructure tend towards the lower luminosity envelope of the relations. Partly this is due to the increase in normalisation of the total sample due to the presence of cool core systems. The scatter about the X-ray luminosity-temperature relation of morphologically disturbed systems is 65 per cent, compared to a scatter of 48 per cent for cool core systems; within the uncertainties, the variance is in fact identical. Furthermore, the scatter in morphologically disturbed systems is identical to that for morphologically relaxed systems, due partly to the preponderance of cool core systems in the relaxed subsample.  \\item Simple exclusion of the emission interior to $0.15\\,\\Rv$ results in a reduction of scatter in all relations. For the X-ray luminosity- temperature relation, the natural logarithmic scatter is 30 per cent, a reduction of more than a factor of two. Similarly significant reductions are seen in  other relations. After exclusion of the core, the scatter in luminosity-temperature and luminosity-$Y_X$ relations is well described with a Gaussian distribution at $> 85$ per cent confidence. A reduction in scatter can also be achieved by considering the central gas density, $n_{e,0}$, as a third parameter in the scaling relations. The $L-T-n_{e,0}$ relation has a natural logarithmic scatter of 47 per cent; the $L-Y_X-n_{e,0}$ relation has a scatter of 22 per cent.  \\item Using $Y_X$ as a mass proxy, a Malmquist bias corrected luminosity mass relation for \\rexcess\\ is steeper than the raw relation due to the under-representation, for a given mass, of low luminosity clusters in the sample.  \\end{itemize}  The behaviour of the observed luminosity scaling relations thus appears to be driven principally by a mass dependence of the total gas content. Plausible physical explanations for the dependence are a variation with mass in the efficiency of conversion of baryons into stars or {\\it in situ} heating after accretion. Greater understanding of the source of the dependence will require deep observations of a similarly representative sample of group scale haloes, to measure accurate luminosities and probe the underlying physical causes; furthermore, precise calibration of the evolution of the scaling relations is needed, ideally with a similarly-selected distant cluster sample, to probe the effect over time."
    },
    "0809/0809.2782_arXiv.txt": {
        "abstract": "We report the discovery of variability in the X-ray emission from the Wolf-Rayet type star WR\\,65. Using archival \\ch\\ data spanning over 5\\,yr we detect changes of the X-ray flux by a factor of 3 accompanied by changes in the X-ray spectra. We believe that this X-ray emission originates from wind-wind collision in a massive binary system. The observed changes can be explained by the variations in the emission measure of the hot plasma, and by the different absorption column along the binary orbit. The X-ray spectra of \\wr\\ display prominent emission features at wavelengths corresponding to the lines of strongly ionized Fe, Ca, Ar, S, Si, and Mg.  WR\\,65 is a carbon rich WC9d star that is a persistent dust maker. This is the first investigation of any X-ray spectrum for a star of this spectral type.  There are indications that the dust and the complex geometry of the colliding wind region are pivotal in explaining the X-ray properties of \\wr. ",
        "introduction": "{\\changed More than 54\\% } of all Wolf-Rayet (WR) stars are in binary systems that consist of a WR and an OB-type star {\\changed \\citep{wal07}}. The collision between two stellar winds in these systems produces characteristic signatures in different wavelength bands that {\\changed can} include non-thermal radio emission \\citep{ei93}, copious and variable X-rays \\citep{stev92}, and dust emission in the infrared (IR) \\citep{wil87}.  Subject of this {\\it Letter} is the WC9d star WR\\,65. \\citet{wil87} reported the presence of warm dust in this object. They discussed that in early-type WC stars an increase in the wind density provided by {\\changed e.g.}\\ shock compression in colliding winds in binary systems may be sufficient to result in grain formation. \\citet{zub98} found no need to invoke binarity to explain the presence of dust in late WC stars. He fitted the observed IR spectrum of \\wr\\ assuming a single star with stellar mass-loss in {\\changed the} form of dust with $\\dot{M_{\\rm d}}=5.4\\times 10^{-10}\\,M_\\odot\\,{\\rm yr}^{-1}$ and surrounded by a dust shell with an inner radius of 520\\,$R_\\ast$. \\citet{lei97} detected radio emission from \\wr, but could not constrain its spectral index.  In a previous paper \\citep{osk03} we demonstrated that single WC type stars are X-ray quiet and proposed that all X-ray active WC stars must be in binary systems. X-ray emission from \\wr\\ was detected and tentatively explained as a result of wind-wind collision. In this paper we report {\\changed on the X-ray light curve of \\wr\\ } and changes in its X-ray spectra that unambiguously confirm its status as a colliding-wind binary (CWB).  \\begin{table} \\caption{Coordinates of \\wr\\ from different sources} \\label{tab:coor} \\centering \\begin{tabular}{ccccc} \\hline \\hline Ref           & {\\changed Band}   & RA J2000    & DEC J2000 \\\\  \\hline Simbad          &  opt.    & 15$^{\\rm h}$\\,13$^{\\rm m}$\\,41\\fs 68 & $-59\\degr\\,11\\arcmin\\,43\\farcs 3$ \\\\ \\citet{chap99}  & radio & 15$^{\\rm h}$\\,13$^{\\rm m}$\\,41\\fs 49 & $-59\\degr\\,11\\arcmin\\,43\\farcs 8$ \\\\ \\ch\\ HRC-I      & X-ray & 15$^{\\rm h}$\\,13$^{\\rm m}$\\,41\\fs 76 & $-59\\degr\\,11\\arcmin\\,44\\farcs 1$ \\\\ Catalog USNO-B1.0 & opt.& 15$^{\\rm h}$\\,13$^{\\rm m}$\\,41\\fs 73 & $-59\\degr\\,11\\arcmin\\,43\\farcs 4$ \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "}  The main findings from the analysis of \\wr\\ X-ray spectra obtained at different epochs are: (i) the spectra display emission features at the wavelengths of the lines of strongly ionized Fe, Ca, Ar, S, Si, and Mg; (ii) the star shows strong spectral X-ray variability; (iii) the spectra can be fitted with a two-temperature model, consisting of a soft component with $0.3 \\lsim kT_1 \\lsim 1.5 $\\,keV and a hard component with $2\\lsim kT_2 \\lsim 20 $\\,keV; (iv) the absorption for the hard component is strong and variable, while the absorption for the soft component is less severe and nearly constant; (v) the changes in the EM are stronger for the soft component.  Can these observations be explained in the framework of a colliding wind model? \\citet{cor05} compare the 2.0-10.0\\,keV light curves of WR\\,140 and $\\eta$ Car (both are systems with high eccentricity and high inclination) and note that they are very similar. \\citet{pit98} explain the light-curve (2.0\\,--\\,10.0\\,keV) of $\\eta$\\,Car using a hydrodynamic simulation of colliding winds.  They show that the data are broadly consistent with the X-ray emission being moderately low throughout most of the orbit. However, a sharp peak occurs near periastron, when a greater fraction of the stellar winds collides, and the higher pre-shock densities and the lower pre-shock velocities result in stronger radiative cooling. The eclipse occurs immediately after periastron passage when the companion of $\\eta$\\,Car swings behind the LBV star with its dense wind.  It is hard to see how a similar model can be used to explain the X-ray variability of WR\\,65.  Strong variations in the absorbing column indicate that the orbit of \\wr\\ has high inclination too. Considering the light-curve (Fig.\\,\\ref{fig:lc}) we can rule out that \\wr\\ has a period of a few days: the two data points obtained in April 2003 indicate constant count rates over at least a week. Moreover, two data points in the light-curve separated by two months (12/2004 and 02/2005) are similar, while there is an increase of emission by nearly 70\\% in the next two months. Therefore, \\wr\\ has most likely an orbital period of a few years.  In this case the long duration of the minimum is inconsistent with its interpretation as attenuation by the wind of the WC star immediately after periastron passage. Quite different from the short-duration rise of column density observed in WR\\,140 and $\\eta$ Car, the column density in \\wr\\ is persistently very high, $N_{\\rm H,2}\\msim 10^{23}\\,{\\rm cm}^{-2}$, sharply decreasing only at the maximum X-ray light.  {\\changed The temporary variations of column density in \\wr\\ are similar to those measured in the  WC8+O binary $\\gamma^2$\\,Vel \\citep{gr00}.  Its hard X-ray emission increases by a factor of $\\sim 4$ when column density decreases by factor of $\\sim$few during the passage of weak-wind inner side of the shock cone across the line of sight \\citep{aw95,gr00,sch04}.}  In \\wr\\ the soft X-ray emission is present at all observed epochs. Its variation can be explained by changing emission measure, while the absorption is nearly constant and comparably low.  A soft spectral component in the low state is also observed in $\\gamma^2$\\,Vel, V444\\,Cyg and $\\eta$\\,Car.  The origin of soft X-rays in these stars is attributed either to individual stellar winds or to a circumstellar nebula. Accordingly, in these objects the absorbing column to the soft component is similar to the interstellar absorbing column. In contrast, we find for \\wr\\ that $N_{\\rm H, ism}\\,<\\,N_{\\rm H,1}\\,<\\,N_{\\rm H, 2}$, and the EM of the soft component changes with the X-ray light-curve.  Thus it is unlikely that the intrinsic wind emission of a companion is solely responsible for the soft X-rays from WR\\,65.  WR\\,65 is a persistent dust maker \\citep{vdh}. \\citet{mon02} notice that binary stars that are persistent dust makers have {\\changed nearly} circular orbits. We can assume that the same holds for \\wr. There is a growing number of WC9d binary stars that constitute the ``pinwheel'' class of objects \\citep{tut08, mar07}. In these binary systems the dust is formed in the wind collision region and is distributed along a wide (compared to the orbital separation) Archimedean spiral. The peak of dust production occurs at some distance downstream from the colliding wind stagnation point.  So far there is no direct evidence that \\wr\\ is a pinwheel. However, the topology and structure of circumstellar matter in a pinwheel could help to explain X-ray spectral variability in \\wr. \\citet{wil00} presented a model for the absorption of X-rays that includes a treatment of dust grains. {\\changed Their work confirmed that large-grain absorption for soft X-rays ($< 2$\\,keV) is} reduced due to the self-blanketing of grains.  Interestingly, from the point of view of radiative transfer this is identical to the reduction of opacity for the X-rays in clumped stellar winds \\citep{feld03}. {\\changed Winds of WC stars are extremely rich in C and O and could have enhanced abundance of Mg and Ne.}  C and O are dominant absorbers for the softer X-rays ($< 1$\\,keV), while {\\changed Mg and Ne strongly} contribute to the absorption of harder X-rays ($> 1$\\,keV). In the {\\changed dusty} shell around a WC9d star, C would be depleted from the gaseous phase and form grains, reducing the opacity for the softer X-rays.  This may help to explain the large differences between the column densities that we determine from soft and hard X-rays.  Recently, \\citet{lem07} presented 3-D hydrodynamical simulations of colliding winds. Scaling the results of their simulation to the parameters of \\wr, soft X-ray emission could originate from a large volume extending above the orbital plane.  This emission would show little variation with orbital phase and suffer comparably mild absorption in the outer regions of the WC star wind, similarly to what {\\changed is} observed in WR\\,65. On the other hand, a high postshock temperature of $kT\\approx 10$\\,keV is predicted for a confined region near the point where the winds collide head-on. Densities up to $\\rho\\msim 10^{-14}$\\,g\\,cm$^{-3}$ can be expected in the dusty spiral arms.  Assuming the wind opacity for hard X-rays as $\\kappa\\approx 5$\\,cm$^2$\\,g$^{-1}$ \\citep{a04}, a spiral arm width of $10^{14}$\\,cm would explain the observed high absorption.  This width is consistent with the observed width of the dusty spiral arms in WR\\,104 \\citep{tut08}. The lowest absorbing column can be expected at the phase when the line of sight crosses only the outer low density part of the spiral arms. It is possible that the \\ch\\ observations of \\wr\\ at high state have caught this orbital phase.  To conclude, our first brief study of the variable X-ray emission from \\wr\\ indicates that it is a colliding wind system, where the dust and complex geometry of its colliding wind region are pivotal in explaining its X-ray properties. {\\changed At present, \\wr\\ is the only known WC9d star where X-ray emission is detected.} Future observations are needed to constrain the orbital parameters of \\wr\\ and allow for detailed modeling."
    },
    "0809/0809.0863_arXiv.txt": {
        "abstract": "We present precise optical and near-infrared ground-based photometry of two Globular Clusters (GCs): $\\omega$ Cen and 47~Tuc. These photometric catalogs are unbiased in the Red Giant Branch (RGB) region close to the tip. We provide new estimates of the RGB tip (TRGB) magnitudes---$m_I(TRGB)=9.84\\pm0.05$, $\\omega$ Cen; $m_I(TRGB)=9.46\\pm0.06$, 47~Tuc---and use these to determine the relative distances of the two GCs. We find that distance ratios based on different calibrations of the TRGB, the RR Lyrae stars and kinematic distances agree with each other within one sigma. Absolute TRGB and RR Lyrae distance moduli agree within 0.10--0.15 mag, while absolute kinematic distance moduli are 0.2--0.3~mag smaller. Absolute distances to \\tuc based on the Zero-Age-Horizontal-Branch and on the white dwarf fitting agree within 0.1 mag, but they are 0.1--0.3 mag smaller than TRGB and RR Lyrae distances. ",
        "introduction": "The road to quantitative astrophysics is paved with improvements in the absolute cosmic distance scale.  Accurate primary distances based on trigonometric parallaxes are limited to measurements with the HIPPARCOS satellite (van Leeuwen et al.\\ 2007), with the Fine Guidance Sensor on board the Hubble Space Telescope (Benedict et al.\\ 2007; Evans et al.\\ 2008) and with near-infrared interferometers on the ground (Kervella et al. 2006). Secondary distance estimates have relied either on period-luminosity relations for classical Cepheids (Sandage \\& Tammann 2006;  Fouqu\\'e, et al.\\ 2007; Caputo 2008), RR Lyrae stars (Szewczyk et al.\\ 2008), type II Cepheids (Feast et al.\\ 2008; Groenewegen et al.\\ 2008); or on fits to the main sequence (Gratton et al.\\ 2003), the white dwarf (WD) cooling sequence (Zoccali et al.\\ 2001; Layden et al.\\ 2005), the Zero Age Horizontal Branch (ZAHB, Caloi \\& D'Antona 2005; Salaris et al.\\ 2007), or the TRGB (Madore \\& Freedman 1995; Salaris \\& Cassisi 1998; Bellazzini 2008). Kinematic distances to GCs offer promise as a new and independent {\\it primary\\/} distance indicator. This method is based on the ratio between the dispersions in proper motion and radial velocity of cluster members.  The accuracy of this method is limited only by the precision of the measurements and the size of the sample and, barring urecognized systematic errors, absolute distances with an accuracy of 1\\% might be provided (King \\& Anderson 2002).  All the above methods have practical limitations, and a definitive consensus on low- and intermediate-mass distance indicators has not yet been reached. The typical problems affecting current distance estimates might be split into two groups:  {\\em i)}-- precision, i.e., the ongoing effort to reduce random observational errors; and {\\em ii)}-- accuracy, i.e., the neverending struggle to identify and eliminate systematic errors.  Problems associated with the former group can be alleviated by increases in sample size and in signal-to-noise ratio.  Usually, problems connected with the latter group can be recognized and addressed only by comparing results obtained by completely different methodologies. The number of GCs for which we have distance estimates based on different methods is quite limited. However, a few GCs do have distances based on multiple methods including the TRGB (Bellazzini et al.\\ 2004), RR Lyraes (Del Principe et al.\\ 2006, DP06; Sollima et al.\\ 2008), and kinematics (van de Ven et al.\\ 2006; McLaughlin et al.\\ 2006). Even with this extremely limited sample, we have evidence that the kinematic distances are significantly and systematically smaller than distances based on the other methods. The difference ranges from 5\\% for M15 (McNamara et al.\\ 2004), to 10--15\\% for both \\omc (van de Ven et al.\\ 2006; DP06) and 47~Tuc (McLaughlin et al.\\ 2006). On the other hand, distances to \\tuc based on the ZAHB and on the WD fitting are on average 10\\% smaller than distances based on other methods (Salaris et al.\\ 2007). ",
        "conclusions": "To identify of the TRGB, we need to distinguish the Asymptotic Giant Branch (AGB) stars.  To do this, we adopted optical--NIR CMDs. Data for \\omc (see top panel of Fig.~1) show that in the $J$, $B-K$ CMD the AGB stars separate well from RGB stars for $11 \\le J \\le 12$, while at the bright end they are, at fixed color, slightly brighter than RGB stars.  \\begin{figure}[!ht] \\begin{center} \\label{fig1} \\includegraphics[height=0.50\\textheight,width=0.45\\textwidth]{f1.ps} \\caption{{\\em a)}-- Optical--NIR $J$, $B-K$ CMD of the bright region $J \\le 14$ of \\omcp. Red circles display candidate AGB stars. The number of stars selected and the adopted selection criteria (intrinsic photometric error, {\\em separation}) are also labeled. The error bars on the right display standard errors in magnitude and color. {\\em b)}-- Same as the top, but for the $I$, $B-K$ CMD. {\\em c)}-- Same as the top, but for the $I$, $V-K$ CMD. {\\em d)}-- Same as the top, but for the $I$, $B-I$ CMD. The red circles plotted in the last three panels are the candidate AGB stars selected from the $J$, $B-K$ CMD. } \\end{center} \\end{figure}  The other panels show the distribution of bright stars in three different optical/NIR CMDs for \\omcp, where the red circles represent the candidate AGB stars selected in the top panel. Overall, this figure suggests that the sample of bright RG stars is minimally contaminated by AGB stars.  Given that, we were able to estimate the RGB luminosity function. Note that the our sample includes $\\sim$220 stars within one $I$-band magnitude of the tip. This is roughly a factor of two larger than the number considered necessary for a robust TRGB detection (Madore \\& Freedman 1995) and $\\approx$20\\% larger than the sample adopted by Bellazzini et al.\\ (2001).   Data plotted in the top panel of Fig.~2 show a well defined jump in the star counts for $m_I\\sim 9.85$, which we take to  mark the position of the TRGB.  The identification is supported by the smoothed luminosity function obtained using a Gaussian kernel with standard deviation equal to the photometric error (Sakai et al. 1996; middle panel). The bottom panel shows the response of the edge detector, a four-point Sobel filter convolved with the smoothed luminosity function; the dashed vertical lines mark the detection of the TRGB at $m_I(TRGB)\\sim 9.84\\pm0.05$. It is worth noting that the current TRGB estimate is in excellent agreement with the value provided by Bellazzini et al.\\ (2001), i.e., $m_I(TRGB)=9.84\\pm0.05$.  To estimate the dereddened magnitude, we adopted a reddening of $E(B-V)=0.11\\pm0.02$ (Kaluzny et al.\\ 2002; Calamida et al.\\ 2005).  This value, combined with the analytical relation for stellar extinction by Cardelli et al.\\ (1989), gives for the Cousins $I$ passband: $A_I=0.59\\times A_V$ = $0.59\\times3.1\\times0.11\\pm0.02$=$0.20\\pm0.06$ mag. Therefore, we end up with $m_{I,0}(TRGB)$=$9.64\\pm0.08$. The error budget includes the uncertainties in the absolute zero-point calibration ($< 0.01$), the cluster reddening, and the reddening law (Fitzpatrick 1999).  \\begin{figure}[!ht] \\begin{center} \\label{fig2} \\includegraphics[height=0.30\\textheight,width=0.45\\textwidth]{f2.ps} \\vspace*{-0.05truecm} \\caption{Top -- $I$-band luminosity function for the bright end of the RGB in \\omcp.  The vertical dashed lines indicate the position of the TRGB. Middle -- Smoothed luminosity function of the same RGB region obtained using a Gaussian kernel with standard deviation equal to the photometric error (Sakai et al.\\ 1996). Bottom -- Response of the Sobel filter to the smoothed luminosity function. } \\end{center} \\end{figure}  We adopted the same approach to estimate the TRGB in 47~Tuc.  In the absence of other NIR observations, we have cross-correlated our optical catalog with the 2MASS NIR catalog, resulting in a catalog including $\\sim 15,000$ stars with measurements in optical and NIR bands. As illustrated by the error bars in Fig.~3, the photometric accuracy is very good from the TRGB down to the red HB ($H$ $\\sim$ 12, $B-K$$\\sim$3).  As was the case for \\omcp, we find that star counts close to the TRGB should be minimally contaminated by AGB stars.  \\begin{figure}[!ht] \\begin{center} \\label{fig3} \\includegraphics[height=0.50\\textheight,width=0.45\\textwidth]{f3.ps} \\vspace*{-0.05truecm} \\caption{Same as Fig.~1, but the top panel shows the $H$, $B-K$ CMD of bright stars in 47~Tuc. Stars were selected according to intrinsic photometric error and {\\em sharpness}. } \\end{center} \\end{figure}  We estimate that the TRGB in 47~Tuc is located at $m_I(TRGB)=9.46\\pm0.06$ (see vertical lines in Fig.~4).  This estimate agrees quite well with the estimate of Bellazzini et al.\\ (2004), i.e., $m_I(TRGB)=9.40\\pm0.13$.  Note that the latter estimate includes a small correction to allow for the limited number of RGs in the last magnitude ($N_{RGB}$=80).  Our sample includes 97 RGs in the same magnitude range--- an increase of $\\sim$25\\%---bringing us closer to the number 100 stars which was proposed as the reference density for the method. We have therefore chosen not to apply this correction.  To estimate the unreddened magnitude, we adopted $E(B-V)=0.04\\pm0.02$ (Salaris et al.\\ 2007) which, with the Cardelli et al. relation $A_I = 0.59\\times A_V= 0.59\\times3.1\\times0.04\\pm0.02=0.07\\pm0.06$ mag, gives $m_{I,0}(TRGB)=9.39\\pm0.08$. The error budget includes the same sources of uncertainty as before.  \\begin{figure}[!ht] \\begin{center} \\label{fig4} \\includegraphics[height=0.30\\textheight,width=0.45\\textwidth]{f4.ps} \\vspace*{-0.05truecm} \\caption{Same as Fig.~2, but for bright RG stars in 47~Tuc. } \\end{center} \\end{figure}  To use the TRGB method to estimate the relative distance independently, we must correct for the difference in the clusters' metal content.  The TRGB calibration provided by Salaris \\& Cassisi (1998) relies on theoretical models, while that provided by Bellazzini et al.\\ (2004) relies on the absolute distance of \\omc based on the eclipsing binary ($13.65\\pm0.11$, $E(B-V)=0.13\\pm0.02$, Thompson et al.\\ 2001) and on an average distance for 47~Tuc ($13.31\\pm0.11$, E(B-V)=0.04). For our purposes we require only the slope of the metallicity-TRGB relation, and in this the two calibrations are in substantial agreement: using the quoted relations and iron abundances of [Fe/H]=-1.7 (\\omcp, metal-poor peak, Johnson et al.\\ 2008) and [Fe/H]=-0.7 (47~Tuc), the differential corrections are +0.09 and +0.06 with an uncertainty of 0.03 if we assume a generous error allowance of 0.2 dex on the metallicity difference between the two clusters. Therefore, splitting the difference between the two corrections, the relative distance becomes $\\delta \\mu=9.64\\pm0.08$(\\omcp) -  $9.39\\pm0.08$ (47~Tuc) + $0.08 = 0.33\\pm0.11$ mag.  For comparison, absolute kinematic distances to \\omc and 47~Tuc have been recently provided by van den Ven et al.\\ (2006) and by McLaughlin et al.\\ (2006). The implied relative distance between the two clusters is $\\delta \\mu = 13.41\\pm0.13$(\\omcp) - $13.02\\pm0.19$(47~Tuc) = $0.39\\pm0.23$ mag.  A similar value results if we use RR Lyrae stars to estimate the relative distances: the variable V9 ($\\log P=-0.1326$) is the only RR Lyrae currently known in 47~Tuc and it is suspected to be slightly evolved (Storm et al.\\ 1994). $K$-band magnitudes have been shown to be minimally affected by evolutionary corrections (Bono et al.\\ 2003; Catelan et al.\\ 2004) and by assuming $E(B-V)=0.04$ and the Cardelli law, we found $<k_0(V9)>=12.66\\pm0.02$ mag. Fortunately, there are five RR Lyrae stars: V90, V127, V136, V87 (Kaluzny et al.\\ 2004) and V141 (Clement et al.\\ 2001) in \\omc with periods ranging from $\\log P= -0.155$ to $-0.119$ and with accurate $K$-band mean magnitudes ($<K>=12.92\\pm0.03$). In order to account for the difference in metal abundances we adopted the metallicity estimates for individual RR Lyraes provided by Rey et al.\\ (2000) and found that the mean metallicity for these five is $<[Fe/H]>=-1.53\\pm0.17$.  The difference in mean magnitude as a function of the iron content was estimated using the theoretical calibration of DP06. The independent theoretical calibration provided by  Catelan et al.\\ (2004) provides very similar distances (DP06). For a difference in metal abundance of 0.8 dex\\footnote{Note that the mean metallicity of the five RR Lyrae stars in \\omc is not exactly the same as the RG metallicity that we employed for the TRGB.} the difference in the $K$ band is 0.07\\footnote{Estimates based on the $K$-band P-L relations for Z=0.0001--0.001--0.004 listed in Table~1 of DP06} mag. Therefore, the relative distance based on RR Lyraes is: $\\delta \\mu=  12.92\\pm0.03$(\\omcp) - $12.66\\pm0.02$(47~Tuc) + 0.07 = $0.33\\pm0.07$~mag.  We thus find that {\\it relative\\/} distances based on three completely independent methods agree within one sigma, or $\\sim 3$\\% in distance, although admittedly the current analysis is based on only two GCs.  We now want take a closer look at the consistency of our relative cluster distances with inferences from various proposed absolute distance scales that have been established on the basis of much larger samples of GCs.  The zero-point of the TRGB calibration provided by Lee et al.\\ (1993) relies on cluster distances obtained with the $M_V$ vs [Fe/H] relation for RR Lyrae stars. On this distance scale \\omc $\\mu = (9.64\\pm0.08) + (4.01\\pm0.05)=13.65\\pm0.09$, and for 47~Tuc $\\mu = (9.39\\pm0.08) + (3.93\\pm0.05)=13.32\\pm0.09$ mag, for a modulus difference of $\\delta \\mu = 0.33\\pm0.13$ mag.  We did not consider the TRGB calibration provided by Rizzi et al.\\ (2007), since 47~Tuc is outside the metallicity range covered by this calibration. The use of the empirical color-metallicity relation for TRGB stars by Bellazzini et al.\\ (2001) and of the Rizzi's TRGB calibration would provide distances very similar to the above ones. The TRGB calibration provided by Salaris \\& Cassisi (1998) is brighter in absolute terms, but new findings (Cassisi et al.\\ 2007) indicate that the discrepancy is significantly reduced.  From absolute distances based on the empirical $K$-band P-L relation for RR Lyrae stars provided by Sollima et al.\\ (2008) we found $\\mu = 13.75\\pm0.11$ for \\omc while for 47~Tuc we found, using the same unreddened magnitude as before for V9, $\\mu = 13.47\\pm0.11$, and therefore $\\delta \\mu = 0.28\\pm0.14$ mag. We note that the quoted absolute distances to \\omc agree quite well with the distances based on the eclipsing binary ($13.71\\pm0.11$, $13.75\\pm0.04$, $E(B-V)=0.11\\pm0.02$, Thompson et al.\\ 2001; Kaluzny et al.\\ 2002). The distances to 47~Tuc agree quite well with the distance based on main sequence fitting ($13.40\\pm0.03$, E(B-V)=0.04, Gratton et al.\\ 2003), but are systematically larger than distances based on fitting the WD cooling sequence ($13.15\\pm0.14$, E(B-V)=0.04, Zoccali et al.\\ 2001) and the ZAHB ($13.18\\pm0.03$, Salaris et al.\\ 2007). Unfortunately, robust distance estimates to \\omc based on these two methods are either not available or hampered by its spread in metal content.  However, the use of these moduli and of the average TRGB and RR Lyrae relative distances would provide distances to \\omc similar to the kinematic distances, but systematically smaller than the other standard candles.  Thus, we are left with the following evidence: {\\em i)} differential distance moduli or---equivalently---distance ratios based upon (a)~the absolute $I$-band relation for the TRGB, (b)~the $K$-band P-L relations for RR Lyrae stars, and (c)~the kinematic method agree with each other to within the precision of each method, or roughly 0.06~mag (see Table~1).  This finding suggests that the estimated metallicity dependences of the different methods are consistent among themselves, since these two clusters differ in metallicity by $\\sim 1$~dex.  However, we obviously cannot completely exclude the possibility that the agreement is either fortuitous or a result of error cancellation. {\\em ii)} Absolute distances based on current calibrations of the TRGB and RR Lyrae methods agree with each other to 0.10--0.15~mag, but the kinematic distance scale is 0.2--0.3~mag shorter (see Table~1). The same outcome applies to \\tuc distances based on ZAHB and WD fitting, they agree with each other within 0.10 mag, but they are 0.1--0.3 mag smaller than TRGB and RR Lyrae distances (Salaris et al.\\ 2007). Taken together, points {\\em i)} and {\\em ii)} suggest that our Population~II distance scale is {\\it not\\/} at present limited by random errors due to data precision or sample size, but rather by systematics:  either a systematic error shared by the TRGB method and the RR Lyrae method, due perhaps to a common rung in the calibration ladder that leads to their respective zero points, or a systematic bias inherent to the kinematic method that is common to the independent studies of \\omc\\ and 47~Tuc.  While we wait for accurate trigonometric parallaxes for GCs from SIM and GAIA, similar analyses of new GCs should help to clarify the relative impact of random and systematic errors among the various distance indicators."
    },
    "0809/0809.4117_arXiv.txt": {
        "abstract": "We present results of a 150~MHz survey of a field centered on $\\epsilon$ Eridani, undertaken with the Giant Metrewave Radio Telescope (GMRT). The survey covers an area with a diameter of $\\sim 2^\\circ$, has a spatial resolution of $\\sim 30''$ and a noise level of $\\sigma=3.1$~mJy at the pointing centre. These observations provide a deeper and higher resolution view of the 150~MHz radio sky than the 7C survey (although the 7C survey covers a much larger area).  A total of 113 sources were detected, most are point-like, but 20 are extended.  We present an analysis of these sources, in conjunction with the NVSS (at 1.4~GHz) and VLSS (at 74~MHz). This process allowed us to identify 5 Ultra Steep Spectrum (USS) radio sources that are candidate high redshift radio galaxies (HzRGs). In addition, we have derived the $dN/dS$ distribution for these observations and compare our results with other low frequency radio surveys. ",
        "introduction": "Over the past 50 years many large radio surveys have taken place, allowing us to build up an understanding of the evolution of the radio source population (see for example, Rowan-Robinson et al. 1993). Various deep surveys have taken place, but many have been at frequencies of 1.4~GHz and higher (for obvious angular resolution reasons) and have reached down to $\\mu$Jy levels (e.g. Hopkins et al. 1998; Huynh et al. 2005). Probably, the most important result from these surveys is the confirmation of the upturn in the Euclidean-normalised differential source counts below $\\sim 2$~mJy (Windhorst et. al. 1985) which is believed to be due to a population of sources that are not AGN (Condon 1989). The ``normal'' population of radio sources (with flux densities larger than a few mJy) can be explained by radio loud AGN which are hosted in elliptical galaxies. Towards $\\mu$Jy flux levels the radio sources counts are increasingly dominated by other populations, such as radio-quiet AGN, starbursts and late-type galaxies (with their radio emission due to massive star formation and associated supernovae).  At lower radio frequencies (below 1.4~GHz) the population of radio sources is less explored, mostly due to the relatively bright flux limits achieved at these frequencies. The Westerbork Synthesis Radio Telescope (WSRT) has been used for this purpose reaching a detection limit of $\\sim 2.5$~mJy at 610~MHz, however this is not sufficiently deep to probe the upturn in source counts (Valentijn 1980). The Giant Metrewave Radio Telescope (GMRT) located near Pune, India, however, has the necessary capabilities (low frequency with good sensitivity) and there have been a few deep surveys using the GMRT at 610~MHz. Examples include Moss et al. (2007), Garn et al. (2007), Bondi et al. (2007) and Garn et al. (2008).  The 150~MHz sky is still less explored. Although the GMRT is capable of observations at this frequency, there have been relatively few results presented in the literature. At 150~MHz the GMRT has a spatial resolution approaching $20''$ and a theoretical {\\sl rms} noise of around $\\sim 1$~mJy. This compares favourably with the Cambridge 7C survey (Hales et al. 2007), which had a resolution of $70\\times 70 \\csc\\delta$ arcsec$^2$, and a typical {\\sl rms} noise of around $\\sim 20$~mJy. Of course, the 7C survey covers a much larger area than is covered in this survey, and the final 7C catalogue contains 43683 sources over an area of 1.7~steradians.  The 7C survey region does not overlap with the field presented here.  A comparison of the point source sensitivities for a range of low frequency surveys is illustrated in Fig.~\\ref{surveys}. The point labelled ``GMRT'' is the survey presented here. Other surveys included the various Cambridge surveys (6C, 7C, 8C) at 38~MHz (Rees 1990) and 151~MHz (Hales et al. 2007), the XMM-LSS survey at 74~MHz and 325~MHz (Cohen et al. 2003), The VLA Low-frequency Sky Survey (VLSS) at 74~MHz (Cohen et al. 2007), The Westerbork Northern Sky Survey (WENSS) at 325~MHz (Rengelink et al. 1997), Molonglo Reference Catalogue (MRC) at 408~MHz (Large et al. 1981), the Sydney University Molonglo Sky Survey (SUMSS) at 843~MHz (Mauch et al. 2003) and the NVSS and FIRST both at 1.4GHz (Condon et al. 1998, White et al. 1997). We also show examples of what can be achieved in deeper smaller area surveys; the ``GMRT 610~MHz'' refer to surveys such as Garn et al. (2007) and Moss et al. (2007), which can achieve point source detections down to around 0.3~mJy, and ``Deep 1.4~GHz'' refers to surveys such as Huynh et al. (2005), with point source detections down to $\\sim 0.07$~mJy.  We also show lines corresponding to sources with spectral indices of $\\alpha=0.8$ and $\\alpha=1.5$, for a flux $S\\propto \\nu^{-\\alpha}$ (which we shall adopt throughout this paper). We can see from Fig.~\\ref{surveys} that the point source sensitivity of this GMRT survey is well matched to the NVSS. By comparing the sources detected at 150~MHz with other surveys (such as the NVSS at 1.4~GHz or the VLSS at 74~MHz) it is possible to identify sources with extreme spectral indices, which are termed Ultra-Steep Spectrum (USS) sources.  In turn, these USS sources are candidates for being high-redshift radio galaxies (HzRGs).  Radio sources can be detected out to high redshift ($z>5$). Empirically, it appears that the majority of high redshift radio sources tend to exhibit steeper radio spectra (see for example, De Breuck et al. 2000; Miley \\& De Breuck 2008). The origin for this correlation is unclear. Part of it may be due to a Malmquist bias, but it now seems likely that the higher gas densities in the galaxy environments at high redshift (constricting the development of radio bubbles from such AGN) plays a role (see Miley \\& De Breuck 2008 for a discussion of this). Having identified candidate HzRGs via their radio spectrum it is then necessary to identify the galaxies in the optical/IR and subsequent determine the redshift. HzRGs seem to have a relatively small scatter in a plot of $K$-band magnitude versus $z$, and they seem to be among the most luminous galaxies at high redshift (De Breuck et al. 2002). Apart from HzRGs there are a number of other sources that exhibit steep spectral indices. These include, but are not limited to, fossil radio galaxies, cluster halos or pulsars (Parma et al. 2007, Manchester et al. 2005).  The primary science driver behind these observations was to detect the magnetospheric radio emission from the exoplanet orbiting $\\epsilon$ Eridani. No detection was made, however, tight constraints on the possible radio emission were determined and these results are summarised in George \\& Stevens (2007). The wide field, good sensitivity and reasonable spatial resolution of the GMRT 150~MHz observations does give an excellent view of the radio sky. From these observations, we show that deep imaging with the GMRT at 150~MHz enables us to identify candidate high redshift radio sources.  Some relevant GMRT results at 150~MHz have been published by Ishwara-Chandra \\& Marathe (2007), although it is clear that this is still an under-exploited capability.  The observations and data analysis are described in Section~2. We discuss the generation of the catalogue in Section~3 and present an analysis of the radio source counts at 150~MHz in Section~4. We then investigate the nature of the sources, including spectral indices and identify the candidate high redshift objects in Section~5. We conclude in Section~6.  \\begin{figure} \\includegraphics[scale=0.3]{surveys.ps} \\caption{A comparison between the point source sensitivity of the GMRT 150~MHz results presented here and other low frequency radio surveys. To illustrate the respective point source sensitivities of the surveys, the two lines are for spectral indices of $\\alpha=0.8$ (solid line) and $\\alpha=1.5$ (dashed line), normalised to the sensitivity of the GMRT 150~MHz observations presented here. See text for a more detailed description of this diagram.} \\label{surveys} \\end{figure} ",
        "conclusions": "We have presented the results of a deep 150~MHz low frequency radio map covering an area with a diameter of $>2^{\\circ}$ in the region around $\\epsilon$ Eridani, and described the production of a catalogue of 113 sources. The radio spectral index analysis of the sources in the field made use of the NVSS (1.4~GHz) and VLSS (74~MHz).  Between this 150~MHz dataset and the NVSS we have made reliable identifications with the NVSS of 81 objects.  We have calculated the 150~MHz source flux distribution and find it to be broadly comparable to that at other wavelengths.  We have made use of these observations, in conjunction with the NVSS and VLSS to identify a number of USS sources. These observations have shown that GMRT 150~MHz observations would can be effectively used to search for USS sources. Of course LOFAR, when fully on-line, will be considerably more effective at finding such sources than the GMRT.  These GMRT 150~MHz observations are clearly capable of detecting USS sources, and we are in the process of following up some of the HzRG candidates. It is also clear that a well designed 150~MHz survey, in conjunction with observations at higher frequencies, could be powerful in detecting high redshift objects."
    },
    "0809/0809.3737_arXiv.txt": {
        "abstract": "We present {\\it B} and {\\it V} photometry of the outlying SMC star cluster BS\\,196 with the 4.1-m SOAR telescope. The photometry is deep (to {\\it V}$\\approx$25) showing $\\approx$3 mag below the cluster turnoff point (TO) at {\\it M$_{\\it V}$}=2.5 (1.03\\,M$_{\\odot}$). The cluster is located at the SMC distance. The colour-magnitude diagram (CMD) and isochrone fittings provide a cluster age of 5.0$\\pm$0.5\\,Gyr, indicating that this is one of the 12  oldest clusters so far detected in the SMC. The estimated metallicity  is [Fe/H]$=-1.68\\pm0.10$. The structural analysis gives by means of King profile fittings a core radius {\\it R}$_{\\rm c}$=8.7$\\pm$1.1\\,arcsec (2.66$\\pm$0.14\\,pc) and a tidal radius {\\it R}$_{\\rm t}$=69.4$\\pm$1.7\\,arcsec (21.2$\\pm$1.2\\,pc). BS\\,196 is rather loose with a concentration parameter {\\it c}=0.90. With {\\it M}$_{\\it V}$=$-1.89\\pm0.39$, BS\\,196 belongs to the class of  intrinsically fainter  SMC clusters, as compared to the well-known populous ones, which starts to be explored. ",
        "introduction": "Luminous old SMC star clusters (we will refer as such to the red SMC clusters with age $>$ 1\\,Gyr -- including intermediate age clusters (IAC) to globular clusters) have been systematically observed  in the last decade,  using mainly {\\it CT$_1$} photometry at CTIO  \\citep[e.g.][]{psc01,psg07}, and Hubble photometry with different cameras \\citep[e.g.][]{msf98,ggg08}. The age and age-metallicity distributions in these studies reveal essential features of the star formation history of the SMC system. Any new star cluster added to these distributions provides fundamental pieces of information to them.  The star cluster BS\\,196 was discovered on Sky Survey plates and  first reported by \\cite{bs95}. It was also listed by \\cite{bd00} in a update of the SMC catalogue. Recently, two star clusters from  Bica \\& Schmitt's  catalogue were found  to be old SMC clusters. The 4.3\\,Gyr old \\citep{ssn07} cluster BS\\,90 is unique, being  projected inside the large H\\,II region complex NGC\\,346 and BS\\,121 with 2.3\\,Gyr \\citep{psg05} is projected on  the SMC main body.  BS\\,196 appears  to be  a less populous cluster on the Sky Survey plates, with a diameter somewhat smaller than 1\\,arcmin. It is  located at $\\alpha_{2000}$ = 1h48m02s and $\\delta_{2000} = -70^\\circ$00\\arcmin15\\arcsec. It is projected  5.1$^\\circ$ to the Northeast  of  the SMC bar and the optical centre at  $\\alpha_{2000}$=00h52m45s, $\\delta_{2000}=-72^\\circ$49\\arcmin43\\arcsec~ \\citep{csp01}, being one of the angularly farthest SMC clusters known \\citep{bs95}. It is more distant from the bar than Lindsay\\,1, an extreme cluster to the West, together with AM-3  \\citep{bs95}.  It is fundamental to carry out deep photometry of the cluster  BS\\,196 to unveil its nature, whether associated to the SMC or a far globular cluster in the Galactic halo.  This paper is organized as follows. In Sect. 2 {\\it BV} photometry of BS\\,196 obtained with the SOAR telescope is presented. In Sect. 3  we determine the cluster astrophysical parameters by means of  Padova isochrone fittings, finding that it belongs to the SMC. In Sect. 4, cluster structural properties are derived. In Sect. 5 we discuss  the cluster age and metallicity  in the context of the SMC  enrichment and star formation histories. Concluding remarks  are given in Sect. 6. ",
        "conclusions": "Since by now most of the populous  old clusters in the SMC have been studied, we turned our attention to fainter clusters in similar environments. We carried out {\\it B} and {\\it V} photometry with the SOAR 4.1-m telescope of the outlying SMC star cluster BS\\,196. It is located in projection 5.1$^\\circ$ (5.6\\,kpc) from the Northeast of the SMC optical centre, thus one of the farthest clusters from the main body.  We confirmed that BS\\,196 corresponds to the SMC distance modulus and it is thus not a far halo Galactic globular cluster. It is intrinsically faint with {\\it M}$_{\\it V}$=$-1.89\\pm0.39$, a new class of objects to be explored. We derive an age of 5.0$\\pm$0.5\\,Gyr and a metallicity [Fe/H]$=-1.68\\pm0.10$. This lower mass cluster is at the lower envelope of values in the early chemical enrichment in the SMC.  This might suggest that some outer clusters share this property like L\\,38 \\citep{psc01}. The cluster structure, fitted by  2 and 3 parameter King profiles, is well described by a core radius {\\it R}$_{\\rm c}$=2.66$\\pm$0.14\\,pc and a tidal radius {\\it R}$_{\\rm t}$=21.2$\\pm$1.2\\,pc. The cluster is rather loose with a concentration parameter {\\it c}=0.90.  The systematic study of such less populous clusters may reveal age gap clusters in the LMC, while in the SMC searches for faint additional genuine (by age) globular clusters are worth conducting to better understand the formation and evolution of star clusters in them. The knowledge of the early structure of the Clouds can also benefit from such studies."
    },
    "0809/0809.1732_arXiv.txt": {
        "abstract": "{The prediction of stellar angular diameters from broadband photometry plays an important role for different applications. In particular, long-baseline interferometry, gravitational microlensing, extrasolar planet transits, and many other observing techniques require accurate predictions of the angular size of stars. These predictions are based on the surface brightness-color (SBC) relations.} {Our goal is to calibrate general-purpose SBC relations using visible colors, the most commonly available data for most stars.} {We compiled the existing long-baseline interferometric observations of nearby dwarf and subgiant stars and the corresponding broadband photometry in the Johnson $BV$ and Cousins $R_cI_c$ bands. We then adjusted polynomial SBC models to these data.} {Due to the presence of spectral features that depend on the effective temperature, the SBC relations are usually not linear for visible colors. We present polynomial fits that can be employed with $BVR_cI_c$ based colors to predict the limb-darkened angular diameters (i.e. photospheric) of dwarf and subgiant stars with a typical accuracy of 5\\%.} {The derived polynomial relations provide a satisfactory approximation to the observed surface brightness of dwarfs and subgiants. For distant stars, the interstellar reddening should be taken into account.} ",
        "introduction": "The surface brightness-color (hereafter SBC) relations link the emerging flux per solid angle unit of a light-emitting body to its color. They express the Stefan-Boltzmann relation between total brightness and effective temperature in measurable quantities such as angular diameter, magnitude and color. It basically assumes that temperature and bolometric correction can be translated into a color measurement. These relations have many astrophysical applications, such as the estimation of Cepheid distances (through the Baade-Wesselink method), extrasolar planet transit studies or the characterization of microlensing sources. Different colors produce tighter or more dispersed relations, simply because color also depends on other variables such as gravity or metallicity, at a different level depending on the adopted photometric bands. It has been shown that a combination of visible and near-IR bands, such as $(V-K)$, is probably optimal (see, e.g., Fouqu\\'e \\& Gieren~\\cite{fouque97}, in the case of Cepheids). However, near-infrared photometry is not always available, for instance due to catalogue incompleteness or field crowding of IR atlases such as 2MASS in the Galactic disk. As a consequence, SBC relations in the visible domain remain very useful.  Although they are relatively similar for stars of different luminosity classes, the literature gives appropriate relations for Cepheids (Kervella et al.~\\cite{kervella04c}), giants (van Belle~\\cite{vanbelle99a}, Nordgren et al.~\\cite{nordgren02}), M giants and supergiants (van Belle et al.~\\cite{vanbelle99b}, Groenewegen~\\cite{groen04}), and dwarf or subgiant stars (Kervella et al.~\\cite{kervella04b}, hereafter K04). Here we limit our discussion to dwarf stars and subgiants, and revise and extend K04's work, which was based on Johnson photometric bands and was limited to linear SBC relations. In the present Note, we calibrate the SBC relations based on Cousins $R_c$ and $I_c$  photometric bands, of more common use than Johnson's $R$ and $I$, and we adjust non-linear SBC relations that are a better fit to the observations than linear laws. In this process, we use exclusively direct interferometric angular diameter measurements, thus avoiding cross-correlations with previously calibrated SBC relations. ",
        "conclusions": "We computed polynomial SBC relations in $BVR_cI_c$ that can be used to predict the angular size of individual stars for which only broadband photometry is available. The visible-infrared linear relations established by K04 usually provide more accurate angular size estimates. However, in many cases the $JHK$ band magnitudes of field stars are not available (due to crowding in 2MASS for instance), and the present $BVR_cI_c$ relations will provide a reliable photospheric angular diameter prediction with a typical uncertainty of $\\approx 5$\\%. Their compatibility with visible-IR relations gives further confidence in their accuracy."
    },
    "0809/0809.1959_arXiv.txt": {
        "abstract": "{ Short-period comet 2P/Encke is a favorable target for studies of light scattering by bare cometary nuclei. These studies enable assessment of the nucleus size, the material albedo, and the light-scattering properties of the cometary surface. } { We characterize the activity of the cometary coma and the physical properties of the nucleus of 2P/Encke such as its size, albedo, colors, spectral slope, surface roughness, porosity, and single-particle properties. } { Broadband imaging photometry and broadband and narrowband linear polarimetry is measured for the nucleus of 2P/Encke over the phase-angle range 4 -- 28 deg. An analysis of the point spread function of the comet reveals only weak coma activity, corresponding to a dust production of the order of 0.05 kg/s. The nucleus displays a color independent photometric phase function of almost linear slope ($\\beta$ = 0.050$\\pm$0.004 mag/deg). The absolute $R$ filter magnitude at zero phase angle is 15.05$\\pm$0.05, and corresponds to an equivalent radius for the nucleus of 2.43$\\pm$0.06km (adopted albedo of 0.047). The nucleus color $V - R$ is $0.47 \\pm 0.07$, suggesting a spectral slope $S'$ of $11\\pm 8$\\,\\%/100nm. The phase function of linear polarimetry in the $V$ and $R$ filter shows a widely color independent linear increase with phase angle ($0.12\\pm0.02$\\,\\%/deg).  We find discrepancies in the photometric and polarimetric parameters between 2P/Encke and other minor bodies in the solar system, which may indicate significant differences in the surface material properties and light-scattering behavior of the bodies. } { The linear polarimetric phase function of 2P/Encke presented here is the first ever measured for a cometary nucleus, and its analysis encourages future studies of cometary nuclei in order to characterize the light-scattering behavior of comets on firm empirical grounds and provide suitable input to a comprehensive modeling of the light scattering by cometary surfaces. } ",
        "introduction": "After its discovery in 1786, the short-period comet 2P/Encke was observed during almost 60 perihelion passages and is one of the most significantly (more than 40 papers in refereed journals) and longest (over 200 years) studied comets. With an orbital revolution period of about 3.3 years, it is closer to the Sun than typical Jupiter family comets (JFC). Levison et al. (2006) suggested that 2P/Encke may be a rare example of a comet in an orbit detached from Jupiter's immediate control (like JFCs are) due to gravitational interaction with the terrestrial planets. They also argued that the nucleus may have been in a quiescent state for a long part of the transition period into its current orbit.  A series of scientific papers dealt with the nucleus properties of 2P/Encke. Fernandez et al. (2000) measured an effective nuclear radius of 2.4$\\pm$0.3km, an axial ratio of at least 2.6, a geometric albedo of $0.047\\pm0.023$, which is in the typical range for comets, and a rather steep photometric phase function with a linear slope of 0.06 mag/deg which was later confirmed by Jewitt (2002). A wider range of values for the effective radius (2.4--3.7 km) was determined from radar echoes of the comet (Harmon and Nolan 2005). Colors of {\\it V-R} = 0.39$\\pm$0.06 and B-V = 0.73$\\pm$0.06 (Lowry and Weissman 2007) indicated only moderate spectral slope S' of the nuclear surface as found by Jewitt (2002) to be 8.9$\\pm$1.6 \\%/100nm. The rotation period was a matter of debate in various papers without a final conclusion (see Belton et al. 2005 and references therein). Attempts to determine the orientation of the rotation axis were made by Sekanina (1991) and Festou and Barale (2000). Jewitt (2004) measured a slightly positive linear polarization of the nucleus at a single phase angle (polarization of 1.0--1.9 $\\pm0.5$ \\% at 22\\degr) in the visible. The activity profile of the nucleus appears to be unusual with a coma being visible around aphelion (Fernandez et al. 2005), but no activity during its inbound arc as close as 1.4 AU solar distance (Jewitt 2004).  The likely absence of activity combined with the predicted nucleus brightness enable us to attempt successfully (i.e. with sufficient signal-to-noise and without significant coma contamination) polarization measurements of 2P/Encke. We report below on the results of the phase-angle-resolved linear polarimetry and quasi-simultaneous broadband photometry  of this cometary nucleus in the visible wavelength range. The main aim of these measurements is to constrain the light-scattering properties of the nuclear surface, i.e. the single-scattering albedo and mean free path as a proxy for the typical grain size to be obtained through modeling. Furthermore, we could test the albedo-polarization relationship known from asteroids (Zellner and Gradie 1976, Zellner et al. 1977, Lupishko and Mohamed 1996, Cellino et al. 1999), for a cometary nucleus. Our results are the first ever linear polarimetric observations of a cometary nucleus over a wider phase-angle range. ",
        "conclusions": "\\label{discussion}  We have presented photometric and linear polarimetric phase functions of 2P/Encke measured during the 2006 approach of the comet to the Sun. The weak coma found about the comet does not affect significantly the photometric and polarimetric light-scattering phase functions measured, i.e. the phase functions reflect the light-scattering behavior of the nucleus surface. Both phase functions show an almost linear behavior with phase angle $\\phi$ (over the measured range from 4 to 28 deg). Trends in the phase function with wavelength are not obvious for the nucleus photometry, but may exist for the linear polarization. A small opposition surge is tentatively suggested in the photometric data, but needs verification by further observations at smaller phase angles. Similarly, the value of minimum polarization P$_{min}$ in the polarization phase curve should be reassessed by new observations covering the range to zero phase angle that is of immediate interest for modeling the opposition effects in the light-scattering by the cometary nucleus. Our radius estimation and spectral slope S' agree with earlier measurements, in particularly well with those obtained from IR observations. The low activity level of the nucleus at solar distances of 2.75 to 2.1AU suggests that the cometary nucleus is covered by a crust that prevents major outgassing. This appears to be true during this part of the orbit, even though 2P/Encke has entered the distance range where water sublimation usually generates significant cometary activity.  Table \\ref{comparison} compares light-scattering properties of the nucleus of comet 2P/Encke with those of other solar-system bodies, that may have similar surfaces and were observed in similar conditions (phase angle, wavelength).  The majority of the data are for V filter - if not indicated otherwise in the table - at the phase angles 5-30\\degr. The analysis of Table \\ref{comparison} indicates that comet 2P/Encke displays photometric properties typical of cometary nuclei. As for other cometary nuclei, the slope of the phase function differs from that measured for cometary dust, which might be due to the difference between light scattering on individual dust particles and their compact layers at the nuclear surfaces.  Among asteroids, the light scattering behavior of C-type objects resembles most closely that of comet 2P/Encke, at least its photometric properties. There are major differences between the polarization phase curve (Figure \\ref{pol_phasefunction}) of the nucleus of 2P/Encke and the polarization phase curves of asteroids of different taxonomy. Primitive asteroids of type C, P, G, and B and evolved asteroids of type V, S, M, E, and A have, in general, larger inversion angles than 2P/Encke (cf. Muinonen et al, 2002; Fornasier et al., 2006), while only primitive F-type asteroids exhibit similar inversion angles (Belskaya et al., 2005). However, other polarimetric characteristics of F-type asteroids differ significantly: The angle of the polarization minimum is at least twice as large, and the slope of polarization at the inversion angle is almost three times larger than that of comet 2P/Encke. Moreover, F-type asteroids do not exhibit red colors, but display neutral or even bluish surface colors. Unfortunately, we cannot compare the properties of comet 2P/Encke with icy bodies such as Transneptunian objects or the small satellites of Saturn, Uranus, and Neptune since these objects can be observed only at small phase angles. The Centaur 2060/Chiron may be the object most appropriate for the comparison. Compared to comet 2P/Encke, it has however neither similar photometric (too high albedo, too red color) nor polarimetric data (too small inversion angle). The polarimetry therefore reveals unique properties of the nucleus of comet 2P/Encke that may be typical for cometary nuclei but differ significantly from the properties of other solar-system bodies. This uniqueness does not manifest itself photometrically. It may indicate a specific composition or structure of the surface layer of cometary nuclei. Whether the narrow negative polarization branch can be explained qualitatively by the single-scattering interference mechanism (Muinonen et al. 2007) remains to be answered by a future study. Based on that mechanism, negative polarization branches of icy objects which corresponds to smaller real parts of refractive indices, can be expected to be narrower than those of silicate-rich stony objects, which represents larger real parts of refractive indices.  The polarization minimum of 2P/Encke, which is not well constrained by the present data set, is less than or equal to -1\\% and, as mentioned above, it is possibly within the range of data for F-type asteroids that have polarization minima between about -1.0 and -1.5\\% (Belskaya et al. 2005). However, we reiterate that the polarization phase curve of 2P/Encke is the only data set for comets available to ourselves. The number of F-type asteroids with measured polarization phase curves is also small such that, in terms of polarization properties, it is difficult to decide whether comets and F-type asteroids are similar.  Our 2P/Encke results have shown that the empirical albedo-polarization relationships of asteroids cannot easily be applied to cometary nuclei, possibly because the surface constitution and light-scattering properties are different. Whether such a relationship can be established for cometary nuclei remains undecided until more objects are studied and a self-consistent modeling of photometric, spectroscopic, and polarimetric light-scattering properties for comets is available.  \\acknowledgement{We would like to thank everyone at the European Southern Observatory at Cerro Paranal in Chile and in Garching/Germany, who were involved in the implementation and execution of this service mode program at the Very Large Telescope VLT.}"
    },
    "0809/0809.1397_arXiv.txt": {
        "abstract": "The properties of Galactic molecular clouds tabulated by Solomon et~al. (1987) (SRBY) are re-examined using the Boston University-FCRAO Galactic Ring Survey of \\coa\\ J=1-0 emission. These new data provide a lower opacity tracer of molecular clouds and improved angular and spectral resolution compared to previous surveys of molecular line emission along the Galactic Plane. We calculate GMC masses within the SRBY cloud boundaries assuming LTE conditions throughout the cloud and a constant \\htwo\\ to \\coa\\ abundance, while accounting for the variation of the $^{12}$C/$^{13}$C with Galacto-centric radius.  The LTE derived masses are typically five times smaller than the SRBY virial masses.  The corresponding median mass surface density of molecular hydrogen for this sample is 42 \\mpcsq, which is significantly lower than the value derived by SRBY (median 206 \\mpcsq) that has been widely adopted by most models of cloud evolution and star formation.  This discrepancy arises from both the extrapolation by SRBY of velocity dispersion, size, and CO luminosity to the 1~K antenna temperature isophote that likely overestimates the GMC masses and our assumption of constant \\coa\\ abundance over the projected area of each cloud.  Owing to the uncertainty of molecular abundances in the envelopes of clouds, the mass surface density of giant molecular clouds could be larger than the valued derived from our \\coa\\ measurements. From velocity dispersions derived from the \\coa\\ data, we find that the coefficient of the cloud structure functions, $v_\\circ=\\sigma_v/R^{1/2}$, is not constant, as required to satisfy Larson's scaling relationships, but rather systematically varies with the surface density of the cloud as $\\sim\\Sigma^{0.5}$ as expected for clouds in self-gravitational equilibrium. ",
        "introduction": "Giant molecular clouds (GMCs) are the exclusive sites of star formation in galaxies.  Their evolution and conversion of interstellar material into stars are governed by the interplay between self-gravity, magneto-turbulent pressure, and feedback processes from newborn stars (McKee 1989). The configuration of a GMC, parameterized by its observed size, velocity dispersion, and mass surface density, offers a snapshot view of its dynamical state. Larson (1981) identified scaling relationships between these observable quantities that have provided the basic, grounding point for all subsequent descriptions of interstellar molecular clouds and star formation These scaling relationships are: 1) a power relationships  between the velocity dispersion, $\\sigma_v$, and the spatial scale of the emitting volume, $L$, $\\sigma_v \\sim L^{0.38}$, 2) self-gravitational equilibrium, $2\\sigma_v L^2/GM\\sim 1$, and 3) an inverse relationship between the mean density, $n$, and size of the cloud, $n \\sim L^{-1.1}$.  The last of these relationships implies that all molecular clouds have comparable gas surface density. The original compilation of these scaling relationships used molecular line data available from earlier studies. Many of these data were collected with the earliest millimeter wave telescopes and instrumentation  and included spatially undersampled maps of molecular line emission from a limited number of nearby ($<$ 2.2 kpc) interstellar clouds with poor sensitivity, compared to currently available data.  The Larson scaling relationships have been supplemented with additional observations with improved sensitivity and larger samples.  The most significant study to confirm Larson's results is described by Solomon et~al. (1987) (hereafter, SRBY). They used the University of Massachusetts-Stony Brook (UMSB) Galactic Plane Survey (Sanders et~al. 1985) to identify 273 GMCs in the first quadrant. Each cloud was defined as a closed surface within the longitude-latitude-velocity data cube at a given threshold of antenna temperature. For most entries in the catalog, the threshold was 4 K (T$_R^*$).  This high threshold, relative to the noise of the data, was necessary to avoid the blending of emission at lower intensity levels from unrelated clouds that are densely distributed within the $l-b-V_{LSR}$ domain of the spectroscopic observations. The blending is particularly severe near the tangent points at each Galactic longitude.  Realizing that such a high threshold would not fully account for the bulk of the emission from a GMC, SRBY extrapolated the sizes, velocity dispersions, and CO luminosities to the 1 K isophote. After accounting for this low level contribution and its effect on the tabulated properties, SRBY identified a size-linewidth relationship with a steeper power relationships index (0.5) than derived by Larson (1981) and concluded that GMCs are self-gravitating objects in virial equilibrium. An algebraic consequence of these two results is that the molecular gas surface density is constant for all clouds with $\\Sigma(H_2)=200$ \\mpcsq. This value is corrected from 170 \\mpcsq\\ quoted in SRBY to account for the difference in the values adopted for the Galactocentric radius of the Sun (10 kpc vs 8.5 kpc) that affects the virial mass. We have similarly corrected other values in the SRBY catalog (CO luminosity, distance, galactocentric radius, sizes) using the rotation curve of Clemens (1985) with $R_\\odot$=8.5 kpc.  The data used by Larson (1981) and SRBY are not optimal for deriving properties of interstellar clouds.  In  both cases, the observed cloud fields were spatially undersampled with respect to the angular resolution of the telescope used to gather the CO data. Owing to poor sensitivity or the need for identifying clouds at a high intensity threshold, the target clouds are biased towards the brightest interstellar clouds, the brightest regions within the clouds,  or clouds that happen to be nearby or associated with conspicuous star formation. In the case of SRBY, the tracer of molecular gas was \\co\\ J=1-0 line emission that is optically thick under most prevailing conditions in molecular clouds. Given these limitations, it is reasonable to inquire whether the properties of GMCs and the corresponding Larson scaling relationships can withstand scrutiny with vastly superior data available today.  For example, a recent study by Bolatto et~al. (2008) examined the properties of resolved giant molecular clouds in dwarf galaxies and those within Local Group spiral galaxies.  They found the properties of extragalactic GMCs similar to those determined by SRBY for Galactic GMCs.  The Boston University-FCRAO Galactic Ring Survey (GRS) imaged the \\coa\\ J=1-0 emission between Galactic longitudes 18$^\\circ$ and 56$^\\circ$ and latitudes, $|b| \\le 1^\\circ$ with the FCRAO 14m telescope (Jackson et~al. 2006). The advantages of the GRS over the UMSB survey includes higher angular sampling and spectral resolution and the use of a mostly optically thin tracer of molecular gas.  The lower opacity of \\coa\\ reduces, but does not fully eliminate, the effect of velocity crowding. It also enables a more direct measure of molecular hydrogen column density and mass under the  assumption of local thermodynamic equilibrium (LTE) (Dickman 1978). These advantages are illustrated in Figure~\\ref{fig1}, which compares the integrated emission from several clouds identified by SRBY derived from each survey. The UMSB \\co\\ data identifies the location, angular extent and velocity of the molecular cloud but is unable to discern any substructure within the cloud. Owing to improved resolution afforded by fully sampling optically thin \\coa\\ emission, the GRS data provide a more precise distribution of molecular material {\\it within} the field  with reduced confusion of signal from unrelated clouds along the line of sight. In this study, we examine the properties of the SRBY sample of GMCs using the \\coa\\ J=1-0 data of the GRS. ",
        "conclusions": "The GMC masses derived in this study are systematically lower than those estimated by SRBY. These lower values have significant implications to the global molecular content of the Galaxy. Solomon \\& Rivolo (1989) provide a detailed summary of the biases and completeness inherent in their cloud definition.  They demonstrate that the SRBY catalog is complete to a limiting mass of 2.5$\\times$10$^5$ \\msun\\ (corrected for current Galactic distances). Accounting for Malmquist bias effects, they estimate that the SRBY clouds account for 40\\% of the total CO luminosity and 40\\% of the flux of the UMass-Stony Brook Survey. The remaining 60\\% of the CO luminosity and flux is presumed to originate from cold and/or small clouds that fall below their GMC identification threshold. The mass distribution, dN(M)/dM, of GMCs follows a power relationships, $ dN/dM \\sim M^{-\\alpha_M} $ with $\\alpha_M$ ranging from 1.5 (SRBY) to 1.8 (Heyer, Carpenter, \\& Snell 2001).  The measured slopes of the GMC mass function imply that most of the molecular mass in the Galaxy resides within the largest clouds that are included in the SRBY sample. Yet, our new calculations of cloud masses imply that the true masses are smaller than the SRBY values by factors of 2-5 depending on the correction for subthermally excited gas and varying abundances within a cloud. Such a rescaling of GMC masses implies a reduced molecular mass content of the Galaxy by these same factors.  Could there be a significant molecular gas mass component residing within smaller, cooler clouds that were not included in the SRBY catalog?  Owing to higher angular sampling than the UMSB survey, the GRS is more sensitive to smaller clouds. A comparison of the GRS field with SRBY clouds does indeed show both discrete and diffuse features that are not included within the SRBY catalog. Yet, we find that 32\\% of the \\coa\\ emission over the GRS field originates within SRBY cloud boundaries, which is comparable to the value (40\\%) estimated by Solomon \\& Rivolo (1989) for \\co.  The emission from volumes external to the SRBY boundaries may provide a significant and unaccounted reservoir of molecular material within the Galaxy.  The molecular gas fraction depends on both the surface density of the cloud and the ambient radiation field (Elmegreen 1989; van Dishoeck \\& Black 1988).  Significant abundances of \\htwo\\ and CO require sufficient column densities to self-shield. It is possible that the fraction of LTE-derived mass to the SRBY-derived virial masses, $M_{LTE}/M_{v,SRBY}$, and the higher SRBY surface densities reflect a more intense ambient radiation field in the inner Galaxy owing to higher star formation activity relative to the Solar neighborhood.  In this case, one would expect to find a dependence of this fraction on Galactocentric radius.  However, we find no evidence for any variation of this fraction over the range of radii (4.1-8.1 kpc) of the cloud sample.  \\subsection{Re-examining Larson's Scaling Relationships} The scaling relationships identified by Larson (1981) have provided a fundamental, observational constraint to descriptions of cloud dynamics and star formation.  The study by SRBY seemingly confirmed these scaling relationships for a larger sample of molecular clouds distributed throughout the Galaxy.  Given the results in the previous section, which demonstrate that the SRBY GMC masses and surface densities are likely overestimates to the true values, it is useful to re-examine the Larson (1981) scaling relationships with the new GRS data within the SRBY defined cloud boundaries.  As has been demonstrated by many previous studies, the Larson scaling relationships are algebraically linked.  Here, we resummarize this coupling to derive the coefficients that are critical to the interpretation of cloud dynamics. Gravitational equilibrium (Larson's second relationship) implies that the observed mass of the cloud, $M_{obs}$, is equal to the virial mass, \\begin{equation} M_{obs} = 5\\sigma_v^2R/G \\end{equation} The cloud size, R, is defined as the radius of a circle with the equivalent area of the cloud. The molecular gas surface density, $\\Sigma$, of a cloud is simply the \\htwo\\ mass divided by the projected area, \\begin{equation} \\Sigma = {{M_{obs}}\\over{\\pi R^2}} \\end{equation} Eliminating $M_{obs}$ and solving for $\\sigma_v$, \\begin{equation} \\sigma_v = (\\pi G/5)^{1/2} \\Sigma^{1/2} R^{1/2} \\end{equation} If $\\Sigma$ is approximately constant for all clouds (Larson's third relationship), then one recovers the size-line width scaling (Larson's first relationship), $\\sigma_v=v_{\\circ,G} R^{1/2}$, with the normalization coefficient \\begin{equation} v_{\\circ,G}=(\\pi G \\Sigma/5)^{1/2} \\end{equation} \\begin{equation} v_{\\circ,G} = 0.52 \\Bigl( { {\\Sigma} \\over {10^2\\; M_\\odot pc^{-2}}}\\Bigr)^{1/2} \\;\\; km s^{-1} pc^{-1/2} \\end{equation} The coefficient, $v_{\\circ,G}$, parameterizes the scaling of velocities within a cloud such that non-thermal pressure balances the self-gravity of the cloud with surface density, $\\Sigma$, and radius, $R$.  More generally, the velocity field of an interstellar cloud is described by the structure function that measures the variation of velocity differences of order p, with spatial scale, $\\tau$, \\begin{equation} S_p(\\tau) = \\langle |v(\\mathbf{x})-v(\\mathbf{x}+\\mathbf{\\tau})|^p\\rangle \\end{equation} where the angle brackets denote a spatial average over the observed field. Within the inertial range, the structure function is expected to vary as a power relationships with $\\tau$. For p=1, \\begin{equation} S_1(\\tau) = {\\delta}v = v_\\circ \\tau^\\gamma \\end{equation} where $\\gamma$ is the scaling exponent and $v_\\circ$ is the scaling coefficient.  These parameters correspond to Type 4 size-line width relationships described by Goodman et~al. (2004). The velocity dispersion \\footnote{Cloud-to-cloud size-velocity dispersion relationships use the full velocity dispersion of the cloud but scaled to the cloud radius, $R \\sim L/2$.  Therefore, the respective definitions for the coefficient may differ by a factor of $\\sim 2^\\gamma$.} of an individual cloud is simply the structure function evaluated at its cloud size, $L$, such that  $\\sigma_v=S_1(L)=v_\\circ L^\\gamma$. Cloud-to-cloud size-velocity dispersion relationships, defined as Type 2 by Goodman et~al. (1998), are constructed from the end-points of each cloud's velocity structure function. The existence of a cloud-to-cloud size-velocity dispersion relationship identified by SRBY necessarily implies narrow distributions of the scaling exponent and coefficient respectively for all clouds (Heyer \\& Brunt 2004).  Large variations of $v_\\circ$ and $\\gamma$ between clouds would induce a large scatter of points that is inconsistent with the observations. From Monte Carlo modeling of the scatter of the SRBY size-velocity dispersion relationship, Heyer \\& Brunt (2004) constrained the variation of $\\gamma$ and $v_\\circ$ between clouds to be less than 20\\% about the mean values that is indicative of a universal structure function.  This universality is also reflected in the structure functions of individual clouds as derived by Brunt (2003) and Heyer \\& Brunt (2004) using Principal Component Analysis.  The Larson scaling relationships are concisely represented within the plane defined by the gas surface density, $\\Sigma$, and the quantity, $\\sigma_v/R^{1/2}$, for a set of GMC properties (see equation 10).  This representation assumes a scaling exponent of 1/2 for the structure function of each cloud so that the ordinate, $\\sigma_v/R^{1/2}$, is equivalent to the scaling coefficient, $v_\\circ$. Absolute adherence to universality and all three of Larson's scaling relationships for a set of clouds would be ideally represented by a single point centered at $\\sigma_v/R^{1/2} = ({\\pi}G\\Sigma/5)^{1/2}$ for a constant value of $\\Sigma$. Given uncertainties in distance and deriving surface densities, one more realistically expects a cluster of points at this location. In Figure~7, we show the corresponding points derived from the GRS data within the SRBY boundaries (area $A_1$) and the area within the half-power isophote of $N_{H_2}$ (area $A_2$). The vertical error bars displayed in the legend reflect a 20\\% uncertainty in the distance to each cloud. As a reference point, the large triangle denotes the location of the SRBY median values ($\\sigma_v/R^{1/2} = 0.72$ \\kms; $\\Sigma(H_2) = 206$ \\mpcsq). The solid line shows the loci of points assuming gravitationally bound clouds, $\\sigma_v/R^{1/2}=(\\pi G/5)^{1/2} \\Sigma^{1/2}$ that is nearly identical to the coefficients used by SRBY. For both considered cloud areas, the \\coa\\ data points are displaced from this loci of virial equilibrium. The median virial parameter, $\\alpha_G=M_{v,13}/M_{LTE}$, is 1.9, where $M_{v,13}$ is the virial mass derived from \\coa\\ data within $A_1$. However, the LTE-derived mass could underestimate the true cloud mass by factors of 2-3 as suggested in \\S2, so the derived properties are consistent with a virial parameter of unity for this sample of clouds.  Figure~7 reveals a systematic variation of $v_\\circ=\\sigma_v/R^{1/2}$ with $\\Sigma$. This trend is separately evident for each area, $A_1$ (open circles) and $A_2$ (filled circles) with Pearson correlation coefficients 0.48 and 0.65 respectively.  For these sample sizes, it is improbable that these data sets are drawn from a random population. The dependence of $\\sigma_v/R^{1/2}$ on $\\Sigma$ signals a departure from the universality of the velocity structure function of clouds.  It implies a necessary modification to Larson's scaling relationships but one that is compatible with the rather basic premise of gravitational equilibrium as described in equation 10. The measured variation of $v_\\circ=\\sigma_v/R^{1/2}$ is larger than the values derived by Heyer \\& Brunt (2004) owing to the larger intrinsic scatter in the size-velocity dispersion relationship determined from the GRS data.  The dependence of $\\sigma_v/R^{1/2}$ on $\\Sigma$ may not have been recognized in previous studies owing to a limited range of surface densities in the observed samples, or the use of a less reliable tracer of molecular gas column density, or simply not considered given the long-standing acceptance of Larson's scaling relationships.  The fidelity of the GRS data provides an excellent relative, if not absolute, measure of gas surface density that allows this relationship to be recognized.  We note that this relationship is algebraically imposed when deriving surface densities from the virial mass, $\\Sigma = M_{vir}/\\pi R^2 \\propto \\sigma_v^2/R$, as calculated by SRBY. However, as shown in Figure~8, the relationship is even evident in the SRBY defined properties when using the mean \\co\\ surface brightness and CO to H$_2$ conversion factor as a measure of gas surface density, $\\Sigma=X_{CO}L_{CO,SRBY}/\\Omega_1 D^2$, where $\\Omega_1$ is the solid angle of the cloud corresponding to $A_1$ and $D$ is the distance.   Moreover, the scaling between $\\sigma_v/R^{1/2}$ and $\\Sigma$ is also present in the sample of extragalactic GMCs tabulated by Bolatto et~al. (2008) (filled squares in Figure~8).  The presence of this scaling within these independent data sets offers a powerful verification that the velocity dispersion of a cloud depends on both the spatial scale of the emitting area and the mass surface density.  \\subsection{GMC Dynamics} Descriptions of cloud dynamics must consider the nature and origin of the observed supersonic motions in GMCs.  While much of the theoretical effort has focused on the scaling exponent of the power spectrum or structure function of the velocity field, the normalization, $v_\\circ$, provides an important measure of the amplitude of these motions as it is evaluated at a fixed scale of 1 pc.  Figure~7 illustrates an additional constraint to these descriptions  -- that for a given cloud, the amplitude of the motions depends on the mass surface density.  It is not evident from the present measurements whether this variation of $v_\\circ$ with $\\Sigma$ is due to varying evolutionary states of the sample clouds or one that reflects different cloud conditions owing to the environmental diversity of the ISM.  For sub-Alfv\\'enic clouds whose neutral gas component is dynamically coupled to the interstellar magnetic field through ion-neutral collisions, the observed motions could arise from the propagation of large amplitude, long wavelength Alfv\\'en waves through the cloud (Arons \\& Max 1975). In fact, a simple model of magnetically supported clouds in which the cloud surface density equals the magnetic critical surface density, $\\Sigma_c=B/(63G)^{1/2}$, predicts the trend observed in Figure~7 (Mouschovias 1987; Myers \\& Goodman 1988a; Mouschovias \\& Psaltis 1995; Mouschovias, Tassis, \\& Kunz 2006). The observed variation of the coefficient with surface density simply reflects plausible differences of the magnetic field strength between clouds owing to Galactic environments and the support of the cloud by the magnetic field such that $\\Sigma \\approx \\Sigma_c$. The vertical scatter of values for a given surface density would arise from varying flux-to-mass ratios owing to ambipolar diffusion.  A definitive test of the role of the interstellar magnetic field in the support of the cloud and the origin of observed internal motions requires measures of the magnetic field strength.  Such measurements are not available for this set of GMCs.   Myers \\& Goodman (1988b) demonstrate that the magnetic field strength derived from thermal OH Zeeman measurements is comparable to the predicted field strength assuming equipartition between the magnetic, kinetic, and gravitational energies, $B_{eq}\\approx(45/G)^{1/2}\\sigma_v^2/R$ for a set of nearby clouds."
    },
    "0809/0809.4267_arXiv.txt": {
        "abstract": "We present high resolution ($R$=25,000--35,000) $K$-band spectroscopy of two young stars, MWC 480 and V1331 Cyg. Earlier spectrally dispersed ($R$=230) interferometric observations of MWC 480 indicated the presence of an excess continuum emission interior to the dust sublimation radius, with a spectral shape that was interpreted as evidence for hot water emission from the inner disk of MWC 480. Our spectrum of V1331 Cyg reveals strong emission from CO and hot water vapor, likely arising in a circumstellar disk. In comparison, our spectrum of MWC 480 appears mostly featureless. We discuss possible ways in which strong water emission from MWC 480 might go undetected in our data. If strong water emission is in fact absent from the inner disk, as our data suggest, the continuum excess interior to the dust sublimation radius that is detected in the interferometric data must have another origin. We discuss possible physical origins for the continuum excess. ",
        "introduction": "CO overtone emission and hot water emission in the $K$-band are both recognized to probe the hot ($>1500$\\,K) inner gaseous disk surrounding young stars. High spectral resolution studies of CO overtone emission have been used to probe disks surrounding low mass T Tauri stars, intermediate mass Herbig Ae stars, and high mass stars (e.g., Carr et al.\\ 1993; Chandler et al.\\ 1993; Najita et al.\\ 1996; Thi \\& Bik 2005; Blum et al.\\ 2004; Thi et al.\\ 2005; Berthoud et al.\\ 2007). High resolution spectroscopy of water emission, both in the $K$-band (Carr et al.\\ 2004; Najita et al.\\ 2000; Thi \\& Bik 2005) and at longer infrared wavelengths (Salyk et al.\\ 2008; Knez et al.\\ 2007), has also been used to probe the properties of inner gaseous disks.  More recently, it has been argued based on spectrally-dispersed interferometric observations that strong water emission is responsible for a significant continuum excess inward of the dust sublimation radius in a disk surrounding a Herbig Ae star (Eisner 2007). Several recent interferometric studies have reported similar evidence for a bright compact source of near-infrared emission from Herbig AeBe stars (Eisner et al.\\ 2007; Kraus et al.\\ 2007; Tannirkulam et al.\\ 2008a,b; Isella et al.\\ 2008). The compact continuum emission is attributed by these authors to emission from a gaseous disk interior to the dust sublimation radius. Such a gaseous disk is expected to show significant spectral structure due to emission lines of CO and water and other opacity sources (e.g., Muzerolle et al.\\ 2004; Eisner 2007).  Here we report high resolution $K$-band spectroscopy of two young stars, MWC 480 and V1331 Cyg. MWC 480 (HD 31648; A3pshe+) is a Herbig Ae star with a stellar mass of $2.3\\Msun$ that is surrounded by a rotating disk of gas and dust (Mannings et al.\\ 1997). CO fundamental emission is detected from the source (Blake \\& Boogert 2004). Spatially resolved millimeter and infrared interferometric studies of the emission from MWC 480 derive an inclination of $i=$26--38 degrees for the disk  (Simon et al.\\ 2000; Eisner et al.\\ 2004). Spectrally dispersed interferometric observations find evidence for an emission excess located interior to the dust sublimation radius with a spectral shape that suggests emission from hot water vapor (Eisner 2007). Evidence for active accretion in this source includes the presence of associated HH objects and emission lines of Si IV (1394\\AA, 1403\\AA) that arise from hot accreting gas (Sitko et al.\\ 2008; also Valenti et al.\\ 2000 and Muzerolle et al.\\ 2004).  The other source in our study, V1331 Cyg (LkHa 120), is a pre-main-sequence star in the L988 dark cloud complex. The distance to L988 (see Herbig \\& Dahm 2006 for a summary) has been estimated as 700-800\\,pc based on photometry and spectroscopy of nebulous sources in the L988 region (Chavarria 1981; Chavarria-K \\& de Lara 1981). Studies of the extinction in the region toward the clouds determine a closer distance for the clouds of 550\\,pc (Shevchenko et al.\\ 1991) and $500\\pm100$\\,pc Alves et al.\\ (1998).  Various studies suggest a spectral type for V1331 Cyg (A8-G5) that falls between those of Herbig Ae stars and typical T Tauri stars (Kuhi 1964; Chavarria 1981; Hamann \\& Personn 1992). V1331 Cyg has an associated molecular outflow (Levreault 1988) and a massive circumstellar disk ($0.5\\Msun$ assuming 550\\,pc distance; McMuldroch et al.\\ 1993; see also Weintraub et al.\\ 1991) that is surrounded by a flattened gaseous envelope and a ring of reflection nebulosity viewed at an inclination of $i \\sim 30$ degrees. V1331 Cyg is known to show emission in CO overtone bands (Carr 1989). High resolution $L$-band spectroscopy that shows water and OH emission from V1331 Cyg has also been presented previously (Najita et al.\\ 2005). ",
        "conclusions": "V1331 Cyg shows a rich spectrum of water emission in the $K$-band. In comparison, the $K$-band spectrum of MWC 480 shows little emission from water or any other spectral lines. The non-detection of water emission is consistent with the absence of CO overtone emission from MWC 480: all sources of $K$-band water emission reported in the literature also show CO overtone emission (Carr et al.\\ 2004; Thi \\& Bik 2005; Najita et al.\\ 2000; this paper; see also van Boekel 2007).  The spectral shape of the low-resolution Keck interferometric spectrum of MWC 480 (Eisner 2007) is consistent with the lack of water emission features in the high-resolution spectrum of MWC 480. One difference between the low-resolution spectrum of MWC 480 and the spectrum of a source like SVS-13, which is known to show strong water emission, is that at low spectral resolution SVS-13 shows an upturn in the continuum at both the long- and short-wavelength ends of the $K$-band.  The spectral shape is interpreted as a signature of water emission from the $1.9\\micron$ and $2.7\\micron$ water bands (Carr et al.\\ 2004). In contrast, the spectrum of MWC 480 shows a significant excess over the continuum at the short wavelength end, while the long-wavelength spectrum is consistent with continuum emission only. To fit the data, Eisner (2007) used hot (2300\\,K) water emission, which produces relatively more flux at the short-wavelength end of the $K$-band than would cooler material.  Although MWC 480 shows neither CO overtone nor water emission, there is evidence for disk gas inward of the dust sublimation radius, as would be expected for a source with a significant outer disk ($0.02-0.2\\Msun$; Mannings \\& Sargent 1997; Hamidouche, Looney, \\& Mundy 2006) that is believed to be accreting (Sitko et al.\\ 2008). The high-$J$ CO fundamental emission from MWC 480 has line wings extending to $\\sim 80 \\kms$ (Blake \\& Boogert 2004). At an inclination of $i=26-38$  and a stellar mass of $2.3\\Msun$, this implies that the CO fundamental emission extends to within $0.06-0.12$\\,AU for Keplerian rotation, well within the dust sublimation radius of $\\sim 0.28$\\,AU (Eisner 2007). While there is clearly gas within the dust sublimation radius of MWC 480, the CO and water column densities and/or the temperature there are perhaps too low to produce CO overtone or water emission.  The presence of CO fundamental emission, accompanied by the absence of CO overtone and $K$-band water emission, is the situation commonly found for classical T Tauri stars. CO fundamental emission is detected from nearly all accreting T Tauri stars studied (Najita et al.\\ 2003), whereas CO overtone and $K$-band water emission are restricted to the high accretion rate systems (Najita et al.\\ 2000, 2007). CO overtone emission is also rare among high mass stars (Hanson et al.\\ 1997).  In contrast, models of gaseous inner disks that assume LTE molecular abundances and include NIR opacity sources such as water, CO, and H$^-$ predict prominent near-infrared spectral structure from Herbig Ae disks. In Muzerolle et al.\\ (2004), a dust-free gaseous inner disk is found to be capable of producing a significant $K$-band flux.  At a nominal Taurus distance of 140\\,pc, an excess of $\\sim 1$\\,Jy in the $K$-band, the flux of the hot compact continuum in MWC 480, can be produced in disks with accretion rates of $\\sim 10^{-8}\\Msunperyr$. Although such a model can plausibly account for the magnitude of the compact excess that is observed in MWC 480, the excess is also predicted to show strong $K$-band emission features of CO and water, which are not observed.  The models further predict prominent CO and water features over a wide range in accretion rate ($10^{-8}-10^{-5}\\Msunperyr$; Muzerolle et al.\\ 2004; Calvet et al.\\ 1991). Emission is predicted at lower accretion rates and absorption at higher accretion rates. The relative rarity of CO overtone and water features (in emission or absorption) from Herbig Ae stars suggests that something is missing from the assumed physical-thermal-chemical structure of the disk in these models.  Possibilities include non-LTE abundances, photodissociation, and missing sources of continuum opacity and ionization.    If water emission does not appear to be responsible for the excess detected inward of the dust sublimation radius in MWC 480, what is the origin of the excess? The lack of molecular emission might be explained by the presence of a competing source of continuum opacity that produces the hot compact excess. Possible additional opacity sources are high temperature condensates, H$^-$, and free-free emission. Materials such as corundum (Al$_2$O$_3$), hibonite (CaAl$_{12}$O$_{19}$), perovskite (CaTiO$_3$), and gehlenite (Ca$_2$Al$_2$SiO$_7$) are found in CI chondrites and have high sublimation temperatures of 1640--1800 K (Posch et al.\\ 2007). However, they have poor $K$-band emission efficiencies $Q_{\\rm abs} \\lesssim 10^{-3}$ (Henning et al.\\ 1991\\footnote{http://www.astro.uni-jena.de/Laboratory/OCDB/oxsul.html}; also B. Sargent, personal communication), which makes it less likely that they will contribute significantly to the $K$-band opacity.  Graphite grains, often invoked to explain the 2175\\AA\\ bump in the interstellar extinction curve, also have a high sublimation temperature ($> 2000$\\,K) at interstellar pressures (Krugel 2003; Salpeter 1977).  However, graphite may be destroyed by processes such as chemisputtering at much lower temperatures $\\sim 1200$\\,K (Lenzuni, Gail, \\& Henning 1995; Duschl, Gail, \\& Tscharnuter 1996). So it is unclear whether it can explain the observed hot, compact excess in MWC 480. Another possibility is that solids disappear through sublimation over a finite range in temperature (and disk radius). For example, in considering the balance between the solid and gas phases for magnesium-iron silicates, Duschl et al.\\ (1996) found that the gas and dust phases coexist over a range in temperature, with a fraction of silicates surviving to temperatures $\\sim 2100$\\,K at the densities of the inner disk region ($\\sim 0.1$\\,AU). Aluminum-calcium-silicates may be sustained to somewhat higher temperatures. Detailed modeling of this kind may be needed to understand the conditions under which dust grains may contribute a modest residual continuum opacity in the high temperature inner regions of disks.   A significant free-free continuum flux requires a significant electron abundance in the inner disk. For optically thin free-free emission from a $10^4$\\,K gas in a cylindrical volume that is $r =0.1$\\,AU in radius and $h = 0.1r$ in height, an electron density of $n_e \\sim 3\\times 10^{11}\\percc$ is needed to produce a continuum flux of $\\sim 1$\\,Jy at a distance of 140\\ pc. In comparison, for a steady accretion disk accreting at the rate $\\Mdot$, the disk column density is $\\Sigma = \\Mdot/3\\pi\\alpha c_s H$ where $\\alpha$ is the viscosity parameter, and $c_s$ is the sound speed. The quantity $H = c_s/\\Omega$ is the disk scale height, where $\\Omega$ is the Keplerian angular velocity. The average disk density can be estimated as $\\bar n_{\\rm H} = \\Sigma/\\mu m_H H \\simeq 5\\times 10^{15}\\percc$ for an accretion rate of $10^{-7}\\Msunperyr,$ a stellar mass $\\Mstar = 2.3\\Msun$, and a viscosity parameter $\\alpha = 0.01$. Comparing these densities, we require a high average electron fraction in the disk $\\bar x_e = n_e / \\bar n_{\\rm H} \\sim 10^{-4}$ in order to produce a $\\sim 1$\\,Jy free-free continuum flux from a disk accreting at $10^{-7}\\Msunperyr$.  Continuum opacity from H$^-$ is another possibility. At a column density $N_{\\rm H} > 10^{26}\\persqcm$ and temperature $T > 4000$\\,K, the H$^-$ continuum optical depth $\\tau > 1$ in chemical equilibrium (e.g., Najita et al.\\ 1996). Such warm temperatures are expected for inner disks with accretion rates $>10^{-7}\\Msunperyr$; the midplane H$^-$ opacity is expected to become significant under these conditions (Muzerolle et al.\\ 2004). The possibility of a high H$^-$ opacity at a column density $N_{\\rm H} > 10^{26}\\persqcm$ is also noted by Thi \\& Bik (2005). Further studies of the gas temperature, electron fraction, and H$^-$ abundance expected for the specific conditions in the inner region of the MWC 480 disk would be useful to sort out these possibilities."
    },
    "0809/0809.2677_arXiv.txt": {
        "abstract": "We use the UKIDSS Ultra-deep survey (UDS), currently the deepest panoramic near infra-red survey, together with deep Subaru optical imaging to measure the clustering, number counts and luminosity function of galaxies at $z\\sim 2$ selected using the BzK selection technique. We find that both star-forming (sBzK) and passive (pBzK) galaxies, to a magnitude limit of $K_{AB} < 23$, are strongly clustered.  The passive galaxies are the most strongly clustered population, with scale lengths of $r_0 = 15.0^{+1.9}_{-2.2}$h$^{-1}$Mpc compared with $r_0 = 6.75^{+0.34}_{-0.37}$h$^{-1}$Mpc for star-forming galaxies. The direct implication is that passive galaxies inhabit the most massive dark-matter halos, and are thus identified as the progenitors of the most massive galaxies at the present day.  In addition, the pBzKs exhibit a sharp flattening and potential turn-over in their number counts, in agreement with other recent studies. This plateau cannot be explained by the effects of incompleteness. We conclude that only very massive galaxies are undergoing passive evolution at this early epoch, consistent with the downsizing scenario for galaxy evolution. Assuming a purely passive evolution for the pBzKs from their median redshift to the present day, their luminosity function suggests that only $\\sim 2.5 \\%$ of present day massive ellipticals had a pBzK as a main progenitor. ",
        "introduction": "There is growing evidence to support the view that the most massive objects in the Universe were the first to assemble and complete their star formation (\\citealt{Kodama_etal:2004, Thomas_etal:2005, DeLucia_etal:2004, Bundy_etal:2006, Stott_etal:2007}); this phenomenon has become known as {\\it downsizing} \\citep{Cowie_etal:1999}. A number of key issues remain unresolved however. There are indications that the build up of the galaxy colour bimodality occurs around $z\\sim2$ (e.g. \\citealt{Cirasuolo_etal:2007}), but the precise evolutionary path from the distant Universe to the present day is still undetermined. The mechanism that terminates the major episode of star-formation is  poorly understood \\citep{Benson_etal:2005}, and it is now clear that massive galaxies undergo significant size evolution from $z\\sim2$ to the present day (e.g. \\citealt{Cimatti_etal:2008} and references therein).  Galaxy clustering is an important tool for investigating these populations, since the amplitude of clustering on large scales ($>1$Mpc) can provide a measurement of the dark matter halo mass \\citep{Mo_and_White:1996, Sheth_and_Tormen:1999}. In principle, therefore, one can relate galaxy populations from the distant past to the present day by tracing the evolution of dark matter halos within the context of a  framework for structure formation.  A full exploration of these issues will require large spectroscopic surveys of IR-selected galaxies over a representative volume of the distant Universe. Such surveys are currently at an early stage, so recent work has focused on the photometric colour selection of passive vs. star-forming galaxies at the crucial $z\\sim 2$ epoch. The two key methods to date are the BM/BX selection \\citep{Erb_etal:2003}, which is an extension of the Lyman-break dropout technique, and the BzK technique \\citep{Daddi_etal:2004} which is based on $K$-band selection.  Recent studies have shown that the BM/BX technique is reasonably efficient at selecting actively star-forming galaxies at $z\\sim 2$, but largely misses the passive galaxy population at these epochs \\citep{Quadri_etal:2007, Grazian_etal:2007}. In contrast, the BzK method appears to be the most complete of the broad-band techniques for the selection of both star-forming and passive galaxies \\citep{Daddi_etal:2004, VanDokkum_etal:2006, Grazian_etal:2007}. Based on initial $K$-band selection, in principle it does not suffer the same heavy biases caused by dust or the age of the stellar populations.  An alternative is to use a larger set of filters over a range of wavelengths and infer a galaxy's redshift from comparison of the calculated magnitudes with a set of templates (\\citealt{Cirasuolo_etal:2007} and references therein). This work makes use of such photometric redshifts and in principle one can derive the stellar age of a galaxy using such templates. However due to the uncertainties in determining stellar ages from the template fits, and the desire for comparison with the literature, the bulk of the present analysis is based on simple BzK selection.   The BzK technique was first developed using the K20 survey \\citep{Cimatti_etal:2002} by \\cite{Daddi_etal:2004}, and preferentially selects star-forming and passive galaxies in the redshift range $1.4 < z < 2.5$ by using the B-, z- and K$_s$-band broadband filters. Recent work has shown that the number counts of these populations differ markedly \\citep{Kong_etal:2006, Lane_etal:2007}, with star-forming galaxies being far more abundant at all magnitudes. The abundance of sBzK galaxies allowed \\cite{Kong_etal:2006} (henceforth K06) to perform a detailed clustering analysis, segregated by limiting K-band luminosity in the range $18.5 < K_{vega} < 20$. They found that the clustering of sBzKs is strongly dependent on $K$-band luminosity. \\cite{Hayashi_etal:2007} studied the clustering of sBzKs over a smaller area (180 arcmin$^{2}$) but to much greater depth (K$_{AB} < 23.2$) and confirmed this strong luminosity dependence. The sBzK population therefore appear to inhabit a range of halo masses, and are therefore likely to be the progenitors of a wide range of present day galaxies.  There have been fewer studies of {\\em passive} galaxy clustering at high redshift, in part because of their relatively low surface density.  K06 found early evidence that pBzKs are more strongly clustered than sBzKs, although the limited depth and area of this survey ($320 {\\rm arcmin}^2$ to K$_{Vega}=20$) gives rise to significant statistical uncertainly, particularly given the relatively low surface density of pBzKs at these magnitudes ($0.38$ arcmin$^{-2}$). \\cite{Blanc_etal:2008} measured the clustering and number counts of BzK-selected galaxies to the same depth as K06, but over a much wider area (0.71 deg$^2$ total between two fields). Their BzK clustering amplitudes are smaller than those of K06 but consistent at the $1\\sigma$ level. Analysis of deeper survey data over a similar area is required to confirm these results, and we present such an analysis in this work.  \\cite{Lane_etal:2007} (henceforth L07) used the Early Data Release (EDR) from the UKIDSS UDS to investigate the number counts and overlap in colour-selected galaxies, including those selected by the BzK technique. They found that the flattening in pBzK number counts observed by K06 extends to become a plateau at the EDR depth of $K_{AB} < 22.5$ and noted that if this were to continue to turn over it could imply an absence of high-redshift passive galaxies at low luminosities.  In this work we investigate the number counts and clustering of BzK-selected galaxies over a substantially greater depth and area than any previous study. Section 2 describes the data and galaxy number counts. In Section 3 we present an angular clustering analysis, which is extended by de-projection to infer the real-space correlation lengths. In Section 4 we use photometric redshifts to derive a luminosity function for these populations, and compare with the present-day galaxy luminosity function. A discussion and conclusions are presented in Section 5. Throughout this paper we assume a flat $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$ and $h = 0.73$.   \\section[]{Source identification and number counts}  \\subsection{Data set and sample definitions}  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/UDS_DR1_BzK_v1.3_daddifilters.ps} \\caption{BzK colour-colour diagram for sources in the UDS DR1. The sBzK and pBzK selection regions are marked accordingly. Also visible are the passive galaxy track identified in \\protect\\cite{Lane_etal:2007}, and the stellar locus used to match our photometric system to that of \\protect\\cite{Daddi_etal:2004} (see text). We show mean error-bars for the galaxies in the BzK selection regions.} \\label{bzk} \\end{center} \\end{figure}  The UKIRT Infrared Deep Sky Survey (UKIDSS) has been underway since spring 2005 and comprises 5 sub-surveys covering a range of areas and depths \\citep{Lawrence_etal:2007}. This survey has been made possible by the advent of the UKIRT Wide-Field Camera \\citep{Casali_etal:2007}.  We base the present study on data from the deepest component of UKIDSS, the Ultra Deep Survey (Almaini et al. in preparation). This aims to reach final depths of $K_{AB} = 25.0, H_{AB} = 25.4, J_{AB} = 26.0$ ($5\\sigma$, point source, 2'') over an area of 0.8 deg$^2$. The size of the UDS field significantly reduces the effects of cosmic variance and on this scale is the deepest near-infrared survey to date. For this work we use the UDS DR1 release \\citep{Warren_etal:2007}, which reaches $5\\sigma$ (point source) depths of K$_{AB} = 23.5$ and J$_{AB} = 23.7$. For details of the completeness estimation, image stacking, mosaicing and catalogue extraction procedures see \\cite{Foucaud_etal:2007} and Almaini et al.  (in preparation). In addition to these data, deep $B,V,R,i^{\\prime},z^{\\prime}$ imaging is also available from Subaru Suprimecam to limiting depths of $ B_{AB} = 28.4, ~V_{AB} = 27.8, ~R_{AB} = 27.7, ~i_{AB}^{\\prime} = 27.7$ and $ z_{AB}^{\\prime} = 26.7$ \\citep{Furusawa_etal:2008}. The UDS and Subaru survey areas are not entirely coincident, which reduces our usable area to 0.63 deg$^2$ in this analysis.  \\subsection{Construction of the BzK colour-colour diagram}   The original BzK selection technique was introduced in \\citet{Daddi_etal:2004} based on photometry in the Bessel $B$-band, Gunn $z$-band and $K_s$-bands as defined on the FORS1, FORS2 and ISAAC instruments at the Very Large Telescope (VLT).  To correct our colours to these filters we used the stellar track of K06, who in turn used published empirical stellar spectra from \\citet{Pickles:1998} and synthetic stellar spectra from \\citet{Lejeune_etal:1997}, convolved with the filter responses of the VLT instruments used in \\citet{Daddi_etal:2004}. These derived colours provide a convenient reference for matching stars on the well-defined BzK stellar locus. We use the adjusted stellar locus of K06 in the following way.  The stellar locus can be split into a main branch and a ``knee'' feature, clearly visible in the lower part of figure \\ref{bzk}. The intersection of these branches provides a fixed point in the BzK plane from which we can derive a quantitative adjustment. Using a least squares fit to each section, we derived the following offset in $z-K$ and $B-z$ colours to this fixed point to convert to the photometric system of \\citet{Daddi_etal:2004}: \\begin{equation} (z-K)_{Daddi} = (z^{\\prime}-K)_{UDS} - 0.26 \\end{equation} \\begin{equation} (B-z)_{Daddi} = (B-z^{\\prime})_{UDS} + 0.06 \\end{equation}  This correction is then applied to all sources in the BzK plane. We henceforth refer to BzK photometry defined in this way.  The standard BzK definitions of \\citet{Daddi_etal:2004} were then used to construct our BzK sample.  K-band sources ($> 5 \\sigma$) from the UDS were used as the primary catalogue, with optical magnitudes extracted directly from the Subaru imaging data after careful matching of the Subaru and UDS astrometric frames. Aperture magnitudes were then extracted using a 2-arcsec diameter. Only sources outside the contaminated halos of saturated optical stars were used.  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/DR1_bzk_K_number_counts_worstcase.ps} \\caption{Differential number counts of passive and star-forming BzK galaxies with apparent K-band magnitude, compared with the full-sample of $K$-selected galaxies. The error-bars shown are standard Poisson errors on the counts. Two sets of values are shown for pBzKs: the standard selection of pBzKs (filled, red triangles); and the worst case sample described in the text (open, green triangles). Numerical values are reproduced in table \\protect\\ref{tab-counts}.  Literature values from K06 and \\protect\\cite{Blanc_etal:2008} are also shown (dot-dashed lines and dashed lines respectively). The plateau identified in previous works (K06, L07, \\protect\\citealt{Blanc_etal:2008}) and subsequent turn-over prior to our magnitude limit are apparent in the pBzK number counts.} \\label{counts} \\end{center} \\end{figure}  All K-band sources with $K_{AB} < 23.5$ were used for BzK selection unless both the B- and z$^{\\prime}$-band magnitudes were fainter than the $3 \\sigma$ limits measured on those images ($0.29\\%$ of sources outside of contaminated regions), since these cannot be constrained within the BzK plane.  Overall this procedure results in the selection of 15177 sBzKs (21.7\\% of the full sample) and 742 pBzKs (1.06 \\% of the sample), 11551 and 702 respectively at $K_{AB} < 23.0$.  \\subsection{Number counts}  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/Kvszk.ps} \\caption{$(z - K)_{AB}$ colour versus K-band magnitude. The dashed line is our completeness limit of $K_{AB} = 23.5$ for the full sample, while the solid line shows how the completeness limit in z affects $(z - K)$ colour. We begin to be incomplete above $(z - K) > 2.94$. The impact of this incompleteness is discussed in the text.} \\label{kzk} \\end{center} \\end{figure}  \\begin{table} \\caption{Differential Number Counts in log(N/deg$^{2}$/0.5mag) bins for sBzK and pBzK in the UDS DR1.} \\begin{tabular} {|l|l|l|l|l|} \\hline K bin&all&sBzK&pBzK&Worst\\\\ centre&sources& & &case pBzK\\\\ \\hline 17.41&2.832&0.198&-&-\\\\ 17.91&2.803&0.676&-&-\\\\ 18.41&2.957&0.198&-&-\\\\ 18.91&3.203&0.500&-&-\\\\ 19.41&3.380&1.403&0.676&0.676\\\\ 19.91&3.556&2.025&1.429&1.429\\\\ 20.41&3.720&2.504&1.914&1.923\\\\ 20.91&3.872&2.982&2.181&2.185\\\\ 21.41&3.998&3.315&2.386&2.394\\\\ 21.91&4.121&3.592&2.386&2.394\\\\ 22.41&4.229&3.778&2.377&2.431\\\\ 22.91&4.285&3.849&2.158&2.394\\\\ 23.41&4.258&3.536&1.560&2.158\\\\ 23.91&3.942&0.801&-&-\\\\ \\hline \\end{tabular} \\label{tab-counts} \\end{table}  The differential K-band number counts are illustrated in Figure \\ref{counts}, and tabulated in Table 1, for both the full sample of galaxies and those selected as passive and star-forming BzKs. We find a steep rise in the counts of star-forming sBzKs toward fainter magnitudes.  In contrast, pBzK number counts exhibit an apparent flattening at $K\\sim 21$, consistent with the findings of K06 and L07. Since these galaxies are sampled over a relatively narrow redshift range, this may imply that the passive population consists largely of luminous galaxies at this early epoch (see Section 4).  To investigate the reality of the turn-over, we note that the galaxies around and beyond this feature have very faint B-band magnitudes, with a substantial fraction (62\\%) below the $3\\sigma$ detection limit in this band ($B_{lim} = 28.4$). A non-detection in the $B$-band does not affect their classification, however, and merely pushes them red-ward in B-z colour in the BzK diagram illustrated in Figure 1. The more worrying class are the objects which are below the $3\\sigma$ limit in $z'$ {\\em and} $B$, which cannot be assigned BzK classification (164 objects). Figure \\ref{kzk} shows that even within our K-band limit, we become incomplete above $(z - K) \\simeq 3$. We study the effect of this incompleteness by considering the extreme case in which all of the objects with B- and $z^{\\prime}$-band magnitudes fainter than the $3\\sigma$ limits (28.4 and 26.7 respectively) are pBzKs. In this 'worst case' sample the number counts still exhibit a plateau, as before, but with the absence of a turn-over.  At our conservative limit of $K_{AB}< 23$ our completeness is $> 95$\\% for compact galaxies (Almaini et al., in prep.), so the feature is unlikely to be due to  $K$-band incompleteness unless there is unprecedented size-evolution in the pBzK population towards fainter magnitudes compared to the general population.  The extreme compactness of passive galaxies at these redshifts observed by \\cite{Cimatti_etal:2008} suggests that of the two classes, the sBzKs should suffer more from such incompleteness. There is no evidence from their number counts to suggest that they are adversely affected in this way, so we expect that the pBzKs are likewise unaffected below this magnitude.  Photometric errors are another source of concern.  Adjusting the boundaries for pBzK selection by the mean photometric errors for sources close to the boundaries ($0.1$ magnitudes), we found that such errors had no noticeable impact on the decline in the pBzK population. Additionally, since the pBzK region of the colour-colour plane is sparsely populated, photometric errors are likely to scatter fainter objects {\\em into} pBzK selection, rather than the opposite.  The faint pBzK counts are thus likely to be slightly {\\em overestimated} because of photometric errors.  We conclude that the flattening in the pBzK population is likely to be real and the subsequent turn-over is probable but not certain. We note that the turnover corresponds to absolute magnitude $ M_K = -23.6$ at z = 1.4, which is close to the value of $L^*$ determined for the K-band luminosity function at these redshifts \\citep{Cirasuolo_etal:2007}.  Taken at face value, these results are suggestive of a sharp decline in the number density of pBzK galaxies towards lower luminosities, consistent with expectation from downsizing.  As correctly pointed out in \\cite{Grazian_etal:2007} and predicted by \\cite{Daddi_etal:2004}, this is not necessarily equivalent to a true decline in the number density of passive galaxies, since the efficiency and completeness of the pBzK technique is largely untested at such faint magnitudes. In particular, there is some evidence that passive galaxies can also be found among the redder sBzK galaxies (at large values of $z-K$), with a possible incompleteness as high as $34\\%$ \\cite{Grazian_etal:2007}. We find that randomly adding an additional $34$\\% of galaxies to the pBzK sample from this part of the diagram is indeed sufficient to remove the apparent turn-over, since these are predominantly the fainter sBzK galaxies.  The sBzK number counts are significantly higher than those of K06, \\cite{Blanc_etal:2008} and \\cite{Imai_etal:2008} (not plotted); while the pBzK number counts are lower. The UDS field is more than 6 times larger than that of the K20 survey and 4 times larger than the one used in \\cite{Imai_etal:2008}. The combined field used in \\cite{Blanc_etal:2008} is of a similar size to the UDS, cosmic variance is the most likely explanation for the difference in number counts in this case.   \\section[]{Clustering properties}  \\subsection{Angular clustering}  The 2-point angular correlation function, $w(\\theta)$ is defined by the joint probability of finding two galaxies in solid angle elements $\\delta\\Omega_1$ and $\\delta\\Omega_2$ at a given separation \\citep{Peebles:1980}, \\begin{equation} \\delta P = n^2 \\delta\\Omega_1 \\delta\\Omega_2 (1 + w(\\theta_{12})). \\end{equation}  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/clustering_deconv_new_BzK.ps} \\caption{The angular correlation function for our BzK-selected galaxy samples. The best-fitting power laws are shown, with slopes fixed to the fiducial value of $\\delta=0.8$. The pBzK galaxies are very strongly clustered, much more so than sBzKs, indicating that they occupy the most massive dark matter halos at their epoch.} \\label{wth} \\end{center} \\end{figure}  To estimate the correlation function we use the estimator of \\cite{Landy_and_Szalay:1993},  \\begin{equation} w(\\theta) = \\frac{DD - 2DR + RR}{RR} \\end{equation}  where DD, DR and RR are the counts of data-data, data-random and random-random pairs respectively at angular separation $\\theta$, normalised by the total number possible. Although this estimator is relatively robust against systematic errors, there remains a small bias due to the finite field size, which is corrected for an integral constraint by a constant, C. We follow the method of \\cite{Roche_and_Eales:1999} by using the random--random counts to estimate the size of this bias:  \\begin{equation} C=\\frac{\\Sigma N_{RR}(\\theta)\\theta^{-0.8}}{\\Sigma N_{RR}(\\theta)}, \\end{equation}  where the sums extend to the largest separations within the field.  The sBzK and pBzK samples were selected to a limit of $K_{AB} < 23$. We adopted this conservative magnitude limit to ensure the minimum contamination by spurious sources. We fit a single power law of the form $w(\\theta) = A(\\theta^{-\\delta} - C)$ to the corrected data over the separation range $0.01 - 0.1$ degrees, fixing the slope to the fiducial $\\delta=0.8$ and minimising the $\\chi^2$. The errors on the measurements are a combination of those due to shot noise (estimated by bootstrap resampling) and an estimate of those due to cosmic variance. The estimates for cosmic variance were found by splitting the field into 4 and computing the variance of $w(\\theta)$.  In order to fit the clustering reliably, we wish to avoid small-scale excesses due to multiple galaxy occupation of a single halo. The lower bound was chosen to correspond to $\\sim 0.9$Mpc (co-moving) at z$\\sim 2$ for this reason. The upper bound is a conservative estimate of the limit to which our data have enough signal for a reliable fit to be obtained.  For sBzK galaxies the amplitude was found to be $A=1.79^{+0.17}_{-0.17} \\times 10^{-3}$ (deg.$^{0.8}$)and for the pBzK population $A = 6.37^{+1.58}_{-1.54} \\times 10^{-3}$ (deg.$^{0.8}$). The clustering amplitude of the pBzKs is inconsistent with the clustering of sBzK galaxies at the 3-sigma level. Figure \\ref{wth} shows the clustering measurements corrected for the integral constraint for the sBzK and pBzK galaxies.  \\subsection{De-projected clustering amplitude}  The real space clustering and projected clustering are linked by the relativistic Limber equation \\citep{Limber:1954}.  If the redshift distribution of a sample is known, the Limber equation can be inverted and the correlation length, $r_0$, can be calculated in a robust manner \\citep{Peebles:1980, Magliocchetti_and_Maddox:1999}  To estimate the redshift distribution we use photometric redshifts based on the Subaru and UDS bands previously introduced, with the addition of Spitzer data taken as part of the SWIRE survey \\citep{Lonsdale_etal:2003}. The method is described fully in \\cite{Cirasuolo_etal:2007}, and consists of minimising the $\\chi^2$ of synthetic galaxy templates. The distributions are shown in figure \\ref{nz}. The photometric redshift distribution was then used directly in the inverted Limber's equation during the calculation of $r_0$. The values obtained in this manner are r$_0 = 15.8^{+2.0}_{-2.2}h^{-1}$ Mpc and $7.69^{+0.40}_{-0.41}h^{-1}$ Mpc for $K < 23$ pBzKs and sBzKs respectively. The quoted errors are due to the error in the fit to the clustering amplitude and therefore take into account the shot noise and cosmic variance.  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/BzK_nz_deconv.eps} \\caption{Photometric redshift distributions for passive (left) and star-forming (right) BzK-selected galaxies (solid line histogram). The over-plotted dashed lines are distributions with photometric errors de-convolved (see text).} \\label{nz} \\end{center} \\end{figure}  The photometric redshifts are subject to errors ($\\sigma/(1+z) = 0.095$ and $0.105$ for pBzKs and sBzKs respectively), however, and assuming that the true redshift distribution is highly peaked, these errors are likely to broaden the measured distribution. A broader distribution will result in a larger inferred clustering scale length. It is therefore important that we take such errors into account. We do so by assuming the errors are Gaussian, and then deconvolve the errors from the redshift distribution using a Fourier-based Wiener filter. A more detailed description of this method is provided in the appendix. The resulting de-convolved redshift distributions are shown in figure \\ref{nz}.  The scaling lengths recovered using the corrected redshift distribution are as follows: $15.0^{+1.9}_{-2.2}h^{-1}$ Mpc and $6.75^{+0.34}_{-0.37}h^{-1}$ Mpc for pBzKs and sBzKs respectively.  However, we note that our photometric redshift code is largely untested for pBzK and sBzK galaxies. This work will be refined by the use of an ongoing ESO Large Programme using VIMOS and FORS2 on the VLT to target one sixth of the UDS DR1 galaxies with photometric redshifts $> 1$.  We also confirm previous claims for a dependence of sBzK clustering on apparent magnitude \\citep{Hayashi_etal:2007}. Figure \\ref{sBzK} and Table \\ref{sBzKtab} show the values for sBzKs with varying limiting magnitude. The dependence is much stronger at magnitudes of $K < 22$, indicating a strong correlation between halo mass and sBzK luminosity for these objects.  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=240pt]{PSfiles/Lit_sBzK_K_deconv.ps} \\caption{The dependence of clustering strength of sBzK-selected galaxies on limiting K$_{AB}$-band magnitude. Our measurements (open squares) are shown together with literature values from \\protect\\cite{Blanc_etal:2008} and \\protect\\cite{Hayashi_etal:2007} (open triangle and asterisks respectively). Our values confirm the magnitude dependence of sBzK clustering strength.} \\label{sBzK} \\end{center} \\end{figure}  \\begin{table} \\caption{Values of r$_0$ for sBzKs in the UDS by limiting K$_{AB}$-band magnitude along with the number of objects brighter than the relevant limit. As for the full samples, the error ranges are derived from estimates of the shot noise and cosmic variance.} \\begin{tabular} {|l|l|l|} \\hline K$_{AB,lim}$&N&r$_0$\\\\ \\hline 20.0&92&$39.0^{+11.8}_{-15.9}$\\\\ 20.5&250&$16.4^{+6.5}_{-10.0}$\\\\ 21.0&689&$12.6^{+2.7}_{-3.2}$\\\\ 21.5&1724&$8.71^{+1.49}_{-1.68}$\\\\ 22.0&3789&$5.00^{+0.49}_{-0.55}$\\\\ 22.5&7199&$6.23^{+0.64}_{-0.69}$\\\\ 23.0&11551&$6.75^{+0.34}_{-0.37}$\\\\ \\hline \\end{tabular} \\label{sBzKtab} \\end{table}  \\subsection{BzK galaxies selected at$1.4 < z < 2.5$ }  The BzK selection was defined by \\cite{Daddi_etal:2004} to isolate galaxies in the redshift range $1.4 < z < 2.5$. Clearly from figure \\ref{nz}, and as expected by \\cite{Daddi_etal:2004}, there are contaminating objects from outside of this range.  Using our photometric redshifts we can asses how successful the BzK selection technique is in reproducing the clustering of objects within the desired range.  Following the same method outlined for the full samples, we compute clustering amplitudes find correlation lengths of $11.1^{+1.7}_{-1.8}h^{-1}$ Mpc and $5.46^{+0.36}_{-0.37}h^{-1}$ Mpc for $1.4 < z < 2.5$ pBzKs and sBzKs respectively. These values are slightly lower than those for the full samples and indicate that the high redshift tails in the full sample are at least as highly clustered as those within the $1.4 < z < 2.5$ range. However, the conclusions drawn from the full sample are still valid, namely that the pBzK galaxies are significantly more strongly clustered than the sBzK galaxies. ",
        "conclusions": ""
    },
    "0809/0809.0996_arXiv.txt": {
        "abstract": "{} {We investigated in detail nine sources in the direction of the young $\\sigma$~Orionis cluster, which is considered a unique site for studying stellar and substellar formation. The nine sources were selected because of some peculiar properties, such as extremely red infrared colours or too strong H$\\alpha$ emission for their blue optical colours.} {We took high-quality, low-resolution spectroscopy (R $\\sim$ 500) of the nine targets with ALFOSC at the Nordic Optical Telescope. We also re-analyzed [24]-band photometry from MIPS/{Spitzer} and compiled the best photometry available at the $V i J H K_s$ passbands and the four IRAC/{Spitzer} channels for constructing accurate spectral energy distributions covering from 0.55 to 24\\,$\\mu$m.} {The nine targets were classified into: one Herbig Ae/Be star with a scatterer edge-on disc, two G-type stars, one X-ray flaring, early-M, young star with chromospheric H$\\alpha$ emission, one very low-mass, accreting, young spectroscopic binary, two young objects at the brown dwarf boundary with the characteristics of classical T~Tauri stars, and two emission-line galaxies, one undergoing star formation, and another one whose spectral energy distribution is dominated by an active galactic nucleus. Besides, we discover three infrared sources associated to overdensities in a cold cloud in the cluster centre.} {Low-resolution spectroscopy and spectral energy distributions are a vital tool for measuring the physical properties and the evolution of young stars and candidates in the $\\sigma$~Orionis cluster.} ",
        "introduction": "\\label{introduction}  \\begin{figure*} \\centering \\caption{False-colour composite images roughly centred on $\\sigma$~Ori at different wavelengths. From left to right: optical photographic $B_J + R_F + I_N$, 2MASS $J + H + K_{\\rm s}$, and {\\em Spitzer} $[3.6] + [8.0] + [24]$ (blue + green + red). Approximate size is $30 \\times 30$\\,arcmin$^2$; north is up, east is left. The nine analyzed targets are encircled. Note the relative dimming of the Trapezium-like star system and the appearence of a large dust cloud eastwards of it at the reddest wavelengths (see Sect.~\\ref{sigOriIRSs}). Colour version of all our figures are available in the electronic publication. {\\bf Note: the images are available only in the A\\&A version.}} \\label{im_all} \\end{figure*}  The \\object{$\\sigma$~Orionis cluster} is one of the very few star-forming regions within 1\\,kpc with massive stars (B2 and earlier -- Zinnecker \\& Yorke 2007). In particular, the Trapezium-like system \\object{$\\sigma$~Ori}, that gives the name to the cluster, contains five known stars with masses above $\\sim$7\\,$M_\\odot$ (A, B, D, E, and the recently confirmed F component -- Caballero 2007a and references therein; D.~M. Peterson et~al., in~prep.). The nearby \\object{Horsehead Nebula}, which {\\em mane} is illuminated by the Trapezium-like system, acts as a retaining wall to the east, while approximately 80\\,\\% of the cluster members are contained in a circular pool of only 20\\,arcmin radius, where the extinction is very low (Lee 1968; B\\'ejar et~al. 2004; Caballero 2008a). Besides, although there are still uncertainties on its actual heliocentric distance ($d \\sim$ 350--440\\,pc -- Caballero 2008b; Mayne \\& Naylor 2008; Sherry et~al. 2008), $\\sigma$~Orionis has been classified as the closest massive-star forming region (Zinnecker \\& Yorke 2007). Furthermore, with $\\sim$3\\,Ma (Zapatero Osorio et~al. 2002a; Oliveira et~al. 2006; Caballero 2008b), the cluster has an age intermediate between those of extremely young, highly embedded ($\\sim$1\\,Ma: \\object{Orion Nebula Cluster}, \\object{NGC~2024}) and relatively cleared, still very young clusters ($\\sim$5--10\\,Ma: \\object{$\\lambda$~Orionis}, \\object{NGC~2023}) at similar heliocentric distances ($d \\sim$~400\\,pc).  The low absorption, relative closeness, and youth make $\\sigma$~Orionis to be one of the most fruitful hunting grounds for substellar objects (B\\'ejar et~al. 1999, 2001, 2004; Scholz \\& Eisl\\\"offel 2004; Kenyon et~al. 2005). Notably, the cluster accounts for more than a half of all known {\\em confirmed} isolated planetary-mass objects in the literature (Zapatero Osorio et~al. 2000, 2002b, 2002c, 2007; Barrado y Navascu\\'es et~al. 2001; Mart\\'{\\i}n et~al. 2001a; Gonz\\'alez-Garc\\'{\\i}a et~al. 2006; Caballero et~al. 2007; Scholz \\& Jayawardhana 2008; Luhman et~al. 2008). The reader can find in Caballero (2008c) additional cluster data and useful references for going deeply into all subjects that have been investigated in $\\sigma$~Orionis (from X-ray emission, through Herbig-Haro objects, to brown dwarfs with~discs).  For a deeper understanding of the processes that take place in the cluster, Caballero (2008c) tabulated, in his Mayrit catalogue, 241 $\\sigma$~Orionis stars and brown dwarfs with known feaures of youth (e.g. OB spectral types, Li~{\\sc i} $\\lambda$6707.8\\,{\\AA} in absorption, abnormal strength of other alkali lines because of low surface gravity, H$\\alpha$ $\\lambda$6562.8\\,{\\AA} in strong emission due to accretion, infrared flux excess due to surrounding disc), 97 candidate cluster members, and 115 back-  and foreground sources. Of the 338 cluster members and member candidates, 54 are fainter than the stellar/substellar boundary at $J \\approx$ 14.5\\,mag for null extinction (Caballero et~al. 2007). The Mayrit catalogue is to date the most comprehensive database of $\\sigma$~Orionis members, and is being used as an input to investigate the cluster spatial distribution, the frequency of discs at different mass intervals, the X-ray emission, the  occurrence of wide binaries, or the initial mass function from about 18\\,$M_\\odot$ to a few Jupiter masses (Caballero 2008b; Luhman et~al. 2008; L\\'opez-Santiago \\& Caballero 2008). Any systematic error in the catalogue (e.g. incompleteness, contamination) may have a pernicious influence on the previous results. It is therefore necessary to confirm or discard membership in the $\\sigma$~Orionis cluster of many Mayrit sources that lack spectroscopy. Some of these disputable sources have been spectroscopically followed-up in the very recent works by Maxted et~al. (2008), Gatti et~al. (2008), and Sacco et~al. (2008). However, there still remain dozens cluster members and member candidates without spectroscopy. ",
        "conclusions": "We have classified our nine targets into three categories, depending on their astrophysical nature. The classes are early-type and solar-like stars (3), late-type young stars and brown dwarfs (4), and emission-line galaxies at intermediate redshifts (2). In the last columns of Table~\\ref{observed.targets}, we list the spectral types, equivalent widths of H$\\alpha$, and cluster membership measured by us of the seven young object candidates. Actually, we measured the pseudo-equivalent width (pEW) of the lines with respect to the pseudo-continuum of the four M-type stars and brown dwarfs. This definition extends to other lines different from the Balmer series (e.g. Li~{\\sc i}). Hereafter, we use the term equivalent width (EW) for simplicity. The error in the EWs comes from the uncertainty in the determination of the (pseudo-)continuum.  The seven young object candidates and the two galaxies are discussed in detail in the following sections.  \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\caption[]{Names, coordinates, and basic spectroscopic data of observed targets.} \\label{observed.targets} \\centering \\begin{tabular}{ll cc l lc lc l} \\hline \\hline \\noalign{\\smallskip} Name$^{a}$\t\t\t& Alternative\t\t\t& $\\alpha$ \t& $\\delta$      & Reference(s)$^{b}$\t & Date \t & $t_{\\rm exp}$\t& Sp.  \t \t\t& EW(H$\\alpha$)\t& Cluster \\\\ % & name\t\t\t\t& (J2000) \t& (J2000)       &\t\t       \t &\t\t & (s)  \t\t& type \t \t\t& [\\AA]\t\t& member  \\\\ % \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\object{Mayrit~459340}\t\t& StHA~50\t\t\t& 05 38 34.45 \t& --02 28 47.6  & St86, DK88, Ca07a      & 2006 Oct 15   & 30\t\t\t& A2--6 V e\t\t& --10$\\pm$1\t& Yes     \\\\ % \\object{Mayrit~968292}\t\t& GSC~04771--00962\t\t& 05 37 44.92\t& --02 29 57.3  & Ca06, Ca07a \t       \t & 2006 Oct 14   & 60\t\t\t& G8--K0 V\t\t& +2.5$\\pm$0.5 \t& Yes?    \\\\ % Mayrit~1064167\t\t& \\object{2MASS~J05390088--0253164}& 05 39 00.88& --02 53 16.4  & ...\t\t       \t & 2006 Oct 12   & 150  \t\t& G2--5 V\t\t& +3.5$\\pm$0.5 \t& No?     \\\\ % \\object{Mayrit~797272}\t\t& [W96]~rJ053751--0235\t\t& 05 37 51.61 \t& --02 35 25.7  & Sh04, Fr06, Ca06       & 2006 Oct 13   & 300  \t\t& M1--3 V e\t\t&--4.5$\\pm$0.5 \t& Yes     \\\\ % \\object{Mayrit~102101}AB\t& [W96]~rJ053851--0236\t\t& 05 38 51.45 \t& --02 36 20.6  & Wo96, Fr06, He07, Sa08 & 2006 Oct 11, 13& 300+500\t\t& M2--4 V e\t\t& --32$\\pm$2\t& Yes     \\\\ % \\object{Mayrit~264077}\t\t& S\\,Ori~J053902.1--023501\t& 05 39 01.94 \t& --02 35 02.9  & He07, Ca07\t       \t & 2006 Oct 11, 12& 900+900\t\t& M2--4 V e\t\t& --72$\\pm$4\t& Yes     \\\\ % \\object{Mayrit~358154}\t\t& S\\,Ori~J053855.4--024121\t& 05 38 55.42 \t& --02 41 20.8  & He07, Ca07\t       \t & 2006 Oct 13, 15& 1200+3$\\times$1200  & M4--6 V e \t\t&--190$\\pm$20  \t& Yes     \\\\ % \\object{UCM~0537--0227}\t\t& Mayrit~926051\t\t\t& 05 39 32.70 \t& --02 26 15.4  & Ca08  \t       \t & 2006 Oct 11, 12& 1200+1200\t\t& ...$^{c}$\t\t& ...$^{c}$\t& No\t   \\\\ % \\object{UCM~0536--0239}\t\t& [HHM2007] 668\t\t\t& 05 38 41.24 \t& --02 37 37.7  & He07, Ca07b\t       \t & 2006 Oct 11, 12& 1200+1200\t\t& ...$^{c}$\t\t& ...$^{c}$\t& No\t   \\\\ % \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{a}$] Mayrit and UCM designations follow the nomenclatures introduced in Caballero (2007b) and Zamorano et~al. (1994), respectively (truncated coordinates in the UCM~HHMM+DDMMW designation are for the 1950.0 equinox). \\item[$^{b}$] References -- St86: Stephenson (1986); DK88: Downes \\& Keyes (1988); Sh04: Sherry, Walter \\& Wolk (2004); Fr06: Franciosini, Pallavicini \\& Sanz-Forcada (2006); He07: Hern\\'andez et~al. (2007); Ca07a: Caballero (2007a); Ca07b: Caballero (2007b). Ca07: Caballero et~al. (2007); Ca08: Caballero (2008c); Sa08: Sacco et~al. (2008). \\item[$^{c}$] Not applicable for galaxies UCM~0537--0227 and UCM~0536--0239. \\end{list} \\vfill \\end{minipage} \\end{sidewaystable*}  \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\caption[]{Photometry of observed targets.} \\label{photometry.targets} \\centering \\begin{tabular}{l c c ccc cccc c c} \\hline \\hline \\noalign{\\smallskip} Name\t\t\t& $V \\pm \\delta V$ \t& $i \\pm \\delta i$ \t& $J \\pm \\delta J$      & $H \\pm \\delta H$      & $K_{\\rm s} \\pm \\delta K_{\\rm s}$ & [3.6]\t& [4.5]\t\t\t& [5.8]\t\t\t& [8.0]\t\t\t& [24.0]\t\t& Sources of \\\\ & [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t        & [mag] \t        & [mag] \t\t& [mag] \t\t& [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t\t& photometry$^{a}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Mayrit~459340\t\t& 11.28$\\pm$0.05      \t& 11.004$\\pm$0.040      & 10.666$\\pm$0.024      & 10.614$\\pm$0.024      & 10.545$\\pm$0.024\t& 10.40$\\pm$0.03\t& 10.29$\\pm$0.03\t& 10.07$\\pm$0.03\t&  9.88$\\pm$0.04\t& 8.90$\\pm$0.55\t\t& A122244440\t\\\\ % Mayrit~968292\t\t& 11.06$\\pm$0.07      \t& 10.747$\\pm$1.000      & 10.202$\\pm$0.026      &  9.960$\\pm$0.022      &  9.905$\\pm$0.022\t&  9.91$\\pm$0.03\t&  9.90$\\pm$0.03\t&  9.91$\\pm$0.03\t&  9.89$\\pm$0.03\t& $\\gtrsim$9.2$^b$\t& A122233330\t\\\\ % Mayrit~1064167\t\t& 12.65$\\pm$0.12\t& 11.912$\\pm$0.030      & 11.451$\\pm$0.024      & 11.184$\\pm$0.023      & 11.101$\\pm$0.024\t& 11.04$\\pm$0.04\t& 10.97$\\pm$0.04\t& 11.01$\\pm$0.05\t& 10.95$\\pm$0.07\t& $\\gtrsim$9.2\t\t& B122244440\t\\\\ % Mayrit~797272\t\t& 15.64$\\pm$0.01      \t& 13.446$\\pm$1.000      & 11.894$\\pm$0.026      & 11.172$\\pm$0.023      & 10.977$\\pm$0.022\t& 10.81$\\pm$0.03      \t& 10.80$\\pm$0.03      \t& 10.79$\\pm$0.03      \t& 10.72$\\pm$0.03      \t& $\\gtrsim$9.2\t\t& 3122233330\t\\\\ % Mayrit~102101AB\t\t& 17.68$\\pm$0.09      \t& 14.248$\\pm$0.030      & 12.438$\\pm$0.027      & 11.813$\\pm$0.023      & 11.552$\\pm$0.025\t& 11.23$\\pm$0.03      \t& 11.16$\\pm$0.03      \t& 11.12$\\pm$0.03      \t& 10.96$\\pm$0.03      \t& 3.44$\\pm$0.27$^c$\t& 3122233330\t\\\\ % Mayrit~264077\t\t& 19.08$\\pm$0.13      \t& 16.380$\\pm$0.080      & 14.445$\\pm$0.035      & 13.378$\\pm$0.030      & 12.607$\\pm$0.030\t& 11.45$\\pm$0.04\t& 10.99$\\pm$0.04\t& 10.52$\\pm$0.04\t&  9.75$\\pm$0.04\t& 6.37$\\pm$0.14\t\t& C122255550\t\\\\ % Mayrit~358154\t\t& $\\gtrsim$20.0\t\t& 18.063$\\pm$0.250      & 15.622$\\pm$0.096      & 14.842$\\pm$0.050      & 13.968$\\pm$0.061\t& 12.88$\\pm$0.08\t& 12.52$\\pm$0.08\t& 12.08$\\pm$0.09\t& 11.15$\\pm$0.08\t& $\\gtrsim$9.2$^d$\t& D122255550\t\\\\ % UCM~0537--0227\t\t& $\\sim$18.5\t\t& 17.622$\\pm$0.250      & 16.141$\\pm$0.098      & 15.099$\\pm$0.071      & 14.345$\\pm$0.092\t& 13.40$\\pm$0.10\t& 13.21$\\pm$0.11\t& 12.47$\\pm$0.10\t&  9.72$\\pm$0.04\t& 6.23$\\pm$0.12\t\t& D122244440\t\\\\ % UCM~0536--0239\t\t& 18.90$\\pm$0.10\t& 17.693$\\pm$0.170      & 16.375$\\pm$0.094      & 15.253$\\pm$0.080      & 14.040$\\pm$0.058\t& 11.88$\\pm$0.03\t& 11.11$\\pm$0.03\t& 10.44$\\pm$0.03\t&  9.39$\\pm$0.03\t& 5.89$\\pm$0.12\t\t& E122233330\t\\\\ % \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{a}$] Sources of photometry -- 0: this work; 1: DENIS; 2: 2MASS; 3: Hern\\'andez et~al. (2007); 4: Caballero et~al. (2008); % 5: Caballero et~al. (2007); % A: transformed from $V_T$ Tycho-2 (H{\\o}g et~al. 2000) to $V$ Johnson using relations in ESA (1997); B: ASAS-3 (Pojma\\'nski 2002); C: Wolk (1996); D: estimated from photographic $B_J$ and $R_F$ magnitudes; E: measured using unpublished $V$ data obtained with the Wide Field Camera at the Isaac Newton Telescope, calibrated into the Johnson system through Tycho-2 comparison stars in an overlapping short exposure $V$ image in Caballero~(2007b). \\item[$^{b}$] Hern\\'andez et~al. (2007) measured $[24]$ = 9.80$\\pm$0.05\\,mag for Mayrit~968292; it is fainter than the 10\\,$\\sigma$ threshold in our data (note their under estimated errors). \\item[$^{c}$] The $[24]$ emission of Mayrit~102101AB actually comes from a spatially-coincident intracluster cloud described in the text. \\item[$^{d}$] Hern\\'andez et~al. (2007) measured $[24]$ = 8.58$\\pm$0.03\\,mag for Mayrit~264077; this emission actually comes from a nearby source (which magnitude we re-determine at $[24]$ = 8.32$\\pm$0.37\\,mag). \\end{list} \\vfill \\end{minipage} \\end{sidewaystable*}   \\subsection{Early-type and solar-like stars}  \\subsubsection{Mayrit~459340} \\label{M459340}  Mayrit~459340 (StHA~50) was a new H$\\alpha$ emission star found above 10\\,deg galactic latitude in Stephenson (1986), who investigated a large number of red-sensitive objective prism plates. Soon afterwards, Downes \\& Keyes (1988) classified Mayrit~459340 as a Be star from a spectrum taken with an image tube scanner. It went unnoticed until Caballero (2007a) listed it as one of the (30) brightest stars of the $\\sigma$~Orionis cluster (Mayrit~459340, with $M \\gtrsim$ 1.4\\,$M_\\odot$, is also one the most massive cluster stars). According to him, Mayrit~459340 deviates at the level of 3\\,$\\sigma$ to the blue of the spectro-photometric cluster sequence in a $V$ vs. $B-V$ colour-magnitude diagram from Tycho-2 magnitudes. Although its early (Be) spectral type matches, within uncertainties, its blue colours in the optical ($B_T - V_T$ = --0.06$\\pm$0.11\\,mag), % Mayrit~459340 is $\\Delta B_T$ = 1.96$\\pm$0.06\\,mag % fainter than the faintest B-type star in $\\sigma$~Orionis (HD~294275 [\\object{Mayrit~1227243}, B9V]; Caballero 2007a). Its $V_T$ magnitude better matches that of an early- or mid-F spectral-type star in the cluster, or a late-B spectral-type star at an unlikely heliocentric distance thrice larger than to $\\sigma$~Orionis. The first hypothesis would easily explain its H$\\alpha$ emission, while the second hypothesis would locate Mayrit~459340 at about 300\\,pc below the Galactic plane (the Galactic vertical scale height of OBA-type stars is only about 100\\,pc; e.g. Bahcall \\& Soneira~1980).  Likewise, Caballero (2007a) also reported that Mayrit~459340 has a disc based on 3.6--8.0\\,$\\mu$m IRAC/{\\em Spitzer} photometry from Caballero et~al. (2008). Our MIPS/{\\em Spitzer} $[24]$ photometry (tabulated in Table~\\ref{photometry.targets}) corroborated his statement: we derived colours $[3.6]-[24] = 1.5 \\pm 0.6$\\,mag and $[3.6]-[8.0] = 0.52 \\pm 0.05$\\,mag, that deviates about 3\\,$\\sigma$ and 10\\,$\\sigma$ with respect the expected values for a disc-less dwarf  ($[3.6]-[24] \\sim [3.6]-[8.0] \\sim 0$\\,mag). From the ALFOSC spectrum, shown in the top of Fig.~\\ref{spAtoG}, we measured a very intense H$\\alpha$ emission of --10$\\pm$1\\,{\\AA} and derived a spectral type in the interval A2--6V (that corresponds to a most probable mass in the interval 3--4\\,$M_\\odot$). This spectral type approximately matches the blue $B_T - V_T$ colour of Mayrit~459340 and the previous determination by Downes \\& Keyes (1988). For the classification, we investigated the depth of the blended doublets He~{\\sc i}~D$_3$ $\\lambda\\lambda$5875.6,5876.0\\,{\\AA} and Na~{\\sc i} $\\lambda\\lambda$5890.0,5895.6\\,{\\AA} and of the line Ca~{\\sc i} $\\lambda$5857.5\\,{\\AA} and compared the slope of the continuum of its spectrum with those of early-type standard stars from Valdes et~al. (2004). The strong H$\\alpha$ emission seems to persist in time, since 1978--1984 (Stephenson 1986), through 1986--1987 (Downes \\& Keyes 1988), to the present (this work). Besides, the absorption feature of the diffuse interstellar band at about 6280\\,{\\AA}, typical in early-type stars in young clusters, is quite clear. To our knowledge, there is only one scenario that simultaneously explain the A spectral type, the H$\\alpha$ emission, the abnormal blue colour for its magnitude (or, alternatively, its abnormal faintness for its colour), and the flux excess at 8--24\\,$\\mu$m: the {\\em blueing} effect in a Herbig Ae/Be (HAeBe) star with a scatterer edge-on disc (Th\\'e et~al. 1994 and references therein).  There are grounds to consider Mayrit~459340 a (long-scale) photometric variable, which is another characteristic typical of HAeBe stars. From data of the All Sky Automated Survey (ASAS-3; Pojma\\'nski 2002), we measured $\\overline{V} = 11.312$\\,mag, $\\sigma_V = 0.045$\\,mag, and $\\overline{\\delta V} = 0.032$\\,mag ($N$ = 248), consistent with the Tycho-2 $V_T$ measurement and with Mayrit~459340 lacking large-amplitude photometric variations ($\\Delta V \\gg$ 0.045\\,mag) during years 2000 to 2005. There is not evidence of photometric variability in the USNO-B1.0 catalogue, either (Monet et~al. 2003). However, from the spectral energy distribution, the DENIS $i$ magnitude, measured in 1996 Feb, seems to be $\\sim$0.6--0.7\\,magnitudes brighter than expected. The value $i = 11.00 \\pm 0.04$\\,mag is still far away from the $i \\sim 10.0$\\,mag threshold below which brighter DENIS stars in $\\sigma$~Orionis are affected by saturation and non-linear effects (Caballero 2008c). From the ratio between our flux-calibrated spectra, we also estimated that during the observations (2006 Oct), Mayrit~459340 was $\\sim$0.9\\,mag {\\em brighter} at 6200--6300\\,{\\AA} than the solar-like, young star candidate Mayrit~968292. However, from the compiled magnitudes in Table~\\ref{photometry.targets}, Mayrit~459340 is only brighter than Mayrit~968292 at $\\lambda \\gtrsim$ 8\\,$\\mu$m (where the flux excess by the disc gets more important). Furthermore, the tabulated $V_T$ magnitudes are contemporaries (early 1990s), and Mayrit~459340 was 0.17$\\pm$0.12\\,mag {\\em fainter} than Mayrit~968292 at the corresponding effective wavelength of 5320\\,\\AA. This apparent brightening at recent epochs, if accompanied by a suitable colour change, would make Mayrit~459340 to lie on the right spectro-photometric sequence of $\\sigma$~Orionis. Recalling the scatterer edge-on disc scenario of Th\\'e et~al. (1994), ``the variations in brightness and in colour [of a HAeBe star] are caused by a dust cloud in a Keplerian orbit revolving in the outer part of a star's circumstellar disc''. Most of the time, the disc of Mayrit~459340 is oriented almost edge on with respect to the observer, and the star is well obscured by the dust cloud. It is at this stage when the dimming is maximum, and scattered blue light is more easily detected. It is expected that when the dust cloud passes away from the front of the HAeBe star, it gets brighter and redder. This effect has been clearly detected in the A3e-type, variable star (with rapid variations) UX~Ori (Bibo \\& Th\\'e 1991). Mayrit~459340 is about 2.6\\,mag fainter than UX~Ori, which probably had made it to get unnoticed until now. High-resolution spectroscopy would enable better constraints on membership through measurments of radial velocity.   \\subsubsection{Mayrit~968292 and Mayrit~1064167}  On the one hand, the Tycho-2 star Mayrit~968292 (GSC~04771--00962) was the brightest (in the optical) $\\sigma$~Orionis member candidate in the Mayrit catalogue {\\em without} spectroscopy (Caballero 2008c). This was a {\\em per~se} reason for first obtaining an optical spectrum with ALFOSC. With $V \\approx$ 11.1\\,mag, it was necessary to exposure only during 1\\,min. Besides, Mayrit~968292 is the brightest component of a visual binary with a relatively small angular separation ($\\rho \\approx$ 9.8\\,arcmin, $\\theta \\approx$ 97\\,deg, $\\Delta J = 0.72 \\pm 0.03$\\,mag). The secondary, \\object{Mayrit~958292}, displayed an apparent lithium absorption feature in a spectrum taken by Caballero (2006), which confirmed the star to be a very young member of the Ori~OB1b association. If both stars formed a physical pair, then it would be one of the very few wide binaries ($r \\approx$ 3800\\,AU) detected so far in the $\\sigma$~Orionis cluster (several wide binary candidates were listed in Caballero [2007b]). See also Caballero (2007a) for a consideration on the possible X-ray emission of Mayrit~968292.  On the other hand, Mayrit~1064167 (2MASS~J05390088--0253164) was an anonymous star whose position in the $i$ vs. $i-K_{\\rm s}$ diagram made it to fall in the side of the candidate fore- and background stars, but still very close to the sequence of the cluster (see fig.~4 in Caballero~2008c). It is the brightest one in a trio of stars separated by 20--40\\,arcsec. One of the other two components of the asterism has also optical colours and magnitudes typical of solar-like cluster members (e.g. Sherry et~al. 2004). However, Caballero (2006) found no trace of lithium in a high signal-to-noise, mid-resolution spectrum of the target. The remaining star in the trio is a foreground late-type dwarf or giant based on a relatively large proper motion and quite red optical/near-infrared colours for its magnitude ($J-K_{\\rm s} = 1.30 \\pm 0.04$\\,mag; Caballero 2008c). Far for extrapolating the non-membership also to Mayrit~1064167, we found three X-ray events tabulated by the Second {\\em ROSAT} Position-Sensitive Proportional Counter (PSPC) Catalogue, 2RXP (ROSAT 2000), and the White-Giommi-Angelini version of the {\\em ROSAT} PSPC Catalogue, WGACAT (White et~al. 2000), at 10--30\\,arcsec to the star. Although the X-ray events could also be associated to a background extragalactic source, we decided to test photometric selection criteria for cluster members in the literature (in particular: Sherry et~al. 2004; Caballero~2008c) with the possible X-ray emitter Mayrit~1064167.  At a first glance at the ALFOSC spectra and the spectral energy distributions of Mayrit~968292 and Mayrit~1064167, the two stars seemed to be G-type stars with no indication of accretion or infrared flux excess. We could only impose lower limits to the $[24]$-band magnitudes. A careful comparison of the blended Na~{\\sc i}~D$_1$D$_2$ doublets and H$\\alpha$ lines indicated that Mayrit~968292 is cooler than Mayrit~1064167 (the absorption of the Na~{\\sc i} doublets and the H$\\alpha$ line increases and decreases, respectively, towards later spectral types). By comparison with several spectral-type standard stars taken from Jacoby et~al. (1984), Montes et~al. (1997), and Valdes et~al. (2004), we classified Mayrit~968292 as G8--K0V and Mayrit~1064167 as G2--G5V. We obtained a similar classification using the lines ratio H$\\alpha$/Fe~{\\sc i} $\\lambda$6495.0\\,{\\AA} (Danks \\& Dennefeld 1994). Other lines typical of solar-like stars that we observed in the spectra of Mayrit~968292 and Mayrit~1064167 were the calcium features Ca~{\\sc i} $\\lambda$6163.8\\,{\\AA} and, especially, the infrared triplet Ca~{\\sc ii} $\\lambda\\lambda\\lambda$8498.0,8542.1,8662.1\\,{\\AA} (partially embedded in the fringing region). Unfortunately, our spectra have not enough spectral resolution to clearly detect any other spectroscopic feature of youth (especially lithium).  Although Mayrit~1064167 has an earlier spectral type than Mayrit~968292, the first star is about 1.2\\,mag fainter than the second one at all passbands. Since such a large reversal of brightness is improbable, we consider that at least one of the two stars is actually a contaminant. The interloper is probably Mayrit~1064167, because Mayrit~968292 follows the cluster sequence and Mayrit~1064167 does not (it has roughly the same magnitudes of mid-K-type young stars in $\\sigma$~Orionis). Albeit Mayrit~968292 remains as a good solar-like cluster member candidate, it is necessary to obtain higher-resolution spectroscopy to confirm its membership. If its youth is verified, its membership in the Mayrit~968292--958292 binary will even so remain unknown for a longer time: because of of the location of $\\sigma$~Orionis in the solar antapex, its cluster members have very low proper motions ($\\mu <$ 10\\,mas\\,a$^{-1}$; Caballero 2007a)   \\subsection{Late-type young stars and brown dwarfs}  We mostly used the PC3 index (Mart\\'{\\i}n, Rebolo \\& Zapatero Osorio 1996; Mart\\'{\\i}n et~al. 1999) for the spectral type classification of the four M-type objects, although we also took into account the information provided by the I1, I2, and I3 indices (indicators of the strength of several CaH and TiO absorption bands; Mart\\'{\\i}n \\& Kun 1996) and some indices associated to metallic hydrides and oxides (CrH, FeH, H$_2$O, TiO, and VO). We also compared the indices values with those of spectral standards in the literature and of GJ~3517, observed with the same instrumental configuration.   \\subsubsection{Mayrit~797272}  Mayrit~797272 ([W96]~rJ053751--0235) is a moderate X-ray emitter, early-M star in $\\sigma$~Orionis. % It was first discovered by Wolk (1996), who classified it as a {\\em ROSAT} M0 star with lithium in absorption (EW(Li~{\\sc i}) = +0.32\\,\\AA) and H$\\alpha$ in faint chromospheric emission (EW(H$\\alpha$) = --6.4\\,\\AA). A decade later, Franciosini et~al. (2006) measured its X-ray count rate with {\\em XMM-Newton} and estimated an M2 spectral type from the $R-I$, $V-I$, and $I-J$ colours using the relations by Kenyon \\& Hartmann (1995) and Leggett et~al. (2001). Almost simultaneously, Caballero (2006) took a low signal-to-noise, optical spectrum where the H$\\alpha$ line appeared in faint, broadened emission and the Li~{\\sc i} line was undetected. Very recently, L\\'opez-Santiago \\& Caballero (2008) have found Mayrit~797272 to be an X-ray flaring star from new {\\em XMM-Newton} data. Finally, Mayrit~797272 forms a visual binary ($\\rho \\approx$ 10.5\\,arcsec, $\\theta \\approx$ 5\\,deg) with a slightly brighter star that does not seem to belong to the $\\sigma$~Orionis cluster.  We have derived the most probable spectral type of Mayrit~797272 in the range M1--3V from the ALFOSC spectrum shown in Fig.~\\ref{spAtoG}. Our determination is consistent with previous analyses. The H$\\alpha$ line in our spectrum of Mayrit~797272 is sharp and with an equivalent width similar to that one measured by Wolk (1996). The broad emission detected by Caballero (2006) was possibly due to the low signal-to-noise ratio and background contamination by a nearby fibre (he used a multifibre spectrograph). The chromospheric H$\\alpha$ emission and the absence of flux excess at $\\lambda \\lesssim$ 8\\,$\\mu$m suggest that Mayrit~797272 lacks a disc, which would make the star to rotate slower. The X-ray variability is easier to explain for a very young, fast-rotating, disc-free, almost completely-convective object like Mayrit~797272 (G\\\"udel et~al. 1995; Feigelson et~al. 2002; Stelzer et~al. 2004; Robrade \\& Schmitt~2005).   \\subsubsection{Mayrit~102101AB}  Mayrit~102101AB ([W96]~rJ053851--0236) was, like Mayrit~797272, first detected as a {\\em ROSAT} X-ray source by Wolk (1996). He derived a K5 spectral type and measured EW(H$\\alpha$) $\\approx$ --86\\,{\\AA} in a faint spectrum where he did not detect the Li~{\\sc i} line. The star has later been detected with {\\em XMM-Newton} and {\\em Chandra} by Franciosini et~al. (2006), Caballero (2007b), and Skinner et~al. (2008). Using IRAC and MIPS/{\\em Spitzer} data, Hern\\'andez et~al. (2007) claimed that Mayrit~102101AB harbours an ``evolved disc'' (i.e. the star exhibits smaller IRAC excesses than optically thick disc systems). Remarkably, Mayrit~102101AB is one of the very few low-mass, spectroscopic binaries (SB2) in the $\\sigma$~Orionis cluster whose orbital parameters have been derived yet ($P$ = 8.72$\\pm$0.02\\,d; Sacco et~al. 2008 -- they also derived a later spectral type, K9.5, and found lithium in both spectra).  Our observations confirm and contradict some of the observables above. First, using the same technique as for Mayrit~797272, we derived an M3$\\pm$1 spectral type for Mayrit~102101AB, which matches better the spectro-photometric cluster sequence than previous determinations. The H$\\alpha$ emission in our ALFOSC spectrum (EW(H$\\alpha$) = --32$\\pm$2\\,{\\AA}) is not so strong as measured by Wolk (1996), although still satisfying the accreting empirical criterion of Barrado y Navascu\\'es \\& Mart\\'{\\i}n (2003 -- the White \\& Basri [2003]'s criterion cannot be applied because of our poor wavelength resolution). The strong (variable?) H$\\alpha$ emission might be originated by material interchange between the two binary components or accretion from a disc. As already noticed by Hern\\'andez et~al. (2007), the IRAC $[3.8] - [8.0]$ colour of Mayrit~102101AB is only of 0.27$\\pm$0.05\\,mag. Furthermore, the apparent flux excess at the MIPS $[24]$ passband is due to an overdensity in a large intracluster dust cloud, and not to the star itself (Sect.~\\ref{sigOriIRSs}). This overdensity may also be responsible of the slight excess at $[8.0]$. The close separation between Mayrit~102101A and~B, of $(a_1 + a_2) \\sin{i}$ = 11.6$\\pm$0.3\\,$R_\\odot$ (i.e. a few objects radii\\footnote{The objects are still contracting and have radii much larger than measured for field dwarf of the same spectral type.}; Sacco et~al. 2008), prevents the formation of inner disc(s) surrounding one (or both) components. From our point of view, the origin of the X-ray and H$\\alpha$ emissions is better explained by the mutual interaction of the two binary components rather than by the presence of an ``evolved disc''. In any case, we do not completely discard the presence of a cool, {\\em circun-binary} disc at several tens solar~radii.   \\subsubsection{Mayrit~264077}  Mayrit~264077 (S\\,Ori~J053902.1--023501) is an object at the stellar/substellar boundary in $\\sigma$~Orionis, with a most probable mass of 0.061\\,$M_\\odot$ (Caballero et~al. 2004, 2007). It had, by far, the reddest near-infrared colours among the objects in the Caballero et~al. (2007)'s survey ($J-K_{\\rm s}$ = 1.84$\\pm$0.05\\,mag [2MASS], $[3.6] - [8.0]$ = 1.70$\\pm$0.06\\,mag) and fell in a differentiated location in the $[3.6] - [8.0]$ vs. $J-K_{\\rm s}$ colour-colour diagram. Simultaneously, Hern\\'andez et~al. (2007) measured $[24]$ = 6.45$\\pm$0.03\\,mag (consistent with our photometry) and sorted Mayrit~264077 as a class~II object. There was no spectroscopic information on it.  The spectral energy distribution in Fig.~\\ref{sedAtoI} complements that one shown in Caballero et~al. (2007). The presence of a circumstellar disc is obvious only from photometric data; the ALFOSC spectrum powerfully reinforces the disc existence. We determined an M3$\\pm$1 spectral type for Mayrit~264077, identical to that of Mayrit~102101AB. However, the H$\\alpha$ emission was stronger: EW(H$\\alpha$) = --72$\\pm$4\\,{\\AA}. Besides, we also found [O~{\\sc i}]~$\\lambda$6300.3\\,{\\AA} and [S~{\\sc ii}]~$\\lambda\\lambda$6716.4,6730.8\\,{\\AA} in forbidden emission (the [N~{\\sc ii}]~$\\lambda$6583.4\\,{\\AA} is blended with the strong H$\\alpha$~line).  It is worthy of notice that the $I$-band light curve of Mayrit~264077 displayed short-, mid-, and long-scale variations in the original data of Caballero et~al. (2004). Since it was one of the brightest targets in their sample and its photometry was, therefore, close to the non-linear limit of the detector, Mayrit~264077 was conservatively discarded from the detailed analysis. However, coming back to the original light curve, it is quite possible that those variations were intrinsic to the object and not due to seeing variations and non-linear effects. As shown by Caballero et~al. (2006), T~Tauri-analog brown dwarfs in $\\sigma$~Orionis with strong H$\\alpha$ emission, forbidden lines, and near-infrared flux excess also display large amplitude photometric variations at all considered time scales.  Mayrit~264077 might actually be a very low-mass T~Tauri star with a large extinction bluewards of $\\sim$1.5\\,$\\mu$m accompanied by a flux excess redwards of it (i.e. the central object is a star brighter than $J$ = 14.44$\\pm$0.04\\,mag, and not a brown dwarf).   \\subsubsection{Mayrit~358154}  Mayrit~358154 (S\\,Ori~J053855.4--024121) was another class~II, brown dwarf candidate without spectroscopy in the works by Caballero et~al. (2007) and Hern\\'andez et~al. (2007). It went unnoticed during the optical/near-infrared survey in Caballero et~al. (2004) because it fell in the glare of the 3.8\\,mag-$V$ multiple star $\\sigma$~Ori. Since Mayrit~358154 is 1.2\\,mag fainter than Mayrit~264077, we expect the first object to be a {\\em bona fide} brown dwarf, even accounting for a possible edge-on disc. To date, there is only one $\\sigma$~Orionis {\\em star} fainter than $J \\sim$ 14.5\\,mag, and it is the source of a well, long-time known Herbig-Haro object (\\object{HH~446} [Mayrit~633105] -- Reipurth et~al. 1998; Andrews et~al.~2004). However, Mayrit~358154 is even fainter in $J$ than the Herbig-Haro star. Besides, Mayrit~358154 is at 8.8\\,arcsec to the north of a brighter photometric cluster member candidate from Sherry et~al. (2004). The hypothetical primary was not, however, in the comprehensive cluster catalogue by Caballero~(2008c).  The ALFOSC spectrum and the spectral energy distribution strenghten the evidence of Mayrit~358154 having a circum(sub)stellar disc. The flux excess at the [8.0] passband is clear, although the spectral energy distribution seems to drop again at redder wavelengths (our upper limit at the [24] passband is still restricting). The distribution has a peculiar feature at the $H$ band (Mayrit~633105 should be brighter at 1.6\\,$\\mu$m) whose true origin we fail to ascertain; the feature is significative and cannot be explained by photometric error. The H$\\alpha$ emission is even larger than for the previous objects (EW(H$\\alpha$) = --190$\\pm$20\\,\\AA) and we could identify the [O~{\\sc i}], [S~{\\sc ii}], and [N~{\\sc ii}] forbidden emission lines in the studied wavelength interval. The derived spectral type, M5$\\pm$1, is consistent with Mayrit~358154 being a high-mass brown dwarf in $\\sigma$~Orionis (the earlier brown dwarf candidates known to date in the cluster have M5.5 spectral types). Although it lies on the substellar domain, Mayrit~358154 displays most of the phenomena observed in classical T~Tauri stars.   \\subsection{Background galaxies}  UCM0537--0227 was the faintest new $\\sigma$~Orionis photometric member candidate in the Mayrit catalogue of Caballero (2008c). He gave the DENIS $i$ and 2MASS $JHK_{\\rm s}$ magnitudes for the object, from where one could derive a mass in the middle of the brown dwarf domain, {\\em if} it belonged to the cluster. He also noticed that UCM0537--0227 possibly has an extended, non-stellar point spread function. Its extremely red 2MASS ($J-K_{\\rm s}$ = 2.34$\\pm$0.11\\,mag) and IRAC/{\\em Spitzer} colours measured by Caballero et~al. (2008) made us to obtain spectroscopy of it. The low-resolution optical spectrum (top panel of Fig.~\\ref{spHandI}) and the spectral energy distribution of UCM0537--0227 correspond to an emission-line, star-forming galaxy at a spectroscopic redshift z$_\\mathrm{sp}$ = $0.1009 \\pm 0.0002$ with strong PAH features and a moderate star formation rate.  UCM0536--0239 ([HHM2007]~668) and its odd spectral energy distribution from the $I$ band to [8.0] was first identified by Caballero (2006). Just afterwards, Hern\\'andez et~al. (2007) measured its [24]-band magnitude and presented it as one of the very few class~I object {\\em candidates} in $\\sigma$~Orionis. Caballero (2007b) noticed that UCM0536--0239 (incorrectly named ``Mayrit~111208'' in his work) had quite blue $I-J$ and extremely red $J-K_{\\rm s}$ colours for its faint magnitude ($J$ = 16.14$\\pm$0.10\\,mag) and found it to be an X-ray emitter detected with the ACIS and HRC-I instruments onboard {\\em Chandra}. The (sub)stellar nature of UCM0536--0239 was put into question by Caballero (2008c): he measured a possible extended point spread function in public optical images. Finally, Skinner et~al. (2008) measured the mean photon energy of UCM0536--0239 at $\\overline{E}$ = 3.40\\,keV from ACIS/{\\em Chandra} data, which is a quite high value for a normal star, but typical of active galactic nuclei. The ALFOSC spectrum (bottom panel of Fig.~\\ref{spHandI}) and the spectral energy distribution of UCM0536--0239 favour its classification as a Type~1 obscured quasi-stellar object at z$_\\mathrm{sp}$ = $0.2362 \\pm 0.0005$."
    },
    "0809/0809.5057_arXiv.txt": {
        "abstract": "All galaxies are thought to reside within large halos of dark matter, whose properties can only be determined from indirect observations.  The formation and assembly of galaxies is determined from the interplay between these dark matter halos and the baryonic matter they host.  Although statistical relations can be used to approximate how massive a galaxy's halo is, very few individual galaxies have direct measurements of their halo masses.  We present a method to directly estimate the total mass of a galaxy's dark halo using its system of globular clusters.  The link between globular cluster systems and halo masses is independent of a galaxy's type and environment, in contrast to the relationship between galaxy halo and stellar masses.  This trend is expected in models where globular clusters form in early, rare density peaks in the cold dark matter density field and the epoch of reionisation was roughly coeval throughout the Universe.  We illustrate the general utility of this relation by demonstrating that a galaxy's supermassive black hole mass and global X-ray luminosity are directly proportional to their host dark halo masses, as inferred from our new method. ",
        "introduction": "Our paradigm of galaxy formation is built upon the foundation that all galaxies form within halos of dark matter (White \\& Rees 1978; Blumenthal~et al. 1984). The evolution and growth of galaxies is governed by the merging of their dark matter halos and the conversion of gas into stars. The mass of a halo is therefore a fundamental parameter that influences a galaxy's form.  While the physics of halo growth in a cold dark matter (CDM) universe can be explored with computer simulations (e.g. Springel al.~2005), actually measuring the total dark mass content of galaxies is problematic.  Recently, great strides have been made in this area by utilising the technique of weak gravitational lensing to infer halo masses at very large radii (Tyson et al. 1984; Hoekstra~et al. 2005; Mandelbaum et al.~2006; Parker et al.~2007). This technique statistically combines the lensing signal from a large number of halos and can yield general relationships between galaxy stellar and total halo masses (Hoekstra~et al. 2005; Mandelbaum et al.~2006).  Such a relation also follows from analysis of the ``conditional luminosity function'', which is an observationally-based model connecting the stellar mass of galaxies to their hosting halos (Yang, Mo \\& van den Bosch 2003, 2008; van den Bosch et al.~2007).  Individual galaxies are expected to scatter about these general relations.  A number of other techniques exist for measuring the halo mass of {\\it individual} galaxies: rotation curves in disk galaxies (Sofue \\& Rubin 2001), X-ray gas in giant elliptical galaxies (O'Sullivan \\& Ponman 2004), kinematics of globular clusters and planetary nebulae (Romanowsky et al. 2003; Romanowsky et al.~2008) and strong gravitational lensing (Ferreras, Prasenjit \\& Williams 2005). However, such techniques all have their limitations, e.g. they are often only applicable to certain galaxy types, require equilibrium conditions, can be observationally very expensive, probe only a limited radial extent and, in the case of strong lensing, is only available for galaxies along random lines of sight.  As a result, only a very small number of individual galaxies have direct measures of their total halo mass available.  In this {\\it Letter} we present a new method for estimating the total halo mass surrounding individual galaxies based on the total mass contained in their globular cluster (GC) systems.  All large galaxies, including the Milky Way, are known to host a system of GCs.  They are among the oldest ($>11$ Gyrs) stellar populations known (Strader et al. 2005; Puzia et al. 2005; Proctor et al. 2007) and therefore trace the first stages of galaxy formation (West et al. 2004; Brodie \\& Strader 2006).  The total mass of a GC system should remain relatively constant over time because the dominant constituent of the GC system mass budget are the massive GCs, whose estimated lifespan is much longer than the age of the Universe (Mc~Laughlin 1999).  Furthermore, the overall properties of GC systems are essentially unaffected by galaxy evolutionary processes (e.g. quiescent, on-going star formation), which can sometimes dramatically change the global properties of a hosting galaxy.  Here we provide evidence for a direct correlation between the mass of a GC system and its host galaxy halo mass.  The direct proportionality implies that GC system masses can be used to approximate the halo masses of individual galaxies. ",
        "conclusions": "We have presented evidence for a direct correlation between the total GC system mass and the host galaxy halo mass.  While this result appears to be robust for intermediate-to-massive galaxies (and possibly galaxy clusters), the halo masses of low-mass galaxies are not well-constrained, and thus the correlation at low galaxy masses must be confirmed with better halo mass estimates.  The direct correlation implies GC system masses can be used to directly measure the total halo mass of individual galaxies.  This technique has the advantage that it can be applied irrespective of the host galaxy type and environment.  It is relatively inexpensive in terms of observing time, requiring only wide-field imaging under reasonable observing conditions (e.g. Spitler et al.~2008).  A few examples have been demonstrated here to help illustrate the promising astrophysical applications of this new technique."
    },
    "0809/0809.2441_arXiv.txt": {
        "abstract": "{ Weakly Interacting Massive Particles (WIMPs) are one of the leading candidates for Dark Matter.  For understanding the properties of WIMPs and identifying them among new particles produced at colliders (hopefully in the near future), determinations of their mass and their couplings on nucleons from direct Dark Matter detection experiments are essential.  Based on our method for determining the WIMP mass model--independently from experimental data, we present a way to also estimate the spin--independent (SI) WIMP--nucleon coupling by using measured recoil energies directly.  This method is independent of the as yet unknown velocity distribution of halo WIMPs. In spite of the uncertainty of the local WIMP density (of a factor of $\\sim$ 2), at least an upper limit on the SI WIMP--nucleon coupling could be given, once two (or more) experiments with different target nuclei obtain positive signals.  In a background--free environment, for a WIMP mass of 100 GeV its SI coupling on nucleons could in principle be estimated with a statistical error of only \\mbox{$\\sim$ 15\\%} with just 50 events from each experiment.  }  \\FullConference{Identification of dark matter 2008 \\\\ August 18-22, 2008 \\\\ Stockholm, Sweden}  \\begin{document} ",
        "introduction": "There is strong evidence that more than 80\\% of all matter in the Universe is dark (i.e., interacts at most very weakly with electromagnetic radiation and ordinary matter). The dominant component of this cosmological Dark Matter must be due to some yet to be discovered, non--baryonic particles.  Weakly Interacting Massive Particles (WIMPs) $\\chi$ with masses roughly between 10 GeV and a few TeV are one of the leading candidates for Dark Matter (for reviews, see Refs.~\\cite{SUSYDM}).  Currently, the most promising method to detect many different WIMP candidates is the direct detection of the recoil energy deposited in a low--background laboratory detector by elastic scattering of ambient WIMPs on the target nuclei \\cite{deta}.  The differential rate for elastic WIMP--nucleus scattering is given by \\cite{SUSYDM}: \\beq \\label{eqn:dRdQ} \\dRdQ = \\calA \\FQ \\int_{v_{\\rm min}}^{v_{\\rm max}} \\bfrac{f_1(v)}{v} dv\\, . \\eeq Here $R$ is the direct detection event rate, i.e., the number of events per unit time and unit mass of detector material, $Q$ is the energy deposited in the detector, $F(Q)$ is the elastic nuclear form factor, $f_1(v)$ is the one--dimensional velocity distribution function of the WIMPs impinging on the detector, $v$ is the absolute value of the WIMP velocity in the laboratory frame.  The constant coefficient $\\calA$ is defined as \\beq \\label{eqn:calA} \\calA \\equiv \\frac{\\rho_0 \\sigma_0}{2 \\mchi m_{\\rm r,N}^2}\\, , \\eeq where $\\rho_0$ is the WIMP density near the Earth and $\\sigma_0$ is the total cross section ignoring the form factor suppression. The reduced mass $m_{\\rm r,N}$ is defined by \\beq \\label{eqn:mrN} m_{\\rm r,N} \\equiv \\frac{\\mchi \\mN}{\\mchi+\\mN}\\, , \\eeq where $\\mchi$ is the WIMP mass and $\\mN$ that of the target nucleus.  Finally, $\\vmin = \\alpha \\sqrt{Q}$ is the minimal incoming velocity of incident WIMPs that can deposit the energy $Q$ in the detector with \\beq \\label{eqn:alpha} \\alpha \\equiv \\sfrac{\\mN}{2 m_{\\rm r,N}^2}\\, , \\eeq and ${v_{\\rm max}}$ is relared to the escape velocity from our Galaxy at the position of the Solar system. ",
        "conclusions": "In this paper we have extended our method for determining the WIMP mass \\cite{DMDDmchi} to estimate the spin--independent WIMP--nucleon coupling from the elastic WIMP--nucleus scattering experiments.  This method is independent of the velocity distribution of halo WIMPs as well as (practically) of the as yet unknown WIMP mass.  The only information needed is the measured recoil energies from at least two experiments with different target nuclei and the local Dark Matter density as the unique assumption.  These information combined with the reconstructed WIMP mass will allow us not only to constrain the parameter space in different extensions of the Standard Model of particle physics, but also to identify WIMPs among new particles produced at colliders (hopefully in the near future).  Once one is confident of this identification, one can use further collider measurements of the mass and couplings of WIMPs.  Together with the reconstruction of the velocity distribution of halo WIMPs \\cite{DMDDf1v}, this will then yield a new determination of the local WIMP density.  On the other hand, knowledge of the WIMP couplings will also permit prediction of the WIMP annihilation cross section.  Together with information on the WIMP density, this will allow to predict the event rate in the indirect Dark Matter detection as well as to test our understanding of the early Universe.   \\subsubsection*"
    },
    "0809/0809.0734_arXiv.txt": {
        "abstract": "Cosmic ray (CR) acceleration at the shock created by the expanding cocoons around active galactic nuclei (AGNs) is studied. It is shown that above the energy $10^{18}$~eV the overall energy spectrum of CRs, produced during the AGN evolution and released in the intergalactic space, has the form $N\\propto \\epsilon^{-\\gamma}$, with $\\gamma\\approx 2.6$, which extends up to $\\epsilon_\\mathrm{max}\\sim 10^{20}$~eV. It is concluded that cocoons shocks have to be considered as a main source of extragalactic CRs, which together with Galactic supernova remnants provide the observed CR spectrum. ",
        "introduction": "The overall origin of cosmic rays (CRs) is still an unresolved problem in astrophysics. The understanding of CR origin requires the determination of the astrophysical objects, that are the CR sources, and of the appropriate acceleration processes, that form the CR spectrum in these objects. In this regard, during the last several years considerable progress has been achieved in this field, both experimentally and theoretically. Recently the steepening of the CR spectrum above $3\\times 10^{19}$~eV was established in the HiRes \\citep{hires07} and Auger \\citep{augerJ07} experiments. It presumably corresponds to the so-called Greizen-Zatsepin-Kuzmin (GZK) cutoff, caused by CR energy losses in their interactions with the microwave background radiation. This is evidence that the highest energy part of the CR spectrum is of extragalactic origin. It was also recently demonstrated \\citep{bv07} that the CRs with energies up to $\\epsilon \\sim 10^{17}$~eV can be produced in supernova remnants (SNRs).  The main reason why SNRs are usually considered as the CR source is a simple argument about the energy required to sustain the Galactic CR population against loss by escape, nuclear interactions and ionization energy loss. Supernovae have enough power to drive the CR acceleration. The high velocity ejecta produced in the supernova explosion interacts with the ambient medium to produce a strong blast wave, which may accelerate a small suprathermal fraction of the ambient plasma to high energies. The only theory of particle acceleration which is sufficiently well developed and specific to allow quantitative model calculations is diffusive acceleration \\citep{krym77,bell78} applied to the strong outer shock associated with SNRs \\citep[e.g. see][for review]{bek88,ber05,ber07}. Considerable efforts have been made during the last years to empirically confirm the theoretical expectation that the main part of CRs indeed originates in SNRs. Theoretically, progress in the solution of this problem has been due to the development of a kinetic nonlinear theory of diffusive shock acceleration \\citep{byk96}. This theory includes all the most relevant physical factors, essential for the SNR evolution and CR acceleration in a SNR and it is able to make quantitative predictions of the expected properties of CRs produced in young SNRs and their nonthermal radiation. Applied to individual young SNRs \\citep[see][for reviews]{ber05, ber07} this theory has successfully explained many observed SNR properties. Making use of the observed synchrotron emission spectrum from radio to X-ray frequencies it permits the determination of the degree of magnetic field amplification, a process advocated earlier from plasma simulations \\citep{lb00,belll01}, which leads to the concentration of the highest-energy electrons in a very thin shell just behind the shock. Such observed filamentary structures have been observed in hard X-rays \\citep{vl03,long03,bamba03,par06} and are nowadays used as a second independent method to infer the magnitude of the amplified field \\citep[see][for interpretation]{bkv03}.  It has been demonstrated, that magnetic field amplification leads to considerable increase of the maximal energy of CRs accelerated in SNRs and that the observed CR energy spectrum can be well represented by two components \\citep{bv07}. The first one, $J^{g}(\\epsilon)$, dominant up to $10^{17}$~eV, consists of CRs produced in Galactic SNRs, whereas the second, $J^{eg}(\\epsilon)$, is produced in extragalactic sources. This so called dip scenario \\citep{aloisio06} requires a relatively steep CR spectrum produced in extragalactic sources, $J_\\mathrm{s}^{eg}\\propto \\epsilon^{-\\gamma}$ with $\\gamma =2.55- 2.75$ at energies $\\epsilon>10^{18}$~eV. It has been demonstrated that peculiarities in the CR spectrum -- the so-called \"dip\" structure at $\\epsilon\\sim 10^{19}$~eV and a GZK cutoff at $\\epsilon>3\\times 10^{19}$~eV \\citep{aloisio06} -- produced due to CR interaction with the cosmic microwave background radiation on their way from the source to the observer are very consistent with the experiment.  Below we consider CR acceleration at the shock created by the expanding cocoon around active galactic nuclei (AGNs). Compared with previous considerations \\citep{norman95} magnetic field amplification is taken into account, which provides CR acceleration with an appropriate spectrum up to the energy $~ 10^{20}$~eV. Therefore the cocoon shocks have to be considered as a main source of extragalactic CRs, which together with Galactic supernova remnants presumably provide the observed CR spectrum. ",
        "conclusions": "Astrophysical objects, which are considered as potential extragalactic sources of ultra high energy CRs, should fulfill a number of requirements. First, they should have sufficiently high energy output, not less than $L_\\mathrm{p}\\sim 2\\times 10^{45}-3.5\\times 10^{46}$~erg~Mpc$^{-3}$yr$^{-1}$ depending on the form of CR spectrum that they produce \\citep{berez06}.  Given the energy requirements, AGNs \\citep[e.g.][]{berez06} and gamma-ray bursts (GRBs) \\citep{milgrom95,waxman95,vietri95} are considered as potential extragalactic sources of the ultra-high-energy CRs. However, as \\citet{berez06} argued, the energy output of GRBs has a serious problem if they are considered as the main source of the extragalactic CRs.  The second requirement to the extragalactic CR sources is their ability to produce a power-law CR spectrum up to the maximal energy, which for protons should be at least as large as $\\epsilon_\\mathrm{max}=10^{20}$~eV, which is well above the GZK cutoff. Particle acceleration at relativistic shocks in AGNs \\citep{achterberg00,rb93} and GRBs \\citep{milgrom95,waxman95,vietri95} are discussed as appropriate acceleration processes. However, simulations of CR shock acceleration \\citep{niemiec06} have demonstrated a low efficiency in the case of relativistic shocks.  The relativistic jet in AGNs is surrounded by the hot cocoon, which expands into the surrounding intergalactic medium with supersonic speed \\citep[e.g.][]{ferrari}. Since the essential fraction of the jet energy is transformed into the internal energy of the background medium due to the outermost nonrelativistic shock, driven by the cocoon overpressure, it is natural to suppose that a  good part of this energy is represented by effectively accelerated CRs, like what takes place in SNRs. Since CR acceleration by nonrelativistic shocks has been very well studied one can obtain reliable estimates of the expected spectrum of CRs. First of all the maximal energy of CRs, accelerated at the expanded shock of size $R(t)$ and speed $V=dR/dt$, is determined by the expression \\citep{ber96} \\begin{equation} \\epsilon_\\mathrm{max}\\approx ZeBRV/c, \\end{equation} where $B$ is the upstream magnetic field, $c$ is the speed of light, $e$ is the proton charge and $Z$ is the charge number of CR nuclei. This expression consistently takes into account the influence of all the most essential physical factors restricted the maximal CR energy, which are the adiabatic cooling of CRs in the downstream region and the shock deceleration followed by the CR escape from the shock vicinity into outer space \\citep{ber96}.  Magnetic field near the shock front, as was established for all young SNRs \\citep{ber05,ber07}, is strongly amplified nonlinearly by CRs up to the level \\begin{equation} B^2/(8\\pi)\\approx 3\\times 10^{-3}\\rho V^2, \\end{equation} which is presumably also appropriate for extragalactic shocks. Here $\\rho$ is the external gas density. Note that the relation $B\\propto V$ following from Eq.(2) is indeed theoretically expected for the shock speeds $V<10^4$~km s${-1}$ \\citep{pell06}. Due to the nonresonantly amplified magnetic field, as was originally shown by \\citet{bell04}\\citep[see also][]{pell06}, at higher shock speed $V>10^4$~km s${-1}$ the expected relation is $B\\propto V^{3/2}$, which leads to larger magnetic field strength compared with the value given by Eq.(2). Due to this fact CR maximal energy $\\epsilon_\\mathrm{max}$ can be a few times larger compared with our estimate below. However this does not play a critical role because even our estimate gives $\\epsilon_\\mathrm{max}\\approx 10^{20}$~eV, which is well above GZK cutoff.  The outer cocoon shock expands with the speed \\citep[e.g.][]{ferrari} \\begin{equation} V\\approx [L_\\mathrm{j}/(\\rho V_\\mathrm{h})]^{1/4}t^{-1/2}, \\end{equation} where $L_\\mathrm{j}$ is the mechanical luminosity of the jet and $V_\\mathrm{h}$ is the hot spot or the jet head speed. Substituting the magnetic field value following from Eq.(2), the shock speed (3) and the shock size $R=2Vt$ into the equation (1), we have \\[ \\epsilon_\\mathrm{max}\\approx 10^{20} Z \\left( \\frac{L_\\mathrm{j}}{10^{46}~\\mbox{erg/s}}\\right)^{3/4} \\left( \\frac{N_\\mathrm{g}}{10^{-4}~\\mbox{cm}^{-3}}\\right)^{-1/4} \\] \\begin{equation} \\hspace{0.5cm}\\times \\left( \\frac{V_\\mathrm{h}}{10^{10}~\\mbox{cm/c}}\\right)^{-1/4} \\left( \\frac{l_\\mathrm{j}}{1~\\mbox{kpc}}\\right)^{-1/2}~\\mbox{eV}, \\end{equation} where $l_\\mathrm{j}=V_\\mathrm{h}t$ is the jet length. According to this expression, during the evolutionary epochs corresponding to the jet length from $l_\\mathrm{j}=1$~kpc to $l_\\mathrm{j}=10$~Mpc the cocoon shock produces a power-law proton spectrum which extends up to the maximal energy from $\\epsilon_\\mathrm{max}=10^{18}$~eV to $\\epsilon_\\mathrm{max}=10^{20}$~eV. Thus we conclude, that AGNs fulfill also the second requirement for extragalactic CR sources: protons, the main kind of ions in the space plasma, are effectively accelerated at least up to the energy $\\epsilon_\\mathrm{max}=10^{20}$~eV. Heavier ions are expected to be accelerated to $Z$ times higher maximal energy. However, since the GZK cutoff energy of nuclei $\\epsilon_\\mathrm{GZK}^A<\\epsilon_\\mathrm{GZK}$ for the atomic number $A>1$ is lower then that for protons $\\epsilon_\\mathrm{GZK}\\approx 5\\times 10^{19}$~eV \\citep{berez06}, the predicted dependence $\\epsilon_\\mathrm{max}\\propto Z$ is difficult to establish experimentally.  The most realistic possibility of finding the signature of the acceleration process of ultra-high-energy CRs is the experimental study of CR composition at energies $\\epsilon>10^{17}$~eV. If extragalactic CRs are indeed produced by large-scale nonrelativistic shocks their composition will have the following peculiarity. Inside and near the source, CR composition corresponds to the composition of the intergalactic gas with substantial enrichment by heavy elements, like with Galactic CR versus interstellar medium compositions \\citep[e.g.][]{byk96}. If compositions of intergalactic and interstellar plasmas are not very different, then the expected source CR composition is characterised by the mean atomic number $<A>\\approx 5$, like Galactic CR components. Compared with this source composition, the observed composition is significantly distorted at almost all energies due to propagation effects in the intergalactic medium and in the Galactic wind. Only at the energy $\\epsilon_*\\approx 2\\times 10^{18}$~eV are these compositions roughly the same because it is just below the lowest GZK cutoff energy $\\epsilon_\\mathrm{GZK}^\\mathrm{He}\\approx 4\\times 10^{18}$~eV and the modification of the proton spectrum at this energy is still not very large \\citep{berez06}. Since CR diffusion mobility for given energy $\\epsilon$ is lower for larger $Z$, heavier CRs undergo stronger modification (depression) during their propagation from the sources to the observer. Due to this effect composition of extragalactic CRs is expected to become progressively lighter with the decrease of their energy at $\\epsilon<\\epsilon_*$. At higher energies, $\\epsilon>\\epsilon_*$, due to the dependence of the GZK cutoff energy $\\epsilon_\\mathrm{GZK}\\propto A$ on the nuclear mass number $A$, CRs become progressively heavier so that the peak value of $<A>$ is expected at the beginning of the GZK cutoff for iron, which is at $\\epsilon\\approx 10^{19}$~eV \\citep{berez06}. At such a high energies only protons and iron nuclei survive and the iron-to-proton ratio has a maximum value at the energy $\\epsilon\\approx 10^{19}$~eV, which provides a peak of CR mean mass very similar to the peak at $\\epsilon\\approx 10^{17}$~eV. To make a quantitative prediction of the expected ultra-high-energy CR composition, one has to perform a detailed consideration, which will be done in a subsequent paper.  Note that the expression (4) is based on the amplified magnetic field value (2) and therefore does not depend on the assumed value of intergalactic magnetic field in contrast to the earlier estimate of $\\epsilon_\\mathrm{max}$ \\citep{norman95}. Note also that as is seen from the Eq.(4), maximal CR energy only slightly depends on external gas density and the jet head speed. The values of mechanical luminosity $L_\\mathrm{j}$ and the jet length $l_\\mathrm{j}$, which determines the maximal CR energy, can be directly measured in the experiment.  The third requirement to CR sources is related to the form of the CR spectrum. The form of the resulting CR spectrum produced during the whole evolution of the expanding shock is determined by three physical factors: (1) nonlinear shock modification due to the CR back-reaction, (2) adiabatic energy losses in the downstream region, and (3) diffusive CR escape from the shock vicinity into outer space. The existence of the CR escape phenomenon makes it possible to estimate the shape of the spectrum of the most energetic CRs produced during the source evolution. As was demonstrated analytically \\citep{bek88,ber96} and confirmed numerically \\citep{byk96}, since the maximal energy $\\epsilon_\\mathrm{max}(t)$ of CRs accelerated by the expanding shock during its most active phase decreases with time $t$ due to shock deceleration, the most energetic particles with energy $\\epsilon\\sim \\epsilon_\\mathrm{max}(t)$ undergo diffusive escape from the shock vicinity into the surrounding space. If such behavior happens during the time period from $t_1$ to $t_2$, then at any given epoch $t_1<t<t_2$ all accelerated particles with energies $\\epsilon >\\epsilon_\\mathrm{max}(t)$ have already escaped and inside the source there are only particles with $\\epsilon \\le\\epsilon_\\mathrm{max}(t)$.  Note that the particle escape phenomenon, has two important aspects. First, the energy of escaped CRs determined for any evolutionary epoch by Eq.(1) does not depend any more on the adiabatic cooling inside the parent expanding source, like what happens with CRs of lower energies $\\epsilon <\\epsilon_\\mathrm{max}(t)$. Second, the spectrum of the escaped CRs can be essentially different compared with the canonical spectrum $N\\propto \\epsilon^{-2}$, produced by strong nonrelativistic shock. To estimate the spectrum of escaped CRs one can use the simple relation \\citep{bek88} \\begin{equation} N(\\epsilon)\\epsilon d\\epsilon \\propto \\rho V^2R^{2 -\\beta}dR, \\end{equation} which determines at any given phase $t$ the overall number of CRs $N(\\epsilon)$ with highest energy $\\epsilon\\sim \\epsilon_\\mathrm{max}(t)$. Due to the hard self-consistent spectrum of CRs, produced by the strong shock, the main contribution to the CR energy content is given by the particles with highest energy $\\epsilon\\sim \\epsilon_\\mathrm{max}(t)$ \\citep{byk96}. Therefore their energy content scales as the shock ram pressure $\\rho V^2$ times the shock volume, as it is given by Eq.(5). Factor $R^{-\\beta}$ describes the progressive slow decrease of CR acceleration efficiency due to the shock weakening. For the shock expansion law $R\\propto t^{-\\nu}$ relation (5) gives \\begin{equation} N(\\epsilon) \\propto \\epsilon^{-\\gamma}~~~\\mbox{with}~~~\\gamma=1+(2-\\beta)/(2/\\nu-3). \\end{equation} In our case $\\nu=1/2$ this gives $\\gamma=3-\\beta$. For the constant acceleration efficiency we have $\\beta=0$ which gives $\\gamma=3$. However, the shock deceleration is accompanied by the decrease of the acceleration efficiency. This effect is described by the amount of shock modification, which is characterized by the shock compression ratio, which for the case of a strong shock depends on the shock speed as $\\sigma \\propto V^{3/8}$ \\citep{byk96}. Since in our case $V\\propto 1/R$ we have $\\beta=3/8$ that gives $\\gamma\\approx 2.6$. Such a spectrum of extragalactic CRs at energies $\\epsilon >10^{18}$~eV corresponds very well to the experiment \\citep{berez06}.  Particles with energies $\\epsilon \\le\\epsilon_\\mathrm{max}(t)$ at any given epoch are contained inside the source and have a spectrum close to $N\\propto \\epsilon^{-2}$. At a very late epoch when the outer shock becomes weak these particles will leave the source. Note however that this is already not important for the Galaxy, because the contribution of extragalactic CRs is expected to be low at energies $\\epsilon <10^{18}$~eV \\citep{berez06}.  Note that based on the data collected at the Auger experiment, a correlation between the arrival directions of CRs with energy above $6\\times 10^{19}$~eV and the position of nearby AGNs has been found \\citep{augercol07}, which strongly supports AGNs as prime candidates for the source of ultra-high-energy CRs."
    },
    "0809/0809.2507_arXiv.txt": {
        "abstract": "I discuss three different topics concerning the chemical evolution of the Milky Way (MW). 1) The metallicity distribution of the MW halo; it is shown that this distribution can be analytically derived in the framework of the hierarchical merging scenario for galaxy formation, assuming that the component sub-haloes had chemical properties similar to those of the progenitors of satellite galaxies of the MW. 2) The age-metallicity relationship (AMR) in the solar neighborhood; I argue for caution in deriving from data with important uncertainties (such as the age uncertainties in the Geneva-Kopenhaguen survey) a relationship between average metallicity and age: derived relationships are shown to be systematically flatter than the true ones and should not be directly compared to models. 3) The radial mixing of stars in the disk, which may have important effects on  various observables (scatter in AMR, extension of the tails of the metallicity distribution, flatenning of disk abundance profiles). Recent SPH + N-body simulations find considerable radial mixing, but only comparison to observations will ultimately determine the extent of that mixing. ",
        "introduction": "The regular shape of the metallicity distribution of the Milky Way (MW) halo can readily be explained by the simple model of galactic chemical evolution (GCE) with outflow, as suggested by Hartwick (1976). However, that explanation lies within the framework of the monolithic collapse scenario for the formation of the MW (Eggen, Lynden-Bell and Sandage, 1962). Several attempts to account for the metallicity distribution of the MW halo  in the modern framework (hierarchical merging of smaller components, hereafter sub-haloes) were undertaken in recent years,  (Bekki and Chiba 2001;  Scanapieco and Broadhurst 2001,   Font et al. 2006,  Tumlinson 2006,  Salvadori, Schneider and Ferrara 2007).  Independently of their success or failure in reproducing the observations, those recent models provide little or no physical insight into the physical processes that shaped the metallicity distribution of the MW halo. Indeed, if the halo was built from a  large number of sucessive mergers of sub-haloes, why is its metallicity distribution so well described by the simple model with outflow (which refers to a single system)? And what determines the peak of the metallicity distribution at [Fe/H]=$\\sim$--1.6, which is (successfully) interpreted in the simple model by a single parameter (the outflow rate) ?  In this work we present an attempt to built the halo metallicity distribution analytically, in the framework of the hierarchical merging paradigm. Preliminary results have already been presented in Prantzos (2007b) and a detailed account is given in Prantzos (2008). It is assumed that the building blocks were galaxies with properties similar to (but not identical with) those of the Local group dwarf galaxies that we observe today. In that way, the physics of the whole process becomes (hopefully) more clear.    \\subsection{The halo metallicity distribution  and the simple model}  The halo metallicity distribution is nicely described by the simple model of GCE, in which the metallicity $Z$ is given as a function of the gas fraction $\\mu$ as $ Z \\ = \\ p \\ ln(1/\\mu) \\ + Z_0 $, where $Z_0$ is the initial metallicity of the system  and $p$ is the {\\it yield} (metallicities and yield are expressed in units of the solar metallicity \\zs).  If the system evolves at a constant mass (closed box), the yield is called the {\\it true yield}, otherwise (i.e. in case of mass loss or gain) it is called the {\\it effective yield}. The {\\it differential metallicity distribution} (DMD)  is: \\begin{equation} {{d(n/n_1)}\\over{d(logZ)}} \\ = \\ {{{\\rm ln10}}\\over{1-{\\rm exp}\\left(-{{Z_1-Z_0}\\over{p}}\\right)}}  \\ {{Z-Z_0}\\over{p}} \\ {\\rm exp}\\left(-{{Z-Z_0}\\over{p}}\\right) \\end{equation} where $Z_1$ is the final metallicity of the system and $n_1$ the total number of stars (having metallicities $\\le Z_1$). This function has a maximum for $Z-Z_0=p$, allowing one to evaluate easily the effective yield $p$ if the DMD is determined observationally.  The DMD of field halo stars peaks at [Fe/H]$\\sim$--1.6, suggesting an effective halo yield $p_{Halo} \\sim$1/40 \\zs \\ for Fe. This has to be compared to the true yield obtained in the solar neighborhood, which is $\\sim$ \\zs \\ (from the peak of the local MD), but $\\sim$2/3 of that originate in SNIa, so that the true Fe yield of SNII during the local disk evolution was $p_{Disk}\\sim$0.32 \\zs; consequently, the {\\it effective Fe yield} of SNII during the halo phase (where they dominated Fe production) was $p_{Halo}\\sim$0.08 $p_{Disk}$.  Hartwick (1976) suggested that {\\it outflow} at a rate $F \\ = \\ k \\ \\Psi$ (where $\\Psi$ is the Star Formation Rate or SFR) occured during the halo formation. In the framework of the simple model of GCE such outflow reduces the true yield $p_{True}$ to its effective value $p_{Halo}=p_{True}(1-R)/(1+k-R)$ (e.g. Pagel 1997) where $R$ is the Returned Mass Fraction  ($R\\sim$0.32-0.36 for most of the modern IMFs, e.g. Kroupa 2002). The halo data suggest then that the ratio of the outflow to the star formation rate was \\begin{equation} k \\ = \\ (1-R) \\ \\left({\\frac{p_{True}}{p_{Halo}}} -1 \\right) \\end{equation} and, by replacing $p_{True}$ with $p_{Disk}$ one finds that $k\\sim$7-8.  \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth]{F1.ps} \\caption{Differential metallicity distribution of field halo stars in linear ({\\it top}) and logarithmic ({\\it bottom}) scales. Data are from Ryan and Norris (1991, {\\it open symbols}), and the ongoing Hamburg/ESO project  ( {\\it filled symbols}). The two data sets are normalised at [Fe/H]-2.2; above that value, the Hamburg/ESO data are incomplete, while below that metallicity the Ryan and Norris (1991) data set is incomplete. The {\\it dotted curve} is a simple model with instantaneous recycling (IRA) and outflow rate equal to 8 times the star formation rate. The solid curve is obtained as in Prantzos (2003), from a model without IRA, an early phase of rapid infall and a constant outflow rate equal to 7 times the SFR. All curves and data are normalised to max=1.} \\label{fig:1}       % \\end{figure}    Although the siple model of GCE with outflow fits extremely well the bulk of the halo DMD, the situation is less satisfactory for the low metallicity tail of that DMD. Prantzos (2003), noting that the situation is reminiscent of the local G-dwarf problem,  suggested that a similar solution should apply, namely an early phase of rapid infall (in a time scale of less than 0.1 Gyr) forming the Milky Way's halo, as illustrated in Fig. 1.  \\subsection{The halo DMD and hierarchical merging}  In this work we assume that the MW halo was formed by the merging of smaller units (\"sub-haloes\"), as implied by the hierarchical merging scenario for galaxy formation. In order to calculate semi-analytically the resulting DMD, it is further  assumed that each of the sub-haloes had a DMD described well by the simple model, i.e. by Eq. (1.1). It remains then to evaluate the effective yield $p(M)$ of each sub-halo as a  function of its mass, as well as the mass function of the sub-haloes $dN/dM$.  It is assumed here that each one of the merging sub-haloes has a DMD described by the simple model with an appropriate effective yield. This assumption is based on recent observations of the dwarf spheroidal (dSph) satellites of the Milky Way, as will be discussed below. It is true that the dSphs that {\\it we see today} cannot be the components of the MW halo, because of  their observed abundance patterns (e.g. Shetrone, C\\^ot\\'e and Sargent 2001; Venn et al. 2004): their $\\alpha$/Fe ratios are typically smaller than the   [a/Fe]$\\sim$0.4$\\sim$const. ratio of halo stars. This implies that they evolved on longer timescales than the Galactic halo, allowing SNIa to enrich their ISM with Fe-peak nuclei and thus to lower the $\\alpha$/Fe ratio by a factor of $\\sim$2-3 (as evidenced from the [O/Fe]$\\sim$0 ratio in their highest metallicity stars). However, the shape of the DMD of the simple model does not depend on the star formation history or the evolutionary timescale, only on gas flows into and out of the system, as well as on the initial metallicity (e.g. Prantzos 2007a).  \\begin{figure} \\includegraphics[width=0.46\\textwidth]{F2.ps} \\qquad \\includegraphics[width=0.46\\textwidth]{F3.ps} \\caption{{\\bf left:} Metallicity distributions  of dwarf satellites of the Milky Way. Data are in {\\it histograms} (from Helmi et al. 2006). {\\it Solid curves} indicate the results of simple GCE  models with outflow proportional to the star formation rate;  the corresponding effective yields (in \\zs) appear on top right of each panel. {\\it Dashed curves} are fits obtained with an early infall phase, while {\\it dotted curves} are models with an initial metallicity log$(Z_0)\\sim$--3; both modifications to the simple model (i.e. infall and initial metallicity) improve the fits to the data. {\\bf Right:} Stellar metallicity vs stellar mass for nearby galaxies; data and model ({\\it upper curves}) are from Dekel and Woo (2003), with {\\it dI} standing for dwarf irregulars and  {\\it dE} for dwarf ellipticals. The {\\it thick dotted} line represents the effective yield of the sub-haloes that formed the MW halo according to this work (i.e. with no contribution from SNIa, see Sec. 3.2). The MW halo, with average metallicity [Fe/H]=--1.6 and estimated mass $\\sim$4x10$^8$ \\ms \\ falls below both curves.  } \\label{fig:2}       % \\end{figure}     The DMDs of four neraby dSphs (Helmi et al. 2006) are displayed as histograms in Fig. 2, where they are compared to the simple model with appropriate effective yields ({\\it solid curves}). The effective yield in each case was simply assumed to equal the peak metallicity (Eq. 1.1). It can be seen that the overall shape of the DMDs is quite well fitted by the simple models. This is important, since i) it strongly suggests that {\\it all} DMDs of small galaxian systems can be described by the simple model  and ii) it allows to determine {\\it effective yields} by simply taking the peak metallicity of each DMD (see below).  Before proceeding to the determination of effective yields, we note that the fit of the simple model to the data of dSphs fails in the low metallicity tails. Helmi et al. (2006) attribute this  to a pre-enrichment of the gas out of which the dSphs were formed. However, early infall is another, equally plausible, possibility, as argued in Prantzos (2008). The recent simulations of Salvadori et al. (2008) find both pre-enrichment and early infall for local dSphs.  If the DMDs of all the components of the Galactic halo are described by the simple model, then their shape is essentially  described by the corresponding effective yield $p$ (and, to a lesser degree, by the corresponding initial metallicity $Z_0$). Observations suggest that the effective yield is a monotonically increasing function of the galaxy's stellar mass $M_*$ (Fig. 3).  In the case of the progenitor systems of the MW halo, however, the effective yield must have been lower, since SNIa had not time to contribute (as evidenced by the high $\\alpha$/Fe$\\sim$0.4 ratios of halo stars), by a factor of about 2-3. We assume then that the effective yield of the MW halo components (accreted satellites or sub-haloes) is given (in \\zs) by \\begin{equation} p(M_*) \\ = \\ 0.005 \\ \\left({\\frac{M_*}{10^6 M_{\\odot}}}\\right)^{0.4} \\end{equation} i.e. the thick dotted curve in Fig. 3.  Obviously, the stellar mass $M_*$ of each of the sub-haloes should be $M_*< M_H$ where $M_H$ is the stellar mass of the MW halo ( $M_H$= 4$\\pm$0.8 10$^8$ \\ms, e.g. Bell et al. 2007).   Hierarchical galaxy formation scenarios predict  the mass function of the dark matter sub-haloes which compose a dark matter halo at a given redshift. Several recent simualtions find $ dN/dM_D \\ \\propto \\ M_D^{-2}  $ (Diemand et al. 2007, Salvadori et al. 2007,   Giocoli et al. 2008). In our case, we are interested in the mass function of the {\\it stellar sub-haloes}, and not of the dark ones. Considering the effects of outflows on the baryonic mass function, Prantzos (20008) finds that \\begin{equation} \\frac{dN}{dM_*} \\ \\propto \\ M_*^{-1.2} \\end{equation} i.e. {\\it the distribution function of the stellar sub-haloes is flatter than the distribution function of the dark matter sub-haloes}. The normalisation of the stellar sub-halo distribution function is made through \\begin{equation} \\int_{M_1}^{M_2} \\frac{dN}{dM_*} \\ M_* \\ dM_* \\ = \\ M_H \\end{equation}  \\begin{figure}  \\includegraphics[width=0.46\\textwidth]{F4.ps} \\qquad \\includegraphics[width=0.46\\textwidth]{F5.ps} \\caption{{\\bf Left:}Properties of the sub-haloes as a function of their stellar mass,  empirically derived as discussed in Sec. 3. From top to bottom: Outflow rate, in units of the corresponding star formation rate; Effective yield, in solar units; Distribution function; Cumulative fraction of stellar mass contributed by the sub-haloes. The total mass of the MW halo is 4 10$^8$ \\ms.  {\\bf Right}:{\\it Top} and {\\it middle} panels: Differential metallicity distribution (in lin and log scales, respectively) of the MW halo, assumed to be composed of a population of smaller units (sub-haloes). The individual DMDs of a few sub-haloes, from 10$^6$ \\ms \\ to 4 10$^7$ \\ms, are indicated in the middle panel, as well as the sum over all haloes ({\\it solid upper curves} in both panels, compared to observations). {\\it Dotted curves} in top and middle panels indicate the results of the simple model with outflow (same as in Fig. 1). Because of their large number, small sub-haloes with low effective yields contribute the largest fraction of the lowest metallicity stars, while large haloes contribute most of the high metallicity stars ({\\it bottom panel}).  } \\label{fig:4} \\end{figure}     The lower mass limit $M_1$ is adopted here to be $M_1$=1-2 10$^6$ \\ms, in agreement with the lower mass bound of observed dSphs in the Local group. Such galaxies have internal velocities $V>$10 km/s. Dekel and Woo (2003) argue that the gas in haloes with $V<$10 km/s cannot cool to form stars at any early epoch and that dwarf ellipticals form in haloes with 10$<V<$30 km/s, the upper limit corresponding to a stellar mass $M_*\\sim$2 10$^8$ \\ms. This is the upper mass limit $M_2$ that we adopt here. The main properties of the sub-halo set constructed in this section appear in Fig. 3 (left) as a function of the stellar sub-halo mass $M_*$. The resulting total DMD is obtained as a sum over all sub-haloes: \\begin{equation} \\frac{d(n/n_1)}{d(logZ)} \\ = \\ \\int_{M_1}^{M_2} \\ \\frac{d[n(M_*)/n_1(M_*)]}{d(logZ)} \\ \\frac{dN}{dM_*} \\ M_* \\ dM_* \\end{equation} The result appears in Fig. 4 (right, with top panel in linear and middle panel in logarithmic scales, respectively). It can be seen that it fits the observed DMDs at least as well as the simple model \\`a la Hartwick. In summary, under the assumptions made here, the bulk of the DMD of the MW halo results naturally as the sum of the DMDs of the component sub-haloes. It should be noted that all the ingredients of the analytical model are taken from observations of local satellite galaxies, except for the adopted mass function of the sub-haloes (which results from analytical theory of structure formation plus a small modification to account for the role of outflows). Obviously, by assuming different values for the slope of the dark matter haloes  and for the mass limits $M_1$ and $M_2$ one may modify the peak of the resulting composite DMD, thus allowing for differences between the halo DMDs of different galaxies. ",
        "conclusions": ""
    },
    "0809/0809.2513_arXiv.txt": {
        "abstract": "We present a new measurement of the space density of high redshift ($z \\simeq $ 3.0-4.5), X--ray selected QSOs obtained by exploiting the deep and uniform multiwavelength coverage of the COSMOS survey. We have assembled a statistically large (40 objects), X--ray selected ($F_{0.5-2 keV} >$  10$^{-15}$ \\cgs), homogeneous sample of $z >$ 3  QSOs for which spectroscopic (22) or photometric (18) redshifts are available. We present the optical (color-color diagrams) and X--ray properties, the number counts and space densities of the $z >$ 3 X--ray selected quasars population and compare our findings with previous works and model predictions. We find that the optical properties of X--ray selected quasars are not significantly different from those of optically selected samples. There is evidence for substantial X--ray absorption (log$N_H >$ 23 cm$^{-2}$) in about 20\\% of the sources in the sample. The comoving space density of luminous ($L_X \\gsimeq 10^{44}$ erg s$^{-1}$) QSOs declines exponentially (by an e--folding per unit redshift) in the $z \\sim$ 3.0--4.5 range, with a behavior similar to that observed for optically bright unobscured QSOs selected in large area optical surveys. Prospects for future, large and deep X--ray surveys are also discussed. ",
        "introduction": "The well known correlations between Super Massive Black Holes (SMBH) and galaxy properties such as bulge luminosity and stellar velocity dispersion (Gebhardt et al. 2000; Ferrarese \\& Merrit 2000), which are by now firmly established at least in the local Universe, point towards a tight link between the assembly of the bulge mass and SMBH. The high redshift AGN luminosity function (LF) and the source counts represent key observational constraints for theoretical models of galaxy and SMBH formation and evolution. While most models are reasonably successful in reproducing several observables in the high-luminosity regime (Di Matteo et al. 2005; Volonteri \\& Rees 2006; Hopkins et al. 2005; Lapi et al. 2006; Menci et al. 2008), they have to face the paucity of data, especially at high redshifts and low luminosities. Given that theoretical predictions are used to determine key physical parameters, such as the QSOs duty cycle, the BH seed mass function, and the accretion rates,  a reliable observational estimate of the QSOs LF and evolution at high redshift is extremely important.  Large optical surveys, most notably the 2dF Quasar Redshift Survey (2QZ; Croom et al. 2004) and the Sloan Digital Sky Survey (SDSS; i.e. Richards et al. 2006) were able to accurately measure the shape of the LF up to $z \\sim 2.0-2.5$ and to constrain the bright (M$_{\\rm  B}<$ --25.5 or M$_{\\rm  I}<$ --27.6, corresponding to bolometric luminosities of log$L_{bol}\\gsimeq 46$ erg s$^{-1}$) QSOs LF up to z $\\sim$ 6.5: the space density of optically selected QSOs peaks at $z \\sim 2-3$, then decreases approximately by a factor 3, per unit redshift, in the range $z\\simeq 3-6$ (Schmidt, Schneider \\& Gunn 1995, Fan et al. 2001, 2004; Richards et al. 2006). The high--redshift quasar surveys mentioned above are relatively shallow and probe only the bright end of the LF. Recently, deep optical surveys were able to probe the LF at significantly fainter magnitudes using different selection criteria (e.g. Wolf et al. 2003, Hunt et al. 2004, Fontanot et al. 2007, Bongiorno et al. 2007, Siana et al. 2008). However, the statistics on the faint QSOs population at $z>3$ is still limited  given the typically small area surveyed.  Moreover, optically selected quasars are known to make up only a fraction of the entire population of accreting SMBH: in fact, most of the accretion power in the Universe is obscured by large amounts of dust and gas (see, e.g., Fabian \\& Iwasawa 1999); moreover, AGN and QSOs span a broad range in accretion rate (see, e.g., Merloni \\& Heinz 2008), and, therefore, of intrinsic luminosities. Deep X-ray surveys with {\\it Chandra} and XMM--{\\it Newton} have proven to be very efficient in revealing obscured accretion (except for Compton Thick AGN) down to relatively low luminosities (see Brandt \\& Hasinger 2005 for a review). According to the most recent models for the synthesis of the X--ray Background (e.g. Gilli,Comastri \\& Hasinger 2007), the obscured AGN population (including heavily obscured Compton Thick AGN) outnumbers the unobscured ones by a luminosity dependent factor ranging from $\\sim2$ at L$_{\\rm X} \\simeq 10^{45}$ erg s$^{-1}$ to $\\sim8$ at L$_{\\rm X} \\simeq 10^{42}$ erg s$^{-1}$. The evidence for a decreasing obscured AGN fraction towards high luminosities is supported by many investigations (e.g. Ueda et al. 2003, Treister \\& Urry 2005). Moreover, there is also increasing evidence that the fraction of obscured AGN grows towards high redshifts (La Franca et al. 2005; Treister \\& Urry 2006;  Hasinger 2008).  By combining deep and shallower X--ray surveys (Ueda et al. 2003; La Franca et  al. 2005; Hasinger et al. 2005; Silverman et al. 2008) it has been possible to address the issue of the evolution of the X-ray LF at high redshifts. Previous works presented evidence for a decline of the comoving space density of X--ray selected AGN at $z >$ 3 in both  the soft (0.5--2 keV; Hasinger et al. 2005; Silverman et al. 2005) and hard (2--10 keV; Silverman et al. 2008) bands with a rate similar to that observed for optically selected quasars. At face value, these findings would imply that the space density of obscured AGN, which are likely to be missed by optical surveys, declines toward high redshifts ($z>$ 3) with a behavior similar to that of the unobscured AGN, at least at the relatively high luminosities probed by present X--ray surveys (log$L_{X}\\gsimeq 44$). However, the results available so far are based on rather heterogeneous and relatively small samples, also affected by a significant level of spectroscopic incompleteness due to the faint magnitudes of the optical counterparts. To cope with this limitation, several photometric selection techniques have been proposed in the recent years. Though differing in the details, the search for high redshift, obscured AGN is mainly based on a suitable combination of X--ray, optical and infrared flux ratios (see e.g. Fontanot et al. 2007). More specifically, hard X--ray emission associated with extremely faint or even undetected optical counterparts and extremely red colors (from optical wavelengths up to mid--infrared) is considered a reliable proxy for a high redshift obscured AGN (see, e.g., Koekemoer et al. 2004; Brusa et al. 2005; Fiore et al. 2008). As an example, at the X-ray fluxes probed by the COSMOS survey, an X-ray to optical flux ratio f$_X$/f$_{opt}>10$ efficiently selects high z, Compton thin obscured AGN (see Fiore et al. 2003). Most of them have red or extremely red (R-K$>5$) optical to near infrared colors and indeed f$_X$/f$_{opt}$ and R-K are tightly correlated among X-ray selected sources (Brusa et al. 2005). Unfortunately, due to the widely different selection criteria, it is difficult to combine the above mentioned approaches to obtain a reliable estimate of their space density. Given the importance of the AGN LF at high redshift for our understanding of SMBH growth and evolution, the search for and the census of high redshift ($z >$ 3) AGN warrant further efforts. The excellent multiwavelength coverage of the COSMOS field (Scoville et al. 2007) offers the unique opportunity to assemble the first statistically large, homogeneous, and  well defined sample of X--ray selected, high redshift AGN, which is almost completely unbiased against obscuration up to column densities of the order of $N_H \\simeq 10^{23}$ cm$^{-2}$.  The paper is structured as follows. In Section 2 we describe the sample selection  which is mainly based on the optical and multiwavelength identification  of the XMM--COSMOS survey (Hasinger et al. 2007, Cappelluti et al. 2007, 2008, Brusa et al. 2007; 2008), complemented  by {\\it Chandra} positions when available (Elvis et al. 2008, Puccetti et al. 2008, Civano et al. 2008) to secure the correct  optical counterpart, and, most importantly, by spectroscopic (Trump et al. 2008, Lilly et al. 2007) and high-quality photometric  redshifts (Salvato et al. 2008). The optical and X--ray properties are presented in Section 3 and 4, respectively. In Section 5 the number counts and space density of the sample are discussed and compared with model predictions. The results are summarized in Section 6, where the perspectives for future observations are also briefly outlined. A $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.3 and  $\\Omega_{\\Lambda}$ = 0.7 cosmology is adopted. Optical magnitudes are in the AB system. ",
        "conclusions": "The analysis of optical colors and X--ray spectra indicates that high redshift, X--ray selected QSOs have optical properties not significantly different from those of optically selected, bright QSOs, even if a non-negligible fraction ($\\sim 20$\\% i.e. the outliers of fig.~4) would not have been selected by the Casey et al. (2008) optical color-color criteria. Not surprisingly, these sources are optically faint and make up a significant fraction (30\\%) of the spectroscopically unidentified objects for which the observed HR does not allow to constrain the column density.   Even if the  bulk of the population of obscured AGN responsible for the X--ray background is not fully sampled at the limiting flux of the XMM--COSMOS survey, the relative fraction ($\\sim$ 20\\%) of obscured (log$N_H > $ 23) QSOs is not inconsistent, given the small number statistics, with that ($\\sim$ 10 \\%) predicted by Gilli et al. (2007). We note, moreover, that the log$N_H>23$ sample might be contaminated by less obscured sources (log$N_H=22-23$) given the relatively large errors associated to the measured HR. The depth of the XMM--COSMOS survey does not allow us to further investigate the absorption distribution for column densities lower than log$N_H \\simeq$ 23, nor to firmly establish whether optically faint sources are X--ray obscured. The fraction of outliers in the optical color--color diagram is very similar to that of X--ray obscured sources and both are close to 20\\%. However, the two subsamples are marginally overlapping (only two optical outliers are also X--ray obscured) suggesting that gas absorption and dust reddening are not tightly correlated (see Maiolino et al. 2001).   The X--ray luminosity function of hard X--ray selected AGN as determined by combining several {\\it Chandra} and XMM surveys (Ueda et al. 2003; La Franca et al. 2005)  is well constrained up to relatively low redshifts ($z\\sim$ 2--3). At higher redshifts, the number statistics has so far prevented a robust estimate of their space density. Evidence for a decline at $z >$ 3 was reported by Hasinger et al. (2005) and Silverman et al. (2005), but is limited to bright, unobscured QSOs. The $z >$ 3 QSOs sample drawn from the XMM--COSMOS survey allowed us, for the first time, to firmly address the issue of the evolution of high redshift QSOs thanks  to a homogeneous and sizable sample of X--ray sources much less biased with respect to mildly obscured AGN and with an almost complete redshift information. The results indicate that the comoving space density of X--ray luminous ($L_X \\gsimeq 10^{44}$ erg s$^{-1}$) QSOs at $z\\sim 3$ is ($3.7\\pm1.2)\\times10^{-6}$ Mpc$^{-3}$ and declines exponentially (by an e--folding per unit redshift) in the $z \\sim$ 3.0--4.5 range. These results appear to be robust, despite the still remaining (small) uncertainties on the photometric redshifts discussed in Sect. 2.2. Moreover, if all the sources undetected in the optical band were at $z>3$ equally populating the different redshifts bins, the shape of the space density as a function of redshift    would not change significantly. Only in the case that all the 10 optically undetected, soft X--ray sources turn out to be at $z>$3.5-4, the analytical parameterization of the observed decay should be modified.  The high--redshift decline is similar to that of luminous ($M_I < -27.6$), optically bright, unobscured QSOs as well established by SDSS observations (Richards et al. 2006) and can be satisfactorily described by the Schmidt et al. (1995) parameterization. Assuming an X--ray to optical spectral index appropriate for these luminosities ($\\alpha_{\\rm ox}=-1.65$, e.g.Vignali et al. 2003, Steffen et al. 2006), the absolute magnitude SDSS limit corresponds to an X--ray luminosity  log$L_X \\sim 45$. Therefore, the observed decline in the space density for the XMM--COSMOS sample strongly suggests that the evolution of mildly obscured (Compton thin) AGN is very similar to that of unobscured, optically luminous QSOs, provided that the shape of the declining function holds also for luminosities which are about an order of magnitude lower than those probed by SDSS.  Silverman et al. (2008) computed the XLF for the optically bright ($i<24$) X-ray population. Even though there is a good agreement at $z <$ 3.5 between their LDDE XLF and the observed XMM-COSMOS space densities (for the same limit in the optical magnitude), the extrapolation of their XLF at higher redshift is larger, by a factor of $\\sim2$, with respect to our data. A better fit would be obtained by tuning the Silverman et al. (2008) LF parameters responsible for the high redshift ($z > $ 3.0--3.5) behavior.     The observed decline in the comoving space density of X--ray selected QSOs, at least for luminosities larger than $\\sim10^{44}$ erg s$^{-1}$, has a significant impact on the predictions of the QSOs number counts expected from future large area X--ray  surveys, and, more in general, on predictions at all the wavelengths based on the current available XLF as representative of the high-redshift, radio quiet population (e.g. Wilman et al. 2008). The present results may also provide a benchmark for theoretical models of SMBH growth. For example, the $z >$ 4 QSOs number counts predicted by the Rhook \\& Haehnelt (2008) models (see central and right panels in their Fig.~6) are about a factor 7--8 higher than what is actually observed, suggesting that some of the model parameters should be revised.  The predicted number counts for a model with and without an exponential decline, at different limiting fluxes and redshift ranges are listed in table 4. The Chandra-COSMOS survey provides a homogeneous coverage with high-resolution (HPD $<2''$) over the central $\\sim0.5$ deg$^2$ in COSMOS (Elvis et al. 2008) down to a limiting flux of $\\sim4\\times10^{-16}$ \\cgs\\ in the 0.5-2 keV band (see vertical line in Fig.~8). It will therefore roughly double the number of $z>3$ quasars in that area, exploring the flux regime where the contribution from most obscured sources is higher.  As a practical example we also report the expectations for the eROSITA (extended ROentgen Survey with an Imaging Telescope Array) survey with the SRG (Spectrum R\\\"ontgen Gamma) mission , which will be launched in the next years. eROSITA will survey the entire extragalactic sky down to a limiting flux of $\\sim 10^{-14}$ \\cgs in the 0.5-2 keV band. The number of expected $z>3$ ($z>4$) QSOs in the all sky survey, when the decline described in $\\S$5.2 is incorporated in our current knowledge of the XLF, is 2.5$\\times10^4$ (2100).  eROSITA will also perform a deeper survey over an area of $\\sim 400$ deg$^2$ down to a depth of $4\\times10^{-15}$ \\cgs.  We expect to reveal $\\sim 2500$ QSOs at $z>3$, about 200 of them at $z>4$ and better constrain their redshift distribution.  The study of QSOs space density and evolution at lower X--ray luminosities and higher redshifts requires much deeper observations. While a few moderately luminous (log$L_X \\sim$ 43--44), high redshift $z >$ 4 quasars are detected in deep {\\it Chandra} fields (i.e. Vignali et al. 2002) the present sample size is not such to constrain their space density. Additional high--z objects are expected to be revealed by ongoing  ultra-deep {\\it Chandra} (Luo et al. 2008) and XMM--{\\it Newton} surveys."
    },
    "0809/0809.2039_arXiv.txt": {
        "abstract": "The recently discovered magnetic Herbig Ae and Be stars may provide qualitatively new information about the formation and evolution of magnetic Ap and Bp stars. We have performed a detailed investigation of one particularly interesting binary system with a Herbig Ae secondary and a late B-type primary possessing a strong, globally ordered magnetic field.  Twenty high-resolution Stokes $V$ spectra of the system were obtained with the ESPaDOnS instrument mounted on the CFHT. In these observations we see clear evidence for a magnetic field in the primary, but no evidence for a magnetic field in the secondary. A detailed abundance analysis was performed for both stars, revealing strong chemical peculiarities in the primary and normal chemical abundances in the secondary. The primary is strongly overabundant in Si, Cr, and other iron-peak elements, as well as Nd, and underabundant in He. The primary therefore appears to be a very young Bp star. In this context, line profile variations of the primary suggest non-uniform lateral distributions of surface abundances. Interpreting  the $0.63995 \\pm 0.00009$ day variation period of the Stokes $I$ and $V$ profiles as the rotational period of the star, we have modeled the magnetic field geometry and the surface abundance distributions of Si, Ti, Cr and Fe using Magnetic Doppler Imaging. We derive a dipolar geometry of the surface magnetic field, with a polar strength $B_{\\rm d}=1230$~G and an obliquity $\\beta=57\\degr$. The distributions Ti, Cr and Fe are all qualitatively similar, with an elongated patch of enhanced abundance situated near the positive magnetic pole. The Si distribution is somewhat different, and its relationship to the magnetic field geometry less clear. ",
        "introduction": "Strong, globally organised magnetic fields have recently been reported in a few Herbig Ae and Be (HAeBe) stars \\citep{Donati1997-major, Hubrig2004-HAeBe, Wade2005-HAeBe_Discovery,Wade2007-HAeBe_survey,Catala2007-HD190073,Alecian2008-HD200775,Alecian2008-ClusterHAeBeLetter}.  HAeBe stars are pre-main sequence stars of intermediate-mass which evolve to become main sequence A and B stars. HAeBe stars have A or B spectral classes, display emission lines and infrared excesses, and are usually found with nebulosity nearby \\citep{Vieira2003-HAeBe_ID}. The detection of magnetic fields in HAeBe stars is of particular interest as it hints at a connection to the main sequence magnetic, chemically peculiar Ap and Bp stars.  Ap and Bp stars display strong magnetic fields, with typical strengths of a few hundred gauss up to a few tens of kilogauss, and globally ordered geometries of approximately dipolar topology.  These stars also display strong and distinctive chemical peculiarities, particularly overabundances of Si and iron peak elements, occasionally in excess of 2 dex, and even greater overabundances of some rare earth elements.  The source of the magnetic fields observed in Ap/Bp stars is not well understood. The two major competing field origin theories propose, on the one hand, that the field is a relic of an earlier stage of stellar evolution, now frozen into the plasma of the star or, on the other hand, that the field is generated contemporaneously by a dynamo. Also unknown are the details of the formation of the observed chemical peculiarities.  Chemical peculiarities are believed to be the result of atomic diffusion leading to a chemically stratified stellar envelope \\citep{Michaud1970-diffusion,Michaud1981-diffusion_magneticApBp}. However, the details of this process, particularly for individual elements, and the impact of magnetic fields are not fully understood \\citep[e.g.][]{AlecianG2007-el-distributions-atmos}. In this context it is of great interest to identify and characterize the evolutionary progenitors of Ap/Bp stars, as these objects can provide critical information on both the origin of magnetic fields and the formation of chemical peculiarities.  The recently discovered magnetic HAeBe stars have been proposed to be pre-main sequence progenitors of Ap/Bp stars \\citep{Wade2005-HAeBe_Discovery}.  Thus a detailed investigation of these stars can shed light on these questions, as well as potentially provide some information about magnetic braking and magnetospheric accretion in intermediate-mass stars. ",
        "conclusions": "We have analyzed 20 high resolution spectropolarimetric observations of the HD 72106 system.  In these observations we see clear evidence for a magnetic field in HD 72106A.  We also confirm that HD 72106B is a HAeBe star, based on emission in H$\\alpha$ and the O~{\\sc i}~7771~\\AA\\, triplet, and that it displays no magnetic field.  There is strong evidence that the HD 72106 system is a true binary system. Both stars have the same proper motions, the same radial velocities, and Hipparcos solution to the system places both stars at the same distance. Although the separation of the components has remained constant in the pas 90 years, there has been a slow systematic increase in the position angle observed. Additionally there is a wide range of possible bound orbits consistent with all available observations.  We find the age of the system to be between 6 and 13 Myr, based on the H-R diagram position of the secondary (and assuming that, as a HAeBe star, the secondary is a pre-main sequence star). Thus the system is fairly evolved, for a pre-main sequence system. In the youngest limit, the primary would within 1 Myr of the ZAMS. However, in the oldest limit the primary would have reached the main sequence 9 Myr ago, but it would have only passed about 1.5\\% of its main sequence lifetime.  Thus, while it may not be on the pre-main sequence, HD 72106A is certainly very young. While it is possible that the system is not coeval and instead was form by capture, the data available (consistent H-R diagram positions and the young age of the secondary) suggests that the system truly coeval. Even with the large uncertainty in age, only a few known Ap/Bp stars approach the maximum fractional age of HD 72106A \\citep{Bagnulo2004-B_NGC2244-334, Landstreet2007-surveyFORS1}. Given its evolutionary status, HD 72106A appears to represent a link between magnetic HAeBe stars and the Ap/Bp stars.  HD 72106A possesses a strong, predominantly dipolar, magnetic field.  We find that a centered dipole with a polar field strength of $1230 \\pm 80$ G and an obliquity angle of $57 \\pm 5\\degr$ is sufficient to provide a good fit to our observations.  In HD 72106B we find no evidence of a magnetic field, with an upper limit on the longitudinal field of about 200 G. HD 72106A is one of the very youngest stars for which detailed modeling of the magnetic field geometry has been performed.  The only clearly younger object is the Herbig Be star HD 200775A, which \\citet{Alecian2008-HD200775} find to have an age of $0.1 \\pm 0.05$ Myr. This star also displays a dipolar magnetic field geometry, with a magnetic field strength at the pole of $1000 \\pm 150$ G and an obliquity of $125 \\pm 8\\degr$, slightly offset at $0.05 \\pm 0.04\\,R_*$ from the star's center \\citep{Alecian2008-HD200775}. HD 200775A and HD 72106A have similar magnetic field characteristics, both of which are analogous to most main sequence Ap/Bp stars. This provides further evidence that there is a continuum of magnetic A and B stars from the pre-main sequence through the main sequence.  We see strong chemical peculiarities in HD 72106A, particularly an underabundance of He and overabundances of Si, Ti, Cr, Fe, and Nd.  Thus, from the point of view of abundance analysis, HD 72106A appears to be a Bp star.  This implies that chemical peculiarities can form very early in a star's lifetime, possibly even on the pre-main sequence. This conclusion is supported by the recent results of \\citet{Alecian2008-ClusterHAeBeLetter} who report chemical peculiarities in the spectrum of the HAeBe star W601 in NGC 6611.  The only other star with a comparably young fractional age in which strong chemical peculiarities have been measured in detail is the main sequence Bp star NGC 2244 334 \\citep{Bagnulo2004-B_NGC2244-334}. This star has been on the main sequence for $2.3 \\pm 0.3$ Myr, but is more massive and hence probably more evolved than HD 72106A, with a fractional age $\\tau = 0.02 \\pm 0.01$ \\citep{Hensberge2000-Age_of_NGC2244,Bagnulo2004-B_NGC2244-334}. \\citet{Bagnulo2004-B_NGC2244-334} find that the star has a temperature of $15000 \\pm 1000$ K and an observed longitudinal magnetic field strength of 9000 G. The authors find a strong underabundance of He by $\\sim1$ dex, as well as strong overabundances of Si and Fe by $\\sim1$ dex, and Ti and Cr by $\\sim2$ dex. These peculiarities are similar to those of HD 72106A, suggesting that there are important similarities between these stars.  We find clear evidence for surface abundance inhomogeneities in HD 72106A. Magnetic Doppler Imaging suggests that there is a large spot of overabundance in Si, Ti, Cr, and Fe near the positive magnetic pole. These inhomogeneities are similar to those seen in Ap/Bp stars, further supporting a link between magnetic HAeBe stars and Ap/Bp stars. The Doppler reconstructions presented here are the earliest stage of intermediate-mass stellar evolution ever mapped.  HD 72106B, in contrast to HD 72106A, appears to be chemically normal. While this is not the first binary system containing a magnetic chemically peculiar and a chemically normal star \\citep[e.g.][]{Carrier2002-binarity_in_Ap}, it does raise the question of how such systems are produced.  Being a young binary, it is likely that the stars of HD 72106 formed at the same time from approximately the same material. The temperatures of these stars differ by only $2250 \\pm 1120$ K, and the \\vs\\, values of the stars differ by only $\\sim10$ \\kms.  However, one star is magnetic and chemically peculiar while the other is not.  This suggests that, whatever mechanism gave rise to the difference between the stars, it must be rather sensitive to the particulars of the star's initial conditions.  It is instructive to compare the characteristics of HD 72106A to those of the magnetic HAeBe stars HD 104237 and HD 190073. HD 104237 was observed to possess a magnetic field by \\citet{Donati1997-major} and confirmed by \\citet{Donati2000-misc} and \\citet{Alecian2008-CpAp_workshop}, with a longitudinal field strength of $\\sim 50$ G (Donati, private communication). HD 104237 has a mass of about 2.3 $M_{\\odot}$ and an age of about 2 Myr \\citep{van_den_Ancker1998-HAeBe-photo/atmo}. \\citet{Acke2004-HAeBe_Abun} performed an abundance analysis of HD 104237 using equivalent widths, and found approximately solar abundances for a range of elements, including Si, Cr, and Fe. The star HD 190073 was reported to possess a magnetic field by \\citet{Catala2007-HD190073}, with a longitudinal magnetic field strength of $74 \\pm 10$ G.  \\citet{Catala2007-HD190073} derive a mass of $2.85 \\pm 0.25$ \\Msun\\, and an age of $1.2 \\pm 0.6$ Myr (measured from the birth line) for this star. \\citet{Acke2004-HAeBe_Abun} also studied the surface chemistry of this star and found roughly solar abundances.  Thus it appears that the majority of known magnetic HAeBe stars are chemically normal, though some peculiar stars seem to exist \\citep[such as NGC 6611 W601,][]{Alecian2008-ClusterHAeBeLetter}. This is in contrast to main sequence Ap/Bp stars, in which magnetic fields are nearly always found with chemical peculiarities. More analysis of chemical abundances in magnetic HAeBe stars must be performed, with an eye to identifying chemically peculiar objects. Interestingly, the young chemically peculiar stars HD 72106A and NGC 2244 334 display no emission, while HD 104237 and HD 190073 both display significant emission in their spectra. This suggests that HD 104237 and HD 190073 may still be undergoing significant accretion or mass loss while HD 72106A and NGC 2244 334 are not. Thus it may be that accretion or mass loss mixes the stellar atmosphere, inhibiting the buildup of chemical peculiarities through diffusion. However, once accretion halts chemical peculiarities may arise quickly."
    },
    "0809/0809.5105_arXiv.txt": {
        "abstract": "We describe the implementation of slitless radial velocity measurements of extragalactic planetary nebulae (PNs) with the 8.2 m Subaru telescope and its Cassegrain imaging spectrograph, FOCAS. As a first application, we have extended a previous search for PNs in NGC 4697 to larger angular distances from its center. A total of 218 PNs were detected, and their radial velocities were measured. We have added 56 new PN detections to the existing sample of 535, observed previously with the ESO VLT + FORS imaging spectrograph; 36 of these new 56 PNs are located at angular distances larger than 230 arcsec from the center of NGC 4697. We compare the new FOCAS velocities with the earlier FORS velocities, for 158 of the 162 reobserved sources, finding good agreement. We now have kinematic information extending out to 5 effective radii from the center. The outer line-of-sight velocity dispersion is a bit lower than estimated earlier. This result is compatible with the existence of a dark matter halo plus some degree of radial anisotropy, but the dark matter halo is rather inconspicuous, and it is still unclear how massive it can be. A more detailed global dynamical study of the whole set of PN velocities will be required to decide if they permit to narrow down the range of possible dark matter distributions in NGC 4697. The new radial velocities reveal no evidence of rotation at 5 effective radii. ",
        "introduction": "Planetary nebulae (PNs) in the outskirts of elliptical galaxies offer valuable kinematic information that can be used to trace the angular momentum distribution and mass distribution in these galaxies. The classic example is the study of PNs in NGC 5128 by Hui et al. (1995), later extended by Peng et al. (2004) to 780 PNs. A dark matter halo and significant outer rotation were detected. A recent dynamical study of 340 globular clusters in NGC 5128 has shown additional evidence of a dark matter halo (Woodley et al. 2007). Another case where a dark matter halo was immediately apparent from PN kinematics is NGC 1344 (Teodorescu et al. 2005).  In a previous work (M\\'endez et al. 2001), hereafter Paper I, we presented photometry and slitless radial velocity information for 535 PNs in the flattened elliptical galaxy NGC 4697. These PNs provided kinematic information up to a distance of 3 effective radii from the center of NGC 4697. The most important result in Paper I was a substantial decline of line-of-sight velocity dispersion (losvd) as a function of angular distance from the center, which can be interpreted as indicating the absence of a dark matter halo around this galaxy, if we assume an isotropic velocity distribution. However, some dark matter can be present if the velocity distribution is radially anisotropic; we explicitly warned about this alternative interpretation in Paper I, and the point has been made in much more detail by Dekel et al. (2005) and De Lorenzi et al. (2008a).  It is important to clarify which of the two interpretations, lack of dark matter or pronounced radial anisotropy, is correct. In the current context, with dark matter firmly established as one of the basic components of the universe, the existence of one or several substantial, medium-size elliptical galaxies without a dark matter halo (which is expected to have made possible its very formation) is surprising. On the other hand, if we are able to confirm radial anisotropy, then we have learned about an important constraint concerning the formation of elliptical galaxies. For example, Athanassoula (2005) has made numerical experiments in which a quick succession of disk mergers is more likely to produce significant radial anisotropy.  De Lorenzi et al. (2008a) made a detailed dynamical analysis of NGC 4697, using the PNs of Paper I and long-slit absorption line kinematic data. Their best-fitting models were radially anisotropic, and included potentials with a variety of dark matter halos. The models with no dark matter were not consistent with the PN losvds. However, more PN velocities at larger angular distances from the galaxy's center were required to narrow down the range of acceptable dark matter models. In this paper we provide more such PN velocities.  Romanowsky et al. (2003) showed similarly declining losvds in NGC 3379, NGC 4494 and NGC 821, claiming that their ``orbit library analysis'' ruled out extreme orbital anisotropies at least in the case of NGC 3379. More recent modeling (De Lorenzi et al. 2008b) indicates again that there are mass-shape-anisotropy degeneracies that need to be resolved to reach a firm conclusion. In addition, a study of two dozens of globular clusters in NGC 3379 (Pierce et al. 2006) suggests that some dark matter may be present there. A recent long-slit spectroscopic study by Forestell \\& Gebhardt (2008) seems to indicate the presence of a dark matter halo around NGC 821, again in contradiction to what Romanowsky et al. (2003) concluded from the PNs in that galaxy. In summary, the case for no dark matter appears to be weakening, but it is still quite possible that less massive ellipticals have, on average, lower dark matter fractions, i.e. less easily detected dark matter halos (e.g. Napolitano et al. 2005). We seem to be in a situation that will require more careful studies of several galaxies before we can arrive to more solid conclusions.  The need to collect large PN samples means that we need to detect faint PNs. This can be better done using the largest telescopes available; so we had a strong motivation to extend the slitless radial velocity method, used in Paper I with VLT+FORS, to other large telescopes in the northern hemisphere. The FOCAS spectrograph of Subaru proved to be a good choice. As a first application of the slitless technique with FOCAS and Subaru, we have extended our search for PNs around NGC 4697 to larger angular distances. In this paper we present the resulting data and a preliminary analysis. In Section 2 we review the basic idea of slitless radial velocities and describe the Subaru + FOCAS observations in the outskirts of NGC 4697. Section 3 describes the reduction and calibration procedures. Section 4 deals with quality control of the slitless radial velocities using observations of the local PN NGC 7293. In Section 5 we give a catalog of 218 PNs detected with FOCAS, and compare the FOCAS with FORS slitless radial velocities for 158 PNs. Having verified the quality of the data, in Section 6 we discuss the line-of-sight velocity dispersion, and we analyze a plot of the escape velocity as a function of projected distance from the center of the galaxy. We adopt a distance of 10.5 Mpc for NGC 4697, taken from Paper I. Section 7 deals with rotation in the outskirts of NGC 4697. Section 8 gives a summary of results. ",
        "conclusions": "We have demonstrated how Subaru and FOCAS can be used to obtain accurate slitless radial velocities of extragalactic PNs. Working in two fields in the outskirts of NGC 4697, we discovered and provided radial velocities for 218 PNs, of which 162 had been previously observed with VLT+FORS. Where possible, we compared the FOCAS radial velocities with those previously obtained with FORS, finding good agreement. We now have PN kinematic information extending out to 5 effective radii from the center of NGC 4697 (its effective radius is 66 arcsec). In particular, we have significantly improved the kinematic information at angular distances larger than 230 arcsec from the center of NGC 4697.  We have found the following: (1) the losvd at radii between 300 and 400 arcsec from the center of NGC 4697 is below 100 km s$^{-1}$, i.e. a bit lower than obtained from the much smaller FORS outer PN sample by De Lorenzi et al. (2008a). This makes it difficult to narrow down the range of acceptable dark matter models. At least we can say that the dark matter halo around NGC 4697 is not as conspicuous as the one around NGC 5128. (2) the rotational signal, which is obvious close to the center, is no longer visible at about 5 effective radii from the center. By comparison to NGC 5128, it seems that much less angular momentum resides in the outskirts of NGC 4697. We need similar kinematic information about more elliptical galaxies.  A more quantitative interpretation of these results concerning NGC 4697 will require careful global modeling of the full, enlarged, FOCAS+FORS PN data set, plus existing absorption-line kinematic data, using NMAGIC or similar codes, following procedures described by De Lorenzi et al. (2008a)."
    },
    "0809/0809.0720_arXiv.txt": {
        "abstract": "We investigate the structure of magnetic field amplified by turbulent velocity fluctuations, in the framework of the kinematic Kazantsev-Kraichnan model. We consider Kolmogorov distribution of velocity fluctuations, and assume that both Reynolds number and magnetic Reynolds number are very large. We present the full numerical solution of the model for the spectra and the growth rates of magnetic fluctuations. We consider astrophysically relevant limits of large and small magnetic Prandtl numbers, and address both helical and nonhelical cases. ",
        "introduction": "Magnetic fields are found everywhere in the universe, in planets and stars, in galaxies and galaxy clusters. Magnetic fields in astrophysical systems are usually generated in a broad interval of scales, ranging from small resistive scales to large scales exceeding the correlation length of plasma motions.  One of the most important, challenging and still open questions in astrophysics is how cosmic magnetic fields have been generated and what is their structure. The prevailing theory for the origin of magnetic fields is dynamo action, which is stretching of magnetic field lines by random motion of highly conducting plasmas or fluids in which these lines are frozen \\citep[e.g.,][]{VZ_72,parker,lynden-bell,zweibel,brandenburg2,kulsrud1,schekochihin,kulsrud2}. Small-scale magnetic fields, that is fields at the scales smaller than the velocity field scales, are generally expected under this mechanism, while  large-scale magnetic fields, correlated at scales larger than the correlation scale of a velocity field, can be generated if some additional conditions are met, such as the condition that the velocity field ${\\bf v}({\\bf x, t})$ is not mirror symmetric, say, possesses nonzero kinetic helicity $H=\\int {\\bf v}\\cdot ({\\nabla \\times {\\bf v}})\\, d^3 x\\neq 0$ \\citep{skr,moffatt78}.  Assume that the velocity fluctuations have the typical correlation scale~$l_0$, and the typical rms value~$v_0$, and the plasma has kinematic viscosity~$\\nu$ and magnetic diffusivity~$\\eta$ (the latter is proportional to electrical resistivity). The range of scales available for velocity and magnetic fluctuations can be characterized by the Reynolds number ${\\rm Re}\\sim l_0v_0/\\nu$ and the magnetic Reynolds number ${\\rm Rm}\\sim l_0v_0/\\eta$, respectively. In astrophysical applications both numbers are very large (for example, in a protogalaxy, where dynamo action is believed to operate, ${\\rm Re}$ and ${\\rm Rm}$ reach values $\\sim 10^{5}$ and $\\sim 10^{26}$, respectively). Their ratio, the magnetic Prandtl number ${\\rm Pm=Rm/Re}$ can be either large or small. For example, in galaxies and galaxy clusters ${\\rm Pm}\\gg 1$, while in planets and stellar interiors ${\\rm Pm}\\ll 1$. Because of the vast range of scales available for magnetic and velocity fluctuations, and generally strongly disparate magnetic and velocity dissipation scales, present-day direct numerical simulations cannot directly address astrophysical magnetohydrodynamic (MHD) regimes. Indeed, maximal Reynolds and magnetic Reynolds numbers accessible with numerical simulations are hopelessly small, of the order of $10^3-10^4$, in which case the typical magnetic Prandtl numbers are not significantly different from ${\\rm Pm}\\sim 1$. A physical picture of magnetic dynamo action and an effective analytic framework for investigating astrophysical dynamo action are therefore in demand.  The first step in understanding dynamo action is to understand the initial, kinematic stage of magnetic field amplification. In this regime, the magnetic field is weak and does not affect the velocity fluctuations. Magnetic field evolution is therefore fully described by the induction equation in which the velocity field is prescribed independently of the magnetic field. An effective framework in this case is provided by the so-called Kazantsev-Kraichnan model, where the velocity field is assumed to be a random Gaussian short-time-correlated field. This formal simplification allows for analytic solutions of the model while capturing the essential physics of the phenomenon. We should note however that even with this simplification the model is nontrivial and its general solution is not known.  Only certain special cases have been solved so far, which reveal a good agreement with numerical simulations in the parameter range accessible to numerics \\citep[e.g.,][]{maron,haugen,boldyrev-cattaneo}.  The Kazantsev-Kraichnan dynamo model allows one to answer the fundamental questions concerning the possibility of turbulent dynamo action for given Reynolds and magnetic Reynolds numbers, the spectrum of growing magnetic fluctuations, the conditions for large-scale magnetic field amplification, etc. In many instances, the results obtained in high-resolution numerical simulations were predicted by the model well before such simulations become available. We therefore believe that the model can provide a valuable insight into astrophysical dynamo regimes that can hardly be accessed through direct numerical simulations in foreseeable future.  Our preliminary results aimed at the full numerical characterization of the kinematic dynamo action in the Kazantsev-Kraichnan framework  were presented in~\\citet{malyshkin}. In particular, we investigated the growth rates of magnetic fluctuations in the velocity field with the Kolmogorov energy spectrum. In the present paper, we explore the spatial structure of the growing magnetic eigenmodes. We assume the Kolmogorov spectrum of velocity fluctuations and address extremely large Reynolds numbers, up to~${\\rm Re}\\sim 10^9$. This allows us to study astrophysically relevant limits of very large and very small magnetic Prandtl numbers. For completeness, we perform our analysis for the cases of both large and small kinetic helicity. In the helical case we also discuss the relevance of the conventional $\\alpha$-model for the description of large-scale dynamo action, in the case when small-scale magnetic fluctuations are amplified as well.  In the next section, we describe the Kazantsev-Kraichnan model of kinematic dynamo action. In Section~3, we present our results, and in Section~4 we give our conclusions. ",
        "conclusions": "We have presented the full numerical characterization of the kinematic Kazantsev-Kraichnan dynamo model in the case of the Kolmogorov scaling of the velocity field. Our main conclusion is that the structure and the characteristic correlation scales of the magnetic field amplified by helical kinematic dynamo action are determined by both bound (localized) and unbound (nonlocalized) growing eigenmodes. In particular, the large-scale component of the field is defined by the unbound eigenmodes and by the shallow bound eigenmodes. Because these shallow bound eigenmodes have growth rates higher than those of all unbound eigenmodes, at any given scale the former may rapidly become dominant over the latter. In practical applications, this means that the shallow bound modes, rather than the unbound modes, are likely to become essential in the large-scale magnetic field configurations in astrophysical systems. In this case the conventional $\\alpha$-dynamo model~\\citep{skr,moffatt78,kulsrud1} gives an inadequate description of the large-scale magnetic field even at the kinematic stage of dynamo action.  The $\\alpha$-model becomes inapplicable in this case because it uses a critical assumption that small-scale fluctuations of the velocity and magnetic fields are much weaker and concentrated at the scales much smaller than the scales of the growing large-scale field. Under this assumption the $\\alpha$-model is obtained by averaging the induction Equation~(\\ref{induction}) over these small-scale fluctuations to obtain a linear and homogeneous differential equation for the large-scale mean magnetic field.\\footnote{The $\\alpha$-model equation for the mean field $\\overline{\\bf B}$ is $\\partial_t \\overline{\\bf B}=\\alpha \\nabla \\times \\overline{\\bf B}+\\beta \\nabla^2 \\overline{\\bf B}$. This is a linear and homogeneous differential equation with constant coefficients $\\alpha\\sim \\overline{{\\bf v}\\cdot ({\\nabla \\times {\\bf v}})} \\tau_0$ and $\\beta \\sim v_0 l_0$, where $v_0$ is a characteristic velocity, $l_0$ is a characteristic scale, and $\\tau_0\\sim l_0/v_0$ is a characteristic decorrelation time of fluid fluctuations. } Thus, the $\\alpha$-model misses all growing bound magnetic eigenmodes, including the essential shallow bound eigenmodes that determine the eventual large-scale configuration of the magnetic field.  We thank Fausto Cattaneo for many useful and stimulating discussions. This work was supported by the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas at the Universities of Chicago and Wisconsin-Madison. S.B.~is supported by the U.S. Department of Energy under the grant no.~DE-FG02-07ER54932."
    },
    "0809/0809.5043_arXiv.txt": {
        "abstract": "Non-electromagnetic emission from cosmic ray particles accelerated in extreme environments has been studied using different variations of semi-classical formalisms. As the energy loss mechanisms of such particles is of great interest, one must improve in some sense the previously intuitive description of classical sources in the presence of fields. In this brief note we evaluate the role played by the classical and dimensional \\emph{proper acceleration} of a particle in radiation processes. One intends to show that the introduction of the acceleration as a measure parameter gives an adequate scale for considering the processes of massive particle emission. If the acceleration threshold is not attained one is considering a regime where the processes of massive particle production are suppressed. The analysis is performed in a \\emph{semiclassical} formalism once applied in different contexts and already checked for consistency in a serie of papers. As a further application of the results we evaluate the possibility of a non negligible meson production by protons accelerated in the framework of polar cap models operating in electromagnetic fields of pulsars. It is shown that inside a systems endowed with magnetic fields $B\\geq 10^{12}$ G the meson emission by protons must not be disregarded ",
        "introduction": "\\label{Int}   \\par The rates and emitted powers of processes in the presence of (classical) $F_{ab}$ fields are usually obtained in the context of classical or semiclassical approachs. It is argued (based on the calculations of the electromagnetic spectra of radiation) that to disregard quantum effects it is mandatory that the source of mass $m$ obeys \\cite{1965ARA&A...3..297G, Zharkov, Erber:1966vv}, \\begin{equation} \\chi = \\frac{B^{\\star}}{B_{\\rm cr}} \\ll 1,\\label{chifactor} \\end{equation} \\noindent where $B^\\star\\equiv \\gamma B$ is the magnitude of the magnetic component of the electromagnetic field as measured in the instantaneous reference frame of the source; $\\gamma$ is the usual Lorentz factor and $B_{\\rm cr}=m^2/e$ is the critical field, which gives the value of the magnetic field where the quantum effects must be taken into account. If $\\chi \\geq 1$, the backreaction effects become important and a full quantum formalism must be employed. Otherwise if condition (\\ref{chifactor}) is satisfied the classical calculations are accurate and the associated emitted power depends only on the proper acceleration of the source \\cite{jackson}. For instance in the presence of a dipolar magnetic field the proper acceleration of the (ultra-relativistic) emitting source is $a\\approx \\gamma eB/m$ and consequently the limit given by the parameter (\\ref{chifactor}) can be rewritten as \\begin{equation} \\chi = \\frac{a}{m} \\ll 1 \\qquad \\longleftrightarrow \\qquad a\\ll m.\\label{accel} \\end{equation} \\noindent  The $\\chi$ invariant is usually defined in a frame independent way as $\\chi\\equiv \\sqrt{(eF^{ab}p_a)^2}/m^3$ and depends on the electromagnetic field ($F_{ab}$) and on the momentum ($p_a$) of the source. \\par In this brief research note we explore the role played by the classical acceleration on processes of emission of massive particles. Essentialy, we specify the ranges of the proper acceleration of the source where a given emission channel of massive fields becomes relevant. At the same time we introduce a semi-quantitative argument for the validity of the classical approximation when one is interested in such a class of processes. Furthermore we expect that the present framework can give some insight about the previous arguments on the applicability of effective approachs. As a further application we analyze the possiblility of extra mechanisms of particle production inside a particular engine of ultra-high energy cosmic rays (UHECR) acceleration that may be useful for some astrophysical observations. The range of proper acceleration where the process becomes favored is a natural consequence of the basic assumptions of the formalism. It is shown that if the threshold for the massive particle production is not taken as a measure parameter it can lead to different conclusions about the relevance of such class of process. \\par We adopt natural units where $c = \\hbar = 1$ unless stated otherwise and  Minkowski signature of the metric $\\eta_{ab}=(+,-,-,-)$ ",
        "conclusions": "\\par In this brief note we have evaluated the role played by the proper acceleration of a source when one is considering the classical limit of field methods. One can see from Eq.(\\ref{final}) and Fig. (\\ref{bpplane}) that the respective threshold for meson production is realized for pulsars endowed with magnetic fields of order $B = 10^{12}-10^{15}$ Gauss and respective periods, $P \\approx (1 - 10)$ ms; for higher periods the corresponding magnetic fields would be $B > 10^{15}$ G. Those statements follows directly from the lower limit for acceleration of the sources, $a\\geq m_x$; it has not been taken into account in the previous works on emission of massive meson fields \\cite{2008JCAP...08..025H,1538-4357-525-2-L117}. The absence of this lower limit for favouring the process has lead to the conclusion that those processes could never take place in an astrophysical environment due the dominance of electromagnetic losses. Then if one is dealing with classical or semiclassical treatment of the emission rates the lower (threshold for massive particle emission) $a_{\\rm source} > m_{\\rm emitted}$ and the upper (non-recoil condition) limit $a_{\\rm source} < m_{\\rm source}$ must constrain the range of calculations. \\par It is desirable to account for the consequences of this extra mechanism of energy loss in models of particle acceleration where high magnetic fields of pulsars are needed. \\ack Thanks to Dr. C. O. Escobar, A. Saa, G. E. A. Matsas, R. Opher and V. Pleitez for suggestions. The author is indebted to O. A. T. Dias for help with the manuscript and to S. Razzaque for illuminating discussion  and for calling my attention to the Ref. \\cite{2008JCAP...08..025H}."
    },
    "0809/0809.3665_arXiv.txt": {
        "abstract": "The X-ray Imaging Spectrometer (XIS) on board the Suzaku satellite is an X-ray CCD camera system that has superior performance such as a low background, high quantum efficiency, and good energy resolution in the 0.2--12~keV band. Because of the radiation damage in orbit, however, the charge transfer inefficiency (CTI) has increased, and hence the energy scale and resolution of the XIS has been degraded since the launch of July 2005.  The CCD has a charge injection structure, and the CTI of each column and the pulse-height dependence of the CTI are precisely measured by a checker flag charge injection (CFCI) technique. Our precise CTI correction improved the energy resolution from 230~eV to 190~eV at 5.9~keV in December 2006.  This paper reports the CTI measurements with the CFCI experiments in orbit.  Using the CFCI results, we have implemented the time-dependent energy scale and resolution to the Suzaku calibration database. ",
        "introduction": "% After the first successful space flight use of the X-ray charge coupled device (CCD) of the SIS (\\cite{Burke93}) on board ASCA, the CCD has been playing a major role in imaging spectroscopy in the field of X-ray astronomy.  However, the charge transfer inefficiency (CTI) of X-ray CCDs increases in orbit due to the radiation damage; the CTI is defined as the fraction of electrons that are not successfully moved from one CCD pixel to the next during the readout.  Since the amount of charge loss depends on the number of the transfers, the energy scale of X-ray CCDs depends on the location of an X-ray event. Furthermore, there is a fluctuation in the amount of the lost charge. Therefore, without any correction, the energy resolution of X-ray CCDs in orbit gradually degrades. In the case of the X-ray Imaging Spectrometer (XIS)~\\citep{Koyama07} on board the Suzaku satellite~\\citep{Mitsuda07} launched on July 10, 2005, the energy resolution in full width at half maximum (FWHM) at 5.9~keV was $\\sim$140 eV in August 2005, but had degraded to $\\sim$230 eV in December 2006.  The increase of the CTI is due to an increase in the number of charge traps at defects in the lattice structure of silicon made by the radiation. Since the trap distribution is not uniform, it would be best if we could measure the CTI of each pixel as Chandra ACIS~\\citep{Grant04}. In the case of the XIS, however, it is impossible to measure the CTI values of all the pixels, mainly because the onboard calibration sources do not cover the entire field of view of the XIS. Therefore, we use the CTI of each column to correct the positional dependence of the energy scale.  The XIS is equipped with a charge injection structure~\\citep{Prigozhin04,Bautz04,LaMarr04,Prigozhin08} which can inject an arbitrary amount of charge in arbitrary positions. Using this capability, we can precisely measure the CTI of each column~\\citep{Nakajima08}.  By applying the column-to-column CTI correction, the positional dependence of the CTI corrected energy scale is greatly reduced, and the over-all energy resolution is also improved~\\citep{Nakajima08}.  In \\citet{Nakajima08}, the results of the CTI correction was mainly based on the ground-based charge injection experiments. In-orbit measurements were limited within one year after the launch. This paper reports more comprehensive and extended  in-orbit experiments up to two years after the launch. The results are based on the data with the normal full window mode \\citep{Koyama07} without a spaced-row charge injection\\footnote{After October 2006, the spaced-row charge injection \\citep{Uchiyama09} has been a normal observation mode, and hence we should use different correction method.} \\citep{Uchiyama09}, and have been implemented to the Suzaku calibration database and applied to all the data obtained with the same mode. All the errors are at the 1$\\sigma$ confidence level throughout this paper unless mentioned. ",
        "conclusions": "We have conducted the CFCI experiments six times in orbit. The CTI correction has been done with the CFCI results. We calibrated the energy scale of the XIS precisely using the onboard calibration sources and 1E$\\,0102.2-7219$. Our calibration results have been applied to all the data  obtained with the normal full window mode without the spaced-row charge injection. The results of the CFCI experiments and the current calibration status are summarized as follows:  \\begin{enumerate} \\item We determined the CTI\\,1 and CTI\\,2 values of each column precisely based on the data of the CFCI experiments.  We also found that the pulse height dependence of the CTI does not change with time.  \\item After the column-to-column CTI correction, we determined the $PH_{\\rm o}$-$E$ relation. We also modeled the time-dependent energy resolution.  \\item The uncertainty of the energy scale is $\\pm$ 0.2 \\% for the O\\emissiontype{VIII} K$\\alpha$ line ($\\sim$ 0.65~keV) of 1E$\\,0102.2-7219$, and $\\pm$ 0.1 \\% for the Mn\\emissiontype{I} K$\\alpha$ line ($\\sim $5.9~keV) of the calibration sources.  \\item With the column-to-column CTI correction, the energy resolution at 5.9~keV in December 2006 was greatly improved from $\\sim$230~eV to $\\sim$ 190~eV.  \\end{enumerate}  \\bigskip The authors thank all the XIS members for their support and useful information. This work was supported by the Grant-in-Aid for the Global COE Program \"The Next Generation of Physics, Spun from Universality and Emergence\" from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. M.O., H.U., and H.N. are financially supported by the Japan Society for the Promotion of Science.  H.M. is also supported by the MEXT, Grant-in-Aid for Young Scientists~(B), 18740105, 2008, and by The Sumitomo Foundation, Grant for Basic Science Research Projects, 071251, 2007. H.T. and K.H. were supported by the MEXT, Grant-in-Aid 16002004.  \\appendix"
    },
    "0809/0809.4759_arXiv.txt": {
        "abstract": "This chapter describes the Galactic star forming regions in the constellation Cassiopeia, in the Galactic coordinate range $120\\deg \\la l \\la 130\\deg$, $-5\\deg \\la b \\la 15\\deg$. At $|b| > 10\\deg$ the nearby clouds L\\,1333 and L\\,1340 are found in this region. The local arm of the Galaxy in Cassiopeia contains only a few star forming regions, smaller and less active than the OB~associations of the neighboring Cepheus. Five members of this system, LkH$\\alpha$\\,198 and its environment, L\\,1287, L\\,1293, L\\,1302/NGC\\,255, and S\\,187 are discussed.  Several more distant OB~associations and giant star forming regions in Cassiopeia are associated with the Perseus arm at 2.0--3.0\\,kpc. Among these, the Herbig~Be star MWC\\,1080 is discussed in this chapter. ",
        "introduction": "The large-scale distribution of the dark clouds in Cassiopeia, adopted from the Atlas and Catalog of Dark Clouds \\citep{DUK}, is displayed in Fig.~\\ref{Fig1}. The most prominent clouds and young stellar objects are marked.  The high Galactic latitude region of the constellation Cassiopeia contains a few nearby star forming molecular clouds. The nearest cloud is Lynds~1333 at ($l,b$)=(129\\deg,\\linebreak +15\\deg) and at $d \\approx 180$\\,pc, associated with a few low-mass pre-main sequence stars. Lynds~1340, a region of intermediate and low mass star formation, can be found at\\linebreak ($l,b$)=(130\\deg,+11\\fdg5) at $d \\approx 600$\\,pc. Towards lower latitudes, the cloud complex\\linebreak  L\\,1355/L\\,1358 is located at ($l,b$) $\\sim$ (133\\fdg5,+8\\fdg6). This complex is illuminated by a loose group of B and A stars and the classical Cepheid SU~Cas \\citep{Turner}. The average distance of these stars, determined by \\citet{Turner} is $d = 258 \\pm 3$\\,pc. No star formation has been observed in these clouds. The starless cores of L\\,1355 and L\\,1358 have been targets of several studies aimed at observing initial conditions of low-mass star formation and protostellar infall \\citep*{Lee99,Lee01,Park04}.  Star forming regions at various distances can be found at latitudes ($|b| < 10\\deg$). The catalog of CO clouds associated with IRAS point sources, published by \\citet{Kerton}, contains numerous likely star forming regions which deserve follow-up studies. L\\,1265, L\\,1287, and S\\,187, located between 600--1000\\,pc, are associated with the Orion arm of our Galaxy. The most distant star forming regions, associated with the Perseus arm, include several OB associations \\citep[e.g.][]{Humphreys,GS92} located between 1.8--2.8\\,kpc, and  the Herbig~Be star MWC\\,1080, one of the best studied members of this class of objects. The most prominent member of this system, the giant star forming complex W3/W4/W5 is discussed in the chapter by Megeath et~al. The regions discussed in the present chapter are  listed in Table~\\ref{Tab_cloud}.  \\begin{figure*}[ht!] \\centerline{ \\includegraphics[width=5.25in,draft=False]{cas_large4.SR.eps} } \\caption{Distribution of visual extinction in the Cassiopeia region \\citep[adopted from][]{DUK}. The prominent clouds and the young stellar objects associated with the clouds discussed in this chapter are marked. Filled circles are outflow sources, plusses are classical T Tauri stars, asterisks mark Herbig Ae/Be stars, and diamonds weak-line T Tauri stars. Rectangles mark the nominal boundaries of the OB associations \\citep{Humphreys}. Dashed lines indicate the boundaries of the constellation. Large circles drawn by radial dashes show the approximate positions of the Cepheus Flare Shell and Cep~OB4 Shell \\citep{Olano}.} \\label{Fig1} \\end{figure*}  \\begin{table} \\caption{Star-forming molecular clouds in Cassiopeia} \\label{Tab_cloud} \\smallskip \\begin{center} {\\footnotesize \\begin{tabular}{llcl} \\tableline Dark Cloud$^a$ & CO cloud$^b$~  & d\\,(pc) & Associated YSO \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} L\\,1238 & 111.7+00.0 & 2200 & MWC\\,1080 \\\\ L\\,1265 & 117.8$-$03.6  & 600  &  LkH$\\alpha$ 198 \\\\ $\\cdots$ &  120.1+03.0$^c$  & 850 & IRAS 00213+6530, IRAS 00259+6510 \\\\ L\\,1287, TDS463 & 121.3+00.6 & \t850 & RNO\\,1, IRAS 00338+6312 \\\\ L\\,1293 & 121.7+00.2 & 850 & IRAS 00379+6248 \\\\ L\\,1302 & 122.0$-$01.1 & 600 & BD\\,+61\\deg154, LkH$\\alpha$201 \\\\ L\\,1317 & 126.6$-$00.6 & 800 & S\\,187  \\\\ L\\,1333 & 129.0+13.8 & 180 & IRAS 02086+7600 \\\\ L\\,1340 & 131.1+11.6 & 600 & RNO\\,7,8,9 \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{4}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^a$ \\citet{Lynds,TDS};  }}\\\\[1ex] \\multicolumn{4}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^b$\\citet{Yonekura}; }}\\\\[1ex] \\multicolumn{4}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^c$This cloud lies within the region studied in this chapter, but, as  it is associated with Cep~OB\\,4 in the literature, it is discussed in the Cepheus chapter.}}\\\\  \\end{tabular}} \\end{center} \\smallskip \\end{table} ",
        "conclusions": ""
    },
    "0809/0809.1394.txt": {
        "abstract": "{  \\begin{abstract}  We report on the results from the first six months of the Catalina Real-time Transient Survey (CRTS). In order to search for optical transients with timescales of minutes to years, the CRTS analyses data from the Catalina Sky Survey which repeatedly covers twenty six thousand of square degrees on the sky. %The CRTS uses data from the Catalina Sky Survey, which repeatedly covers thirty five thousand of square %degrees of the sky, in order to search for stationary optical transients with timescales of minutes to years. The CRTS provides a public stream of transients that are bright enough to be followed up using small telescopes.  Since the beginning of the survey, all CRTS transients have been made available to astronomers around the world in real-time using HTML tables, RSS feeds and VOEvents.  As part of our public outreach program the detections are now also available in KML through Google Sky.  The initial discoveries include over 350 unique optical transients rising more than two magnitudes from past measurements.  Sixty two of these are classified as supernovae, based on light curves, prior deep imaging and spectroscopic data.  Seventy seven are due to cataclysmic variables (only 13 previously known), while an additional 100 transients were too infrequently sampled to distinguish between faint CVs and SNe.  The remaining optical transients include AGN, Blazars, high proper motions stars, highly variable stars (such as UV Ceti stars) and transients of an unknown nature.  Our results suggest that there is a large population of SNe missed by many current supernova surveys because of selection biases. These objects appear to be associated with faint host galaxies.  We also discuss the unexpected discovery of white dwarf binary systems through dramatic eclipses.  %} ",
        "introduction": "Time-domain astronomy is one of the most rapidly emerging areas of astronomy (Paczy\\'nski 2000). Several large experiments (LSST, Ivezic et~al. 2008; PanSTARRS, Hodapp et~al. 2004; %\\citep{Hod04}, PTF\\cite{}, SkyMapper, Keller et~al. 2007) are set to make a major impact in the near future by covering thousands of square degrees with targets ranging from distant supernovae to near earth asteroids. Past time-domain surveys have mainly concentrated on areas of tens to hundreds of square degrees and have typically searched for specific types of astronomical transients such as microlensing (MACHO, Alcock et~al. 1993; OGLE, Udalski et~al. 1994; EROS, Aubourg et~al. 1995), GRB afterglows (ROTSE; Akerloff et~al. 2000) and supernovae (KAIT; Filippenko et~al. 2001). %such as MACHO\\cite{Alc01}, OGLE \\cite{} (Microlensing), DLS \\cite{Becker} (Lensing), ROTSE (GRBs)  Deep surveys for variability have been carried out over small areas ($< 25$ $\\rm deg^2$) by the Subaru/XMM-Newton Deep Survey (SXDS; Morokuma et al. 2008), the Deep Lensing Survey (DLS; Becker et al. 2004) and the Faint Sky Variability Survey (FSVS; Huber et al. 2006) wheras large-area surveys, such as the Sloan Digital Sky Survey (SDSS; York et~al. 2000), the Two Micron All Sky Survey (2MASS; Skrutskie et~al. 2006), and the Galaxy Evolution Explorer (GALEX; Martin et~al. 2005) have generally not been synoptic. One survey intermediate between past deep surveys, and wide future surveys was carried out by the SDSS consortium to discover type Ia supernovae. This survey covered 300 $\\rm deg^2$ in 5 optical bands to $r \\sim 23$ during nine months spaced over 2005 to 2007 (Sesar et~al. 2007).   \\subsection{Current Transient Surveys}  Current wide-field transient surveys include the Robotic Optical Transient Search Experiment (ROTSE-III; Akerlof et al. 2003) and the All Sky Automated Survey (ASAS-3; Pojmanski 2001). However, like many other surveys, both are targeted to the discovery and characterization of GRB afterglows. ASAS-3 surveys $\\sim$ 30,000 square degrees of the southern sky to  $V\\sim13.5$, every 2 nights (Pojmanski 2001), while the four telescopes of the ROTSE-III survey routinely cover 1260 square degrees to $R\\sim17.5$ (Rykoff et al. 2005; Yost et al. 2007).  In the time between targeted surveys, and future wide, deep, transient surveys, two experiments have been searching tens of thousands of square degrees of the sky for transients; the Palomar-Quest digital synoptic sky survey (PQ\\footnote{http://palquest.org}, Djorgovski et al. 2008a), and the Catalina Real-time Transient Survey (CRTS).  The PQ survey started analyzing driftscan data for optical transients in real-time in August 2006, while the CRTS began in November 2007. The transients discovered in PQ data will be presented in Mahabal et~al. (2008a).  Both CRTS and PQ experiments use a purpose-built pipeline for real-time transient detection, analysis and distribution (Drake et~al. 2008a).  These surveys are now set to provide estimates of the rates and types of optical transients that can be expected from future synoptic surveys.  In these experiments, unlike most current and past optical transient surveys, the results are made public within minutes of discovery using the IVOA standard VOEvent protocol and the VOEventNet astronomical event distribution network (Drake et al. 2007a). In addition, VOEventNet, in collaboration with Google, makes transient astronomy available to the world in real-time via a dedicated layer in Google Sky\\footnote{http://earth.google.com/sky}. Current events available through VOEventNet also include MOA and OGLE microlensing, GCN gamma ray burst alerts, and PQ transients.  The CRTS routinely searches for Optical Transients (OTs) in data from the Catalina Sky Survey Schmidt telescope (CSS). In this paper we present results from the first six months.  Firstly, we will discuss the observations and data analysis ($\\S 2$). We will then list the current results ($\\S 3$) and discuss some of the types of transients discovered. Finally, we will summarize the current findings ($\\S 4$). ",
        "conclusions": "We have presented a few of the types of optical transients found in the first six months of CRTS analysis of CSS observations. With the observational parameters of this survey the transient discoveries are dominated by two types of objects, DNe and SNe.  Apart from Miras, few of the many kinds of highly variable stars were discovered. We believe the main reasons for this are our high detection threshold (2 mags), use of catalogs based on co-added images (where cyclic variability on short timescales are averaged out), and the natural limit to the number of stars observed by only observing fields with $|b| > 10\\arcdeg$.  While photometric follow-up was taken for many of the OTs in the first six months (Mahabal et al. 2008c), spectroscopic follow-up was carried out to classify a small number of the candidates. Clearly, one of the most important next steps is that of classifying these and other such transients, such that the expensive spectroscopic resources are used to a minimum. Given that surveys tend to provide a small number of data points, this requires optimizing photometric follow-up in combination with probabilistic classification techniques. We have been making steady progress to tackle this hard problem using advanced mathematical and statistical methodologies (see e.g. Mahabal et al. 2008d).  As the survey progresses we will also optimize transient discovery by employing machine learning techniques (Borne 2008). This will enable unsupervised rapid follow-up to be carried out on the most interesting objects.  The CRTS has a firm commitment to making discoveries and the associated meta-data public as quickly as possible. In this way we hope all aspects and types of OTs can be explored. Many of the CRTS OTs were posted in ATels and a number of these were independently observed after these announcements. We chose to announce most of the obvious SNe with bright hosts in CBETs and these were usually spectroscopically followed and thus confirmed. Although significant follow-up would be required to characterize all the OTS discovered by CRTS, this is possible as the SDSS-II Supernova Survey acquired $\\sim180$ nights of spectroscopic time on 2.4 to 11 meter telescopes, to follow SNe found during $\\sim70$ nights of imaging (Frieman et~al.  2008).  All OTs will continue to be made publicly available from VOEventNet and those events that appear the most interesting will be announced in ATels and CBETs. In addition, we will be using the two additional CSS survey telescopes to search for transients.  These will allow us to find similar numbers of transients in the Southern sky as well as fainter transients in the North.  In conclusion, while there is little doubt that many of the ``major advances in our understanding of the Universe have historically arisen by improvements in our ability to see'' (Ivezic et~al. 2008),  %\\cite{Ive08}, we believe that significant advances have, and will continue to be made by also changing how and what is observed. The upcoming next-generation of transient surveys (PanSTARRs and LSST) promise a unique ability to discover rare types of optical transients in the faint time-domain sky. However, it is clear that the full characterization of many kinds of transients will require follow-up observations and these are most easily achieved when the transients are both accessible to the entire astronomical community, and bright enough to be followed by the large numbers of existing small telescopes. %The CRTS delivers both."
    },
    "0809/0809.0934.txt": {
        "abstract": "Several accurate analyses have revealed a statistically significant North-South ecliptic asymmetry in the large-angle correlations strength of the Cosmic Microwave Background (CMB) radiation temperature field data from the Wilkinson Microwave Anisotropy Probe (WMAP). This asymmetry is inconsistent with the statistical isotropy expected in the concordance cosmological model $\\Lambda$CDM. It has been suggested that a possible cause-effect relationship exists between this large-angle anisotropy and the anomalous CMB quadrupole-octopole planes alignment. In turn, this later phenomenon (or both) would be a consequence of one or more of the following undesired effects in CMB data: a systematic error in the data processing or in the instrument characterization, residual foregrounds, and large-angle correlations induced by the incomplete sky CMB data (cut-sky masks are needed to reject Galactic foregrounds). Here, it is proved that the North-South asymmetry is unrelated to the quadrupole ($\\ell=2$) and the octopole ($\\ell=3$) properties because we find, at high confidence levels, such large-angle anisotropy in three- and five-year WMAP CMB maps containing only the multipole components $4 \\leq \\ell \\leq 10$. The statistical significance depends on both, the CMB map analyzed as well as the cut-sky mask applied to exclude foregrounds. In general, we obtain that the significance level of the North-South asymmetry is less in five-year WMAP data with KQ75 ($\\gtrsim 90\\%$ CL) than it is in three-year data with Kp0 ($\\gtrsim 96\\%$ CL). For instance, in the WMAP ILC-five-year map with the KQ75 mask (a sky cut of 28.4\\%) this phenomenon is observed at 92.7\\% CL, whereas for the WMAP ILC-three-year map with the Kp0 mask (a sky cut of 23.5\\%) this phenomenon appears at 96.5\\% CL. Moreover, it is also shown that this hemispherical asymmetry is unlikely due to systematics or foreground contaminants, because it is present in single-frequency, multifrequency, and cleaned ILC-type CMB maps. Additionally, robustness tests show that the statistical significance of the North-South asymmetry is affected by the use of the KQ85 and KQ75 masks in five-year single-frequency WMAP CMB maps, whereas it is insensitive with respect to application of the Kp2 and Kp0 masks in three-year single-frequency WMAP CMB maps. The confidence levels were obtained using sets of Monte Carlo CMB maps produced according to the $\\Lambda$CDM model. ",
        "introduction": "\\label{Introduction} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The five-year data release from the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{Hinshaw08,Gold08,Nolta08,Dunkley08,Komatsu08} have recently confirmed the validity of the concordance cosmological model $\\Lambda$CDM, which successfully describes the contents and evolution of the Universe. These, almost full-sky, precise cosmological data of\\/fer the unusual possibility to scrutinize the large-scale properties of the Universe, in particular those expected within this concordance model. One of these features concerns the set of Cosmic Microwave Background (CMB) radiation temperature f\\/luctuations which is assumed to be a stochastic realization of a Gaussian random field, implying that their angular distribution on the celestial sphere is statistically isotropic at all angular scales. Examinations of the CMB maps from the first-year~\\cite{Bennett03a,Bennett03b,Hinshaw03} and three-year WMAP data~\\cite{Hinshaw06,Jarosik06,Spergel06} have revealed highly significant departures from statistical isotropy, at large angular scales. Evidences of anomalous large-angle anisotropies come from the power asymmetry of the CMB angular correlations between the northern and southern ecliptic hemispheres (hereafter NS-asymmetry), with indications of a preferred axis of maximum hemispherical asymmetry~\\cite{Hansen04a,% Hansen04b,Bielewicz04,Eriksen04a,Eriksen05,Land04,Land05b,Wiaux06,Copi06,Huterer06,% Hansen06,PPG,Vielva06,Vielva07,Land07,Eriksen07,Copi05,BVWLF,Samal,BMRT,BMRT07,% Bunn08,Monteserin,Lew} (for a dif\\/ferent point of view see, e.g.,~\\cite{Hajian03,% Hajian04,Souradeep04,Souradeep05,Souradeep06,Hajian06}). Furthermore, other manifestations of CMB large-angle anomalies include the unlikely quadrupole-octopole planes alignment (referring both to the strong planarity of these multipoles as well as to the alignment between such planes~\\cite{TOH,OTZH,Weeks04,Wiaux06,Copi06,Abramo06,PPG}). Other unexpected results include deviations from Gaussianity~\\cite{Chiang03,% Vielva04,Copi04,Cruz05,Cruz07,McEwen06,BTV07,Martinez08}.  Attempts to explain anomalous CMB phenomena have also considered globally axisymmetric space-times~\\cite{Aurich05,Land05a,% Jaffe06,Hipolito,Cresswell,Ghosh}, motivated by the fact that CMB data seem to indicate a preferred direction in the space. Other studies, instead, have investigated physical mechanisms that break statistical isotropy, or during the epoch of inf\\/lation~\\cite{Gordon05,Contaldi07,Pullen,Ackerman,% Dutta,Dvorkin} or during the decoupling era~\\cite{Demianski,Pereira,Campanelli2,% Morales,BH,Kahniashvili1,Kahniashvili2}, in order to account for two (or more) large-angle CMB anomalies. The motivation for searching a scenario that correlates two CMB anomalies is to comprehend whether these phenomena have dif\\/ferent causes or instead they are manifestations of a unique ef\\/fect. For this, a number of studies have searched for a possible relationship between the NS-asymmetry and the CMB quadrupole-octopole alignment phenomena (see, e.g.,~\\cite{Wiaux06,Rakic07,Land07,BH}).  Although the origin of these CMB anomalies remains unknown, possible sources include CMB residual foregrounds~\\cite{Eriksen04b,WW,% OT,Cruz06,ASW,Lopez07,Chiang07}, and systematic errors~\\cite{Copi05,Copi06,% Helling06,Bunn07}. In particular, there are indications that unremoved contaminants~\\cite{Naselsky07} or omitted systematics~\\cite{Schwarz04} could be inf\\/luencing the reconstruction of the CMB quadrupole and octopole. Moreover, the WMAP cut-sky masks used to remove residual foregrounds probably induce spurious large-angle correlations, producing forged low-order CMB multipoles~\\cite{Bielewicz05}. Thereby ef\\/forts are being done by the WMAP science team to improve the data processing by minimizing the ef\\/fects caused by foregrounds and systematic errors~\\cite{Jarosik06,Gold08}.  Here we show that the NS-asymmetry phenomenon is not related to the CMB's quadrupole or octopole components. To obtain this result we perform a two-point angular correlations analysis in an array of large-angle spherical caps that scan the CMB sky, where the WMAP maps investigated have their corresponding quadrupole and octopole components removed. The outline of this paper is the following. In Sec.~\\ref{method}, we present the geometrical-statistical method, termed {\\em sigma map}, that led us to investigate the directional angular correlations strength in a set of WMAP CMB maps. In Sec.~\\ref{CMB analyses} we first discuss the maps to be analyzed, after that we apply our method to these CMB data, where the statistical confidence is evaluated according to Monte Carlo maps produced according to the $\\Lambda$CDM model. Finally, in Sec.~\\ref{conclusions}, we formulate our conclusions and final remarks.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We investigated a set of foregrounds-reduced and systematics-cleaned single-frequency, multifrequency, and ILC-type CMB maps from three- and five-year WMAP data. We performed the sigma-maps analyses of these CMB maps at large-angles corresponding to the multipoles range $4 \\leq \\ell \\leq 10$, and for dif\\/ferent three- and five-year WMAP cut-sky masks, that is Kp2, Kp0, KQ85, and KQ75.  We found that the single-frequency Q-3yr, V-3yr and W-3yr WMAP maps show, independently of the mask (Kp2 or Kp0) used to obtain the multipoles $4 \\leq \\ell \\leq 10$, the same sigma-map angular power spectra with confidence level $\\sim 96\\%$ in all the cases (see in the top panel of Fig.~\\ref{fig3}, symbols plotted close to the vertical line). This result indicates, at least concerning the sigma-map analysis, that the single-frequency Q-3yr, V-3yr, and W-3yr maps are reasonable foreground cleaned. As a consequence, the multifrequency CMB maps, that is VW, QVW, and ILC-type, produced with these single-frequency maps behaves accordingly. For the WMAP-5yr-KQ85 data (see in the top panel of Fig.~\\ref{fig3}, symbols plotted slighty shifted to the right) we found in all, but one, of these CMB maps, at 95\\% $-$ 99\\% CL, an uneven hemispherical distribution in the intensity of the two-point large-angle temperature correlations. The exceptional case refers to the Q-5yr map. On the contrary the single-frequency Q-5yr, V-5yr and W-5yr WMAP maps exhibit a disparate behavior as seen in their corresponding sigma-map angular power spectra, in both KQ85 and KQ75 mask cases. This probably suggests that these CMB maps have residual foregrounds that diminish but are not fully eliminated with the use of the more severe KQ75 mask.  We also notice, despite that the ILC-5yr-KQ75 and the ILC-3yr-Kp0 CMB maps are fully correlated (Pearson's correlation coef\\/ficient: 0.985180 and 0.985466, for the CMB data outside the KQ75 and Kp0 masks) their corresponding sigma maps are not so much. This fact motivated a detailed examination of the masking ef\\/fects on the sigma-map angular power spectra. To estimate the inf\\/luence of the dif\\/ferent masks employed we calculate the sigma maps for the ILC-3yr CMB map obtained after applying the five-year KQ75 mask, and for the ILC-5yr CMB map after applying the three-year Kp0 mask. The angular power spectra of sigma-map-3yr-KQ75, sigma-map-3yr-Kp0, sigma-map-5yr-KQ75, and sigma-map-5yr-Kp0 data are shown in the top panel of Fig.~\\ref{fig5}. We conclude that the KQ75 mask, as compared with Kp0, has a significant impact on the detection of the NS-asymmetry phenomenon decreasing its statistical significance. In fact, the KQ75 mask cuts a sky patch that turns out to be critical for the computation of the large-scale 2PACF in the sigma maps (as revealed by the dif\\/ference map: sigma-map-ILC-5yr Kp0 minus sigma-map-ILC-5yr KQ75, plotted at the bottom panel of Fig.~\\ref{fig5}).  Summarizing, first, our results prove that the NS-asymmetry phenomenon is present at the large angular scales defined by the CMB multipoles $4 \\le \\ell \\le 10$, at a high significance level, in a set of single and multifrequency CMB maps from three- and five-year WMAP data, where the statistical significance depends on both, the CMB map analyzed as well as the cut-sky mask applied to exclude foregrounds. In general, we obtain that the NS-asymmetry is statistically less significant in five-year WMAP data with KQ75 ($\\gtrsim 90\\%$ CL) than it is in three-year data with Kp0 ($\\gtrsim 96\\%$ CL). For instance, in the ILC-5yr-KQ75 map the NS-asymmetry is observed at 92.7\\% CL, whereas in the ILC-3yr-Kp0 map this phenomenon appears at 96.5\\% CL. Second, these results also show that the NS-asymmetry in WMAP maps is unrelated to the quadrupole and octopole CMB components because they were removed from the CMB maps in our analysis. Third, due to the whole set of sensitivity tests performed our results strongly suggest that this hemispherical power asymmetry, is unlikely correlated with spurious phenomena like residual foregrounds, masking applications, or known systematic ef\\/fects. Additionally, we found that the statistical significance of the NS-asymmetry phenomenon is af\\/fected by the use of the KQ85 and KQ75 masks in  five-year single-frequency WMAP CMB maps, whereas it is insensitive with respect to application of the Kp2 and Kp0 masks in three-year single-frequency WMAP CMB maps. Moreover, we also examined satisfactorily the robustness of our results under several parameters used in the application of the sigma-map method. The statistical significance of the sigma-maps-WMAP angular power spectra were calculated by comparing them against those obtained from $1\\,000$ Monte Carlo CMB maps produced in agreement with the WMAP $\\Lambda$CDM concordance model~\\cite{Nolta08}. Finally, as an illustrative case, we also performed the sigma-map analysis for the full-sky ILC-5yr-$\\ell_{4-10}$ CMB map, where we find the NS-asymmetry phenomenon at 95.8\\% CL for the angular scales $4 \\le \\ell \\le 10$, a result that is independent of the quadrupole and octopole CMB components. For completeness, we plot in Fig.~\\ref{fig6} the sigma-map-ILC-5yr-$\\ell_{4-10}$ (top panel) and its corresponding dipole term (bottom panel), which reveals that the NS-asymmetry axis for the ILC-5yr-$\\ell_{4-10}$ map points along the direction $(l, b) \\simeq (220^{\\circ},\\, 120^{\\circ})$.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 1 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\vspace{-1cm} \\includegraphics[width=15cm,height=24cm]{fig1.ps}  \\vspace{-3cm} \\caption{\\label{fig1} The top panel shows the CMB map ILC-5yr-KQ85-$\\ell_{4-10}$. The second, third, and fourth panels are the sigma-maps calculated from this CMB map considering spherical caps with apertures $\\gamma_0 = 30^{\\circ}, 45^{\\circ}, 60^{\\circ}$, respectively. The NS-asymmetry is clearly seen in the uneven hemispherical distribution of the angular correlations strength, depicted as red and blue pixels (large and small sigma-values, respectively). All the maps are plotted in Galactic coordinates.} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%% FIGURA 1 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%% FIGURA 2 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\vspace{-1cm} \\includegraphics[width=15cm,height=24cm]{fig2.ps}  \\vspace{-3cm} \\caption{\\label{fig2} The top panel shows the CMB map ILC-5yr-KQ75-$\\ell_{4-10}$. The second, third, and fourth panels are the sigma-maps calculated from this CMB map considering spherical caps with apertures $\\gamma_0 = 30^{\\circ}, 45^{\\circ}, 60^{\\circ}$, respectively. Again, the NS-asymmetry is clearly seen in the uneven hemispherical distribution of the angular correlations strength, depicted as red and blue pixels (large and small sigma-values, respectively). All the maps are plotted in Galactic coordinates. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%% FIGURA 2 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%% FIGURA 3 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\includegraphics[width=15cm,height=24cm]{fig3.ps}  \\vspace{-9cm} \\caption{\\label{fig3} Top: Angular power spectra $\\mbox{\\sc S}_{\\ell}$ of the sigma-maps-WMAP-KQ85-$\\ell_{4-10}$ compared with the mean (solid line) and their corresponding 95\\% CL (dashed line) of $1\\,000$ sigma-maps MC, obtained from an equal number of MC-KQ85-$\\ell_{4-10}$ CMB maps. The symbols representing data from three-year WMAP maps are plotted (in red) close to the vertical axis, while the symbols representing data from the five-year WMAP maps are plotted (in blue) slightly shifted to the right. The Q,  V,        W,    VW,       QVW,    OT,   PPG,  KNC, and ILC CMB maps are represented by the bullet, asterisk, star, triangle, circle, nabla, {\\bf\\large $\\times$}, diamond, and {\\bf +} symbols, respectively. The NS-asymmetry phenomenon in WMAP CMB maps is evidenced, except for the Q-5yr map, through the anomalously large dipole value of the sigma-maps-WMAP KQ85, at more than 95\\% CL. Bottom: This is the dipole term of the sigma-maps-ILC-5yr KQ85 obtained with $\\gamma_0 = 45^{\\circ}$ (actually the $\\gamma_0 = 30^{\\circ}, 60^{\\circ}$ cases are similar), where we observe that it points in the direction $(l, b) \\simeq (190^{\\circ},\\, 130^{\\circ})$. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%% FIGURA 3 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%% FIGURA 4 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\includegraphics[width=15cm,height=24cm]{fig4.ps}  \\vspace{-9cm} \\caption{\\label{fig4} Top: Angular power spectra $\\mbox{\\sc S}_{\\ell}$ of the sigma-maps-WMAP-KQ75-$\\ell_{4-10}$ compared with the mean (solid line) and their corresponding 95\\% CL (dashed line) of $1\\,000$ sigma-maps MC, obtained from an equal number of MC-KQ75-$\\ell_{4-10}$ CMB maps. The symbols representing data from three-year WMAP maps are plotted (in red) close to the vertical axis, while the symbols representing data from the five-year WMAP maps are plotted (in blue) slightly shifted to the right. The Q,  V,        W,    VW,       QVW,    OT,   PPG,  KNC, and ILC CMB maps are represented by the bullet, asterisk, star, triangle, circle, nabla, {\\bf\\large $\\times$}, diamond, and {\\bf +} symbols, respectively. We observe that the NS-asymmetry phenomenon in WMAP CMB maps is revealed at more than 98\\% CL, in all the single, multifrequency, and ILC-type WMAP-3yr-Kp0 maps and also in the KNC-5yr-KQ75 foreground-cleaned map. However, in all the other sigma-maps-WMAP-5yr-KQ75 maps the statistical significance is lower, and this matter has been suitably analyzed in Sec.~\\ref{results2}. Bottom: This is the dipole term of the sigma-maps-ILC-5yr KQ75 obtained with $\\gamma_0 = 45^{\\circ}$ (actually the $\\gamma_0 = 30^{\\circ}, 60^{\\circ}$ cases are similar), where we observe that it points in the direction $(l, b) \\simeq (180^{\\circ},\\, 130^{\\circ})$. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 4 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 5 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\vspace{-1cm} \\includegraphics[width=15cm,height=24cm]{fig5.ps}  \\vspace{-9cm} \\caption{\\label{fig5} Top: Sigma-map angular power spectra of the analyze CMB maps: ILC-3yr-KQ75, ILC-3yr-Kp0, ILC-5yr-KQ75, and ILC-5yr-Kp0, represented as circle, filled-circle (or bullet), square, and filled-square symbols, respectively. These results show that $\\mbox{\\sc S}_{\\ell}^{_{\\mbox{\\footnotesize ILC-3yr-KQ75}}} \\simeq \\mbox{\\sc S}_{\\ell}^{_{\\mbox{\\footnotesize ILC-5yr-KQ75}}} $ and $\\mbox{\\sc S}_{\\ell}^{_{\\mbox{\\footnotesize ILC-5yr-Kp0}}} \\simeq \\mbox{\\sc S}_{\\ell}^{_{\\mbox{\\footnotesize ILC-3yr-Kp0}}} $, demonstrating that the main reason for the dif\\/ferent statistical significance found in Fig.~\\ref{fig4} for the sigma-map-ILC-3yr and sigma-map-ILC-5yr is due to the distinct cut-sky area defined by the KQ75 and Kp0 masks. In fact, the NS-asymmetry phenomenon is present, at more than 96\\% CL, in those maps where the Kp0 mask has been applied (i.e., the ILC-3yr-Kp0 and ILC-5yr-Kp0 CMB maps). Bottom: This is the dif\\/ference map: sigma-map-ILC-5yr-Kp0 minus sigma-map-ILC-5yr-KQ75, which clearly reveals through its red-intense spots that the application of dif\\/ferent masks has a significant impact on the detection of the NS-asymmetry phenomenon: the KQ75 mask cuts a sky patch that turns out to be critical for the computation of the large-scale 2PACF in the sigma-maps. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 5 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 6 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\includegraphics[width=15cm,height=24cm]{fig6.ps}  \\vspace{-10cm} \\caption{\\label{fig6} For illustration we show the sigma-map-ILC-5yr-$\\ell_{4-10}$ computed from the WMAP-ILC-5yr full-sky CMB map (top panel) and its corresponding dipole term (bottom panel). Observe that this dipole term is pointing in the direction $(l, b) \\simeq (220^{\\circ},\\, 120^{\\circ})$. } \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURA 6 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.1456_arXiv.txt": {
        "abstract": "Type I planetary nebulae (PNe) have high He/H and N/O ratios and are thought to be descendants of stars with initial masses of $\\sim 3$ -- 8$\\Msun$. These characteristics indicate that the progenitor stars experienced proton-capture nucleosynthesis at the base of the convective envelope, in addition to the $slow$ neutron capture process operating in the He-shell (the $s$-process). We compare the predicted abundances of elements up to Sr from models of intermediate-mass asymptotic giant branch (AGB) stars to measured abundances in Type I PNe. In particular, we compare predictions and observations for the light trans-iron elements Se and Kr, in order to constrain convective mixing and the $s$-process in these stars.  A partial mixing zone is included in selected models to explore the effect of a \\iso{13}C pocket on the $s$-process yields.  The solar-metallicity models produce enrichments of [(Se, Kr)/Fe] $\\lesssim~0.6$, consistent with Galactic Type I PNe where the observed enhancements are typically $\\lesssim 0.3$ dex, while lower metallicity models predict larger enrichments of C, N, Se, and Kr. O destruction occurs in the most massive models but it is not efficient enough to account for the $\\gtrsim 0.3$~dex O depletions observed in some Type~I PNe. It is not possible to reach firm conclusions regarding the neutron source operating in massive AGB stars from Se and Kr abundances in Type I PNe; abundances for more $s$-process elements may help to distinguish between the two neutron sources. We predict that only the most massive ($M \\gtrsim 5\\Msun$) models would evolve into Type I PNe, indicating that extra-mixing processes are active in lower-mass stars (3--$4\\Msun$), if these stars are to evolve into Type~I PNe. ",
        "introduction": "After the thermally-pulsing AGB (TP-AGB) phase is terminated, low to intermediate mass stars ($\\sim 0.8$ to 8$\\Msun$) evolve to become post-AGB stars and possibly planetary nebulae (PNe), if the ejected envelopes have sufficient time to become ionized by the hot central stars before dissipating into the interstellar medium. For recent reviews of TP-AGB and post-AGB stars see \\citet{herwig05} and \\citet{vanwinckel03}, respectively. The illuminated PN is comprised of material from the deep convective envelope, hence nebular abundances should reveal information about the efficiency of mixing events and chemical processing that took place during previous evolutionary phases, in addition to the initial composition of the parent star \\citep{dopita97,karakas03a}.  Briefly, during the TP-AGB phase the He-burning shell becomes thermally unstable every $10^{4}$ years or so, depending on the core mass. The energy from the thermal pulse (TP) drives a convective zone in the He-rich intershell, that mixes the products of He-nucleosynthesis within this region. The energy provided by the TP expands the whole star, pushing the H-shell out to cooler regions where it is almost extinguished, and subsequently allowing the convective envelope to move inwards (in mass) to regions previously mixed by the flash-driven convective zone. This inward movement of the convective envelope is known as the third dredge-up (TDU), and is responsible for enriching the surface in \\iso{12}C and other products of He-burning, as well as heavy elements produced by the $s$ process. Following the TDU, the star contracts and the H-shell is re-ignited, providing most of the surface luminosity for the next interpulse period. In intermediate-mass AGB stars with initial masses $\\gtrsim 4\\Msun$, the base of the convective envelope can dip into the top of the H-burning shell, causing proton-capture nucleosynthesis to occur there \\citep[hot bottom burning, HBB; see][for a review]{lattanzio96}.  Abundances can be obtained from PN spectra for a number of elements including C, N, O, S, Cl and the noble gases He, Ne, and Ar \\citep{aller83,kingsburgh94,dopita97,stanghellini00,leisy06}; and heavy elements synthesized by the $s$ process, including Ge, Se, Kr, Xe, Ba, and Br \\citep[][hereafter SD08]{pequignot94, dinerstein01a, sterling02, sharpee07, sterling07, sterling08}. Whereas the metallicity of a star is usually taken from the abundance of iron, for ionized nebulae oxygen is generally the primary index of metallicity. However, O is not an ideal metallicity index for PNe, because it can be destroyed by CNO cycling and/or synthesized during helium burning. Consequently, Ar is sometimes used in place of O \\citep[][SD08]{leisy06}, under the assumption that it remains unchanged by AGB nucleosynthesis. Zinc is a better indicator of iron-group element abundances in nebulae than Fe itself, because Zn does not condense into dust as easily as Fe \\citep{york82,sembach95,welty99} and can be observed in PNe \\citep{dinerstein01,dinerstein04}. However, Zn is at the beginning of the $s$-process chain and might be produced to some extent in AGB stars. In order to interpret the elemental abundances in PNe it is important to determine which elements are produced, destroyed or left unaltered by AGB nucleosynthesis.  The set of elements heavier than Fe that can be detected in PNe include all even-numbered elements from $Z$ = 30 -- 36, which lie at the light-element end of the $s$-process distribution, just below the first $s$-process peak at Sr, Y, and Zr ($Z$ = 38 -- 40). In \\citet{karakas07a} we attempted to use Ge abundances derived from a small sample of PNe \\citep{sterling02} to try and constrain the amount of TDU mixing after the final few TPs. There a hint was found that efficient TDU at the tip of the AGB may be required in order to match the abundances; however, stronger conclusions were precluded by model uncertainties, the small PNe sample size, and the large uncertainties of the Ge abundance determinations. In the present study we proceed to analyze Se and Kr abundances from a much larger sample of 120 Galactic PNe from SD08, where the $s$-process enrichments were investigated for correlations with PN morphology and other nebular and stellar characteristics, as outlined in detail in \\S\\ref{obs}. We also use PNe abundances of C, N, and O along with the abundances of post-AGB stars for comparison to the stellar models.  The composition of Type I PNe provides important constraints on nucleosynthesis and mixing in intermediate-mass AGB stars. These PNe tend to exhibit bipolar morphologies and have high He/H and N/O ratios characteristic of proton-capture nucleosynthesis \\citep{stanghellini06}, suggesting they are descendants of AGB stars with initial masses $\\sim$3 -- 8 $\\Msun$ \\citep{peimbert78,kingsburgh94,stanghellini06}. Low-mass AGB stars, on the other hand, are defined as having initial masses $\\sim$1 -- 3$\\Msun$ and tend to show carbon and $s$-process element enrichments \\citep{busso01,travaglio04}. In the Milky Way, the spatial distribution and kinematics of Type I PNe indicate that they are a young population \\citep[e.g.,][]{corradi95}. This supports their identification with intermediate-mass stars, which have relatively short evolutionary lifetimes. In the Galaxy, intermediate-mass AGB stars are difficult to identify owing to a lack of reliable distances. This has resulted in a paucity of observational evidence for constraining stellar models, and especially the efficiency of the TDU in intermediate-mass AGB stars, which is still not well determined \\citep{karakas02,ventura02,herwig04a,stancliffe04b}. \\citet{garcia06} identified several Galactic intermediate-mass AGB stars within a sample of OH/IR stars (i.e., bright O-rich giants with large infrared excesses). The large enhancements of Rb found by these authors, combined with the fact that these stars are O-rich, support the prediction that HBB and an efficient TDU has occurred in these stars. At lower metallicities, constraints are provided by observations of luminous O-rich AGB stars in the Large and Small Magellanic Clouds (LMC and SMC, respectively) that are rich in Li and $s$-process elements \\citep{wood83,smith89,plez93}.  We begin in \\S\\ref{obs} with a summary of observational results for $n$-capture elements in PNe. We review the numerical method and present the stellar models in \\S\\ref{models}, and outline the results in \\S\\ref{results}. In \\S\\ref{discussion}, we discuss our findings and their implications, and in \\S\\ref{conclusions} we summarize our conclusions. ",
        "conclusions": "\\label{conclusions}  Type I PNe have high He/H and N/O ratios, indicating that they are the descendants of intermediate-mass AGB stars with initial masses between $\\sim 3$ to 8$\\Msun$. \\citet{sterling08} found that Type I PNe exhibit significantly smaller enrichments of Se and Kr ($\\lesssim 0.3$~dex) on average than other PNe. We calculated $s$-process enrichments in a set of AGB models covering a range in mass from 2.5 to 6.5$\\Msun$, and metallicity from 0.2$Z_{\\odot}$ to $Z_{\\odot}$. The 2.5$\\Msun$, $Z=0.008$ model was included to show an example of a model that would produce a clear non-Type~I PN abundance signature, with [Se, Kr/Fe] $>0.5$ and C/O $\\sim 4$. The main conclusion is that the results for the 3--6.5$\\Msun$ are a good match to the observed abundances. The only real exception out of the HBB models is the low-metallicity 5$\\Msun$, $Z = 0.004$ model that produced much larger enhancements in Se and Kr than observed. This suggests that Galactic Type I PNe do not descend from such low-metallicity objects.  We also compare calculated abundances for selected intermediate-mass (Ne, P, S, Ar) and iron-peak (Fe, Zn) elements to observations of post-AGB stars and PNe. We find that the elemental abundances of P, S, Cl, Ar, Fe, and Zn are essentially unchanged by AGB nucleosynthesis, although there are isotopic shifts caused by neutron captures in the He-shell. These results justify use of elements such as S, Cl, and Ar in PNe as tracers of Galactic chemical evolution.  It is difficult to reach firm conclusions about the neutron source operating in massive AGB stars from Se and Kr abundances in Type I PNe. Certainly it seems that at the least the \\iso{22}Ne source, with efficient TDU, is required to produce enhancements in $s$-process elements in massive Type I PN progenitors. It is less clear whether a \\iso{13}C pocket is also required in these stars. Increases in the [Kr/Fe] abundance ratio were observed in the models when a \\iso{13}C pocket was included,  but not beyond the amounts observed in Type~I PN spectra. Obtaining abundances for more $n$-capture elements, particularly Rb and elements beyond the first $s$-process peak, from a large data set of OH/IR stars and PNe will help distinguish among the possibilities.  Finally, only the models with HBB ($M \\gtrsim 5\\Msun$, depending on $Z$) show the high He/H and N/O ratios that define the Type I PN class.  Given that $3 \\Msun$ stars are more common than $6.5 \\Msun$ stars (according to initial mass function and evolutionary time considerations), it may be necessary that another mixing process other than HBB is active in stars of the lower-mass range (3--$4\\Msun$), if these stars do in fact evolve into Type~I PNe. Rapid stellar rotation in single $\\sim 3\\Msun$ stars may, for example, be able to account for the increased He and N abundances in the progenitor stars. The problems of convection and mixing, and of other modeling uncertainties (e.g., $n$-capture cross sections) requires further study in the context of $s$-process nucleosynthesis in intermediate-mass AGB stars."
    },
    "0809/0809.1726_arXiv.txt": {
        "abstract": "Several works have reported changes of the Sun's subsurface stratification inferred from $f$-mode or $p$-mode observations. Recently a non-homologous variation of the subsurface layers  with depth and time has been deduced from $f$-modes. Progress on this important transition zone between the solar interior and the external part supposes a good understanding of the interplay between the different processes which contribute to this variation. This paper is the first of a series where we aim to study these layers from the theoretical point of view. For this first paper, we use solar models obtained with the CESAM code, in its classical form, and analyze the properties of the computed theoretical $f$-modes. We examine how a pure variation in the calibrated radius influences the subsurface structure and we show also the impact of an additional change of composition on the same layers. Then we use an inversion procedure to quantify the corresponding $f$-mode variation and their capacity to infer the radius variation. We deduce an estimate of the amplitude of the 11-year cyclic photospheric radius variation. ",
        "introduction": "\\label{S-Introduction}  Helioseismology has been extremely useful to probe the internal structure of the Sun but two crucial regions need to be improved for a good understanding of the time evolution of the solar activity: (1) the solar core for a proper description of the transport of momentum implied by  rotation,  gravity waves and magnetic field along the evolution and (2) the subsurface layers which correspond to the transition region between large  and small scale dynamics evolution.  This transition zone is important to study because it couples the internal dynamics to the dynamics of  the solar atmosphere and it plays  a crucial role in the emergence of the space weather science (see for example \\citet{Rozelot06}). Figure \\ref{F-lepto} sketches a schematic view of these subsurface layers above 0.96 $\\Rs$, which are called the \\textit{leptocline} region (from the greek ``leptos''= thin, ``klino\"= tilt). This term was proposed by \\citet{Godier01}, who computed the solar oblateness and showed a curvature change in this zone which was interpreted as the presence of a double layer. Just below the surface, there is a change in radial and latitudinal rotation \\citep{Basu99,Corbard02} and the treatment of the superadiabiatic region supposes a proper description of the convection (presently described by the mixing length parameter)  and of detailed molecular and atomic opacity calculations. One needs also to add the turbulent pressure and the emergence of local magnetic field to these processes, then below, hydrogen and helium pass from neutral to partially ionized and then totally ionized in a region where the magnetic pressure cannot be ignored. The mean and varying magnetic field amplitudes have been tentatively extracted by \\citet{Nghiem06} from the analysis of the low degree acoustic modes.% The proper interaction between these different physical processes is not yet included in stellar evolution models and could lead to some variation of the solar radius along the solar cycle.  The observation of these subsurface layers is not so easy. The difficulty comes from the fact that the high-degree $p$-modes \\citep{Korzennik04} have a small lifetime and are largely perturbed by the turbulent motions and the emerging magnetic field. But important progress has appeared recently  thanks to the analyses of long series of seismic data. \\citet{Rabello08}  estimate the variations of the high-degree $p$-modes frequencies over the solar cycle and  using ring diagram analysis  \\citet{Basu07}  extract now latitudinal and temporal variations of  the sound speed or the density along the Hale solar cycle (at least its last half cycle).  \\clearpage \\begin{figure*}[htbp] \\centerline{\\includegraphics[angle=0,width=16cm]{f1.eps}} \\caption{Schematic view of the leptocline (non-scale scheme).} \\label{F-lepto} \\end{figure*} \\clearpage   In parallel, \\citet{Lefebvre05} and \\citet{Lefebvre07} reported changes with the solar cycle of the solar subsurface stratification, inferred from inversion of SOHO/MDI $f$-mode frequencies. They notice a different behavior for the layers around 0.99 $\\Rs$. Between 0.97 and 0.99 $\\Rs$, it seems that the position of the layers varies in phase with the solar cycle, whereas it appears opposite in the upper part, above 0.99 $\\Rs$, where the variability is in antiphase. So,  it seems that the most external layers of the Sun shrink during the ascending phase of the solar cycle and relax after the maximum, but these changes are not uniform with depth. At the surface, this result  is coherent with the observation of the photospheric radius variations. The observed series of ground-based measurements  \\citep{Laclare96,Reis03} and the measurements aboard stratospheric balloons \\citep{Sofia94} suggest a reduction of the solar photospheric radius at the maximum of the cycle. The balloon variation estimates are much smaller than the ground-based observations, polluted by the atmospheric turbulence. They are consistent with results from \\citet{Kuhn04} who reported no evidence of solar-cycle visible radius variations between 1996 and 2004 larger than 7 mas. But they disagree with computations of \\citet{Sofia05}. These authors predict also non-homologous subsurface stratification changes with the 11 year solar cycle but with an amplification up to a factor 1000 from the depth at 5 Mm to the surface. It leads to a variation of the solar radius up to 600 km along the cycle  as shown in Fig. 2 of \\citet{Lefebvre07}. Such a variation seems extremely large.  So the study of the detected  $f$-modes is very useful  in the present context even the interpretation of $f$-modes has appeared puzzling \\citep{Dziem01,Antia03,Dziem04} due to some apparent lack of sensitivity of these modes at the real surface. We develop here a theoretical approach  to validate  the $f$-modes inversion procedure used previously and applied here to some specific cases that we know perfectly. This paper is the first of a series where we estimate the impact of a change of radius and composition on the subsurface layers and on the $f$-mode frequencies using classical solar models including a detailed microscopic description of these layers. In the present work, we do not try to justify the origin of the variation of the radius. The next step will be to develop more complex models including magnetic field and differential rotation. Our general aim  is to properly qualify the variabilities of $f$-modes and solar radius that we will study with the SDO and PICARD missions (Kosovichev et al. 2007 and Thuillier, Dewitte \\& Schmutz 2006, respectively).% We will also investigate the capability of the $f$-modes to estimate the solar photospheric radius variability. This study will contribute also to clarify the notion of ``solar radius\" to bridge two different communities and obtain without ambiguity a unique definition of this fundamental quantity \\citep{Haberreiter08}.  After a description of our study (Section \\ref{S-Context}), we describe in Section \\ref{S-model}, the models used and examine the subsurface layer changes on different variables produced by a change in the solar radius and composition and we calculate the corresponding $f$-mode frequencies. Section \\ref{S-inversion} is devoted to the use of these frequencies to infer the change in the position of the subsurface layers. The validation of the procedure, an estimate of the present solar cyclic radius variation and the perspectives of this work will conclude the paper in Section \\ref{S-discussions}. ",
        "conclusions": "\\label{S-discussions}  Equation 8 of \\citet{Dziem04} supposes that the radial position $r$ used in Equation \\ref{Eq-eq_radius} is the lagrangian radius. $f$-modes are surface oscillation trapped waves, so each mode oscillates around an equilibrium radius. We relate in this section this radial position to that of the structure.  \\subsection{ Validity of the procedure of inversion}  The extraction of the solar subsurface structure and the photospheric radius variation along the solar cycle is not so easy to determine and several papers have been published with different conclusions \\citep{Dziem01, Antia03, Sofia05}. The most recent work of \\citet{Dziem04} shows the difficulty to obtain a general expression for the inversion, so this work allows to test such a procedure. We verify in this section if the displacement $\\Delta r$ of Equation \\ref{Eq-eq_radius} corresponds to the radial displacement at a given mass for the different cases studied.  \\clearpage \\begin{figure*}[htbp] \\centerline{\\includegraphics[width=8cm]{f7a.eps} \\includegraphics[width=8cm]{f7b.eps}} \\caption{\\textit{Left panel}: Radial displacement $\\Delta r$ of a given m(r), p(r), $c$,$H_p$, $\\Gamma_1$ and $\\kappa$ defined in the reference model ($M_1$) for models of set 2 (the curves of  set 1 models superpose the previous ones). These curves have to be compared to the left panel of Figure \\ref{F-inversion1}. \\textit{Right panel}: Idem with models of set 3.} \\label{F-radius_param} \\end{figure*}  \\begin{figure}[htbp] \\centerline{\\includegraphics[angle=-90,width=8cm]{f8.eps}} \\caption{Radial displacement $\\Delta r$ issued from the inversion (same curves than in the left panel of Figure \\ref{F-inversion2} added with the associated displacement at $H_p$ constant as plotted in the right panel of Figure \\ref{F-radius_param} (here dashed curves associated at the same color for the corresponding model). Note the good superposition of the different curves.} \\label{F-superposition} \\end{figure} \\clearpage  With this purpose, we calculate for each set (sets 1 and 2 lead to the same result) the radial displacement of a layer corresponding to a given value of the structural quantities, i.e. $\\Delta r = (r_N-r_1)_{q_N=q_1}$ where $q$ represents a model structure quantity like $m$, $c$, $H_p$ or $\\Gamma_1$, and $N$ is the model number.  Figure \\ref{F-radius_param} shows the panels representing $\\Delta r$ at constant $m$, $c$, $H_p$ and $\\Gamma_1$. For sets 1 or 2, we have the same behaviour for all the quantities and Equation \\ref{Eq-eq_radius} can be applied without any doubt. For the set 3, the radial difference $\\Delta r$, computed by the procedure described above, differs from one quantity to another. The only curve similar in amplitude and in shape with the one computed by inversion of $f$-mode frequencies in Figure  \\ref{F-inversion2} is the curve computed at fixed pressure height $H_p$. Effectively Figure \\ref{F-superposition} shows an excellent superposition of the corresponding curve and the radial displacement $\\Delta r$ issued from the inversion. This agreement, in shape and amplitude, demonstrates that the validity of equation 1 in this case requires to associate the radial displacement to this quantity. It is in fact not surprising because the $f$-mode frequencies are naturally sensitive to the pressure scale height.   We note also that the inversion is not able to reproduce precisely the behavior very near the surface due to the lack of sensitivity of  the $f$-modes. But we can deduce from this study an estimate of the uncertainty on the determination of the photospheric radius variation: we estimate it of the order of 15\\%.  \\subsection{A new prediction on the solar radius variation along the solar cycle}  We show in this paper that a pure variation in the solar radius produces changes in the subsurface layers that are in the same zone than that studied with the observed $f$-modes by \\citet{Lefebvre05}. Such a change is characterized by changes in the subsurface stratification and more exactly variations in the computed position of the layers. We show that the changes detected by $f$-modes are physically related to the variation of the pressure. In our modelling we are able to produce non-uniform changes by slightly modifying the chemical composition below the solar surface. This  is not  excluded if one introduces dynamical processes in keeping the same constraints on luminosity and radius. The two studies show different stratifications of the outlayers but with similar effects on the $f$-mode frequencies. Considering larger effects on the $f$-mode frequencies than in the solar case, we propose to deduce from  this study an estimate of the variation of the solar radius along the cycle. From $\\Delta\\nu/\\nu$ of about $3\\times10^{-4}$ for a $\\Delta r$ of about 140 km, we lead to a solar radius variation along the solar cycle of about $7\\pm 1$ km for the observed variations in frequency of about $1.5\\times10^{-5}$. This result is consistent with the radius extrapolated by \\citet{Lefebvre05} (see respectively Figures 1 and 3 of this article). It is interesting to note that this radius variation estimate is also compatible (in order of magnitude) with the observed low degree acoustic mode variation of about 0.4 $\\mu Hz$ along the solar cycle of the low degree $p$-modes \\citep{Fossat87,Chaplin01}.  \\subsection{Perspectives}  This study represents the first step on the way to explain the variations of the subsurface stratification along the solar cycle and reinforces the interest of the $f$-modes. The next step will be the development of models including magnetic field and other dynamical processes, which will permit a more realistic study.  Nevertheless it will suppose also an improved expression for the inversion of $f$-modes which is not yet obtained. We believe that the introduction of a magnetic field will influence the stratification of the subsurface layers. The study by \\citet{Nghiem06} pointed out non-uniform radial changes, depending on the importance of the magnetic pressure. Moreover, it will be also interesting to take into account the differential rotation and consequently some asphericity to confront them to the subsurface latitudinal stratification over the 11-year cycle.  We would like to emphasize that a better knowledge of these subsurface layers will contribute to our understanding (i) of the dynamics of the 11-year solar cycle and (ii) of the Sun-Earth relationship for space climate. This supposes in parallel  the simultaneous measurement of the radius and frequencies change. This study is in the framework of coming future space missions: SDO (see \\url{http://sdo.gsfc.nasa.gov/}) and PICARD (see \\url{http://smsc.cnes.fr/PICARD/} and \\citet{Thuillier06}).  The DynaMICCS/HIRISE perspective \\citep{Turck06,Turck08}, proposed in the framework of the ESA Cosmic Vision 2015-2025, could bring even more data on all sources of internal variability by putting together instruments which will follow seismically all the layers down to the core together with instruments measuring the variability of the above atmosphere."
    },
    "0809/0809.4335_arXiv.txt": {
        "abstract": "We present a novel cosmological model in which scalar field matter in a biaxial Bianchi IX geometry leads to a non-singular `pancaking' solution: the hypersurface volume goes to zero instantaneously at the {`Big Bang'}, but all physical quantities, such as curvature invariants and the matter energy density remain finite, and continue smoothly through the Big Bang. We demonstrate that there exist geodesics extending through the Big Bang, but that there are also incomplete geodesics that spiral infinitely around a topologically closed spatial dimension at the Big Bang, rendering it, at worst, a quasi-regular singularity. The model is thus reminiscent of the Taub-NUT vacuum solution in that it has  biaxial Bianchi IX geometry and its evolution exhibits a dimensionality reduction at a quasi-regular singularity; the two models are, however, rather different, as we will show in a future work.  Here we concentrate on the cosmological implications of our model and show how the scalar field drives both isotropisation and inflation, thus raising the question of whether structure on the largest scales was laid down at a time when the universe was still oblate (as also suggested by {\\cite{PitrouUzan2007CosmoPerturbationsAnisotropic, Pitrou2008PredictionsAnisotropicEra, Contaldi2007InflationaryPertubationsInBianchi}}).  We also discuss the stability of our model to small perturbations around biaxiality and draw an analogy with cosmological perturbations. We conclude by presenting a separate, bouncing solution, which generalises the known bouncing solution in closed FRW universes. ",
        "introduction": "}  At the core of most of theoretical cosmology lie the assumptions of homogeneity and isotropy. These were originally motivated by the cosmological principle and the mathematical tractability of the resulting FRW models. However, the observed universe is obviously neither homogeneous nor isotropic, so these symmetries can only ever be approximate. Thus the question arises as to which assumptions we can relax whilst maintaining analytical tractability. We choose here to consider more general, anisotropic cosmologies, whilst keeping the assumption of (large-scale) spatial homogeneity. The resulting class of cosmologies is collectively known as the Bianchi models.  Scalar fields are ubiquitous in theories of high energy physics. The Standard Model of particle physics postulates the scalar Higgs particle, and Superstring Theory and string-inspired models motivate a plethora of moduli fields from compactifications, dilatons, radions etc. Scalar fields are also of interest in cosmology, as they can drive periods of accelerated expansion of the universe in a relatively straightforward and plausible manner. Cosmological scalar fields could therefore account for the posited period of inflation and the apparent present period of $\\Lambda$-domination.  It is therefore natural to consider the effects of a scalar field dominating the dynamics of a Bianchi model. However, the phenomenological success of FRW models suggests that we should only consider the Bianchi classes that allow FRW universes as special cases. This narrows the phenomenologically interesting Bianchi types down to $\\mathrm{I}$, $\\mathrm{V}$, $\\mathrm{VII_{h}}$ and $\\mathrm{IX}$. From previous work \\cite{LasenbyDoran2005ClosedUniversesdSandInflation, LasenbyDoran2004ConformalModelsICsCMB} we are particularly interested in closed models, which are still consistent with observations \\cite{Spergel2007WMAP3Cos}.  The only closed Bianchi type that also allows for FRW subclasses is Bianchi IX, so we shall consider the dynamics of a scalar field in a Bianchi IX model \\cite{Toporensky1999BianchiIXMassiveScalarField, Fay2005BianchiIXMassiveScalarFieldChaos, ColeyGoliath2000}. Bianchi IX models are known to have very complicated dynamics, exhibiting oscillatory singularities and chaos (mixmaster) \\cite{Misner1969Mixmaster,BelinskyKhalatnikovLifshitz1972,BelinskyKhalatnikovLifshitzKL1969, Ellis1969ClassOfHomogeneousCosmologicalModels, Novikov1973ProcessesNearSingularities, Ellis2006BianchiModelsThenAndNow, UgglaHeinzle2009mixmaster, UgglaHeinzle2009newproof}. In this light, it is even more intriguing that the biaxial model we will consider is so well-behaved.  {It is interesting that our solution {provides} another example of a regular solution with axial symmetry; for others, see Pitrou et al \\cite{PitrouUzan2007CosmoPerturbationsAnisotropic, Pitrou2008PredictionsAnisotropicEra} and Senovilla and collaborators \\cite{Senovilla1990NewClassOfInhomogeneous,Senovilla1992GeneralClass,Senovilla1992Singularity-free,Fernandez2002AWideFamily}. In the different context of inhomogeneous {but axially symmetric} cosmologies, \\cite{Senovilla1990NewClassOfInhomogeneous,Senovilla1992GeneralClass,Senovilla1992Singularity-free,Fernandez2002AWideFamily} also consider the important concept of geodesic completeness of a spacetime.  Completeness is a criterion for the physical significance of a solution that is more stringent than mere geometric regularity, which we will also need to address later.}  {Often inflation is invoked to justify the assumptions of (acausal) homogeneity and isotropy. However, this then raises the question of how sensitive inflation itself is to the rather unnatural initial conditions of a homogeneous initial state. For these considerations, the interested reader may refer to, for instance, \\cite{GoldwirthPiran1992ICsforinflation}. Here, however, we consider a Bianchi geometry as our starting assumption,  without invoking inflation in order to justify it. For recent related literature concerning the case where a vector field is present to drive a phase of anisotropic inflation, see for example \\cite{Soda2008AnisotropicInflationVector, Contaldi2008Instability}.}  This paper is organised as follows. We begin with a brief introduction to Bianchi models in Section \\ref{cla}. In Section \\ref{tri} we construct a generic Bianchi IX model dominated by scalar field matter and then specialise to the biaxial case in Section \\ref{axi}. ({An alternative form of the Einstein field equations for this system, using the $3+1$ {covariant} approach, is presented in the appendix.}) We demonstrate scaling behaviour of solutions, and present a particular series solution, in which one radius goes to zero at the Big Bang (pancaking). We also consider geodesics through the Big Bang. We then present a realistic cosmology based on our model in Section \\ref{cos} and, in Section \\ref{pert}, the stability of such a model about biaxiality. We finally present a separate, bouncing solution in Section \\ref{bounce}, before we conclude in Section \\ref{con}. ",
        "conclusions": "}  We have presented a novel scenario in which the effect of a scalar field and biaxial Bianchi IX geometry is to render the Big Bang much better behaved.  Physical quantities remain finite at the Big Bang, and there is no curvature singularity at the time of pancaking. This is in contrast with singularities in flat and open FRW-models, and essential singularities in the full triaxial Bianchi IX case.  In addition, the investigation of the behaviour of geodesics suggests that some physical observers can cross over and move smoothly into a pre-Big Bang phase. Some observers, however, may not be able to cross and this would appear to make the model incomplete, with the pancake being a quasiregular singularity (the type that also occurs in Taub-NUT). However, in the light of periodic identification of angular coordinates on the 3-sphere and the fact that the winding direction is the collapsing dimension, this might not be such a problem.  The model also exhibits stability under sufficiently small perturbations around biaxiality.  Though the model isotropises at late times, we show that structure on the largest scales could in fact stem from a time at which the universe was significantly anisotropic.  Such behaviour is very desirable for a satisfactory cosmological model, as a closed universe can be matched to the observed asymptotic de Sitter phase of $\\Lambda$- domination, whilst the scalar field can drive isotropisation, so as to account for the observed degree of isotropy of the universe, and inflation, thought to be needed in order to explain the observed perturbation spectrum. We predict a spectral index and a tensor-to-scalar ratio compatible with current constraints, as well as a dip at low multipoles of the CMB power spectrum, which is consistent with an exponential cutoff that has also been argued for on phenomenological grounds.  We explain the apparent qualitative similarity between the two very different kinds of perturbations, namely curvature perturbations and perturbations around biaxiality, from their evolution equations.  A separate, bouncing solution is also presented, and will be considered in more detail elsewhere. A more thorough analysis of many subtleties of the model, in particular the issue of geodesic completeness and a detailed comparison with Taub-NUT is in preparation."
    },
    "0809/0809.1510_arXiv.txt": {
        "abstract": "We present numerical simulations and laboratory experiments on an eight-octant phase-mask (EOPM) coronagraph. The numerical simulations suggest that an achievable contrast for the EOPM coronagraph can be greatly improved as compared to that of a four-quadrant phase-mask (FQPM) coronagraph for a partially resolved star. On-sky transmission maps reveal that the EOPM coronagraph has relatively high optical throughput, a small inner working angle and large discovery space. We have manufactured an eight-segment phase mask utilizing a nematic liquid-crystal device, which can be easily switched between the FQPM and the EOPM modes. The laboratory experiments demonstrate that the EOPM coronagraph has a better tolerance of the tip-tilt error than the FQPM one. We also discuss feasibility of a fully achromatic and high-throughput EOPM coronagraph utilizing a polarization interferometric technique. ",
        "introduction": "% An attempt to discover and characterize extrasolar Earth-like planets is one of the most challenging issues in modern astronomy. For direct detection of the extrasolar planets, many concepts of high-contrast imagers have been proposed, especially in the context of the Terrestrial Planet Finder Coronagraph (TPF-C) mission. It is necessary for the high-contrast imagers to realize extremely high dynamic range over a broad spectral bandwidth, together with high optical throughput, a small inner working angle (IWA), and large discovery space. For a practical aspect, in addition, technically simple instrumentation would be preferable because of a severe environment of space observations.   A classical Lyot coronagraph is a precursory work for the high-contrast imager aiming at observing a solar coronal prominence by using an opaque mask on a focal plane \\citep{Lyot39}. Recently, more advanced focal-plane masks, such as a four-quadrant phase mask (FQPM), a band-limited mask, and an optical vortex mask, have been proposed for planet detection \\citep{Rouan00,Kuchner02,Kuchner05,Foo05}. These techniques have the ability to strongly suppress both stellar diffraction core and halo, and tend to have relatively small IWAs. Pupil-apodization techniques, such as an apodized square aperture and binary shaped masks, have been proposed to suppress only the stellar diffraction halo \\citep{Nisenson01,Kasdin03,Enya07}. An apodized pupil Lyot coronagraph (APLC), a combination of the pupil-apodization and the classical Lyot coronagraph, has also been proposed for realizing the high-contrast imaging with arbitrary aperture shapes of telescopes \\citep{Soummer05}. The pupil-apodizations are less sensitive to tip-tilt errors and/or finite stellar angular sizes than the focal-plane mask coronagraphs. However, these techniques tend to have relatively large IWAs and low optical throughputs. A phase-induced amplitude apodization (PIAA) has been proposed to achieve a small IWA without a loss of the optical throughput \\citep{Guyon03}.   A comprehensive evaluation of theoretical capabilities for extrasolar planet detection has been conducted for several promising concepts \\citep{Guyon06}. For direct detection of the extrasolar Earth-like planets, the high-contrast imagers must be tolerant of the tip-tilt error and the stellar angular size problem. For the FQPM coronagraph, however, residual stellar noise increases rapidly with $\\delta^2$, where $\\delta$ is the tip-tilt error. This second-order behavior for the tip-tilt error would be insufficient for direct detection of the Earth-like planets around nearby stars. For this reason, the FQPM coronagraph has been excluded from the candidates for the TPF-C mission. Note that it is possible to realize the higher-order insensitivities to the tip-tilt error by using the other focal-plane mask (the band-limited mask and the optical vortex mask) coronagraphs. Nevertheless, the FQPM coronagraph is still attractive because of its ability to realize a perfect stellar suppression, a high optical throughput, a large discovery space and a small IWA together with technically simple instrumentation. \\cite{Rouan07} point out that it is possible to realize more efficient stellar suppressions by multiplying the number of sectors of the phase mask. We have also been paying attention to an eight-octant phase mask (EOPM) because of its acceptable optical throughput, IWA, and discovery space. We expect that it is easy to manufacture and achromatize the high-performance EOPM, because it has neither continuous phase retardation like the optical vortex masks nor complex and fine structure like the band-limited masks.   We have manufactured a chromatic version of the phase mask by using a nematic liquid crystal device, which can be easily switched between the FQPM and the EOPM modes. In this paper, we present numerical simulations and preliminary laboratory experiments on the EOPM coronagraph comparing it with the FQPM one. In \\S. \\ref{sec:EOPM}, we show the results of the coronagraphic numerical simulations for a partially resolved star. On-sky transmittances of both the coronagraphs are also shown for evaluating their optical throughputs, IWAs, and discovery spaces. In \\S. \\ref{sec:Lab}, we describe the manufactured phase mask, and show results of the laboratory experiments on both the coronagraphs for various tip-tilt errors. In \\S. \\ref{sec:Achromatization}, we discuss feasibility of a fully achromatic and high-throughput EOPM coronagraph utilizing a polarization interferometric technique. Finally, our conclusions are summarized in \\S. \\ref{sec:Conc}. ",
        "conclusions": "% \\label{sec:Conc} In this paper, we report the numerical simulations and the laboratory experiments on the EOPM coronagraph. The numerical simulations show that the EOPM coronagraph greatly improves the achievable contrast as compared to the FQPM one for partially resolved stars. In addition, on-sky transmission maps reveal that the EOPM coronagraph has a relatively small IWA and a large discovery space. We manufactured the eight-segment phase mask by using the NLC device, which can be easily switched between the FQPM and EOPM modes. Our laboratory experiments confirm that the EOPM coronagraph has the fourth-order response to the tip-tilt error, which is better than that of the FQPM one (expected to be second-order). This higher-order behavior of the EOPM coronagraph could enable us to detect extrasolar Earth-like planets around nearby stars with relatively large apparent sizes.   We also discuss feasibility of a fully achromatic and high-throughput EOPM coronagraph. We expect that the polarization interferometric phase mask by using the FLC device (so-called EOPoM) and the two-channel optical configuration to improve the achromaticity and the optical throughput. The two-channel configuration is also useful for obtaining polarization differential images to suppress residual stellar speckle noise and to carry out polarimetric measurements of extrasolar planets \\citep{Baba03}. For our future work, we are planning to design and manufacture the achromatic EOPoM by using the FLC (or possibly by the TNLC) device, and conduct on-sky high-contrast observations of circumstellar disks and faint companions around nearby stars."
    },
    "0809/0809.2934_arXiv.txt": {
        "abstract": "{Understanding grain-surface processes is crucial to interpreting the chemistry in many regions of the interstellar medium. However, accurate surface chemistry models are computationally expensive and are difficult to integrate with gas-phase simulations.} {A new modified-rate method for solving grain-surface chemical systems is presented. The purpose of the method is to trade a small amount of accuracy, and certain excessive detail, for the ability to accurately model highly complex systems that can otherwise only be treated using the sometimes inadequate rate-equation approach.} {In contrast to previous rate-modification techniques, the functional form of the surface production rates was modified, and not simply the rate coefficient. This form is appropriate to the extreme ``small-grain'' limit, and can be verified using an analytical master-equation approach. Various further modifications were made to this basic form, to account for competition between processes, to improve estimates of surface occupation probabilities, and to allow a switch-over to the normal rate equations where these are applicable.} {The new method was tested against a number of systems solved previously using master-equation and Monte Carlo techniques. It is found that even the simplest method is quite accurate, and a great improvement over rate equations. Further modifications allow the master-equation results to be reproduced exactly for the methanol-producing system, within computational accuracy. Small discrepancies arise when non-zero activation energies are assumed for the methanol system, which result from complex reaction-competition processes that cannot be resolved easily without using exact methods. Inaccuracies in computed abundances are never greater than a few tens of percent, and typically of the order of one percent, in the most complex systems tested.} {The new modified-rate approach presented here is robust to a range of grain-surface parameters, and accurately reproduces the results of exact methods. Furthermore, it may be derived from basic approximations, making the behaviour of the system understandable in terms of physical processes rather than time-dependent probabilities or other more abstract quantities. The method is simple enough to be easily incorporated into a full gas-grain chemical code. Implementation of the method in simple networks, including hydrogen-only systems, is trivial, whilst the results are highly accurate.} ",
        "introduction": "Chemical models consisting of hundreds of species and thousands of reactions are frequently used to interpret the morphologies and evolutionary histories of interstellar clouds and star-formation regions. Towards such ends, simulations dealing solely with gas-phase chemistry have achieved much success; however, the explicit consideration of grain-surface processes in chemical models is becoming increasingly important to our understanding of interstellar chemistry. From the formation of the most basic and abundant interstellar molecule, H$_2$ (e.g. Gould \\& Salpeter 1963; Hollenbach \\& Salpeter 1971; Duley \\& Williams 1984), to some of the most chemically complex organic species observed in star-forming regions (e.g. Charnley et al. 1992, 1995; Cazaux et al. 2003; Horn et al. 2004; Garrod \\& Herbst 2006; Garrod et al. 2008), the action of grain-surface chemical processes appears to be crucial.  Unfortunately, the accurate coupling of gas-phase and grain-surface processes in chemical models is difficult, due to the different nature of the chemistry in each phase. Gas-phase chemistry may be accurately modelled by describing chemical abundances as averaged concentrations, or number densities. By assuming an arbitrarily large ``cell'' of gas, the absolute number of particles of any species present in the system can be assumed to be statistically large, making stochastic fluctuations unimportant. Such a ``deterministic'' treatment allows the chemistry to be described by a set of ordinary, first-order differential equations (rate equations), which are solved numerically, allowing the time-dependent behaviour to be traced.  The simplicity of this rate-equation approach naturally encourages its use in solving grain-surface problems, especially in coupled gas-phase/grain-surface (i.e. gas-grain) models; however, its accuracy in these applications is in some cases questionable. Grain-surface chemistry takes place on finite surfaces, where the populations of certain chemical species can become very small, of the order of 1. If surface reactions occur quickly in such a regime, then stochastic effects may come into play; this renders the rate-equation treatment inaccurate, yielding production rates that are faster than physically possible.  These failures do not arise in all regimes; \\cite{katz} showed that hydrogen diffusion at low temperature takes place via thermal hopping, rather than fast quantum tunnelling. For grains of canonical size ($0.1\\mu$m), at temperatures of $\\sim$10 K, the reaction rates are typically no faster than the accretion or evaporation rates of the reactants, putting the system into the deterministic limit. It is in the consideration of grains that are smaller or hotter than this that stochastic effects must be considered.  Monte Carlo methods have been developed (e.g. Tielens \\& Hagen 1982; Charnley et al. 1997; Charnley 1998, 2001) to take account of stochastic behaviour. The most recent techniques (Cuppen \\& Herbst 2005) trace the behaviour of individual atoms and molecules on a grain surface; however, even the simplest of such schemes is difficult to integrate with a rate equation-based gas-phase code (Chang et al. 2007). Alternatively, the master-equation method describes the system using the time derivatives of the probabilities of specific population states, which are easily solved in tandem with gas-phase rate equations (Biham et al. 2001). Cut-offs are imposed, representing the maximum populations of surface species. Unfortunately, with large networks the number of possible population-state combinations becomes unmanageable, even with low cut-offs.  Stantcheva et al. (2002) developed hybrid schemes that mix deterministic and stochastic methods, treating abundant, ``deterministic'' species such as CO using rate-equations, but treating ``stochastic'' species using a master-equation method that employs low cut-offs. These schemes are successful for the small networks upon which they have been tested, but require prior knowledge from the exact methods in order to distinguish ``deterministic'' from ``stochastic'' species.  Green et al. (2001) made approximations within the master-equation method to give analytic expressions for production rates in some simple chemical systems; but the approximations are valid only in the stochastic regime, ruling out an extension to complex networks involving deterministic species such as CO.  The ``method of moments'' recently developed by Barzel \\& Biham (2007a,b) involves the calculation of second moments of population, beginning from a master-equation standpoint. This allows the calculation of accurate production rates. However, in the deterministic regime, the method gives accurate production rates, but can produce inaccurate populations (Barzel \\& Biham, 2007b; B. Barzel, private comm.). It is currently unclear how this would affect the behaviour of large chemical networks. Also, the scheme is untested against networks that include activation energies, such as are necessary in the methanol-producing system.  Attempts have been made previously (Caselli et al. 1998; Shalabiea et al. 1998; Stantcheva et al. 2001) to rectify the inaccuracies of the rate-equation method by modifying the rate coefficients of grain-surface production rates according to a set of simple rules; \\cite{caselli02} also devised a more complex modification scheme for reactions involving activation energy barriers. These modifications were largely empirical, and were not universally successful; convergence with exact techniques was achieved only for low temperatures, where atomic hydrogen is the only mobile reactant. The comparison of techniques conducted by Rae et al. (2003), using a standardised two-reactant system, found that the method was in general no more accurate than the standard rate-equation approach.  Here is presented a new method of modification of the production rates, different from previous techniques. This method modifies not the rate {\\em coefficient}, but the functional form of the production rate itself, adopting production rates derived for the extreme ``small-grain'' regime. The method switches smoothly to a rate-equation treatment where appropriate. The scheme retains its functional dependence on averaged population values, $\\langle N \\rangle$, rather than analysing individual states, $N$, allowing it to be easily implemented in standard gas-grain codes. The purpose of the new method is not primarily to produce a {\\em perfect} match to the results of exact methods, but to improve accuracy in large chemical networks that cannot be modelled with exact techniques.  Section 2 outlines the rates for surface processes. Basic modified rates are formulated in Section 3. Sections 4 and 5 detail further modifications to the basic rates. Section 6 applies the method to reaction systems with activation energies. A discussion follows in Section 7, and conclusions in Section 8. ",
        "conclusions": "The production-rate modification schemes presented here substitute the deterministic production rates with so-called ``small-grain'' rates, applicable when populations are very small and reaction rates are fast. These basic rates are easily formulated, and have been more rigorously derived elsewhere using a master-equation approach (e.g. Lipshtat et al. 2004). A simple switching/rate-limiting system is employed for each reaction, allowing the deterministic rates to be used when appropriate. Even this simple scheme provides a highly accurate reproduction of the hydrogen system examined by Barzel \\& Biham (2007a,b).  The consideration of competition between surface processes allows the water- and methanol-producing systems of Barzel \\& Biham also to be very accurately reproduced. A further refinement is made to better estimate the probability of the presence of a reactant on the grains, improving accuracy in regimes in which average populations approach a value of 1. Excellent agreement with the results of the Monte Carlo models of Stantcheva et al. (2002) is achieved with this scheme, labelled method C.  The modification method is useful in that it trades off a small reduction in accuracy for a scheme that does not require the explicit treatment of population states. This allows the methods to be easily implemented in gas-grain chemical models. The emphasis on competition processes also renders the systems understandable in terms of physical processes, rather than probability distributions or other more abstract quantities.  The greatest practical obstacle to achieving a perfect reproduction of the results of exact methods is the existence of activation-energy barriers for certain key reactions. However, the method C results never fall more than a few tens of percent away from the true populations of species produced by these routes. Critically, the inaccuracies necessarily fall in a parameter space close to the stochastic--deterministic threshold, and are therefore limited by the switch-over to the standard rate-equations. Even these inaccuracies may be overcome by judicious consideration of competition processes; however, a comprehensive approach of this sort would be difficult, and computationally expensive.  Perfect accuracy over all parameter space can of course only be achieved with exact techniques such as the Monte Carlo or master-equation methods. The fundamental cause of inaccuracies in the modification schemes is the break-down of the two-particle approximation that is applied, explicitly or implicitly, throughout this paper. The modification schemes, culminating in method C, act to diminish the effects of this break-down, thereby bridging more closely the stochastic and deterministic regimes.  It should be stated again that the purpose of the methods developed here is not to achieve perfect accuracy, but rather to achieve acceptable accuracy in the simulation of systems that cannot (currently) be treated using exact methods. In this aim, the modification schemes presented here appear promising. In many regimes, some of the modifications may not even be neccessary; competition between evaporation and reaction should only be important if large $E_{b}$:$E_{D}$ ratios are assumed, or if extremely small grains or extremely high temperatures are considered. The most important modification to the basic (continuous) method is that applied in method C. Even the use of just the basic (continuous) method would be an improvement over rate equations, in any regime. Table 7 details the equations necessary to implement each method.  Certain elements of the modification method are not rigorously tested by the systems used here, and should be further tested in future. This includes the switch-over to deterministic rates when populations become large; deterministic rates typically become valid before populations of 1 or more are achieved. Also, the ``full competition'' scheme first implemented in method B shows only a limited effect here. The accuracy of most simple reaction networks should not suffer greatly by its omission.  The methods of production-rate modification presented are simple enough that they may be implemented in significantly larger systems than those explored here, sufficient for a full treatment of gas-phase and grain-surface chemistry. In such networks it may also be necessary to treat photodissociation of surface species, being both a source of ``formation'' in equation (18), and a source of ``destruction'' for the evaluation of competition efficiencies in equations (17), (20) and (21). Photodissociation is, like evaporation, a deterministic process, rendering its implementation in the relevant equations trivial."
    },
    "0809/0809.2796_arXiv.txt": {
        "abstract": "Weak gravitational lensing provides a potentially powerful method for the detection of clusters.  In addition to cluster candidates, a large number of objects with possibly no optical or X-ray component have been detected in shear-selected samples.  Determining the nature of these so-called ``dark\" lenses is an important step towards understanding the reliability of shear-selection techniques.  We develop an analytic model to investigate the claim of \\citet{WK2002} that unvirialised protoclusters account for a significant number of dark lenses.  In our model, a protocluster consists of a small virialised region surrounded by in-falling matter.  We use a simple model for the density profile that assumes the Navarro-Frenk-White form inside of the virial radius and a power law $\\rho \\sim r^{-\\alpha}$ outside.  We find that, in order for a protocluster to simultaneously escape X-ray detection and create a detectable weak lensing signal, it must have a small virial mass ($\\sim 10^{13}~\\Msun$) and large total mass ($\\sim10^{15}~\\Msun$), with a relatively flat density profile outside of the virial radius ($\\alpha \\sim 0 - 1$).  Such objects would be characterized by rising tangential shear profiles well beyond the virial radius.  We use a semi-analytic approach based on the excursion set formalism to estimate the abundance of lensing protoclusters with a low probability of X-ray detection.  We find that they are extremely rare, accounting for less than $0.4$ per cent of the total lenses in a survey with background galaxy density $n = 30 ~ \\mathrm{arcmin^{-2}}$ and intrinsic ellipticity dispersion $\\sigma_{\\epsilon} = 0.3$.  Their  abundance decreases significantly if flat density profiles outside of the virial radius are not common.  We conclude that lensing protoclusters with undetectable X-Ray luminosities are too rare to account for a significant number of dark lenses. ",
        "introduction": "The abundance of collapsed dark matter haloes in the Universe yields a wealth of information on the dynamics of structure formation.  The number density of galaxy clusters and its time-evolution can be used to probe the normalization of the linear power spectrum, $\\sigma_8$, and the density parameters $\\Omega_m$ and $\\Omega_{\\Lambda}$ \\citep[e.g.,][]{Lilje1992,WEF1993,CenOstriker1994,ECF1996,Henry1997,BahcallFan1998,VianaLiddle1999,Holder2001,ReiprichBohringer2002,Dahle2006}.  In addition, since the linear growth of overdensities and comoving volume-element are sensitive to the dark energy equation-of-state parameter $w = P/\\rho$, the cluster abundance can be used to constrain dark energy \\citep[e.g.,][]{Haiman2001,MajumdarMohr2004,Mantz2008}.  In order to effectively use the cluster number counts as a cosmological tool, one must be able to create a complete catalog out to high redshifts.  To date, most studies aiming to complete this task have relied on X-ray based selection and constraints on cluster masses \\citep[see however][]{Dahle2006}.  Mass estimates derived from their X-ray emission require the additional assumption of hydrostatic equilibrium for the gas, which is not robust.  Galaxy clusters are the most recently assembled structures in the Universe.  One should therefore expect to find a large number of them in a dynamically unrelaxed state.  In recent years, much effort has been devoted to developing techniques that utilize weak-lensing in cluster surveys.  Gravitational lensing offers a method to measure masses that is independent of their dynamical state.  In addition to improving the accuracy of mass measurements, the weak distortion in the shapes of background galaxies may also provide a powerful method of cluster detection.  In this spirit, \\citet{Schneider1996} introduces the aperture mass measure as a way to systematically search for mass concentrations using weak lensing.  In this approach, a weighted sum of image ellipticities is used as a proxy for the projected mass contained within an aperture.  Since gravitational lensing probes mass concentrations in a way that is independent of their baryonic content, shear-selected samples can potentially provide a new and exciting view of large-scale structure.  In fact, there have been several shear-selected cluster candidates that appear to lack the characteristic galaxy over-density and X-ray luminosity. \\citep[e.g.,][]{Erben2000,UmetsuFutamase2000,Miralles2002,Dahle2003,Schirmer2007,Maturi2007}.  The first detection of these so-called dark lenses was reported by \\citet{Erben2000}.  Their initial analysis indicates significant tangential alignment roughly 7 arcminutes South of the cluster Abell 1942.  Assuming the presence of a mass concentration with roughly the same redshift as Abell 1942 ($z \\sim 0.2$), they obtain a lower-bound mass estimate of $M \\sim 10^{14} h^{-1}~\\Msun$ inside of a sphere of radius $r = 0.5h^{-1}~\\Mpc$.  However, a follow-up analysis by \\citet{vonderLinden2006} using HST observations detects the tangential alignment with a lower significance, casting doubt on the hypothesis that the object is a true mass concentration.  The HST data contains a larger number of distant galaxies and should therefore increase the significance of the detection in the case of a dark mass concentration.  Among the most recent detections, \\citet{Schirmer2007} find 86 dark lenses out of 158 possible mass concentrations in a wide field with average galaxy density $n \\sim 12~\\mathrm{arcmin}^{-2}$.  These objects, which show no detectable optical component, were identified using the aperture mass measure and a variant of it.  Similarly, using the linear filter of \\citet{Maturi2005} to minimize spurious signals created by large-scale structure (LSS), \\citet{Maturi2007} detect 7 dark lenses with no optical or X-ray component out of 14 identified lenses.  The exact nature of these detections remains an open question.  There are several possible explanations for the appearance of dark lenses in shear-selected surveys.  One possibility is that these objects truly correspond to dark matter concentrations that are abnormally deficient in baryons.  A significant abundance of such objects is not expected and would require a rethinking of current structure formation scenarios \\citep{vonderLinden2006,Maturi2007}.  A second possibility is that dark lenses are spurious signals in a lensing map resulting from the alignment of intrinsic galaxy ellipticities or LSS projected along the line-of-sight.  Both scenarios are likely to be significant problems for weak-lensing surveys and have been investigated by several authors \\citep{ReblinskyBartelmann1999,White2002,Hamana2004,Hennawi2005,Pace2007,Fan2007}.  Another interesting possibility was proposed by \\citet{WK2002} (WK2002).  They suggest that dark lenses may be cluster progenitors that are not fully virialised.  They argue that these objects should have a low galaxy over-density and X-ray luminosity compared to fully virialised clusters of the same mass.  If some of these objects are sufficiently massive and over-dense to create a detectable weak-lensing signal, then they might have the same observable signatures as dark lenses.  For clarity, we will refer to these objects as lensing protoclusters (LP).  Using an analytic approach based on the Press-Schechter formalism, WK2002 estimate that $10-20$ per cent of weak lenses should be dark LPs.  Of course, it is entirely possible that the dark lens phenomenon is due to a combination of the above scenarios.  Determining the extent to which each contributes to dark lens abundances, if at all, is an important step towards understanding the reliability of shear- selection techniques.  In this paper, we explore the scenario proposed by WK2002 in further detail.  We create a simple analytic model to investigate the likelihood that LPs can have the observational properties of dark lenses.  The remainder of this paper is organized in the following manner.  In section \\ref{APMASSMEASURE}, we briefly review the aperture mass technique.  In section \\ref{AMLPs}, we present our analytic model of LPs.  In section \\ref{LPsDL}, we calculate the properties that LPs must have to be detected as dark lenses.  We calculate the dark LP mass function in section \\ref{ABUNDANCES} and explore what the abundance of dark LPs implies for weak-lensing surveys.  Finally, we discuss our results and other potential dark lenses in section \\ref{DISCUSSION}.  In what follows, we assume a $\\Lambda$CDM cosmology with parameters $\\Omega_m = 0.26$, $\\Omega_{\\Lambda} = 0.74$, $\\Omega_b = 0.044$, $H = 100h~\\mathrm{km s^{-1}}~\\Mpc^{-1}$ (with $h = 0.72$), $n = 0.96$, and $\\sigma_8 = 0.8$, consistent with five-year WMAP results \\citep{Dunkley2008}.  Unless otherwise noted, all distances and volumes are reported in physical units. ",
        "conclusions": "\\label{DISCUSSION}  We have developed an analytic model to determine whether LPs can account for a significant fraction of dark lenses.  In our model, a protocluster consists of a small virialised central region surrounded by infalling matter.  A dark LP corresponds to a cluster-scale mass concentration with group-sized ($\\sim 10^{13}~\\Msun$) virial mass.  The small virial mass results in a low probability of X-ray detection, while the large total mass yields a high aperture mass $S/N$.  As initially suggested by WK2002, these objects can potentially share the same observational properties as dark lenses.  For the LP mass distribution we used an idealized model consisting of an NFW profile inside of the virial radius and power-law, $\\rho \\sim r^{-\\alpha}$, extending from the virial radius to the truncation radius.  In this case, the total $S/N$ is the sum of contributions from inside and outside of the virial radius.  In the case of small virial mass, the $S/N$ is dominated by the infall region.  Dark objects in shear-selected samples would likely be followed up with deep X-Ray searches.  In order to quantify the likelihood for a LP to display detectable X-ray emission, we used the analysis of \\citet{Nord2008}.  We found that LPs with a low probability of being detected via their X-Ray luminosities (or equivalently low virial masses) must have large total masses ($M \\sim 10^{15} \\Msun$) and $\\alpha$-values $\\la 1$ to meet the aperture mass detection threshold.  Such infall regions may exist given the recent findings of \\citet{Tavio2008}, who showed that the density profiles of haloes deviate from the NFW form beyond the virial radii, and display considerable scatter.  The abundance of these objects in N-body simulations has yet to be investigated.  Objects with the above characteristics would display rising $\\left< \\gamma_t \\right>$ at larger radii.  A comparison of our results with the shear profiles of detected dark clumps is difficult due to the fact that their redshifts are unmeasurable by definition.  Hence, it is impossible to determine whether features in the shear profile occur at the appropriate radii.  We have used the excursion set formalism to calculate the abundance of dark LPs.  In our approach, the number density of mass concentrations that are sufficiently overdense to meet the $S/N$ threshold is multiplied by the fraction that are unlikely to be detected in the soft X-ray band.  This subset of objects contains zero virialised sub-haloes with a probability $\\ge f_{\\mathrm{det}}$ of being detected via X-Ray emission.  In most cases of interest this fraction is extremely small, resulting in a suppression of dark LP abundances.  These results appear to be consistent with the average profiles derived in \\citet{Tavio2008}, which indicate that infall regions typically do not contain enough mass to create dark lenses.  In section \\ref{IMPLICATIONS}, we compared the differential number counts of dark LPs to ordinary clusters in a hypothetical shear-selected survey with source density $n = 30~\\mathrm{arcmin^{-2}}$ and intrinsic ellipticity dispersion $\\sigma_{\\epsilon} = 0.3$.  In both cases, we found that most detections originate from objects at $z_l \\sim 0.1 - 0.6$.  The dark LP number counts are generally $2-3$ orders of magnitude smaller than the cluster number counts.  We varied the minimum allowed power-law index, $\\alpha_{\\mathrm{min}}$, to simulate scenarios in which flat infall power-laws are dynamically unlikely.  We found that dark LP abundances are highly sensitive to $\\alpha_{\\mathrm{min}}$, dropping rapidly with increasing $\\alpha_{\\mathrm{min}}$.  If infall regions typically fall off steeper than $r^{-1}$, then lensing contributions from the outskirts of dark LPs may be insufficient to meet the detection threshold.  We also varied $f_{\\mathrm{det}}$ to explore the remote possibility that objects with large virialised sub-haloes could be detected as dark lenses.  We found that the differential number counts are relatively insensitive to $f_{\\mathrm{det}}$, varying only by $\\sim 20-30 $ per cent between $f_{\\mathrm{det}} = 0.3$ and $0.7$.  Finally, we have calculated the fraction of lenses that correspond to dark LPs.  We found that they constitute $\\la 0.4$ per cent of lenses at any given redshift.  Moreover, dark LPs account for $\\la 0.4$ per cent of the total number of lenses in our hypothetical shear-selected survey.  We therefore concluded that dark LPs are too rare to be considered a plausible dark lens candidate.  Our approach adds to initial work by WK2002 in several important ways.  The first is our use of the density profile (\\ref{LPprofileEQ}), which provides a physical model for lensing protoclusters that takes into account deviations from the NFW form beyond the virial radius.  Such deviations have been recently pointed out in high resolution N-body simulations by \\citet{Cuesta2008} and \\citet{Tavio2008}.  An additional advantage of (\\ref{LPprofileEQ}) is that we are able to quantify the $S/N$ contributions from virialised and unvirialised matter.    In contrast, the model of WK2002 does not allow one to quantify the lensing contribution from the central regions that meet the virialization over-density threshold.  Hence, in their approach it is possible to consider cases where the $S/N$ is dominated by the virial mass of the LP.  These cases typically occur when the total over-density of the LP is close to the virialization threshold.  Note that this difference is one of the reasons that our approach generally yields higher LP masses compared to the WK2002 results.  By forcing the virial region to be smaller in order to simultaneously minimize its lensing contribution and avoid X-ray detection, larger total masses are required to meet the $S/N$ threshold.  We also point out that the fraction of mass in lensing overdensities does not correspond to the fraction contained in dark objects.  Many of these overdensities contain large virialised sub-haloes.  In practice, these cases would correspond to true cluster detections since follow up X-ray searches would be sensitive to these sub-haloes.  By incorporating the halo statistics of \\citet{CasasMiranda2002}, our approach only counts overdensities without large virialised sub-haloes.  This key difference accounts for the lower dark LP abundances that we obtained compared to WK2002.  Following the work of \\citet{ReblinskyBartelmann1999,White2002,Hamana2004,Hennawi2005,Pace2007,Fan2007}, it is more likely that dark detections will correspond to false peaks resulting from: 1) LSS along the line-of-sight.  In this case, the $S/N$ is due to projected mass; it cannot be associated with a single isolated structure.  2) the random or correlated alignment of intrinsic galaxy ellipticities.  These alignments alone can lead to spurious detections, especially in shallow surveys.  However, it is also possible they can boost peaks that correspond to smaller mass concentrations \\citep{vonderLinden2006,Fan2007}.  Owing to an artificially high $S/N$, these detections can be misinterpreted as dark lenses.  Using ray-tracing through stacked snapshots of cosmological N-Body simulations, \\citet{Pace2007} tested the performance of the \\citet{Schirmer2004} filter used above (referred to as OAPT in their paper).  By removing individual lens planes of haloes that may be associated with a particular $S/N$ peak, \\citet{Pace2007} were able to separate true detections from spurious ones.  A true detection is associated with a cataloged cluster in the N-body simulation; a spurious detection remains when lens planes of individual candidates are removed.  They found that the OAPT filter yields spurious detection fractions $\\la 20~(25)$ per cent at $S/N = 4$ for source redshifts of $z_s = 1~(2)$.  For larger aperture sizes, this fraction decreases only mildly at higher $S/N$ thresholds.  In addition, they point out that these spurious detections are indistinguishable from true detections in a $S/N$ map.  Therefore, it is likely that LSS accounts for at least some of the dark lenses reported in the literature.  Note that a dark LP would likely be counted as a ``true\" detection with the algorithm of \\citet{Pace2007}.  The removal of the lens plane containing the dark LP would significantly diminish the signal observed in the $S/N$ map.  In addition, since the CVR would be cataloged in the N-body simulation, the detection might be associated with this small-mass halo.  Therefore, it would be instructive to determine what causes the lensing enhancement of small-mass detections in studies such as \\citet{Pace2007}.  Finally, we point out that the intrinsic galaxy ellipticities were randomly oriented in \\citet{Pace2007}.  The effect of \\emph{correlated} alignment of intrinsic ellipticites on the number of false detections was not taken into account.  As \\citet{Fan2007} points out, galaxy formation is sensitive to the local environment.  One would therefore expect the  orientation of a galaxy to at least be correlated with its closest neighbors.  Such alignments can increase the number of false detections in convergence $\\kappa$-maps significantly.  \\citet{Fan2007} showed that including this source of noise can increase the likelihood of false detection due to intrinsic ellipticities in a given field.  This increase can affect whether intrinsic ellipticities can be ruled out in a dark lens detection.  Future numerical studies on false peaks in weak lensing surveys should investigate this important possibility.  While the analytic model presented in this paper provides important insight into why dark LPs should be extremely rare, it is limited by several key issues.  The first is the simplistic density profile (\\ref{LPprofileEQ}), which neglects the effects of anisotropy and substructure on the $S/N$.  A more detailed analysis should incorporate these properties, which are expected to have a significant effect on the lensing signal.  Secondly, since the excursion set formalism does not yield any information about the density profiles of individual trajectories, it is impossible to rigorously determine whether objects meet the $S/N$ threshold.  The best we can do in our analytic approach is to assume that objects above the derived over-density threshold can be detected.  In addition, we have used a simple model for the X-ray luminosities of protoclusters in order to estimate the probability of detection.  In reality this is a highly complicated problem with many caveats that can only be addressed numerically.  Lastly, our model does not include galaxy overdensities.  We assume that objects with low virial masses also escape optical selection.  Future studies should focus on whether LPs display low galaxy overdensities as well.  Each of the above issues would be ideally addressed in a high-resolution N-body simulation containing a baryonic component.  Our model provides a starting point for more detailed investigations on the characteristics of simulated protoclusters and their impact on shear-selected samples."
    },
    "0809/0809.3759_arXiv.txt": {
        "abstract": "{Deuterated ions, especially H$_2$D$^+$ and N$_2$D$^+$, are abundant in cold ($\\rm\\sim 10~K$), dense ($\\rm\\sim 10^5~cm^{-3}$) regions, in which CO is frozen out onto  dust grains. In such environments, the N$_2$D$^+$/N$_2$H$^+$ ratio can exceed the elemental abundance ratio of D/H by a factor of $\\simeq 10^4$.} {We use the deuterium fractionation to investigate the evolutionary state  of Class 0 protostars. In particular, we expect the N$_2$D$^+$/N$_2$H$^+$ ratio to decrease as temperature (a sign of the evolution of the protostar) increases.} {We observed N$_2$H$^+$~1-0, N$_2$D$^+$~1-0, 2-1 and 3-2, C$^{18}$O~1-0 and HCO$^+$~3-2 in a sample of 20 Class~0 and borderline Class~0/I protostars. We determined the deuteration fraction and searched for correlations between the N$_2$D$^+$/N$_2$H$^+$ ratio and well-established evolutionary tracers, such as T$\\rm _{Dust}$ and the CO depletion factor. In addition, we compared the observational result with a chemical model.} {In our protostellar sample, the N$_2$H$^+$~1-0 optical depths are significantly lower than those found in prestellar cores, but the N$_2$H$^+$ column densities are comparable, which can be explained by the higher temperature and larger line width in protostellar cores. The deuterium fractionation of N$_2$H$^+$ in protostellar cores is also similar to that in prestellar cores. We found a clear correlation between the N$_2$D$^+$/N$_2$H$^+$ ratio and evolutionary tracers. As expected, the coolest, i.e. the youngest, objects show the largest deuterium fractionation. Furthermore, we find that sources with a high N$_2$D$^+$/N$_2$H$^+$ ratio show clear indication for infall (e.g. $\\rm \\delta v<0$). With decreasing deuterium fraction the infall signature disappears and $\\rm \\delta v$ tends to be positive for the most evolved objects. The deuterium fractionation of other molecules deviates clearly from that of N$_2$H$^+$. The DCO$^+$/HCO$^+$ ratio stays low at all evolutionary stages, whereas the NH$_2$D/NH$_3$ ratio is $>0.15$ even in the most evolved objects.} {The N$_2$D$^+$/N$_2$H$^+$ ratio is known to trace the evolution of prestellar cores; we show that this ratio can be used to trace core evolution even after star formation. Protostars with an N$_2$D$^+$/N$_2$H$^+$ ratio above 0.15 are in a stage shortly after the beginning of collapse. Later on, deuterium fractionation decreases until it reaches a value of $\\sim 0.03$ at the Class~0/I borderline.} ",
        "introduction": "\\label{intr}  Early attempts to describe an evolutionary sequence of protostars classified them according to their near- and mid-infrared spectra.  Lada \\& Wilking~(\\cite{class}) divided a sample of protostellar sources in the $\\rho$~Ophiuchi cloud into three classes (Class I to III with progressive evolutionary stage). Andr\\' e et al.~(\\cite{Class0}) introduced the Class~0 stage. The strong submillimetre emission of Class~0 objects relative to their bolometric luminosity suggests that the mass of the circumstellar envelope is larger than the mass of the central star. Class~0 objects are the youngest among the protostars. In the subsequent years, more sensitive indicators for protostar evolution have been established, such as the bolometric temperture (Myers \\& Ladd~\\cite{tbol}), the L$\\rm _{BOL}$/L$\\rm _{smm}$ ratio (Andr\\' e et al.~\\cite{Class0}) and the correlation between the bolometric luminosity and the distance normalized 1.3~mm flux (Saraceno et al.~\\cite{sar}). All these evolutionary indicators are based on the change of the spectral energy distribution (SED) of the protostar. As it heats up, an increasing fraction of the radiation is emitted at shorter wavelengths, leading to an increase of L$\\rm _{BOL}$ and T$\\rm _{BOL}$ with time.   Several theoretical models connect physical properties of the protostars with the time since gravitational collapse (e.g. Smith~\\cite{mod1}, Myers et al.~\\cite{mod2}). However, although the evolutionary stages determined by different models is the same (i.e. the time sequence is the same), the absolute ages vary significantly.  The absolute age for Class~0/I borderline objects, for example, varies between $10^4$  and a few times $10^5$  years (Froebrich~\\cite{fro}). In this paper, we show that it is possible to use the deuterium fractionation of N$_2$H$^+$ as an evolutionary tracer.   The chemistry of cold ($\\rm T<20~K$), dense ($\\rm n=10^5~cm^{-3}$) environments in space is peculiar. Common molecular tracers, such as CO and CS, appear reduced in abundance because they freeze out onto dust grains (e.g. Willacy et al.~\\cite{wlv98}, Caselli et al.~\\cite{cwt99}, Kramer et al.~\\cite{kal99}, Bergin et al.~\\cite{ted}, Redman et al.~\\cite{rrn02}, Tafalla et al.~\\cite{taf},~\\cite{tmc04}). By comparison, molecules such as N$_2$H$^+$ and NH$_3$ suffer less from depletion. The persistence of nitrogen-bearing species is likely due to the fact that N$_2$, the progenitor of these species, remains in the gas phase longer than CO, despite the fact that recent laboratory measurements have found that the binding energy, as well as the sticking coefficient of N$_2$ and CO on dust, are quite similar ({\\\"O}berg et al.~\\cite{ovf05}, Bisschop et al.~\\cite{bis}).  Recent theoretical work suggests that the resistance of nitrogen-bearing species to depletion may be due to the fact, that a significant fraction of interstellar nitrogen is in atomic (N) rather than molecular (N$_2$) form (Maret et al.~\\cite{mar}, Flower et al.~\\cite{flo}). Atomic nitrogen has a smaller binding energy than molecular nitrogen and CO.  One consequence of the low temperatures and freeze-out of CO in pre- and protostellar cores is that the abundance of deuterated molecules is greatly enhanced. In such regions, column density ratios for N$_2$D$^+$/N$_2$H$^+$ of 0.24 (Caselli et al.~\\cite{cas2}), D$_2$CO/H$_2$CO $\\simeq$ 0.01-0.1 (Bacmann et al. \\cite{bac03}), NH$_2$D/NH$_3$ up to 0.33 (Hatchell~\\cite{nh3})  and DCO$^+$/HCO$^+$ in the order of 0.01 (J\\o rgensen et al.~\\cite{jor}) are found. These ratios are three to four orders of magnitudes larger than the average ISM D/H ratio of 1.5$\\cdot10^{-5}$ (Oliveira et al.~\\cite{dh}). Models show that the abundance of H$_2$D$^+$, a fundamental molecule in deuterium chemistry, rises steeply as temperature decreases (e.g. Flower et al.~\\cite{fpfw}, Roberts \\& Millar~\\cite{rob}). Consequently, the deuterium fractionation of many other molecules also increases (Millar et al.~\\cite{tom}).  In a recent study, Crapsi et al.~(\\cite{crap}) investigated the N$_2$D$^+$/N$_2$H$^+$ ratio in 31 low-mass $starless~cores$ and found ratios in the range of a few percent to 0.44. They also found a tight correlation between the N$_2$D$^+$/N$_2$H$^+$ ratio and the CO depletion factor; the ratio is greater where CO is depleted. Furthermore, they attempted to find correlations between the N$_2$D$^+$/N$_2$H$^+$ ratio and various chemical and dynamical evolutionary indicators (e.g. core density, molecular column density, line width, and line asymmetry). Although the dependencies on particular indicators are not always clear, they identified a recognizable trend. The N$_2$D$^+$/N$_2$H$^+$ ratio was generally found to increase as a cloud core evolves towards the moment of protostellar collapse. Another results from their work is that the trends seen in the subsample of Taurus cores were more homogeneous than trends in the total sample. This homogeneity led them to speculate that a core's external environment plays an important role in its evolution. In this work we follow the trend of the N$_2$D$^+$/N$_2$H$^+$ ratio to the subsequent, early protostellar stages.     One recent study that also sought to determine the behavior of deuterated molecules in protostellar environments is from Roberts \\& Millar~(\\cite{RM}). They determined the N$_2$D$^+$/N$_2$H$^+$ ratio in five low mass $protostellar~cores$, using the University of Arizona 12m telescope (HPBW of 70$''$ for N$_2$H$^+$~1-0). Furthermore they measured the D$_2$CO/H$_2$CO and the HDCO/H$_2$CO ratio in these cores. The HDCO/H$_2$CO ratio was similar in most cores ($\\sim 0.06$), with the exception of the object with the highest N$_2$D$^+$/N$_2$H$^+$ ratio,  where HDCO/H$_2$CO was lowest. For four sources, the N$_2$D$^+$/N$_2$H$^+$ ratio appeared anti-correlated with the D$_2$CO/H$_2$CO ratio. Only in L~1527 were both N$_2$D$^+$ and D$_2$CO abundances anomalously low. They did not find a correlation between deuterium fractionation and bolometric temperature for any of the observed molecules. However, due to the small number of protostars, this work cannot rule out any trend of the deuterium fractionation in protostellar cores.  We observed 20 Class~0 protostars in N$_2$H$^+$ and N$_2$D$^+$, using the IRAM~30m telescope (HPBW of 27$''$ for N$_2$H$^+$~1-0, see section~\\ref{obs}). The goal of our study is to investigate the trend of deuterium fractionation after the moment of star formation. Spectra and line parameters are presented in Sect.~\\ref{res}. We expect that the N$_2$D$^+$/N$_2$H$^+$ ratio (determined in section~\\ref{ana1}) should decrease during protostellar evolution due to the internal heating of the core. As the core environment warms up, H$_2$D$^+$ and its more highly deuterated isotopologues are destroyed.  Consequently, all other deuterated molecules formed by reactions with H$_2$D$^+$ (such as N$_2$D$^+$) are destroyed as well. In Sect.~\\ref{Tdustsec}-~\\ref{ana4}, the N$_2$D$^+$/N$_2$H$^+$ ratio is compared with evolutionary indicators, namely dust temperature, CO depletion factor, and L$\\rm _{BOL}$/F$_{1.3}$. Possible correlations with kinematical parameters of the gas are investigated in section~\\ref{anav}. In addition, we compare the N$_2$D$^+$/N$_2$H$^+$ ratio with the deuterium fractionation of other molecules (section~\\ref{anam}). Finally, we compare our results with the predictions of a simple chemical model (section~\\ref{model}). ",
        "conclusions": "We observed 20 Class~0 protostars in N$_2$H$^+$~1-0, N$_2$D$^+$~1-0, 2-1 and 3-2, C$^{18}$O~1-0 and HCO$^+$~3-2. The integrated intensities of the N$_2$H$^+$~1-0 lines in  our sample are  comparable to those found in prestellar cores (Crapsi et al.~\\cite{crap}). The optical depths are typically 67\\% lower in the Class~0 sources, but because the excitation temperature is 2-3~K higher and the line width is about 2.5 time lager, the N$_2$H$^+$ column densities ($\\rm \\sim 10^{13}~cm^{-2}$) are also comparable. The presence of a protostar also affects the kinematics of the cores. N$_2$H$^+$ lines are significantly wider in Class~0 sources than in prestellar cores (on average, 0.61~km/s and 0.26~km/s, respectively). A comparison of our HCO$^+$ observations with observations conducted with a $\\sim$three times larger beam (Gregersen et al.~\\cite{greg}) leads us to the conclusion that HCO$^+$~3-2 stems  mainly from a compact region of about 4000~AU in size.  The dependence of the N$_2$D$^+$/N$_2$H$^+$ ratio on dust temperature, CO depletion factor, and the L$\\rm _{BOL}$/F$\\rm ^{220~pc}_{1.3}$ ratio clearly indicates a close relation between deuterium fractionation and evolutionary stage of a Class 0 protostar. For the sub-sample of sources in the Perseus cloud, the correlation is striking. This might be an indication for the influence of the environment on the deuterium fractionation.  The correlation of the deuterium fractionation and the CO depletion factor looks qualitatively like the one found in prestellar cores (Crapsi et al.~\\cite{crap}). However, whereas in prestellar cores, the N$_2$D$^+$/N$_2$H$^+$ ratio stays low until f$\\rm _{D}(CO)\\sim 10$, in protostars, the ratio starts to rise at f$\\rm _{D}(CO)\\sim 3$. This difference occurs because CO is highly depleted at the centre of prestellar cores, but not at the center of protostellar cores.  There is a weak correlation of N$_2$D$^+$ line width with N$_2$D$^+$/N$_2$H$^+$ ratio, but line broadening by multiple sources makes this conclusion suspect. A clear correlation with other kinematical tracers, especially with $\\delta$v, is found. Sources with a high N$_2$D$^+$/N$_2$H$^+$ ratio clearly show infall motion ($\\rm\\delta v\\sim -0.4$). The lower the deuterium fractionation gets, the more $\\delta$v increases. This might reflect the fact that outflow activity increases with evolutionary stage (Richer et al.~\\cite{pp5}).   A model with a power law density and temperature profile reproduces the correlation between the N$_2$D$^+$/N$_2$H$^+$ ratio and both T$\\rm _{Dust}$ and the CO depletion factor very well. In these models, the protostellar envelopes are chemically stratified. The inner part is too warm for CO to freeze out, so that the N$_2$H$^+$ abundance and the deuterium fractionation are low. At radii where the temperature drops below $\\sim 20$~K, CO is mostly depleted and the N$_2$D$^+$/N$_2$H$^+$ ratio increases by several orders of magnitude. At even larger radii, the lower density decreases the degree of CO depletion, and thus the N$_2$D$^+$/N$_2$H$^+$ ratio also decreases.   The results of this work, in combination with the results of Crapsi et al.~(\\cite{crap}) show that the deuterium fractionation is highest at the moment of collapse. The primary use of the deuterium fractionation of N$_2$H$^+$ is to identify very young protostellar objects, because the decline of the   N$_2$D$^+$/N$_2$H$^+$ ratio is rather steep.   Comparisons of N$_2$H$^+$ fractionation with the deuterium fractionation of HCO$^+$ and NH$_3$ show clear differences, which are due either to the influence of freeze-out on dust grains or differences in chemical reactions or both. The low DCO$^+$/HCO$^+$ ratio observed by J\\o rgensen et al.~(\\cite{jor}) is a prediction of our model. In young, i.e. cool objects, the NH$_2$D/NH$_3$ ratio, observed by Hatchell~(\\cite{nh3}) is comparable to the deuterium fractionation of N$_2$H$^+$. As the temperature increases, the N$_2$D$^+$/N$_2$H$^+$ ratio falls off much faster than the NH$_2$D/NH$_3$ ratio, which stays above 0.15 in all observed objects.  This cloud be due to desorption of deuterium enhanced ammonia from the dust grains.    \\begin{acknowledgement}  We would like thank the IRAM~30m staff for their assistance during the observations. Furthermore, we gratefully thank L. Pagani for providing us the Monte-Carlo code. M.~Emprechtinger, N.~H.~Volgenau, and M.~C.~Wiedner are member of the Nachwuchsgruppe of the Sonderforschungsbereich 494, which is funded by the Deutsche Forschungsgemeinschaft (DFG)  \\end{acknowledgement}"
    },
    "0809/0809.3959.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are a mixed class of sources consisting of, at least, the long duration and short-hard subclasses, the X-ray flashes, and the low-luminosity GRBs. In all cases, the release of enormous amounts of energy on a short timescale makes an energetic, relativistic or mildly relativistic fireball that expands until it reaches a coasting Lorentz factor determined by the amount of baryons mixed into the fireball. Radiation is produced when the blast wave interacts with the surrounding medium at an external shock, or when shell collisions dissipate kinetic energy at internal shocks. This series of notes is organized as follows: (1) The observational situation of GRBs is summarized; (2) Progenitor models of GRBs are described; (3) An overview of the the blast-wave physics used to model leptonic emissions is given; (4) GRB physics is applied to hadronic acceleration and ultra-high energy cosmic ray production; (5) Prospects for GRB physics and $\\gamma$-ray astronomy with the Fermi Gamma-ray Space Telescope (FGST, formerly GLAST), and space-based and ground-based observatories are considered. Also included are exercises and problems. ",
        "introduction": "A GRB may flare up from any direction in space. The classical, long-duration GRB (LGRB) releases most of its energy  in hard X-ray and soft $\\gamma$-ray (X/$\\gamma$) energies during (to the observer) a fraction of a second to tens of seconds. There is no compelling evidence that LGRBs are recurrent events. Therefore, a wide field-of-view instrument is necessary for serendipitous detection.  \\subsection{Discovery of Gamma Ray Bursts}  %[1-4] GRBs were discovered in data returned between 1967 and 1973 by the Vela series of satellites used to monitor compliance with the nuclear test ban treaty. The Vela spacecraft carried non-imaging CsI detectors that were sensitive in the $\\approx 200$ keV -- 1 MeV range (Klebesadel, Strong, \\& Olson 1973). Above the large background, coincident events were identified in the light curves. Timing studies and triangulation were used to give an approximate direction to the GRBs, revealing a cosmic/non-terrestrial origin. [{\\it Exercise:} Perform a simple timing analysis from synthetic satellite data to show how to reconstruct arrival direction information. Give uncertainty analysis.]  %[1-5] The basic observational data in GRB studies are the spectral photon fluxes $\\phi(\\e;t)$ measured at time $t$ and at photon energy $h\\nu = m_ec^2 \\epsilon$. From this quantity, one can derive, after subtracting backgound flux, the $\\nu F_\\nu$ flux (cgs units of ergs cm$^{-2}$ s$^{-1}$) $f_\\e(t) = m_e c^2 \\int_{\\e_1}^{\\e_2} d\\e \\; \\e\\; \\phi(\\e;t) \\approx m_ec^2 \\e_d^2 \\phi(\\e_d;t)$, where $\\e_d$ is the typical photon energy at which the detector is most sensitive. The fluence between times $t_1$ and $t_2$ and within the energy range $\\e_1$ and $\\e_2$ (cgs units of ergs cm$^{-2}$) is given by ${\\cal F}(t_1,t_2,\\e_1,\\e_2) = m_e c^2 \\int_{t_1}^{t_2}dt\\;\\int_{\\e_1}^{\\e_2} d\\e \\; \\e\\; \\phi(\\e;t)$. Flux and fluence distributions can be constructed from observations of many GRBs.  \\subsection{BATSE Observations: GRBs are Cosmological}  %[1-6] The Burst and Transient Source Experiment BATSE on CGRO consisted of an array of large area detectors (LADs) most sensitive in the 50-300 keV band, in addition to smaller spectroscopy detectors. The BATSE has given the most extensive data base of GRB observations during the prompt phase. It searched for GRBs by examining strings of data for $> 5.5\\sigma$ enhancements above background on the 64 ms, 256 ms, and 1024 ms time scales, and triggers on GRBs as faint as $\\approx 0.5$ ph cm$^{-2}$ s$^{-1}$, corresponding to energy flux sensitivities $\\lesssim 10^{-7}$ ergs cm$^{-2}$ s$^{-1}$. %[1-7,1-8] At hard X-ray and soft $\\gamma$-ray energies, the peak flux may reach hundreds of photons cm$^{-2}$ s$^{-1}$ in rare cases. Empirical morphological studies of LGRBs give various phenomenological relations, including hardness-intensity correlation, generic hard-to-soft evolution of $\\e_{pk}(t)$, and variability-distance correlation.  %[1-9] Expressing the sensitivity of a high-energy radiation detector in terms of a threshold energy flux $\\Phi_{thr}$ (same units as $f_\\e$) imposes the condition that $\\Phi \\geq \\Phi_{thr}$. For unbeamed sources with luminosity $L_*$ and distance $d$, $\\Phi = L_*/4\\pi d^2$, and the maximum source distance for a give source flux $\\Phi$ is $$d(\\Phi) = \\sqrt{L_*\\over 4\\pi\\Phi}\\;.$$ (The luminosity distance which includes cosmological effects is defined by $d_L = \\sqrt {L_*/4\\pi \\Phi}$, and $d\\cong d_L$ at low redshifts $z \\ll 1$. [{\\it Exercise.} Relate energy to fluence, including redshift.]) Hence the well-known $-3/2$ result for sources uniformly distributed with density $n_0$ in Euclidean space, namely \\begin{equation} N(>\\Phi) = N(< d) = 4\\pi n_0 \\int_0^{d(\\Phi)}dx\\;x^2 \\propto \\Phi^{-3/2}\\;\\;. \\label{Nphi} \\end{equation}  The $\\langle V/V_{max}\\rangle$ statistic \\begin{equation} \\big\\langle V/V_{max}\\big\\rangle\\;=\\;{1\\over N}\\sum_{i=1}^N\\; \\big({\\Phi_i\\over \\Phi_{thr}}\\big)^{-3/2}\\; \\label{vovervmax} \\end{equation} expresses the deviation from 0.5 expected for a uniform Euclidean distribution of sources (Schmidt 1968). Here $V$ stands for volume, and $V_{max}$ is the maximum volume from which a source with flux $\\Phi$ could be detected. Values of $\\langle V/V_{max}\\rangle > 0.5$ represent positive evolution of sources, that is, either more sources and/or brighter sources at large distances or earlier times.  Values of $\\langle V/V_{max}\\rangle < 0.5$ represent negative source evolution, i.e., fewer or dimmer sources in the past. Detailed treatments of the statistical properties of black hole sources must consider cosmological effects and evolution of source properties.  %[1-10] The integral size distribution of BATSE GRBs in terms of peak flux $\\phi_p$ is very flat below $\\sim 3$ ph cm$^{-2}$ s$^{-1}$, and becomes steeper than the $-3/2$ behavior expected from a Euclidean distribution of sources at $\\phi_p \\gtrsim 10$ ph cm$^{-2}$ s$^{-1}$. The directions to the BATSE GRBs are isotropically distributed in the sky and display no clustering toward the Galactic plane. When coupled with the flattening of the peak flux distribution, the implication is that we are at the center of an isotropic though bounded distribution of GRB sources. A cosmological distribution of sources is most compatible with these observations.   %[1-11] The duration of a GRB is defined by the time during which the middle 50\\% ($t_{50}$),  90\\% ($t_{90}$), or $i$\\% ($\\langle t_i\\rangle $) of the counts above background are measured. %[1-12] A bimodal duration distribution is measured, irrespective of whether the $t_{50}$ or $t_{90}$ durations are considered (Kouveliotou et al.\\ 1993) %[1-13] About two-thirds of BATSE GRBs are long-duration GRBs with $t_{90}\\gtrsim 2$ s, with the remainder comprising the short duration GRBs. The short duration GRBs tend to have harder spectra, so that they are referred to as the short, hard class of GRBs. They are also much weaker in average fluence in the BATSE range.  %[1-14] GRBs typically show a very hard spectrum in the hard X-ray to soft $\\gamma$-ray regime, with a photon index breaking from $\\approx -1$  at photon energies $E_{ph}\\lesssim 50$ keV to a $-2$ to $-3$  spectrum at $E_{ph} \\gtrsim$ several hundred keV. In BATSE studies,  the distribution of the peak photon energies $E_{pk}$ of the time-averaged $\\nu F_\\nu$ spectra of BATSE GRBs are typically found in the  100 keV - several MeV range. The time-averaged or, for very bright GRBs, time-sliced GRB spectrum is usually well-described by the ``Band function\" $N_B(\\e)$ (Band et al.\\ 1993), a power-law times an exponential that smoothly connects to a steeper power-law, given by $$ N_{_B}(\\epsilon) = k_{_B} \\; \\epsilon^{\\alpha} \\exp[-\\epsilon (\\alpha-\\beta)/\\epsilon_{_{\\rm br}}] \\; H(\\epsilon; \\e^{\\rm B}_{\\rm min},\\epsilon_{_{\\rm br}}) $$ % \\begin{equation} + \\;  k_{_B} \\; \\epsilon_{_{\\rm br}}^{\\alpha-\\beta} \\exp(\\beta-\\alpha) \\; \\epsilon^\\beta \\; H(\\e;\\e_{\\rm br},\\e^{\\rm B}_{\\rm max}) \\; .% \\label{Nband-eq} % \\end{equation} % Here $\\alpha$ and $\\beta$ are the low and high energy photon number indices, and $E_{\\rm br} = m_ec^2\\epsilon_{_{\\rm br}}$ is the ``break energy.'' The Heaviside function $H(x;y,z)$ vanishes everywhere except at $y\\leq x<z$, where it equals unity. The term $k_{_B}$ is the constant normalizing the number fluence to the $> 20$ keV BATSE energy fluence $\\Phi_{_B} (> 20 {\\rm \\; keV})$ of a particular GRB [{\\it Exercise:} Derive the form of the Band function and normalizing constant, and convolve with nontrivial model detector response.] %[1-15] Typical values of Band alphas $\\alpha \\approx -1$ and Band betas $\\beta \\approx -2.2$ -- $-2.5$. Deviations of these values give valuable information about radiation processes and existence of separate radiative components.    \\subsection{Beppo-SAX and the Afterglow Revolution}  %[1-16] Beppo-SAX GRB observations reveal that essentially all long-duration GRBs have fading X-ray afterglows. Beppo-SAX, launched April 30, 1996, carried three instruments. The Gamma Ray Burst Monitor was sensitive in the range 60 -- 600 keV to GRBs brighter than $\\approx 10^{-6}$ ergs cm$^{-2}$ s$^{-1}$. The Wide Field Camera was sensitive in the range 2 -- 30 keV down to $\\approx 10^{-10}$ ergs cm$^{-2}$ s$^{-1}$ and provided $\\lesssim 10^\\prime$ error boxes. %[1-17] The spacecraft was then slewed, requiring 6 -- 8 hours, but was fast enough for the Narrow Field Instruments, sensitive  to $> 0.1$ keV GRB emissions as faint as $\\sim 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$, to give error boxes $\\lesssim 0.5^\\prime$. The first X-ray afterglow was obtained from GRB 970228 (Costa 1997), which revealed an X-ray source that decayed according to a power-law, $\\phi_X(t) \\propto t^{\\chi}$, with $\\chi \\sim -1.33$. Typically, $\\chi \\sim -1.1$ to $-1.5$ in Beppo-SAX $\\sim 2$ -- 10 keV X-ray studies.  The small X-ray error boxes allowed deep optical and radio follow-up studies. %[1-18] GRB 970228 was the first GRB from which an optical counterpart was observed (van Paradijs et al.\\ 1997), and GRB 970508 was the first GRB for which a redshift was measured. %[1-19] Redshifts are provided by detection of optical emission lines from the host galaxy and absorption lines in the fading optical afterglow due to the presence of intervening gas. %[1-20,1-21] Host galaxies of long duration GRBs are bluish star-forming galaxies, primarily consisting of dwarf irregular and spiral galaxies. No optical counterparts were detected from approximately one-half of Beppo-Sax GRBs with well-localized X-ray afterglows, and are termed ``dark\" bursts. These sources may be undetected in the optical band because of dusty media, %\\cite{gro98} or intrinsically faint afterglows. These results give compelling evidence that LGRBs are associated with star-forming galaxies %[1-22] and the deaths of massive stars, especially given the detection of supernova emissions a few weeks after the GRB in a few, nearby faint GRBs %[1-23] which, however, may not be fully representative of the LGRBs (but rather the low luminosity GRBs, LLGRBs.  %[1-24] Apparent isotropic energy releases of LGRBs are enormous, exceeding $10^{54}$ ergs, and the redshift distribution of Beppo-SAX, BATSE, HETE-II and INTEGRAL (pre-Swift) GRBs is peaked near $\\langle z \\rangle \\approx 1$. %[1-25] Achromatic breaks in the optical curves of GRB afterglows gives evidence for a beamed/jetted geometries, reducing the apparent energy release to a beaming-corrected energy release by a beaming factor ${\\cal F}_{bm}$. %[1-26,1-27] Approximately 40\\% of GRBs have radio counterparts, and the transition from a scintillating to smooth behavior in the radio afterglow of GRB 980425 provides evidence for an expanding source. %[1-28] For BATSE/Beppo-SAX type GRBs, the lion's share $\\sim 65$\\% of the energy is released in the form of $>25$ keV X-rays and soft $\\gamma$ rays,  $\\sim 7$\\% is softer X-rays, $\\sim 0.1$\\% in optical, during the prompt phase $t \\lesssim 2\\langle t_i\\rangle$.  %[1-29] A class of X-ray rich GRBs, with durations on the order of seconds to minutes and X-ray fluxes in the range $10^{-8}$ -- $10^{-7}$ ergs cm$^{-2}$ s$^{-1}$ in the 2-25 keV band, was detected with many X-ray satellites, including Ariel V, HEAO-1, ROSAT, and Ginga, but conclusively established with Beppo-SAX (Heise et al.\\ 2001). %[1-30] These X-ray flashes (XRFs) are ``$\\gamma$-ray challenged,\" as indicated by Band-function model fits to the prompt emission spectrum. Several phenomenological correlations of XRFs and LGRBs have been reported. One goal is to establish a pseudo-redshift indicator, another a ``pulse\" paradigm, %\\cite{nor00} and correlations between the duration of quiescent and subsequent pulse periods in separated pulses. %\\cite{ram01}. [1-31] A quantitative relation between integrated fluence and $E_{pk}$ in well-defined pulses is reported. %\\cite{lk96}  %[1-32] The Amati relation (Amati et al.\\ 2002) correlates the $\\nu F_\\nu$ peak photon energy $E_{pk}$  with apparent isotropic energy release ${\\cal E}_{iso}= 10^{54}{\\cal E}_{54}$ ergs according to \\begin{equation} E_{pk} \\propto {\\cal E}_{iso}^{1/2}\\;. \\label{Amati} \\end{equation} The Ghirlanda relation (Ghirlanda et al.\\ 2004) correlates $E_{pk}$ with the collimation-corrected absolute $X/\\gamma$ energy release ${\\cal E}_{abs} = 10^{51}{\\cal E}_{51}$ ergs  according to \\begin{equation} E_{pk} \\propto {\\cal E}_{abs}^{0.7}\\;. \\label{Ghirlanda} \\end{equation} The Amati and Ghirlanda relations are challenging to explain and potentially useful for cosmolgical studies.    \\subsection{Swift and Various Classes of Bursting Sources}  %[1-33] The Swift Observatory is a NASA-ASI--supported MidEx launched November 20, 2004. Its main scientific payload consists of 3 instruments: the Burst Alert Telescope (BAT), triggering between 15 and 150 keV, and a very sensitive X-ray Telescope (XRT) operating between $\\sim 0.3$ -- few keV, and the Ultra-violet optical telescope (UVOT). The spacecraft slews autonomously in 20 -- 75 s in response to triggers from the BAT.  \\subsubsection{Long duration GRBs (LGRB).}  %[1-34] The redshift distribution of the GRBs detected with Swift, with average Swift GRB redshift $\\langle z \\rangle \\approx 2$, differing markedly from $\\langle z \\rangle \\approx 1$ for pre-Swift GRBs. This can be explained by the different triggering energy range of BAT vs.\\ BATSE, and $E_{pk}$-flux correlations. %[1-35] Knowledge of the early X-ray afterglow phase, an important goal realized by the Swift mission, provided surprising unpredicted behavior, most notably rapid X-ray declines and X-ray flares to late times (Zhang et al.\\ 2006; Nousek et al.\\ 2006). The physical meaning of the rapid declines in the X-ray flux between $\\sim 100$ -- $10^3$ s is in dispute, including muti-component jet models, refreshed jets, and hadronic signatures.  \\subsubsection{Short Hard GRBs (SGRB).}  %[1-36,1-37] As expected for an old population of host galaxy progenitors, as would be the case for coalescing neutron stars and black holes, SGRBs are expected to be associated with ellipticals as well as spirals. %[1-38] No evidence for supernova emissions has been found in SGRB afterglows, meaning that the progenitors are not associated with a young, massive stellar population. That this expectation received spectacular confirmation with Swift, and also HETE-II,  only opens up the surprises, including delayed X-ray afterglow emissions, rapid X-ray declines, and X-ray flares.  %[1-39] The heterogeneous Swift/HETE sample has a redshift distribution broadly distributed around $\\langle z \\rangle \\approx 0.4$, and differs in important ways from the LGRBs in terms of host galaxies, offsets from host galaxies, apparent energy releases, and lag-luminosity relation. Absolute energy determination for SGRBs is compromised by the difficulty of finding beaming breaks in light curves of SGRBs. ({Exercise: Beppo-SAX was not sensitive to SGRBs. Why?)   \\subsubsection{Low Luminosity GRBs (LLGRBs).}  %[1-40] The LLGRB class is the long tail to the LGRBs to low apparent energy release (reviewed by Zhang 2007). %[1-41] LLGRBs could very well be a separate and distinct population from LGRBs, possibly associated with magnetar activity, as indicated by GRB 060218, or extended black-hole fueling activity. A possible new class of GRBs was found in relation to nearby GRB 060614 with no supernova emissions, though it could possibly be a nearby SGRB. Identification of LLGRBs as a separate class, known since GRB 980425/SN 1998bw, makes us rethink conclusions about the relation of LGRBs to supernovae (SNe), which mostly derived from nearby LLGRBs. ({\\it Exercise:} Report on GRB 98025.)  \\subsubsection{Soft Gamma Repeaters (SGRs).}  %[1-42] The giant flare of December 27, 2004 was the most intense cosmic transient observed historically. It was detected by over 20 spacecraft from the Earth to Saturn, and apparently released $\\approx 1000$ times more energy than all the Milky Way's stars, $\\gtrsim 1$ erg cm$^{-2}$ in hard X-rays and $\\gamma$ rays at Earth (Hurley et al.\\ 2005). ({\\it Exercise.} Estimate the bolometric radiation flux from the stars in the Milky Way. Estimate the total stellar energy flux of the universe.) The event, releasing $\\sim 10^{47}$ ergs, lowered the level of the Earth's ionosphere.  %[1-43] SGR 041227 began with an $\\sim 0.2$ s long, hard spectrum spikes with $E\\sim 10^{46}$ -- $10^{47}$ erg. The spike is followed by a pulsating tail with $\\approx 1/1000^{th}$ of the energy. Viewed from a large distance, only the initial spike would be visible, and would resemble a short GRB. It could be detected out to 100 Mpc GRB 050906 at $z=0.03$ could be a magnetar flare. Giant flares like this must occur in other galaxies, and could comprise $\\sim 10$\\% of the the SGRB population.   \\subsection{ %[1-44] Summary}  \\begin{enumerate} \\item Pioneering phase (1967 \u0096- 1991): era of confusion \\item BATSE/CGRO era (1991 --2000): GRBs are cosmological \\item Beppo-SAX afterglow era (1996 -- 2006): distance and energy scale for classical LGRBs   (HETE-II/INTEGRAL) \\item XRFs, first seen with Ginga, clarified with Beppo-SAX \\item Low luminosity GRBs class (first example in 1998: GRB 980425/SN 1998bw) \\item Swift era (2004 -- ): early afterglows and the distance and energy scale for short-hard SGRBs (HETE-II) \\item Fermi GBM (GLAST Burst Monitor)/Fermi LAT (Large Area Telescope) era (2008 -- ) \\end{enumerate} ",
        "conclusions": ""
    },
    "0809/0809.3045_arXiv.txt": {
        "abstract": "Cosmochemical evidence for the existence of short-lived radioisotopes (SLRI) such as $^{26}$Al and $^{60}$Fe at the time of the formation of primitive meteorites requires that these isotopes were synthesized in a massive star and then incorporated into chondrites within $\\sim 10^6$ yr. A supernova shock wave has long been hypothesized to have transported the SLRI to the presolar dense cloud core, triggered cloud collapse, and injected the isotopes. Previous numerical calculations have shown that this scenario is plausible when the shock wave and dense cloud core are assumed to be isothermal at $\\sim 10$ K, but not when compressional heating to $\\sim 1000$ K is assumed. We show here for the first time that when calculated with the FLASH2.5 adaptive mesh refinement (AMR) hydrodynamics code, a 20 km/sec shock wave can indeed trigger the collapse of a 1 $M_\\odot$ cloud while simultaneously injecting shock wave isotopes into the collapsing cloud, provided that cooling by molecular species such as H$_2$O, CO$_2$, and H$_2$ is included. These calculations imply that the supernova trigger hypothesis is the most likely mechanism for delivering the SLRI present during the formation of the solar system. ",
        "introduction": "The discovery of evidence for live $^{26}$Al (half-life of $0.7 \\times 10^6$ yr) at the time of the formation of refractory inclusions in the Allende meteorite (Lee et al. 1976) led quickly to the suggestion that the $^{26}$Al was synthesized in a supernova, then transported by the supernova's shock wave to a dense molecular cloud, where the shock wave triggered the collapse of a cloud core and injected the $^{26}$Al into the collapsing presolar cloud (Cameron and Truran 1977). Detailed numerical calculations showed that such simultaneous triggering and injection was possible, provided that the shock wave had slowed to speeds of $\\sim$ 10 to $\\sim$ 40 km/sec (Boss 1995; Foster \\& Boss 1996, 1997) and that the shock wave and cloud were effectively isothermal at $\\sim$ 10 K. This isothermal assumption is appropriate for the radiative phase of a supernova shock (Chevalier 1974), which occurs when the shock has travelled $\\sim$ 10 pc and swept up a cool shell of gas and dust. Recent adaptive mesh refinement (AMR) studies (Nakamura et al. 2006; Melioli et al. 2006) have shown that shock-triggered star formation is likely to occur when the supernova shock has evolved into a radiative shock, i.e., the shock wave is able to cool so rapidly by radiation that the shock wave is effectively at the same temperature as the ambient medium, confirming the results of Boss (1995) and Foster \\& Boss (1996, 1997).  Vanhala \\& Cameron (1998; hereafter VC) used a smoothed particle hydrodynamics (SPH) code and detailed equations of state to confirm the previous results for isothermal shocks and clouds. However, when VC allowed their models to become nonisothermal by including compressional heating and molecular, atomic, and dust cooling, they found that they could not achieve simultaneous triggering of cloud collapse and injection of shock wave material, raising serious doubts about the supernova trigger hypothesis.  Alternative explanations for the origin of the SLRI in chondritic meteorites have also arisen. Shu et al. (1997) suggested that the SLRI were produced {\\it in situ} by spallation reactions involving energetic particles emanating from protosolar flares striking dust grains near the X-wind point, from whence the grains would be lofted upward by the X-wind and transported outward to the solar nebula. However, the SLRI $^{60}$Fe cannot be produced by spallation, and requires a stellar nucleosynthetic source (Tachibana \\& Huss 2003) such as a supernova. The fact that the results of $^{26}$Al dating agree well with the independent Pb-Pb dating system (Connelly et al. 2008) implies that the $^{26}$Al was spatially homogeneous in the solar nebula, which is inconsistent with the bulk of its production by spallation reactions at the X-wind point. Regardless of whether spallation reactions contributed to the SLRI inventory, then, a supernova source for the $^{26}$Al and $^{60}$Fe appears to be necessary.  Given a supernova origin for the $^{26}$Al and $^{60}$Fe, it has also been suggested that a nearby ($\\sim$ 0.1 pc) supernova may have injected these SLRI into the already-formed solar nebula (Ouellette, Desch \\& Hester 2007). While such hot shock front gas is unable to penetrate into the much denser disk gas (Chevalier 2000), Ouellette et al. (2007) suggested that SLRI residing in dust grains micron-sized and larger shot through the stalled shock-wave gas and into the disk. However, supernova dust grains are essentially all smaller than 0.1 micron, and are sputtered to even smaller sizes in the shock (Bianchi \\& Schneider 2007). Gounelle \\& Meibom (2008) noted that disks formed within $\\sim$ 0.3 pc of a massive star would be photoevaporated away prior to its supernova explosion. Furthermore, Krot et al. (2008) argued that injection into a late-phase solar nebula would have led to the injection of stable oxygen isotopes as well as SLRI into the disk, leading to an oxygen isotope distribution that would be inconsistent with the observed values and their explanation by mass-independent fractionation. Krot et al. (2008) and Thrane et al. (2008) argued that injection must have occurred instead into the presolar cloud, so that the sun and the solar nebula shared a common reservoir of oxygen isotopes.  Here we return to the question of whether nonisothermal shock fronts are fatal for the supernova triggering hypothesis, by using the FLASH2.5 AMR code and an improved cooling law to reinvestigate this basic question of solar system origin. ",
        "conclusions": "These models show that when cooling by molecular species is included, a 20 km/sec shock wave is able to trigger the gravitational collapse of an otherwise stable, solar-mass dense cloud core, as well as to inject appropriate amounts of supernova shock wave material into the collapsing cloud core. This injected material consists of shock wave gas as well as dust grains small enough to remain coupled to the gas, i.e., sub-micron-sized grains, which are expected to characterize supernova shock waves (Bianchi \\& Schneider 2007) and to carry the SLRI whose decay products have been found in refractory inclusions of chondritic meteorites (Lee et al. 1976). Evidently a radiative-phase supernova shock wave is able to cool sufficiently rapidly to behave in much the same way as a shock wave that is assumed to remain isothermal with the target cloud (Boss 1995; Foster \\& Boss 1996, 1997; Vanhala \\& Boss 2000, 2002). These models thus lend strong support to the hypothesis first advanced by Cameron \\& Truran (1977) that a supernova shock wave carrying SLRI may have triggered the formation of the solar system."
    },
    "0809/0809.4179_arXiv.txt": {
        "abstract": "We report the results of our project devoted to study the chemical enrichment history of the field population in the Magellanic Clouds using Ca II triplet spectroscopy. ",
        "introduction": "The study of the ages and metallicities of the resolved stars in a nearby galaxy provides very detailed information on its evolutionary history. The Magellanic Clouds are examples of galaxies where this method can be applied very successfully. However, their vastness and our limitations in observing sizable samples of stars imply that there are particularly large gaps of knowledge in this area, compared with others. For example, their chemical enrichment histories have mainly been studied from their cluster systems. However, these objects have some drawbacks, like the age-gap of the LMC clusters or the single old cluster known in the SMC. Although the SMC is more metal-poor than the LMC, the chemical evolution of the cluster systems has been qualitatively similar in both galaxies. They show a first episode of chemical enrichment followed by a period of slow chemical evolution until around 4 Gyrs ago. Then, both galaxies took off again (\\cite[Olszewski \\etal\\ 1991]{ol91};\\cite[Piatti \\etal\\ 2001]{piatti01}). However, as mentioned above, there are some epochs in which the lack of clusters makes it difficult to extract definite conclusions. The chemical evolution of the Magellanic Clouds has also been investigated from their planetary nebulae (\\cite[Dopita \\etal\\ 1991]{dopita97};\\cite[Idiart \\etal\\ 2007]{idiart07}), which show a behaviour similar to the clusters one. Finally, \\cite[Cole \\etal\\ (2005)]{c05} have studied the chemical evolution of the LMC bar based on a sample of red giant branch (RGB) stars. The age-metallicity relationship (AMR) of the bar stars is similar to the cluster's one, particularly for the older ages. However, the increase of metallicity observed in the clusters in the last 3 Gyrs is not observed in the bar.  What about the field populations in the LMC disk and in the SMC? Do they share the cluster and planetary nebulae behaviour in each galaxy? Does chemical evolution show a global pattern in each galaxy or on the contrary, does it depends on the position? To address these questions, and as part of a large project devoted to study the stellar content of the Magellanic Clouds, we have obtained metallicities and ages for stars in 4 large LMC and 13 SMC fields. In this paper we will discuss the procedure used to derive metallicities and ages, and describe our main conclusions on the chemical enrichment history of both Magellanic Clouds. ",
        "conclusions": ""
    },
    "0809/0809.1106_arXiv.txt": {
        "abstract": "We compute the mid-infrared extinction law from $3.6-24\\:\\mu$m in three molecular clouds: Ophiuchus, Perseus, and Serpens, by combining data from the ``Cores to Disks'' \\spitzer{} Legacy Science program with deep \\jhks{} imaging. Using a new technique, we are able to calculate the line-of-sight extinction law towards each background star in our fields. With these line-of-sight measurements, we create, for the first time, maps of the \\chisq{} deviation of the data from two extinction law models.  Because our \\chisq{} maps have the same spatial resolution as our extinction maps, we can directly observe the changing  extinction law as a function of the total column density.  In the \\spitzer{} IRAC bands, $3.6-8\\:\\mu$m, we see evidence for grain growth.  Below $A_{K_s} = 0.5$, our extinction law is well-fit by the \\citet{weingartner01} $R_V = 3.1$ diffuse interstellar medium dust model.  As the extinction increases, our law gradually flattens, and for $A_{K_s} \\ge 1$, the data are more consistent with the Weingartner \\& Draine $R_V = 5.5$ model that uses larger maximum dust grain sizes.  At $24\\:\\mu$m, our extinction law is $2-4\\times$ higher than the values predicted by theoretical dust models, but is more consistent with the observational results of \\citet{flaherty07}.  Lastly, from our \\chisq{} maps we identify a region in Perseus where the IRAC extinction law is anomalously high considering its column density.  A steeper near-infrared extinction law than the one we have assumed may partially explain the IRAC extinction law in this region. ",
        "introduction": "Dust is a ubiquitous component of the interstellar medium. It is important to understand the dust because it attenuates, or extincts, the light from background objects seen through the dust. The amount of extinction varies with both wavelength and position. Because of the constancy of the gas-to-dust ratio, the quantity of dust gives you a direct measure of the cloud's mass \\citep{bohlin78}. Note, however that the \\citet{bohlin78} equation does depend on the $E(B-V)$ color excess, which can potentially vary depending on interstellar medium properties such as metallicity.  The dust extinction must also be accounted for when computing the luminosities for protostars within molecular clouds. The wavelength dependence of the extinction, known as the extinction law, directly relates to dust grain properties such as size, composition, and the presence or lack of icy mantles on the grains.  Many authors have found that the near-infrared extinction law has the form of a power law, $A_\\lambda \\propto \\lambda^{-\\beta}$, between $\\lambda \\sim 1\\:\\mu$m and $\\sim4\\:\\mu$m with $\\beta = 1.6-1.8$ \\citep[and references therein]{draine03}. The extinction law in the mid-infrared wavelengths has not been as well studied. \\citet{rieke85} found the near-infrared power law extended out to $\\sim7\\:\\mu$m before the extinctions increased again towards the $10\\:\\mu$m silicate absorption peak. More recent results have found a different behavior. \\citet{lutz99} used ISO to measure hydrogen recombination lines towards the Galactic center and \\citet{indebetouw05} and \\citet{flaherty07} used the Spitzer Space telescope to measure the mid-infrared extinction law in several regions. These authors all find a similar extinction law in the mid-infrared which is much flatter than the \\citet{rieke85} result.  To date no one has studied the changes in the extinction law within a large region with high spatial resolution. In this paper we use \\spitzer{} and deep ground-based \\jhks{} observations to probe changes in the dust properties within three molecular clouds: Ophiuchus, Perseus, and Serpens. We compare our results with different dust models to investigate the relationship between dust grain properties and total extinction (column density). In \\S\\,\\ref{sec:obs} we describe our observations. Then, in \\S\\,\\ref{sec:clouddatareduce} we discuss our data reduction method, including how we identified background stars and computed line-of-sight (LOS) extinctions. In \\S\\,\\ref{sec:cloudext}-\\ref{sec:dustprop-clouds} we create extinction and \\chisq{} maps and compare them to each other. We also compute the average extinction law from $3.6-24\\:\\mu$m for different ranges of extinction. Finally, we summarize our results in \\S\\,\\ref{sec:cloud-discuss}. ",
        "conclusions": "Conclusions }  In this paper we took deep \\jhks{} observations of regions within three molecular clouds: Ophiuchus, Perseus, and Serpens. These were combined with \\spitzer{} c2d data to create band-merged catalogs in eight wavebands from $1.2-24\\:\\mu$m. Our final catalogs contain 2,365 stars in Ophiuchus, 11,280 in Perseus, and 49,485 in Serpens. From these catalogs we created uniform maps of the extinction and \\chisq. We also computed the average extinction law in four extinction bins: $0 < A_{K_s} \\le 0.5$, $0.5< A_{K_s} \\le 1$, $1 < A_{K_s} \\le 2$, and $A_{K_s} > 2$.  Our results in the IRAC bands, $3.6-8\\:\\mu$m, show that grain growth occurs in the dense regions of our clouds. The \\chisq{} maps show the deviations from two dust models: \\citet{weingartner01} with $R_V = 3.1$ and $R_V = 5.5$ (abbreviated as WD3.1 and WD5.5). There is a strong correlation between high extinction and high WD3.1 \\chisq{}, but not with WD5.5 \\chisq. This suggests denser regions are more consistent with the WD5.5 extinction model, which has grain sizes up to $10\\times$ larger than those in the WD3.1 dust model. The same behavior is seen in the average extinction law. For $0 < A_{K_s} \\le 0.5$, it is most consistent with WD3.1 ($94\\%$ probability of the data fitting this model).  As \\ak{} increases, the extinction law gradually flattens and becomes more consistent with WD5.5 for $1 < A_{K_s} \\le 2$, with a $76\\%$ probability.  For $A_{K_s} > 2$, the WD5.5 model is not as good a fit to the data, primarily due to the $A_{5.8}/A_{K_s}$ relative extinction, which is higher than the model.  This may indicate the presence of water ice in the densest regions.  Not all areas in our clouds can be explained using grain growth. Our \\chisq{} map of Perseus reveals a region that has only a modest average extinction of $A_{K_s} = 0.60$, yet has a peak WD3.1 \\chisq{} of 39. The extinction law in this region is much flatter than WD5.5. Because of the low average extinction here, it doesn't seem likely that grain growth is altering the extinction law, nor that ices have formed on the dust grains. It is possible the background sources in this region are not really stars, or that the near-infrared extinction law is steeper in this region. More work is needed to understand the physical cause of the extinction law in this region.  Unlike in the IRAC bands, the $24\\:\\mu$m extinction law is $2-4\\times$ higher than current dust models predict. There is also a negative correlation between \\ak{} and the value of $A_{24}/A_{K_s}$. Both of these results maybe be partially mitigated by considering the small numbers of stars detected at $24\\:\\mu$m or slight changes to our average stellar model flux. However, even with the most conservative assumptions, our average value for $A_{24}/A_{K_s}$ is still $\\sim2\\times$ larger than that predicted by dust grain models. This results agrees with the measured values from \\citet{flaherty07}. Future dust models will need to reproduce the relatively  large value of the $24\\:\\mu$m extinction."
    },
    "0809/0809.0701_arXiv.txt": {
        "abstract": "A low frequency stochastic background of gravitational waves may be detected by pulsar timing experiments in the next five to ten years. Using methods developed to analyze interferometric gravitational wave data, in this paper we lay out the optimal techniques to detect a background of gravitational waves using a pulsar timing array. We show that for pulsar distances and gravitational wave frequencies typical of pulsar timing experiments, neglecting the effect of the metric perturbation at the pulsar does not result in a significant deviation from optimality. We discuss methods for setting upper limits using the optimal statistic, show how to construct skymaps using the pulsar timing array, and consider several issues associated with realistic analysis of pulsar timing data. ",
        "introduction": "The search for gravitational waves is at the forefront of current fundamental physics research. The direct detection of gravitational waves will usher in a new era in astronomy and astrophysics. Gravitational waves will reveal information about black holes, supernovae, and neutron stars that cannot be gleaned from electromagnetic observations. Furthermore, the detection of a gravitational wave background will open an observational window onto a time in the early universe before recombination, prior to which the universe is opaque to electromagnetic waves. The scientific rewards for such a detection would be truly exceptional.  Several international efforts are underway to detect gravitational waves and two of these efforts are expected to result in the detection of gravitational waves in the next 5 to 10 years: Interferometric ground-based gravitational wave detectors and pulsar timing observations.   Neutron stars radiate powerful beams of radio waves from their magnetic poles. If a neutron star's magnetic poles are not aligned with its rotational axis, the beams sweep through space like the beacon on a lighthouse and the neutron star is said to be a pulsar. If the Earth lies within the sweep of a pulsar's beams, the star is observed as a point source in space emitting short, rapid bursts of radio waves~\\cite{Hulse:1974eb}.  Due to their enormous mass, neutron stars have a very large moment of inertia and the radio pulses we observe arrive at a very constant rate.  Pulsar timing experiments exploit this regularity~\\cite{1975GReGr...6..439E,Sazhin78,Detweiler,% LorimerKramer2005}. Fluctuations in the time of arrival of radio pulses, after all known effects have been subtracted out, could be due to the presence of gravitational waves. Since the 1970s, when these ideas were first conceived, pulsar timing precision has improved dramatically. Several known pulsars can now be timed with a precision of about 1 micro-second, and a handful can be timed with a precision around several hundred nanoseconds \\cite{Hobbs:2005yx}. Recent work~\\cite{Jenet:2005pv} has shown that the presence of nanohertz gravitational waves could be detected by observing 20 pulsars with timing precisions of 100 nanoseconds over a period of 5 to 10 years. Non-detection would still improve current bounds on the low frequency stochastic gravitational wave background~\\cite{Jenet:2006sv}.   Gravitational waves from supermassive black hole binary systems could be detected via pulsar timing observations~\\cite{Wyithe:2002ep,Jaffe:2002rt,Enoki:2004ew,% Kramer:2004hd,alberto}. In addition, pulsar timing has the potential to measure the polarization properties of gravitational waves which could confirm (or even change) the current theory of gravity~\\cite{PhysRevLett.30.884,PhysRevD.8.3308}. Gravitational wave observations in the nanohertz band could also yield information about the early universe~\\cite{Boyle:2007zx}. Cosmic strings, line-like topological defects, could produce gravitational waves in the nanohertz band. Cosmic strings can form during phase transitions in the early universe due to the rapid cooling that takes place after the Big Bang~\\cite{Kibble:1976sj,Hindmarsh:1994re,% alexbook}. Cosmic string production is generic in supersymmetric grand unified theories~\\cite{Jeannerot:2003qv}. Additionally, in string theory motivated cosmological models cosmic strings may also form (dubbed cosmic superstrings to differentiate them from field theoretic cosmic strings)~\\cite{Jones:2002cv,% Sarangi:2002yt,Dvali:2003zj,Jones:2003da,Copeland:2003bj,% Jackson:2004zg}. Cosmic strings and superstrings are expected to produce a background of stochastic gravitational waves and bursts of gravitational waves~\\cite{Berezinsky:2000vn,% Damour:2000wa,Damour:2001bk,Damour:2004kw,Siemens:2006vk,% Siemens:2006yp} that could be detected using pulsar timing observations. Pulsar timing observations are already producing some of the most interesting constraints on cosmic string models and a detection would have profound implications~\\cite{Siemens:2006vk,% Siemens:2006yp}.  Recently Lommen~\\cite{Lommen:2002je} produced an upper limit on the stochastic gravitational wave background using observations of three millisecond pulsars spanning 17 years. The methods used by Lommen were based on those developed by Kaspi and collaborators~\\cite{Kaspi:1994hp} and have been the subject of some criticism in the literature~\\cite{Damour:2004kw,Jenet:2006sv}.  More recently Jenet and collaborators~\\cite{Jenet:2005pv,Jenet:2006sv} developed a new technique for gravitational wave stochastic background searches in pulsar timing data and applied it to Parkes Pulsar Timing Array data~\\cite{Manchester:2007mx,PPTA}.  In this paper we consider optimal strategies for extraction of a gravitational wave stochastic background signal using data from a pulsar timing array. Our methods are based on those developed for and used in ground based interferometric gravitational-wave detectors such as LIGO and Virgo~\\cite{Michelson,Christensen:1992wi,Flanagan:1993ix,% AllenandRomano,Abbott:2003hr,% Abbott:2005ez,Abbott:2006zx,Abbott:2007tw,Abbott:2007wd}, and improve on existing methods in several ways. In Section~\\ref{sigsec} we write the redshift of pulsar signals induced by passing gravitational waves, first derived by Detweiler~\\cite{Detweiler}, in a coordinate-independent way more suitable for our analysis, and discuss its form in the frequency domain including the long-wavelength limit.  In Section~\\ref{DSsec} we construct the optimal cross-correlation filter for a pulsar pair by maximizing the signal to noise. We find that the overlap reduction function is well approximated by a constant, or equivalently, that the metric perturbation at the pulsar can be neglected for values of pulsar distances and gravitational wave frequencies typical of pulsar timing experiments, without significant losses in sensitivity. In Section~\\ref{DSsec} we also show how to construct the optimal combination of cross-correlations of pulsar pairs in a pulsar timing array and include a more sophisticated derivation of the optimal detection statistic based on the likelihood ratio. In Section~\\ref{ULDsec} we discuss upper limit and detection methods. In Section~\\ref{SMsec} we show how to construct skymaps using pulsar timing data---a pulsar timing radiometer. In Section~\\ref{issues} we discuss several important issues relating to the realistic analysis of pulsar timing data, including the Lomb-Scargle periodogram for power spectrum estimation of unevenly sampled data, and optimal procedures for computing Fourier transforms. We conclude in Section VIII. Lommen, Romano and Woan~\\cite{Lommenea} will extend our work using a likelihood based approach developed in~\\cite{Allen:2002jw}, and consider the case of stochastic backgrounds that are loud compared to the noise, closely examine time-domain implementations of the optimal statistic, and provide a detailed comparison of the optimal statistic described here with the methods of Jenet and collaborators~\\cite{Jenet:2005pv,Jenet:2006sv}. ",
        "conclusions": ""
    },
    "0809/0809.5124_arXiv.txt": {
        "abstract": "In this work we study how the cosmological parameter, the Hubble constant $H_0$, can be constrained by observation of very high energy (VHE) $\\gamma$-rays at the TeV scale. The VHE $\\gamma$-rays experience attenuation by background radiation field through $e^+e^-$ pair production during the propagation in the intergalactic space. This effect is proportional to the distance that the VHE $\\gamma$-rays go through. Therefore the absorption of TeV $\\gamma$-rays can be taken as cosmological distance indicator to constrain the cosmological parameters. Two blazars Mrk 501 and 1ES 1101-232, which have relatively good spectra measurements by the atmospheric Cerenkov telescope, are studied to measure $H_0$. The mechanism measuring the Hubble constant adopted here is very different from the previous methods such as the observations of type Ia supernovae and the cosmic microwave background. However, at $2\\sigma$ level, our result is consistent with which given by other methods. ",
        "introduction": "The modern cosmology achieves great progress in recent years. A concordance $\\Lambda$CDM cosmology has been built thanks to the precise observations of the ``distance indicators'' type Ia supernovae \\citep[SNe Ia,][]{1998AJ....116.1009R,2004ApJ...607..665R, 1999ApJ...517..565P} and the anisotropy of cosmic microwave background \\citep[CMB,][]{2000Natur.404..955D,2003ApJS..148..175S}. There are also other cosmological probes such as the large scale structures \\citep{2004ApJ...606..702T}, galaxy clusters \\citep{2008MNRAS.383..879A}, observational Hubble parameters \\citep{2005PhRvD..71l3001S} and the weak gravitational lensing \\citep{2008PhR...462...67M} further supporting this scenario. Different methods are roughly consistent with each other within the observation uncertainties. It is very important to develop additional complementary observational evidence to test this model and measure the cosmological parameters.  It has been pointed out that the observations of VHE $\\gamma$-rays at the energy scale of hundred GeV to TeV scale are possible to provide another independent constraint on the cosmological parameters \\citep{1994ApJ...423L...1S}. Thanks to the rapid technical development of VHE $\\gamma$-ray detection, especially the atmospheric Cerenkov telescopes, great progress of VHE $\\gamma$-ray astronomy is achieved in recent years and a large number of VHE $\\gamma$-ray sources are detected. Even $\\gamma$-ray sources at cosmological distances, such as Mrk 501 at $z=0.034$ and 1ES 1101-232 at $z=0.186$ are observed. More importantly the spectra of these sources have been measured with relatively high precision, which provide us the possibility to untangle the effect of attenuation when $\\gamma$-rays propagate in the intergalactic space.  The attenuation of VHE $\\gamma$-rays is induced by the electron-positron pair production $\\gamma+\\gamma_{bk} \\rightarrow e^+ + e^-$ during its propagation in the background radiation field \\citep{Nikishov1962, 1967PhRv..155.1404G,1966PhRvL..16..252G}. This process is actually complex. The observed spectra of extragalactic sources are related with several issues: the intrinsic spectra at sources, the cross section of $\\gamma\\gamma$ interaction, the intensity of the cosmic infrared background (CIB), and the physical distance the VHE $\\gamma$-rays cross. Even before the first detection of the VHE $\\gamma$-rays from distant extragalactic sources, the perspective to explore the CIB using the attenuation effect was proposed \\citep{1992ApJ...390L..49S}. The discovery of the first extragalactic VHE $\\gamma$-ray source, an active galactic nuclei (AGN) Mrk 421, was performed by Whipple in 1992 \\citep{1992Natur.358..477P}. Till now more than 20 extragalactic sources, most of which are AGNs, are discovered by ground-based observatories\\footnote{See the VHE source web by Wagner, http://www.mppmu.mpg.de/$\\sim$rwagner/sources/}. The observations of these sources provide us valuable information in understanding the $\\gamma$-ray production mechanism and give useful implication or constraint on the CIB intensity \\citep[e.g.,][]{1999APh....11...35C,2000A&A...353...97K, 2001A&A...371..771R,2006Natur.440.1018A}. On the other hand, once the primary spectra and the CIB intensity are specified, the distance-redshift relation (accordingly the cosmological model parameters) of the sources can be derived from the absorption effect.  In this work we try to constrain the Hubble constant from the absorption effect of distant  VHE $\\gamma$ sources by the CIB. By a global fitting to the observational spectra of two TeV blazars, Mrk 501 and 1ES 1101-232, we get the Hubble constant with larger errors compared with other methods. In our work the $\\Lambda$CDM universe with matter component $\\Omega_M=0.28$ and dark energy $\\Omega_{\\Lambda}=0.72$ is adopted \\citep{2008arXiv0803.0547K}. We find that the best-fitting to the data of the two sources intend to give similar Hubble constant, although they have very different intrinsic spectra and redshifts. This is very encouraging that the attenuation may indeed give implications on the cosmological parameters. We noticed in a previous work \\cite{2008arXiv0804.3699B} adopt the similar effect to derive the lower limit of the Hubble constant from observation of Mrk 501. In their work, the direct measurements of CIB intensity was adopted and the intrinsic spectrum of the source was required to be concave.  The outline of this paper is as follows. Sec. 2 describes the absorption of VHE $\\gamma$ photons by CIB. In Sec. 3 we present an introduction to the observations of the two TeV blazars. The implication on Hubble constant is given in Sec. 4. Finally we give conclusion and some discussion in Sec. 5. ",
        "conclusions": "In this work we constrain the cosmological parameters, especially the Hubble constant $H_0$, by observations of extragalactic VHE $\\gamma$-ray sources at cosmological distances. The VHE $\\gamma$-rays experience attenuation by background radiation field through $e^+e^-$ pair production. This attenuation is proportional to the distance that the VHE $\\gamma$-rays go through. Therefore the absorption of VHE $\\gamma$-rays can be used to determine the distance of VHE $\\gamma$-ray sources, accordingly to get constraints on the cosmological parameters.  By fitting the spectra of two blazars Mrk 501 and 1ES 1101-232 we get the best-fitting Hubble constant is $H_0\\sim 100$ km s$^{-1}$ Mpc$^{-1}$. A large Hubble constant implies that the absorption of VHE $\\gamma$-rays is not as significant as we usually expected \\citep{2000PhLB..493....1P,2006Natur.440.1018A}. Since the observations of VHE $\\gamma$-rays and CIB are still rough, the errors of the fitting parameters are also very large. The mechanism constraining the Hubble constant adopted here is very different from previous methods, however, our results are consistent with the recent combined analysis of CMB, SNe Ia and baryon acoustic oscillation data on Hubble constant $h=0.701\\pm0.026$ at $2\\sigma$ level \\citep{2008arXiv0803.0547K}.  In fact, for each single method measuring the Hubble constant there are relatively large uncertainties, including both the statistical and the systematic ones. The HST Key Project measured the Hubble constant from several secondary distance indicators using Cepheid as calibration \\citep{2001ApJ...553...47F}. They gave the results that $h=0.71\\pm0.06$ for SNe Ia, $h=0.71\\pm0.08$ for the Tully-Fisher relation of spiral galaxies, $h=0.70\\pm0.08$ for the surface brightness fluctuations of galaxies, $h=0.72\\pm0.11$ for Type II supernovae and $h=0.82\\pm0.11$ for the fundamental plane method of elliptical galaxies. The combined result of HST Key Project suggested $h=0.72\\pm0.08$. While after the CMB observations are involved, the result is greatly improved \\citep{2008arXiv0803.0547K}, as shown in Fig. \\ref{combine}. It shows that the cross check and combination of different methods are very helpful to find the right answer and improve the accuracy.  Before concluding we would like to briefly comment the simple assumptions adopted in our work. We assume the intrinsic spectrum to be a single power law. In fact the spectral energy distribution of many AGNs can be well described by the so-called synchrotron self-Compton (SSC) model. In the SSC scenario, the VHE $\\gamma$-rays are produced through the inverse Compton (IC) scatterings between the electrons and synchrotron photons generated themselves. The VHE spectrum generally has an ``IC peak'' originates from the transition from Thomson regime to Klein-Nishima regime \\citep{1970RvMP...42..237B}. However, if the energy range is narrow, e.g., within a decade of energy, it can be described approximately by a power law.  The CIB model and its uncertainties are also too simplified. Because of the contamination of foreground radiation, it is usually difficult to get reliable CIB from measurements. The galaxy evolution models to predict CIB also have large uncertainties \\citep{2001ARA&A..39..249H,2006ApJ...648..774S}. In \\cite{2006ApJ...648..774S} the results between the ``baseline'' and ``fast'' evolution models differ by about $20\\% \\sim 40\\%$. While the results given by \\cite{2005AIPC..745...23P} differ from \\cite{2006ApJ...648..774S} by a factor of 2 at some wavelengths.  Anyway we think the present work is only a prototype of such studies, since the observation of VHE $\\gamma$ rays from extragalactic sources is achieved only in the recent years. This field is actually immature and at its early stage. With the next generation of space and ground-based instruments more extragalactic sources will be observed with high precision. Especially the space observatory Fermi\\footnote{See the homepage of Fermi, http://www-glast.stanford.edu} (energy range from MeV to hundred GeV) can explore $\\gamma$-ray sources to redshift $z\\sim 1$, due to an estimate of the ``$\\gamma$-ray horizon'' $\\log(z)\\sim 1-0.7\\log(E/1{\\rm GeV})$ \\citep{2007AIPC..921...24H}. The ``$\\gamma$-ray horizon'' is the redshift corresponding to absorption depth $\\tau\\approx 1$ for energy $E$. With larger sample of data, higher precision of spectra and higher redshift sources from Fermi we can even explore more cosmological parameters besides the Hubble constant, such as the cosmological component or the equation of state of dark energy. The development of ground-based instruments also aims to lower threshold energy and improve sensitivities. We anticipate the field of VHE $\\gamma$-ray will develop quickly in the near future. Our work shows that the VHE $\\gamma$-rays may become an more important field, not only to astrophysics but also to the cosmology."
    },
    "0809/0809.2829_arXiv.txt": {
        "abstract": "We investigate the X-ray origin in {FR\\,Is} using the multi-waveband high resolution data of eight FR I sources, which have very low Eddington ratios. We fit their multi-waveband spectrum using a coupled accretion-jet model. We find that X-ray emission in the source with the highest $L_{\\rm X}$ ($\\sim 1.8 \\times 10^{-4}L_{\\rm Edd}$) is from the advection-dominated accretion flow (ADAF). Four sources with moderate $L_{\\rm X}$ ($\\sim$ several $\\times 10^{-6}L_{\\rm Edd}$) are complicated. The X-ray emission of one {FR\\,I} is from the jet, and the other three is from the sum of the jet and ADAF. The X-ray emission in the three least luminous sources ($L_{\\rm X}\\leq1.0\\times 10^{-6}L_{\\rm Edd}$) is dominated by the jet. These results roughly support the predictions of Yuan and Cui (2005) where they predict that when the X-ray luminosity of the system is below a critical value, the X-radiation will not be dominated by the emission from the ADAF any longer, but by the jet. We also find that the accretion rates in four sources must be higher than the Bondi rates, which implies that other fuel supply (e.g., stellar winds) inside the Bondi radius should be important. ",
        "introduction": "Radio galaxies are usually classified as {FR\\,I} or {FR\\,II} sources depending on their radio morphology. {FR\\,I} radio galaxies (defined by edge-darkened radio structure) have lower radio power than {FR\\,II} galaxies (defined by edge-brightened radio structure due to compact jet terminating hot spots) (Fanaroff \\& Riley 1974). Different explanations of division of {FR\\,I} and {FR\\,II} radio galaxies invoke either the interaction of the jet with the ambient medium or the intrinsic nuclei properties of accretion and jet formation processes (e.g., Bicknell 1995; Reynolds \\textit{et al} 1996a; Hardcastle \\textit{et al} 2007). Accretion mode in low power {FR\\,Is} may be different from that in powerful {FR\\,IIs}. There is growing evidence to suggest that {FR\\,Is} type radio galaxy nuclei possess advection-dominated accretion flow (ADAF; or ``radiative inefficient accretion flows''; see Narayan \\& McMlintock 2008 for a review). In fact, we now have strong observational evidence that ADAFs may be powering various types of low-luminosity active galactic nuclei (LLAGNs; e.g., Yuan 2007; Ho 2008 for reviews), not only {FR\\,Is}. It was found that the hard X-ray photon indices of both X-ray binaries (XRBs) and AGNs are anti-correlated with the Eddington ratios when the Eddington ratio is less than a critical value, while they become positively correlated when the Eddington ratios higher than the critical value (e.g., Wu and Gu 2008). This results provide evidence for the accretion mode transition near the critical Eddington ratio.  One of the uncertainties in {FR\\, Is} is the respective contribution of ADAFs and jets at various wavebands. The least controversial nuclear emission are the radio and also optical emission, which is believed to be dominated by the jet (e.g., Wu and Cao 2005; Chiaberge \\textit{et al} 2006). It is still an open question for the origin of the core X-ray emission in {FR\\,Is} (and also LLAGNs). One possibility for the X-ray emission is dominated by the ADAF (e.g., Reynolds et al. 1996b; Quataert \\textit{et al} 1999; Merloni 2003). The other possibility is come from the jet (e.g., Falcke \\textit{et al} 2004; Markoff \\textit{et al} 2008), and recent observation on some individual extremely low luminosity sources supported this possibility (e.g., Garcia et al. 2005 for M\\,31; Pellegrini \\textit{et al} 2007 for NGC\\,821, etc.). Then an important question is \\emph{systematically} in what kind of condition the radiation from the jet or ADAF will be important in X-ray band. ",
        "conclusions": "Two possibilities exist for the X-ray origin in {FR\\,Is}: ADAF or jet. We use a coupled ADAF-jet model to fit the multi-waveband spectrum to try to investigate this problem. We find that the jet can well describe the radio and optical spectra for most {FR\\,Is} in our sample. The soft X-ray flux at $\\sim$1\\,keV of all {FR\\,Is} is roughly consistent with the predictions of the jet. This result indicates why a tight correlation is found among radio, optical and soft X-ray (Evans \\textit{et al} 2006; Balmaverde \\textit{et al} 2006). For 3C 346, the highest X-ray luminosity source in our sample which has $L_{\\rm X}=1.8 \\times 10^{-4} L_{\\rm Edd}$, its X-ray spectrum is dominated by the ADAF and the jet contribution is negligible. However, for the four sources with ``intermediate X-ray luminosities'' ({B2\\, 0755+37}, 3C 31, 3C 317, and B2\\,0055+30; $L_{\\rm X}=(2.4-5.2) \\times 10^{-6} L_{\\rm Edd}$), their X-ray origin is complicated. The X-ray spectrum of {B2\\,0755+37 is dominated by the jet, while for 3C 31 and B2\\,0055+30} spectra are dominated by the ADAF. For the other one (3C 317), the contributions of the ADAF and jet are comparable. For the three least luminous sources (3C 66B, 3C 449, and 3C 272.1) which have $L_{\\rm X}=(6.8\\times 10^{-8}-1\\times 10^{-6}) L_{\\rm Edd}$, their X-ray spectra are dominated by the jet. The X-ray emission of 3C 449 is also interpreted by the sum of a jet and an ADAF, which requires higher quality data to further constrain it. Our results are roughly consistent with the predictions of Yuan and Cui (2005). The ``intermediate luminosity'' here in our sample corresponds to the critical luminosity in Yuan and Cui (2005). However, the former is about 10 times higher than the latter. The value of the critical luminosity depends on the ratio of the mass loss rate in the jet to the mass accretion rate in the ADAF. This ratio is adopted in Yuan and Cui (2005) from fitting the data of a black hole X-ray binary---XTE\\,J1118+480. The current result indicates that the ratio in FR Is is about 10 times higher than in XTE\\,J1118+480. One possible reason could be that systematically the black holes in FR Is are spinning more rapidly. We note that the critical luminosity (or Eddington ratio) in this work is for the transition of \\emph{X-ray emission} dominated by the ADAF or by the jet, which is different from the critical value for the transition of the \\emph{jet power dominated systems} and a\\emph{ccretion power dominated systems} (Wu and Cao 2008 for more details and references therein).  The kinetic luminosity, $L_{\\rm kin}=\\Gamma_{j}(\\Gamma_{j}-1) \\dot{M}_{\\rm jet}c^{2}$, can be derived from our modeling results. We use $\\eta_{\\rm jet}=L_{\\rm kin}/\\dot{M}(10R_{S})c^{2}$ to describe the efficiency of the jet power converted from the accretion power, where $\\dot{M}(10R_{S})$ is mass accretion rate at $10R_{S}$. We find that $\\eta_{\\rm jet}=0.03-0.44$ for the sources in this sample (see Table 1), and most of them (six of eight) have $\\eta_{\\rm jet}$ significantly higher than 0.057, which is the largest available accretion energy at the innermost stable circular orbit (ISCO) for a non-rotating black hole. It implies that the black holes in these sources are spinning rapidly (so ISCO is smaller thus more accretion energy is available).  \\ack Wu Q W thanks the local organizing committee of the conference for partial financial support. This work is supported by the One-Hundred-Talent Program of China, the National Science Fund for Distinguished Young Scholars (grant 10325314), the NSFC (grants number 10773024, 10773020, 10703009 and 10633010), and a postdoctoral fellowship of the KASI."
    },
    "0809/0809.3969_arXiv.txt": {
        "abstract": "HH~262 is a group of emitting knots displaying an ''hour-glass'' morphology in the H$\\alpha$ and \\sii\\ lines, located 3.5 arcmin to the northeast of the young stellar object  L1551-IRS5, in Taurus. We present new results of the kinematics and physical conditions of HH~262 based on  Integral Field Spectroscopy covering a field of $\\sim$ 1.5 $\\times$ 3 arcmin$^2$, which includes all the bright knots in HH~262. These data show  complex kinematics and  significant variations  in physical conditions  over the mapped region of HH~262 on a  spatial scale of $\\leq$ 3 arcsec. A new result derived from the IFS data is  the weakness of the \\nii\\ emission (below detection limit in most of the mapped region of HH~262), including the brightest central knots. Our data  reinforce the association of HH~262  with the redshifted lobe of the evolved molecular outflow L1551-IRS5. The interaction of this outflow with a younger one, powered by L1551 NE, around the position of HH~262 could give rise  to the complex morphology and kinematics of HH~262. ",
        "introduction": "HH~262 is the designation of a group of emitting knots, surrounded by  more diffuse emission, located  $\\sim$ 3.5 arcmin to the  northeast of the source L1551-IRS5  in Taurus. The two brightest HH~262 condensations (named GH9 and GH10) were first reported by \\citet{Rod89} and \\citet{Gra90}. In the H$\\alpha$+\\nii\\ and \\sii\\ narrow-band images, HH~262 shows a rather chaotic, ''hour-glass''-shaped morphology that is quite different from the collimated structures found in most of the Herbig--Haro (HH) jets (\\citealt{Gar92}, \\citealt{Lop95}). Apparently, the morphology of HH~262 seems to  trace a wind-blown cavity from an embedded source located close to the GH10 knot. However, no  point source  suitable for powering the HH~262  emission has been currently reported at this location.  Since its discovery, the origin of HH~262, as well as its relationship  with other outflows in the L1551 star-forming region, has been  matter of discussion. \\citet{Rod89} first suggested that the GH9 and GH10 knots were associated  with the redshifted lobe of the  molecular outflow powered by the source L1551-IRS5. They proposed that GH9 and GH10 may be the redshifted counterpart of a chain of HH bright knots  (which includes the well-known HH objects 28 and 29) extending from L1551-IRS5 to the southwest up to $\\sim$ 5 arcmin, and suggested that  all these optical knots  traced the walls of a wind cavity evacuated by the CO outflow powered by L1551-IRS5.  Later on, \\citet{Gra92} obtained a long-slit spectrum across the GH9 and  GH10 knots. From the H$\\alpha$ line, they derived redshifted radial velocities covering a range extending over $\\sim$ 100~\\kms  and measure peak  velocities of +40~\\kms\\ and +90~\\kms\\ and average velocities of +20~\\kms and +45~\\kms in the GH9 and GH10 knots, respectively.  These results reinforced the suggestion that GH9 and GH10 were associated with the redshifted lobe of the CO outflow powered by L1551-IRS5.  The kinematics of HH~262, obtained by \\citet{Lop98} from proper motion measurements of the brightest knots, and from radial velocities derived from Fabry-Perot observations in the \\sii\\ $\\lambda$ 6716 \\AA\\ line, is also consistent with the interpretation of the HH~262 origin mentioned before. However, \\citet{Dev99} were uncertain whether the exciting source of HH~262 is L1551-IRS5 or  L1551 NE, another younger, embedded proto-binary source located $\\sim$ 150 arcsec northeast of IRS5 (\\citealt{Eme84}, \\citealt{Rod95}).  L1551 NE also drives a molecular CO outflow, so that it is likely that the outflows powered by these two sources are interacting.  Further observations from \\citet{Sto06} showed evidence of blueshifted molecular gas arising close to L1551 NE that  extends to the northeast of the source and encompasses HH~262. As pointed out by \\citet{Dev99} and \\citet{Mor06}, HH~262 could be related to either the L1551-IRS5 or L1551 NE outflows, its rather chaotic  morphology originating through the interaction  between  the  molecular outflows powered by these two sources (see diagram of Fig.\\ref{diagram1}).  More recent sub-mm continuum observations at 850 $\\mu$m by \\citet{Mor06}  detected emission around HH~262 that were interpreted as produced by warm dust entrained by the molecular outflows. HH~262 is thus a suitable place to search for  optical  signatures of the interactions among supersonic outflows. In addition, because of its morphology, HH~262 is better suited for performing Integral Field Spectroscopy (IFS) than long-slit observations when the kinematics and physical conditions of the  emission need to be analysed.  With the aim of advancing in the study of HH~262, we included this target within a program of Integral Field Spectroscopy of Herbig--Haro objects using the Potsdam Multi-Aperture Spectrophotometer (PMAS) in the wide-field IFU mode PPAK.  We present in this paper new results on the kinematics and physical conditions of HH~262 obtained from this IFS observing program. The paper is organized as follows. The observations and data  reduction are described in \\S\\ 2.  Results are given in \\S\\ 3: The integrated maps in \\S\\ 3.1, the channel maps in \\S 3.2, and the spectral atlas of HH~262, including discussions of the characteristic of the more relevant knots in \\S\\ 3.3. A global discussion and the main conclusions are given  \\S\\ 4.  \\begin{figure} \\centering \\includegraphics[width=\\hsize]{Fig1_diagram1.eps} \\caption{ Diagram of the HH~262 region showing the location of HHs and  YSOs. The contours correspond to the blueshifted and redshifted CO emission from the three molecular outflows identified in the region (two outflows in the NE-SW direction, associated with the YSOs L1551 IRS 5 and L1551 NE, and the east-west outflow; \\citealt{Mor06}). The straight lines are drawn along the outflow axes. \\label{diagram1}} \\end{figure}      \\section[]{Observations and Data Reduction}  Observations of HH~262 were made on 22 November 2004 with the 3.5-m  telescope of the Calar Alto Observatory (CAHA). Data were acquired with the  Integral Field Instrument Potsdam Multi-Aperture Spectrophotometer PMAS \\citep{Rot05} using the PPAK configuration that has 331 science fibres, covering an hexagonal FOV of 74$\\times$65 arcsec$^2$ with a spatial sampling of 2.7 arcsec per fibre, and 36 additional fibres to sample the sky (see Fig.\\ 5 in \\citealt{Ke06}). The I1200 grating was used, giving an effective sampling  of 0.3 \\AA\\ pix$^{-1}$ ($\\sim15$  km~s$^{-1}$ for H$\\alpha$) and covering the wavelength range $\\sim6500$--7000 \\AA, thus including characteristic HH emission lines in this wavelength range (H$\\alpha$, \\nii\\ $\\lambda\\lambda$6548, 6584 \\AA\\  and \\sii\\ $\\lambda\\lambda$6716, 6731 \\AA). The spectral resolution (\\ie\\ instrumental profile) is $\\sim2$ \\AA\\  FWHM ($\\sim90$ km~s$^{-1}$) and the accuracy in the determination of the position of the line centroid is $\\sim0.2$ \\AA\\  ($\\sim10$ km~s$^{-1}$ for the  strong observed emission lines). Eight overlapped pointings of 1800-s exposure time each were observed to obtain a mosaic of $\\sim$ 1.5 $\\times$ 3 arcmin$^2$  to cover almost the entire emission from HH~262.  Note, however, that  this mosaic does not cover the extended, arc--shaped emission known  as HH~262E, which is most probably related  to HH~262.  In Table 1 we list the centre positions of each  pointing and the corresponding HH~262 knot identification following the nomenclature of \\citet{Gar92},  extended by \\citet{Lop95}.   Data reduction was performed using a preliminary version of the {\\small R3D} software \\citep{Sa06}, in combination with {\\small IRAF}% \\footnote{{\\small IRAF} is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} and the {\\small Euro3D} packages \\citep{Sa04}. The reduction consists of the standard steps for fibre-based integral field spectroscopy. A master bias frame was created by averaging all the bias frames observed during the night and subtracted from the science frames. The location of the spectra in the CCD was determined using a continuum-illuminated exposure taken before the science exposures. Each spectrum was extracted from the science frames by co-adding the flux  within an aperture of 5 pixels along the cross-dispersion axis for each pixel in the dispersion axis, and stored in a row-stacked-spectrum (RSS) file \\citep{Sa04}. At the date of the observations we lacked of a proper calibration unit for PPAK. Therefore the wavelength calibration was performed iteratively, following all the steps described in \\citet{Lop08} for HHL~73,  another target of our HH  IFS survey.  The final wavelength solution was set by using the sky emission lines found in the observed wavelength range. The  accuracy achieved for the wavelength calibration was better than  $\\sim0.1$  \\AA\\ ($\\sim5$ km~s$^{-1}$). Observations of a standard star were used to perform a relative flux calibration. A final datacube  containing the 2D spatial plus the spectral information of HH~262 was then created from the 3D data by using {\\small Euro3D} tasks  to interpolate the data spatially until reaching a final grid of 2.16 arcsec of spatial sampling and with a spectral sampling of 0.3 \\AA\\ . Further manipulation of this datacube, devoted to obtain  integrated emission line maps, channel maps and position--velocity maps, were made using several common-users tasks of {\\small STARLINK}, {\\small IRAF} and {\\small GILDAS} astronomical packages and {\\small IDA} \\citep{Ga02}, an specific {\\small IDL} software to analyse 3D data.      \\begin{table} \\begin{minipage}{130mm} \\caption{Positions of the single pointings of the HH~262 mosaic. \\label {tjcmt}} \\begin{tabular}{@{}lcc@{}} \\hline \\multicolumn{2}{c}{Position} &knot\\\\ \\hline 04$^{\\rm h}$31$^{\\rm m}$58$\\fs7$&+18$\\degr$12$\\arcmin$07$\\arcsec$& GH9\\\\ 04$^{\\rm h}$32$^{\\rm m}$01$\\fs1$&+18$\\degr$12$\\arcmin$35$\\arcsec$& K18\\\\ 04$^{\\rm h}$32$^{\\rm m}$00$\\fs5$&+18$\\degr$12$\\arcmin$05$\\arcsec$& K19\\\\ 04$^{\\rm h}$32$^{\\rm m}$00$\\fs9$&+18$\\degr$11$\\arcmin$53$\\arcsec$& K20\\\\ 04$^{\\rm h}$32$^{\\rm m}$00$\\fs1$&+18$\\degr$11$\\arcmin$33$\\arcsec$& GH10W\\\\ 04$^{\\rm h}$32$^{\\rm m}$01$\\fs1$&+18$\\degr$11$\\arcmin$24$\\arcsec$& GH10E\\\\ 04$^{\\rm h}$32$^{\\rm m}$00$\\fs0$&+18$\\degr$11$\\arcmin$17$\\arcsec$& K26\\\\ 04$^{\\rm h}$31$^{\\rm m}$59$\\fs7$&+18$\\degr$10$\\arcmin$30$\\arcsec$& K22\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "We have carried out a study of the kinematics and physical conditions of HH~262 based on IFS observations. From the analysis of these data we find:  \\begin{itemize}  \\item HH~262 shows a non-collimated, ''hour-glass'' morphology in the H$\\alpha$ and \\sii\\ lines. In contrast, the emission from the \\nii\\ lines is  below detection limit except in the northwestern knot GH9. This work therefore presents the first narrow-band \\nii\\ image currently obtained of HH~262 to show its surprisingly different morphology as compared with the morpholgy found in the other characteristic HH emission lines. \\item In all the HH~262 knots, radial velocities derived from the line centroids appear redshifted, relative to the cloud velocity. The highest velocities and FWHM values are found in the central  knots. \\item A more detailed picture of the kinematics was derived from the channel maps, which shows  that the emission spreads over a wide range of velocities,  the range slightly varying for the different knots. Furthermore, inspection of the HH~262 extracted spectra shows that there are  broad,  asymmetric line profiles at the southern (K22) and central knots (GH10E-W) that are  double-peaked at several positions, which would suggest contribution from two velocity components. The line profiles appear  more symmetric and single-peaked away from the centre, in the northern (K18--K20)  and northwestern (GH9) knots. \\item Concerning the physical conditions in HH~262, we found variations on a spatial scale of $\\le$3 arcsec (\\ie\\ of the order of our spatial resolution). As a general trend, the electron densities, derived from the \\sii\\ line ratios, are  low. In particular, only an upper limit could be set for the brightest, central HH~262 knots.  Because $n_\\rmn{e}$ is referred to the density of the ionized gas, we propose that the only a small fraction of the atomic gas  in HH~262 is  ionized, the remainder being mostly neutral gas. It is also plausible that the total density of the  gas  is lower in this region because the evolved molecular  outflow has had enough time to clear the medium. Over the entire velocity range of the emission, the gas excitation, derived from the \\sii/H$\\alpha$ line ratios, increases as we move towards the western region of knot GH10W and to knot GH9  \\ie\\ moving towards the edges of HH~262 which are  closer to the higher extinction region of the cloud. The general trend found is thus that the densest knots have the highest gas excitation  and the lower centroid velocities. \\end{itemize}  All these  characteristics of HH~262 derived from the IFS data reinforce the causal relationship of the  optical line emission with the two molecular  outflows (one powered by L1551-IRS5 and the other, by L1551 NE) found in the region were HH~262 is located. In fact, is widely accepted the association between HH~262 and the redshifted lobe of the L1551-IRS5 molecular outflow, first proposed by \\citet{Rod89} and confirmed later on by the Fabry-Perot observations \\citep{Lop98}.  Detailed PV maps of the L1551-IRS5 outflow performed by \\citet{Sto06} from high-sensitivity CO observations show a pattern that is characteristic of an evolved molecular outflow, according to current model simulations of the entrainment outflow mechanisms  (\\citealt{Vel98}, \\citealt{Ric00}). In the \\citet{Sto06} PV CO maps,  the HH objects are found associated with the molecular emission at the highest velocities, highlighting the interaction regions of the shocks at the head of the shock interface between the  supersonic gas and the environment. In contrast, the CO emission at lower velocities appears as a limb-brightened shell-like emission.  Thus, the result derived from the \\citet{Sto06} work support the proposed  scenario, in which the HH~262 emission traces the walls of a conical  cavity evacuated by an evolved molecular outflow (L1551-IRS5). In such a scenario, one would expect that the HH~262 radial velocities to decrease (due to projection effects) as we move away from the axis outflow (\\ie\\ from the central GH10 knots, where we found the highest integrated velocities, to the other knots). It  is also expected to find double-peaked line profiles along the outflow axis (corresponding to material  moving in the forward and rear faces of the cone) and a single-peaked profile away from the axis. Such a trend is observed in HH~262 from our IFS data, where the lines of the central knots appear with a complex, asymmetric line profiles, being double-peaked for most positions, while the line profiles of the knots farther away from the outflow axis (\\eg\\ K18, K19 or GH9) appear more symmetric and single-peaked.  In addition to the L1551-IRS5 outflow, the blueshifted lobe of another younger, dimmer CO outflow, powered by L1551 NE, also encompases HH~262 \\citep{Mor06}. Thus, HH~262 is located in a region where two molecular outflows are interacting (see  Fig.\\ref{diagram1}) and it is expected that this interaction   gives signatures on the jet turbulence. This interaction could in part be  responsible for the complex velocity pattern and the wide line profiles found in the optical emission lines.  In a turbulent jet scenario,  the highest FWHM values are expected  to be found close to the axis, where there are contributions from gas within a wider range of velocities because we intersect the entire jet beam, while FWHM values would be lower at the edge of the jet, away from the axis. This is the trend found in the IFS HH~262 data for the FWHMs , where the highest values correspond to the central GH10 knots and decrease moving towards the edges of the emission.  The association of HH~262 with an evolved molecular outflow near the edge of a dark cloud would in part justify the low density values derived for most of the knots, as the outflow has cleared the medium.  Note however that the   density of the gas is probably higher than the  density derived from the \\sii\\ line ratio, which is an indicator of only the density of the ionized fraction of the gas rather than of the total density (\\citealt{San07}).  Finally, let us discuss the new finding  presented here concerning  the low intensity of the \\nii\\ emission in HH~262. Such a behaviour has  been found in slow-velocity shocks as those modelled by  \\citet{Har87}. In particular, the \\nii/H$\\alpha$ and \\sii/H$\\alpha$ line ratios derived for the knot GH9 are fully compatible with those derived in shock models with $V_\\rmn{s}$ $\\simeq$ 30--40~\\kms\\ occurring in a low density, partially ionized medium. Lower \\nii/H$\\alpha$ line ratios ($\\simeq$ 0.01) are found in models with lower $V_\\rmn{s}$ ($\\le$ 20 \\kms\\ ), thus being compatible, within the errors, with the marginal detection of \\nii\\ in the northern knots and with the  line ratios derived in those knots. The HH~262 line ratios are also compatible with those predicted by the shock models of \\citet{Har94} for low magnetic field strength and low ionization fraction values. Note however that the  morphology of HH~262 is too complex and rather different from the morphologies found in most known HH jets. Thus, it should not be expected that the radiative planar shock models of \\citet{Har87} and \\citet{Har94} succeeded in reproducing all the HH~262 properties derived from the current IFS data. Further IFS observations covering a wavelength range blueward of \\nii\\ + H$\\alpha$ are needed to search for emission in the \\nin\\ $\\lambda$ 5200 \\AA\\ line whose strength could justify the weakness of the \\nii\\ emission found in HH~262 provided  the nitrogen abundance were not far from the solar value."
    },
    "0809/0809.4122_arXiv.txt": {
        "abstract": "This paper makes the first systematic attempt to determine using perturbation theory the positions of images by gravitational lensing due to arbitrary number of coplanar masses without any symmetry on a plane, as a function of lens and source parameters. We present a method of Taylor-series expansion to solve the lens equation under a small mass-ratio approximation. First, we investigate perturbative structures of a single-complex-variable polynomial, which has been commonly used. Perturbative roots are found. Some roots represent positions of lensed images, while the others are unphysical because they do not satisfy the lens equation. This is consistent with a fact that the degree of the polynomial, namely the number of zeros, exceeds the maximum number of lensed images if N=3 (or more). The theorem never tells which roots are physical (or unphysical). In this paper, unphysical ones are identified. Secondly, to avoid unphysical roots, we re-examine the lens equation. The advantage of our method is that it allows a systematic iterative analysis. We determine image positions for binary lens systems up to the third order in mass ratios and for arbitrary N point masses up to the second order. This clarifies the dependence on parameters. Thirdly, the number of the images that admit a small mass-ratio limit is less than the maximum number. It is suggested that positions of extra images could not be expressed as Maclaurin series in mass ratios. Magnifications are finally discussed. ",
        "introduction": "Gravitational lensing has become one of important subjects in modern astronomy and cosmology (e.g., Schneider 2006, Weinberg 2008). It has many applications as gravitational telescopes in various fields ranging from extra-solar planets to dark matter and dark energy at cosmological scales. This paper focuses on gravitational lensing due to a N-point mass system. Indeed it is a challenging problem to express the image positions as functions of lens and source parameters. There are several motivations. One is that gravitational lensing offers a tool of discoveries and measurements of planetary systems (Schneider and Weiss 1986, Mao and Paczynski 1991, Gould and Loeb 1992, Bond et al. 2004, Beaulieu et al. 2006), compact stars, or a cluster of dark objects, which are difficult to probe with other methods. Gaudi et al. (2008) have recently found an analogy of the Sun-Jupiter-Saturn system by lensing. Another motivation is theoretically oriented. One may be tempted to pursue a transit between a particle method and a fluid (mean field) one. For microlensing studies, particle methods are employed, because the systems consist of stars, planets or MACHOs. In cosmological lensing, on the other hand, light propagation is considered for the gravitational field produced by inhomogeneities of cosmic fluids, say galaxies or large scale structures of the universe (e.g., Refregier 2003 for a review). It seems natural, though no explicit proof has been given, that observed quantities computed by continuum fluid methods will agree with those by discrete particle ones in the limit $N \\to \\infty$, at least on average, where $N$ is the number of particles.  Related with the problems mentioned above, we should note an astronomically important effect caused by the finiteness of $N$. For most of cosmological gravitational lenses (both of strong and weak ones), a continuum approximation can be safe and has worked well. There exists an exceptional case, however, for which discreteness becomes important. One example is a quasar microlensing due to a point-like lens object, which is possibly a star in a host galaxy (for an extensive review, Wambsganss 2006). A galaxy consists of very large number $N$ particles, and light rays from an object at cosmological distance may have a chance to pass very near one of the point masses. As a consequence of finite-$N$ effect in large $N$ point lenses, anomalous changes in the light curve are observed. For such a quasar microlensing, hybrid approaches are usually employed, where particles are located in a smooth gravitational field representing a host galaxy. It is thus likely that $N$ point-mass approach will be useful also when we study such a finite-$N$ effect at a certain transit stage between a particles system and a smooth one. Along this course,  An (2007) investigated a $N$ point lens model, which represents a very special configuration that every point masses are located on regular grid points.  For a $N$ point-mass lens at a general configuration, very few things are known in spite of many efforts. Among known ones is the maximum number of images lensed by $N$ point masses. After direct calculations by Witt (1990) and Mao, Petters and Witt (1997), a careful study by Rhie (2001 for N=4, 2003 for general N) revealed that it is possible to obtain the maximum number of images as $5(N-1)$. This theorem for polynomials has been extended to a more general case including rational functions by Khavinson and Neumann (2006). (See Khavinson and Neumann 2008 for an elegant review on a connection between the gravitational lens theory and the algebra, especially the fundamental theorem of algebra, and its extension to rational functions).  \\noindent {\\bf Theorem} (Khavinson and Neumann 2006): \\\\ Let $r(z)=p(z)/q(z)$, where $p$ and $q$ are relatively prime polynomials in $z$, and let $n$ be the degree of $r$. If $n>1$, then the number of zeros for $r(z)-z^* \\leq 5(n-1)$. Here, $z$ and $z^*$ denote a complex number and its complex conjugate, respectively.  Furthermore, Bayer, Dyer and Giang (2006) showed that in a configuration of point masses, replacing one of the point deflectors by a spherically symmetric distributed mass only introduces one extra image. Hence they found that the maximum number of images due to $N$ distributed lensing objects located on a plane is $6 (N-1)+1$.  Global properties such as lower bounds on the number of images are also discussed in Petters, Levine and Wambsganss (2001) and references therein.  In spite of many efforts on $N$ lensing objects, functions for image positions are still unknown even for $N$ point-mass lenses in a general configuration under the thin lens approximation. Hence it is a challenging problem to express the image positions as functions of lens and source locations. Once such an expression is known, one can immediately obtain magnifications via computing the Jacobian of the lens mapping (Schneider et al. 1992).  Only for a very few cases such as a single point mass and a singular isothermal ellipsoid, the lens equation can be solved by hand and image positions are known, because the lens equation becomes a quadratic or fourth-order one (For a singular isothermal ellipsoid, Asada et al. 2003). For the binary lens system, the lens equation has the degree of five in a complex variable (Witt 1990). It has the same degree also in a real variable (Asada 2002a, Asada et al. 2004). This improvement is not trivial because a complex variable brings two degrees of freedom. This single-real-variable polynomial has advantages. For instance, the number of real roots (with vanishing imaginary parts) corresponds to that of lensed images. The analytic expression of the caustic, where the number of images changes, is obtained by the fifth-order polynomial (Asada et al. 2002c). Galois showed, however,  that the fifth-order and higher polynomials cannot be solved algebraically (van der Waerden 1966). Hence, no formula for the quintic equation is known. For this reason, some numerical implementation is required to find out image positions (and magnifications) for the binary gravitational lens for a general position of the source. Only for special cases of the source at a symmetric location such as on-axis sources, the lens equation can be solved by hand and image positions are thus known (Schneider and Weiss 1986). For a weak field region, some perturbative solutions for the binary lens have been found (Bozza 1999, Asada 2002b), for instance in order to discuss astrometric lensing, which is caused by the image centroid shifts (for a single mass, Miyamoto and Yoshii 1995, Walker 1995; for a binary lens, Safizadeh et al. 1999, Jeong et al. 1999, Asada 2002b).  If the number of point masses $N$ is larger than two, the basic equation is much more highly non-linear so that the lens equation can be solved only by numerical methods. As a result, observational properties such as magnifications and image separations have been investigated so far numerically for $N$ point-mass lenses. This makes it difficult to investigate the dependence of observational quantities on lens parameters.  This paper is the first attempt to seek an analytic expression of image positions without assuming any special symmetry. For this purpose, we shall present a method of Taylor-series expansion to solve the lens equation for $N$ point-mass lens systems. Our method allows a systematic iterative analysis as shown later.  Under three assumptions of weak gravitational fields, thin lenses and small deflection angles, gravitational lensing is usually described as a mapping from the lens plane onto the source plane (Schneider et al. 1992). Bourassa and Kantowski (1973, 1975) introduced a complex notation to describe gravitational lensing. Their notation was exclusively used to describe lenses with elliptical or spheroidal symmetry (Borgeest 1983, Bray 1984, Schramm 1990). For $N$ point lenses, Witt (1990) succeeded in recasting the lens equation into a single-complex-variable polynomial. This is in an elegant form and thus has been often used in investigations of point-mass lenses. An advantage in the single-complex-variable formulation is that we can use some mathematical tools applicable to complex-analytic functions, especially polynomials (Witt 1993, Witt and Petters 1993, Witt and Mao 1995). One tool is the fundamental theorem of algebra: Every non-constant single-variable polynomial with complex coefficients has at least one complex root. This is also stated as: every non-zero single-variable polynomial, with complex coefficients, has exactly as many complex roots as its degree, if each root is counted as many times as its multiplicity. On the other hand, in the original form of the lens equation, one can hardly count up the number of images because of non-linearly coupled properties. This theorem, therefore, raises a problem in gravitational lensing. The single-variable polynomial due to $N$ point lenses has the degree of $N^2+1$, though the maximum number of images is $5(N-1)$. This means that unphysical roots are included in the polynomial (for detailed discussions on the disappearance and appearance of images near fold and cusp caustics for general lens systems, see also Petters, Levine and Wambsganss (2001) and references therein). First, we thus investigate explicitly behaviors of roots for the polynomial lens equation from the viewpoint of perturbations. We shall identify unphysical roots. Secondly, we shall re-examine the lens equation, so that the appearance of unphysical roots can be avoided.  This paper is organised as follows. In Section 2, the complex description of gravitational lensing is briefly summarised. The lens equation is embedded into a single-complex-variable polynomial in Section 3. Perturbative roots for the complex polynomial are presented for binary and triple systems in sections 4 and 5, respectively. They are extended to a case of $N$ point lenses in section 6. In section 7, we re-examine the lens equation in a dual-complex-variables formalism and its perturbation scheme for a binary lens for its simplicity. The perturbation scheme is extended to a $N$ point lens system in section 8. Section 9 is devoted to the conclusion. ",
        "conclusions": "Under a small mass-ratio approximation, this paper developed a perturbation theory of N coplanar (in the thin lens approximation) point-mass gravitational lens systems without symmetries on a plane. The system can be separated into a single mass lens as a background and its perturbation due to the remaining point masses.  First, we investigated perturbative structures of the single-complex-variable polynomial, into which the lens equation is embedded. Some of zeroth-order roots of the polynomial do not satisfy the lens equation and thus are unphysical. This appearance of correct but unphysical roots is consistent with the earlier work on a theorem on the maximum number of lensed images (Rhie 2001, 2003). However, the theorem  never tells which roots are physical (or unphysical). What we did is that unphysical roots are identified.  Next, we re-examined the lens equation in the dual-complex-variables formalism to avoid inclusions of unphysical roots. We presented an explicit form of perturbed image positions as a function of source and lens positions. As a key tool for perturbative computations, Eq. ($\\ref{zz*root}$) was also found. For readers' convenience, the perturbative roots are listed in Table $\\ref{list}$. If one wishes to go to higher orders, our method will enable one to easily use computer algebra softwares such as MAPLE and MATHEMATICA. This is because it requires simpler algebra (only the four basic operations of arithmetic), compared with vector forms which need extra operations such as inner and outer products.  There are numerous possible applications along the course of the perturbation theory of N point-mass gravitational lens systems. For instance, it will be interesting to study lensing properties such as magnifications by using the functional form of image positions. Furthermore, the validity of the present result may be limited in the weak field regions. It is important also to extend the perturbation theory to a domain near the strong field.  Our method considers only the images which exist in the small mass limit as $\\nu_i \\to 0$. The number of the images that admit the small mass-ratio limit is less than the maximum number. This suggests that the other images do not have the small mass limit. Therefore, it is conjectured that positions of the extra images could not be expressed as Maclaurin series in mass ratios. This may be implied also by previous works. For instance, the appearance of the maximum number of images for a binary lens requires a finite mass ratio and the caustic crossing (Schneider and Weiss 1986). Regarding this point, further studies will be needed to determine positions of all the images with the maximum number as a function of lens and source parameters."
    },
    "0809/0809.0317_arXiv.txt": {
        "abstract": "{Evolutionary trends in the surface abundances of heavier elements have recently been identified in the globular cluster NGC~6397 ([Fe/H]$=-2$), indicating the operation of atomic diffusion in these stars. Such trends constitute important constraints for the extent to which diffusion modifies the internal structure and surface abundances of solar-type, metal-poor stars.} {We perform an independent check of the reality and size of abundance variations within this metal-poor globular cluster.} {Observational data covering a large stellar sample, located between the cluster turn-off point and the base of the red giant branch, are homogeneously analysed. The spectroscopic data were obtained with the medium-high resolution spectrograph FLAMES/GIRAFFE on VLT-UT2 ($R\\sim27\\,000$). We derive independent effective-temperature scales from profile fitting of Balmer lines and by applying colour-$T_{\\rm eff}$ calibrations to Str\\\"{o}mgren $uvby$ and broad-band $BVI$ photometry. An automated spectral analysis code is used together with a grid of MARCS model atmospheres to derive stellar surface abundances of Mg, Ca, Ti, and Fe.} {We identify systematically higher iron abundances for more evolved stars. The turn-off point stars are found to have $0.13$\\,dex lower surface abundances of iron compared to the coolest, most evolved stars in our sample. There is a strong indication of a similar trend in magnesium,  whereas calcium and titanium abundances are more homogeneous. Within reasonable error limits, the obtained abundance trends are in agreement with the predictions of stellar structure models including diffusive processes (sedimentation, levitation), if additional turbulent mixing below the outer convection zone is included.} {} ",
        "introduction": "A globular cluster (GC) constitutes, in many respects, a homogeneous stellar population and is therefore an excellent laboratory for testing and constraining models of stellar structure and evolution. Chemical abundances inferred from observations of GCs can give important clues to the physical processes at work in individual stars. Recently a discovery was made pointing to the existence of systematic differences in surface abundances between stars in different evolutionary stages in a metal-poor GC. \\citet{Korn07} (hereafter Paper I) observe four samples of stars located between the main sequence (MS) turn-off (TO) point and the red giant branch (RGB) of NGC 6397, using FLAMES/UVES on the VLT, and found significant variations in abundances with effective temperature and surface gravity. In particular, they conclude iron to be under-abundant by $0.16\\pm0.05$\\,dex in the atmospheres of the TO stars compared to the RGB stars in their sample. Proposed as an explanation to this phenomenon is a continuous, long-term depletion of iron and other heavy elements from the surface layers, at work during the stars' life-time on the MS. The mechanism responsible for the depletion is generally referred to as atomic diffusion -- an umbrella term accounting for several diffusive processes. High ages and thin convective envelopes are factors that make metal-poor TO stars particularly prone to atomic diffusion. Once a star leaves the MS and begins to develop into the giant stage, its outer convection zone gradually reaches deeper layers, and elements that had previously been drained from the surface layers are mixed up again, eventually restoring the initial chemical composition (with fragile elements such as Li as notable exceptions). This is predicted to result in a gradual rise in the surface abundances of the depleted elements as the star evolves along the subgiant and red giant branches (hereafter, we write SGB for subgiant branch). GC stars presumably share the same original composition of elements such as iron, titanium, and calcium, which in combination with their high ages and low metal content make them suitable test cases for stellar structure models including atomic diffusion.\\\\ The abundance trend for iron presented in Paper I is directly contradictory to result by \\citet{Gratton01}, who by studying TO vs.\\ base-RGB (bRGB) stars conclude that there are no significant difference in iron abundance between the groups (magnesium is, however, found to be 0.15\\,dex less abundant in the TO stars). The cause for the differing result is twofold: lower effective temperatures are found for the TO stars in Paper I and the analysis include stars on the RGB.\\\\ Here we present a follow-up analysis of spectroscopic data from FLAMES/GIRAFFE collected simultaneously to the observations presented in Paper I. A large sample of stars, $\\approx 100$, covering the range from the TO to the bRGB is homogeneously analysed, with the aim of further constraining possible variations in surface abundance in the cluster. In particular, we investigate whether the results presented in Paper I are robust for a larger sample of stars, analysed with an independent set of tools. We start by describing the observations and the data-reduction procedure in Sect.\\ 2. Sect.\\ 3 is devoted to the determination of fundamental stellar parameters and a description of the spectrum analysis code. In Sect.\\ 4 we present and discuss our abundance results. Sect.\\ 5 compares predictions from stellar-structure models including atomic diffusion to the abundances we have inferred from observations. In Sect.\\ 6 we summarise our conclusions.\\\\ ",
        "conclusions": "We have presented a homogeneous analysis of medium-high resolution spectroscopic data covering a large sample of stars, located between the TO point and bRGB of the globular cluster NGC 6397. The obtained iron abundances show a significant trend with evolutionary stage. The magnesium abundances also indicate the presence of a trend of similar size, but the high sensitivity of this element to stellar parameters combined with the possibility of differential NLTE corrections prevent an unambiguous detection. The calcium and titanium abundances can be reconciled with a flat behaviour. The difference in iron abundances between the stars at $T_{\\rm eff,H\\alpha}=5\\,450\\,$K and stars at $T_{\\rm eff,H\\alpha}=6\\,250\\,$K amounts to $0.13$\\,dex in $\\log{\\epsilon_{\\rm Fe}}$, a trend that is robust to realistic errors in stellar parameters. Raising the effective temperatures of the hottest stars by $\\sim 350$\\,K would remove the trend in iron abundance. Considering that the total range in effective temperature between the considered points is $800$\\,K, this would imply an error in our $T_{\\rm eff}$-scale of $43\\,$\\%, which we regard as unlikely. Alternatively, the same effect may be achieved by lowering the effective temperatures of the coolest stars by $\\sim200$\\,K, correspoding to a systematic error of almost 25\\,\\%. Even that is an improbably large correction, given the good agreement between spectroscopy, through profile-fitting of H$\\alpha$ and the different photometric calibrations that were exploited. The excitation equilibrium of Fe\\,I points to cooler $T_{\\rm\\,eff}$-values overall, but on the relative scale the agreement is satisfactory. Overall, the results of this study independently support the conclusions of \\citet{Korn07}, indicating that atomic diffusion significantly affects the surface abundances of metal-poor stars near the main turn-off point.\\\\ The abundance analysis is based on traditional LTE analysis using 1D, hydrostatic model atmospheres. Although it seems that corrections to the obtained abundances of Fe and Mg by 1D, NLTE modelling have no large differential impact in comparison to the sizes of the found trends, we cannot rule out that our results are artifacts from insufficient modelling techniques. More sophisticated models, implementing NLTE line formation in 3D, dynamical model atmospheres may provide a more definite answer. Some explorations with LTE line formation using 3D models were performed in Paper I, indicating that the results are robust in this respect. \\\\ The obtained abundance trends with effective temperature are, for magnesium, calcium, and iron, well reproduced by stellar-structure models including atomic diffusion and turbulent mixing. Especially, the good relative agreement seen at the cooler half of the SGB, where additional turbulent transport is irrelevant, indeed lends support to the notion that atomic diffusion shapes the surface abundances of unevolved metal-poor stars."
    },
    "0809/0809.2638_arXiv.txt": {
        "abstract": "We have searched for pulsation of the anomalous X-ray pulsar (AXP) 4U 0142+61 in the $K^\\prime$ band ($\\lambda_{\\rm eff} = 2.11$ $\\mu$m) using the fast-readout mode of IRCS at the Subaru 8.2-m telescope. We found no significant signal at the pulse frequency expected by the precise ephemeris obtained by the X-ray monitoring observation with RXTE. Nonetheless, we obtained a best upper limit of 17\\% (90\\% C.L.) for the root-mean-square pulse fraction in the $K^\\prime$ band. Combined with $i^\\prime$ band pulsation \\citep{Dhillon+_2005}, the slope of the pulsed component ($F_\\nu \\propto \\nu^\\alpha$) was constrained to $\\alpha > -0.87$ (90\\% C.L.) for an interstellar extinction of $A_{V} = 3.5$. ",
        "introduction": "Anomalous X-ray pulsars (AXPs; see \\citet{Woods_Thompson_2006} for a recent review) have been recognized as being enigmatic X-ray pulsars, for which none of characteristics of binary companions have been detected \\citep{Mereghetti_Stella_1995}. Therefore, the observed X-ray luminosity of AXPs ($L_{\\rm X} \\sim 10^{34} - 10^{36}$ erg s$^{-1}$) cannot be explained by the accretion of matter from companions, as is the case for binary X-ray pulsars. The spin periods of AXPs ($P$) are concentrated in a narrow range (2$-$12 s \\footnote{Recently, the period range was extended down to 2 s by the discovery of the radio-emitting AXP 1E 1547.0$-$5408 \\citep{Camilo+_2007}.}) and the spin-down rates ($\\dot{P}$) are distributed in the range $5 \\times 10^{-13}$ -- $1 \\times 10^{-10}$ s s$^{-1}$. The rate of the spin-down energy loss of neutron stars ($\\dot{E} = - 4 \\pi^2 I \\dot{P}/P^3 \\sim 10^{32.6}$ erg s$^{-1}$, where $I \\simeq 10^{45}$ g cm$^2$ is the moment of inertia of a neutron star) is too small to maintain the X-ray luminosity, suggesting that they are not rotation-powered pulsars.  Two competing models have been proposed for the energy source: accretion from a fossil disk and decay of ultra-strong magnetic fields. In the former scenario, the fossil disks are thought to be produced through a fall-back of matter after a supernova explosion \\citep{Chatterjee_Hernquist_Narayan_2000, Chatterjee_Hernquist_2000, Alpar_2001} or as a remnant of the common-envelope phase in the evolution of a binary \\citep{van_Paradijs_Taam_van_den_Heuvel_1995, Ghosh_Angelini_White_1997}. The advantage of the fossil disk model is that it naturally explains the narrow range of the concentration of the spin periods through the propeller phase in the spin evolution \\citep{Chatterjee_Hernquist_Narayan_2000, Chatterjee_Hernquist_2000, Alpar_2001}. In the latter scenario AXPs are magnetars, which are strongly magnetized neutron stars with emissions powered by the dissipation of magnetic energy \\citep{Thompson_Duncan_1995, Thompson_Duncan_1996}. Assuming that their spins are braked by magnetic dipole radiation, a strong magnetic field of $10^{14}$ -- $10^{15}$ G at their surface is implied. AXPs are similar to soft gamma repeaters (SGRs), which are thought to be magnetars (e.g. see \\cite{Thompson_Duncan_1995}). The observation of SGR-like bursts in AXPs suggests the unification of AXPs and SGRs by the magnetar scenario \\citep{Gavriil_Kaspi_Woods_2002}.  4U 0142+61 is a typical object among AXPs. It is the nearest and least absorbed object among AXPs and SGRs and is hence the most intensively observed of these objects in multi-wavelength studies. The optical counterpart of this object was discovered by \\citet{Hulleman_van_Kerkwijk_Kulkarni_2000}. \\citet{Kern_Martin_2002} discovered the optical pulsation. The pulsed fraction of $PF_{\\rm opt} = 27^{+8}_{-6}$\\% was found to be larger than that of soft X-rays, suggesting that the optical pulsation cannot be explained by thermal reprocessing of a disk illuminated by X-rays. \\citet{Dhillon+_2005} also measured the optical pulsation in a $i^\\prime$ band and obtained a consistent pulsed fraction of $PF_{i^\\prime} = 29 \\pm 8$\\%.  We showed that the near-infrared/optical spectrum consists of two components, near-infrared excess and optical pulsating components, by first near-simultaneous multi-band photometry \\citep{Morii+_2005, Tanaka+_2008}. \\citet{Wang_Chakrabarty_Kaplan_2006} discovered the mid-infrared emission with Spitzer, and showed that the near-infrared excess component extends to the mid-infrared region. They proposed that the infrared excess component can be modeled by a dust disk around a neutron star, where the spectrum is similar to that of the dust disk of a proto-planetary system. Their model is different from the original fossil disk model, in that the disk does not power the X-ray emission. In their model, the flux of the $K$ band is mainly composed of the dust disk component. Therefore, the pulsed flux of the $K$ band is expected to be small. \\citet{Durant_van_Kerkwijk_2006b} collected data spanning seven years and showed that the near-infrared/optical flux and the spectral shape are variable. The correlations between the X-ray total/pulsed fluxes and the near-infrared/optical fluxes were uncertain, although these correlations are expected for a fossil disk. Therefore, they suggested that the near-infrared flux is composed of components other than the disk component. As described above, the origin of the infrared emission is not yet clear.  Some authors have proposed theoretical models for the optical/infrared emission of AXPs, based on the magnetar scenario or fossil disk scenario. In the former scenario, \\citet{Eichler_Gedalin_Lyubarsky_2002, Lu_Zhang_2004} applied mechanisms of coherent radio emission of rotation-powered pulsars to a magnetar, where the emission frequency is boosted to the infrared, optical, or even UV region by the strong magnetic field of the magnetar. \\citet{Heyl_Hernquist_2005} proposed an exotic quantum electrodynamics model for the broad-band emission of magnetars. \\citet{Ertan_Cheng_2004} applied the outer gap model of rotation-powered gamma-ray pulsars on magnetars. In the latter scenario, \\citet{Perna_Hernquist_Narayan_2000, Perna_Hernquist_2000, Hulleman_van_Kerkwijk_Kulkarni_2000} discussed the optical/infrared spectrum generated by an irradiated fall-back disk and found that the sizes of the inner and outer radii of the disk become implausible. In addition, hybrid models of both scenarios have also been proposed. \\citet{Ertan_Cheng_2004} showed that the pulsed optical emission can be explained by the disk-star dynamo gap model, where the fossil disk plays a key role in the emission mechanism. \\citet{Ertan_Caliskan_2006, Ertan+_2007} showed that an irradiated and viscously heated disk model can account for the observed infrared/optical spectrum of both pulsed and un-pulsed components, by magnetospheric optical emission \\citep{Ertan_Cheng_2004} and disk emission, respectively.  In any of the above models, the origin of the emission can be clearly determined by the pulsation. The emission from the co-rotating magnetosphere must be pulsed, while that from the disk is unlikely to be pulsed. In this paper, we report a search for the pulsation of the AXP 4U 0142+61 in the $K^\\prime$ band, using the fast-readout mode of IRCS mounted on the Subaru 8.2-m telescope. This is the first search for the pulsation in the infrared band.  In the following, the dereddened spectrum of 4U 0142+61 is obtained using the interstellar extinction of $A_V = 3.5 \\pm 0.4$, which was determined by using red clump stars \\citep{Durant_van_Kerkwijk_2006a}. ",
        "conclusions": "\\begin{center} \\begin{tabular}{llll} \\hline Date (UT)   &  2004-07-31 \\\\ Start time (UT) & 12:57:14 \\\\ End time (UT)   & 15:21:07 \\\\ Duration    &  8633 s  \\\\ Filter      & $K^\\prime$ ($\\lambda_{\\rm eff}: 2.11$ $\\mu$m) \\\\ NDR         &   16        \\\\ Co-adds     &   1         \\\\ Exposure    &  0.84  s    \\\\ \\# of frames & $300 \\times 22 = 6600$ \\\\ Net exposure &  5.5 ks  \\\\ Seeing (FWHM) & 0.$^{\\prime\\prime}$37 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}     To discuss the emission mechanism of the AXP, we generated the spectral energy distribution ($\\nu F_\\nu$) of 4U 0142+61 for the phase-averaged fluxes and pulsed fluxes (figure \\ref{fig: 4u_nufnu_av3.5+pulse.ps}), using all the published data (see caption of figure \\ref{fig: 4u_nufnu_av3.5+pulse.ps}) and the $K^\\prime$ band fluxes obtained in this work.  For the Subaru observation on Sep. 2003 \\citep{Morii+_2005}, we analysed the data in our own way and obtained different results from that shown in \\citet{Durant_van_Kerkwijk_2006b}. We will present the details of the photometry in our companion paper \\citep{Tanaka+_2008}. In the phase-averaged spectrum, a peak at the $H$ band is clearly seen. As discussed in \\citet{Hulleman_van_Kerkwijk_Kulkarni_2004}, this feature may be a proton cyclotron feature. Although other cyclotron features in other bands are expected, they are not yet clear. Further simultaneous multi-band observations are necessary, because the spectrum is variable.  To understand the emission mechanism of the pulse component, we calculated the slope of the spectral shape. Combined with the pulsed flux of the $K^\\prime$ band (this work; $PF_{\\rm rms}$) and that of the $i^\\prime$ band \\citep{Dhillon+_2005}, the slope of the pulsed component ($F_\\nu \\propto \\nu^\\alpha$) was found to be constrained to $\\alpha > -0.87$ (90\\% C.L.) for the interstellar extinction of $A_V = 3.5$.  Because the constraint on the slope of the pulsed flux is not so stringent, any emission mechanisms are still allowed for the pulse component. For example, the Rayleigh-Jeans tail of the blackbody is $\\alpha = +2$ and synchrotron self-absorption is $\\alpha = +2.5$. In the quantum electrodynamics model of \\citet{Heyl_Hernquist_2005}, $\\alpha = 0$ for the optical-infrared region. In the disk-star dynamo gap model of \\citet{Ertan_Cheng_2004}, $\\alpha = -0.5$; while in the magnetar model by the same authors, $\\alpha$ is slightly smaller than $-0.5$.  In the curvature-drift-induced maser mechanism proposed by \\citet{Lu_Zhang_2004}, the maser frequency of 4U 0142+61 is about $\\nu_M = 1.39 \\times 10^{14}$ Hz, which is near the $K$ band. Then, the maser amplification ratio equating to the pulse fraction in their model is maximum near the $K$ band. Since the upper limit for the pulse fraction of the $K^\\prime$ band is similar to the pulse fraction of the $i^\\prime$ band, this model is a possible mechanism. The $H$ band peak in the phase-averaged spectrum may be the maser frequency peak of this emission mechanism. If so, then a larger pulse fraction at the $H$ band is expected.  Viewed differently, the dust disk model is also a possible mechanism \\citep{Wang_Chakrabarty_Kaplan_2006}, because the fraction of un-pulsed flux is dominant in the $K^\\prime$ band.   In the near future, we can make use of the adaptive optics (AO) technique for the sensitive detection of the AXP. The Subaru Telescope plans to install an upgraded 188-element AO system with a laser guide star \\citep{Hayano+_2003}. Using this technique, we will be able to constrain the pulse fraction to $PF < 3\\%$ and constrain the spectral slope of the pulsed component to $\\alpha > 1$ for a feasible observation time. Furthermore, the $K$ band flux varies about 40\\% over the time scale of years and hence it is possible that the $K$ band pulsation would grow above the detection limit. Further monitoring of the infrared pulsation should prove interesting.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){fig/figure3.ps} \\end{center} \\caption{Energy spectrum ($\\nu F_\\nu$) of 4U 0142+61. The horizontal and vertical axes are shown in units of Hz and Jy$\\cdot$Hz, respectively. Phase-averaged fluxes derived from \\citet{Hulleman_van_Kerkwijk_Kulkarni_2000, Hulleman_van_Kerkwijk_Kulkarni_2004, Morii+_2005, Durant_van_Kerkwijk_2006b, Wang_Chakrabarty_Kaplan_2006, Wang_Kaspi_2008, Tanaka+_2008} are marked as small circles. The flux obtained by this work (see subsection \\ref{subsection: Phase-averaged flux}) is marked as a large circle. The pulsed flux in the $i^\\prime$ band is derived from \\citet{Dhillon+_2005} and marked as a large triangle. The $90$\\% C.L. upper limit for the pulsed flux in the $K^\\prime$ band obtained in this work ($PF_{\\rm rms}$) is marked as a large square and downward arrow. The pulsed fluxes are calculated as the product of the phase-averaged flux and the pulse fraction. The vertical bars show the 1$\\sigma$ level error range. Fluxes obtained on the same day are connected by solid lines. For the Subaru Telescope data, we used the results of \\citet{Morii+_2005} and \\citet{Tanaka+_2008}. To convert the magnitudes to fluxes, we used the zero fluxes of the Johnson-Cousins photometric system for the $I$$R$$V$$B$ bands of the KeckI(II)/LRIS, KeckII/ESI and UH88/Optic data, the Ellis system for the $R_E$ band of the KeckII/ESI data, the SDSS system for the $i^\\prime$$g^\\prime$ bands of the UHT/ULTRACAM data \\citep{Fukugita_Shimasaku_Ichikawa_1995} and the MKO-NIR system for the $K$$K^\\prime$$Ks$ bands of the KeckI/NIRC, CFHT/AOB, Subaru/IRCS and Gemini/NIRI data \\citep{Tokunaga_Vacca_2005}. We corrected the observed flux by the interstellar extinction of $A_V = 3.5$ \\citep{Durant_van_Kerkwijk_2006a}. Extinction other than that of the $V$ band was calculated by the relative extinction listed in \\citet{Schlegel_Finkbeiner_Davis_1998} and the extinction law given in \\citet{Chiar_Tielens_2006}. The systems or bands not listed in \\citet{Schlegel_Finkbeiner_Davis_1998} were corrected using values of the Landolt and UKIRT systems for optical and near-infrared bands, respectively. }\\label{fig: 4u_nufnu_av3.5+pulse.ps} \\end{figure}  \\bigskip  M.~M. is supported by the Post-doctoral Researchers Program of Rikkyo University and a Grant-in-Aid for Young Scientists(B)(18740120)."
    },
    "0809/0809.5197_arXiv.txt": {
        "abstract": "The hydrodynamics  of the interaction of pulsar and stellar winds in  binary systems harboring a pulsar and its impact  on the nonthermal  radiation  of the binary pulsar PSR~B1259-63/SS2883 is discussed.  The collision of an ultrarelativistic pulsar wind with a nonrelativistic stellar outflow results in significant bulk acceleration of the shocked material from the pulsar wind. Already at distances comparable to the size of the binary system, the  Lorentz factor of the shocked flow can be  as large  as $\\gamma\\sim4$. This results in  significant anisotropy of the inverse Compton radiation of accelerated electrons. Because of the Doppler boosting of the produced radiation, one should expect a variable gamma-ray signal from the system.  In particular, this effect may naturally explain the reported tendency of a decrease of  TeV gamma-ray flux close to the periastron. The modeling of the interaction of pulsar and stellar winds allows  self-consistent calculations of  adiabatic losses. Our results show that adiabatic losses dominate over  the radiative losses. These results  have direct impact  on the orbital variability of radio, X-ray and gamma-ray signals detected from the binary pulsar  \\psr{}. ",
        "introduction": "Binary systems represent an important population of very high energy (VHE) sources. Presently this category comprises  four objects \\cite{hess,ls5039, ls50392,lsi,cyg}. At least one of them,  PSR~B1259-63/SS2883, is a binary pulsar system consisting of a  $48\\, \\rm ms$ pulsar in  an elliptic orbit around a massive B2e optical star \\cite{johnston1}. The multiwavelength observations of the system \\cite{johnston,connors,johnston05,cominsky,kaspi,hirayama,chernyakova} indicate correlations between different energy bands \\cite{neronov}. One of the interesting  features of the high energy  emission of  PSR~B1259-63/SS2883 is the  behavior of the gamma-ray light-curve which deviates from the earlier  theoretical predictions \\cite{kirk99,kawachi04}.    The understanding of the nature/origin of the  high energy radiation of this object  is a quite complex issue.  It requires not only proper modeling  of the acceleration and  radiation process in the context of multiwavelength properties of the source, but also detailed studies  of the hydrodynamics of the interaction of the pulsar and stellar winds.  Indeed, the particle acceleration regime and formation of  energy distribution of relativistic particles after termination of the pulsar wind strongly depend on the dynamics  of the flow. In this regard one should note that the recent attempts at explanation  of the gamma-ray light-curve  of  PSR~B1259-63/SS2883 \\cite{neronov,khangulyan07} contain assumptions  which are closely linked to the  hydrodynamics of the system. For example, the density changes along current lines determine  the adiabatic loss rate in the flow. The flow bulk  motion is also responsible for the particle advection (escape) from the system. The nonradiative losses caused by these processes may have a strong impact on the energy spectrum and  light-curve of gamma-rays \\cite{khangulyan07}. Therefore, any self-consistent treatment of  the high energy radiation of the system requires hydrodynamic calculations of key parameters characterizing the interaction of two winds. This concerns, in particular,  the magnetic field. The magnetic field is expected to be nonhomogenous in the pulsar wind nebulae \\cite{kc}. The calculations of the spatial distribution of the strength of the magnetic field generally requires a  proper magnetohydrodynamical (MHD) approach.  However,  in the case when the the impact of the magnetic field on the  dynamics of  plasma is not dramatic,  it can be treated  hydrodynamically   within a  ``froze-in''  approach.  Below we assume the influence of the magnetic field to be weak and we discuss the impact of hydrodynamics  characterizing the collision of pulsar and stellar  winds on the high energy radiation of the system. The results are based on the recent hydrodynamic  calculations of Bo\\-go\\-va\\-lov et al. \\cite{bogovalov08} conducted through two different approaches using  a relativistic code  in the ``pulsar zone''  and a nonrelativistic code  in  the ``optical star'' region. It is assumed that  initially both the pulsar and stellar winds expand radially. ",
        "conclusions": "the signal from the source was detected only during  epochs close to the periastron passage.   The effect of the flow bulk acceleration strongly modifies the  relationship between the synchrotron X-ray and inverse Compton gamma-ray fluxes produced  by the same population of relativistic electrons. Indeed, as long as the magnetic field is frozen in the plasma, its strength is determined by the structure of the flow and may significantly vary across the flow. The plasma bulk motion affects the IC radiation production as well, but in a different way. Namely, in the co-moving frame, where electrons are distributed isotropically, the target radiation field will be Doppler boosted. This implies that hydrodynamical effects  can lead to a significant deviation from the standard relations between  the X-ray  and VHE gamma-ray fluxes (see e.g. \\cite{atoyan,khangulyan2}).  Moreover, due to the large  bulk Lorentz factors (for details see \\cite{bogovalov08}), the observed fluxes should have strong orbital phase dependence. Indeed, the direction of the post shock flow varies with the motion of the pulsar along the orbit around the  star. This implies significant changes of the Doppler factor $\\delta$.  Namely, $\\delta \\gg 1$ for small viewing angles (when the flow bulk velocity is directed towards the observer);  and $\\delta \\ll 1$ for large viewing angles.  Such a geometry is expected, in particular,  at the  periastron passage. Correspondingly, this will have a strong impact on the light-curve of nonthermal radiation of electrons, $F_\\gamma \\propto \\delta^{n}$ where typically $n \\geq 3$. Interestingly, in such a case one should expect  a direct correlation between different energy bands (from radio to VHE), even though the electrons responsible for these energy intervals of electromagnetic radiation may not have the same origin. This conclusion is supported by multiwavelength observation of \\psr{}\\cite{neronov}.  Moreover, this effect could be the reason for the  deep in the  VHE light-curve   close to the periastron as reported by the HESS collaboration  \\cite{hess}.  \\balance"
    },
    "0809/0809.0584_arXiv.txt": {
        "abstract": "The High Energy Stereoscopic System (HESS) survey of the Galactic plane has established the existence of a substantial number ($\\sim$40) of Galactic TeV $\\gamma$-ray sources, a large fraction of which remain unidentified. HESS\\,J0632+057 is one of a small fraction of these objects which is point-like in nature ($<2'$ rms), and is one of only two point-like sources that remain unidentified.  Follow-up observations of this object with \\emph{XMM-Newton} have revealed an X-ray source coincident with the TeV source and with the massive star MWC\\,148, of the spectral type B0pe. This source exhibits a hard spectrum, consistent with an absorbed power law with $\\Gamma=1.26\\pm0.04$, and shows significant variability on hour timescales. We discuss this spatial coincidence and the implied spectral energy distribution of this object and argue that it is likely a new $\\gamma$-ray binary system with a close resemblance to the three known members of this class and, in particular, to LS\\,I\\,+61\\,303.  Further X-ray, radio and optical observations of this system are needed to firmly establish HESS\\,J0632+057 as a new member of this rare class of Galactic objects. ",
        "introduction": "The emerging class of $\\gamma$-ray binaries has three well-established members, PSR\\,B1259$-$63/SS\\,2883 \\citep{HESS:psrb1259} LS\\,5039 \\citep{HESS:ls5039p1,HESS:ls5039p2} and LS\\,I\\,+61\\,303 \\citep{ MAGIC:lsi61,MAGIC:lsi61b, VERITAS:lsi61}, all high-mass X-ray binary systems \\footnote{There is evidence of a single TeV \\emph{flare} from Cyg X-1, but $\\gamma$-ray emission from this object is not firmly established~\\citep{MAGIC:cygx1}}.  All these objects show variable or periodic TeV $\\gamma$-ray emission with peak fluxes around $10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ (1--10\\,TeV). As such fluxes are well above the sensitivity achieved by the HESS Galactic plane survey \\citep{HESS:scanpaper2}, the serendipitous discovery of new $\\gamma$-ray binaries might be expected in HESS observations. Indeed, the TeV emission of LS\\,5039 was discovered in this way. Binaries are one of the few classes of sources expected to appear \\emph{point-like} for TeV instruments.  Of the $\\sim$40 unidentified TeV sources \\citep{Jim:NJP}, only HESS\\,J1745$-$290 \\citep{HESS:gcprl} (at the gravitational centre of the Galaxy) and HESS\\,J0632+057 \\citep{HESS:0632} are unresolved. For this reason alone HESS\\,J0632+057 can be considered as a candidate for a new binary system.  HESS\\,J0632+057 was discovered in HESS observations of the region of apparent interaction between the Monoceros Loop supernova remnant and the star-forming regions of the Rosette Nebula. Three possible associations were suggested by the HESS collaboration for this object: the unidentified \\emph{ROSAT} source 1RXS J063258.3+054857, the unidentified \\emph{EGRET} source 3EG\\,0634+0521 and the massive star MWC\\,148 (HD\\,259440, spectral type B0pe).  To resolve the question of the association of these three objects, and to probe the greater than 3 keV emission of the X-ray source, an \\emph{XMM-Newton} observation of this region was conducted in 2007, the results and implications of which are described below. ",
        "conclusions": "A fundamental question in the interpretation of XMMU\\,J063259.3+054801 is the physical process responsible for the X-ray emission. Given the existence of TeV emission from this region, it seems natural to attribute the power-law X-ray emission as synchrotron emission from ultrarelativistic electrons. However, spectra resembling that observed for source \\#1 (consistent with two temperature blackbody models with $kT_{1}\\sim$0.5~keV and $kT_{1}\\sim$2~keV) have been observed for isolated magnetic Bp stars, such as $\\sigma$\\,Ori\\,E, and been attributed to shock heating of magnetically confined winds \\citep[see][and references therein]{Townsend}. It seems possible that MWC\\,148 is a system of this type, and that the observed X-ray emission is thermal in origin. However, it is difficult to explain the TeV $\\gamma$-ray emission in such a scenario. In the discussions that follow we assume that the observed X-ray emission is nonthermal in origin. If this is not the case, then the observed flux provides an upper limit on the nonthermal X-ray flux of HESS\\,J0632+057.  The probabilities of chance associations of HESS\\,J0632+057 with MWC\\,148 and 1RXS J063258.3+054857 have been estimated as $10^{-4}$ and $10^{-3}$ \\citep{HESS:0632} respectively. The chance coincidence of a massive star with the 1$''$ error box of the brightest (and the hardest spectrum) source in an \\emph{XMM-Newton} observation is even less probable ($\\sim10^{-6}$). The least secure association is that with 3EG\\,0634+0521, as $\\sim$10\\% of all Galactic TeV sources lie within the 99\\% confidence contour of a 3EG source \\citep{GeVTeV}. Furthermore, the binary system SAX\\,J0635+0533 (composed of a 34~ms period pulsar and a Be star) has been suggested to power most or all of the flux of 3EG\\,0634+0521~\\citep{SAX0635}. However, as 3EG\\,0634+0521 is flagged as extended and/or confused it seems plausible that HESS\\,J0632+057 may contribute to this emission. The relative contributions of these and other sources to 3EG\\,0634+0521 should be established by the first months of data from \\emph{GLAST}. We note that there is no evidence of TeV emission from SAX\\,J0635+0533 \\citep{HESS:0632}.  If all the apparent associations of HESS\\,J0632+057 are correct then its spectral energy distribution (SED) bears a close resemblance to that of the known TeV binaries. All have hard ($\\Gamma<2$) X-ray spectra and significantly softer TeV spectra.  Furthermore, X-ray variability with timescales comparable to those observed here for source \\#1 has been observed in LS\\,I\\,+61\\,303 \\citep{Esposito}. Associations with GeV $\\gamma$-ray sources also exist in the cases of LS\\,5039 (3EG\\,J1824$-$1514) and LS\\,I\\,+61\\,303 (3EG\\,J0241+6103).  These systems are powered by either a relativistic pulsar wind or a wind-accretion powered jet \\citep[see e.g.][]{BinariesReview}.  An assessment of the possible power source of HESS\\,0632+057 requires an estimate of the distance to the system, Using the apparent visual magnitude $m_{V}=9.16$ of MWC\\,148 \\citep{Tycho2}, assuming an absolute magnitude $M_{V}=-4.0$, and accounting for reddening assuming an intrinsic $M_{B}-M_{V}$ of -0.3, we find a distance of $\\approx$1.5~kpc, consistent with that of the Rosette Nebula \\citep{RosetteDist} \\footnote{Following \\citet{ryter}, the implied optical extinction of $E_{B-V} = 0.75$, is compatible with the best fit $N_{H}$ for source \\#1, supporting the physical association of the two sources.}.  At this distance, the $1''$ limit on the intrinsic size (and the similar limit on the projected displacement from MWC\\,148) of the X-ray source implies an origin of the emission within $\\approx2\\times10^{16}$~cm of the star and hence a radiation density of greater than 10$^{6}$~eV~cm$^{-3}$ \\footnote{Causality arguments limit the size of the X-ray emission region (rather than the distance from the star) to $r_{X}<c\\Delta t$, where $\\Delta t$ is the variability timescale of $\\sim$30~ks, leading to $r_{X}<10^{15}$ cm.}.  In the presence of such a strong, high-energy ($kT$\\,$\\sim$\\,3~eV) photon field, inverse Compton (IC) scattering of $\\sim$TeV electrons must occur in the deep Klein--Nishina (KN) regime. IC cooling is dominant if $f_{\\rm KN}(E_{e})U_{\\rm rad} > U_{\\rm mag}$, where $f_{\\rm KN}(E_{e})$ is the suppression factor due to the KN effect. Following \\citet{Moderski}, $f_{\\rm KN}(E_{e}) \\approx (1.0 + (4E_{e}\\times2.8kT)/(m_{e}c^{2})^{2})^{-1.5}$, and for $kT=3$~eV and $E_{e}\\,=\\,1$~TeV, $f_{\\rm KN}\\sim10^{-3}$.  If IC cooling in the KN regime is dominant, then a hard X-ray spectrum and softer TeV $\\gamma$-ray spectrum, as observed in this case, can be expected. Such a scenario has been recently discussed for a pulsar wind nebula in the central parsec of our galaxy~\\citep{GCPWN} and for binary systems such as PSR\\,B1259$-$63/SS\\,2883~\\citep{Khangulyan07}.  Assuming for the moment a distance of the emission region from the star of a few AU, radiation densities of $U_{\\rm rad} \\sim$1 erg~cm$^{-3}$ are expected. The relative flux of X-rays and TeV $\\gamma$-rays then implies $B\\sim 5 (f_{KN}\\,F_{X}/F_{\\gamma})^{0.5}\\,\\mathrm{G}\\,\\sim\\,70$ mG.  For these radiation and magnetic fields 1 TeV electrons produce X-rays of a few keV and $\\gamma$-rays of almost TeV energy (if the IC scattering proceeds in the deep KN regime). The cooling time ($E/\\mbox{d}E/\\mbox{d}t$) for these electrons is then $\\sim$20~ks, compatible with the X-ray variability observed for source \\#1. Furthermore, in contrast to synchrotron-dominated cooling, the cooling timescale is almost constant with energy in the region probed by synchrotron X-rays and no spectral variability is expected, again compatible with this observation. We note that in this scenario correlated X-ray and TeV $\\gamma$-ray variability is expected.  An illustrative one-zone model with these parameters ($B=70$~mG, $U_{rad}=1$ erg~cm$^{-3}$) is shown in Figure~\\ref{fig_sed}. Particles are injected with a spectrum $dN/dE \\propto E^{-\\alpha} \\exp{(-E/E_{\\rm max})}$, with $\\alpha=2$ and $E_{\\rm max} = 10$ TeV, and cool continuously. The SED reaches an equilibrium state (as shown) for timescales $\\gg 20$ ks. Lower values of $E_{\\rm max}$ are difficult to reconcile with the highest-energy X-ray and $\\gamma$-ray data points. The effect of changing the minimum energy of injected electrons is shown in Figure~\\ref{fig_sed}. We note that values as high as 100~GeV (long-dashed curve) might be expected in the binary PWN scenario: \\emph{GLAST} is well suited to probe this region (see sensitivity curve in Figure~\\ref{fig_sed}).  If genuinely associated, the flux level of 3EG\\,0634+0521 may represent a higher state of this variable object, or indicate a softer electron spectrum in the GeV range.  \\begin{figure*}[t] \\epsscale{1.04} \\plotone{f4.eps} \\caption{Spectral energy distribution of HESS\\,J0632+057 / XMMU\\,J063259.3+054801 (red symbols/regions) compared to that of the $\\gamma$-ray binary system LS\\,I\\,+61\\,303 \\citep[grey symbols/regions, see][and references therein]{Chernyakova} and the estimated 1 year GLAST sensitivity for this region (thick solid line). The solid, short-dash and long-dash curves show the synchrotron and IC components of a simple time-dependent one-zone model with $E_{\\rm min}$ set to 1 GeV, 10 GeV and 100 GeV, respectively - see text for details.\\label{fig_sed}} \\end{figure*}  The power requirements of the observed emission place limitations on the possible energy source in the system. The model shown in Fig.~\\ref{fig_sed} requires a power of $10^{34}$ erg s$^{-1}$. The luminosities in individual wavebands are $\\approx1.3\\times10^{32}$ erg s$^{-1}$ (1--10 keV), $\\approx$3$\\times10^{34}$ erg s$^{-1}$ (0.2--2~GeV), and $\\approx$5$\\times10^{32}$ erg s$^{-1}$ (0.4--4~TeV). These luminosities should be compared to the available kinetic power of the Be-star wind \\citep[$\\cal O$($10^{34}$) erg s$^{-1}$, see e.g.][]{Waters}, the wind accretion luminosity on to an (unseen) compact companion \\citep[$\\sim$$10^{35}$ erg s$^{-1}$, see ][]{AccretionBook}, and the spin-down power of an (unseen) young pulsar companion \\citep[$\\dot{E}\\approx$$8\\times10^{35}$ erg s$^{-1}$ e.g. in PSR\\,B1259$-$63, ][]{Johnston}."
    },
    "0809/0809.2853_arXiv.txt": {
        "abstract": "We report the observations of 12 globular clusters with the Far-Infrared Surveyor (FIS) on-board AKARI infrared satellite. Our goal is to search for emission from the cold dust within clusters. We detect diffuse emissions toward NGC~6402 and 2808, but the IRAS 100~$\\mu$m maps show the presence of strong background radiation. They are likely emitted from the galactic cirrus, while we cannot rule out the possible association of a bump of emission with the cluster in the case of NGC~6402. We also detect 28 point-like sources mainly in the WIDE-S images ($90~\\micron$). At least several of them are not associated with the clusters but background galaxies based on some external catalogs. We present the spectral energy distributions (SEDs) by combining the near-and-mid infrared data obtained with the Infrared Camera (IRC) if possible. The SEDs suggest that most of the point sources are background galaxies. We find one candidate of the intracluster dust which has no mid-infrared counterpart unlike the other point-like sources, although some features such as its point-like appearance should be explained before we conclude its intracluster origin. For most of the other clusters, we have confirmed the lack of the intracluster dust. We evaluate upper limits of the intracluster dust mass to be between $10^{-5}$ and $10^{-3}$~M$_\\odot$ depending on the dust temperature. The lifetime of the intracluster dust inferred from the upper limits is shorter than 5~Myr ($T_{\\mathrm{d}}=70$~K) or 50~Myr (35~K). Such short lifetime indicates some mechanism(s) are at work to remove the intracluster dust. We also discuss its impact on the chemical evolution of globular clusters. ",
        "introduction": "Red giants in globular clusters are known to have low initial masses of about 1~M$_\\odot$ considering their high ages of 10--13~Gyr. Observations revealed that some of the red giants have the mass loss accompanied with dust formation (e.g.\\ \\cite{Frogel-1988}, \\cite{Matsunaga-2005}). Circumstellar dust has been detected in on-going space missions ({\\it Spitzer Space Telescope}, SST, Boyer \\etal\\ \\yearcite{Boyer-2006}, \\yearcite{Boyer-2008}, \\cite{Origlia-2007}; AKARI, \\cite{Ita-2007}). Terminal velocities of the stellar winds from red giants, 10--15 km~s$^{-1}$ \\citep{McDonald-2007}, are lower than typical escape velocities from massive globular clusters, $\\sim$~30~km~s$^{-1}$ \\citep{Gnedin-2002}. For this reason, one would expect that the released materials are trapped and accumulated within clusters in the halo region. On the other hand, the intracluster matter is removed by ram pressure when clusters cross the galactic plane. It was expected that gas of the order of 10~$M _\\odot$ exists in a cluster depending on its size and the time since the last passage through the galactic plane \\citep{Tayler-1975}.  In spite of the early expectations, most previous attempts failed to detect intracluster matter and placed upper limits which were lower than the expected amounts. The surveys range from the ones of dust to those of different forms of gas (dust, \\cite{Hopwood-1999}; molecular gas, \\cite{Smith-1995}; atomic gas, \\cite{Faulkner-1991}; ionized gas, \\cite{Knapp-1996}; and references therein). Among the challenges to search for intracluster dust, the most secure detection was found in the globular cluster M~15 (\\cite{Knapp-1995}, \\cite{Evans-2003}). Recently, \\citet{Boyer-2006} confirmed the extended emission of the intracluster dust, which they named IR1a, with SST. They also obtained its temperature ($\\sim 70$K) and dust mass ($9 \\times 10^{-4}$~M$_\\odot$) based on the spectral energy distribution (SED). The temperature was found to be higher than that of the interstellar dust (\\cite{Draine-1985}; \\cite{Sodroski-1987}). The obtained mass is still 4 times smaller than the expected mass released via the mass loss of red giants. Intracluster gas has been also detected in a couple of cases (e.g.\\ \\cite{Faulkner-1991}, \\cite{Freire-2001}).  The deficiency of the intracluster dust suggests that some mechanism(s) are at work to remove or destroy the dust even in the halo region. The discrepancy between the expected amount of the released dust and the upper limit is larger than ten for some clusters whose intracluster dust has not been detected \\citep{Knapp-1995}. It remains to be understood what mechanism(s) are efficient to what extent. For example, \\citet{Lynch-1990} proposed sputtering of dust grains by hot halo gas to explain the deficiency. On the other hand, Hopwood \\etal\\ (\\yearcite{Hopwood-1998}, \\yearcite{Hopwood-1999}) rejected the sputtering as the mechanism responsible for the removal of the intracluster dust. \\citet{Penny-1997} presented a list of possible mechanisms. Very recently, \\citet{Umbreit-2008} demonstrated that stellar collisions can generate the kinetic energy to remove the intracluster dust and gas.  The fate of the released matter within the clusters may have an effect on evolution of globular clusters. Although it was once considered that stars in a globular cluster are chemically homogeneous, more and more studies revealed evidence of inhomogeneity. Readers are referred to \\citet{Smith-1987}, for example, for early investigations mainly about carbon and nitrogen abundances. Recently, the existence of stars enriched in helium have attracted a lot of attention (\\cite{Norris-2004}, \\cite{Lee-2005}). Many authors have argued that those enrichment results from the matter released from intermediate-mass asymptotic giant branch (AGB) stars (e.g.\\ \\cite{Tsujimoto-2007}, \\cite{Ventura-2008}). \\citet{Suda-2007} proposed that low-mass AGB stars can also play an important role in generating stars with enriched helium contents via intracluster-matter accretion. However, it is not well studied whether the ejecta of AGB stars are actually retained within clusters in the context of dynamics of intracluster gas and dust. The ejecta may disappear before they are used for the generation of enriched stars. The connection to such evolutionary scenarios strengthens the importance  to understand how the intracluster matter decreases faster than expected if it is true.  The Far-Infrared Surveyor (FIS; \\cite{Kawada-2007}) on-board AKARI satellite \\citep{Murakami-2007} provides us with a new opportunity to investigate the intracluster dust. The FIS can take much deeper images in the far infrared with much higher angular resolution than before. Observations in the far infrared have high sensitivities for the cold dust, $\\sim$30~K. We used the FIS to observe a dozen of globular clusters as a part of the Mission Program to investigate the mass-loss phenomenon in evolved stars in globular clusters and nearby dwarf galaxies (AGBGA; PI, Y. Nakada).  We describe the observation and data reduction in section \\ref{sec:Observation}. As we will present, we detect a few extended emissions and several point-like sources. Analyses of the extended emissions and the point-like sources are presented in section \\ref{sec:diffuse} and \\ref{sec:points}, respectively. We argue that most of them are not associated with the clusters, while one point-like source remains to be a candidate of the dust within the cluster (section \\ref{sec:candidate}). In section \\ref{sec:Discussion}, we present discussions on the mass of the intracluster dust (for both the possible detection and upper limits) and discuss the impact on evolution of globular clusters. Section \\ref{sec:Summary} summarizes this work. ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{Upper limit of the dust mass from the background fluctuation} \\label{sec:MassLimit}  In this section, we assess the upper limit of undetected intracluster dust within the half-mass radius $r_{\\mathrm{h}}$ based on the background fluctuation. We follow the equation of \\citet{Lynch-1990}, \\begin{equation} f_\\nu = 6 \\sqrt{N} A_{\\mathrm{p}} \\sigma _\\nu, \\end{equation} where $N$ is the number of pixels corresponding to the area, $A_{\\mathrm{p}}$ is the area of the pixel, and $\\sigma _\\nu$ is the standard deviation of the background, to estimate the flux limits (eq. \\ref{eq:Mdust}). We list the limits of $90~\\micron$ flux in table \\ref{tab:limits} with some parameters of the clusters.  According to \\citet{Boyer-2006}, the flux $f_\\nu$ [mJy] can be converted into the dust mass as, \\begin{equation} M_{\\mathrm{d}} = 4.79 \\times 10 ^{-17} ~f_\\nu \\frac{D_{\\mathrm{sun}} ^2}{\\kappa _\\nu B_\\nu (T_{\\mathrm{d}})} ~[{\\rm M}_\\odot], \\label{eq:Mdust} \\end{equation} under the assumption that the dust cloud is optically thin. $D_{\\mathrm{sun}}$ is the distance in kilo-parsec, $\\kappa _\\nu$ is the dust absorption coefficient in cm$^2$~g$^{-1}$, $B_\\nu (T_{\\mathrm{d}})$  is the Planck function in cgs units, and $T_{\\mathrm{d}}$ is the dust temperature. We calculate $\\kappa _\\nu = 40$ cm$^2$~g$^{-1}$ from \\citet{Ossenkopf-1994} assuming a standard Mathis-Rumpl-Nordsieck dust distribution as \\citet{Boyer-2006} did. The temperature has a significant effect on the estimate of the dust mass through the Planck function $B _\\nu (T _{\\mathrm{d}})$. At $90~\\micron$, its dependency on the temperature is as $B_{\\mathrm{100K}}:B_{\\mathrm{70K}}:B_{\\mathrm{35K}} = 2.2:1:0.09$, which directly changes the mass estimate according to equation (\\ref{eq:Mdust}). \\citet{Angeletti-1982} performed realistic estimates of the temperature of intracluster dust, and obtained temperatures between 35 and 90~K dependent on the assumed chemical composition of dust grain. Here we use two assumed dust temperatures of 35~K and 70~K. The value of 70~K is adopted from the temperature of M~15 IR1a\\citep{Boyer-2006}. The $M_{\\mathrm{d}}$ gets smaller in case of $T_{\\mathrm{d}}=90$~K, which accentuates the lack of intracluster dust. The results are listed in table \\ref{tab:limits}.   \\begin{table*} \\begin{minipage}{0.99\\hsize} \\caption{ Parameters of the observed globular clusters and obtained upper limits. The updated version of \\citet{Harris-1996} is used to take metallicities [Fe/H], integrated magnitudes $\\mathrm{M_V}$, distances $D_{\\mathrm{sun}}$, and heights $Z$ above or below the galactic plane. Escape velocities at half-mass radii are taken from \\citet{Gnedin-2002}. If available, the velocities $V_{\\mathrm{Z}}$ perpendicular to the plane are taken from \\citet{Dinescu-1999} and \\citet{Casetti-Dinescu-2007}. The $3 \\sigma$ upper limits of the emission $f_{90\\micron}$ and the mass $M_{\\mathrm{d}}$ of the intracluster dust are obtained within the half-mass radii $r_{\\mathrm{h}}$, while $\\tau _{\\mathrm{d}}$ is the upper limit of the lifetime of the dust obtained by equation (\\ref{eq:lifetime}). Two values are listed based on the two assumed dust temperatures, $T_{\\mathrm{d}}=35 and 70$~K, for both $M_{\\mathrm{d}}$ and $\\tau _{\\mathrm{d}}$. The $M_{\\mathrm{d}}$ values for N5024 FIS2 (this work) and for M~15 and NGC~6356 (from early papers) are also listed and their lifetimes are obtained in the same manner. \\label{tab:limits}} \\begin{center} \\begin{tabular}{crrrrrrrrrrr} \\hline Object  & \\multicolumn{1}{c}{[Fe/H]} & \\multicolumn{1}{c}{M$_{\\mathrm{V}}$} & \\multicolumn{1}{c}{$V_{\\mathrm{esc}}^{\\mathrm{h}}$} & \\multicolumn{1}{c}{$D_{\\mathrm{sun}}$} & \\multicolumn{1}{c}{$Z$} & \\multicolumn{1}{c}{$V_{\\mathrm{Z}}$} & \\multicolumn{1}{c}{$f_{90\\micron}$} & \\multicolumn{2}{c}{$M_{\\mathrm{d}}$} & \\multicolumn{2}{c}{$\\tau _{\\mathrm{d}}$} \\\\ & \\multicolumn{1}{c}{dex} & \\multicolumn{1}{c}{mag} & \\multicolumn{1}{c}{km/s} & \\multicolumn{1}{c}{kpc} & \\multicolumn{1}{c}{kpc} & \\multicolumn{1}{c}{km/s} & \\multicolumn{1}{c}{mJy} & \\multicolumn{2}{c}{$10^{-5}$~M$_\\odot$} & \\multicolumn{2}{c}{$10^6$~yr} \\\\ & & & & & & & & 70~K & 35~K & 70~K & 35~K \\\\ \\hline NGC1261 & $-1.35$ & $ -7.81$ & 20.2 & 16.4 & 12.9 &  ---   &  16.9 &  8.8 & 96 &  1.9 & 20 \\\\ NGC1851 & $-1.22$ & $ -8.33$ & 24.3 & 12.1 &  6.9 & $-109$ &  10.6 &  3.0 & 32 &  0.3 &  3 \\\\ NGC1904 & $-1.57$ & $ -7.86$ & 26.1 & 12.9 &  6.3 & $   6$ &  14.5 &  4.6 & 50 &  1.6 & 17 \\\\ NGC2808 & $-1.15$ & $ -9.39$ & 45.8 &  9.6 &  1.9 & $  63$ &  25.8 &  4.6 & 50 &  0.1 &  2 \\\\ NGC5024$^*$ & $-1.99$ & $ -8.70$ & 20.9 & 17.8 & 17.5 & $ -83$ &  20.1 & 12.3 &133 &  5.0 & 55 \\\\ NGC5139 & $-1.62$ & $-10.29$ & 44.0 &  5.3 &  1.4 & $  -9$ & 122.8 &  6.7 & 72 &  0.3 &  3 \\\\ NGC5272 & $-1.57$ & $ -8.93$ & 22.7 & 10.4 & 10.2 & $-154$ &  20.2 &  4.2 & 46 &  0.5 &  6 \\\\ NGC5634 & $-1.88$ & $ -7.69$ & 17.0 & 25.2 & 19.1 &  ---   &  11.0 & 13.5 &146 & 10.9 &118 \\\\ NGC5904 & $-1.27$ & $ -8.81$ & 29.2 &  7.5 &  5.4 & $-222$ &  33.4 &  3.6 & 39 &  0.3 &  3 \\\\ NGC6205 & $-1.54$ & $ -8.70$ & 26.9 &  7.7 &  5.0 & $-119$ &  26.9 &  3.1 & 34 &  0.4 &  5 \\\\ NGC6341 & $-2.28$ & $ -8.20$ & 29.2 &  8.2 &  4.7 & $  35$ &  19.7 &  2.6 & 28 &  3.2 & 35 \\\\ NGC6402 & $-1.39$ & $ -9.12$ & 26.2 &  9.3 &  2.4 &  ---   &  72.9 & 12.2 &132 &  0.9 &  9 \\\\ \\hline N5024 FIS2 & $-1.99$ & $ -8.70$ & 20.9 & 17.8 & 17.5 & 1.11 & \\multicolumn{1}{c}{114} & 50 & 500 &  10 & 100 \\\\ NGC~7078 & $-2.26$ & $-9.17$ & 27.4 & 10.3 & 4.7 & 1.06 & \\multicolumn{1}{c}{---} & 90$^\\dagger$ & \\multicolumn{1}{c}{---} & 20 & \\multicolumn{1}{c}{---}  \\\\ NGC~6356 & $-0.50$ & $-8.52$ & 31.7 & 15.2 & 2.7 & 0.74 & \\multicolumn{1}{c}{---} & 860$^\\ddagger$ & --- & 7 & --- \\\\ \\hline \\end{tabular} \\end{center} $^*$In the event that N5024 FIS2 is not associated with the cluster.\\\\ $^\\dagger$\\citet{Boyer-2006},~ $^\\ddagger$\\citet{Hopwood-1998}. \\end{minipage} \\end{table*}  In figure \\ref{fig:Mdust}, we plot the estimated upper limits of the intracluster dust against some parameters of the clusters which are taken from \\citet{Harris-1996} and \\citet{Gnedin-2002}\\footnote{We used their machine-readable tables on Web:\\\\ http://www.physics.mcmaster.ca/Globular.html\\\\ http://www.astro.lsa.umich.edu/\\%7Eognedin/gc/vesc.dat}. We used the Harris catalog available online which was updated in 2003. We also include the mass of the intracluster-dust candidate, N5024 FIS2 (see section \\ref{sec:N5024FIS2}), and those of the dust clouds in M~15 and NGC~6356 reported in previous studies. No parameter can be an exclusive factor to determine the accumulated mass of the intracluster dust. Although \\citet{Hopwood-1999} suggested that the metallicity may be an influencing factor based on the fewer samples with shallower observations, the detection of M~15 IR1a disclaims such possibility. It is unclear what parameter(s) determine the amount of the accumulated intracluster dust.  \\begin{figure} \\begin{minipage}{80mm} \\begin{center} \\FigureFile(70mm,105mm){fig11.ps} \\end{center} \\caption{ The mass and upper limits of the intracluster dust plotted against (a) $Z$-height, (b) metallicity [Fe/H], (c) integrated magnitude $M_{\\mathrm{V}}$, and (d) escape velocity at the half-mass radius $V_{\\mathrm{esc}}^{\\mathrm{h}}$. The $y$-axis value is based on the assumption of $T_{\\mathrm{d}}=70$~K. In case of $T_{\\mathrm{d}}=35$~K, all the points shift upward by a factor of 11. The arrows indicate the upper limits obtained in this work, and star symbols indicate the estimated mass of the intracluster-dust candidate N5024 FIS2. Two filled circles indicate M~15 (NGC~7078) and NGC~6356 based on the estimate by \\citet{Boyer-2006} and \\citet{Hopwood-1998} respectively. \\label{fig:Mdust}} \\end{minipage} \\end{figure}  The estimated dust mass of M~15 IRa1 is $(9 \\pm 2) \\times 10^{-4}$~M$_\\odot$, much larger than the upper limits in table \\ref{tab:limits}. \\citet{Boyer-2006} used a box aperture of $130\\arcsec.5 \\times 130\\arcsec.5$, which is larger than a circle with $r_{\\mathrm{h}}=1.06\\arcmin$. Nonetheless, the size of the far-infrared emission of the detected intracluster dust, IR1a, is smaller than the circle so that we can make a direct comparison of the values for IR1a with the upper limits obtained in this work. It is clear that the intracluster dust in M~15, which is still smaller than the estimated mass released from red giants \\citep{Boyer-2006}, is exceptionally large compared with possible intracluster dust in other globular clusters if any. The parameters such as $V_{\\mathrm{esc}}^{\\mathrm{h}}$ and $Z$-height for at least some of our samples are comparable with those of M~15.   \\subsection{The lack of intracluster dust}\\label{sec:LackDust}  Here we discuss the lack of intracluster dust quantitatively by estimating the lifetime from the expected amount of gas released from red giants. We can calculate the gas amount corresponding to that of the intracluster dust with the dust-to-gas mass ratio $\\psi$. The ratio is considered to be dependent on metallicity. \\citet{Marshall-2004} concluded that $\\psi$ is proportional to the metallicity by comparing OH/IR stars in the solar neighborhood and the large Magellanic Cloud (also see \\cite{vanLoon-2006}). It is worthwhile noting uncertainties on $\\psi$ especially in the low-metallicity environment. \\citet{Marshall-2004} found the proportionality between $\\psi$ and metallicity based on the objects between $\\sim 0.1~Z_\\odot$ and $\\sim 1~Z_\\odot$. It is not clear if the extrapolation can be applied below [Fe/H]~$= -1$~dex. \\citet{Boyer-2006} found more than 20 stars with dusty wind in M15. Although they did not give any estimate on the dust mass-loss rate, some of them have large mid-infrared excesses comparable with the mass-losing stars in the metal-rich cluster 47~Tuc \\citep{Origlia-2007}. Such dusty mass-loss was not expected in a metal-poor cluster like M~15. If $\\psi$ is not so small as expected from the proportionality, the amount of gas corresponding to the upper limits of dust gets smaller. The $\\psi$ value can be changed by non-canonical evolution of stars such as interacting binaries and/or blue stragglers. In the following estimations, we simply assume $\\psi = 0.01 \\times 10^{\\mathrm{[Fe/H]}}$.   It is possible to evaluate the number of mass-losing stars in a cluster by using the specific evolutionary flux, \\begin{equation} B \\sim 2 \\times 10^{-11} {\\rm ~stars ~yr^{-1} ~L_\\odot^{-1}}, \\end{equation} which is the number of stars entering or leaving any post-main-sequence stage per year and per solar luminosity of the stellar population \\citep{Renzini-1986}. The average dust mass released per year within a cluster can be obtained as, \\begin{equation} F = \\Delta M_{\\mathrm{d}} B L_{\\mathrm{cl}}, \\label{eq:MassLoss} \\end{equation} where $\\Delta M_{\\mathrm{d}}$ is the amount of dust to be lost during the post-main-sequence stage and $L_{\\mathrm{cl}}$ is the total luminosity of the cluster. We assume that each red giant in a cluster loses the gas of 0.2~M$_\\odot$, then $\\Delta M_{\\mathrm{d}}$ is calculated as $0.2 \\psi$~M$_\\odot$ $= 2 \\times 10^{-3+{\\mathrm{[Fe/H]}}}$~M$_\\odot$. We obtain the upper limits of the lifetime of the intracluster dust: \\begin{equation} \\tau _{\\mathrm{d}} = \\frac{M_{\\mathrm{d}}}{F} = \\frac{M_{\\mathrm{d}}}{\\Delta M_{\\mathrm{d}} B L_{\\mathrm{cl}}/2}. \\label{eq:lifetime} \\end{equation} $L_{\\mathrm{cl}}$ is divided by two because the upper limit $M_{\\mathrm{d}}$ has been obtained for the intracluster dust within the half-mass radius. This factor may not be correct if the dust released from stars outside the half-mass radius falls into the central region of the cluster, in which case the estimated lifetime $\\tau_{\\mathrm{d}}$ gets even smaller.  We list the obtained lifetimes in table \\ref{tab:limits}. Two values are obtained based on the assumed dust temperatures of 35~K and 70~K. The value with $T_{\\mathrm{d}}=35$~K is 11 times larger than the one with $T_{\\mathrm{d}}=70$~K from the ratio of Planck function. In both cases, they are significantly shorter than the typical time interval, $\\sim 10^8$~yr, between passages through the galactic plane for most of the clusters. Please note that the $\\tau _{\\mathrm{d}}$ gets even smaller if $\\psi$ is larger than the value expected from the proportionality between $\\psi$ and [Fe/H].  The kinematics information of the clusters suggests that most of the clusters have actually spent a significant time as large as $10^8$~yr since the last passage. Here we consider the orbits obtained by \\citet{Dinescu-1999} and \\citet{Casetti-Dinescu-2007}. The velocity components perpendicular to the galactic plane, i.e. $V_{\\mathrm{Z}}$ listed in table \\ref{tab:limits}, indicate that five clusters are now approaching to the plane. For NGC~1904 and 5139, $V_{\\mathrm{Z}}$ almost equals to zero, so that they are in the middle of the passage through the galactic plane. Time since the most recent passages is about the half of the orbital periods or longer for these seven clusters. NGC~1851 and 6341 are still moving away from the plane. NGC~1851 is well separated from the plane by 6.9~kpc. The obtained lifetime, $\\tau_{\\mathrm{d}} < 3$~Myr, is certainly shorter than the time since its last passage. For NGC~6341, on the other hand, the discrepancy between the obtained lifetime and the time since the last passage is not large to conclude the lack of the intracluster dust in case of $T_{\\mathrm{d}}=35$~K. There are no proper motions available for NGC~1261, 5634, and 6402.  We confirmed that the estimated lifetimes are shorter than the duration in which the intracluster dust can be accumulated since the cluster crossed the galactic plane for at least seven clusters in our sample. The conclusion for NGC~5024 depends on whether N5024 FIS2 belongs to the cluster or not. The short lifetimes indicate that the intracluster dust is not accumulated as was expected and disappears in a short time scale. Thanks to the high sensitivity of the AKARI/FIS, we could conclude the lack of the intracluster dust even with the lower dust temperature of 35~K.  \\subsection{Estimated mass of the intracluster-dust candidate} \\label{sec:N5024FIS2}  We have found a candidate of the intracluster dust, N5024 FIS2. Here, we estimate its mass assuming that it is associated with the cluster. The diffuse emission around NGC~6402 can be associated with the cluster as discussed in section \\ref{sec:NGC6402}. However, the clumpy background prevents us from a robust mass estimation, so that we do not include it in quantitative discussions below.  The temperature for N5024 FIS2 cannot be determined well based on our data because of the small coverage of wavelength (only in N60 and WIDE-S) and the limited photometric accuracy ($\\pm~20$\\%). The flux ratio of $F_{90}/F_{65} = 1.33\\pm 0.4$ corresponds to $T=55\\pm 15$~[K]. We perform two estimates assuming 70~K and 35~K again. Then, the dust mass of N5024 FIS2 is estimated at $5 \\times 10^{-4}$~M$_\\odot$ or $5 \\times 10^{-3}$~M$_\\odot$ for 70~K and 35~K respectively from equation (\\ref{eq:Mdust}). The former value agrees within a factor of 2 with that of M~15 IR1a \\citep{Boyer-2006}. Under the assumption of the proportionality between $\\psi$ and [Fe/H], the corresponding gas mass is 5~M$_\\odot$ (or 50~M$_\\odot$ if $T_{\\mathrm{d}}=35$~K) so that it is made up of the dust released from a number of stars. The obtained lifetime $10^7$~yr (or $10^8$~yr) is longer than the upper limits obtained above. If we assume the same dust temperature as the cases of M~15 and NGC~6356, $\\tau_{\\mathrm{d}}$ for N5024 FIS2 is comparable with the lifetimes for the intracluster dust in these clusters.   There remain a few features to be explained if N5024 FIS2 is actually a dust cloud  within the cluster. It is offset from the cluster center, i.e. the bottom of the gravitational potential, by $52\\arcsec$ $(\\sim 0.8 r_{\\mathrm{h}})$. M~15 IR1a is located about $17\\arcsec$ $(\\sim 0.3 r_{\\mathrm{h}})$ to the west of the cluster core. \\citet{Boyer-2006} argued that the intracluster dust is short-lived and has not been dynamically relaxed. This can be true for N5024 FIS2 too. Another intriguing feature for our case is that N5024 FIS2 looks like a point source. The estimated mass suggests that a couple of dozen mass-losing stars contributed to N5024 FIS2. The dust clouds released from mass-losing stars aggregated into one blob and it is constrained in a rather small region providing N5024 FIS2 is associated with the cluster. On the other hand, if the dust-to-gas ratio $\\psi$ does not follow the proportionality (see section \\ref{sec:LackDust}), the estimated dust mass of N5024 FIS2 does not necessarily exceed the amount released from one object. With $\\psi = 0.01$ and $T_{\\mathrm{d}}=70$~K assumed, the estimated amount of the corresponding gas reduces to 0.05~M$_\\odot$. The size evolution of ejecta from a single AGB star is discussed by \\citet{Wareing-2007}. In the interaction with the interstellar medium, the size of circumstellar shell can be restricted to within a few pc, which is as small as the FWHM ($\\sim 40\\arcsec$) of the FIS at the distance of NGC~5024 (17.8~kpc). It is essential to explain these features in order to conclude the intracluster-dust origin of N5024 FIS2. Theoretical work may be useful to study how the materials released from red giants evolve in the intracluster space and interact with each other.   \\subsection{ Impacts on evolution of globular cluster }  As mentioned in the Introduction, intracluster matter may have played an important role in the chemical evolution of globular clusters. \\citet{Tsujimoto-2007} calculated that a significant amount of the mass released from intermediate-mass AGB stars can be accreted to surrounding stars in NGC~5139 during 2~Gyr. However, it is possible that no significant mass remains to be accreted if the intracluster matter disappears with the short lifetime mentioned above. Of course, it is not dust but gas that should be considered as the source of the mass to be accreted because \\citet{Tsujimoto-2007} are interested in the enrichment of helium. The fate of the intracluster gas can be rather different from that of the intracluster dust. Previous observations, however, suggest that the intracluster gas is also deficient in globular clusters as mentioned in the introduction. The deficiency was also reported for NGC~5139 \\citep{Smith-1990}. The conclusion of \\citet{Tsujimoto-2007} can be altered if the intracluster gas is removed before it is accreted to the stars within the cluster. Likewise, the fate of the intracluster dust can have an effect on the possible chemical enrichment via the ejecta of evolved stars especially for heavy elements depleted in dust. For example, aluminum abundance pattern has been also explained by the ejecta of AGB stars (\\cite{Denissenkov-1998}, \\cite{Ventura-2008}).  In any case, it is of interest to investigate the evolution, destruction, and/or removal of the intracluster gas and dust in the context of their effects on the evolution of globular clusters. It is possible that globular clusters were embedded in dark matter halos once in the early universe and the intracluster-matter evolution was different from that in current clusters \\citep{Bekki-2006}. Nonetheless, observation of current ones can provide important constraints to be compared with detailed modeling.  \\subsection{Future prospects}\\label{sec:future}  Besides pointed observations as those presented here, AKARI accomplished the All-Sky Survey. The expected sensitively is lower than those of the pointed observations. The 5~$\\sigma$ detection limit in the WIDE-S filter is 0.55 Jy for a point source \\citep{Kawada-2007}. Nonetheless, it is similar to that in previous observations made with the ISO which resulted in a couple of detections (\\cite{Hopwood-1998}; \\cite{Evans-2003}). The All-Sky Survey data will enable us to search for the intracluster dust in all the globular clusters with a reasonable detection limit.  On the other hand, with the high sensitivity of the FIS pointed observations, a significant number of background galaxies appear above the detection limit. In the future surveys, it can become a problem to discriminate the intracluster dust from the background galaxies unless they are more extended than expected for galaxies. A potentially useful method is to combine the radio observation for the sources. It is well known that the FIR-radio correlation holds for various kinds of galaxies \\citep{Dickey-1984}. \\citet{Oyabu-2005} explored the correlation for the far-infrared sources in the Lockman Hole region based on the ISO's survey (90 and $170~\\micron$) and 1.4~GHz observations with Very large Array (VLA). We can obtain $\\log (F_{90\\micron} / F_{\\mathrm{1.4GHz}}) \\sim 2.8$ based on their data. This value increases to about 6.2 if we consider the blackbody radiation of 70~K. Therefore, the radio flux clearly discriminates between the intracluster dust and background galaxies assuming that intracluster dust has a temperature of 70~K as M~15 IR1a. As for the source we detected, $F_{90\\micron}$ of about 100 mJy predicts $F_{\\mathrm{1.4GHz}} \\sim 700$ $\\mu$Jy for the galaxies and $\\sim 60$~nano-Jy for the blackbody radiation of 70~K. Though the expected $F_{\\mathrm{1.4GHz}}$ value is lower than the detection limits of the wide-field survey catalog currently available even for the case of galaxies, deep follow-up observations may tell which is the case.     We obtained far-infrared images for 12 globular clusters with the AKARI/FIS, which allowed us to do deep surveys of the cold intracluster dust. While we found one candidate of the intracluster dust, we confirmed the paucity of such dust in most clusters.  In the FIS images, we detected extended emissions toward two clusters, namely NGC~2808 and NGC~6402. They are separated from the galactic plane by less than 15 degrees, and large-scale emission of the galactic cirrus is dominant in their fields. We suggest that the detected bumps of emission are merely fluctuations of the background emission. In case of NGC~2808, the emission is largely away from the cluster center, while the bump coincides with the center of NGC~6402. We cannot rule out the possible dust component associated with NGC~6402.  On the other hand, we detected 28 point-like sources, but concluded that most of them are background galaxies. For those with the near-to-mid infrared data by the AKARI/IRC, the obtained SEDs resemble those of galaxies. We found one source, N5024 FIS2, without any clear counterpart in the mid infrared. There is no galaxy found in the HST/ACS image at around the coordinate. It is a possible candidate of the intracluster dust. If so and the temperature of 70~K is assumed, the dust mass is estimated at $5 \\times 10^{-4}$~M$_\\odot$. However, some intriguing features such as the point-like appearance remain to be explained if N5024 FIS2 actually belongs to the cluster. Besides the candidate, our results give upper limits of 2--9 $\\times 10^{-5}$~M$_\\odot$ within the half-mass radii of the 12 clusters from the background fluctuation.  Some mechanism(s) are clearly efficient to remove the intracluster dust in most clusters, if not all. The lifetime of the intracluster dust is as short as $10^{5-6}$ yr. On the other hand, we found that the intracluster dust detected in M~15 and NGC~6356 is exceptionally large compared with other clusters. Our result strengthens the mystery of the lack of the intracluster dust. It urges us to study physical processes which determine the fate of the intracluster dust. Together with that of the intracluster gas, it has a potential to affect the evolution of globular clusters.  \\bigskip  We acknowledges the anonymous referee whose helpful comments gave us a chance to improve the paper. We thank Shinki Oyabu for his suggestions about the usefulness of radio follow-up observation. We are grateful to Iain McDonald for helpful comments based on his analysis on the Spitzer Space Telescope data."
    },
    "0809/0809.1899_arXiv.txt": {
        "abstract": "We consider the generation of gravitational waves by primordial helical inverse cascade magnetohydrodynamic (MHD) turbulence produced by bubble collisions at the electroweak phase transition. We extend the previous study~\\cite{kgr08} by considering both currently discussed  models of MHD turbulence. For popular electroweak phase transition parameter values, the generated gravitational wave spectrum is only weakly dependent on the MHD turbulence model. Compared to the unmagnetized electroweak phase transition case, the spectrum of MHD-turbulence-generated gravitational waves peaks at lower frequency with larger amplitude and can be detected by the proposed Laser Interferometer Space Antenna. ",
        "introduction": "The direct detection of a gravitational wave background will provide us with a probe of physical conditions in the early Universe at the epoch when the gravitational radiation was generated. This is because gravitational waves are massless and after generation weakly-coupled gravitational radiation propagates freely. Hence, once generated, any gravitational wave spectrum retains its shape, with all wavelengths simply scaling with the expansion of the Universe. Gravitational wave astronomy holds the potential for helping construct a picture of the Universe at energy scales even higher than those the Large Hadron Collider will reach, to probe in detail physics at the electroweak energy scale.  There are various mechanisms that might generate gravitational waves in the early Universe, for reviews see~\\cite{source,m00}. A well-known one is parametric amplification of quantum fluctuations~\\cite{inflation}, which can also take place shortly after inflation~\\cite{shortafter}.  Other mechanisms include bubble wall motion and collisions during phase transitions; gravitational wave generation during the electroweak and QCD phase transitions are discussed in Refs.~\\cite{bubble,kos1,huber,nicolis,cds06}. Cosmic strings and other defects can produce relic gravitational waves~\\cite{strings}. Cosmological magnetic fields~\\cite{magnet,cdk04,cd06} and hydrodynamical or magnetohydrodynamical turbulence can also induce primordial gravitational waves~\\cite{kmk02,dolgov,kgr05,gkk07,kgr08}.  Gravitational waves will also provide information on symmetries present in the Universe. In particular, the relic gravitational wave background carries information about parity asymmetry, that might be generated at the electroweak phase transition, or earlier. Parity symmetry testing based on the direct detection of gravitational waves in the near future (with currently planned missions) seems to be promising for gravitational waves with frequencies around the Laser Interferometer Space Antenna (LISA) sensitivity band at $0.1 - 100$ mHz~\\cite{lisa}. The Standard Model electroweak phase transition, when parity symmetry breaking and magnetic helicity production might be expected~\\cite{helicity}, produces gravitational waves with a significant amplitude in this frequency range. Recently, some of us have investigated the generation of gravitational waves through inverse cascade MHD turbulence~\\cite{kgr08} and found that even a small amount of initial helicity enlarges the range of phase transition parameters for which gravitational waves are potentially detectable by LISA. This gravitational wave signal has an important feature: it is circularly polarized. For direct cascade hydrodynamical turbulence some of us have previously estimated the polarization degree~\\cite{kgr05}, which is obviously turbulence model dependent. Polarized gravitational waves are present in other models~\\cite{gw-pol}, and the polarization of the gravitational wave background is in principle observable, either directly~\\cite{seto} or through the cosmic microwave background~\\cite{cdk04,a}.  The detection prospects of gravitational radiation from the early Universe depend on the energy scale when gravitational waves were generated and even more  so on the source duration time and efficiency of the generation process itself. In particular, to produce a detectable signal the electroweak phase transition must be strong enough~\\cite{detection,gs06,grojean}. The currently popular standard model of particle physics does not have a strong enough electroweak phase transition. Also, all currently discussed electroweak phase transition models do not predict an observable gravitational radiation signal if one assumes only bubble collisions~\\cite{grojean}. It is of particular interest to study gravitational wave production during phase transitions in the minimal and next to minimal supersymmetric standard models (MSSM and nMSSM \\cite{models}), \\cite{huber}. The gravitational wave signal produced in the MSSM model is below the LISA sensitivity, while the nMSSM results in a stronger electroweak phase transition \\cite{models} and as a consequence a detectable  gravitational wave signal \\cite{huber}.  On the other hand, sources associated with stochastic vector fields in plasma, i.e.\\ kinetic velocity or magnetic fields, significantly change gravitational wave detection prospects~\\cite{kkgm08}.  In the case of unmagnetized hydrodynamical turbulence the peak frequency of the gravitational wave power spectrum is determined by the characteristic time of turbulence, i.e. the inverse turn-over time of the largest eddy. Another important characteristic is the energy scale when gravitational radiation is generated. In recent modifications of the standard electroweak phase transition model where the transition is moved to a higher energy scale~\\cite{gs06,rs06}, the gravitational wave power spectrum peak frequency is shifted to a higher frequency, which, since the gravitational wave spectrum is sharply peaked, reduces the possibility of detection by LISA. The presence of a cosmological magnetic field affects the dynamics of turbulence. In the case of MHD turbulence the presence of an energy inverse-cascade leads to an increase in the effective size of the largest eddy (now associated with a helical magnetic field), and can result in the gravitational wave power spectrum peaking in the LISA band, with amplitude large enough to be detected by LISA~\\cite{kgr08}.  In this paper we extend the study of Ref.~\\cite{kgr08} by also considering a different model of MHD turbulence. In particular, Ref.~\\cite{kgr08} adopts the inverse cascade model of Refs.~\\cite{BM99,CHB05}, and in this paper we also include in the consideration the model of Refs.~\\cite{son,jedamzik,campanelli}. In both cases,  gravitational wave generation is considered assuming non-zero magnetic helicity during the phase transition.  For our analysis in this paper we adapt the technique developed in Ref.~\\cite{gkk07}. We model MHD turbulence and obtain the gravitational wave spectrum by using an analogy with the theory of sound wave production by hydrodynamical turbulence~\\cite{L52,P52,G,my75}. We perform the computation of the gravitational wave spectrum in real space, instead of using conventional Fourier space techniques as in Refs.~\\cite{kmk02,cds06}. This makes the physical interpretation of all quantities straightforward. We use natural units with $\\hbar = c = k_B = 1$ throughout.  The paper is organized as follows. In Sec.~II we summarize the general formalism of gravitational wave generation. In Sec.~III we describe the turbulence models we consider. In Sec.\\ IV we derive the gravitational wave spectra, and we conclude in Sec.\\ V. ",
        "conclusions": "We have analyzed the generation of gravitational waves at the electroweak phase transition, from hydrodynamic turbulence in the presence of a helical magnetic field. It is convenient to consider this as a two-stage process. During the first stage the effects of helicity are negligible when studying the production of gravitational waves from turbulence. The first stage ends when quasi-conservation of magnetic helicity starts to trigger an inverse cascade of the magnetic field, that is, a transfer of magnetic energy from small to large scales. The details of the inverse cascade process in helical MHD are still under debate and there exist in the literature two different models. In this paper we have examined gravitational wave generation in both models of helical MHD turbulence. We have found that, for realistic values of the parameters defining the electroweak phase transition, the generated gravitational wave spectrum  is independent of turbulence modeling (within a factor of order 2).  The inverse-cascade associated peak frequency always coincides with the Hubble frequency at the moment of gravitational wave production. Moreover, as clear from Fig.~1, the main contribution to the total gravitational waves energy density is from the second, inverse-cascade stage, even for small values of magnetic helicity.  Gravitational waves produced {\\it via} helical turbulence are strongly polarized, since magnetic helicity is a parity-odd quantity and is maximal at the end of the first stage~\\cite{kgr05}. LISA should be able to detect such a gravitational wave polarization~\\cite{seto}. Also, in contrast to the unmagnetized case, the contribution to the gravitational wave amplitude coming from the inverse-cascade stage is large enough at $0.1$ mHz to be detectable by LISA. If the electroweak phase transition occurs at higher energy (see Fig.~2) the peak frequency is higher, closer to the LISA sensitivity peak, which leads to a stronger signal. Our formalism is also applicable to gravitational wave production at an earlier QCD phase transition, assuming the presence of colored magnetic fields~\\cite{QCD}, or to any other phase transition~\\cite{bubble}: the peak frequency will shift according to  changes in $T_\\star$ and $g_\\star$.  Finally, we stress that the gravitational wave signal arising from helical MHD turbulent motion exceeds that from bubble collisions~\\cite{kos1,detection,gs06} and that  from hydrodynamical turbulence alone~\\cite{nicolis,gkk07}. Of course, the strong signal estimated here (see also Ref.~\\cite{kgr08}) assumes initial non-zero (although small) magnetic helicity, so detection of polarized gravitational waves by LISA will indicate parity violation during the phase transition, as proposed in Refs.~\\cite{helicity}."
    },
    "0809/0809.3252_arXiv.txt": {
        "abstract": "Several lines of evidence suggest that the galaxy cluster Cl0024+17, an apparently relaxed system, is actually a collision of two clusters, the interaction occurring along our line of sight. Recent lensing observations suggest the presence of a ring-like dark matter structure, which has been interpreted as the result of such a collision. In this paper we present $N$-body simulations of cluster collisions along the line of sight to investigate the detectability of such features. We use realistic dark matter density profiles as determined from cosmological simulations. Our simulations show a ``shoulder'' in the dark matter distribution after the collision, but no ring feature even when the initial particle velocity distribution is highly tangentially anisotropic ($\\sigma_\\theta/\\sigma_r >> 1$).  Only when the initial particle velocity distribution is circular do our simulations show such a feature.  Even modestly anisotropic velocity distributions are inconsistent with the halo velocity distributions seen in cosmological simulations, and would require highly fine-tuned initial conditions.  Our investigation leaves us without an explanation for the dark matter ring-like feature in Cl 0024+17 suggested by lensing observations. ",
        "introduction": "Clusters of galaxies have proven to be interesting laboratories for studying the dynamics of the different kinds of matter in the universe. Early measurements of galaxy velocity dispersions in clusters demonstrated that most of the mass of galaxy clusters was in the form of a non-luminous component \\citep{zwy37}. Modern cosmological constraints (e.g. Turner 2000) suggest that this non-luminous material must be mostly non-baryonic. In addition, X-ray observations of clusters of galaxies reveal that most of the baryonic mass in clusters is in the form of a hot, diffuse, X-ray emitting gas, the intracluster medium (ICM). The presence of these different types of matter (stellar, gaseous, and dark) in these systems provides not only insights into the nature of the clusters themselves but also the nature of the different components of matter and their dynamics.  Galaxy clusters are themselves the product of many mergers between galaxies, galaxy groups, and smaller clusters of galaxies \\citep{dav85}. Observations indicate that this process is still ongoing in many systems. A famous example is the so-called ``Bullet Cluster'' (1E 0657-56), a merging system in which X-ray and weak lensing observations demonstrate a clear separation between a collisionless dark matter component and the intracluster medium \\citep{mar02, clo06}.  A system that has also garnered attention recently is Cl 0024+17, a cluster at $z = 0.395$ with weak and strong lensing. Early attempts at reconstructing the mass profile of this system using lensing \\citep{tys98, bro00, com06} suggested a conflict with the predictions of the standard cold dark matter (CDM) model due to the flattening of the density profile in the inner regions of the cluster. Cosmological simulations assuming CDM indicate that galaxy cluster mass profiles should exhibit a logarithmic slope in the inner regions (e.g. Navarro, Frenk, \\& White 1997). This apparent discrepancy (and others) led to suggestions that the CDM paradigm would need to be modified \\citep{spr00,hog00,moo00}, for example to include self-interaction of the dark matter.  \\citet{czo01, czo02} demonstrated that the redshift distribution of the cluster galaxies in Cl~0024+17 is bimodal and suggested that the system is composed of two clusters undergoing a collision along the line of sight. They also performed a simulation demonstrating that such a scenario reproduces not only the bimodal redshift distribution but also the flattening of the central density profile. Observations of the cluster by Chandra \\citep{ota04} and XMM-Newton \\citep{zhg05} revealed that the surface brightness profile is better fit by a superposition of two ICM models rather than one. They suggested on the basis of the isothermal temperature profile that the system had returned to equilibrium after the collision and that consequently the encounter must have occured several Gyr ago.  Recently, a weak lensing analysis presented by \\citet{jee07} revealed a ring-like structure in the projected matter distribution. They proposed that dark matter from the cores of the clusters had been disrupted and ejected from the systems by the collision, and they demonstrated that such features could be reproduced in a simulation of a collision of two pure dark matter halos. On the basis of this result they suggested the current state of the system is $\\sim$1-2 Gyr past the pericentric passage of the cluster cores.  We wish to shed additional light on the formation of ring structures within the dark matter distribution. ``Ring galaxies'' provide another astrophysical situation in which ``particle'' rings are produced due to collisions (e.g., the Cartwheel Galaxy). We look to the previous $N$-body simulations of these phenomena to guide our understanding of what might lead to such a feature in the dark matter in clusters of galaxies. Two major differences exist between the dark matter particles in galaxy clusters and star ``particles'' in disk galaxies. In the case of disk galaxies, the stars are concentrated in a disk, and the velocities are essentially circular. In galaxy clusters, dark matter particles form a roughly spheroidal distribution and have isotropic to radially anisotropic velocity dispersions. To investigate the effects of the second of these differences we perform a set of pure dark matter ($N$-body only) simulations with a physically motivated dark matter profile and variations on the velocity distribution of the dark matter particles.  Throughout this paper we assume a flat $\\Lambda$CDM cosmology with $h = 0.7$ and $\\Omega_{\\rm m} = 0.3$. ",
        "conclusions": "A recent lensing analysis of the cluster Cl 0024+17 revealed possible evidence of a ring-like dark matter structure. This has been interpreted as the result of a collision between two clusters of galaxies along the line of sight, a scenario for which there are other lines of evidence in this system. We have performed simulations of collisions between galaxy cluster dark matter halos to test this hypothesis, using density profiles for the dark matter halos motivated by observations and simulations and investigating a parameter space of initial velocity distribution where we vary the initial velocity anisotropy. Our simulations show that a more pronounced ``shoulder'' feature appears in the projected dark matter density when we make the velocities of the dark matter particles more tangentially anisotropic, by analogy with the phenomena of \"ring galaxies.\"  Our simulations show that, although the collision ejects a large amount of mass from the dark matter cores, a ring-like features does not form, even for initial velocity distributions that are highly tangentially anisotropic.  Only when the initial velocity distribution is circular does a ring form.  Since we do not expect dark matter particles in clusters of galaxies to have tangentially anisotropic velocities, our investigation of this parameter space leaves us without an explanation for the ring-like feature that appears in the dark matter distribution of Cl~0024+17."
    },
    "0809/0809.4228_arXiv.txt": {
        "abstract": "After a brief  non technical introduction of the basic properties of strange quark matter (SQM) in compact stars, we consider some of the late important advances in the field, and  discuss some recent astrophysical observational data that could shed new light on the possible presence of SQM in compact stars. We show that above a threshold value of the gravitational mass a neutron star (pure hadronic star) is metastable to the decay (conversion) to an  hybrid neutron star or to a strange star. We explore the consequences of the metastability of ``massive'' neutron stars and of the existence of stable compact ``quark'' stars (hybrid neutron stars or strange stars) on the concept of limiting  mass of compact stars, and we give an extension of this concept with respect to the {\\it classical} one given in 1939 by Oppenheimer and Volkoff. ",
        "introduction": "\\label{intro} In 1964 Murray Gell-Mann\\cite{gell-mann} and George Zweig\\cite{zweig}  proposed that the proton, the neutron, and all the other hadrons ({\\it i.e.} particles which feel the {\\it strong interaction}) are not elementary, but are composed of smaller (point-like) constituents which were called {\\it quarks}. \\footnote{Gell-Mann borrowed this name from a line of the novel {\\it Finnegans Wake} by James Joyce.} In the original model there are three fundamental quarks nicknamed {\\it up} ({\\it u}) {\\it down } ({\\it d}) and {\\it strange} ({\\it s}). The premiere evidence for the existence of quarks came already at the end of 1967 from the first high energy electron-proton scattering experiment\\cite{slac} performed at the Stanford Linear Accelerator Center (SLAC),  and  a growing body of direct evidence for their  reality \\footnote{Isolated quarks have never been observed. This has led to the quark confinement hypothesis, as one of the basic features of Quantum Chromodynamics (QCD), which is the fundamental theory of strong interactions.} was accumulated in the following years in numerous experiments in different laboratories all around the world.  It was then natural to speculate that when nuclear matter is compressed to densities so high that nucleons substantially overlap, a new phase of matter, in which quarks are the relevant  degrees of freedom, could be formed. Already in 1965, Ivanenko e Kurdgelaidze\\cite{ik65} suggested the possible existence of a ``quarkian core'' in very massive stars.  Two years later, at the end of 1967,  Jocelyn  Bell discovered\\cite{PSR}  the first pulsar. Pulsars  were soon interpreted\\cite{pac-gold}  as strongly magnetized rotating neutron stars,  the compact stellar remnants of supernova explosions hypothesized\\cite{bz} in 1934  by Baade and Zwicky \\footnote{In 1937 Landau suggested the existence of a ``neutron core'' (neutron matter core) in normal stars, to explain the energy source of stars as due to the release of gravitational energy via matter accretion  into the  stellar  neutron core.} soon after the Chadwick's discovery of the neutron. The first calculation of the structure on a neutron star was performed in 1939  by Oppenheimer and Volkoff\\cite{ov39}.  They assumed the star to be composed by pure neutron matter described as a non-interacting relativistic Fermi gas. More sophisticated calculations, with the inclusion of the effects of nuclear interactions on the equation of state (EOS) of neutron matter, were available\\cite{cam59} at the end of 1950s. A more refined model for the internal composition of neutron stars was introduced in 1960 by Ambartsumyan and Saakyan\\cite{amb60}, which suggested the possible presence of hyperons ($\\Lambda$, $\\Sigma^{-}$, $\\Sigma^{0}$, $\\Sigma^{+}$, $\\Xi^{-}$ and $\\Xi^{0}$ particles) in the inner core of neutron stars. All these early  models gave,  with a reasonable accuracy,  the gross properties ({\\it i.e.}  maximum mass, radii, and central densities) of neutron stars. Particularly, it was clear that the central density of the maximum mass configuration could be as high as 10 times the central density ($\\rho_0 = 2.8 \\times 10^{14} {\\rm g/cm}^3$)  of heavy atomic nuclei. Thus it was quite evident at the end of the 1960s and in the early 1970s, that neutron stars were the best candidate in the universe where quark matter could be found. It was in those years that the idea of pure quark stars\\cite{ito70} or hybrid hadronic-quark stars\\cite{ik69}  (termed ``baryon-quarkian stars'' by the authors of ref. \\cite{ik69}) was conceived. The earliest studies\\cite{cp75,bc76,kk76}, however  indicated that it was unlikely that quark matter could be found in stable neutron stars, or to have a third family of quark compact stars. This conclusion was mainly due to some simplification in the treatment of the quark-deconfinement phase transition, which was considered as one at constant pressure. Further investigations\\cite{gle92,gle_book} have established that,  since neutron star matter is a multicomponent system with  two conserved ``charges'' ({\\it i.e.,} electric charge and baryon number), the quark-deconfinement phase transition proceeds through a mixed phase over a finite range of pressures and densities according to the Gibbs' criterion for phase equilibrium. These \\cite{gle92,gle_book} and subsequent studies have established that neutron stars may, very likely, contain quark matter in their interiors. Neutron  stars  which possess a quark matter core either as a mixed phase \\footnote{see however ref.s\\cite{heisel93,voskr02} for later studies on the mixed hadron-quark phase, where the effects of Coulomb, surface energies, and charge screening are taken into account.} of deconfined quarks and hadrons or as a pure quark matter phase are called {\\it hybrid neutron stars} or shortly {\\it hybrid stars}\\cite{gle_book}  (HyS). The  more {\\it conventional} neutron stars  in which no fraction of quark matter is present, are currently referred to as  {\\it pure hadronic stars} (HS).  A further crucial step in the study of quark matter in astrophysical contest was made by Arnold R. Bodmer in 1971.  In his pioneering work\\cite{bod71}, Bodmer suggested the possibility of {\\it collapsed nuclei}. He conjectured that collapsed nuclei could have a lower energy than normal ({\\it i.e.} made of protons and neutrons) nuclei, and he speculated that normal nuclei are very long-lived  {\\it isomers} against collapse because of a ``saturation '' barrier between normal and collapsed nuclei. Among other possibilities, Bodmer discussed collapsed nuclei as composite states of 3A  {\\it u}-{\\it d}-{\\it s} quarks (being A the baryon number of the nucleus). These many-body {\\it u}-{\\it d}-{\\it s} quark states are what we call, with modern terminology,  {\\it strangelets} or {\\it quark nuggets}. Bodmer also speculated that {\\it collapsed nuclei may have been copiously produced in the initial extremely hot and dense stages of the universe},  and proposed that collapsed nuclei may condensate as peculiar very  compact massive ``black''  objects being possible candidates  for dark matter.  The feasible stability of multi-quark systems was also investigated\\cite{ter79}  by Hidezumi Terazawa, who termed them {\\it super-Hypernuclei}. However, these ideas brought a wide attention by the scientific community only after Edward Witten published, in 1984,  his widely known paper\\cite{witt84} on  the ``cosmic separation of phases'', where he examined  the cosmological and astrophysical consequences of the {\\it absolute stability} of strange quark matter (SQM){\\footnote{see next section for the definition of strange quark matter.}, and calculated the bulk properties of {\\it Strange Stars} within a simple model for SQM inspired to the MIT bag model for hadrons. In the following years the properties of SQM and the properties of strange stars were calculated by numerous groups\\cite{afo86,hzs86,bh89}.  A recent valuable advance  in the comprehension of quark matter properties, that is having a substantial impact on the study of neutron star physics, is the  discovery of color superconductivity. Superconductivity, as it is well known, is a general feature of degenerate Fermi systems, which become unstable if there exist any attractive interaction at the Fermi surface.  As recognized by Bardeen, Cooper and Schrieffer (BCS) this instability leads to the formation of a condensate of Cooper pairs and to the appearance of superconductivity. In Quantum Chromodynamics (QCD) the quark-quark (qq) interaction is strongly attractive in many channels. This will lead to quark pairing and color superconductivity. This possibility was already  pointed out\\cite{cp75,bar77,fra78} in the mid 1970s (see also ref.\\cite{bl84}).  Very recently the study of color superconductivity has become a research subject full of activity, since the realization that the typical superconducting energy gaps in quark matter may be of the order of $\\Delta \\sim 100$ MeV, thus much higher than those predicted in the early works.  The phase diagram of QCD has been widely analyzed in the light of color superconductivity\\cite{Raj-Wil00,Nard01,Alf01,Cas-Nard04,sho05,bub05,Schaf05,Rust-etal06,mann07}. Quarks, unlike electrons, have different flavors  ({\\it u}, {\\it d}, {\\it s} for the quark matter relevant for neutron star physics) and have color as well as spin degrees of freedom, thus many different quark pairing schemes are expected. At asymptotic densities (in the asymptotic freedom regime where perturbative QCD is valid) the ground state of QCD is the so called color-flavor locked (CFL) phase, in which all three quark flavors participate to pairing symmetrically {\\footnote{This phase of SQM is electrically charge neutral without any need for electrons.}. In the density regime relevant for neutron star physics (well below the asymptotic density regime) the ground state of quark matter is uncertain and many possible patterns for color superconductivity have been predicted.  \\begin{figure}[t] \\begin{center} \\psfig{file=hybrid_star_sect.eps,width=3.4in} \\end{center} \\caption{Schematic cross section of a hybrid star (see text for more details). } \\label{aba:cross_sect_hybr} \\end{figure} ",
        "conclusions": ""
    },
    "0809/0809.3593_arXiv.txt": {
        "abstract": "It is established that the formation of rotationally supported disks during the main accretion phase of star formation is suppressed by a moderately strong magnetic field in the ideal MHD limit. Non-ideal MHD effects are expected to weaken the magnetic braking, perhaps allowing the disk to reappear. We concentrate on one such effect, ambipolar diffusion, which enables the field lines to slip relative to the bulk neutral matter. We find that the slippage does not sufficiently weaken the braking to allow rotationally supported disks to form for realistic levels of cloud magnetization and cosmic ray ionization rate; in some cases, the magnetic braking is even enhanced. Only in dense cores with both exceptionally weak fields and unreasonably low ionization rate do such disks start to form in our simulations. We conclude that additional processes, such as Ohmic dissipation or Hall effect, are needed to enable disk formation. Alternatively, the disk may form at late times when the massive envelope that anchors the magnetic brake is dissipated, perhaps by a protostellar wind. ",
        "introduction": "Disk formation is an integral part of star formation that has been studied for a long time. Early works concentrated on the collapse of rotating, non-magnetic cores \\citep[e.g.,][]{1984ApJ...286..529T}. \\citet{1995ARA&A..33..199B} reviewed these works, and noted a number of unresolved problems \\citep[see also][]{1998ASPC..148..314B}. Topping the list is the effect of magnetic braking on disk formation.  The magnetic properties of star-forming dense cores are reasonably well constrained. There is now ample evidence for ordered magnetic fields on the core scale from polarization of dust emission \\citep[e.g.,][]{2000ApJ...537L.135W}. A spectacular recent example is the Submillimeter Array (SMA) observation of NGC 1333 IRS 4A, which shows a pinched magnetic configuration on the $10^3$~AU scale \\citep{2006Sci...313..812G}. The magnitude of the magnetic field is harder to determine. \\citet{2008ApJ...680..457T} carried out an extensive Zeeman survey of dark cloud cores in OH, and determined a mean mass-to-flux ratio of $\\lambda \\sim 4.8$ in units of the critical value $(2\\pi G^{1/2})^{-1}$ \\citep{1978PASJ...30..671N}; correction for uncertain projection effects may bring this value to $\\sim 2$. The implication is that dense cores are typically moderately strongly magnetized, with a dimensionless mass-to-flux ratio of a few to several. A relatively low value of $\\lambda \\sim 2-3$ is expected for dense cores formed out of magnetically subcritical clouds (with $\\lambda < 1$) through ambipolar diffusion \\citep[e.g.,][]{1989ApJ...342..834L,1994ApJ...432..720B,2008arXiv0804.4201N,2007ApJ...671..497A}. The values of $\\lambda$ are expected to have a wider spread for cores formed out of turbulent compression. Nevertheless, the cores tend to be more magnetized relative to their masses than the cloud as a whole, because only a fraction of the mass along a flux tube that threads the cloud is compressed into the core \\citep[e.g.,][]{2007ApJ...661..262D,2005prpl.conf.8473T}. It is unlikely for the cores to have a value of $\\lambda$ more than several, unless the cloud as a whole is magnetized to an unrealistically low level.  There have been a number of recent calculations that included magnetic fields in the collapse of rotating cores, focusing mostly on the early phase before the mass of the central object reaches the stellar range \\citep[e.g.,][]{1998ApJ...502L.163T,2006ApJ...647L.151M,2006ApJ...641..949B,2007Ap&SS.311...75P}. Other studies have concentrated on the angular momentum evolution and disk formation in the main accretion phase (\\citealp{2003ApJ...599..363A,2006ApJ...647..374G}; \\citealp[hereafter Paper I]{2008ApJ...681.1356M}; \\citealp{2008A&A...477....9H}). These studies find that magnetic braking efficiently removes angular momentum from the collapsing matter, preventing a large, 10$^2$-AU scale rotationally supported disk from forming in realistic cores in the ideal MHD limit. However, such disks are routinely observed around young stars, at least in relatively late, Class I and II phases. How can the disk be saved?  Molecular cloud cores are lightly ionized, so perfect coupling between the magnetic field and matter assumed in ideal MHD is not expected. As the density increases, the coupling is weakened first by ambipolar diffusion, then the Hall Effect, then Ohmic dissipation. In this paper we will concentrate on ambipolar diffusion as a first step. The effect of ambipolar diffusion on magnetic braking in the formation \\citep{1994ApJ...432..720B} and collapse \\citep{2002ApJ...580..987K} of rotating cores has been examined previously using the thin disk approximation. \\citet{2002ApJ...580..987K} parametrized the braking strength and found that disk formation can be suppressed in the ambipolar diffusion limit when the braking parameter is large. We use two-dimensional models which allow direct calculation of the strength of the magnetic braking. We find that ambipolar diffusion does not weaken the magnetic braking enough to allow rotationally supported disks to form under realistic conditions. The implication is that additional processes must be found to save the disk; two possibilities are discussed in section 5. ",
        "conclusions": "We have studied the collapse of rotating molecular cloud cores magnetized to moderate degrees, concentrating on a dimensionless mass-to-flux ratio $\\lambda=4$, as suggested by recent Zeeman observations \\citep{2008ApJ...680..457T}. In the ideal MHD limit, we have previously shown that a slight twist of the field lines in such cores is sufficient to remove essentially all of the angular momentum of the collapsing matter and suppress the formation of rotationally supported disks during the protostellar accretion phase completely \\citep[Paper I ; see also][]{2006ApJ...647..374G,2008A&A...477....9H}. One may expect ambipolar diffusion to reduce the efficiency of magnetic braking and potentially enable the disks to form, since the twisting of field lines by the rotation of neutrals will be reduced by diffusion. However, for realistic cosmic ray ionization rates, the small amount of field twisting needed for angular momentum removal drives a drift speed between the ions (and the field lines tied to them) and neutrals in the azimuthal direction that is much smaller than the neutral rotation speed (see the lower panel of Fig.~\\ref{standard}). By itself, the azimuthal slippage of field lines relative to neutrals does not reduce the strength of the toroidal magnetic field (and the braking rate) significantly. The ambipolar diffusion is expected to modify the poloidal magnetic field as well, especially in regions where the magnetic force is large. The modification is most evident at small radii, where magnetic diffusion has allowed the magnetic flux that would have been dragged onto the central object to occupy an extended region \\citep{1996ApJ...464..373L,1998ApJ...504..257C,2007ApJ...660..388T}. The increased poloidal field strength tends to make the magnetic braking more efficient \\citep{2002ApJ...580..987K}. Given these competing effects in opposite directions, it not at all obvious whether ambipolar diffusion would increase or decrease the overall efficiency of magnetic braking. Our detailed calculations have shown that, for realistic rates of cosmic ray ionization, the overall efficiency is increased, because the azimuthal field slippage is small and the poloidal field trapped at small radii is strong. The efficiency is reduced only for unrealistically low ionization rates, which speed up the azimuthal field slippage and lower the poloidal field strength by enabling the trapped flux to escape to large distances.  There is, however, ample evidence for large, rotationally supported disks on $10^2$~AU scale or more around young stars, at least in Class I and II, and possibly in Class 0, sources. This begs the question: what is the origin of these disks?  The most obvious possibility is that additional non-ideal MHD effects may weaken the magnetic braking further. Ohmic dissipation is expected to be important at high densities \\citep{2006ApJ...647..382S}. \\citet{2002ApJ...573..199N} estimates that it dominates other processes at densities above $\\sim 10^{12}$~cm$^{-3}$ \\citep[see also][]{2001ApJ...550..314D}. Since the density in our standard model never exceeds this value, we believe that Ohmic dissipation would not significantly weaken the magnetic braking inside our computation domain (from $10^{14}$~cm$ < r < 2\\times10^{17}$~cm). However, Ohmic dissipation may still enable a rotationally supported disk to form at a smaller radius where the density is higher, as demonstrated explicitly by \\citet{2008ApJ...676.1088M} using a resistive MHD core. They found a disk of $\\sim10^{12}$~cm in radius shortly ($\\sim 11$~days) after the formation of the so-called ``second core''; the disk was not present in the ideal MHD counterpart. Such Ohmic dissipation-enabled disks can form inside our inner boundary (of $10^{14}$~cm). Since the material crossing the inner boundary of our simulation is typically braked to a rotational speed well below the local Keplerian speed, any disk that may form later should be relatively small in the absence of angular momentum redistribution.  Besides Ohmic dissipation, another non-ideal mechanism of magnetic flux diffusion is the Hall Effect. When dust grains abundances are low, the Hall Effect becomes important when collisions with neutrals decouple the ions (but not electrons) from the magnetic field. This occurs when the ion Hall parameter $\\beta_i$ is less than unity \\citep[e.g.,][]{1984FCPh....9..139N}. We have checked that $\\beta_i>1$ everywhere inside our computational domain for the standard model. We do not expect the Hall Effect to change our results significantly, although this may be changed by the inclusion of dust grains. The exact effect depends on the grain size distribution \\citep{1999MNRAS.303..239W}, which is uncertain. Further study is required to determine the importance of the Hall Effect. If the additional non-ideal MHD effects do not allow the observed large scale disks to form, other means of weakening the magnetic braking must be sought.  One possibility is protostellar outflow. Protostellar outflows are observed at all stages of star formation. They are thought to clear away the envelope material as the source ages \\citep{2006ApJ...646.1070A}. As the envelope is the anchor for the magnetic brake in our models, removing it may reduce the strength of braking sufficiently to allow large scale rotationally supported disks to form. This is particularly true for the more evolved class II and perhaps class I sources, where little envelope is left to brake the disk. It may not work for class 0 sources where the outflows tend to be confined to the polar regions (e.g., HH 211, \\citealp{1999A&A...343..571G}). If this is the case, we would not expect these deeply embedded objects to harbor large scale, rotationally supported disks. They could still have relatively small disks enabled by, for example, Ohmic dissipation. It is difficult to determine the exact amount of rotational support in the equatorial region from current observations of Class 0 sources. The situation will improve when ALMA comes online."
    },
    "0809/0809.2595_arXiv.txt": {
        "abstract": "We use Sloan Digital Sky Survey (SDSS) data to investigate galaxy cluster properties of systems first detected within DPOSS. With the high quality photometry of SDSS we derived new photometric redshifts and estimated richness and optical luminosity. For a subset of low redshift ($z \\le 0.1$) clusters, we have used SDSS spectroscopic data to identify groups in redshift space in the region of each cluster, complemented with massive systems from the literature to assure the continuous mass sampling. A method to remove interlopers is applied, and a virial analysis is performed resulting in estimates of velocity dispersion, mass, and a physical radius for each low-$z$ system. We discuss the choice of maximum radius and luminosity range in the dynamical analysis, showing that a spectroscopic survey must be complete to at least M$^*+1$ if one wishes to obtain accurate and unbiased estimates of velocity dispersion and mass. We have measured X-ray luminosity for all clusters using archival data from RASS. For a smaller subset (twenty-one clusters) we selected temperature measures from the literature and estimated mass from the M-T$_X$ relation, finding that they show good agreement with the virial estimate. However, these two mass estimates tend to disagree with the caustic results. We measured the presence of substructure in all clusters of the sample and found that clusters with substructure have virial masses higher than those derived from T$_X$. This trend is not seen when comparing the caustic and X-ray masses. That happens because the caustic mass is estimated directly from the mass profile, so it is less affected by substructure. ",
        "introduction": "Galaxy clusters represent the largest and latest systems to form under the influence of their own gravity. The study of the evolution of the large scale structure through the cluster mass function and how that varies with time, is one of the most important tools in unveiling the global properties of the Universe. The way clusters evolve strongly depend on several cosmological parameters like $\\Omega_m$, $\\Omega_{\\lambda}$, and $\\sigma_8$ \\citep {eke98, bah97, bah03, don98, rei99, bla98}.  The main difficulty encountered in these studies, lies in the exact estimate of the cluster mass. A variety of techniques are in use today, but accurate estimates for large data sets are still impractical, as all methods are expensive from the observational viewpoint.  A simple way to estimate mass for a large sample of clusters is to find an unbiased photometric proxy such as X-ray luminosity (L$_X$), richness (N$_{gal}$), or optical luminosity (L$_{opt}$). L$_X$ and L$_{opt}$ correlate well with mass, though with larger scatter than relations involving X-ray temperature (T$_X$) and velocity dispersions ($\\sigma_P$). The measurement and calibration of these relations are essential for reliably establishing the cluster mass function. As pointed out by \\citet {voi05}, underestimation of the scatter in the mass$-$observable relation leads to overestimation of $\\sigma_8$. Furthermore, contradictory results appear when one relates mass to cluster properties measured at different wavelengths. This maybe reflecting  the selection function of the cluster catalogs created from data in different wavelength regimes, or the physical processes affecting the cluster galaxy population and the intra-cluster gas.  The study of cosmology using galaxy clusters has as a key ingredient the proper recognition of the biases present in multi-wavelength surveys. Some of these issues have been addressed in the recent literature, either by comparing X-ray and optical cluster catalogs or conducting joint X-ray/optical surveys of galaxy clusters \\citep {don01, don02, gil04, pop04, pop05, lop06, dai07, gal08}. Some of these works investigate the complementary properties of clusters selected in other wavelength regimes (optical or X-ray; \\citealt {don02, gil04, pop04, lop06, dai07, gal08}).  Information on the distribution of galaxies, hot gas and dark matter within clusters can play an important role to understand which observational biases contribute to increasing the scatter of the scaling relations. The presence of substructure is a clear sign of incomplete relaxation in a cluster and can be seen from the X-ray emission from the intra-cluster medium \\citep{jon99} as well as in the distribution of cluster galaxies \\citep {pin96, mat99, kol01}. Recent works have investigated the possibility that substructures in clusters increase the scatter of the scaling relations, but the results have been contradictory. \\citet {smi05} find that the observed scaling relations for relaxed and not relaxed clusters are significantly different. However, \\citet {oh06} reach the opposite conclusion, basing their analysis on simulated and observed data. \\citet {lop06} find that the T$_X$-N$_{gals}$ (or T$_X$-L$_{opt}$) relation is severely affected by the presence of substructures, although the relations involving L$_X$ are not.  The main goals of this paper are: (i) collect higher quality (SDSS) data for the NoSOCS supplemental catalog to redshift $z \\sim$ 0.5 (ii) define a low-redshift subset of the NoSOCS sample ($z \\le$ 0.1); (iii) derive unbiased measurements of velocity dispersion and masses for the clusters in this low$-z$ sample; (iv) compare optical and X-ray mass estimates, allowing checks for possible biases in the mass estimates. This paper gives a description of the estimation of cluster properties for the NoSOCS catalog using SDSS. New photometric redshifts, richness and optical luminosity are derived for the full sample. Velocity dispersion, mass, and physical radius are obtained for the low-$z$ clusters. This data set will be used for future cluster studies.  This work relies on the supplement version of NoSOCS (\\citealt {lop03, lop04}), which contains 9,956 cluster candidates over 2,700 square degrees. This catalog, based on the Digitized Second Palomar Observatory Sky Survey (DPOSS, \\citealt {djo03}), includes clusters up to $z \\sim 0.50$ and is the largest catalog to this redshift limit to date. Here, we discuss the use of higher quality photometric data from the SDSS (\\citealt {yor00}) to assess NoSOCS cluster properties. We also used spectroscopic data from the SDSS to study the dynamics of the low redshift systems ($z \\le 0.10$).  The presentation of the paper is as follows: In $\\S$2 we briefly describe the NoSOCS. Selection of SDSS data for the NoSOCS clusters is described in $\\S$3, where we also discuss the estimates of richness and optical luminosity, as well as the effect of different k$-$corrections and evolution in M$^*$ to richness. In $\\S$ 4 we detail the way we characterize the properties of the low-$z$ systems (velocity dispersion, characteristic radius, and mass). Specific issues regarding the method for interloper rejection, as well as the radial cutoff and magnitude limit of the spectroscopic survey are also discussed. We estimate X-ray luminosity for a subset of these local clusters and discuss the estimation of mass for a subset of systems with accurate T$_X$ available. In $\\S,$ 5 we present the substructure tests employed for the clusters at low-$z$. We compare optical and X-ray mass estimates in $\\S$ 6 (as well as optical masses from the caustic method and the virial analysis). We summarize our main conclusions in $\\S$ 7. Throughout this work we assumed a cosmology with $\\Omega_m = 0.3, \\Omega_{\\Lambda} = 0.7$ and H$_0 = 100$ $h$ $km$ $s^{-1}$ Mpc$^{-1}$, with $h$ set to 0.7. ",
        "conclusions": ""
    },
    "0809/0809.2076_arXiv.txt": {
        "abstract": "A catalogue of 14453 radio--loud AGN with 1.4 GHz fluxes above 3.5 mJy in the redshift range $0.4<z<0.8$, has been constructed from the cross--correlation of the NVSS and FIRST radio surveys with the MegaZ--LRG catalogue of luminous red galaxies derived from Sloan Digital Sky Survey imaging data. The NVSS provides accurate flux measurements for extended sources, while the angular resolution of FIRST allows the host galaxy to be identified accurately. New techniques were developed for extending the cross--correlation algorithm to FIRST detections that are below the nominal 1 mJy S/N limit of the catalogued sources. The matching criteria were tested and refined using Monte--Carlo simulations, leading to an  estimated reliability of $\\sim$98.3\\% and completeness level (for LRGS) of about 95\\% for our new catalogue.  We present a new determination of the luminosity function of radio AGN at $z \\sim 0.55$ and compare this to the luminosity function of nearby ($z \\sim 0.1$) radio sources from the SDSS main survey. The comoving number density of radio AGN with luminosities less than $10^{25}$ W Hz$^{-1}$  increases by a factor  $\\sim 1.5$ between $z =0.1$ and $z=0.55$. At higher lumiosities, this factor increases sharply, reaching values of more than 10 at radio luminosities larger than $10^{26}$ W Hz$^{-1}$. We then study how the relation between radio AGN and their host galaxies evolves with redshift. Our main conclusion is that the fraction of radio--loud AGN increases towards higher redshift in all massive galaxies, but the evolution is particularly strong for the lower mass galaxies in our sample. These trends may be understood if there are two classes of radio galaxies (likely associated with the ``radio\"  and ``quasar mode\" dichotomy) that have different fuelling/triggering mechanisms and hence evolve in different ways. ",
        "introduction": "Ever since the first studies of \\citet{longair}, it has been widely accepted that evolution is required to explain the radio source counts. Strong cosmological evolution of the high luminosity radio source population has been established out to at least $z=2-3$ (e.g. \\citealt{dunpeack}), with the space--densities of the most powerful radio sources increasing by a factor $\\sim10^3$ relative to their densities at $z\\sim0$. However, the evolution of the low--luminosity radio galaxy population remained more uncertain. \\citet{clewley} found that the low luminosity population of radio sources with $L_{\\rm 325 MHz} < 10^{25}$ W Hz$^{-1}$sr$^{-1}$ exhibits no evolution out to $z\\sim0.8$. In contrast, \\citet{brown} found significant cosmic evolution out to $z \\sim 0.55$, which they modelled using a pure luminosity evolution model as $P(z) = P(0)(1+z)^{k_L}$, with $3<{k_L}<5$. More recently, similar results were obtained by \\citet{sadler07}, who analyzed radio galaxies in the `2dF--SDSS Luminous Red Galaxy and Quasar' (2SLAQ) survey (\\citealt{cannon}). They found that the number density of radio sources below $\\rm P_{\\rm 1.4 GHz}=10^{26}$ W Hz$^{-1}$ grows by a factor of $\\sim 2$ between $z=0$ and $z=0.55$. These results confirm that the cosmic evolution of low power radio sources is considerably weaker than that of their higher radio power counterparts. However, because of the relatively small sizes of samples, none of these studies was able to accurately constrain the luminosity dependence of the evolution.  Over the same period, our physical understanding of radio galaxies has developed steadily. \\citet{fanaroff} showed that radio galaxies exhibit a dichotomy in radio morphology. Low radio luminosity sources are usually `edge--darkened', with radio jets that quickly decelerate and flare as they advance through the interstellar medium (Fanaroff--Riley class I sources, or FRIs). In contrast, most higher luminosity sources have jets that remain relativistic over kpc and up to Mpc scales, ending in bright hotspots (`edge--brightened' sources, or FRIIs). The dividing line between FRIs and FRIIs occurs around a monochromatic power of log$_{10}$(P$_{\\rm 1.4 GHz}$)$\\sim25$ W Hz$^{-1}$, although \\citet{ledlowowen} showed that it occurs at higher radio luminosity in more optically--luminous galaxies. This division is close to the break in the radio luminosity function, and has led to the hypothesis these two populations of radio galaxies correspond to the different populations of cosmologically--evolving sources, FRIIs being strongly evolving and FRIs much less so (e.g. \\citealt{jacksonwall}).  A more important division in the radio source population, however, appears to relate to the radiative properties of the AGN. Powerful FRIIs and a small subset of the more powerful FRIs (e.g. see \\citealt{heywood}) show strong emission lines and either a visible or hidden quasar--like nucleus (\\citealt{barthel}), indicative of radiatively efficient accretion. In contrast, most FRIs and a sub--population of weak--emission--lined low luminosity FRIIs show no evidence for a powerfully radiating nucleus (see \\citealt{hardcastle} and references therein) suggesting radiatively inefficient accretion. This implies that a fundamental difference exists between high and low power radio sources that is distinct from the radio morpological dichotomy. Instead, it is likely to be related to different origins or mechanisms of gas fuelling.  Understanding this low luminosity radio source population, and its cosmic evolution, is not only important for understanding the physical origin of the radio activity, but also for understanding galaxy evolution. There is growing evidence that low--luminosity radio--loud AGN play a critical role in regulating the masses and star formation rates of galaxies. Semi--analytic models of galaxy formation  (for example see \\citealt{whitefrenk} and \\citealt{cole}) successfully reproduce many of the  the observed properties of galaxies. However, these models over--predict the abundance of galaxies at the bright end of the luminosity function. This problem arises because many massive galaxies are predicted to sit at the centre of hot hydrostatic haloes, in which cooling flows are expected to develop. If the gas cools and forms stars, central group and cluster galaxies are too massive by the present day and have much bluer colours than observed. \\citet{tabor} first suggested that radio galaxies could in principle regulate these cooling flows, preventing significant accretion of gas and limiting the mass of galaxies. This concept of AGN feedback has now been successfully incorporated into the semi--analytic models of \\citet{bower} and \\citet{croton}.  Observational support for a scenario in which radio AGN play an important role in regulating the growth of massive galaxies has slowly been accumulating. Both luminous FRII and low power FRI sources can in principle inject a significant amount of energy into the surrounding gas, as their radio lobes expand and interact with the surrounding medium ($\\sim10^{42}$ ergs$^{-1}$ for the jets of a borderline FRI/FRII object, see \\citealt{bicknell}). This has been directly observed in clusters of galaxies, where the expanding radio lobes create buoyant bubbles and drill cavities in the hot intracluster medium (\\citealt{bohrin}; \\citealt{mcnam}; \\citealt{fabian}). In nearby systems, which can be studied in detail, the energy supplied by the radio source is well--matched to that required to balance the cooling (e.g. \\citealt{fabian}).  With the availability of large redshifts surveys like the 2--degree--Field Galaxy Redshift Survey (2dFGRS, \\citealt{colless}) and the Sloan Digital Sky Survey (SDSS, \\citealt{york}) it has become possible to obtain the sky and redshift coverage needed to define sufficiently large samples of nearby radio sources to study population statistics and global energetics. \\citet{best05a} constructed a sample of 2215 radio--loud AGN brighter than 5 mJy with redshifts $0.03<z<0.3$ using the SDSS Data Release 2 (DR2) in combination with the National Radio Astronomy Observatory Very Large Array Sky Survey (NVSS, \\citealt{condon}) and the Faint Images of the Radio Sky at Twenty centimeters (FIRST, \\citealt{becker}). \\citet{best05b}  found that the fraction of radio--loud AGN, $f_{RL}$, is a strong function of black hole and stellar mass, with dependencies $f_{RL}\\propto M_{*}^{2.5}$ and $f_{RL}\\propto M_{BH}^{1.6}$. Twenty--five percent of the very most massive galaxies host radio sources more powerful than $10^{23}$ W Hz$^{-1}$. By combining these results with estimates of the energetic output of these sources, \\citet{best06} showed that the time--averaged mechanical energy output of radio sources in elliptical galaxies is comparable to the radiative cooling losses from the hot halo gas, for ellipticals of all masses. Radio sources can therefore provide sufficient energy to suppress cooling flows and regulate the growth of massive galaxies in the nearby Universe.  The role of powerful (radiatively--efficient) FRII objects could not be addressed by Best et~al. due to the scarcity of such systems in the low redshift catalogue. Another  open question is how the heating--cooling balance in massive galaxies evolves with redshift. In order to address these issues, we require a large radio galaxy sample at higher redshifts. In this paper we extend the work of Best et~al. by constructing a comparable catalogue of radio--loud AGN  over the redshift range $0.4 < z < 0.8$. The catalogue contains a large number of luminous radio sources. By comparing with a sample of local radio sources selected  from SDSS DR4, we aim to  study not only of the evolution of the radio galaxies, but also the evolution of the duty cycle of radio AGN activity in galaxies as a function of their stellar mass.  This paper is organized as follows. In Section 2 we will describe the surveys and samples used in this work, as well as the matching procedure. In Section 3 we will present results on the evolution of the radio luminosity function, radio--loud AGN fractions and the bivariate radio luminosity--stellar mass function, and the stellar mass and colour distributions of radio AGN host galaxies. Finally, we will summarize our results in Section 4 and discuss the implications of our work. ",
        "conclusions": "The main results of the present work can be summarized as follows:\\\\  A catalogue of 14453 radio--loud AGN with 1.4 GHz fluxes above 3.5 mJy in the redshift range $0.4<z<0.8$, has been constructed from the cross--correlation of NVSS and FIRST radio source catalogues with the MegaZ--LRG catalogue of luminous red galaxies. The vast majority of the radio AGN are single component sources in both NVSS and FIRST. However, there is a signifiant fraction of objects without high S/N FIRST detections, which are required for accurate identification of the optical counterpart. We thus introduced a method for analyzing the FIRST radio maps around the candidate positions of these sources. This allowed us to dig deeper into the FIRST survey and use lower S/N detections to pinpoint the location of the host galaxy.  We have presented a new determination of the luminosity function of radio--loud AGN at $z \\sim 0.55$ that is in excellent agreement with other results in the literature, but with notably smaller error bars. By comparing our radio luminosity function at $z \\sim 0.55$ to that derived for a large sample of nearby radio AGN from the SDSS DR4, we find compelling evidence for strong cosmic evolution of radio sources. The comoving number density of radio AGN with luminosities less than $10^{25}$ W Hz$^{-1}$ increases by a factor of $\\sim 1.5$ between $z=0.14$ and $z=0.55$. At higher lumiosities, this factor increases sharply, reaching values $\\sim 10$ at a radio luminosity of $10^{26}$ W Hz$^{-1}$. Neither a pure luminosity evolution nor a pure density evolution scenario provides a good description of the data.  We then turn to an analysis of how the relation between radio AGN and their host galaxies evolves with redshift. The fraction of galaxies with radio luminosities above a given threshhold is a steeply increasing function of stellar mass at both $z \\sim 0.1$ and at $z \\sim 0.55$. The fraction of radio loud AGN increases with redshift and this increase is largest at the highest radio luminosities and also for lower mass galaxies. We have also calculated the bivariate radio luminosity--mass function at $z \\sim 0.14$ and $z \\sim 0.55$. Its shape does not appear to depend on mass at either redshift, but there is strong evolution in the shape at the {\\em bright--end} of the function. The fraction of galaxies brighter than  $10^{25}$ W hz$^{-1}$ declines with significantly shallower slope at higher redshifts. Within the range $10^{25-27}$ W hz$^{-1}$, simple power law fits to the DR4 bivariate function in Figure 8 ($\\rm frac \\propto \\rm L_{\\rm 1.4 GHz}^{\\gamma}$), produce slopes $\\gamma$ of $-1.84\\pm0.21$, $-1.66\\pm0.02$ and $-1.83\\pm0.04$ for the first, second and third mass bins. In MegaZ, we obtain $-1.05\\pm0.02$, $-1.16\\pm0.05$ and $-1.20\\pm0.02$ for the same luminosity range and mass bins.  In short, two main conclusions have emerged from our analysis: \\begin{itemize} \\item There is a characteristic luminosity  of $\\sim 10^{25}$ W Hz$^{-1}$ below which the radio source population appears to evolve only very weakly with redshift. Above this characteristic luminosity, there is strong evolution, with the most powerful radio sources undergoing the largest increase in co--moving number density. These results are in broad agreement with past studies. \\item The strongest evolution in the fraction of galaxies that host radio--loud AGN takes place in the lower mass galaxies in our sample. \\end{itemize}  The most plausible explanation of these trends is that there are two classes of radio galaxy, likely associated with the high excitation/low excitation ``dichotomy\") that have different fuelling/triggering mechanisms and hence evolve in different ways. This has also been argued by \\citet{tasse}, whose analysis of a sample 1\\% of the size of ours in the XMM-LSS region provided hints of similar evolutionary properties. As discussed before, it has been hypothesized that the class of low luminosity AGN is likely associated with massive ($ \\sim 10^{13} M_{\\odot}$) dark matter halos with quasistatic hot gas atmospheres. The abundance of such halos is predicted to remain approximately constant out to redshifts $\\sim 1$ (\\citealt{mowhite}) and this is broadly consistent with the very weak evolution we see in this population. These ``hot--halo\" triggered sources at $z\\sim0.5$ would then produce radio emission that roughly compensates the radiative losses of the hot gas.  The reason for the remarkably strong evolution in the comoving abundance of high--luminosity radio AGN is somewhat more difficult to understand. As shown in Figure 8, this evolution is most dramatic in galaxies at the low stellar mass end of our sample. It is thus tempting to link this population with the strongly accreting (and evolving) population of luminous quasars, whose space densities also increase by factors of more than 10 over the redshift range studied in this paper. Recent work has shown that quasars are hosted in dark matter halos of masses $\\sim 10^{12} M_{\\odot}$, independent of redshift or the optical luminosity of the system (e.g. \\citealt{croom}; \\citealt{shen}). Only $\\sim 10$\\% of optical quasars are radio--loud, however, and this raises the question as to whether there is simply a short radio--loud phase during the lifetime of every optical quasar, or whether optically luminous AGN and powerful radio galaxies are triggered under different physical conditions. If, as in quasars, these sources are not regulated by cooling from a surrounding hot gaseous halo, but by other mechanisms like interactions/mergers, then we would not expect to find them in an equilibrum state like in low excitation AGNs.  One way to shed further light on these matters is to compare the clustering properties of quasars to those of the luminous radio galaxies in our sample. If their clustering properties are identical, this would favour the hypothesis that all quasars experience a radio--loud phase. Recently, \\citet{kbh} found that nearby radio-loud AGN were more strongly clustered than matched samples of  radio-quiet AGN with the same black hole masses and extinction corrected [OIII] line luminosities. It will be interesting to see if the same conclusion holds at higher redshifts and for more powerful systems. This will be the subject of a future paper."
    },
    "0809/0809.0629_arXiv.txt": {
        "abstract": "Non-axisymmetric oscillations of rapidly rotating relativistic stars are studied using the Cowling approximation. The oscillation spectra have been estimated by  Fourier transforming the evolution equations describing the perturbations. This is the first study of its kind and provides information on the effect of fast rotation on the oscillation spectra while it offers the possibility in studying the complete problem by including spacetime perturbations. Our study includes both axisymmetric and non-axisymmetric perturbations and provides limits for the onset of the secular bar mode rotational instability. We also present approximate formulae for the dependence of the oscillation spectrum from rotation. The results suggest that it is possible to  extract  the relativistic star's parameters from the observed gravitational wave spectrum. ",
        "introduction": "Relativistic stars during  their evolution may undergo oscillations which can become unstable under certain conditions. Newly born neutron stars are expected to oscillate wildly during their creation shortly after the supernovae collapse \\cite{Ott:2006eh}. They are expected also to oscillate if they are members of binary systems  and there is tidal interaction \\cite{Kokkotas:1995xe,Flanagan:2006sb} or mass and angular momentum transfer from a companion star and also when they undergo phase transitions \\cite{2002PhRvD..66f4027M,Ferrari:2007ur} which might be responsible for the observed glitches in pulsars. Rotation  strongly affects these oscillations and perturbed stars can become unstable if they rotate faster than some critical velocities. During these oscillatory phases of their lives compact stars emit copious amounts of gravitational waves which together with viscosity tend to suppress the oscillations. The oscillations are divided into different families according to the restoring force \\cite{1988ApJ...325..725M,1999LRR.....2....2K}. If pressure is the main restoring force then these modes are called (pressure) \\emph{p-modes}, buoyancy is the restoring force of another class of oscillation mode, the \\emph{g-modes} while Coriolis force is the restoring force for the \\emph{inertial modes}. Spacetime induces another family of oscillations which couple only weakly to the fluid, these are the so-called \\emph{w-modes} \\cite{Kokkotas:1992ks}. There are more families of modes if one assumes the presence of crust \\cite{1983MNRAS.203..457S,1985ApJ...297L..37M,1988ApJ...325..725M,2008MNRAS.384.1711V} or magnetic fields \\cite{1986ApJ...305..767C}. For a complete description of the relativistic star perturbation theory one may refer to \\cite{1999LRR.....2....2K,2001IJMPD..10..381A,lrr-2003-3}.  The study of stellar perturbations in the framework of general relativity (GR) dates back to mid 1960s with the seminal works by Thorne and his collaborators \\cite{Thorne:1967th,Price:1969pt,Thorne:1969to,Thorne:1969th}. Since then the study of stellar oscillations and possible instabilities was a field of intensive work in  relativistic astrophysics \\cite{Detweiler:1985dl, 1985ApJ...297L..37M}. Moreover, during the last two decades, these studies become even more important due to the relations of the oscillations and instabilities to the emission of gravitational waves and the possibility of getting information about the stellar parameters (mass, radius, equation of state) by the proper analysis of the oscillation spectrum \\cite{Andersson:1996ak,Andersson:1998ak,Benhar:1998au,2001MNRAS.320...307K,benhar-2004-70,2004PhRvD..69h4008S,2004PhRvD..70h4026S}. Still, all these studies were mainly dealing with non-rotating stars, because the combination of rotation and general relativity made both the analytic and numerical studies extremely involved. This led to certain approximations in studying rotating stars in GR. The most obvious of them include the so called ``slow-rotation\" and the ``Cowling\" approximation. Actually, both approximation were known and have been used extensively in Newtonian theory of stellar oscillations \\cite{1989nos..book.....U}. In the \\emph{slow rotation approximation} one expands the perturbation equations in terms of a small parameter $\\varepsilon=\\Omega/ \\Omega_K$, where $\\Omega$ is the angular velocity of the star and $\\Omega_K$ stands for the ``Kepler angular velocity\" which is the maximum velocity that can be attained before the star splits apart due to rotation. In the \\emph{Cowling approximation} one typically neglects the perturbations of the Newtonian potential or the spacetime in the case of GR. This is a quite good approximation, both qualitatively and quantitatively, for the higher order \\emph{p-modes}, for the \\emph{g-modes} and the inertial modes while  it is only qualitatively  good  for the \\emph{f-modes}, see for example \\cite{1968AnAp...31..549R}. Although, most of our understanding on the oscillations of relativistic stars is due to perturbative studies, recently,  it has become possible to study stellar oscillations using evolutions of the non-linear equations of motion for the fluid \\cite{Font:1999wh,2001PhRvL..86.1148S,Font:2001eu,Stergioulas:2003ep,Dimmelmeier:2005zk,2006PhRvD..74l4024K,2008arXiv0804.1151K,2008PhRvD..77d4042B,2008arXiv0803.3804B}. Finally, differential rotation is another key issue that is believed to play an important role in the dynamics of nascent neutron stars. Actually, it is associated with dynamical instabilities both for fast and slowly rotating neutron stars \\cite{2001ApJ...550L.193C,2002MNRAS.334L..27S} and affects the onset of secular instabilities; still it has not yet been studied extensively  \\cite{Dimmelmeier:2002bm,2007PhRvD..75f4019S,2008PhRvD..77b4029P} and remains an open issue.   Since rotational instabilities are typically connected with fast rotating stars, it is of the great importance the study of oscillations of this type of objects. As a first step one can drop the ``slow rotation\" approximation but still maintain the Cowling approximation i.e. to freeze the spacetime perturbations or in the best case to freeze the radiative part of these perturbations. This was the approach used up to now for most of the studies of the oscillations of fast rotating relativistic stars either in perturbative approaches \\cite{2002ApJ...568L..41Y,2007PhRvD..75d3007B} or in non-linear (but axisymmetric) approaches \\cite{2006MNRAS.368.1609D,2008arXiv0804.1151K}.  In this article, we present the first results of a new approach based on 2D time evolution of the perturbation equations which seems to be promising for the study of the oscillations and instabilities of fast rotating neutron stars.  This the first study of its kind, since earlier 2D perturbative studies have been done either in Newtonian theory \\cite{2002MNRAS.334..933J} (see also \\cite{Passamonti:2008nx} for a very recent work) or in reduction to eigenvalue problem \\cite{2002ApJ...568L..41Y,2007PhRvD..75d3007B,2007PhRvD..76j4033F}. The advantage of this method is that it can be easily extended to include differential rotation or the perturbations of the spacetime. On the other hand it provides a robust tool in studying the onset of rotational instabilities in fast rotating neutron stars, while as it has been demonstrated here, one can get easily results for realist equations of state which is vital for developing gravitational wave asteroseismology. Finally, via this approach one may answer the question of the existence or absence of a continuous spectrum for the inertial modes as it has been suggested by the 1D studies in the slow rotation approximation \\cite{1998MNRAS.293...49K,Beyer:1999te,2002MNRAS.330.1027R,2003MNRAS.339.1170R,2008PhRvD..77b4029P}.  In the next section we describe in detail the derivation of the perturbation equations and the conventions that we have adopted. In section 3, we present the numerical tools that have been developed in order to study the problem while section 4  describes the results for axisymmetric and non-axisymmetric perturbations. In the last section we discuss the results, their application to astrophysics and the possible extensions of this work. ",
        "conclusions": "In this work we present a study of the oscillation properties of fast rotating relativistic stars for bot both axisymmetric and non-axisymmetric perturbations. The study was based on the  Cowling approximation  using a 2D version of the perturbation equations. The results for axisymmetric perturbations are in excellent agreement with earlier ones while the non-axisymmetric results are the first of their kind in the literature. We demonstrated the neutral points for the onset of the CFS instability and suggested possible normalizations which can be used in order to extract the parameters of the rotating star.  The method  presented here (evolution of the 2D linearized equations) can be extended by including the perturbations of the spacetime. This will offer the possibility in testing the earlier results \\cite{Stergioulas:1997ja} for the onset of the secular instability in fast rotating stars. Moreover, the  dependence of oscillations frequencies on rotation will be based on the exact results and will not rely on the Cowling approximation.  Finally, the effect of differential rotation on the spectra can be studied both for testing earlier, mainly Newtonian, results \\cite{PhysRevD.64.024003,  2002ApJ...568L..41Y} but more importantly in finding the correct dependence of the frequencies on rotation as well as the neutral points for the onset of the secular instabilities. This is actually an important step since newly born neutron stars are expected to rotate differentially at least during the stages that they will be in an oscillatory phase or the even when they are secularly unstable.    \\[ \\] {\\bf Acknowledgments:} We thank D. Petroff and N. Stergioulas  for helpful discussions. This work was also supported by the German Science Foundation (DFG) via SFB/TR7."
    },
    "0809/0809.4300_arXiv.txt": {
        "abstract": "{The \\roh\\ and \\rdoh\\ isotope ratios are measured in the Oort-Cloud comet C/2002 T7 (LINEAR) through ground-based observations of the OH $\\, A\\,^{2}\\Sigma^{+} - X\\,^{2}\\Pi_{i}$ ultraviolet bands at 3063~\\AA\\ (0,0) and 3121~\\AA\\ (1,1) secured with the Very Large Telescope (VLT) feeding the Ultraviolet-Visual Echelle Spectrograph (UVES).  From the \\roh\\ ratio, we find \\ro\\ = 425 $\\pm$ 55, equal within the uncertainties to the terrestrial value and to the ratio measured in other comets, although marginally smaller.  We also estimate \\rdoh\\ from which we derive D/H = 2.5 $\\pm$ 0.7 10$^{-4}$ in water. This value is compatible with the water D/H ratios evaluated in other comets and marginally higher than the terrestrial value. } ",
        "introduction": "\\label{sect:intro}  The determination of the abundance ratios of the stable isotopes of the light elements in different objects of the Solar System provides important clues into the study of their origin and history.  This is especially true for comets which carry the most valuable information regarding the material in the primitive solar nebula.  The \\ro\\ isotopic ratio has been measured from space missions in a few comets. In-situ measurements with the neutral and ion mass spectrometers onboard the Giotto spacecraft gave \\ro\\ = 495$\\pm$37 for H$_2$O in comet 1P/Halley (Eberhardt et al. \\cite{Eberhardt}). A deep integration of the spectrum of the bright comet 153P/2002 C1 (Ikeya-Zhang) with the sub-millimeter satellite Odin led to the detection of the H$_2^{18}$O line at 548 GHz (Lecacheux et al. \\cite{Lecacheux}).  Subsequent observations resulted in the determination of \\ro\\ = 530$\\pm$60, 530$\\pm$60, 550$\\pm$75 and 508$\\pm$33 in the Oort-Cloud comets Ikeya-Zhang, C/2001 Q4, C/2002 T7 and C/2004 Q2 respectively (Biver et al. \\cite{Biver}).  Within the error bars, these measurements are consistent with the terrestrial value (\\ro\\ {\\small(SMOW\\footnote{Standard Mean Ocean Water})} = 499), although marginally higher (Biver et al. \\cite{Biver}).  More recently, laboratory analyses of the silicate and oxide mineral grains from the Jupiter family comet 81P/Wild~2 returned by the Stardust space mission provided \\ro\\ ratios also in excellent agreement with the terrestrial value. Only one refractory grain appeared marginally depleted in $^{18}$O (\\ro\\ = 576$\\pm$78) as observed in refractory inclusions in meteorites (McKeegan et al.~\\cite{McKeegan}).  The D/H ratio has been measured in four comets. In-situ measurements provided D/H = 3.16$\\pm$0.34 10$^{-4}$ for H$_2$O in 1P/Halley (Eberhardt et al. \\cite{Eberhardt}, Balsiger et al. \\cite{Balsiger}), a factor of two higher than the terrestrial value (D/H {\\small (SMOW)} = 1.556 10$^{-4}$).  The advent of powerful sub-millimeter telescopes, namely the Caltech Submillimeter Observatory and the James Clerck Maxwell telescope located in Hawaii, allowed the determination of the D/H ratio for two exceptionally bright comets.  In comet C/1996 B2 (Hyakutake), D/H was found equal to 2.9$\\pm$1.0 10$^{-4}$ in H$_2$O (Bockel\\'ee-Morvan et al. \\cite{Bockelee}), while, in comet C/1995 O1 (Hale-Bopp), the ratios D/H = 3.3$\\pm$0.8 10$^{-4}$ in H$_2$O and D/H = 2.3$\\pm$0.4 10$^{-3}$ in HCN were measured (Meier et al. \\cite{Meier1,Meier2}), confirming the high D/H value in comets. Both Hyakutake and Hale-Bopp are Oort-Cloud comets.  Finally, bulk fragments of 81P/Wild~2 grains returned by Stardust indicated moderate D/H enhancements with respect to the terrestrial value. Although D/H in 81P/Wild~2 cannot be ascribed to water, the measured values overlap the range of water D/H ratios determined in the other comets (McKeegan et al.~\\cite{McKeegan}).  Among a series of spectra obtained with UVES at the VLT to measure the \\rn\\ and \\rc\\ isotope ratios in various comets from the 3880$\\,$\\AA\\ CN ultraviolet band (e.g. Arpigny et al. \\cite{Arpigny}, Hutsem\\'ekers et al. \\cite{Hutsemekers}, Jehin et al. \\cite{Jehin2}, Manfroid et al. \\cite{Manfroid2}), we found that the spectrum of \\com\\ appeared bright enough to detect the $^{18}$OH lines in the $A\\,^{2}\\Sigma^{+} - X\\,^{2}\\Pi_{i}$ bands at 3100~\\AA\\ allowing --for the first time-- the determination of the \\ro\\ ratio from {\\it ground-based} observations.  We also realized that the signal-to-noise ratio of our data was sufficient to allow a reasonable estimate of the \\rdoh\\ ratio from the same bands.  The possibility of determining the \\ro\\ ratio from the OH ultraviolet bands has been emphasized by Kim (\\cite{Kim}). Measurements of the OD/OH ratio were already attempted by A'Hearn et al. (\\cite{AHearn}) using high resolution spectra from the International Ultraviolet Explorer and resulting in the upper limit D/H $< 4 \\, 10^{-4}$ for comet C/1989 C1 (Austin). These observations now become feasible from the ground thanks to the high ultraviolet throughput of spectrographs like UVES at the VLT. ",
        "conclusions": "We have measured the oxygen isotopic ratio \\ro\\ = 425 $\\pm$ 55 from the OH \\ $A\\,^{2}\\Sigma^{+} - X\\,^{2}\\Pi_{i}$ ultraviolet bands in comet \\com .  Although marginally smaller, our value do agree within the uncertainties with \\ro\\ = 550 $\\pm$ 75 estimated from observations by the Odin satellite (Biver et al. \\cite{Biver}), with the \\ro\\ ratios determined in other comets, and with the terrestrial value (Sect.~\\ref{sect:intro}).  To explain the so-called ``oxygen anomaly'' i.e. the fact that oxygen isotope variations in meteorites cannot be explained by mass-dependent fractionation, models of the pre-solar nebula based on CO self-shielding were proposed, predicting enrichments, with respect to the {\\small SMOW} value, of $^{18}$O in cometary water up to \\ro\\ $\\sim$ 415 (Yurimoto \\& Kuramoto \\cite{Yurimoto}).  Recently, Sakamoto et al. (\\cite{Sakamoto}) found evidence for such an enrichment in a primitive carbonaceous chondrite, supporting self-shielding models. The value of \\ro\\ we found in \\com\\ is also marginally smaller than the terrestrial value and compatible with these predictions.  On the other hand, the measurement of \\ro\\ = 440 $\\pm$ 6 in the solar photosphere (Ayres et al. \\cite{Ayres}; cf. Wiens et al. \\cite{Wiens} for a review of other, less accurate, measurements) indicates that solar ratios may deviate from the terrestrial ratios by much larger factors than anticipated, requiring some revision of the models.  More observations are then critically needed to get an accurate value of \\ro\\ in comets, assuming that cometary water is pristine enough and can be characterized by a small set of representative values. Namely, if self-shielding is important in the formation of the solar system, it is not excluded that significant variations can be observed between comets formed at different locations in the solar system, like Oort cloud and Jupiter-family comets.  We also detected OD and estimated D/H = 2.5 $\\pm$ 0.7 10$^{-4}$ in water. Our measurement is compatible with other values of D/H in cometary water and marginally higher than the terrestrial value (Sect.~\\ref{sect:intro}).  Our observations were not optimized for the measurement of OD/OH (neither for \\roh ) and one of our best spectra was obtained at airmass $\\sim$ 2 with less than 20 min of exposure time for a comet of heliocentric magnitude m$_r \\simeq$ 5 (for comparison, comet Hale-Bopp reached m$_r \\simeq$ $-$1). All these observing circumstances can be improved, including observations at negative heliocentric velocities to increase the OD/OH fluorescence efficiency ratio (cf. figure~1 of A'Hearn et al. \\cite{AHearn}). This opens the possibility to routinely measure both the \\ro\\ and D/H ratios from the ground, together with the \\rc\\ and \\rn\\ ratios, for a statistically significant sample of comets of different types (e.g. Oort-cloud, Halley-type, and hopefully Jupiter-family comets although the latter are usually fainter). The measurement of D/H is especially important since it allows to limit the contribution of comets to the terrestrial water, the high D abundance implying that no more than about 10 to 30\\% of Earth's water can be attributed to comets (e.g. Eberhardt et al. \\cite{Eberhardt}, Dauphas et al. \\cite{Dauphas}, Morbidelli et al. \\cite{Morbidelli}).  However, only a full census of D/H in comets could answer this question. In particular, if Jupiter-family comets, thought to have formed in farther and colder places in the Solar System, are characterized by an even higher D/H, closer to the ratio measured in the interstellar medium water, then the fraction of cometary H$_2$O brought onto the Earth could be even smaller."
    },
    "0809/0809.3185_arXiv.txt": {
        "abstract": "We present a measurement of metallicity in the Galactic center Quintuplet Cluster made using quantitative spectral analysis of two Luminous Blue Variables (LBVs). The analysis employs line-blanketed NLTE wind/atmosphere models fit to high-resolution near-infrared spectra containing lines of H, \\HeI, \\SiII, \\MgII, and \\FeII.  We are able to break the H/He ratio vs. mass-loss rate degeneracy found in other LBVs and to obtain robust estimates of the He content of both objects. Our results indicate solar iron abundance and roughly twice solar abundance in the $\\alpha$-elements. These results are discussed within the framework of recent measurements of oxygen and carbon composition in the nearby Arches Cluster and iron abundances in red giants and supergiants within the central 30~pc of the Galaxy. The relatively large enrichment of $\\alpha$-elements with respect to iron is consistent with a history of more nucleosynthesis in high mass stars than the Galactic disk. \\end {abstract} ",
        "introduction": " ",
        "conclusions": "}  Our results suggest solar Fe abundances and approximately twice-solar $\\alpha$-element abundances for the Quintuplet LBVs.  Presumably, these abundances were the same in the gas that condensed to form these stars and the other stars in the Quintuplet cluster and indeed in the whole of the present-day Galactic center. The results can be discussed in the context of similar measurements of Galactic center objects and with respect to the trend one might expect if the region is an inward extension of the disk or the bulge. In addition, the ratio of Fe to $\\alpha$-elements might be used to decipher the star formation history in the Galactic center.  \\input table2.tex  Table~2 displays a number of recent determinations of stellar metal abundances in the Galactic Center together with the above mentioned three reference patterns for solar abundances.  The values derived for Fe abundances in cool stars agree with our result \\citep{car00,ram97,ram99,ram00}.  \\citet{cun07} find a very narrow range of Fe abundances clustered around the solar value for a population of cool stars in the central 30~pc. Of particular interest is star VR5-7 from \\citet{cun07} sample which is located in the Quintuplet Cluster and shows A(Fe/H)=7.60 and A(Ca/H)=6.41  There are relatively few measurements of the $\\alpha$-element abundances ([$\\alpha$/Fe]) in GC stars.  \\citet{naj04} find solar abundances (as defined in this paper) for hot stars in the Arches cluster based on the oxygen abundance and, to a lesser degree, carbon abundance, and adopting the canonical solar value of A(O/H)=8.93 \\citep{and89}. Those estimates assume that nitrogen has reached its maximum surface abundance value. Evolutionary models indicate that 95\\% of that value is already attained by the time that H/He$<$2 (by number). \\citet{naj04} followed the metallicity patterns from the Geneva evolutionary models and assumed no selective enrichment of CNO or $\\alpha$-elements vs Fe, in concluding that the stars in the Arches Cluster have solar $\\alpha$-element abundances. However, estimates of solar abundances have varied considerably over the past 15 years \\citep[e.g.][, see also Table~2]{alle08}. Thus, depending on the assumed solar CNO composition, the derived nitrogen abundance by \\citet{naj04} could imply solar \\citep[][]{and89}, 1.2~$\\times$~solar \\citep[][]{gre93} or 2.0~$\\times$~~solar \\citep[][]{asp05} CNO composition.  Recently \\citet{mar07a,mar07b} have analyzed a larger sample of hot stars in the Arches and Central Parsec clusters and find similar results (see Table~2). Interestingly, if one considers only the objects in \\citet{mar07b} with He/H$>0.1$ and those with Z(C)$<0.05$, \\ie\\ fulfilling the condition to be close enough to Z(N)$_{max}$, the average value of Z(N) is 1.7.  \\citet{geb06} estimate roughly solar oxygen abundance, A(O/H)=8.91, in IRS~8, an OIf supergiant near the central parsec.  \\citet{cun07} find $<$A(O/H)$>$=9.04 ([0.37]) and $<$A(Fe/H)$>$=7.59 ([0.14]) for their sample of cool stars, where the numbers in brackets are the ratio with respect to the solar value in dex.  This implies [O/Fe]=0.22, i.e.\\ a clear enhancement over the solar ratio. Again, the \\citet{cun07} measurements could be interpreted as indicating solar ratios in O over Fe if the solar O abundance in evolutionary models is used. It is crucial to have accurate solar abundances, and that values used in stellar evolution calculations should be consistent with these. An excellent example is attempting to determine whether the possible oxygen enhancement is due to a top-heavy IMF favoring $\\alpha$-elements vs Fe enrichment, or simply an overall CNO and metal enhancement. Thus, taking the CNO abundances for the GC objects from \\citet{cun07} and assuming C/N equilibrium values one can interprete their results either as solar CNO with mildly enhanced (30\\%) oxygen \\citep{and89} or a clearly supersolar environment with a factor of 1.7 enhancement for C and N and 2.5 for oxygen \\citep{asp05}.   Fortunately, there are other $\\alpha$-elements whose adopted solar abundances have suffered basically no major revision.  Thus, we believe that the enhanced values obtained in this work for Mg and Si, roughly a factor of two solar, together with the enhancement of Ca found by \\citet{cun07}, are a strong indication of the enrichment of $\\alpha$-elements compared to Fe.  Our results run counter to the trend in the disk \\citep{roll00,sma01,mar03}, and are more consistent with the values found for the bulge \\citep{fro99,fel00}. This may imply that the ISM in the disk does not extend inward to the GC, so that material is dragged into the central molecular zone from the bulge rather than from the disk. Another possibility is that the GC stars are forming out of an ISM that has an enrichment history distinctly different from that of the disk. At this point, further studies of the $\\alpha$-elements vs Fe would be useful. Future high S/N and high resolution spectroscopy of the \\OI\\ lines in LBVs and K-band spectroscopy of WNL stars in the same cluster (Najarro et al. in prep.) will provide two independent measurements of the original oxygen content, and thus set definite constraints on metallicity.  The modest enrichment in $\\alpha$-elements versus Fe that we find in the two Quintuplet LBVs is consistent with a top-heavy IMF in the GC \\citep{fig99a}. In such a scenario, enhanced yields of $\\alpha$-elements compared to Fe are expected through a higher than average ratio of the number of SNII vs SNIa events \\citep{whe89,cun07}."
    },
    "0809/0809.4593_arXiv.txt": {
        "abstract": "Our \\textit{Swift} observations of RS Oph form an unprecedented X-ray dataset to undertake investigations of both the central source and the interaction of the outburst ejecta with the circumstellar environment. Over the first month, the XRT data are dominated by emission from rapidly evolving shocks. We discuss the differences in derived parameters from those found for \\textit{RXTE} at early times and the evolution of the X-ray emission to much later times. It is apparent that at late times several emission components are present. We find no strong evidence of the proposed shock break-out in our data. ",
        "introduction": "Prior to 1985, evidence for the interaction of outburst ejecta with the circumstellar environment in RS~Oph came from observations of coronal lines in optical spectra, the narrowing of initially broad emission lines, and the evolution of superimposed narrow emission features \\citep[see e.g.][and references therein]{mas87_bs}. \\textit{EXOSAT} observations from $t = 55$ to 251 days post-outburst in 1985 showed that RS Oph was a bright, but rapidly declining, X-ray source \\citep{mas87_bs}. Bode \\& Kahn (1985) formulated an analytical model and concluded that RS Oph evolved like a supernova remnant, but on timescales around $10^5$ times faster. Subsequently, O'Brien et al. (1992) developed detailed hydrodynamical models of the interaction of the ejecta with the red giant wind to derive parameters of the explosion and the circumstellar environment. ",
        "conclusions": ""
    },
    "0809/0809.5020_arXiv.txt": {
        "abstract": "{ In the framework of hierarchical structure formation ellipticals can form from merging of smaller disk galaxies. The nearby interacting 'Antennae' galaxy pair (NGC 4038/39) is one of the best-studied local systems of merging spirals, thus presenting us with an ideal laboratory for the study of galaxy evolution models. The Antennae are believed to be in a state prior to their final encounter with rapid subsequent merging, which puts them in the first position of the Toomre (1977) merger sequence. Here we present first numerical high-resolution, self-consistent, smoothed particle hydrodynamics (SPH) simulations of the Antennae system, including star formation and stellar feedback, and compare our results to VLA HI observations by Hibbard et al.~(2001). We are able to obtain a close, but not yet perfect match to the observed morphology and kinematics of the system. } ",
        "introduction": "The 'Antennae' galaxies (NGC 4038/39, Arp 244, VV245) are a well-known, archetypal example of a spiral-spiral\\linebreak galaxy merger (see Fig. \\ref{FigHibbHIoptical}). They are nicknamed after their spectacular appearance with a prominent pair of tidal tails, which was already noted in early optical images (e.g. Duncan 1923, Schweizer 1978). At present there is a huge\\linebreak amount of data collected for the Antennae both from\\linebreak ground- and space-based telescopes, covering a wide range of wavelength regimes. Recent examples include: optical HST WFPC data (Whitmore et al.~1999), near-IR WIRC (Brandl et al.~2005) and mid-IR IRAC (Wang et al.~2004) imaging, and Chandra ACIS-S X-ray observations (Baldi et al. 2006). The elongated tails are a clear sign of tidal interaction of nearly equal-mass galaxies (Toomre \\& Toomre 1972). They form kinematically by gravitational tides exerted during the interaction process and their morphology and velocity fields give strong hints on the encounter geometry and history. This makes the Antennae galaxies a key system for investigating galaxy interactions and their associated physical phenomena. Starting from the first numerical model by Toomre \\& Toomre (1972), where they used the restricted N-body method to model the evolution of the disks, there have been many attempts to obtain numerical 'look-alikes' for the Antennae galaxies. Barnes (1988) was the first to use self-consistent N-body models with multiple components consisting of a bulge, a disk, and dark halo with a mass ratio of 1:3:16 and a total mass of $2.75 \\cdot 10^{11}\\, \\mathrm{M}_\\odot$. Dubinski, Mihos \\& Hernquist~(1996) used the extent of the tidal features in the Antennae to probe the amount of dark matter in these galaxies. Mihos, Bothun \\& Richstone~(1993) were the first to include gas dynamics and star formation into their dynamical model of the Antennae galaxies. For the latest review of the Antennae models we would also like to refer to the results presented by Hibbard~(2003). However, most of the Antennae models are still based on the orbital parameters given in Toomre \\& Toomre (1972). Here we want to extend the probed parameter space in order to find orbital parameters which allow for a closer fit to the observed morphological and kinematical data of the Antennae. We use medium-resolution VLA HI mappings of the Antennae ($\\sim 20$'', $\\Delta v=5.21 \\mathrm{km}\\,\\mathrm{s}^{-1}$) by Hibbard et al.~(2001) for comparison with our models. In this case it is advantageous to use the cold atomic gas as a sensitive tracer of the large-scale dynamics of the system as it is unlikely to be disturbed by star formation. Recently, there has been a debate about the exact distance to the Antennae, ranging from a modest $13.3 \\pm 1.0$ Mpc (Saviane et al. 2008), based on photometry of the tip of the red giant branch, to $22 \\pm 3$ Mpc (Schweizer et al.~2008), based on observations of a supernova of type Ia. Sometimes even higher and lower values have been quoted (e.g. Zezas \\& Fabbiano 2002, Rubin, Ford \\& D'Odorico 1970). A widely-used, intermediate distance to the Antennae is $D = 19.2$ Mpc, which is derived from the systemic recession velocity relative to the Local Group assuming a Hubble constant of $H_0 = 75 ~\\mathrm{km}~\\mathrm{s}^{-1}~\\mathrm{Mpc}^{-1}$ (Whitmore et al.~1999). In this paper the distance will be adopted as part of the model matching process. \\begin{figure} \\centering \\includegraphics[width=50mm,height=41.25mm,angle=90, scale=1.]{Hibbard2001_HIoptical2a.ps} \\caption{NGC 4038 (south) and NGC 4039 (north). Distribution of HI and optical light. HI data are shown in blue, together with a combined B + V + R optical image in white and green. Taken from Hibbard et al.~(2001). North is pointing upwards.} \\label{FigHibbHIoptical} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=50mm,height=50mm,scale=1.]{FigPlaneOfSky_GSD2008.eps} \\caption{'Best match' to the Antennae at $t=0.60$ Gyr in the plane-of the sky. Disk and bulge particles are shown in green, gas in blue, and stars, which have formed during the simulation, in red. Note that for the old stellar component (green) we plot only every fifth particle.} \\label{FigPlaneOfSky} \\end{figure} ",
        "conclusions": "In this work we have presented first steps towards a good numerical model for the Antennae galaxies. The simulations include star formation and stellar feedback using sub-grid physics as proposed by Springel \\& Hernquist (2003). We are able to obtain a close match to the observed morphology of the system, while a perfect fit to the kinematical HI data is still to be achieved. Our goal is to investigate in detail the history and distribution of newly formed stars in the Antennae galaxies, in particular in the 'overlap' region of the two progenitor disks. Close comparison to observations then enables us to adjust the star formation prescriptions used in numerical simulations and figure out the important physical properties controlling star formation in interacting galaxies. In this spirit, it is worth examining alternative formulations for the star formation laws in numerical simulations, see e.g. Barnes'~(2005) shock-induced star formation model for the 'Mice' galaxies. There is a number of possible processes relevant for the enhanced star formation in the Antennae merger, specifically in the 'overlap' region. One idea is that collisions of giant molecular clouds (GMCs) during the merger process could trigger starbursts (Noguchi 1991) in interacting galaxies. As the disk filling factor of GMCs is rather low ($f < 0.01$), Jog \\& Solomon~(1992) proposed that a hot ionized, high-pressure gas medium, resulting from fast collisions between HI clouds, could cause a radiative shock compression in the outer layers of pre-existing GMCs, thereby inducing bursts of star formation. Elmegreen \\& Efremov~(1997) advocate high-pressure regions, which originate in the large-scale shocks of interacting galaxies, for being the driver of bound star cluster formation, while, lately, Renaud et al.~(2008) raised the idea that compressive tidal modes in galaxy mergers could play an important role for the formation of globular clusters. It would also be interesting to study whether one can form candidates for tidal dwarf galaxies in the tails of the merger as is observed in the southern tidal arm of the Antennae (Schweizer 1978, Mirabel, Dottori \\& Lutz 1992, see also Wetzstein, Naab \\& Burkert 2007)."
    },
    "0809/0809.2600_arXiv.txt": {
        "abstract": "We report the discovery of a third white dwarf hosting a gaseous debris disc, SDSS\\,J084539.17+225728.0. The typical double-peaked \\Lines{Ca}{II}{8498,8542,8662}\\,\\AA\\ emission lines can be modelled in terms of a Keplerian gas disc with a radial extent from $\\sim0.5$\\,\\Rsun\\ to $\\sim1.0$\\,\\Rsun.  The effective temperature of SDSS\\,0845+2257, $\\Teff\\simeq18600\\pm500$\\,K, is comparable to the two other white dwarfs with gaseous discs, SDSS\\,1043+0855 and SDSS\\,1228+1040, and hence substantially hotter than the bulk of white dwarfs where dusty debris discs were identified through the presence of infrared excess flux. This may suggest that the conditions to produce emission lines from debris discs in the optical wavelength range are only met for a relatively narrow range in $\\Teff$.  The observed asymmetry in the line profiles indicates a substantial eccentricity in the disc. Two spectra obtained four years apart reveal a significant change in the shapes and equivalent widths of the line profiles, implying that the circumstellar disc evolves on relatively short time scales. In contrast to SDSS\\,1043+0855 and SDSS\\,1228+1040, SDSS\\,0845+2257 has a helium-dominated atmosphere. We detect photospheric absorption lines of He, Ca, Mg, and Si in the Sloan Digital Sky Survey spectrum, and hence classify SDSS\\,0845+2257 as DBZ white dwarf. The abundances for the three metals determined from model atmosphere fits are $\\mathrm{Ca/He}\\simeq1.3\\times10^{-7}$, $\\mathrm{Mg/He}\\simeq6.0\\times10^{-6}$, and $\\mathrm{Si/He}\\simeq8.0\\times10^{-6}$. From the non-detection of H$\\alpha$ we derive $\\mathrm{H/He}<3\\times10^{-5}$, which implies that the hydrogen-to-metal abundance ratio of the circumstellar material is $\\ga1000$ times lower than in the Sun. This lends strong support to the hypothesis that the gaseous and dusty debris discs found around roughly a dozen white dwarfs originate from the disruption of rocky planetary material. ",
        "introduction": "Recent years have seen a surge of interest in the evolution of extra-solar planetary systems through the late phases in the evolution of their host stars.  \\citet{ignace01-1}, \\citet{chuetal01-1}, and \\citet{burleighetal02-1} suggested that Jovian planets around white dwarfs may be detected either through direct imaging or spectroscopy. While the survival of planets seems plausible \\citep{villaver+livio07-1}, there is as yet no unambiguous detection of a planet around a white dwarf (\\citealt{burleighetal08-1}, but see also \\citealt{mullallyetal08-1} for a good candidate).  \\begin{figure} \\centerline{\\includegraphics[angle=0,width=0.75\\columnwidth]{caii.ps}} \\caption{\\label{f-caii} Comparison of the \\Lines{Ca}{II}{8498,8542,8662}\\,\\AA\\ emission lines in the SDSS spectra of the three white dwarfs SDSS\\,1228+1040 (top, \\citealt{gaensickeetal06-3}), SDSS\\,0845+0855 (middle) and SDSS\\,1043+0855 (bottom, \\citealt{gaensickeetal07-1}). The spectra of SDSS\\,0845+0855 and SDSS\\,1228+1040 are offset by one and two units, respectively. } \\end{figure}  However, as part of a search for cool companions to white dwarfs, an infrared excess was discovered around the cool white dwarf G29--38 \\citep{zuckerman+becklin87-1}. HST observations of G29--38 revealed a large photospheric abundance of metals, implying that the star is accreting at a relatively high rate \\citep{koesteretal97-1}.  Deep imaging and asteroseismological studies of G29--38 ruled out a brown dwarf companion \\citep{grahametal90-1, kleinmanetal94-1, kuchneretal98-1} and led to the hypothesis of a cool dust cloud around the white dwarf. The presence of dust near G29--38 has been verified by infrared observations with the \\textit{Spitzer Space Telescope} \\citep{reachetal05-1}. Recent infrared surveys have boosted the number of cool white dwarfs harbouring dust discs to $\\sim10$ (e.g. \\citealt{becklinetal05-1, kilicetal05-1, kilicetal06-1, kilic+redfield07-1, vonhippeletal07-1, juraetal07-1, farihietal08-1}). These dust discs originate from the tidal disruption of either comets \\citep{debesetal02-1} or asteroids \\citep{jura03-1}. Asteroids appear to be more likely as they can explain the large metal abundances in the material accreted by the white dwarfs from the dusty environment.  Hydrogen-rich white dwarfs with radiative atmospheres ($\\Teff\\ga11500$\\,K) have short diffusion time scales for heavy elements, and hence their photospheric abundances should closely reflect the chemical abundances of the circumstellar debris discs from which they are accreting. \\citet{zuckermanetal07-1} suggested that such systems offer potential insight into the chemical composition of extra-solar planetary systems, once more detailed calculations of diffusion time scales for a wide range of elements and atmospheric parameters are available.  We have recently identified two hydrogen-dominated (DA) white dwarfs, SDSS\\,J122859.93+104032.9 and SDSS\\,J104341.53+085558.2 (henceforth SDSS\\,1228+1040 and SDSS\\,1043+0855), that exhibit double-peaked emission lines in the $I$-band \\Ion{Ca}{II} triplet \\citep{gaensickeetal06-3, gaensickeetal07-1}. From the morphology of the \\Ion{Ca}{II} lines profiles we inferred the presence of gaseous circumstellar discs close ($\\sim R_\\odot$) to the white dwarfs.  Like the majority of the white dwarfs with dusty debris discs, both SDSS\\,1228+1040 and SDSS\\,1043+0855 have hydrogen-rich metal-polluted (DAZ) atmospheres, indicating that they are accreting from the circumstellar material. The absence of photospheric helium absorption lines, as well as of Balmer emission lines from the gaseous discs, strongly suggest that circumstellar material around SDSS\\,1228+1040 and SDSS\\,1043+0855 is depleted in volatile elements. A dusty extension to the gaseous disc in SDSS\\,1228+1040 has been detected with \\textit{Spitzer} \\citep{brinkworthetal08-1}.  Here we present a third white dwarf with a metal-rich gaseous debris disc, the helium-dominated (DB) SDSS\\,J084539.17+225728.0 (henceforth SDSS\\,0845+2257).  \\begin{figure*} \\includegraphics[angle=-90,width=12.8cm]{fit.ps} \\caption{\\label{f-opt_fit} Model atmosphere fit to the spectrum of SDSS\\,0845+2257. Main panel: best-fit to the SDSS spectrum with $\\Teff=18621$\\,K and $\\log g=8.175$. The positions of the \\Ion{Ca}{II} and \\Ion{Fe}{II} emission lines are indicated. Top panels: close-ups of the average WHT spectrum regions featuring photospheric metal absorption lines Ca, Mg, and Si. The model shown has been computed for the following abundances: $\\mathrm{Ca/He}\\simeq1.3\\times10^{-7}$, $\\mathrm{Mg/He}\\simeq6.0\\times10^{-6}$, and $\\mathrm{Si/He}\\simeq8.0\\times10^{-6}$. The non-detection of H$\\alpha$ absorption implies $\\mathrm{H/He}<3\\times10^{-5}$. } \\end{figure*} ",
        "conclusions": "We have identified SDSS\\,0845+2257 as the third white dwarf known to host a metal-rich gaseous circumstellar disc. SDSS\\,0845+2257 has a similar temperature as SDSS\\,1043+0855 and SDSS\\,1228+1040, and all three are substantially hotter ($\\Teff\\sim20\\,000$\\,K) compared to the majority of the white dwarfs with dusty discs identified in surveys for infrared excess fluxes. The total roster of white dwarfs with gaseous and/or dusty debris discs stands at 13, of which $\\sim25$\\% were identified through the presence of \\Ion{Ca}{II} emission lines in their SDSS spectra. Taking the masses of the three white dwarfs with gaseous debris discs at face value, we find $<\\Mwd>=0.73\\pm0.07\\,\\Msun$, which is marginally higher than the mean mass of field white dwarfs \\citep{liebertetal05-1}. A speculative line of thought is that the potentially higher occurence of debris discs around more massive white dwarfs is related to the higher frequency of debris discs around A-stars compared to Sun-like stars \\citep{suetal96-1, trillingetal08-1}.  In contrast to SDSS\\,1228+1040 and SDSS\\,1043+0857, SDSS\\,0845+2257 has a helium-dominated atmosphere and the detection of Ca, Mg, and Si absorption lines qualifies it as a DBZ white dwarf. The presence of a circumstellar disc strongly suggests that ongoing accretion is the origin of the observed photospheric metals. The non-detection of H$\\alpha$ absorption in the SDSS spectrum of SDSS\\,0845+2257 implies that the hydrogen-to-metal ratio in the circumstellar material is at least a factor 1000 smaller than the solar value \\citep{dupuisetal93-2, friedrichetal99-1}. The most likely origin of the metal-rich material is the planetary debris, such as the tidal disruption of a rocky asteroid \\citep{jura03-1}.   Modelling the broad double-peaked \\Ion{Ca}{II} emission lines observed in SDSS\\,0845+2257 with a Keplerian gas disc confines the region of the \\Ion{Ca}{II} emission to a radial extension of $\\sim0.5$\\,\\Rsun\\ to $\\sim1.0$\\,\\Rsun. The line profiles display a large asymmetry, which we interpret as a significant asymmetry in the gaseous disc. The line profile shapes and equivalent widths of the \\Ion{Ca}{II} emission lines varied substantially in between two observations taken $\\sim4$ years apart. Following the detection of variability in the equivalent widths of the photospheric \\Ion{Ca}{H,K} absorption lines lines seen in G29--38 \\citep{vonhippel+thompson07-1}, this is the second piece of evidence indicating that the structure of circumstellar debris around white dwarfs undergo changes on relatively short time scales. \\citet{jura08-1} suggested that rather than a single tidal disruption event, white dwarfs with remnants of planetary systems may experience repeated accretion of small asteroids. Such recurrent events may potentially be related to variability such as observed in SDSS\\,0845+2257. Whatever the cause of the changes in the \\Ion{Ca}{II} emission lines, long-term monitoring of SDSS\\,0845+2257 is likely to offer direct dynamical insight into the evolution of debris discs around white dwarfs.      \\vspace*{-3ex}"
    },
    "0809/0809.0605_arXiv.txt": {
        "abstract": "We have made multi-epoch VLBI observations of $\\rmn{H_2O}$ maser emission in the massive star forming region {\\it IRAS\\/} 06061$+$2151 with the Japanese VLBI network (JVN) from 2005 May to 2007 October. The detected maser features are distributed within an 1\\arcsec$\\times$1\\arcsec\\ (2000$\\:$au$\\times$2000$\\:$au at the source position) around the ultra-compact \\,H\\,{\\small\\bf II} region seen in radio continuum emission. Their bipolar morphology and expanding motion traced through their relative proper motions indicate that they are excited by an energetic bipolar outflow. Our three-dimensional model fitting has shown that the maser kinematical structure in {\\it IRAS\\/} 06061$+$2151 is able to be explained by a biconical outflow with a large opening angle ($>$ 50\\degr). The position angle of the flow major axis coincides very well with that of the large scale jet seen in 2.1$\\:\\mu\\rmn{m}$ hydrogen emission. This maser geometry indicates the existence of dual structures composed of a collimated jet and a less collimated massive molecular flow. We have also detected a large velocity gradient in the southern maser group. This can be explained by a very small (on a scale of several tens of au) and clumpy (the density contrast by an order of magnitude or more) structure of the parental cloud. Such a structure may be formed by strong instability of shock front or splitting of high density core. ",
        "introduction": "The studies of massive star formation are essentially important for astrophysics. Because of their powerful protostellar or stellar winds and radiation, their formation is deeply concerned with activities of local star formation in galaxies (e.g., \\citealt{Elmegreen1977}). Once massive young stellar objects (MYSOs) are formed, their strong activities immediately interact with surrounding clouds. This interaction often results in fragmentation and dispersion of the parental cloud, and sometimes stops following star formation. However, another star formation is triggered frequently in such compressed fragment. In spite of numerous observations and numerical simulations (e.g., \\citealt{Elmegreen1998}, and references therein), quantitative and detailed studies on such interaction are still insufficient.  Direct observations of local kinematics around MYSOs would give intrinsic decision for such studies. If detailed information of the motion of gas in these interactions is available, we can examine the triggering effect quantitatively. However, observational studies of their kinematics have some serious difficulties. First, all of MYSOs are still deeply embedded in the parental molecular clouds because of their rapid evolution. Second, they are usually seen in crowded star clusters. Third, they are far distant from the Sun ($>\\:$500$\\:$pc). In such deep and crowded massive star forming regions, star formation process is complex since several effects, such as stellar winds, radiations and stellar or interstellar magnetic fields, etc, are mixed. Consequently, high angular resolution is essential to resolve into individual effects and local kinematics.  Observations of $\\rmn{H_2O}$ maser emission with a very long baseline interferometre (VLBI) is very suitable for such studies. Because the maser excitation requires a high density ($n_{\\rmn{H_2}}\\sim10^9$$\\:\\rmn{cm}^{-3}$) and a temperature ($\\sim$600$\\:$K), $\\rmn{H_2O}$ maser emission in a massive star forming region is usually associated with strong shock front induced by MYSOs (e.g., \\citealt{Elitzur1992}). Each maser source is composed of small gas clumps called maser feature, which has typically a size of 1$\\:$au and line of sight velocity width of a few $\\rmn{km}\\:\\rmn{s}^{-1}$ (e.g., \\citealt{Reid1981}). This scale is enough to resolve the spatial and kinematical structure mentioned above. Moreover, extremely high spatial resolution of VLBI obsevations enable us to detect small proper motions of a few milli-arcseconds (mas) per year, which reflects the local kinematics projected onto the celestial plane. For these reasons, $\\rmn{H_2O}$ maser emission is a good tracer of three-dimensional kinematics around MYSOs. In fact, many VLBI observations have shown highly systematic and directional proper motions of $\\rmn{H_2O}$ maser features (e.g., \\citealt*{Genzel1981,Imai2000,Seth2002,Goddi2005,Moscadelli2007}), which are thought to trace propagation of shock fronts induced by a bipolar outflow, a protosteller wind, or an expansion of ultra-compact (UC)\\,H\\,{\\small\\bf II} region.  In this paper, we report results of the VLBI observations of $\\rmn{H_2O}$ maser source WB755 in the massive star forming region {\\it IRAS\\/} 06061+2151. {\\it IRAS\\/} 06061+2151 is one of the most massive star forming region in the Gemini OB1 cloud complex \\citep*{Carpenter1995b}. Its photometric distance is 2.0$\\:$kpc \\citep*{Carpenter1995a}. Several observations have indicated active star formation in this region (e.g., \\citealt*{Hanson2002,Hill2005,Kumar2006,Thompson2006}). This region is deeply embedded in the extended ($\\sim$1$\\:$pc) and massive ($\\sim$7000$\\:\\rmn{M}_\\odot$) molecular cloud complex observed in C$^{18}$O \\citep{Saito2007}. There is a compact cluster seen in near-infrared (NIR), AFGL 5182 at the centre \\citep{Carpenter1995b}. Five bright {\\it Ks}-band sources are identified in this cluster \\citep{Anandarao2004}. They have concluded these sources are early B-type protostars (S1-S5 in \\citealt{Anandarao2004}). They also have detected knot-like structure which suggests a protostellar jet. In addition, three weak centimetre continuum sources have been detected with the Very Large Array (VLA) in B configuration, and all of them are thought to be UC\\,H\\,{\\small\\bf II} regions \\citep*{Kurtz1994}. Only the most southern and brightest UC\\,H\\,{\\small\\bf II} region is coincident with the sourthern {\\it Ks}-band source, S5. The other two UC\\,H\\,{\\small\\bf II} regions are also very close to S4 source, but they are clearly individual sources. The geometry between these five {\\it Ks}-band sources and three centimetre continuum sources is clearly described in \\citet{Bik2006}. The maser source, WB755, is associated with the most northern UC\\,H\\,{\\small\\bf II} region, G188.794+1.031, in \\citet{Kurtz1994} (see section 3). This region is too small to resolve with 1\\arcsec\\ resolution in the VLA observation at 8.4$\\:$GHz. This scale corresponds to about 2000$\\:$au at 2.0$\\:$kpc and about ten times smaller than the typical size of UC\\,H\\,{\\small\\bf II} region \\citep{Hoare2007}. It is probable that this UC\\,H\\,{\\small\\bf II} region is extremely young and its ionization source is still in the accretion phase.  We describe the observations and data reduction in section 2, and the observational result in section 3. Discussions and model fitting are described in section 4. We summarize our conclusions in section 5. ",
        "conclusions": "We made VLBI observations of H$_2$O masers in massive star forming region {\\it IRAS\\/} 06061+2151 with JVN from 2005 May to 2007 October. The H$_2$O masers are clustered within 1\\arcsec$\\times$1\\arcsec\\ area around extremely young UC\\,H\\,{\\small\\bf II} region. They are divided into several groups based on their radial velocities, and show expanding motions. Their alignment and detected three-dimensional motion indicate they are associated with a bipolar outflow from the source in the UC\\,H\\,{\\small\\bf II} region. The southern blue-shifted group shows an arc-like alignment within several tens of au. Their continuous alignment indicate they are excited in a single gas layer. They also show two clearly different kinematics, which are the vertical part (V-wing) of 90--180$\\:$km$\\:$s$^{-1}$ and the horizontal part (H-wing) of 30--65$\\:$km$\\:$s$^{-1}$. This may be caused by far small (several tens of au) and clumpy (10$^7$--$10^8\\:$cm$^{-3}$) structure of parental cloud core. Such a clumpy structure is expected to be induced by strong instability of shock front (e.g., expansion of UC\\,H\\,{\\small\\bf II} region). Further high-resolution MHD simulation is desirable to examine this effect.  The biconical flow has an opening angle of 54\\degr\\ and the axis along NW-SE direction based on our model fitting. This NW-SE axis coincides very well with the large scale jet seen in 2.1$\\:\\mu\\rmn{m}$ hydrogen emission. These geometry and large opening angle are thought to indicate the dual structure of highly collimated jets and massive and largely opened molecular flow in this region. We estimate density of the outflow from the deceleration of the V-wing of $\\sim$10$^{6}\\:$cm$^{-3}$, although such a massive outflow has not been detected in this region yet. This may mean such a flow is still compact and deeply embedded in the dense clump, and we can only see an collimated jet which can penetrates surrounding cloud. In this point of view, we conclude that these two types of flow is associated with the ionizing source of G188.794+1.031. There are, in addition, two B-type protostars located on the jet axis. These two sources are seemed to be younger than the ionizing source from the absence of any ionized region and comparison of a spectral type. This time sequance and their geometry may indicate that the formation of two protostars is affected by the jet and outflow. With additional observations which has higher resolution and accuracy, they can be a good example of sequential star formation in a massive cluster."
    },
    "0809/0809.0269_arXiv.txt": {
        "abstract": "We present here additional results of a spectropolarimetric survey of a small sample of stars ranging from spectral type M0 to M8 aimed at investigating observationally how dynamo processes operate in stars on both sides of the full convection threshold (spectral type M4).  The present paper focuses on early M stars (M0--M3), i.e.\\ above the full convection threshold. Applying tomographic imaging techniques to time series of rotationally modulated circularly polarised profiles collected with the NARVAL spectropolarimeter, we determine the rotation period and reconstruct the large-scale magnetic topologies of 6 early M dwarfs. We find that early-M stars preferentially host large-scale fields with dominantly toroidal and non-axisymmetric poloidal configurations, along with significant differential rotation (and long-term variability); only the lowest-mass star of our subsample is found to host an almost fully poloidal, mainly axisymmetric large-scale field ressembling those found in mid-M dwarfs.  This abrupt change in the large-scale magnetic topologies of M dwarfs (occuring at spectral type M3) has no related signature on X-ray luminosities (measuring the total amount of magnetic flux);  it thus suggests that underlying dynamo processes become more efficient at producing large-scale fields (despite producing the same flux) at spectral types later than M3.  We suspect that this change relates to the rapid decrease in the radiative cores of low-mass stars and to the simultaneous sharp increase of the convective turnover times (with decreasing stellar mass) that models predict to occur at M3;  it may also be (at least partly) responsible for the reduced magnetic braking reported for fully-convective stars. ",
        "introduction": "Activity and magnetic fields are ubiquitous to cool stars of all spectral types \\citep[e.g.,][]{Saar85, Donati97}.  The analogy with the Sun suggests that these fields are produced through dynamo mechanisms operating in a thin interface layer at the base of the convective envelope, where angular rotation gradients are strongest.  Starting from a weak poloidal configuration, differential rotation progressively winds the field around the star, building a strong toroidal belt at the base of the convective zone;  cyclonic turbulence then restores a weak poloidal field (with a polarity opposite to that of the initial one) once the toroidal field has grown unstable \\citep{Parker55}. However, there is still considerable controversy (even for the Sun itself) on how and where exactly the field is amplified and on which physical processes (differential rotation, meridional circulation) are mainly controlling the magnetic cycle \\citep[e.g.,][]{Charbonneau05}.  Observing stars other than the Sun is of obvious interest for this question as they provide a direct way of studying how dynamo processes depend on fundamental stellar parameters such as mass (and thus convective depth) and rotation rate.  In this respect, low-mass fully-convective stars are particularly interesting as they host no interface layer (where dynamo processes presumably concentrate in the Sun), and nevertheless show both strong magnetic fields \\citep[e.g.,][]{Johns96} and intense activity \\citep[e.g.,][]{Delfosse98}.  Many theoretical models were proposed to attempt resolving this issue \\citep[e.g.,][]{Durney93, Dobler06, Browning08} but observations of the large-scale magnetic fields at the surfaces of fully-convective stars, and in particular of their poloidal and toroidal components, are still rare.  Thanks to time-resolved spectropolarimetric observations of cool stars, we are now able to recover information on how magnetic fields distribute at (and emerge from) the surfaces of cool active stars \\citep[e.g.,][]{Donati03a}.  With the advent of new generation instruments optimised in this very purpose \\citep[ESPaDOnS at the 3.6m Canada-France-Hawaii Telescope and NARVAL at the 2m T\\'elescope Bernard Lyot in France, ][]{Donati03c}, this technique can now be also applied to faint M dwarfs.  The first attempt focussed on the fully-convective rapidly-rotating M4 star V374~Peg and revealed that, against all theoretical expectations, fully-convective M dwarfs are apparently very efficient at producing strong large-scale mainly-axisymmetric poloidal fields \\citep{Donati06} despite very small levels of differential rotation.  A follow-up study confirmed this point and further demonstrated that the magnetic configuration of V374~Peg is apparently stable on timescales of $\\simeq1$~yr \\citep{Morin08}.  To investigate this issue in more details, we embarked in a spectropolarimetric survey of a small sample of M dwarfs, located both above and below the full-convection threshold (corresponding to a mass of 0.35~\\msun\\ and to a spectral type of M4).  The first results of this survey, focussing mainly on active M4 dwarfs, confirms the results obtained on V374~Peg and demonstrate that strong large-scale mainly-axisymmetric poloidal fields are indeed fairly common in fully-convective mid-M active dwarfs \\citep[][hereafter M08]{Morin08b}. In this new paper, we present the results for the 6 early M dwarfs that we observed in this survey, namely DT~Vir, DS~Leo, CE~Boo, OT~Ser, GJ~182 and GJ~49, with spectral types ranging from M0 to M3 (see Table~\\ref{tab:sample}). After briefly describing the observations and the magnetic modelling method that we use, we detail the results obtained for each star and discuss their implication for our understanding of dynamo processes in cool active stars.  \\begin{table*} \\caption{Fundamental parameters of our sample early-M dwarf stars. Columns 1 to 6 respectively list the name, the spectral type, the mass \\mstar\\ (derived from the Hipparcos distance and the J, H and K magnitudes using the mass-luminosity relations of \\citealt{Delfosse00} except for GJ~182; for which we used the evolutionary models of \\citealt{Baraffe98}, see text), the logarithmic bolometric luminosity \\logLb\\ (derived from the mass and the models of \\citealt{Baraffe98}), the logarithmic relative X-ray luminosity \\logRx\\ (i.e., \\logLX, from \\citealt{Kiraga07} or from the NEXXUS data base, \\citealt{Schmitt04}, or from \\citealt{Wood94} for GJ~49), the projected rotation velocity (this paper, accuracy $\\simeq1$~\\kms), the rotation period (this paper),  the convective turnover time $\\tau_{\\rm c}$ (from \\citealt{Kiraga07}), the effective Rossby number $Ro=\\prot/\\tau_{\\rm c}$, the radius \\rstar\\ (predicted by the theoretical models of \\citealt{Baraffe98}) and the assumed inclination angle of the rotation axis to the line-of-sight (this paper). } \\begin{tabular}{lccccccccccc} \\hline Star & ST & \\mstar & \\logLb & \\logRx & \\vsini & \\prot & $\\tau_{\\rm c}$ &  $Ro$ & \\rstar & $i$ \\\\ &    & (\\msun) & (\\eps)&        & (\\kms) & (d) & (d) & & (\\rsun) & (\\degr) \\\\ \\hline GJ 182           & M0.5 & 0.75 & 32.7 & --3.1 & 10 & 4.35 & 25 & 0.174 & 0.82 & 60 \\\\ DT Vir / GJ 494A & M0.5 & 0.59 & 32.3 & --3.4 & 11 & 2.85 & 31 & 0.092 & 0.53 & 60 \\\\ DS Leo / GJ 410  & M0   & 0.58 & 32.3 & --4.0 &  2 & 14.0 & 32 & 0.438 & 0.52 & 60 \\\\ GJ 49            & M1.5 & 0.57 & 32.3 & $<-4.3$ &  1 & 18.6 & 33 & 0.564 & 0.51 & 45 \\\\ OT Ser / GJ 9520 & M1.5 & 0.55 & 32.2 & --3.4 &  6 & 3.40 & 35 & 0.097 & 0.49 & 45 \\\\ CE Boo / GJ 569A & M2.5 & 0.48 & 32.1 & --3.7 &  1 & 14.7 & 42 & 0.350 & 0.43 & 45 \\\\ \\hline \\end{tabular} \\label{tab:sample} \\end{table*} ",
        "conclusions": "\\label{sec:disc}  We report in this paper the results of our spectroscopic survey of M dwarfs; following M08 (concentrating on mid-M dwarfs), we describe here the Zeeman signatures and the large-scale magnetic topologies we observed on 6 early-M dwarfs (from M0 to M3). We also determined or confirmed the rotation period of all stars (ranging from 2.8 to 18.6~d), and detected significant surface differential rotation in 4 of them (with a strength comparable to that of the Sun).  \\begin{figure} \\center{\\includegraphics[scale=0.55,angle=-90]{fig/mdw_49fit.ps}} \\caption[]{Same as Fig.~\\ref{fig:494fit} for GJ~49.  } \\label{fig:49fit} \\end{figure}  \\begin{figure*} \\center{\\includegraphics[scale=0.65]{fig/mdw_49map.ps}} \\caption[]{Same as Fig.~\\ref{fig:494map} for GJ~49.  } \\label{fig:49map} \\end{figure*}  The magnetic fields we detect in early-M dwarfs are weak (typically a few tens of G), smaller in particular than those found in mid-M dwarfs (M08) by typically a factor of 5 (see Fig.~\\ref{fig:diag}) with a sharp transition occuring at 0.4~\\msun\\ (see Fig.~\\ref{fig:ross}, left panel).  The large-scale magnetic topologies we derive are also significantly different, involving a much larger fraction of toroidal fields and a lower axisymmetric degree of poloidal fields whenever $\\mstar>0.5$~\\msun\\ (5 stars in the present sample); below 0.5~\\msun, the poloidal field is largely dominant and axisymmetric and its strength is increasing rapidly as mass decreases.  We also observe that the typical lifetime of the large-scale magnetic topology is very different on both sides of the 0.4--0.5~\\msun\\ threshold, with lifetimes smaller than a few months on the hot side and longer than 1~yr on the cool side (M08). This threshold is very sharp and well defined, with little apparent dependence with the rotation period.  At this stage, it is interesting to consider the effective Rossby number $Ro$, defined as $Ro=\\prot/\\tau_{\\rm c}$ where $\\tau_{\\rm c}$ is the convective turnover time\\footnote{The $\\tau_{\\rm c}$ values that we use here are those of \\citet{Kiraga07}, determined empirically from relative X-ray luminosities of stars with different masses and rotation periods.  At masses of $\\simeq1$~\\msun, they match the usual value of $\\simeq15$~d;  they steeply increase with decreasing mass below masses of $\\simeq0.6$~\\msun.};  in particular, $Ro$ is a convenient parameter for comparing the strength of dynamo action (and, e.g., \\logRx\\ that indirectly reflects dynamo action through coronal heating) in stars with different masses.  Figure~\\ref{fig:ross} (right panel) illustrates how \\logRx\\ smoothly varies with $Ro$ for both early- and mid-M dwarfs studied here and in M08;  we find that \\logRx\\ increases steeply with decreasing $Ro$ until $Ro\\simeq0.1$ where \\logRx\\ saturates at a level of about --3.0, in good agreement with previous studies \\citep[e.g.,][]{James00}. Only GJ~182 (at $Ro=0.17$) lies slightly above the overall trend, as a likely consequence of its extreme youth and the related differences in its internal structure.  \\begin{figure*} \\includegraphics[scale=0.6,angle=-90]{fig/mdw_diag.ps} \\caption[]{Basic properties of the large-scale magnetic properties of early- and mid-M dwarfs as a function of stellar mass and rotation period.  Symbols size indicates magnetic energies, symbol colour illustrates the field configuration (blue and red for purely toroidal and purely poloidal fields respectively) while symbol shape depicts the degree of axisymmetry of the poloidal field component (decagon and stars for a purely axisymmetric and purely non-axisymmetric poloidal fields respectively).  Results for early-M stars are from this paper and results for mid-M stars are from M08.  The full, dashed and dash-dot lines respectively trace the location of the $Ro=1$, 0.1 and 0.01 contours, while the dotted line shows the theoretical full-convection threshold ($\\mstar\\simeq0.35$~\\msun).} \\label{fig:diag} \\end{figure*}  \\begin{figure*} \\center{\\includegraphics[scale=0.37,angle=-90]{fig/mdw_ross.ps}\\hspace{5mm} \\includegraphics[scale=0.37,angle=-90]{fig/mdw_lx.ps}} \\caption[]{Reconstructed magnetic energy (left) and relative X-ray luminosity (with respect to the bolometric luminosity, right) as a function of $Ro$ for stars studied in this paper and in M08.  Stars with masses larger and smaller than 0.4~\\msun\\ are shown as green squares and red circles respectively.  In the left panel, measurements at different epochs (whenever available) are shown for each star to illustrate the typical scatter expected from temporal variability.   } \\label{fig:ross} \\end{figure*}  We note that the abrupt step in the large-scale magnetic energy between early- and mid-M dwarfs (see Fig.~\\ref{fig:ross}, left panel) apparently correlates better with stellar mass than with $Ro$;  at $Ro=0.05$, AD~Leo exhibits a magnetic flux compatible with that of all other early M dwarfs, but significantly weaker than that of EV~Lac (at $Ro=0.07$) and YZ~CMi (at $Ro=0.04$).  More data are needed to confirm this point, especially for high-mass rapid rotators (i.e., having $Ro\\simeq0.01$ and $\\mstar>0.5$~\\msun) and low-mass slow rotators (with $Ro>0.1$ and $\\mstar<0.4$~\\msun).  We also note that this abrupt step does not show up in the \\logRx\\ vs $Ro$ plot of Fig.~\\ref{fig:ross} (right panel); this is presumably because X-rays are sensitive to overall magnetic energies while we are only sensitive to the largest scales.  Our result therefore suggests that, at some specific stellar mass ($\\simeq0.4$~\\msun, rather than at some specific $Ro$), dynamo processes become suddenly much more efficient at triggering large-scale magnetic fields;  we also observe that, at more or less the same mass ($\\simeq0.5$~\\msun), large-scale topologies of M dwarfs become dominantly poloidal and axisymmetric.  Note however that, even in the case of mid-M dwarfs, the large-scale fields we derive are significantly smaller than the corresponding equipartition field (a few kG, M08).  \\begin{table*} \\caption{Properties of the large-scale field topologies derived in the present study.  For each star, observations at different epochs are listed separately.  The table lists sequentially the name of the star, the mass, the rotation period, the effective Rossby number and the logarithmic relative X-ray luminosity (taken from Table~\\ref{tab:sample}), the surface angular rotation shear between the equator and pole \\dom\\ (whenever detected), the derived filling factor $f$ (whenever applicable), the reconstructed magnetic energy and flux (i.e., $<B^2>$ and $<B>$), and the fractional energies in the poloidal field, the poloidal dipole ($\\ell=1$) modes, the poloidal quadrupole ($\\ell=2$) modes, the poloidal octupole ($\\ell=3$) modes, and the poloidal axisymmetric ($m<\\ell/2$) modes.  All field components with $\\ell\\leq3$ are required to fit the data at noise level.  } \\begin{tabular}{rccccccccccccc} \\hline Star  & \\mstar & \\prot & $Ro$ & $\\logRx$ & \\dom & $f$ & $<B^2>$ & $<B>$ & pol & dip & qua & oct & axi \\\\ & (\\msun) & (d)  &      &            &  (\\rpd) &  & ($10^4$~G$^2$) & (G) &  &  &  &  & \\\\ \\hline GJ 182 2007 & 0.75 & 4.35 & 0.174 & --3.1 & $0.06\\pm0.03$   &      & 3.95 & 172 & 0.32 & 0.48 & 0.18 & 0.14 & 0.17 \\\\ DT Vir 2007 & 0.59 & 2.85 & 0.092 & --3.4 &                 &      & 2.78 & 145 & 0.38 & 0.64 & 0.17 & 0.08 & 0.12 \\\\ 2008 &      &      &       &       & $0.060\\pm0.006$ &      & 3.21 & 149 & 0.53 & 0.10 & 0.17 & 0.17 & 0.20 \\\\ DS Leo 2007 & 0.58 & 14.0 & 0.438 & --4.0 &                 &      & 1.24 & 101 & 0.18 & 0.52 & 0.37 & 0.08 & 0.58 \\\\ 2008 &      &      &       &       & $0.076\\pm0.020$ &      & 1.05 &  87 & 0.20 & 0.52 & 0.31 & 0.07 & 0.16 \\\\ GJ 49  2007 & 0.57 & 18.6 & 0.564 &$<-4.3$&                 &      & 0.09 &  27 & 0.48 & 0.71 & 0.20 & 0.07 & 0.67 \\\\ OT Ser 2007 & 0.55 & 3.40 & 0.097 & --3.4 &                 & 0.05 & 2.28 & 136 & 0.80 & 0.47 & 0.19 & 0.18 & 0.86 \\\\ 2008 &      &      &       &       & $0.12\\pm0.02$   & 0.10 & 2.38 & 123 & 0.67 & 0.33 & 0.17 & 0.21 & 0.66 \\\\ CE Boo 2008 & 0.48 & 14.7 & 0.350 & --3.7 &                 & 0.05 & 1.48 & 103 & 0.95 & 0.87 & 0.06 & 0.03 & 0.96 \\\\ \\hline \\end{tabular} \\label{tab:syn} \\end{table*}  Significant surface toroidal fields are detected even in DS~Leo and GJ~49, i.e., the two slowest rotators with masses larger than 0.5~\\msun;  it suggests that the transition between mainly poloidal and mainly toroidal fields in $\\mstar>0.5$~\\msun\\ stars occurs at $Ro\\simeq0.5-1.0$, with the Sun located on the other side of this boundary (at $Ro\\simeq1.5-2.0$).  Note that this boundary coincides with the sharp onset of photometric variability in convective stars (occuring below $Ro\\simeq0.7$, \\citealt{Hall91}). With a poloidal field concentrating 70--80\\% of the reconstructed magnetic energy, OT~Ser is off this trend;  we suspect that this is due to its proximity with the 0.5~\\msun\\ sharp threshold below which magnetic topologies become dominantly poloidal.  Early-M dwarfs are found to show significant differential rotation;  the values we obtain for the surface angular rotation shear \\dom\\ ranges from 0.06 to 0.12~\\rpd, i.e., from once to twice the strength of the surface latitudinal shear of the Sun.  Our detection is further confirmed by the small (but significant) differences between the rotation periods we measure and the values reported in the literature (derived from photometric fluctuations) and by the short lifetimes of the large-scale field topologies (quickly distorted beyond recognition by differential rotation). Previous Doppler imaging studies of early-M dwarfs with very fast rotation ($Ro\\simeq0.01$) report that differential rotation is very small \\citep{Barnes05};  our study suggests that the situation may significantly differ in moderate rotators like those we considered. Our result is also different from what is observed in mid-M dwarfs where differential rotation is very small (a few \\mrpd\\ at most, i.e., more than 10 times smaller than that of early-M dwarfs, M08) and large-scale magnetic topologies long-lived (M08).  It is not clear yet what this difference is due to;  while small $Ro$ may contribute at freezing differential rotation, this is likely not the only relevant parameter for this problem (e.g., with DT~Vir and EV~Lac showing respectively significant and no differential rotation despite their similar $Ro$).  The sharp transition that we report between the magnetic (and differential rotation) properties of early- and mid-M dwarfs is surprising at first glance;  naively, one would expect the properties of large-scale magnetic fields to change smoothly with stellar mass as the radiative core gets progressively smaller.  From the evolutionary models of \\citet{Siess00}, we however note that the outer radius of the radiative core of early-M dwarfs is changing very quickly with stellar mass, from about 0.5~\\rstar\\ for a 0.5~\\msun\\ star to a negligible fraction for a 0.4~\\msun\\ star. We speculate that this sharp transition is the main reason for the abrupt magnetic threshold that we report here.  The rapid increase in empirical convective turnover times occuring at about the same location \\citep{Kiraga07} also likely contributes at making the transition between both dynamo regimes very sharp.  The most recent dynamo simulations of fully convective M dwarfs \\citep{Browning08} (carried out for $Ro\\simeq0.01$) are successful at reproducing the frozen differential rotation that we observe (M08);  they however predict the presence of strong toroidal fields that we do not see in mid-M dwarfs with similar $Ro$. We speculate that the abrupt change in the large-scale magnetic topology of M dwarfs that we report here to occur at spectral type M3 may also be (at least partly) responsible for the reduced magnetic braking observed for stars later than M3 \\citep[e.g.,][]{Delfosse98};  MHD simulations of magnetic winds are necessary to estimate quantitatively whether the observed change in the large-scale magnetic topology can indeed explain the longer spin-down timescales.  In the two stars having \\vsini\\ measured with sufficient precision (i.e., $\\vsini\\geq10$, GJ~182 and DT~Vir), we find that $\\rstar \\sin i$ (equal to $0.86\\pm0.09$~\\rsun\\ and $0.62\\pm0.06$~\\rsun\\ respectively) is already larger than the predicted radius from theoretical models \\citep{Baraffe98};  while this could result from overestimating the true age (and hence underestimating the true radius) of GJ~182, this explanation does not apply for DT~Vir, for which we conclude that the observed radius is truly larger (by at least 10\\% and potentially as much as 30\\%) that what theoretical models predict. A similar conclusion is reached for V374~Peg \\citep{Donati06, Morin08};  furthermore, M08 obtains that $\\rstar\\sin i$ is equal to the predicted theoretical radius (within the error bars) for 4 other active stars, suggesting again that \\rstar\\ is larger than expected. Following \\citet{Chabrier07}, we propose that this effect is a direct consequence of magnetic fields getting strong enough (and hence saturating the dynamo, see Fig.~\\ref{fig:ross}) to affect the energy transport throughout the convective zone and hence the radius. Our results therefore independently confirm the report that cool low-mass active stars, either single \\citep{Morales08} or within close eclipsing binaries \\citep{Ribas06}, usually have oversized radii with respect to inactive stars of similar spectral types.  We also detect significant RV fluctuations (with a full amplitude of up to 0.40~\\kms) in the 3 very active stars of our sample (with rotation periods smaller than 5~d).  For the most active ones (DT~Vir and GJ~182, showing the largest RV modulation), the RV variations correlate reasonably well (though not perfectly) with longitudinal fields, suggesting that the origin of the variations is indeed the magnetic field (and the underlying activity).  It also suggests that spectropolarimetric observations should be carried out simultaneously with RV measurements of active stars to enable filtering out efficiently the activity jitter from the RV signal;  this technique may prove especially useful when looking at Earth-like habitable planets orbiting around M dwarfs in the future, e.g., with a nIR spectropolarimeter such as SPIRou (a nIR counterpart of ESPaDOnS, proposed for CFHT).  Our spectropolarimetric survey is an on-going study;  we are now concentrating on late-M dwarfs (M5-M8) to derive similar observational constraints about the large-scale magnetic topologies of stars in the yet unexplored 0.08--0.20~\\msun\\ region of Fig.~\\ref{fig:diag} to investigate how dynamo processes operate down to the brown dwarf threshold, i.e., when stellar atmospheres get so cool that they start to decouple from their magnetic fields."
    },
    "0809/0809.2746_arXiv.txt": {
        "abstract": "{\\lb\\ (\\aador) is an eclipsing, close, post common-envelope binary (PCEB) consisting of an sdOB primary star and an unseen secondary with an extraordinarily low mass ($M_2 \\approx 0.066\\,\\mathrm{M_\\odot}$) -- formally a brown dwarf. A recent NLTE spectral analysis shows a discrepancy with the surface gravity, which is derived from analyses of radial-velocity and lightcurves. } {We aim at precisely determining of the photospheric parameters of the primary, especially of the surface gravity, and searching for weak metal lines in the far UV. } {We performed a detailed spectral analysis of the far-UV spectrum of \\lb\\ obtained with FUSE by means of state-of-the-art NLTE model-atmosphere techniques. } {A strong contamination of the far-UV spectrum of \\lb\\ by interstellar line absorption hampers a precise determination of the photospheric properties of its primary star. Its effective temperature (\\Teffw{42}) was confirmed by the evaluation of new ionization equilibria. For the first time, phosphorus and sulfur have been identified in the spectrum of \\lb. Their photospheric abundances are solar and 0.01\\,times solar, respectively. From the \\Ionww{C}{3}{1174-1177} multiplet, we can measure the rotational velocity $v_\\mathrm{rot}=35\\pm 5\\,\\mathrm{km/sec}$ of the primary of \\lb\\ and confirm that the rotation is bound. From a re-analysis of optical and UV spectra (analogue to Rauch 2000), we determine a slightly higher surface gravity \\loggw{5.3\\pm 0.1} compared to Rauch (2000, \\loggw{5.2\\pm 0.1}). } {The rotational velocity of the primary of \\lb\\ is consistent with a bound rotation. The higher \\logg\\ reduces the discrepancy in mass determination in comparison to analyses of radial-velocity and lightcurves. However, the problem is not completely solved. } ",
        "introduction": "\\label{sect:introduction} The eclipsing binary system \\lb\\, (\\aador) is a blue foreground object of the LMC at a spectroscopic distance of \\mbox{$d=396\\,\\mathrm{pc}$} \\citep{Rauch00}. It consists of aa sdOB with an effective temperature of \\Teffw{42} and a low-mass companion from which no direct spectroscopic information has been obtained yet. Due to its mass of about $0.066\\,\\mathrm{M_\\odot}$ it formally lies in the brown-dwarf range. The orbital period is about 0.26\\,d and the inclination is $i=90\\degr$.  The system was analyzed several times, starting with \\citet{Kilkenny79,Kilkenny81} and \\citet{Paczynski80}, who established the basic parameters. First NLTE analyses of the primary were performed by \\citet{Kudritzki82} and \\citet{Lynas-Gray84}. Details on the history of this object can be found in \\citet{Rauch00, Rauch04}.  Recent spectral analyses by means of NLTE model-atmospheres that were based on optical (ESO CASPEC\\footnote{European Southern Observatory, Cassegrain Echelle Spectrograph}) and ultraviolet (IUE\\footnote{International Ultraviolet Explorer}) observations \\citep{Rauch00} have shown a discrepancy with analyses of radial-velocity and lightcurves \\citep{Hilditch96,Hilditch03}. \\citet{Rauch00} determined a surface gravity of \\loggw{5.21\\pm 0.1}, while \\citet{Hilditch03} obtained \\loggw{5.45 - 5.51} from lightcurve and mass function. Because the analysis of \\citet{Rauch00} suffered from the long exposure times (some hours) hence smearing, due to the orbital movement of the available spectra, \\citet{RauchWerner03} measured the radial-velocity curve from optical spectra with short (180\\,sec) exposure times obtained with ESO's VLT\\footnote{Very Large Telescope} and UVES\\footnote{Ultraviolet and Visual Echelle Spectrograph}. They obtained a radius $r_1=0.169\\,\\mathrm{R_\\odot}$ which is smaller than the $r_1=0.236\\,\\mathrm{R_\\odot}$ found by \\citet{Rauch00}. Since the stellar radius depends on $g$ this may be a hint of a higher value than \\loggw{5.21}. However, data reduction of the Balmer lines in the UVES spectra used by \\citet{RauchWerner03} was not very accurate, so a precise determination of \\logg\\ of the primary by means of NLTE modeling techniques is still lacking.  Consequently, high-resolution and high-S/N observations (exposure times of 200\\,sec each) in the far-UV range were performed with the FUSE\\footnote{Far Ultraviolet Spectroscopic Explorer} satellite. The FUSE wavelength range ($904\\,\\mathrm{\\AA} < \\lambda < 1187\\,\\mathrm{\\AA}$) covers the hydrogen Lyman series except for Ly\\,$\\alpha$. The series decrement is a sensitive indicator for \\logg.  A detailed spectral analysis of the FUSE observation of \\lb\\ is described in \\sK{sect:analysis}. Since the far-UV spectrum of \\lb\\ turned out to be strongly contaminated by interstellar absorption, we modeled the ISM line absorption \\sK{sect:observations} in order to distinguish weak lines of iron-group elements (here: Ca, Sc, Ti, V, Cr, Mn, Fe, Co, Ni; \\se{sect:analysis}). ",
        "conclusions": "\\label{sect:results}  We performed an NLTE spectral analysis of FUSE observations of the post common-envelope binary \\lb. The short FUSE exposure times (200\\,sec) allowed us to measure the rotational velocity $v_\\mathrm{rot}=35 \\pm 5\\,\\mathrm{km/sec}$ of the primary star of \\lb. This is in good agreement with the $v_\\mathrm{rot}=34.7 - 38.6\\,\\mathrm{km/sec}$ given by \\citet{Hilditch03}, and it confirms that the rotation of \\lb\\ is bound.  A re-analysis of optical spectra \\citep[cf\\@.][]{Rauch00} has shown that \\logg\\ is about 0.1 higher than given by \\citet[][\\loggw{5.2}]{Rauch00}, who assumed bound rotation because the system has been classified to be a post common-envelope binary consisting of an sdOB and a main-sequence star \\citep{deKoolRitter93} where the common-envelope phase is much longer than the synchronization time. Since the primary's radius was assumed to be larger ($r_1=0.236\\,\\mathrm{R_\\odot}$) due to the lower \\logg, a higher $v_\\mathrm{rot}=45\\,\\mathrm{km/sec}$ was adopted then for the analysis. From $v_\\mathrm{rot}=35\\pm 5\\,\\mathrm{km/sec}$ and an orbital period of $P=22\\,597.033\\,\\mathrm{sec}$ \\citep{Kilkenny00}, we can calculate a stellar radius of $r_1=0.181\\pm 0.025\\,\\mathrm{R_\\odot}$.  The spectral analysis is hampered by a strong ISM contamination. We used \\owens\\ to model the ISM line absorption qualitatively in order to identify and model stellar lines. Our NLTE models consider opacities of 18 species from H -- Ni. It is obvious that the iron-group elements (here: Ca -- Ni) contribute to the opacity. However, we were not able to identify any individual iron-group line in the FUSE observation.  We identified phosphorus and sulfur lines in the FUSE spectrum. For phosphorus we can determine a solar abundance. The sulfur abundance is surprisingly low, and we determined a 0.01 times solar abundance. The photospheric abundances of \\lb\\ are summarized in \\ab{fig:abundances}. The derived abundance pattern reflects the interplay of gravitational settling and radiative levitation \\citep[cf\\@.][]{Rauch00}.   \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics[]{0738_f14.eps}} \\caption[]{Photospheric abundances of \\lb\\ determined by \\citet{Rauch00} and in this work (Si, P, S, Ni). [X] denotes $\\log \\mathrm{(mass~fraction / solar~mass~fraction)}$ of species X.} \\label{fig:abundances} \\end{figure}   The ``gravity problem'' is still unsolved although the surface gravity is slightly higher: \\loggw{5.3} rather than the 5.2 found by \\citet{Rauch00}. If we take the radius of the primary as $r_1=0.181 R_\\odot$ as noted above, and its mass $M_1=0.330M_\\odot$ \\citep{Rauch00}, then we obtain {\\boldmath $\\log g = \\log \\left(G M_1/R_1^2\\right) = 5.44 \\pm 0.13$\\unboldmath} (where G is the gravitational constant). Note that $\\log g \\approx 5.5$ is necessary for good agreement \\citep[cf\\@.][]{Rauch00}.  An additional uncertainty \\citep[cf\\@.][]{Rauch00} is the lack of appropriate evolutionary models for post common-envelope binaries to compare so as to derive the primary's mass.  High-resolution and high S/N observations in the near-UV wavelength range is highly desirable when searching for the weak lines of iron-group elements. \\aba{fig:COS} demonstrates that lines of all elements from Ca -- Ni but Sc and V should be detectable. The search for signatures of the ``heated-up'' secondary in IR observations in order to further investigate the nature of this star also appears possible, and its contribution to the continuum at $\\lambda = 12\\,\\mu$ is a few percent.   \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics[]{0738_f15.eps}} \\caption[]{Flux ratios (in the far UV) of synthetic spectra for \\lb\\ (calculated with abundances from \\ab{fig:abundances} and Kurucz's LIN lines), which include only line opacities of one of the elements Ca -- Ni and a synthetic spectrum, which does not include any Ca -- Ni line.} \\label{fig:COS} \\end{figure}"
    },
    "0809/0809.1745_arXiv.txt": {
        "abstract": "This paper presents an analysis of the optical night sky brightness and extinction coefficient measurements in $UBVRI$ at the Indian Astronomical Observatory (IAO), Hanle, during the period 2003$-$2008. They are obtained from an analysis of CCD images acquired at the 2 m Himalayan Chandra Telescope (HCT) at IAO. Night sky brightness was estimated using 210 HFOSC images obtained on 47 nights and covering the declining phase of solar activity cycle-23. The zenith corrected values of the moonless night sky brightness in mag arcsec$^{-2}$ are 22.14 $\\pm$ 0.32 ($U$), 22.42 $\\pm$ 0.30 ($B$), 21.28 $\\pm$ 0.20 ($V$), 20.54 $\\pm$ 0.37 ($R$) and 18.86 $\\pm$ 0.35 ($I$) band. This shows that IAO is a dark site for optical observations. No clear dependency of sky brightness with solar activity (implied by the 10.7 cm solar flux) is found. Extinction values at IAO are derived from an analysis of 1325 images over 58 nights. They are found to be 0.36 $\\pm$ 0.07 in $U$-band, 0.21 $\\pm$ 0.04 in $B$-band, 0.12 $\\pm$ 0.04 in $V$-band, 0.09 $\\pm$ 0.04 in $R$-band and 0.05 $\\pm$ 0.03 in $I$-band. On average, extinction during the summer months is slightly larger than that during the winter months. This might be due to an increase of dust in the atmosphere during the summer months. No clear evidence for a correlation between extinction in all bands and the average night time wind speed is found. Also presented here is the low resolution moonless optical night sky spectrum for IAO covering the wavelength range 3000 $-$ 9300 \\AA. Features from O, OH,  N and Na are seen in the spectra. Hanle region thus has the required characteristics of a good astronomical site in terms of night sky brightness and extinction, and could be a natural candidate site for any future large aperture Indian optical-infrared telescope(s). ",
        "introduction": "A good astronomical site is characterised by various atmospheric conditions (which includes atmospheric transparency, seeing, meteorological parameters such as wind, snowfall, surface temperature, rainfall etc. and sky brightness) and geographical conditions (such as local topography, seismicity, source availability i.e, latitude etc.). Sites having minimum cloud coverage, very low frequency of snowfall/rainfall, low relative humidity, low nocturnal temperature variation, high atmospheric transparency and low night sky brightness are good for ground based optical and infrared observations. Two main characteristics of the night sky are the night sky brightness and atmospheric extinction.   Even in the absence of artificial light, the moonless night sky is not dark. This is because the atmosphere scatters into the sky, light emitted by the following processes (cf. Krisciunas 1997), (i) zodiacal light (caused by sunlight scattered off interplanetary dust), (ii) faint unresolved stars and diffuse galactic light due to atomic processes within our galaxy, (iii) diffuse extragalactic light (due to distant, faint unresolved galaxies) and (iv) airglow and aurorae (produced by photochemical reactions in the Earth's upper atmosphere). Of these, (i)$-$(iii) are extraterrestrial in origin and thus independent of the site, whereas (iv) depends on the site and time of observation. These are the  natural processes which produce the night sky brightness in any astronomical site.  In addition to the above, night sky can be affected by light pollution due to scattering of street lights in the Earth's lower atmosphere. Humans do not have control on any of the natural sources causing the brightness of the night sky, but do have a control on the brightness caused by artificial lights scattered on to the sky. It is thus possible to maintain the night sky brightness at any observatory site to its natural level by minimising light pollution in the immediate vicinity of the observatory.  Apart from the natural and artificial sources affecting the night sky, the light coming from any celestial source being observed suffers from scattering by air molecules and aerosols as well as absorption by water vapour and ozone while passing through the Earth's atmosphere. This leads to attenuation of their light and is referred to as the atmospheric extinction. This depends on the constituents of the atmosphere, the wavelength of the incoming light, and the altitude of the site. Precise knowledge of the extinction coefficient of each site is essential to compare observations of the same object taken from different locations of the globe. A good astronomical site needs low extinction values. Apart from low extinction, its stability during a night is also equally important.  In this article we present the moonless night sky brightness and atmospheric extinction in $UBVRI$ passbands at IAO, Hanle. IAO is located at the Himalayan range in Northern India (longitude = 78$^{\\circ}$57$^{\\prime}$51.2$^{\\prime\\prime}$ E, latitude = 32$^{\\circ}$46$^{\\prime}$46.5$^{\\prime\\prime}$ N and altitude = 4467 m) and run by the Indian Institute of Astrophysics, Bangalore. This is a thinly populated, cold and dry desert region. The sky at IAO is thus not much affected by dust and light pollution due to human activities.  A 2 m telescope, the Himalayan Chandra Telescope (HCT) is operational at IAO since May 2003. The data used in this study for night sky brightness span the period 2003$-$2008 which correspond to a major part of the declining phase of solar activity (which could affect night sky brightness) cycle-23, while the data for extinction estimates span the period 2000$-$2008. Preliminary estimates of the night sky brightness and extinction at IAO have been reported by Parihar et al. (2003) and an analysis of the meteorological parameters at IAO is presented in Stalin et al. (2008). The structure of this paper is as follows. Section 2 describes the data set used in this study, Section 3 presents the analysis of extinction and Section 4 presents an analysis of night sky brightness. A spectrum of the night sky at IAO is presented in Section 5 and the results are summarized in the final section. ",
        "conclusions": "We have studied the nature of the night sky at IAO. The results are summarized below: \\begin{enumerate} \\item The measured average extinction coefficient at IAO during the period 2003$-$2008 are 0.36 $\\pm$ 0.07 in $U$, 0.21 $\\pm$ 0.04 in $B$, 0.12 $\\pm$ 0.04 in $V$, 0.09 $\\pm$ 0.04 in $R$ and 0.05 $\\pm$ 0.03 in $I$. However, the average extinction during summer months is slightly larger than that of winter months. There is no clear evidence for a correlation between the measured extinction coefficient and the average night time wind speed. \\item The moonless night sky brightness at zenith are 22.14 $\\pm$ 0.32 in $U$, 22.42 $\\pm$ 0.30 in $B$, 21.28 $\\pm$ 0.20 in $V$, 20.54 $\\pm$ 0.37 in $R$ and 18.86 $\\pm$ 0.35 in $I$. Except for the $I$ band, the sky brightness in other bands at IAO are similar to those of other dark sites in the world. The bright nature of the sky in $I$ band at IAO, might be due to the presence of strong OH rotation vibrational Meinel bands in the red region of the optical spectrum.  We find no dependence of the night sky brightness with 10.7 cm solar flux, probably due to insufficient data coverage during the solar cycle. \\item The moonless night sky spectrum covering the wavelength range 3000 to 9300 has been presented. Features from OI, OH, Na and N are seen in the spectra. Lines due to light pollution are not seen in the spectrum, and thus Hanle is free from any man made light pollution. \\end{enumerate}  We conclude that IAO has a night sky similar to the best sites in the world. In addition to the good skies, IAO also has a longitudinal advantage in covering the longitudinal gap between observatories in the West and the East. Hanle region could thus be a potential site for any future large Indian optical-infrared telescope(s)."
    },
    "0809/0809.3789_arXiv.txt": {
        "abstract": "We use the observed distribution of Eddington ratios as a function of supermassive black hole (BH) mass to constrain models of quasar/AGN lifetimes and lightcurves. Given the observed (well constrained) AGN luminosity function, a particular model for AGN lightcurves $L(t)$ or, equivalently, the distribution of AGN lifetimes (time above a given luminosity $t(>L)$) translates directly and uniquely (without further assumptions) to a predicted distribution of Eddington ratios at each BH mass. Models for self-regulated BH growth, in which feedback produces a self-regulating ``decay'' or ``blowout'' phase after the AGN reaches some peak luminosity/BH mass and begins to expel gas and shut down accretion, make specific predictions for the lightcurves/lifetimes, distinct from e.g.\\ the expected distribution if AGN simply shut down by gas starvation (without feedback) and very different from the prediction of simple phenomenological ``light bulb'' scenarios. We show that the present observations of the Eddington ratio distribution, spanning nearly 5 orders of magnitude in Eddington ratio, 3 orders of magnitude in BH mass, and redshifts $z=0-1$, agree well with the predictions of self-regulated models, and rule out phenomenological ``light bulb'' or pure exponential models, as well as gas starvation models, at high significance ($\\sim5\\,\\sigma$). We also compare with observations of the distribution of Eddington ratios at a given AGN luminosity, and find similar good agreement (but show that these observations are much less constraining). We fit the functional form of the quasar lifetime distribution and provide these fits for use, and show how the Eddington ratio distributions place precise, tight limits on the AGN lifetimes at various luminosities, in agreement with model predictions.  We compare with independent estimates of episodic lifetimes and use this to constrain the shape of the typical AGN lightcurve, and provide simple analytic fits to these for use in other analyses. Given these constraints, the average local BH must have gained its mass in no more than a couple of bright, near peak-luminosity episodes, in agreement with models of accretion triggering in interactions and mergers. ",
        "introduction": "\\label{sec:intro}  Quasars and active galactic nuclei (AGN) are among the most luminous, energetic, and distant objects in the Universe. Simple integral arguments \\citep{soltan82} make it clear that the supermassive black hole (BH) population was grown primarily through accretion in luminous AGN phases, and that the accretion luminosity released in these phases dominates the X-ray background and constitutes a large fraction of the bolometric energy production of the Universe. Comparison of e.g.\\ the clustering \\citep{croom:clustering,porciani2004, hopkins:clustering} and host galaxy properties \\citep{bahcall:qso.hosts,canalizostockton01:postsb.qso.mergers, dunlop:qso.hosts,hopkins:red.galaxies,zakamska:qso.hosts, zakamska:mir.seds.type2.qso.transition.at.special.lum} of high and low redshift AGN and galaxies demonstrates that AGN are the progenitors of modern-day spheroids.  Moreover, with the discovery of tight correlations between the masses of black holes and the velocity dispersion \\citep{FM00,Gebhardt00}, masses \\citep{magorrian}, and perhaps most fundamentally binding energy or potential well depth \\citep{hopkins:bhfp.obs,hopkins:bhfp.theory,aller:mbh.esph} of the host demonstrates a fundamental link between the growth of supermassive black holes and galaxy formation. A number of models have been developed arguing that the energy or momentum released from an accreting supermassive black hole, even if only a small fraction couples to the surrounding ISM, is sufficient to halt further accretion onto the black hole and drive away gas, self-regulating growth by shutting off the quasar\\footnote{In what follows, we use the term ``quasar'' somewhat loosely, as a proxy for high-Eddington ratio accretion activity, rather than as a reference to specific optical properties. We use the term AGN to refer to BH accretion at all levels.} and quenching star formation in the galaxy and therefore allowing it to redden rapidly \\citep[see e.g.][]{ciottiostriker:cooling.flow.selfreg.1, ciottiostriker:cooling.flow.selfreg.2,silkrees:msigma,burkertsilk:msigma, dimatteo:msigma,hopkins:lifetimes.letter,hopkins:qso.all, murray:momentum.winds,sazonov:radiative.feedback, springel:models,springel:red.galaxies}.  One of the most basic aspects of black hole growth, and a powerful test of these self-regulating models for AGN evolution, is the quasar/AGN lifetime. Given a sufficiently well-known lifetime distribution (lifetime as a function of e.g.\\ luminosity, black hole mass, and other properties), the well-constrained quasar luminosity function (QLF) can be empirically translated (without invoking any models or additional assumptions) into the triggering rate of AGN as a function of e.g.\\ luminosity, BH/host galaxy/dark matter halo mass, redshift, and other properties, as well as the active BH mass function, Eddington ratio and duty cycle distributions, and differential (mass-dependent) rate of buildup of the BH mass function.  Observations generally constrain {\\em quasar} lifetimes to the range $\\approx 10^{7}-10^{8}$ yr \\citep[for a review, see][]{martini04}.  These estimates are primarily based on demographic or integral arguments which combine observations of the present-day population of supermassive black holes and accretion by the high-redshift quasar population \\citep[e.g.,][]{soltan82,haehnelt:bh.synthesis.model,yutremaine:bhmf,yulu:bhmf, haiman:bhmf,marconi:bhmf,shankar:bhmf}, or incorporate quasars into models of galaxy evolution \\citep[e.g.,][]{kh00,wyitheloeb:sam, dimatteo:cosmo.bh.growth.sim.1,dimatteo:qso.host.metals, dimatteo:cosmo.bhs,granato:sam,scannapieco:sam,lapi:qlf.sam, hopkins:qso.all,hopkins:groups.qso,sijacki:radio} or reionization of HeII \\citep{sokasian:heII.reion.epoch, sokasian:heII.reion,faucher:heII.reion.qso.constraints,faucher:ion.background.evol, mcquinn:helium.reionization.model}. Results from clustering in quasar surveys \\citep[e.g.,][]{porciani2004,grazian:local.qso.clustering, croom:clustering,myers:clustering.old,myers:clustering, lidz:clustering,porciani:clustering,shen:clustering, daangela:clustering,hopkins:clustering}, the proximity effect in the Ly$\\alpha$ forest \\citep{bajtlik:gunnpetersen.qso.age.est, haiman:gunnpetersen.qso.age.est,yulu:stromgren,faucher:proximity} \\citep[but see also][]{lidz:proximity}, and the transverse proximity effect \\citep{jakobsen:heII.ion.transverse.proximity,schirber:transverse.proximity, worseck06:indirect.transverse.proximity,worseck07:indirect.transverse.proximity, goncalves:transverse.proximity}  similarly suggest lifetimes $\\sim{\\rm a\\ few}\\,\\times 10^{7}\\,$yr.  These observations, however, pertain in particular to the {\\em quasar} lifetime -- i.e.\\ the characteristic time spent at high Eddington ratios/accretion rates, where much of the mass of a BH is accreted. Unsurprisingly, the observations suggest a lifetime similar to the \\citet{salpeter64} time (the $e$-folding time for Eddington-limited black hole growth) $t_{S}=4.2\\times10^{7}\\,(\\epsilon_{r}/0.1)\\,$yr for accretion with radiative efficiency $\\epsilon_{r}=L/\\dot{M}_{\\rm BH} c^{2}\\sim0.1$. AGN, however, are not a homogeneous population, and so the lifetime is not a single number -- in general, it should be a function of luminosity and other parameters such as BH or host mass and (possibly) various physical effects. As advocated by \\citet{hopkins:lifetimes.letter,hopkins:lifetimes.methods}, the AGN lifetime should properly be thought of as a {\\em luminosity-dependent lifetime} distribution -- i.e.\\ some time $t(>L)$ as a function of $L$ or differential $\\dtdlogL$.  \\citet{hopkins:lifetimes.letter,hopkins:lifetimes.methods, hopkins:qso.all,hopkins:faint.slope,hopkins:seyferts} study this luminosity-dependent lifetime/duty cycle distribution in hydrodynamic simulations and analytic models of feedback-regulated BH growth, and show that such self-regulation leads to a generic and unique predicted form for the lifetime distribution. After some initial trigger that fuels gas inflows \\citep[such as e.g.\\ major and minor mergers or disk instabilities; for discussion see][]{dimatteo:msigma,hopkins:seyferts, hopkins:cusps.mergers,hopkins:cusps.ell,hopkins:cores,hopkins:cusps.fp, hopkins:cusps.evol,hopkins:disk.survival,hopkins:seyfert.limits,younger:minor.mergers}, AGN grow in approximately Eddington-limited fashion until reaching some critical mass where, if some small fraction of the radiant energy/momentum can couple to the surrounding ISM, the feedback is sufficient to halt inflows and/or expel gas and shut down future accretion. This ``upper limit'' to growth is essentially an Eddington limit effective at the scales where the host galaxy, rather than the BH, dominates the gravitational potential, and therefore is set not by the details of fueling mechanisms, but instead by (relatively generic) global parameters of the accretion physics such as the host galaxy mass and AGN luminosity.  That the resulting lifetime distribution is independent of fueling mechanism (i.e.\\ is not specifically related to, for example, merger-induced fueling, but is a generic consequence of models where BH growth is self-regulated by feedback) has been demonstrated in the nearly identical lifetime distributions obtained by models of fueling in major mergers, minor mergers, flyby events, bar instabilities, and random ``stochastic'' encounters with nuclear molecular clouds, under similar feedback-regulated conditions \\citep{hopkins:qso.all,hopkins:faint.slope,hopkins:seyferts, johansson:mixed.morph.mbh.sims, younger:minor.mergers}. These various fueling mechanisms do result in other important differences (in e.g.\\ host properties and spectral properties of observed AGN), and they will lead to different evolution of BHs in a cosmological sense (global triggering rates and their evolution varying significantly for the mechanisms above) and produce BHs of different masses (to the extent that they produce bulges across a large range of masses, their self-regulating nature and the existence of the BH-host correlations ensures this to be the case). These are discussed in more detail in other papers \\citep[see e.g.][]{hopkins:seyfert.limits}, but the important point is that, for a given triggering rate at some redshift and BH mass interval (set by some cosmological or galactic processes), self-regulated models predict a similar effective Eddington limit and lightcurve.  Just as growth at the traditional Eddington limit leads to a self-similar solution for the AGN lightcurve -- exponential growth -- the expulsion of gas in this analogous limit leads to a self-similar lightcurve once the gas begins to be removed from the vicinity of the BH -- a power-law decay of the form $L\\propto t^{-(1.5-2.0)}$. In turn, this translates into a lifetime distribution $\\dtdlogL$ with a characteristic faint-end (low-$L$) power-law like behavior: i.e.\\ the time spent above a given luminosity/Eddington ratio (at fixed final BH mass) scales $\\propto L^{-\\beta}$ with a $\\beta\\sim0.6$ at low-$L$, with a cutoff near the Eddington limit/peak luminosity of the system. The normalization is set by the characteristic timescales of the system -- the Salpeter time and the (very similar) characteristic dynamical times in the central regions of the galaxy -- naturally yielding a robust prediction of the observed quasar lifetime as well as the complete AGN lifetime distribution.  This is significantly different from what is assumed in commonly adopted phenomenological models, as well as other physical prescriptions. Quasars are often treated crudely as ``light-bulbs'' -- i.e.\\ assumed to be ``on'' for a fixed time (the ``quasar lifetime'') at fixed luminosity or Eddington ratio, and otherwise ``off.'' In such a case, the lifetime/Eddington ratio distribution does not increase towards lower luminosities, but instead is strongly peaked about a characteristic high Eddington ratio/luminosity (strictly speaking, a $\\delta$-function; or in a Schechter-function parameterization with some scatter, $\\beta\\ll 0$). Similar results are obtained if one assumes that quasar lightcurves are pure exponentials (corresponding to growth at fixed Eddington ratio with either an instantaneous cutoff or time-mirrored exponential luminosity decay), which yields equal time spent per logarithmic interval in luminosity (i.e.\\ a luminosity-independent lifetime, or $\\beta=0$ Schechter function).  More physically motivated but distinct models make their own predictions for the lifetime distribution.  For example, if one assumes that the quasar accretion is regulated by a standard \\citet{shakurasunyaev73} thin disk, and the fuel supply is immediately removed but there is no feedback (i.e.\\ a gas starvation scenario), one obtains a similarity solution for the accretion rate versus time that yields a lifetime distribution more akin to, but still significantly distinct from that predicted in self-regulated models \\citep[see e.g.][]{yu:mdot.dist}.  A number of indirect tests have been proposed to break the degeneracies between these models and constrain the quasar lifetime distribution, and the observations in these scenarios have thus far supported the predictions of self-regulated models. These include: the dependence of quasar clustering on luminosity \\citep{lidz:clustering,hopkins:clustering}, where observations finding e.g.\\ a weak dependence of clustering amplitude on luminosity (at fixed redshift) support lifetime models with more time spent at lower $L$ \\citep{adelbergersteidel:lifetimes,coil:agn.clustering, myers:clustering,daangela:clustering}; the shape of the active BH mass function in various luminosity-selected samples \\citep{hopkins:lifetimes.interp, hopkins:qso.all,hopkins:groups.qso}, including more massive systems at lower $L$ (in extreme cases even being peaked) rather than tracing an identical shape to the QLF \\citep{heckman:local.mbh,greene:active.mf}; the evolution in the faint-end QLF slope with redshift \\citep{hopkins:faint.slope}, flattening (weakly) with redshift \\citep{ueda03:qlf,hasinger05:qlf,lafranca:hx.qlf,silverman:hx.spacedensity.ldde, hopkins:bol.qlf,fontanot:highz.qlf,silverman:hx.lf,siana:z3.swire.qlf} as predicted in some self-regulated models owing to a (weak) dependence of quasar lifetime distributions on BH mass/peak luminosity; the shape of the distribution in quasar host galaxy masses/luminosities \\citep{hamilton:qso.host.lf,hopkins:merger.lfs}, similar in nature to the active BH mass function as a test of lifetime models; the mass functions and clustering as a function of mass of quasar ``remnants'' \\citep[i.e.\\ bulges/spheroids, given the observed $M_{\\rm BH}-\\sigma$ relation; see][]{hopkins:red.galaxies, hopkins:groups.ell,hopkins:transition.mass,bundy:mtrans,haiman:qlf.from.ell.ages, shankar:implied.msig.from.gal.ages,yulu:lightcurve.constraints.from.bhmf.integration}; and the relation between observed luminosity functions in different bands and active BH masses \\citep{hopkins:lifetimes.obscuration,shankar:bol.qlf, merloni:synthesis.model,bundy:agn.lf.to.mf.evol,yulu:lightcurve.constraints.from.bhmf.integration}.  Although these observations are consistent with the predictions from self-regulated models, they are indirect, and as such are not able to rule out alternative interpretations (in the case of e.g.\\ clustering or the evolution of the faint-end QLF slope) or depend on additional (albeit observationally and physically-motivated) assumptions. Moreover, many are restricted to relatively bright Seyfert/quasar populations, where the models are all similar (near these luminosities, they all predict that the entire population must be dominated by relatively massive BHs at high Eddington ratios $\\sim 0.1-1$, and given the rapid growth at these Eddington ratios, the lifetime in this regime must be similar, comparable to the \\citet{salpeter64} time). Biases introduced (selecting for specific Eddington ratios, for example) will also be important in samples selected by e.g.\\ broad emission lines or optical/UV/IR colors/spectral shape \\citep[see][]{hopkins:seyfert.bimodality}.  The observed Eddington ratio distributions -- specifically, the complete distribution of Eddington ratios for {\\em all} BHs of a given mass (obscured or unobscured, luminous or under-luminous), however, represent a direct and powerful test of these models, with the predicted behavior at lower Eddington ratios/luminosities being strongly divergent. Given the QLF, a lifetime model directly and uniquely translates (without any additional assumptions) into a distribution of Eddington ratios at each BH mass, and vice versa. The requirements are demanding: breaking these degeneracies necessitates large samples, complete to all objects of a given BH mass, in a large volume (to constrain rare high-Eddington ratio objects), but sufficiently deep to measure even faint levels of AGN activity in those objects. Furthermore, a probe of AGN activity such as X-ray or narrow-line emission, robust to obscuration effects (or the possible disappearance of the broad-emission line region and/or thin disk at low Eddington ratios) is important both to obtain a complete census of AGN activity and to avoid biases from e.g.\\ luminosity or Eddington ratio-dependent obscuration/dilution/SED shapes. Fortunately, with the advent of wide area spectroscopic surveys such as the SDSS, it has become possible to constrain the Eddington ratio distribution at low redshifts, over a sufficiently large dynamic range to break the degeneracies between these models, as a function of various AGN and galaxy properties.  Here, we combine a large number of observations of AGN Eddington ratio distributions as a function of BH and host galaxy mass, AGN luminosity, and redshift, in order to test these models and tightly constrain critical quantities such as the quasar lifetime as a function of luminosity, and show how the present observations are already sufficient to rule out, at high significance, a wide variety of alternative, simplified physical and phenomenological AGN lifetime/lightcurve models.  In \\S~\\ref{sec:compare.dist} we compare these model predictions with observations of the Eddington ratio distribution measured directly at $z=0$ over a range of BH mass, and inferred indirectly at $z=0-1$. In \\S~\\ref{sec:compare.L} we similarly compare with Eddington ratio distributions measured not at fixed BH mass, but at fixed AGN luminosity, again over the range $z=0-1$. In \\S~\\ref{sec:compare.models} we show how these observations tightly constrain physical and phenomenological models for the quasar lifetime/accretion rate distribution, relate to a possible dependence of quasar feedback efficiency on mass, and rule out a number of alternative lifetime/lightcurve models. We show how the observations tightly constrain even general, parameterized lifetime models to a narrow range about the physical models, and can directly be converted to yield the cosmologically integrated AGN lifetime and $z=0$ duty cycles as a function of Eddington ratio. In \\S~\\ref{sec:lightcurves} we translate these Eddington ratio/lifetime distribution constraints to limits on the form of the ``typical'' AGN lightcurve, and discuss constraints on the ``episodic'' quasar lifetime and how, combined with the duty cycle constraints, this can give a bound on the number of accretion episodes per AGN and shape of the typical lightcurve. In \\S~\\ref{sec:implications} we demonstrate the constraints from these observations on how much mass present-day BHs accreted in various intervals in Eddington ratio and luminosity. Finally, in \\S~\\ref{sec:discussion} we discuss our results, the implications of the observations for a broad range of AGN properties, and the prospects for future observational tests.  For ease of comparison, we convert all observations to bolometric luminosities given the appropriate bolometric corrections from \\citet{hopkins:bol.qlf} \\citep[see also][]{elvis:atlas,richards:seds}. We adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology and normalize all observations and models appropriately \\citep[note that this generally affects only the exact normalization of quantities here, not the qualitative conclusions, and differences are negligible within the range of cosmologies allowed by present constraints; e.g.][]{komatsu:wmap5}. ",
        "conclusions": "\\label{sec:discussion}   \\subsection{Overview} \\label{sec:discussion:overview}  We have compared observations of the Eddington ratio distribution as a function of BH mass, redshift, and luminosity with various theories. We find good agreement between the observations and the predictions of the self-regulated feedback models for BH growth and evolution from \\citet{hopkins:lifetimes.letter,hopkins:lifetimes.methods, hopkins:lifetimes.interp,hopkins:qso.all,hopkins:faint.slope,hopkins:seyferts}. The agreement covers the entire observed dynamic range, with observations extending down $5-6$ orders of magnitude in Eddington ratio ($L/L_{\\rm Edd} \\sim 10^{-6}-1$), over three orders of magnitude in BH mass ($M_{\\rm BH}\\sim10^{6}-10^{9}\\,\\msun$), corresponding to $\\sim8$ orders of magnitude in luminosity, and (given indirect observational constraints) from redshifts $z\\sim0-1$. The models are not fitted to these observations -- there are no free parameters adjustable, and the range of uncertainty between different versions of the models constrained to reproduce the observed quasar luminosity function is small (within the observational error bars).  The lifetime -- already as constrained in a purely empirical sense from the observations, is clearly {\\em not} a delta function -- i.e.\\ quasars are not ``light bulbs'' -- rather, it is a smooth, continuous, and relatively steep function of luminosity/Eddington ratio, with more time spent at lower luminosities/Eddington ratios.  We also show agreement with observed Eddington ratio distributions as a function of AGN luminosity, but these are less constraining, because of the inherent Eddington ratio limits implied by the selection \\citep[wavelength-dependent effects further complicate comparisons; discussed in][]{hopkins:seyfert.bimodality}.    \\subsection{Implied AGN Lifetimes and Growth Histories} \\label{sec:discussion:constraints}  The observationally implied lifetime distribution can be generally parameterized (in model-independent fashion) as a Schechter function, with a characteristic normalization lifetime, a turnover at Eddington ratios near $\\sim1$ (reflecting a physical limit around the Eddington limit), and a faint-end slope $\\beta$ (such that the luminosity-dependent lifetime $t(>\\mdot)\\propto \\mdot^{-\\beta}$ at small $\\mdot$). The observations favor a best-fit $\\beta\\approx0.6\\pm0.05$ (in agreement with feedback-regulated models, discussed below) for typical $\\sim L_{\\ast}$ galaxies and BHs. Combined with constraints from the quasar luminosity function, there is evidence for weak BH mass-dependence of $\\beta$, as predicted by hydrodynamic simulations \\citep{hopkins:faint.slope}.  The observations directly yield quasar duty cycles and lifetimes as a function of Eddington ratio and BH mass. In terms of the integrated time at a given Eddington ratio interval (around some final BH mass), the ``quasar'' lifetime -- i.e.\\ lifetime at high Eddington ratios $\\gtrsim0.1$, is similar to that expected from the \\citet{salpeter64} time and other observational constraints, $\\sim10^{8}\\,$yr. However, the time at lower Eddington ratios rises rapidly -- with $t(\\mdot\\gtrsim0.01)\\sim 0.5-1\\,$\\,Gyr and $t(\\mdot \\gtrsim 0.001)\\sim1-5\\,$Gyr.  This allows us to quantify the fractional contribution to present-day BH mass from various intervals in $\\mdot$; growth is dominated by large $\\mdot\\sim0.2$, with a small -- but non-negligible -- contribution $\\sim20\\%$ from accretion at low Eddington ratios $\\mdot\\lesssim0.01$, in agreement with integral arguments \\citep{soltan82} and various independent constraints \\citep[][and references therein]{hopkins:old.age}.  Comparison of these integral constraints with the $z=0$ observations (comparing ``integrated'' and ``effective'' AGN lifetimes/duty cycles) reflects the increasingly established AGN downsizing trend: the low-mass BH population is still growing today, the high-mass population has shut down since its earlier epoch of peak activity.   \\subsection{Constraints on Models of Lightcurves and Lifetimes} \\label{sec:discussion:models}  There is enormous constraining power in the Eddington ratio distribution at low $\\mdot$ in BH mass-limited samples. Already, the observations at $z=0$ are sufficient to limit the Eddington ratio/quasar lifetime distribution to a narrow range around the theoretical predictions from recent hydrodynamic simulations incorporating a self-consistent model for accretion and feedback. Those models predict that the self-regulated nature of BH accretion should lead to a relatively self-similar decay phase in AGN luminosity or Eddington ratio (regardless of e.g.\\ triggering mechanisms, or the exact details of feedback physics), $L\\propto t^{-(1.5-2.0)}$, which gives rise to the observed power-law-like faint behavior in the lifetime distribution with predicted slope $\\beta\\approx0.6$.  The observations strongly rule out (at $>5\\,\\sigma$ significance) simplified models for quasar lightcurves, including: ``light-bulb'' models in which quasars turn ``on'' for some time at a fixed or relatively narrow range of accretion rates, or ``exponential'' models in which AGN grow at fixed Eddington ratio and then ``shut off'' rapidly.  Models where BH accretion simply traces stellar mass loss are also ruled out ($\\gtrsim4\\,\\sigma$). Stellar mass loss may still be a fuel source; however, the observations argue that some process must further regulate AGN activity, shutting down accretion more efficiently than the simple slow starvation expected if BH growth directly traced stellar mass loss.  At $\\sim3-5\\,\\sigma$ significance, the observations also rule out isolated accretion disk models \\citep[][and references therein]{yu:mdot.dist} - i.e.\\ accretion disks fueled rapidly but then cut off from a future gas supply, without feedback (such that they evolve by gas exhaustion). These are a considerable improvement on the light bulb model, but are still ruled out formally -- with the sense again that some process must shut down growth more efficiently (possibly the addition of even mild outflows to the isolated disk solution) -- especially for high-mass BH populations.  Local observations of the Eddington ratio distribution alone do not strongly constrain the mass or redshift dependence of these properties. Combining the observations here with constraints from the quasar luminosity function, observations do favor a weak mass-dependence in the shape of the lifetime distribution, with the sense that $\\beta$ decreases with increasing BH mass ($\\beta\\sim0.6 - 0.2\\{\\log{[M_{\\rm BH}/10^{7}\\,\\msun]}\\}$). This is equivalent to the statement that more massive systems ``shut off'' more abruptly than low-mass systems; a trend predicted by hydrodynamic models, as quasar feedback becomes relatively more dominant (and bulge-to-disk ratios increase while global disk gas fractions decrease) in more massive systems.  Recently, \\citet{kauffmann:new.mdot.dist} have shown that the $\\lambda$ distribution depends on the stellar population age of the system. We show that the observed trends naturally arise from a {\\em single} lightcurve of the form constrained here, as a consequence of different triggering histories. Systems which are still active/growing (i.e.\\ have not ``quenched'' their gas supply, BH growth, or star formation) remain at higher accretion rates reflecting the median lightcurve directly, while systems which quenched at some earlier epoch have decayed down to the lower $\\lambda$ power-law-like tail of the lightcurve distribution. The existence of these trends, and observational constraints on the Eddington ratio distribution at $z=1$ \\citep{merloni:mdot.dist.fits.prep}, set some constraints on the evolution of typical lightcurves with redshift.  We find that $\\beta$ and the episodic quasar lifetime $t_{Q}$ cannot evolve strongly (parameterized as $\\beta\\propto(1+z)^{\\beta^{\\prime}_{z}}$ and $t_{Q}\\propto(1+z)^{t^{\\prime}_{z}}$; the allowed range is $-0.5\\le \\beta^{\\prime}_{z} \\le 0.05$, $-0.7\\le t^{\\prime}_{z} \\le 0.25$). Evolution towards a stellar wind or isolated accretion disk solution is strongly ruled out -- rather, the sense of (mild) evolution allowed is towards more sharply peaked lightcurves at higher redshift, which may be expected if fueling is more violent and/or feedback is more efficient in this regime. The evolution of duty cycles (and correspondingly ``effective'' AGN lifetimes) at the high-luminosity end ($\\lambda\\gtrsim0.1$) is directly constrained by evolution of the quasar luminosity function -- this limits the combination of $t_{0}$ and $\\eta$ (but not the faint-end slope $\\beta$), and (given some constraint on the evolution of episodic lifetimes $t_{Q}$) the triggering rate of quasars to evolve according to Equations~\\ref{eqn:z.evol}-\\ref{eqn:z.evol.alpha}, increasing with redshift more rapidly in higher-mass BHs.  These constraints on the AGN can be transformed into constraints on the average AGN lightcurve, suggesting a characteristic power-law like decay ($L\\propto t^{-(1.5-2.0)}$) similar to that predicted in models, after rapid growth to some peak luminosity. The observations of the Eddington ratio distribution tightly limit the shape of such lightcurves, but only weakly constrain the characteristic timescale for a single such ``event'' (i.e.\\ the width in time of a single ``peak'') -- the episodic AGN lifetime -- by setting an upper limit.  Observational constraints from e.g.\\ the proximity effect in bright quasars can independently set lower limits on the episodic lifetime, however, and suggest that (at least over the observed range for the proximity effect, namely quasar activity with $\\mdot\\sim0.1-1$) the episodic lifetime is comparable to the integrated lifetime, in the range $t_{Q}\\sim0.3-1.0\\,t_{\\rm int}$. This implies that {\\em the typical massive BH has experienced no more than $\\sim$ a couple of bright, high Eddington ratio quasar episodes while near its current ($z=0$) mass}. This is in agreement with predictions from cosmological models that associate the brightest quasar activity with fueling in violent major mergers \\citep{hopkins:groups.qso}, for which there are only a couple of events expected since high redshift. Improved constraints on the episodic lifetime will allow the observed quasar luminosity function to be more robustly translated into the distribution of AGN triggering rates as a function of BH and host mass. Extending the observational constraints on the episodic lifetime to lower luminosities/Eddington ratios is also important, as one might expect that although a typical object only experiences a couple of triggers to near-peak activity, it could have many more triggers for lower-level activity.   \\subsection{Other Tests and Future Work} \\label{sec:discussion:future}  Our findings agree with other (less direct) independent constraints. Recently, for example, \\citet{yulu:lightcurve.constraints.from.bhmf.integration} showed that the joint evolution of AGN luminosity functions with redshift favors similar lightcurves. These results (from the QLF and/or BH mass function and integral/continuity arguments) provide independent support for the constraints here, but are primarily sensitive to luminous (high-$\\lambda$) behavior and cannot distinguish between similar lightcurves (with slightly different $\\beta$), such as the isolated accretion disk or feedback-regulated predictions. The lightcurve shape can also be probed by the dependence of AGN clustering on luminosity and shape of the QLF or ``active'' BH mass function in deep samples; the observations at present appear to favor feedback-regulated models over simplified light-bulb or exponential lightcurve models \\citep[see][]{adelbergersteidel:lifetimes, myers:clustering,greene:active.mf,daangela:clustering}, but again deeper observations are needed to distinguish between the proposed physically motivated models.  Extending the observations of the Eddington ratio distribution to higher redshift will greatly improve these constraints as well as limit possible redshift evolution in quasar light curves. More massive black holes will be closer to their peak growth at higher redshifts, allowing observations to probe the times of greatest interest. At $z\\sim1-3$, the $\\mdot\\gtrsim0.1$ end of the distribution, in broad-line luminous quasars, is already constrained by the observed QLF and application of the virial BH mass estimators \\citep{kollmeier:mdot,fine:broadline.distrib}. Combined with these constraints, smaller volume but deep redshift surveys can be used to construct samples which are complete to a given BH/bulge mass, and similar narrow-line searches in these hosts could limit the Eddington ratio distribution at much lower ratios. If coverage is sufficient, X-ray data can be used in the same manner. At $z\\ge2$, complete, large-volume spectroscopic or X-ray samples are not available at present. Here, indirect tests will remain important for the near future, but improved constraints on what (if any) evolution is seen at lower redshifts in the shape of the lifetime distributions will considerably inform future models and observational efforts."
    },
    "0809/0809.3740_arXiv.txt": {
        "abstract": "{% We present an extinction map of a $\\sim 1\\,700 \\mbox{ deg sq}$ region that encloses the Ophiuchus, the Lupus, and the Pipe dark complexes using 42 million stars from the Two Micron All Sky Survey (2MASS) point source catalog.  The use of a robust and optimal near-infrared method to map dust column density (\\textsc{Nicer}, described in \\citealp{2001A&A...377.1023L}) allow us to detect extinction as low as $A_K = 0.05 \\mbox{ mag}$ with a 2-$\\sigma$ significance, and still to have a resolution of $3 \\mbox{ arcmin}$ on our map.  We also present a novel, statistically sound method to characterize the small-scale inhomogeneities in molecular clouds. Finally, we investigate the cloud structure function, and show that significant deviations from the results predicted by turbulent models are observed. ",
        "introduction": "\\label{sec:introduction}     In this paper we present an extinction map of the Ophiuchus and Lupus complexes covering $\\sim 1\\,700 \\mbox{ deg sq}$, computed by applying an optimized multi-band technique dubbed Near-Infrared Color Excess Revisited (\\textsc{Nicer} \\citealp{2001A&A...377.1023L}, hereafter Paper~I) to 42 million JHK photometric measurements of stars from the Two Micron All Sky Survey (2MASS; \\citealp{1994ExA.....3...65K}). This paper is the second of a series where we apply the \\textsc{Nicer} method to point sources from the 2MASS database.  The main aim of this coordinated study is to investigate in detail the large-scale structure of molecular clouds.  In addition, the use of a uniform dataset and of a consistent and well tested pipeline allows us to characterize many properties of molecular clouds, such as their reddening law, and to identify cloud-to-cloud variations in such properties.  The region considered here encloses in addition the Pipe nebula, considered in the first paper of this series (\\citealp{2006A&A...454..781L}; hereafter Paper~II), and allows us to study in detail the relationships among these molecular clouds.  Lupus is a well studied complex, composed of several subclouds showing different modes of star formation: Lupus 1 shows isolated star formation, while cluster formation is observed in Lupus 3 \\citep{1977ApJS...35..161S, 1994AJ....108.1071H, 2000AJ....119..873N}; no evident star formation activity is observed in Lupus 5.  The distance of this cloud complex is still highly controversial.  The conventional value, $150 \\mbox{ pc}$ \\citep{1991ESOSR..11....1R}, has been challenged by \\citet{1998A&A...338..897K}, who reported $100 \\mbox{ pc}$, and by \\citet{1998MNRAS.301L..39W}, $190 \\mbox{ pc}$, both using Hipparcos observations.  In addition, \\citet{2001A&A...373..714K} suggested that Lupus~2 is physically disconnected from the other Lupus subclouds, and is located at a distance of $360 \\mbox{ pc}$; similarly, \\citet{1999A&A...352..574B} suggested Lupus~3 to be disconnected from the rest of the Lupus complex.  However, surprisingly no velocity difference is found among the various Lupus subclouds in radio observations.  Recently, we used a novel maximum-likelihood technique based on Hipparcos data obtaining $d = (155 \\pm 8) \\mbox{ pc}$ \\citep{2008A&A...480..785L}; this value will be adopted in this paper as the actual distance of Lupus.  Due to its large area on the sky, millimeter observations have mostly covered limited regions of this cloud.  The first extensive cloud survey in ${}^{12}$CO ($J=1 \\rightarrow 0$) was made by \\citet{1986A&A...167..234M} over $\\sim 170 \\mbox{ sq deg}$ with an effective resolution of $30 \\mbox{ arcmin}$.  More recently, \\citet{2001PASJ...53.1081T} have performed a complete survey of the Lupus complex (also in ${}^{12}$CO, $J=1 \\rightarrow 0$) using the NANTEN sub-millimeter telescope.  Their data covered $~ 550 \\mbox{ sq deg}$ on a grid of $8 \\mbox{ arcmin}$, and smaller areas on a $4 \\mbox{ arcmin}$ grid.  Finally, recently \\citet{2005ApJ...629..276T} studied the structure of Lupus~3 in the NIR.  The $\\rho$ Ophiuchi molecular cloud is a filamentary complex located at the edge of the Upper Scorpius subgroup in the Sco-Cen OB association.  It is composed of a series of dark clouds that extend eastward from cores of dense molecular gas \\citep{1992A&A...262..258D}.  With an estimated distance of $(119 \\pm 6) \\mbox{ pc}$ \\citep{2008A&A...480..785L}, it is one of the nearest star-forming regions, and has thus been the target of numerous investigations.  The main cloud, L1688, is situated about one degree south of the star $\\rho$ Ophiuchi and has been studied in the near and far infrared, in the millimeter continuum, as well as in the X-ray and radio continuum, and hosts an embedded infrared cluster with $\\sim 200$ young stellar objects (YSO; e.g., \\citealp{1992ApJ...395..516G, 2001A&A...372..173B, 1993ApJ...416..185C, 2000PASJ...52.1147T}). L1688 is known to host a $1 \\mbox{ pc} \\times 2 \\mbox{ pc}$ central molecular core with visual estimated visual extinction exceeding $50 \\mbox { mag}$ \\citep{1983ApJ...274..698W}.  This ``main cloud'' of Oph covers an area of roughly $480 \\mbox{ sq arcmin}$ and has been dissected into about a dozen cloud components.  The entire star formation complex extends over several degrees on the sky, containing a few major clouds.  \\citet{1998A&A...336..150M} have conducted an extensive $1.3 \\mbox{ mm}$ continuum mapping of the central region using the IRAM 30-m telescope.  \\citet{2006A&A...447..609S} presents SIMBA $1.2 \\mbox{ mm}$ observations of the cloud covering an area of $4600 \\mbox{ arcmin}^2$ with a resolution of $24''$; \\citet{2000ApJ...545..327J} have published a SCUBA $850 \\ \\mu\\mbox{m}$ thermal emission map of the $\\rho$ Ophiuchi cloud.  This paper is organized as follows.  In Sect.~\\ref{sec:nicer-absorpt-map} we briefly describe the technique used to map the dust and we present the main results obtained.  A statistical analysis of our results and a discussion of the bias introduced by foreground stars and unresolved substructures is presented in Sect.~\\ref{sec:statistical-analysis}. Section~\\ref{sec:mass-estimate} is devoted to the mass estimate of the cloud complex.  Finally, we summarize the results obtained in this paper in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The main results of this paper can be summarized as follows: \\begin{itemize} \\item We used approximately $42$ million stars from the 2MASS point source catalog to construct a $1\\,672$ square degrees \\textsc{Nicer} extinction map of the Ophiuchus and Lupus dark nebul\\ae.  The map has a resolution of $3\\mbox{ arcmin}$ and an average $2\\sigma$ detection level of $0.5$ visual magnitudes. \\item We considered in detail the effect of sub-pixel inhomogeneities, and derived an estimator useful to quantify them.  We also showed that inhomogeneities play a significant role only in the densest cores with $A_K > 6$--$8 \\mbox{ mag}$. \\item We derived the structure functions of both dark clouds and compared them with several theoretical models.  We could not find any reasonable fit of the Ophiuchus data with models, while the Lupus scaling index ratio is well described in terms of the \\citet{1994PhRvL..72..336S} turbulent model. \\end{itemize}"
    },
    "0809/0809.1868_arXiv.txt": {
        "abstract": "We perform Halo Occupation Distribution (HOD) modeling to interpret small-scale and intermediate-scale clustering of 35,000 luminous early-type galaxies and their cross-correlation with a reference imaging sample of normal $L_*$ galaxies in the Sloan Digital Sky Survey. The modeling results show that most of these luminous red galaxies (LRGs) are central galaxies residing in massive halos of typical mass $M \\sim$ a few times $10^{13}$ to $10^{14} \\hMsun$, while a few percent of them have to be satellites within halos in order to produce the strong auto-correlations exhibited on smaller scales. The mean luminosity $L_c$ of central LRGs increases with the host halo mass, with a rough scaling relation of $L_c \\propto M^{0.5}$. The halo mass required to host on average one satellite LRG above a luminosity threshold is found to be about 10 times higher than that required to host a central LRG above the same threshold. We find that in massive halos the distribution of $L_*$ galaxies roughly follows that of the dark matter and their mean occupation number scales with halo mass as $M^{1.5}$. The HOD modeling results also allows for an intuitive understanding of the scale-dependent luminosity dependence of the cross-correlation between LRGs and $L_*$ galaxies. Constraints on the LRG HOD provide tests to models of formation and evolution of massive galaxies, and they are also useful for cosmological parameter investigations. In one of the appendices, we provide LRG HOD parameters with dependence on cosmology inferred from modeling the two-point auto-correlation functions of LRGs. ",
        "introduction": "The clustering of galaxies depends on their properties, such as morphology (e.g., \\citealt{hubble36,zwicky68,davis76,dressler80,postman84,guzzo97,willmer98, zehavi02,goto03}), luminosity (e.g., \\citealt{davis88,hamilton88,white88,park94,loveday95,guzzo97, benoist96,norberg01,zehavi02,Zehavi05b,Coil06,Coil08}), color (e.g., \\citealt{willmer98,brown00,zehavi02,Zehavi05b,Coil08}), and spectral type (e.g., \\citealt{norberg02,budavari03,madgwick03}). Galaxy clustering thus provides important clues to the physics of galaxy formation. Often found to reside in galaxy groups and clusters, luminous red galaxies (LRGs) constitute the bright end of the galaxy luminosity function. Clustering of LRGs encodes information about their environments, which are typically the central regions of groups and clusters. The LRG redshift sample \\citep{Eisenstein01} of the Sloan Digital Sky Survey (SDSS; \\citealt{York00}) provides an enormous data set with which to measure the clustering of LRGs. In this paper, we perform theoretical modeling of auto-clustering of LRGs and cross-clustering of LRGs with other types of galaxies to understand the origin of their clustering properties, to learn how they are distributed among massive dark matter halos, and to aid the study of formation and evolution of massive galaxies.  The theoretical understanding of galaxy clustering has been greatly enhanced through the framework of the halo occupation distribution (HOD, see, e.g., \\citealt{Jing98a,Ma00,Peacock00,Seljak00,Scoccimarro01, Berlind02}) and the closely related approach of the conditional luminosity function (CLF, \\citealt{Yang03}). The HOD formalism describes the bias relation between galaxies and matter at the level of individual virialized dark matter halos, whose distribution and properties can be readily predicted by numerical simulations or analytic models given a cosmological model. The key ingredients of this formalism are the probability distribution $P(N|M)$ that a halo of mass $M$ contains $N$ galaxies of a given type and spatial and velocity distributions of galaxies within halos. In the CLF approach, the dependence of $P(N|M)$ on galaxy luminosity is implicitly derived by inferring the conditional luminosity distribution of galaxies as a function of halo mass $M$. Given the HOD/CLF and the cosmology, any statistics of galaxy clustering on any scales can be predicted, therefore the HOD/CLF method provides a complete description of the bias relation between galaxies and dark matter. HOD/CLF modeling has been applied to interpret galaxy clustering measurements in several galaxy surveys (e.g., \\citealt{Jing98b,Jing02, Bullock02,Moustakas02,Bosch03,Magliocchetti03,Yan03,Zheng04,Yang05,Zehavi05b, Lee06,Hamana06,Cooray06,Conroy06,White07,Zheng07,Blake08,Wake08}). HOD/CLF analysis recasts galaxy clustering measurements into a form that is more physically informative and conducive for testing galaxy formation theories (see, e.g., \\citealt{Berlind02,Berlind03,Bosch03,Zheng05}). \\citet{Zehavi05b} present clustering measurements and HOD modeling for galaxies in the SDSS main galaxy spectroscopic sample \\citep{Strauss02}. The HOD modeling results of the luminosity and color dependence of galaxy clustering are found to be in good qualitative agreement with predictions from galaxy formation models (\\citealt{Zheng05}).  In this paper, we apply HOD modeling to interpret clustering in the SDSS LRG spectroscopic sample, which uses color and magnitude cuts to effectively select LRGs out to redshift $z\\sim 0.45$ \\citep{Eisenstein01}. The large volume probed by the LRG sample has led to the detection of the baryon acoustic peak in the two-point correlation function (\\citealt{Eisenstein05b}). \\citet{Zehavi05a} report the measurements of two-point correlation functions of 35,000 LRGs on scales of 0.3--40$\\hMpc$. They find that LRGs are highly clustered (correlation length $\\sim$10$\\hMpc$) and that more luminous LRGs are more clustered. Clear deviations from a power law are seen in the correlation functions, with a dip at $\\sim$2$\\hMpc$. \\citet{Eisenstein05a} measure the two-point cross-correlation between 32,000 spectroscopic LRGs and 16 million galaxies in the SDSS imaging sample. Since these reference galaxies have luminosities around $L_*$, the characteristic luminosity of the \\citet{Schechter76} luminosity function, they are denoted as $L_*$ galaxies. \\citet{Eisenstein05a} find a strong luminosity dependence of the LRG--$L_*$ cross-clustering amplitude. The form of the luminosity dependence is itself dependent on scale, with more variation in the clustering amplitude on small scales. Understanding all these auto- and cross-clustering features is one of the goals of our HOD modeling. Since LRGs trace massive halos, the cross-correlation between LRGs and $L_*$ galaxies also allows us to study the HOD of $L_*$ galaxies in these massive halos.  The structure of this paper is as follows. In \\S~2, we briefly describe the SDSS LRG and $L_*$ galaxy samples we use and our modeling method. In \\S~3, we describe our HOD parameterization for LRG samples and $L_*$ samples. In \\S~4, we perform HOD modeling for LRG two-point auto-correlation functions and the two-point cross-correlation functions between LRGs and $L_*$ galaxies. We show how LRGs occupy dark matter halos and how $L_*$ galaxies occupy massive halos. Based on the modeling results, we interpret the luminosity dependence of the cross-clustering. Finally, we summarize and discuss our results in \\S~5. We also include three appendices in the paper. In Appendix~\\ref{sec:appendixA}, we investigate which HOD parameter plays the major role in the departures of the galaxy two-point correlation function from a pure power law. In Appendix~\\ref{sec:appendixB}, we present HOD modeling results for two luminosity threshold LRG samples for different cosmological models. In Appendix~\\ref{sec:appendixC}, we have a brief discussion on the mass function of the most massive halos and the fluctuation of the number of massive halos in the volume probed by the SDSS LRG samples. ",
        "conclusions": "\\label{sec:discuss}  We have modeled the two-point auto-correlation functions of LRGs and the two-point cross-correlation functions between LRGs and $L_*$ galaxies in the SDSS, within the HOD framework, obtaining results on the mean relation between central LRG luminosity and halo mass, the dispersion about this relation, the slope and amplitude of the satellite occupation function, and the abundance of $L_*$ galaxies in massive halos.  The continuous rise toward small scales of two-point auto-correlation functions of LRGs implies that not all LRGs are the bright central galaxies in galaxy groups or clusters; a fraction of LRGs must be satellites to produce small scale, one-halo pairs. However, the satellite fraction is small and decreases with the LRG luminosity, e.g., $\\sim$5--6\\% for $M_g<-21.2$ and $\\sim$2--3\\% for $M_g<-21.8$ based on the HOD modeling. The characteristic minimum host halo mass of central LRGs ($\\Mmin$, at which 50\\% of halos host a galaxy above the luminosity threshold) is a few times $10^{13}\\hMsun$ and increases with LRG luminosity.  \\citeauthor{Zehavi05b} (\\citeyear{Zehavi05b}; see also \\citealt{Zheng07}) found a ratio $M_1/\\Mmin\\sim 20$ between the halo mass required to host a satellite above a luminosity threshold and the mass required to host a central galaxy above the same threshold. For these LRG samples, which populate higher mass halos and have a median redshift $z\\sim 0.3$, we find a smaller ratio, $M_1/\\Mmin\\sim 10$. A similar drop is seen for the brightest sample ($M_r<-22$) in \\citet{Zehavi05b}, which has a mean redshift $\\sim0.16$. The decrease of the scaling factor reflects the balance between accretion and destruction of satellites (\\citealt{Zentner05}) --- massive halos assemble more recently and their satellites have less time to merge with the central galaxy. The relatively higher redshift of the samples further strengthens this effect.  The HOD of LRGs has been inferred using different methods in several recent investigations. To compare our results with others, one needs to pay attention to the differences in the sample definition, the underlying assumptions, and the modeling details.  \\citet{Mandelbaum06} present a mass determination of host halos for two SDSS LRG samples based on galaxy-galaxy lensing measurements. The construction of their LRG samples is not identical to ours, but the halo masses determined from their two samples appear to be consistent with those of the two luminosity-threshold LRG samples we model. The galaxy lensing directly measures the halo masses, while our results come from the mass distribution of halos and the galaxy assignment required to reproduce the observed clustering. The agreement between the two results is therefore encouraging. \\citet{Wake08} model the projected two-point correlation functions for LRGs in the 2dF-SDSS LRG and QSO survey with a three-parameter model. For the SDSS $z=0.21$ LRG sample, their inferred mean occupation function is in general agreements with ours. \\citet{Blake08} perform HOD modeling of the two-point angular correlation function of $0.4<z<0.7$ SDSS LRGs with photometric redshifts. Their parameterization is a slight variation of our five-parameter model presented in Appendix~\\ref{sec:appendixB}. For samples with similar number densities, their inferred halo mass scales, cutoff widths of central galaxy occupation function, and high mass slopes of satellite occupation function are close to what we obtain. \\citet{Kulkarni07} constrain the HOD of SDSS LRGs with redshift-space two-point and three-point correlation functions, by comparing the measurements to those from mock catalogs generated through populating halos identified in $N$-body simulations. They use a three-parameter HOD description and assume no velocity bias. They find that redshift-space three-point correlation functions favor a lower high mass slope ($\\sim 1.4$) for the satellite occupation function. Since the $z=0$ outputs of $\\sigma_8=0.9$ simulations are used in their modeling while the median redshift of LRGs is about 0.3, the effective $\\sigma_8$ in their modeling is about 1.05. From Appendix~\\ref{sec:appendixB} (eq.~[B3]), we see that the high mass slope from our modeling is 1.5 for such a high $\\sigma_8$, which is close to the value favored by \\citet{Kulkarni07}. However, we note that they use quite a different halo definition than ours (with a much lower overdensity threshold), which complicates the comparison. A more detailed comparison of the LRG modeling results with different methods and samples can be found in \\citet{Brown08}.  The inferred high mass slopes of the LRG occupation functions tend to be substantially larger than unity, either from our results or others (e.g., \\citealt{Blake08,Kulkarni07,Wake08}). This appears to differ from observational inferences and theoretical predictions for low luminosity samples (e.g., \\citealt{Zehavi05b,Zheng07,Kravtsov04,Zheng05,Conroy06}). Is the steep slope of the inferred LRG occupation function a true feature or merely a result of modeling imperfections? First of all, the steep slope should not be a result of any restriction in our HOD parameterizations. Our five-parameter HOD model (Appendix~\\ref{sec:appendixB}) introduces a cutoff in $\\Nsat$ at the low mass end so that the connection between the high mass end slope and that at $\\Nsat\\sim$ a few is broken. For the fainter LRG sample, the constraint on the slope up to $M\\sim 2\\times 10^{15}\\hMsun$ remains tight even with our more flexible parameterization [see Fig.~\\ref{fig:wp_LRG}({\\it b})]. In the LRG survey volume, there are not many halos that are more massive than $10^{15}\\hMsun$. The true halo mass function in this volume can therefore deviate from the theoretical one used in the modeling. In Appendix~\\ref{sec:appendixC}, we show that this fluctuation in the halo mass function does not seem to introduce any systematic biases in model fitting as it is already reflected in the covariance matrix of the data. To determine which features of the data drive the steep slope, we fit the faint LRG sample with the five-parameter model after excluding some data points and find that the two data points at $r_p\\sim 1.7\\hMpc$ and $2.7\\hMpc$ play a large role --- if they are excluded, the slope drops from 1.8 to 1.6 (for $\\sigma_8=0.8$). These scales are in the one-halo to two-halo transition region where the model is sensitive to the treatments of halo exclusion and scale-dependent halo bias. A more accurate scale-dependent halo bias with halos defined by spherical over-density (SO) could lead to a somewhat lower value of the high mass end slope (J.~L. Tinker, private communication). \\citet{Reid09} constrain the LRG HOD with the counts-in-cylinders multiplicity function and the correlation function through populating halos in a simulation. They obtain a good fit by using SO halos and the high-mass end slope is found to be close to unity. Compared to friends-of-friends (FoF; \\citealt{Davis85}) finder, the SO halo finder does not have the problem of linking two halos by a thin bridge of particles. We plan to pursue analytic models of galaxy clustering based on SO halo properties \\citep{Tinker08} in future work.  Because of the steep high mass slope, our model fits predict that massive clusters should host multiple LRGs. For example, the LRG occupation number of a $2\\times 10^{15}\\hMsun$ cluster would be about ten. Such a prediction can be tested with a cluster catalog if the mass can be determined. Using a sample of X-ray selected galaxy clusters at $0.2<z<0.6$, \\cite{Ho07} assess cluster membership for LRGs based on their photometric redshifts and assign halo mass based on X-ray luminosity. They define halos as objects with mean density of 200 times the critical density rather than the mean background density as we do, and they assume $\\Omega_m=0.238$. After correcting the differences in the halo definition and cosmology, our HOD result for the faint LRG sample matches theirs in the regime of $\\NavgM \\sim$ a few. (However, near the cutoff mass they find a much larger occupation number, while we have $\\NavgM <1$.) In their catalog, there are only two very massive clusters with masses of $\\sim 8\\times10^{14}\\hMsun$ and $\\sim 1.1\\times10^{15}\\hMsun$ (corrected to be consistent with our halo definition), and the (corrected) numbers of LRG members are about 8 and 14, respectively. We also performed a rough calculation to associate LRGs in the faint sample to MaxBCG clusters \\citep{Koester07} in the overlapped sky region and redshift range. The mass of each cluster is estimated from the total number of MAIN galaxies inside the virial radius, with calibration by weak lensing \\citep{Sheldon07}. For each cluster, we infer the halo radius and velocity dispersion from the halo mass corrected to match our halo definition. LRGs that fall in the projected halo radius and three times the velocity dispersion along the line-of-sight direction are assigned as cluster members (with no completeness and edge correction applied). The number of LRG members for clusters more massive than $6\\times 10^{14}\\hMsun$ is found to range from 0 to 7. The predicted number of LRGs in high mass halos (Fig.~\\ref{fig:wp_LRG}) appears approximately consistent with the estimates from clusters. However, given the uncertainties in the mass estimator and the small number statistics, more work is needed to test our derived values of the high mass slope.  \\begin{figure*} \\epsscale{1.0} \\plotone{f6.eps} \\epsscale{1.0} \\caption[]{ \\label{fig:samcmp} Modified projected two-point correlation function and six-parameter HOD fit. Panel ($a$) shows the original $w_p(r_p)$ data points (filled circles) and the modified ones (open circles), which approximate those predicted by the semi-analytic model of \\citet{Bower06} (see \\citealt{Almeida08}). The thick curve is the best fit to the original data from the five-parameter HOD model and the dotted curve is that from the HOD model presented in \\S~3.1. The thin curve is the best fit to the modified data from the six-parameter HOD model (see the text). The two curves are best HOD model fits. Panel ($b$) shows the best fit mean occupation functions. The thick curves are from fitting the original $w_p(r_p)$ with the five-parameter model and the thin ones are the $\\Delta\\chi^2<1$ envelopes of mean occupation functions from fitting the modified $w_p(r_p)$ with a six-parameter model (see the text) and and those for all galaxies are shaded. Dashed, dotted, and solid curves are for central, satellite, and all galaxies, respectively. } \\end{figure*}  By modeling clustering of LRGs with different luminosities, we infer how the mean luminosity of central galaxies changes with the mass of their host halos, $L_c\\propto M^p$ with $p\\sim$ 0.46--0.51 for $M\\sim 10^{14}\\hMsun$. This is consistent with the relation $M\\propto L^2$ found by \\citet{Mandelbaum06} based on galaxy-galaxy lensing measurements. \\citet{Padmanabhan09} model the clustering of a sample of photometrically selected SDSS LRGs that are fainter than those in the spectroscopic samples we use. Combining their HOD modeling results with ours for LRGs at $z\\sim 0.3$, we find that the value of $p$ is in the range of 0.4--0.5 over a larger range of LRG luminosity. From a study of the halo occupation statistics of galaxies using galaxy groups identified in the Two Degree Field Galaxy Redshift Survey (2dFGRS; \\citealt{Colless01}), \\citet{Yang05} infer a $L_c$--$M$ relation that is well described by a broken power law with $p\\sim$ 2/3 and 1/4 below and above $10^{13}\\hMsun$, respectively. The inferred $L_c$--$M$ relation by \\citet{Vale04} from matching the luminosity function from the 2dFGRS with the theoretical subhalo mass function has a similar high-mass slope $p\\sim 0.28$. \\citet{More09} constrain the high-mass slope to be $0.28_{-0.09}^{+0.07}$ for $z<0.072$ SDSS galaxies based on satellite kinematics. The high-mass slope extrapolated from the HOD modeling results of SDSS MAIN galaxies \\citep{Zheng07} is also about 0.3. The value we obtain at the high-mass end from modeling the LRG clustering differs significantly from these results. It is interesting to see whether such a difference can be explained by the difference in the galaxy samples. The 2dFGRS luminosity is in the $b_J$ band, while we use the SDSS $g$-band luminosity. Since the wavelength coverages of these two bands are close, it seems unlikely that the band difference can cause the apparent discrepancy. The other thing to notice is that the 2dFGRS galaxies and the SDSS MAIN galaxies have a mean redshift $\\sim 0.1$ and the LRG galaxies in our analysis are located around redshift 0.3. Could the discrepancy indicate an evolution effect over the intervening $\\sim$2 billion years? Through fitting restframe $B$-band galaxy luminosity functions at different redshifts using a CLF approach, \\citet{Cooray05} finds that the data are compatible with a halo-mass-dependent central galaxy luminosity evolution, with the high mass slope $p$ increasing with redshift. The fitting results of \\citet{Cooray05} imply that $p$ could be as high as 0.5 at $z\\sim 0.3$, close to our inferred value.  If we take the inferred values of $p$ at $z\\sim 0.3$ from our analysis and at $z\\sim 0.1$ from the 2dFGRS studies at face value, they indicate that the luminosity evolution of central bright galaxies depends on halo mass in the sense that either galaxies in more massive halos fade more or the fraction of stars that were assembled into central galaxies from mergers between $z\\sim 0.3$ and $z\\sim0.1$ is smaller in more massive halos.  \\citet{Bernardi06}'s study of the properties of early-type galaxies in the SDSS as a function of local environment and redshift suggests that star formation in early-type galaxies happens earlier in dense regions and lasts over a shorter time-scale, which implies that central galaxies in more massive halos on average experience star formation at an earlier epoch. Since younger stellar populations fade faster than older ones, this seems to rule out the possibility that luminosity of central galaxies in more massive halos fades more between $z\\sim 0.3$ and $z\\sim 0.1$. We are left with the possibility that mergers in more massive halos are less efficient in adding stars to the central galaxies, which appears to be consistent with our finding of the drop of $M_1/\\Mmin$ and its interpretation based on the competition between accretion and destruction as a function of halo mass. This is also in line with the results of LRG evolution in \\citet{Brown08} from HOD modeling of their clustering from $z\\sim 0.2$ to $z\\sim 1.0$ in the NOAO Deep Wide-Field Survey.  Since LRGs naturally separate out massive halos, our HOD modeling of the two-point cross-correlation functions between LRGs and $L_*$ galaxies circumvents the usual challenge in constraining the HOD of $L_*$ galaxies in massive halos from their two-point auto-correlation function. We show that the cross-correlation data are consistent with the case where the distribution of $L_*$ galaxies inside massive halos roughly follows that of the matter at distances greater than $\\sim 0.2\\hMpc$ and that the mean number of $L_*$ galaxies scales with halo mass as $M^{1.5}$. The slope of the mean occupation function is steeper than what is found for $z\\sim 0$ galaxies, which is $\\sim 1.10$. There may not be inconsistency between the results. The constraint for the slope for $z\\sim 0$ galaxies from auto-correlation function is not sensitive to the occupation in very massive halos, while here from cross-correlation with LRGs, we have a better constraint on the $L_*$ occupation function in massive halos. On the other hand, the selection of $L_*$ galaxies and its redshift dependence make it hard to do a precise comparison between results at $z\\sim 0$ and at $z\\sim 0.3$.  The luminosity dependence of the LRG--$L_*$ galaxy cross correlation depends on scale in a rather complex way \\citep{Eisenstein05a}. By separating contributions from pairs of $L_*$ galaxies with central and satellite LRGs in common halos and in different halos, our HOD modeling results explain these trends, in a relatively transparent way. At a fixed scale, the luminosity dependence in the cross correlation reflects the luminosity-dependent HOD of LRGs. As the scale in consideration changes, the relative contributions of the one-halo central-satellite, the one-halo satellite-satellite, and the two-halo LRG-$L_*$ pairs to the cross-correlation function vary, which leads to the scale-dependent luminosity dependence of the cross correlation.  Our LRG modeling results establish the relation between massive galaxies and dark matter halos at $z\\sim 0.3$, which itself provides useful tests to models of formation of massive galaxies. \\citet{Almeida08} present predictions for properties of LRGs in semi-analytic galaxy formation models. They show that the \\citet{Bower06} model, which is based on the Millennium simulation \\citep{Springel05}, predicts a $z=0.24$ LRG luminosity function that is in good agreement with the observation, although it fails at $z=0.5$. Compared to the measurement, the \\citet{Bower06} model seems to predict the $z=0.24$ LRG two-point correlation function remarkably well on both small and large scales (see their Fig.13). However, the LRG HOD predicted in this model differs significantly from our modeling results presented in this paper: (a) in the \\citet{Bower06} model, LRGs in the fainter sample can reside in halos of $10^{12}\\hMsun$ (with mean occupation number of $\\sim$0.01), much lower than the mass scale we infer; (b) the mass of halos that can on average host one LRG is about $3\\times 10^{14}\\hMsun$, higher than what we find (see Fig.11 in \\citealt{Almeida08} and Fig.16 in \\citealt{Wake08}); (c) the probability distribution of LRG host halo masses is broad (see  Fig.16 in \\citealt{Wake08}), ranging from $10^{12}\\hMsun$ to $10^{15}\\hMsun$, rather than narrowly peaked around a few times $10^{13}\\hMsun$ as we find (Fig.~\\ref{fig:wp_LRG}$d$).  Does the discrepancy between the theoretically predicted HOD and our observationally inferred HOD imply that our HOD parameterization is not generic enough to model LRG clustering? To investigate the problem, we modify the five-parameter HOD model (Appendix~\\ref{sec:appendixB}) to mimic the shape of the mean occupation function predicted in \\citet{Bower06} model. The parameterization for the satellite mean occupation function remains unchanged (with three parameters). For the central galaxy occupation function, we model it as a linear (in logarithmic space) ramp going from $\\langle N_{\\rm min} \\rangle$ at $\\Mmin$ to unity at $M_u$ and staying at unity for $M>M_u$. In total, this parameterization has six free parameters. With the measured $w_p(r_p)$, we find that the best-fit HOD from this six-parameter model (not shown in Figure~\\ref{fig:samcmp}$b$) closely follows the result from the five-parameter model (shown as thick curves in Figure~\\ref{fig:samcmp}$b$). Therefore, change in the HOD parameterization does not solve the discrepancy and the five-parameter model is not inadequate in modeling LRG clustering. A close look at Figure~13 in \\citet{Almeida08} shows that the predicted two-point correlation function does not match the data perfectly --- it is $\\sim$15\\% higher on scales of 0.1--2$\\hMpc$ and $\\sim 30\\%$ lower on scales larger than 2$\\hMpc$ (note that the vertical range of the plot is over eight orders of magnitude). We therefore modify the amplitude of the observed $w_p(r_p)$ data points to mimic the predicted correlation function and perform an HOD fit with the six-parameter model (by adopting the predicted number density, which is 10\\% lower than the observed one). The results on the mean occupation functions are shown in panel ($b$) of Figure~\\ref{fig:samcmp} as thin curves. Note that the shaded region represents the $\\Delta\\chi^2<1$ envelopes and the linear ramp for the central galaxy mean occupation number (thin dashed curves) has not reached unity at the highest mass in the plot. The mean occupation function can extend to halos of mass as low as $2\\times 10^{12}\\hMsun$ and reaches unity around $4\\times 10^{14}\\hMsun$, which appears to be approximately consistent with the prediction of the \\citet{Bower06} model. From the above investigations, we conclude that the discrepancy between the theoretically predicted LRG HOD and the observationally inferred LRG HOD reflects the imperfection of the semi-analytic galaxy formation model --- the $15-30\\%$ discrepancies with observed clustering are real and physically significant --- rather than limitations in our HOD parameterization. The results suggest that the mechanism of turning blue galaxies to red in the semi-analytic model is too efficient in halos of a few times $10^{12}\\hMsun$ and not efficient enough in more massive halos. Our HOD modeling results thus provide important tests to galaxy formation theory.  In combination with passive evolution of LRGs and a halo merger history (e.g., \\citealt{White07,Seo08}), our modeling results can be used to predict the HOD of LRGs and the clustering of LRGs at lower or slightly higher redshifts. Supplemented with corresponding observations at these redshifts, we would be able to test our understanding of the formation and evolution of massive galaxies. Constraints on the HOD of LRGs are also useful for cosmological parameter investigations based on LRG clustering. For example, the most precise measurements of the large scale galaxy power spectrum have come from the SDSS LRG sample (\\citealt{Tegmark06,Percival07}), and the principal limitation in interpreting these measurements is the uncertain level of scale-dependent bias between galaxy and matter power spectra in the mildly non-linear regime. With HOD constraints like those derived here, this scale-dependent bias can be calculated and corrected \\citep{Yoo09}. As another example, HOD constraints could be combined with galaxy-galaxy lensing measurements of the LRG-mass cross-correlation (R. Mandelbaum et al., in preparation) to improve determinations of $\\sigma_8$ and $\\Omega_m$ \\citep{Yoo06}. The role of LRGs in cosmological studies seems destined to grow with surveys that target large numbers of LRGs to measure baryon acoustic oscillations, including AAOmega LRG survey \\citep{Ross08} and the Baryon Oscillation Spectroscopic Survey (BOSS), part of a proposed successor to SDSS-II. Understanding the evolving relation between LRGs and dark matter halos will be crucial to exploiting their power as cosmological probes and to revealing the physics that governs the formation of the most massive galaxies in the universe."
    },
    "0809/0809.4160_arXiv.txt": {
        "abstract": "To minimize instrumentally induced systematic errors, cosmic microwave background (CMB) anisotropy experiments measure temperature differences across the sky using paires of horn antennas, temperature map is recovered from temperature differences obtained in sky survey through a map-making procedure. To inspect and calibrate residual systematic errors in  recovered temperature maps is important as most previous studies of cosmology are based on these maps. By analyzing pixel-ring couping and latitude dependence of CMB temperatures, we find notable systematic deviation from  CMB Gaussianity in released Wilkinson Microwave Anisotropy Probe (WMAP) maps. The detected deviation can not be explained by the best-fit $\\Lambda$CDM cosmological model at a confidence level above $99\\%$ and can not be ignored for a precision cosmology study. ",
        "introduction": "To minimize instrumentally induced systematic errors, the WMAP mission measures temperature differences across the sky using paires of horn antennas with a fixed separation angle $\\theta$, temperature maps are recovered from raw time-ordered temperature differences obtained in sky survey with a map-making algorithm$^{[1]}$. When an antenna points to a sky pixel $i$, the scan path of the other one will draw a ring $R_\\theta(i)$ in the sky with angular radius  $\\theta$ to the center pixel $i$. The map temperature $t(i)$ of a pixel $i$ should be in some extent correlated to measured temperatures in its scan-ring $R_\\theta(i)$ through the map-making procedure. The beam separation angle of WMAP radiometers  is $\\theta\\sim141\\degree$. The off-diagonal terms of noise covariance matrixes have been inspected by the WMAP team  with two-point correlation functions and small positive blips of order $0.3\\%$ of the diagonal elements have been found at $141\\degree$ pixel separation angle$^{[2,3]}$.  Instead of the pixel-pixel noise coupling, in this work we inspect the pixel-ring temperature coupling, and find in released WMAP CMB maps that notable temperature distortions exist on scan-rings of hot foreground sources: scan-rings of hot sources are significantly cooled (presented in \\S2) and strongest anti-correlations between pixel and scan-ring temperatures appear at a separation angle $\\theta\\sim141\\degree$ (\\S3). In \\S4 we find a systematic deviation from CMB isotropy in Galactic latitude distributions of released WMAP maps, which can be understood by pixel-ring temperature coupling found in \\S2 and \\S3. Finally we give a brief discussion in \\S5. ",
        "conclusions": "The magnitude of temperature distortion of the hottest Galactic sources upon their scan-rings can be estimated to be $>10\\mu$K from Table~1 (the corresponding region covers about $15\\%$ of the entire celestial sphere after removing KP0 mask region), and the average distortion for the region of $|b|<55\\degree$ to be $\\sim 5\\mu$K from Table~4 (the corresponding region covers $>80\\%$ of the entire celestial sphere or $>60\\%$ after removing KP0 mask region). Such a wide-spread effect with considerable strength can not be ignored for a precise cosmology study and need to be further investigated.  The detected large scale distortion from CMB Gaussianity is closely connected with the Galactic hot sources and the WMAP beam separation angle, which is hard to explain by the best-fit $\\Lambda$CDM cosmological model and most possibly comes from foreground effect and map-making. To mitigate errors produced by effects from the hot foreground sources, the WMAP team used the Kp8 mask as a \"processing mask\" during the map-making process$^{[2,7]}$. However, the detected temperature coupling between hot sources and their scan-rings indicates that the Kp8 mask might be not wide enough to exclude all unwanted effect in the map-making process, a broader mask, e.g. Kp0, might be better. The future CMB observation project Planck is designed to measure the CMB anisotropy with completely different mode to WMAP, therefore, it is expected to be totally unaffected by such distortion."
    },
    "0809/0809.1435_arXiv.txt": {
        "abstract": "{ High precision radial velocity (RV) measurements in the near infrared are on high demand, especially in the context of exoplanet search campaigns shifting their interest to late type stars in order to detect planets with ever lower mass or targeting embedded pre-main-sequence objects.  ESO is offering a new spectrograph at the VLT -- CRIRES -- designed for high resolution near-infrared spectroscopy with a comparably broad wavelength coverage and the possibility to use gas-cells to provide a stable RV zero-point.  We investigate here the intrinsic short-term RV stability of CRIRES, both with gas-cell calibration data and on-sky measurements using the absorption lines of the Earth's atmosphere imprinted in the source spectrum as a local RV rest frame. Moreover, we also investigate for the first time the intrinsic stability of telluric lines at 4100\\,nm for features originating in the lower troposphere.  Our analysis of nearly 5 hours of consecutive observations of MS~Vel, a M2II bright giant centred at two SiO first overtone band-heads at 4100\\,nm, demonstrates that the intrinsic short-term stability of CRIRES is very high, showing only a slow and fully compensateable drift of up to 60\\,m/s after 4.5 hours. The radial velocity of the telluric lines is constant down to a level of approx. $\\pm10$\\,m/s (or 7/1000 of one pixel). Utilising the same telluric lines as a rest frame for our radial velocity measurements of the science target, we obtain a constant RV with a precision of approx. $\\pm20$\\,m/s for MS~Vel as expected for a M-giant.} ",
        "introduction": "CRIRES\\footnote{CRIRES stands for \\underline{CR}yogenic \\underline{I}nfra\\underline{R}ed \\underline{E}chelle \\underline{S}pectrograph} is a newly commissioned adaptive optics (AO) - fed spectrograph at UT1 (Antu) of the Very Large Telescope on Paranal, Chile. The spectrograph provides a resolution of up to 100\\,000 at 960--5200\\,nm with an instantaneous wavelength coverage of $\\sim\\lambda/50$ on a 4096$\\times$512 pixel detector mosaic \\citep{kaufl04, kaufl06b}.  % A MACAO\\footnote{MACAO stands for \\underline{M}ultiple \\underline{A}pplication \\underline{C}urvature \\underline{A}daptive \\underline{O}ptics} % adaptive optics system \\citep{arsen2004} is providing diffraction limited images at the entrance slit of CRIRES with Strehl ratios of up to 65\\% in the K band, thus roughly doubling the throughput under reasonable seeing conditions for the nominal 0.2\\arcsec~slit when compared to seeing limited performance \\citep{jerome06}.  The spectrum, being either a full order at the shortest wavelength in the J band and about 1/5 of an order at the longest wavelength in M band, is imaged in long slit mode onto a mosaic of four Aladdin III detectors with a gap of $\\sim 250$ pix between the chips. Thus, the instantaneous wavelength coverage is much shorter than for optical cross-dispersed spectrographs.  However, the gain in efficiency for RV measurements, primarily dependent on spectral coverage and S/N, will shift in favour of CRIRES as soon as the effective SED of the science target peaks longwards of 900\\,nm where other optical spectrographs commonly used for RV measurements, are either not operating (e.g. HARPS) or loose dramatically in sensitivity (e.g. UVES).  For precise RV measurements the spectrograph has to be well characterised since we deal with various sources of uncertainties that could induce wavelength shifts in the measured spectra.  Since CRIRES is a cryogenic, thus actively stabilised spectrograph, temperature drifts are small but certainly not negligible. The same is true for vibrations and flexure, especially due to changing thermal gradients in the structure. Being mounted on a Nasmyth platform of the VLT, the system is intrinsically very stable but vibrations induced by the cryogenic cooling system can be suspected for causing instability.  A second effect is the imperfect repeatability of dispersive parts on moving functions suffering slip-stick effects: the prism (pre-disperser), the intermediate slit and the grating. It is forseen that functions are left untouched when the same instrumental setup is used in subsequent observations. However, for observations that span days to months the repeatability of the setup could be an issue, if it can't be well calibrated.  Last but not least a long-slit spectrograph suffers the natural problem of alignment uncertainties of the source with respect to the slit centre. For spectrographs fed by an adaptive optics system the situation is worse in some aspects:\\label{DCR} \\begin{enumerate} \\item The FWHM of the source is usually smaller than the slit width. \\item The alignment of the adaptive optics in respect to the entrance slit of the spectrograph suffers instabilities, requiring an active slit-viewer guiding in order to keep the source centred on the slit. \\item The algorithm for slit-viewer guiding can only work on the wings of the stellar PSF, since its core is masked by the slit. \\item The refraction caused by the Earth's atmosphere has to be measured and compensated simultaneously as observations are performed at three different wavelengths: in the optical for the AO wave front sensor, in J, H or K band for the slit-viewer (used for active guiding) and at the respective wavelength of the observation, in our case at 4100\\,nm. \\end{enumerate} All these effects will cause a degraded stability and a possible drift of the wavelength scale with time. We thus aimed to characterise the system and the means to calibrate it and test the on-sky performance in terms of RV stability during a single night.  For CRIRES a RV precision of 15~m/s would require pointing and tracking with a precision of one milli~arcsec (mas), given that one pixel in dispersion direction (1.5 km/s) equals 100~mas at the entrance slit of the spectrograph. The tracking precision is limited by the amount of light in the wings of the PSF on which the guiding algorithm of the slitviewer operates. Tracking the core of the PSF on the seeing-limited wings is a general limitation. Even for bright sources, the performance of the guiding algorithm is not sufficient to stabilise the source in the slit with a precision better than approx. 25 mas (approx. 200~m/s), which makes a reference source (telluric or gascell absorption lines) indispensable when aiming for high precision RV measurements. ",
        "conclusions": "We have performed the first radial velocity stability test of CRIRES with 4.5~h of on-sky measurements of the giant MS~Vel at 4100\\,nm. The dispersive parts of the spectrograph have been electronically stabilised, thus decoupling the problems of functional repeatability, spectrograph stability and tracking errors. Besides this non-standard settings we used the spectrograph as offered to the community, which demonstrates that even at this very early stage of its life, CRIRES is a stable instrument, ready to measure radial velocities with high precision including a potential for even higher performance when using an appropriate gas as local RV zero point.  We could demonstrate that CRIRES shows only a slow drift in its intrinsic radial velocity zero point of $\\sim$ 60\\,m/s in 4.5\\,h. This drift is fully compensatable using a N$_2$O gas-cell in calibration mode or telluric lines as a local rest frame. Moreover, if our assumption holds that the drift is due to a small change in temperature of the grating (by approx. 1/20~K), it will not appear in any later measurements as the grating is now actively stabilised to approx. 1~mK.  The telluric lines in emission, observed with a resolution of R$\\sim$50\\,000 are stable in RV to a level of $\\leq \\pm$ 10\\,m/s and are therefore suitable as a local rest frame for high precision RV measurements. For the relative radial velocity of our science target, MS~Vel, we could recover the 300\\,m/s drift induced by rotational and barycentric Earth's motion. The remaining residuals (rms) in the stellar spectra are $\\leq \\pm$ 22\\,m/s.  We can conclude that the telluric lines imprinted in the science object are not a limiting factor for precise RV measurements in the near infrared but can instead be used as a substitute for a gas-cell providing the necessary rest frame in cases where no suitable gas can be found or the interesting stellar features are in regions of high telluric line density. First scientific results using this method are presented in \\citet{TWHya}.  The influence of mismatches between the diffraction limited FWHM of the PSF and the typical slit width in combination with dynamic (i.e. tracking induced) misalignments of the PSF in respect to the slit centre are substantial. However, these problems can be compensated when telluric or gas-cell lines are measured simultaneously with the source and used as a local rest frame. This method is well known and widely used in the optical, where Iodine lines are used as the local restframe. Given the high rate of close blends of telluric lines and photospheric lines from our science target and taken the limitation of our analysis into account, we see a clear margin for a gain in performance. Especially the use of a gas-cell, well characterised by a high resolution laboratory spectrum (R$\\gg$100\\,000), should provide a better local rest frame when aiming at RV precision below $\\approx$20\\,m/s. The gas has to exhibit a high and stable line density in regions where the science target shows also a reach spectrum. To avoid blends with telluric lines, rare isotope enhanced gases could be used which would elegantly avoid problems with blending telluric features.  In the example presented here - as for all conceivable RV applications of CRIRES - rotational vibrational transitions of molecules in the vibrational ground-state are used.  These gases behave very close to ideal gases. This implies that the density of molecules in the beam is invariable and relatively easy controllable. This situation is fundamentally different in the Iodine gas-cells which are being used for radial velocity work. There iodine vapour is used in a heated cell and condensation on cooler parts can result in strong changes of the vapour pressure. Therefore, in our case it is safe to assume, that the transition frequencies are very close to those, which have been determined by NIST using heterodyne-techniques relative to the Caesium time standard. The spectra are relatively simple and lines which may consist of a blend can be positively excluded. The only conceivable process which may change transition frequencies in the gas-cell in the case of an ideal gas is pressure shift. The pressure may change due to the external temperature, which will be monitored. It has been estimated that this effect is less than ~1m/s. A more serious problem is pressure shift by air, in case the gascell has a leak. Fig. 2b in \\citet{Glenar} gives an example. A PH3-gascell line in the presence of a leak is measured relative to a gain-stabilised CO2-laser in a heterodyne setup. The kinks in the PH3-trace correspond to refillings of the cell. To control this problem, the pressure in the CRIRES gas-cells will be monitored as part of regular operations.  A gascell filled at low pressure with species, for which NIST has done a heterodyne reference measurement relative to the time standard (Caesium-clock), in principle provides for an absolute secondary frequency standard. To that end CRIRES measurements are linked in an absolute and traceable way to the time standard. This is, strictly speaking, not an actual experimental proof, that this indeed allows for a \"cm/s\"-type calibration of the spectrograph. Indeed e.g. a spurious blend of gas-cell lines with weak telluric lines may produce artificial shifts exceeding this value. Experts in the field indeed stress that the frequency calibration of spectrographs breaks down at a precision of $10^{10}$ and only heterodyne techniques can provide for this precision (Theodor Haensch, priv. communication).  We finally note that this paper aimed at demonstrating the short-term stability of CRIRES and the implications when observing with an underfilled entrance slit. The long-term stability of the instrument, especially important for RV studies, will be considered in a later paper."
    },
    "0809/0809.4789_arXiv.txt": {
        "abstract": "We discuss the role of mass loss for the evolution of the most massive stars, highlighting the role of the predicted {\\it bi-stability} jump that might be relevant for the evolution of rotational velocities during or just after the main sequence. This mechanism is also proposed as an explanation for the mass-loss variations seen in the winds from Luminous Blue Variables (LBVs). These might be relevant for the quasi-sinusoidal modulations seen in a number of recent transitional supernovae (SNe), as well as for the double-throughed absorption profile recently discovered in the H$\\alpha$ line of SN 2005gj. Finally, we discuss the role of metallicity via the $Z$-dependent character of their winds, during both the initial and final (Wolf-Rayet) phases of evolution, with implications for the angular momentum evolution of the progenitor stars of long gamma-ray bursts (GRBs). ",
        "introduction": "Mass loss has a major effect on the evolution of stars of all initial masses, however its effect is most prominent for the more massive stars due to their large luminosities -- in close proximity to the Eddington limit. Mass loss is relevant both in terms of evolutionary pathways as well as the properties of the pre-supernova (SN) circumstellar environments. In this contribution, we first discuss mass-loss predictions that are relevant for predicting the forward evolution of massive stars. There are actually two aspects that need to be accounted for: (i) the loss of {\\it mass} as winds ``peel off'' the star's outer layers (Conti 1976), but as massive stars start their evolution as rapid rotators, also (ii) the associated loss of {\\it angular momentum} (e.g. Langer 1998).  Towards the end of the main sequence massive stars encounter the so-called bi-stability jump, for which we discuss the implications of the loss of angular momentum. For the final stages, the evolution of angular momentum is particularly relevant for our understanding of the long gamma-ray burst (GRB) phenomenon, as the popular collapsar model (MacFadyen \\& Woosley 1999) requires the core of the progenitor star to be rapidly rotating before collapse (but see also Lee \\& Ramirez-Ruiz 2006). A key parameter in the story is that of metallicity, due to the $Z$-dependence of radiation-driven winds during both the main sequence and evolved Luminous Blue Variable (LBV) and Wolf-Rayet (WR) phases of massive star evolution.  In the canonical scenario of massive stars, LBVs represent a transitional phase between the main sequence and the WR phase, however we also discuss the possibility for a new evolutionary paradigm in which the variable winds of LBVs might betray themselves as the {\\it direct} progenitors of SNe. ",
        "conclusions": "\\label{sec:concl}  We presented theoretical mass-loss rates and their implications for the peeling off and angular momentum loss of massive stars during the various evolutionary stages from the main sequence, to the LBV, and WR phases. The role of the bi-stability jump at an effective temperature of $\\sim$25\\,000 K was discussed in the context of the observed drop in rotational velocities in this part of the Hertzsprung-Russell diagram as well as for the variable winds of LBVs, suggesting these objects could be in a {\\it direct} pre-SN state of massive star evolution."
    },
    "0809/0809.2098.txt": {
        "abstract": "We present photometric  redshifts and spectral energy  distribution (SED) classifications for  a sample  of 1542 optically identified sources detected with {\\it XMM} in the COSMOS field. Our template fitting classifies 46 sources as stars and 464 as non-active galaxies, while the remaining 1032 require templates with  an AGN contribution. High  accuracy   in  the  derived photometric redshifts was accomplished as the result of  1) photometry in  up to 30  bands with high significance  detections, 2) a  new set of  SED  templates   including 18   hybrids covering  the   far-UV  to mid-infrared, which have been  constructed by the combination of AGN and non-active  galaxies templates, and  3) multi-epoch observations that have been  used to correct for variability  (most important for type  1  AGN).  The  reliability  of  the  photometric redshifts  is evaluated  using the  sub-sample   of  442  sources  with  measured spectroscopic  redshifts.  We achieved an accuracy  of $\\sigma_{\\Delta   z/(1+z_{spec})}   =  0.014$ for i$_{AB}^*<$22.5 ($\\sigma_{\\Delta z/(1+z_{spec})} \\sim0.015$  for  i$_{AB}^*<$24.5). The high accuracies  were accomplished for both type 2  (where the SED  is often dominated  by the host  galaxy) and type 1 AGN and QSOs out to $z=4.5$. The number of outliers  is a large improvement over  previous photometric redshift estimates for X-ray selected sources (4.0\\% and 4.8\\% outliers for i$_{AB}^*<$22.5 and i$_{AB}^*<$24.5, respectively). We show that the intermediate band photometry is vital to achieving accurate photometric redshifts for AGN, whereas the broad SED coverage provided by mid infrared ({\\it Spitzer}/IRAC) bands is important to reduce the number of outliers for normal galaxies. ",
        "introduction": "It is now well established that  AGN play a significant role in cosmic evolution -- tracing the growth  of super massive black holes in galaxy nuclei and providing important energetic feedback to the host galaxies and  the  IGM.  Characterization  of  the  AGN  properties  and  their redshifts  is  therefore  of  paramount  importance  to  understanding cosmological evolution.   The mass of the central  super massive black hole  (SMBH)  directly  correlates  with  the mass  and  the  velocity dispersion     of     the     bulge     of     its     host     galaxy \\citep{Gebhardt:2000jt,Ferrarese:2000hb}    and   these   correlations apparently persist  both for normal, less-active galaxies  and for AGN \\citep{Mclure:2002la}. This  is strongly suggesting  a co-evolution of galaxies and their nuclear BHs.  Unfortunately, given the low fraction of galaxies with high luminosity AGN and the fact that those with high luminosity may give a biased view of the overall importance of the AGN feedback,  the co-dependent  evolution  of galaxies  and  BH has  been pursued  very little  at high  redshift,  where the  majority of  both galaxy  and AGN  evolution occurs.  Here  we address  this problem  -- attempting to  identify a much  larger sample of  AGN over a  range of luminosity  and  to provide  classifications  and  redshifts for  much fainter samples.  At present, no criterion exists to identify all AGN in a field using a single photometric  band. Optical,  infrared, radio, narrow  bands and X-ray  techniques   each  yield  incomplete,   partially  overlapping, sub-samples of the overall AGN  population.  A complete census can only be  achieved  by  combining  the  results  of  the  various  selection techniques (Zamorani et~al.,  in preparation). Nevertheless, the X-ray band  appears to  be  very  efficient and  rather  complete for  most classes of AGN \\citep{Brandt:2005fj}.  The COSMOS  survey provides full multi-wavelength  data-sets from X-ray to  radio and  we focus  here  on {\\it  XMM}--COSMOS selected  sources \\citep{Hasinger:2007dn} using spectral  energy distributions (SEDs) to disentangle  between  non-active  galaxies  and  AGN  and  to  provide accurate  photometric  redshifts.   Photometric redshifts  for  normal (inactive) galaxies can reach  an accuracy of $\\sigma_{\\Delta z/(1+z)} $    =    0.05    \\citep{Grazian:2006kx,Ilbert:2006vl}    or    better \\citep{Gabasch:2004qq,Wolf:2004yq} and most  recently better than 0.02 (Ilbert et~al. 2008).  In contrast, for  AGN the most accurate          results         have          been         0.06--0.1 \\citep{Zheng:2004gd,Polletta:2007cs,Rowan-Robinson:2008ph}          for redshift z$<$1.  Uncertainties in the  SED fitting of AGN come  from two major sources. First,  AGN can vary  with time  and quasi-simultaneous  photometry is required to  obtain a snapshot  SED. Second, the observed  emission is often  a  superposition of  AGN  and  host  galaxy. Depending  on  the properties  of the host  and the  relative power  of the  nucleus, the different  spectral  bands can  be  dominated  by  either of  the  two components.  For example, NGC3079 is a Seyfert 2 galaxy, identified on the  basis of  its  optical nucleus  and  a radio  jet,  yet its  {\\it Spitzer}/IRS spectrum  is often used  as a template for  non active, starburst  galaxies in  the infrared  \\citep{Weedman:2006ve}.   On the other  hand, Mrk231, a  Seyfert 1  and ultra-luminous  infrared galaxy (ULIRG), exhibits  strong silicate absorption in the  IRS spectrum and is therefore highly obscured  with much AGN (and starburst) luminosity re-emitted in the far infrared.  Clearly, to analyze samples which, as {\\it XMM}-COSMOS, contain both galaxies  and AGN, one needs: 1) a fine sampling and full spectral  coverage with photometry spanning from the UV  to  mid-infrared and  2)  spectral  energy distribution  templates representing the  full range  of AGN and  QSOs together  with variable mixing fractions of AGN and host galaxy emission.  Only recently, with  {\\it Spitzer} and {\\it GALEX}  providing full SED templates      of      AGN      beyond     the      optical      bands \\citep[e.g.][]{Polletta:2007cs} such  analysis has become  possible at moderately high redshift.  In this paper we derive photometric redshifts for the {\\it XMM}-COSMOS sources -- mostly dominated  by an AGN (see $\\S$~\\ref{sec:results}) -- achieving  an  accuracy  comparable   that  obtained  for  non  active galaxies. The  COSMOS field  \\citep{Scoville:2007uo} with its  size (2 square degrees) and deep photometry  in 30$^+$ bands from the radio to the X-ray, is an excellent testbed for such a census.  In $\\S$~\\ref{sec:dataset}, we describe the sample and the photometric data set.  The variability analysis is presented $\\S$~\\ref{sec:variability}.  In $\\S$~\\ref{sec:fitting} the procedure used for the AGN classification and photometric redshift determination is discussed.  The results are given in $\\S$~\\ref{sec:results} and the description on the released catalog are described in $\\S$~\\ref{sec:catalog}.  The discussion and the conclusions end the paper in $\\S$~\\ref{sec:discussion} and $\\S$~\\ref{sec:conclusions}.  Throughout this work we use AB magnitudes and assume H$_0$=70\\,kms$^{-1}$Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.7, and $\\Omega_{M}$=0.3.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions} In   this   paper  we   presented   the   photometric  redshifts   and classifications  for  1542   {\\it  XMM}-COSMOS  sources  with  optical counterpart.   Using SED  matching to  up to  30 photometric  bands we identified 46 Galactic  foreground stars, 464 sources best  fit with a non-active galaxies template and 1032 best fit by a template which has an  AGN contribution.  A high  reliability in  recognizing  stars from galaxies and AGN was demonstrated using a spectroscopically identified sub-set of  470  sources.  Overall, we  reach a photometric redshift accuracy of   $\\sigma_{\\Delta z/(1+z_{\\rm spec})}=0.014$ (i$^+_{\\rm AB}<22.5$)  with  4.0\\,\\% outliers.  However, the most important result  is perhaps  the accuracy for  the sub-sample  of AGN dominated  sources ($\\sigma_{\\Delta z/(1+z_{\\rm  spec})}=0.011$) which marks  an significant improvement  in photometric redshift estimation for this kind of sources. The number of  outliers, however, is still non-negligible (6.3\\% when considering  the \"qsov\" sample alone).  The excellent photometric redshifts and classifications have been made possible by a number of factors.  Here, the most important was likely the large spectral coverage from UV to 8\\,$\\mu$m, densely sampled with broad, intermediate and narrow filters.  The 25 optical bands essentially mimic low resolution spectra and are powerful in constraining the positions of spectral features.  The near/mid-infrared coverage provides the spectral range required to better distinguish stars from galaxies and helps to constrain SED templates, especially for passive galaxies.  Also important was the recognition of and correction for variability. This is especially relevant for AGN dominated sources, for which photometry collected over a large time span may dramatically vary from a snapshot SED.  We applied a correction for variability by re-scaling the optical photometry to a common epoch close to that of IRAC observations.  The correction decreased the number of outliers by a factor of 2.  The significant improvement obtained in our sample with this correction demonstrates the importance of quasi-simultaneous observations for projects aiming for good photometric redshift estimates of AGN. Future surveys will have to consider this point very carefully.  Deep imaging in a common filter in each observing epoch would provide the possibility of a valuable first order correction.  Finally, we used a revised set of SED templates to fully represent the entire range of different contributions of active galactic nuclei and their host galaxies. Empirical SED templates are often obtained from objects in the nearby Universe where the nuclear and extended fluxes are easily resolved.  Their application to high-redshift objects, where the two contributions are mixed together, may not always be appropriate.  Our hybrid templates already showed a significant improvement, although further advancements are required.  The catalog of {\\it XMM}-COSMOS selected sources with photometric redshifts and classifications, together with the X-ray properties (Brusa et~al.  2008) offers an unprecedented possibility to study the luminosity function and redshift distribution of AGN. In combination with the immense multi-band COSMOS data set and the new photometric redshifts for optically selected sources (Ilbert et~al. 2008), studies of the environment of AGN will also allow to probe AGN fuelling models, clustering and evolution."
    },
    "0809/0809.1679_arXiv.txt": {
        "abstract": "A simple model of quintessential inflation with the modified exponential potential $e^{-\\alpha \\phi} \\left[A+\\left(\\phi-\\phi_0\\right)^2\\right]$ is analyzed in the braneworld context. Considering reheating via instant preheating, we conclude that the model exhibits transient acceleration at late times for $0.96 \\lesssim A \\alpha^2 \\lesssim 1.26$ and $271 \\lesssim \\phi_0\\, \\alpha \\lesssim 273$, while permanent acceleration is obtained for $2.3\\times10^{-8} \\lesssim A \\alpha^2 \\lesssim 0.98$ and $255 \\lesssim \\phi_0\\, \\alpha \\lesssim 273$. The steep parameter $\\alpha$ is constrained to be in the range $5.3 \\lesssim \\alpha \\lesssim 10.8$. ",
        "introduction": "The Universe seems to exhibit an interesting symmetry with regard to the accelerated expansion, namely, it underwent inflation at early epochs and is believed to be accelerating at present.It is then natural to ask whether one can build a model to join these two ends without disturbing the thermal history of the Universe. Attempts have been made to unify both these concepts  in which a single scalar field plays the role of the inflaton and quintessence - the so-called quintessential inflation.  On the other hand, in recent years there has been increasing interest in the cosmological implications of a certain class of braneworld scenarios where the Friedmann equation is modified at very high energies. In particular, in the Randall-Sundrum type II (RSII) model~\\cite{Randall:1999vf} the square of the Hubble parameter, $H^2$, acquires a term quadratic in the energy density, \\begin{equation} H^2 = {1 \\over 3\\, M_4^2}\\, \\rho\\, \\left[1 + {\\rho \\over 2 \\lambda}\\right]~. \\label{eq:Friedmann} \\end{equation} allowing slow-roll inflation to occur for potentials that would be too steep to support inflation in the standard Friedmann-Robertson-Walker (FRW) cosmology. where $M_4$ is the 4D reduced Planck mass and  $\\lambda$ is the brane tension.   In this paper we show that the modified exponential potential (hereafter we adopt natural units, $M_4=1$, unless stated otherwise) \\begin{equation} V(\\phi)=e^{-\\alpha \\phi} \\left[A+\\left(\\phi-\\phi_0\\right)^2\\right] \\label{eq:albrecht} \\end{equation} also leads to a successful quintessential inflation model.  In the context of quintessence, this potential was first analyzed by Albrecht and Skordis (AS)~\\cite{Albrecht:1999rm}.  The model displays an interesting feature: it can lead to both permanent and transient acceleration regimes.  These models belong to the category of nonoscillating models in which the standard reheating mechanism does not work. In this case, one can employ instant preheating. This mechanism is quite efficient and robust, and is well suited to nonoscillating models~\\cite{Felder:1999pv}.    \\begin{figure*}[t] \\centerline{\\psfig{file=fig4.eps,width=11cm}} \\caption{Parameter space consistent with all the observational constraints considered, for the permanent acceleration case.} \\label{fig:permanent} \\end{figure*}  \\begin{figure*}[t] \\centerline{\\psfig{file=fig5.eps,width=11cm}} \\caption{Same as in Fig. \\ref{fig:permanent}, but for the transient acceleration case.} \\label{fig:transient} \\end{figure*} ",
        "conclusions": "We have analyzed a simple model of quintessential inflation in the RSII braneworld context with a  modified exponential potential. Assuming that the Universe was reheated via the instant preheating mechanism, we have shown that the evolution of the scalar field from inflation till the present epoch is consistent with the observations in a wide region of the parameter space. Requiring that the model meets various cosmological constraints at the different stages of the evolution, we were able to constrain tightly its parameters, as summarized in Eqs.~(\\ref{eq:boundtrans1})-(\\ref{eq:boundNstarf})."
    },
    "0809/0809.1954_arXiv.txt": {
        "abstract": "We incorporate a contribution to reionization from X-rays within analytic and semi-numerical simulations of the 21-cm signal arising from neutral hydrogen during the epoch of reionization. The relatively long X-ray mean free path (MFP) means that ionizations due to X-rays are not subject to the same density bias as UV ionizations, resulting in a substantive modification to the statistics of the 21-cm signal. We explore the impact that X-ray ionizations have on the power spectrum (PS) of 21-cm fluctuations by varying both the average X-ray MFP and the fractional contribution of X-rays to reionization.  In general, prior to the epoch when the intergalactic medium is dominated by ionized regions (H\\,{\\sevensize II} regions), X-ray-induced ionization enhances fluctuations on spatial scales smaller than the X-ray MFP, provided that X-ray heating does not strongly supress galaxy formation.  Conversely, at later times when \\H2 regions dominate, small-scale fluctuations in the 21-cm signal are suppressed by X-ray ionization. Our modelling also shows that the modification of the 21-cm signal due to the presence of X-rays is sensitive to the relative scales of the X-ray MFP, and the characteristic size of \\H2 regions.  We therefore find that X-rays imprint an epoch and scale-dependent signature on the 21-cm PS, whose prominence depends on fractional X-ray contribution.  The degree of X-ray heating of the IGM also determines the extent to which these features can be discerned. We further show that the presence of X-rays smoothes out the shoulder-like signature of \\H2 regions in the 21-cm PS. For example, a 10\\% contribution to reionization from X-rays translates to a 20--30\\% modulation in the 21-cm PS across the scale of \\H2 regions. We show that the MWA will have sufficient sensitivity to detect this modification of the PS, so long as the X-ray photon MFP falls within the range of scales over which the array is most sensitive ($\\sim0.1$\\,Mpc$^{-1}$). In cases in which this MFP takes a much smaller value, an array with larger collecting area would be required. As a result, an X-ray contribution to reionization has the potential to substantially complicate analysis of the 21-cm PS. On the other hand, a combination of precision measurements and modelling of the 21-cm PS promises to provide an avenue for investigating the role and contribution of X-rays during reionization. ",
        "introduction": "\\label{sec:intro} The process of hydrogen reionization started with ionized (\\H2) regions around the first galaxies, which later grew to surround groups of galaxies. Reionization completed once these \\H2 regions overlapped and occupied most of the volume between galaxies. The inhomogeneous nature of the reionization process suggests that the spatial dependence of the statistics of redshifted 21-cm fluctuations will provide the most powerful probe of the reionization epoch \\citep{FOB06}. With this as motivation, several experiments are currently under development that aim to detect the 21-cm signal during reionization, including the Low Frequency Array\\footnote{http://www.lofar.org/} (LOFAR), the Murchison Widefield Array\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} (MWA) and the Precision Array to Probe Epoch of Reionization\\footnote{http://astro.berkeley.edu/$\\sim$dbacker/eor/} (PAPER), and more ambitious designs are being planned such as the Square Kilometre Array\\footnote{http://www.skatelescope.org/} (SKA). Detection of the redshifted 21-cm signal will not only probe the astrophysics of reionization, but also the matter power spectrum (PS) during the epoch of reionization \\citep{McQ+06,B+06}.  In anticipation of forthcoming 21-cm observations, much recent theoretical attention has focused on the prospects of measuring the PS of 21-cm emission \\citep[e.g.][]{ZFH04,F+04,morales2005,bowman2006}. In particular, the ionization structure of the intergalactic medium (IGM) owing to UV emission associated with star formation has been studied in detail using analytic~\\citep[e.g.][]{F+04,B07} and numerical simulations~\\citep[e.g.][]{McQ+06, Il08}. These studies describe a scenario in which very large \\H2 regions form around clustered sources within overdense regions of the IGM. The formation of these \\H2 regions has a significant effect on the shape of the 21-cm PS, moving power from small to large scales and thus leaving a shoulder-shaped feature on the PS at the characteristic scale of the \\H2 regions \\citep{F+04}. The detailed morphology of the ionization structure, which is encoded in the PS, will yield information about the ionizing sources, as well as the structure of absorbers in the IGM on small scales~\\citep{McQ+06}. More recently, semi-numerical models have also been developed to study the formation of \\H2 regions during the epoch of reionization (EoR). These have again focused on the role of UV emission from star-forming regions \\citep{MF07} and quasars \\citep{GW08}.         In addition to UV photons from the first galaxies and quasars, it has been suggested that a background of X-ray photons at very high redshift may be important in explaining the large Thomson optical depth~\\citep{D+08} of the IGM measured by \\textit{WMAP} \\citep{RO04a,PF07}. Indeed, several authors have interpreted this observation by proposing an initial phase of preheating and partial ionization of the IGM by X-rays, resulting primarily from black hole accretion \\citep{SVS85,ME99,HAR00,VGS01,RO04a}. An X-ray background could also have originated from X-ray binaries and supernova remnants, through accretion in the former case and inverse Compton scattering in the latter \\citep[see, e.g.][for a review of these processes and the global evolution of the IGM]{FOB06}.   The process by which this early X-ray background ionizes and heats the IGM begins with the photoionization of hydrogen and helium. The resulting high energy free electrons then deposit their energy through collisional ionization and excitation of hydrogen and helium, and electron-electron scattering with cooler free electrons. Thus X-rays ionize hydrogen both directly and through secondary ionizations by photoelectrons from ionized helium, with the latter dominating so that many hydrogen atoms can be ionized by a single X-ray photon. As reionization proceeds, however, the proportion of each X-ray's energy that is deposited into the IGM as heat, increases. In particular, for ionization fractions $\\gtrsim 10\\%$, ionization by secondary electrons becomes inefficient \\citep{RO04a}, as the electron-electron scattering cross-section becomes large. As a result, X-ray ionization is self regulating at the level of $\\sim10-20\\%$. Further details of the X-ray emission mechanisms are discussed in \\cite{O01}. In addition, observational limits on the possible contribution of X-rays to reionization are provided by the unresolved soft X-ray background. \\cite{D+04} show that the hard X-ray background produced by a population of black holes at high redshift would be observed as a present-day soft X-ray background. They find that accreting black holes with a hard spectrum could not reionize the Universe to greater than the 50\\% level without saturating the unresolved component of the soft X-ray background.  \\cite{PF07} study the effect of the inhomogeneous preheating of the IGM \\citep{O01,VGS01,MRVHO04} that results from fluctuations in the X-ray background, and calculate its impact on 21-cm fluctuations at $z\\geqslant 10$. The effect of this X-ray heating has also been calculated within numerical simulations~\\citep{S+07}. Our work is distinct from these studies, which do not include an X-ray component of reionization. In our modelling we neglect temperature fluctuations, and instead concentrate on the possible effect of X-ray ionization late in the reionization era when \\H2 regions dominate the ionization structure, and the IGM is already heated above the cosmic microwave background (CMB) temperature.  The \\H2 region dominated ionization structure of the IGM seen in numerical simulations forms as a result of the biased clustering of sources, combined with the short mean free path (MFP) of UV photons. In addition to the relatively large amount of energy deposited into the IGM, X-rays differ from UV photons in their effect on the IGM due to their comparatively long MFP. This long MFP means that X-rays provide a weakly fluctuating radiation background, rather than the highly biased and strongly nonlinear ionizing structure thought to arise from UV photons. The ionization structure of the IGM would therefore be modified if X-rays contributed significantly to reionization, because the long X-ray MFPs imply that their ionizing effect is not confined to the vicinity of large scale overdensities in the matter field \\citep{O01}. For this reason, the inclusion of an X-ray component of the emission responsible for the reionization of the IGM could modify fluctuations in the 21-cm PS by reducing the influence of strongly biased ionized regions.   In this paper we extend existing analytic and semi-numerical simulations of reionization to include an X-ray contribution to the ionizing flux. The structure of our paper is as follows: we begin by summarising estimates of the relative X-ray background (\\S\\,\\ref{sec:Xflux}) arising from stellar and accreting stellar mass black hole sources. \\S\\,\\ref{sec:model} describes our analytic model of the 21-cm signal arising from the EoR. In \\S\\,\\ref{sec:model_X} we describe how this model is extended to include a flux contribution from X-rays, and present results for the effect of X-ray ionization on the fluctuation statistics. \\S\\,\\ref{sec:model_num_X} summarizes the semi-numerical scheme used to generate ionization maps of the IGM arising from both UV and X-ray emission. Our semi-numerical results pertaining to the impact of X-rays on the 21-cm PS are discussed in \\S\\,\\ref{sec:results}. In \\S\\,\\ref{noise} we consider the sensitivity of low frequency arrays to X-ray signatures in the 21-cm PS.  We summarize our results in \\S\\,\\ref{sec:summ}.  Throughout this paper we adopt a concordance cosmology for a flat $\\Lambda$CDM Universe, $(\\Omega_{\\rm m},\\Omega_{\\Lambda},\\Omega_{\\rm b},h,\\sigma_8,n) = (0.27,0.73,0.046,0.7,0.8,1)$, consistent with the \\textit{WMAP} constraints from \\cite{komatsu2008}. All distances are in comoving units unless stated otherwise. ",
        "conclusions": "\\label{sec:model} We model fluctuations in the 21-cm signal from neutral hydrogen during reionization using a semi-numerical scheme following the work of \\cite{GW08}, which is based on the analytic model of \\cite{WL07}. The analytic model encodes the preference for galaxy formation in overdense regions \\citep{MW96}, in concert with the suppression of low mass galaxies in ionized regions resulting from ionization feedback. The robust prediction of this model is that overdense regions reionize first. Individual \\H2 bubbles surrounding clustered sources then percolate into less overdense regions, resulting in the eventual overlap of ionized regions throughout the IGM. This behaviour mimics that seen in sophisticated numerical simulations \\citep[e.g.][]{MLZDHZ07}. In this paper we consider a model in which the mean IGM is reionized at $z \\sim 6$ \\citep{F06,GF06,WBFS03}. The model differentiates between star formation in neutral regions, which proceeds in halos above the hydrogen cooling threshold $T_{\\rm min}\\sim 10^4\\,\\rm{K}$, and star formation in ionized regions, which is suppressed below a higher threshold $T_{\\rm ion}\\sim 10^5\\,\\rm{K}$ due to radiative feedback.  The model predicts the time dependent ionization fraction by mass $Q_{\\delta,R}$ of a particular region on a scale $R$ with overdensity $\\delta_R$. On scales much greater than the ionising photon MFP, this ionization fraction evolves according to the following differential equation \\citep{WL07} \\begin{eqnarray} \\label{eq:dqdt} \\frac{dQ_{\\delta,R}}{dt} &=& \\frac{N_{\\rm ion}}{0.76} \\left[ Q_{\\delta,R}\\frac{dF_{\\rm coll}(\\delta_R,R,z,M_{\\rm ion})}{dt}\\right.\\nonumber\\\\ & & \\left.+ (1-Q_{\\delta,R})\\frac{dF_{\\rm coll}(\\delta_R,R,z,M_{\\rm min})}{dt}\\right]\\nonumber\\\\ & &  -\\alpha_{\\rm B}Cn_{\\rm H}^0 \\left[1+\\delta_R\\frac{D(z)}{D(z_{\\rm obs})}\\right](1+z)^3 Q_{\\delta,R}\\,, \\end{eqnarray} where $N_{\\rm ion}$ is the number of ionizing photons per baryon in galaxies, $\\alpha_{\\rm B}$ is the case-B recombination coefficient, $C$ is the clumping factor (taken to be $2$ in the work presented here) and $D(z_{\\rm obs})$ is the growth factor between redshift $z$ and the present. The halo masses $M_{\\rm min}$ and $M_{\\rm ion}$ correspond to virial masses of temperature $T_{\\rm min}$ and $T_{\\rm ion}$ respectively. The last term in equation (\\ref{eq:dqdt}) describes the recombination rate of ionized gas. Ionizing photons are assumed to be produced at a rate proportional to the collapsed fraction of mass $F_{\\rm coll}$ in halos above the minimum thresholds in neutral ($M_{\\rm min}$) and ionized ($M_{\\rm ion}$) regions. In a region of comoving radius $R$ and mean overdensity $\\delta_R(z)$, the relevant collapsed fraction is obtained from the extended Press-Schechter \\citep{PS74} formalism \\citep{B91} \\begin{equation} \\label{eq:Fcol} F_{\\rm coll}(\\delta_R,R,z) = {\\rm erfc}\\left[\\frac{\\delta_{\\rm c}-\\delta_R(z)} {\\sqrt{2(\\sigma^2_{\\rm gal}-\\sigma_R^2)}}\\right]\\,. \\end{equation} From the ionization fraction $Q_{\\delta,R}$ we can calculate the corresponding brightness temperature $T$ using \\begin{equation} T(\\delta_R,R) = 22\\,{\\rm mK}\\left(\\frac{1+z}{7.5}\\right)^{1/2}(1-Q_{\\delta,R})(1+\\delta_R)\\,. \\end{equation}     \\label{sec:summ}  The presence of an ionizing background of X-rays imprints a scale dependent signature on the 21-cm power spectrum during reionization. The impact of X-rays on the topology of reionization depends on the redshift at which one observes, as well as the mean energy of the X-rays and their fractional contribution relative to UV emission. The findings of this investigation can be summarized as follows: \\begin{itemize} \\item  The effect of X-rays on the statistics of 21-cm fluctuations can be prominent on any scale, rather than only on scales comparable to the MFP as might be na\\\"{i}vely expected. \\item  In a mostly neutral IGM, X-ray ionizations enhance 21-cm fluctuations. Conversely, in a mostly ionized IGM, X-rays suppress fluctuations, provided that X-ray heating does not suppress low mass galaxy formation. \\item  ionization by X-rays modifies, and in some cases smoothes out, the signature of \\H2 regions in the 21-cm PS. If reionization of the IGM included a non-negligible contribution due to X-rays, then the modification of the PS shape introduces a feature at a level that is comparable to that introduced by \\H2 regions. \\item The details of PS modification depend on the relative scales of the \\H2 regions and MFP, and on the overall ionization fraction of the IGM. \\item The increase in source bias due to suppression of low mass galaxy formation by X-ray heating counterpoints the decorrelation of ionizing photons and density due to the long X-ray MFP. \\item The MWA will have sufficient sensitivity to detect the modification of the PS due to a 10-30\\% contribution to reionization by X-rays, so long as the MFP falls within the range of scales over which the array is most sensitive ($\\sim0.1$\\,Mpc$^{-1}$). In cases where the MFP takes a much smaller value, an array with larger collecting area would be required. \\item In some scenarios, the requirements of foreground removal will limit the observability of the effect of X-ray ionization on the PS at scales larger than the MFP. However, in all cases, the magnitude of the signature imprinted on the 21-cm PS at small scales is sensitive to the relative X-ray fraction, and will fall within the observable range. \\end{itemize}  In conclusion, our study shows that the details of the PS shape depend on the typical X-ray photon energy, while the magnitude of any modification to the PS is determined by the relative contribution of X-rays to reionization. The possible contribution of X-rays therefore has the potential to substantially complicate analysis of the 21-cm PS. On the other hand, precision measurements and modelling of the 21-cm PS promise to provide a method for investigating the role and contribution of X-rays during reionization.  \\bigskip  {\\bf Acknowledgements} We wish to thank Steve Furlanetto for detailed comments on an earlier draft of this paper.  We also thank the anonymous referee, whose report served to improve this paper. LW and PMG acknowledge the support of Australian Postgraduate Awards. PMG is grateful for the hospitality of the Institute for Theory and Computation at the Harvard-Smithsonian Centre for Astrophysics. The research was supported by the Australian Research Council (JSBW).   \\appendix"
    },
    "0809/0809.3483_arXiv.txt": {
        "abstract": "We identify compact groups of galaxies (CGs) within mock galaxy catalogues from the Millennium Simulation at z=0 with the semi-analytic models (SAMs) of galaxy formation of  Bower et al., Croton et al. and De Lucia \\& Blaizot. CGs are identified using the same 2D criteria as those visually applied by Hickson (1982) to his CGs (HCGs), but with a brightest galaxy magnitude limit, and we also add the important effect of observers blending close projected pairs. Half of the mock CGs identified in projection contain at least 4 accordant velocities (\\emph{mvCG}s), versus 70\\% for HCGs. In comparison to \\emph{mvCG}s, the HCGs are only 8\\% complete at distances $< 9000 \\, \\rm km \\, s^{-1}$, missing the CGs with small angular sizes, a strongly dominant galaxy, and (for the second SAM) the \\emph{mvCG}s that are fainter and those with lower surface brightness. 10\\% of the mock \\emph{mvCG}s are identical to the parent virialized group, meaning that they are isolated, while the remainder are embedded in their parent virialized groups. We explore different ways to determine the fraction of physically dense groups given the data from the simulations. Binding energy criteria turn out to be inapplicable given the segregation between galaxies and dark matter particles. We rely instead on the combination of the three-dimensional length of the CGs (maximum real space galaxy separation) and their elongation along the line-of-sight (ratio of maximum line-of-sight to maximum projected separations), restricting ourselves in both cases to smallest quartets within the CGs. We find that between 64\\% and 80\\% (depending on the SAM) of the \\emph{mvCG}s have 3D lengths shorter than $200 \\, h^{-1} \\, \\rm kpc$, between 71\\% and 80\\% have line-of-sight elongations less than 2, while between 59\\% and 76\\% have either 3D lengths shorter than $100 \\, h^{-1} \\, \\rm kpc$ or both lengths shorter than $200 \\, h^{-1} \\, \\rm kpc$ and elongations smaller than 2. Therefore, chance alignments (CAs) of galaxies concern at most 40\\% of the \\emph{mvCG}s. These CAs are mostly produced from larger host groups, but a few have galaxies extending a few Mpc beyond the host group. The \\emph{mvCG}s built with the Hickson selection (respectively without the close projected pair blending criterion) have 10\\% higher (lower) fractions of physically dense systems. ",
        "introduction": "Compact Groups (CGs) are small, relatively isolated systems of typically four or five luminous galaxies in close proximity to one another. The first example of a CG was found by \\cite{Stephan1877}. Several catalogues of CGs are now available: \\cite{Rose77} and \\cite{Hickson82} visually identified CGs on POSS~I photographic plates. After the Hickson compact group (HCG) catalogue, several CG catalogues have been automatically extracted from galaxy catalogues, themselves automatically extracted from photographic plates: from the COSMOS/UKST Southern Galaxy Catalog \\citep*{PIM94,Iovino02}, from the DPOSS catalogue \\citep{Iovino03,decarvalho05} or CCD frames from the Sloan Digital Sky Survey (SDSS) photometric catalogue \\citep{Lee+04}. CG catalogues have also been extracted from galaxy catalogues in redshift space: from the CfA2 \\citep{Barton+96}, Las Campanas \\citep{AT00}, and SDSS \\citep{Deng+07,McCPES09} surveys, as well as from the 3D UZC Galaxy Catalog \\citep{FK02}. CGs are so compact that the median projected galaxy separation in HCGs is only $39 \\, h^{-1} \\, \\rm kpc$ \\citep{HMdOHP92}.  HCGs have been studied in detail, in particular their internal structures, shapes, morphologies, luminosities, and environments \\citep{Hickson82,HNHM84,Mamon86,HKH88,HR88,MdOH91,Zepf93,Moles+94,PIM94,KF04,TPT06}. To summarise, these studies indicate that CG galaxies have star formation properties, colours and morphological mixes that lie in between binary galaxies and isolated ones.  The nature of the CGs has been a puzzling matter for quite some time. How can a few bright galaxies coexist within less than 100 kpc, given that galaxies are expected to merge fast in such systems \\citep*{CCS81,Barnes85,Mamon87,BCL93}? There are three schools of thought on this matter. One view is that compact groups are recently formed dense systems that are about to coalesce into a single galaxy \\citep{HR88}. The galaxies lost in the merger may be replenished by galaxies in the loose group environment \\citep*{DGR94}, and the predicted rate of formation of CGs appears to be sufficient to explain the observed frequency of HCGs \\citep{Mamon00_IAU174}. The second view states that CGs may be transient unbound cores of looser groups \\citep{Rose77,RDGH94,TYT01}. And the third scenario places CGs as chance alignments of galaxies along the line of sight within larger loose groups (\\citealp{Rose77} for CGs elongated in projection; \\citealp{Mamon86} and \\citealp{WM89} in general), clusters \\citep{WM89} and cosmological filaments \\citep*{HKW95}. In this scenario, the numerous signs of interaction and star formation is explained by the frequent occurrence of binaries and triplets in the chance alignments \\citep{Mamon92_DAEC}.  If CGs are physically dense, their dynamical times should be short (1\\% of the age of the Universe), and the hot intra-group gas should trace a smooth gravitational potential. The launch of X-ray observatories with good sensitivity in the soft X-ray band (ROSAT, ASCA, Chandra and XMM-Newton) has led to the detection of hot X-ray emitting gas from many CGs. Since the X-ray emissivity scales as the square of the gas density, X-ray emission is less prone to projection effects than optical surveys (but see \\citealp*{OLH95}). However, although 22 HCGs were detected out of 32 pointed observations \\citep{PBEB96}, it is not clear what the global fraction of detections would be on the full sample of HCGs (69 [92] groups with at least 4 [3] accordant velocities, according to \\citealp{HMdOHP92}). Moreover, some of the detected groups appear clumpy (e.g. HCG~16 according to \\citealp{DSM99}), which strongly suggests their unvirialized state.  The distinction between compact groups that are dense in 3D, or chance alignments within loose groups or longer filaments is difficult, because redshift space distortion introduces uncertainties in the computation of the line of sight coordinate which might result in misidentified compact configurations. For a group with line of sight velocity dispersion $\\sigma_v$, the redshift distortion will amount to a spread of $\\delta r_z = \\sigma_v/H_0$ in the line of sight coordinate. Assuming that the square velocity dispersion is half the square circular velocity at the virial radius, $\\sigma_v^2 = (1/2)\\,GM(r_v)/r_v$ (appropriate for an $\\rho \\propto 1/r^2$ density profile), one finds $\\delta r_z/r_v = (1/2)\\,\\sqrt{100} = 5$ if the virial radius is defined where the mean density at that radius is 100 times the critical density of the Universe. So redshift space distortions prevent measuring distances within virial systems (see also the Introduction of \\citealp{WM89}).  Nevertheless, there is one CG meeting the HCG criteria discovered by one of us \\citep{Mamon89} that is so close (within the Virgo cluster) that surface brightness fluctuation distance measurements by \\cite{Mei+07} are able to settle the issue of its nature: \\cite{Mamon08} concludes that this CG is a chance alignment of galaxies along the line of sight, at least 440 kpc and most probably 2 Mpc long.  In summary, even though many efforts have been devoted to look for an explanation about the nature of CGs, the debate is still wide open.  The advent of increasingly realistic cosmological simulations now allow one to distinguish whether CGs are truly dense in 3D, or caused by chance alignments within looser groups, filamentary structures, or the general field. In an early pioneering attempt, \\cite*{HKW95}, who identified galaxies as dense knots of cold gas in their $N$-body + hydrodynamical simulation, and searched for CGs in redshift space in many viewing directions. They found four CGs with at least 4 accordant velocities, all of which were longer than $2 \\, h^{-1} \\, \\rm Mpc$ along the line of sight (one was as long as $4 \\, h^{-1} \\, \\rm Mpc$), and yet presented accordant velocities, despite the (Hubble law) stretching of velocities caused by their elongation along the line of sight. The analysis of \\citeauthor{HKW95} suffers from several drawbacks (according to present-day standards for cosmological simulations): the simulation box was small ($44 \\, h^{-1} \\, \\rm Mpc$ wide), the mass resolution was poor (their simulation had $32^3$ dark matter particles and $32^3$ gas particles, and their galaxies were identified with as few as 8 gas particles), and the spatial resolution was poor (the dark matter particles had a softening length of $10 \\, h^{-1} \\, \\rm kpc$). Furthermore, the identification of galaxies with knots of dense gas was not optimal, especially that feedback from supernovae and active galactic nuclei were not incorporated.  In this work, we  quantify the fraction of CGs that can be considered as physically dense entities in samples of automatically identified CGs, based upon more realistic cosmological $N$ body simulations. At present, one can build realistic CGs in two ways: 1) from dissipationless cosmological simulations on top of which galaxies are painted using fairly complex semi-analytical galaxy formation/evolution models; 2) from hydrodynamical codes that resolve galaxies. We have chosen the first approach and use for this purpose the largest cosmological $N$ body simulation ever performed (in 2006, when the present study began), the Millennium Run \\citep{Springel+05}, on which galaxies were identified in three ways, using three different state-of-the-art semi-analytic models (SAMs) of galaxy formation \\citep{Bower+06,Croton+06,dLB07}.  These three galaxy samples provide an opportunity to test both, projection effects and the real nature of systems identified using standard algorithms like that proposed by \\cite{Hickson82}. The CGs are identified in mock redshift-space catalogues constructed from the real-space galaxy sample derived with the semi-analytical model, from the Millennium Run.  In comparison with the analysis of \\citeauthor{HKW95}, our study is based upon a simulation in a box whose volume is over 1 million times greater, with 30 thousand times as many particles, 25 times finer  mass resolution and a softening scale 4 times smaller. However, the simulation we use does not contain gas particles, so the galaxy parameters are highly dependent on the physics of galaxy formation and evolution of the three SAMs that we analyse.  We focus here on the HCG catalogue, which is by far the best studied sample of Compact Groups.  The layout of this paper is as follows. In Section~\\ref{sample}, we present the different steps for the construction of the mock CG catalogue and Sect.~\\ref{classes} describes how the resulting CGs are classified. The conclusions are summarised and discussed in Section~\\ref{discus}. Once our analysis was well advanced, we learnt about the work of \\cite{McCEP08}, who performed a similar analysis of the properties of Hickson-like CGs from the \\cite{dLB07} galaxy catalogue, and found that 70\\% of the mock CGs selected in projection were caused by chance alignments of galaxies. We highlight in Sect.~\\ref{compare} the similarities and several important differences between our two studies. ",
        "conclusions": "\\label{discus} \\subsection{Summary of results}  The aim of this paper is twofold: 1) predicting the nature of automatically identified CGs, and 2) predicting the nature of the well-studied but highly incomplete and biased HCGs. We identify CGs in three mock galaxy catalogues, constructed from the Millennium Simulation at $z=0$ combined with three semi-analytical models \\citep{Bower+06,Croton+06,dLB07}. Several thousand mock CGs are identified using a two dimensional automated algorithm similar to that applied by \\cite{Hickson82} plus a restriction in the brightest galaxy magnitude. We also allowed for CGs that contain isolated and compact subgroups and furthermore considered the important effects of extended galaxies causing confusion for close projected pairs.  Our main results are: \\begin{enumerate} \\item Among the observable \\emph{mpCG}s, $\\sim 60\\%$ have at least 4 galaxies within $1000\\,\\rm km \\, s^{-1}$ of the median velocity of the group, regardless of the SAM. Since our study is carried out on a sample with a redshift cut-off ($z \\sim 0.17$), we tested whether our results might be bias by this fact. We find that identifying in our (magnitude+)volume limited catalogue produces 1.6 times more \\emph{mvCG}s than identifying on an only-magnitude limited catalogue, because our box sample misses distant galaxies that spoil the isolation of mock Compact groups. Correcting for this effect, we deduce that the fraction of \\emph{mpCG}s that survive the velocity filter is 50\\%. In comparison, there are as many as 52 \\emph{vHCG}s among 72 \\emph{pHCG}s, meaning that 71\\% of observed HCGs survive the velocity filter. Given binomial statistics the probability that we find as many as 52 given that 36 (half of 72) are expected is negligible ($P=0.0\\%$).  \\item Comparing the space densities of the \\emph{mvCG}s and the \\emph{vHCG}s, we deduce that the HCG catalogue is only 8\\% complete. A comparison of the parameter distributions between the \\emph{mvCG}s and \\emph{vHCG}s indicates that the HCG is incomplete in groups of small angular sizes, high fraction of light in the brightest galaxy, as well as (for the C06 model) faint brightest galaxy magnitudes and low surface brightness.  \\item We find that the velocity filtering of \\emph{mpCG}s does not necessarily imply that the resulting accordant-velocity CGs are physically dense. We tested different criteria to classify the accordant-velocity CGs according to the maximum 3D galaxy separations and the line of sight elongations. We find that, with the most conservative criterion, \\emph{at least 3/5 of the mock accordant-velocity CGs are physically dense}, although the precise fraction depends on the galaxy formation model used.  \\item The large majority of non-Real mock accordant-velocity CGs are caused by chance alignments within larger groups, rather than within larger regions such as large-scale filaments.  \\item We find that the fraction of chance alignments decreases from 28--47\\% for the particle-based \\emph{pmvCG}s to 24--41\\% for the sample with close pairs removed (\\emph{mvCG}s) to only 14--30\\% once we fold in the biases of the HCG (\\emph{mvHCG}s). This explains in part why simulation studies \\citep{Mamon86,WM89,HKW95} predict more chance alignments than one infers from actual observations.  \\end{enumerate}  Table~\\ref{samples} indicates that typically half of the groups are lost once we apply the criterion to blend close projected pairs of galaxies. While the discarded groups no longer satisfy the selection threshold of 4 galaxies within 3 magnitudes from the brightest, there are still many blended pairs within the groups that survived the blending criterion (in 1/3, 1/4, and 1/8 of the mock CGs for B06, DLB, and C06, respectively, with similar fractions for \\emph{mpCG}s and \\emph{mvCG}s). Could this mean that a significant fraction of the \\emph{pHCG}s and \\emph{vHCG}s contain blended galaxies?  A comparison of the lists of \\cite{Hickson82} (selected from photographic plates) and \\cite{HKA89} (CCD-based) indicates that the latter found 10 HCGs with extra galaxies (24, 26 [+3], 27, 43, 51, 70, 72, 76[+2], 83 and 99), plus one with one galaxy less (HCG 40). % So 10\\% of the original HCG groups contained blended galaxies as discovered with the better CCD photometry. It therefore appears that our fractions of 1/8 to 1/3 of \\emph{mpCG}s with blended pairs is high. But if the truth is in between our particle case and our extended case, then we should have of order 6--17\\% of groups with blended galaxies, which is in rough agreement with what we see in the HCGs.  Now, three HCGs have been observed at much higher resolution with HST imaging. One shows no extra galaxies (HCG 87), while the other two show more interacting units than counted by \\cite{HKA89}: HCG 31 seems to have 7 galaxies and not just 4, while HCG 90 has 5 galaxies and not 4. Binomial statistics suggest that, with 95\\% confidence, the view of 2 groups out of 3 with extra galaxies implies that the fraction of such groups with blended pairs is between 14\\% and 86\\%. But HST is probably biased towards dense interacting HCGs. Still, there could be interacting pairs showing galaxies that have been blended even with the CCD images of \\cite{HKA89}. So, with the high resolution of the HST, the fraction of \\emph{pHCG}s and \\emph{vHCG}s with blended pairs may be considerably higher than 10\\% and in agreement with the fractions found in the mock CGs.  We can also estimate the number of HCGs that are physically dense groups of at least 4 galaxies. As discussed in Sect.~\\ref{hcgsample}, the HCG catalogue has 100 members, among which 99 are compact groups (since HCG~54 is a collection of \\hii regions), of which only 83 actually fulfil the original magnitude concordance criterion ($R$-band magnitude range less than 3). Among the 72 \\emph{pHCG}s whose brightest and faintest galaxies are brighter than $R=14.44$ and $R=17.44$, respectively, only 52 (72\\%) have at least 4 accordant velocities, and among these, we expect roughly between 36 and 44 HCGs that are physically dense groups of at least 4 galaxies. Extrapolating to the 68 accordant-velocity HCGs (including those with magnitude range greater than 3 mags), \\emph{we expect no more than $\\sim$ 58 physically dense HCGs with at least 4 galaxies}.  In comparison, \\cite{Mamon86} had predicted that 47 out of what he thought would be 78 accordant velocity HCGs are caused by chance alignments (60\\%), while the remaining 40\\% are physically dense (but he predicted that half of these dense groups were unbound systems). We are therefore less pessimistic than \\cite{Mamon86} on the fraction of chance alignments polluting the HCG catalogue, since chance alignments appear to represent between 14\\% and 30\\% of the \\emph{mvHCG}s (we were not able to check what fraction of the Real ones are unbound: see Appendix~\\ref{appbind}). Part of this discrepancy is caused by Mamon's (1986) reliance on simulations without consideration of selection effects such as observers blending close projected pairs of galaxies. Still, the percentage of chance alignments in the particle mock velocity-filtered compact groups (\\emph{pmvCG}s) is only 28--47\\% (depending on the SAM). Nevertheless, given the wide range of chance alignments fractions among the three SAMs, one cannot rule out that a more realistic galaxy formation model would lead to as much as 60\\% of chance alignments.   \\subsection{Comparison with McConnachie et al.} \\label{compare}  \\cite{McCEP08} (MEP) have published a study very similar to ours: they also extracted \\emph{pmpCG}s from the DLB model obtained from the Millennium dark matter simulations. Their sample extended to $r = 18$, which corresponds to roughly one-quarter of a magnitude fainter than our limit of $R < 17.44$. Their other \\emph{pmpCG} criteria appear to be almost exactly the same as ours (following the criteria of \\citealp{Hickson82}), although their algorithm works differently \\citep{McCPES09}. MEP found a total of over 15\\,000 \\emph{pmpCG}s over $4\\pi\\,\\rm sr$.  Using precisely the same input galaxy catalogue (from  \\citealp{Blaizot+05}) as MEP, we find 25\\,000 \\emph{pmpCG}s, so, our algorithm is nearly 1.6 times more efficient than MEP's in finding \\emph{pmpCG}s. Surprisingly, if we build a light cone as we did in Sect.~\\ref{testing}, with apparent magnitude limit of $R=17.67$ ($\\simeq r_{\\rm SDSS}=18$), we obtain a mock galaxy catalogue that is 3 times denser than the \\citeauthor{Blaizot+05} mock galaxy  catalogue used by MEP. From our mock galaxy catalogue, we extract 15\\,191 \\emph{pmpCG}s in 1.2693 sr, which means that, with the data used in this work, we are $\\sim 10$ times more efficient than MEP in finding CGs, principally by differences in the parent samples of galaxies, but also by differences in the CG detection algorithm.   MEP and us agree that a significant fraction of \\emph{mpCG}s are caused by chance alignments: MEP found 71\\% of their \\emph{pmpCG}s are CAs while we find 80\\% (with our hybrid classification, see Table~\\ref{nattab}).  There are, however, several important differences between our two studies: \\begin{itemize} \\item MEP build their sample from a mock that extends beyond the box size of the Millennium simulation --- using the output of the Mock Map Facility (MoMaF) code of \\cite{Blaizot+05}, while our mock galaxy catalogues are limited to the size of the simulation box. However, as shown in Sect.~\\ref{testing}, working on a light cone or working on a single simulation box leads to similar numbers of mock CGs (we identify a factor of 1.2 fewer \\emph{mpCG}s, and 1.5 more \\emph{mvCG}s). \\item MEP consider CGs with a faint magnitude limit, while we also tie in a bright magnitude limit to ensure that all mock CGs were built from galaxies that spanned a range of over 3 magnitudes. \\item We have analysed the galaxies from 3 different SAMs, while MEP have only considered the DLB sample, which we found to produce the smallest fraction of physically dense \\emph{mpCG}s. \\item MEP only provide statistics for the mock CGs defined in projection (\\emph{pmpCG}s, which they refer to as `HA's) but do not consider the subset of accordant-velocity groups (\\emph{pmvCG}s). We think this would have been worthwhile because ever since \\cite{HMdOHP92} published the HCG galaxy redshifts, most analyses have thrown out the discordant velocity HCGs. \\item MEP did not consider selection effects, while we considered both the galaxy confusion from close, blended, projected pairs, as well as the biases that we determined for the Hickson's visual selection of the HCGs.  \\item MEP only considered those \\emph{mpCG}s with $k\\geq 3$  galaxies that lie very close in real space, while we considered $k \\ge 4$ (to be consistent with Hickson's initial motivation to have at least 4 galaxies per HCG). \\item MEP define the Real \\emph{mpCG}s using a Friends-of-Friends linking length in real space, while we use a maximum real-space separation and the elongation along the line-of-sight. Structures built from small numbers of components with Friends-of-Friends algorithms tend to be more filamentary (e.g. \\citealp*{MFW93}). For \\emph{mpCG}s that are CALGs or CAFs, the most distant outlier will determine a similar maximum length and critical linking length. However, for \\emph{mpCG}s without outliers (e.g. Real \\emph{mpCG}s), the linking length will be smaller than the 3D length. In other words, selecting Real groups with a linking length of $200 \\, h^{-1} \\, \\rm kpc$ will result in group 3D lengths considerably greater. Moreover, for \\emph{mpCG}s with both foreground and background galaxies, MEP's $\\ell$ must be compared to our \\emph{half}-maximum size $s_{\\rm cut}/2$, and there is here a discrepancy of a factor two. Worse, for those (admittedly rare) cases of, say 4, galaxies aligned along the line of sight at roughly equal separations just below $\\ell$, one will end up with a group that spans up to $3\\,\\ell = 600 \\, h^{-1} \\, \\rm kpc$, which is now three times our maximum 3D length, but will still be called \\emph{Compact Association} (Real) by MEP, although it clearly is a chance alignment. In summary, this point and the previous one imply that MEP's criterion for calling an \\emph{mpCG} Real is much more liberal than ours.  \\end{itemize}   \\subsection{Perspectives} In forthcoming papers, we will analyse the distribution of and correlations between the physical characteristics of the \\emph{mvCG}s, and show how they depend on their classification, in view of optimising the probability that a CG selected in redshift-space is physically dense. It would be worthwhile to probe the formation of the physically dense CGs by analysing the merger trees of galaxies in the mock CGs. Finally, the analysis presented here will need to be confirmed with increasingly realistic simulations of galaxy catalogues, for example constructed from future galaxy formation models run on the recent high resolution Millennium-II dark matter simulation \\citep{BoylanKolchin+09}, and also on future high-resolution cosmological hydrodynamical simulations, with realistic prescriptions for feedback from AGN and supernovae."
    },
    "0809/0809.4606_arXiv.txt": {
        "abstract": "We study the generation of vorticity and velocity dispersion by orbit crossing using cosmological numerical simulations, and calculate the backreaction of these effects on the evolution of large-scale density and velocity divergence power spectra. We use Delaunay tessellations to define the velocity field, showing that the power spectra of velocity  divergence and vorticity measured in this way are unbiased and have better noise properties than for standard interpolation methods that deal with mass weighted velocities. We show that high resolution simulations are required  to recover the correct large-scale vorticity power spectrum, while poor resolution can spuriously amplify its amplitude by more than one order of magnitude. We measure the scalar and vector modes of the stress tensor induced by orbit crossing using an adaptive technique, showing that its vector modes lead, when input into the vorticity evolution equation, to the same vorticity power spectrum obtained from the Delaunay method. We incorporate orbit crossing corrections to the evolution of large scale density and velocity fields in perturbation theory by using the measured stress tensor modes. We find that at large scales ($k\\simeq 0.1 \\kMpc$) vector modes have very little effect  in the density power spectrum, while scalar modes (velocity dispersion) can induce percent level corrections at $z=0$, particularly in the velocity divergence power spectrum.  In addition, we show that the velocity power spectrum is smaller than predicted by linear theory until well into the nonlinear regime, with little contribution from virial velocities. ",
        "introduction": "The evolution of cosmological perturbations is determined, at scales larger than those where baryonic physics becomes important, by the gravitational clustering of cold dark matter, which can be taken as collisionless to a very good approximation.   Therefore, in this regime the Vlasov equation, i.e. the collisionless limit of the Boltzmann equation, describes the dynamics of cosmological perturbations~\\cite{1980lssu.book.....P}.  At large scales, where orbit crossing may be neglected, the Vlasov equation reduces to the dynamics of a pressureless perfect fluid (hereafter PPF). The PPF approximation has been used extensively in analytic approaches such as standard perturbation theory (hereafter PT; see~\\cite{2002PhR...367....1B} for a review) and the more recent renormalized perturbation theory~\\cite{2006PhRvD..73f3519C} (hereafter RPT) and related techniques~\\cite{2007PhRvD..75d3514M,2007A&A...465..725V,2007JCAP...06..026M,2008ApJ...674..617T,2007arXiv0711.2521M,2008arXiv0806.0971P}.  At small scales, as the first nonlinear structures are formed, orbit crossing generates a nontrivial stress tensor (the second cumulant of the phase space distribution function), which leads to velocity dispersion and vorticity in the dark matter distribution. Both of these effects would not be significantly present otherwise; vorticity corresponds to vector modes which are not produced primordially (at least in the simplest models of inflation) and even if they were they decay due to the expansion of the universe, velocity dispersion does get generated primordially e.g. during thermal equilibrium of dark matter in the early universe but for cold dark matter candidates typical values are vanishingly small ($\\sim 10^{-6}$ km/s for WIMPs and at most $\\sim10^{-10}$ km/s for axions~\\cite{1997PhRvD..56.1863S}) compared to typical velocity flows generated during structure formation.  Although a number of works have attempted to go beyond the PPF approximation~\\cite{1993ApJ...416L..71W,1998A&A...335..395B,2000PhRvD..62j3501D,2001A&A...379....8V,2002MNRAS.334..435D,2003ApJ...583L...1S,2005A&A...438..443B}, there has been no quantitative estimate in the literature at what scale corrections to the PPF become important. This is the main goal of this paper.  In order to achieve this, one has to compare PPF solutions with solutions to the Vlasov equation. N-body simulations of collisionless cold dark matter attempt to solve the latter by discretizing the distribution function using particles that follow the characteristics of the Vlasov equation~\\footnote{For direct solutions of the Vlasov equation see~\\cite{2005MNRAS.359..123A} are references therein.}. The N-body solution will differ from the PPF in regions where particle orbits cross, also known as ``caustics\" or ``shell-crossing\" in the context of the spherical collapse. This generates a nontrivial stress tensor and higher-order cumulants of the distribution function in the dark matter. We use the N-body solution to measure the induced stress tensor generated by orbit crossing and calculate from it the corrections to the PPF predictions for the density and velocity divergence power spectra at large scales. Recent work on orbit crossing has concentrated on enhancement of dark matter annihilation in caustics~\\cite{2006MNRAS.366.1217M,2008MNRAS.385..236V,2008arXiv0809.0497W}. We are instead interested on the impact of orbit crossing on the large-scale dynamics, outside dark matter halos.   Along the way we provide a number of results regarding the nonlinear evolution of peculiar velocities, which compared to the density field has not been studied in as much depth. The generation of vorticity and velocity dispersion impacts the reconstruction of primordial fluctuations from peculiar velocities~\\cite{1989ApJ...336L...5B,1993ApJ...412....1D} which assume a cold (single-stream) potential flow, as do other reconstruction methods based on galaxy positions~\\cite{1997MNRAS.285..793C,2006MNRAS.365..939M,2007ApJ...664..675E}. The understanding of the nonlinear evolution of the {\\em volume weighted} (as opposed to mass or galaxy weighted) velocity field is important, since the velocity difference PDF is one of the building blocks that contributes to the redshift space galaxy power spectrum, independent of galaxy bias and calculable from first principles (\\cite{2004PhRvD..70h3007S}, see also~\\cite{2008arXiv0808.0003P} for recent discussion). The study of the peculiar velocity field has been recently highlighted as a means to constrain dark energy and large-distance modifications of gravity~\\cite{2007PhRvD..76j3523F,2008arXiv0802.3897S,2008arXiv0807.0810S}. Finally, as old and new techniques to measure the peculiar velocity power spectrum improve, some of the issues we study here should be important for making predictions that model nonlinear effects accurately for future observations~\\cite{1995PhR...261..271S,2007MNRAS.379..343W,2008JCAP...02..022H,2007PhRvL..99h1301G,2008MNRAS.388..884Z,2008arXiv0802.1935A}.  This paper is organized as follows. In section~\\ref{secdelresults}, we present results for the power spectrum of velocity divergence and vorticity that follow from applying the Delaunay method to our N-body simulations. We discuss how vorticity and velocity dispersion get generated by orbit crossing in section~\\ref{orbitcross}, where we also propose an estimator of the stress tensor induced by orbit crossing based on an adaptive method.  In section~\\ref{PPFcorrect} we extend PT to include vorticity and velocity dispersion. Finally we present our conclusions in section~\\ref{conclude}.   \\begin{figure} \\begin{center} \\includegraphics[width=75mm]{RPTvsDel.eps} \\end{center} \\caption{Velocity divergence power spectrum at $z=0$ from 50 realizations of the LR1280 simulations. The symbols with error bars denote the velocity divergence power spectrum measured with the Delaunay method, normalized as in Eq.~(\\ref{norm}). The  blue solid line is the RPT prediction, and the dotted line is the linear power  spectrum.} \\label{rptvsdel} \\end{figure}  \\begin{figure*} \\begin{center} \\hspace*{-20mm} \\includegraphics[width=200mm]{vortilus_hor.eps} \\end{center} \\vspace*{-10mm} \\caption{Illustration of overdensity, divergence and vorticity in a  $1 \\Mpc$ thick cross-section of the simulation box at $z=0$. The divergence and   vorticity components panels correspond to the dimensionless quantities  $\\nabla\\cdot\\uv/{\\cal H}f$ and $\\nabla\\times\\uv/{\\cal H}f$. The panel labeled ``Halo particle overdensity'' shows the overdensity  of particles belonging to dark matter halos with mass $m \\protect\\ga 2.4 \\times 10^{10} h^{-1}M_\\odot$.} \\label{vortilus} \\end{figure*} ",
        "conclusions": "\\label{conclude}  We studied the impact of orbit crossing in the large-scale power spectra of density and velocity divergence fields, which are usually described in the pressureless perfect fluid (PPF) approximation. We presented a method to  extend perturbation theory (PT) beyond the PPF approximation, based on measuring the stress tensor induced by orbit crossing in numerical simulations. The stress tensor, when decomposed into scalar and vector modes leads to corrections associated with velocity dispersion and the effects of vorticity. We found the effects due to the scalar modes to be small, but not negligible at large scales ($k\\simeq 0.1 \\kMpc$), particularly for the velocity divergence power spectrum (see Fig.~\\ref{dispcorrfig}). The impact of vorticity on large scales is much smaller, see Fig.~\\ref{vort_corr}. These two effects appear at different orders in PT and have been included separately as we are interested in large scales where the induced corrections are small. Both lead to suppressions of the power spectra predicted by the PPF approximation, as expected physically since velocity dispersion and vorticity should inhibit collapse. In this regard we emphasize that neglecting orbit crossing has {\\em opposite} effects on Eulerian compared to Lagrangian PT. For Lagrangian PT, neglecting orbit crossing leads to (much more severe) {\\em underestimates} of the density power spectrum (see e.g.~\\cite{2007arXiv0711.2521M} for a recent example), since neglecting self-gravity in caustics leads to artificial thickening of such structures when trajectories cross without interacting.  A novel aspect of our calculation is the estimation of the stress tensor and the vorticity and divergence power spectra from numerical simulations. To estimate velocity fields, we applied the Delaunay tessellation method, which we have shown to be a more reliable estimator than traditional mass weighting schemes. While estimates of the velocity divergence are robust, we found that measurements of the vorticity power spectrum are significantly more difficult, due to aliasing during the measurement process and most importantly lack of resolution in the simulations. For the latter we have found that low resolution simulations can overestimate the vorticity power spectrum by an order of magnitude. This maybe be due to insufficient spatial resolution in multistreaming regions, with the overestimate perhaps related to aliasing effects during the PM part of the force calculation, which may generate a vector mode. In any event, for high enough resolution we find that the vorticity power spectrum converges to a stable answer. On the other hand, care must be taken that these spurious effects are not present when using numerical simulations to study nonlinear velocities, since artificial vorticity can amplify the velocity power spectrum at small scales.  A nontrivial check of our numerical calculation of the stress tensor, which we have done using an adaptive method independent of the Delaunay tesselation, is that its vector modes source the growth of vorticity. Therefore, using linear PT from this vector source one should recover at large scales the vorticity power spectrum measured by the Delaunay method, as we do (see Fig.~\\ref{vort_disp}). This does not test the scalar mode of the stress tensor though, which ends up inducing the largest correction to the PPF approximation. In this respect, it would be interesting to test how robust the scalar part of the stress tensor is to details of the numerical simulations, as spurious effects due to discreteness may amplify velocity dispersion in simulations~\\cite{1997ApJ...479L..79M,1998ApJ...497...38S} (see also~\\cite{2008arXiv0805.1357J}). As far as we know, our work is the first to make a quantitative connection between the growth of velocity dispersion and that of the density power spectrum, which will be useful to probe more in order to make sure that simulations can correctly reproduce the matter power spectrum to percent level, as required for the next generation of weak lensing surveys designed to probe cosmic acceleration~\\cite{2005APh....23..369H}.   The deviations we found from the PPF approximation at large scales are small but not negligible, in particular for the velocity divergence power spectrum, for which corrections are a factor of about three larger than for the density power spectrum. Our estimate, being based on numerical simulations, corresponds to fixed cosmological parameters (e.g. $\\sigma_8=0.9$, $\\Omega_m=0.27$ and $n_s=1$). Given the strong dependence on the growth factor of the correction ($\\propto D_+^{2.25}$ relative to PPF) we expect it to be smaller for lower normalization amplitudes, as well for cosmological parameters that correspond to less power at small scales (i.e. $\\Omega_m<0.27$ and $n_s<1$).  In section~\\ref{orbitcross} we sketched what must be done to include these effects from first principles into analytic calculations such as RPT that usually start from the PPF approximation, instead of using the hybrid approach we develop here partially based on numerical simulations.  Including velocity dispersion in a self-consistent manner should cure divergences that appear in PT for scale-free models with initial power spectra $P(k)\\sim k^n$ for $n>-1$, as well as regulate the divergences that appear in the resummation of the Lagrangian space propagator~\\cite{2008arXiv0805.0805B}. In addition, we expect velocity dispersion and vorticity to be crucial to describe the virial turnover in the density power spectrum.  Another interesting application of the ideas presented in section~\\ref{orbitcross} is to use the cumulant hierarchy to describe nonlinear effects in a massive neutrino component, to improve on recent calculations~\\cite{2008PhRvL.100s1301S,2008arXiv0809.0693W} that assume linearity. We hope to report on some of this in the near future."
    },
    "0809/0809.3518_arXiv.txt": {
        "abstract": "We discuss the non-adiabatic or entropy perturbation, which controls the evolution of the curvature perturbation in the uniform density gauge, for a scalar field system minimally coupled to gravity with non-canonical action. We highlight the differences between the sound and the phase speed in these systems, and show that the non-adiabatic pressure perturbation vanishes in the single field case, resulting in the conservation of the curvature perturbation on large scales. ",
        "introduction": "\\label{sec:thermo}  Considerable physical insight can be gained by studying the thermodynamic properties of a system. In order to keep the discussion as simple as possible, we focus here on a single fluid system. For an analysis including an arbitrary number of interacting fluids and canonical fields see e.g.~Ref.~\\cite{MW2004}.    \\subsection{Entropy and non-adiabatic pressure} \\label{sec:entropy}  A general thermodynamic system is fully characterised by three variables, of which only two are independent. Here we choose the energy density $\\rho$, and the entropy $S$, as independent variables, with the pressure $P$ being $P\\equiv P(\\rho,S)$. The pressure perturbation can then be expanded into a Taylor series as \\be \\label{eq:dP} \\dP = \\frac{\\p P}{\\p S}\\delta S + \\frac{\\p P}{\\p \\rho}\\drho \\,. \\ee This can be cast in the more familiar form \\bea \\label{defdPnad} \\dP=\\dPn+\\cs\\delta\\rho\\,, \\eea by introducing the adiabatic sound speed \\bea \\label{defcs2} \\cs\\equiv\\left.\\frac{\\p P}{\\p \\rho}\\right|_{S}\\,, \\eea and by identifying the non-adiabatic pressure perturbation as $\\dPn\\equiv\\left.\\frac{\\p P}{\\p S}\\right|_{\\rho}\\delta S$ \\cite{WMLL}. Equations (\\ref{eq:dP}) and (\\ref{defdPnad}) provide an intuitive and convenient definition for the adiabatic and the entropy, or non-adiabatic, perturbations in our system: the single adiabatic degree of freedom is proportional to the energy density, all other degrees of freedom will contribute to the non-adiabatic perturbation. This definition can easily be extended to more than two degrees of freedom as we shall show below, and can also be extended to higher order in the perturbations, as we shall briefly demonstrate at the end of this section.   With the above definitions, we see from \\eq{defdPnad} that the adiabatic sound speed $\\cs$ is of zeroth order in the perturbations, and since all background quantities depend only on time we get from \\eq{defcs2} \\be \\label{cs2back} \\cs=\\frac{\\dot P_0}{\\dot\\rho_0}\\,, \\ee where subscript ``0'' denotes a background quantity. Together with the behaviour of $\\delta P$ and $\\delta\\rho$ under a gauge-transformation $\\tilde t= t-\\delta t$, namely $\\wt{\\delta\\rho} =\\delta\\rho+\\dot\\rho_0\\delta t$ and $\\wt{\\delta P} =\\delta P+\\dot P_0\\delta t$, see e.g.~Ref.~\\cite{MM2008}, only the definition for the sound speed $\\cs$ given in \\eq{cs2back} renders the non-adiabatic pressure perturbation $\\dPn$ gauge-invariant.   We now focus on a system with two degrees of freedom. In order to obtain the non-adiabatic pressure perturbation in terms of different variables, we simply change from $\\rho$ and $P$ to the new variables denoted here $\\vp$ and $X$, which we shall identify in the following section with the amplitude of the scalar field and its kinetic term, respectively. The energy density and the pressure are then functions of $\\vp$ and $X$, that is $P=P(\\vp,X)$ and $\\rho=\\rho(\\vp,X)$, and we can therefore write \\bea \\label{eq:dP1} \\dP = \\frac{\\p P}{\\p \\vp}\\dvp +\\frac{\\p P}{\\p X}\\delta X \\,, \\eea and \\bea \\label{eq:drho} \\drho = \\frac{\\p \\rho}{\\p \\vp}\\dvp+\\frac{\\p \\rho}{\\p X}\\delta X \\,. \\eea Substituting \\eqs{eq:dP1} and (\\ref{eq:drho}) into \\eq{eq:dP} we obtain \\be \\label{eq:nad} \\delta P_{\\rm{nad}} =\\rho_{,\\vp} \\left(\\frac{ P_{,\\vp}}{\\rho_{,\\vp}}-\\cs\\right)\\dvp{} +\\rho_{,X}\\left(\\frac{P_{,X}}{\\rho_{,X}}-\\cs\\right)\\delta X\\,, \\ee where, for example, $P_{,\\vp}\\equiv\\p P/\\p \\vp$. Note that in \\eq{eq:nad} the terms in brackets are of zeroth order, and hence $\\delta P_{\\rm{nad}}$ can be evaluated once $P(\\vp,X)$ has been specified using background quantities only, which we shall do below in Section \\ref{sec:conditions}.  Equation (\\ref{eq:nad}) can readily be extended to more than two degrees of freedom.  For example, if the energy density and the pressure are functions of $N$ fields $\\vp_I$ and $X$, that is $P=P(\\vp_I,X)$ and $\\rho=\\rho(\\vp_I,X)$, then \\be \\label{eq:nad:mult} \\delta P_{\\rm{nad}} =\\sum_K\\rho_{,{\\vp_K}} \\left(\\frac{ P_{,{\\vp_K}}}{\\rho_{,{\\vp_K}}}-\\cs\\right)\\dvpK{} +\\rho_{,X}\\left(\\frac{P_{,X}}{\\rho_{,X}}-\\cs\\right)\\delta X\\,. \\ee    To extend Eq.~(\\ref{eq:dP}) to higher order we simply do not truncate the Taylor expansion at the linear order, that is \\bea \\label{eq:dP2} \\dP = \\frac{\\p P}{\\p S}\\delta S + \\frac{\\p P}{\\p \\rho}\\drho  +\\frac{1}{2}\\Bigg[ \\frac{\\p^2 P}{\\p S^2}\\delta S^2 +\\frac{\\p^2 P}{\\p \\rho\\p S}\\drho \\delta S +\\frac{\\p^2 P}{\\p \\rho^2}\\drho^2\\Bigg]+\\ldots\\,. \\eea The entropy, or non-adiabatic pressure perturbation at second order, for example, is then found from \\eq{eq:dP2}, and expanding $\\delta P\\equiv\\delta P_1+\\frac{1}{2}\\delta P_2$, and similarly $\\delta\\rho$, we get \\cite{Malik:2003mv} \\be \\delta P_{2\\rm{nad}} =\\delta P_2-\\cs\\delta\\rho_1-\\frac{\\p \\cs}{\\p\\rho}\\ \\delta\\rho_1^2\\,. \\ee After this higher order excursion, we return to linear theory below.     \\subsection{Sound speed and phase speed} \\label{sec:soundphase}   Many oscillating systems can be described by a wave equation, that is a second order evolution equation of the form \\be \\label{eq:wave_equ} \\frac{1}{\\cph}\\ddot\\phi-\\nabla^2\\phi+F(\\phi,\\dot\\phi)=0\\,, \\ee where $\\phi$ is, for example, the velocity potential or the scalar field amplitude, $F(\\phi,\\dot\\phi)$ is the damping term, and $\\cph$ is the phase speed, i.e.~the speed with which perturbations travel through the system \\cite{arnold}.   There is some confusion in the literature on the meaning of adiabatic sound speed and phase speed. Although this seems not to affect the results of previous works, and the adiabatic sound speed is often simply used as a convenient shorthand for $\\dot{\\Pb}/\\dot{\\rhob}$, as defined in \\eq{cs2back}, it is often confusingly used to denote the phase speed. We take the opportunity to clarify these issues here.   The adiabatic sound speed defined above describes the response of the pressure to a change in density at constant entropy \\footnote{More intuitive is the introduction of the adiabatic sound speed using the compressibility, $\\beta\\equiv\\frac{1}{\\rho}\\frac{\\p\\rho}{\\p P}\\Big|_{S} =\\frac{1}{\\rho\\cs}$, which describes the change in density due to a change in pressure while keeping the entropy constant.}, and is the speed with which pressure perturbations travels through a classical fluid. A collection of scalar fields can also be described as fluid, but the analogy is not exact. Whereas in the fluid case phase speed $\\cph$ and adiabatic sound speed $\\cs$ are equal, this is not true in the scalar field case and the speed with which perturbations travel is given \\emph{only} by $\\cph$. We shall illustrate the difference in $\\cph$ and $\\cs$ for a concrete example in Section \\ref{sec:speeds} below.    \\section {Governing equations} \\label{sec:equations}  We now derive the governing equations for a scalar field system with general non-canonical action. Consider a general Lagrangian for a single scalar field $\\vp$, \\be \\label{eq:lagrangian} \\mathcal{L}=f(X,\\vp)\\,, \\ee where $X=-\\frac{1}{2}g^{\\mu\\nu}\\vp_{,\\mu}\\vp_{,\\nu}$.   The energy-momentum tensor is defined as \\cite{LLBook} \\be \\label{eq:energymomentum} T_{\\mu\\nu}=-2\\frac{ \\p \\mathcal{L}} {\\p g^{\\mu\\nu}}+g_{\\mu\\nu}\\mathcal{L} \\,, \\ee and substituting the Lagrangian Eq.~(\\ref{eq:lagrangian}) into Eq.~(\\ref{eq:energymomentum}) we get \\be \\label{eq:EM2} {T^{\\mu}}_{\\nu} =f_{,X}g^{\\mu\\lambda}\\vp_{,\\lambda}\\vp_{,\\nu}+{\\delta^\\mu}_\\nu f\\,. \\ee The energy-momentum tensor for a perfect fluid is \\be \\label{eq:EMfluid} {T^\\mu}_\\nu=(\\rho+P)u^\\mu u_\\nu+P{\\delta^\\mu}_\\nu\\,, \\ee where $P$ and $\\rho$ are the pressure and energy density, respectively, and $u^\\mu$ is the four-velocity, subject to the constraint $u^\\mu u_\\mu=-1$.  We consider here only scalar perturbations in a flat FRW background with line element \\cite{Bardeen80,KS,MFB} \\bea \\label{eq:metric} ds^2=-(1+2A)dt^2+2aB_{,i}dx^idt +a^2\\left[(1-2\\psi)\\dij+2E_{,ij}\\right]dx^idx^j\\,, \\eea where $a=a(t)$ is the scale factor, $A$ is the lapse function, $B$ and $E$ make up the shear ($\\sigma=a^2\\dot{E}-aB$), and $\\psi$ is the dimensionless curvature perturbation. The four-velocity is then given by \\be \\label{eq:4velocity} u^{\\mu}=\\left[(1-A),\\frac{1}{a} v_,^i\\right]\\,, \\ee where $v$ is the scalar velocity potential of the fluid.    We split quantities into a time-dependent background and a time- and space-dependent perturbation, and get e.g.~for the scalar field $\\vp=\\vpb(t)+\\dvp(t,x^i)$. The energy density and pressure of the scalar field are then, in the background \\be \\label{rho_back} \\rhob=2f_{,X}\\Xb-f\\,, \\qquad \\Pb=f\\,, \\ee where $\\Xb=\\frac{1}{2}\\dot{\\vpb}^2$, and to first order in the perturbations \\bea \\label{eq:perturbed} \\drho&=&\\left(f_{,X}+2\\Xb f_{,XX}\\right)\\delta X +\\left(2\\Xb f_{,X\\vp}-f_{,\\vp}\\right)\\dvp{}\\,, \\nonumber \\\\ \\dP&=&f_{,X}\\delta X+f_{,\\vp}\\dvp{}\\,, \\eea where we have defined $\\dX=\\dot{\\vpb}\\dot{\\dvp}-A\\dot{\\vpb}^2$. It is convenient to work with the covariant velocity perturbation, $V=a(v+B)$, which in terms of the scalar field is given from ${T^i}_0$ component of the energy-momentum tensor as \\be \\label{Vrel} V=-\\frac{\\dvp}{\\dot{\\vpb}}\\,. \\ee    \\subsection{Background} \\label{sec:background}  The conservation of energy-momentum is ${T^\\mu}_{\\nu;\\mu}=0$, which in the background reduces to \\be \\label{eq:bkgd} \\dot{\\rhob}=-3H(\\rhob+\\Pb)\\,. \\ee In terms of the scalar field, this becomes the evolution equation \\be \\label{eq:bkgdfield} \\ddot{\\vpb}\\left(f_{,XX}\\dot{\\vpb}^2+f_{,X}\\right)+f_{,X\\vp}\\dot{\\vpb}^2 -f_{,\\vp}+3Hf_{,X}\\dot{\\vpb}=0 \\,. \\ee The Friedmann constraint equation is given by \\be H^2=\\frac{8\\pi G}{3}\\rhob\\,, \\ee where $H\\equiv\\dot a/a$ is the Hubble parameter.  \\subsection{First order} \\label{sec:firstorder}  Conservation of energy-momentum to linear order gives the perturbed energy conservation equation \\be \\label{eq:fluidevolution} \\dot{\\drho}=-3H(\\drho+\\dP)+(\\rho_0+P_0)\\left[3\\dot{\\psi} -\\frac{\\nabla^2}{a^2}(\\sigma+V)\\right]  \\,. \\ee This can be readily rewritten in terms of the curvature perturbation on uniform density hypersurfaces $\\zeta\\equiv-\\psi- H\\delta\\rho/\\dot\\rho$, and working in the uniform density gauge, where $\\delta\\rho\\equiv0$ and hence $\\dP=\\dPn$ and $\\zeta\\equiv -\\psi$, we get \\cite{WMLL} \\be \\label{eq:zetadot} \\dot{\\zeta}=\\frac{-H}{(\\rhob+\\Pb)}\\dPn-\\frac{\\nabla^2}{3a^2}(\\sigma+V) \\,. \\ee Equation (\\ref{eq:zetadot}) reduces on large scales, where gradient terms can be neglected, to \\eq{eq:zetadotlarge} and shows that if the pressure perturbation is adiabatic, then the curvature perturbation is conserved. See Ref.~\\cite{WMLL} for a more detailed discussion.    In the following we will also need the constraint equation (see e.g.~Ref.~\\cite{MW2004}), \\be \\label{deltarho_const} \\delta\\rho-3H(\\rhob+\\Pb)V+\\frac{H}{4\\pi G a^2}k^2\\sigma=0\\,, \\ee where we have replaced spatial Laplacians with the wave-numbers of their respective eigenvalues, that is $\\nabla^2\\rightarrow -k^2,$ and have chosen the flat gauge, where $\\psi=0=E$. On large scales, where gradient terms can be neglected, we can rewrite \\eq{deltarho_const} in terms of the scalar field, using \\eqs{eq:perturbed} and (\\ref{Vrel}), as \\be \\label{X_vp_rel} \\left(f_{,X}+\\dot{\\vpb}^2 f_{,XX}\\right)\\dX +\\left(\\dot{\\vpb}^2 f_{,X\\vp}-f_{,\\vp} +3Hf_{,X}\\dot{\\vpb}\\right)\\dvp{}=0 \\,, \\ee giving a convenient relation between $\\dX$ and $\\dvp$.   \\subsection {Speeds} \\label{sec:speeds}  The scalar field system we are considering here can be described by ``borrowing'' terminology from fluid dynamics, though the analogy is not exact. If we are interested in the speed with which perturbations travel through the system, we have to calculate the phase speed, which can be read off from the perturbed Klein-Gordon equation governing the evolution of the scalar field. This is just a damped wave equation, like \\eq{eq:wave_equ}, and so by comparing the coefficients of the second order temporal, and second order spatial derivatives, we obtain the phase speed.   For the general non-canonical Lagrangian, (\\ref{eq:lagrangian}), the evolution equation for the first order scalar field perturbation is found by substituting Eq.~(\\ref{eq:perturbed}) into (\\ref{eq:fluidevolution}), and is given here already in the form of Eq.~(\\ref{eq:wave_equ}) as \\be \\label{eq:pertevolution} \\frac{1}{\\cph}\\ddot{\\dvp}+\\Big[\\frac{3H}{\\cph}+C_1\\Big]\\dot{\\dvp} +\\Big[\\frac{k^2}{a^2}+C_2\\Big]\\dvp=0\\,, \\ee where the coefficients $C_1$ and $C_2$, both functions of $\\vp$ and $X$, are given in Eq.~(\\ref{eq:pertevolution2}) in the Appendix. We can therefore read off the phase speed as \\be \\label{eq:cphgeneral} \\cph=\\frac{f_{,X}}{f_{,X}+2\\Xb f_{,XX}}\\,. \\ee Using \\eq{rho_back}, this reduces to \\cite{Garriga:1999vw} \\be \\label{def_cph_final} \\cph=\\frac{P_{0,X}}{\\rho_{0,X}}\\,. \\ee The adiabatic sound speed for the Lagrangian (\\ref{eq:lagrangian}) is given by \\be \\label{eq:csgeneral} \\cs=\\frac{f_{,X}\\ddot{\\vpb}+f_{,\\vp}} {f_{,X}\\ddot{\\vpb}-f_{,\\vp}+f_{,XX}\\dot{\\vpb}^2\\ddot{\\vpb}+ f_{,X\\vp}\\dot{\\vpb}^2} \\,. \\ee  The Lagrangian for a canonical scalar field is obtained by setting $f(X,\\vp)=X-U(\\vp)$. In this case the adiabatic sound speed reduces to \\be \\label{eq:csstandard} \\cs=1+\\frac{2U_{,\\vp}}{3H\\dot{\\vpb}} \\,, \\ee and becomes $\\cs=-1$ in slow-roll. The phase speed for a canonical scalar field is $\\cph=1.$ ",
        "conclusions": "\\label{sec:conclusions}   We have studied whether the curvature perturbation in a system with a scalar field minimally coupled to gravity with non-canonical action is constant on super-horizon scales. Deriving the adiabatic and the non-adiabatic or entropy perturbation we found that imposing the strict super-horizon limit $k \\to 0$ is sufficient to guarantee conservation of $\\zeta$ and hence of the primordial power spectrum, just as in the canonical case.  The non-adiabatic pressure, or entropy, perturbation is in general non-zero, if there is more than one degree of freedom in the system, as is the case in multi-field systems. In this case the curvature perturbation evolves, even on large scales, and it is necessary to solve the evolution equations of the scalar fields to construct the non-adiabatic pressure. However, it is then more efficient to construct $\\zeta$ directly using its definition and substituting in the energy densities in terms of the fields, instead of first constructing $\\dPn$ and then solving the evolution equation for $\\zeta$, \\eq{eq:zetadot}, see \\cite{Gordon:2000hv,Langlois}.    We have also highlighted the difference between the phase speed and the adiabatic sound speed of a system in Section \\ref{sec:speeds}. Although these are the same in classical fluid systems, they are different in the scalar field models studied here, with only the phase speed describing the speed with which perturbations travel. We emphasise that only the definition for the adiabatic sound speed, \\eq{cs2back}, enters the gauge-transformation of the pressure perturbation, as pointed out in Section \\ref{sec:entropy}. Similarly, in the adiabatic case when $\\dPn=0$, \\eq{defdPnad} reduces to $\\dP=\\cs\\delta\\rho$. Again, this convenient relation between the pressure and the energy density perturbation is only correct when using the adiabatic sound speed, as defined in \\eq{cs2back}."
    },
    "0809/0809.1037_arXiv.txt": {
        "abstract": "We report the detection of the first two planets from the N2K Doppler planet search program at the Magellan telescopes. The first planet has a mass of $M \\sin i = 4.96 M_{Jup}$ and is orbiting the G3 IV star HD154672 with an orbital period of 163.9 days. The second planet is orbiting the F7 V star HD205739 with an orbital period of 279.8 days and has a mass of $M \\sin i = 1.37 M_{Jup}$. Both planets are in eccentric orbits, with eccentricities $e$ = 0.61 and $e$ = 0.27, respectively. Both stars are metal rich and appear to be chromospherically inactive, based on inspection of their Ca II H and K lines. Finally, the best Keplerian model fit to HD205739b shows a trend of 0.0649 m s$^{-1}$ day$^{-1}$, suggesting the presence of an additional outer body in that system. ",
        "introduction": "\\label{sec:intro} Ongoing Doppler radial velocity surveys of nearby stars have detected over 200 extrasolar planets in the past decade \\citep{Butler06}. Those surveys focus on late F, G, K, and M dwarfs within 50 pc and most of the planets they have found to date are more massive than Saturn, and are presumably gas giants. Recently, several Neptune-mass and lower mass planets have been detected, most of them with orbital periods of a few days \\citep{Butler04, McArthur04, Santos04, Bonfils05, Bonfils07, Rivera05, Udry06, Udry07, Endl08}  As the search for new extrasolar planets continues, Doppler surveys now look through a broad parameter space, including long-period Jupiter analogs, very low mass planets in short-period orbits, multiple planetary systems, and new planets around stars with spectral types that extend beyond the ones traditionally been searched, i.e.  K0V to F8V. The N2K program \\citep{Fischer05} is a Doppler survey with distributed observing campaigns at the Keck, Magellan and Subaru telescopes, and is primarily aimed at increasing the number of known hot Jupiters.  Because of their proximity to the host stars, the atmospheres of hot Jupiters can be as hot as 2000 K, resulting in detectable emission at IR wavelengths.  This makes these short period planets ideal targets for spaceborn follow up to observe exoplanet atmospheres \\citep{Harrington07}, especially when they transit their host stars and undergo a secondary eclipse so that small differential changes in emission from the star-planet system can be measured \\citep[e.g.][]{Charbonneau05, Knutson07}.  It is also important to have a sample of hot Jupiters that is large enough to provide meaningful constraints on the formation, migration, and evolution mechanisms of these planets \\citep{Ford06}.  The N2K program searches a fresh sample of the ``next 2000'' stars, not included in other current Doppler surveys, by extending the search radius to distances of 100--110 parsecs from the Sun. This project is intentionally biased towards metal rich stars to exploit the correlation between formation of gas giant planets and high stellar metallicity \\citep{Santos04, Fischer05b}. The search strategy is optimized for the detection of Jupiter-mass planets with orbital periods shorter than 14 days by obtaining radial velocity measurements on three consecutive nights. Variability in the velocity of the host star over this short timescale flags the star as a candidate host for a short-period planet. However, gas giant planets in longer orbits are also identified and detected with additional observations. As an example, the N2K programs at the Keck and Subaru telescopes have already detected seven planets with periods between 21 days and 3.13 years \\citep{Fischer07, Robinson07}.   Here, we report the detection of the first two planets from the Magellan N2K program. This program has been underway at the Magellan telescopes at Las Campanas Observatory in Chile since 2004, and includes $\\sim$ 300 FGK metal rich stars at distances between 50 and 100 parsecs.  The new planets are orbiting the stars HD 154672 and HD 205739. Both host stars showed significant radial velocity scatter in the first year of observations and were initially flagged as hot Jupiter candidates. However, follow-up observations over subsequent years revealed the presence of longer period planets and a trend on the radial velocity curve of HD 205739 that continues to increase after over 3.5 years. ",
        "conclusions": "\\label{sec:discussion}  We present two new Jovian--mass planets orbiting metal rich stars.  HD 154672b is a fairly massive planet with a mass of M$\\sin i$ = 4.96 $M_{Jup}$ and a very pronounced orbital eccentricity of $e$ = 0.61, that causes the planet to move from 0.23 to 0.96 AU between periastron and apastron.  The planet will therefore experience surface temperature changes of about 300K along its orbit, reaching a maximum temperature at periastron of about 600K, assuming the albedo of the planet is low. If water is present in the atmosphere of HD 154672b, it could transition between gaseous and liquid phases along the planet's orbit.  When placed in the eccentricity vs. orbital period parameter space diagram of known exoplanets illustrated in Figure~\\ref{fig:eccvsp}, HD 154672b shows an orbital eccentricity larger than 90\\% of the discovered planets, and is only the seventh planet found with an orbital period shorter than 300 days and an eccentricity larger than 0.6. Of the other six planets, four have been found to be either in multiple-planet systems \\citep[HD 74156;][]{Naef04, Bean08}, to have a brown dwarf companion \\citep[HD 3651;][]{Mugrauer06}, to be part of a wide stellar binary system \\citep[HD 80606;][]{Naef01}, or to present a large radial velocity trend induced by a distant body associated to that system \\citep[HD 37605;][]{Wittenmyer07}. The high eccentricies in those cases can be explained by either Kozai oscillations \\citep{Kozai62} or chaotic evolution of planetary orbits in multiple systems. Another planet in this subgroup, HD 17156, has been recently reported to have a large orbital axis misalignment with respect to the stellar rotation axis, that can be best explained by gravitational interactions with other planets \\citep{Narita07}. There is however no evidence for additional objets associated with the last planet in this subgroup, orbiting HD 89744.  The radial velocity curve of HD 154672 in Figure~\\ref{fig:154672} shows no trend nor significant residuals to the fit that indicate the presence of other massive objects in that system.   HD 205739b has a mass of M$\\sin i$ =1.37 $M_{Jup}$, with an average relative separation of 0.896 AU and an eccentricity of $e$ = 0.27. In the eccentricity vs. orbital period diagram in Figure ~\\ref{fig:eccvsp}, the parameters of this planet do not seem atypical. For planets with orbital periods longer than 20 days, the mean eccentricity is 0.29, therefore HD205739b has an orbital eccentricity that is typical of detected planets that have not experienced tidal circularization.  The separation of HD205739b from its host star changes from 0.65 AU to 1.14 AU between periastron and apastron. The maximum surface temperature of this planet is expected to be of the order of 400K, and the amplitude of its surface temperature change will only be of about 100K along the entire orbit. One peculiarity of the radial velocity curve of HD 205739 is the presence of a pronounced trend of 0.0649 m s$^{-1}$ per day, indicating the presence of an additional outer body in the system with an orbital period longer than the 3.5 year time span covered by our observations, and a radial velocity semi-amplitude greater than 35 m s$^{-1}$. Finally, the residuals of our best fit are a factor of two larger than expected, what hints the possible presence of other bodies in this system. However, more observations are needed to confirm this hypothesis."
    },
    "0809/0809.4338_arXiv.txt": {
        "abstract": "We present spectroscopic and photometric results of Nova V1280 Sco which was discovered in outburst in early 2007 February. The large number of spectra obtained of the object leads to one of the most extensive, near-infrared spectral studies of a classical nova. The spectra evolve from a P-Cygni phase to an emission-line phase and at a later stage is dominated by emission from the dust that formed in this nova. A detailed model is computed to identify and study characteristics of the spectral lines. Inferences from the model address the vexing question of which novae have the ability to form dust. It is demonstrated, and strikingly corroborated with  observations,  that the presence of lines in the early spectra of low-ionization species like Na and Mg - indicative of low temperature conditions - appear to be reliable indicators that dust will form in the ejecta. It is theoretically expected that mass loss  during a nova outburst is a sustained process. Spectroscopic evidence for such a sustained mass loss, obtained by tracing the evolution of a P-Cygni feature in the Brackett $\\gamma$ line, is presented here allowing a lower limit of 25-27 days to be set for the mass-loss duration. Photometric data recording the nova's extended 12 day climb to peak brightness after discovery is used to establish an early fireball expansion and also show that the ejection began well before maximum brightness. The JHK light curves indicate the nova  had a fairly strong second outburst $\\sim$ 100 days after the first. ",
        "introduction": "Nova V1280 Sco was discovered on 2007 February 4.8  independently by Sakurai and Nakamura  (2007), within  a short time of each other, with a reported visual magnitude  in the range  9.4 - 9.9.  The nova brightened quickly to its maximum in visual light on Feb. 16.19 ($m_{vmax}$= 3.79, Munari et al., 2007) to become one of the brightest novae in recent times ( Schmeer 2007); V1500 Cyg in 1975 had a $m_{vmax}$=2.0 and nova V382 Vel 1999 had $m_{vmax}$ =2.5. V1280 Sco has been observed in different wavelength regimes since its discovery. Early stage pre-maxima spectra was described by Munari et al. (2007) and Naito $\\&$ Narusawa (2007) in the visible region and by Rudy et al. (2007a,b) in the infrared. The early spectrum was dominated by absorption lines of hydrogen, neutral nitrogen and carbon, and displayed deep P-Cygni profiles. After passing the maxima the spectrum turned to that typically shown by a classical nova; the lines started appearing in emission which included strong emission lines of Ca II H and K, the hydrogen Balmer series, the Na I D doublet, and Fe II 42, 49 and 74 in visible region (Buil, 2007; Munari et al., 2007). Early near-IR spectra by  Das et al. (2007a) reported the presence of HI Brackett and Paschen lines and several strong C I lines. No X-ray detection was reported from the source from RXTE and SWIFT observations respectively (Swank et al. 2007; Osborne et al., 2007) .  The nova has an interesting lightcurve with  the possibility of more than one outburst. After passing the maxima, a  smooth and slow decline followed. This was interrupted, about 24 days after discovery, by the formation of dust which was evidenced from the sharp decline in visible light curve and a rise in the infrared continuum (Das et al., 2007b). The dust shell has been directly detected by interferometric techniques  and its subsequent spatial expansion has also been tracked (Chesneau et al. 2008). Apart from providing the first direct detection of a dust shell around a classical nova, the Chesneau et al. (2008) study emphasises the  important, emerging role of interferometry in the study of dust formation and dust properties in novae. We present here spectroscopic and photometric results based on  observations between 14 and 125 days after the discovery. A large number of spectra were recorded, sampling the nova's evolution at regular intervals, thereby leading to one of the most comprehensive studies of a  classical nova in the near-infrared. ",
        "conclusions": ""
    },
    "0809/0809.4879_arXiv.txt": {
        "abstract": "Dynamical models of prototype gravastars are constructed and studied. The models are the  Visser-Wiltshire three-layer gravastars, in which  an infinitely thin spherical shell of a perfect  fluid with the equation of state $p = (1-\\gamma)\\sigma$ divides the whole spacetime into two regions, where the internal region is  de Sitter, and  the external is  Schwarzschild. When $\\gamma < 1$ and $\\Lambda \\not= 0$, it is found that in some cases the models represent stable gravastars, and in some cases they represent ``bounded excursion\"  stable gravastars, where the thin shell is oscillating between two finite radii, while in some other cases they collapse until the formation of black holes. However, when $\\gamma \\ge 1$, even with $\\Lambda \\not= 0$, only black holes are found. In the phase space, the region for both stable gravastars and  ``bounded excursion\"  gravastars is very small in comparison to that of black holes, although it is not completely empty. ",
        "introduction": "\\renewcommand{\\theequation}{1.\\arabic{equation}} \\setcounter{equation}{0}   As alternatives to black holes,  gravastars have received some attention recently \\cite{grava}, partially due to the tight connection between the cosmological constant and a currently accelerating universe \\cite{DEs}, although very strict observational constraints on the existence of such stars may exist \\cite{BN07}. In the original model of Mazur and  Mottola (MM) \\cite{MM01}, gravastars consist of five layers: an internal core $0 < r < r_{1}$, described by the de Sitter universe, an intermediate thin layer of stiff fluid $r_{1} < r < r_{2}$, and an external region $ r > r_{2}$, described by the Schwarzschild solution, \\bq \\lb{1.1} ds^{2} = -f(r) dt^{2} + \\frac{1}{f(r)} dr^{2} + r^{2}\\left(d\\theta^{2} + \\sin^{2}\\theta d\\varphi^{2}\\right), \\eq where the function $f(r)$ is given by $f(r) =  1 - {2{\\cal{M}}}/{r}$, in the units where $c = 1 = G$. In addition, in such a setup,  two infinitely thin shells also  appear, respectively, on the hypersurfaces $r = r_{1}$ and $r = r_{2}$. By properly choosing the free parameters involved, one can show that the two shells can have only tensions but with opposite signs \\cite{MM01}. Visser and Wiltshire (VW) argued that such five-layer models can be simplified to three-layer ones \\cite{VW04}, in which the two infinitely thin shells and the intermediate region are replaced by one infinitely thin shell, so that the function $f(r)$ in the metric (\\ref{1.1}) is given by \\bq \\lb{1.3} f(r) = \\begin{cases} 1 - \\frac{2{\\cal{M}}}{r}, &  r > a(\\tau),\\\\ 1 - \\left(\\frac{r}{l}\\right)^{2}, &  r < a(\\tau), \\end{cases} \\eq where $r = a(\\tau)$ is a timelike hypersurface, at which the infinitely thin shell is located, and $\\tau$ denotes the proper time of the thin shell. The constant $l \\equiv \\sqrt{3/\\Lambda}$ denotes  the de Sitter radius. On the hypersurface $r = a(\\tau)$ Israel junction conditions yield \\bq \\lb{1.4} \\frac{1}{2}\\dot{a}^{2} + V(a) =  0, \\eq where an overdot denotes the derivative with respect to the proper time $\\tau$ of the thin shell. % Therefore, in the region $ r > a(\\tau)$ the spacetime is locally Schwarzschild, while in the region $ r < a(\\tau)$ it is locally de Sitter.  These two different regions are connected through a dynamical infinitely thin  shell located at $r = a(\\tau)$ to form a new spacetime of gravastar.   Two different types of stable gravastars are identified by VW, stable gravastars and ``bounded excursion\" gravastars.  {\\bf Stable gravastars}: In this case,  there must exist a radius $a_{0}$ such that \\bq \\lb{1.5} V\\left(a_{0}\\right) = 0, \\;\\;\\; V'\\left(a_{0}\\right) = 0, \\;\\;\\; V''\\left(a_{0}\\right) > 0, \\eq where a prime denotes the ordinary differentiation with respect to the indicated argument. If and only if there exists such an $a_{0}$ for which the above conditions are satisfied, the model is said to be stable. Among other things, VW found that there are many equations of state for which the gravastar configurations are stable, while others are not \\cite{VW04}. Carter studied  the same problem and found new equations of state for which the gravastar is stable \\cite{Carter05}, while De Benedictis {\\em et al} \\cite{DeB06} and Chirenti and Rezzolla \\cite{CR07} investigated the stability of the original model of  Mazur and  Mottola against axial-perturbations, and found that gravastars are stable to these perturbations. Chirenti and Rezzolla also showed that their quasi-normal modes differ from those of a black hole of the same mass, and thus can be used to discern a gravastar from a black hole. Early work on dynamical thin shells with de Sitter interior and (A)dS-RN exterior can be found in \\cite{ALT99}.  {\\bf ``Bounded excursion\" gravastars}: As VW noticed, there is a less stringent notion of stability, the so-called ``bounded excursion\" models, in which there exist two radii $a_{1}$ and $a_{2}$ such that \\bq \\lb{1.6} V\\left(a_{1}\\right) = 0, \\;\\;\\; V'\\left(a_{1}\\right) \\le 0, \\;\\;\\; V\\left(a_{2}\\right) = 0, \\;\\;\\; V'\\left(a_{2}\\right) \\ge 0, \\eq with $V(a) < 0$ for $a \\in \\left(a_{1}, a_{2}\\right)$, where $a_{2} > a_{1}$.  Lately, we studied this type of gravastars \\cite{JCAP}, and found that, among other things,  such configurations can indeed be constructed, although   the region for the formation of this type of gravastars is very small in comparison to that of black holes.  In this paper, we generalize our previous work \\cite{JCAP} to the case where the equation of state of the infinitely thin shell is given by $p= (1-\\gamma)  \\sigma$ with $\\gamma$ being a constant. When $\\gamma = 0$ it reduces to the case  we studied in \\cite{JCAP}. We shall first construct three-layer dynamical models, in analogy to the VW models, and then show both  stable gravastars of the both types and black holes exist for $\\gamma < 1$ and $\\Lambda \\not= 0$. However, when $\\gamma \\ge 1$ even  with $\\Lambda \\not= 0$, only black holes are found. In the phase space, the region of gravastars and the one of black holes  are non-zero, although the former is much smaller than the latter. The rest of the paper is  organized as follows: In Sec. II we shall study various cases, in which all the possibilities of forming black holes, gravastars, de Sitter, and Minkowski spacetime exist. In Sec. III we present our main conclusions. ",
        "conclusions": "In this paper, we have generalized our previous work on the problem of stable gravastars by constructing dynamical three-layer VW models  \\cite{VW04}, which consists of an internal de Sitter space, a dynamical infinitely thin shell of a perfect fluid with the equation of state $p = (1-\\gamma)\\sigma$, and an external Schwarzschild spacetime. We have shown explicitly that the final output can be a black hole, a  ``bounded excursion\"  gravastar, a stable gravastar, a Minkowski, or a de Sitter spacetime, depending on the total mass $m$ of the system,  the cosmological constant $\\Lambda$, and the initial position $R_{0}$ of the dynamical thin shell. All these possibilities have non-zero measurements in the phase space of $m, \\; \\Lambda (\\not= 0)$, $\\gamma (< 1)$ and $R_{0}$, although the region  of gravastars is very small in comparison with that of black holes. When $\\gamma \\ge 1$ even with $\\Lambda \\not= 0$, only black holes are found. Therefore, although the existence of gravastars cannot be completely excluded in these dynamical models, our results show that, even if gravastars indeed exist, they do not exclude the existence of black holes."
    },
    "0809/0809.4562_arXiv.txt": {
        "abstract": "Clouds seem like an every-day experience. But -- do we know how clouds form on brown dwarfs and extra-solar planets? How do they look like? Can we see them?  What are they composed of?  Cloud formation is an old-fashioned but still outstanding problem for the Earth atmosphere, and it has turned into a challenge for the modelling of brown dwarf and exo-planetary atmospheres.  Cloud formation imposes strong feedbacks on the atmospheric structure, not only due to the clouds own opacity, but also due to the depletion of the gas phase, possibly leaving behind a dynamic and still supersaturated atmosphere. I summarise the different approaches taken to model cloud formation in substellar atmospheres and workout their differences. Focusing on the phase-non-equilibrium approach to cloud formation, I demonstrate the inside we gain from detailed micro-physical modelling on for instance the material composition and grain size distribution inside the cloud layer on a Brown Dwarf atmosphere.  A comparison study on four different cloud approaches in Brown Dwarf atmosphere simulations demonstrates possible uncertainties in interpretation of observational data. ",
        "introduction": "Brown Dwarfs are object much more massive than planets, but they can become as cold as planets during their lifetime. Young gas giants and also so-called close-in giant gas planets can be as hot as Brown Dwarfs (e.g. \\cite{ca2008}, \\cite{heb2008}). Brown Dwarfs can only nucleosynthesis energy during a short period of their lift time, which is enough to bust their luminosity to a level where they can be much easier observed than any planet at a similar distance. The largest similarity between Brown Dwarfs and planets is the rich, multi-phase atmospheric chemistry.  Descending the main sequence makes Brown Dwarfs and possibly late-type M-dwarfs the only objects, that form like stars and that contain dust (i.e. small solid particles) in their atmosphere. After this discovery \\cite{Jones1997}, it became immediately clear that Brown Dwarfs were not the perfect example for a static atmosphere but feedback of the dust clouds on the atmospheric structure, the chemistry, and the radiative transfer would need attention. Other places of efficient dust formation are the circulstellar envelopes of AGB stars which are extremely dynamic systems due to the presence of dust (see review by A.C. Andersen, and contributions by e.g. S.H\\\"ofner, N. Thureau this volume). \\cite{cu2008}\\, have presented spectra for L-type Brown Dwarfs showing broad absorption features between 8$\\mu$m and 11$\\mu$m which is the wavelength range were silicate dust would absorb. Although the idea of dust absorbing in the atmosphere is thermodynamically correct, it is hard to proof its existence directly since we are basically searching for the absorption signature of a thin transparent layer on top of an optically thick wall (\\cite{hel2006}). \\cite{gel2002}\\, concluded that inhomogeneities in cloud decks and the evolution thereof can plausibly produce observed photometric $I$-band variations.  Extensive studies searching for uncorrelated, time-dependent variability due to a variable cloud coverage of the atmosphere, suggest a variability of 2--3\\% (\\cite{gold2008}) or even 2--10\\% (\\cite{bj2008}) depending in the wavelength studied.  Richardson et al. (2007) infer the presence of silicate haze based on secondary-eclipse Splitzer observations on giant gas-planets, a conclusion which was also put forward by \\cite{po2008} based on HST transmission spectra for HD\\,189733b (see also reviews by G. Tinetti and by J. Harrington, this volume).   \\begin{figure} \\includegraphics[height=.3\\textheight]{CircuitofDust.ps} \\caption{Interaction of multi-scale processes during cloud formation in Brown Dwarf atmospheres.Dust formation and turbulence act on small scales, and gravitational settling and element replenishment act on large scales. } \\label{f:circ} \\end{figure} ",
        "conclusions": "The modelling of Brown Dwarf atmospheres has turned into an unexpected challenge as these atmosphere contain dust which causes strong feedbacks as opacity source and element sink. Different approaches are applied to model dust in Brown Dwarf atmospheres, and obeying its nature as coupled system of nonlinear equation, atmosphere simulations applying different cloud models do produce different results although the over-all cloud-structure appears robust. A fundamental difference amongst the models is the consideration of phase-equilibrium. Kinetic models did show that in particular compounds made or rare elements like Ti, Al, and Ca never achieve phase equilibrium, i.e. thermal stability, in the Brown Dwarf atmosphere. Furthermore, the transition from the L- into the T-dwarf regime is not understood in enough detail that consistent atmosphere simulations are available. This is a sever challenge for evolutionary models.  What else needs to be done? A still outstanding problem is the modelling of turbulent dust formation in large scale simulations which is also of significance for planetary atmospheres or protoplanetary disks. The challenge is here not to lose the small-scale interactions between the chemistry and the fluid field which has been shown to support the turbulent nature of the fluid field and to even cause the appearance of medium-scale cloud structures. A second challenge is the possible coupling of the dust-ionised atmosphere to the object's magnetic field. \\cite{mo2002} and \\cite{gel2002} argue against an ionisation potential in Brown Dwarf atmospheres. However, drifting grains in a turbulent environment might make a re-consideration necessary."
    },
    "0809/0809.3174_arXiv.txt": {
        "abstract": "{Over the last decade there have been a series of results supporting the hypothesis of  the existence of a long thin bar in the Milky Way with a half-length of 4.5 kpc and a position angle of around 45$^\\circ$. This is apparently a  very different structure from the triaxial bulge of the Galaxy, which is thicker and shorter and dominates the star counts at $|l|<10^\\circ$.} {In this paper, we analyse the stellar distribution in the inner Galaxy to see if there is clear evidence for two triaxial or bar-like structures in the Milky Way.} {By using the red-clump population as a tracer of Galactic structure, we determine the apparent morphology of the  inner Galaxy. Deeper and higher spatial resolution NIR photometry from the  UKIDSS Galactic Plane Survey allows us to use in-plane data even at the innermost Galactic longitudes, a region where the source confusion is a dominant effect that makes it impossible to use other NIR databases such as 2MASS or TCS-CAIN.} {We show that results previously obtained with using  the red-clump giants  are confirmed with the in-plane data from UKIDSS GPS. There are two different structures coexisting in the inner Galactic plane: one with a position angle of 23$\\fdg$60$\\pm$2$\\fdg$19 that can be traced from the Galactic Centre up to $\\sim$10$^\\circ$(the Galactic bulge), and other with a larger position angle of 42$\\fdg$44$\\pm$2$\\fdg$14, that ends around l=28$^\\circ$  (the long Galactic bar).} {} ",
        "introduction": "The nature of the morphology of the inner Galaxy is being continously unveiled as new and precise data on its stellar content is being accumulated over the years. There are now little doubts, if any, of there being a stellar bar in the Milky Way. From the pioneering work in this topic of de Vaucouleurs (1964), using gas velocity data, to date many insights in this direction have been produced using different methods; from the analysis of infrared (IR) surface brightness maps (e.g.\\ Blitz \\& Spergel 1991; Dwek et al. 1995) to studies of the asymmetry in the number counts  (Weinberg 1992; Hammersley et al.\\ 1994; Stanek et al.\\ 1994; L\\'opez-Corredoira et al. 2001),  which  both show far  more stars at positive galactic longitudes than negative longitudes for  $l<30^\\circ$, close to the Galactic plane.  However, the exact morphology  of the inner Galaxy is still subject of some controversy, which is also being resolved as the precise stellar distribution is being delineated by a succesion of progessively larger scale, depper sensitivity and higher spatial resolution NIR surveys (TMGS, Garz\\'on et al. 1993; DENIS, Epchtein et al. 1997; 2MASS, Skrutskie et al. 2006; TCS-CAIN, Cabrera-Lavers et al. 2006; GLIMPSE,  Benjamin et al. 2003; UKIDSS, Lawrence et al. 2007).  While some authors have refered to the bar as a thick structure, around 2.5 kpc in length with a position angle of 15--30 degrees with respect to the Sun--Galactic Centre direction (Dwek et al. 1995; Nikolaev \\& Weinberg 1997; Stanek et al. 1997; Binney et al. 1997; Freudenreich 1998; L\\'opez-Corredoira et al. 1999; Bissantz \\& Gerhard 2002; Babusiaux \\& Gilmore 2005), other researchers suggest that there is a long bar with a half length of 4 kpc and a position angle of around 45 degrees (Weinberg 1992; Hammersley et al. 1994, 2000; Sevenster et al. 1999; Van Loon et al. 2003; Picaud et al. 2003; Benjamin et al.\\ 2005; L\\'opez-Corredoira et al. 2007).  It should be noted, however, that the references supporting the shorter bar all examine the region  at $|l|<12^\\circ$ and normally off the plane, partly due to the interstellar extinction, whereas those mentioning the long bar with the larger angle are trying to explain counts for $10^\\circ<|l|<30^\\circ$. This suggests that there are probably two different Galactic components coexisting in the inner Galaxy, instead of a single component alone with such a complex geometry.  In this picture, the structure present in the inner Galaxy ($|l|<12^\\circ$) might be interpreted as a triaxial bulge with axial ratios of 1:0.5:0.4 (L\\'opez-Corredoira et al.\\ 2000, 2005) rather than a bar, and certainly not a long straight bar. Furthermore L\\'opez-Corredoira et al. (2001) showed that a triaxial bulge + disc model cannot reproduce the observed counts in the Galactic plane for $10^\\circ<l<30^\\circ$, something  also suggested by NIR photometry of red-clump stars (Nishiyama et al.\\ 2005).    Along this line, Cabrera-Lavers et al. (2007a) presented a detailed analysis of the distribution of red-clump stars in the inner Galaxy by using deep near infrared (NIR) photometry from the TCS-CAIN survey (Cabrera-Lavers et al. 2006). They showed that there are evidences for a bar + bulge scenario in the innermost Galaxy. While the Galactic bulge dominates the counts at higher latitudes, the long bar can be traced from the in-plane data at least from l=18$^\\circ$ up to l=27$^\\circ$. However, up to date, no evidences of the presence of the bulge component were found in the Galactic plane due to incompleteness effects in the available NIR data at those coordinates.  The intention of this paper is to repeat and complete the analysis in Cabrera-Lavers et al. (2007a) to confirm or refute the above stated conclusions, using the best available dataset on stellar distribution of the central Galaxy in the NIR, the \\emph{United Kingdom Infrared Deep Sky Survey} (UKIDSS, Lawrence et al. 2007). The use of the red-clump source distribution in the inner Galactic plane to trace the 3D morphology of the stellar population has been performed by several authors (Hammersley et al 2000; Stanek et al. 1994; L\\'opez-Corredoira et al. 2002, 2004; Babusiaux \\& Gilmore 2005; Nishiyama et al.\\ 2005, 2006b). In this paper, we will pay close attention to the potential differences on the structural parameters derived fron the two datasets, TCS-CAIN (Cabrera-Lavers et al. 2007a) and UKIDSS (this work), which could then be assigned to different stellar distribution being traced by both surveys, due to their differences in sensitivity and spatial resolution. As it will be shown in secs. 4.2 and 5, no significant differences are found, pointing to the fact that the majority, or at least the most prominent section, of the stellar content of the inner Milky Way have been already measured. Perhaps not surprisingly, a similar results on the main structural parameter of the stellar galactic bar can also be found in Hammersley et al. (1994), where they analised data from the TMGS (Garz\\'on et al. 1993), with limiting magnitude of K=10.5, extended with data from DIRBE (Boggess et al. 1992). ",
        "conclusions": "By using the deepest wide area NIR data available to date for analyzing the inner Galactic Plane, we find that there are two different structures coexisting in the inner 4 kpc of the Milky Way: From the Galactic Centre to l=10$^\\circ$, the distribution of red-clump sources traces a thicker structure with a position angle with respect to the Sun-Galactic Centre line of 23$\\fdg$60$\\pm$2$\\fdg$19, while for l$\\le$10$^\\circ$ the red-clump sources follow a more inclined structure that can be traced up to l=28$^\\circ$ with a position angle of 42$\\fdg$44$\\pm$2$\\fdg$14, that ends around l=28$^\\circ$.  This result corroborates the previous hypothesis of Hammersley et al. (2000) that the Milky Way presents a inner triaxial bulge together with a long thinner structure, that appears  as a long Galactic bar, with semi-length of about 4.5 kpc from the Galactic Centre. There are also evidences that the inner triaxial bulge has a 'boxy' morphology (L\\'opez-Corredoira et al. 2000, 2005), but nothing in that sense can be said from the analysis presented here.  Previous works in the same direction (CL07 and L\\'opez-Corredoira et al. 2007) could not analyze the inner Galactic Plane due to the shallowness in their NIR data. For this reason, they had to use a combination of on and off plane data to reach the same conclusions as ours. Some other researchers also pointed to this double morphology in the Milky Way, and suggested that the inner Galaxy is dominated at higher latitudes by the Galactic Bulge, while the star counts in the plane are dominated by the long bar who dominates the counts (Sevenster et al. 1999). The result presented here confirms that the same picture can be observed close to the Galactic Plane.  This suggested  scenario of a triaxial bulge + long bar for the  Galaxy is plausible as many galaxies have triaxial bulges contained within a primary stellar bar (Friedli et al. 1996), and it is well supported by recent N-body simulations about secular galactic evolution\\footnote{In fact, the ratio of sizes between the semi-major axis of the long bar and the major axis of the triaxial bulge in the simulations is 1.5, nearly coincident what it is observed in Our Galaxy.}. (Athanassoula 2005; Athanassoula \\& Beaton 2006) and also by NIR observations in external galaxies (Beaton et al. 2005). Boxy bulges are in fact a part of the long bar (Athanassoula 2005), and this is what is seen in the  Milky Way."
    },
    "0809/0809.3342_arXiv.txt": {
        "abstract": "I discuss the use of neutrinos from the CN cycle and pp chain to constrain the primordial solar  core abundances  of C  and N  at an interesting level of  precision.  A comparison of  the  Sun's  deep interior and surface compositions would test  a  key assumption of the standard solar  model (SSM), a homogeneous zero-age Sun.  It would   also  provide   a   cross-check  on   recent  photospheric   abundance determinations that have altered the  once excellent agreement between the SSM and  helioseismology.   Motivated by the discrepancy between convective-zone abundances and helioseismology, I discuss the possibility that a two-zone Sun could emerge from late-stage metal differentiation in the solar nebula connected with formation of the gaseous giant planets. ",
        "introduction": "One of the initial motivations  for pursuing solar neutrino physics was to test our understanding of main-sequence stellar evolution and  the Standard Solar Model (SSM).   In the past two decades this goal was put aside, as difficulties in understanding the pattern of solar neutrino fluxes led to the discovery of solar neutrino oscillations.  But because of the precision with which the relevant flavor physics is now known -- and because the solar neutrino problem also spurred progress in the nuclear physics of the Sun and the development of high-statistics detectors such as Super-Kamiokande (SK) and SNO -- the use of neutrinos as quantitative solar probe is now a practical possibility.  Specifically, this talk summarizes recent arguments by Aldo Serenelli and me \\cite{whas} that a measurement of the CN neutrino flux could test a key assumption of the SSM, the homogeneity of the zero-age main sequence Sun.  This  assumption, the  basis for equating the SSM's  primordial core metal abundances to  today's surface metal abundances, may now  be in some degree of  conflict with observation:  recent 3D  modeling of photospheric absorption  lines has  led to a  downward revision in  the metal content  of  the solar  convective  zone  \\cite{ags05}, altering helioseismology  in  the upper radiative zone, where the temperature $\\sim$ 2-5 $\\times$ 10$^6$ K \\cite{bbps05,bsb05,antia,montalban}. In this region C, N, O, Ne, and Ar are partially ionized  and  particularly  O and  Ne  have  a  significant influence  on  the radiative opacity.  A  quantitative comparison between  the Sun's  surface and  core abundances could prove  useful in understanding  the chemical evolution of  other gaseous bodies in  our solar system,  whose interiors are  not as readily  probed.  The Galileo and Cassini  missions found significant metal enrichments  in the H/He atmospheres of Jupiter and Saturn, e.g.,  abundances of C and N of $\\sim$ four times solar for Jupiter and $\\sim$ 4-8 for Saturn \\cite{guillot}.  Planetary models that  take account of these  data show that the  gaseous giants are very significant  solar-system metal reservoirs.   The metal excess in these planets, estimated to be $\\sim~40-90M_\\oplus$, is very similar to the depletion that would be needed in the convection zone to bring the photospheric abundances and helioseismology into agreement. Serenelli and I raised the possibility that a two-zone Sun resulted from a process in which the last $\\sim$ 1\\% of nebular gas was scoured of its metals by the process of planetary formation, then deposited on the early Sun to dilute the convective zone. ",
        "conclusions": ""
    },
    "0809/0809.3668_arXiv.txt": {
        "abstract": "We present a study of the chemistry of a dense photon-dominated region (PDR) using a time-dependent chemical model.  Our major interest is to study the spatial distribution of complex molecules such as hydrocarbons and cyanopolyynes in the cool dense material bordering regions where star formation has taken place. Our standard model uses a homogeneous cloud of density $2\\times10^4$\\cmt\\ and temperature $T=40$ K, which is irradiated by a far-ultraviolet radiation field of intermediate intensity, given by $\\chi=100$. We find that over a range of times unsaturated hydrocarbons (e.g., \\cdh, \\cqh, \\cthd) have relatively high fractional abundances in the more external layers of the PDR, whereas their abundances in the innermost layers are several orders of magnitudes lower. On the other hand, molecules that are typical of late-time chemistry are usually more abundant in the inner parts of the PDR. We also present results for models with different density, temperature, intensity of the radiation field and initial fractional abundances. Our results are compared with both high- and moderate-angular resolution observations of the Horsehead nebula.  Our standard model is partially successful in reproducing the observations. Additional models run with different physical parameters are able to reproduce the abundance of many of the observed molecules, but we do not find a single model that fits all the observations at the same time.  We discuss the suitability of a time-dependent model of a dense PDR such as ours as an estimator of the age of a PDR, provided that enough observational data exist. ",
        "introduction": "Far-ultraviolet (FUV) photons (6 eV$<h\\nu<13.6$ eV) can drastically affect the structure, chemistry, thermal balance, and evolution of portions of the interstellar medium, even if only the mean galactic interstellar radiation field (ISFR) is present. For this reason, the simplest chemical models of dense interstellar clouds should include a treatment of the radiative transfer of this radiation through the outer and intermediate layers of material, because the field affects observations that also sample the interior gas. Indeed, most of the molecular gas and certainly all the atomic gas of the Galaxy is found in regions where the FUV photon flux is several times that of the mean ISRF \\citep{Hollenbach97}.  Photon-dominated regions (PDRs) are special regions of neutral gas exposed to intense fluxes of FUV radiation, which dominate the physical (heating) and chemical processes that take place in their interiors. PDRs can be observed in very different environments, including dense molecular clouds such as the Orion Bar or the Horsehead nebula, planetary nebulae, winds of red giants and AGB stars, the interstellar medium of starburst galaxies such as M82, low metallicity dwarf galaxies, etc. PDRs are studied through atomic spectral lines such as [C\\,{\\small II}] at 158$\\mu$m, [O\\,{\\small I}] at 63 and 146 $\\mu$m, [C\\,{\\small I}] at 370 and 609 $\\mu$m; many H$_2$O lines; millimetre and sub-millimetre rotational transitions of molecules such as CO, $^{13}$CO, or C$^{18}$O; ro-vibrational and pure rotational lines of H$_2$; and many broad mid-IR features attributed to PAHs.  Physical conditions of PDRs are thus very diverse.  They can be diffuse, with gas density $n\\sim 10-100$~cm$^{-3}$, or dense, with $n > 10^4$ cm$^{-3}$, while the incident FUV flux may range from the ISRF to $10^6$ times the ISRF very near to an O star.  An important characteristic of PDRs is that they have a layered structure, as a result of the interaction of the radiation with the gas and dust. Typically, they contain an outer layer of partially ionised gas, where H is atomic and the predominant form of carbon is C\\,{\\small II}; a layer of neutral gas, where H$_2$ is predominant and carbon is present mainly as C\\,{\\small I}; and a molecular layer, where CO is the main form of carbon. The ratio between the flux of the incident radiation field and the gas density determines the size of these layers. The distinct layers should be observable if the PDR is seen edge-on, because then the geometry of the cloud is more apparent. One of the main difficulties of the study of PDRs is, however, the determination of the geometry of the object.  The gas temperature is also dependent on the strength of the radiation field and the position inside the cloud, and can range from 10-100 K in the molecular layer to more than 10$^3$ K in the surface layers.  \\begin{table} \\caption{Standard initial (dark cloud) fractional abundances with respect to $n_{\\rm H}$} \\label{iniabund} \\begin{center} \\begin{tabular}{llll} \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} H$_2$  & 0.5                  & C$_2$H & $1.0 \\times 10^{-8}$ \\\\ H      & $7.5 \\times 10^{-5}$ & CO$_2$ & $1.3 \\times 10^{-8}$ \\\\ He     & 0.14                 & H$_2$O & $3.5 \\times 10^{-8}$ \\\\ C      & $2.8 \\times 10^{-8}$ & HCN    & $1.0 \\times 10^{-8}$ \\\\ O      & $1.0 \\times 10^{-4}$ & HNC    & $1.0 \\times 10^{-8}$ \\\\ N      & $1.3 \\times 10^{-5}$ & NH$_3$ & $1.0 \\times 10^{-8}$ \\\\ S      & $7.2 \\times 10^{-8}$ & SO$_2$ & $5.0 \\times 10^{-10}$ \\\\ Si     & $7.8 \\times 10^{-9}$ & C$_3$H & $5.0 \\times 10^{-9}$ \\\\ Cl     & $4.0 \\times 10^{-9}$ & C$_4$H & $4.5 \\times 10^{-8}$ \\\\ Fe     & $3.9 \\times 10^{-10}$ & C$_3$H$_2$ & $5.0 \\times 10^{-9}$ \\\\ Mg     & $1.9 \\times 10^{-9}$ & HC$_3$N & $1.0 \\times 10^{-8}$ \\\\ Na     & $4.7 \\times 10^{-10}$ & C$^+$  &  $4.7 \\times 10^{-9}$ \\\\ P      & $3.0 \\times 10^{-9}$ & H$^+$  &  $4.2 \\times 10^{-10}$ \\\\ CH     & $1.0 \\times 10^{-8}$ & He$^+$  & $3.5 \\times 10^{-10}$ \\\\ CN     & $2.5 \\times 10^{-9}$ & Fe$^+$ &  $2.6 \\times 10^{-9}$ \\\\ CO     & $7.3 \\times 10^{-5}$ & Mg$^+$ &  $5.1 \\times 10^{-9}$ \\\\ CS     & $2.0 \\times 10^{-9}$ & Na$^+$ &  $1.5 \\times 10^{-9}$ \\\\ N$_2$  & $4.2 \\times 10^{-6}$ & S$^+$  &  $1.2 \\times 10^{-9}$ \\\\ NO     & $1.5 \\times 10^{-8}$ & Si$^+$ &  $2.5 \\times 10^{-10}$ \\\\ O$_2$  & $8.1 \\times 10^{-8}$ & H$_3^+$ & $1.4 \\times 10^{-9}$ \\\\ OH     & $1.0 \\times 10^{-7}$ & HCO$^+$ & $4.0 \\times 10^{-9}$ \\\\ S$_2$  & $1.8 \\times 10^{-9}$ & HCS$^+$ & $2.0 \\times 10^{-10}$ \\\\ SO     & $1.0 \\times 10^{-9}$ & N$_2$H$^+$ & $2.0 \\times 10^{-10}$ \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{table}  Dense PDRs are of particular interest in our attempt to understand the process of star formation. Newly formed stars produce high intensity photon fields deep inside clouds that lead to the formation of H\\,{\\small II} regions and, further removed from the exciting stars, PDRs. The gas-phase chemistry that occurs under these conditions is in most layers greatly different from the chemistry found in quiescent regions. The abundances of some atoms or molecules are greatly enhanced by the presence of the FUV field, while other species are destroyed by the radiation. For instance, the abundances of several hydrocarbons (C$_2$H, c-\\cthd, C$_4$H, C$_6$H) have been found to be almost as high in relatively unshielded regions as in dark clouds \\citep{Fuente03, Teyssier04}. However, some of these high abundances are difficult to explain with normal gas-phase chemical models and a link between small hydrocarbons and PAHs has been proposed \\citep{Teyssier04}.  In this paper, we study the chemistry of a dense PDR created after a star has been formed inside a dark cloud. We are particularly interested in the spatial distributions of hydrocarbons and other complex molecules, and whether or not time-dependent chemistry can explain some of the discrepancies found between the observed abundances and the results of steady-state models \\citep{Teyssier04,Pety05} for the Horsehead nebula, as well as providing us with some estimate of the age of this PDR. Another use of time-dependent PDR chemistry is described by \\citet{Bell05}. The structure of the remainder of the paper is as follows: in Section~\\ref{sct_model}, we describe the time-dependent chemical model used.  Our results obtained with our standard model are then considered in Section~\\ref{sct_results}, while in Section~\\ref{sct_disc} we discuss the dependence of our results on assorted physical parameters.  The agreement between model results and observations of the Horsehead nebula is probed as a function of time in Section~\\ref{hhnebula}.  Finally, Section~\\ref{sectsum} contains a summary of the paper. ",
        "conclusions": "\\label{sectsum}  We present the results of a time-dependent model of a dense photon-dominated region created after a star has formed inside a dark cloud. We have assumed for our standard model an homogeneous cloud with density $n=2\\times10^4$\\cmt\\ and gas temperature $T=40$ K, on one side of which impinges an FUV radiation field $\\chi=100$ times the ISRF.  We have found that the more unsaturated and/or simpler hydrocarbons, in particular \\cdh, \\cqh, or C$_6$H, have peak fractional abundances in the outer layers of the PDR, at A$_V\\sim$1--2, at times as early as a few $\\times 10^4$ yr, while more saturated hydrocarbons, such as CH$_4$, have their peaks deeper into the cloud at these times, around A$_V\\sim5$.  The outer and intermediate layers reach steady-state rapidly, at times shorter than a few $\\times 10^5$ yr, while for the inner layers it takes almost an order of magnitude longer.  The molecule CS and the cyanopolyynes also peak nearer to the surface of the cloud, at A$_V\\sim1.7$--2, at times following a few $\\times10^4$ yr, while HCN, HNC, HCO$^+$, \\nthp\\ and \\nh\\ peak deeper into the clouds, from A$_V=5.5$ onwards. This result is very interesting because it reproduces in space a similar distribution of molecules to the one found in time for a homogeneous gas-phase model. The external radiation field seems to keep the chemistry `young' in the outer layers of the cloud, while only the more shielded slabs develop a  `late-time' chemistry.  We have run several models with modifications of some of the parameters of the standard model, varying the intensity of the radiation field, the temperature and density of the cloud, and the initial distribution of fractional abundances in the gas. For the three models run with non-standard scaling factors of the radiation field of 10, $10^3$, and $10^4$, the effects on the peak fractional abundances of hydrocarbons are not large, typically about an order-of-magnitude.   At large values of the radiation field, all molecules are swept out of the outer regions, while only the gas at A$_V> 9-10$ is completely shielded from variations in the radiation field.  One salient result of increasing the density is that for hydrocarbons that peak in the outer layers, the abundance profile when plotted against distance has a smaller region of large fractional abundance.  Another salient feature is that the peak fractional abundances of selected hydrocarbons and cyanopolyynes increase with density in the range 10$^{3}$ - 10$^{6}$ cm$^{-3}$ while other species show little change or even a decrease. Variations in peak fractional abundance as temperature is changed tend to be less significant than with density for temperatures in the range 10-200 K.  We also explored the effects that different initial fractional abundances have on the results of the model by running a model that starts with atoms in the gas, finding that the evolution of fractional abundances in time is very different from the standard model, although the end result (steady state) is, of course, the same.  In a comparison of the results of our standard model with recent high-angular resolution observations of the Horsehead nebula, we were able to reproduce, within the estimated observational uncertainties, the peak fractional abundance and spatial dependence of \\cdh; not so with \\cqh\\, and \\cthd\\ which show much worse agreements at the outermost positions. For the rest of the molecules, studied at moderate- and high-angular resolution, the agreement to the observations go from good (for CS and \\hcsp) to poor (especially for \\hctn). We have also used additional models to explore how different physical conditions can improve the agreement with observations. Even if some models did reproduce better some of the molecular abundances, we could not find a single model that was able to fit all the observations at the same time.  We also ran two models in which the initial CO abundance was changed from its  `standard'value of $7.3\\times10^{-5}$ with respect to total H. First, we reduced this value to that measured in TMC1-CP $(4\\times10^{-5})$, while conserving the total elemental abundances of C and O by increasing the atomic abundances of carbon and oxygen. Secondly, we used the reduced CO abundance with no attempt to maintain our standard elemental abundances of C and O. In both instances, little change was seen for carbon-chain species from the standard model at times near steady state. At earlier times, however, the first modification produces strong enhancements in carbon-chain molecules for the inner layers while the second modification produces smaller diminutions of these abundances that are even smaller than at steady state. The status of the agreement with observations in the Horsehead nebula is unaffected by changes in the initial CO abundance, since the changes affect the inner layers, which have not been observed.  Recently, there have been some new results that might help to solve the disagreement between observation and theory. First, Rimmer and Herbst (in preparation)  have re-calculated the cosmic-ray ionization rate as a function of depth into clouds. We found in our models that a different value of $\\zeta$ may dramatically change the abundances of the species. Secondly, there is evidence that calculations including negative ions may indeed increase the calculated abundances of neutral carbon chain species. These calculations are currently being done by R.~Ni Chiumin, Tom Millar, Nanase Harada, Eric Herbst and others. Thirdly, \\citet{parise} have confirmed that the Horsehead Nebula is rather clumpy by their analysis of deuterated molecules. We intend to use these new calculations in further work to find out if they are able to reproduce the abundances of the molecules observed in the Horsehead nebula.  From the currently available observational data, it is difficult to obtain a definitive age of the Horsehead nebula. We would need more observations, especially of layers deep into the cloud (A$_V\\ga4$), where the abundances of species such as hydrocarbons and cyanopolyynes are more sensitive to changes in time."
    },
    "0809/0809.0341_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  A large number of X-ray binaries show X-ray flux periodicities at their orbital period. The current comprehensive list of such objects is given by Wen et al.\\ (2006). Using strict criteria for the statistical significance of the presence of a modulation, they find 33 secure orbital periods in the \\xte/ASM data. Due to the limitation of the ASM sampling, they were unable to detect the 685-s X-ray binary period of the ultracompact low-mass X-ray binary (LMXB) 4U 1820--303, which, however, is clearly detected, e.g., in \\xte/PCA data (e.g., Zdziarski et al.\\ 2007b, hereafter Z07, and references therein). Even faster X-ray modulations, at 569 s and 321 s, are found in \\rosat\\/ data for RX J1914.4+2456 (Cropper et al.\\ 1998) and RX J0806.3+1527 (Israel et al.\\ 1999), respectively. These two systems contain pairs of white dwarfs orbiting each other at the above respective period (Cropper et al.\\ 1998; Israel et al.\\ 2002; Ramsay et al.\\ 2002).  In addition, a number of X-ray binaries show X-ray modulations at periods longer than their orbital periods, so-called superorbital periodicity. Sood et al.\\ (2007) state that 35 such systems have been reported in literature. However, many of those claims remain uncertain, and Sood et al.\\ (2007) selected 19 of them as secure. However, even some of those appear, in our opinion, uncertain. For example, they list the high-mass X-ray binary (HMXB) Cen X-3 as having a superorbital period of 140 d. On the other hand, the same object is presented by Raichur \\& Paul (2008) as a case of the {\\it lack\\/} of stable superorbital modulation. Another example is 4U 1916--053, which is listed as having a 199-d period in Sood et al.\\ (2007) whereas Homer et al.\\ (2001) (and later Wen et al.\\ 2006) find no trace of such periodicity. Wen et al.\\ (2006), using a uniform analysis method, find 6 objects as having coherent superorbital modulations in the \\xte/ASM data, and further 5 with quasi-periodicities. Using their method with no pre-selection of data, they did not find a superorbital periodicity in the black-hole HMXB Cyg X-1. However, its $\\sim$150-d period manifests itself mostly in the hard spectral state, where its statistical significance is very high (e.g., Brocksopp et al.\\ 1999a; Karitskaya et al.\\ 2001; {\\\"O}zdemir \\& Demircan 2001; Lachowicz et al. 2006; Ibragimov et al.\\ 2007; Poutanen et al.\\ 2008, hereafter PZI08). Another black-hole binary in which the superorbital modulation manifests itself only in the hard state appears to be the LMXB GX 339--4, in which it shows a $\\sim$220-d periodicity, whereas that modulation is absent in the soft state (Zdziarski et al.\\ 2004). (See, e.g., Zdziarski \\& Gierli\\'nski 2004; Done et al.\\ 2007 for reviews and physical interpretation of the spectral states of black-hole binaries.)  The MAXI, an X-ray all-sky monitor mission on the Japanese Experiment Module of the {\\it International Space Station}, is scheduled to be launched in 2009 March. Due to its unprecedented sensitivity of a few mCrab in a day covering most of the sky, the MAXI will monitor variability of a large number of X-ray sources at much lower flux levels than is now possible with the \\xte/ASM. It will thus be able to greatly enhance our knowledge of the variability of X-ray binaries, in particular of their orbital and superorbital periodicities. ",
        "conclusions": "We have reviewed orbital and superorbital modulation of X-ray binaries, concentrating on the X-ray spectral band. The physical causes of the orbital modulation are relatively well understood, see Section \\ref{physical} The most common cause of the superorbital modulation appears be precession of accretion disc. However, details remain unclear, in particular what drives the precession, whether it is prograde or retrograde, and how a viscous disc can coherently precess. On the other hand, some objects, especially 4U 1820--303, show very strong evidence of the superorbital variability being due to periodic changes of $\\dot M$, and not due to precession. However, physical details also here remain unclear. One possibility is the effect of a third star in a hierarchical triple (Chou \\& Grindlay 2001; Zdziarski et al.\\ 2007a).  Then, the properties of the orbital modulation may depend on the superorbital phase. We have discussed two specific examples. In Cyg X-1, the orbital modulation is strongest at the lowest observed flux, which appears to be due to the presence of an accretion bulge at the outer edge of the precessing disc. The bulge goes into the line of sight when the disc is most edge-on, which then corresponds to the lowest observed flux. An opposite effect is observed in 4U 1820--303. Here, the orbital modulation is weakest at the lowest observed flux. The reason for that appears that the size of the obscuring bulge increases with the increasing $\\dot M$, periodically changing in that system (unlike Cyg X-1, where most likely there is no periodic $\\dot M$ modulation).  \\vspace{1pc} \\noindent AAZ has been supported by the Polish MNiSW grant NN203065933 (2007--2010) and the Polish Astroparticle Network 621/E-78/SN-0068/2007. This study has been partially supported by the Academy of Finland grants 110792 and 112986, and the International Space Science Institute (Bern). AI has been supported by the Finnish Graduate School in Astronomy and Space Physics, the V\\\"ais\\\"ala Foundation and the Russian President's program for support of leading science schools (grant NSH-4224.2008.2). LW has been supported by the Alexander von Humboldt Foundation's Sofja Kovalevskaja Programme funded by the German Federal Ministry of Education and Research."
    },
    "0809/0809.5278_arXiv.txt": {
        "abstract": "{Magnetic fields are believed to play a crucial role in the process of star formation.} {We compare high-angular resolution observations of the submillimeter polarized emission of NGC~1333~IRAS~4A, tracing the magnetic field around a low-mass protostar, with models of the collapse of magnetized molecular cloud cores.} {Assuming a uniform dust alignment efficiency, we computed the Stokes parameters and synthetic polarization maps from the model density and magnetic field distribution by integrations along the line-of-sight and convolution with the interferometric response.} {The synthetic maps are in good agreement with the data. The best-fitting models were obtained for a protostellar mass of 0.8~$M_\\odot$, of age $9\\times 10^4$~yr, formed in a cloud with an initial mass-to-flux ratio $\\sim 2$ times the critical value.} {The magnetic field morphology in NGC~1333~IRAS~4A is consistent with the standard theoretical scenario for the formation of solar-type stars, where well-ordered, large-scale, rather than turbulent, magnetic fields control the evolution and collapse of the molecular cloud cores from which stars form.} ",
        "introduction": "\\label{intro}  NGC~1333~IRAS~4A (herafter IRAS~4A) is one of the prototypical low-mass young stellar systems in the earliest stages of evolution, and it is still deeply embedded in an infalling dense molecular and dusty envelope (Sandell et al.~1991; Di Francesco et al.~2001) and powering a well-collimated outflow (Blake et al.~1995; Choi~2005; Choi et al.~2006). The BIMA spectropolarimetric observations have detected and partially resolved the polarization in both the dust (at 1.3mm) and line (CO $J=2$--1) emission (Girart et al.~1999). Recent observations with the submillimeter Array (SMA) at 877~$\\mu$m (see Fig.~\\ref{iras4a}) show that the magnetic field associated with the infalling envelope has a well-defined hourglass morphology on scales of a few hundred AU (Girart et al.~2006, hereafter GRM06). In this paper, we perform a quantitive comparison of the observed submillimeter polarization data with models of the collapse of magnetized molecular cloud cores and show that the data support the theoretical scenario where the ordered, mean component of the interstellar magnetic field controls the evolution and collapse of the molecular cloud cores from which stars form (see e.g. Shu et al.~1987, 1999; Machida et al.~2005a,b; Banerjee \\& Pudritz~2006).  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{GGG08_fig1.ps}} \\caption{Map of NGC~1333~IRAS~4A, from GRM06. Contours show the continuum emission at 877~$\\mu$m, bars indicate the direction and degree of polarization (magnetic field vectors), and the color map shows the polarized intensity. At the distance of 300~pc, $1\\arcsec$ corresponds to 300~AU.} \\label{iras4a} \\end{figure} ",
        "conclusions": "\\label{discussion}  From the discussion of Sect.~\\ref{results}, it is clear that the agreement of magnetic collapse models with the continuum polarization data for IRAS~4A is very good, supporting the standard theoretical scenario for the formation of low-mass stars from magnetized molecular cloud cores.  The residuals distributions indicate that the best agreement with the data is obtained with GS93 models with early ages ($\\sigma=14.8^\\circ$ for $\\tau=0.3$) or for S06 models with low values of the Ohm's radius ($\\sigma=11.7^\\circ$ for $r_{\\rm Ohm}=5$~AU). The low value of $r_{\\rm Ohm}$ obtained from the best-fitting S06 model ($r_{\\rm Ohm}=5$--50~AU) might also indicate that the steady-state conditions assumed in that model have not yet been reached, and the size of the region of strong magnetic dissipation is below the resolution limit of the observation. At any rate, both models suggest that the magnetic field lines in a region of radius $\\sim 500$~AU from the central source(s) are almost radial, as expected for collapse under ideal MHD conditions.  The almost radial geometry of magnetic field lines in IRAS~4A is also supported by the distribution of polarized intensity shown in Fig.~\\ref{iras4a}. In fact, only models with strong central concentration of magnetic field lines reproduce the observed two-lobe structure of the observed polarized intensity, with two large ``polarization holes'' in the midplane of the pseudo-disk, on both sides of the central protostar.  Rather than a variable efficiency alignement of dust grains, the presence and extent of these ``holes'' seem to indicate a combination of geometrical effect (many field lines point to the observer in a strongly pinched geometry) combined with beam dilution effects (the pinched region is smaller than the beam), resulting in a cancellation of Stokes parameters along the line-of-sight.  Some amount of field dissipation must have already occurred in the central regions of IRAS~4A. The presence of two continuum peaks in IRAS~4A separated by $1.8\\arcsec$ (Looney et al~2000; Reipurth et al~2002),  corresponding to 540~AU at the assumed distance of 300~pc, suggests that some fragmentation of the pseudo-disk has already occurred, leading to the formation of a (possibly bound) protobinary system.  Without field dissipation, the catastrophic magnetic braking associated to the concentration of magnetic fields in ideal-MHD collapse calculation provides a fierce opposition to fragmentation and to the formation of circumstellar disks (Galli et al.~2006; Hennebelle \\& Fromang~2008; Mellon \\& Li~2008).  From the distribution of residuals of the fit, it is possible to estimate the relative magnitude of the turbulent vs. the ordered component of the magnetic field using the formula $|\\delta{\\bf B}|/|{\\bf B}|\\approx (\\sigma^2-\\sigma_{\\rm instr.}^2)^{1/2}$ (GRM06). With the $\\sigma$ obtained for the best-fitting GS93 and S06 models, we derive $|\\delta{\\bf B}|/|{\\bf B}|\\approx 20\\%$. This is probably an upper limit on the intensity of the turbulent field. Thus, The magnetic field morphology in IRAS~4A is consistent with the standard theoretical scenario for the formation of low-mass stars from cores threaded by ordered rather than turbulent magnetic fields."
    },
    "0809/0809.2943_arXiv.txt": {
        "abstract": "{  We present here the discovery of rapid, large amplitude intraday variability in the compact flat-spectrum radio quasar 1156+295. The detection of 40\\% flux density variations at 15\\,GHz on a timescale of only 2.7\\,hours was serendipitously made when the source was observed with the Very Long Baseline Array as a part of the MOJAVE survey programme on February 5, 2007. Intraday variability on timescales of a few hours or less is rare, and there exist very few sources that show large-amplitude variations on a timescale as short as what is now observed for 1156+295. The shape of the visibility function of the source changes very little during the observation, although the correlated flux density changes by 40\\%. This suggests that the variability occurs in a single dominant compact component. The observed variability characteristics are consistent with interstellar scintillation in nearby, highly turbulent medium. The rms amplitude of modulation at 15\\,GHz is unusually large and it implies a rather high scattering measure along the line-of-sight towards 1156+295.  }  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Flux density variations on timescales of $\\lesssim2$ days -- so-called intraday variability (IDV) -- are seen at centimetre wavelengths in a significant fraction of compact, flat-spectrum AGN \\cite{wag95,lov03}. There is now good evidence that, at least in the case of the most extreme sources showing large-amplitude variations on intra-hour timescales, the IDV is due to interstellar scintillation (ISS) in the turbulent, ionised interstellar medium\\footnote{We note that there are some sources likely exhibiting also intrinsic variability on a timescale of $\\sim1$\\,day \\cite{qui91}.}  (see e.g. \\cite{big07} and references therein). This evidence is mainly based on the detection of time delays in the variability pattern arrival times between widely separated telescopes and on the observations of an annual modulation of the variability timescale. By detailed analysis of the scintillation it is possible to probe both the ISM and the structure of compact radio sources at microarcsecond resolution, which is far higher than what can be achieved by the present day VLBI \\cite{mac02}. The short variability timescale of the most extreme sources facilitates such studies by allowing well-sampled observations of the flux density fluctuations to be made within a single observing run. Unfortunately, however, while IDV in flat-spectrum radio sources is common, variability on timescales of a few hours or less and with an rms amplitude of modulation larger than 10\\% is extremely rare and only a handful of such sources are known. We report here the serendipitous discovery of one new source, QSO 1156+295, showing IDV on a timescale of less than 3\\,h and with 13\\% rms amplitude of modulation. This discovery is unusual in two ways. Firstly, it was made during a VLBI experiment, and secondly, the modulation index of 13\\% is atypically high at 15\\,GHz.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{1156+295.u.2007_02_05.icn_color.ps} \\caption{Naturally weighted 15\\,GHz false colour image of 1156+295 observed with the VLBA on February 5, 2007. The map peak brightness is 1.46 Jy beam$^{-1}$ and the low flux density cut-off is at 0.7 mJy beam$^{-1}$.} \\label{map} \\end{figure} ",
        "conclusions": "Through light-travel time arguments, the variability timescale of 2.7\\,hours translates into a brightness temperature of $\\gtrsim 2 \\times 10^{19}$\\,K for 1156+295 if the observed flux density variations are source-intrinsic. This is far in excess of the inverse Compton (IC) catastrophe limit of $10^{12}$\\,K \\cite{kel69}. Therefore, a very high Doppler factor would be needed in order to explain the observations in the standard framework of incoherent synchrotron radiation from a jet of relativistic electrons. A Doppler factor of $\\gtrsim 270$ would be needed in order to avoid the IC catastrophe in the case of a power-law electron energy distribution, and even in the case of the quasi-monoenergetic electron population \\cite{kir06}, a Doppler factor of $\\gtrsim200$ would be needed. Such fast jets are unlikely to exist on both theoretical \\cite{beg94} and observational \\cite{coh07} grounds. Thus, we consider it probable that the fast variability observed in 1156+295 is due to source-extrinsic effects.  A sufficiently compact radio source will scintillate due to the wave propagation effects in the ionised interstellar medium of our Galaxy (for a review, see e.g. \\cite{ric90,goo97}). The scattering strength is determined by the strength of the electron density fluctuations along the line of sight. Below a certain critical frequency $\\nu_\\mathrm{s}$, both narrow-band diffractive and broad-band refractive scattering phenomena can be observed and the scattering is referred to as ``strong''. Above the critical frequency, in the so-called ``weak'' scattering regime, the refractive and diffractive scattering length scales become equal and flux density variations arise due to slight focusing and defocusing over the Fresnel scale.  The variability timescale and modulation index can be used to constrain the properties of the scattering medium. IDV has been previously observed in 1156+295 at 5\\,GHz in 2002, albeit with less extreme characteristics: $m=5.8$\\% and $t_\\mathrm{var}\\sim20$\\,h have been reported \\cite{lov03}. For a point source, $m \\propto \\nu^{17/30}$ and $t_\\mathrm{var} \\propto \\nu^{-11/5}$ in the case of strong refractive ISS, while in the weak regime $m \\propto \\nu^{-17/12}$ and $t_\\mathrm{var} \\propto \\nu^{-1/2}$ \\cite{ric90}. Therefore, the difference between timescales observed at 5\\,GHz and 15\\,GHz suggests that $\\nu_\\mathrm{s}$ is above 5\\,GHz. This implies a rather large scattering measure towards 1156+295, which is somewhat unexpected considering the high galactic latitude ($b=78.4^\\circ$) of the source \\cite{cor02}. However, if the scintillation occurs in the very local ISM, there is not necessary a strong correlation between the scattering measure and the galactic latitude \\cite{bha98}. There are also 5 years between these observations, and the source could have been in a more compact stage at our epoch than in 2002. Simultaneous multifrequency measurements are needed to settle this.  Setting a lower limit to the intrinsic source size, $\\theta_\\mathrm{s}^{FWHM}$, by the IC catastrophe argument, it is possible to constrain the maximum distance to the scattering screen. If the maximum Doppler factor $\\delta$ is assumed to be $<50$ \\cite{coh07}, the lower limit to $\\theta_\\mathrm{s}^{FWHM}$ is about 17\\,$\\mu$as. Together with the short variability timescale of 2.7\\,h and the modulation index of $13$\\%, this implies a maximum screen distance of about 300\\,pc for typical screen velocities \\cite{goo97,sav08}. If equipartition conditions in the source are assumed, we can estimate that $\\theta_\\mathrm{s}^{FWHM} \\sim 270 \\cdot \\delta^{-7/17}$\\,$\\mu$as. Again assuming $\\delta < 50$, this would place the screen at the distance of $\\lesssim 100$\\,pc. The above calculations constrain also the scattering measure \\cite{goo97}, which indeed turns out to be rather large: {\\it SM} $\\gtrsim 0.5 $\\,m$^{-20/3}$\\,pc for a screen distance of $\\lesssim 300$\\,pc, and {\\it SM} $\\gtrsim 4$\\,m$^{-20/3}$\\,pc for a screen distance of $\\lesssim 100$\\,pc (an uncertainty of 3\\% in $m$ has been assumed when calculating these limits; a more detailed analysis is presented in \\cite{sav08}).  The rapid, large amplitude IDV in 1156+295 is in principle consistent with interstellar scintillation due to a nearby, localised region of highly turbulent ionised gas. The above-derived lower limits for {\\it SM} in the direction of 1156+295 are, however, 5-40 times larger than what is predicted by Cordes \\& Lazio \\cite{cor02} model of the distribution of free electrons in the Milky Way. Therefore, further observations should be carried out to search for other signs of increased turbulence in the direction of 1156+295. Finally, we note that the case of 1156+295 is also an important reminder that large amplitude IDV within the timescale of a VLBI experiment violates the basic assumption made in Earth rotation synthesis, and can significantly degrade VLBI imaging results."
    },
    "0809/0809.2491_arXiv.txt": {
        "abstract": "I review our current understanding of positron sources in the Galaxy, on the basis of the reported properties of the observed 511 keV annihilation line. It is argued  here that most of the disk positrons propagate away fom the disk (due to the low density environment) and the resulting low surface brightness annihilation emission is currently undetectable by SPI/INTEGRAL. It is also argued that a large fraction of the disk positrons may be transported via the regular magnetic field of the Galaxy into the bulge and annihilate there. These ideas may alleviate current difficulties in interepreting INTEGRAL results in a \"conventional\" framework. ",
        "introduction": "The origin of the Galactic electron-positron annihilation radiation remains problematic ever since the original detection of its characteristic 511 keV line (e. g. Diehl et al. 2006 and references therein). In particular, recent  observations of the line intensity and spatial morphology with the SPI instrument aboard INTEGRAL put severe constraints on its origin, since it appears that 1.5$\\pm$0.1 10$^{43}$ \\ee/s are annihilated in the bulge alone and  0.3$\\pm$0.2 10$^{43}$ \\ee/s  in the disk, i.e. that the bulge/disk   ratio of annihilating positrons is $B/D\\sim$3-9 (Kn\\\"odlseder et al. 2005).   In a recent work, Weidenspointer et al. (2008) find a significant asymmetry in the disk emission (factor 1.7 between positive and negative latitudes).\\footnote{Note that this finding is not confirmed in a recent analysis by Bouchet et al. 2008. }They also find that the observed distribution of low-mass X-ray binaries in the hard state has a remarkable similarity to the 511 keV longitude profile; they suggest then that those objects may be the main sources of positrons in the disk, with a sizeable (but certainly insufficient) contribution to the bulge emission.  In this short review, I discuss the Galactic positron sources in  the light of recent developments. I argue that, {\\it if} it is assumed  that {\\it positrons annihilate near their sources},  then none of the proposed positron production sites, either conventional or ``exotic''ones, completely satisfies the observational constraints. Furthermore, I argue that a  large fraction of the disk positrons, produced by thermonuclear supernovae (SNIa) or other sources, may be transported via the regular magnetic field of the Galaxy into the bulge, where they annihilate. This increases both the bulge positron annihilation rate and the bulge/disk ratio, alleviating considerably  the constraints imposed by the  SPI/INTEGRAL data analysis. In fact, I argue that the SPI data are compatible with values of $B/D$ as low as 0.5, because positrons can propagate away from their sources and fill a rather large volume (of low surface brightness), much larger than the relatively thin disks adopted in the analysis of  SPI data. This property is crucial to the success of the scenario proposed here, which depends also on the poorly known properties of the Galactic  magnetic field. In any case, the issue of positron propagation in the ISM and the magnetic field of the Galaxy appears to be crucial for our understanding of the 511 keV emission. ",
        "conclusions": ""
    },
    "0809/0809.2805_arXiv.txt": {
        "abstract": "Cosmological tests based on cluster counts require accurate calibration of the space density of massive halos, but most calibrations to date have ignored complex gas physics associated with halo baryons.   We explore the sensitivity of the halo mass function to baryon physics using two pairs of gas-dynamic simulations that are likely to bracket the true behavior.   Each pair consists of a baseline model involving only gravity and shock heating, and a refined physics model aimed at reproducing the observed scaling of the hot, intracluster gas phase.   One pair consists of billion-particle re-simulations of the original $500 \\hinv\\Mpc$ Millennium Simulation of Springel \\ea (2005), run with the SPH code Gadget-2 and using a refined physics treatment approximated by preheating (PH) at high redshift.   The other pair are high-resolution simulations from the adaptive-mesh refinement code ART, for which the refined treatment includes cooling, star formation, and supernova feedback (CSF).   We find that, although the mass functions of the gravity-only (GO) treatments are consistent with the recent calibration of  \\cite{tinker:08}, both pairs of simulations with refined baryon physics show significant deviations.  Relative to the GO case, the masses of $\\sim 10^{14} \\hinv\\msol$ halos in the PH and CSF treatments are shifted by averages of $-15 \\pm 1$ percent and $+12 \\pm 5$ percent, respectively.   These mass shifts cause $\\sim \\! 30\\%$ deviations in number density relative to the Tinker function, significantly larger than the $5\\%$ statistical uncertainty of that calibration. ",
        "introduction": "Deep cluster surveys offer the promise of tightly constraining cosmological parameters, including the nature of dark energy  \\citep{holder:01,levine:02,majumdar:03,lima:04,lima:05,younger:06,sahlen:08}.  Realizing this promise requires accurate calibration of the expected counts and clustering of massive halos, along with a  careful treatment of how halo mass relates to the signals observed by such surveys.  This logical division is reflected by two long-standing threads of effort, one focused on the emergence of massive structures from gravity and the other focused on scaling relations of multiple signals within the population of massive halos.  The fact that $17\\%$ of clustered matter in the universe is baryonic ties these threads together.  Non-gravitational physics is required in massive halos, not simply to create galaxies \\citep{white:78} but also to reproduce scaling behavior of the hot, intracluster medium (ICM)  observed in X-rays \\citep{evrard:91,borgani:01,reiprich:02,stanek:06,nagai:07}.  If a significant fraction of halo baryons become spatially segregated from the dark matter, either condensed within galaxies or disbursed from non-gravitational heating, then the gravitational development of massive structures will be altered, perhaps at the $\\sim 10\\%$ level, under strong baryonic effects.  The spatial number density of halos, or mass function, expected from Gaussian random initial conditions was originally derived using a mix of analytic arguments and numerical simulations \\citep[e.g.][]{ps:74,bond:91,st:99}.  Modern efforts focus on providing fitting functions of increasing statistical precision  \\citep{jenkins:01,warren:06,tinker:08}.  The recent \\citet{tinker:08} mass function (hereafter TMF), calibrated to a wide range of cosmological simulations that include gas-dynamic, Marenostrum simulations \\citep{yepes:07,gottlober:07}, has pushed statistical errors to the level of $5\\%$.  To date, however, there have been few gas-dynamic simulations that include a non-gravitational treatment of baryonic processes in volumes large enough to provide good statistics for high-mass halos.  Calibration of the mass function at the level of the \\citet{tinker:08} using hydrodynamic simulations is too expensive to be feasible in the near-term.  A less computationally expensive technique is to compare realizations of fixed initial conditions evolved with different baryonic physics.  \\citet{jing:06} and \\citet{rudd:08} employ this approach to study baryonic effects on the matter power spectrum,  finding 2-10\\% modifications of the matter power spectrum at scales $k \\sim 1 h \\Mpc^{-1}$.  \\cite{rudd:08} finds a halo mass function that is enhanced by $\\sim10\\%$ relative to the dark-matter only case.  Neither set of simulations were sufficiently large to properly probe rich cluster scales, however.  In this letter, we take a similar approach to examine the effect of non-gravitational, baryonic physics on the cluster mass function.  Specifically, we consider two pairs of gas-dynamic simulations, each comprised of a treatment of the gas with gravity and shock heating only and a second, more complicated treatment.   One pair are Millennium Gas Simulations \\citep[MGS hereafter]{hartley:08}, with SPH gas dynamics under Gadget-2 \\citep{gadget2}, and the second pair are adaptive-mesh ART simulations from \\cite{rudd2007phd}.    As the two simulations in each pair have the same initial conditions, we infer the effects of baryonic physics by comparing halos in the more detailed simulations with their gravity-only counterparts. The outline of this paper is as follows: in Section \\ref{sec:sims}, we discuss the various simulations and their bulk cluster properties, comparing them to observations.  Section \\ref{sec:mf} compares halo masses between corresponding halos in each pair of simulations, and discusses the effect on the total mass function.  All halos masses are identified as $M_{500}$, the mass of a spherical halo with radius $r_{500}$ and mean density $500\\rho_c(z)$, where $\\rho_c(z)$ is the critical density of the universe. ",
        "conclusions": "\\label{sec:conc}   Calibrations of the halo space density from ensembles of N-body and dissipationless gas dynamic simulations now have very small statistical uncertainties, $\\sim 5\\%$ in number \\citep{tinker:08}. At the high-mass end, this level of precision in number is equivalent to a precision in halo mass at the $2\\%$ level.   Since baryons represent $17\\%$ of the matter density, complex gas dynamics associated with galaxy formation physics could plausibly lead to effects on halo masses of more than a few percent. In this letter, we demonstrate that shifts approaching $10\\%$ in mass are possible, and that the sign of this effect is not yet understood.  We use two extreme treatments of gas physics that are likely to bracket the range of behavior due to astrophysical processes in galaxy clusters.  A simple assumption of preheating reduces the local baryon fraction in halos, thereby suppressing their mass at levels ranging from $15\\%$ at $10^{14} \\hinv\\msol$ to $5\\%$ at $10^{15} \\hinv\\msol$.  A more complete physics treatment with cooling and star formation increases the local baryon fraction and deepens the halo potential, thus enhancing halo mass, by an average of $12\\%$ at  $10^{14.5} \\hinv\\msol$.  The effects of cooling and star formation on halo mass are qualitatively consistent with the  systematic enhancement in small-scale power seen in previous simulations \\citep{jing:06,rudd:08}.    In both of the complex physical treatments we consider, the shifts in mass lead to statistically significant offsets in cluster counts from the TMF expectations.  These shifts in mass depend on the choice of scale used in defining halos: in both treatments, the mass shifts are larger when identifying halos via higher density contrasts.  Although both the \\PH and \\CSF simulations provide fair matches to the mean observed $L-T$ relation, implying the structure of the hot gas phase is nearly correct, neither describes well the stellar content  of clusters.   The \\PH simulation ignores galaxies while  the \\CSF simulation converts nearly $\\sim 50\\%$ of baryons into a large stellar component.   Although it is tempting to dismiss the \\PH model due to its lack of detailed physics, a growing body of observations, particularly the ubiquity of strong winds in moderate redshift DEEP2 galaxies \\citep{weiner:08} and the remarkably simple evolution to $z=1.4$ of the color of red sequence galaxies seen in the Spiter/IRAC Shallow Survey \\citep{eisenhardt:08}, provide supporting evidence for a scenario in which the fireworks associated with galaxy formation in clusters is both rapid and effective.  Improvements in the physical and computational modeling of cooling and star formation are needed to match the full set of observational constraints on the baryonic mass components of cluster halos.  We have shown here that varying these treatments can affect total halo masses at levels up to ten percent. Improving the accuracy of the halo mass function calibration will therefore entail a suite of sophisticated  gas dynamic simulations, not more or larger N-body simulations."
    },
    "0809/0809.4032_arXiv.txt": {
        "abstract": "We investigate the impact of both slow and fast polarization modulation strategies on the science return of upcoming ground-based experiments aimed at measuring the $B$-mode polarization of the cosmic microwave background. Using detailed simulations of the \\clover\\, experiment, we compare the ability of modulated and un-modulated observations to recover the signature of gravitational waves in the polarized CMB sky in the presence of a number of anticipated systematic effects.  The general expectations that fast modulation is helpful in mitigating low-frequency detector noise, and that the additional redundancy in the projection of the instrument's polarization sensitivity directions onto the sky when modulating reduces the impact of instrumental polarization, particularly for fast modulation, are borne out by our simulations. Neither low-frequency polarized atmospheric fluctuations nor systematic errors in the polarization sensitivity directions are mitigated by modulation. Additionally, we find no significant reduction in the effect of pointing errors by modulation. For a \\clover-like experiment, pointing jitter should be negligible but any systematic mis-calibration of the polarization coordinate reference system results in significant $E$-$B$ mixing on all angular scales and will require careful control.  We also stress the importance of combining data from multiple detectors in order to remove the effects of common-mode systematics (such as un-polarized $1/f$ atmospheric noise) on the measured polarization signal. Finally we compare the performance of our simulated experiment with the predicted performance from a Fisher analysis.  We find good agreement between the (optimal) Fisher predictions and the simulated experiment except for the very largest scales where the power spectrum estimator we have used introduces additional variance to the $B$-mode signal recovered from our simulations. In terms of detecting the total $B$-mode signal, including lensing, the Fisher analysis and the simulations are in excellent agreement. For a detection of the primordial $B$-mode signal only, using an input tensor-to-scalar ratio of $r=0.026$, the Fisher analysis predictions are $\\sim 20$~per cent. better than the simulated performance. ",
        "introduction": "\\label{sec:intro} There is currently a great deal of interest in the rapidly evolving field of observations of the polarization of the cosmic microwave background (CMB). This interest stems from the fact that such observations have the potential to discriminate between inflation and other early-universe models through their ability to constrain an odd-parity $B$-mode polarization component induced by a stochastic background of gravitational waves at the time of last scattering~\\citep{kamionkowsky97,seljak97}.  From an observational point of view, we are still a long way off from a detection of this $B$-mode signature of inflation. However, much progress has been made recently with the detection of the much stronger $E$-mode polarization signal on large scales by the \\wmap\\, experiment \\citep{page07,nolta08}. On smaller scales, a growing number of balloon-borne and ground-based experiments have also measured $E$-mode polarization including \\dasi\\, \\citep{leitch05}, \\cbi\\, \\citep{sievers07}, \\boomerang\\, \\citep{montroy06}, \\maxipol\\, \\citep{wu07}, \\capmap\\, \\citep{bischoff08} and \\quad\\, \\citep{pryke09}.  Most recently, the high precision measurement of small-scale polarization by the \\quad\\, experiment has, for the first time, revealed a characteristic series of acoustic peaks in the $E$-mode spectrum and put the strongest upper limits to date on the small-scale $B$-mode polarization signal expected from gravitational lensing by large-scale structure.  Building on the experience gained from these pioneering experiments, a new generation of experiments is now under construction with the ambitious goal of observing the primordial $B$-mode signal. Observing this signal is one of the most challenging goals of modern observational cosmology. There are a number of reasons why these types of observations are so difficult. First and foremost, the sought-after signal is expected to be extremely small -- in terms of the tensor-to-scalar ratio\\footnote{% Our normalisation conventions follow those adopted in the \\camb\\, code~\\citep{lewis00}, so that $r$ is the ratio of primordial power spectra for gravitational waves and curvature perturbations. Explicitly, for slow-roll inflation in a potential $V(\\phi)$, $2 r\\approx M_{\\rm Pl}^2 [V'(\\phi)/V(\\phi)]^2$ where $M_{\\rm Pl}= 2.436 \\times 10^{18}\\, \\mathrm{GeV}/c^2$ is the reduced Planck mass.}, the RMS polarization signal from primordial $B$-modes is $0.4 \\sqrt{r}\\, \\mu \\mathrm{K}$, and the current 95 per cent limit $r< 0.22$ from \\wmap\\, temperature and $E$-mode polarization plus distance indicators~\\citep{komatsu09} implies an RMS $<180\\,\\mathrm{nK}$. Secondly, polarized emission from our own galaxy and from extra-Galactic objects act as a foreground contaminant in observing the CMB polarized sky. Although our knowledge of such polarized foregrounds is currently limited, particularly at the higher frequencies $\\gsim 100\\,\\mathrm{GHz}$ of relevance to bolometer experiments, models suggest that such contamination could be an order of magnitude larger than the sought-after signal on the largest scales (e.g. \\citealt{amblard07}). Thirdly, gravitational lensing by large-scale structure converts $E$-modes into $B$-modes on small to medium scales (see~\\citealt{lewis06} for a review) and acts as a source of confusion in attempts to measure the primordial $B$-mode signal \\citep{knox02, kesden02}. Note however that the lensing $B$-mode signal is a valuable source of cosmological information in its own right and can be used to put unique constraints on dark energy and massive neutrinos (e.g. \\citealt{kaplinghat03,smithchallinor06}). Unfortunately these latter two effects (foreground contamination and weak gravitational lensing) contrive in such a way as to render $B$-mode polarization observations subject to contamination on all angular scales (the primordial $B$-mode signal is dominated by foregrounds on large scales whilst on smaller scales it is swamped by the lensing signal). Last, but not least, exquisite control of systematic and instrumental effects will be required, to much better than $100\\, \\mathrm{nK}$, before any detection of $B$-modes can be claimed. The sought-after signal is so small that systematic and instrumental effects considered negligible for an exquisitely precise measurement of $E$-modes say, could potentially ruin a detection of $B$-modes, if left uncorrected. One possible approach to mitigating some of these systematics in hardware is to modulate the incoming polarization signal such that it is shifted to higher frequency and thus away from low-frequency systematics which would otherwise contaminate it. There are a number of techniques for achieving this including the use of a rotating half-wave plate (HWP; e.g.~\\citealt{johnson07}), phase-switching, or Faraday rotation modulators~\\citep{keating03}.  In this paper, we investigate the ability of modulation techniques that are either slow or fast with respect to the temporal variation of the signal to mitigate a range of possible systematic effects. Our analysis is based on simulations, and the subsequent analysis, of data from a ground-based CMB $B$-mode polarization experiment. Previous investigations of the impact of systematic effects on $B$-mode observations include the analytic works of \\citet{hu03}, \\citet{odea07} and~\\citet{shimon08}, as well as the simulation-based analysis of \\cite{mactavish08} who based their study on signal-only simulations of the \\spider\\, experiment. The simulation work presented here is complementary to these previous analyses but we also take our analysis further by including realistic noise in our simulations --- we are thus able to quantify not only any bias found, but also any degradation of performance due to the presence of systematic effects. Our work makes use of a detailed simulation pipeline which we have created in the context of the \\clover\\, experiment. Although the precise details of our simulations are specific to \\clover\\,, our general conclusions regarding the impact of modulation on a variety of systematic effects are relevant to all upcoming ground-based $B$-mode experiments and many of them are also relevant for both balloon-borne and space-based missions.  The paper is organised as follows. In Section \\ref{sec:cmb_expts}, we review the relevant upcoming $B$-mode polarization experiments. Section \\ref{sec:sims} describes our simulation technique and the systematic effects we have considered. In Section \\ref{sec:analysis}, we describe the map-making and power spectrum estimation techniques that we use to analyse the simulated data. Section \\ref{sec:results} presents the results from our main analysis of systematic effects. We discuss our results in Section \\ref{sec:discussion} where, for clarity, we group the possible systematic effects considered into those that are, and those that are not, mitigated by a modulation scheme. In this section, we also demonstrate the importance of combining information from multiple detectors during analysis and compare the simulated performance of \\clover\\, to the predicted performance from a Fisher analysis. Our conclusions are summarised in Section \\ref{sec:conclusions}. Finally, Appendix~\\ref{app:pointing} develops a simple model of the spurious $B$-mode power produced by pointing jitter for experiments with a highly redundant scan strategy. ",
        "conclusions": "\\label{sec:conclusions} We have performed a detailed investigation of the ability of both slow and fast polarization modulation schemes to mitigate possible systematic effects in upcoming CMB polarization experiments, targeted at measuring the $B$-mode signature of gravitational waves in the early universe. To do this we have used a simulation pipeline developed in the context of the \\clover\\, experiment, which includes realistic instrument and observation parameters as well as $1/f$ detector noise and $1/f$ atmospheric noise correlated across the \\clover\\, focal-plane array. Using this simulation tool, we have performed simulations of \\clover\\, operating with no explicit modulation, with a stepped HWP and with a HWP rotating continuously at $3$~Hz. We have analysed the resulting time-stream simulations using the technique of detector differencing coupled with a na\\\"{\\i}ve map-making scheme, and finally have reconstructed the $E$ and $B$-mode power spectra using an implementation of the near-optimal ``pure'' pseudo-C$_\\ell$ power spectrum estimator.  As expected, we find that fast modulation via a continuously rotating HWP is extremely powerful in mitigating a correlated $1/f$ component in the detector noise but that a stepped HWP is not. In addition, we have demonstrated that a polarized $1/f$ component in the atmosphere is not mitigated by any amount of modulation and if present, would need to be mitigated in the analysis using a sophisticated map-making technique. We have further verified with simulations that fast modulation is very effective in mitigating instrumental polarization that is fixed relative to the instrument basis, for example the $T \\rightarrow Q$ leakage caused by systematic gain errors and mis-matches between detectors, in agreement with the conclusions of~\\citet{odea07}. We have also demonstrated that modulation does not mitigate a systematic mis-calibration of polarization angles and that these angles will need to be measured accurately in order to avoid a systematic leakage between $E$ and $B$-modes. The other systematics which we have investigated (pointing errors, mis-estimated time-constants) have a negligible impact on the recovered power spectra for the parameters adopted in our simulations.  In addition to our investigation of systematic effects, we have stressed the importance of combining data from multiple detectors and have demonstrated the superior performance of a differencing technique as opposed to one based on measuring all three Stokes parameters from single detectors in isolation. We suggest that the latter technique, although possible in the presence of rapid modulation, is likely a poor choice of analysis technique, at least in the presence of a common-mode systematic effect such as atmospheric $1/f$ noise.  Finally, we have compared the simulated performance of the \\clover\\, experiment with the expected performance from a simplified Fisher-matrix analysis. For all but the very lowest multipoles, where the simulations fail to match the Fisher predictions, we find excellent agreement between the predicted and simulated performance. In particular, despite the highly anisotropic noise distribution present in our simulated maps, our measurement of the total $B$-mode signal matches closely with the Fisher matrix prediction (the latter assuming isotropic noise).  On the other hand, the measurement of the large scale $B$-mode signal (and thus of the tensor-to-scalar ratio, $r$) from the simulations is around 20 per cent worse than the Fisher prediction.  This is almost certainly due to the sub-optimality of our power spectrum analysis on large scales. It is possible that the Fisher matrix predictions could be recovered from the simulations by using a more optimal weighting scheme in the pure pseudo-$C_\\ell$ analysis, or, more likely, by using a maximum-likelihood $C_\\ell$ estimator for the low multipoles.  One important class of systematic effects which we have not considered in this paper are those associated with imperfect optics. Additionally, we have considered only two effects associated with an imperfect HWP. The efficacy of fast modulation to mitigate systematic effects from imperfect optics, for example instrumental polarization due to beam mis-match, is expected to depend critically on the optical design (such as the HWP location).  We are currently working to include such optical systematics in our simulation pipeline, along with a more detailed physical model of the atmosphere and models of the expected polarized foreground emission. In future work, in addition to investigating further systematic effects, we will extend our simulations to multi-frequency observations and will use these to test alternative foreground removal techniques. We will also apply the ``destriping'' map-making technique of \\cite{sutton09} to our simulations to assess the relative merits of destriping in analysis as opposed to a hardware based approach for mitigating $1/f$ noise."
    },
    "0809/0809.3424_arXiv.txt": {
        "abstract": "The super massive black hole candidate, Sagittarius~A*, exhibits variability from radio to X-ray wavelengths on time scales that correspond to < 10 Schwarzschild radii.  We survey the potential of millimeter-wavelength VLBI to detect and constrain time variable structures that could give rise to such variations, focusing on a model in which an orbiting hot spot is embedded in an accretion disk.  Non-imaging algorithms are developed that use interferometric closure quantities to test for periodicity, and applied to an ensemble of hot-spot models that sample a range of parameter space.  We find that structural periodicity in a wide range of cases can be detected on most potential VLBI arrays using modern VLBI instrumentation.  Future enhancements of mm/sub-mm VLBI arrays including phased array processors to aggregate VLBI station collecting area, increased bandwidth recording, and addition of new VLBI sites all significantly aid periodicity detection.  The methods described herein can be applied to other models of Sagittarius~A*, including jet outflows and Magneto-Hydrodynamic accretion simulations. ",
        "introduction": "\\label{introduction}  Observations of the compact radio/IR/X-ray source Sagittarius~A* (Sgr~A*) make the most compelling case for the existence of super massive black holes.  Both speckle imaging and adaptive optics work in the near-infrared (NIR) band shows that multiple stars orbit the position of Sgr~A* \\citep{schodel03,ghez05} These orbits are consistent with a central mass of $\\sim4\\times10^6 ~ M_\\sun$ contained within 45~AU - the closest approach of any star.  Radio interferometric proper motion measurements \\citep{backer99,reid04} limit the motion of Sgr~A* to $<15$~km\\,s$^{-1}$, implying that Sgr~A* must trace at least 10\\% of the mass determined from stellar orbits.  Very long baseline interferometry (VLBI) at 1.3~mm wavelength \\citep{doeleman08} has resolved Sgr~A*, and measures an intrinsic size of $<0.3$~AU \\citep[assuming a distance of 8.0~kpc, from][]{reid93}, or 4 times the Schwarzschild radius of the central mass ($\\Rs \\approx 10~\\mu$as). Assuming Sgr~A* marks the position of the black hole, the implied density, using the proper motion lower limit on the mass and the VLBI size, is $>9.3\\times10^5~M_\\sun\\,{\\mathrm{AU}}^{-3}$.  Almost any conceivable aggregation of matter would, at this density, quickly collapse to a black hole \\citep{maoz98}.  The 1.3~mm VLBI result confirms the existence of structures within Sgr~A* on size scales commensurate with the innermost accretion region, and matches size scales inferred from light curve monitoring over a broad wavelength range. Sgr~A* exhibits variability on time scales of minutes to hours in the radio, millimeter, NIR, and X-ray bands \\citep[e.g.,][]{baganoff01, aschenbach04, genzel03, ghez04, belanger06, meyer06, yusefzadeh06, hornstein07, marrone07}, and flare rise times in the X-ray and NIR correspond to light-crossing times of $<12~\\Rs$.  Models that produce flaring X-ray flux density via Synchrotron Self Compton scattering of IR photons require emission regions of diameter $<10~\\Rs$ \\citep{marrone07}.  The resolution of millimeter- and submillimeter-wavelength VLBI is well matched to the scale of inner disk physics.  Baselines from Hawaii or Western Europe to Chile provide fringe spacings as small as $30~\\mu$as ($3~\\Rs$) at 230~GHz and $20~\\mu$as at 345~GHz.  Millimeter VLBI thus has the potential to detect signatures of hot spot and jet models proposed to explain the rapid variability of Sgr~A* as well as strong general relativistic effects, such as the black hole silhouette or shadow \\citep{falcke00,broderick06b,huang07,markoff07}.  Extending the VLBI technique to short (0.85~mm, 1.3~mm) wavelengths is essential for this work due to the interstellar scattering towards Sgr~A*, which broadens radio images with a $\\lambda^2$ dependence \\citep{backer78}.  VLBI at 7~mm and 3~mm wavelengths \\citep{rogers94,doeleman01,bower04,shen05} limits the intrinsic size of Sgr~A* to be $<2$~AU and $<1$~AU respectively, but VLBI at these wavelengths is strongly influenced by scattering effects.  For $\\lambda < 1.3$~mm the scattering size is less than the fringe spacings on the longest possible baselines.  Recent 1.3~mm VLBI results \\citep{doeleman08}, coupled with ongoing technical advances to reach 0.85~mm, strongly suggest that it is not a question of \\emph{if} but of \\emph{when} VLBI will directly probe Sgr~A* on event horizon scales.  Claims of observed periodicity in IR and X-ray light curves \\citep{belanger06, meyer06, eckart06} during Sgr~A* flares, can potentially be explained in the context of hot spots orbiting the black hole at a few times $\\Rs$.  It has been proposed that the fastest periodicity can be used to constrain the spin of the black hole, since the period of the innermost stable circular orbit (ISCO) is much shorter for a maximally rotating Kerr black hole than for a nonrotating Schwarzschild black hole \\citep{genzel03}.  Indeed, several authors have argued that the black hole must be rotating based on observed rapid X-ray and infrared periodicities \\citep{aschenbach04,belanger06,meyer06}.  However, it has also been proposed that the flaring activity can be explained by magnetohydrodynamic turbulence along with density fluctuations \\citep{goldston05,chan06}.  Rossby wave instabilities may naturally produce periodicities on the order of tens of minutes, in which case it is not necessary to appeal to a nonzero black hole spin to explain the fast quasi-periodic flares seen at multiple wavelengths \\citep{tagger06,falanga07}.  Regardless of the source of variability at millimeter wavelengths, VLBI has the potential to confirm conclusively its association with the inner disk of Sgr~A*, to probe the size of the region of variability (since an interferometer acts as a spatial filter on the emission), and to extract periodicity.  Preliminary studies involving the analysis of simulated data from expected models are important for a number of reasons.  Such studies will highlight the abilities and limitations of millimeter VLBI in regards to detecting the signatures of the physical processes in the accretion disk surrounding the black hole in Sgr~A*.  Critical resources (such as stable frequency standards, high-bandwidth recording equipment, and phased array processors) are likely to be limited initially, and telescope upgrades (such as surface accuracy improvement, expanded IF bandwidth, additional receiver bands, and simultaneous dual-polarization capability) must necessarily be prioritized. Simulated observational data can help assess the tradeoffs that must be considered for optimization of Galactic Center VLBI observations.  The ultimate goal is to explore the potential of black hole parameter estimation by present and future millimeter VLBI observations.  In this manuscript, we explore the observational signatures of an orbiting hot spot embedded in a quiescent disk around Sgr~A*.  We consider a non-imaging approach to analyzing millimeter VLBI data for several reasons.  First, one fundamental assumption of Earth rotation aperture synthesis is that the source structure is not changing.  Since orbital periods in the hot spot models are much shorter than the rotation of the Earth, this assumption is clearly violated.  Second, phase-referencing the data is presently not feasible at millimeter wavelengths since the phase path through the atmosphere changes on a time scale which is faster than the time needed to move the antennas between the reference source and Sgr A*.  We note, though, that this problem could be circumvented at connected-element arrays where some antennas could be dedicated to simultaneously observing a reference source while others observe Sgr A*, but this is also currently limited by the low SNR that can be achieved in the coherence time of the atmosphere.  Potential arrays for millimeter VLBI will have a small number of telescopes, initially as few as three, possibly with vastly different sensitivies.  Low expected signal-to-noise ratios (SNRs) on some baselines combined with the few antennas available will prevent adequate self-calibration via closure relations.  Third, even if the visibility data could be properly calibrated and the source structure were not changing over the observation, the $(u,v)$-plane would be sparsely populated with noisy data points, resulting in poor image fidelity, as shown by the simulations of \\citet{miyoshi04}.  At least initially, it will be more productive to analyze the data by model fitting in the visibility domain rather than in the image domain. ",
        "conclusions": ""
    },
    "0809/0809.1617_arXiv.txt": {
        "abstract": "We simulate black hole binary interactions to examine the probability of mergers and black hole growth and gravitational radiation signals using a specific initial distribution of masses for black holes in globular clusters and a simple semi-analytic formalism for dynamical interactions.  We include 3-body recoil and the latest results in numerical relativity for gravitational radiation recoil.  It is found that while 99\\% of binaries are ejected from low metallicity, low mass clusters; metal rich massive clusters retain 5\\% of their binaries.  An interesting fraction of the ejected binaries, especially those from high mass, high metallicity systems, merge on timescales short enough to be gravitational radiation sources during their mergers with rates approaching those expected for galactic field black hole binaries.  While the merger rates are comparable, the much larger mass of these binaries and their localization will make them appealing targets for advanced LIGO.  We single out two possible Milky Way clusters (NGC 6441 and NGC 6388) as having the properties for a good probability of retention. ",
        "introduction": "Observed black hole masses occupy two regimes, $M_{BH}\\lesssim100M_{\\sun}$ for black holes formed from core collapse supernova, and supermassive black holes with $M_{BH}\\gtrsim10^6M_{\\sun}$ which reside in the centers of galaxies.  Observations of some objects, however have suggested that a middle regime of intermediate mass black holes (IMBH, review by \\citet{Miller04}) could exist with masses between stellar and supermassive black holes.  Ultraluminous X-ray sources ($L_X>10^{39}$ ergs sec$^{-1}$) have been found in intense star forming regions outside the nuclei of some galaxies \\citep{Kaaret01,Matsumoto01,Fabbiano01}, and recently even in one globular cluster \\citep{Maccarone07}.  The lower limit on the mass for these objects, assuming isotropic emission at the Eddington limit, is a few hundred solar masses.  A stellar mass black hole would require special geometry of its accretion disk for sufficient beaming to occur and the accretion to be sub-Eddington \\citep{King01}, or need special conditions on the gas to provide a super-Eddington accretion rate \\citep{Begelman01}.  Conversely, a supermassive black hole ($M\\gtrsim10^6 M_{\\sun}$) would experience dynamical friction and sink to the center of its galaxy in too short a time to be plausibly observed at locations within the host galaxies where the IMBH candidates are projected to be seen today \\citep{Kaaret01}.  An intermediate mass for these objects would seem to be indicated.  In the nearly bulgeless galaxy NGC4395, there has been found \\citep{Fillipenko03} an AGN the black hole mass for which seems to be $\\lesssim 10^5M_{\\sun}$, which is in the upper range of IMBH masses.  An object with a similar mass may also exist in the galaxy POX 52 \\citep{Barth04}.  Another possible place to look for IMBHs besides in starburst regions in galaxies is in the centers of globular clusters.  Observations have seen an increase in the mass-to-light ratio towards the centers of two globular clusters that might be consistent with a massive object, a $M>10^4M_{\\sun}$ object in the Andromeda Galaxy cluster G1 \\citep{GRH02, GRH05}, a $4\\times10^4M_{\\sun}$ object found by \\citet{Noyola08} in the cluster $\\omega$ Cen, and a few thousand solar mass object in M15 \\citet{vanderMarel02, Gerssen02}, although in the case of M15, \\citet{Baumgardt03} are able to simulate the observations without an IMBH using smaller compact objects.  The velocity dispersion of the central stars in the cores of these globular clusters as compared to the conjectured mass of the IMBH put these clusters on the same M-$\\sigma$ relation as the bulges of galaxies with supermassive black holes \\citep{GRH02, vanderMarel02, Gerssen02}.  While the theory of the origin of the M-$\\sigma$ relation for supermassive black holes would probably not apply to globular clusters, it is intriguing that, at least in these two cases, the IMBHs in these clusters are consistent with it.  Two possible formation scenarios for the formation of IMBHs within a globular cluster have been proposed recently.  The first involves the process of core collapse, by which the heavier stars in a cluster first sink to the middle through mass segregation.  The stellar density goes very high, and might be sufficient that several stars collide forming a very large star (mass a few hundred solar masses), which collapses directly to an IMBH \\citep{Begelman78,PortZ02,Freitag07}.  Otherwise, stellar evolution causes the high mass stars to form black holes, which can become binaries through exchange into existing binaries of lower mass main sequence or neutron stars \\citep{Sigurdsson93b}.  We assume that formation of an IMBH by runaway merger does not occur in this case.  Three-body interactions, which in the cores of clusters are dominated by interactions with all three objects being black holes, can then begin to work to harden binaries to the point where they merge.  It is this scenario we intend to investigate.  Previous studies of dynamical formation of a IMBH from stellar mass black holes have been performed.  Black holes have been shown to be dynamically important in such aspects as the radius-age relation of Magellanic Cloud clusters \\citep{Mackey07,Mackey08}, and even in galactic nuclei \\citep{Lee95}.  \\citet{PortZ00} include a study of how important black hole binaries in clusters are to gravitational wave research.  Recently, \\citet{HB07} did a study of mergers of black holes in a system already containing a few hundred solar mass black hole, which represents the next step in IMBH formation after our work.  \\citet{Kulkarni93} and \\citet{Sigurdsson93a} used only 10 M$_{\\sun}$ black holes, and determined that the formation of 10$^3$ M$_{\\sun}$ objects is possible.  On the other hand, \\citet{Miller02} showed that these binaries tend to be ejected before reaching a size at which recoil becomes unimportant, precluding further growth.  Since black holes are produced from progenitors with a wide range of masses (20-100 M$_{\\sun}$) and have varied evolution histories (wind losses and mass transfer) just prior to becoming black holes, a distribution in masses may better reflect the actual situation in globular clusters.  \\citet{OLeary06} did a study using the distribution of black hole masses and binary periods as given in \\citet{Belczynski04} and a more complicated method of computing interactions than the semi-analytic model we use.  They use the old prescription for gravitational radiation recoil similar to that found in \\citet{Favata04}.  Numerically simulating black hole mergers through ringdown has now been done \\citep{Gonzalez06} and definitive recoil velocities determined, so that the sole remaining uncertainty in determining the circumstances of black hole mergers is the initial distributions of their masses.  Our semi-analytic method can more quickly respond to updates in stellar population synthesis models than direct many-body integration.  We describe the conditions under which the simulations were done, including initial conditions of the binary and the analytical form of the 3-body interactions and relevant time scales, then report results from several simulation runs, with a few parameters (e.g. metallicity) adjusted after each one.  Finally, we discuss what the results imply for observed systems and suggest two systems that fall in the higher probability category for harboring an IMBH. ",
        "conclusions": "The event rate from merging black hole binaries can be calculated from the fraction of systems that merge within a Hubble time and the relative contributions from low and high metallicity systems and light or massive clusters.  To conservatively estimate the event rate, we assume 100 globular clusters per galaxy (e.g. the Milky Way is currently thought to have about 150) and 100 BH per globular cluster ($N_{BH}\\sim10^{-4}N_{\\star}$).  We assume that the break between high versus low metallicity is at an [Fe/H] of -1.3 and that light globulars have $v_{esc}$ of less than 30 km sec$^{-1}$ and heavy globulars above this value.  The escape velocity for a globular cluster is given by $v_{esc}=\\sqrt{2\\Phi_0}$, where $\\Phi_0$ is the central potential of the cluster, $W=\\Phi_0/\\sigma^2$ is the King parameter and is correlated with the cluster concentration, and $\\sigma$ is approximately equal to the velocity dispersion except in the case for shallow globulars.  We determine the concentration and thus $W$ from the catalog of \\citet{Harris96}, while 1D velocity dispersion data were obtained from \\citet{Pryor93}.  For the 56 Milky Way clusters for which we could determine the escape velocity, we find that the percentage of clusters in each of our models is as follows:  A 45\\% (25), B 21\\% (12), C 20\\% (11), and D 14\\% (8).  Including data from table~\\ref{restab} on the number of mergers within a Hubble time, we find that ejected binaries account for $\\sim$640 mergers per galaxy in a Hubble time, with another 335 coming from those binaries that merge while still in the cluster.  Over a Hubble time this gives a rate of 10$^{-7}$ per year per galaxy.  These are very conservative estimates for the rate, as the Milky Way is assumed to have about 150 globular clusters, and giant ellipticals can have on the order of $10^3$.  Assuming a value of 300 globular clusters per galaxy and 300 black hole binaries per cluster, there would be an order of magnitude jump in the rate to 10$^{-6}$ per year per galaxy.  We also expect a further increase in rate from the additional mergers produced by black holes that are retained after their first merger.  Galactic binary BH merger rates are estimated at $~10^{-6}$ per year \\citep{OShaun05a}.  The mergers are expected to be delayed from the formation of the clusters, which in the case of globulars is close to the beginning of the universe.  The last interaction before the binary is ejected typically happens when the semimajor axis is 0.1-1 AU, giving a $t_{enc}$ of $~10^8-10^9$ years.  The timescales for the gravitational merger of the ejected binaries spans a wide range of values ($5<\\log t_{GW}<20$).  Figure \\ref{tGW} shows the distribution of merger timescales for ejected binaries for each of the models.  We find that the percentage of binaries which merge between 1 and 10 Gyr is 2.6\\% for model A, 2.3\\% for model B, 4.3\\% for model C, and 8.2\\% for model D.  While we have used a thermal distribution ($P(e)=2e$, $\\langle e\\rangle=0.67$) for the eccentricity after an exchange, the 3-body study by \\citet{Sigurdsson93b} found that this works for equal mass exchanges, but for non-equal masses, the eccentricities may be higher ($\\langle e\\rangle\\approx 1-1.3(m_3/m_2)$).  This does not affect recoil velocities, but the $t_{GW}$ would be shortened, and the expected rates of black hole mergers would increase by a factor of a few.  If we choose black hole binaries as we do for the simulation runs and determine their $t_{GW}$ without any interactions, we find that for metal rich systems, only 0.1\\% merge in less that a Hubble time, whereas 7\\% of low metallicity binaries do so.  This we explain as a model dependent result, the low metallicity period distribution includes systems which have periods shorter than 100 days while the metal rich distribution does not.  Interactions are of great importance in metal rich systems for producing observable mergers, while they are ambivalent in metal poor systems.  The chirp masses for cluster binary mergers are much higher due to exchanges undergone while the binary was in the cluster. We find that, while galactic mergers have chirp masses of 3-8M$_{\\sun}$ \\citep{Belczynski07}, the chirp masses for the ejected binaries are 15-25M$_{\\sun}$ in the metal rich case and 20-40M$_{\\sun}$ in metal poor clusters.  The higher chirp masses, while dependent on the models used for the initial mass function of the black holes, is a distinct prediction characteristic of the globular cluster binaries, and easily observable by gravitational radiation instruments.  Since the strain due to gravitational radiation scales as $M/r$, the factor of 4-6 increase in mass of the cluster binaries makes them visible over a factor of 60-200 larger volume, which makes them almost as important source as galactic binaries for the conservative values of GC/galaxy and binaries/GC.  If we assume the less conservative numbers, the cluster binary inspirals would dominate the signal.  LIGO will have an abundance of targets from the ejected binaries.  Our work in this paper provides a first step from population synthesis to the possible formation of an IMBH in the center of a globular cluster.  Examining a second merger once the merged object has exchanged into a new binary is beyond the scope of this work, but has been studies by \\citet{HB07}.  To connect our theoretical models to observed clusters, a plot of metallicity versus $v_{esc}$ is given in figure~\\ref{metvesc} using metallicity data from the catalog by \\citet{Harris96} and the escape velocity as described above.  Two clusters that have both high metallicity and $v_{esc}>50$ km sec$^{-1}$ are NGC 6388 and NGC 6441.  Both of these clusters lie within 4 kpc of the galactic center.  These clusters are most well known for their contribution to the \"second parameter\" problem in that they have more extended blue horizontal branches than their metallicity would indicate.  It is speculated that this might be due to dynamical interactions in the clusters \\citep{Rich97,Miocchi07}.  They note that the M31 cluster G1, a cluster suspected of having an IMBH by \\citet{GRH02}, also shows an extended blue horizontal branch.  These clusters may have been at one point the nuclei of dwarf galaxies, as a couple of other suspected nuclei appear in interesting regions of the metallicity-$v_{esc}$ plot.  Other clusters suspected of being dwarf galaxy nuclei are M54 (due to its association with the Sagittarius dwarf galaxy) by \\citet{Ibata94} and $\\omega$ Cen \\citep{Norris96,Norris97}.  The presence of extended blue horizontal branch stars in the metal-rich clusters NGC 6388 and NGC 6441 is thought to give a similar argument for their being formed in a similar manner \\citep{Piotto97}.  A couple others which are less outstanding but still in the upper right part of the diagram are NGC 2808 and M62.  Observations of variability in the recently discovered ULX in a globular cluster of NGC 4472 by \\citet{Maccarone07} lead to estimates of a 300M$_{\\sun}$ IMBH, though their other solution gives a mass of 30M$_{\\sun}$.  They find a metallicity of the cluster of -1.7 from color-metallicity relations, and the luminosity gives it a absolute magnitude of -9.2.  When compared to analogous clusters in the Milky Way (e.g. NGC 6273), this cluster fits into category C, low metallicity high mass.  Figure~\\ref{metvesc} places NGC 6388 and NGC 6441 in context with other massive, well studied globular clusters.  In conclusion, we find that within our simplified model assumptions, most black hole binaries are ejected through gravitational 3-body interaction from the cluster into the general potential of the galaxy.  Of those binaries that survive to merge by gravitational radiation, about 2/3 to half are ejected through gravitational radiation recoil.  Between 0.5\\% and 3.5\\%, depending on metallicity and cluster escape velocity, of all black hole binaries in clusters are predicted to be retained upon merger of the binary, with typical final masses of 20-50 M$_\\sun$, but in some instances over 100M$_{\\sun}$.  Of course if other formation channels dominate, or there is significant gas accretion after the dynamical interaction phase, then the final black hole masses may be very different (higher if there is significant accretion).  We find that the rate per galaxy of black hole binary mergers from the globular cluster population, through gravitational radiation is competitive with the total merger rate from the parent galaxy, but biased towards higher masses.  While most globular clusters in massive galaxies probably form at high redshift, this suggests that black hole binary coalescence from clusters in low mass, nearby star forming galaxies may be a significant contributor to the total high frequency gravitational radiation signal in the local universe.  The current results are dependent on the exact form of the initial conditions of mass distributions and period distributions obtained from population synthesis.  As the formation mechanisms for black holes become more well understood, it would be appropriate and easy to refine the results in this paper.  We thank the Center for Gravitational Wave Physics for support, and Ben Owen and Richard O'Shaughnessy for helpful discussions in making this paper.  We would also like to thank the referee for his helpful suggestions.  This research was funded by NSF grant PHY 0114375."
    },
    "0809/0809.4204_arXiv.txt": {
        "abstract": "{ Approximately 20\\% of very metal-poor stars ([Fe/H] $<$ --2.0) are strongly enhanced in carbon ([C/Fe] $>$ +1.0). Such stars are referred to as carbon-enhanced metal-poor (CEMP) stars. A variety of abundance patterns are found among CEMP stars. Strong overabundances of nitrogen are common, and overabundances of neutron capture elements are often, however not always, present. The variety of abundance patterns among CEMP stars strongly suggests that this population of stars comprises several astrophysical origins.  We are conducting a high-resolution follow-up of candidate EMP stars extracted from the Sloan Digital Sky Survey (SDSS; York et al.~2000) using UVES at the VLT. Three of the programme  stars, SDSS J0912+0216, SDSS J1036+1212 and SDSS J1349-0229, where deliberately targetted as CEMP stars since a strong $G$ band was evident from the SDSS spectra and the weakness of the Ca\\,{\\sc ii} K line testified their very low metallicity. The UVES high resolution follow-up confirmed the original findings ([Fe/H] $<-2.50$ ) and allowed a more detailed investigation of their chemical composition.  We determined the carbon abundance from molecular lines which form in the outer layers of the stellar atmosphere. It is known that convection in metal-poor stars induces very low temperatures which are not predicted by classical 1D stellar atmospheres. To obtain the correct temperature structure, one needs full 3D hydrodynamical models. 3D carbon abundances were determined for all three stars, using CO$^5$BOLD 3D hydrodynamical model atmospheres. 3D effects on the carbon abundance are found to be quite significant for these stars, with 3D corrections of up to --0.7 dex.  Two of the stars, SDSS J0912+0216 and SDSS J1349-0229 exhibit an overabundance of neutron capture elements which classifies them as CEMP-s. Star SDSS J1036+1212, instead belongs to the elusive class of CEMP-no/s stars, with enhanced Ba, but deficient Sr, of which it is the third member discovered to date. }  \\FullConference{10th Symposium on Nuclei in the Cosmos\\\\ July 27 - August 1, 2008\\\\ Mackinac Island, Michigan, USA}  \\begin{document} ",
        "introduction": "The EMP candidates have been selected using a version of our automatic analysis code (Bonifacio \\& Caffau 2003) tailored to SDSS spectra of turn-off stars which provides an estimate of [M/H]. Follow up high resolution spectra were acquired with UVES at the VLT, with a 1\\farcs{4} slit and $2\\times2$ on-chip binning, providing a resolution of $\\sim 30\\,000$. Figure~\\ref{fig1} displays the CH {\\it G} band of all three stars, showing clearly the C enhancement.  \\begin{figure}[ht] \\centering \\includegraphics{GbandsCH.ps} \\caption{Observed spectra of the CH {\\it G} band in the three CEMP stars, compared to the same spectral region of a dwarf metal-poor star, HD84937, showing no carbon enhancement. Spectra have been offset for clarity.} \\label{fig1} \\end{figure} ",
        "conclusions": "The 3D effects on the carbon abundance determined using CH lines were found to be significant, with 3D corrections of up to --0.7 dex. This suggests that all the published C abundances in dwarf CEMP stars, based on molecular lines and 1D models should be reconsidered in the light of hydrodynamic simulations. Star SDSS J1349-0229 exhibits a strong C$_2$ Swan band which classifies it as a carbon star (C/O$>1$), while the other two stars are carbon-enhanced stars.  In Figure~\\ref{fig4} we plot the Sr and Ba abundances of the three stars (red symbols) together with a sample of 24 CEMP stars from Table 4 of Sivarani et al.~(2006). It is now widely accepted that the CEMP class includes objects which have rather diverse astrophysical origin. One sub-classification of such stars was proposed by Beers \\& Christlieb~(2005) based on the abundances of neutron-capture elements in CEMP stars. One therefore speaks of CEMP-no stars, when no enhancement of neutron capture elements is observed; CEMP-s when only the s-process elements are enhanced, and CEMP-r/s when both r-process and s-process elements are enhanced. The class CEMP-no/s was introduced by Sivarani et al.~(2006) to contain the stars CS~29527-041 and CS~31080-095 which where overabundant in Ba but underabundant in Sr. Star SDSS J1036+1212 is the third star known to fall into this category. Interestingly all three known CEMP-no/s stars are warm dwarf/TO stars. The vast majority of CEMP dwarf stars is understood as the result of mass-transfer from a presently unseen companion which has undergone the AGB phase. The peculiar Sr/Ba ratios of the CEMP-no/s stars should provide useful constraints on the properties of the companion star while on the AGB.  \\begin{figure} \\includegraphics{SrFe_BaFe.ps} \\caption{[Ba/Fe] vs. [Sr/Fe] for our three CEMP stars (red filled symbols) plotted alongside values for 24 CEMP stars (open squares for dwarfs, circles for giants) listed in Table 4 of Sivarani et al.~(2006). } \\label{fig4} \\end{figure}  The other two stars, SDSS J0912+0216 and SDSS J1349-0229, exhibit both high abundances of Sr and Ba, indicating that they are CEMP-s stars. Our single epoch observations offer no indications on the possible existence of binary companions, however radial velocity monitoring of all three would be highly desirable."
    },
    "0809/0809.4945_arXiv.txt": {
        "abstract": "% Near-infrared polarimetry of point sources reveals the presence of a toroidal magnetic field in the central 20' x 20' region of our Galaxy. Comparing the Stokes parameters between high extinction stars and relatively low extinction ones, we have obtained a polarization originating from magnetically aligned dust grains at the central region of our Galaxy of at most 1-2 kpc. The derived direction of the magnetic field is in good agreement with that obtained from far-infrared/submillimeter observations, which detect polarized thermal emission from dust in the molecular clouds at the Galactic center. Our results show that by subtracting foreground components, near-infrared polarimetry allows investigation of the magnetic field structure at the Galactic center. The distribution of the position angles shows a peak at around 20$^{\\circ}$, nearly parallel to the direction of the Galactic plane, suggesting a toroidal magnetic configuration. ",
        "introduction": "The magnetic field configuration at the Galactic center (GC) has been investigated with a wide variety of methods. Recent far-infrared (FIR) and sub-millimeter (sub-mm) polarimetric observations point out that the magnetic field is generally parallel to the Galactic plane \\citep{Novak00,Novak03,Chuss03}. This is contrast to the poloidal field traced by non-thermal radio filaments, most of which align nearly perpendicular to the Galactic plane.   Previous near-infrared (NIR) polarization measurements of the GC were discussed in terms of selective absorption by the intervening interstellar dust grains in the Galactic disk. To our knowledge, no one has studied the magnetic field configuration {\\it at} the GC with NIR polarimetry. In this paper, we present results of NIR polarimetric observations toward the GC. We demonstrate that NIR polarization of point sources can provide information on the magnetic field structure at the central region of our Galaxy. ",
        "conclusions": "The direction of the magnetic field at the GC has been investigated from polarized dust emission in the FIR and sub-mm wavelengths. As seen in Fig. 2, the magnetic field configuration we obtained at the GC shows a good agreement globally with those obtained by \\citet{Dotson00}, \\citet{Novak00}, and \\citet{Chuss03}, which are the highest angular resolution polarimetry data sets in the FIR/sub-mm wavelengths. The polarized FIR/sub-mm emission comes from molecular clouds, which are known to be located in the GC. Therefore we conclude that the position angles derived from our NIR polarimetry represent the direction of the magnetic field {\\it in} the GC.   \\begin{figure}[!ht] \\plotone{fig2.eps} \\caption{$K_S$-band vector map derived from the Galactic center component ($P_{\\mathrm {R-B}}$ \\& $\\theta_{\\mathrm {R-B}}$, thick bars). The length of the bars is proportional to the measured degree of polarization, and their orientation is drawn parallel to the inferred magnetic field direction. The polarization map derived from FIR/sub-mm observations (thin bars) is also shown. The data sets of FIR/sub-mm wavelengths are from 60 $\\mu$m \\& 100 $\\mu$m polarimetry by \\citet{Dotson00}, and 350 $\\mu$m polarimetry by \\citet{Novak00} and \\citet{Chuss03}. } \\end{figure}   From $H-K_S$ color differences between peaks in the $H-K_S$ histograms and mean colors of the blue and red stars, we tried to estimate the depth of the region where we have mapped the magnetic field configuration. Using $A_{K_S}/E_{H-K_S}=1.44$ \\citep{Nishi06} and the model of dust distribution by \\citet{Davies97}, we obtained the average distances of 0.5 kpc from the GC for the blue stars, and 1.0 kpc for the red stars. This suggests that the polarization shown in Fig. 2 occurs between the average distances of $(R_0 - 0.5)$ kpc and $(R_0 + 1.0)$ kpc from the Sun, arising probably from the central 1$-$2 kpc region of our Galaxy (where $R_0$ is the distance between the GC and the Sun).    We have shown that the polarization of starlight can be a probe of the magnetic field near the GC. \\citet{Morris98} enumerated five different ways in which the magnetic field near the GC has been studied: morphology, polarization angle, Faraday rotation of the radio continuum, Zeeman effect, and polarized dust emission in FIR/sub-mm wavelengths. The wide field-of-view of the NIR polarimeter SIRPOL, and the statistical treatment of tens of thousands of stars enable us to study the magnetic field near the GC; that is, the NIR polarimetry of starlight is a new way to investigate the magnetic field {\\it in} the GC.   NIR polarimetry has the advantage of providing information about the magnetic field at locations where FIR/sub-mm emission is weak. NIR polarization of starlight is attributed to extinction along the line of sight by aligned dust grains, while FIR/sub-mm polarization is due to emission from the aligned dust. Hence, NIR polarimetry can investigate the magnetic field in regions where FIR/sub-mm polarimetry is absent, if background stars exist. This advantage is clearly shown in Fig. 2. We have detected polarization at positions where thin bars are not shown. The distribution of the position angles in most of the observed regions including such low emission regions shows a globally toroidal magnetic configuration at the GC."
    },
    "0809/0809.2305.txt": {
        "abstract": "We provide estimates of volcanism versus time for planets with Earth-like composition and masses 0.25 - 25 {\\Me}, as a step toward predicting atmospheric mass on extrasolar rocky planets. Volcanism requires melting of the silicate mantle. We use a thermal evolution model, calibrated against Earth, in combination with standard melting models, to explore the dependence of convection-driven decompression mantle melting on planet mass. We show that (1) volcanism is likely to proceed on massive planets with plate tectonics over the main-sequence lifetime of the parent star; (2) crustal thickness (and melting rate normalized to planet mass) is weakly dependent on planet mass; (3) stagnant lid planets live fast (they have higher rates of melting than their plate tectonic counterparts early in their thermal evolution), but die young (melting shuts down after a few Gyr); (4) plate tectonics may not operate on high mass planets because of the production of buoyant crust which is difficult to subduct; and (5) %Per unit mass, rates of volcanism vary only by a factor of three of each other on planets $<$ 3 Gyr %old and $>$ 1 {\\Me}, but a stronger mass dependence develops for older planets. Melting in plate-tectonic %mode ceases when potential temperature falls below the zero-pressure solidus, or when plates arrive at %subduction zones with densities far lower than the mantle. The second condition disfavors plate tectonics on %planets $\\gtrsim$ 4 {\\Me}. In the stagnant-lid mode, internal temperatures are higher, and melting initiates at higher %pressures, but melt production rates can be lower because convecting mantle does not approach the surface. %Decompression melting of non-plume %mantle will cease on stagnant-lid planets after $<$ 9 Gyr, regardless of size. melting is necessary but insufficient for efficient volcanic degassing -- volatiles partition into the earliest, deepest melts, which may be denser than the residue and sink to the base of the mantle on young, massive planets. Magma must also crystallize at or near the surface, and the pressure of overlying volatiles must be fairly low, if volatiles are to reach the surface. If volcanism is detected in the $\\tau$ Ceti system, and tidal forcing can be shown to be weak, this would be evidence for plate tectonics. ",
        "introduction": "Theory predicts the existence of rocky planets having 1-10 Earth masses (e.g. \\citet{ida04}). Planets in this mass range are now being detected \\citep{riv05}, and next-decade observatories, such as \\emph{James Webb Space Telescope} and \\emph{Giant Magellan Telescope}, may be able to detect any atmospheres. %It is even possible that near-term coronographic instruments (Cheops at ESO, Gemini Planet Imager) could succeed in this goal. A planet's atmosphere will consist of gas (1) accreted from the nebula, (2) degassed during impact accretion, and (3) degassed during subsequent geologic activity. (1) will depend on the lifetime of the nebula, (2) and (3) will depend on the volatile abundance of material \\citep{elk08a}, and all will be modified by atmospheric escape. Very large planets far from their parent star will retain primitive gas, but smaller planets closer to their parent star will not. %Such worlds, if massive enough, may retain a nebular (H-He rich) atmosphere. Degassing of an accretion-heated magma ocean %probably envelops the growing planet in a steam atmosphere. %But what if this %initial volatile envelope is lost, or never accumulates? Then atmospheric composition will reflect the balance %between volcanic degassing and atmospheric loss, as on Venus, Earth and Mars --\u0096 perhaps modified %by biology. Loss rates vary between gases, so planetary atmospheres could be a mixture of gases left over from the initial atmosphere, and those replenished by volcanism. %Arguably, we have a nearby example in Titan, whose nitrogen ($\\sim$95\\% of atmospheric mass) was %released early, but whose methane ($\\sim$5\\% of atmospheric mass) may have been %supplied by cryo-volcanoes over geological time \\citep{tob06}. Thermal emission phase-curves gathered from extrasolar planets can set bounds on atmospheric mass. Spectral observations of atmospheric constituents with short photochemical lifetimes, such as SO$_2$, would require an ongoing source --\u0096 most likely, volcanic degassing.  Volcanism results from partial melting of upper mantle. (Planetary mantles can cool convectively without volcanism - present-day Mercury is almost certainly an example). Partial melting occurs when the adiabat crosses the solidus. Assuming that the adiabat is steeper than the solidus, this requires that the potential temperature of the mantle $T_{p}$ exceed the zero--pressure solidus of mantle rock (e.g., peridotite) (Figure \\ref{PEDAGOG}a): \\begin{equation} T_p = T_{m,r} - P_{r}\\frac{\\partial V}{\\partial S} \\ge  T_{sol}(0), \\end{equation} where $T_{m,r}$ is the mantle temperature evaluated at some reference pressure $P_{r}$, \\(\\frac{\\partial V}{\\partial S}\\) is the adiabat, potential temperature is defined as the temperature a parcel of solid mantle would have if adiabatically lifted to the surface, and $T_{sol}$ is the solidus, evaluated here at zero pressure. On planets with plate tectonics, the thickness of the crust (the crystallized melt layer) is a convenient measure of the intensity of volcanism. The rate of crust production is the product of crustal thickness, plate spreading rate, and mid-ocean ridge length. The pressure at the base of the crust is the product of crustal thickness, the planet's surface gravity, and crustal density. A reasonable approximation to the pressure at the base of the crust is the integral of the fractional-melting curve from great depth to the surface (Figure \\ref{PEDAGOG}a).  Venus and Mars lack plate tectonics: their mantles are capped by largely immobile, so-called `stagnant lid' lithospheres, which cool conductively. Because mantle cannot rise far into the stagnant lid, melting can only occur if the temperature at the base of the stagnant lid exceeds the local solidus of mantle rock: \\begin{equation} T_{m,r} - (P_{r} -P_{lith}) \\frac{\\partial V}{\\partial S} \\ge  T_{sol}( P_{lith} ), \\end{equation} where $P_{lith} = \\rho_{lith} g Z_{lith}$ is the pressure at the base of the stagnant lid, $\\rho _{lith}$ is lithospheric density, $g$ is gravity, and $Z_{lith}$ is stagnant lid thickness. This is a more stringent condition than (1) if the adiabat is steeper than the solidus (Figure \\ref{PEDAGOG}b). Io shows yet another style of rocky-planet mantle convection: magma pipe cooling. %However, we find that this style is not important for massive Earth-like planets $\\ge$ 1 Gyr old.  Previous studies have examined both rocky planet atmospheres and massive-earth geodynamics. \\citet{elk08a} estimate the mass and composition of super-Earth atmospheres degassed during accretion. Complementary to that study, we emphasize long-term geological activity. There is disagreement over whether plate tectonics will operate on massive planets. One previous study uses scaling arguments to argue that higher gravity favors subduction \\citep{val07}. Another study shows that subduction might never begin if the yield stress of old plate exceeds the stresses imposed by mantle convection, which fall steeply with increasing planet mass \\citep{one07}. We do not consider yield stresses in this paper. Instead, we analyze five possible volcanism-related limits to plate tectonics, tracing the consequences in more detail than the short paper of \\citep{val07}. The approach of \\citet{pap08} is most similar to that taken here. We differ from \\citet{pap08} in that we neglect the pressure dependence of viscosity, use more realistic melting models, and account for energy advected by magma.  Here we examine (1) the history of partial melting on planets with either plate tectonics or stagnant lid convection, and (2) the effect of melting and crust production on the stability of plate tectonics. Our method is given in \\S 2. We use three different melting models to compute the intensity of volcanism for massive Earth-like planets of different ages. %These models involve many assumptions and uncertainties, listed in . Our results are given in \\S 3, for both plate tectonic (\\S 3.2) and stagnant lid (\\S 3.3) modes of mantle convection. We also trace the implications of galactic cosmochemical evolution for heat production and planetary thermal evolution (\\S 3.4). We find that results differ greatly depending on the mode of convection. Massive Earths with plate tectonics will produce melt for at least as long as the age of the Galaxy, stagnant lid planets will not. In \\S 4, we analyse the effect of melting on the style of mantle convection. We show that plate buoyancy is likely to be a severe problem, and may be limiting for plate tectonics. In \\S 5, we relate our results to atmospheric degassing, and discuss the possible suppression of degassing (and, perhaps, melting) by the higher ocean pressures expected on massive Earth-like planets. Finally, in \\S 6, we summarize our results; justify our approximations and model limitations; and compare our results to solar system data. ",
        "conclusions": "\\subsection{Summary of results}  \\begin{list}{$\\bullet$} { \\setlength{\\labelwidth}{2cm} } \\item \\emph{Modest effect of planet mass on temperature and crustal thickness:}  Scaling analysis predicts $T_{m}$ should only be weakly dependent on planet mass \\citep{ste03}. Because of the exponential dependence of viscosity on temperature, modest variation in $T_{m}$ can accommodate the range of heat flows generated by planets of varying masses. Our more detailed study confirms this, and although the nonlinearity of mantle melting leads to  greater variations in crustal thickness, these are still of order unity (\\S 3.2; \\S 3.3; Figure \\ref{CRUSTVSTIME}). However, planet mass can determine which mode of mantle convection the planet is in (\\S 4), which in turn determines whether melting is possible at all (\\S 3.3). \\item \\emph{Challenges for plate tectonics on massive Earths:} Plate buoyancy is an increasingly severe problem for plate tectonics as planet mass increases (\\S 4.5). Because Early Earth faced the same challenge, geologic fieldwork may determine the threshhold beyond which plate tectonics shuts down (\\S 4.5). Middle-aged super-Earths may suffer from continental spread, which could choke off plate tectonics (\\S 4.4). Shutdown of plate tectonics could place a planet in stagnant lid mode. \\item \\emph{Shutdown of melting on old, stagnant lid planets:}  All our melting models predict that stagnant lid planets planets older than Earth cease melting and volcanic activity (\\S 3.3-5, but see caveats in \\S 6.2). Therefore, observed volcanism on rocky planets $>$ 8 Gya would provide some support for plate tectonics -- if and only if tidal heating could be shown to be small. A candidate planetary system has been identified around $\\tau$ Ceti \\citep{mar02}, a 10 Gya star at 3.65 parsec, and may provide an early test of this prediction. \\item \\emph{Implications for the stability of atmospheres, climate, and habitability:} The atmospheres of Venus, Earth and Mars were produced by release of gases dissolved in partial melts, as well as temperature-dependent exchange of volatiles between the atmosphere and surface rocks. Planets lacking dynamos and on close orbits around their parent stars may experience high rates of atmospheric erosion due to stellar winds and coronal mass ejections \\citep{kho07}, and their atmospheres would persist on main sequence timescales only if maintained by mantle melting. Under more benign conditions, crust formation may establish the long-term carbon dioxide content of Earth-like atmospheres; CO$_2$ released in volcanism is sequestered as carbonates produced during the low-temperature weathering of that crust \\citep{wal81}. Volcanism may underpin the long-term habitability that the fossil record teaches us is a prerequisite for advanced life \\citep{but07}.  \\end{list} %Because of the exponential dependence of viscosity on temperature, modest variation in $T_{m}$ can %accommodate the range of heat flows generated by planets of varying masses. If plate tectonics %is assumed to be possible on massive Earth-like planets, %our model  confirms the expected, %relatively weak dependencies: crustal thickness decreases slightly with increasing planet mass, and %vertical tectonics is slightly favored over plate tectonics with increasing planet mass. For real planets, %these effects are convolved with the effects of varying cosmochemistry and initial radiogenic element %complement.  \\subsection{Overview of approximations and model limitations}  In this paper we assume whole-mantle convection. Layered mantle convection can alter the volcanic history of a planet by introducing long-term sensitivity to initial thermal conditions.  For a phase boundary to form a barrier to flow its Clapeyron slope must be strongly endothermic. Historically, the $sp \\to pv$ transition was thought to have this character but this is now known not to be the case. Deeper phase transitions are either exothermic ($pv \\to ppv$) or very gradual \\citep{sea07}. In spite of these arguments against layered mantle convection, barriers to flow are possible if stable deep mantle layers arise during or shortly after fractional crystallization of the primordial magma ocean. %\\citet{but02} recover Earth's thermal history %using a parameterization of $Ra$ dependent mantle layering from (2D) numerical models of %mantle convection. However, more complete (3D) models show that rare mantle avalanches overcome %the apparently high degree of layering in high-Ra models, which invalidates the proposed lower-mantle %buffering of upper-mantle temperature. Although CMB pressures reach 14 Mbar for 10 \\Me and 40 Mbar for 25 \\Me, we have not considered metalized silicates \\citep{ume06}, nor the possibility that lower-mantle convection is extremely sluggish (even isoviscous; \\citet{fow83}) due to low homologous temperatures, nor pressure dependent viscosity \\citep{pap08}. Our parameterized treatment of mantle convection assumes the solid state and is inappropriate for very high temperatures, when the greater part of the lithosphere is underlain by a magma ocean. For almost all cases, however, the transition to a magma ocean takes place at temperatures beyond the range of validity of our melting models, so magma ocean development is not the limiting consideration in interpreting our results. %In \\S 3-4 %we assumed that the pressure exerted by volatiles overlying the silicate layer is small. %Even for rock-dominated planets in the range 1-10 {\\Me}, late-stage accretion may deliver far %greater quantities of volatiles than are found in the accessible reservoirs of Earth \\citep{ray06}. %We explored the consequences of relaxing the small-overburden-pressure constraint in \\S 5.2.  Our model does not consider tidal heating, which may be important for Earthlike planets on close eccentric orbits about low-mass stars \\citep{jac08}, nor does it take into account the energetics of the core \\citep{nim07}. Because the magnitude of core cooling is set by mantle cooling, core energetics only dominate mantle thermal evolution in Earthlike planets as $t\\to\\infty$ and $H\\to$ 0. Surface temperature is assumed to be constant, which is likely to be a good approximation except for close-in exoplanets with thin atmospheres in a 1:1 spin-orbit resonance \\citep{gan08}.  All our melting models show `solidus rollover' at high P - that is, $\\frac{T_{sol}}{P} \\to \\frac{\\partial V}{\\partial S}$. As a consequence, plots of $X$ versus $P$ at high $T_p$ have a long, thin high-pressure tail, and the pressure of first melting diverges at finite $T_p$. Although solidus rollover is a real effect, resulting from the greater compressibility of silicate melts versus solids \\citep{ghi04}, infinite pressures of first melting are unphysical. In particular, phase transitions near 14 GPa introduce unmodelled kinks of the solidus. To sidestep the solidus rollover problem, we truncate our melting integration at 8 GPa, and do not plot model output where melting models predict nonzero $X$ at 8 GPa. But this solution has costs: 1) we cannot model hot planets (grayed-out regions in Figures \\ref{CRUSTVSTIME} \\& \\ref{CRUSTVSTIMESL}), and 2) because pressure at first melting is so sensitive to solidus rollover, we have little confidence in our model output for the rate of processing of the mantle through the melt zone. Melting models calibrated to higher pressures and temperatures would remove this impediment to progress. One such model (xMELTS), based on a high-pressure equation-of-state for silicate liquids \\citep{ghi04}, will soon become available \\citep{ghi07}. Because crustal production rate is proportional to the integral under the curve of $X$ versus $P$, rather than the pressure at which $X = 0$, it is much less sensitive to this problem. Representing the solidus as a straight line in T-P space \\citep{pap08} is geologically not realistic, but has the advantage that both crustal thickness and the pressure of first melting are finite for all $T_p$.  We calculate crustal thickness as the integral of melting fraction from the surface to great depth. In plate tectonic mode this corresponds to the observable (seismically-defined) crustal thickness. But in stagnant lid mode, the observable crustal thickness need not be the same everywhere (because melting will only occur above mantle upwellings). Also, the observable crustal thickness may be everywhere greater than the integral of melting fraction from the surface to great depth, because crust generated in previous melting events is located directly underneath `fresh' crust. This situation is described in the context of Io by \\citet{moo01}. %Therefore, the rate of volcanism (\\S 5) is a more meaningful quantity than crustal thickness for comparing plate-tectonic and stagnant-lid planets.  On stagnant lid planets, volcanism may be episodic (as on Mars; \\citet{wil01}) and absent for long periods. This means that an observed lack of compounds with a short photochemical lifetime would not preclude a high level of volcanic activity averaged over a sufficiently long period ($\\ge 10^{\\sim 8}$ yr). Volcanism is not the only possible source of degassing. Metamorphic decarbonation \\citep{bic96}, and the episodic release of deep-seated volatiles as on the Moon \\citep{gor73}, could permit (low) rates of degassing on non-volcanic worlds.  By assuming that all melt reaches the surface, we have ignored the distinction between extrusion and intrusion. However, extrusion can alter the reflectance properties of the planet without sustaining an atmosphere -- consider the Moon -- and it may eventually be possible to discriminate between the two possible fates for melt.  \\subsection{Comparison with Solar System data}  \\subsubsection{Thermal evolution, with an emphasis on Earth}  Four decades since plate tectonics became widely accepted, geologists have not determined the detailed thermal history of the Earth, neither have we even produced a model that satisfyingly accounts for the few data points in hand. The model of Korenaga (2006, taking a parameterized approach) may be an exception, but this model assumes that present-day observed averages are valid for the past, and is probably not portable to other Earth-like planets. Complicating the picture are two sharp-edged lower mantle anomalies mapped by seismic tomography under Africa and the S. Pacific, which are likely distinct in composition and radiogenic-element density from the rest of the mantle, are probably not sampled in our collections of mantle rocks brought up by volcanoes and which may be persistent, ancient features that have interfered in mantle convection since more than 2.5 Gyr ago. Their origin is unknown \\citep{gar08}. Geoneutrino detectors will break some of these degeneracies before 2020, and continued field mapping of ancient tectonic features will sharpen the history of past dynamics of plate deformation \\citep{con08}. But using simple models and making changes along one dimension only -- the approach taken here -- is still a meaningful approach because the more complex models introduce many more free parameters, and it is not clear how these free parameters vary with mass. By contrast, our approach involves only one `tuning' parameter, the offset $T_m - T_p$.  The situation for the other terrestrial planets -- single plate planets with stagnant lithospheres - is in many ways better. Parameterized cooling fits all the observations. This is not just because of a lack of data: stagnant lid planets seem to be truly simpler, and the relevant fluid physics is probably better understood than the `viscoelastoplastic plus damage' phenomenology needed to accommodate the observed geological motions of the Earth's plates. On Earth, mantle heat loss is largely by conduction aided by hydrothermal convection at mid-ocean ridges, but mantle heat loss on stagnant lid planets should be less sensitive to the details of surface deformation.  The degree to which the mantle cools over geological time depends on the activation energy for deformation of mantle rock, which is known only imprecisely \\citep{kar08}. We take $T_{\\nu} \\cong$  43 K, for which $A_0$ = 7 x $10^4$ K, giving an order-of-magnitude viscosity change for each 100K temperature increment (although $T_{\\nu}$ could be as high as 100K) \\citep{sle07}. This is conservative (but not neccessarily more correct) in terms of the effect of mass on temperature, because small values of $T_{\\nu}$ allow mass (heat flux) variations to be accommodated by small changes in temperature (Figure \\ref{TNU}).  The range of mantle states that permit plate tectonics to continue will differ from the range of states permitting plate tectonics to initiate, and ongoing plate tectonics will alter the state of the mantle.  %Because in plate tectonics mode melting depends on $T_{p}$ alone, we can relate the behaviour of large planets (of comparable age to the Earth) %to that of earlier versions of the Earth. For example, a 4.5 Gyr-old 10 {\\Me} planet %behaves like a 1.5 Gyr-old 1 {\\Me} planet, apart from corrections for gravity (which thins the melting zone and thus the crust) %and surface area (the more massive planet will have a larger gross rate of plate production). This means that early Earth %data is useful in understanding the geodynamics of massive Earths (\\S 6.2).  Continents affect thermal evolution (\\S 4.4). However, continent growth and survival is not well understood. Hydration of the oceanic crust at mid-ocean ridges, and subduction of this H$_2$O, is probably required for continental growth \\citep{cam83}. There is geological and isotopic evidence for episodic continental growth \\citep{con06}. Durable continents may require photosynthetic life both in the oceans \\citep{ros06} and on land \\citep{aeo08}. None of these effects are straightforward to model. So, though excessive continental surface area and radiogenic-element sequestration both have the potential to throttle our predicted increase in mantle melting on high-mass Earths, we consider this only a tentative prediction.  \\subsubsection{Melting: constraints from Venus, Mars, Io, and the Moon} If massive Earth-like planets are in heat-pipe mode for $\\sim$ 1 Gya (\\S 4.2), a substantial fraction of the census of nearby massive Earth-like planets would be in heat-pipe mode. And if Io is any guide (e.g., \\citet{lop07}), dramatic spatial and temporal variations in thermal emission should be present. However, the relatively slow cooling in our thermal evolution model is inappropriate for magma oceans \\citep{sle07}, because our parameterization of viscosity does not include the fifteen-order-of-magnitude decrease associated with melting. Planets with $Mo > 1$ are likely to cool quickly to $Mo < 1$; we suspect that Io-type cooling is unsustainable for geologically significant periods without a nonradiogenic source of energy.  Our melting model is tuned to Earth only, but fits Venus observations reasonably well. With the MB88 model, our calculations show that volcanism on a planet in stagnant lid mode and with Venus' mass ({\\Me} = 0.85) will cease after 3.9 Gyr, which is consistent with observations \\citep{sch92}.  However, when decompression melting of mantle material at background temperatures cannot sustain volcanism on Earth-like planets undergoing whole-mantle convection, other mechanisms may take over. Plumes rising from the core-mantle boundary, or from a compositional interface within the mantle, can have temperatures up to several hundred degrees greater than that of passively upwelling mantle. This requires a more detailed treatment of thermal evolution than we have taken in this paper, and could be the basis of a more detailed study.  In particular, Mars' continuing --\u0096 although fitful and low-rate \u0096-- volcanic activity \\citep{bor05} indicates that volcanic activity can continue on planets in stagnant-lid mode for longer than our simple models predict, perhaps as the result of plumes in a compositionally-layered mantle \\citep{wen04} or a thick, thermally insulating crust \\citep{sch07}. On Venus, mantle fluxing by sinking delaminated lithosphere has been proposed as a mechanism that would allow volcanism to continue through to the present day \\citep{elk07}, although there is no robust evidence of ongoing volcanism. Continued Solar System exploration is needed if we are to fully exploit nearby data points to understand the geodynamic window for life.  %%It may be possible for otherwise habitable planets to enter an entirely ice-covered (snowball) state as a consequence %%either of evolutionary innovations \\citep{kop05} or stochastic variations in climate. The ability of a planet to %%eventually exit the snowball state depends on the rate of degassing \\citep{pie05}.  %On the Earth, the relative abundances of the rare gases and their %isotopes indicate the loss of an earlier (although not necessarily %primordial) atmosphere and its replacement by an atmosphere sustained % by slow degassing of the mantle over time.  The persistence %of this atmosphere and its modification by chemical reaction %with surface rocks and fluids are considered critical to the %maintenance of ``habitable'' conditions on the Earth over a %significant fraction of the age of the Sun. Since the appearance %of free molecular oxygen in the atmosphere 2.4 Ga it is thought that %carbon dioxide has been the leading greenhouse gas. A negative %feedback involving weathering of silicate rocks and sequestration of %carbon dioxide in carbonate rocks has been invoked to explain the (almost?) continuous stability of surface liquid water, % which would otherwise be hard to maintain in the face of a brightening Sun \\citep{Walker81}. % (The alternative to a negative feedback is that Earth's apparent climate stability is the result of chance.) % Partial melting is required for the silicate-weathering feedback both by providing the CO$_2$ that is degassing %and by producing fresh volcanic crust for weathering.  %On the Earth, plate tectonics, via degassing at active plate margins, subduction of carbon-bearing rocks, %and exposure of crustal rocks in mountain belts, is intimately associated with %the long-term geologic carbon cycle \\citep{hay06}.  However, the absence of plate %tectonics on other, otherwise Earth-like planets need not %preclude the operation of a carbonate-silicate %feedback nor habitable conditions, i.e. stable liquid water at the %surface.  At the minimum, such a condition requires that the rate of %silicate weathering and production of alkalinity at surface %temperatures above freezing balances outgassing of CO$_2$ from all %sources, and that the temperature dependence of weathering is positive. %On a planet with plate tectonics, the first requirement is expressed as: %\\begin{equation} %f_{mantle} + f_{arc} + f_{meta} = f_{weathering} %\\end{equation} %where the terms are, respectively:-- the flux of carbon dioxide to/from the %atmosphere/ocean reservoir due to mantle melting and degassing; %arc volcanism, partially recycling subducted carbon; metasomatism of carbonate rocks; and carbonate deposition due to the %weathering of silicate rocks and production of alkalinity.  In the %absence of plate tectonics and any mountain-building, the only non-zero terms are the right hand side, and the first term on the LHS. %A stagnant lid planet is less capable of %recycling carbon dioxide back into the atmosphere from its crustal %reservoir.  We also predict that SL planets the same mass and age of %the Earth will have less mantle melting - or none at all.  (We have also %shown that younger SL planets will have {\\it higher} rates of mantle %melting.).  %For example, estimated CO$_2$ fluxes on the present Earth are 2.2, %2.5, and $<3$ for mid-ocean ridge, arc, and plume volcanism, %respectively, in units of $10^{12}$ mol a$^{-1}$ \\citep{Marty98}. %%Metamorphic degassing from orogenic belts is $3 \\times 10^{12}$ mol %%a$^{-1}$ \\citep{Kerrick98}. % On an SL planet of identical age and %mass, mantle degassing should either have ceased entirely (using the MB88 model) or %be ~1/3 that of the present Earth (using the K03 model; Figure \\ref{RATERATIO}). (However, earlier in their evolution, stagnant lid planets have %rates of volcanism/degassing greater than that of contemporary plate tectonic planets; Figure \\ref{RATERATIO}). % %We schematically show the balance between degassing of CO$_2$ and %weathering of silicates in Figure \\ref{CARBONCYCLE}.  Weathering rates in the %laboratory have an exponential temperature dependence. %%, with some %%activation energy that depends on the specific reaction. %The mean %surface temperature is set by the balance of degassing and weathering; %if temperatures increase (decrease), weathering exceeds (falls below) %degassing and the CO$_2$ content of the atmosphere and oceans %decreases (increases).  Ultimately, silicate weathering requires the %production and exposure of igneous rocks.  Rocks formed from the %weathering and alteration of igneous rocks are less fertile sources of %alkalinity (e.g., quartz sandstone) during subsequent weathering.  At %sufficiently high rates of degassing and sufficiently high surface temperatures on a wet %planet, weathering may be limited by the availability of fresh igneous %rock, rather than chemical rates of reaction. In this case, the rate %of weathering becomes temperature-independent, the silicate weathering %``thermostat'' fails, and climate may become unstable and enter a %runaway wet greenhouse state. However, seafloor weathering may %provide a base level of weathering that is relatively independent of surface %temperature, but proportional to the rate of mid-ocean ridge volcanism %\\citep{Brady97}. % %Therefore, all else being equal, massive stagnant-lid planets of comparable age to Earth will experience lower %rates of silicate weathering (Figure \\ref{CARBONCYCLE}).  Because of a lower rate of %igneous rock production, the weathering curve saturates at a lower %temperature (Figure \\ref{CARBONCYCLE}).  These considerations suggest that massive stagnant-lid planets that are Earth's contemporaries %will have a colder climate than those with plate tectonics, but that %the silicate weathering feedback can still operate. Once volcanic degassing ceases, however, no process can balance long-term CO$_2$ drawdown, and plants %would fast become extinct \\citep{cal92}.  %Stagnant-lid planet climate systems operating at a lower temperature will %necessarily have less CO$_2$ in their atmosphere.  As stellar %luminosity rises the point at which there is no CO$_2$, and at which surface %temperatures begin to rise, is reached earlier.  However, the %evolutionary trajectories of the atmospheres of SL and PT planets will %converge once both planets have no atmospheric CO$_2$.  \\bigskip \\bigskip  Many of the ideas discussed here are drawn from the conversation of Dave Stevenson, whose support for the early stages of this project was invaluable. We thank Nick Butterfield, Rhea Workman, Brook Peterson, Norm Sleep, and Itay Halevy for their productive suggestions. We acknowledge support from the NASA Astrobiology Institute. E.K. acknowledges support from the Berkeley Fellowship and a Caltech Summer Undergraduate Research Fellowship.  %% Appendix material should be preceded with a single \\appendix command. %% There should be a"
    },
    "0809/0809.2087_arXiv.txt": {
        "abstract": "In recent years several hypervelocity stars (HVSs) have been observed in the halo of our Galaxy. Such HVSs have possibly been ejected from the Galactic center and then propagated in the Galactic potential up to their current position. The recent survey for candidate HVSs show an asymmetry in the kinematics of candidate HVSs (position and velocity vectors), where more outgoing stars than ingoing stars (i.e. positive Galactocentric velocities vs. negative ones) are observed. We show that such kinematic asymmetry, which is likely due to the finite lifetime of the stars and Galactic potential structure, could be used in a novel method to probe and constrain the Galactic potential, identify the stellar type of the stars in the survey and estimate the number of HVSs. Kinematics-independent identification of the stellar types of the stars in such surveys (e.g. spectroscopic identification) could further improve these results. We find that the observed asymmetry between ingoing and outgoing stars favors specific Galactic potential models. It also implies a lower limit of $\\sim54\\pm8$ main sequence HVSs in the survey sample ($\\gtrsim648\\pm96$ in the Galaxy), assuming that all of the main sequence stars in the survey originate from the Galactic center. The other stars in the survey are likely to be hot blue horizontal branch stars born in the halo rather than stars ejected from the Galactic center. ",
        "introduction": "Hypervelocity stars (HVSs) are stars with extremely high peculiar velocities relative to the velocity distribution of their parent population. In recent years several HVSs have been observed in the Galactic halo, some of them unbound to the Galaxy (with velocities beyond the escape velocity;\\citealp{bro+07b}). From these a Galactic population of $96\\pm10$ such HVSs was inferred \\citep{bro+07b}, up to the $100$ kpc distance limit of the survey. Many similar bound HVSs (with velocities lower than the escape velocity) have been observed at larger numbers. Most of the observed HVSs are B-type stars (\\citealt{bro+05,bro+06a,bro+06b,bro+07a,bro+07b,ede+06}; future observations of HVSs of other stellar types are discussed in \\citealt{kol+07,bro+08,ken+08}). Given the color selection of the targeted survey for these stars \\citep{bro+06a}, such stars could be either main sequence (MS; or blue straggler) B stars or hot blue horizontal branch (BHB) stars. Currently only three of the stars in the survey have specific unambiguous identification and were found to be MS stars \\citep{fue+06,lop+08,prz+08b}.  Extreme velocities as found for these stars most likely suggest a dynamical origin from an interaction with or close to the massive black hole (MBH) in the Galactic center \\citep[GC; ][]{hil88,yuq+03,lev05,ole+07,per+07}. In the following we discuss only HVSs ejected from the GC% \\footnote{\\label{fn:Disk HVSs}Note that recent observations possibly suggest a different origin for two of HVSs observed serendipitously (\\citealp{bon+08,prz+08,heb+08}, but see \\citealp{per08b}). Throughout this paper we assume that most (if not all) of the MS stars in the HVSs survey sample have a GC origin. We note however that the methods we describe here to constrain the Galactic potential could similarly be used, in principle, for the analysis of high velocity stars ejected from the Galactic disk. % } and observed in the Galactic halo, $>10$ kpc from the GC, which would require ejection velocities from the GC of $\\gtrsim800$ km s$^{-1}$. Such HVSs could serve as probes of the GC environment, stellar population and dynamics (see e.g. \\citealp{ses+07a,ken+08,per08a}) and serve as an independent evidence for the existence of a MBH in the Galactic center \\citep{hil88}. Most of the B-type stars observed through the HVSs survey have lower velocities and are either bound HVSs \\citep{bro+07a} or are just halo stars. The survey shows an asymmetry in the kinematics of the stars, where more stars have positive radial velocities in Galactocentric coordinates (i.e. outgoing stars) than negative ones (ingoing or returning stars). As we show this asymmetry is dependent on both the absolute velocity and the distance of the stars from the GC, and could be used in a novel method to probe and constrain the Galactic potential, identify the stellar type of the stars in the survey and estimate the number of HVSs.  This paper is organized as follows. We first briefly describe the HVSs survey (\\S \\ref{sec:vel-dist}), and then suggest a novel method to probe the Galactic potential using such surveys (\\S \\ref{sec:HVS-potential}). In \\S \\ref{sec:HVS-lifetime} we use a similar analysis to infer the statistics of the stellar type of stars in the HVSs survey and estimate a lower limit to the number of HVSs in the Galaxy. ",
        "conclusions": "\\label{sec:summary}  In this paper we studied the characteristics, the origins and the use of stars observed in the survey for HVSs in the Galactic halo. The kinematics of currently observed HVSs ejected from the Galactic center depend strongly on their lifetimes and their propagation in the Galactic potential. We suggest a novel method to probe the Galactic potential up to large distances using the kinematics and the spectral identification of HVSs. We also use a reverse method, where a specific Galactic potential model is assumed, to give lower limit estimates on the number of HVSs ejected from the Galactic center. Future observations of HVSs in M31 \\citep{she+07} and in other Galaxies (Perets et al., in preparation) could have a similar use for studies of Galactic structures and potentials."
    },
    "0809/0809.5172_arXiv.txt": {
        "abstract": "Motivated by the observed shortfall of baryons in the local universe, we investigate the ability of high resolution cosmic microwave background (CMB) experiments to detect hot gas in the outer regions of nearby group halos. We construct hot gas models with the gas in hydrostatic equilibrium with the dark matter and described by a polytropic equation of state. We also consider models that add entropy to the gas in line with constraints from X-ray observations. We calculate the thermal Sunyaev-Zel'dovich (SZ) signal in these halos and compare it to the anticipated sensitivities of forthcoming SZ survey experiments such as ACT, PLANCK and SPT. Using a multi-frequency Wiener filter we derive SZ detectability limits as a function of halo mass and redshift in the presence of galactic and extragalactic foregrounds and the CMB. We find that group-sized halos with virial masses below $10^{14} M_\\odot$ can be detected at $z ~\\lsim~ 0.05$ with the threshold mass dropping to $3-4\\times 10^{13} M_\\odot$ at $z~\\lsim~0.01$. The SZ distortion of nearby group-sized halos can thus be mapped out to the virial radius by these CMB experiments, beyond the sensitivity limits of X-ray observations. These measurements will provide a unique probe of hot gas in the outer regions of group halos, shedding insight into the local census of baryons and the injection of entropy into the intragroup medium from non-gravitational feedback. ",
        "introduction": "\\label{section:intro}  Through the distinctive signatures that cosmological parameters, such as the baryon density, matter density and spatial curvature have on the cosmic microwave background (CMB) anisotropy, recent measurements of the CMB temperature and polarization spectra \\citep[][and references therein]{nolta_2008} have placed precise constraints on these parameters, thereby establishing a cosmological model that is consistent with a wide range of astronomical observations \\citep[][and references therein]{dunkley_2008}. Moreover, because the physics underlying the CMB anisotropy is to a good approximation linear on these scales, the CMB provides a powerful probe into the physics of the early universe and the primordial perturbations \\citep[][and references therein]{komatsu_2008}.  Whereas on large angular scales the CMB anisotropy is fairly well described, apart from a late-time integrated Sachs-Wolfe (ISW) contribution \\cite{sachs_wolfe_1967}, by the geometrical projection of inhomogeneities at the last scattering surface onto our celestial sphere, the so-called primary anisotropies, on smaller angular scales of around an arcminute CMB photons are more significantly affected by gravitational and scattering processes along the line of sight. These secondary anisotropies provide an effective probe of the physics of the low-redshift universe.  The most significant secondary anisotropy on scales of around a few arcminutes comes from inverse Compton scattering of CMB photons off hot gas in galaxy clusters along the line of sight, the so-called thermal Sunyaev-Zel'dovich (SZ) effect \\cite{sunyaev_zeldovich_1970}. The population of hot electrons in the cluster gas imparts energy to photons in the Rayleigh-Jeans part of the spectrum pushing them into the Wien tail. This distorts the thermal CMB spectrum by creating a deficit of photons at low frequencies and an excess of photons at high frequencies with no distortion at the SZ null frequency, $\\nu \\simeq 218$ GHz. The galaxy cluster appears as a cold spot on the sky at frequencies below the SZ null and as a hot spot above the SZ null. The unique spectral signature of the SZ effect will allow multi-frequency, high-resolution CMB experiments to make a relatively clean separation of the thermal SZ effect from the primary CMB and other contaminants such as point sources and galactic foregrounds.  The utility of the thermal SZ effect as a cosmological probe arises from the fact that the Compton distortion of the CMB is a scattering effect, so that the central decrement towards a cluster is independent of the redshift of the cluster. Moreover the angular diameter distance flattens out at $z\\sim 1$ so that the angular size of the cluster is approximately constant at high redshift. This means that the SZ effect does not suffer from the strong redshift dimming of its optical and X-ray counterparts with the consequence that a microwave background SZ cluster survey can detect a larger proportion of lower mass clusters out to higher redshifts. The thermal SZ effect can thus be used to construct galaxy cluster catalogs with a well defined selection function, thereby providing a potentially powerful probe of cosmology through the evolution of the cluster abundance (for a review see \\cite{carlstrom_2002}).  While most work has focussed on the detection of the SZ effect in galaxy clusters and its cosmological application, there is an intriguing possibility that high resolution CMB experiments with sufficient sensitivity may be able to detect hot gas in lower mass galaxy or galaxy group halos. Whereas the central SZ distortion scales roughly in proportion to the mass of the halo, the total SZ flux, $S_{\\rm SZ} \\propto M^{5/3}/d_A^2,$ has an additional dependence on the angular diameter distance of the halo. This means that nearby low mass halos with a large angular size could produce a significant SZ flux. This fact was exploited in \\cite{taylor_2003} to study the possibility of detecting the SZ effect in the halo of our neighbouring galaxy, M31, with the upcoming PLANCK surveyor. In this paper we focus on how well current and forthcoming CMB experiments will detect the SZ effect in nearby galaxy and group halos. Such a detection will provide unique insights into the distribution of baryons and the properties of hot gas in the outskirts of galaxy and group halos.  Indeed, in a census of baryons in the local universe \\cite{fukugita_1998, fukugita_peebles_2004} it has been argued that most of the uncertainty in the baryon budget comes from the uncertainty in the mass in ionized plasma associated with groups of galaxies.  Whereas the baryon fraction inferred from light element abundances (e.g., \\cite{steigman_2007}), the CMB anisotropy \\cite{dunkley_2008} and the high redshift Ly$\\alpha$ forest \\cite{kirkman_2003} give a consistent baryon fraction of $\\Omega_b \\approx 0.044,$ only about ten percent of these baryons are observed to be in stars, hot gas in clusters, and neutral and molecular gas in galaxies at low redshift. Cool plasma in and around groups of galaxies indirectly detected through quasar absorption lines \\cite{penton_2000, penton_2004} makes up twenty to twenty five percent of the baryon budget, while the warm-hot gas residing in the filamentary large-scale structure is believed to account for another thirty to thirty five percent of the baryons \\cite{dave_2001, cen_ostriker_2006}. The remaining thirty to forty percent of baryons is likely to be tied up in galaxy groups in a low density, warm-hot plasma.  Recent X-ray observations \\cite{Sun_2008} that report average gas fractions, $f_{\\rm gas} \\approx 0.08,$ out to $r_{500}$ in galaxy groups may have detected roughly half of these baryons. When combined with the fact that the gas fraction profiles inferred from these measurements are still rising at $r_{500}$ (see also \\cite{Vikhlinin_2006}) this indicates that a large fraction, perhaps as much as fifteen to twenty percent, of baryons in the universe may be in the outskirts of galaxy halos and in the intragroup medium beyond $r_{500}$. Recent evidence for this warm-hot component has come from quasar absorption lines in the ultraviolet (O VI; e.g., \\cite{tripp_2000, sembach_2003, tumlinson_2005}) and in X-rays (OVII and OVIII; e.g., \\cite{fang_2003}) observed around galaxies and groups of galaxies but these detections only map the two-dimensional gas distribution along a few lines of sight to allow a complete study of its spatial properties.  More detailed observations of this warm-hot component will yield important clues into the impact of galactic feedback on the distribution of gas in the intragroup medium. The relative dearth of hot gas observed in the central regions of galaxies and galaxy groups results in these halos being less X-ray luminous for a given temperature as compared to more massive galaxy clusters. The resulting break in the luminosity-temperature relation observed in X-ray clusters and groups \\cite{McCarthy_2004} suggests that non-gravitational processes that modify the entropy structure in the central regions of low mass clusters and groups are responsible for departures from the cluster scaling relations expected in self-similar models. Several models have been proposed to explain the excess entropy (for a recent review see \\cite{Voit_2005}), including pre-heating of infalling gas before it was shock-heated to the virial temperature, radiative cooling of low entropy gas to form stars, and feedback from supernovae and active galactic nuclei that increased the temperature and reduced the density of gas in the central regions of galaxies and groups.  While a combination of radiative cooling and feedback seems to reproduce the central entropy excess in low mass clusters and groups, it is clear that improved measurements of the properties of hot gas in these halos are necessary to understand the exact details of entropy injection. Whereas X-ray observations lack sensitivity to the hot gas beyond $r_{500}$ because the X-ray luminosity scales as the square of the (decreasing) gas density, the SZ effect scales linearly with the density and thus provides a potentially more sensitive probe of the hot gas in the outskirts of galaxy and group halos. Moreover, the SZ effect measures the projected gas pressure which is a complementary probe to the X-ray luminosity and can provide new constraints on models of the intracluster and intragroup medium.  While the SZ effect has been detected in a number of galaxy clusters by single-dish and interferometric experiments (for a review see \\cite{carlstrom_2002}) it has not not yet been convincingly detected in any galaxy groups. However, a new generation of high resolution CMB experiments with superior detector sensitivity offers the hope of detecting and constraining the properties of the hot gas in galaxy groups.  The outline of this paper is as follows. In Section 2 we construct models of the hot gas in virialized halos of a given mass and redshift, over the mass range $M_{\\rm vir}=10^{13} M_\\odot - 10^{15} M_\\odot$ and redshift range $z=0 - 1.5$.  The models take into account constraints on the temperature, entropy injection and gas fraction from X-ray observations. In Section 3 we utilise multi-frequency filtering techniques to determine the detectability of hot gas in nearby halos with current and forthcoming microwave background experiments such as the Atacama Cosmology Telescope (ACT~ \\cite{act_kosowsky_2006}), the South Pole Telescope (SPT~ \\cite{spt_ruhl_2004}) and the PLANCK mission \\cite{planck_collab_2006}. In a final section we summarize our findings and discuss how detection of the SZ effect in group halos will constrain gas models and allow a measurement of physical parameters such as the entropy injection and baryon fraction in these halos. Throughout this paper we assume a flat $\\Lambda$CDM cosmological model with parameters $h=0.72,~\\Omega_m=0.26, ~\\Omega_b=0.044, ~\\Omega_{\\rm de}=0.74, ~w_{\\rm de}=-1.0,$ and $\\sigma_8 = 0.8$ that provide a best fit to the WMAP 5 year data \\cite{dunkley_2008}. ",
        "conclusions": "We have investigated the detectability of SZ groups using an analytic prescription for the hot gas in these halos. The models that we studied were based on hot gas being in hydrostatic equilibrium with the dark matter halo, and described by a polytropic equation of state, or an equation of state modified to include an entropy injection term. We have found that the entropy models are distinguishable from the polytropic models via measurements of their SZ distortion, even in low mass clusters and galaxy groups. While these models provide a useful starting point to evaluate the detectability of SZ groups, an improved analysis will include a more realistic treatment of the gas distribution as provided by high resolution cosmological simulations, which we intend to pursue in a forthcoming paper.  Another issue that we have only partially addressed here, through the inclusion of an SZ background contaminant, is the confusion caused by the superposition of SZ distortions from hot gas in halos along the line of sight (see for e.g., \\cite{holder_2007, hallman_2007b}). Here again a large volume cosmological simulation will help to quantify the impact of the SZ background on the detection of group halos and their recovered flux. By taking advantage of the fact that nearby group halos produce a more extended SZ signal, we aim to mitigate the impact of the SZ background by devising algorithms to separate the group signal from the SZ emission at smaller angular scales produced by higher redshift clusters. Finally the combination of maps of the various foreground contaminants with simulated SZ maps will allow us to undertake a more accurate treatment of the foreground contamination. While we have been relatively conservative in our modelling of the foreground contaminants, we have not included effects such as the clustering of infrared point sources, which could turn out to be a significant contaminant in the extraction of SZ halos  \\cite{righi_2008}. We have also not included the kinetic SZ effect as a possible contaminant because it has the same frequency dependence as the primary CMB but a much smaller amplitude on the relevant angular scales. For the same reason we have not attempted to detect the halo via its kinetic SZ signal, though we note that a detection of the kinetic SZ signal could be enhanced by cross-correlation with optical or X-ray observations of the halo or the thermal SZ signal.  \\begin{figure}[t!] \\includegraphics[width=0.5\\textwidth]{\\plotsdir m_z_cut_ACT_with_data.ps} \\caption{The minimum detectable mass, $M_{\\rm vir},$ of SZ halos as a function of redshift, $z,$ for the ACT experiment compared to a compilation of groups from the USGC \\cite{ramella_2002} (plotted as crosses) and a galaxy group catalog based on the Sloan Digital Sky Survey (SDSS DR4)\\cite{yang_2007} (plotted as dots). The SDSS groups are only shown up to a redshift of $z=0.1$. The minimum detectable mass is plotted for an entropy model with $S/N=5$ (solid curve) and $S/N=3$ (dashed curve).} \\label{halo_detection_limits_vs_data} \\end{figure}  Prospects for detection of SZ clusters have been studied previously in the case of PLANCK \\cite{schaefer_2006, melin_2006}, ACT \\cite{pace_2008, sehgal_2007} and SPT \\cite{melin_2006}, but these studies have mainly focused on the statistics of SZ detections above the mass completeness limit of the respective surveys. Pace et al. \\cite{pace_2008} found that ACT could detect SZ halos down to $6\\times 10^{13} h^{-1} M_\\odot$ fairly independent of redshift, whereas we have found that these halos become hard to detect at higher redshifts. The analysis of \\cite{pace_2008} only included the CMB as a foreground so it is conceivable that the inclusion of point source foregrounds would degrade their forecasts for low mass halos at high redshift, as we have found to be the case in our analysis. We also note that the analysis of \\cite{schaefer_2006} found that PLANCK could detect halos of mass $6\\times 10^{13} M_\\odot/h$ below $z=0.1$ when they included the CMB and all galactic foregrounds. In our analysis we have emphasised that smaller halos, with $M \\simeq 3-4\\times 10^{13} M_\\odot,$ could be detected at $z~\\lsim~0.01$.  The detection of hot gas in these galaxy groups and low mass clusters with the upcoming set of SZ survey experiments will provide an interesting probe of galaxy formation and its effect on the distribution and state of the hot gas, which is most prominent in these halos. A measurement of the SZ distortion at the virial radius of these halos will set a joint constraint on the level of entropy injection and the baryon fraction, which can be compared to the predictions from galaxy formation models. In particular a measurement of the baryon fraction in the outskirts of galaxy groups and low mass clusters will provide a unique update to the baryon census in the local universe.  While the upcoming generation of SZ experiments have been designed to carry out blind surveys of galaxy clusters, a targeted survey of galaxy groups already detected in optical and X-ray observations would yield a very interesting set of objects to study. The measurement of a diffuse signal on the scale of tens of arcminutes will be challenging though, and carefully planned and executed observations will be necessary to control systematic effects and produce high fidelity maps of the SZ distortion in these halos. In combination with existing optical and X-ray observations of these halos, these measurements will enhance our knowledge of the physics of galaxy formation and its effect on the intragroup medium.  \\begin{figure}[!t] \\includegraphics[width=0.45\\textwidth]{\\plotsdir density_profiles.ps} \\caption{Electron number density profiles, $n_e$ ($10^{-3}\\mbox{cm}^{-3}$), plotted against $r/\\Rvir$ for the polytropic model (solid curves) and entropy model (dot-dashed curves). The panels from top to bottom show the density profiles for halo masses $M_{\\rm vir}=10^{13},10^{14}$ and $10^{15}~M_\\odot$ respectively.} \\label{fig:density_profiles} \\end{figure}  \\begin{figure}[!t] \\includegraphics[width=0.45\\textwidth]{\\plotsdir temperature_profiles.ps} \\caption{Electron temperature profiles, $T_e$ (keV), plotted against $r/\\Rvir$ for the polytropic model (solid curves) and the entropy model (dot-dashed curves). The panels from top to bottom show the temperature profiles for halo masses $M_{\\rm vir}=10^{13},10^{14}$ and $10^{15}~M_\\odot$ respectively.} \\label{fig:temperature_profiles} \\end{figure}"
    },
    "0809/0809.0427_arXiv.txt": {
        "abstract": "{The PICsIT detector onboard the INTEGRAL satellite was designed to provide information about emission in the soft $\\gamma$-ray band for many bright sources. Due to strong and variable instrumental background, only 4 objects have been detected so far using standard software.} {The moderate sensitivity of PICsIT can be compensated for in the case of many objects by adopting a long exposure time, thanks to INTEGRAL's large field of view and an observing strategy focused on the Galactic plane. With angular resolution far higher than that of all other instruments operating in a similar energy band, PICsIT is suitable for fields too crowded or too significantly affected by Galactic diffuse emission. Therefore, it is desirable to improve the spectral extraction software to both obtain more reliable results and enlarge the number of objects that can be studied.} {The new PICsIT spectral extraction method is based on three elements: careful modelling of the background, an energy-dependent pixel-illumination function, and the computation of the probability density  of the source count rate. Background maps of the detector plane are prepared for short periods and relatively narrow energy bands to insure that the background dependence on time and energy is modelled well. The most important element of the new spectral extraction method is the proper treatment of the Poisson-distributed data, developed within a Bayesian framework.} {The new method was tested extensively on both a large true data set and simulated data. Results assumed in simulations were reproduced perfectly, without any bias and with high precision. Count rates measured for Crab were far more stable than those obtained with the standard software. For weaker sources, the new method produced spectra of far higher quality and allows us to detect at least  8  additional objects. Comparison with other INTEGRAL instruments demonstrated that PICsIT is well calibrated and provides valuable information about the continuum emission in the 250 keV -- 1 MeV band, detectable currently only by INTEGRAL.} {} ",
        "introduction": "The low-energy $\\gamma$-ray band (100 keV -- 10 MeV) is extremely difficult to observe, mainly because of the intense nuclear-decay emission induced by cosmic rays inside the detector and surrounding materials. As the flux of the observed objects decline significantly with increasing energy from X-rays to $\\gamma$-rays, the strong increase in internal background of the detector makes the ratio of observed signal to background very small. This ratio usually does not exceed 0.01 even for sources as bright as the Crab nebula, despite the passive and active shielding applied. Moreover,  the fact that  soft $\\gamma$-rays cannot be focused as well as soft X-rays reduces the angular resolution in this energy range. Therefore, observations of crowded Galactic fields with $\\gamma$-ray telescopes are often affected by contamination from nearby sources or from Galactic diffuse emission.  There have been many $\\gamma$-ray observatories, but due to their generally limited life times and sensitivities only a few have provided a significant amount of data. More detailed studies of the soft $\\gamma$-ray spectra became possible in the last decade of the 20th century, thanks to the SIGMA detector \\citep{Paul91} onboard the GRANAT satellite, and the OSSE and BATSE detectors \\citep{Johnson93} onboard the CGRO satellite. SIGMA was a coded-mask instrument with a good angular resolution of 13$\\arcmin$, and its nominal energy range was 35--1300 keV. The OSSE telescope has been the most sensitive soft $\\gamma$-ray instrument sent into orbit until now, due to its large size. The high sensitivity of OSSE was in practice slightly reduced by limitations related to its low orbit. In addition, the large field of view of OSSE \\rm made it \\rm difficult to resolve closeby objects at the Galactic centre and to distinguish point-source emission from Galactic diffuse emission. BATSE was an all-sky monitor, used mainly for detecting GRBs and monitoring the activity of brighter $\\gamma$-ray sources. The nominal energy range of the PDS detector \\citep{Frontera97} onboard the BeppoSAX satellite was 15--300 keV but its sensitivity above 150 keV was insufficient to provide more precise data.  There are four missions presently in operation with soft $\\gamma$-ray instruments, namely RXTE (HEXTE), INTEGRAL, Swift (BAT), and Suzaku (HXD/GSO). The HEXTE \\citep{Rothschild98} and BAT \\citep{Gehrels04} energy ranges do not exceed $\\approx$150 keV, and these instruments have therefore limited capabilities of studying the high energy emission. GSO onboard Suzaku was declared to be the detector with the lowest internal background in the 40--600 keV band, due to a more advanced active-shielding technique \\citep{Takahashi07}. This low background should correspond to high sensitivity, but 3 years after Suzaku's launch almost no results above 200 keV have been published, apart from Cyg X-1 observed up to 400 keV \\citep{Makishima08}. Therefore, taking into account the limited angular resolution of GSO above 100 keV ($\\approx$ 4.5$\\degr$), current interest in the soft $\\gamma$-ray research is focused on the data acquired by INTEGRAL detectors.  The INTEGRAL satellite \\citep{Winkler03} has onboard three coded-mask $\\gamma$-ray detectors. ISGRI \\citep{Lebrun03} is the upper detector of INTEGRAL's imager IBIS \\citep{Ubertini03}, which operates nominally in the 13--1000 keV range. Owing to its relatively high sensitivity and angular resolution of 12$\\arcmin$, ISGRI remains the INTEGRAL instrument of choice for studies of point sources, particularly Galactic ones. The lower layer of the IBIS imager, the PICsIT detector \\citep{Labanti03}, is described in the next section. The spectrometer SPI \\citep{Vedrenne03} is a high spectral resolution $\\gamma$-ray telescope operating in the energy range 20 keV -- 8 MeV. Despite its lower sensitivity and limited angular resolution ($\\approx$ 2.5$\\degr$), SPI spectra extracted for brighter sources complement the ISGRI spectra well, especially above 150 keV where the ISGRI sensitivity declines rapidly. Due to its reliable, ground-based efficiency calibration, SPI serves also as a reference instrument for the INTEGRAL cross-calibration tests. Observations of crowded fields, however, are difficult, as for OSSE, and special care is required in analysing such data \\citep{Roques05}.  Notwithstanding their unique capabilities, SPI and ISGRI provide rather limited information about the continuum emission from point sources above $\\approx$200 keV. In the case of SPI, besides its limited angular resolution, there are some difficulties in modelling the instrumental background at higher energy, where the number of photons is low. However, by applying a more sophisticated analysis to SPI data, it was possible to detect 20 objects above 200 keV \\citep{Bouchet08}. The sensitivity of ISGRI declines dramatically above 150 keV, and the efficiency calibration remains uncertain in that range \\citep{Jourdain08}.  Initially, the INTEGRAL imager sensitivity was planned to exceed that of OSSE by about an order of magnitude \\citep{Winkler94}. Unfortunately, due to a reduction in mission funds, this plan became impossible, when the foreseen three layers of the high energy detector were reduced to two layers of smaller volume and a limited telemetry was ascribed to them. The effect is that the PICsIT detector, being still the largest soft $\\gamma$-ray detector on orbit, has the thickness of only 3 cm compared to almost 18 cm of OSSE. The sensitivity of PICsIT is reduced further because the pixellated structure and data-acquisition logic produce an additional loss in the fraction of the Compton-scattered photons. On the other hand, due to the elongated orbit, both the Earth occultations and passages through the radiation belts affect only a small part of the INTEGRAL revolution period. This and the observing strategy focused on the Galactic centre and Galactic plane ensures that many objects are observed by PICsIT with a long exposure time, compensating to some extent the moderate sensitivity of the instrument.  Shortly after launch, it appeared that the  PICsIT background level was about two times lower than the pre-launch conservative estimate \\citep{DiCocco03}. Only the energy band below 300 keV showed a higher background, with an excess originating in track events induced by high-energy cosmic rays \\citep{Segreto03}. Despite the lower background, the sensitivity achieved after the first year of detector operation was low, a factor between 5 and 10 lower than the statistical limit. The situation was improved considerably when the long-exposure background maps were included in the OSA 4.0 software release \\citep{Foschini04}, allowing a sensitivity comparable to that of SPI to be achieved.  Nonetheless, almost six years after the INTEGRAL launch, there remain only three objects for which results with PICsIT have been reported: Crab \\citep{DiCocco03}, Cyg X-1 \\citep{Cadolle06}, and XTE J1550-564 during its 2003 outburst \\citep{Foschini05}. The PICsIT spectra presented did not provide significantly more information than the SPI spectra. More spectacular PICsIT results were instead obtained for the observations of gamma-ray bursts (GRBs) \\citep{Malaguti03b}, in particular those  appearing outside the INTEGRAL field of view \\citep{Marcinkowski06}, because these events are detectable at high energy with the use of IBIS Compton-mode data.  Since the possibilities of improving the spectral extraction performance for PICsIT using the standard OSA software were limited, a completely novel approach was developed for testing purposes. Preliminary results demonstrated that this new approach produced very good results, by for example allowing us to detect several objects not detected by the standard procedure. This paper presents all the elements of the new method, which is now fully developed. Since we decided after discussions with the IBIS Team that the new method should not be implemented in the  standard software for INTEGRAL data analysis, we provided a detailed presentation here to allow a potential user to reproduce the results. In Sects. \\ref{instrument}--\\ref{ppdf}, basic ingredients of the spectral extraction technique are described. Section \\ref{limits} presents the results of testing the PICsIT detection limits. In Sect. \\ref{spectra}, our main results, PICsIT spectra, are widely presented and discussed. After a summary given in Sect. \\ref{summary}, an extended Appendix A presents the results of tests completed for different methods that can be used to extract the source count rates from $\\gamma$-ray instruments. ",
        "conclusions": "\\label{summary}  We have developed a comprehensive method for extracting PICsIT spectra from spectral-imaging single-event data, which enables us to derive information to the physical limitations of the instrument. This was achieved mainly by application of a correct description of the Poisson-distributed data and careful modelling of the detector background.  We adopted a technique that handles the Poisson probability density functions in a Bayesian manner for the first time in extracting the count rates from a $\\gamma$-ray detector.  The effectiveness of this technique in the `low number of counts' regime and its superiority over other methods was proven with exhaustive tests.  In our spectral extraction technique we employed a background model that was as close as possible to the data because it consisted of a map of the mean count rate measured for each pixel during a given object observation. This approach is justified for a coded-mask instrument operating in a dithering mode, especially when the background is orders of magnitude stronger than the source emission.  The detection reliability was verified by detailed studies of fake source spectra extracted from empty field observations. This approach provided the most reliable estimate of the noise level possible, which accounted for true background fluctuations and other systematic effects related to the instrument and spectral extraction method. We recommend similar tests for other instruments, because the standard analytical calculation of the detection limit can often underestimate the true noise limit.  We tested the performance of the new PICsIT spectral extraction method by comparing our results with those from other INTEGRAL instruments. Using the standard PICsIT response files, we found good agreement between the spectral fitting results for data from PICsIT and both SPI and ISGRI, obtained for the Crab, Cen A, and Cyg X-1. PICsIT appears to be well calibrated, although some tuning of the response files may be needed below 300 keV and above 1 MeV.  Confronted with the standard OSA software results, count rates computed with the new method are far more stable and about a factor of two higher, in close agreement with results derived for data from other INTEGRAL instruments. Crab spectra extracted with both OSA and the new method are described well by a model of the same spectral slope. However, for weaker objects and for observations made at large off-axis angles, we were unable to obtain high quality spectra with the standard software.  The new spectral extraction method was applied to date for about 30 objects observed with INTEGRAL. Eight new sources were added to the four already reported, which were detected with the standard PICsIT software. A refined analysis and an expanding data set should provide at least several detections more. Due to a very long exposure time, the spectrum obtained for the microquasar GRS 1758-258 reached an unprecedented quality and could be extended to higher energies up to about 700 keV. Preliminary spectra extracted for several other sources, in particular those observed during outbursts, also provide information unachievable for any other instrument currently operating.  Despite its unique advantages, PICsIT has a limited sensitivity when compared to the OSSE or BATSE detectors. The main reason for this is the smaller thickness of the detector layer, 3 cm for PICsIT, compared to 17.8 cm for OSSE and 7.6 cm for BATSE. In addition, the PICsIT ability to detect photons is decreased by the pixellated structure of the detector and the logic discriminating single and multiple events. Nevertheless, it seems that a coded-mask instrument is the most appropriate solution for soft $\\gamma$-ray astronomy, because this technique offers at the same time a large field of view, good angular resolution and direct measurement of the background when dithering observations are used."
    },
    "0809/0809.2939_arXiv.txt": {
        "abstract": "Radiation-hydrodynamical simulations of surface convection in low-mass stars can be exploited to derive estimates of i) the efficiency of the convective energy transport in the stellar surface layers; ii) the convection-related photometric micro-variability. We comment on the universality of the mixing-length parameter, and point out potential pitfalls in the process of its calibration which may be in part responsible for the contradictory findings about its variability across the Hertzsprung-Russell digramme. We further comment on the modelling of the photometric micro-variability in HD\\,49933 -- one of the first main \\textit{COROT} targets. ",
        "introduction": "Radiation-hydrodynamical (RHD) simulations of surface convection in low-mass stars have reached a high level of maturity. Such simulations provide quantitative predictions about the spatial and temporal statistics of the flows taking place in the stellar surface layers. We used the 3D RHD code \\COBOLD\\ to study convective flows in late-type stars at different metalicity. This contribution deals with two distinct applications of \\COBOLD\\ models. The first one is related to the efficiency of the convective energy transport in convective envelopes, the second one to the low-level photometric variability related temporal evolution of the surface granulation pattern. We shall rather discuss problems than solutions with a focus on dwarfs. ",
        "conclusions": ""
    },
    "0809/0809.4761.txt": {
        "abstract": "The northern Milky Way in the constellation of Cepheus ($100\\deg \\le l \\le 120\\deg; 0\\deg \\le b \\le 20\\deg$) contains several star forming regions. The molecular clouds of the Cepheus Flare region at $b > 10\\deg$, are sites of low and intermediate mass star formation located between 200 and 450~pc from the Sun. Three nearby OB~associations, Cep~OB2, Cep~OB3, Cep~OB4, located at 600--800 pc, are each involved in forming stars, like the well known high mass star forming region S\\,140 at 900~pc. The reflection nebula NGC\\,7129 around 1~kpc harbors young, compact clusters of low and intermediate mass stars. The giant star forming complex NGC\\,7538 and the young open cluster NGC\\,7380, associated with the Perseus arm, are located at $d > 2$\\,kpc. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0010_arXiv.txt": {
        "abstract": "We apply spherical needlets to the Wilkinson Microwave Anisotropy Probe 5-year cosmic microwave background (CMB) data\\_set, to  search for imprints of nonisotropic features in the CMB sky. We use the needlets' localization properties to resolve peculiar features in the CMB sky and to study how these features contribute to the anisotropy power spectrum of the CMB. In addition to the now well-known ``cold spot'' of the CMB map in the southern hemisphere, we also find two hot spots at greater than 99\\% confidence level, again in the southern hemisphere and closer to the Galactic plane. While the cold spot contributes to the anisotropy power spectrum  in the multipoles between $\\ell=6$ to $\\ell=33$, the hot spots  are found to be dominating the anisotropy power in the range between $\\ell=6$ and $\\ell=18$. Masking both the cold and the two hot spots results in a reduction by about 15\\% in the amplitude of the angular power spectrum of CMB around $\\ell=10$. The resulting changes to the cosmological parameters when the power spectrum is estimated masking these features (in addition to the WMAP team's KQ85 mask) are within the 1$\\sigma$ errors published with the WMAP mask only. We also study the asymmetry between the angular power spectra evaluated on the northern and southern hemispheres. When the features detected by needlets are masked, we find that the difference in the power, measured in terms of the anisotropy variance between $\\ell=4$ and $\\ell=18$, is reduced by a factor $2$. We make available a mask related to needlet features for more detailed studies on asymmetries in the CMB anisotropy sky. ",
        "introduction": "Beyond the angular power spectrum of the cosmic microwave background (CMB) anisotropies, the high-sensitivity all-sky CMB maps produced by the Wilkinson Microwave Anisotropy Probe (WMAP)\\footnote{http://lambda.gsfc.nasa.gov/product/map/current/} \\cite{WMAP52008} have enabled detailed statistics studies to  extract higher-order information. These studies include characterizations of asymmetries in the CMB sky \\cite{Eriksen2008} and the search for anomalous features in the anisotropy field. Previous analyses have shown evidences for alignments \\cite{Copi2007}, asymmetry in the CMB statistics between northern and southern Galactic hemispheres \\cite{Eriksen2007}, and features such as the ``cold spot,'' a significant negative feature in the CMB map first identified with wavelets \\cite{Cruz2005ColdSpot}.   Wavelets can be constructed both in real space on the sphere, for instance by means of the so-called stereographic projection \\cite{AntoineVandergheynst1999,Wiaux2005,McEwen2006b,McEwen2007}, and in harmonic space, for instance following the prescription described in \\cite{Holschneider1996,Mcewen2006a}, see also \\cite{Wiaux-st}. There have been already several applications of wavelets to CMB data analysis, including tests for non-Gaussianity and asymmetries \\cite{Vielva2004,Cabella2004,Mcewen2008,Wiaux-ng,Wiaux-glo}, polarization analysis \\cite{CabellaNatoliSilk2007}, foreground subtraction \\cite{Hansen2006}, foreground component separation \\cite{Moudden2005,Starck2005}, and point source detection in CMB anisotropy maps \\cite{Sanz2006}. The reason for such a wide success can be motivated as follows: the comparison between models and CMB data are primarily undertaken in the Fourier domain, where each multipole can be measured separately with all-sky experiments. However, in the presence of sky cuts data analysis must take into account the coupling among multipoles induced by the mask; statistical studies consequently become much more challenging. Wavelets address these issues by providing the possibility to combine sharp localization properties in pixel space with a peaked window function in the harmonic space.   In the latest two years, needlets have been proposed as a new wavelet system which enjoys several advantages over existing constructions. The introduction of needlets into the mathematical literature is due to \\cite{NarcowichPetrushevWard2006}, where their localization properties in pixel and harmonic space are established. The analysis of their statistical properties for spherical random fields is first provided by \\cite{Baldi2006}, where uncorrelation is discussed and it is also shown how needlets can be used to implement estimators for the angular power spectrum or tests of non-Gaussianity. The first application to WMAP data is due to \\cite{Pietrobon2006}, where needlets are used to estimate (cross-)angular power spectra in order to search for dark energy imprints on the correlation between large-scale structures and CMB. A full description of needlets and their potentiality for CMB data analysis is then provided by \\cite{Marinucci2008}; Ref.~\\cite{Guilloux2007} investigates the effect of different window functions in their constructions, whereas Ref.~\\cite{Baldi2007} provides further mathematical results on their behavior for partially observed sky-maps. The previous results have been extended to angular power spectrum estimation in the presence of noise \\cite{Fay2008,Fay2008b}, estimation of the bispectrum \\cite{Lan2008NeedBis}, foreground component separation \\cite{Delabrouille2008maps}, and analysis of directional data \\cite{Baldi2008Adaptivedensity}; see also Refs.~\\cite{Geller2007,Lan2008,Mayeli2008} for further developments.   Among the advantages of needlets compared to previously adopted wavelets, we recall that they are compactly supported in harmonic space, so that they allow to focus the analysis on a specific set of multipoles, which can be completely controlled \\cite{NarcowichPetrushevWard2006}. Needlet coefficients can be shown to be (approximately) uncorrelated both across different frequencies and, at high multipoles, at different locations in pixel space, which makes statistical analysis much more efficient; it is possible to provide an exact analytic expression for their variance and correlation, in terms of the underlying angular power spectrum \\cite{Baldi2006,Pietrobon2006}. Needlets do not rely on any tangent plane approximation, but they are directly embedded into the manifold structure of the sphere; they allow for direct reconstruction formulas (a consequence of the fact that they make up a tight frame system) and they are computationally very convenient \\cite{Marinucci2008}.   In this paper, we make use of needlets to further study features in the WMAP CMB maps. We focus on the large angular scales or, equivalently, on multipoles smaller than 200. We recover the cold spot that was previously detected in WMAP data with wavelets, and we also detect other features, including two hot spots, which have so far received less attention. By masking these features, we study how the angular power spectrum of CMB anisotropies is modified. Given that these features are located in the southern hemisphere, we also discuss the extent to which these features could be responsible for the north-south asymmetry in WMAP data \\cite{Lew2008,Eriksen2008,Hansen2004Asym}. As is well known, this asymmetry has also drawn much interest in the theoretical community, since it could entail strong implications on the physical nature of primordial perturbations, including inflation \\cite{Erickcek2008}. By masking the low-$\\ell$ features, we find that the difference in the CMB anisotropy variance between the two hemispheres is reduced by a factor of 2, reducing the significance of previous detections.  We also explore the evidence that statistically significant bumps and dips in the CMB anisotropy power spectrum at multipoles of 20 and 40 could be related to features in the CMB sky. We did not locate any particular feature in the sky that dominates the power spectrum at these multipoles; masking the significant features detected by needlets tuned to these multipoles did not change the power spectrum more than $\\sim$5\\%.  The paper is organized as follows: In the following section we describe briefly the needlet formalism, then we apply it to the WMAP 5-year temperature map describing the features we observed and their significance in Sec.~\\ref{sec:gauss_check}. In Sec.~\\ref{sec:cls} we address the angular power spectrum modification induced by the specific features we measure in the temperature map, computing the effect on the cosmological parameters. We then draw our conclusion in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We apply spherical needlets to the Wilkinson Microwave Anisotropy Probe 5-year cosmic microwave background dataset, to  search for imprints of nonisotropic features in the CMB sky.    After calibration by means of a large set of mock simulations, the analysis of needlet coefficients highlights the presence of the now well-known ``cold spot'' of the CMB map in the southern hemisphere, and in addition two hot spots at significance greater than  99\\% confidence level, again in the southern hemisphere and closer to the Galactic plane. While the cold spot primarily contributes to the anisotropy power spectrum  in the multipoles between $\\ell=6$ to $\\ell=33$, the hot spots  are found to be dominating the anisotropy power in the range between $\\ell=6$ and $\\ell=18$.  We also studied the effect the two spots have on the CMB power spectrum, by building $1000$ mock CMB simulations. We conclude that, especially at low multipoles, the effect is measurable: masking both the cold and the two hot spots results in an increase in the quadrupole amplitude of 10\\%, while at $\\ell=10$ power is reduced by 12\\%. To investigate the effect of this difference on the value of cosmological parameters, we modified the WMAP 5-year CMB fiducial power spectrum and the KQ$85$ mask used to perform the analysis by cutting out the contribution of the two spots, and we repeated the parameter estimation analysis. The results we obtain are slightly different, but fully consistent within the 1$\\sigma$ errors on parameters published by the WMAP team.  Since all three spots appear in the southern hemisphere, we also studied the power spectrum asymmetry between the two hemispheres, which has been previously found to be statistically significant. When the features detected by needlets are masked, we find that the difference in the power, measured in terms of the anisotropy variance between $\\ell=4$ and $\\ell=18$, is reduced by a factor of $2$. This decreases the significance of the previously claimed north-south asymmetry. We make the mask resulting from needlet features available for future, more detailed studies on the asymmetries in the CMB anisotropy sky\\footnote{http://www.fisica.uniroma2.it/$\\sim$cosmo/masks}.  \\begin{center} {\\bf Acknowledgments} \\end{center}  DP thanks Marcella Veneziani for useful discussions and the UCI Physical Sciences Department where this work has been performed. DM is grateful to P.Baldi, G.Kerkyacharian and D.Picard for many useful discussions.  \\newpage"
    },
    "0809/0809.1625.txt": {
        "abstract": "We analyze stellar convection with the aid of 3D hydrodynamic simulations, introducing the turbulent cascade into our theoretical analysis. We devise closures of the Reynolds-decomposed mean field equations by simple physical modeling of the simulations (we relate temperature and density fluctuations via coefficients); the procedure (CABS, Convection Algorithm Based on Simulations) is terrestrially testable and is amenable to systematic improvement. We develop a turbulent kinetic energy equation which contains both nonlocal and time dependent terms, and is appropriate if the convective transit time is shorter than the evolutionary time scale. The interpretation of mixing-length theory (MLT) as generally used in astrophysics is incorrect; MLT forces the mixing length to be an imposed constant. Direct tests show that the damping associated with the flow is that suggested by Kolmogorov ($\\varepsilon_K \\approx \\rho \\ubar^{3}/\\ld$, where $\\ld$ is the size of the largest eddy and $\\ubar$ is the local rms turbulent velocity). This eddy size is approximately the depth of the convection zone $\\lcz$ in our simulations, and corresponds in some respects to the mixing length of MLT. New terms involving the local heating due to turbulent dissipation should appear in the stellar evolutionary equations, and are not guaranteed to be negligible. The enthalpy flux (stellar ``convective luminosity'') is directly connected to the buoyant acceleration, and hence to the scale of convective velocity. MLT tends to systematically underestimate the velocity scale, which affects estimates of chromospheric and coronal heating, mass loss, and wave generation. Quantitative comparison with a variety of 3D simulations reveals a previously unrecognized consistency. Extension of this approach to deal with rotational shear and MHD is indicated. Examples of application to stellar evolution will be presented in subsequent papers in this series. ",
        "introduction": "More than fifty years ago,  the version of the ``mixing length theory'' of convection (MLT) which became the preferred basis for subsequent study of stellar evolution was introduced \\citep{bier51,vitense53,bv58}. Despite much effort (still ongoing), MLT is still the standard choice for the field.  In this paper we develop a new procedure, \"Convective Algorithms Based on Simulations\" (CABS), in which we close the Reynolds-decomposed, angular and time averaged equations by simple physical models based upon analysis of fully three-dimensional, time dependent turbulent stellar convection \\citep{ma07b}. These simulations include convective boundaries within the computational volume (as far from the edges of the grid as feasible), and allow interface physics to be examined. The resulting theoretical formalism allows us to incorporate content from other simulations (especially \\cite{cs89,cs96,kim95,kim96,pw00,pwj00,rob04}), and from research in other fields, such as terrestrial fluids \\citep{turner73}, oceanography \\citep{gill}, and meteorology \\citep{dutton}, which have a firmer empirical basis. We develop a simple description of the convective velocity field as seen in our simulations. This effort brings some startling suggestions for revision of our interpretation of MLT, and suggests how our approach may be generalized to include rotation and magnetic fields \\citep{balbus,pessah}. This is timely, considering recent success in simulating turbulent plasma with magnetic fields \\citep{brown08,schu08}.  Since the formulation of MLT, there has been a considerable development in understanding the nature of chaotic behavior in nonlinear systems; see \\cite{cvit} for a review and reprints of original papers, and \\cite{frisch,gleick,thomp}. \\cite{lorenz} presented a solution to the Rayleigh problem of thermal convection \\citep{chandra} which captured the seed of chaos in the Lorenz attractor, and contains a representation of the fluctuating aspect of turbulence not present in MLT. \\cite{kolmg} and \\cite{obukhov} developed the modern version of the turbulent cascade. Although already formulated, the original theory \\citep{kolmg41} was not used in MLT.  We derive our approximate theory from a consideration of the full equations of 3D compressible hydrodynamics for a multiple component fluid. These are close to the corresponding equations for a high beta (matter dominated) plasma. This approach will allow us to incorporate a variety of phenomena in a coherent way, and to evaluate their relative importance, rather than to patch together various bits of physics piecemeal. Here we focus on the dominant features of non-rotating, non-magnetic, turbulent, compressible fluid flow. The results are applicable to almost all stages of stellar evolution (any stage having convection or shear)\\footnote{The shear from convection is similar to the shear from differential rotation; fluid experiments may use either to investigate the physics of shear \\citep{turner73}. Although different in some details, there are deep connections between convective mixing as described here and the rotational mixing investigated by \\cite{mm00}.}.   In Section~2 we examine the physical aspects of stellar convection, and show  that the velocity scale is set by the balance between buoyant driving, and damping in the Kolmogorov cascade. Appropriate averaging allows us to deal with mild time dependence and reveals a robust underlying behavior. We develop a kinetic energy equation describing the average properties of turbulent convection, which shows how turbulent motion is created, transported, and destroyed. In Section~3 we incorporate this theoretical development into the equations of stellar evolution with turbulent convection, including new terms. Section~4 uses the theoretical development to compare our simulations to others, and finds previously unrecognized similarities. Section~5 indicates some important implications of this work, including how the effects of burning, rotation and magnetic fields may be included. Section~6 summarizes our results and conclusions. Quantitative treatment of the dynamics of fluctuations and a comprehensive algorithm for stellar evolutionary calculations will be developed in subsequent papers in this series. ",
        "conclusions": "We find that our three-dimensional time-dependent simulations of compressible stellar turbulence (ILES) are well represented by a master equation (Eq.~\\ref{a12}) for kinetic energy which includes dissipation implied by the Kolmogorov cascade. The damping length is found in three independent ways, with reasonable consistency, and is the size of the largest eddy, which is approximately the geometric linear dimension (depth) of the convective zone. Unlike the mixing length of MLT, it is not a free parameter, but a mathematical consequence of the turbulent flow. Balancing turbulent buoyant driving with Kolmogorov damping provides a reasonable estimate of the damping length. The divergence of kinetic energy and acoustic fluxes is nonzero, and provides the mechanism to spread turbulence through the convective zone, and make the Kolmogorov damping a good approximation. The turbulent flow is highly dissipative and must be maintained by continual driving; this ``frictional heating'' term is missing from the standard stellar evolution equations, as are the kinetic energy and acoustic fluxes.  Fluctuations in kinetic energy are significant (of order 50 percent), and damping lags driving by about a transit time for the convective zone.  Comparison with some other simulations, which were dramatically different in many respects, gives a consistent picture. Turbulent convection is more vigorous for deeper (more stratified) convection zones. Turbulent kinetic energies are lower for nonideal equations of state, such as partially ionized plasmas and electron-positron plasmas, to the extent that their specific heat at constant volume is less that their specific heat at constant pressure. The average flow structure changes in a simple way depending upon the depth of the convection zone; deep convection zones have downwardly directed flows of kinetic energy, cancelling some of the upward enthalpy flux.  It appears that extension of this approach, using simulations to define stellar convection algorithms, can establish a theoretical model of turbulent convection that does not require astronomical calibration, but can be based upon a combination of computer simulations, terrestrial observations, and  experiments. Efforts to implement this general theory in a stellar evolution code are underway."
    },
    "0809/0809.4225_arXiv.txt": {
        "abstract": "We explore the implications of the QCD phase transition during the postbounce evolution of core-collapse supernovae. Using the MIT bag model for the description of quark matter and assuming small bag constants, we find that the phase transition occurs during the early postbounce accretion phase. This stage of the evolution can be simulated with general relativistic three-flavor Boltzmann neutrino transport. The phase transition produces a second shock wave that triggers  a delayed supernova explosion. If such a phase transition happens in a future galactic supernova, its existence and properties should become observable as a second peak in the neutrino signal that is accompanied by significant changes in the energy of the emitted neutrinos. In contrast to the first neutronization burst, this second neutrino burst is dominated by the emission of anti-neutrinos because the electron-degeneracy is lifted when the second shock passes through the previously neutronized matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.3319_arXiv.txt": {
        "abstract": "We are developing a new photon detector with micro pattern gaseous detectors. A semitransparent CsI photocathode is combined with 10cm$\\times$10cm GEM/$\\mu$PIC for the first prototype which is aimed for the large liquid Xe detectors. Using Ar+C$_2$H$_6$ (10$\\%$) gas, we achieved the gas gain of $10^5$ which is enough to detect single photoelectron. We, then, irradiated UV photons from a newly developed solid scintillator, LaF$_3$(Nd), to the detector and successfully detected single photoelectron. ",
        "introduction": "% In the last two decades, many gaseous photomultipliers with CsI photocathodes operated in both flushed gas mode and sealed gas mode have been developed and tested\\cite{charpak,berskin,carlson}. Moreover, recently, large area micro pattern gaseous detectors with avalanche multiplication structures, such as Micromegas\\cite{megas}, GEM\\cite{gem}, and $\\mu$PIC\\cite{upic} have been developed and successfully operated. These devices with photocathodes can realize a low cost large area photon detector with position sensitivity and can be applied to large astroparticle detectors. In particular, the quantum efficiency of the CsI photocathode matches the liquid Xe scintillator, thus dark matter search via Xe is one of the first targets of this photon detector.  In this work, for the first step, we investigated the feasibility of manufacturing a large size gaseous photon detector with CsI photocathode and evaluated its performance. ",
        "conclusions": "A UV photon detector based a semitransparent CsI photocathode combined with large area GEMs and a $\\mu$PIC has been reported. It is demonstrated that the detector has the ability of detecting single photoelectron. The quantum efficiency is about 1$\\%$ as expected.  In order to increase the quantum efficiency, a reflective type CsI photocathode (200nm-thick CsI evaporated to one side of the GEM) is being developed and tested now."
    },
    "0809/0809.4539_arXiv.txt": {
        "abstract": "% We investigate correlations between the direction of the optical linear polarization and the orientation of the host galaxy/extended emission for type1 and type2 radio-loud and radio-quiet quasars. We have used high resolution Hubble Space Telescope data and a deconvolution process to obtain a good determination of the host galaxy/extended emission (EE) position angle. With these new measurements and a compilation of data from the literature, we find a significant correlation, different for type1 and type2 objects, between the linear polarization position angle and the orientation of the EE, suggesting scattering by an extended UV/blue region in both types of objects. Our observations support the extension of the Unification Model to the higher luminosity AGNs like the quasars, assuming a two component scattering model. ",
        "introduction": "The discovery of type1-like broad emission lines in the polarized spectrum of the type2 Seyfert galaxy NGC1068 \\citep{borguet_anto85} led to the foundation of the so-called Unification Model (UM) of AGNs \\citep[e.g.][]{borguet_anto93} in which the orientation of a dusty torus plays a crucial role in the spectroscopic classification of the AGN for a given observer. While this UM is now well accepted for the lower luminosity AGNs, a key question is whether it applies to higher luminosity objects like quasars since the presence of the dusty torus may be affected by the higher radiation flux.  In order to get some insights into the quasar inner structure, we investigate the possible existence of a correlation between the direction of the linear optical polarization ($\\theta_{pola}$) and the orientation of the major axis ($PA_{host}$) of the host galaxy/extended emission surrounding RQ and RL quasars. ",
        "conclusions": ""
    },
    "0809/0809.3775_arXiv.txt": {
        "abstract": "The Magellanic Clouds (MCs) present a rich system of stellar clusters that can be used to probe the dynamical and chemical evolution of these neighboring and interacting irregular galaxies. In particular, these stellar clusters (SCs) present combinations of age and metallicity that are not found for this class of objects in the Milky Way, being therefore very useful templates to test and to calibrate integrated light simple stellar population (SSP) models applied to unresolved distance galaxies. On its turn, the age and metallicity for a cluster can be determined spatially resolving its stars, by means of analysis of its colour-magnitude diagrams (CMDs). In this work we present our method to determine self-consistent physical parameters (age, metallicity, distance modulus and reddening) for a stellar cluster, from CMDs modelling of relatively unstudied SCs in the Small Magellanic Cloud (SMC) imaged in the BVI filters with the 4.1 m SOAR telescope. Our preliminary results confirm our expectations that come from a previous integrated spectra and colour analysis: at least one of them (Lindsay~2) is an intermediate-age stellar cluster with $\\sim$~2.6 Gyr and [Fe/H]~$\\sim$~-1.3, being therefore a new interesting witness regarding the reactivation of the star formation in the MCs in the last 4 Gyr. ",
        "introduction": "SCs are useful tools to study the complex stellar content observed in nearby galaxies, as they may be modelled as SSP of a fixed age and metallicity. In particular, the MC SCs form a rich system ($>$~3700 objects) (Bica et al. 2008), therefore they can be used to test and to calibrate SSP models for combinations of age and metallicity that are not found in the Milky Way (Santos Jr. \\& Piatti 2004). Such information can be used to probe the dynamical and chemical evolution of these neighboring and interacting dwarf irregular galaxies.  The age distribution based on clusters is likely distinct from the star formation history (SFH) as inferred from field stars (Holtzman et al. 1999). More recently, Rafelski \\& Zaritsky (2005) analysed a sample of 195 clusters and shows that the populations of field stars are similar to the populations of SMC stars. The large period of quiescent star formation in the MCs between $\\sim$~4 and 10~Gyr (Harris \\& Zaritsky 2001, 2004) seems to be imprinted in the low number of populous SCs with these ages (Rich et al. 2000, 2001). Almost all of them are in the SMC as well (Mighell et al. 1998; Piatti et al. 2005, 2007).  In spite of that the SMC cluster system is comparatively much less studied than the LMC, e.g. Piatti et al. (2001) list only 16 SMC clusters with known ages and metallicities. A few more have been recently added by the same authors (Piatti et al. 2005). Detailed CMDs for SMC clusters are still very scarce as well, which prevents more reliable age, metallicity and structural information from being derived for them.  Therefore we are studying SMC clusters by using photometry and CMD analysis, based on a combination of CMD modelling and statistical tools. The confirmation of some of these clusters as intermediate or old age ones will significantly improve the poor census in the age range corresponding to the age gap for the SMC clusters. Our sample data are for Lindsay~2, and Lindsay~72, for which preliminary estimates in the literature indicate ages of 3 to 8~Gyr and $\\sim$~200~Myr for Lindsay~72 (Chiosi et al. 2006), respectively. An age indication for Lindsay~2 was obtained from low resolution integrated spectra with the 1.5m telescopes at the ESO and LNA observed by us and based on integrated magnitudes and colours from Rafelski \\& Zaritsky (2005). ",
        "conclusions": "Using SOAR/SOI photometry we can objectively determine accurate and self-consistent physical parameters for SMC clusters by means of CMD modelling, as Lindsay~2 results showed.  Soon we will combine results from analysis of V,B-V and V,V-I CMDs and we will also apply our method to Lindsay~72 and other far west SMC clusters that are suspect to be of intermediate or even old age. In case we confirm that, we would be able to combine our results with those of Crowl et al. (2001) and discuss the possibility that these clusters may be tidally stripped from the SMC into the Magellanic stream."
    },
    "0809/0809.3069_arXiv.txt": {
        "abstract": "We study the effects of the tidal interaction with the companion, via orbital separation and binary mass ratio, on the global one-armed oscillation modes in disks around binary Be stars. Our model takes into account the three-dimensional effect that contributes to the mode confinement, which was recently found by \\citet{Ogi08}. We find that the one-armed oscillations are well confined in systems with disks larger than a few tens of stellar radii. In such systems, the oscillation period depends little on the binary parameters. On the other hand, in systems with smaller disks, where the mode confinement is incomplete, the oscillation period increases with increasing orbital separation and/or decreasing binary mass ratio. The eigenmode is insensitive to the spectral type of the central star. Our results suggest that the dependence of V/R oscillation period on the orbital separation and binary mass ratio should be observed only in short period binary systems, and that, for systems with a similar orbital period, those with higher mass ratios will show shorter V/R variations. ",
        "introduction": "\\label{sec:intro} Be stars are non-supergiant B-type stars whose spectra show, or have at some time shown, one or more Balmer lines in emission (\\cite{Col87}). A Be star is a rapidly rotating, B-type star as the central star with a two-component circumstellar envelope, a polar wind and an equatorial disk. The polar wind is a low-density, fast outflow emitting UV radiation. The wind structure is well explained by the line-driven wind model (\\cite{CAK75}; \\cite{FO86}). On the other hand, the equatorial disk is a geometrically thin, high-density plasma in nearly Keplerian rotation (e.g., \\cite{Por03}), from which the optical emission lines and the IR excess arise. Although there are still several competing scenarios for the Be disk formation, the most promising one is the scenario where the disk is formed by viscous decretion of gas ejected from the central star (\\cite{Lee91}; see also \\cite{Por03}, and references therein).  Many Be stars show long-term variations in their spectra over years to decades. One of these phenomena is called long-term V/R variations, which are variations of the ratio of relative intensity of violet (V) and red (R) peaks of a double-peaked emission line profile. The period of the long-term V/R variations is typically in the range 5-10\\,yr for isolated Be stars. It is widely accepted that the long-term V/R variations are attributed to global one-armed (i.e., the azimuthal wave number $m=1$) oscillations in the Be disk (e.g., \\cite{Por03}). It is based on the fact that in a nearly Keplerian disk a one-armed mode shows up as a very slowly revolving perturbation pattern (e.g., \\cite{Kat83}). The early versions of the one-armed oscillation model (\\authorcite{Oka91} \\yearcite{Oka91}, \\yearcite{Oka97}; \\cite{Pap92, Sav93}) qualitatively well explained the observed characteristics of the V/R variations. However, owing to the lack of the mechanism for confining the modes to the inner part of the disk and the high sensitivity of the oscillation period on various parameters, the model had little predictive power \\citep{FH06}.  Recently, the model was advanced greatly. \\citet{Ogi08} has solved the mode confinement problem by investigating a three-dimensional effect on the mode characteristics. While previous versions of the model had assumed motions to be independent of height since the disk is geometrically thin, he argued that the variation of the vertical gravitational acceleration around an elliptical orbit excites an oscillatory vertical motion in an eccentric disk that should not be neglected. \\citet{Ogi08} showed that the three-dimensional effect alone allows confined prograde modes.  Hitherto, theoretical studies of one-armed oscillation modes in Be disks are mostly limited to isolated Be stars. However, according to the current census, the fraction of Be stars found in binary systems is $\\sim 1/3$ (\\cite{Por03}). The presence of the companion is most likely to change some features of the global one-armed oscillations in Be disks. Actually, in Be/X-ray binaries, which consist of a neutron star and an early-type Be star, the time-scale of V/R variations is frequently about 1\\,yr, which is much shorter than the typical V/R time-scale for isolated Be stars (e.g., \\cite{Rei05}).  Moreover, there are observationally two types of long-term V/R variations in binary Be stars (\\cite{Ste07}). One is quasi-periodic V/R variations, where V/R peak ratio varies quasi periodically, which shows no correlation with their orbital phase. This type of V/R variations are similar to those usually observed for V/R variations in isolated Be stars. The other type is periodic V/R variations locked to the orbital period. This characteristic is found only in binary Be star systems. Unfortunately, no theoretical explanation is available yet for the phase-locked V/R variations.  In this paper, we investigate the effects of the companion on the characteristics of one-armed oscillation modes in binary Be stars, taking into account the tidal effects in addition to the three-dimensional effect. For simplicity, we assume the binary orbit to be circular. We find that the mode is well confined in disks larger than a few tens of stellar radii, which is consistent with \\citet{Ogi08}. In smaller disks, however, the mode confinement is incomplete and the oscillation period depends on the binary parameters significantly.  The structure of this paper is as follows. In section 2, we describe the unperturbed disk and derive the basic equations for linear, one-armed perturbations, following the formulation by \\citet{Ogi08}. The boundary conditions to be adopted are also discussed. In section 3, we present our numerical results, and in section 4, we compare our results with the observational features of V/R variations in binary Be stars. The final section is devoted to conclusions. ",
        "conclusions": "We have studied the tidal effect of the companion on the global oscillation modes in equatorial disks around binary Be stars. For simplicity, we assumed {the binary orbit to be circular and} the {Be} disk to be inviscid, isothermal, and truncated at the tidal radius. We {solved} linearized {equations for global} $m=1$ {perturbations in a} three-dimensional {Be} disk {with a power-law density distribution,} with the rigid wall inner boundary condition, which is applicable to systems where material is ejected continuously, so there is no gap between the star and the disk.  We {have} obtained prograde {fundamental} modes {even when the quadrapole contribution to the potential is negligible.} This confirms the results of \\citet{Ogi08}}. In our study of binary Be star systems, the modes are well confined when the disk is larger than a few tens of stellar radii.  We have found that the oscillation period increases with increasing binary separation {and/or {decreasing} binary mass ratio. The effect is, however, only appreciable for small binary separation. In systems with large binary separations, the fundamental mode is well confined to the inner part of the disk, so the eigenfrequency no longer depends on the binary parameters.  Be stars sometime{s} show {a} cavity {between the star and} the disk when the disk is being dissipated. {In order to study the global oscillation modes in such disks, we have solved the perturbation equations with a free inner boundary condition, without changing the density distribution. We have obtained much higher eigenfrequencies with the free inner bounday condition than with the rigid one, as in \\citet{Pap06} and \\citet{Ogi08}. The result implies an interesting possibility that the V/R variability time-scale is shorter when the disk is dissipating than when it is forming or persistent.}  {In this paper we have assumed a power-law density distribution. In the disk dissipation stage, however, it is likely that} the density distribution is far from {a} power-law {form} and {is} highly peaked near the inner {radius}. In a subsequent paper, we will study {the} effect {of density distribution on the global $m=1$ oscillations}.  \\bigskip  FO thank{s} Masayuki Fujimoto for helpful discussions. She also acknowledges the scholarship from Ministry of Education, Culture, Sports, Science and Technology. {We thank Gordon Ogilvie for kindly showing us his latest result on the three-dimensional effect on the $m=1$ mode confinement.}"
    },
    "0809/0809.2290_arXiv.txt": {
        "abstract": "{ } { % We create a catalogue of \\emph{simulated} fossil groups and study their properties, in particular the merging histories of their first-ranked galaxies. We compare the \\emph{simulated} fossil group properties with those of both \\emph{simulated} non-fossil and \\emph{observed} fossil groups. } {Using simulations and a mock galaxy catalogue, we searched for massive ($>$ 5 $\\times$ 10$^{13} \\ h^{-1} {\\cal M}_\\odot$) fossil groups in the Millennium Simulation Galaxy Catalogue. In addition, attempted to identify observed fossil groups in the Sloan Digital Sky Survey Data Release 6 using identical selection criteria. } { Our predictions on the basis of the simulation data are: (a) fossil groups comprise about 5.5\\% of the total population of groups/clusters with masses larger than 5 x 10$^{13} \\ h^{-1} {\\cal M}_\\odot$. This fraction is consistent with the fraction of fossil groups identified in the SDSS, after all observational biases have been taken into account; (b) about 88\\% of the dominant central objects in fossil groups are elliptical galaxies that have a median R-band absolute magnitude of $\\sim -23.5-5 \\ log \\ h$, which is typical of the observed fossil groups known in the literature; (c) first-ranked galaxies of systems with $ {\\cal M} >$   5 x 10$^{13} \\ h^{-1} {\\cal M}_\\odot$, regardless of whether they are either fossil or non-fossil, are mainly formed by gas-poor mergers; (d) although fossil groups, in general, assembled most of their virial masses at higher redshifts in comparison with non-fossil groups, \\emph{first-ranked galaxies in fossil groups merged later, i.e. at lower redshifts}, compared with their non-fossil-group counterparts. } {We therefore expect to observe a number of luminous galaxies in the centres of fossil groups that show signs of a recent major merger.} ",
        "introduction": "\\cite{Jones03} identified fossil groups as spatially extended X-ray sources with an X-ray luminosity $L_X>10^{42} \\ h_{50}^{-2} \\ erg \\ s^{-1}$ whose optical counterpart was a bound system of galaxies with $\\Delta M_{12}>2$ mag, where $\\Delta M_{12}$ was the difference in absolute magnitude in R-band between the brightest and the second brightest galaxies in the system within half the projected virial radius ($r_{vir}$). The dynamical masses of the systems studied so far are comparable to those of rich clusters ($\\sim 10^{13}-10^{14} \\ h^{-1} {\\cal M}_\\odot$) \\citep{Mendes06,cyp06,Mendes08, KPJ06}. Fossil groups may be of considerable importance as the place of formation of a significant fraction of all giant ellipticals. Beside minor differences in the definition of fossil groups, their incidence rate was estimated by observational, analytical, numerical, and semi-analytical analyses. \\cite{Vik99} and \\cite{Jones03} stated that fossil groups represent (8-20)\\% of observed systems in the same mass range. \\cite{vandenbosch07} used the 2dF Galaxy Redshift Survey data to measure a fossil fraction of 6.5\\% among groups with masses $(10^{13}-10^{14}) \\ h^{-1} {\\cal M}_\\odot$. \\cite{M06} estimated analytically that fossil groups represent 5-40\\% of groups with masses in the range $\\sim 10^{13}-10^{14} \\ h^{-1} {\\cal M}_\\odot$, while the percentage decreased to 1-3\\% for groups of mass larger than $10^{14} \\ h^{-1} {\\cal M}_\\odot$, the  latter result having been confirmed using the photometric Sloan Digital Sky Survey Data Release 2 (SDSS DR2). Numerical simulations by \\cite{Donghia} predicted a higher fraction of fossil systems (33\\%) amongst groups of mass $\\sim 10^{14}  \\ h^{-1} {\\cal M}_\\odot$. \\cite{vonbenda07}, also using numerical simulations, found that 24\\% of groups with masses in the range $(1-5) \\times 10^{13} \\ h^{-1}{\\cal M}_\\odot$ were fossil groups. From semi-analytical models, \\cite{sales07} estimated that fossil groups represent $(8-10)\\%$ of groups with masses $(10^{13}-10^{15}) \\ h^{-1}{\\cal M}_\\odot$, while \\cite{Dariush07} stated that $\\sim 13  \\%$ of groups in that mass range were fossil systems and predicted that this percentage decreased to 3-4 \\% for X-ray rich systems.  A natural question is whether the large magnitude difference between the first and second ranked galaxies ($\\Delta$M$_{12} > $ 2 in the R-band), characteristic of these groups, implies that they are a distinct class of objects or if they instead represent a tail of the cluster (hereafter non-fossil group) distribution. To investigate this question, \\cite{Donghia} used high-resolution N-body/hydrodynamical simulations to compare four simulated fossil and eight non-fossil groups, all of virial masses close to $1 \\times 10^{14} \\ h^{-1} \\ {\\cal M}_\\odot$. They found that the values of the magnitude gap, $\\Delta$M$_{12}$, for the 12 systems, were correlated with the halo assembly time, such that fossil groups assembled earlier than non-fossils. Similarly, \\cite{Dariush07} concluded, by the study of fossil groups in the Millennium Simulation, that fossils assemble a higher fraction of their masses at higher redshifts than non-fossil groups. The most accepted scenario for fossil groups is then that they are not a distinct class but, are instead, examples of groups/clusters that collapsed early.  Although fossil groups in general do not appear differ from galaxy clusters of similar masses, except for their earlier times of formation, it was realised, observationally, that the first-ranked galaxies in fossil and non-fossil groups differ in some respects. First, their shapes differ: while fossil, first-ranked, galaxies are often disky, brightest cluster galaxies are often boxy \\citep{KPJ06}. Second, their stellar populations are dissimilar: fossil group first-ranked galaxies are not as old as brightest cluster galaxies  \\citep{delaRosa08}. These findings motivated us to revisit the study of \\cite{Dariush07} of fossil groups in the Millennium Simulation, but now focusing on the properties of the first-ranked galaxies (as opposed to the complete system). By searching in the Millennium Simulation,  we created two samples, one of fossil groups with $\\Delta$M$_{12} > 2$ (in the R-band) and a second control sample of non-fossil groups  with systems with $\\Delta$M$_{12} < 0.5$.  Both the simulations themselves and a mock catalogue were used to complete these searches. In addition, we searched for fossil groups in the Sloan Digital Sky Survey Data Release 6 (SDSS DR6) \\citep{AMSDSS08} using the same criteria, to compare with the mock catalogue as a test of the semi-analytic model and also to examine observationally the results of our searching algorithm.  The layout of this paper is as follows. In Sect.~\\ref{construction} we briefly describe the galaxies in the Millennium Simulation, and the search for simulated fossil and non-fossil groups. Section \\ref{differences} contains a comparison between fossil and non-fossil groups, and simulated and observed fossils. In particular, we discuss the implications of these results for the evolution of first-ranked galaxies in fossil groups. In Sect.~\\ref{samples} we describe the construction of a mock catalogue and the procedures of group identification. We also perform an identification of fossil groups in SDSS DR6 and compare the results with those obtained from the mock catalogue. Finally, we summarise the paper in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} To help interpret observations, we have analysed the properties of the first-ranked galaxies of simulated fossil groups and their main differences with respect to the brightest galaxies of non-fossil systems. To perform this analysis, we have studied a sample of fossil groups obtained from the largest simulated galaxy catalogue at present, the Millennium simulation combined with a semi-analytic model of galaxy formation \\citep{deLucia07}. The sample of fossil groups identified was sufficiently large to enable robust statistical analyses to be completed. A control sample of non-fossil groups was also selected. The goal of this research was twofold: to predict results based on semi-analytical models about fossil groups and test the semi-analytical model of formation and evolution by comparing with observational results.  Our analysis was performed in both the Millennium Galaxy Catalogue (MSGC) and a mock catalogue. The fraction of fossils predicted by the mock catalogue was compared with the fraction of fossils found in the SDSS data set by applying the same algorithm of identification. If the two main sources of incompleteness affecting the spectroscopic sample of SDSS were taken into account, the fraction of fossil groups found in the mock catalogue was consistent with that observed in the SDSS. We have presented six candidate fossil systems in SDSS DR6. One has been rejected after visually analysing the surroundings of the brightest galaxy. The remaining groups await X-rays observations or deeper photometric studies to confirm their nature as fossil groups.  Our main results can be summarised as follows: we confirm the old age feature of fossil systems and that fossil groups with masses larger than $5 \\times 10^{13} \\ h^{-1} {\\cal M}_\\odot$ represent $\\sim 5.5\\%$ of massive systems in the same mass range, in the Millennium Galaxy Catalogue and a lower percentage (of 3\\%) of similar systems in a mock galaxy catalogue. Fossils in the Mock Catalogue were identified in redshift space (just as achieved in observations), where galaxies from the outskirts of the groups affect the $\\Delta M_{12}$ selection criterion and reduce the resulting number of fossils. We found that $88\\%$ of fossil systems have central galaxies that are ellipticals. In addition, we found that the first-ranked galaxies of fossil groups and non-fossil groups have the same morphological mixtures,  which can be considered to be the results of gas poor mergers. If the central galaxy of a fossil group is always elliptical, the likelihood that its progenitor galaxies merge in a dry merger is higher. Although this result disagrees with the observational studies of \\cite{KPJ06} for seven elliptical galaxies in fossil groups, further observational data is required to fully resolve this issue. Finally, we investigated the nature of the central galaxies by keeping track of their merging histories. On the one hand, the fossil groups in general assembled earlier than non-fossil groups, whilst their central galaxies assembled later. We also found that first-ranked galaxies in fossil groups have undergone a major merger later than their counterparts in non-fossil systems. We expect this result to be confirmed by future observational catalogues."
    },
    "0809/0809.2403_arXiv.txt": {
        "abstract": "A representative sample of unevolved early B-type stars in nearby OB associations and the field is analysed to unprecedented precision using NLTE techniques. The resulting chemical composition is found to be more metal-rich and much more homogeneous than indicated by previous work. A rms scatter of $\\sim$10\\% in abundances is found for the six stars (and confirmed by six evolved stars), the same as reported for ISM gas-phase abundances. A cosmic abundance standard for the present-day solar neighbourhood is proposed, implying mass fractions for hydrogen, helium and metals of $X$\\,$=$\\,0.715, $Y$\\,$=$\\,0.271 and $Z$\\,$=$\\,0.014. Good agreement with solar photospheric abundances as reported from recent 3D radiative-hydrodynamical simulations of the solar atmosphere is obtained. As a first application we use the cosmic abundance standard as a proxy for the determination of the local ISM dust-phase composition, putting tight observational constraints on dust models. ",
        "introduction": "The Sun is unique among the stars because independent~in\\-dicators allow its chemical composition to be constrained with a precision unmatched for any other star. This can be done by spectroscopic analysis of its photosphere and by measurement of solar wind and solar energetic particles. Solar nebula abundances can be determined from CI chon\\-drites, which are unaltered since the formation of the system. The wealth of information established the Sun~as~the principal standard for the chemical composition of cosmic matter (e.g. Grevesse \\& Sauval 1998, GS98; Holweger 2001; Asplund et al 2005, AGS05). % However, is a 4.6\\,Gyr old star indeed representative of the cosmic matter in its neighbourhood\\footnote{ We consider the region at distances shorter than $\\sim$1\\,kpc (and $\\pm$500 pc in Galactocentric direction) as solar neighbourhood in order to minimize bias due to Galactic abundance gradients, see Fig.~\\ref{evol} for a schematic overview.} at present?  Ideal indicators for pristine abundances are unevolved early B-stars of spectral types B0--B2. Slowly rotating stars are preferred as their photospheres should be essentially unaffected by mixing of CN-processed material \\citep[][]{maeder00}. The atmospheres of early B-stars are also unaffected by atomic diffusion that gives rise to peculiarities of metal abundances in many later-type stars \\citep[e.g.][]{smith96}. A major practical advantage is also their relatively simple photospheric physics, which is represented well by classical model atmospheres, unaffected by complications such as stellar winds or convection.  As a consequence, early B-stars in the solar neighbourhood were subject of several NLTE studies in the past \\citep[e.g.][]{gies92,kilian92,kilian94,cunha94,daflon99,daflon01a,daflon01b,daflon03,lyubimkov04,lyubimkov05}. Overall, they found a wide range of abundances, by about a factor $\\sim$10, and an average metallicity of only $\\sim$2/3 solar (GS98). Hence, the impression arose that the solar neighbourhood is chemically highly heterogeneous, and the Sun anomalously metal-rich compared to young stars.  Both findings are problematic in terms of Galactic chemical evolution. Dispersal of stellar nucleosynthesis products increases the metallicity over time \\citep[e.g.][]{chiappini03} and hydrodynamic mixing tends to homogenize the interstellar medium (ISM) locally \\citep{edmunds75}. Characteristic timescales for homogenization are short, ranging from 10$^6$--10$^8$\\,yrs on scales of 100-1000\\,pc \\citep{roy95}.  In contrast to the young stars the interstellar gas shows a high degree of chemical homogeneity in the solar neighbourhood \\citep{sofia04}, with the rms scatter of mean abundances being $\\sim$10\\%. However, the ISM gas phase is not suitable as a tracer for cosmic abundances because of selective depletion of elements onto dust grains. Here we reinvestigate the conundrum of inhomogeneous stellar vs. homogeneous ISM gas-phase abundances in the solar neighbourhood, motivated by the finding of homogeneous B-star abundances for carbon \\citep[see detailed analysis by][NP08]{nieva08}. ",
        "conclusions": ""
    },
    "0809/0809.0138_arXiv.txt": {
        "abstract": "Using RHESSI and some auxiliary observations we examine possible connections between spatial and temporal morphology of the sources of non-thermal hard X-ray (HXR) emission which revealed minute quasi-periodic pulsations (QPPs) during the two-ribbon flares on 2003 May 29 and 2005 January 19. Microwave emission also reveals the same quasi-periodicity. The sources of non-thermal HXR emission are situated mainly inside the footpoints of the flare arcade loops observed by the TRACE and SOHO instruments in the EUV range. At least one of the sources moves systematically both during the QPP-phase and after it in each flare that allows to examine the sources velocities and the energy release rate via the process of magnetic reconnection. The sources move predominantly parallel to the magnetic inversion line or the appropriate flare ribbon during the QPP-phase whereas the movement slightly changes to more perpendicular regime after the QPPs. Each QPP is emitted from its own position, which does not coincide with the origin of the previous pulsations. It is also seen that the velocity and the energy release rate don't correlate well with the flux of the HXR emission calculated from the sources. The sources of microwaves (observed by NoRH during the May 29 flare only) and thermal HXRs are situated near the apex of the loop arcade and are not stationary either. Almost all QPPs and some spikes of HXR emission during the post-QPP-phase reveal the soft-hard-soft spectral behavior indicating separate acts of electrons acceleration and injection, rather than modulation of emission flux by some kinds of magnetohydrodynamic (MHD) oscillations of coronal loops. In all likelihood, the flare scenarios based on the successively firing arcade loops are more preferable to interpret the observations, although we can not conclude exactly what mechanism forces these loops to flare up. ",
        "introduction": "\\label{S-Introduction}  Quasi-oscillatory phenomena are ubiquitous in nature and always attractive for researchers. This is also true for solar and stellar dynamic processes such as flares, filament oscillations (\\eg, \\opencite{Ballester06}), and magnetic loop oscillations (\\eg, \\opencite{Nakariakov05}; \\opencite{Aschwanden06}). Sometimes, flux of non-thermal flaring emission reveals QPPs with periods in the range from milliseconds to minutes and in different frequencies from radio waves to HXRs (see review by \\opencite{Aschwanden87}; and more recent \\opencite{Aschwanden03}; \\opencite{Nakariakov05}). It seems difficult to explain all variety of such QPPs by a unified physical process.  In this paper we concentrate only on long-periodic pulsations (with periods $>5$ \\textrm{s} according to the classification given by \\opencite{Aschwanden03}) simultaneously observed in the HXR and microwave ranges which can be manifestations of single population of accelerated electrons in the solar atmosphere, because of similarity of these emissions' light curves during many flares (\\eg{}, see for details a review by \\opencite{Bastian98}). It's accepted that a bulk of non-thermal HXRs in solar flares is the thick-target bremsstrahlung radiation of non-thermal electrons ($20-100$ \\textrm{keV}) precipitating along foots of magnetic loops into the chromosphere, whereas the microwave emission is produced via the gyrosynchrotron process of accelerated electrons ($>100$ \\textrm{keV}) interacting with magnetic field. QPPs with such periods often assumed to be generated by some kind of MHD oscillations of a flare coronal loop with fixed footpoints which contains already accelerated electrons (apparently, the idea have appeared firstly in \\opencite{Rosenberg70}, although it was used for interpretation of the fast $\\sim 1$ \\textrm{s} broad band fluctuations). Such MHD oscillations can cause periodic oscillations of the magnetic loop cross section. Consequently, the magnetic field strength and magnetic mirroring can be modulated and the flux of already accelerated electrons, precipitating towards the loop's footpoints, becomes modulated too, causing QPPs of the HXR emission (\\eg{}, \\opencite{Zaitsev82}; \\opencite{Zaitsev89}). QPPs of the microwave emission can naturally appear in this scenario, because of magnetic field variations. \\inlinecite{Brown75} developed a bit different model of the QPPs based on a betatron action of the flaring vibrating magnetic bottle on the already accelerated and trapped electrons changing their spectrum and flux.  Another class of models based on the idea, that MHD oscillations of coronal loops can modulate efficiency of accelerating process. Thus, \\inlinecite{Asai01} suggested under the \\inlinecite{Tsuneta98} flare model, that QPPs can be caused by quasi-periodic modulation of the fast shock's length (that can accelerate electrons by the first-order Fermi process) via oscillations of an underlying flaring loop. Be guided by their own observations of the large transequatorial loop, which places near the flaring region of thermal X-ray emission with QPPs, \\inlinecite{Foullon05} suggested another scenario according to which QPPs can be caused by periodic pumping of electrons into an acceleration region by the linked oscillating loop. \\inlinecite{Nakariakov06} proposed the similar model but with modifications, according to which a modulation of electric current density in an acceleration region can be produced by penetrating fast magnetoacoustic oscillations from a non-flaring oscillating loop. In its turn oscillations of current density can produce oscillations of the magnetic reconnection process via anomalous resistivity and, as a result, an oscillatory acceleration of electrons. Unfortunately, to the best of our knowledge, there were not yet simultaneous observations of EUV loop oscillations and QPPs of microwave and non-thermal HXR emission, although spatially resolved oscillations of a flare loop in the microwave range were found \\cite{Nakariakov03}. This remains possibilities for QPPs to be originated not only by a non-reconnecting oscillatory coronal loop with fixed footpoints.  According to the standard model of erupting solar flares (the CSHKP model; \\opencite{Carmichael64}; \\opencite{Sturrock66}; \\opencite{Hirayama74}; \\opencite{Kopp76}), an emerging filament can stretch overlying magnetic field lines and form a quasi-vertical current sheet, where magnetic field lines can reconnect releasing the stored magnetic energy that can be partially converted into accelerated charged particles and consequently in the HXR and microwave emission. The released energy rate, $dW/dt$, can be estimated as a product of the area of the reconnection region and the Poynting flux, $\\left|\\vec{S}\\right| \\sim \\left|\\vec{E_{cor}}\\times\\vec{B_{cor}}\\right|$ (\\opencite{Isobe02}), if one assume that all energy is released (this is a serious assumption). Coronal electric field, $E_{cor}$, can be estimated as $v_{in}B_{cor}$, where $v_{in}$ is the inflow velocity of coronal magnetic field lines with $B_{cor}$ drawing into the reconnection region. Thus, \\begin{equation}  \\label{Eq_1} dW/dt \\sim v_{in}B_{cor}^{2}, \\end{equation} if one assumes the constant area of the reconnection region.  The inflow velocity can be estimated if one observes an upward motion of flaring loops with a high cadence and spatial resolution (say, in EUV or in soft X-rays), that can be successively involved in the moving reconnection site. This process reveals itself via EUV or H$\\alpha$ flare ribbons separation preferably perpendicular to the magnetic inversion line and via growing loop arcades. \\inlinecite{Sui04} using \\textit{Reuven Ramaty High-Energy Solar Spectroscopic Imager} (RHESSI; \\opencite{Lin02}) found correlation between the rate of the flare loops upward motion and HXR flux. Theoretically, motions of the reconnection site can have an oscillatory behavior, as it was shown by numerical MHD modeling \\cite{Chen99}. But one exciting observation with the \\textit{Transition Region and Coronal Explorer} (TRACE; \\opencite{Handy99}) of the oscillatory shrinkage of flaring loops in 195 \\textrm{\\r{A}} \\cite{Li06} did not reveal clearly an oscillatory flux of the RHESSI HXR emission, as it could be expected. Nevertheless, theoretical evidences of an oscillatory regime of the magnetic reconnection under the coronal conditions were shown in quantity (\\eg{}, \\opencite{Tajima87}; \\opencite{Kliem00}; \\opencite{Ofman06}) and we have to wait more precise observations to examine this possibility.  Another possibility of a QPP-generating mechanism consists in successive acts of energy release in different places of flaring arcade (\\eg, \\opencite{Vorpahl76}). As multiple and quite random non-periodic pulsations of non-thermal emission are often observed during ribbon flares, but QPPs are rare phenomena, QPPs may be just a special case of non-periodic pulsations under a unique configuration of magnetic field. An impressive example of pulsatory RHESSI HXR emission (but without obvious periodicity) from a two-ribbon flare was presented by \\inlinecite{Grigis05}, were analysing motions of footpoint-like HXR sources. Authors showed that each observed ``elementary flare bursts'' were not modulated by a single oscillating flare loop with stationary footpoints or by an oscillating reconnection process in a confined flare region, rather by a process of reconnection site motion along the flare arcade. Since detailed observations of coronal source motions are succeeded rarely, but motions of HXR footpoints are observed in quantity (using Yohkoh and RHESSI), $v_{in}$ was related to a component of velocity of footpoint-like HXR sources which is perpendicular to the magnetic inversion line, $v_{ftp}^{\\bot}$, in the frame of the simple 2D reconnection model (under the assumption of the translation simmetry along a flare arcade; see review by \\opencite{Priest02}). Three extra assumptions were made for this: (a) just reconnected field lines have footpoints close to the footpoints of previous lines; (b) the magnetic flux conservation is satisfied, \\ie{}, $v_{in}B_{cor}=v_{ftp}^{\\bot}B_{ftp}$, where $B_{ftp}$ is the magnetic field in the footpoints of a reconnecting coronal loop; (c) magnetic field in the recconection site in the corona is proportional to the magnetic field in the footpoints of the loops. Thus, the relation~(\\ref{Eq_1}) can be rewritten as \\begin{equation}  \\label{Eq_2} dW/dt \\sim v_{ftp}^{\\bot}B_{ftp}^{2}. \\end{equation} \\inlinecite{Grigis05} did not find correlation between the rate of footpoints motion perpendicular to flare ribbons and the HXR flux as it can be expected from the simple 2D models. Authors suggested an interesting idea amending these models: a filament may erupt unevenly along its length.  Several flares accompanied by motions of the footpoint-like HXR sources along the magnetic inversion line or flare ribbons were really found in quantity (\\eg{}, \\opencite{Fletcher02}; \\opencite{Bogachev05}; \\opencite{Jing07}; \\opencite{Gan08}) using HXR data from Yohkoh and RHESSI. Motions of H$\\alpha$ bright kernels along the flare ribbons were also found, \\eg{}, by \\inlinecite{Qiu02}. Authors of the papers also concluded that it was difficult to interpret their observations in terms of the simple 2D reconnection model. Most probably, the source of primary energy release can somehow propagate mainly along the sheared flaring arcade in some ribbon flares. Quasi-pulsatory behavior of non-thermal emission may be a consequence of sequentially bursting adjacent arcade loops with quite similar physical properties (size, magnetic field, density).  One can see that the problem of solar flare QPPs in non-thermal range is quite intricate and is still far from its complete solving. Many possibilities were supposed to interpret  quasi-periodicities of the observed emission light curves. Recently, \\inlinecite{Li08} analysing one solar flare event accompanied by QPPs of HXR and microwave emission came to the conclusion that their event ``cannot be easily explained with the existing mechanisms''. Thus, it seems reasonable to find and analyse another flares with QPPs of non-thermal emission using imaging observations.  The main goal of the paper is to examine connections between spatial and temporal behavior of non-thermal HXR emission observed by RHESSI during a quite rare phenomenon - the clearly marked phase of minute QPPs during the solar flares on 2003 May 29 and 2005 January 19, and to verify if some changes occur in morphology of the HXR sources under transition from the QPP-phase to the post-QPP-phase. Unfortunately, it is impossible to examine dynamics of coronal magnetic structures in the EUV range, because of the instrumental problems. One extra goal is to examine validity of the simple 2D reconnection model for the concerned events using the relation~(\\ref{Eq_2}). RHESSI observations of the January 19 flare, but without a detailed analysis of spatial evolution of the HXR sources, was reported by \\inlinecite{Nakariakov06} and \\inlinecite{Ofman06} as an impressive example of the flare with QPPs. Authors proposed different mechanisms to generate the observed QPPs. \\inlinecite{Saldanha08} and \\inlinecite{Grigis08} reported the detailed HXR spectral analysis and also some imaging analysis of the same event but without giving consideration to the QPP-problem. Detailed spectral analysis of non-thermal emission during the QPP-phase of the 2003 May 29 flare was reported by \\inlinecite{Minoshima08} and by \\inlinecite{Ning07}. Thus, we will restrict ourself only by the detailed imaging analysis of these QPP-events. ",
        "conclusions": "% \\label{S-Discussion} Analysis of the RHESSI observations of two solar flares which were accompanied by minute QPPs of non-thermal HXR and microwave emission reveals motions of thermal and non-thermal HXR sources. Microwaves and the thermal HXR emission were emitted mainly from the arcade apex, but non-thermal HXRs were emitted successively from the footpoints of the adjacent coronal loops stacked into the arcades. This is especially seen for the flare on January 19, when the HXR source was systematically moving along the flare ribbons about 20 minutes. The general velocity component of the moving sources was directed parallel to the magnetic inversion line (the May 29 flare) or parallel to the flare ribbon (the January 19 flare). The peak-to-peak correlation between the velocity of the HXR sources and the flux of non-thermal emission was not found. Moreover, anticorrelation between the mentioned parameters is observed during some moments of the flare. The sources movement looks like abrupt jumps mainly along the arcade and sitting a bit during the QPPs' maxima. These may indicate a non-validity of the simple 2D model of the eruptive flares based on the magnetic reconnection process in the observed events. What useful information could be obtained from these imaging observations to understand the observed quasi-periodicity of non-thermal emission and its relation to the processes of primary energy release?  First, we can conclude, that the separate acts of energy conversion from accelerated electrons into the neutral emission have occurred by piecemeal in different flaring loops rather than in a single one (we note that as before we can't observe the sites of primary energy release). Thus, that models of the flare QPPs, based on a single oscillatory flaring loop, don't work here. Apparently, this is the only evident conclusion from our observations. The following is just a discussion.  Probably, an interacting loop model proposed by \\inlinecite{Emslie81} also does not work in its direct form in the observed events. According to this model the separation distance between two interacting and bursting loops, $D$, can be roughly estimated as \\begin{equation}  \\label{Eq_3} D \\sim v_{A}T\\approx\\frac{2.18\\times{10^{11}}\\times{B}\\times{T}}{\\sqrt{n}} \\end{equation} in cm, where $v_{A}$ is the Alfven velocity, $B$ is magnetic field in \\textrm{G}, $T$ is the time between two consecutive bursts in \\textrm{s}, $n$ is the electron density in \\textrm{cm$^{-3}$}. In the case of the May 29 flare $T\\approx{60}$ \\textrm{s} and $B\\approx{500}$ \\textrm{G}  near the arcade apex as it was estimated from the calculated potential field. The emission measure is $\\approx(1.6 - 5.4)\\times{10^{49}}$ \\textrm{cm$^{-3}$} from the GOES observations during the QPP-phase. One can roughly interpret the emitting volume, $V$, as the entire flare arcade which can be approximated by a half of cylinder with the length of $\\approx40^{\\prime\\prime}$ or $\\approx3\\times10^9$ \\textrm{cm} and the diameter of $\\approx25^{\\prime\\prime}$ or $\\approx1.8\\times10^9$ \\textrm{cm}. Hence, $V\\approx3.8\\times10^{27}$ \\textrm{cm$^3$} and $n\\approx{(0.7 - 1.2)\\times10^{11}}$ \\textrm{cm$^{-3}$}. Thus, $D\\approx{(1.9-2.5)\\times{10^{10}}}$ \\textrm{cm} that is one order of magnitude larger than the length of the entire flare arcade. Even if we assume $B\\approx{50}$ \\textrm{G}, $D$ will be equal to the length of the arcade. But $D$ must be less than the arcade length, because observations show that non-thermal HXRs are emitted from the arcade foots. The same we can expect for the January 19 flare, which has $T\\approx{180}$ \\textrm{s} and $n\\approx{(0.4-0.6)\\times10^{11}}$ \\textrm{cm$^{-3}$}. Unfortunately, we don't have information about magnetic field, but it is reasonably to suppose $B\\approx100$ \\textrm{G}, that is less than it was supposed for the May 29 flare, because the sunspots are also less in this active region. Thus, $D\\approx{(1.6-2)\\times10^{10}}$ \\textrm{cm}, that is also a very large value. Therewith, the evaluated Alfven velocity is $\\sim{1000}$ \\textrm{km s$^{-1}$}, that is several times larger than the observed velocity of HXR sources (the same was observed by \\opencite{Grigis05}). From here we can also conclude that a magnetosonic wave with the velocity $V \\approx{\\sqrt{v_{A}^{2}+v_{S}^{2}}}$, where $v_{A}<<c$ and $v_{S}$ is the sound velocity, can not also be considered as the propagating trigger of the primary energy release as it was suggested by \\inlinecite{Vorpahl76}. As far as the standard 2D reconnection model of the flare does not work properly in the observed events and the potential magnetic field lines don't conjugate the observed HXR sources well, one can expect that the observed arcades are sheared. So, different modes of the MHD waves may be presented, but it is beyond the scope of this paper to discuss them.  A scenario based on an oscillatory loop arcade as a whole, which is sometimes observed (\\eg{}, \\inlinecite{Verwichte04} have observed standing kink oscillations of post-flare arcade loops), can not be ruled out with certainty, although there were not direct observations of such loops behavior in the examined events. We can imagine two principally different possibilities of such models to be without going into details in this work:\\\\ \\textbf{(a)} consecutive involvement of similar and adjacent arcade loops into the processes of primary energy release, acceleration, and precipitation of electrons, like in an interacting loop model but via, say, transverse oscillations in arcade loops and further via the mechanism proposed by \\inlinecite{Nakariakov06};\\\\ \\textbf{(b)} simultaneous trapping of the bulk of accelerated electrons by an entire flaring arcade and consecutive pulsed precipitation in different loops, because of that sort of MHD waves which can propagate along the arcade and decrease the magnetic mirror ratio in each loops. The \\textbf{(b)}-type models seem quite doubtful for the explanation of the observed events. It is difficult to imagine a long effective trapping of many electrons without a significant decrease of their amount as it was in the case of the QPPs 5 and 6 of the January 19 flare about $15-20$ minutes after the first non-thermal spike. Therewith, imaging analysis of the X-class flare on 2005 January 17 from the same active region with the 19$^{th}$ January flare made by \\inlinecite{Temmer07} in H$\\alpha$ and HXR ranges shows that the bulk of electrons was mainly accelerated and injected into a certain set of arcade loops and only a small fraction of electrons was injected to the entire flare arcade. Though, ribbon-like HXR sources were observed, although it is a rare phenomenon (\\inlinecite{Masuda01}; \\inlinecite{Jing07}).  Evidence in favor of the \\textbf{(a)}-type scenario can be found in the work of \\inlinecite{Minoshima08}. Authors compared the observed behaviors of the HXR and microwave emission with their numerical simulations of the static ``trap-plus-precipitation'' model (\\eg{}, \\opencite{Melrose76}) for two adjacent pulsations of the May 29 flare. In the frame of their model the observed behaviors of the separate pulsations can be qualitatively explained by the separate injections of accelerated electrons into a magnetic loop having the observed parameters. Authors considered the same loop parameters for the both spikes of emission. Such approach could occur even in the case of the separate traps, because the loops are quite similar in the concerned part of the arcade (Figure~\\ref{F-5}). Also, the soft-hard-soft spectral behavior of the HXR emission was successfully explained by an electron injection function with the same spectral property.  For all that, from our observations it is impossible to finally conclude what physical mechanism forces primary energy to be released quasi-periodically in different places of the active region. We have to wait new simultaneous imaging observations in the microwave, EUV, and HXR ranges of solar flares with quasi-periodic pulsatory non-thermal emission.  \\clearpage{}"
    },
    "0809/0809.2417_arXiv.txt": {
        "abstract": "We study the evolution of the first stars in the universe (Population III) from the early pre--Main Sequence (MS) until the end of helium burning in the presence of WIMP dark matter annihilation inside the stellar structure. The two different mechanisms that can provide this energy source are the contemporary contraction of baryons and dark matter, and the capture of WIMPs by scattering off the gas with subsequent accumulation inside the star. We find that the first mechanism can generate an equilibrium phase, previously known as a {\\it dark star}, which is transient and present in the very early stages of pre--MS evolution. The mechanism of scattering and capture acts later, and can support the star virtually forever, depending on environmental characteristic of the dark matter halo and on the specific WIMP model. ",
        "introduction": "Within the scenario of a $\\Lambda$CDM cosmology, it has been recently recognized that if the dark matter (DM) dominant component is a weakly interacting massive particle (WIMP), its annihilation could play a relevant role on the formation and evolution of the first baryonic objects in our universe. Spolyar, Freese, and Gondolo (2008) noticed that during the proto--stellar phase, the cooling of the baryonic gas could be overcome by the energy deposition following the annihilation of DM concentrated in the star formation site. This is a consequence of the peculiar formation characteristics of the first stars, at the center of a minihalo whose gas cooling is dominated by the little efficient primordial chemistry; the authors suggested that this phase (called a {\\it dark star}) could prevent the formation of the first stars, and be a new phase of stellar evolution. Iocco (2008) and Freese, Spolyar, and Aguirre (2008) noticed that, if a star does eventually form, the process of WIMP capture by scattering could be so efficient that their subsequent accumulation and annihilation inside the celestial object may provide an energy source comparable or even exceeding its nuclear luminosity. Motivated by these works, Iocco et al. (2008) and Freese et al. (2008b,c) have studied the early pre--MS phase, in which the baryonic structure is sustained by the annihilation of the DM accreted inside it by gravitational contraction, with the help of numerical codes. Both the groups find this phase is transient, although the techniques adopted are different and the details of the treatment lead to different duration estimates; they both conclude, however, that the collapse must continue at the end of this process, which for the sake of simplicity we call the Adiabatic Contraction ({\\it AC}) phase. In Iocco et al. (2008) (hereafter, I08), we have also studied the pre--MS phase of these stars in presence of annihilating scattered and captured DM ({\\it SC} phase) and followed the evolution of stellar models of different mass until the end of the helium burning, in different DM environments. We find the duration of the MS is dramatically prolonged, up to a potentially everlasting phase (depending  on the choice of parameters, see later): the energy released by the annihilating DM can support the core at temperatures low enough that nuclear reactions are never ignited. Yoon, Iocco, and Akiyama (2008) and Taoso et al. (2008) also studied the {\\it SC} phase, confirming our results with different codes and carrying their analysis further. These proceedings are based on the results obtained by I08, to which we address the reader for detailed referencing and more quantitative details: here we aim to a more qualitative description of the physical processes at the basis of this class of objects. ",
        "conclusions": "We have studied the effects of WIMP dark matter annihilation from the late stages of stellar formation through the end of helium burning. Addressing the reader to Iocco et al. (2008) for quantitative details, here we wish to stress the fundamental physical difference between the {\\it SC} mechanism, which needs a weak process to mediate the capture and thermalization of DM particles (and therefore the DM density profile inside the star), and the {\\it AC} one, where the DM density profile is dictated by gravitational contraction only. The difference in the two physical processes dictates diverse characteristics for the two phases: the {\\it AC} is transient, and takes place early in the pre--Main Sequence evolution of the star, when it is almost at the proto--stellar stage. The {\\it SC} phase becomes active when the star is at the bottom of the Hayashi track or later in the pre--MS evolution, and it can dramatically extend the stellar life, up to orders of magnitude more than its standard lifespan, depending on parameters.   As a note added in proof, we address the reader to a paper very similar to this one, Freese et al. (2008d) in these proceedings; and also to the analysis of Natarajan, Tan, and O' Shea (2008) who find results in good agreement with Spolyar et al. 's, studying gas and dark matter profiles from cosmological simulations of first star formation."
    },
    "0809/0809.2598_arXiv.txt": {
        "abstract": "We introduce a modified detection method for measuring the luminosity of the tip of the red giant branch (TRGB) by introducing the composite magnitude $T \\equiv I - \\beta [(V-I)_{\\circ} - 1.50]$, where $\\beta$ is the slope of the tip magnitude as a function of color (or metallicity).  The method is specifically designed to account for known systematics due to metallicity.  In doing so, this simple transformation does away with arbitrary color selections in measuring the tip, and thereby significantly boosts the population of resolved stars that go into defining the TRGB distance.  Moreover this method coincidentally reduces the impact of reddening on the true modulus as well as its final uncertainty. ",
        "introduction": "The luminosity of the tip of the red giant branch (TRGB) has been profitably used by many groups to determine high-precision distances to nearby galaxies (for example, see the compilation of Ferarrese et al.  2001). And, by all measures this method continues to grow in popularity (e.g., Rizzi et al. 2007 and references therein). Its success is driven, in all probability, by its low cost in observing time, its conceptual simplicity, and its wide range of application (see Madore \\& Freedman 1998 for a review).  In our earliest paper on quantifying the procedures for measuring the magnitude of the TRGB tip (Lee, Freedman \\& Madore 1993) we first recommended the use of a digital (Sobel) filter operating on a binned histogram representation of the RGB luminosity function. Being aware of the slight dependence of the RGB tip on metallicity (as manifest by a monotonic decline of the TRGB I-band magnitude with increasing color) we also outlined (a regretably convoluted) means of correcting the observed tip magnitude back to a fiducial metallicity/color-corrected tip magnitude based largely on theory. With the clarity of hindsight (drawing on better data, with continuing theoretical support) we now correct that shortcoming and suggest a means of measuring a TRGB distance modulus that is explicitly corrected for metallicity at the filtering level.  In correcting for metallicity the new methodology introduced below is simple, but its additional utility and power comes from the fact that it allows the use of a much enhanced sample of tip stars (independent of their metallicity) to bolster the statistical certainty of the distance modulus determination.  The method outlined below is, at the same time, both systematically and statistically more powerful and robust than previously-used methods. ",
        "conclusions": "As relatively insensitive to metallicity as the I-band tip of the red giant branch is, there is still observed to be a slight dependence of the TRGB magnitude on the metallicity as tracked by the intrinsic color. We have the means to account for it, since that dependence has been found to be extremely stable from galaxy to galaxy and is well calibrated (and indeed well understood). We have incorporated that new information into a very simply modified TRGB edge detector that makes distance determinations using the TRGB explicitly independent of the mean metallicity and/or metallicity distribution of the stars in a given halo population. In the process this new method allows all tip stars to add to the statistical significance of the detection of the core helium ignition discontinuity.  As an added bonus the modified method is also slightly less sensitive (by 10\\%) to reddening corrections (and/or uncertainties in those same reddening corrections) as compared to the standard (fixed-metallicity) option.  The method does, of course, now require that both the V and I magnitudes be obtained for the color-correction to be applied. Previous applications measuring the tip discontiuity from only an I-band luminosity function could not be corrected for metallicity effects, were always met with justifiable skepticism, and are not to be recommended.  We thank the referee, Brent Tully, for his careful reading of the paper and his useful suggestions and comments."
    },
    "0809/0809.2883_arXiv.txt": {
        "abstract": "The solar convective zone, or SCZ, is nearly adiabatic and marginally convectively unstable.  But the SCZ is also in a state of differential rotation, and its dynamical stability properties are those of a weakly magnetized gas.  This renders it far more prone to rapidly growing rotational baroclinic instabilities than a hydrodynamical system would be. These instabilities should be treated on the same footing as convective instabilites.  If isentropic and isorotational surfaces coincide in the SCZ, the gas is marginally (un)stable to {\\em both} convective and rotational disturbances.  This is a plausible resolution for the instabilities associated with these more general rotating convective systems.  This motivates an analysis of the thermal wind equation in which isentropes and isorotational surfaces are identical.  The characteristics of this partial differential equation correspond to isorotation contours, and their form may be deduced even without precise knowledge of how the entropy and rotation are functionally related.  Although the exact solution of the global SCZ problem in principle requires this knowledge, even the simplest models produce striking results in broad agreement with helioseismology data.  This includes horizontal (i.e. quasi-spherical) isorotational contours at the poles, axial contours at the equator, and approximately radial contours at midlatitudes.  The theory does not apply directly to the tachocline, where a simple thermal wind balance is not expected to be valid.  The work presented here is subject to tests of self-consistency, among them the prediction that there should be good agreement between isentropes and isorotational contours in sufficiently well-resolved large scale numerical MHD simulations. ",
        "introduction": "The problem of differential rotation in the solar convective zone (SCZ) has vexed theorists for many years.  At the simplest level, one is faced with the fact that the only region of the sun characterized by significant turbulence is also the only region showing any significant differential rotation.  This is a cogent reminder of the hazards of modeling a rotating, turbulent gas by an effective viscosity parameter. The SCZ is far too complex to yield to such an approach.  Indeed, large scale numerical simulations have been in place for many years, trying to elicit the rotation profile of the sun's outer layers (Brummell, Cattaneo, \\& Toomre 1995; Thompson et al. 2003). While these studies are now at a stage where they can begin to make contact with the observational data, they have not yet been able to reproduce the salient features of the observed solar rotation profile. This is especially true for the isorotational contours (hereafter ``isotachs''), which tend to be cylindrical in simulations, but are decidedly noncylindrical in the sun.  In this paper, we examine an effect that could be important for understanding the SCZ rotation profile, but has not been sufficiently emphasized in previous investigations: the extreme sensitivity of ionized differentially rotating gas to the presence of even very weak magnetic fields of arbitrary geometry (Balbus \\& Hawley 1994, Ogilvie 2007).  Far from behaving as a passive vector field, a weak magnetic field triggers rapidly growing unstable local modes in rotating systems that would be hydrodynamically unstable.  The field endows the gas with degrees of freedom that have no hydrodynamical counterpart. The classical manifestation of this behavior takes the form of what is known as the magnetorotational instability (MRI), widely regarded as the process responsible for turbulence and enhanced transport in accretion disks (Balbus \\& Hawley 1991, 1998).   But destabilization by a weak magnetic field is a process that extends beyond the domain of rotationally supported disks.  Such fields are agents of dynamical destabilization in baroclinic and convective flows as well (Balbus 1995), with potentially interesting applications to the SCZ.  It is essential to understand two key counterintuitive points at the outset.  The first is that destabilization by a magnetic field can occur even when the field plays essentially no role in the dynamical equilibrium: a weakly magnetized gas does not behave ``almost hydrodynamically'' with respect to its stability properties.  The second is that the destabilization is independent of the field strength and field geometry.  In MRI simulations, the disruption remains vigorous and self-sustaining even in highly convoluted turbulent flow.  This is what makes MHD processes potentially so important for understanding rotating, stratified flows.  It is the emphasis on these points that sets the present approach apart from earlier work that also addressed the stability of rotating, stratified, magnetized flows (e.g. Acheson 1983, Ogden \\& Fearn 1995).  In these earlier studies, the focus is upon toroidal fields, nonaxisymmetric disturbances, and instability driven directly by the magnetic field.  By contrast, what is crucial here are poloidal fields and axisymmetric disturbances.  Moreover, while the instability of interest relies on the presence of a magnetic field, the (nonconvective) seat of free energy is the differential rotation, not the field itself. The role of the field is to provide the critical degrees of freedom to the fluid required to tap into the (destabilizing) differential rotation.  We shall begin our presentation, however, not with magnetic fields, but with a theoretical observation that may be taken at a purely phenomenological level, independently of deeper MHD considerations. This is the remarkable agreement with the overall pattern of solar isorotation contours produced by a simple calculation in which a dominant thermal wind balance from the vorticity equation is combined with the assumption that entropy and angular velocity gradients are counteraligned. Whatever the cause of the counteralignement may be, the analysis on its own is highly suggestive.  (A connection between entropy and angular velocity gradients has in fact been advocated on purely hydrodynamical grounds [Miesch, Brun, \\& Toomre 2006].)  This material is presented in \\S 2.  In \\S 3, we discuss in some detail the case for an MHD coupling between angular velocity and entropy gradients.  The final section is a critical discussion of unresolved issues provoked by this work, and a summary.     \\section {Solving the thermal wind equation}  \\subsection {Review}\\label{revsec}  Let $(R, \\phi, z)$ be a standard cylindrical coordinate system, and $(r, \\theta, \\phi)$ a standard spherical coordinate system.  Unit vectors are denoted $\\bb{e_R}$, $\\bb{e_\\theta}$, etc.  The angular velocity $\\Omega$ is assumed to be independent of $\\phi$, but otherwise general.  Our notation for the fluid variables is likewise standard: $\\bb{v}$ is the velocity, $P$ is the gas pressure, $\\rho$ is the mass density, and $\\bb{B}$ is the magnetic field.  We adopt a fiducial value of $2.5\\times 10^{-6}$ rad s$^{-1}$ for the angular velocity $\\Omega$ in the SCZ, corresponding to rotation velocities between 1 and 2 km s$^{-1}$.  This is well in excess of the $\\sim 30$ m s$^{-1}$ expected of convective velocities, but such a direct comparison is not necessarily the most relevant one. A more telling comparison is between the squared \\BV frequency \\beq\\label{fiduN} \\left| N^2\\right|  = \\left| -{1\\over \\rho\\gamma}{\\dd P\\over \\dd r} {\\dd\\ln P\\rho^{-\\gamma} \\over \\dd r}\\right| \\sim 3.8\\times 10^{-13}\\ {\\rm s}^{-2} \\eeq (we have adopted a value of $10^{-6}$ for $\\dd\\ln P\\rho^{-\\gamma}/\\dd\\ln r$ [Schwarzschild 1958]) and the rotational parameter \\beq\\label{fiduOm} {d\\Omega^2\\over d\\ln R} \\sim 1.8\\times 10^{-12} \\ {\\rm s}^{-2}. \\eeq This is evidently about 5 times as large as $N^2$. (Whether $R$ or $z$ is used in the angular velocity gradient is not critical at midlatitudes.)  Taking this estimate at face value, the rotational parameter is significantly larger than $N^2$.  The significance of this will shortly become apparent.  Let us first consider systems whose equilibrium velocity is differential rotation, $\\bb{v} = R\\Omega(R, z)\\bb{e_\\phi}$. The thermal wind equation follows directly from the time-steady form of the vorticity equation, \\beq\\label{vort2} R{\\dd\\Omega^2\\over \\dd z} = {1\\over \\rho^2}\\left( \\del P\\btimes \\del\\rho \\right)\\bcdot\\bb{e_\\phi}. \\eeq Expressing the right side in $r,\\theta,\\phi$ spherical coordinates: \\beq R {\\dd\\Omega^2\\over \\dd z} = {1\\over r \\rho^2} \\left( {\\dd\\rho\\over\\dd \\theta}{\\dd P\\over \\dd r} - {\\dd\\rho\\over\\dd r}{\\dd P\\over \\dd \\theta} \\right). \\eeq Now, rewrite this in terms of the entropy gradients: \\beq R {\\dd\\Omega^2\\over \\dd z} = {1\\over C_P  \\rho r} \\left( {\\dd P \\over\\dd \\theta}{\\dd S\\over \\dd r} - {\\dd P \\over\\dd r}{\\dd S\\over \\dd \\theta} \\right), \\eeq where $S$ is the specific entropy, \\beq S={k\\over \\gamma -1} \\ln P\\rho^{-\\gamma} + {\\rm constant}, \\eeq $k$ is the Boltzmann constant, and $C_P$ is the constant pressure specific heat, \\beq C_P =\\left(\\gamma\\over \\gamma - 1\\right) k \\eeq In the SCZ, the $r$ gradient of $P$ clearly dominates over the $\\theta$ gradient, whereas the $r$ gradient of $S$ is unlikely to greatly exceed the $\\theta$ gradient. (In fact, we shall presently argue just the opposite.) We may then conclude that \\beq\\label{thwnd} R {\\dd\\Omega^2\\over \\dd z} = {g\\over \\ C_P r } {\\dd S\\over\\dd\\theta} ={g\\over \\gamma r } {\\dd(\\ln P\\rho^{-\\gamma}) \\over  \\dd\\theta} \\eeq where $g=-(1/\\rho) {\\dd P/\\dd r}$ is the gravitational field in hydrostatic equilibrium, ignoring the small effects of rotation. Significant latitudinal entropy gradients are required to avoid cylindrical isotachs.  Equation (\\ref{thwnd}) is known as the thermal wind equation (Kitchatinov \\& R\\\"udiger 1995, Thompson et al. 2003).  This equation holds in our MHD analysis, because we explicitly assume that the magnetic fields are sufficiently weak in the equilibrium state that they do not affect the large scale rotation profile.  (For a 10 G field and a density of $0.05$ g cm$^{-3}$, the Alfv\\'en velocity is 13 cm per second.) In full spherical coordinates, the thermal wind equation is \\beq\\label{thw2} \\left( \\cos\\theta{\\dd\\Omega^2\\over \\dd r}  - {\\sin\\theta\\over r} {\\dd\\Omega^2\\over \\dd \\theta} \\right) = {g\\over C_P r^2\\sin\\theta } {\\dd S\\over\\dd\\theta} \\eeq  Even before a detailed solution is developed, a simple scaling argument reveals something of interest here.  The $r$ gradient of $S$ is generally determined by the need to transport the solar luminosity by thermal convection (Schwarzschild 1958, Clayton 1983).  On the other hand, the $\\theta$ gradient is, from equation (\\ref{thwnd}), linked directly to differential rotation.  If the fiducial numbers (\\ref{fiduN}) and (\\ref{fiduOm}) are reasonably accurate, at midlatitudes equation (\\ref{thwnd}) implies that the $\\theta$ gradient of $S$ will significantly exceed the $r$ gradient.  The point of interest here is that this anisotropic feature is directly seen in the $\\Omega$ gradient deduced from the helioseismology data.  This suggests that there is a deeper dynamical coupling present between $S$ and $\\Omega$, beyond just the general trend that one goes up as the other goes down.  We will develop this idea more fully in the next two sections.   \\subsection{Analysis}\\label{anall}  The thermal wind equation is a familiar tool to practitioners of solar rotation theory.   The gross features of the sun's angular velocity profile ($\\Omega$ decreasing polewards from the equator) can be understood relatively simply with the aid of this equation and some reasonable assumptions of the efficiency of convection in the presence of Coriolis forces (Thompson et al. 2003).  The idea is that convection in the equatorial direction is impeded by Coriolis forces, resulting in a more efficient transport of heat along the rotation axis. The poles are then regions of higher specific entropy compared with the equatorial zone, and the resulting $\\theta$ gradient in $S$ drives axial gradients in $\\Omega$.  This is a well-motivated and physically plausible scenario.  It seems natural therefore, to search for explicit model solutions to the thermal wind equation incorporating this idea.  Moving upward from the equator, we expect the angular velocity to decrease as the entropy increases. Clearly the gradients of these quantities are in broadly opposite senses. But the final paragraph of \\S \\ref{revsec} suggests that ``counteraligned'' may be a better description then ``broadly opposite.''  This raises the question of what the solutions to equation (\\ref{thw2}) would look like if the two gradients were in fact {\\em precisely} oppositely aligned.  Following the notion that isentropic and isorotational surfaces coincide, we assume that $S=S(\\Omega^2)$. ($S$ should not change if $\\Omega$ changes sign, hence the dependence on $\\Omega^2$.) Equation (\\ref{thw2}) then becomes \\beq\\label{maineq} {\\dd \\Omega^2\\over \\dd r} - \\left( {g S' \\over C_P r^2\\sin\\theta \\,\\cos\\theta}+ {\\tan\\theta\\over r} \\right) {\\dd \\Omega^2 \\over\\dd\\theta} =0, \\eeq where $S'\\equiv dS/d\\Omega^2$. The solution of equation (\\ref{maineq}) is that $\\Omega^2$ is constant along the characteristic \\beq\\label{char1} {d\\theta\\over dr} =  - {\\tan\\theta\\over r} - {g S'\\over C_P r^2\\sin\\theta\\ \\cos\\theta } \\eeq But if $\\Omega^2$ is constant along this characteristic, then so must be $S'$. This is therefore a self-contained, ordinary differential equation for $\\theta$ as a function of $r$, precisely the isorotational contours that we seek. To solve equation (\\ref{char1}), let $y=\\sin\\theta$.  Then, our differential equation simplifies to \\beq {dy^2\\over dr} +{2y^2\\over r} = - {2g S' \\over C_P r^2} \\eeq Multiplying by $r^2$ and regrouping, \\beq {d(r^2y^2)\\over dr} = - {2g S' \\over C_P }= - {2GM_\\odot S' \\over C_P  r^2 } \\eeq where $G$ is Newton's constant and $M_\\odot$ is a solar mass. (We have ignored the local self-gravity of the SCZ.)  Since $S'$ is a constant along the characteristic, this integrates immediately to \\beq\\label{conts} r^2\\sin^2\\theta = R^2 = A - {B\\over r} \\eeq where $A$ is a constant of integration and \\beq B= - {2GM_\\odot S' \\over C_P   }. \\eeq We have inserted a minus sign since $S$ will generally be a decreasing function of $\\Omega^2$.  Indeed, as will become very clear, the characteristics make no sense if $S'$ is positive, but a very great deal of sense if it is negative.  In this model, the solar isotachs are given by a remarkably simple formula.  To estimate the magnitude of $B$, note that it may be written \\beq {B\\over r_\\odot^3}= \\left( 2GM_\\odot/r_\\odot\\over \\gamma r_\\odot^2\\Omega^2\\right) \\left( -{d\\ln P\\rho^{-\\gamma}\\over d\\ln r}\\right) \\left({d\\ln r\\over d\\ln \\Omega^2}\\right) \\eeq The first factor is large, of order $10^5$.   The second, we have already estimated at $10^{-6}$.  The helioseismology data suggest that the third factor ranges between 1 and 10, and is larger near the equator. Crudely speaking, we expect $B/r_\\odot^3$ to be of order unity or less.  There should not be a large difference in scale between $A/r_\\odot^2$ and $B/r_\\odot^3$.  \\subsection {Alternative isotach solution}\\label{2pt3}  By way of comparison, consider the solution of the thermal wind equation that would obtain were the {\\em angular momentum} counteraligned with the entropy, rather than the angular velocity.   (In hydrodynamic baroclinic turbulence, one might expect the coupling to be of this nature since entropy and angular momentum tend to be retained by a displaced fluid element.)  Then $S=S(l^2)$, where $l=R^2\\Omega$ is the specific angular momentum.  We denote $dS/dl^2$ by $S_{l^2}$. Our analysis proceeds along lines identical to \\S\\ref{anall}, and $l^2$ satisfies the PDE \\beq {\\dd l^2\\over\\dd r} -\\tan\\theta \\left( {1\\over r} +{gr^2\\sin^2\\theta S_{l^2}\\over C_P} \\right) {\\dd l^2\\over\\dd\\theta}=0. \\eeq The contours of constant $l^2$ are found to be of the form \\beq\\label{contsl} {1\\over R^2}={1\\over r^2\\sin^2\\theta} = A_{l} +{B_{l}\\over r} \\eeq where $A_{l}$ is an integration constant, and now \\beq B_{l}= - {2GM_\\odot S_{l^2}\\over C_P} \\eeq Note that $r_\\odot B_{l}$ is a dimensionless constant of order unity. We will use this solution as a point of comparison in the next section.  \\subsection {An explicit solution}\\label{anexpl}  Let us turn to the data to see how our solutions fare, beginning with our first, $\\Omega$ based solution. As an illustrative example, consider the reduced problem in which $S'$ is constant, not just along a particular characteristic, but everywhere.  Then $B$ is also everywhere constant.  If $\\Omega$ is now specified on some particular radial shell as a function of $\\theta$, we may write down the solution everywhere.  For this purpose, it is easiest to choose the solar surface radius $r=r_\\odot$.  Let the angular velocity at $r=r_\\odot$ be $\\Omega_\\odot(\\cos^2\\theta_0)$. We use the fit of Ulrich et al. (1988): \\beq\\label{fit} \\Omega_\\odot(\\cos^2\\theta_0)=2\\pi\\left( 451.5 - 65.3 \\cos^2\\theta_0 -66.7\\cos^4\\theta_0\\right)\\ {\\rm nHz.} \\eeq Note that $\\theta_0$ carries a subscript to indicate that it is the particular value of $\\theta$ along each trajectory characteristic (\\ref{conts}) that intersects the surface shell $r=r_\\odot$.  Equation (\\ref{conts}) becomes \\beq\\label{cont1} r^2\\sin^2\\theta = r_\\odot^2\\sin^2\\theta_0 +B\\left( {1\\over r_\\odot} -{1\\over r} \\right), \\eeq or \\beq\\label{cont2} \\cos^2\\theta_0 =  1- x^2 + \\left(B\\over r_\\odot^3\\right) \\left( 1 -{1\\over \\varpi} \\right) \\eeq where \\beq \\varpi=r/r_\\odot, \\qquad x^2 = \\varpi^2\\sin^2\\theta . \\eeq (Equations (\\ref{cont1}) and (\\ref{cont2}) are also valid if $B$ is a function of $\\cos\\theta_0$.) Substituting equation (\\ref{cont2}) for $\\cos^2\\theta_0$ into equation (\\ref{fit}) then generates the solution everywhere in the shell.  \\begin{figure} \\epsfig{file=e1pt5e2pt0.eps, width=8 cm}% \\epsfig{file=e1pt6e2pt0.eps, width=8 cm}% \\caption {Contour plot of $\\Omega(r,\\theta)$ in solar interior, using surface fit of Ulrich et al. (1988).  SCZ boundaries marked in white.  Calculation based on eqs. [\\ref{fit}] and characteristic equations of isorotational contours, [\\ref{cont2}].  $B/r_\\odot^3=0.5\\,{\\rm (left)},\\,  0.6\\, {\\rm (right)}$.} \\end{figure}  The result of this procedure is shown in figure (1) for the cases $B/r_\\odot^3=0.5, 0.6$.  The detailed fit of the contours to the actual data is not perfect---there is no tachocline in these simple models, and the true high latitude isotachs show stronger curvature, following spherical shells, before turning upward.   But the overall trend of the isotachs being predominantly quasi-spherical at high latitudes, increasingly radial at midlatitudes, and axial at small latitudes is unmistakable. Moreover, fitting our solution to the observed surface data, while convenient, is unlikely to show off its best form: thermal wind balance probably breaks down near the solar surface (Thompson et al. 2003). Given the simplicity of our direct approach, the qualitative agreement is both striking and encouraging.  The iso-angular-momentum contours of equation (\\ref{contsl}) can also be used to construct an explicit solution.  In this case, the surface angular momentum $l$ fit is \\beq\\label{ll} l(\\cos^2\\theta_0) =2\\pi r_\\odot^2 \\sin^2\\theta_0 \\left( 451.5 - 65.3 \\cos^2\\theta_0 -66.7\\cos^4\\theta_0\\right). \\eeq Instead of equation (\\ref{cont2}), we have \\beq\\label{cosl} \\cos\\theta_0^2 = 1-\\sin^2\\theta_0= 1 - \\left[ {1\\over x^2} +r_\\odot B_{l} \\left(1-{1\\over \\varpi} \\right) \\right]^{-1} \\eeq Substitution of (\\ref{cosl}) into (\\ref{ll}) generates the full solution for the specific angular momentum $l(r,\\theta)$, and the angular velocity solution follows immediately from \\beq\\label{cosll} \\Omega(\\cos^2\\theta_0) = (2\\pi/x^2)\\sin^2\\theta_0 \\left( 451.5 - 65.3 \\cos^2\\theta_0 -66.7\\cos^4\\theta_0\\right) \\eeq  \\begin{figure} \\epsfig{file=ange1pt2.eps, width=8 cm}% \\epsfig{file=ange1pt5.eps, width=8 cm}% \\caption {Contour plot of $\\Omega(r,\\theta)$ in solar interior, based on counter aligned entropy and angular momentum gradients, using surface fit of Ulrich et al.\\ (1988).  SCZ boundaries marked in white.  Calculation based on eq. [\\ref{cosll}] and characteristic equation of iso-angular-momentum contours, [\\ref{ll}]. $r_\\odot B_{l^2}=0.2\\,{\\rm (left)},\\,  0.5\\, {\\rm (right)}$. } \\end{figure}  In figure 2, we show two representative diagrams of the isorotational contours taken from this alternative angular momentum based approach. In general the contours are too cylindrical, to some extent exhibiting the same syndrome often seen in numerical SCZ simulations.  The contrast between figures 1 and 2 is very apparent.  There seems to be a real linkage between $S$ and $\\Omega$, and it matters very much that the coupling is between $S$ and $\\Omega$, not $S$ and $l$.  It is possible that the refractory nature of the cylindrical contours of the simulations is due to an $S-l$ coupling that remains too strong, as noted in in \\S \\ref{2pt3}.  While it is possible that this may be cured by a more highly resolved treatment of the turbulent fluid, it is also possible that magnetic fields may be playing a non-negligible role, enforcing an $S-\\Omega$ coupling by field line tethering of the fluid elements. We pursue this possibility in \\S 3.  In the remainder of the paper, we focus exclusively on our original $S-\\Omega$ solution.  \\subsection {Tightening the contours}  By allowing the $B$ parameter to vary from one isotach to another, the isorotational contours we have found can become more tightly spaced near the poles. In this sense, our solutions admit, but do not demand, something reminscent of tachoclinic structure. Equation (\\ref{conts}) may be written \\beq\\label{tach1} r= {B/A\\over 1 -(R^2/A)} \\eeq If the variation of $\\Omega$ leads to a nearly constant ratio for $B/A$ close to the outer radius of the radiative zone $r_{rad}$, but very different $A$ values, tachoclinic structure results.  In the polar regions near the axis of rotation, the contours would all converge to $r_{rad}$, while corresponding to very different $\\Omega$ values. This would look very much like a tachocline.  Does this make physical sense? Equation (\\ref{cont1}) implies $A=r_\\odot^2 \\sin^2\\theta_0 +B/r_\\odot$, or \\beq {B\\over A} =  {r_\\odot\\over 1 + (r_\\odot^3\\sin^2\\theta_0/B)}, \\eeq assuming that the data are initially specified on $r=r_\\odot$. Hence, $B\\propto \\sin^2\\theta_0$ would produce tachoclinic structure.   Physically , this would mean that $B\\propto dS/d\\Omega^2$ is small and negative near the pole, and that the entropy decreases toward the equator as $\\Omega$ increases.  In the vicinity of the equator the entropy drops sharply.  This is not unreasonable behavior: near the pole, unencumbered by Coriolis deviations, convection is most effective, whereas at the equator, the opposite is true.  At the same time, the observations suggest that $\\Omega$ is changing rapidly near the pole, and only slowly at the equator.  This is consistent with our simple picture.  We return to equation (\\ref{cont2}) and replace $B/r_\\odot^3$ by \\beq\\label{eta1} B/r_\\odot^3 = \\eta_1 +\\eta_2\\sin^2\\theta_0. \\eeq Mathematically, this corresponds to the first two terms in a Taylor series expansion of $S'$ as a function of $\\sin^2\\theta_0$ (or equivalently $\\cos^2\\theta_0$). By varying $\\eta_1$ and $\\eta_2$ we can go between a singular tachocline $\\eta_1=0$ and the previous case $\\eta_2=0$.  Substituting (\\ref{eta1}) in (\\ref{cont2}) and solving for $\\cos^2\\theta_0$ gives \\beq\\label{tacos} \\cos^2\\theta_0 = 1 - {\\varpi x^2  -\\eta_1(\\varpi-1) \\over \\varpi+\\eta_2(\\varpi-1)}. \\eeq Using equation (\\ref{tacos}) in (\\ref{fit}) now produces the interior structure $\\Omega(r,\\theta)$.  With two parameters available, one can of course produce somewhat better fits to the rotation profile.  In practice, the improvement over \\S\\ref{anexpl} is noticeable, but not dramatic.  In figure 3, we show on the left the isotachs for $\\eta_1=0.3$, $\\eta_2=0.2$.  This gives a very respectable fit to the data away from the tachocline, $r\\ge 0.75 r_\\odot$, say.  If we wish to include an explicit tachocline in our modeling, our earlier considerations suggest that we should restrict ourselves to small values of $\\eta_1=0$.  On the right side of figure 3, the interesting case of $\\eta=0.12$ and $\\eta_2=0.8$ is presented.  A striking ``tachocline'' structure appears, though formally it lies just beneath the SCZ.  The solar midlatitude radial contours and equatorial cylinders regions are, however, rather well-represented.  Note as well the gentle nonmonotonic behavior of $\\Omega(R)$ near the equator, a feature seen in the helioseismology data. Increasing the value of $eta_2$ to bring the tachocline to larger radii (while keeping the surface layers fixed) appears to cause too much global distortion, though an exhaustive parameter search has not been performed.   \\begin{figure} \\epsfig{file=e1pt3e2pt2.eps, width=8 cm}% \\epsfig{file=e1pt12e2pt8.eps, width=8 cm}% \\caption {As in figure 1, but calculation now based on equations [\\ref{fit}] and [\\ref{tacos}] with $\\eta_1=0.3$, $\\eta_2=0.2$ (left), and $\\eta_1=0.12$, $\\eta_2=0.8$ (right).  The two-parameter fit on the left is a slight improvement over the earlier single parameter models.  The fit on the right is striking in its overall resemblance to the SCZ, though the ``tachocline'' formally lies below the convective zone lower boundary.  Equatorial and midlatitude contours are well-represented.  } \\end{figure}  At this stage, these results are suggestive, but not more than that. It seems likely that more complex choices for $S'(\\Omega^2)$ could improve the contour fits, but the agreement is impressive even in the simplest models. A truly compelling explanation of the tachocline will involve more than just the thermal wind equation and the outer convective zone layers.  The solar tachocline arises from the complex coupling of the rigidly rotating radiative core and the overlying strongly shearing convective zone, not the demands of the surface rotation and vorticity conservation, and very different dynamical processes are likely to be involved.  Our simple, prescription valid above the tachocline ($r\\gtsim 0.75 r_\\odot$) may represent a sort of outer SCZ solution that asymptotically matches onto an inner solution in which tachocline dynamics become locally dominant.  \\subsection {Generic features of isotachs}  The isotachs we have found have a very distinctive ``viking helmet'' structure.  This unusual feature is also characteristic of solar isorotation contours, and worth examining in isolation.  To understand the general structure of the isotachs, rewrite eqation (\\ref{tach1}) as \\beq {r\\over r_\\odot} = {B/Ar_\\odot \\over 1 - (r_\\odot^2/A)(R^2/r_\\odot^2) } ={\\alpha\\over 1-\\beta R^2/r_\\odot^2} \\eeq which defines our two parameters $\\alpha$ and $\\beta$, \\beq \\alpha=B/r_\\odot A, \\qquad \\beta=r_\\odot^2/A \\eeq We now transform to Cartesian coordinates in a meridional plane, \\beq y= (r/r_\\odot)\\cos\\theta, \\quad x= (r/r_\\odot)\\sin\\theta. \\eeq The Cartesian contour structure is slightly easier to visualize: \\beq\\label{ycont} y = \\left[ {\\alpha^2\\over (1-\\beta x^2)^2 } -x^2 \\right]^{1/2} \\eeq  We must of course view equation (\\ref{ycont}) through the convective zone slot $0.7<\\sqrt{x^2+y^2}<1$.  There are three types of contours. The first is a typical polar contour, emerging perpendicular to the axis before bending upward, seen at high latitudes in figure 1.  The second class, visible at midlatitudes in figure 1, is hidden in the radiative zone at small $x$, and does not emerge into the SCZ until the contour is well separated from the axis, at which point the isotach has a quasi-radial character.  The third contour class is deeply buried, running along a very small radius in the core (not seen), completely disappearing in the bulk of the sun ($y$ is imaginary). It then makes a sudden leap upward into the equatorial region of the SCZ, where it appears nearly axial.  These are the low latitude regions of figure 1, corresponding physically to Taylor columns of constant rotation on cylinders.  Equation (\\ref{ycont}) thus displays the three key traits one associates with isotachs: horizontal near the poles, radial at midlatitudes, and cylindrical near the equator. ",
        "conclusions": ""
    },
    "0809/0809.4188_arXiv.txt": {
        "abstract": "{In this work we perform the first test of a stellar spectral classification method (Stock \\& Stock 1999) by applying it to early type stars. The sample of stars are the members of the open cluster IC~2391 that have high-resolution spectra available in the UVES Project of Paranal Observatory. We show that, in general, absolute magnitudes $M_V$ and intrinsic colors $(B-V)_0$ can be recovered within the uncertainties stated in the original calibration ($\\sim 0.4$ for the magnitudes and $\\sim 0.03$ for the colors). This accuracy allows us to estimate distances and to infer membership of individual stars to obtain an average distance to the cluster of $156\\pm 24$ pc, which is in good agreement with previous reported determinations. Finally, we identify and discuss the real strengths and limitations of this method and we suggest how it can be improved for future studies.}    \\addkeyword{Methods: data analysis} \\addkeyword{Stars: fundamental parameters, classification} \\addkeyword{Open cluster: individual (IC~2391)}  \\begin{document} ",
        "introduction": "The absolute magnitude and the intrinsic color are the main quantities to describe stars. The first provides the amount of energy emitted, and combined with the apparent magnitude gives an estimate of the distance to the object. The second is directly related to the temperature of the atmosphere of the star. These properties are expected to imprint themselves on the optical band of the spectrum, specially trough the absorption lines of different species. The spectral classification methods are advocated to infer these parameters by the inspection of those absorption profiles, and it is often made by visual inspection, as is the case for the MK system. New {\\it quantitative} classification methods are needed to improve the results and provide consistency in the process. Many authors have pointed out the need of a fully automated system to classify stellar spectra \\citep[see for example the reviews by][]{Bai02,All04,Gir06}. Automated systems are more homogeneous (and therefore more appropriate) for the classification of a large number of stars on the basis of a self-consistent system.  Due to the progress in astronomical instrumentation and to the increase in the data acquisition rate, new automated and efficient systems are constantly being developed and improved to classify stellar spectra and/or to infer stellar atmospheric parameters from spectral profiles \\citep[recent examples include][]{Chr02,All03,Bai05,Zha05,Rec06,Ref07}. An example of this kind of systems is the method proposed by \\citet[][hereafter SS99]{SS99}. This method is based on the determination of a spectral index defined by two external bands placed at each extreme of a central band \\citep{Wor94}. This index is related with the stellar parameters through second order polynomials. The calibration of this (and any) method is performed by comparing the defined indices with physical parameters of interest, such as the absolute magnitude $M_V$ or the color index $(B-V)_0$. The errors in these parameters will propagate in the resulting calibration, so that in the best case one would expect the results to be as accurate as the data used for calibration. Using the SS99 method is possible, in theory, to derive absolute magnitudes and intrinsic colors with average uncertainties of $\\sim 0.3$ and $\\sim 0.02$ magnitudes, respectively. This is a reasonably good accuracy taking into account the simplicity of the method, which does not require large computational resources. Another advantage of this method is its relatively low sensitivity to the spectral type. In general, the morphological differences in the spectra of early-type and late-type stars can lead to problems when defining robust indices. The results found in SS99 showed that, on average, the residuals did not vary significantly among stars of different spectral types (in the range A-K). This is because the method does not need to determine the underlying true continuum. The extension of the method to B-type stars was done by \\citet[][hereafter SS02]{SS02}. In a subsequent paper, \\citet{Mol04} applied the same technique to a sample of stars within the spectral types K-M. The effects of variations in the spectral resolution on the sensitivity of the method was analyzed by \\citet{Gar05}. They showed that the method can be applied to low resolution spectra without loss of accuracy and that small variations in the spectral resolution would not affect the determination of the physical parameters for early-type stars. However, their results also established that for late-type stars the method need to be improved.  To date the proposed method (SS99) has not been used to derive physical parameters of any sample of stars, except for the same stars used for the calibration. In this paper we perform the first test of the method by its application to the stars in an open cluster. Our main goal is to evaluate the method in order to identify strengths and limitations. The open cluster chosen for this study is IC~2391, since high resolution spectra for $\\sim 50$ possible members are available from the UVES Paranal Observatory Project \\citep{Bag03}. Another reason to choose this cluster is the large amount of studies published in the last years, which will allow us to verify the results obtained using the SS99 method and to compare them with previous works.  This paper is organized in the following way. In Section~\\ref{metodo} we briefly explain the method used to determine the stellar parameters. In Section~\\ref{muestra} we describe the sample of stars over which the method is applied. The results obtained are shown in Section~\\ref{resultados}, where a detailed analysis lead us to propose some improvements to the method which allow us to determine the stellar parameters and to compare them with literature values. Finally, the main conclusions are summarized in Section~\\ref{conclusiones}. ",
        "conclusions": "  \\begin{itemize} \\item In the method described by SS99 and SS02, the combinations of lines having the smallest average residuals are not necessarily the ones yielding the best results for $M_V$ and $(B-V)_{0}$. In fact, the average uncertainties may vary significantly among the different combinations, mainly in late type stars. In this work we propose to use the averages of the results derived from all the available combinations instead of from only one combination. We have verified that this procedure reduces considerably the overall uncertainty.  \\item The method used does not provide reliable results (it generates relatively large uncertainties) when applied to type A stars. The problem seems to lie in the fact that this spectral type was not adequately represented in the original calibration (SS99, SS02). It becomes necessary, therefore, a new calibration of the method using a more extensive spectral library.  \\item From the results obtained for this sample of stars, we conclude that it is possible to derive absolute magnitudes with uncertainties of $0.59$ or $0.35$ magnitudes, depending on whether the comparison is done with values derived from the spectral type or values reported in the literature, respectively. It is also possible to derive intrinsic colors with average uncertainties of $0.023$ or $0.029$ magnitudes, depending again on the source used for the comparison. These uncertainties are in agreement with the average residuals associated to $M_V$ and $(B-V)_{0}$ according to SS99 and SS02.  \\item Stars numbered 1, 17, 36, 45 and 46 are probably not members of IC~2391. The rest of the stars allowed us to estimate a distance of $156 \\pm 24\\ \\mathrm{pc}$ for the cluster, which is in good agreement with published values. \\end{itemize}  Summarizing, this work shows that the method proposed in SS99 and SS02 can be applied in a reliable way to calculate the physical parameters of stellar systems. Both the simplicity in the definition of the indices used and the relatively small uncertainties in magnitude and color values make this method a suitable tool for the analysis of a large number of stellar spectra. Nevertheless, a new calibration is necessary in order to overcome the current limitations of the method, mainly those associated with type A stars."
    },
    "0809/0809.1078_arXiv.txt": {
        "abstract": "{ Measuring amplitudes of solar-like oscillations and the granulation power spectral density constitute two promising sources of information to improve our understanding and description of the convection in outer layers of stars. However, different instruments, using different techniques and different band passes, bring measurements which cannot be directly compared neither to each other nor to theoretical values. } { In this work, we define simple response functions to derive intrinsic oscillation amplitudes and granulation power density, from photometry measurements obtained with a specific instrument on a specific star. } {We test this method on different photometry data sets obtained on the Sun with two different instruments in three different band passes. } {We show that the results are in good agreement and we establish reference intrinsic values for the Sun in photometry. We also compute the response functions of the CoRoT instrument for a range of parameters representative of the Main Sequence solar-like pulsators to be observed with CoRoT. We show that these response functions can be conveniently described by simple analytic functions of the effective temperature of the target star. } {} ",
        "introduction": "Solar-like oscillations are being detected in a rapidely growing number of stars \\citep[see e.g.][]{Bedding07}. The excitation of these oscillations first observed in the Sun is attributed to the acoustic noise generated by convection in the outer layers of stars and the measurement of their amplitude is a source of information on the convection process \\citep[see e.g.][]{Samadi07,Samadi07b}.  The existing theoretical works generally bring parametric scaling laws calibrated on the Sun. However, as noticed by \\citet{Kjeldsen05}, measurements made on different stars with different instruments using different techniques in velocimetry or photometry, in different spectral lines or band passes, have different sensitivity to the oscillations. They cannot be compared directly to each other, nor to theoretical values. The comparison to the Sun is not straightforward either, since the different existing data sets obtained on the Sun have not been translated into a proper standard reference suited for comparison with stars. \\citet{Kjeldsen05} initiated such a normalization work and a comparison between several stars. Then, very recently, \\citet{Kjeldsen08} measured the solar oscillation amplitude with stellar techniques, aiming at setting up a consistent reference for stellar oscillation measurements.  This was done in velocimetry, since till now the vast majority of solar-like oscillations measured in other stars has been obtained with this technique. However, CoRoT \\citep{Baglin06}  has started bringing photometric measurements of oscillation in solar-like pulsators which will need to be measured quantitatively and compared with those of the Sun and with those obtained in velocimetry. In addition to oscillations, rapid photometry might allow measuring, in approximately the same domain of frequency, the power density spectrum contribution associated with the stellar granulation. Granulation being a manifestation of the convective motions at the photosphere level, the profile of its power density spectrum is expected to reflect characteristic time scales and geometrical scales associated with the convection process as described by heavy 3D numerical simulations \\citep[see e.g.][]{Ludwig06,Trampedach98} or by parametrized models \\citep[see e.g.][]{Baudin06}.  In the present work, we consider measurements of solar photometrical variations obtained with two different instruments in four different band passes (SOHO/VIRGO/PMO6 and SPM three channels). In the corresponding instrumental power density spectra, we fit contributions from the solar background and from the acoustic oscillations (Sect.2). Then, in Sect.3, we establish a simple instrumental response function relating the instrumental power density measurement to the intrinsic bolometric luminosity relative variation. These response functions can be applied to infer intrinsic (bolometric) power density of solar background from specific photometry measurements. They also can be used to derive intrinsic amplitude of solar radial oscillations from the same data.  We discuss how they can be adapted for non radial modes. Following \\citet{Kjeldsen05} and \\citet{Kjeldsen08}, we also propose to relate the oscillation mean power density measurement to an intrinsic amplitude chosen here to be the bolometric amplitude for radial modes. In Sect.4, we show that the results obtained with the different data sets considered here are consistent to a good approximation and allow us to produce a reference value of bolometric radial oscillation amplitude for the Sun observed as a star, and a reference bolometric power density spectrum for Solar granulation. Then (Sect.5), we compute the response functions adapted to the CoRoT instrument for stars representative of potential solar-like pulsators on the Main Sequence in terms of effective temperatures, log g values and chemical compositions. We show that to a great extent, the dependency with log g and chemical composition can be neglected and that the CoRoT response functions can be conveniently described with a good precision by analytic functions of $T_{\\rm eff}$. ",
        "conclusions": "Measurement of stellar oscillations or granulation brings instrumental values which depend on the instrumental technique and bandpass and on the star considered. In this work, with the pourpose of helping future comparisons between stars observed in photometry, \\begin{enumerate} \\item we propose a simple expression for response functions connecting specific instrumental photometric measurements with intrinsic bolometric values for oscillation amplitudes and granulation power density. \\item we test and validate this expression on four sets of data obtained on the Sun, in four different band passes and with two different instrumental techniques. \\item we establish reference bolometric measurements for the Solar oscillation amplitudes ($ 2.53 \\pm 0.11$ ppm) and for the Solar granulation power density. \\item we compute the response functions for the CoRoT instrument and give an analytic expression for it. \\end{enumerate}"
    },
    "0809/0809.3885_arXiv.txt": {
        "abstract": "We have discovered small whirlpools in the Sun, with a size similar to the terrestrial hurricanes ($\\la$~0.5~Mm). The theory of solar convection predicts them, but they had remained elusive so far. The vortex flows are created at the downdrafts where the plasma returns to the solar interior after cooling down, and we detect them because some magnetic bright points (BPs) follow a logarithmic spiral in their way to be engulfed by a downdraft. Our disk center observations show 0.9$\\times 10^{-2}$~vortexes per Mm$^{2}$, with a lifetime of the order of 5~min, and with no preferred sense of rotation. They are not evenly spread out over the surface, but they seem to trace the supergranulation and the mesogranulation. These observed properties are strongly biased by our type of measurement, unable to detect vortexes except when they  are engulfing magnetic BPs. ",
        "introduction": "A proper understanding of the solar surface convection  has impact on many different areas of astrophysics. For example, (a) the solar metallicity is used as a universal reference, and it has been recently modified by fifty percent after considering realistic convective motions in the analysis \\citep{asp05b}, and (b) turbulent dynamo action can produce magnetic fields in many astrophysical environments \\citep[e.g.,][]{mar04,ber05}, but its observational study is so far confined to the turbulent dynamo driven by the solar convective motions \\citep[e.g.,][]{cat99a,vog07}.    We have discovered small whirlpools in the Sun, with a size similar to the terrestrial hurricanes. This discovery confirms a specific prediction of the current theory of solar convection that had remained elusive so far. The energy produced in the solar interior is transported to the surface firstly by radiation exchange and, during the last thirty percent of the way, by convective motions \\citep[e.g.,][]{sti91}. The hot buoyant plasma rises, releases energy, and then falls down. According to the current theory, these convective motions are driven by strong highly localized downdrafts at the solar surface \\citep{spr90,stei98,ras98}. The downdrafts are sinks where the cold plasma returns to the solar interior. Since the matter has angular momentum with respect to the draining point, it must spin up when approaching the sink, giving rise to a whirl flow ({\\em bathtub effect}). Although the vortex motions predicted by the theory have been repeatedly sought, they had not been found so far \\citep{nov89,spr90,rou97,hoe98,nis03}. (The exception in \\citealt{bra88} represents the same phenomenon at much larger scale.) We have detected these vortexes as proper motions of magnetic bright points (BPs), which follow spiral paths in the way to be engulfed by a downdraft. The BPs are easily blurred by seeing \\citep{tit96}, thus the unique 90~km spatial resolution provided by the Swedish Solar Telescope at La Palma turned out to be necessary \\citep{sch02}. Magnetic BPs are present almost everywhere on the Sun, also in the quiet internetwork regions \\citep{san04a,dew05,dew08}. Precisely looking for properties of these  quiet Sun BPs, we found out that some of them spiral when heading towards the supposed location of a sink, which is the result reported in this Letter. We note that the vortical motion of photospheric magnetic concentrations may have significant impact on its own, e.g., on the heating of the solar corona \\citep[][]{bal98,gud05,ama08}. Previous searches for photospheric vortexes were often motivated by this interest \\citep{bal98,nis03}. ",
        "conclusions": "Observing these convectively driven vortex flows has been elusive because of the size of the whirlpools. Motions take place at sub-arcsec scales  and, therefore, close to the best angular resolution achieved at present. In addition, the case analyzed in  \\S~\\ref{whirlpool} seems to be extreme in the sense that the ratio between the azimuthal velocity and the radial velocity is particularly large ($A\\simeq 6.4$), which renders closed spiral paths and clear swirling motions. Assuming that the velocity of the plasma that feeds the whirlpools has an important stochastic component set by the complex granular dynamics, then one should expect the azimuthal and the radial velocities to be similar, rendering $A\\sim 1$. In this case the swirling motions are moderate and so easy to overlook. The BP trajectories corresponding to these more typical cases are hard to distinguish from straight lines within the angular resolution of the observations (see Fig.~\\ref{twocases}). \\begin{figure} \\includegraphics[width=0.5\\textwidth]{f4.ps} \\caption{ BP trajectories like the one analyzed in \\S~\\ref{whirlpool} (the solid line, corresponding to $A=6.4$), and what we expect to be the most probable case, with $A\\sim 1$ (the dashed line).  The swirling motions of the latter are difficult to detect since the deviations from straight lines are of the order of the angular resolution of the observation (represented by the hashed disk at the core of the whirlpool). } \\label{twocases} \\end{figure}  Our work should be complemented with a spectroscopic confirmation of the vortex motions. It requires observing the whirlpools out of the disk center so that the horizontal circular motions have a component along the line-of-sight (LOS). Such spectroscopic observation is challenging. Achieving the required spatial resolution is even more difficult in spectroscopy. Moreover, one would like to maximize the velocity component along the LOS by observing well outside the solar disk center. However, the solar surface is corrugated and, outside the disk center, the granules protrude hiding the intergranular lanes that harbor the vortexes \\citep[e.g.,][]{car04,kel04}."
    },
    "0809/0809.3761_arXiv.txt": {
        "abstract": "% The idea that we live near the centre of a large, nonlinear void has attracted attention recently as an alternative to dark energy or modified gravity.  We show that an appropriate void profile can fit both the latest cosmic microwave background and supernova data. However, this requires either a fine-tuned primordial spectrum or a Hubble rate so low as to rule these models out.  We also show that measurements of the {\\em radial\\/} baryon acoustic scale can provide very strong constraints.  Our results present a serious challenge to void models of acceleration. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.3411_arXiv.txt": {
        "abstract": "We present several corrections for point source photometry to be applied to data from the Infrared Array Camera (IRAC) on the Spitzer Space Telescope.  These corrections are necessary because of characteristics of the IRAC arrays and optics and the way the instrument is calibrated in-flight.  When these corrections are applied, it is possible to achieve a $\\sim$2\\% relative photometric accuracy for sources of adequate signal to noise in an IRAC image. ",
        "introduction": "The Infrared Array Camera (IRAC) was built at the NASA Goddard Space Flight Center under the direction of a team led by Giovanni Fazio at the Smithsonian Astrophysical Observatory \\citep{fazio04}. IRAC is the mid-infrared camera on the Spitzer Space Telescope \\citep{werner}, with four arrays, or ``channels'', simultaneously taking data in two separate fields of view. The four channels are referred to in this paper with their standard labels of 3.6, 4.5, 5.8, and 8.0 $\\mu$m for channels 1, 2, 3, 4, respectively, although as described by the IRAC documentation\\footnote{http://ssc.spitzer.caltech.edu/irac/} and by \\cite{fazio04,reach05} the nominal wavelengths differ from these labels. The absolute calibration of the camera was performed in flight by comparison to a set of stars that had been selected and characterized before launch \\citep{megeath03,cohenxiii}. \\cite{reach05} presented the IRAC in-flight calibration results, including the observing strategy, the predictions and measurements, and an assessment of the calibration accuracy and stability of the instrument and the pipeline-processed data provided by the Spitzer Science Center (SSC) to observers.  There are several characteristics of the IRAC instrument that affect the accuracy of the photometry obtained from the images.  The effects considered in this paper are:  \\begin{enumerate} \\item The IRAC science pipeline generates images in units of surface brightness, but because of distortion, the pixels do not subtend constant solid angle.  This causes errors in point source photometry that vary over the field of view (FOV) in each channel. \\item The IRAC spectral response varies over the field of view, and therefore the color corrections are field dependent. \\item The electron rates in the 3.6 $\\mu$m and 4.5 $\\mu$m channels depend slightly on pixel phase (the position of the star relative to the nearest pixel center).  The pipeline calibration factors are correct on average, as is appropriate for sources observed multiple times at multiple dither positions.  For the most precise photometry, however, pixel phase should be taken into account. \\item A correction must be applied for the size of the aperture and background region used in aperture photometry, if different from that used by \\citet{reach05} to derive the calibration from standard stars. \\end{enumerate}  The corrections described in this paper should be applied to the photometry in a manner consistent with those applied by \\citet{reach05} (which includes using the centroiding technique and the aperture and annulus background sizes described by Reach et al., applying the point source gain correction described in Section 2.2 below, and the pixel phase correction described in Section 3 below), in order for the absolute calibration to remain valid and to achieve $<$2\\% photometric accuracy reported. ",
        "conclusions": "The IRAC camera has FOV-dependent transmission characteristics that affect the measurement of astronomical sources.  After correction of these effects for standard stars, \\citet{reach05} found that the calibration has a relative accuracy of 1.8\\%, 1.9\\%, 2.0\\%, and 2.1\\% in channels 1 (3.6 $\\mu$m), 2 (4.5 $\\mu$m), 3 (5.8 $\\mu$m), and 4 (8 $\\mu$m), respectively.  To measure fluxes at this level of accuracy requires several photometric corrections: array position dependence (due to changing spectral response and pixel solid angle over the camera of view), pixel phase dependence (due to nonuniform quantum efficiency over a pixel), color correction (due to the different system response integrated over the passband for sources of different color), and aperture correction (due to the fractions of light included within the measurement aperture and lost in the background aperture).  The same accuracies are possible for sources with spectra similar to the A-type standards used in the absolute calibration with the array position dependence, pixel phase, and aperture corrections.  For sources with spectra different from A stars, a knowledge of the source spectrum is necessary to make the necessary color corrections to achieve the same photometric accuracy."
    },
    "0809/0809.0718_arXiv.txt": {
        "abstract": "We report the results of a local stability analysis for a magnetized, gravitationally stratified plasma containing cosmic rays. We account for cosmic-ray diffusion and thermal conduction parallel to the magnetic field and allow $\\beta = 8\\pi p/B^2$ to take any value, where $p$ is the plasma pressure and $B$ is the magnetic field strength.  We take the gravitational acceleration to be in the $-z$-direction and the equilibrium magnetic field to be in the $y$-direction, and we derive the dispersion relation for small-amplitude instabilities and waves in the large-$|k_x|$ limit. We use the Routh-Hurwitz criterion to show analytically that the necessary and sufficient criterion for stability in this limit is $n k_B dT/dz + dp_{\\rm cr}/dz + (1/8\\pi) dB^2/dz > 0$, where $T$ is the temperature, $n$ is the number density of thermal particles, and $p_{\\rm cr}$ is the cosmic-ray pressure.  We present approximate analytical solutions for the normal modes in the low- and high-diffusivity limits, show that they are consistent with the derived stability criterion, and compare them to numerical results obtained from the full, unapproximated, dispersion relation. Our results extend earlier analyses of buoyancy instabilities in galaxy-cluster plasmas to the $\\beta\\lesssim 1$~regime. Our results also extend earlier analyses of the Parker instability to account for anisotropic thermal conduction, and show that the interstellar medium is more unstable to the Parker instability than was predicted by previous studies in which the thermal plasma was treated as adiabatic. ",
        "introduction": "\\label{sec:intro}  Convection plays an important and well-known role in the transport of energy in stellar interiors. It has also been argued that convection is important in a number of low-density astrophysical plasmas, such as the intracluster medium in clusters of galaxies (Chandran \\& Rasera 2007) and accretion flows onto compact objects (Quataert \\& Gruzinov 2000).  Although convection in stellar interiors has been thoroughly studied over the course of several decades, the theory of convection in low-density plasmas is still being developed, and investigations carried out during the last several years have led to some interesting surprises.  For many years, it was widely assumed that the convective stability criterion for a low-density, non-rotating, weakly magnetized plasma is the Schwarzchild criterion, $ds/dz>0$, where $s$ is the specific entropy of the plasma and the gravitational acceleration is in the~$-z$ direction.  However, Balbus~(2000, 2001) showed that even weak magnetic fields strongly modify the convective stability criterion by causing heat to be conducted almost exclusively along magnetic field lines.  This anisotropy in the thermal conductivity arises when the electron gyroradius is much less than the electron mean free path, a condition that is easily satisfied for realistic magnetic fields in most cases of interest.  Balbus considered an equilibrium in which the magnetic field is in the~$xy$-plane, and in which $\\beta = 8\\pi p/B^2 \\gg 1$, where $p$ is the pressure and~$B$ is the magnetic field strength.  He showed that near marginal stability, the temperature of a rising fluid parcel is almost constant, essentially for two reasons. First, the parcel remains magnetically connected to material at its initial height. Second, near marginal stability a fluid parcel rises very slowly, so thermal conduction has enough time to approximately equalize the temperature along the perturbed magnetic field lines. As a result, the stability criterion becomes~$dT/dz>0$.  When this criterion is satisfied, a rising fluid parcel is cooler than its surroundings, and hence denser at the same pressure, so that it falls back down to its initial height. Parrish \\& Stone (2005, 2007) carried out numerical simulations that validated Balbus' analysis and extended it to the nonlinear regime.  An immediate question arises, namely, why doesn't Balbus's stability criterion apply to stars? Although stellar plasmas are magnetized, heat is conducted in stellar interiors primarily by photons. As discussed by Balbus (2000, 2001), the conductivity is thus almost isotropic, so it is the Schwarzchild criterion that applies. The reason that the Schwarzchild criterion applies even if the isotropic conductivity is large is somewhat subtle. If $ds/dz<0$, then adiabatic expansion would cause a slowly rising fluid parcel to be hotter and lighter than its surroundings, and hence buoyant.  The effect of isotropic conductivity is then to relax the temperature in the parcel towards that of the immediately surrounding fluid. However, because the conductivity is finite, the rising fluid parcel's temperature is never decreased all the way to the temperature of its surroundings.  The fluid parcel thus remains slightly hotter than its surroundings, and hence slightly lighter at the same pressure, and the fluid is convectively unstable. Although isotropic conductivity does not modify the Schwarzchild stability criterion, it does reduce the convective heat flux and the ``efficiency of convection'' in a convectively unstable fluid by decreasing the temperature difference $\\delta T$ between rising fluid parcels and their surroundings (Cox \\& Guili 1968).  Balbus's analysis has been extended in two ways by recent studies. First, Chandran \\& Dennis (2006), (hereafter referred to as \\paperi), investigated how the stability criterion is affected by the presence of cosmic rays that diffuse primarily along magnetic field lines. Like Balbus (2000, 2001), they assumed that~$\\beta \\gg 1$ and took the equilibrium magnetic field to be in the~$xy$-plane They showed that near marginal stability, the cosmic-ray pressure is nearly constant within a rising fluid element. This is because the fluid element remains connected to material at its initial height, and because fluid elements rise very slowly near marginal stability, so that there is plenty of time for cosmic-ray diffusion to approximately equalize the cosmic-ray pressure~$p_{\\rm cr}$ along the perturbed magnetic field lines. \\paperi\\ showed analytically that the stability criterion in the presence of cosmic rays is $n k_B dT/dz + dp_{\\rm cr}/dz > 0$, where $n$ is the total number density of thermal particles.  More recently, Quataert (2007) considered buoyancy instabilities in a low-density, high-$\\beta$ plasma in the absence of cosmic rays, but allowing the equilibrium magnetic field to have a component in the~$z$ direction, parallel or antiparallel to the direction of gravity. Since the temperature is a function of~$z$, the $z$ component of the equilibrium magnetic field leads to an equilibrium heat flux. Quataert (2007) showed that this heat flux causes the plasma to become convectively unstable even if~$dT/dz>0$, so that the plasma is always convectively unstable if the magnetic field has a nonzero $z$ component and a nonzero component in the $xy$ plane.  This heat-flux-buoyancy instability arises because of the geometry of the perturbed magnetic field lines in the plasma. For example, when the magnetic field is in the~$z$ direction and a fluid element is displaced upwards at a 45-degree angle with respect to the $z$ axis, field lines converge as they enter the fluid element from ``above'' (i.e., from larger~$z$).  As a result, if $dT/dz>0$ then the parallel heat flux converges within the fluid element, causing the fluid element to become hotter than its surroundings, and thus less dense at the same pressure. Buoyancy forces then cause the upwardly displaced fluid element to rise unstably. (Quataert~2007) The nonlinear development of this instability was investigated numerically by Parrish \\& Quataert (2007).  One of the open questions in this area of research is whether the buoyancy instabilities identified in these previous studies for the~$\\beta \\gg 1$ regime still operate when the magnetic field strength is increased to the point that~$\\beta \\lesssim 1$. We address this question in this paper. We consider the equilibrium geometry investigated by Balbus (2000, 2001) and \\paperi, in which the magnetic field is in the~$xy$-plane - in particular, we set $\\Bvec_0 = B_0 \\yhat$. We also allow for cosmic rays that diffuse along magnetic field lines, but now we allow~$\\beta$ to take any value.  We focus on wave vectors in the ``quasi-interchange'' limit, in which $|k_x| \\gg |k_y|$, $|k_x| \\gg |k_z|$, and $|k_x H| \\gg 1$, where~$H$ is the density scale height. This is the most unstable wave-vector regime for stratified adiabatic plasmas, because a small~$k_y$ reduces the stabilizing effects of magnetic tension and a large~$k_x$ allows a rising fluid element to easily get out of the way of the next rising element beneath it by moving just a small distance in the~$x$ direction. (Parker 1967, Shu 1974, Ferri\\`ere et al. 1999) We show analytically that the stability criterion in this limit is \\begin{equation} nk_B \\frac{dT}{dz} + \\frac{d p_{\\rm cr}}{d z} + \\frac{1}{8\\pi} \\frac{dB^2}{dz} > 0, \\label{eq:stabcrit0} \\end{equation} and we present a heuristic derivation of this stability criterion from physical arguments.  We also derive approximate analytical solutions to the dispersion relation for small-amplitude perturbations to the equilibrium in different parameter regimes, and compare these solutions to numerical solutions of the full dispersion relation.  Our results are important for determining the convective stability of galaxy-cluster plasmas, in which cosmic-rays are often produced by central radio sources.  Convection in intracluster plasmas is of interest because it may provide a mechanism for regulating the temperature profiles of galaxy-cluster plasmas and offsetting radiative cooling, thereby solving the so-called ``cooling-flow problem.''  (Chandran 2004, 2005; Parrish \\& Stone 2005, 2007; Chandran \\& Rasera 2007). Our results are also important for determining the conditions under which the Parker instability can operate in the interstellar medium (ISM). Previous treatments of the Parker instability assume an adiabatic thermal plasma (Parker 1966, 1967, Shu 1974, Ryu 2003).  Our results show that anisotropic thermal conductivity makes the ISM more unstable to the Parker instability, so that the instability can operate under a wider range of equilibrium profiles than was previously recognized.  The remainder of this paper is organized as follows. In section \\ref{sec:disprel} we outline the derivation of the general form of the dispersion relation.  In section \\ref{sec:qil} we specialize to the quasi-interchange limit, present our derivation of the necessary and sufficient condition for convective stability, and describe the properties of the unstable eigenmodes in plasmas that are very close to marginal stability. In section~\\ref{sec:physical} we present a heuristic, physical derivation of the stability criterion. We discuss the implications of our work for galaxy-cluster plasmas and the interstellar medium in sections~\\ref{sec:gc} and~\\ref{sec:parker}, respectively.  In section~\\ref{sec:conc} we summarize our results, and in appendix~\\ref{ap:eigenmodes} we present approximate analytic solutions and numerical solutions to the dispersion relation. ",
        "conclusions": "\\label{sec:conc}  In this paper we derive the stability criterion for local buoyancy instabilities in a stratified plasma, with the equilibrium magnetic field in the~$\\yhat$ direction and gravity in the~$-\\zhat$ direction. We take into account cosmic-ray diffusion and thermal conduction along magnetic field lines and focus on the large-$|k_x|$ limit, which is the most unstable limit for adiabatic plasmas.  Our work extends the earlier work of Balbus (2000, 2001) and \\paperi\\ by allowing for arbitrarily strong magnetic fields.  Applying our work to galaxy-cluster plasmas, we find that increasing the magnetic field to the point that~$\\beta = 8\\pi p/B^2 \\lesssim 1$ would shut off buoyancy instabilities at wavelengths along the magnetic field that are much shorter than the equilibrium scale height.  Our analysis also extends our understanding of the Parker instability by allowing the equilibrium values of $T$, $\\beta$, and $p_{\\rm cr}/p$ to vary with~$z$, and by accounting for anisotropic thermal conduction. We find that the interstellar medium is more unstable to the Parker instability than was predicted by earlier studies, which treated the thermal plasma as adiabatic."
    },
    "0809/0809.0697_arXiv.txt": {
        "abstract": "%   A search of the Two Micron All Sky Survey and Sloan Digital Sky Survey reveals 36 previously unknown high proper motion objects with $J<17$.  Their red-optical colors indicate that 27 are M dwarfs, 8 are early-type L dwarfs, and 1 is a late-type T dwarf.  The L dwarfs have $J-K_s$ colors near the extrema of known L dwarfs indicating that previous surveys for L dwarfs using color as a selection criterion may be biased.  Followup near-infrared spectroscopy of 6 dwarfs confirm they are all late-type with spectral types ranging from M8 to T4. Spectroscopy also shows that some of the L dwarf spectra exhibit peculiar features similar to other peculiar \"blue\" L dwarfs which may indicate that these dwarfs have a relatively condensate free atmosphere or may be metal poor.  Photometric distance estimates indicate that 22 of the new M, L and T dwarfs lie within 100 pc of the Sun with the newly discovered T dwarf, 2MASS J10595185+3042059, located at $\\sim$25 pc.  Based on the colors and proper motions of the newly identified objects, several appear to be good subdwarf candidates.  The proper motions of known ultracool dwarfs detected in our survey were also measured, including, for the first time, SDSS J085834.42$+$325627.6 (T1), SDSS J125011.65$+$392553.9 (T4) and 2MASS J15261405$+$2043414 (L7). ",
        "introduction": "Stars that exhibit high proper motions (HPMs: $\\mu > 0\\farcs4$ yr$^{-1}$) are likely to be within the solar neighborhood (less than $\\sim$100 pc).  Indeed most of our nearest stellar neighbors have been identified through their proper motions using shallow optical surveys such as the Digital Sky Survey (DSS) (Luyten 1979; Wroblewski and Costa 2001; Ruiz et al. 2001; Oppenheimer et al. 2001; Lepine et al. 2003; Teegarden et al. 2003; Hambly et al. 2004; Scholz et al. 2004; Pokorny et al. 2004; Subasavage et al. 2005, Deacon et al. 2005; Lodieu et al. 2005; Finch et al. 2007; Lepine 2008). However, these surveys are not sensitive to the very red ($I-J > 4$), low mass ultracool ($T_{\\mathrm{eff}}<2400$ K) L and T dwarfs.  The L and T dwarfs are extremely faint in the optical and to date have only efficiently been identified based on their colors at red optical (0.6$-$1.0 $\\mu$m) wavelengths in the Sloan Digital Sky Survey (SDSS: e.g., Fan et al. 2000, Leggett et al. 2002, Chiu et al. 2006) and the Canada France Brown Dwarf Survey (CFBDS: e.g., Delorme et al. 2008) and at near-infrared (1$-$2.5 $\\mu$m) wavelengths in the Two Micron All Sky Survey (2MASS: e.g., Kirkpatrick et al. 1999, 2000; Burgasser et al. 2002, 2003b; 2004a; Cruz et al. 2003; Tinney et al. 2005; Kendall et al. 2007; Looper et al. 2007), the Deep Near Infrared Southern Sky Survey (DENIS: Phan-Bao et al. 2008; Kendall et al. 2004) and the UKIRT Infrared Deep Sky Survey (UKIDSS: e.g., Lawrence et al. 2007; Chiu al. 2008, Pinfield et al. 2008).  Because these surveys select ultracool dwarf candidates based on color alone, they may be biased against objects with unusual colors.  Metchev et al. (2008) have been able to partially relax this color constraint for low mass objects by comparing the colors of objects in the 2MASS catalog to colors of objects in the SDSS catalog.  Although such surveys have proven very successful in identifying L and T dwarfs, it is likely, based on the luminosity function of stellar and substellar objects, that the census of the solar neighborhood within 25 pc remains roughly 20 to 60\\% deficient (Henry et al. 1997; Henry et al. 2002; Reid et al. 2004; Cruz et al. 2007; Reid et al. 2007).   A search for high proper motions stars at red optical and near infrared wavelengths could 1) potentially identify brown dwarfs in the solar neighborhood via their large proper motions, and 2) remove any potential color bias inherent in most of the surveys conducted to date.  Two surveys covering up to a few hundred square degrees have been successful in identifying ultracool dwarfs via proper motion in the near infrared (Artigau et al. 2006; Looper et al. 2008).  In this paper, we present a large area survey ($\\sim$ 8000 square degrees) to search for faint, high proper motion stars in the near infrared using the 2MASS and SDSS catalogs. ",
        "conclusions": "The original goal of this survey was to identify brown dwarfs that were close to the Sun and to understand whether or not there is a color bias in previous surveys that use color as a selection criterion.  Although no brown dwarfs located very close to the Sun were identified, our results suggest that a color bias does exist in current surveys.  As shown in Figure ~\\ref{fig:ijjkcolors}, six of the eight L dwarfs identified in this survey have $J-K_s$ colors that lie at the extrema of the distribution of L dwarfs colors.  Followup spectroscopy shows that 2MASS J1431$+$1436 and 2MASS J1434$+$2202 appear to have anomalous spectral features given their (tentative) spectral types.  Based on their $J-K_s$ colors alone, both dwarfs would have a spectral type of earlier than M9 V (Kirkpatrick 2008) indicating that these dwarfs are much bluer than typical L dwarfs. Indeed Burgasser et al. (2004a, 2008a) have noted that the spectra of 2MASS J0041$-$3547 and SIPS J0921$-$2104, the best fitting library spectra for 2MASS J1431$+$1436 and 2MASS J1434$+$2202 respectively, appear anomalous with stronger FeH absorption, weaker CO absorption, and a bluer near-infrared continuum than typical L dwarfs.  Such spectral characteristics have been previously ascribed to subsolar metallicities (e.g., Cruz et al. 2007), high surface gravities or variations in the condensate clouds properties (e.g., Burgasser et al. 2008a).  A sub solar metallicity may explain the peculiar nature of 2MASS J1434$+$2202 given that the spectral type of its best fitting library spectrum, 2MASS J0041$-$3547, is sdM9.  Most of the newly discovered objects are relatively distant (Table 4). It is thus likely that most of these objects have large tangential velocities in order for them to show such high proper motions at relatively large distances.  Metal poor subdwarfs are usually associated with large tangential velocities.  In order to determine if any of the objects reported in this survey are subdwarfs we calculated their Reduced Proper Motions, RPM.  The RPM of each object was determined using the relation $RPM = r' + 5 + 5 \\log \\mu$ (Salim and Gould 2002; Subasavage et al. 2005; Marshall 2008).  The RPM can be thought of as a rough luminosity class (or absolute magnitude) using only proper motion ($\\mu$) and apparent magnitude ($r'$) by assuming they are directly related.  This assumption is usually valid though many factors make the scatter in the relationship large.  The RPM is a simple technique to allow identification of object types that differ from the normal main sequence.  This technique is useful when looking for a crude way to identify white dwarfs and subdwarfs since they will have bluer colors or fainter RPMs for a given color.  In Figure \\ref{fig:reducedPM} we plot the $r'-J$ colors versus the RPM for the newly discovered high proper motion objects discovered in this survey along with some white dwarfs found in this survey, the L4 subdwarf 2MASS J1626$+$3925 also found in this survey and several confirmed M subdwarfs seen in both the SDSS and 2MASS from Marshall (2008) and Lepine and Scholz (2008).  It can clearly be seen that the white dwarfs occupy the far left in the figure while the subdwarfs occupy the center and lower portions of Figure \\ref{fig:reducedPM}. Several of our newly identified HPM objects are near the center and lower portions of Figure \\ref{fig:reducedPM} and are thus good subdwarf candidates (see Table 4).  In fact, 2MASS J1434$+$2202, which is near the known subdwarf 2MASS J1626$+$3925 in Figure \\ref{fig:reducedPM}, has been identified as a likely subdwarf from our near-infrared spectroscopy (see Table 7).  Most of our newly found HPM objects are on the far right of Figure \\ref{fig:reducedPM} and thus not likely subdwarfs and are definitely not white dwarfs.  Because the scatter is rather large a more detailed spectroscopic observational campaign will be required to determine which of these objects are subdwarfs."
    },
    "0809/0809.2692_arXiv.txt": {
        "abstract": "We search for stars with proper motions in a set of twenty deep Subaru images, covering about 0.28 square degrees to a depth of $i' \\simeq 25$, taken over a span of six years. In this paper, we describe in detail our reduction and techniques to identify moving objects. We present a first sample of 99 stars with motions of high significance, and discuss briefly the populations from which they are likely drawn. Based on photometry and motions alone, we expect that 9 of the candidates may be white dwarfs. We also find a group of stars which may be extremely metal-poor subdwarfs in the halo. ",
        "introduction": "The basic structure of our Milky Way galaxy seems clear: a thin disk of young stars, gas and dust circles the center quietly, immersed within a thicker disk of older stars. Both disks sit inside a nearly spherical halo of very old, metal-poor stars which do not share the overall rotation of the disks. Surrounding everything is an extended distribution of dark matter. Our knowledge of the details within this big picture, on the other hand, is not so clear. Recent large-scale projects, such as the Sloan Digital Sky Survey, have measured the properties of high-luminosity stars throughout the halo (see, for example, \\cite{Yanny2000}, and \\cite{Juric2008}), while the rapid development of infrared detectors has allowed projects such as the Two Micron All Sky Survey (\\cite{Skrutskie2006}) and Spitzer Space Telescope (\\cite{Patel2004}) to pierce the dusty disks and measure the properties of their stars. Nonetheless, some portions of the Milky Way remain largely unexplored.  In particular, we know little of the low-mass, low-luminosity stars of the halo. The white dwarfs and metal-poor subdwarfs of the halo glow so faintly, from so great a distance, that they are rarely seen and more rarely recognized. As a recent review (\\cite{Reid2005}) points out, however, these shy and elusive stars may dominate the microlensing events observed towards the Galactic bulge and the Magellanic Clouds.  There are two ways to search for these stars: cover a very large area on the sky to a shallow depth, or use a ``pencil-beam'' survey to examine a tiny region much more deeply. The first approach (see, for example, \\cite{Oppenheimer2001} and \\cite{Carollo2006}) will find objects in many directions, but only out to a small distance from the Sun; the second approach (see, for example, \\cite{Mendez2002}, \\cite{Nelson2002}, and \\cite{Kalirai2004}) probes farther into the halo, but only in a specific direction. One way to characterize surveys is to combine their area with the distance out to which they would detect some specific star to generate an ``effective volume'' for that type of star. In Table \\ref{table:survey_volume}, we compare the projects mentioned above by this metric, using a star of absolute magnitude $M_V = 16.5$, appropriate for a cool white dwarf. We assumed a color $(V-R) = 0.5$ to convert limiting magnitudes for $R$-based surveys to $V$-band.  \\begin{longtable}{l l l r} \\caption{Effective volumes, for $M_V = +16.5$} \\label{table:survey_volume} \\hline \\hline Survey & Area (sq.deg.) & limiting mag $V$ & volume (pc$^3$) \\\\ \\endfirsthead \\hline \\hline \\endhead \\hline \\endfoot \\hline \\endlastfoot  \\hline Oppenheimer et al. & 4165                & \\quad 19.8     & 80000 \\\\ Carollo et al.     & 1150                & \\quad 20       & 15000 \\\\ Nelson  et al.     & \\phantom{416}0.021  & \\quad 26.5     &  2100 \\\\ Mendez             & \\phantom{416}0.0013 & \\quad 26.0     &    64 \\\\ Kalirai et al.     & \\phantom{416}0.0031 & \\quad 29       &  9800 \\\\ this work          & \\phantom{416}0.28   & \\quad 26       & 14000 \\\\  \\end{longtable}   Our project could be described as a ``pencil-beam'' survey, but it uses a very thick pencil. We examine images in the area of the Subaru Deep Field (SDF) (\\cite{Kashikawa2004}) acquired over a period of six years to search for moving objects. These images were acquired primarily to study high-redshift galaxies (\\cite{Nagao2004}, \\cite{Shimasaku2006}), but have also been used to find high-redshift supernovae (\\cite{Poznanski2007}). The images are nearly as deep as those in some HST-based surveys, but cover a significantly wider area, yielding a large effective volume. We can refer to the SDF catalogs compiled by \\citet{Kashikawa2004} and \\citet{Richmond2005} for information on our candidates in multiple passbands ($B$, $V$, $R_c$, $i'$ and $z'$). Unlike most other pencil-beam surveys, we have measurements at many epochs: our dataset contains images taken on 20 nights. We can therefore measure the proper motion of our candidates very well, and place strong constraints on the uncertainties in our measurements.  This paper is the first in a series on proper motions in several small fields studied with Subaru. We will concentrate on techniques, leaving detailed analysis of the results for later papers. In Section 2, we describe the observations, their reductions, and the combination of individual frames into a single combined image for each night. In Section 3, we walk through our procedure for finding moving objects, and discuss our criteria for separating good candidates from bogus ones; we end up with a first sample of stars which have very well measured proper motions. In Section 4, we compute the reduced proper motions for objects in this sample, and compare their properties to those of objects drawn from a simulated survey of our field. Finally, in Section 5, we list our plans for future work on this dataset, and in other fields with multiple epochs of deep Subaru imaging.   Astronomers at the Observatoire de Besan\\c{c}on have created a model of the stellar populations in the Milky Way (\\cite{Robin2003}) with a very convenient web-based interface\\footnote{http://bison.obs-besancon.fr/modele}. The model consists of four populations of stars -- thin disk, thick disk, spheroid, outer bulge (the innermost portions of the Milky Way are poorly constrained) -- plus white dwarfs added to each component separately. The parameters of each component are adjusted to produce the best fit to the observed stellar populations and their dynamics. Recent work by \\citet{Ibata2007} finds that the Besan\\c{c}on model does a very good job of reproducing observed star counts in two of three deep fields down to $i_{0} = 24$. Those authors criticized the Besan\\c{c}on populations for being too sharply defined, forcing them to smooth the model colors by small amounts ($\\sim 0.10$ mag) in order to match observed color-magnitude diagrams. However, since our main concern is to classify objects very broadly using a mixture of kinematics, magnitudes and colors, we will adopt the Besan\\c{c}on model and use it as a reference throughout this paper to help us interpret properties of our sample. ",
        "conclusions": ""
    },
    "0809/0809.0973_arXiv.txt": {
        "abstract": "We use the combined {\\sc galform} semi-analytical model of galaxy formation and {\\sc grasil} spectrophotometric code to investigate the properties of galaxies selected via their sub-millimeter (sub-mm) emission.  The fiducial model we use has previously been shown to fit the properties of local ULIRGs, as well as the number counts of faint sub-mm galaxies.  Here we test the model in more detail by comparing the SEDs and stellar, dynamical, gas and halo masses of sub-mm galaxies against observational data.  We precisely mimic the sub-mm and radio selection function of the observations and show that the predicted far-infrared properties of model galaxies with $S_{850}>5$\\,mJy and $S_{1.4}>30\\mu$Jy are in good agreement with observations.  Although the dust emission model does not assume a single dust temperature, the far-infrared SEDs are well described by single component modified black-body spectrum with characteristic temperature $32\\pm5$\\,K, in good agreement with observations.  We also find evidence that the observations may have uncovered evolution in the far-infrared--radio relation in ULIRGs out to $z\\sim 2$.  We show that the predicted redshift distribution of sub-mm galaxies provides a reasonable fit to the observational data with a median redshift $z=2.0$.  The radio-selected subset of sub-mm galaxies are predicted to make up approximately 75\\% of the population and peak at $z=1.7$, in reasonable agreement with the observed radio detected fraction and redshift distribution.  However, the predicted $K$-band and mid-infrared (3--$8\\mu$m) flux densities of the sub-mm galaxies (and LBGs) are up to a factor $10\\times$ fainter than observed.  We show that including the stellar TP-AGB phase in the stellar population models does not make up for this deficit. This discrepancy may indicate that the stellar masses of the sub-mm galaxies in the model are too low: M$_{\\star}\\sim10^{10}$\\,M$_{\\odot}$, while observations suggest more massive systems, M$_{\\star}\\gsim10^{11}$\\,M$_{\\odot}$.  However, if the predicted $K$- and 3--8$\\mu$m extinctions in the model could be dramatically reduced, then this would reduce, but not eliminate, this discrepancy. Finally we discuss the potential modifications to the models which may improve the fit to the observational data, as well as the new observational tests which will be made possible with the arrival of new facilities, such as SCUBA2. ",
        "introduction": "The discovery a decade ago of a population of faint submm-selected galaxies (SMGs) revolutionised our view of the cosmic star formation history of the Universe \\citep{Smail97,Hughes98,Barger98}.  These galaxies (originally discovered with the SCUBA instrument at 850$\\mu$m) appear to be high redshift galaxies with star-formation rates exceeding 1000$\\Msolyr$ \\citep{Ivison00,Smail02}.  These Ultra-Luminous Infrared Galaxies (ULIRGs) peak around $z\\sim2$ \\citep{Chapman03a,Chapman05a} and show a thousand fold increase in their abundance between $z=0$ and $z=2$.  If this far-infrared emission arises solely from star-formation with a standard (``solar neighborhood'') IMF, then they could potentially dominate the star-formation activity in the early Universe, dwarfing the contribution of galaxies selected in the rest-frame ultraviolet. The apparent intensity of these starbursts, the resulting high metallicity, along with their large dynamical masses, high gas fractions and inferred strong clustering \\citep{Blain04a,Greve05,Tacconi06,Tacconi08,Swinbank04,Swinbank06b} are all suggestive of a close link to the formation phase of the most massive spheroids and black holes \\citep{Lilly99,Smail02,Smail04,Webb03,Genzel03,Smail04,Alexander03a,Alexander05a,Swinbank06b}.  Reproducing the sub-mm galaxy population has been a major challenge for theoretical models \\citep{Granato00,Baugh05}.  In part this is because (quite reasonably) many of the recipes and constraints used to develop the models are based on the formation and evolution of ``normal'' galaxies rather than the extreme populations.  Hence these models have had difficulty reproducing extremely luminous galaxies with sufficient cool dust at high redshift without over-predicting the abundance of bright galaxies in the local Universe.  Initial attempts to match the basic properties of sub-mm galaxies (whilst maintaining the match to the present day $K$-band luminosity function and {\\it IRAS} 60$\\mu$m luminosity function) under-predicted the 850$\\mu$m counts by a factor $30\\times$ \\citep{Baugh05} despite $\\Lambda$CDM producing enough baryons in massive halos at $z\\sim2$ to match observations of gas masses in sub-mm galaxies \\citep{Genzel03,Greve05}.  One solution to this problem was to alter the Initial Mass Function (IMF): \\citet{Blain99b} suggested that a Salpeter IMF \\citep{Salpeter55} with a low-mass threshold of 3\\,M$_{\\odot}$ was required in far-infrared luminous galaxies around $z\\sim2$ since the implied star-formation rate for a Salpeter IMF integrated to the canonical value of 0.07\\,M$_{\\odot}$ would over-predict the integrated stellar density at $z=0$.  Invoking a top-heavy (or ``flat'') IMF in the semi-analytic framework has been shown to provide a much improved fit to the sub-mm galaxy counts and redshift distribution as well as the Lyman-break galaxy luminosity function whilst still maintaining a good fit to the properties of the present day galaxy populations (\\citealt{Baugh05}, hereafter B05).  In this model, a standard IMF is adopted in quiescently star-forming galaxies, whilst the star-formation induced by galaxy mergers produces stars with a flat IMF: $dn/dln(m)=m^{-x}$ with $x=0$ (compared to a Salpeter IMF which has $x=1.35$; \\citealt{Salpeter55}).  With a larger proportion of high mass stars, the energy radiated in the ultra-violet per unit mass of stars produced is increased, thus increasing the amount of radiation to heat the dust.  Moreover, the flat IMF produces a higher yield of metals from type II supernovae, thus increasing the dust content of the galaxy and boosting the luminosity in the sub-millimeter waveband.  This assumption of a flat IMF in bursts is controversial and so it is essential to explore the predictions of this model in more detail.  In particular, the model can be (i) tested against the growing observational multi-wavelength data on sub-mm galaxies to test whether the choice of parameters is suitable and (ii) provide useful constraints on observational data, in particular for understanding possible selection biases.  For example, since much of what is known about sub-mm galaxies is based on the radio-detected sub-sample, one outstanding issue is how much this radio-selection (which does not benefit from the negative K-correction experienced in the sub-mm waveband) has affected the conclusions being drawn about SMGs.  Are the radio-undetected fraction of SMGs at either significantly higher redshift \\citep{Younger07} or do they instead have far-infrared colours which represent much colder systems \\citep{Chapman05a}?  In \\S2 we describe the basic properties of the galaxy formation model we employ, but refer the reader to \\citet{Cole2k}, \\citet{Benson02,Benson03}, \\citet{Baugh05} and \\citet{Lacey08} for a detailed description.  In \\S3 we test the predicted far-infrared properties of sub-mm-selected galaxies against the available observational data.  In \\S4 we discuss the masses and evolution of model SMGs compared to those estimated from observational data and in \\S5 we give our conclusions and future prospects for constraining the properties of sub-mm galaxies both theoretically and observationally.  We use a flat, $\\Lambda$CDM cosmology with $\\Omega_{m}=0.3$, $\\sigma_8=0.93$ and $H_{o}=100h$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ with $h=0.7$. ",
        "conclusions": "In this paper we have used the {\\sc galform} semi-analytic model for galaxy formation together with the {\\sc grasil} spectrophotometric code to investigate in detail the properties of sub-mm galaxies in the model of \\citet{Baugh05}. We mimic the observational selection in our model catalogue and identify galaxies both via their 850$\\mu$m and 1.4\\,GHz (radio) emission. We find that the radio-identified SMGs (rSMGs) are predicted to make up approximately 75\\% of the whole population. This is in good agreement with the fraction found in observational surveys (eg.  \\citealt{Ivison05,Ivison07}). The main difference between the radio-identified and radio-unidentified sub-mm galaxies is the characteristic dust temperature, where the radio undetected SMGs are $\\sim10$\\% cooler.  This results in model SMGs having a median redshift of $z=2.0$ whilst model rSMGs (which include a radio flux density threshold S$_{1.4}>30\\mu$Jy) have a median redshift of $z=1.7$.  We show that the far-infrared colours of model rSMGs are consistent with observational data. Fitting the 350 and 850$\\mu$m photometry we show that the model far-infrared SEDs are well described by single component modified ($\\beta=1.5$) blackbodies with temperature of $32\\pm5$\\,K (although {\\sc grasil} does not assume a single dust temperature either within or between galaxies). Integrating the model SED and the black-body fit, the inferred and true bolometric luminosities for the model SEDs also agree, with the single modified blackbody fit recovering an average of 0.94$\\pm$0.25 of the bolometric luminosity.  However, the median 850$\\mu$m/1.4\\,GHz flux density ratio for SMGs predicted by the model is systematically lower than that observed by a factor $\\sim1.26\\pm0.24\\times$, even after the model has been normalized to match the far-infrared--radio correlation for ULIRGs at $z=0$.  This could be due to evolution in the far-infrared--radio correlation, or due to modest contributions to the observed radio flux density from AGN.  We show that the predicted velocity dispersions are in good agreements with observations, with $\\sigma_{1D}\\sim160\\kms$ for the model SMGs, compared to observational constraints of $\\sigma=170$--$200\\kms$ inferred from the kinematics of the redshifted H$\\alpha$ and CO emission.  We also show that the predicted gas masses are agree well with those derived from observations using $\\alpha=0.8$.  Turning to the halo masses, we find that the model SMGs reside in halos with masses of M$_{halo}\\sim10^{12.4}$M$_{\\odot}$, and a correlation length of $r_{o}=8.8\\pm0.3$\\,Mpc, in excellent agreement with the tentative observational estimate of $r_{o}=9.8\\pm3.0$\\,Mpc from \\citet{Blain04a}. We find no evidence for a significant merger bias in the clustering of SMGs in the model.  However, we find that the predicted $K$-band and mid-infrared ($5.8\\mu$m) flux densities of the model SMGs are up to a factor $10\\times$ fainter than observed. We discuss the possible reasons for this. The stellar masses of the model SMGs, M$_{\\star}\\sim10^{10}$\\,M$_{\\odot}$, are a factor 10$\\times$ smaller compared to observational estimates based on the same mid-infrared data (M$_{\\star}\\gsim10^{11}$\\,M$_{\\odot}$), but the observational estimates are very sensitive to the assumed IMF, burst age and dust extinction. The simplest explanation for this discrepancy in the near- and mid-infrared fluxes is that it is due to the stellar masses being too low in the model. However, other factors may be contributing, for example in the model SMGs, the dust extinction is quite large even at $5.8\\mu$m, resulting in a $4\\times$ reduction in the observed flux. If the rest-frame visible and near-infrared dust extinctions in the model SMGs were removed, this would improve (but not eliminate) this discrepancy (an offset of a factor 2$\\times$ would remain).  However, we note that there is support for these high extinctions in the $K$-band, at least for the emission-line gas.  We also discuss possible routes for increasing the near-infrared luminosities and stellar masses, both within the context of the current \\citep{Baugh05} model, which assumes a top-heavy IMF in bursts, and using the prescriptions described in \\citet{Bower06}, which uses AGN feedback to prevent too many baryons cooling in massive halos but assumes a normal IMF throughout. The \\citet{Bower06} model predicts more high-mass galaxies at high redshift, and a $K$-band luminosity function in better agreement with observations at high redshift, but under-predicts the sub-mm counts by a factor 20. We will discuss the properties of sub-mm galaxies in models which combine features from the \\citet{Baugh05} and \\citet{Bower06} models in a future paper.  We also investigate the evolution of the sub-mm galaxy population in the \\citet{Baugh05} model by comparing the halo and stellar masses for model SMGs at $z=1$--3 to those at $z>3.5$.  The $z>3.5$ SMG population are predicted to make up approximately 2\\% of the total SMG population (with the $z>4$ rSMGs contributing approximately 0.5\\% of the total population).  Nevertheless, we find only a mild increase in the halo and stellar masses between $z\\sim4$ and $z\\sim2$, and therefore attribute the strong evolution seen in the redshift distribution primarily to the evolution in the space density of massive halos hosting SMGs.  As well as the need for the models to provide a better match to observations, a number of upcoming surveys and facilities will provide valuable observational constraints which can be used to further test them.  In particular, deep SCUBA2 surveys will probe both 450 and 850$\\mu$m, providing much tighter constraints on the SEDs of galaxies with bolometric luminosities a factor $10\\times$ lower than currently achievable. When combined with results from {\\it Herschel} (which will probe 60--670$\\mu$m), the temperatures, bolometric luminosities and dust masses of SMGs will be constrained with much higher accuracy than currently possible from the sparse and low signal-to-noise 350 and 850$\\mu$m photometry alone.  Interestingly, the model predicts that there is a significant tail of faint 450$\\mu$m selected galaxies at high redshift (Lacey et al.\\ 2009 in prep).  Selecting galaxies with S$_{450}>2.5$\\,mJy (the expected 5-$\\sigma$ flux density limit for SCUBA2) the model suggests that there should be $\\geq8000$, 850 and 150 galaxies selected at 450$\\mu$m per square degree at $z>2$, 4 and 5 respectively ($\\sim50$\\%, 5\\% and 1\\% of the total population respectively), suggesting that the fainter confusion limit at 450$\\mu$m may provide the most efficient route to get an unbiased census for the far-infrared selected galaxies at the highest redshifts, providing insights into the physical processes of star-formation occurring at the peak epoch of galaxy formation, $z\\sim2$."
    },
    "0809/0809.2855_arXiv.txt": {
        "abstract": "We simulate planet migration caused by interactions between planets and a planetesimal disk. We use an N-body integrator optimized for near-Keplerian motion that runs in parallel on a video graphics card, and that computes all pair-wise gravitational interactions. We find that the fraction of planetesimals found in mean motion resonances is reduced and planetary migration rates are on average about 50\\% slower when gravitational interactions between the planetesimals are computed than when planetesimal self-gravity is neglected. This is likely due to gravitational stirring of the planetesimal disk that is not present when self-gravity is neglected that reduces their capture efficiency because of the increased particle eccentricity dispersion. We find that migration is more stochastic when the disk is self-gravitating or comprised of more massive bodies. Previous studies have found that if the planetesimal disk density is below a critical level, migration is ``damped'' and the migration rate decays exponentially, otherwise it is ``forced'' and the planet's migration rate could accelerate exponentially. Migration rates measured from our undamped simulations suggest that the migration rate saturates at a level proportional to disk density and subsequently is approximately power law in form with time. ",
        "introduction": "The field of solar system dynamics, has a timeline of discoveries that is related to the computational power available (e.g., \\citealt{morby01}). As computational power has increased over time, so to has our ability to more accurately simulate more complex systems. The interaction between planets and planetesimals subsequent to formation is a well posed n-body problem but displays remarkable complexity even in the absence of collisions, including resonance capture, planetary migration, and heating and scattering of planetesimals. In this paper we attempt to more accurately simulate this system by using the increased computational power and on board memory recently available on video graphics cards.  Planet-planetesimal scattering can lead to planet migration (e.g., \\citealt{fernandez84,malhotra95,hahn99,ida00,levison03,gomes04,gomes05,hahn05, levison08}). Recent work on dusty circumstellar disks with clearings, suggest that planets are required to account for the morphology of the disk \\citep{quillen06,wyatt06}.  The planetesimal masses suggested by collisional models (e.g., \\citealt{wyatt02,DD03,quillen07}) and proximity of the hypothetical planet to the disks for systems such as Fomalhaut intimate that planetary migration could be taking place.  To simulate planet migration due to planetesimal scattering, planets must gravitationally interact with planetesimals. However, because of the $O(N^2)$ required computational intensity previous simulations have necessarily neglected the gravitational interactions between the planetesimals (e.g., \\citealt{hahn99,gomes04,hahn05,thommes08}). This means that collective gravitational effects and gravitational self stirring in the disk have been ignored.  Here we take all interparticle forces into account to explore what differences might be seen in simulations of migrating planets with and without planetesimal interactions. ",
        "conclusions": "In this paper we have described an implementation of the heliocentric democratic 2nd order symplectic integration scheme \\citep{duncan98} written in CUDA and designed to work on a graphics processing unit. This allows us to efficiently use the large number of processors available on the card in parallel. The code is designed to compute a large number of gravitational interactions as well as compute Keplerian advances in parallel. Because all memory resides on the device this parallel computation platform has advantages over message passage interfaces that must pay a penalty to share information between processors.  The biggest drawback with our current computational platform is the low precision, a problem that is now being resolved with the recent availability of double precision capable video cards.  We have used the new code to simulate planetary migration into a planetesimal disk with $10^4$ bodies and have computed all gravitational interactions. We have explored the difference between simulations of planets migrating into self-gravitating planetesimal disks and those lacking computed interactions between planetesimals. We find that the fraction of particles present in mean motion resonances is reduced in the simulations with self-gravitating disks compared to those lacking planetesimal interactions.  We attribute this to a reduction in the resonance capture rate due to the larger eccentricity dispersion caused by gravitational stirring in the self-gravitating planetesimal disks. This suggests that disks in proximity to migrating planets could be featureless (as is true for Fomalhaut, \\citealt{kalas05}). A dust disk lacking resonant structure could still host migrating planets.  We find that planet migration is smoother but somewhat faster in the simulations lacking self-gravity.  Migration rates in both cases are approximately linearly dependent on planetesimal disk density. The planet can undergo rapid changes in migration in the simulations with the most massive disks. However for the most of the disk densities considered no exponential increase in migration rate is seen implying that the migration rate saturates at a level approximately proportional to the disk density.  \\vskip 1.0 truein We thank Matt Holman for suggesting how to reduce the radial drift caused by single precision computation during the Keplerian computation step. Support for this work at University of Rochester and Rochester Institute of Technology was also provided by NASA through an award issued by JPL/Caltech, by NSF grants AST-0406823 \\& PHY-0552695 and and HST-AR-10972 to the Space Telescope Science Institute. This work is based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA."
    },
    "0809/0809.0850_arXiv.txt": {
        "abstract": "The problem of a test body in the Schwarzschild geometry is investigated in a Keplerian limit.  Beginning with the Schwarzschild metric, a solution to the limited case of approximately elliptical (Keplerian) motion is derived in terms of trigonometric functions.  This solution is similar in form to that derived from Newtonian mechanics, and includes first-order corrections describing three effects due to general relativity:  precession; reduced radial coordinate; and increased eccentricity.  The quantitative prediction of increased eccentricity may provide an additional observational test of general relativity.  By analogy with Keplerian orbits, approximate orbital energy parameters are defined in terms of a relativistic eccentricity, providing first-order corrections to Newtonian energies for elliptical orbits.  The first-order relativistic equation of orbit is demonstrated to be a limiting case of a very accurate self-consistent solution.  This self-consistent solution is supported by exact numerical solutions to the Schwarzschild geometry, displaying remarkable agreement.  A more detailed energy parameterization is investigated using the relativistic eccentricity together with the apsides derived from the relativistic effective potential in support of the approximate energy parameters defined using only first-order corrections.  The methods and approximations describing this Keplerian limit are applied to more general static spherically-symmetric geometries.  Specifically, equations of orbit and energy parameters are also derived in this Keplerian limit for the Reissner-Nordstr\\\"{o}m and Schwarzschild-de\\,Sitter metrics. ",
        "introduction": "The problem of a test body in the Schwarzschild geometry is of fundamental importance in understanding orbital characteristics due to general relativity.   Detailed analyses and approximate analytical solutions of this geometry exist \\cite{MTW,weinberg,rindler,PleKra,OR,dinverno,carroll1,ovanesyan,schild,BroCle,hartle,wang,stump,nobili,magnan,deliseo,brillgoel,dean}, emphasizing a first order approximate rate of precession of perihelia of Mercury and the unstable circular orbit and relativistic capture representative of extreme astrophysical environments.  The general problem has been solved in terms of Weierstra\\ss\\ elliptic functions by Hagihara \\cite{hagihara}, and more recently by Kraniotis and Whitehouse \\cite{KW1}.  Ashby \\cite{ashby} has solved the limited case of approximately Keplerian motion in terms of Jacobian elliptic functions, ``...which are closer in spirit to trigonometric functions, with which most are familiar.''  In this paper, beginning with the Schwarzschild geometry, a solution to the limited case of approximately Keplerian motion is derived in terms of trigonometric functions.  This approximate solution lends itself to easy comparison with the familiar Keplerian orbits (ellipses) of Newtonian mechanics and is consistent with both the reduced radius of circular orbit, commonly derived from the Schwarzschild effective potential, and the observed rate of precession of perihelia of the inner planets, commonly derived perturbatively.  Additional insights regarding the sizes and shapes of bound relativistic orbits are provided, including a relativistic correction to eccentricity which may be subjected to observational tests.  By analogy with Keplerian orbits, a relativistic eccentricity is used to define a Schwarzschild energy parameter, providing corrections to Newtonian energies for Keplerian orbits.  The relativistic equation of orbit together with the energy parameterization comprise a simple model that is useful for a qualitative and quantitative understanding of the corrections to Keplerian orbits due to general relativity.  This model is easily extended to include more general static spherically-symmetric geometries.  Specifically, models are also derived for the Reissner-Nordstr\\\"{o}m and Schwarzschild-de\\,Sitter metrics.  The methods and approximations describing a Keplerian limit are detailed in Sec.~\\ref{Sec_S_KepLim}.  Beginning with the Schwarzschild geometry, an approximate equation of orbit is derived in terms of trigonometric functions.  When compared to that of a corresponding Keplerian orbit, this equation of orbit clearly displays three characteristics of relativistic orbits:  precession; reduced radial coordinate; and increased eccentricity.  These characteristics arise as first-order relativistic corrections to the familiar equation of orbit describing Keplerian orbits.  This solution is found to be valid for near-circular orbits requiring only small relativistic corrections.  Predictions of relativistic precession and reduced radius of circular orbit are in agreement with known results.  This provides confidence in a new relativistic correction to eccentricity.  (The possibility of observing relativistic corrections to eccentricity is discussed briefly in Sec.~\\ref{SubSec_S_Char}.)  In addition, first-order relativistic corrections to Keplerian apsides are predicted, resulting in the conclusion that the overall size of a relativistic orbit is smaller than a corresponding Keplerian orbit.  By analogy with Newtonian mechanics, Schwarzschild energy parameters are defined using the virial theorem for circular orbits, and using a relativistic eccentricity for noncircular orbits, resulting in first-order relativistic corrections to Newtonian energies for Keplerian orbits.  This simple model is substantiated by a more detailed investigation in Sec.~\\ref{Sec_S_self_cons}.  Again beginning with the Schwarzschild geometry, the derivation of a very accurate self-consistent relativistic equation of orbit is detailed in Sec.~\\ref{Sec_S_self_cons}.  This self-consistent equation of orbit is identical in form to the more approximate equation of orbit derived in Sec.~\\ref{Sec_S_KepLim} and predicts the same orbital characteristics.  However, corrections to Keplerian orbits due to general relativity are described more accurately.  For example, the predicted radius of circular orbit is identical to that derived by minimizing the relativistic effective potential.  This solution is accurate in predicting long-term behavior of Schwarzschild orbits for a large parameter space, as is demonstrated by comparisons with exact numerical solutions.  The more approximate equation of orbit derived in Sec.~\\ref{Sec_S_KepLim} is shown to be a limiting case of this self-consistent solution.  A more detailed energy parameterization is investigated for comparison with the simple parameterization of Sec.~\\ref{Sec_S_KepLim}.  This parameterization is constructed using the relativistic eccentricity derived from the self-consistent equation of orbit together with the relativistic apsides derived from the Schwarzschild effective potential.  The energy parameters of Sec.~\\ref{Sec_S_KepLim} approximate well the results from this more detailed parameterization, lending value to the simpler approach and more approximate results.  The methods and approximations describing a Keplerian limit to the Schwarzschild geometry are applied to a more general class of static spherically-symmetric geometries in curvature coordinates \\cite{PleKra,OR,dinverno,carroll2,SH2,GaoZha}.  Specifically, the path of a small test mass in a static spherical spacetime is taken to be described by the metric \\begin{gather} \\mathrm{d}s^2 \\!= \\mathrm{e}^{2\\nu(r)} c^2 \\mathrm{d}  t^2 - \\mathrm{e}^{-2\\mu(r)} \\mathrm{d}r^2 - r^2 \\mathrm{d}\\varOmega^2, \\label{eq_metric_general} \\end{gather} where $\\mathrm{d}\\varOmega^2 = \\mathrm{d}\\theta^2 + \\sin^2{\\!\\theta}\\,\\mathrm{d}\\varphi^2$, and \\begin{gather} \\mathrm{e}^{2\\nu(r)} = \\mathrm{e}^{2\\mu(r)} = 1 - \\frac{2M}{r} + \\frac{Q^2}{r^2} - \\frac{1}{3}\\Lambda r^2 > 0. \\end{gather} The Reissner-Nordstr\\\"{o}m $(\\Lambda = 0)$ geometry is considered in Sec.~\\ref{Sec_RN_KepLim}, followed by the Schwarzschild-de\\,Sitter $(Q = 0,\\Lambda>0)$ geometry in Sec.~\\ref{Sec_SdS_KepLim}.  These geometries have been shown to meet all of the conditions for physical acceptability \\cite{delgaty,semiz,bronnikov,maartens,skmhh}.  In each of these cases the relativistic correction due to matter is considered to be the dominant contribution.  The resulting orbital equations are of the same form as that derived for the Schwarzschild geometry, including additional corrections for each geometry.  For examples: there is an additional contribution to relativistic precession due to charge opposite in direction to the contribution due to matter; there is also an additional contribution to relativistic precession due to $\\Lambda$, but in the same direction as the contribution due to matter.  Numerical studies are provided only for the Schwarzschild geometry in order to establish the self-consistent approach.  A self-consistent treatment of the Schwarzschild-de\\,Sitter geometry is intractable and is not pursued.  A concise summary of the results for the three geometries is given in Sec.~\\ref{Sec_Summary}. ",
        "conclusions": "\\label{Sec_Summary}  The relativistic central-mass problem is investigated in a Keplerian limit.  Beginning with the Schwarzschild metric, a relativistic equation of orbit is derived that is similar in form to that describing Keplerian orbits.  This equation of orbit includes three relativistic corrections to Keplerian orbits: precession; reduced radial coordinate; and increased eccentricity.  The prediction for the relativistic contribution to precession is in agreement with existing calculations and observations.  The predicted reduction in size of a circular orbit is also in agreement with existing calculations.  These agreements provide confidence in a new quantitative prediction of increased eccentricity,  which may be subjected to observational tests.  The methods and approximations describing this Keplerian limit to the Schwarzschild geometry are also applied to the Reissner-Nordstr\\\"{o}m and Schwarzschild-de\\,Sitter metrics.  The resulting equations of orbit are identical in form to that derived for the Schwarzschild metric and include relativistic corrections due to charge (RN) and cosmological constant (SdS).  In every case the relativistic equation of orbit has a form that is easily compared to that describing Keplerian orbits (\\ref{eq_N_eoo}) of Newtonian mechanics, \\begin{gather} \\frac{\\tilde{r}_\\text{c}}{r} = 1 + \\tilde{e} \\cos{\\tilde{\\kappa}\\varphi}. \\label{eq_summary1} \\end{gather} The coefficients $\\tilde{r}_\\text{c}$, $\\tilde{e}$, and $\\tilde{\\kappa}$ provide relativistic corrections for each geometry.  These corrections are given to first order in Table~\\ref{tab_summary1}. \\begin{table}[ht] \\label{tab_summary1} \\caption{Summary of first-order relativistic corrections to Keplerian (K) orbits for the Schwarzschild (S), Reissner-Nordstr\\\"{o}m (RN), and Schwarzschild-de\\,Sitter (SdS) geometries.  The coefficients of the relativistic equation of orbit, Eq.~(\\ref{eq_summary1}), are given for each of the three geometries.  In the bottom two rows the subscript $X$ represents mass $M$, charge $Q$, or cosmological constant $\\Lambda$.  The parameter $r_\\text{c} \\equiv \\ell^2/GM$ is the radius of a circular Keplerian orbit with angular momentum $\\ell$.} \\vspace*{-4ex} \\begin{equation*} \\begin{array}{c||c|c|c|c} \\hline\\hline {} & \\text{\\,K\\,} & \\text{S} & \\text{RN} & \\text{SdS} \\\\ \\hline\\hline \\tilde{\\kappa} \\vphantom{\\Bigr]} & 1 & 1 - 3\\epm & 1 - 3\\epm + \\epq & 1 - 3\\epm - 3\\epL \\\\ \\hline \\tilde{r}_\\text{c}/r_\\text{c} \\vphantom{\\Bigr]} & 1 & 1 - 3\\epm & 1 - 3\\epm + 2\\epq & 1 - 3\\epm + 2\\epL \\\\ \\hline \\tilde{e}/e \\vphantom{\\Bigr]} & 1 & 1 + 3\\epm & 1 + 3\\epm & 1 + 3\\epm + 8\\epL \\\\ \\hline \\epsilon^{}_{\\!X} \\vphantom{\\Biggr]^2_2} & \\text{--} & \\epm \\equiv \\dfrac{\\rm}{2 r_\\text{c}}  & \\epq \\equiv \\dfrac{\\rqsquared}{\\rm r_\\text{c}} & \\epL \\equiv \\dfrac{r_\\text{c}^3}{\\rm \\rLsquared} \\\\ \\hline r^{}_{\\!X} \\vphantom{\\Biggr]^2_2} & \\text{--} & \\rm \\equiv \\dfrac{2GM}{c^2} & \\rq \\equiv \\Bigl( \\dfrac{GQ^2}{4\\pi\\varepsilon_0 c^4}\\Bigr)^{\\!1/2} & \\rL \\equiv \\Bigl( \\dfrac{3}{\\Lambda}\\Bigr)^{\\!1/2} \\\\ \\hline\\hline \\end{array} \\end{equation*} \\end{table} \\noindent The first-order shift in apside (precession) per revolution $\\Delta\\varphi = 2\\pi (\\tilde{\\kappa}^{-1} -1)$ is consistent with known results for the Schwarzschild,  Reissner-Nordstr\\\"{o}m, and Schwarzschild-de\\,Sitter geometries.  (See Table~\\ref{tab_summary2}, left column.)  The RN geometry predicts an additional contribution to precession in the opposite direction to that due to matter, while the SdS geometry predicts an additional contribution to precession in the same direction as that due to matter.   For each of the three geometries, the radius of circular orbit is consistent to first order with that determined by minimizing the relativistic effective potential.  The Schwarzschild geometry predicts a reduced radius of circular orbit.  Both the RN and SdS geometries predict an radius of circular orbit that is larger than that predicted by the Schwarzschild geometry, but still smaller than that for a corresponding Keplerian orbit.  Schwarzschild orbits are predicted to be more eccentric than corresponding Keplerian orbits, and a cosmological constant (SdS) serves to further increase the eccentricity.  The presence of electric charge (RN) does not result in any additional contribution to eccentricity.  Relativistic corrections to eccentricity may serve as additional tests of general relativity. \\begin{table}[ht]\\label{tab_summary2} \\caption{Summary of additional first-order relativistic corrections to Keplerian (K) orbits for the Schwarzschild (S), Reissner-Nordstr\\\"{o}m (RN), and Schwarzschild-de\\,Sitter (SdS) geometries.  First-order angular shifts in apsides (precession) per revolution are listed in the left column.  First-order relativistic corrections to apsides are listed in the right column.  The relativistic correction parameters $\\epsilon^{}_{\\!X}$ are listed in Table~\\ref{tab_summary1}, bottom two rows.} \\vspace*{-4ex} \\begin{equation*} \\begin{array}{c||c|c} \\hline\\hline \\vphantom{\\biggr]} & (2\\pi)^{-1}\\Delta\\varphi & \\tilde{r}_\\pm/r_\\pm \\\\ \\hline\\hline \\text{K} & 0 & 1 \\\\ \\hline \\text{S} \\vphantom{\\biggr]^2_2} & 3\\epm & 1 - 3\\epm\\dfrac{1 \\mp 2e}{1 \\mp e} \\\\ \\hline \\text{RN} \\vphantom{\\biggr]^2_2} & 3\\epm - \\epq & 1 - 3\\epm\\dfrac{1 \\mp 2e}{1 \\mp e} + 2\\epq \\\\ \\hline \\text{SdS} \\vphantom{\\biggr]^2_2} & 3\\epm + 3\\epL & 1 - 3\\epm\\dfrac{1 \\mp 2e}{1 \\mp e} + 2\\epL \\dfrac{1 \\pm 3e}{1 \\mp e} \\\\ \\hline\\hline \\end{array} \\end{equation*} \\end{table}  This model and the resulting first-order corrections are valid for near-circular orbits $(e \\ll 1/2)$ that require only small relativistic corrections $(\\epm \\ll 1/12; \\epq \\ll \\epm; \\epL \\ll \\epm)$.  In addition to the properties listed in Table~\\ref{tab_summary1},  the overall size of a Schwarzschild orbit is predicted to be smaller than a corresponding Keplerian orbit.  This is determined not only by the radius of circular orbit, but also by a comparision of relativistic apsides to those for corresponding Keplerian orbits.  Both the apocenter and pericenter distances are found to be smaller for Schwarzschild orbits.  (See Table~\\ref{tab_summary2}, right column.)  Both the RN and SdS geometries predict apsides that are larger than the Schwarzschild apsides, but still smaller than the corresponding Keplerian apsides.  Long-term orbital behavior is predicted very accurately using a self-consistent Keplerian limit.  This is demonstrated by comparing a self-consistent equation of orbit to the exact numerical solution for Schwarzschild orbits.  The self-consistent equation of orbit is also given by Eq.~(\\ref{eq_summary1}), but with more accurate expressions for the coefficients $\\tilde{r}_\\text{c}$, $\\tilde{e}$, and $\\tilde{\\kappa}$.  For examples, the self-consistent solution predicts: a radius of circular orbit that is identical to that calculated by minimizing the relativistic effective potential; and a relative error in angular frequency of $~10^{-10}$ over 1600\\,cycles.  (See Fig.~\\ref{fig_moderate}.)  This solution is accurate in describing Schwarzschild orbits with eccentricities as large as $e=1/2$ and requiring large relativistic corrections.  The coefficients in Table~\\ref{tab_summary1} are found to be limiting cases of those derived for the more accurate self-consistent equation of orbit, lending value to the simpler approach.  A self-consistent equation of orbit is also derived for the Reissner-Nordstr\\\"{o}m geometry, foregoing numerical studies.  A self-consistent model for the Schwarzschild-de\\,Sitter geometry is intractable and is not pursued. \\begin{table}[ht]\\label{tab_summary3} \\caption{Summary of relations between Newtonian energies for Keplerian (K) orbits and Schwarzschild (S), Reissner-Nordstr\\\"{o}m (RN), and Schwarzschild-de\\,Sitter (SdS) energy parameters.  First-order relativistic corrections to energies for circular orbits are listed in the left column.  First-order relativistic corrections to energies for near-Keplerian orbits are listed in the right column.  The relativistic correction parameters $\\epsilon^{}_{\\!X}$ are listed in Table~\\ref{tab_summary1}, bottom two rows.} \\vspace*{-4ex} \\begin{equation*} \\begin{array}{c||c|c} \\hline\\hline \\vphantom{\\biggr]} & \\tilde{E}_\\text{c}/E_\\text{c} & (\\tilde{E} - E)/|E_\\text{c}| \\\\ \\hline\\hline \\text{K} & 1 & 0 \\\\ \\hline \\text{S} \\vphantom{\\Bigr]} & 1 + 2\\epm & -2\\epm (1 - 4e^2) \\\\ \\hline \\text{RN} \\vphantom{\\Bigr]} & 1 + 2\\epm - 2\\epq & -2\\epm (1 - 4e^2) + 2\\epq (1 - e^2) \\\\ \\hline \\text{SdS} \\vphantom{\\Bigr]} & 1 + 2\\epm + 2\\epL & -2\\epm (1 - 4e^2) - 2\\epL (1 - 9e^2) \\\\ \\hline\\hline \\end{array} \\end{equation*} \\end{table} Very accurate relativistic energy parameters for circular orbits are derived using the virial theorem.  This is useful for comparing the energy of a circular relativistic orbit $\\tilde{E}_\\text{c}$ to the energy of a corresponding circular Keplerian orbit $E_\\text{c}$.  This relation is summarized for the three geometries in Table~\\ref{tab_summary3}, left column.  (See also Fig.~\\ref{fig_dE_circular}.)  Because the relativistic orbits are taken to be very near-Keplerian, an energy parameterization for noncircular bound orbits analogous to that for Newtonian mechanics is investigated.  The total energy of a relativistic orbit $\\tilde{E}$ is defined by simple ansatz, \\begin{gather} \\tilde{E} = (1 - \\tilde{e}^2)\\tilde{E}_\\text{c}, \\label{eq_summary2} \\end{gather} where $\\tilde{E}_\\text{c}$ is the energy of a relativistic circular orbit, and $\\tilde{e}$ is a new relativistic eccentricity derived in the context of the relativistic equation of orbit (\\ref{eq_summary1}).  Then, using the relation between $\\tilde{e}$ and $e$ (Table~\\ref{tab_summary1}, middle row), together with the relation between $\\tilde{E}_\\text{c}$ and $E_\\text{c}$ (Table~\\ref{tab_summary3}, left column), a relation between the total energy for a noncircular relativistic orbit and total energy for a Keplerian orbit is derived.  This relation is summarized for the three geometries in Table~\\ref{tab_summary3}, right column.  (See also Fig.~\\ref{fig_dE}.)  The virial theroem and simple ansatz (\\ref{eq_summary2}) provide first-order relativistic corrections to Newtonian energies for bound orbits.  Finally, this simple energy parameterization is compared to a more detailed parameterization constructed using the intersection of the relativistic effective potential with a line of constant energy.  The results are similar, lending value to the simple ansatz (\\ref{eq_summary2}) and resulting approximate relations.  A more detailed energy parameterization for the Schwarzschild-de\\,Sitter geometry results in a quintic equation for $\\tilde{E}$ and is not pursued.  The additional unstable circular orbit and relativistic capture described in the standard Newtonian limit to general relativity \\cite{MTW,wald,hartle,carroll2,rindler,HobEfsLas,OR} are absent in the present treatment.  However, the present approach to the relativistic central-mass problem results in an equation of orbit that exhibits several characteristics of relativistic orbits at once.  In this Keplerian limit characteristics of general-relativistic orbits are provided as corrections to Keplerian orbits of Newtonian mechanics, providing a qualitative and quantitative understanding of the effects of general relativity on bound systems.  It should also be possible to adapt these results to relative motion of binary systems \\cite{JS,iorio1}.  Perhaps more general statements may be made concerning larger systems and more extreme environments.  Globular clusters are expected to be biased toward high eccentricity in galaxies with larger cores.  Individual stars orbiting near blackholes are expected to be in anomolously small and eccentric orbits.  These characteristics could also serve as an indicator of dark matter and dark energy.  Consider two galaxies, each having approximately the same amount of visible matter.  A large difference in the amount of coexisting dark matter should be apparent in the eccentricities of the orbits of individual stars and star clusters.  Individual outlying members of galaxy clusters are expected to be biased toward high eccentricity and large precession rates due to both the large central mass and the cosmological constant.  Although effects of the cosmological constant on planetary orbits have been ruled out, they may be observable in galaxy clusters with larger radii, for which $\\epL \\sim \\Lambda r_\\text{c}^2$ becomes non-negligible.  The methods and approximations describing this Keplerian limit may be applied to other static spherical spacetimes.  The results summarized in Tables~\\ref{tab_summary1}, \\ref{tab_summary2} and \\ref{tab_summary3} already describe related geometries using simple replacements.  The results for the Reissner-Nordstr\\\"{o}m geometry are extended to include magnetic charge with the replacement $Q^2/\\varepsilon_0 \\rightarrow Q^2/\\varepsilon_0 + P^2\\mu_0 c^2$, where $P$ is the magnetic charge \\cite{carroll2,PleKra}.  The results for the Schwarzschild-de\\,Sitter geometry become those for the Schwarzschild-anti-de\\,Sitter $(\\Lambda < 0)$ with the replacement $\\epL \\rightarrow -\\epL$ \\cite{KW2,SH,SH2,COV}.  It may also be possible to apply this Keplerian limit to more exotic objects such as wormholes \\cite{SR}, naked singularities, and Boson and Fermion stars \\cite{perlick}."
    },
    "0809/0809.0311_arXiv.txt": {
        "abstract": "We study the efficiency and dynamics of supermassive black hole binary mergers driven by angular momentum loss to small-scale gas discs. Such binaries form after major galaxy mergers, but their fate is unclear since hardening through stellar scattering becomes very inefficient at sub-parsec distances. Gas discs may dominate binary dynamics on these scales, and promote mergers. Using numerical simulations, we investigate the evolution of the semi-major axis and eccentricity of binaries embedded within geometrically thin gas discs. Our simulations directly resolve angular momentum transport within the disc, which at the radii of interest is likely dominated by disc self-gravity. We show that the binary decays at a rate which is in good agreement with analytical estimates, while the eccentricity grows. Saturation of eccentricity growth is not observed up to values $e \\simgt 0.35$. Accretion onto the black holes is variable, and is roughly modulated by the binary orbital frequency. Scaling our results, we analytically estimate the maximum rate of binary decay that is possible without fragmentation occuring within the surrounding gas disc, and compare that rate to an estimate of the stellar dynamical hardening rate. For binary masses in the range $10^5 \\msun \\simlt M \\simlt 10^8 \\msun$ we find that decay due to gas discs may dominate for separations below $a \\sim 10^{-2} \\ {\\rm pc} - 0.1 \\ {\\rm pc}$, in the regime where the disc is optically thick. The minimum merger time scale is shorter than the Hubble time for $M \\simlt 10^7 \\msun$. This implies that gas discs could commonly attend relatively low mass black hole mergers, and that a significant population of binaries might exist at separations of a few hundredths of a pc, where the combined decay rate is slowest. For more massive binaries, where this mechanism fails to act quickly enough, we suggest that scattering of stars formed within a fragmenting gas disc could act as a significant additional sink of binary angular momentum. ",
        "introduction": "The assembly of present-day galaxies occurs via the hierarchical merger of lower mass progenitors, many of which seem likely to have supermassive black holes (SMBHs) in their nuclei. Following a merger of two galaxies, each harbouring a black hole, we expect the SMBHs to sink toward the centre of the merger product and eventually form a black hole binary. What happens subsequently is less clear. There is no doubt that at sufficiently small separations ($a \\simlt 10^{-3} \\ {\\rm pc}$) loss of angular momentum to gravitational radiation occurs rapidly enough to effect coalescence. However, at larger separations an as-yet undetermined mix of angular momentum loss to stars and to gas is required to bring the black holes into the gravitational radiation inspiral regime. Observations provide scant information. Detecting black hole binaries on pc or sub-pc scales (where the black holes' gravity dominates the galactic potential) is extremely difficult, with the closest confirmed binary in the radio galaxy 0402+379 \\citep{Rodriguez06} having a projected separation of 7.3~pc. A variety of circumstantial evidence -- of which the most compelling is probably the small observed scatter in the relationships between black hole mass and galaxy properties \\citep{Ferrarese06,Gebhardt00} -- suggests that binaries do merge, though these arguments do not constrain the time scale or the mechanism that facilitates mergers.  One longstanding motivation for considering the role of gas in hastening binary mergers has come from the fact that stellar dynamical processes may simply fail. Although dynamical friction against the stellar background is rapid for large separations \\citep{Milosavljevic03}, for pc-scale binaries the rate slows dramatically as the black holes eject the finite number of stars that have orbits that come close enough to the centre to interact with the binary. Replenishment of stars on these loss cone orbits occurs on the relaxation time, which typically exceeds the Hubble time \\citep{Begelman80,Yu02b,Merritt06}. This `last parsec' problem occurs for spherical or axisymmetric nuclei, and is particularly severe for massive galaxies which are observed to possess rather low density stellar cores \\citep{Faber97}.  More recently, interest in this long standing problem has been further piqued by the realization that {\\em if} gas is responsible for driving mergers, then it is likely that some of that material will survive in the immediate vicinity of the holes at the moment of coalescence. Plausibly, some small fraction of the energy released during the merger may go into heating the gas, producing an electromagnetic precursor \\citep{Armitage02} or afterglow \\citep{Milosavljevic05,Lippai08,Schmittman08,Shields08} to the gravitational wave event. Detection of an electromagnetic counterpart would greatly improve the accuracy with which space-based gravitational wave observatories such as {\\em LISA} can localize mergers, and allow much more information to be gathered about the properties of the host galaxy \\citep{Kocsis07}. It may also be possible to determine whether gas discs typically attend mergers by searching for periodic electromagnetic signals produced by binaries at larger separations, independent of any gravitational wave information \\cite{Haiman08}.  Theoretically, it is well established that the dissipative nature of gas allows inflow down to pc scales in galactic nuclei following galactic mergers \\citep{Mihos96,DiMatteo05b,Mayer07b,Levine08}. Perhaps more surprisingly, smaller amounts of gas appear to be able to penetrate deep within the black hole's sphere of influence even in galaxies such as the Milky Way that are far removed from any significant merger activity. Evidence for such inflows comes from observations of young stars with disc-like kinematics close to Sgr~A* \\citep{Paumard06,Lu08}, which may well have formed in situ from gas at sub-pc scales \\citep{Levin03,NC05}. Taken together, these lines of argument suggest that it is plausible to invoke the existence of gas discs at small radii as a common feature of galactic nuclei, though the typical masses and structure of such discs are subject to considerable uncertainty. Prior theoretical work has established that gas discs are likely to be important for SMBH binary mergers, provided that the mass of gas in the disc is roughly of the same order of magnitude as the mass in black holes \\citep{Gould00,Armitage02,Escala05,Dotti07,MacFadyen08,Hayasaki08}.  In this paper we revisit the interaction between a SMBH binary and a relatively low mass circumbinary gas disc. Our main innovation is to study this interaction using numerical simulations that directly resolve the physical mechanism of angular momentum transport within the gas. At the radii of greatest interest for the SMBH binary merger problem (typically tenths of a pc, where the stellar dynamical time scales are longest for many galaxies) it is expected that self-gravity will dominate the dynamics of surrounding gas discs \\citep{Shlosman90,Goodman03}. The onset of self-gravity in an accretion disc can result either in a quasi-stable disc with outward angular momentum transport, or in fragmentation of the disc into stars. Either possibility is of great interest for its effect of SMBH binary mergers. Here, we study directly the first regime in which the disc remains stable and angular momentum is transported by the action of self-gravitating spiral arms. Physically this occurs if the cooling time in the disc exceeds the local dynamical time \\citep{Gammie01}. We use the simulations to study the rate at which the gas drives the binary toward merger, and the effect that the gas disc has on the eccentricity of the binary. Previous simulations that have used either artificial viscosity or a Navier-Stokes formulation\\footnote{Although the use of a Navier-Stokes viscosity is an improvement over a reliance on purely numerical effects to transport angular momentum, it too may lead to unphysical behaviour. In particular, it is known that a disc subject to only a Navier-Stokes shear viscosity is unstable to the growth of eccentricity even in the absence of perturbations \\citep{Ogilvie01}. In general, there is no reason to assume that a `turbulent viscosity' resulting from a physical process such as self-gravity or the magnetorotational instability \\citep{Balbus98,Balbus99} will behave in the same way as a microscopic fluid viscosity.} to model angular momentum transport have found that eccentricity growth (both of the binary, and within the gas disc) attends the decay of the binary semi-major axis \\citep{Artymowicz91,Armitage05,MacFadyen08}. The time dependent accretion signatures that result may permit the identification of black hole binaries within gas rich systems prior to the final coalescence phase.    \\begin{figure*} \\centerline{\\epsfig{file=disc_6.eps,width=.95\\textwidth}} \\caption{Logarithmic maps of the disc column density (in units of $M a_0^{-2}$) at different times during the simulation.  The panel at $t=0$ shows the smooth initial conditions.  Until $t \\approx 350\\Omega_0^{-1}$, material piles up at $R \\approx 3 a_0$ as a result of the torques shown in Fig.~\\ref{fig:torques}, forming a dense ring.  The ring breaks due to its self-gravity (see Fig.~\\ref{fig:toomre}), spreading the gas approximately over the original radial range.  A spiral pattern develops, and the disc stays in that state until at least $t \\approx 1200\\Omega_0^{-1}$, when the simulation ends.} \\label{fig:disc6} \\end{figure*} ",
        "conclusions": "In this paper we studied the interaction of a super-massive binary black hole with a self-gravitating circumbinary gaseous disc. Our focus has been on sub-pc scales, where binary hardening due to stellar dynamical processes is thought to be slow, and where efficient radiative cooling makes it likely that the gas will form a geometrically thin accretion structure. Studies of galaxy mergers, and indirect evidence from our own Galactic Centre, suggest that gas flows to such small scales may be a common feature of galactic nuclei, particularly following mergers. If present, the existence of a gas disc around the binary can be a dominant factor in determining whether the binary merges and the time-scale for this process. The presence of a disc will influence the binary orbit, as it will transport the angular momentum of the binary outward. At the radii of greatest interest the evolution of such gas discs will be dominated by their self-gravity, which will alter the structure of the disc. Our simulations resolve angular momentum transport due to self-gravity physically, without requiring any ad hoc prescription to mimic viscosity.  Our simulations show that the disc, after a transient phase due to the initial conditions, develops a quasi-steady spiral pattern that lasts for at least hundreds of dynamical times. The binary couples gravitationally to the disc, and loses angular momentum. This process makes the orbit of the binary decay at a rate $da/dt \\sim - {\\rm few} \\times 10^{-5} a_0 \\Omega_0$ (for our simulation parameters), in good agreement with theoretical calculations for non-self-gravitating discs. The eccentricity of the binary also increases, at a rate $de/dt \\sim 10^{-4} \\Omega_0$.  Due to the computational expense of the simulations, we have not followed the system for long enough to see a large change in the orbital parameters. Our direct numerical experiments show only a modest (20\\%) decrease in the binary separation, accompanied by a growth of eccentricity to values of around $e \\approx 0.3$. The evolution of both $a,e$ appears to slow down with time, approaching asymptotic values $a \\approx 0.8 a_0 ; e \\approx 0.4$.  However, we attribute this trend to the fact that the gaseous disc absorbs angular momentum and moves away from the binary, making the interaction less efficient.  This is consistent with a viscous time-scale $t_{\\rm visc} \\sim (R/H)^2 t_{\\rm orb} / \\alpha \\sim 3000\\Omega_0^{-1}$ at $R\\sim3a_0$, where most of the disc mass is located early in the simulation. On the other hand, the ratio $de/da$ remains roughly constant throughout our calculations, suggesting that the mechanism for eccentricity growth has not saturated.  We speculate then that, given enough influx of external material (which could easily occur on the viscous time-scale described above), the orbit will continue its decay and the eccentricity will keep growing, perhaps reaching the very large values needed to influence the decay in the relativistic regime.  We also studied the accretion onto the black holes during the orbital decay, and found that it is highly variable, especially on the time-scale of the orbital frequency.  This variability pattern, if observed in AGN, could help us to identify binaries at sub-parsec separations.  Motivated by the good agreement between the decay rate obtained from our simulations, and that predicted analytically by \\cite{Ivanov99}, we combine their formalism with that of \\cite{Levin07} to calculate the maximum decay rate that can be obtained due to a self-gravitating disc.  We find that the decay produced by the disc can dominate over stellar scattering once the binary separation is 0.01--0.1 pc. The time-scale for decay at the critical separation where gas disc and stellar processes are equally important is shorter than the Hubble time for binary masses $M \\simlt 10^7\\msun$. This implies that geometrically thin gas discs could provide a solution to the `last-parsec' problem for binaries in this range, whereas gas discs are unable to drive mergers within galaxies hosting the most massive black holes with masses of $10^8 \\msun$ and above. However, even for masses where gas discs can in principle result in mergers, the time scales are typically rather long. At low redshift, where the typical time scale between major mergers for galaxies hosting $10^7 \\msun$ black holes is several Gyr \\citep{Volonteri03}, we would then expect that many galaxies would host binaries with separations close to the hang-up radius of a few hundredths of a parsec. At higher redshift the interval between mergers is much shorter -- less than a Gyr -- and interactions between three or more black holes in a single system could not be neglected.  If we accept that most galactic mergers result in black hole mergers, the results presented here suggest that geometrically thin gas discs could in principle drive mergers on a short enough time scale for relatively low mass black holes at low redshift. Additional processes, either stellar dynamical or hydrodynamic, are needed if very massive binaries are to merge, and to avert ubiquitous 3-body interactions (that might result in ejections) at high redshift. One logical alternative that we did not explore here is that provided by even more massive discs. Such discs will fragment and form stars very close to the binary, and if those stars can be scattered by the binary prior to exploding as supernovae they will absorb the binary angular momentum almost as efficiently as if the disc had remained fluid. Given slower scattering, only low mass stars and stellar remnants will be able to contribute to binary decay. The stellar mass function is evidently critical in determining whether such a scenario is viable, as is the stellar dynamics of stars formed in a disc close to a (probably eccentric) binary black hole."
    },
    "0809/0809.3915_arXiv.txt": {
        "abstract": "We investigate inflationary scenarios driven by a class of potentials which are similar in form to those that arise in certain minimal supersymmetric extensions of the standard model. We find that these potentials allow a brief period of departure from inflation sandwiched between two stages of slow roll inflation. We show that such a background behavior leads to a step like feature in the scalar power spectrum. We set the scales such that the drop in the power spectrum occurs at a length scale that corresponds to the Hubble radius today---a feature that seems necessary to explain the lower power observed in the quadrupole moment of the Cosmic Microwave Background (CMB) anisotropies. We perform a Markov Chain Monte Carlo analysis to determine the values of the model parameters that provide the best fit to the recent WMAP $5$-year data for the CMB angular power spectrum. We find that an inflationary spectrum with a suppression of power at large scales that we obtain leads to a much better fit (with just one extra parameter, $\\chi_{\\rm eff}^{2}$ improves by $6.62$) of the observed data when compared to the best fit reference $\\Lambda$CDM model with a featureless, power law, primordial spectrum. ",
        "introduction": "Measurements of the Cosmic Microwave Background (CMB) anisotropies---from the early days of the COsmic Background Explorer (COBE) satellite until the most recent observations of the Wilkinson Microwave Anisotropy Probe (WMAP)---have consistently indicated a low value of the quadrupole, below the cosmic variance of the concordant $\\Lambda$CDM cosmological model with a nearly scale invariant, primordial spectrum~\\cite{cobe-2,cobe-4,wmap-1,wmap-3,wmap-5}. While there has been a recurring debate about the statistical significance of the quadrupole and the other outliers (notably, near the multipole moments $22$ and $40$) in the CMB angular power spectrum (see Refs.~\\cite{odwyer-2004,magueijo-2007,park-2007,chiang-2007} and references therein), there has also been constant activity to understand possible underlying physical reasons for the outliers (see, for an inexhaustive list, Refs.~\\cite{starobinsky-1992,jing-1994,feng-2003,efsthathiou-2003,luminet-2003,hajian-2003-05,kawasaki-2003-05,gordon-2004,contaldi-2003,cline-2003,moroi-2004,hunt-2004-07,sriram-2004a,sriram-2004b,sriram-2005,sinha-2006,covi-2006-07,boyanovsky-2006,donoghue-2007,powell-2007,nicholson-2008,joy-2007-08,destri-2008}).  Given the CMB observations, different model independent approaches have been used to recover the primordial spectrum (see, for example, Refs.~\\cite{wmap-3,bridle-2003,mukherjee-2003,hannestad-2004,tarun-2004-08,tv-2005}). While all these approaches arrive at a spectrum that is nearly scale invariant at the smaller scales, most of them inevitably seem to point to a sharp drop in power at the scales corresponding to the Hubble scale today. Within the inflationary scenario, a variety of single and two field models have been constructed to produce such a drop in power at the large scales~\\cite{starobinsky-1992,feng-2003,contaldi-2003,cline-2003,moroi-2004,boyanovsky-2006,powell-2007,nicholson-2008,destri-2008}. However, in single field inflationary models, in order to produce such a spectrum, we find that many of the scenarios either assume a specific pre-inflationary regime, say, a radiation dominated epoch, or special initial conditions for the background scalar field, such as an initial period of fast roll~\\cite{contaldi-2003,cline-2003,powell-2007,nicholson-2008}. Moreover, some of them impose the initial conditions on the perturbations when the largest scales are outside the Hubble radius during the pre-inflationary or the fast roll regime~\\cite{contaldi-2003,powell-2007,nicholson-2008}. Such requirements are rather artificial and, ideally, it would be preferable to produce the desired power spectrum during an inflationary epoch without invoking any specific pre-inflationary phase or special initial conditions for the inflaton. Furthermore, though a very specific pre-inflationary phase such as the radiation dominated epoch may allow what can be considered as natural (i.e. Minkowski-like) initial conditions for the perturbations even at super-Hubble scales, we believe that choosing to impose initial conditions for a small subset of modes when they are outside the Hubble radius, while demanding that such conditions be imposed on the rest of the modes at sub-Hubble scales, can be considered unsatisfactory.  It has long been known that power spectra with large deviations from scale invariance can be generated in inflationary models that admit one or more periods of departure from the slow roll phase (see, for instance, Refs.~\\cite{starobinsky-1992,hodges-1990,mukhanov-1991,tarun-1995,adams-1997,lesgourgues-2000,barriga-2001,adams-2001,jinn-ouk-2005}). The degree of the deviation from a nearly scale invariant spectrum would be determined by the extent and duration of the departure, which are, in turn, controlled by the parameters of the model. A departure from slow roll affects the evolution of modes that leave the Hubble radius just before the departure. Rather than remaining constant, the curvature perturbations, say, ${\\cal R}_{k}$, corresponding to these modes evolve at super-Hubble scales, sourced by the intrinsic entropy perturbations of the inflaton field which, typically, exhibit a rapid growth during the fast roll regime~\\cite{leach-2001,rajeev-2007}. Such an evolution on super-Hubble scales results in dips or bursts of oscillations in the scalar power spectrum. Usually, such a departure is induced by introducing a sharp feature in the potential of the inflaton field, such as a step or a sudden change in the slope~\\cite{starobinsky-1992,covi-2006-07,joy-2007-08,jinn-ouk-2005}. However, this is not necessary, and transitions to fast roll for brief periods can be generated even with smooth and better motivated effective potentials~\\cite{hodges-1990,tarun-1995,leach-2001,rajeev-2007}.  Our purpose in this paper is to present a simple model of inflation that supresses the power spectrum on large scales, a feature---as we discussed above---that seems to be necessary to fit the lower power in the quadrupole (and, to some extent, in the deviant power at other lower multipoles such as the octopole and the multipole $\\ell=22$) of the CMB angular power spectrum, using an effective potential of the canonical scalar field {\\it without}\\/ introducing any ad hoc sharp feature. We find that the form of the potentials motivated by a class of certain minimal supersymmetric extensions of the standard model provide us with the desired behavior~\\cite{allahverdi-2006,allahverdi-2007,sanchez-2007,allahverdi-2008}. These large field models allow a period of fast roll sandwiched between two stages of slow roll inflation\\footnote{Earlier, in the literature, two successive stages of slow roll inflation have often been driven by two scalar fields~\\cite{kofman-1985,silk-1987,polarski-1992-95,langlois-1999,tsujikawa-2003}. Instead, in this paper, we achieve the two stages of slow roll inflation including a brief period of departure from inflation, all with just a single scalar field.}. The first phase of slow roll inflation allows us to impose the standard Bunch-Davies initial conditions on the modes which exit the Hubble radius during the subsequent fast roll regime, an epoch due to which the curvature perturbations on the super-Hubble scales are suppressed. The second slow roll phase lasts for about $50$-$60$ $e$-folds, thereby allowing us to overcome the standard horizon problem associated with the hot big bang model. The advantages of our approach over other single field models mentioned earlier are twofold. Firstly, we do not need to assume any specific pre-inflationary phase. The entire evolution of the inflationary era is described by a single inflaton potential and, therefore, is much simpler. Secondly, the modes which exit the Hubble radius during the fast roll regime are inside the Hubble radius during the first stage of slow roll inflation and, hence, we do not have to impose any special initial conditions on the large scale modes.  This paper is organized as follows. In Sec.~\\ref{sec:bd}, we shall review the essential features of the effective inflaton potential that we shall consider, and describe the background dynamics in situations of our interest. In Sec.~\\ref{sec:ps}, after an outline of the slow roll `expectations' of the scalar spectrum that can arise in such a background, we shall discuss the spectra that we obtain through numerical integration. In Sec.~\\ref{sec:comparison}, using the cosmological Boltzmann code CAMB and the Monte Carlo code COSMOMC, we shall compare the power spectra from the models we consider with the recent WMAP 5-year data. Finally, we shall close with Sec.~\\ref{sec:discussion}, wherein after a brief summary of our results, we shall discuss as to how the results from our model compare with those that have been obtained in another closely related single field model.  In the discussions below, we shall set $\\hbar$ and $c$ as well as $M_{_\\mathrm{Pl}} = (8\\,\\pi\\, G)^{-1/2}$ to unity. Also, throughout, an overdot and an overprime shall denote differentiation with respect to the cosmic and the conformal times, respectively. ",
        "conclusions": "\\label{sec:discussion}  In this section, after a quick summary of our results, we compare the results we obtain with those obtained in another single scalar field model that has been considered earlier. We emphasize the fact that the difference between these models has immediate observational consequences.   \\subsection{Summary}  In this work, we have investigated a two stage slow roll inflationary scenario sandwiching an intermediate period of deviation from inflation, driven by potentials that are similar in shape to certain MSSM potentials~\\cite{allahverdi-2007}. In the MSSM case, inflation occurs when the field values are much smaller than the Planck scale~\\cite{allahverdi-2007,sanchez-2007,allahverdi-2008}. However, in our case, since we demand two epochs of slow roll, it necessarily requires that the initial values of the field (assuming, say, about $60$ $e$-folds of inflation) be greater than $M_{_{\\rm Pl}}$. The period of fast roll period produces a sharp drop in the scalar power spectrum for the modes that leave the Hubble radius just before the second slow roll phase. We choose our scales such that the drop in power corresponds to the largest cosmological scales observable today. We find that the resulting scalar power spectrum provides a much better fit to the recent WMAP data than the canonical, nearly scale invariant, power law, primordial spectrum.   \\subsection{Discussion}  At this stage, it is important that we compare our results with those obtained in another single field model that has been studied before. As we had mentioned in the introduction, an initial kinetic dominated (i.e. fast roll) stage preceding slow roll inflation, driven by a quadratic potential has been considered earlier to provide a sharp drop in the scalar power spectrum at large scales~\\cite{contaldi-2003,nicholson-2008}. At first glance, one may be tempted to conclude that the model we have studied here is equivalent to such a scenario if we disregard the first slow roll stage, since we have a kinetic dominated phase preceding a period of slow roll inflation. However, there are crucial differences between the two models which we have outlined below.  To begin with, in the model we consider, there is no freedom to choose the type of fast roll (say, the equation of state during the epoch of fast roll). It is fixed once we have chosen the parameters so as to fit the observations. Secondly, in the scenario considered earlier, the modes which are outside the Hubble radius during the kinetic dominated phase {\\it would have always remained so}\\/ in the past~\\cite{contaldi-2003,nicholson-2008}. The authors assume that somehow there may have been a previous phase of inflation, during which they were inside the Hubble radius and began life in the Bunch-Davies vacuum. While it is not impossible to think of situations where there may have been a previous inflationary epoch---for instance, it can be achieved by invoking another scalar field~\\cite{polarski-1992-95,langlois-1999,tsujikawa-2003}---the consequences can be quite different. In contrast, in the scenario that we have considered here, the standard, sub-Hubble, Bunch-Davies, initial conditions have been imposed on {\\it all}\\/ the modes. Thirdly, it was argued that, since the suppression of power for the scalar spectrum proves to be sharper than that of the tensor, the tensor-to-scalar ratio~$r$ displays a sharp rise towards large physical scales~\\cite{nicholson-2008}, a feature that may possibly be detected by upcoming missions such as, for instance, PLANCK~\\cite{planck}. However, in the models that we consider, the tensor amplitude on scales of cosmological interest proves to be too small ($r <10^{-4}$) to be detectable in the very near future. In conclusion, we would like to mention that a detection of the $C_{\\ell}^{^{\\rm BB}}$ modes corresponding to, say, $r > 10^{-4}$, can rule out the class of models that we have considered in this work.   \\subsection*"
    },
    "0809/0809.4704_arXiv.txt": {
        "abstract": "{Cosmologists have suggested a number of  intriguing hypotheses for the origin of the ``WMAP cold spot'', the coldest extended region seen in the CMB sky, including a very large void and a collapsing texture.  Either hypothesis predicts a distinctive CMB lensing signal.  We show that the upcoming generation of high resolution CMB experiments such as ACT and SPT should be able to detect the signatures of either textures or large voids.  If either signal is detected, it would have profound implications for cosmology.} ",
        "introduction": "One of the most intriguing features in the WMAP\\footnote{Wilkinson Microwave Anisotropy Probe; http://map.gsfc.nasa.gov/}\\citep{bennett.halpern.ea:2003} maps of the microwave sky is the Cold Spot \\citep{vielva.martnez-gonzalez.ea:2004,cruz.martnez-gonzalez.ea:2005,cruz.tucci.ea:2006,cruz.cayon.ea:2007}. Under the standard assumption of statistically homogenous Gaussian random fluctuations, the \\emph{a posteriori} probability of finding such a feature on the last scattering surface is less than $2\\%$ \\citep{cruz.martnez-gonzalez.ea:2005,cruz.tucci.ea:2006}. The Cold Spot also appears to have a flat frequency spectrum and is located in a region of low foreground emission,  making it unlikely to be caused by Galactic foregrounds or the Sunyaev-Zel'dovich effect \\citep{sunyaev.zeldovich:1970}. This has led some theorists to speculate that the Cold Spot is a secondary effect, generated at some intermediate distance between us and the last scattering surface. One such  model proposes that the Cold Spot may have been caused by the Rees-Sciama effect \\citep{rees.sciama:1968} due to an underdense void of comoving radius $\\sim 200 h^{-1} \\mathrm{Mpc}$ and fractional density contrast $\\delta \\sim -0.3$ at redshift of $z\\lesssim 1$ \\citep{inoue.silk:2006,inoue.silk:2007}. Interestingly, \\cite{rudnick.brown.ea:2007} reported a detection of an underdense region with similar characteristics  in the distribution of extragalactic radio sources in the NRAO VLA Sky Survey in the direction of the Cold Spot, a claim which has recently been challenged \\citep{smith.huterer:2008}. An alternative view \\citep{cruz.turok.ea:2007} proposes that the spot was caused by the interaction of the CMB photons with a cosmic texture, a type of topological defect that can give rise to hot and cold spots in the CMB  \\citep{turok.spergel:1990}. Bayesian analysis by \\cite{cruz.martinez-gonzalez.ea:2008} claims that the texture hypothesis seems to be favored over the void explanation, mainly because such large voids as required by the latter is highly unlikely to form in a $\\Lambda$CDM structure formation scenario. Irrespective of whether the Cold Spot was caused by a void or a texture, the CMB photons interacting with such an entity would have been gravitationally deflected. The deflections would lead to a systematic remapping of the primordial CMB anisotropies in and around the Cold Spot. In this brief report, we use simple analytic models for the void and the texture to address the issue of detectability of the gravitational lensing signature of either model, using upcoming high resolution CMB experiments.\\par For calculations presented in this paper, we assume a WMAP 5-year \\citep{dunkley.komatsu.ea:2008} flat $\\Lambda$CDM cosmology with a total matter density parameter $\\Omega_{m} = 0.258$ and a vacuum energy density  $\\Omega_{\\Lambda} =0.742$. The spectral index of the primordial power spectrum is set to $n_{s} = 0.963$ and the primordial amplitude for curvature perturbations is taken as $A_{s} = 2.41 \\times 10^{-9}$ at a pivot scale of $0.002 \\hmpc$. The present value of the Hubble parameter is taken as $H_{0} = 72 \\mathrm{~km ~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "If either a texture or a void is responsible for the WMAP cold spot, then there should be a distinctive lensing signature seen in the CMB.  We have shown that a void would gravitationally lens the CMB photons appreciably so that its presence should be detectable  with arcminute scale CMB experiments. For  a cosmic texture that collapsed at $z\\sim 6$,  we find that the gravitational lensing effect on the CMB is more subtle than the void, but should be detectable with longer integration. Together with other indicators, like the power spectrum and the bispectrum \\citep{masina.notari:2008} and measurements of the temperature-polarization cross-correlation \\cite{cruz.turok.ea:2007}, CMB lensing appears to be a powerful aid in constraining the theories of the WMAP cold spot anomaly."
    },
    "0809/0809.0891.txt": {
        "abstract": "% E+A galaxies have been interpreted as post-starburst galaxies based on % the presence of strong Balmer absorption lines combined with the % absence of major emission lines ([OII] nor H$\\alpha$). As a population % of galaxies in the midst of the transformation, E+A galaxies has been % subject to an intense research activity. %It has been, however, % difficult to investigate E+A galaxies statistically since E+A galaxies % are an extremely rare population of galaxies ($<$1\\% of all galaxies).   We have performed a spatially-resolved medium resolution long-slit spectroscopy of a nearby E+A (post-starburst) galaxy system, SDSSJ161330.18+510335.5, with the FOCAS spectrograph mounted on the Subaru telescope. This E+A galaxy has an obvious companion galaxy 14kpc in front with the velocity difference of 61.8 km/s. Both galaxies have obviously disturbed morphology. Thus, this E+A system provides us with a perfect opportunity to investigate the relation between the post-starburst phenomena and galaxy-galaxy interaction.   We have found that H$\\delta$ equivalent width (EW) of the E+A galaxy is greater than 7\\AA~ galaxy wide (8.5 kpc) with no significant spatial variation. The E+A galaxy have a weak [OIII] emission (EW$\\sim$1\\AA) by $\\sim$2.6 kpc offset from the peak of the Balmer absorption lines. We detected a rotational velocity in the companion galaxy of $>$175km/s.  The progenitor of the companion may have been a rotationally-supported, but yet passive S0 galaxy. %The progenitor of the E+A may have been pressure-supported but yet (puzzlingly) gas rich. We did not detect significant rotation on the E+A galaxy. Metallicity estimate based on the $r-H$ colour suggests $Z=0.008$ and 0.02, for the E+A and the companion galaxies, respectively. Assuming these metallicity estimates, the age of the E+A galaxy after quenching the star formation is estimated to be 100-500Myr, with its centre having slightly younger stellar population. The companion galaxy is estimated to have older stellar population of $>$2 gigayear of age with no significant spatial variation.  These findings are inconsistent with a simple picture where the dynamical interaction creates infall of the gas reservoir that causes the central starburst/post-starburst. Instead, our results present an important example where the galaxy-galaxy interaction can trigger a galaxy-wide post-starburst phenomena. ",
        "introduction": "\\begin{figure*} \\begin{center} %\\includegraphics[scale=0.6]{f_chart_sk.eps} %\\includegraphics[scale=0.6]{f_chart_sk_2.eps} \\includegraphics[scale=0.6]{f_chart_sk_3.eps} \\end{center} \\caption{The SDSS $g,r,i$-composite image of the J1613+5103. The long-slit positions are overlayed. The E+A galaxy is to the right (west), with bluer colour. The companion galaxy is to the left (east). }\\label{fig:fchart} \\end{figure*}    \\begin{figure*} \\begin{center} \\includegraphics[scale=0.49]{080112_J1613.ps_pages2} \\includegraphics[scale=0.49]{080112_J1613.ps_pages4} \\includegraphics[scale=0.49]{080112_J1613.ps_pages1} \\includegraphics[scale=0.49]{080112_J1613.ps_pages3} \\end{center} \\caption{Spectra of the entire galaxy. Right panels are for the E+A galaxy. Left panels are for the companion galaxy. Bottom panels are from the core. Top panels are through the slit shifted by 2'' to the north.  The spectra are shifted to the rest-frame and smoothed using a 5-pixel box. }\\label{fig:J1613_whole} \\end{figure*}  It has long been known that galaxies have varieties of morphology; Edwin Hubble was the first to classify galaxies into the so-called Hubble tuning fork, which mainly consists of elliptical, lenticular, spiral and barred spirals \\citep{1926ApJ....64..321H}. Because of similarity and gradual change of galaxy properties along the Hubble's tuning fork, people speculated this might be an evolutionary sequence, and extensively studied the galaxy evolutionary sequence since the dawn of the extragalactic astronomy \\citep[e.g.,][]{1994ARA&A..32..115R}. However, it is still not very well understood what physical mechanism drives the galaxy evolution, creating the diverse galaxy properties such as in the  Hubble tuning fork. As a galaxy in a transition phase, one of the key populations in elucidating this galaxy evolution is so-called E+A galaxies.  Galaxies with strong Balmer absorption lines without any emission in [OII] nor H$\\alpha$ are called E+A galaxies. The existence of strong Balmer absorption lines shows that E+A galaxies have experienced starburst recently \\citep[within a gigayear;][]{2004A&A...427..125G}.  However, these galaxies do not show any sign of on-going star formation as non-detection in the [OII] emission line indicates. Therefore, E+A galaxies have been interpreted as post-starburst galaxies, that is,  a galaxy which truncated starburst suddenly \\citep{1983ApJ...270....7D,1992ApJS...78....1D,1987MNRAS.229..423C,1988MNRAS.230..249M,1990ApJ...350..585N,1991ApJ...381...33F,1996ApJ...471..694A}. Recent study found that E+A galaxies have $\\alpha$-element excess \\citep{2007MNRAS.377.1222G}, which also supported the post-starburst interpretation of E+A galaxies. However, the reason why they started starburst, and why they abruptly stopped starburst remains one of the mysteries in the galaxy evolution. Since a post-starburst is an important stage of the overall galaxy evolution in the Universe, it is important to investigate the physical origin of the E+A galaxies.   At first, E+A galaxies are found in cluster regions, especially at higher redshift \\citep{1985MNRAS.212..687S,1986ApJ...304L...5L,1987MNRAS.229..423C,1988MNRAS.235..827B,1991ApJ...381...33F,1995A&A...297...61B,1996MNRAS.279....1B,1998ApJ...498..195F,1998ApJ...507...84M,1998ApJ...497..188C,1999ApJS..122...51D,1999ApJ...518..576P,2003ApJ...599..865T,2004ApJ...617..867D,2004ApJ...609..683T}. Therefore a  cluster specific phenomenon such as the ram-pressure stripping was thought to be responsible for the  violent star formation history of E+A galaxies \\citep{1951ApJ...113..413S,1972ApJ...176....1G,1980ApJ...241..928F,1981ApJ...245..805K,1983ApJ...270....7D,1999MNRAS.308..947A,1999PASJ...51L...1F,2000Sci...288.1617Q,2003ApJ...584..190F,2003ApJ...596L..13B,2004PASJ...56..621F}.  However, \\citet{2004MNRAS.355..713B} found that low redshift E+A galaxies are located predominantly in the field environment, suggesting that a physical mechanism that works in the field region is at least partly responsible for these E+A galaxies. Recently, \\citet{2005MNRAS.357..937G} has shown that E+A galaxies have more close companion galaxies than average galaxies, providing a statistical evidence that the dynamical merger/interaction could be the physical origin of field E+A galaxies. Dynamically disturbed morphologies of E+A galaxies also support this scenario \\citep{2006astro.ph.12053L,2005MNRAS.359.1557Y}.  To understand the origin of E+As is cluster-related or merger-driven (or both), independent evidence on the origin of E+A galaxies has been waited. Previous work mentioned above has been focused on the investigation of the global/external properties of E+A galaxies, such as the environment of E+A galaxies \\citep{2005MNRAS.357..937G}, and the integrated spectra of the E+A galaxies \\citep{2004ApJ...617..867D}. However, if the physical origin of E+A galaxies is merger/interaction or gas-stripping, these mechanisms should leave traces to the spatial distribution of stellar-populations within each E+A galaxy. For example, a centrally-concentrated post-starburst region is  expected to be found in case of the merger/interaction origin \\citep[e.g.,][]{1992ARA&A..30..705B}. In contrast, the gas-stripping would create a more uniform, galaxy-wide post-starburst region. Thus, the spatial-distribution of the post-starburst region inside the E+A galaxy contains important and independent clues on the physical origin of E+A galaxies \\citep[e.g.,][]{2005MNRAS.359.1421P,2005ApJ...622..260S}. In this series of papers \\citep{2006AJ....131.2050Y,YGH,goto_3D}, we try to obtain such independent hints on the origin of E+A galaxies by revealing internal structure of E+A galaxies. In this paper, we focus on a spatially-resolved long-slit scan spectroscopy of a nearby E+A system J1613+5103. This E+A galaxy previously studied by \\citet{YGH} has an obvious companion galaxy in front (Fig.\\ref{fig:fchart}), and thus, provides us with a perfect opportunity to reveal the relation between the post-starburst phenomena and galaxy-galaxy interaction. Color gradient of the E+A galaxy is studied in detail by \\citet{2005MNRAS.359.1557Y}. Compared to our own previous work \\citep{2006AJ....131.2050Y,YGH,goto_3D}, this time we obtained E+A spectra with higher resolution and better signal-to-noise taking advantage of the Subaru telescope.  Unless otherwise stated, we adopt the WMAP cosmology: $(h,\\Omega_m,\\Omega_L) = (0.7,0.3,0.7)$ \\citep{2008arXiv0803.0547K}.    \\begin{figure*} \\begin{center} %\\includegraphics[scale=0.49]{080115_J1613_ap9.ps_pages1} %\\includegraphics[scale=0.49]{080115_J1613_ap9.ps_pages2} %\\includegraphics[scale=0.49]{080115_J1613_ap9.ps_pages3} %\\includegraphics[scale=0.49]{080115_J1613_ap9.ps_pages4} \\includegraphics[scale=0.39]{080220_J1613_ap9.ps_pages4.ps} \\includegraphics[scale=0.39]{080220_J1613_ap9.ps_pages2.ps} \\includegraphics[scale=0.39]{080220_J1613_ap9.ps_pages3.ps} \\includegraphics[scale=0.39]{080220_J1613_ap9.ps_pages1.ps} \\end{center} \\caption{ Spectra in each aperture. Right panels are for the E+A galaxy. Left panels are for the companion galaxy. Bottom panels are from the core. Top panels are shifted by 2'' to the north. Each spectrum is spatially-divided into 9 equally-spaced bins. In each panel, these 9 spectra are normalized by the peak flux and aligned in the vertical direction (thus, the flux scale is arbitrary). The spectra are shifted to the rest-frame and smoothed using a 5-pixel box. }\\label{fig:J1613_ap9} \\end{figure*} ",
        "conclusions": "We have performed a spatially-resolved long-slit spectroscopy of the nearby interacting E+A system J1613+5103 using the FOCAS/Subaru. The E+A galaxy with very strong H$\\delta$ EW of $\\sim$7\\AA~(estimated age of 250Myr) has an obvious companion at 14 kpc in front. Combined with a disturbed morphology, the system is a clear example that a dynamical interaction can actually create an E+A galaxy. The slit was aligned to include cores of both the E+A galaxy and the companion galaxy. Additional spectra were taken at 2 arcsec away in the north from the cores. Our finding based on these spectra are as follows:  \\begin{itemize} \\item The strong Balmer absorption lines of the E+A galaxy are strong in the entire galaxy; Especially, H$\\delta$ EW is $>$7\\AA~ for the entire galaxy region ($>$8.5 kpc). \\item The peak of the weak (EW$\\sim$1\\AA) [OIII] emission of the E+A galaxy is slightly offset from that of the Balmer absorption lines, possibly suggesting that the remaining starburst is going on at a different location from the post-starburst region. \\item For the companion galaxy emission line ([OII]) profile is more extended than that of the H$\\beta$, suggesting on-going star-formation happens in more extended region than the central region of old stellar population. \\item We found a rotational velocity of $>$175 km/s for the companion galaxy (Fig.\\ref{fig:velocity}), implying the companion was a disk galaxy without star-formation activity before the dynamical interaction. No significant rotation was found for the E+A galaxy. %, implying an interesting possibility that the progenitor of the E+A might have been a pressure-supported but yet gas-rich galaxy. \\item Using the $r-H$ colour, we estimated the metallicity of the E+A and the companion were $Z=0.008$ and 0.02, respectively (Fig.\\ref{fig:nir}). \\item Based on the H$\\delta$-D4000 and  H$\\delta$-[OII] EW plots (Fig.\\ref{fig:d4000}), we estimate age of the E+A galaxy is around 250 Myr after the truncation of the burst. Possibly the centre of the E+A has a slightly younger age of $\\sim$100 Myr than its outskirts. \\end{itemize}  These results present an important example that a galaxy-galaxy interaction with equal mass can create a galaxy wide, uniformly-distributed post-starburst regions. In this example, it is obvious that the dynamical interaction is the physical origin of the E+A galaxy. At the same time, it warns us that simple classification of the E+A origin with the H$\\delta$ profile may not be adequate enough."
    },
    "0809/0809.2702_arXiv.txt": {
        "abstract": "Using an SPH simulation of a star-forming region in a molecular cloud, we show that the emergence of a clump mass function (CMF) resembling the stellar initial mass function (IMF) is a ubiquitous feature of molecular cloud structure, but caution against its over-interpretation. We employ three different techniques to extract the clumps in this study. In the first two, we interpolate the SPH particle data to 2 and 3 dimensional grids before performing the clump-find, using position-position (PP) and position-position-velocity (PPV) information respectively. In the last technique, the clump-finding is performed on the SPH data directly, making use of the full 3 dimensional position information. Although the CMF is typically similar to that observed in regions of nearby star formation, the individual clumps and their masses are found to be unreliable since they depend strongly on the parameters and the method of the clump-finding. In particular we find that the resolution and orientation of the data make a significant difference to the resulting properties of the identified clumps in the PP and PPV cases. We conclude that making comparisons between a CMF and the stellar IMF should be done with caution, since the definition of a clump boundary, and hence the number of clumps and their properties, are arbitrary in the extraction method. This is especially true if molecular clouds are truly scale free. ",
        "introduction": "\\label{sec:introduction}  Recent observations \\citep{Motte98,Johnstone00,Johnstone06, Nutter07,Alves07,Ward-Thompson07} of dense gas cores in molecular clouds have shown their distribution resembles the stellar initial mass function (IMF). It is hoped that by studying these dense cores  the earliest stages of star formation can be probed and the origins of the IMF revealed.  \\citet{Motte98} have shown that the clump mass function (CMF) in $\\rho$ Oph can be described by a similar power-law fit as the IMF. Above $\\sim 0.5$ M$_{\\odot}$ they find $dn/dm \\propto m^{-\\alpha}$ is well-fitted by an $\\alpha = 2.5$, while at lower masses they see a turn-over that can be fitted with $\\alpha = 1.5$. These values for $\\alpha$ are broadly consistent with the usual fits to the IMF \\citep{Kroupa02}. Similar results have been confirmed by a number of authors for a variety of nearby star-forming regions, although the range of core masses found and the break mass of the power law fit can vary \\citep{Johnstone00,Johnstone06, Nutter07,Alves07, Testi98}. A more detailed account of the clump/core mass distribution can be found in \\citet{Ward-Thompson07}.  One interpretation of the similarity between the CMF and IMF is that the masses of stars are controlled by the fragmentation of the cloud. \\citet{Myers08} propose that for cores with a high density contrast to their surroundings, the resulting protostellar mass will be the gas mass whose free-fall time equals the core dispersal time due to outflows. Additionally \\citet{Swift08} take a slightly different approach and argue that the resulting shape of an IMF formed from a Saltpeter-like CMF is robust against different core evolution scenarios.  In this paper, we show that the shape of clump mass function seems to be a ubiquitous feature of molecular clouds, but the identified clumps and their masses are sensitive to the method used to find them.  Molecular cloud (MC) structure as a whole is complex, filamentary and extends through all observable scales (\\citealt{Williams00}). It can also be thought of as hierarchical: small dense cores are part of small clumps of gas which are themselves part of larger gas clumps. This continues over the full spatial range of observations \\citep{Bergin07}. MCs also exhibit supersonic turbulence. \\citet{Larson81} showed that the internal velocity dispersion $\\sigma(v)$ of a cloud is related to its size $L$ by the relation $\\sigma(v)\\propto L^{-0.38}$. Such a scale-free relation is a natural outcome of a turbulent cascade, with \\citet{Brunt04} proposing that compressible shock dominated Burgers turbulence is most consistent with observations.  As supersonic flows compress the gas in the molecular cloud they create a complex filamentary density field, suggesting that the observed structure in MCs is a direct consequence of turbulent motions \\citep{Klessen00, Klessen01}. At the smallest scales of the hierarchy, the transonic dense cores formed by the turbulent cascade may be the precursors of star formation (\\citealt{Goodman98}, \\citealt{Ward-Thompson07}). Indeed \\citet{Padoan02} have predicted that mass distribution of Jeans unstable clumps created by super-Alvenic turbulence will resemble the IMF, a prediction that is supported by the observations of the CMFs discussed above.  However there are still questions as to whether the CMF observations can unambiguously support the cloud fragmentation picture. Typically, the observations must rely on two-dimensional column density maps (PP), which are limited by resolution and sensitivity, and the observed features in the column density may be contaminated by integrated emission along the line-of-sight. Three-dimensional information is only available by obtaining position-position-velocity (PPV) cubes of the emission from a particular molecular tracer. Converting between emission and column density relies on many assumptions \\citep{Pineda08} and accounting for the properties of dust is itself challenging \\citep{Schnee06}. Further, the observed emission is itself dependent on the entangled effects of density, temperature, molecular tracer properties, and the observational technique \\citep{DiFrancesco07}.  In this paper, we investigate how the differences in the form of the data -- whether it be PP, PPV or full three-dimensional PPP -- affects the resulting CMF. We explore how the resolution the data and orientation of the cloud affect the mass functions profile. Since we are interested in comparing observational style clump-finds to a full three dimensional clump-find, the data must necessarily be synthetic. As such we use the results from a high-resolution smoothed particle hydrodynamics (SPH) simulation of driven turbulence. A full description of this simulation can be found in Section \\ref{sec:simulation}. To produce the PP and PPV style clump-finds the SPH data is interpolated to a grid, while for the PPP data the clump-find is done directly on the SPH data. The method used to extract the clumps is similar to that described by \\citet{Williams94}, although there are some differences that can be found in Section \\ref{sec:clumpfind}. ",
        "conclusions": "We have shown that when a CLUMPFIND style algorithm is applied to different types of data-sets from the same underlying simulation, clump characteristics vary. This is due to a number of factors. Details of the observational maps or the simulation obviously play their part, particularly parameters  such as resolution, orientation and density range. The more levels of structure resolved within a data-set the greater the number of clumps that can be found. Blending of structure along the line of sight decreases the number of clumps found in 2D clumpfinds and increases clumps masses. A significant amount of material which is in reality not connected to a high density core can be artificially included. Applying a high density cut to PP data yields more distinct objects but can bias the resulting mass function.  Using velocity as a third dimension improves accuracy and provides a closer match to the true 3D structure than just using PP data. The clump boundaries, however, still differ and less mass is assigned to clumps.  Qualitatively the profile of the clump mass function always resembles the Stellar IMF, regardless of the data-set used. However, the quantitative values of the CMF are dependent on the extraction method, and therefore causal comparisons to the IMF should be made with caution."
    },
    "0809/0809.3368_arXiv.txt": {
        "abstract": "At present, ten years after they were first discovered, ten accreting millisecond pulsars are known. I present a study of the aperiodic X-ray variability in three of these systems, which led to the discovery of simultaneous kHz quasi periodic oscillations in XTE~J1807--294 and extremely strong broadband noise at unusually low variability frequencies in IGR~J00291+5934. Furthermore, we classified SWIFT~J1756.9--2508 as an atoll source and measured in its 2007 outburst spectral and variability properties typical of the extreme island state. I also give detailed estimates of the total fluence during the studied outbursts. ",
        "introduction": "During the last decade accreting millisecond pulsars (AMPs) have revealed a number of interesting phenomena and have opened a new window to the physics of accretion onto neutron stars (NSs). The first of such systems was discovered by \\citet[][SAX~J1808.4--3658]{Wijnands98}, presenting the first prove of an accreting neutron star having both millisecond spin period and dynamically important magnetic field. Ten AMPs have been discovered to date, and in three of them millisecond X-ray pulsations have been seen to appear in and disappear from the persistent emission, producing predominant \\citep[HETE~J1900.1-2455;][]{Galloway07} intermittent \\citep[SAX~J1748.9-2021;][]{Altamirano07} or very rare \\citep[Aql~X-1;][]{Casella07} episodes of pulsations. This implies that the AMP within them is only active or visible during a relatively small fraction of the time, which may provide a link with the much more numerous class of non-pulsating neutron star low-mass X-ray binaries (NS-LMXBs).  One way to study accretion onto compact objects is to analyze the aperiodic variability in the X-ray flux coming from these sources, which tells us about processes occuring in the inner accretion flow \\citep{vanderKlis95,vanderKlis06}. Such timing of the accretion flow, combined with a study of the X-ray spectrum, reveals different ``accretion states''. We show in this paper three different AMPs in three different accretion states, we describe their aperiodic variability and quantify their outburst fluence. ",
        "conclusions": ""
    },
    "0809/0809.3642_arXiv.txt": {
        "abstract": "Supergiant fast X-ray transients are a new class of high mass X-ray binaries recently discovered with \\int.\\ Hours long outbursts from these sources have been observed on numerous occasions at luminosities of $\\sim$10$^{36}$-10$^{37}$~erg~s$^{-1}$, whereas their low level activity at $\\sim$10$^{32}$-10$^{34}$~erg~s$^{-1}$ has not been deeply investigated yet due to the paucity of long pointed  observations with high sensitivity X-ray telescopes. Here we report on the first long ($\\sim$32~ks) pointed \\XMM\\ observation of \\j16479,\\ a member of this new class. This observation was carried out in March 2008, shortly after an outburst from this source, with the main goal of investigating its low level emission and physical mechanisms that drive the source activity. Results from the timing, spectral and spatial analysis of the EPIC-PN \\XMM\\ observation show that the X-ray source \\j16479\\ underwent an episode of sudden obscuration, possibly an X-ray eclipse by the supergiant companion. We also found evidence for a soft X-ray extended halo around the source that is most readily interpreted as due to scattering by dust along the line of sight to \\j16479.\\ We discuss this result in the context of the gated accretion scenarios that have been proposed to interpret the behaviour of supergiant fast X-ray transient. ",
        "introduction": "\\label{sec:intro}  Supergiant Fast X-ray transients (SFXTs) are a new class of high mass X-ray binaries (HMXBs), recently discovered with \\int.\\ These sources display sporadic outbursts lasting from minutes to hours with peak luminosities of $\\sim$10$^{36}$-10$^{37}$~erg~s$^{-1}$, and spend long time intervals at lower X-ray luminosities, ranging from $\\sim$10$^{33}$~erg~s$^{-1}$ to $\\sim$10$^{34}$~erg~s$^{-1}$. In some SFXTs a very faint state, with a typical luminosity of $\\sim$10$^{32}$~erg~s$^{-1}$ has also been observed. The outbursts of SFXTs have so far been studied in relatively fairly detail, whereas the lower luminosity states remained poorly known due to the paucity of long pointings with high sensitivity X-ray telescopes. Data collected with both wide field instruments, such as \\int,\\ and imaging X-ray telescopes, such as those on board \\chan,\\ \\swift\\ and \\XMM,\\ showed that the outburst spectra are well described by a power law of photon index $\\Gamma$$\\simeq$1-2 and an absorption column density that is larger than the interstellar Galactic value \\citep[$N_{\\rm H}$$\\sim$10$^{22}$-10$^{23}$~cm$^{-2}$,][]{sguera06,walter06}. Similar spectral properties were inferred at luminosities of $\\sim$10$^{33}$-10$^{34}$~erg~s$^{-1}$ \\citep{walter06,sidoli08}, whereas indications were found that very faint states might be characterized by significantly softer spectra \\citep{zand05}. The short duration of SFXT outbursts likely indicates the accretion flow towards the NS is not mediated by a disk (viscous timescales are of order of weeks to months). Instead, proposed models generally involve accretion onto a neutron star (NS) immersed in the clumpy wind of its supergiant companion \\citep{zand05,leyder,walter07,bozzo08}, and suggest these sources have orbital periods $\\gtrsim$7-10~days (or, equivalently, orbital separations $\\gtrsim$2.5~R$_*$, with R$_*$ the radius of the supergiant star).  \\\\ Different models make different predictions with respect to properties and origin of states with luminosities below the outburst levels. Observations of these states, however, are still sparse and the properties of SFXTs at low luminosities are poorly known at present. It is generally believed that the X-ray luminosity in these states is powered by residual accretion onto the NS, taking place at much reduced rate than in outburst. In fact, large variations in the mass accretion rates in SFXTs are expected due to: i) a centrifugal and/or magnetic gating mechanism acting on a moderately clumpy wind \\citep{bozzo08}; ii) clumps in the wind of the supergiant companion with extreme velocity and/or density contrasts \\citep{walter07}; iii) the presence of two components with different densities in the wind of the supergiant companion \\citep{sidoli07}. According to interpretation (i), the wide dynamic range in the X-ray luminosity (up to $\\sim$5 decades) displayed by SFXTs, if accompanied by a slow spin period ($\\gg$100~s), might indicate that these sources host ``magnetars'' \\citep[i.e., neutron stars with extremely high magnetic fields, $\\sim$10$^{15}$G,][]{duncan92}. During very faint states, an important contribution to the X-ray luminosity may also derive from shocks in the strong wind of the OB supergiant star \\citep{zand05}.  Here we report on the first long high sensitivity observation ($\\sim$32~ks) of the SFXT \\j16479,\\ when the source was in a low activity state. This source was detected in outbursts many times and appears to be characterized by a persistent luminosity of $\\sim$10$^{34}$~erg~s$^{-1}$ \\citep{sguera08,romano08,walter07}. The \\XMM\\ target of opportunity observation was carried out after \\swift\\ discovered a very bright outburst from this source on March 19 \\citep{romano08b}, and was aimed at investigating its low level emission and gaining insight in the physical mechanisms that drive its activity. ",
        "conclusions": "\\label{sec:discussions}  We reported on the first long ($\\sim$32~ks) pointed observation of \\j16479\\ during a period of low X-ray emission. The decay of the light curve and the increase of the iron line EW between states A and B, likely indicated that the X-ray radiation from this source was being obscured by an occulting body. Based on the present knowledge of \\j16479,\\ we suggested that \\XMM\\ might have caught the source entering an X-ray eclipse. The possibility that the obscuration event was caused by a wind clump cannot be excluded at present. However, we regard this interpretation as unlikely, because photoelectric absorption would be expected to cause a more pronounced flux decrease at soft X-ray energies than at hard X-ray energies (at least in the early stages of the obscuration event) contrary to the results presented here. On the other hand, an energy-independent obscuration could be caused by a fully ionized clump that is optically thick to electron scattering, a very unlikely possibility in consideration of the low X-ray luminosity of the source. Obscuration by a tilted accretion disk, which is known to cause similar effects in other HMXBs \\citep[see e.g,][]{zane04}, also does not appear as a very promising possibility here. This is because the outbursts of \\j16479,\\ and SFXT in general, have durations that are incompatible with viscous timescales in a typical accretion disc (see Sect.~\\ref{sec:intro}).  In order to interpret the source flux and spectral variations revealed by the \\XMM\\ observation, we considered the same scenario that has been extensively discussed for eclipsing HMXBs with a dust scattering halo \\citep[see e.g.,][]{audley06,ebisawa96}: during the eclipse ingress source photons reaching us directly dominated the X-ray flux we observed; while in eclipse the residual X-ray flux is due to the sum of a wind-scattered component and a dust scattering halo. Strong evidence for the wind scattered component in the case of \\j16479\\ comes from the marked increase of the EW of the iron fluorescence line at $\\sim 6.5$~keV. A tens of arcsec wide dust-scattering halo around \\j16479\\ could be directly seen in the soft X-ray \\XMM\\ images accumulated during the eclipse. Moreover the residual soft X-ray decreased on a timescale of $\\sim 1$~hr after the eclipse onset, as expected for a dust scattering halo of the size seen by \\XMM\\ at the source distance of 4.9~kpc. \\begin{figure} \\centering \\includegraphics[scale=0.48]{psf.ps} \\caption{\\XMM\\ EPIC-PN PSF as a function of the distance from the source. The upper panels are for images extracted in states A and B in the energy range 2-5~keV. Instead, the PSF in the lower panels are extracted from the same images but in the 5-10~keV energy range. For this analysis we used the {\\sc ximage} tool (version 4.4).} \\label{fig:psf} \\end{figure}  We conclude that \\j16479\\ is the first SFXT that displayed evidence for an X-ray eclipse. In this case, some constraints on the orbital period of this system can be derived by using \\citep{rappaport83} \\begin{equation} R_*=a\\left[\\cos{i}^2+\\sin{i}^2\\sin{\\Theta_{\\rm e}}\\right]^{1/2}. \\label{eq:relation} \\end{equation} Here $i$ is the inclination of the orbit to the plane of the sky, $\\Theta_{\\rm e}$ is the eclipse half-angle, $a$ is the binary separation, and we assume in the following a supergiant star with a mass of $\\sim$30~M$_{\\odot}$ and a radius of R$_*$$\\sim$20~R$_{\\odot}$. We consider circular orbits, as it is generally expected for an HMXB with a supergiant companion \\citep{rappaport83}. By using Kepler's law, Eq.~\\ref{eq:relation} can be solved for the orbital period, as a function of the inclination angle, for any fixed value of the eclipse duration. This is shown in Fig.~\\ref{fig:eclipse}. The solid line in this figure represents the solution obtained by assuming an eclipse duration of $\\sim$28~ks. This value should be regarded as a lower limit on $\\Theta_{\\rm e}$, since in the \\XMM\\ observation discussed here there is no evidence for the eclipse egress. The dashed line is for an eclipse that lasts for 55~ks (chosen for comparison with the previous case), whereas the triple-dot-dashed line is for an eclipse that extends for half of the binary orbital period. The latter provides an upper limit  on $\\Theta_{\\rm e}$. We also include in the figure, the lower limit on the orbital period that was obtained by \\citet{negueruela08}, by assuming the sporadic outbursting activity of SFXTs is related to accretion of clumps from the supergiant stellar wind (see Sect.~\\ref{sec:intro}). If we consider that the clumpy wind model applies to \\j16479,\\ then the region of allowed parameter space in Fig.~\\ref{fig:eclipse} suggest this source is a high inclination SFXT. From this figure we also note that, for this source, an eclipse with a duration much longer than the $\\sim$28~ks observed here is very unlikely: in fact the longer is the duration of the eclipse, the smaller is the allowed parameter region in Fig.~\\ref{fig:eclipse}. In particular, a much longer eclipse would require a very high inclination angle and short orbital period that can be hardly reconciled with the requirements of the clumpy wind model (the NS orbit should be very close to the supergiant surface, and a persistent luminosity of $\\gtrsim$10$^{36}$~erg~s$^{-1}$ would be expected). \\begin{figure} \\centering \\includegraphics[scale=0.5]{orbital.ps} \\caption{Orbital period constraints on \\j16479,\\ as a function of the binary inclination angle. The solid line is given by Eq.~\\ref{eq:relation} assuming an eclipse duration of 28~ks, the dashed line is for an eclipse duration of 55~ks, and the triple-dot-dashed line is for an eclipse lasting half of the binary orbital period. The dot-dashed line is the lower limit on the orbital period obtained by applying the clumpy wind model to \\j16479.\\ According to this model, the region of allowed parameters in this plot indicates \\j16479\\ is a high-inclination SFXT.} \\label{fig:eclipse} \\end{figure}  \\citet{walter07} derived the average quiescent flux of \\j16479\\ as observed by \\int\\ over $\\sim$67~d, which converts to an X-ray luminosity of few 10$^{34}$~erg~s$^{-1}$. This relatively high luminosity is comparable with that of state A, and thus provides support in favor of state B being an unfrequent state, as expected in the eclipse interpretation (eclipses cannot last more than half of the orbital period).  At present, it is not clear whether eclipses should be expected also in other SFXTs. However, we note that, if other SFXTs display X-ray eclipses, their maximum luminosity swing might not be due to genuine changes in the mass accretion rate, impacting on the application of the gating accretion model for SFXT \\citep{bozzo08}. Further observations of \\j16479,\\ as well as other SFXTs in quiescence, will improve our knowledge of the low level emission of these sources, and will help clarifying this issue."
    },
    "0809/0809.4254_arXiv.txt": {
        "abstract": "We report on the results of a long term X-ray monitoring campaign of the galactic binary LS I +61 303 performed by the Rossi X-ray Timing Explorer. This dataset consists of 1 ks pointings taken every other day between 2007 August 28 until 2008 February 2. The observations covered six full cycles of the 26.496 day binary period and constitute the largest continuous X-ray monitoring dataset on LS I +61 303 to date with this sensitivity. There is no statistically strong detection of modulation of flux or photon index with orbital phase; however, we do find a strong correlation between flux and photon index, with the spectrum becoming harder at higher fluxes. The dataset contains three large flaring episodes, the largest of these reaching a flux level of 7.2$^{+0.1}_{-0.2}$$\\times$10$^{-11}$ erg cm$^{-2}$s$^{-1}$ in the 3-10 keV band, which is a factor of three times larger than flux levels typically seen in the system. Analysis of these flares shows the X-ray emission from LS I +61 303 changing by up to a factor of six over timescales of several hundred seconds as well as doubling times as fast as 2 seconds.  This is the fastest variability ever observed from LS I +61 303 at this wavelength and places constraints on the size of the X-ray emitting region. ",
        "introduction": "The galactic binary LS I +61 303 is one of the most heavily studied binary star systems in the Milky Way, and is one of only three so called ``TeV Binaries'' to be regularly detected in the $>$100 GeV gamma-ray regime; the other two being LS 5039 (Aharonian et al. 2005a) and PSR B1259-63 (Aharonian et al. 2005b). Although the subject of many observational campaigns, the fundamental identification of the components of the system remains relatively unclear. From observations detailed in Hutchings and Crampton (1981) and Casares et al. (2005) it is clear that the system can be classified as a high mass X-ray binary (HMXB) located at a distance of 2 kpc; the components of the system consisting of a compact object in a 26.496($\\pm$0.003) day orbit around a massive BO Ve main sequence star. Although the range of masses derived for the compact object favor a neutron star component, a black hole cannot be ruled out (Casares et al. 2005). Periastron passage of the compact object occurs at $\\phi=$0.23 (here $\\phi$ represents the orbital phase ranging from 0.0 to 1, $\\phi$=0.0 set at  JD 2443366.775), with apastron passage occuring at 0.73, and inferior and superior conjunctions occuring at 0.26 and 0.16, respectively (Casares et al 2005).  The two main competing scenarios to explain the system can be summarized as $\\textit{microquasar}$ (i.e. non-thermal emission powered by accretion and jet ejection) or $\\textit{binary pulsar}$ (i.e. non-thermal emission powered by the interaction between the stellar and pulsar winds). The microquasar model (Massi et al. 2001,  Bosch-Ramon et al. 2006) used to describe this system has several drawbacks (for example, lack of blackbody X-ray spectra), however the scenario has not been ruled out, for example see Romero et al. (2007). There is growing evidence for the binary pulsar model of the system (Maraschi and Treves 1981, Dubus 2006), most strongly supported by the result of Dhawan (2006) in which VLBA monitoring of the system revealed a cometary radio structure around LS I +61 303 which was interpretated as the interaction between the pulsar and Be star wind structures. However, there is currently no evidence for the presence of a pulsar within the system either in radio or X-rays.  LS I +61 303 has been historically an object of interest due to its quasi-periodic outbursts at radio (Paredes et al. 1998, Gregory 2002) and X-ray wavelengths (Taylor et al. 1996, Leahy et al. 1997, Paredes et al. 1997, Harrison et al. 2000, Greiner and Rau 2001). The radio outbursts are well correlated with the orbital phase (Gregory 2002),  although the phase of maximum emission can vary between 0.45 and 0.95. The first high energy association of the source was with the COS-B source 2CG 135 +01 (Hermsen et al. 1977). LSI +61 303 has also been identified with the EGRET source 3EG J0241+6103, a source which also shows evidence for a 26.5 day modulation in the GeV band (Massi 2004b). More recently, LS I +61 303 has been detected as a variable TeV gamma-ray source (Albert et al. 2006, Acciari et al. 2007) with high emission appearing near apastron.  The collection of X-ray observations on LS I +61 303 is expansive, however, the exact character of the X-ray emission from this source remains unclear. The X-ray source was originally weakly identified by the Einstein satellite (Bignami et al. 1981). Further observations of the system by ROSAT in 1991 and 1992 (Goldoni and Mereghetti 1995, Taylor et al. 1996) in the 0.01-2.4 keV band showed the source to be variable by a factor of 3 over a single orbital cycle; a peak flux occuring between orbital phases 0.4 and 0.6. The first accurate observations in the $>$5 keV range were performed by ASCA (Leahy et al. 1997) in the 0.5-10 keV band. ASCA performed two separate 30 ks exposures near orbital phases 0.2 and 0.45 showing relatively low flux states (0.6/0.4 $\\times$10$^{-11}$erg cm$^{-2}$s$^{-1}$). The spectra from these observations were well fit by an absorbed powerlaw with photon indices of 1.7$\\pm$0.1 and 1.8$\\pm$0.1, respectively.  RXTE observed LS I +61 303 for an entire orbital cycle in 1996, seeing a clear increase in flux between $\\phi$=0.3 and 0.7, with a peak occurring near phase 0.45 at a flux of 2 $\\times$10$^{-11}$erg cm$^{-2}$s$^{-1}$ (Harrison et al. 2000). There have been three separate published analyses of this dataset (Harrison et al. 2000, Greiner and Rau 2001, Neronov and Chernyakova 2006). All three analyses find that a simple absorbed power law provides a resonable fit to the derived spectra. In 2002, XMM-\\textit{Newton} was used to monitor LS I +61 303 in the 0.2-12 keV range (Neronov and Chernyakova 2006). Four 5-6 ks pointings were taken spread out over a single orbit, with a single 6 ks pointing taken several months later. These pointings showed the 2-10 keV flux to be highly variable, having a minimum near periastron and peaking near phase 0.55. The spectral fitting (again, using an absorbed powerlaw) was consistent with a photon index value of 1.5 for 4 out of 5 of the pointings, with a much softer index of 1.78 resulting for the data point just preceding the transition from a high to low X-ray flux. Neronov and Chernyakova (2006) interpreted this as evidence for correlation between spectral behavior and flux states.  Several years later in 2005, XMM-\\textit{Newton} was again used for an extended pointing (48.7 ks) on LS I +61 303 during orbital phase 0.61 (Sidoli et al. 2006). This observation showed evidence for variation of the hardness ratio (defined as the ratio between 0.3$\\rightarrow$2 keV and 2$\\rightarrow$12 keV emission) over the timescale of hours. These observations also showed a sharp decrease in flux of the order of a factor of three over a time period of a few thousand seconds, with flux decreasing from 1.2$\\rightarrow$0.4 $\\times$10$^{-11}$erg s$^{-1}$cm$^{-1}$, which were the first observations to detail such rapid variability of the X-ray flux from this system. The Chandra satellite carried out a 50 ks pointed observation of LS I +61 303 in April 2006 (Paredes et al. 2007) near orbital phase 0.04. These observations (0.3 to 10 keV) revealed the presence of kilosecond scale miniflares in the source, with emission increasing by a factor of 2 over a timespan of roughly one hour. Similar to the XMM-\\textit{Newton} extended exposure, the Chandra observations also show an implied correlation between harder emission and increased flux. The mean flux was determined to be 0.71$^{+0.18}_{-0.14}$ $\\times$ 10$^{-11}$erg s$^{-1}$cm$^{-1}$ which is also consistent with previous measurements near that orbital phase. The derived photon index, however, was determined at $\\alpha$=1.25$\\pm$0.09, much harder than any previous measurements made. The Chandra exposure also found evidence (3.2$\\sigma$ significance) for extended emission between 5 and 12.5\" towards the North of LS I +61 303, an indication that particle acceleration resulting in X-ray emission may be taking place as far away as 0.05-0.12 parsecs from LS I +61 303.  Taken as a whole, the X-ray observations conducted on LS I +61 303 indicate that the source is an extremely unpredictable one. Although long term X-ray monitoring conducted by the All Sky Monitor aboard RXTE indicate that the system exhibits a long term 26.42 $\\pm$0.05 day period in X-ray emission (Leahy 2001), the X-ray observations  descibed above show that extreme variation in the lightcurve is seen between different orbital cycles. As for spectral behavior it is clear that a simple absorbed powerlaw provides an effective fit to the observational data, however, there is mounting evidence that the photon index may be correlated to the flux level of the system. ",
        "conclusions": "The current RXTE dataset demonstrates the overall variability of the X-ray behavior of the system. Additionally, we find a strong correlation between photon index and flux from the system. While previously believed to be an X-ray source with somewhat dependable modulation, the observations reported here show strong flares which complicate the measurement of the typically expected flux and spectral variations as a function of binary orbital phase. The X-ray behavior (with respect to the orbital phase) is not completely understood.  The three large flares observed in this dataset represent the strongest X-ray activity recorded from this system to date with luminosities exceeding 3$\\times$10$^{34}$erg s$^{-1}$. These flares show that the flux levels within the flare structures rise and fall by factors of up to 6 on the 100s time scale, with doubling times as short as 2 seconds (FL1). Through this variability, the causal size of the emission region can be constrained to be less than R$<c\\Delta$t = 6$\\times$10$^{10}$ cm. It should be noted that this constraint does not account for the scenario in which the emission is originating from a relativistic jet with a Doppler factor which would modify the above calculation. Since the existence of such a jet within the system is not clearly evident (and much less its Doppler factor), this modification is neglected at this time. Although rapid flaring over the kilosecond timescale has been demonstrated in both Chandra and XMM-$\\textit{Newton}$ data, this is the first indication that the X-ray flux from the sytem varies on such a rapid scale.  We note that fast X-ray flaring behavior has been observed in binary X-ray pulsars such as IGR J16465-4507 and AX J1841.0-0536 (Lutivinov et al. 2005, Sguera et al. 2006 ). Following along these lines,  a possible scenario which could be used to explain the available observations is the binary pulsar model of Zdziarski et al. (2008). In this scenario, the radio through TeV emission results from the interactions between the pulsar wind and a two-component Be star wind. The two components consist of a fast ``clumpy'' polar wind and a dense circumstellar disk. When the pulsar is close enough to the Be star to be heavily exposed to the fast polar wind, clumps of relativistic electrons in the polar wind are maintained by magnetic field inhomogeneities and are sufficiently energetic to cause mixing with the pulsar wind. Shocks between the clumps and the pulsar wind accelerate electrons which can then inverse-Compton (IC) scatter off of the Be star photons, giving rise to the highly variable X-ray emission. Clumps such as these have been observed on the scale of 10$^{11}$ cm in 2S 0114+65 (Apparao, Bisht, and Singh 1991), which agrees (within error) with the implied timescale for variability observed in this work. However, even if clumps are much larger than this, it is believed that the interaction between a clump and the pulsar wind will cause a flare in a region with characteristic size equal to the stand-off distance between the pulsar and equatorial stellar wind (Dubus 2008b). At periastron passage the stand-off distance can be as small as 5$\\times$10$^{10}$ cm, which is consistent with the variability reported in this work. On a larger timescale, the pulsar wind moves through region of varying density of the equatorial wind. During traversal of the densest parts of the disk, X-ray emission is quenched due to Coulomb losses, and when the pulsar moves towards the more spare regions of the disk, IC losses can dominate over both Coulomb and synchrotron mechanisms, giving rise to a orbital modulation of the X-ray flux. Although this dataset does not strongly indicate a modulation of the X-ray flux with orbital phase (possibly due to complications from shorter timescale flares), the possibility of low amplitude modulation is still consistent with the observations reported here.  We also note that fast X-ray variability such as that reported here has been observed in good candidates for galactic black hole systems such as Cygnus X-1 (Gierlinski and Zdziarski 2003) and GRS 1915 +105 (Belloni et al. 2000); which are both known accretion driven sources. RXTE observations of the TeV binary LS 5039 (Bosch-Ramon et al. 2005a) show a photon index versus flux correlation and hourly timescale flux variations similar to what we find for LS I +61 303, but no very rapid variability. The authors interpret these results in an accretion driven scenario where a stellar wind feeds a disk around either a black hole or neutron star and argue that the X-ray emission is caused by a population of relativistic electrons accelerated within the jet, which is in turn fed by the wind fed accretion disk. These electrons are responsible for the observed X-ray emission via either the synchrotron process or inverse-Compton scattering of local stellar photons. In this scenario, fast variations in the stellar wind are responsible for the observed ``miniflares'', while sudden increases in accretion (higher flux) would also result in a higher electron acceleration efficiency in the jets. This efficiency increase results in a higher peak of the electron energy distribution,  in turn resulting in a harder emission spectra. Additionally, in the model described in Bosch-Ramon et al. (2005b),  they predict a jet length of $\\sim$10$^{11}$ cm, which is consistent with the several second variability described in this work. It should be noted that no X-ray spectral signatures of an accretion disk (i.e. a blackbody spectral component) are found in LS 5039 either, further reinforcing the link between the X-ray emission of the two sources.  Recently, Swift-BAT has reported the detection of a short (0.23s), extremely powerful X-ray burst with a luminosity of 10$^{37}$ erg in the 15-150 keV energy range within the 90$\\%$ containment radius of LS I +61 303 (Barthelmy et al. 2008). Follow up observations with the Swift-XRT instrument 921 seconds later showed no significantly elevated flux in the 0.3-10 keV energy range (0.92$\\times$10$^{-11}$ erg cm$^{-2}$s$^{-1}$). While the authors of this report note the possibility that this emission was due to an unrelated short gamma-ray burst, they claim that the evidence is in favor of activity from a source within LS I +61 303. Dubus and Giebels (2008) submit that this burst could be evidence in favor of magnetar-like activity from a source within LS I +61 303. Although reconciling the presence of a magnetar within LS I +61 303 with previous X-ray observations is somewhat problematic (Bosch-Ramon 2008, Dubus 2008a), this could be the first magnetar to be discovered within a high mass X-ray binary system. If confirmed with following observations, this type of extremely powerful bursting may resolve the question of the identification of LS I +61 303.  In conclusion, while neither conclusively ruling out or confirming either the microquasar or binary pulsar scenarios, the observations reported here will aid further modeling work of LS I +61 303 by the demonstrated constraint on emission region size, as well as the correlation between flux and photon index."
    },
    "0809/0809.1653_arXiv.txt": {
        "abstract": "The modification of Einstein gravity at high energies is mandatory from a quantum approach. In this work, we point out that this modification will necessarily introduce new degrees of freedom. We analyze the possibility that these new gravitational states can provide the main contribution to the non-baryonic dark matter of the Universe. Unfortunately, the right ultraviolet completion of gravity is still unresolved. For this reason, we will illustrate this idea with the simplest high energy modification of the Einstein-Hilbert action: $R^2$-gravity. ",
        "introduction": "Different astrophysical observations agree that the main amount of the matter content of our Universe is in form of unknown particles that are not included in the Standard Model (SM). Typical candidates to account for the missing matter can be found in well motivated extensions of the electro-weak sector. However, there is a fundamental sector in our model of particles and interactions, where the introduction of new degrees of freedom is not only well motivated, but absolutely necessary.  The non-unitarity and non-renormalizability of the gravitational interaction described by the Einstein-Hilbert action (EHA) demands its modification at high energies. The main idea of this work is to realize that this correction cannot be accomplished without the introduction of new states; these states will typically interact with SM fields through Planck scale suppressed couplings and potentially work as dark matter (DM). ",
        "conclusions": "In conclusion, we have studied the possibility that the DM origin resides in UV modifications of gravity. Although our results may seem particular of $R^2$-gravity, the low energy phenomenology of the studied scalar mode is ubiquitous in high energy corrections of the EHA coming from string theory, supersymmetry or extra dimensions (reproduced in form of dilatons, radions, graviscalars or other moduli fields). The instability of the deduced DM predicts a deviation from EEs at an energy scale: $1\\,\\tev \\lesssim \\Lambda_G\\lesssim 10^5\\,\\tev$. Consequences for hierarchy interpretations, baryogenesis or inflation deserve further investigations."
    },
    "0809/0809.4076_arXiv.txt": {
        "abstract": "We present thick de Sitter brane solutions which are supported by two interacting {\\it phantom} scalar fields in five-, six- and seven-dimensional spacetime. It is shown that for all cases regular solutions with anti-de Sitter asymptotic (5D problem) and a flat asymptotic far from the brane (6D and 7D cases) exist. We also discuss the stability of our solutions. ",
        "introduction": "In recent years, there is growing interest in higher-dimensional cosmological models inspired by the recent developments in string theory. One sort of such models are branes which are submanifolds embedded into the higher-dimensional spacetime. They play an important (even essential) role, for instance, in a description of the confinement of the nongravitational interactions and/or stabilization of the extra dimensions. Braneworld models, in which our universe is exactly one of four-dimensional branes, are the central issues in this research field. Much efforts to reveal cosmology on the brane have been done in the context of five-dimensional spacetime, especially after the stimulating proposals by Randall and Sundrum \\cite{rs} (see, e.g., \\cite{bw}). But of course there is no particular reason to restrict the number of dimensions to be five. In recent years, the focus has turned to higher-dimensional braneworld models.     It is still unclear how the gravity and cosmology on the brane embedded in the higher-dimensional spacetime behave. In higher dimensions, there would be many kinds of classes of possible braneworld models. Here we focus on the generic feature of higher-dimensional braneworld models and from this we deduce the subject to be revealed in this paper.    It is commonly assumed that a brane is an infinitely thin object. In five dimensions, this is a good approximation as long as one is interested in the scales larger than the brane thickness. Then, the cosmology on the brane in five dimensions, i.e., dynamics of the brane in one bulk spacetime, is uniquely determined by the Israel junction conditions once the matter on the brane is specified. But, in higher dimensions the situation is quite different. It happens because self-gravity of an infinitesimally thin three-brane in a higher-dimensional spacetime develops a severe singularity and one cannot put any kind of matter on the brane if the number of codimensions is larger than 2. Also in the case of a codimension-two brane, one cannot put any kind of matter other than the pure tension. The cosmology on such a brane must be investigated under some prescription of the singular structure of the brane. The well-motivated prescription is to regularize the brane by taking its microscopic structure into consideration. However, in the case of six-dimensional models, the cosmology on the brane strongly depends on the way of regularization. This implies that there is no unique brane cosmology in the case of higher codimensions.    Thus, in this paper we take a different point of view. We shall take the brane thickness into account {\\it from the beginning} and look for the regular braneworld solutions with finite thickness, called {\\it thick brane solutions}, in a field theory model. The thickness of the brane may be very close to the length scale of the quantum gravitational theory, e.g., the string length scale. As we mentioned, the brane thickness must be more essential in higher-dimensional spacetimes than in the five-dimensional one. Thus, we look for the thick brane solutions in higher dimensions as well as in five dimensions. Bearing the cosmological applications, we focus on the thick brane solutions which have de Sitter geometry in the ordinary four-dimensions.    The properties of thick branes in five-dimensions, especially localization of gravity, have been discussed in Ref. \\cite{tb_gen}. In five dimensions, the thick brane solutions in five-dimensional spacetime have been explored in various literatures (see, e.g., \\cite{thick}). In this paper, we consider thick brane solutions supported by two interacting scalar fields $\\phi$ and $\\chi$, whose potential is given by \\begin{equation} \\label{pot2} V(\\phi,\\chi)=\\frac{\\lambda_1}{4}(\\phi^2-m_1^2)^2+ \\frac{\\lambda_2}{4}(\\chi^2-m_2^2)^2+\\lambda_3 \\phi^2 \\chi^2+\\Lambda. \\end{equation} A stronger interaction between two scalar fields develops the local metastable vacua, which may be able to support our brane universe, other than the global one. Several works on thick brane solutions in such a setup have already been performed. Five-dimensional Minkowski brane solutions were derived in \\cite{2scalar1, phantom}. Bearing cosmological applications in mind, our goal in this work is to find thick de Sitter brane solutions. Thick brane solutions in higher-dimensional spacetimes with two interacting scalar field were also investigated in \\cite{2scalar2} (see, e.g., \\cite{higher_dim} for other kinds of models). We shall give de Sitter versions of these solutions.   In Ref.~\\cite{gauge} there are some arguments in favor of the fact that the scalar fields, used in \\cite{2scalar2}, are a quantum nonperturbative condensate of a SU(3) gauge field. Briefly these arguments consist of the following: components of the SU(3) gauge field can be divided in some natural way onto two parts. The first group contains those components which belong to a subgroup $SU(2) \\in SU(3)$. The remaining components belong to a factor space $SU(3)/SU(2)$. Following Heisenberg's idea \\cite{heisenberg} about the non-perturbative quantization of a nonlinear spinor field the nonperturbative quantization for the SU(3) gauge field is being carried out as followings. It is supposed that two-point Green functions can be expressed via the scalar fields. The first field $\\phi$ describes two-point Green functions for SU(2) components of a gauge potential, and the second scalar field $\\chi$ describes two-point Green functions for $SU(3)/SU(2)$ components of the gauge potential. It is supposed further that four-point Green functions can be obtained as some bilinear combination of the two-point Green functions. Consequently, the Lagrangian of the SU(3) gauge field takes the form \\eqref{Daction}. It allows us to regard such sort of thick brane solutions as some defect in a spacetime filled by a condensate of the gauge field living in the bulk.    This paper is organized as follows. In Sec. II, we introduce the basic field theory model to find thick brane solutions. In Sec. III, we present thick brane solutions in five-, six- and seven-dimensional models. In Sec. IV, we investigate perturbations in the five-dimensional model and then give speculations about stability of higher-dimensional models. Secion V contains a brief summary. ",
        "conclusions": "In this paper, we have presented five-, six-, and seven-dimensional thick de Sitter brane solutions supported by two interacting ({\\it phantom}) scalar fields. The special form of the potential~\\eqref{pot2} allows us to find regular solutions with finite energy density. It was shown that asymptotically there exist anti-de Sitter spacetime for the five-dimensional case and flat spacetime for the six- and seven-dimensional cases.  Also we have investigated the stability of our thick brane solutions. In the five-dimensional thick de Sitter brane solutions, we have shown the presence of decaying solutions for perturbations. We also could not find any mode with negative eigenvalue. This, of course, implies that our solutions in five-dimensional spacetime are stable. It is also quite reasonable to expect that our thick brane solutions in higher dimensions are also stable since the spacetime and vacuum structures are essentially the same as in the case of five spacetime dimensions."
    },
    "0809/0809.4306_arXiv.txt": {
        "abstract": "Supernova remnant (SNR) \\snr\\ is morphologically characterized by brightened radio, infrared, and X-ray emission on the southeastern rim, with the 1720~MHz OH masers detected in the northeastern and southeastern regions at various local standard rest (LSR) velocities. We have performed a millimeter observation in CO and \\HCOp\\ lines toward \\snr. From the northeastern compact maser region, \\twCO\\ and \\thCO\\ emission's peaks around $65\\km\\ps$ and $85\\km\\ps$, which are consistent with the masers' LSR velocities, are detected. In the southeast, a molecular (\\twCO) arc is revealed at 77--$86\\km\\ps$, well coincident with the partial SNR shell detected in the radio continuum and mid-infrared observations. An $85\\km\\ps$ \\HCOp\\ emission is found to arise from a radio peak on the shell. Both the molecular arc and the \\HCOp\\ emission at $\\sim85\\km\\ps$ seem to be consistent with the presence of extended OH masers along the southeastern boundary of \\snr. The morphology correspondence between the CO arc and other band emission of the \\snr\\ shell provides strong evidence for the association between SNR~\\snr\\ and the $\\sim85\\km\\ps$ component of molecular gas. The multiwavelength emissions along the southeastern shell can be accounted for by the impact of the SNR shock on a dense, clumpy patch of molecular gas. This pre-existing gas is likely to be a part of the cooled debris of the material swept up by the progenitor's stellar wind.  The association of SNR~\\snr\\ with the molecular cloud at the systemic velocity of $\\sim85\\km\\ps$ enables us to place the SNR at a kinematic distance of 5.2~kpc. ",
        "introduction": "The progenitors of core-collapse supernovae are most probably formed in giant molecular clouds (MCs). Due to the short lifetime, they are not far away from their matrices when they explode. Therefore it is common that the supernova remnants (SNRs) are located in the vicinity of MCs and may encounter them in evolution. The association with MCs often results in irregular morphology of the SNRs in multiwavelengths, which indicates sophisticated shock interaction with the inhomogeneous environmental medium. About 20 SNRs have been discovered to be physically interacting with ambient molecular gas based on the detection of the 1720~MHz OH masers (Frail et al.\\ 1996), which are believed to be a tracer of the shock interaction with MCs (Lockett et al.\\ 1999; Frail \\& Mitchell 1998; Wardle \\& Yusef-Zadeh 2002).  SNR~Kesteven~69 is thought to be probably associated with MCs, because of the OH masers detected toward this remnant; however, the masers are found at various local standard rest (LSR) velocities and at various projected locations. Green et al.\\ (1997) detected a compact OH maser using the Very Large Array (VLA) and the Australia Telescope Compact Array (ATCA) at $\\VLSR=69.3\\km\\ps$, which is located projectionally in the northeastern part of the remnant. Recently Hewitt, Yusef-Zadeh, \\& Wardle (2008) not only found that this compact maser also has faint emission at $85\\km\\ps$, but also detected extended OH maser emission at the velocity of $\\sim85\\km\\ps$ with the Green Bank Telescope observation and the VLA archival observation toward the southern bright radio shell. The different LSR systemic velocities imply the MCs at different distances that may be impacted by the SNR shock wave. Therefore investigation is needed to clarify at which systemic velocity the maser emission is the product of the \\snr\\ shock interaction.   \\snr\\ has an irregular X-ray morphology, as observed by {\\sl ROSAT} and {\\sl Einstein}, inside an incomplete radio shell (Seward 1990; Yusef-Zadeh et al.\\ 2003). The {\\sl Spitzer} Infrared Array Camera (IRAC) mid-infrared observation toward \\snr\\ shows an arc at 4.5$\\mu$m in the same location as the southeastern radio shell (Reach et al.\\ 2006). The extended OH emission along the southern incomplete radio and Infrared (IR) shell hints an interaction of the SNR with the dense molecular gas in the south. If the SNR/MC association is established, a big progress can be made toward resolving the open question of the disparate velocity maser components and the distance to Kes~69 can also be determined.  Motivated by the supposed association of \\snr\\ with molecular gas, we have performed millimeter CO and \\HCOp\\ observations toward this remnant. The observations and results are described in \\S2 and \\S3, and the conclusion is summarized in \\S4. ",
        "conclusions": "We have performed a millimeter observation in CO and \\HCOp\\ lines toward SNR~\\snr. From the northeastern compact 1720~MHz OH maser region, the \\twCO\\ and \\thCO\\ emission's peaks around $65\\km\\ps$ and $85\\km\\ps$, which are consistent with the masers' LSR velocities are detected. In the southeast, a molecular (\\twCO\\ J=1--0) arc is revealed at 77--$86\\km\\ps$, well coincident with the partial SNR shell detected in the 1.4~GHz radio continuum and mid-IR observations. An $85\\km\\ps$ \\HCOp\\ emission is found to arise from a radio peak on the shell. Both the molecular arc and the \\HCOp\\ emission at $\\sim85\\km\\ps$ seem to be consistent with the presence of the extended 1720 MHz OH emission along the southeastern boundary of \\snr. The morphology correspondence between the CO emission and other band emission of the \\snr\\ shell provides strong evidence for the association of the SNR with the $\\sim85\\km\\ps$ component of molecular gas. There is another section of molecular arc at 79--$86\\km\\ps$ in the northwest.  Both the molecular arcs, together with the faint northern radio features, seem to be distributed along a circle of radius $8'.7$. The multiwavelength emissions along the southeastern shell can be explained by the impact of the SNR shock on a dense, clumpy patch of molecular gas. This pre-existing gas is likely to be a part of the cooled, clumpy debris of the interstellar molecular gas swept up by the progenitor's stellar wind.  The association between SNR~\\snr\\ and the MC at the systemic velocity $\\sim85\\km\\ps$ enables us to place the SNR at a kinematic distance of 5.2~kpc. \\\\  {\\sl Note added in proof.} In another observation in CO lines toward SNR Kes 75, we similarly discover a molecular shell, a part of which follows the bright partial SNR shell seen in X-rays, mid-IR, and radio continuum, and provide effective evidence for the association between Kes 75 and the adjacent MC (Su et al. 2009). Molecular shells are probably common in a number of SNRs but have rarely been studied."
    },
    "0809/0809.1523_arXiv.txt": {
        "abstract": "Light dark matter annihilating into electron-positron pairs emits a significant amount of internal bremsstrahlung that may contribute to the cosmic gamma-ray background. The amount of emitted gamma-rays depends on the dark matter clumping factor. {Recent calculations indicate that this value should be of order $10^6-10^7$. That allows us to calculate the expected gamma-ray background contribution from dark matter annihilation. We find that the light dark matter model can be ruled out if a constant thermally-averaged cross section is assumed ({s}-wave annihilation). For more massive dark matter candidates like neutralinos, however, cosmic constraints are weaker.} ",
        "introduction": "{Recent observations of the Galactic center in different frequencies have provided increasing motivation for theoretical models that consider dark matter annihilation and/or decay. One finds an excess of GeV photons \\citep{deBoer05}, of microwave photons \\citep{Hooper07}, of positrons \\citep{Cirelli08} and of MeV photons \\citep{Jean06, Weidenspointner06}{, which correlates with the Galactic bulge instead of the disk and cannot be attributed to single sources. It is therefore controversial whether it can be explained by conventional astrophysical sources} alone or if dark matter annihilation models are required{. It is not clear whether these different phenomena are related. Hence, their interpretation is still under discussion \\citep{deBoer08}.} For this reason, it is particularly interesting to also consider constraints which are independent of the Galactic center observations. % }  {Observations in the MeV energy range favor models based on light dark matter. These models suggest that light dark matter particles with masses of {$1-100$~MeV} annihilate into electron-positron pairs \\citep{BoehmHooper04}. In such a scenario, the dark matter mass needs to be larger than $511$~keV {to be able to produce electron-positron pairs through dark matter annihilation}. {It should be smaller than $100$~MeV, as otherwise {pion final states that produce too many gamma rays would be possible} \\citep{BoehmHooper04}.} Electromagnetic radiative corrections to the annihilation process require the emission of internal bremsstrahlung \\citep{Beacom05}. {It has also been proposed that such internal bremsstrahlung emission might explain the observed gamma-ray background in the $10-20$~MeV range \\citep{Ahn05b, Ahn05a} for dark matter masses of $\\sim20$~MeV, though these models appear less favorable in light of stronger upper limits on the dark matter particle mass \\citep{Beacom06, Sizun06}. Supernovae data require dark matter particle masses larger than $10$~MeV, although this limit depends on assumptions made regarding the scattering cross section between dark matter particles and neutrinos \\citep{Fayet06}. }  {Calculations regarding the annihilation of light dark matter in the Milky Way show that predicted and observed fluxes are only in agreement for {p}-wave annihilation models \\citep{BoehmHooper04}. In addition, dark matter models can be constrained by the cosmic backgrounds \\citep{Ahn05c, Ando05, Rasera06, Yuksel07, Mack08, SchleicherBanerjee08c}. Such constraints provide a highly complementary approach {based on the observed cosmic gamma-ray background}. In this letter, we show that such constraints provide a strong independent confirmation that {s}-wave annihilation of light dark matter is ruled out. } ",
        "conclusions": "{As shown by \\citet{Beacom06}, light dark matter scenarios that explain observations of the Galactic center require a mass less than $3$~MeV. {This results from a comparison between $511$~keV emission and higher-energy gamma-rays and does not require assumptions regarding the dark matter profile. For low} masses, clumping factors $C_0$ in the range $10^6-10^7$, which we would expect for the currently-favored Einasto profiles, are excluded due to the gamma-ray background constraints. Even if we adopted clumping factors corresponding to a flat Burket profile with little substructure, corresponding to values $C_0\\ge10^5$ \\citep{Cumberbatch08}, this would violate the constraints from the gamma-ray background. This shows that the light dark matter model can no longer be maintained with a constant thermally-averaged annihilation cross section if constraints from the cosmic background {are combined with the upper mass limits from the spectral constraints{, thus providing an independent confirmation that s-wave annihilation of light dark matter can be ruled out.}}   {These conclusions hold even if we adopted the very conservative upper mass limit of \\citet{Sizun06}, which is $7.5$~MeV if the ISM in the Galactic center would be strongly ionized. This appears unlikely and provides a very strong upper limit. In addition, more recent INTEGRAL observations favor an even smaller emission region, which increases the inflight annihilation intensity by a factor of $1.7$ and pushes the allowed maximum positron injection energy even further down (John Beacom, private communication). }  {Still feasible are models based on {p}-wave annihilation, in which the annihilation cross section becomes much smaller at late times when the velocities of dark matter particles are no longer relativistic \\citep{BoehmFayet04}. It is also interesting to note that the gamma-ray background at $10-20$~MeV might be explained by a power-law component of non-thermal electrons in active galactic nuclei (AGN) \\citep{Inoue08}. It is however unclear whether a sufficient number of non-thermal electrons is actually available. We expect that the FERMI satellite \\footnote{http://www.nasa.gov/mission\\_pages/GLAST/science/index.html} will shed more light on such questions, as the anisotropic distribution of the gamma-ray background may allow one to distinguish between astrophysical sources and dark matter annihilation \\citep{Ando07}.}  {We finally note that we have checked if similar constraints apply for more massive dark matter candidates, based on the recent analysis of \\citet{Mack08} and \\citet{Yuksel07}. For such models, however, clumping factors in the range $10^7-10^9$ are still feasible and well within the allowed parameter space. {The upper limits derived for light dark matter do not apply here, as the annihilation products may be different.} }  { We finally note that if the clumpiness of dark matter was high already at early times, annihilation of light dark matter may provide a significant contribution to the observed reionization optical depth \\citep{SchleicherBanerjee08a, Chuzhoy08, Natarajan08}. The detailed effects of such annihilations on the IGM have been explored by \\citet{Ripamonti07}, and consequences for $21$~cm observations have been explored by \\citet{Furlanetto06, Chuzhoy08}.}"
    },
    "0809/0809.4656_arXiv.txt": {
        "abstract": "We report {\\it Suzaku} observations of the active, Compton-thick Circinus galaxy. Observations were obtained with both the X-ray Imaging spectrometer (XIS) and the Hard X-ray Detector (HXD). Below 10 keV, the nuclear spectrum is dominated by radiation reflected from cold dense gas of high column density, while above 13~keV the radiation is directly transmitted nuclear emission seen through a column density of $\\simeq 4 \\times 10^{24}$~\\pcmsq. In the 0.2--10~keV band, the XIS spectrum is contaminated at 5\\% level by the brightest off-nuclear source in Circinus (CG X-1), but drops to 1\\% in the 5-10~keV and is negligible at higher energies. We find no significant evidence for variability in the hard ($>12$~keV) emission. The Circinus is marginally detected with the HXD/GSO in the 50--100~keV band at 2.5$\\sigma$ level. We model the 3-70~keV band XIS+PIN spectra with a four components: the Compton transmitted nuclear emission, the reflected nuclear emission, a soft power law (representing a combination of scattered nuclear emission, extended emission and contamination by sources in the galaxy below a few keV). The hard nuclear power-law is found to have a photon index $\\Gamma_h \\simeq 1.6$, very similar to the soft power-law. The high energy cut-off is $E_C \\simeq 49$~keV. These results agree with those from BeppoSax. An extrapolation of this model up to the GSO band shows good agreement with the GSO spectrum and supports our detection of the Circinus up to $\\simeq 100$~keV. ",
        "introduction": "The Circinus galaxy is a large, nearby ($4\\pm 1$~Mpc; \\markcite{freeman77}{Freeman} {et~al.} 1977) galaxy that harbors both a circumnuclear star burst and a Seyfert 2 nucleus.  Evidence for an obscured Seyfert 1 nucleus is provided by the finding of a  broad (FWHM $\\simeq$ 3000 \\kmps)  H$\\alpha$ line component in polarized light \\markcite{oliva98}({Oliva} {et~al.} 1998). The picture of an obscured Seyfert 1 nucleus is also supported by both the discovery of highly ionized gas extending along the minor axis of the galaxy, with a morphology that is reminiscent of the ionization cones seen in other Seyfert galaxies \\markcite{marconi94, wilson00}({Marconi} {et~al.} 1994; {Wilson} {et~al.} 2000), and direct X-ray detection of the nucleus through a column density of $\\sim 4 \\times 10^{24}$~\\pcmsq (\\markcite{matt99}{Matt} {et~al.} 1999, hereafter M99; \\markcite{soldi05}{Soldi} {et~al.} 2005).  {\\it ASCA}, {\\it BeppoSax} and {\\it XMM-Newton}  spectra  below 10 keV show both a flat (hard) spectrum and high equivalent width Fe K$\\alpha$ emission, characteristic of Compton reflection from cold gas illuminated by a power-law continuum. A single scattering Compton shoulder was found in {\\it Chandra} HETG observations \\markcite{bianchi02}({Bianchi} {et~al.} 2002) and was confirmed by {\\it XMM-Newton} \\markcite{molendi03}({Molendi}, {Bianchi}, \\& {Matt} 2003), demonstrating conclusively the existence of Compton thick matter.   Because of their very flat hard X-ray spectra, Compton-thick AGNs are a key ingredient in the synthetic model of the hard cosmic X-ray background \\markcite{setti89, madau94}({Setti} \\& {Woltjer} 1989; {Madau}, {Ghisellini}, \\& {Fabian} 1994), which peaks at 30--40~keV and rolls over at higher energies. After  NGC~4945 and  NGC~1068, Circinus galaxy is the third brightest Compton-thick AGN that allows detailed study.  In this paper, we present a broad band observation of the Circinus galaxy using the {\\it Suzaku} telescope. The low background and high sensitivity of {\\it Suzaku} allow improved  constraints on the nuclear continuum and the nature of the absorbing and reflecting matter. In \\S~\\ref{data}, we describe the observation and data reduction. The results are presented in \\S~\\ref{results}.  We discuss the results and their implications in \\S~\\ref{concl}. ",
        "conclusions": "} Our main results from an X-ray observation of the Circinus galaxy with the {\\it Suzaku} satellite are:  \\begin{enumerate} \\item There is minimal contamination ($\\pm 5\\%$) of the 0.2--10~keV radiation by the periodic ULX CG X-1, and less at higher energies. \\item Below 10~keV, we see a soft power-law ( that could come from a combination of extended emission, contaminating point sources in the galaxy, or the Thompson scattered nuclear emission) and a strong cold reflection component (as observed by earlier satellites) from the nuclear emission. \\item The spectrum above 20~keV is well described by a heavily absorbed Compton thick ($\\rm{N_H} = (4-5) \\times 10^{24}$~\\pcmsq) power-law model with a high energy cut-off ($E_C \\simeq 49$~keV). Based on the lack of variability in the PIN light curve, we suggest that the nucleus is obscured by a torus with a large covering fraction, with the reflection component originating from the far side of the  torus. \\item The galaxy is detected in the 50--200~keV band at 2.5$\\sigma$ level. This flux agrees well with an extrapolation to the higher energies of the transmitted component seen in the 13--50~keV band. \\item We have detected  strong emission lines at energies consistent with the florescent lines of neutral Fe K$\\alpha$, K$\\beta$, and  Ni K$\\alpha$. These features are consistent with the reflection from cold dense gas in the torus. However, we were not able to find the Compton shoulder in our XIS spectra. We point out this non-detection can be attribute to a calibration problem. \\item We also found emission features which are consistent with emission lines from He-like Fe~K$\\alpha$ (6.71~keV),  H-like Fe~K$\\beta$ (8.2~keV). The emission line at 7.04~keV  is also consistent with the line energy of H-like Fe K$\\alpha$.  These suggest that highly ionized gas is present in the Circinus galaxy. \\end{enumerate}"
    },
    "0809/0809.3240_arXiv.txt": {
        "abstract": "In both the Sun and the early universe, the \\he3($\\alpha$,$\\gamma$)\\be7 reaction plays a key role.  The rate of this reaction is the least certain nuclear input needed to calculate both the primordial \\li7 abundance in big bang nucleosynthesis (BBN) and the solar neutrino flux.  Taking advantage of several recent highly precise experiments, we analyse modern \\he3($\\alpha$,$\\gamma$)\\be7 data using a robust and minimally model dependent approach capable of handling discrepant data sets dominated by systematic rather than statistical errors. We find $S_{34}(0)=0.580\\pm0.043(0.054)$~keV~b at the 68.3(95.4)\\% confidence level. ",
        "introduction": "\\label{sect:introduction}  As the nearest star, the Sun is the best studied.  Models predict the neutrino fluxes produced by the radioactive decay of \\be7 and \\b8 in the solar core \\cite{bahcall2}.  Over the last decade, various neutrino observatories have measured the flux and flavour composition of the \\b8 neutrinos from the Sun ~\\cite{Super-K,SNO}, allowing constraints to be placed on neutrino mixing angles and mass-squared differences.  New experiments sensitive to the \\be7 flux~\\cite{borexino} are ongoing.  Both the \\be7 and \\b8 neutrino flux predictions are nearly directly proportional to the astrophysical $S$ factor for the $\\he3(\\alpha,\\gamma)\\be7$ reaction, $S_{34},$ at a relative energy of $\\sim 20$~keV; in fact, $\\phi_\\nu(\\be7)\\propto S_{34}(0)^{0.86}$ and $\\phi_\\nu(\\b8)\\propto S_{34}(0)^{0.81}$~\\cite{bahcall1,adelberger}. This sensitivity allowed $S_{34}$ to be constrained using the measured solar neutrino fluxes and laboratory measurements of $S_{17}$, the astrophysical $S$ factor for the $\\be7(p,\\gamma)^8$B reaction \\cite{cfo4}.  The \\he3($\\alpha$,$\\gamma$)\\be7 reaction is important not only in determining the solar neutrino fluxes, but also in other astrophysical environs.  Primordial nucleosynthesis has been a robust prediction of hot big bang cosmology for over 40 years~\\cite{wagoner67,wssok,osw,fo06,fs06}. The theory explains the large universal abundance of \\he4 as well as the origin of trace quantities of the light isotopes D, \\he3, and \\li7.  It is a theory with three free parameters, the cosmic baryon density, the neutron mean lifetime, and the number of active light neutrino species. Via fits to the standard electroweak theory, measurements at the Large Electron Positron Collider have determined the number of active light neutrino species to be $N_\\nu=2.984\\pm0.008$~\\cite{pdg08}, justifying the choice $N_\\nu = 3$.  Many experiments have been performed to determine the mean lifetime of the neutron; the current Particle Data Group recommendation is $\\tau_n = 885.7\\pm0.8$ sec~\\cite{pdg08}.  The cosmic baryon density has been determined by analysing anisotropies in the cosmic microwave background radiation; the latest WMAP satellite results are $\\Omega_bh^2=0.02273\\pm0.00062$~\\cite{wmap5}, where $\\Omega_b$ is the universal mean density of baryons expressed in units of the critical density required to close the universe and the Hubble constant is $100h$~km~sec$^{-1}$~Mpc$^{-1}$.  Using the Friedmann-Robertson-Walker cosmological model, the standard model of particle physics, and nuclear reaction rates one can therefore predict the light element abundances very precisely and compare directly with primordial abundance estimates based on observations of the oldest systems.  When observations of the Li abundances in the oldest stars in the Milky Way were compared with the predictions of BBN, a 2-3$\\sigma$ discrepancy was found \\cite{cyburt04}. Several possibilities exist for resolving this discrepancy, including particle physics beyond the standard model, improvements to the stellar models used to interpret astronomical observations, better approximations of the matter distribution in relativistic cosmology \\cite{wiltshire07,leith08}, modifications of gravitational theory (e.g., \\cite{bekenstein04}), and improved nuclear reaction rates. In this paper, we examine the last possibility. Lithium is made as beryllium in the early universe via the $\\he3(\\alpha,\\gamma)\\be7$ reaction.  Of course the rates of reactions that destroy \\be7 must be known in addition to those that create it. Recent studies of the $\\be7(d,p)2\\alpha$ reaction cross section suggest that this reaction is not the source of the discrepancy \\cite{coc04,angulo05}. The primordial \\li7 abundance prediction is nearly directly proportional to the $\\he3(\\alpha,\\gamma)\\be7$ cross section at a relative energy of $\\sim 300$~keV; the primordial abundance ratio \\li7/H $\\propto S_{34}^{0.96}$~\\cite{cyburt04}.  On account of the obvious importance of this reaction for both solar neutrinos and BBN, it has been subjected to extensive theoretical as well as experimental study.  Since the first potential model calculations \\cite{christy61,tombrello63}, many theoretical studies of the \\he3($\\alpha,\\gamma$)\\be7 reaction have been performed \\cite{kajino84,buck85,kajino86,langanke86,kajino87,mohr93,igmanov97,csoto00,baye00}.  Examining the potential models, a hard sphere potential yields $S_{34}(0)=0.47$ keV b and $S^\\prime(0)/S(0)=-0.60$ MeV$^{-1}$~\\cite{tombrello63}, while a more physical potential yields $S_{34}(0)=0.46$ keV b and $S^\\prime(0)/S(0)=-0.79$ MeV$^{-1}$~\\cite{baye00}.  Though the zero-energy $S$ factor values are in good agreement, the shapes  as measured by the logarithmic derivatives are quite different. Similarly, in the single channel resonating group method (RGM) calculations of Ref.\\ \\cite{kajino86}, $S(0)$ varies widely with the nucleon-nucleon potential. The most commonly used fitting function for experimental data is the microscopic cluster model calculation of Kajino \\cite{kajino84,kajino86,kajino87}. Ref.\\ \\cite{kajino86} reports $S_{34}(0)=0.50\\pm0.03$ keV~b and $S^\\prime(0)/S(0)=-0.548\\pm0.033$ MeV$^{-1}$.  The estimated theoretical uncertainty of $\\pm$6\\%~\\cite{kajino86} on both of these quantities must be considered when fitting this predicted shape to experimental data. The commonly cited potential model and microscopic cluster model calculations are shown in Figure \\ref{fig:thry}. Another RGM calculation that includes the \\li6+p configuration in addition to the \\he3+$\\alpha$ configuration \\cite{csoto00} finds that the energy dependence of the astrophysical $S$ factor is not uniquely determined and yields a range of $S^\\prime(0)/S(0)$ from $-0.70$ to $-0.50$ MeV$^{-1}$. Analyses that renormalize theoretically calculated $S$ factors must include an additional systematic error in order to account for uncertainties in the true shape of the $S$ factor when extrapolating to zero energy. It would be desirable to use the data themselves to determine the shape of the $S$ factor.  \\begin{figure} \\includegraphics[width=\\linewidth]{S34_4} \\caption{Theoretical $S_{34}$ calculations by Tombrello and Parker~\\cite{tombrello63} (solid) and by Kajino~\\cite{kajino87} (dashed) are plotted relative to their values at $E=0$ in the upper graph.  The deviations relative to the Kajino calculation are plotted in the lower graph. The dotted lines delineate the uncertainty due to the $\\pm$ 6\\% theoretical uncertainty in the zero-energy $S$ factor and logarithmic derivative estimated in Ref. \\cite{kajino86}.} \\label{fig:thry} \\end{figure}   In this paper, we evaluate the modern prompt capture $\\gamma$ ray and induced \\be7 activity measurements of the  $\\he3(\\alpha,\\gamma)\\be7$ cross section in a nearly model-independent way.  We use well known physics to constrain the energy dependence of the low energy cross section and disuss the treatment of systematic errors and their propagation into the final uncertainties, providing reliable values for $S_{34}$ at energies relevant to solar neutrino production and BBN with realistic uncertainties derived from Markov Chain Monte Carlo calculations. ",
        "conclusions": "\\label{sect:conclusions}  We have performed a minimally model dependent analysis of modern \\he3($\\alpha$,$\\gamma$)\\be7 data. This analysis properly takes into account the systematic errors of potentially discrepant data and yields reliable central values and uncertainties. At the 68.3\\% confidence level, we find $S_{34}(0)=0.580\\pm 0.043$ keV~b and $S^\\prime(0)/S(0) = -0.92\\pm0.18$ MeV$^{-1}$. We have used this value of $S_{34}(0)$ to compute the thermally averaged rate of the \\he3($\\alpha$,$\\gamma$)\\be7 reaction per particle pair, and have fit this rate with an analytical form accurate within 0.5\\%.  This reaction rate is about 9\\% higher than the recommendation of \\cite{adelberger}, but is quite consistent with it. It is also perfectly consistent with the recommendations of References \\cite{nacre,de04}. Moreover, access to precise modern data has allowed us to minimize the uncertainties associated with extrapolation and arrive at a more precise and better founded value that does not depend on nuclear structure models.  Our recommended reaction rate implies that the predicted \\be7 and \\b8 solar neutrino fluxes should both be increased by 8\\% compared with predictions based on the $S_{34}(0)$ value from Ref.\\ \\cite{adelberger}. Similarly, the present analysis implies that the value of $S_{34}$ in the energy range relevant to big bang nucleosynthesis is 17\\% larger than the value adopted in Ref.\\ \\cite{cyburt04}, resulting in a 17\\% increase in the predicted primordial \\li7 abundance. The uncertainties found here reduce the total error in the \\li7 abundance prediction by a factor of $\\sim2$. The discrepancy between the predicted and observed primordial \\li7 abundances increases from $\\sim3\\sigma$ to $\\sim5\\sigma$~\\cite{cfo5}, implying a rather serious disagreement that must be resolved.  In order to improve our knowledge of the rate of the \\he3($\\alpha$,$\\gamma$)\\be7 reaction, future experiments should collect data at at least six different energies. No modern data exist in the Gamow window for big bang nucleosynthesis, so this would be of primary interest. Information on capture into the ground and first excited states would also be useful, arguing in favour of prompt $\\gamma$ ray measurements rather than those based on induced \\be7 activity. Finally, precise measurements from 1 - 3 MeV would allow better constraints on the shape of the $S$ factor at the low energies of astrophysical interest."
    },
    "0809/0809.3289_arXiv.txt": {
        "abstract": "We present results of a search for a young stellar moving group associated with the star HD 141569, a nearby, isolated Herbig AeBe primary member of a 5$\\pm$3 Myr-old triple star system on the outskirts of the Sco-Cen complex.  Our spectroscopic survey identified a population of 21 Li-rich, $\\lesssim$30 Myr-old stars within 30\\degr\\ of HD 141569 which possess similar proper motions with the star.  The spatial distribution of these Li-rich stars, however, is not suggestive of a moving group associated with the HD 141569 triplet, but rather this sample appears cospatial with Upper Scorpius and Upper Centaurus Lupus.  We apply a modified moving cluster parallax method to compare the kinematics of these youthful stars with Upper Scorpius and Upper Centaurus Lupus.  Eight new potential members of Upper Scorpius and five new potential members of Upper Centaurus Lupus are identified.  A substantial moving group with an identifiable nucleus within 15\\degr ($\\sim$30 pc) of HD 141569 is not found in this sample. Evidently, the HD 141569 system formed $\\sim$5 Myr ago in relative isolation, tens of parsecs away from the recent sites of star formation in the Ophiucus-Scorpius-Centaurus region. ",
        "introduction": "The HD 141569 stellar system initially garnered the interest of \\citet{Rossiter:1943} as a potential triple star system.  The existence of this group was confirmed over half a century later when \\citet{Weinberger:2000} showed that the stars' position angles relative to one another had not changed since Rossiter's 1938 observations, thus confirming common proper motions.  They also noted that all three stars are consistent with being the same age, indicating a comoving and coeval system.  This triumvirate is particularly interesting because all components are quite young; the system was found to be 5$\\pm$3 Myr \\citep{Weinberger:2000}, and recently the A star's age was constrained using surface gravities and effective temperatures to be $\\sim$4.7 Myr \\citep{Merin:2004}. HD 141569 itself is near enough to resolve its large, dusty disk, and near-infrared signatures indicate perturbations to the disk's structure which could be explained by planet formation \\citep{Weinberger:1999}. Were a coherent group present, it would therefore be distinctly possible to similarly observe additional young, disk-bearing stars.  Furthermore, surveying a coeval sample at this distance would be useful for determining disk frequency and mechanisms by which they arise and evolve.  In some cases, seemingly isolated young stars have later been found to have an entourage of other low-mass young stars which constitute a stellar association \\citep[e.g. HD 104237, HR 4796, $\\beta$ Pictoris; cf.][]{Mamajek:2002,Li:2005}. The lower mass members of these associations have been useful, indeed critical, for estimating the age of the massive star in question, as often the massive star is on the main sequence and uncertainty in its HR diagram position is substantial enough to impede an isochronal age estimate.  Any additional association members should be identifiable by their similar space motions.  The youth of these objects makes them useful for studies of circumstellar disks and stellar evolution.  In this region of sky rich with already identified associations such as $\\rho$ Ophiucus, Upper Scorpius (US) and Upper Centaurus Lupus (UCL), placing an ``HD 141569 Association'' into context with surrounding moving groups could aid in the understanding of star formation histories in giant molecular clouds.   So motivated, in this paper we seek to identify a sample of stars associated with HD 141569.  Previous studies \\citep[e.g.][]{Mamajek:2002,Li:2005} have shown that catalog searches based on X-ray activity, proper motions, and distance criteria can yield new members of associations. Here we also identify candidate members with a catalog search for X-ray active stars in the vicinity of HD 141569 possessing proper motions consistent with HD 141569 (\\S\\ref{targets}). From new spectroscopic observations (\\S\\ref{obs}), in \\S\\ref{analysis} we identify a subset of 21 stars which we claim as youthful using Li {\\sc I} equivalent widths; furthermore, we also note the presence of H$\\alpha$ emission as an interesting quantity usually indicative of chromospheric activity. Lacking direct distance measurements for these stars, we derive distances using a modified version of the traditional moving cluster parallax method \\citep{deBruijne:1999, deZeeuw:1999, Mamajek:2005}, with which we further determine the probability that the sample stars are comoving with HD 141569 (\\S\\ref{mcp}).  Placing these stars on a Hertzsprung-Russell diagram (hereafter HRD or HR diagram, \\S\\ref{sec_hrd}), we find that the stars' distances---and hence ages---are most consistent with their strong lithium abundances when we derive their moving cluster parallaxes assuming comovement with Upper Scorpius or UCL associations of young stars. In \\S\\ref{summconc} we summarize our study and present thirteen newly identified Sco-Cen members. ",
        "conclusions": "\\label{summconc}  We have identified a group of 21 PMS stars within 30\\degr\\ of HD 141569 on the basis of strong Li {\\sc I} absorption and, in 9 of those 21 cases, H$\\alpha$ in emission. These stars were selected through a joint catalog search for X-ray sources with spatial and proper-motion characteristics similar to those of HD 141569, a B9.5Ve star at 116 pc that harbors a circumstellar disk and for which two low-mass companions had previously been identified \\citep{Weinberger:2000}.  For these 21 stars, we have applied a moving cluster parallax technique to proper-motion data from the literature.  Table \\ref{Membership} outlines our final membership assessments: we present eight potential new members of Upper Sco and five new potential members of UCL.  These stars possess lithium presence consistent with youth and futhermore appear youthful on the HR diagram.  Primarily we utilize spatial position as the principal criterion for determining membership and supplement that with comovement probabilties from the moving cluster parallax derivation.  Additionally, we examine the motion of the HD 141569 system away from the galactic midplane over its lifetime.  The system, surprisingly, appears to have formed in isolation, well outside of presently known star forming regions and molecular clouds."
    },
    "0809/0809.1245_arXiv.txt": {
        "abstract": "{stars: evolution --  supernovae: general -- stars: Wolf-Rayet -- stars: supergiants}  The study of the stars that explode as supernovae used to be a forensic study, working backwards from the remnants of the star. This changed in 1987 when the first progenitor star was identified in pre-explosion images. Currently there are 8 detected progenitors with another 21 non-detections, for which only a limit on the pre-explosion luminosity can be placed. This new avenue of supernova research has led to many interesting conclusions, most importantly that the progenitors of the most common supernovae, type IIP, are red supergiants as theory has long predicted. However no progenitors have been detected thus far for the hydrogen-free type Ib/c supernovae which, given the expected progenitors, is an unlikely result. Also observations have begun to show evidence that luminous blue variables, which are among the most massive stars, may directly explode as supernovae. These results contradict current stellar evolution theory. This suggests that we may need to update our understanding. ",
        "introduction": "The term supernovae was first introduced by Baade \\& Zwicky (1935) when they separated such events from \\textit{common} novae.  Novae occur frequently in our galaxy and are the result of hydrogen accreting on to the surface of a white dwarf star that ignites after a thick enough layer is deposited. The \\textit{super}-novae are far more luminous and rare, only occurring once or twice a century in our own Galaxy. The last supernovae observed in our Galaxy were the Tycho and Kepler Supernovae in 1572 and 1604 respectively.  Supernovae can be so luminous that they outshine all the other stars in a galaxy. After a sharp rise to maximum light over a few days their luminosity decays and slowly fades over weeks and months. Some have a light curve plateau and remain at a constant brightness for up to four  months before fading. The energy required to power a typical supernova display is similar to the amount of energy our Sun will output over its 10 billion year lifetime. To produce that amount of energy over a few days the progenitor star must be significantly altered in a supernova.  The modern understanding of supernovae divides them into two main types, thermonuclear and core-collapse. Thermonuclear supernovae, (also know as type Ia supernovae) are the detonation of carbon-oxygen white dwarfs. We do not discuss them in this review (see Hillebrandt \\& Niemeyer 2000). Core-collapse supernovae account for 72 percent of all supernovae in a volume limited sample (Smartt et al. 2008). They are the final events in the lives of stars more massive than approximately 8 solar masses ($8M_{\\odot}$). The basic evolution of a star is relatively well understood. A star begins on the main sequence, burning hydrogen to helium. Here it spends the largest fraction of its lifetime. Once a helium core is formed, hydrogen continues to burn in a shell around the core and the stellar radius swells to between 100 and 1000 times that of the Sun. At this stage the star becomes a red giant or, for the most luminous, a red supergiant. Massive stars undergo further burning stages. Helium burns to form carbon and oxygen and then progressively heavier elements burn, preventing stellar collapse until an iron core is formed. Iron fusion is an endothermic reaction. With no further energy source to prevent the core from collapsing a neutron star or a black hole is formed. Figure \\ref{single} shows the evolution of a star in schematic form.  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=100mm]{jje_figure1.eps} \\caption{A massive star's life cycle. Stars begin their lives on the main sequence as compact blue stars. Once the helium core is formed the surrounding hydrogen envelope expands and the star becomes a red supergiant. Further burning stages occur until an iron core is formed. The star will explode as a type-II supernova. } \\label{single} \\end{center} \\end{figure}  Creating a neutron star releases a tremendous amount of energy in neutrinos. A fraction of these neutrinos interact with the stellar envelope, providing the energy to heat and eject it and thus give rise to the supernova. The ejected material is rich in heavy elements synthesized during the star's life and in the supernova. Eventually this material mixes with the interstellar medium and pollutes it. When future generations of stars form they will be more metal rich than the previous generation. Most of the heavy elements in our bodies was formed in the evolution of a massive star and its explosive end.  The exact mechanism of how the energy is transfered to the envelope is uncertain. Most current simulations of supernova do not produce explosions after the core becomes a neutron star. Different additional mechanisms such as acoustic driving by the proto-neutron star or jets from a magneto-rotational instability as material accretes on to the proto-neutron star have been suggested (see Burrows et al. 2007, Dessart et al. 2008 and references therein).  There are two exceptions to the standard picture of iron core collapse. Stars around 7 to $8M_{\\odot}$ are not massive enough to progress past carbon burning they therefore form oxygen-neon-magnesium cores supported by electron degeneracy pressure. If the core mass reaches the Chandrasekhar mass of $1.4M_{\\odot}$ it begins to collapse. Eventually the central density is high enough that electrons are captured by magnesium and neon. This removes the electrons supporting the core and accelerates the collapse to a neutron star and an electron-capture supernova occurs (Eldridge, Mattila \\& Smartt 2007; Poelarends et al. 2008). The other exception is the pair-production instability when photons in the core form electron-positron pairs reducing the radiation pressure that was supporting the star and hence leading to its collapse. The resulting evolution is complex and the entire star can be disrupted in an explosion (Heger et al. 2003). However such supernovae are rare, occurring only in the most metal poor galaxies (Langer et al. 2007).  The appearance of a supernova strongly depends on the structure and composition of the material ejected. This is determined by how much mass is lost from the surface of the star during its evolution. Stars of mass less than 25$M_{\\odot}$ have weak stellar winds and do not lose their hydrogen envelopes before an iron core is formed. Stars more massive than this have strong stellar winds and all hydrogen is lost from the surface before the star explodes. These stars are known as helium stars or Wolf-Rayet stars and Figure \\ref{single2} shows the structure of such a star. The hydrogen envelope can also be removed if the star is in a binary. The stars in a binary interact if their radius is similar to the radius of the orbit and mass-transfer can occur with one star losing mass and the other gaining mass.  \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=100mm]{jje_figure2.eps} \\caption{The details of the figure are similar to Figure \\ref{single}. However after the star becomes a red supergiant strong stellar winds (or binary interaction) remove the hydrogen envelope and the star becomes a Wolf-Rayet star. Once an iron core is formed the star explodes as a type-Ib/c supernova.} \\label{single2} \\end{center} \\end{figure}  Core-collapse supernova progenitors naturally separate into two main groups, those that contain hydrogen when they explode and those that do not. Supernovae are also classified by the same aspect. If hydrogen is detected in the supernova spectrum then it is classified as a type-II supernova otherwise it is a type-I. These two main types can be further subdivided based on the photometric and spectroscopic behaviour of the supernova (Filippenko 1997).  There are two types of type-I core-collapse supernovae: Type Ib with helium lines but no silicon lines in their spectra and Type Ic with neither helium nor silicon lines in their spectra (silicon lines indicate a thermonuclear type Ia supernova).  For type-II supernovae, if the luminosity is constant for a few months after the supernova appeared, a plateau in the light curve, then it is a type IIP. These are the most common type of supernovae, making up 59 percent of core-collapse supernovae in a volume limited survey (Smartt et al. 2008). The luminosity is constant because, as the ejecta expand in radius, the visible surface where the observed light is emitted from (the photosphere) moves inwards in mass and therefore remains stationary.  There are three other rarer subtypes of type-II supernovae. If the light curve decays linearly then the supernova is a type IIL. The progenitors of these stars are thought to have lost hydrogen from their envelopes so that there is not enough to produce the luminosity plateau. If the supernova has narrow hydrogen lines which indicates slow moving ejecta it is a type IIn. This occurs when the supernova ejecta encounter a large amount of material surrounding the progenitor star and are decelerated from a typical ejection velocity of $10^{5} {\\rm km\\,s^{-1}}$ to $1000 \\, {\\rm km\\,s^{-1}}$. The last subtype is the hybrid class, type IIb. These are type-II events which metamorphose into Ib supernovae. The progenitors of these supernovae have only the barest trace of hydrogen left at the time of core collapse.  The deduction of which stars produce which supernova was a series of educated guesses. It was thought that red supergiants produced type-IIP supernovae and that the other types were produced depending on the amount of mass lost before core collapse. The more mass that is lost from a star, the deeper the layers and the heavier elements that are exposed at the surface. This leads to the stellar type of the progenitor changing from a red supergiant to a Wolf-Rayet star and the supernova type progresses from: IIP $\\rightarrow$ IIL $\\rightarrow$ IIb $\\rightarrow$ Ib $\\rightarrow$ Ic. The amount of stripping depends on the initial mass of the star with the most massive stars losing the largest fraction of their initial mass and exposing the deepest interiors at explosion.  The only way to confirm whether this theory is correct is to study the star that exploded in a supernova. This is difficult because the star is destroyed in the supernova. It is possible however that an image may have been taken before the explosion but supernovae are rare. Also only a few galaxies are close enough for individual stars to be resolved in ground-based images. However the Supernovae, 1987A and 1993J, took place in nearby galaxies and it was possible to find the progenitor star in pre-explosion images and determine their nature (see Section 2).  The situation dramatically improved with the launch of the Hubble Space Telescope (HST). Its resolving power and sensitivity were such that it became possible to resolve individual stars in galaxies out to 60 million light-years rather than out to a few million light-years. This increased the number of galaxies observed in detail and for a few supernovae a year pre-explosion images of the progenitor should exist (rather than a few per century before HST). In 2003 two groups simultaneously found the first progenitor of a type IIP supernova (Van Dyk et al. 2003; Smartt et al. 2004). With the concept tested, many discoveries have followed. The success can be reflected in that, out of the 135 supernovae that occurred from 1999 to 2007 within range of HST, 29 had HST pre-explosion images of the supernova site. The results have confirmed some theories and caused problems for others.  In this article we review the study of supernova progenitors, beginning with the first two discovered. This is followed by discussion of the study of progenitors during the HST era, with highlights of the main results of the observations of type-IIP progenitors. We also discuss the progenitors for type-Ib/c supernovae and the growing evidence that type-IIn supernova have progenitors that challenge the preconceptions of stellar theorists. ",
        "conclusions": "The direct study of pre-supernova imaging in the archives has led to some quite remarkable and usually understated successes in the determination that the progenitors of type IIP supernovae are red supergiants. Despite this, it is still a subject in its infancy. The remaining supernova types are still lacking firm constraints. The importance of direct study is highlighted by the current confusion of type-IIn progenitors.  By detecting the progenitors of supernovae, we can make fundamental tests of stellar evolution that were not possible before we were able to observe their progenitors. Each time we see something new it provokes us to reach a new understanding. Our understanding of the lives of stars and the nature of supernovae is evolving dramatically because of this exciting avenue of science."
    },
    "0809/0809.1559_arXiv.txt": {
        "abstract": "{Theory suggests that about 10\\% of \\emph{Swift}-detected gamma-ray bursts (GRBs) will originate at redshifts, $z$, greater than 5 yet a number of high redshift candidates may be left unconfirmed due to the lack of measured redshifts.} {Here we introduce our code, \\emph{GRBz}, a method of simultaneous multi-parameter fitting of GRB afterglow optical and near infrared, spectral energy distributions. It allows for early determinations of the photometric redshift, spectral index and host extinction to be made.} {We assume that GRB afterglow spectra are well represented by a power-law decay and model the effects of absorption due to the Lyman forest and host extinction. We use a genetic algorithm-based routine to simultaneously fit the parameters of interest, and a Monte Carlo error analysis.} {We use GRBs of previously determined spectroscopic redshifts to prove our method, while also introducing new near infrared data of \\object{GRB\\,990510} which further constrains the value of the host extinction.} {Our method is effective in estimating the photometric redshift of GRBs, relatively unbiased by assumptions of the afterglow spectral index or the host galaxy extinction. Monte Carlo error analysis is required as the method of error estimate based on the optimum  population of the genetic algorithm underestimates errors significantly. } ",
        "introduction": "\\label{grbz:intro}   Theory suggests and observations agree that approximately 10\\% of \\emph{Swift}-detected gamma-ray bursts (GRBs) will originate at redshifts, $z \\gtrsim 5$ \\cite{bromm2006:ApJ642,jakobsson2006:A&A447} and about 3\\% at redshifts, $z \\gtrsim 6$ \\cite{daigne2006:MNRAS372}. Furthermore, due to the favourable $K$-correction at a fixed observer time and cosmological time dilation, the observed optical flux is not expected to fade significantly with increasing redshift. Hence GRBs should be detectable out to $z \\gtrsim 10$ \\cite{ciardi2000:ApJ540,lamb2000:ApJ536}, making them the most distant observable objects in the Universe. Yet a number of high redshift candidates may be left unconfirmed due to the lack of observed counterparts or lack of observations \\cite{ruiz2007:ApJ669}.  The most usual way of determining the redshift of a GRB is to spectroscopically measure the redshift of the afterglow at early times if it is bright enough, or to measure that of its host galaxy, either photometrically (e.g. \\citeNP{bolzonella2000:A&A363}) or spectroscopically, once the optical afterglow of the burst has dimmed sufficiently. Neither of these may be possible given the dimness of the afterglows \\cite{fynbo2007:Msngr130} and the host galaxies. In addition, programs to obtain such measurements are frequently only triggered if there is evidence to suggest a high redshift, or a burst of significant interest. There are a number of redshift indicators, or pseudo redshift methods, which rely on correlations between observable properties and redshift (e.g., \\citeNP{guidorzi2005:MNRAS364,amati2006:MNRAS372,pelangeon2006:GCN5004}) that offer approximate redshifts to varying degrees of success.   An alternative method of redshift determination, useful when the source is not bright enough for spectroscopic observations, is to photometrically measure the redshift of the afterglow itself, given enough simultaneous optical/near infrared data points.  Attempts at photometric redshift determination usually assume a spectral index and a value for host extinction, while fitting for redshift, or fit for a limited number of discrete values of spectral index and host extinction (e.g., \\citeNP{andersen2000:A&A364,jakobsson2006:A&A447}). Independently, many attempts have been made to fit host extinction and spectral index to similar data, when the redshift is previously known (e.g., \\citeNP{galama2001:ApJ549,stratta2004:ApJ608,kann2006:ApJ641,schady2007:MNRAS377,starling2007:ApJ661}). This is important as the properties of the circumburst medium and host galaxy can be used to constrain progenitor types and burst models themselves.  In this paper we present our code, \\emph{GRBz}, which, unlike other methods of photometric redshift determination, fits all three parameters simultaneously and allows for determinations of the photometric redshift to be made, relatively unbiased by assumptions of the spectral index or the host extinction. In section \\ref{grbz:method} we describe the method of modelling and fitting that we apply to the data described in section \\ref{grbz:data}, including our previously unpublished near IR data of GRB\\,990510. In section \\ref{grbz:results} we present our results and in section \\ref{grbz:discussion} we discuss the results as they apply to the fitting mechanism. We summarise our findings in section \\ref{grbz:conclusions}. ",
        "conclusions": "\\label{grbz:conclusions}   We have developed a method whereby early time photometric nIR, optical and UV data of the afterglow can be used to estimate the photometric redshift of GRBs.  This has advantages over the more usual methods of spectroscopic redshift determination, in that photometric observations do not require the source to be as bright, and over photometric estimates of the host galaxy as the afterglow is frequently much brighter than the host. In this implementation, we assume that GRB afterglow spectra are well represented by a power-law, and model the effects of absorption due to the Lyman forest and host extinction. We use a genetic algorithm-based routine to simultaneously fit the parameters of interest, relatively unbiased by assumptions of the spectral index or host extinction.  We introduce new nIR data for GRB\\,990510, which we have fitted along with the previously published data, to give new estimates of the host extinction and spectral index. We offer a new photometric redshift for GRB\\,050814, slightly higher than that previously suggested.  The Monte Carlo error analysis, though computationally time consuming, is required as the method of error estimate based on the optimum population width of the genetic algorithm underestimates errors significantly. As the distribution of the best fit parameters obtained via Monte Carlo analysis are not Gaussian, caution is required when interpreting the nominal 68\\% errors of the average parameters."
    },
    "0809/0809.1904_arXiv.txt": {
        "abstract": "{The energy spectrum of nuclear recoils in Weakly Interacting Massive Particle (WIMP) direct detection experiments depends on the underlying WIMP mass (strongly for light WIMPs, weakly for heavy WIMPs).  We discuss how the accuracy with which the WIMP mass could be determined by a single direct detection experiment depends on the detector configuration and the WIMP properties. In particular we  examine the effects of varying the underlying WIMP mass, the detector target nucleus, exposure, energy threshold and maximum energy, the local velocity distribution and the background event rate and spectrum.}  \\FullConference{Identification of dark matter 2008\\\\ August 18-22, 2008\\\\ Stockholm, Sweden}  \\begin{document} ",
        "introduction": "The direct detection of WIMPs in the lab would not only directly confirm the existence of dark matter but would also allow us to probe the WIMP properties, in particular its mass. This would shed light on its nature and probe extensions of the standard model of particle physics. Furthermore definitive detection of the WIMP may well require consistent signals (i.e. with the same inferred WIMP properties) from direct detection, indirect detection and collider experiments. Here we give a brief overview of recent work~\\cite{green} on determining the WIMP mass from direct detection experiments (see also Refs.~\\cite{ls,massall,ds}).  The differential event rate (number of events per unit energy, time and detector mass) has a roughly exponential energy dependence:~\\cite{ls} \\begin{equation} \\frac{{\\rm d}R}{{\\rm d} E} \\approx c_{1} F^2(E) \\left( \\frac{{\\rm d}R}{{\\rm d} E} \\right)_{0} \\exp{\\left(-\\frac{E}{c_{2} E_{\\rm R}}\\right)} \\,, \\end{equation} where $c_{1} $ and $c_{2}$ are fitting parameters of order unity (which depend on the target mass number and energy threshold), $({\\rm d} R/{\\rm d} E)_{0}$ is the event rate in the $E \\rightarrow 0 \\, {\\rm keV}$ limit and $F(E)$ is the form factor. The characteristic energy scale of the exponential, $E_{\\rm R}$, depends on the WIMP mass, $m_{\\chi}$, and is given by \\begin{equation} E_{\\rm R} = \\frac{ 2 m_{A} m_{\\chi}^2 v_{\\rm c}^2}{(m_{\\chi} + m_{A})^2} \\,, \\end{equation} where $m_{A}$ is the target nuclei mass and $v_{\\rm c}$ is the local circular speed. For light WIMPs ($m_{\\chi} \\ll m_{\\rm A}$) $E_{\\rm R} \\propto m_{\\chi}^2$, while for heavy WIMPs ($m_{\\chi} \\gg m_{\\rm A}$) $E_{\\rm R} \\sim {\\rm const}$. In other words, for light WIMPs the energy spectrum is strongly dependent on the WIMP mass while for heavy WIMPs the dependence on the WIMP mass is far weaker. Consequently it should be easier to measure the mass of light (compared with the target nuclei) WIMPs than heavy WIMPs. ",
        "conclusions": "\\label{results}  \\begin{figure} \\includegraphics[width=.8\\textwidth]{Greenfig1.eps} \\caption{The fractional deviation of the WIMP mass limits from the input mass, $(m_{\\chi}^{\\rm lim}-m_{\\chi}^{\\rm in})/m_{\\chi}^{\\rm in}$, for exposures ${\\cal E}= 3 \\times 10^{3}, 3 \\times 10^{4}$ and $3 \\times 10^{5} \\, {\\rm kg \\, day}$ and input cross-section $\\sigma_{\\rm p} = 10^{-8} \\, {\\rm pb}$ for the benchmark SuperCDMS like detector. The solid (dotted) lines are the 95\\% (68\\%) confidence limits.  } \\label{fig1} \\end{figure}    The accuracy with which the WIMP mass could be measured by the benchmark SuperCDMS~\\cite{SuperCDMS} like Ge detector described above is shown in Fig.~\\ref{fig1}. With exposures of ${\\cal E}= 3 \\times 10^{4}$ and $3 \\times 10^{5} \\, {\\rm kg \\, day}$ it would be possible to measure the mass of a light, $m_{\\chi} \\sim {\\cal O}(50 \\, {\\rm GeV})$, WIMP with an accuracy of roughly $25\\%$ and $10\\%$ respectively. For heavy WIMPs ($m_{\\chi} \\gg 100 \\, {\\rm GeV}$) even with a large exposure it will only be possible to place a lower limit on the mass.  For very light WIMPs, $m_{\\chi} < {\\cal O}(20 \\, {\\rm GeV})$, the number of events above the detector energy threshold would be too small to allow the mass to be measured accurately.  The number of events detected is directly proportional to both the exposure and the cross-section, therefore these quantities have the greatest bearing on the accuracy of the WIMP mass determination.   The energy threshold, $E_{\\rm th}$, and the maximum energy, $E_{\\rm max}$, above which recoils are not detected/analysed also affect the accuracy with which the WIMP mass can be determined. Increasing $E_{\\rm th}$ (or decreasing $E_{\\rm max}$) not only reduces the number of events detected, but also reduces the range of recoil energies and the accuracy with which the characteristic energy of the energy spectrum, $E_{\\rm R}$, and hence the WIMP mass, can be measured.  For light WIMPs the small $E_{\\rm R}$ means that the expected number of events decreases rapidly as the energy threshold is increased, while for heavy WIMPs the large $E_{\\rm R}$, and flatter energy spectrum, means that the smaller range of recoil energies reduces the accuracy with which $E_{\\rm R}$ can be measured. Reducing the maximum energy only has a significant effect for heavy WIMPs.  The WIMP and target mass dependence of $E_{\\rm R}$ suggests that heavy targets will be able to measure the mass of a heavy WIMP more accurately, however the rapid decrease of the nuclear form factor with increasing momentum transfer which occurs for heavy nuclei means that this is in fact not the case (see also Ref.~\\cite{ds}).  If the WIMP distribution on the ultra-local scales probed by direct detection experiments is smooth, then the $\\pm 20 \\, {\\rm km \\, s}^{-1}$ uncertainty in the local circular speed~\\cite{klb} leads to a $\\sim 10 \\%$ systematic error in the determination of $m_{\\chi}$. Changes in the detailed shape of the local velocity distribution lead to relatively small changes in the shape of the differential event rate~\\cite{drdens}, and hence a relatively small, ${\\cal O} (5 \\%)$, systematic uncertainty in the WIMP mass. If the ultra-local WIMP distribution consists of a finite number of streams, then the energy spectrum will consists of a number of steps. The positions of the steps will depend on the (unknown) stream velocities, as well as the target nuclei and WIMP masses. With multiple targets it would in principle be possible to constrain the WIMP mass without making any assumptions about the WIMP velocity distribution~\\cite{ds}.  Future experiments aim to have negligible backgrounds, however if the background rate is not negligible compared with the WIMP event rate it will be difficult to disentangle a WIMP signal (and the WIMP mass) from the background if the background spectrum has a similar shape to the WIMP spectrum (i.e. exponential background, or flat background and a heavy WIMP). The uncertainties from backgrounds could be mitigated by using multiple targets and/or using multiple scatter events to measure/constrain the background spectrum."
    },
    "0809/0809.0844_arXiv.txt": {
        "abstract": "The vacuum expectation values of the energy--momentum tensor and the fermionic condensate are analyzed for a massive spinor field obeying the MIT bag boundary condition on a cylindrical shell in the cosmic string spacetime. Both regions inside and outside the shell are considered. By applying to the corresponding mode-sums a variant of the generalized Abel--Plana formula, we explicitly extract the parts in the expectation values corresponding to the cosmic string geometry without boundaries. In this way the renormalization procedure is reduced to that for the boundary-free cosmic string spacetime. The parts induced by the cylindrical shell are presented in terms of integrals rapidly convergent for points away from the boundary. The behavior of the vacuum densities is investigated in various asymptotic regions of the parameters. In the limit of large values of the planar angle deficit, the boundary-induced expectation values are exponentially suppressed. As a special case, we discuss the fermionic vacuum densities for the cylindrical shell on the background of the Minkowski spacetime. ",
        "introduction": "It is well known that different types of topological objects may have been formed in the early universe after Planck time by the vacuum phase transition \\cite{Vile85}. Depending on the topology of the vacuum manifold these are domain walls, strings, monopoles and textures. Among them the cosmic strings are of special interest. Although the recent observational data on the cosmic microwave background radiation have ruled out cosmic strings as the primary source for primordial density perturbations, they are still candidates for the generation of a number of interesting physical effects such as the generation of gravitational waves, gamma ray bursts and high-energy cosmic rays (see, for instance, \\cite{Damo00}). Recently, cosmic strings attract a renewed interest partly because a variant of their formation mechanism is proposed in the framework of brane inflation \\cite% {Sara02}.  In the simplest theoretical model describing the infinite straight cosmic string the spacetime is locally flat except on the string where it has a delta shaped Riemann curvature tensor. In quantum field theory the corresponding non-trivial topology induces non-zero vacuum expectation values for physical observables. Explicit calculations for the geometry of a single cosmic string have been done for different fields \\cite{Hell86}-\\cite% {Spin08}. Vacuum polarization effects by higher-dimensional composite topological defects constituted by a cosmic string and global monopole are investigated in Refs. \\cite{Beze06Comp} for scalar and fermionic fields. Another type of vacuum polarization arises when boundaries are present. The imposed boundary conditions on quantum fields alter the zero-point fluctuations spectrum and result in additional shifts in the vacuum expectation values of physical quantities. This is the well-known Casimir effect (for a review see \\cite{Most97}). In Ref. \\cite{Beze06sc}, we have studied both types of sources for the polarization of the scalar vacuum, namely, a cylindrical boundary and a cosmic string, assuming that the boundary is coaxial with the string and that on this surface the scalar field obeys Robin boundary condition. For a massive scalar field with an arbitrary curvature coupling parameter we evaluated the Wightman function and the vacuum expectation values of the field squared and the energy-momentum tensor. The polarization of the electromagnetic vacuum by a conducting cylindrical shell in the cosmic string spacetime is investigated in \\cite{Beze07El} (For a combination of topological and boundary induced quantum effects in the gravitational field of a global monopole see Refs. \\cite{Saha03Mon,Saha04Mon,Beze06Mon}.)  Continuing in this line of investigation, in the present paper we analyze the polarization of the fermionic vacuum by a cylindrical shell coaxial with the cosmic string on which the field obeys MIT bag boundary condition. We evaluate the fermionic condensate and vacuum expectation values of the energy-momentum tensor in both interior and exterior regions of the shell. The renormalized vacuum expectation value of the energy-momentum tensor for a fermionic field in the geometry of a cosmic string without boundaries is investigated in \\cite{Frol87,Dowk87b,Line95,More95,Beze06}. In addition to describing the physical structure of the quantum field at a given point, the energy-momentum tensor acts as a source of gravity in the Einstein equations and plays an important role in modelling a self-consistent dynamics involving the gravitational field. In the problem under consideration all calculations can be performed in a closed form and it constitutes an example in which the topological and boundary-induced polarizations of the vacuum can be separated in different contributions.  From the point of view of the physics in the region outside the string, the geometry considered in the present paper can be viewed as a simplified model for the non-trivial core. This model presents a framework in which the influence of the finite core effects on physical processes in the vicinity of the cosmic string can be investigated. In particular, it enables to specify conditions under which the idealized model with the core of zero thickness can be used. The corresponding results may shed light \\ upon features of finite core effects in more realistic models, including those used for string-like defects in crystals and superfluid helium. In addition, the problem considered here is of interest as an example with combined topological and boundary induced quantum effects in which the vacuum characteristics such as energy density and stresses can be found in closed analytic form. From the results of the present paper, as a special case, we obtain the fermionic Casimir densities for a cylindrical shell with MIT boundary conditions in the Minkowski background (for the combined effects of a magnetic fluxon and MIT boundary conditions on the vacuum energy of a Dirac field see Refs. \\cite{Lese98}). Note that, in addition to traditional problems of quantum field theory under the presence of material boundaries, the Casimir effect for cylindrical geometries can also be important to the flux tube models of confinement \\cite{Fish87,Barb90} and for determining the structure of the vacuum state in interacting field theories \\cite{Ambj83}.  We have organized the paper as follows. In the next section, the eigenspinors for the region inside the cylindrical boundary are constructed and the eigenvalues of the corresponding quantum numbers are specified. These eigenspinors are the basis for the anlysis of the Casimir densities in the following sections. In section \\ref{sec:FermC}, by using a variant of the generalized Abel-Plana formula, we extract from the mode-sum of the fermionic condensate the part corresponding to the geometry of a cosmic string without the shell. The part induced by the shell is investigated in various asymptotic regions for the parameters. Section \\ref{sec:EMT} is devoted to the investigation of the boundary induced parts in the vacuum expectation value of the energy-momentum tensor inside the cylindrical shell. The vacuum densities in the region outside the cylindrical shell are discussed in section \\ref{sec:Exterior}. The main results of the paper are summarized in section \\ref{sec:Conc}. In appendix \\ref{sec:Appendix} we give an alternative representation for the fermionic condensate and the expectation value of the energy-momentum tensor for a massive fermionic field in the geometry of a cosmic string without boundaries. ",
        "conclusions": "\\label{sec:Conc}  In this paper the vacuum polarization effects are investigated for a fermionic field in the geometry of a cosmic string with a coaxial cylindrical shell. We have assumed that on the shell the field obeys the MIT bag boundary condition. In order to evaluate the fermionic condensate and the VEV\\ of the energy-momentum tensor one needs the complete set of normalized eigenspinors satisfying the boundary condition. This set for the region inside the cylindrical shell is considered in section \\ref% {sec:EigSpinor}. The corresponding mode-sums for both fermionic condensate and the energy-momentum tensor contain series over the zeros of the combination (\\ref{ftilde}) of the Bessel function of the first kind and its derivative. For the summation of these series we used a variant of the generalized Abel-Plana formula previously derived in Ref. \\cite{Saha04Mon}. This formula allows us to extract from the respective VEVs the parts corresponding to the cosmic string geometry without a cylindrical shell and to present the part induced by the shell in terms of exponentially convergent integrals for points away from the boundary. In this way the renormalization procedure for the fermionic condensate and the energy-momentum tensor is reduced to the renormalization of the corresponding quantities in the geometry of the boundary-free cosmic string. The renormalized VEV of the energy-momentum tensor for a fermionic field in the boundary-free geometry is well investigated in literature. In appendix % \\ref{sec:Appendix}, by using the Abel-Plana summation formula, we give alternative integral representations for both fermionic condensate and the energy-momentum tensor in the case of a massive field.  In the region inside the shell, the parts in the VEVs induced by the presence of the cylindrical boundary are given by formula (\\ref{FCcyl}) for the fermionic condensate and by Eq. (\\ref{T00int}), (\\ref{Tnuint}) for the vacuum energy densities and stresses. These formulae are further simplified for a massless fermionic field with the vacuum densities given by Eqs. (\\ref% {FCm0}) and (\\ref{EMTm0}). For points near the cylindrical shell the boundary induced parts in the VEVs dominate over the boundary-free parts and diverge on the cylindrical shell. These type of divergences are well known in quantum field theory with boundaries and are investigated for various bulk and boundary geometries. In the problem under consideration, the leading terms in the asymptotic expansions in powers of the distance from the boundary are given by Eq. (\\ref{FCnear}) for the fermionic condensate and by Eqs. (\\ref{T22NearCyl}) and (\\ref{T00nearCyl}) for the components of the energy-momentum tensor. These leading terms do not depend on the planar angle deficit and are the same as for a cylindrical boundary in the Minkowski bulk. The boundary induced parts in the VEVs vanish on the string axis for $q>1$ and are non-zero in the case of a cylindrical boundary in the Minkowski bulk. Since the boundary-free part diverges on the string axis, for points near the string it dominates. For large values of the parameter $% q $, which corresponds to a large planar angle deficit, the boundary-induced VEVs are suppressed by the factor $(r/a)^{q}$. The boundary induced VEVs have non-trivial dependence on the mass of the field and, as it is illustrated by figure \\ref{fig3}, for a massive field the polarization effects can be stronger than for a massless one.  Fermionic vacuum densities in the region outside a cylindrical shell with MIT bag boundary condition are investigated in section \\ref{sec:Exterior}. Subtracting from the mode-sums the parts corresponding to the geometry of a string without boundaries and by making use of a complex rotation, we have derived explicit expressions for the boundary induced VEVs. The corresponding parts in the fermionic condensate and the energy-momentum tensor are given by Eqs. (\\ref{FCcylext}), (\\ref{T00ext}) and (\\ref{Tnuext}% ). When the cylinder radius goes to zero, for a fixed value of the radial distance, the boundary induced part in the VEV\\ of the energy-momentum tensor vanishes as $a^{q}$ for a massive field and as $a^{q+1}$ for a massless one. At large distances from the cylindrical shell this part is exponentially suppressed for a massive field and decay as $r^{-4}(r/a)^{q+1}$ in the case of a massless field. Note that in the latter case the boundary-free part behaves as $r^{-4}$ and it dominates at large distances. For points near the cylindrical shell the leading terms in the asymptotic expansions in powers of the distance from the boundary are given by the same formulae as for the interior region. In this limit the total VEV is dominated by the boundary induced part. In dependence of the mass, the vacuum stresses can be either positive or negative, whereas the energy density is positive. In the special case $q=1$, from the formulae derived in the present paper we obtain the fermionic Casimir densities for a cylindrical boundary in the Minkowski spacetime.  We have considered the idealized geometry of a cosmic string with zero thickness. A realistic model for cosmic string has a non-trivial structure on a length scale defined by the phase transition at which it is formed. As it has been shown in Refs. \\cite{Alle90,Alle92,Alle96}, the internal structure of the string may have non-negligible effects even at large distances. Here we note that when the cylindrical boundary is present, the VEVs of the physical quantities in the exterior region are uniquely defined by the boundary conditions and the bulk geometry. This means that if we consider a non-trivial core model with finite thickness $b<a$ and with the line element (\\ref{ds21}) in the region $r>b$, the results in the region outside the cylindrical shell will not be changed. As regards to the interior region, the formulae given in this paper are the first stage of the evaluation of the VEVs and other effects could be present in a realistic cosmic string. Note that from the point of view of the physics in the exterior region the cylindrical surface with MIT bag boundary condition can be considered as a simple model of non-trivial string core."
    },
    "0809/0809.0596_arXiv.txt": {
        "abstract": "A new search result of the Tokyo axion helioscope is presented. The axion helioscope consists of a dedicated cryogen-free 4T superconducting magnet with an effective length of 2.3\\, m and PIN photodiodes as x-ray detectors. Solar axions, if exist, would be converted into X-ray photons through the inverse Primakoff process in the magnetic field. Conversion is coherently enhanced even for massive axions by filling the conversion region with helium gas. The present third phase measurement sets a new limit of $g_{a\\gamma\\gamma}<\\mbox{5.6--13.4}\\times10^{-10}\\rm GeV^{-1}$ for the axion mass of $0.84<m_a<1.00\\rm\\,eV$ at 95\\% confidence level. ",
        "introduction": "The existence of axion is implied to solve the so-called strong CP problem \\cite{axion-bible1,axion-bible2,axion-bible3,axion-bible4,axion-bible5}. Axions are expected to be produced in solar core through their coupling to photons with energies of order keV, and the so-called `axion helioscope' technique may enable us to detect such axions directly \\cite{sikivie1983,bibber1989}.  The differential flux of solar axions at the Earth is approximated by \\cite{bahcall2004,raffelt2005} \\begin{eqnarray} \\rm d \\Phi_a/\\rm d E&=&6.020\\times10^{10}[\\mathrm{cm^{-2}s^{-1}keV^{-1}}] \\nonumber\\\\ &&{}\\times\\left(g_{a\\gamma\\gamma}\\over10^{-10}\\mathrm{GeV}^{-1}\\right)^2 \\left( \\frac{E}{1\\,\\mathrm{keV}}\\right)^{2.481} \\exp \\left( -\\frac{E}{1.205\\,\\mathrm{keV}}\\right), \\label{eq:aflux} \\end{eqnarray} where $g_{a\\gamma\\gamma}$ is the axion-photon coupling constant. Their average energy is 4.2\\,keV reflecting the core temperature of the sun. Then, they would be coherently converted into X-rays through the inverse process in a strong magnetic field at a laboratory. The conversion rate in a simple case is given by \\begin{equation} P_{a\\to\\gamma} = \\left( \\frac{g_{a\\gamma\\gamma}B_\\bot L}{2}\\right)^2 \\left[ \\frac{\\sin(qL/2)}{qL/2}\\right]^2, \\label{eq:prob_plain} \\end{equation} where $B_\\bot$ is the strength of the transverse magnetic field, $L$ is the length of the field along the axion path, $q=(m_\\gamma^2-m_a^2)/2E$ is the momentum transfer by the virtual photon, $m_a$ is the axion mass, and $m_\\gamma$ is the effective mass of the photon which equals zero in vacuum.  If one can adjust $m_\\gamma$ to $m_a$, coherence will be restored for non-zero mass axions. This is achieved by filling the conversion region with gas. A photon in the X-ray region acquires a positive effective mass in a medium. In light gas, such as hydrogen or helium, it is well approximated by \\begin{equation} m_\\gamma=\\sqrt{4\\pi\\alpha N_e\\over m_e}, \\end{equation} where $\\alpha$ is the fine structure constant, $m_e$ is the electron mass, and $N_e$ is the number density of electrons. We adopted cold helium gas as a dispersion-matching medium. It is worth noting that helium remains at gas state even at 5--6\\,K, the operating temperature of our magnet. Since the bore of the magnet is limited in space, the easiest way is to keep the gas at the same temperature as the magnet. Moreover, axions as heavy as a few electronvolts can be reached with helium gas of only about one atmosphere at this temperature. ",
        "conclusions": "The axion mass around 1\\,eV has been scanned with an axion helioscope with cold helium gas as the dispersion-matching medium in the $4{\\rm\\,T}\\times2.3\\rm\\,m$ magnetic field, but no evidence for solar axions was seen. A new limit on $g_{a\\gamma\\gamma}$ shown in Fig. \\ref{fig:exclusion} was set for $0.84<m_a<1.00\\rm\\,eV$. It is the first result to search for the axion in the $g_{a\\gamma\\gamma}$-$m_a$ parameter region of the preferred axion models~\\cite{GUT_axion} with a magnetic helioscope. Full description of the present result is published in Ref. \\cite{sumico2008}."
    },
    "0809/0809.4261_arXiv.txt": {
        "abstract": "The stellar masses, mean ages, metallicities, and star formation histories of galaxies are now commonly estimated via stellar population synthesis (SPS) techniques.  SPS relies on stellar evolution calculations from the main sequence to stellar death, stellar spectral libraries, phenomenological dust models, and stellar initial mass functions (IMFs) to translate the evolution of a multi-metallicity, multi-age set of stars into a prediction for the time-evolution of the integrated light from that set of stars. Each of these necessary inputs carries significant uncertainties that have until now received little systematic attention.  The present work is the first in a series that explores the impact of uncertainties in key phases of stellar evolution and the IMF on the derived physical properties of galaxies and the expected luminosity evolution for a passively evolving set of stars.  A Monte-Carlo Markov-Chain approach is taken to fit near-UV through near-IR photometry of a representative sample of low- and high-redshift galaxies with this new SPS model.  Significant results include the following: 1) including uncertainties in stellar evolution, stellar masses at $z\\sim0$ carry errors of $\\sim0.3$ dex at 95\\% CL with little dependence on luminosity or color, while at $z\\sim2$, the masses of bright red galaxies are uncertain at the $\\sim0.6$ dex level; 2) either current stellar evolution models, current observational stellar libraries, or both, do not adequately characterize the metallicity-dependence of the thermally-pulsating asymptotic giant branch phase; 3) conservative estimates on the uncertainty of the slope of the IMF in the solar neighborhood imply that luminosity evolution per unit redshift is uncertain at the $\\sim0.4$ mag level in the $K-$band, which is a substantial source of uncertainty for interpreting the evolution of galaxy populations across time.  Any possible evolution in the IMF, as suggested by several independent lines of evidence, will only exacerbate this problem.  4) Assuming a distribution of stellar metallicities within a galaxy, rather than a fixed value as is usually assumed, can yield important differences when considering bands blueward of $V$, but is not a concern for redder bands.  Spectroscopic information may alleviate some of these concerns, though uncertainties in the stellar spectral libraries and the importance of non-solar abundance ratios have not yet been systematically investigated in the SPS context. ",
        "introduction": "\\label{s:intro}  The spatially-integrated UV, optical, and near-IR light of a galaxy contains a wealth of information regarding its physical properties. The spectrum of this light is governed principly by the star formation and metal enrichment histories of a galaxy in conjunction with stellar evolution and attenuation by interstellar dust.  Combining these ingredients in order to predict the spectrum of a galaxy is known as stellar population synthesis (SPS), and has an extensive history \\citep[e.g.][]{Tinsley76, Tinsley80, Bruzual83, Renzini86, Bruzual93, Worthey94, Maraston98, Leitherer95, Fioc97, Vazdekis99, Yi03a, Bruzual03, Jimenez04, Maraston05, CidFernandes05, Ocvirk06}.  Sophisticated SPS models have allowed for an enormous body of work aimed at constraining physical parameters of galaxies by comparing the observed spectroscopic and/or photometric properties of $z\\sim0$ galaxies to SPS models.  Physical parameters that are constrained by SPS include, but are not limited to, stellar masses \\citep[e.g.][]{Bell03, Kauffmann03a, Panter07}, star formation histories and rates \\citep[e.g.][]{Kauffmann03b, Panter07}, and metallicities \\citep[e.g.][]{Gallazzi05}.  SPS is now routinely used not only at $z\\sim0$ but also at higher redshifts in order to interpret the observed broad-band colors and spectra of galaxies \\citep[e.g.][]{Shapley05, Daddi07a, Kriek06a, Bundy06}.  Key aspects of stellar evolution theory --- an essential ingredient of any SPS model --- have been extensively tested against observations including globular clusters and star counts in the solar neighborhood, stellar halo, and the Large and Small Magellenic Clouds (LMC and SMC, respectively), suggesting that in many respects they are sufficiently reliable to be applied in SPS modeling of observed galaxies.  However, there are significant phases in stellar evolution that are still not well understood, both observationally and theoretically.  Such phases include blue stragglers (BSs), the horizontal branch (HB), the asymptotic giant branch (AGB), and thermally pulsating AGBs (TP-AGBs). Relatedly, the different evolutionary path(s) taken by binary systems, which appear to constitute the majority of all stellar systems, is not well-understood.  These phases are particularly important for SPS models because stars in such phases are more luminous than the main sequence (in some cases by several orders of magnitude), and in the case of AGBs they can dominate the red/IR light in some age ranges of a coeval set of stars, while in the case of the HB and BS they can contribute importantly to the blue and UV light.  These uncertainties will thus impact our ability to derive physical properties of galaxies \\citep[e.g.][]{Charlot96a, Charlot96b, Yi03, LeeHC07}.  The potential importance of uncertain phases of stellar evolution was highlighted recently by the different treatments of TP-AGBs in the SPS models of \\citet{Maraston05} and \\citet{Bruzual03}.  These models, when applied to high-redshift galaxies where light is dominated by $\\sim0.2-2$ Gyr old populations, result in {\\it systematic} factors of $\\sim2$ differences in mass and age \\citep{Maraston06, Bruzual07, Kannappan07}.  This example demonstrates the need for understanding in detail how the uncertainties in stellar evolution propagate into estimates of physical properties of galaxies.  In addition to uncertainties in stellar evolution, there are also substantial uncertainties in our knowledge of the stellar initial mass function \\citep[IMF][]{Kroupa01}.  The IMF effects not only the normalization of the mass-to-light ratio but also, to a lesser degree, the colors of a coeval set of stars \\citep{Tinsley80}.  The IMF thus plays an important role in deriving physical properties of galaxies \\citep[e.g.][]{Papovich01, LeeHC04a}.  The IMF also controls the rate of luminosity evolution for a passively evolving system, and is thus of fundamental importance for interpreting the evolution of galaxy populations across time \\citep[e.g.][]{Tinsley80, Yi03}.  In this paper we present a new SPS code that was developed in order to address the impact of uncertainties in stellar structure and evolution on the derived physical properties of galaxies.  This new SPS model is then applied to a variety of galaxies drawn from the Sloan Digital Sky Survey \\citep[SDSS][]{York00} and the Two Micron All Sky Survey \\citep[2MASS][]{Jarrett00} at $z\\sim0$, and from a sample of luminous red galaxies at $z\\sim2$, in order to investigate the effects of uncertain phases of stellar evolution on derived physical properties such as stellar mass and star formation history.  In addition to the uncertain aspects of stellar evolution, we also explore the effects of the IMF and the potential biases that arise when assuming a single metallicity for all stars in a galaxy rather than the more physical assumption of a distribution of stellar metallicities.  Where necessary, we adopt a flat $\\Lambda$CDM cosmology with $(\\Omega_m, \\Omega_\\Lambda,h)=(0.26,0.74,0.72)$.  Throughout, mass-to-light ratios are quoted in units of $M_\\Sol/L_\\Sol$. ",
        "conclusions": "\\label{s:disc}  \\subsection{The importance of accurate uncertainties}  The importance of accurate and robust estimates for not only the central values of various physical properties of galaxies but also their associated uncertainties cannot be underestimated.  We highlight here several areas where accurate uncertainties play an important, if often neglected role.  With the advent of large galaxy surveys and publically-available SPS codes, the estimation of stellar mass functions has become routine in both the local and distant Universe \\citep[see e.g.][]{Cole01, Bell03, Fontana04, Drory04, Bundy05, Borch06}.  Since the mass function declines exponentially at the massive end, any uncertainty in the masses of individual galaxies can lead to a substantial misestimation of the underlying mass function.  \\citet{Cattaneo08} provide illustrative examples of the extent to which uncertainties can broaden an intrinsically steep mass function.  This important effect is rarely taken into account when either trying to estimate the underlying mass function or when trying to compare to models of galaxy formation. With accurate and reliable uncertainties, one can begin to deconvolve the true mass function from the broadening effects of uncertainties. This task is critical if one is to make accurate comparisons to models of galaxy formation and evolution.  This concern does not reside solely with the mass function.  There are many relations that are effected, including the relation between star formation rate and stellar mass, the star formation rate function, and the relation between stellar mass and large scale clustering strength. This last relation is essential for understanding the connection between galaxies and the underlying dark matter structure --- an accurate accounting of the uncertainties is thus essential.   \\subsection{Why the IMF matters} \\label{s:whyimf}  As discussed in $\\S$\\ref{s:uimf}, it is extremely difficult to constrain the slope of the IMF in the solar neighborhood, let alone in external galaxies.  This is unfortunate given the immense importance of the IMF in many areas of astrophysics.  For example, the IMF is not often considered a source of uncertainty when attempting to constrain the evolution of the luminosity function of galaxies through time \\citep[though see][who discuss this source of uncertainty]{Cool08}.  In recent years it has become common to construct the luminosity function for passively evolving (i.e. red) galaxies at multiple epochs in order to understand how these galaxies grow with time.  Recently, several authors have found evidence for very little intrinsic evolution in the bright end of the luminosity function of passive galaxies, suggesting that massive galaxies have evolved little since $z\\sim1$ \\citep[e.g.][]{Wake06, Brown07, Cool08}.  Passive galaxies are considered relatively easy to study because, in the absence of any appreciable star formation, the time-dependence of the luminosity of such systems is determined simply by stellar evolution and the IMF.  Authors take expected luminosity evolution between two epochs for fiducial stellar evolution models and an IMF and compare to the evolution in the luminosity functions between the same two epochs in order to infer the underlying (intrinsic) change in the population with time.  We found in $\\S$\\ref{s:uimf} however, that the uncertainty in the slope of the IMF in the solar neighborhood translates into an uncertainty in the amount of passive luminosity evolution of $\\sim0.4$ magnitudes per unit redshift in the $K-$band.  The difficulty of accounting for systematics in the observed IMF, and the unique difficulty of measuring the IMF slope near $1\\Msun$ --- precisely where it must be measured accurately for passive evolution corrections --- suggest that this uncertainty in luminosity evolution could be an underestimate.  Any possible evolution in the IMF, as discussed in $\\S$\\ref{s:uimf}, will only exacerbate this problem.  This uncertainty can have serious consequences for understanding the evolution of galaxies.  We highlight one example here.  There has been some tension in the literature between merger rates of luminous passive systems inferred from close-pair counts and from the evolution of the luminosity function.  Taken at face value, constraints from the luminosity function suggest that the merger rate amongst bright, passive galaxies has been very modest since $z=1$ \\citep{Wake06, Brown07, Cool08}, in contrast to the apparently higher rates inferred from the close pair counts \\citep{vanDokkum05, Bell06b, Masjedi06, Lin08}.  If the slope of the IMF were shallower than the canonical value, then from Figure \\ref{fig:lumev} we conclude that luminosity evolution would be stronger than is typically assumed.  If this were the case, then evolution of the luminosity function would instead require a higher merger rate than inferred with the canonical IMF slope. The slope of the IMF is thus intimately related to the merger rate as inferred from evolution of the luminosity function.  A comprehensive comparison between various merger rate indicators has yet to be undertaken, and the apparent tension in the literature between indicators may be largely interpretive.  If however there is a quantitative tension, a non-canonical IMF slope may reconcile different measures.  An additional and important constraint is the evolution of the fundamental plane.  As discussed in detail in \\citet{vanDokkum08}, the evolution in the fundamental plane can provide constraints on the logarithmic slope of the IMF because one can measure and compare mass evolution to luminosity evolution, so long as one is confident that the descendants of high redshift galaxies can be reliably identified at low redshift.  There are additional assumptions that make it difficult in practice to set strong constraints.  For example, this approach requires knowledge of the fraction of mass that is attributed to dark matter.  Nonetheless, use of additional constraints, such as the fundamental plane and its evolution, can provide valuable additional constraints not only on the form of the IMF but also on uncertain aspects of stellar evolution.   \\subsection{Spectroscopy and the SDSS} \\label{s:spec}  We have focused on fitting our SPS model to the broad-band photometry of galaxies in the near-UV through near-IR.  We now discuss the use of observed spectra in constraining SPS models.  In addition to broad-band photometry, the SDSS provides reasonable signal-to-noise, medium resolution ($R\\sim2000$) spectroscopy for $\\sim670,000$ galaxies as of DR4.  These spectra provide a substantial amount of information not available in broad-band photometry.  Indeed, several independent groups have made extensive use of spectroscopy to infer physical properties of SDSS galaxies \\citep[e.g.][]{Kauffmann03a, Panter07}.  There are several concerns related to using spectroscopy for SPS modeling that have not received extensive attention in the SPS literature.  First, the robustness of the stellar spectral libraries is not known over the full luminosity, temperature, and metallicity range required.  The TP-AGB spectra are clearly deficient, as described in $\\S$\\ref{s:tpagb}, and will thus introduce uncertainty in the interpretation of optical and near-IR spectral features.  As discussed in $\\S$\\ref{s:stellib}, there are also known systematic problems with the color-temperature relations in the stellar libraries, indicating again that the libraries probably contain unknown systematics.  Furthermore, the recent results of \\citet{Maraston09} suggest that the theoretical spectral libraries are deficient in important ways.  An example of these complications was highlighted by \\citet{Wild07}, who recently showed that there is a significant offset in spectral indices between the BC03 model and SDSS data.  The authors attribute the mis-match to the stellar spectral libraries used in the BC03 model and caution that systematic errors in SPS models may be significantly larger than quoted statistical errors.  A second source of concern, at least for SDSS galaxies, is that the spectra are taken through fibers that only cover a fraction of the total galaxy.  Averaged over the full spectroscopic sample with $0.01<z<0.15$, the spectra only captures the light from $\\approx30$\\% of the total galaxy.  For example, a galaxy may have both a bulge and a disk but the fiber would only cover the bulge and thus the spectra of this galaxy would not be representative of the total galaxy.  The standard procedure to deal with this fiber effect is to use the observed spectra in the SPS fitting and then make a correction based on the observed color(s) outside of the fiber \\citep[see e.g.][]{Brinchmann04}.  Given the unknown systematics in the spectral libraries, in conjunction with the simple correction that is typically adopted to account for light outside of the fiber, we do not believe that global physical properties of galaxies derived from spectra are more robust than information derived solely from broad-band photometry.  In fact, there is some evidence that the masses derived from spectra contain significant systematic uncertainties.  \\citet{Maller09} showed that the stellar masses derived with spectroscopic information by \\citet{Kauffmann03a} are a function of the inclination of the galaxy in the sense that face-on galaxies are on average $\\approx 0.2$ dex heavier than edge-on galaxies.  Of course, an intrinsic property of a galaxy such as its stellar mass should not be a function of its orientation projected on the sky.  This trend thus suggests a systematic bias in the masses derived by \\citet{Kauffmann03a}.  The stellar masses estimated from broad-band photometry by \\citet{Blanton07} show a weaker dependence on inclination, and the estimates based on broad-band photometry from \\citet{Bell03} show no trend with inclination.  It is not known if the trend in the spectroscopically-derived masses is an intrinsic problem with using spectroscopic data or a systematic in the methodology of \\citet{Kauffmann03a}.  \\subsection{Additional uncertainties not explored herein} \\label{s:other}  There are additional uncertainties not explored herein that may significantly impact the results of SPS modeling.  In this section we briefly discuss some of the more important uncertainties.  In $\\S$\\ref{s:stellib} we discussed several of the limitations of the current generation of stellar spectral libraries.  \\citet{Martins07} has presented a detailed comparison between several theoretical and empirical libraries and found substantial disagreement for both blue and red colors (such as $U-B$ and $V-K$, respectively).  The difficulty in modeling stellar spectra is exacerbated by the fact that empirical libraries span a relatively narrow range in metallicity. Without a well-sampled empirical library it is thus difficult to understand and isolate the limitations of the models \\citep[see also][]{Charlot96a, Charlot96b, Yi03}.  The impact of non-solar abundance ratios has not been discussed in the present work.  Observations consistently show that the $\\alpha$-to-Fe ratio is a function of galaxy mass, with more massive galaxies being more highly $\\alpha$-enhanced \\citep[e.g.][]{Thomas05}.  Despite this well-known observational fact, all SPS models that attempt to determine physical properties of galaxies from either their colors or spectral shapes rely on stellar spectra and evolutionary calculations with solar abundance patterns\\footnote{A considerable amount of work has gone into re-calibrating the basic outputs of solar-scaled SPS models (the SSPs) so that they may be applied more generally to variable abundance ratios.  Such efforts attempt to use a variety of absorption line strengths as a diagnostic to determine the abudance patterns of individual galaxies.  This method is calibrated on the abundance patterns of globular clusters for which detailed abundance patterns can be measured on a star-by-star basis.  For further information see e.g. \\citet{Worthey94, Thomas03}.}.  The effect of this omission on derived physical properties of galaxies has not been adequately quantified, but various calculations suggest that the effect can be substantial, especially for spectral signatures \\citep[see e.g.][]{Coelho07}.  The complication here is two-fold because the abundance patterns effect not only the spectra of individual stars but also the isochrones and their evolution.  Fully self-consistent models of $\\alpha$-enhanced abundance patterns are only just beginning \\citep{Coelho07, Dotter07, LeeHC09}, but early results suggest that the magnitude of this effect on colors can be as high as 0.1 magnitudes \\citep{Coelho07}.  We have not explored uncertainties in either the main sequence turn-off point or the red giant branch.  \\citet{Gallart05} has presented an extensive comparison between the most popular stellar evolution models and found that slightly different treatments in the underlying physics can lead to detectable differences in derived properties such as ages and metallicities \\citep[see also][]{Maraston05}.  For example, current stellar models predict different times for the onset of the red giant branch, owing in part to different treatments of convection in the interiors (i.e. assumptions relating to convective core overshoot and to the mixing length parameter).  These differences are large enough to be detected in star clusters in the LMC \\citep{Ferraro04}, and can impact the broad-band colors of galaxies \\citep{Yi03, LeeHC07}.  Another source of uncertainty that is currently not treated in SPS models is contamination of the galactic light from active galactic nuclei (AGN).  Overall, the fraction of galactic light attributed to AGN is small \\citep[e.g.][]{Hao05}, but there are certain regimes where the contribution may be non-negligible.  The problem can be particularly acute when spectroscopic information is not available to isolate AGN candidates."
    },
    "0809/0809.4943.txt": {
        "abstract": "We report the first direct observations of neutral, molecular gas streaming in the nucleus of NGC1068 on scales of $<30$ pc using SINFONI near-infrared integral field spectroscopy. At a resolution of $0.075\\arcsec$, the flux map of 2.12 $\\mu$m 1--0 S(1) molecular hydrogen emission around the nucleus in the central arcsec reveals two prominent linear structures leading to the AGN from the north and south. The kinematics of the gas in these features are dominated by non-circular motions and indicate that material is streaming towards the nucleus on highly elliptical or parabolic trajectories whose orientations are compatible with that of the disk plane of the galaxy. We interpret the data as evidence for fueling of gas to the central region. The radial transport rate from $\\sim30$ pc to a few parsec from the nucleus is $\\sim$15 M$_\\sun$ yr$^{-1}$. One of the infalling clouds lies directly in front of the central engine. We interpret it as a tidally disrupted streamer that forms the optically thick outerpart of an amorphous clumpy molecular/dusty structure which contributes to the nuclear obscuration. ",
        "introduction": "A detailed description of the distribution and kinematics of the molecular gas in the central region of Seyfert galaxies is crucial for understanding the fueling of the nucleus and the role of gas in obscuring the active galactic nucleus (AGN). NGC1068, at a distance of 14.4 Mpc \\citep{bland97}, is the brightest Seyfert 2 galaxy and therefore one of the best candidates for direct investigations of the morphology and dynamics of the molecular gas neighboring an active nucleus. The idea of a rotating molecular/dusty torus surrounding a supermassive black hole emerged from the interpretation given by \\citet{miller83} of optical spectropolarimetry of precisely this galaxy, which revealed scattered type--I emission from the obscured broad-line region (BLR). Since then, several near- and mid-IR observations have in fact found and begun to elucidate the physical conditions of the concentration of molecular gas and dust in the nucleus of NGC1068 (near-IR: Young et al. 1996; Marco et al. 1997; Rouan et al. 1998, 2004; Alloin et al. 2001; Galliano \\& Alloin 2002; Galliano et al. 2003; Weigelt et al. 2004; Wittkowski et al. 2004; Gratadour et al. 2005, 2006; mid-IR: Bock et al. 2000, Tomono et al. 2001, 2006; Jaffe et al. 2004; Galliano et al. 2005; Mason et al. 2006; Poncelet et al. 2006, 2007), all favouring indirectly the existence of a compact molecular/dusty torus but failing to obtain a clear image of it. Alternatively, by means of mid-IR observations over the central 140 pc (2$\\arcsec$), \\citet{cameron93} proposed a model in which the bulk of the molecular gas and dust is located at large distances (several tens of parsecs) from the AGN. Line-of-sight attenuation of the BLR in this case would be a mere consequence of one or more intervening molecular clouds. More recently, \\citet{jaffe07} analyzed new interferometric mid-IR observations of NGC1068 and found that the fitted Gaussian components to the $(u,v)$ plane of the central 10 $\\mu$m source resemble a disk similar to the H$_2$O masers \\citep{greenhill96}. Both, the dust and H$_2$O maser disks, appear to be oriented neither perpendicular to nor aligned with the radio jet. However, one needs to be cautious when interpreting all of these observations. As they are measurements of the continuum emission, the inferred gas/dust distributions are strongly dependent on temperature, tracing in fact radiation of matter at a given temperature rather than spatial gas/dust distributions.  The 2.122 $\\mu$m H$_2$ rovibrational $\\nu=$1--0 S(1) emission line probes hot ($\\geq10^3$ K) and moderately dense ($\\geq10^3$ cm$^{-3}$) molecular gas and as such, may be an excellent tracer of gas in the nuclear region. Its spatial distribution can expose the potential presence of a molecular/dusty torus, and also provide physical information on fueling or feedback mechanisms. Imaging spectroscopy of this line \\citep{rotaciuc91, blietz94} has indicated the presence of significant amounts of hot, dense, circumnuclear molecular gas extending over the central 4 arcseconds which is associated with the narrow-line clouds. The nuclear region shows a strong peak almost $1\\arcsec$ east of the nucleus and a weaker one $\\sim1\\arcsec$ to the west. Furthermore, the H$_2$ emission at a few arcseconds from the nucleus is more extended along the major axis of the bar \\citep{davies98}. More recently, \\citet{galliano02} obtained 2D spectroscopic observations of the warm molecular gas in the central $4\\arcsec\\times4\\arcsec$ of the galaxy. They did not detect H$_2$ emission at the location of the $K-$band continuum core (which is known already to be coincident with the central engine at these scales). Instead they confirmed the two main regions of H$_2$ emission at about $1\\arcsec$ east and west of the nucleus along a PA=$90\\degree$, and a region with complex line profiles at $\\sim50$ pc north of the AGN. They interpreted the observed H$_2$ emission at these scales in terms of a warped disk consistent with the interpretation given by \\citet{schinnerer00} to millimeter interferometric maps of $^{12}$CO(2--1) emission. It is now apparent from recent SINFONI H$_2$ 1--0 S(1) data at these scales that simple warped disk models of the molecular gas \\citep{schinnerer00, baker00} cannot account for the fantastic variety and detail in the morphological and kinematical structure (see Davies et al. 2006), in particular since there is a considerable amount of gas inside the inner edge of the 100 pc ring.  In order to analyze in detail the physical conditions of the warm molecular gas in the central region of an active nucleus, we have obtained adaptive optics assisted SINFONI integral field spectroscopic data to peer deeply into the energetic core of NGC1068. In this paper we focus on the observations of the molecular gas emission in the central arcsecond of the galaxy done with the smallest spatial scale of the instrument, resulting in a resolution of 75 milli-arcseconds, approximately 7 times better than previous observations. We refer to \\citet{mueller08} for the analysis of the stars, molecular and ionized gas in the central $3\\arcsec$ of the galaxy corresponding to the rest of the results from the SINFONI observations of NGC1068. ",
        "conclusions": "\\label{conclusions}  We have presented in this paper high resolution SINFONI observations of the molecular gas in the nucleus of NGC1068. The distribution of the H$_2$ 1--0 S(1) emission at a resolution of $0.075\\arcsec$ has been resolved. Two bright regions connected by a linear structure extending up to $0.7\\arcsec$ north of the nucleus along a PA of $\\sim-14\\degree$ are distinguished: one lying right in front of the nucleus (the southern tongue) and another one $0.35\\arcsec$ north of the center (the northern tongue). The northern tongue correlates with the mid-IR emission and its tip coincides with a knot of radio continuum emission providing direct evidence of the shock interface between the jet and a molecular cloud that has caused the jet direction to change. The main results of our analysis on the kinematics of these components are summarized as follows:   \\begin{itemize}  \\item  Dynamical modeling shows that material is streaming towards the nucleus. The infalling gas is contained on elliptical/parabolic orbits whose orientation is consistent with that of the plane of the galaxy. We interpret this as strong evidence of how gas, on scales of a few parsecs, is fueling the AGN.  \\item The gas transport is from $\\sim70$ pc to a few parsecs from the nucleus in the plane of the galaxy. The modeling reveals the existence of two streamers: a northern streamer associated with the northern tongue that passes very close to the nucleus (pericenter of 5 pc), and a southern streamer which lies currently in front of the nucleus associated with the southern tongue, and has a pericenter of $\\sim1$ pc.   \\item The mass inflow rate $\\mathrm{d}M/\\mathrm{d}t$ is  $\\sim 15$ M$_\\odot$ yr$^{-1}$ from scales of 30 pc to a few pc. This is about 1000 times that from a few pc to a few times the Schwarzschild radius $R_S$. A similar change in the mass inflow rate with radius is observed in the Galactic Center. This may be a natural consequence of inefficient angular momentum transport.  \\item The geometry, kinematics and high column density of the nuclear concentration of molecular gas (the southern tongue) can be explained by a tidally disrupted streamer consisting of a set of infalling clouds that form the optically thick outerpart of an amorphous clumpy molecular/dusty structure. \\end{itemize}"
    },
    "0809/0809.1819_arXiv.txt": {
        "abstract": "Several recent papers have suggested that the cosmological constant $\\Lambda$ directly influences the gravitational deflection of light. We place this problem in a cosmological context, deriving an expression for the linear potentials which control the cosmological bending of light, finding that it has no explicit dependence on the cosmological constant. To explore the physical origins of the apparent $\\Lambda$-dependent potential that appears in the static Kottler metric, we highlight the two classical effects which lead to the aberration of light. The first relates to the observer's motion relative to the source, and encapsulates the familiar concept of angular-diameter distance. The second term, which has proved to be the source of debate, arises from cosmic acceleration, but is rarely considered since it vanishes for photons with radial motion. This apparent form of light-bending gives the appearance of curved geodesics even within a flat and homogeneous universe. However this cannot be construed as a real lensing effect, since its value depends on the observer's frame of reference. Our conclusion is thus that standard results for gravitational lensing in a universe containing $\\Lambda$ do not require modification, with any influence of $\\Lambda$ being restricted to negligible high-order terms. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1734_arXiv.txt": {
        "abstract": "We investigate the role that dry mergers play in the build-up of massive galaxies within the cold dark matter paradigm. Implementing an empirical shut-off mass scale for star formation, we find a nearly constant dry merger rate of $ \\sim 6 \\times 10^{-5}$ Mpc$^{-3}$ Gyr$^{-1}$ at $z \\leq 1$ and a steep decline at larger z. Less than half of these mergers are between two galaxies that are morphologically classified as early-types, and the other half is mostly between an early-type and late-type galaxy. Latter are prime candidates for the origin of tidal features around red elliptical galaxies. The introduction of a transition mass scale for star formation has a strong impact on the evolution of galaxies, allowing them to grow above a characteristic mass scale of $M_{*,c} \\sim 6.3 \\times 10^{10}$ M$_{\\odot}$ by mergers only. As a consequence of this transition, we find that around $M_{*,c}$, the fraction of 1:1 mergers is enhanced with respect to unequal mass major mergers. This suggest that it is possible to detect the existence of a transition mass scale by measuring the relative contribution of equal mass mergers to unequal mass mergers as a function of galaxy mass. The evolution of the high-mass end of the luminosity function is mainly driven by dry mergers at low z. We however find that only $10\\% -20\\%$ of galaxies more massive than $M_{*,c}$ experience dry major mergers within their last Gyr at any given redshift $z \\le 1$. ",
        "introduction": "Mergers are fundamental to the cold dark matter paradigm of structure formation. They not only drive mass evolution by merging smaller dark matter haloes into larger ones, but they also change the morphology of galaxies from late to early-type \\citep[e.g.][]{1972ApJ...178..623T,1992ARA&A..30..705B,2003ApJ...597..893N}, and drive gas to the centre of the merger remnant that triggers star formation \\citep{1991ApJ...370L..65B,2006MNRAS.372..839N} and AGN activity \\citep{2005Natur.433..604D}. Early structural studies of elliptical galaxies by \\citet{1992ApJ...399..462B} already hinted at the mass-dependent importance of dissipation during their formation. In a first systematic study of the morphology of merging pairs in a CDM galaxy formation model, \\citet{2003ApJ...597L.117K} could show that massive elliptical galaxies are mainly formed from dry mergers of early-type galaxies, while less massive ones show mixed mergers between an elliptical and a spiral galaxy. Only elliptical galaxies well below $L_*$ are predominantly formed by wet mergers from two spiral galaxies. Subsequent work on the role of dry mergers revealed that they can explain the formation of slow rotating boxy ellipticals \\citep{2005MNRAS.359.1379K,2006ApJ...636L..81N}, that they lie on the fundamental plane \\citep{2001ApJ...552L..13C,2003MNRAS.342..501N,2006MNRAS.369.1081B,2006ApJ...641...21R}, follow the $M_{\\bullet}-\\sigma$-relation \\citep{2008arXiv0802.0210J} and that they could possibly explain the formation of a stellar density core in the centre of the remnant due to a binary black hole merger \\citep{2001ApJ...563...34M,2004ApJ...613L..33G,2006ApJ...648..976M}. Furthermore, it has been argued that the strong size evolution of  massive early-type galaxies \\citep{2007MNRAS.382..109T,2007ApJ...671..285T,2008ApJ...677L...5V,2008A&A...482...21C} provides evidence for dry merging \\citep{2006ApJ...648L..21K}. In an attempt to model the size-evolution of early-type galaxies, \\citet{2006ApJ...648L..21K} showed that the amount of dissipation during mergers can account for the observed size evolution. In their model, dry mergers result in remnants with larger sizes than remnants from gaseous mergers of the same mass.  Similar results have been reported from numerical simulations of mergers with varying degrees of gas fractions by \\citet{2006ApJ...650..791C}.  The natural question that immediately arises is, what is the reason for dry merging? The early seminal work of \\citet{1977ApJ...215..483B}, \\citet{1977ApJ...211..638S} and \\citet{1977MNRAS.179..541R} predicts the existence of a characteristic mass scale, below which the cooling time $t_{cool}$ of a collapsing gas cloud is shorter than its dynamical  time $t_{dyn}$, allowing for efficient collapse on a dynamical time scale and subsequent star formation. In massive dark matter halos  with $M_{DM} > 10^{12}$ M$_{\\odot}$, one generally finds $t_{cool} \\gg t_{dyn}$ and that shock heating of the collapsing gas supports the formation of  a hot, static atmosphere at the halo virial temperature \\citep[e.g.][]{2003MNRAS.345..349B,2007MNRAS.380..339B}. From an observational point of view, the existence of a bimodality in the properties of the galaxy population, occurring at a  mass scale of M$_* > 3 \\times 10^{10}$ M$_{\\odot}$ \\citep{2003MNRAS.341...54K}, lends support to the notion of a transition in the mode of galaxy formation \\citep{2006MNRAS.368....2D}. Hence one expects that dry merging will occur when cooling is sufficiently hindered at masses above a transition mass scale and if the reservoir of cold gas in the galaxy is used up by star formation before the merger happens.  The existence of a characteristic shut-off mass scale in galaxy formation seemingly provides a simple way to truncate star formation within galaxy formation models. Such an ad-hoc prescription has been used in earlier work by \\citet{1999MNRAS.303..188K} with the aim of avoiding too massive and too blue galaxies in clusters, and more recently in work by \\citet{2006MNRAS.370.1651C}. These latter  authors assume a shut-down of star formation in halos of mass $\\geq 10^{12}$ M$_{\\odot}$ at $z \\leq 3$, and show  that the colour bimodality and luminosity function can be reproduced accurately in their model. The choice of $10^{12}$  M$_{\\odot}$ draws its support from two main arguments laid out in \\citet{2006MNRAS.368....2D}. One is that at this mass scale,  stable shocks appear that allow for shock heating of gas \\citep{2003MNRAS.345..349B}. The second argument, more important, is that this shock-heated gas is generally so dilute and vulnerable to feedback that it literally stays hot forever and does not cool down to subsequently fuel star formation. While the first argument draws support from various simulations \\citep[e.g.][]{2005MNRAS.363....2K}, the second argument is less clear. One main uncertainty is with regards to the heating source of the hot gas. Several plausible candidates are suggested in the literature such as  AGN-feedback \\citep{1998A&A...331L...1S}, dynamical friction heating \\citep{2004MNRAS.354..169E,2007ApJ...658..710N}, or heating by gravitational potential energy \\citep{2008ApJ...680...54K,2008MNRAS.383..119D}. Neither the relative contribution nor the overall magnitudes with which these processes heat the hot gas are theoretically certain or observationally confirmed to date.  The aim of this letter is twofold. Firstly, we predict the dry merger rate and its evolution by adopting a shut-off mass scale, and secondly, we use these results to propose observational strategies on how to test the existence  of a critical mass scale for the quenching of star formation, that relies on the continued merging activity within CDM-cosmologies and is independent to first order on the underlying baryonic physics involved in the quenching process. \\begin{figure} \\includegraphics[width=0.45\\textwidth]{f1.eps} \\caption{The fraction of galaxies that were formed by a dry major merger within the last Gyr. Lines show results for galaxy mass ranges $5 \\times 10^{10} -10^{11}$, $10^{11} -5 \\times 10^{11}$ M$_{\\odot}$, $5 \\times 10^{11} -10^{12}$ M$_{\\odot}$ and $ >  10^{12}$ M$_{\\odot}$. Galaxies with $M_* > M_{*,c}=6.3 \\times 10^{10}$ M$_{\\odot}$ show similar fractions of dry mergers at $z \\le 1$. } \\label{fig1} \\end{figure} ",
        "conclusions": "In this Letter, we predicted properties of dry mergers in a model that assumes a critical shut-off mass scale for cooling of gas. The impact on the galaxy population and the merger rates can be summarized as follows. The high-mass end of the luminosity function is dominated by continued dry mergers. At any redshift $z \\le 1$,  $10 \\% - 20 \\%$ of massive galaxies have had experienced a dry merger within their last Gyr. We find a dry merger rate of $ \\sim 6 \\times 10^{-5}$ Gyr$^{-1}$ Mpc$^{-3}$ and that the  number density of dry major mergers is significantly increased with respect to a model without shut-off mass scales. The relative fraction of equal mass mergers is enhanced with respect to unequal mass mergers at $M_{*,c}=6.3 \\times 10^{10}$ M$_{\\odot}$ which marks the transition of galaxies from being predominantly formed in gaseous mergers or through star formation in discs to dry mergers.  In a model where the transition between star forming and non-star forming galaxies is regulated by a physical process that does not result in a sharp shut-off mass scale, the relative rates of equal to unequal mass mergers do not show a significant change with mass (e.g. K07). Around the same mass scale, the merger rate is enhanced with respect to the general trend of a decreasing merger rate with mass consistent with recent observations by \\citet{2008arXiv0806.0018P}. All of these features can be explained by considering that galaxies grow through star formation in discs only until their host halos reach $M_{DM,crit}$ and their supply of fuel in the form of cold gas stops. At this mass scale, efficient shock heating kicks in, as well as the gas becoming  prone to efficient heating from various feedback sources \\citep{2006MNRAS.368....2D}. Galaxies on average have masses of $M_{*,c}$ when this is the case, and can only grow to become more massive by mergers. As a result the relation between central galaxy stellar mass and host dark mater halo mass will become shallower, resulting in unequal mass dark halo mergers resulting in similar mass galaxy mergers (see also Hopkins et al. in prep).  The results presented here can be used to test the existence of a shut-off mass scale (see e.g. \\citep{2009ApJ...695..900Y} for a recent study based on galaxy group catalogue, arguing that this mass scale must be larger than $10^{12.5}$ M$_{\\odot}$). If indeed various physical processes conspire to generate a characteristic mass scale,  it should leave its fingerprint in the equal mass  merger rate. Using systematic surveys of galaxies to count pair statistics one can measure the relative fraction of equal to unequal mass mergers and look for a change as a function of mass. This approach is rather insensitive to difficulties with observing changes in luminosity functions, morphologies of galaxies or signs of interactions, and therefore should be able to provide robust results even at larger redshifts. \\\\  We would like to thank the referee for his valuable comments, as well as Shardha Jogee, Gary Mamon, Michael Brown  and Avishai Dekel for helpful comments that improved the manuscript."
    },
    "0809/0809.0218_arXiv.txt": {
        "abstract": "We study the inverse problem of deducing the dynamical characteristics (such as the potential field) of large systems from kinematic observations. We show that, for a class of steady-state systems, the solution is unique even with fragmentary data, dark matter, or selection (bias) functions. Using spherically symmetric models for simulations, we investigate solution convergence and the roles of data noise and regularization in the inverse problem. We also present a method, analogous to tomography, for comparing the observed data with a model probability distribution such that the latter can be determined. ",
        "introduction": "Matter in the universe is usually contained in systems bound together by gravitation: planets and their moons, planetary systems, star clusters, galaxies, and groups of galaxies. Indeed, the modern concepts of gravitation and gravitational potential were brought about by the realization that a mathematically well definable universal force field must keep celestial bodies in their observed orbits. Newton's solution to the inverse problem of ``What kind of a force keeps two pointlike bodies in an elliptic orbit around each other?'' was the inverse-square law of attraction\\footnote{Newton actually solved the direct problem of ``What kind of an orbit is produced by the inverse-square attraction?'' (this is what Halley asked him). The solution of the inverse problem is an instant corollary; it is also unique for motion around the focal point of an ellipse -- the harmonic oscillator creates an elliptic orbit around the centre.}. In such a force field, the equations of motion for a body can be written as \\bee \\frac{d^2 x}{dt^2}=-\\nabla \\Phi(x),\\label{newton} \\eee where $t\\in\\R$ is the time, $x\\in \\R^3$ describes the position of the body, and $\\Phi:\\R^3\\rightarrow \\R$ is the gravitational potential; here $\\Phi(x)\\sim \\Vert x\\Vert^{-1}$.  Outside isolated two-body configurations, the two-body solution only serves as a useful first approximation for systems dominated by one massive body; for such a class of systems, the perturbations from this solution can be examined analytically to a relatively high precision. The behaviour of a system consisting of a moderate number of bodies of more or less equal weight must be numerically integrated assuming the Newtonian potential for each pointlike mass. However, if the system consists of many ($>10^6$ or so) gravitationally interacting bodies collectively forming the potential, we can assume it to obey some principles of statistical mechanics. Accordingly, we no longer analyze the motion of separate particles, but want to define the dynamical characteristics of the system as a whole, such as its global-scale potential field  $\\Phi(x)$, matter density $\\rho(x)$, etc. We define {\\em dynamical tomography} as the inverse problem of deducing these smoothly distributed characteristics from kinematic observations. A more detailed definition of this large-scale generalization of Newton's dynamical inverse problem is given in Section 2.1. The term tomography can be used here in both intuitive and technical senses: we determine density/distribution functions inside a domain, and this can be done by employing the principles of tomographic analysis.  Large gravitationally bound systems have been studied for more than a century, and the corresponding theoretical framework is described in several textbooks, of which \\cite{bt} is the most up-to-date one. However, the studies have mostly concentrated on the direct problem of creating numerically or analytically well tractable dynamical models, or of designing approximate system models to explain various observed phenomena, while the general inverse problem (as posed in sections 2 and 3 here) has received less attention. This is mostly due to the lack of data sufficient for a comprehensive analysis, and the scarcity of efficient methods of dealing with such data. New sky surveys such as LSST (Large Synoptic Survey Telescope, USA) and Gaia (EU), both starting at the beginning of the next decade, will provide a vast amount of kinematic data of the stars in our galaxy, making the construction of a robust machinery for inverse problem analysis essential. We describe here some theoretical and practical aspects of solving problems in dynamical tomography.  The structure of the paper is as follows: in Section 2, we review as well as define some basic concepts of large gravitational systems, and pose the problem in corresponding terminology. In Section 3, we examine the fundamental aspects and uniqueness properties of the inverse problem, and in Section 4 present examples in the simplified case of spherically symmetric systems. In Section 5, we redefine the problem in the sense of probability distributions, which is usually necessary in practice. Finally, in Section 6, we present conclusions and discuss future work and applications. ",
        "conclusions": "We have defined the inverse problem of dynamical tomography and shown that, with suitable assumptions and mathematical tools, the problem is well-posed and solvable. The uniqueness theorems are central to dynamical tomography: they demonstrate that the steady-state assumption allows a unique solution even with fragmentary data, unobservable mass, and observational bias functions.  Another important concept is the possibility to approximate general (or at least near-integrable) systems with integrable ones. Near-integrable systems that are not very old, but old enough to have settled to a quasistable state, appear closer to integrable than old ones as phase-space diffusion (Arnold diffusion) is not extensive yet. As Nekhoroshev's theorem as well as semianalytical approximations and numerical estimates show \\cite{licht}, the time scale for such diffusion grows fast in inverse proportion to the distance of the initial phase-space point from a KAM torus. For example, in \\cite{kaaspre} it was shown that even in chaotic zones it is possible to construct approximate tori such that orbits of the system resemble perturbed motion on these tori for small enough time intervals.  Torus construction \\cite{kaasprl} is a well-defined concept for near-integrable potentials \\cite{poschel}: if a Hamiltonian system is near-integrable in the sense of the KAM theorem, it is integrable on a Cantor set consisting of the surviving KAM tori of the system, i.e., it is possible to construct tori such that their defining action integrals indeed parametrize a global, ordered set. Via torus construction, we find a potential $\\Phi(x)$ for which we can create an approximate set of tori (and determine expressions for their action integrals $J$ and thus for distribution functions $f[J(\\Phi(x);x,v)]$) defining an integrable system such that the steady-state integrability assumption used here agrees with the observations as well as possible. Note that $\\Phi(x)$ itself does not have to be integrable or even near-integrable; the torus set constructed defines an integrable Hamiltonian that is not usually derivable from a potential, so we never get an approximate integrable potential in the first place. This gives the derived $\\Phi(x)$ some additional flexibility; indeed, any integrable potential approximating a real galaxy is probably not a very good representation. The key principle allowing the flexibility is the same as in the uniqueness theorems: we look for isosurfaces in phase space best explaining the observations. The goodness of our solution $\\Phi(x)$ is not directly measured by its closeness to an integrable system.  Our $\\Phi(x)$ should thus be able to mimic the real potential quite well in cases of near-integrable or even somewhat chaotic but not very old systems. The feasible potentials $\\Phi(x)$ that can reproduce, via the torus-construction principle, the approximate isosurface structure of the observed matter distribution can all be expected to be close to each other. In fact, even if $\\Phi(x)$ were integrable, we would have to use torus construction to find it even if we started with an integrable initial $\\Phi_0(x)$, since the iteration procedure (or equivalent) for fitting a model to observations will generally explore non-integrable potentials $\\Phi(x)$.  Let us denote by $H_0$ the integrable Hamiltonian corresponding to the tori constructed for $\\Phi(x)$, and by $H$ the Hamiltonian of $\\Phi(x)$. For optimal $H_0$, the difference (in some chosen norm) \\bee \\Vert \\delta H\\Vert =\\Vert H-H_0\\Vert \\eee is as small as possible over the whole phase space; we call the corresponding tori the optimal tori of $\\Phi$, and $H_0$ the optimal Hamiltonian of $\\Phi$. If $\\Phi$ is integrable, then its optimal tori are its invariant tori: $\\delta H$ vanishes everywhere. For our purposes, $\\Vert\\delta H\\Vert$ does not have to be small (although the smaller it is, the better).  We surmise that: \\begin{enumerate}[i)] \\item The optimal tori constructed for $\\Phi(x)$ form a map $\\Phi\\rightarrow H_0$. With the same construction scheme, a potential $\\Phi(x)+\\epsilon\\phi(x)$ yields an optimal Hamiltonian $H_0+\\epsilon h_0$. \\item The isosurfaces of the distribution function of the system can be approximated by constructing them from a set of 3-tori describable by an integrable Hamiltonian $H_0$ (but not necesssarily by an integrable potential). \\end{enumerate}  Then our $\\Phi$ can be expected to be a good approximation of the real potential of the system (as a regularizing constraint, we can choose, e.g., the smallness of $\\Vert \\delta H\\Vert$).  The possibility of constructing optimal tori and Hamiltonians holds another great advantage: a multitude of deviations from integrability can readily be modelled with a suitably tailored version of Hamiltonian perturbation theory \\cite{kaasham}. This includes both structural detail (resonant orbit families and other details) and timelike irregularities (deviations from steady state).  In a forthcoming study, we will study more general and realistic systems such as axisymmetric and fully three-dimensional St\\\"ackel potentials. We will also introduce more general observational biases and other factors, and investigate their influence. The final goal is the combination of a general torus-construction machinery and an analysis procedure for dynamical tomography, so that we can work with any potentials. With such tools, we can analyze the data from large-scale surveys and construct a consistent mathematical model of the dynamics of our galaxy.  The real galactic problem can be expected to suffer from considerable model noise, so we should employ various forms of modelling the distribution functions and potential. In addition to mapping the density distribution of the dark matter in our galaxy, we should also be able to test whether distributions with alternative theories of gravity are distinguishable from standard models with dark matter."
    },
    "0809/0809.5051_arXiv.txt": {
        "abstract": "We consider gas at densities appropriate to protoplanetary disks and calculate its ability to cool due to line radiation emitted by $\\water$ molecules within the gas.  Our work follows that of Neufeld \\& Kaufman (1993; \\apj, 418, 263), expanding on their work in several key aspects, including use of a much expanded line database, an improved escape probability formulism, and the inclusion of dust grains, which can absorb line photons.  Although the escape probabilities formally depend on a complicated combination of optical depth in the lines and in the dust grains, we show that the cooling rate including dust is well approximated by the dust-free cooling rate multiplied by a simple function of the dust optical depth.  We apply the resultant cooling rate of a dust-gas mixture to the case of a solar nebula shock pertinent to the formation of chondrules, millimeter-sized melt droplets found in meteorites.  Our aim is to assess whether line cooling can be neglected in chondrule-forming shocks or if it must be included.  We find that for typical parameters, $\\water$ line cooling shuts off a few minutes past the shock front; line photons that might otherwise escape the shocked region and cool the gas will be absorbed by dust grains.  During the first minute or so past the shock, however, line photons will cool the gas at rates $\\sim 10^4 \\ {\\rm K} \\, {\\rm hr}^{-1}$, dropping the temperature of the gas (and most likely the chondrules within the gas) by several hundred K.  Inclusion of $\\water$ line cooling therefore must be included in models of chondrule formation by nebular shocks. ",
        "introduction": "\\subsection{Radiative Transfer and Line Cooling}  In many astrophysical settings, emission of radiation via rotational and vibrational transitions of molecules plays an important role in cooling warm gas. Because the emission of such radiation is in sharp spectral lines, this mechanism is referred to as line cooling. Line radiation from the water molecule ${\\rm H}_{2}{\\rm O}$, with its permanent electric dipole and its high cosmochemical abundance, is significant in a variety of settings ranging from molecular clouds to protostellar envelopes (e.g., Cernicharo \\& Crovisier 2005). More recently, line cooling from ${\\rm H}_{2}{\\rm O}$ molecules has been recognized to play a pivotal role in the energetics following the passage of a shock wave through the dense gas in the solar nebula protoplanetary disk (see Desch et al.\\ 2005).  The rate of emission of line radiation from a warm gas containing water molecules is difficult to calculate, because there are so many accessible rotational and vibrational energy levels, and therefore a great many transitions. In dense molecular gas ($n_{\\rm H2} > 10^{10} \\, {\\rm cm}^{-3}$), these energy levels are populated according to Boltzmann statistics (i.e., in local thermodynamic equilibrium, or LTE), but at lower densities the populations must be calculated by balancing transition rates. Finally, the cooling of gas by line radiation relies on the ability of the photons to escape the system without being reabsorbed; to the extent that other, nearby molecules can absorb the emitted photons, the gas does not cool. Photons generated during a molecular transition are emitted and also reabsorbed over a small range of frequencies centered on the line frequency, with an efficiency that depends on the frequency shift from the line center. Rather than integrating the equations of radiative transfer over all frequencies within the line, all standard treatments of the problem instead assume integration over this ``line profile\" and use a frequency-integrated escape probability of photons. This escape probability is a function of the column density of the molecule within the system. This is the approach taken by Neufeld \\& Kaufman (1993; hereafter NK93), in particular.  The calculation of NK93 of the rate of line cooling from ${\\rm H}_{2}{\\rm O}$ molecules has remained the state of the art for 15 years; but, as we discuss below, several aspects of the NK93 calculation are now out of date. Their calculation uses a limited database of transitions, an oversimplified escape probability formulism, and ignores absorption by dust grains. It is one of the goals of the work here to update the ${\\rm H}_{2}{\\rm O}$ line cooling rates to improve on the calculations of NK93. A second goal of this paper is assess whether or not line cooling plays a significant role in cooling the gas following the passage of a shock through the dense gas of the solar nebula. As discussed by Desch et al.\\ (2005), previous modeling has not determined whether or not line cooling can be neglected. This has significant implications for the issue of chondrule formation.  \\subsection{Chondrules}  The parent bodies of the most primitive meteorites, the chondrites, formed at about 2-3 astronomical units (AU) from the Sun, 4.57 billion years ago (Wadhwa \\& Russell 2000). Chondrites are the most primitive meteorites in our collection, in that they have suffered very little alteration since their formation, and therefore contain information about the conditions that existed in the early solar nebula. Chondrites are remarkable for containing calcium-rich, aluminum-rich inclusions (CAIs), the oldest solids in the Solar System, whose formation has been dated to $4567.2 \\pm 0.6$ Mya (Amelin et al.\\ 2002). Also found in abundance within chondrites are sub-millimeter- to millimeter-sized, (mostly ferromagnesian) igneous spheres, called chondrules, from which the chondrites derive their name. Chondrules formed, at most, \\app~2-3 million years after CAIs (Amelin et al. 2002; Kita et al. 2005), as melt droplets that were heated to high temperatures while they were independent, free-floating objects in the early solar nebula. After they were heated, cooled and recrystallized, chondrules were incorporated into the parent bodies from which chondrites originate. Chondrules are capable of providing incredibly detailed information about conditions in the Solar System protoplanetary disk, if the process that led to their heating, melting and recrystallization could be understood. Chondrules make up to 80\\% of the volume of ordinary chondrites (Grossman 1988), and it is estimated that \\app 10$^{24}$ g of chondrules exist in the asteroid belt today (Levy 1988). Such a prevalence of chondrules suggests that chondrule-forming events were widespread in the solar nebula. A process that can melt $10^{24} \\, {\\rm g}$ of rock is surely a dominant process in the solar nebula disk.  The chondrule formation process, despite its obvious importance, has been a mystery in the field of meteoritics for two centuries (Sorby 1877). Proposed mechanisms include interaction with the early active Sun, through jets (Liffman \\& Brown 1995; Liffman \\& Brown 1996) or solar flares (Shu et al. 1996, 1997, 2001), melting by lightning (Pilipp et al.\\ 1998; Desch \\& Cuzzi 2000), and melting by planetesimal impacts (Urey \\& Craig 1953; Urey 1967; Sanders 1996; Lugmair \\& Shukolyukov 2001). The dominant theory, though, is that chondrules were melted in shock waves in the protoplanetary disk (Hewins 1997; Jones et al. 2000; Connolly \\& Desch 2004; Connolly et al. 2006). A shock wave, or shock front, is a sharp discontinuity between supersonic and subsonic gas, over an area only a few molecular mean free paths thick, typically only meters in the solar nebula. The gas is slowed, compressed, and heated by the time it reaches the other side of the shock front. Solids moving with the gas are heated not only by thermal exchange upon entering the shocked region, but also by friction as they are slowed to the post-shock gas speed, and by absorbing the infrared radiation emitted by other solids. The shock model of chondrule formation appears able to resolve the chondrule formation mystery, because it makes several detailed predictions about chondrule formation that are largely borne out by observation and experimentation, especially regarding chondrule thermal histories.  \\subsection{Chondrule Formation by Nebular Shocks}  Heating and cooling rates of chondrules have been determined experimentally. According to furnace experiments, in which melt droplets with chondrule compositions are allowed to cool and crystallize, reproduction of chondrule textures requires particular ranges of cooling rates between the liquidus temperature ($\\approx 1800$ K) and solidus temperature ($\\approx 1400$ K; Hewins \\& Connolly 1996). Porphyritic  chondrules cooled at about $10 -10^3 \\, {\\rm K} \\, {\\rm hr}^{-1}$, and barred-olivine chondrules cooled at about $10^3 \\, {\\rm K} \\, {\\rm hr}^{-1}$ (Hewins et al.\\ 2005; see also Desch \\& Connolly 2002). Additionally, chondrules retain volatile elements such as S, indicating that they did not remain above the liquidus for more than minutes, and cooled quite rapidly ($\\gtrsim 10^4 \\, {\\rm K} \\, {\\rm hr}^{-1}$ while above the liquidus (Yu \\& Hewins 1998)). Finally, there is no indication of the isotopic fractionation that would arise from the free evaporation of alkalis such as Na, which constrains the time spent at high temperature before melting (Tachibana et al. 2004). Modeling of isotopic fractionation has shown that chondrules must heat up from 1300 to 1600 K on the order of minutes or less (Tachibana \\& Huss 2005).  Passage through nebular shocks satisfies nearly all the experimental constraints on chondrule formation in broad brush (Iida et al. 2001, hereafter INSN; Desch \\& Connolly 2002, hereafter DC02; Ciesla \\& Hood 2002, hereafter CH02; Desch et al. 2005). Prior to passage through the shock front, chondrules are moderately heated by absorbing radiation emitted by chondrules which have already passed through the shock front. Upon passage through the shock front, the gas is immediately slowed, and compressed and heated, but the chondrules continue at supersonic speeds through the gas. They achieve their peak heating during this stage, due to absorption of radiation, thermal exchange with the gas, and supersonic frictional drag heating. This stage lasts until the chondrules slow to the gas velocity, which takes an aerodynamic stopping time, $t_{\\rm stop} = \\rho_{\\rm s} a_s / \\rho_{\\rm g} C \\sim 1$ minute (where $\\rho_{\\rm s}$ and $a_s$ are the particle density and radius, and $\\rho_{\\rm g}$ and $C$ are the post-shock gas density and sound speed). At this time, the gas and chondrules are dynamically coupled and chondrules are heated only by absorption of radiation and thermal exchange with the gas. Soon thereafter the solids and gas become thermally coupled as well, and both components achieve similar temperatures. Although the chondrules pass through the first 100 km of the post-shock region more rapidly than the gas, once they do, they will thermally equilibrate to the gas temperature. Once this happens, the gas and chondrules cool together as fast as they can either radiate away energy or leave the source of infrared radiation that heats them. DC02 showed that the cooling rate of gas and chondrules should be \\begin{equation} CR=50 \\; \\left( \\frac{\\rho_{\\rm g}}{10^{-9} {\\rm g} \\, {\\rm cm}^{-3} } \\right) \\left( \\frac{V_{\\rm s}}{7 \\, {\\rm km} \\, {\\rm s}^{-1}}  \\right) \\left( \\delta + \\frac{C}{50} \\right) \\, {\\rm K} \\, {\\rm hr}^{-1}, \\label{eq:cr} \\end{equation} where $\\rho_{\\rm g}$ is the pre-shock gas density, $V_{\\rm s}$ the shock speed, $C$ is the ``concentration\" of chondrules (the chondrules-to-gas mass ratio normalized to $3.75 \\times 10^{-3}$, where a chondrule radius $300 \\, \\mu\\rm m$ has been assumed), and $\\delta$ the concentration of submicron dust (the dust-to-gas mass ratio normalized to $1.25 \\times 10^{-3}$). More precisely, $\\delta$ should be interpreted as the Rosseland mean opacity of the dusty gas, normalized to $1.14 \\, {\\rm cm}^{2} \\, {\\rm g}^{-1}$. For typical parameters, then, cooling rates $\\sim 10^2 \\, {\\rm K} \\, {\\rm hr}^{-1}$ are obtained, consistent with the cooling rates of chondrules.  \\subsection{Estimates of Line Cooling in Shocks}  The effect of line cooling---the cooling of gas by emission of so-called line photons by trace molecules in the gas such as CO and ${\\rm H}_{2}{\\rm O}$--- in solar nebula shocks has been considered by INSN and by Miura \\& Nakamoto (2006). In both cases the cooling rates of NK93 were used. INSN assumed a gas optically thin to the line radiation and found chondrule cooling rates $\\sim 10^{4} \\, {\\rm K} \\, {\\rm hr}^{-1}$ in many cases. Miura \\& Nakamoto (2006) allowed the gas to become optically thick to this line radiation, using the formulation of NK93 described below. They found that for plausible parameters the cooling rates of chondrules and gas cluster around $5000 \\, {\\rm K} \\, {\\rm hr}^{-1}$, with line cooling playing an important role in the cooling of the gas. Chondrules and gas are thermally coupled only a few minutes after passing through the shock front (INSN; DC02; CH02), so these chondrule cooling rates are controlled by the rate at which gas cools by emission of line radiation. To understand the essence of their results without reproducing all of their calculations in detail, we now estimate the cooling rate of chondrules and gas following the shock, using the tabulations of NK93.  We first examine how the cooling rates predicted by NK93 and used by INSN and Miura \\& Nakamoto (2006) are calculated. For a wide range of densities and temperatures, NK93 calculated the steady-state level populations of $\\water$ and CO molecules (denoted `$M$'), and the rates of emission of photons. The escape of these photons from a system was considered under the large-velocity-gradient (Sobolev) approximation as a function of column density $n(M) d$, via use of the parameter \\begin{equation} \\tilde{N} \\equiv \\frac{n(M)}{d v_{z} / d z}, \\end{equation} These results were adapted to a large range of geometries of interest, including those not consistent with the Sobolev approximation, in particular the one of relevance to nebular shocks, a static, plane-parallel slab of thickness $d$. According to NK93, the cooling rate at the center of the slab (at depth $d/2$) will depend on the similar parameter \\begin{equation} \\tilde{N} = \\frac{n(M)\\;d}{\\Delta v}, \\end{equation} where $\\Delta v$ is the thermal velocity of the molecules $M$. An important assumption made by NK93 is that if the line photons can escape to the edge of the slab, then the gas is cooled; if the photons do not escape to the edge of the slab, they are essentially re-absorbed on the spot and there is no cooling. The cooling rate per volume is then $n(m) {\\cal L}_{\\rm LTE} (\\tilde{N})$, where the subscript `LTE' is the notation used by NK93 to denote cooling rates derived under the assumption that the population levels follow a Boltzmann distribution, which is the regime of interest for solar nebula shocks. NK93 have tabulated ${\\cal L}_{\\rm LTE}(\\tilde{N})$ for both vibrational and rotational lines of $\\water$ and CO, for every decade in $\\tilde{N}$.  The plane-parallel slab is equivalent to the nebula shock geometry because of the large velocity jump at the shock front. In a typical chondrule-forming shock, the gas velocity changes nearly instantaneously from $\\approx 7 \\, {\\rm km} \\, {\\rm s}^{-1}$ to $\\approx 1.6 \\, {\\rm km} \\, {\\rm s}^{-1}$, a jump of over $5 \\, {\\rm km} \\, {\\rm s}^{-1}$. This velocity difference is much larger than the thermal velocities of the $\\water$ molecules [$\\Delta v = (2 k T / m_{\\rm H2O})^{1/2}$ $= 1.36 \\, {\\rm km} \\, {\\rm s}^{-1}$ at 2000 K] that control the width of the line profile. Essentially, once a line photon emitted from the post-shock gas escapes to the shock front, it will continue to travel relatively unimpeded by gas. Only one change to the NK93 formulation must be made before it is adopted to the nebular shock geometry. For gas a distance $z$ past the shock front, the equivalent thickness of the slab is $d = 2 z$, and the cooling rate is only half the cooling rate calculated by NK93, who assumed radiation could escape either side of the slab.  The cooling rate of the gas yields \\begin{equation} \\frac{\\partial e}{\\partial t} = \\frac{\\partial}{\\partial t} \\left( \\frac{p}{\\gamma -1} \\right) = 2.8 \\, n_{\\rm H2} \\, k \\, \\frac{\\partial T}{\\partial t} = - \\frac{1}{2} \\; n(\\water)\\;\\mathcal{L_{\\rm LTE}}(\\tilde{N}) , \\end{equation} where we have ignored compression of the gas after passing through the shock front in this toy model, and we have assumed an abundance of He of 10\\% by number. Converting the time derivative to a spatial derivative by assuming a constant gas velocity $V_{\\rm g}$ past the shock front, we find \\begin{equation} \\frac{\\partial T}{\\partial \\tilde{N}} = -\\frac{ \\Delta v }{ 5.6 n_{\\rm H2} V_{\\rm g} k } \\, \\frac{ {\\cal L_{\\rm LTE}}(\\tilde{N}) }{ 2 }, \\label{eq:nkcool} \\end{equation} where we have assumed a constant $\\Delta v$ throughout the post-shock region. In actuality, $\\Delta v$ would decrease past the shock front, but that would merely reduce the total cooling, so our assumption overestimates the degree of cooling. This formula provides the gas temperature as a function of $\\tilde{N}$, the effective column density past the shock front.  Fedkin \\& Grossman (2006) showed that the $n_{\\rm H2O} / n_{\\rm H2}$ ratio in the solar nebula probably varied with temperature, but found that $n_{\\rm H2O} / n_{\\rm H2} = 5 \\times 10^{-4}$ was a good average for the temperatures of interest. We have adopted a canonical ratio $n_{\\rm H2O} / n_{\\rm H2} = 8 \\times 10^{-4}$. This is only slightly higher than some commonly assumed canonical ratios but is consistent with the abundance  of Lodders (2003), $n_{\\rm H2O} / n_{\\rm H2} = 8.88 \\times 10^{-4}$. We note that the exact value of $n_{\\rm H2O} / n_{\\rm H2}$ in the chondrule formation environment is likely variable with time (Ciesla \\& Cuzzi 2006), and consider variations in this ratio later in the paper. Some dissociation of $\\water$ is likely to occur at the peak post-shock temperatures we consider, but we are interested in the maximum cooling possible by line photons, so we neglect dissociation in our calculations.  Interpolating between the tabulated values of $\\mathcal{L}_{\\rm LTE}(\\tilde{N})$ provided by NK93, we have integrated equation~\\ref{eq:nkcool} to find the total drop in temperature by the time the gas is at an equivalent of $\\tilde{N} = 10^{21} \\, {\\rm cm}^{-2} \\, {\\rm km}^{-1} \\, {\\rm s}$ past the shock front. For typical parameters (pre-shock density $n_{\\rm H2} = 2 \\times 10^{14} \\, {\\rm cm}^{-3}$ and velocity $V_{\\rm g} = 7 \\, {\\rm km} \\, {\\rm s}^{-1}$; mixing ratio $n_{\\rm H2O} / n_{\\rm H2} = 8 \\times 10^{-4}$; post-shock density $n_{\\rm H2} = 1.2 \\times 10^{15} \\, {\\rm cm}^{-3}$ and velocity $V_{\\rm g} = 1.2 \\, {\\rm km} \\, {\\rm s}^{-1}$; and post-shock thermal velocity $\\Delta v = 1.6 \\, {\\rm km} \\, {\\rm s}^{-1}$), this equates to $d = 1.3 \\times 10^{4} \\, {\\rm km}$, a distance $z = 6.6 \\times 10^{3} \\, {\\rm km}$, and a time 1.6 hours after passing through the shock front. (Again, this toy model neglects the compression and slowing of the gas. It also assumes wrongly that $\\Delta v$ is constant, whereas it should decrease slightly as the gas cools.  It is nonetheless illustrative.) At these distances past the shock front, the total cooling due to rotational lines of $\\water$ and CO is completely negligible, $< 4 \\, {\\rm K}$. This illustrates that the rotational lines become optically thick so quickly that they are unable to escape the post-shock region and cool the gas. On the other hand, the total cooling from CO vibrational line photons is roughly 72 K. This leads to an average cooling rate of $\\sim 40 \\, {\\rm K} \\, {\\rm hr}^{-1}$, which is important, but not significantly greater than the cooling rates of chondrules ($\\sim 10^{2} \\, {\\rm K} \\, {\\rm hr}^{-1}$). The assumption of INSN that the lines of $\\water$ and CO are optically thin is therefore invalidated, because these lines are indeed optically thick.  The final line cooling mode (not considered by INSN, but considered by Miura \\& Nakamoto 2006) is vibrational line cooling from $\\water$. We find a total cooling due to these line photons on the order of 800 K. This is a significant cooling of the gas: the average cooling over 1.6 hours is $\\sim 500 \\, {\\rm K} \\, {\\rm hr}^{-1}$, with even higher cooling rates obtained early on. The cooling of the gas by vibrational line photons of $\\water$, despite being somewhat optically thick, is the source of the high chondrule cooling rates predicted by Miura \\& Nakamoto (2006).  We confirm the statement made by Miura \\& Nakamoto (2006) that ``...line emission is important for the gas cooling.'' We hereafter focus on cooling due to $\\water$ alone, as the dominant coolant is expected to be water.  Clearly, any model of chondrule cooling rates in solar nebula shocks must account for the emission of line photons from $\\water$. However, incorporating the cooling rates of NK93 directly into a shock code, as Miura \\& Nakamoto (2006) did, is not ideal. First, line cooling due to rotational and vibrational transitions of $\\water$ was calculated by NK93 using roughly 50,000 transitions from the HITRAN database (Rothman et al. 1987). While this was the best available data at the time, much more extensive databases now exist. Second, NK93 calculated escape probabilities under the large-velocity-gradient (Sobolev) approximation, and asserted that the results will apply to the case of a static, plane-parallel slab if $\\tilde{N}$ is defined as described above. Third, NK93 assumed a one-sided escape probability for line photons equal to $0.5 / (1 + 3\\tau)$, where $\\tau$ is the Sobolev optical depth; more exact escape probabilities exist, as described below. Fourth, it is not necessarily the case that if photons do not escape to the edge of the slab that the gas does not cool locally; photons re-absorbed halfway to the edge still cool the gas at the slab center. Finally, NK93 did not consider the absorption of line photons by intervening dust. Inclusion of dust is absolutely mandatory, as the following argument makes clear: For a solar-composition gas with 0.5\\% of the mass in the form of $a = 0.5 \\, \\mu{\\rm m}$ radius grains (similar to the size of matrix dust grains in meteorites), the opacity to short-wavelength radiation is as high as $\\kappa = (\\rho_{\\rm dust} / \\rho_{\\rm gas}) \\, (3 / 4 \\rho_{\\rm s} a) = 30 \\, {\\rm cm}^{2} \\, {\\rm g}^{-1}$, where the internal density of the dust grains is $\\rho_{\\rm s} \\approx 2.5 \\, {\\rm g} \\, {\\rm cm}^{-3}$ . Over the distance for which the optical depth in dust is unity, the parameter $\\tilde{N}$ is therefore less than $10^{19} \\, {\\rm cm}^{-2} \\, {\\rm km}^{-1} \\, {\\rm s}$, assuming a ratio $n_{\\rm H2O} / n_{\\rm H2} = 8 \\times 10^{-4}$. This corresponds to distances $\\sim 10^{3} \\, {\\rm km}$, or times of about 10 minutes, after passage through the shock front. As this is smaller than the $\\tilde{N}$ for which the most significant cooling takes place, dust grains are capable of absorbing photons that would otherwise cool the gas. A thorough treatment of line cooling should improve on the calculations of NK93 in the areas identified above.   \\subsection{Outline}  In this paper we calculate the cooling rate of gas due to emission of line photons by $\\water$ molecules, improving upon the work of NK93 with the use of an expanded database, improved escape probabilities, and the inclusion of dust. First we evaluate escape probability approximations, both in the presence and absence of dust. We then calculate the cooling rates from $\\water$ molecules using the 1.2-million line SCAN-$\\water$ database (J\\\"{o}rgensen et al.\\ 2001). Finally, we incorporate these cooling rates into a toy model of chondrule thermal evolution. We conclude that for canonical parameters, line cooling is significant in the first few minutes past the shock front, but that afterward dust grains absorb the line photons that would otherwise cool the gas. We discuss the effect of the important dust-to-water ratio on these conclusions. ",
        "conclusions": "We have built upon the line cooling work of NK93, by using a much expanded database of spectral lines, improved escape probabilities, and the inclusion of dust grains.  We have found that although the cooling rate due to $\\water$ line cooling is a complicated combination of gas and dust opacity, the cooling rate can be well approximated by  $\\mathcal{L_{\\rm LTE}}(\\tau,\\tau_{\\rm d}) \\approx \\mathcal{L_{\\rm LTE}}(\\tau_{\\rm d}~=~0) \\times E_{2}(0.95\\;\\tau_{\\rm d})$.  This allows us to investigate the maximum cooling effect $\\water$ line cooling will have in chondrule-forming nebular shocks.  Because vibrational line photons from $\\water$ so effectively cool the gas, the high oxygen fugacities and $\\water$ densities inferred for the chondrule-forming environment (Fedkin \\& Grossman 2007) would seem incompatible with cooling rates $<$ 100 K/hr, but we find that dust significantly reduces the photon escape probabilities and the cooling rates. Dust grains absorb the line photons that would escape and cool the gas. Without dust, $\\water$ line cooling reduces the temperature of the gas to 1400 - 1800$\\;$K in $\\leq$ 0.1 hr and cooling continues at this rate ($>$10$^4$ K/hr). With dust, however, the grains absorb line photons and inhibit cooling, leading to a cessation of cooling  after $\\sim$ 0.1 hr ($\\sim$ 8 min).  Line cooling is a very effective and important cooling agent in the first few minutes and must be included in any comprehensive shock model. After the first few minutes, however, the thermally well-coupled gas and chondrules cool together at a slow rate (only as fast as they travel $\\sim$ 1 optical depth through the dust), as in DC02 and CH02.  INSN failed to note this effect;  although they treat the gas and dust as well coupled (T$_{\\rm gas}$ = T$_{\\rm dust}$), their system was effectively optically thin to line emission, allowing all line photons to escape the region and cool the gas.  As we are interested at this time only in the maximum cooling effects of line photons (the ``worst case scenario\"), our calculations of cooling rates made no assumptions about radiative transfer; the chondrule temperatures followed the gas temperature. Thermal histories of chondrules including radiative transfer, line cooling, and other relevant effects will be reported on in future work (Morris et al.\\ 2009b, in preparation). Since we have assumed that line photons either escape completely or are absorbed immediately, the true situation will clearly be bracketed between the two extremes shown in Figure~\\ref{fig:chondrule}. In these calculations, we have assumed that the dust is thermodynamically coupled to the gas. All temperatures (gas, dust, and chondrules), in actuality, will be determined by dust/chondrule opacity and radiative transfer. DC02 showed that the cooling rates of chondrules (without the inclusion of line cooling) were given by the formula that appears above as Equation 1. In the case considered in Figure~\\ref{fig:chondrule}, the chondrule concentration may be considered to be low ($C \\ll 1$), and the dust opacity is characterized by $\\delta = 0.1 \\times 28.76 / 1.14 = 2.5$. This yields a total cooling rate of gas and dust through the crystallization range of about $125 \\, {\\rm K} \\, {\\rm hr}^{-1}$ after line cooling shuts off. Higher chondrule concentrations ($C > 10^2$) appear to have been typical during chondrule formation (Cuzzi \\& Alexander 2006; Alexander et al.\\ 2008). These would increase the cooling rate to several $\\times 10^{2} \\, {\\rm K} \\, {\\rm hr}^{-1}$. Line cooling is significant during the first few minutes past the shock, and more detailed calculations will be necessary to constrain the initial drop in temperature from the peak to below the liquidus. However, our results have shown that dust effectively shuts down line cooling within minutes and it then takes hours for chondrules to crystallize in a shock, consistent with DC02 and CH02.  As mentioned previously, we did not calculate line cooling due to CO, as it is a factor of \\app 20 less effective than $\\water$. CO also cools by emitting line photons, but dust will absorb these line photons as readily as those emitted by $\\water$. Figure~\\ref{fig:chondrule} shows a drop in temperature of \\app 600~K. If CO were included, the temperature would probably drop an extra 5\\% (\\app 630~K), which is within the uncertainties of the $\\water$ abundances we have assumed.  This study has shown that the gas and chondrules will cool much too rapidly to match the experimental constraints on chondrule cooling rates if dust is not present. The question of the dust abundance and its opacity therefore becomes paramount.  But, significantly, we have found that enhancing only the amount of $\\water$ causes higher cooling rates, whereas enhancing the amount of $\\water$ and dust together does not.  What affects the cooling rate the most, therefore, is not the amount of water, but the dust-to-water ratio.  Potential complications, though, include the possibility that dust will evaporate immediately upon passage of the shock, and the possibility that dust vapor may recondense in the post-shock region (e.g., Scott \\& Krot 2005).  In a related vein, we note an application of this work only tangentially related to chondrule formation. Ciesla et al.\\ (2003) suggested that phyllosilicates found in the fine-grained accretionary rims of chondrules in CM meteorites may have resulted from rapid gas-phase reactions of water vapor with silicates. Previous studies indicated that phyllosilicate production would be kinetically inhibited in the solar nebula (Prinn \\& Fegley 1987); however, Ciesla et al. (2003) showed that nebular shocks could result in a local increase in the water vapor pressure, thereby increasing the rate of phyllosilicate formation, as well as increasing the temperature at which phyllosilicates become stable. Chondrule-forming shocks in icy regions of the solar nebula could therefore account for both the formation of chondrules and their fine-grained phyllosilicate rims. A possible objection to this hypothesis is that the high abundance of water vapor in the shocked region would cool the gas too rapidly to allow the production of much phyllosilicates. As our work here shows, however, rapid cooling by line photons will not be associated with shocked, water-rich nebula gas, as long as the dust-to-water ratio is close to the canonical value for the solar nebula. An investigation of this hypothesis is also planned for future work."
    },
    "0809/0809.4488_arXiv.txt": {
        "abstract": "In a universe with a cosmological constant, the large-scale gravitational potential varies in time and this is, in principle, observable. Using an N-body simulation of a $\\Lambda$CDM universe, we show that linear theory is not sufficiently accurate to predict the power spectrum of the time derivative, $\\dot{\\Phi}$, needed to compute the imprint of large-scale structure on the cosmic microwave background (CMB). The linear part of the $\\dot{\\Phi}$ power spectrum (the integrated Sachs-Wolfe effect or ISW) drops quickly as the relative importance of $\\Omega_{\\Lambda}$ diminishes at high redshift, while the non-linear part (the Rees-Sciama effect or RS) evolves more slowly with redshift. Therefore, the deviation of the total power spectrum from linear theory occurs at larger scales at higher redshifts. The deviation occurs at $k\\sim 0.1 $ $h$~Mpc$^{-1}$ at $z=0$. The cross-correlation power spectrum of the density $\\delta$ with $\\dot{\\Phi}$ behaves differently to the power spectrum of $\\dot{\\Phi}$. Firstly, the deviation from linear theory occurs at smaller scales ($k\\sim 1 $ $h$~Mpc$^{-1}$ at $z=0$). Secondly, the correlation becomes negative when the non-linear effect dominates. For the cross-correlation power spectrum of galaxy samples with the CMB, the non-linear effect becomes significant at $l\\sim 500$ and rapidly makes the cross power spectrum negative. For high redshift samples, the cross-correlation is expected to be suppressed by $5-10\\%$ on arcminute scales. The RS effect makes a negligible contribution to the large-scale ISW cross-correlation measurement. However, on arc-minute scales it will contaminate the expected cross-correlation signal induced by the Sunyaev-Zel'dovich effect. ",
        "introduction": "The most intriguing topic in contemporary cosmology is the nature of the dark energy which appears to dominate the energy density of the Universe at late times.  Strong evidence for the existence of dark energy comes from both the combined analysis of the cosmic microwave background radiation (CMB) and the galaxy large-scale structure (LSS) \\citep[e.g.][]{Efstathiou02,Spergel03}, and from high redshift type~Ia supernovae \\citep[e.g.][]{Riess98,Perlmutter99}.  Both of these techniques infer the presence of dark energy from geometrical measures. A complementary probe of dark energy is provided by techniques that measure the dynamical effect of dark energy through its influence on the rate of growth of structure.  Large deep galaxy redshift surveys (like the EUCLID, the ESA Mission to Map the Dark Universe, and the JDEM, the Joint Dark Energy Mission) are being planned that will exploit the redshift space anisotropy of galaxy clustering, caused by coherent flows into overdense regions and outflows from underdense regions, to measure directly the growth rate as a function of redshift.  The Integrated Sachs-Wolfe (ISW) effect \\citep{Sachs67}, in which the decay of the large-scale potential fluctuations induces CMB temperature perturbations, provides another measure of the dynamical effect of dark energy.  In principle, the ISW effect could be detected directly in the CMB power spectrum at very low multiples.  In the $\\Lambda$CDM cosmology, it would boost the plateau in the power spectrum at $l \\sim 10$. However, as the increase of the power is not large in comparison to the cosmic variance, it cannot be unambiguously detected even in the WMAP data \\citep{Hinshaw08}. A more sensitive technique is to search for the ISW signal in the cross-correlation of the LSS with the CMB. As the expected signal is weak and occurs on large scales, a very large galaxy survey is needed to trace the LSS. Currently individual detections based on surveys such as APM, 2MASS, NVSS and SDSS \\citep[e.g.][]{Fosalba03,Afshordi04,Fosalba04, Padmanabhan05, Cabre06,McEwen07, Rassat07, Raccanelli08} are not of very high statistical significance \\citep[but see][]{Granett08}. There are also analyses of the ISW cross-correlation, which combine multiple galaxy survey samples and achieve $\\sim 4\\sigma$ detection of the ISW effect (Ho et al. 2008; Giannantonio et al. 2008). These measurements may be the best that can be obtained before the next generation of surveys (BOSS\\footnote{http://www.sdss3.org/cosmology.php}, Pan-STARRS1\\footnote{http://www.ps1sc.org/}) come to fruition and make redshift tomography possible. If such surveys are to place robust, meaningful constraints on the properties of the dark energy it is important to take full account of other processes beyond the (linear) ISW effect that may contribute to the cross-correlation signal. Here, we focus on deviations caused by non-linear gravitational evolution, the Rees-Sciama effect \\citep{Rees68}.  Other processes are known to contribute to the cross-correlation signal. First, the thermal Sunyaev-Zel'dovich (SZ) effect \\citep{Sunyaev72} caused by hot ionized gas in galaxy clusters induces an anti-cross-correlation signal which can cancel the ISW effect on small scales. Its statistical contribution can be modelled and subtracted given the value of $\\sigma_8$ (the $rms$ linear mass fluctuations within a sphere of 8 $h^{-1}$~Mpc) which determines the abundance of galaxy clusters \\citep[e.g.][]{White93,Fan01,Mei04}. Also, since the thermal SZ effect is frequency dependent, it can be subtracted in frequency space given sufficient spectral coverage.  Second, the redshift dependence of galaxy bias, if not properly taken into account, can introduce systematic effects in the determination of dark energy parameters. Other effects such as lensing magnification and the Doppler redshift effect can also boost the cross-correlation signal, but are only important at high redshift \\citep{Loverde07,Giannantonio07}. These effects are well documented and can be calibrated and removed.  In this paper we will solely explore the contribution of the non-linear terms, or the Rees-Sciama (RS) effect, on the cross-correlation signal.  The RS effect arises from the non-linear evolution of the potential \\citep{Rees68}. It is believed to be much smaller than the CMB signal at all scales \\citep{Seljak96, Puchades06}. Indeed, compared with the CMB power spectrum, the RS effect is orders of magnitude lower. Also, compared with the complete integrated ISW power spectrum, the RS effect has been shown, using the halo model approach \\citep[e.g.][]{Cooray02a, Cooray02b}, to be unimportant at $l<100$. However, the RS effect has not been taken into account in cross-correlation analyses and it is important to assess its importance ahead of the completion of the next generation of large deep galaxy surveys.  We use a large N-body simulation to investigate the effect of the non-linear contribution on the interpretation of the ISW cross-correlation signal.  We use the $488^3$-particle L-BASICC simulation described by \\cite{Angulo08} which, with a box size of 1340 $h^{-1}$~Mpc, is ideal for this purpose because not only does it enable us to extrapolate our analysis to non-linear scales at different redshifts, but it includes the very large scale power necessary to check the agreement with linear theory. The cosmology adopted in the L-BASICC simulation is $\\Lambda$CDM, with $\\Omega_\\Lambda=0.75$, $\\Omega_{\\rm m}=0.25$, $\\Omega_{\\rm b}=0.024$, $\\sigma_8=0.9$ and $H_0=73$~km~s$^{-1}$~Mpc$^{-1}$.  The paper is organised as follows. In \\S2, we compute the power spectrum of the ISW plus RS effects from our simulation and compare them with linear theory. In \\S3, we analyse these two effects in terms of the cross-correlation of the LSS with the CMB. Finally, in \\S4, we discuss our results and present our conclusions. ",
        "conclusions": "We have used an N-body simulation to calculate the non-linear (Rees-Sciama) contribution to the Integrated Sachs Wolfe effect. The comparison of the 3-D and 2-D power spectra measured from the simulation with those given by linear theory reveals a strong nonlinear contribution whose physical scale increase with redshift. We investigated the strength of this effect on the cross-correlation of the CMB with galaxy samples in terms of angular power spectra and in angular coordinates at different redshifts. We find that there is a non-linear contribution to the cross-correlation signal at sub-degree scales. The non-linear effect alters not only the amplitude, but also the shape of the cross-correlation power spectrum. With current galaxy samples which cover relatively small volumes, it is not yet possible to disentangle the contribution of the RS effect from that of the ISW effect within the noise. However, in future surveys like Pan-STARRS and LSST, for which the number of galaxies and the sky coverage will increase dramatically, the error bars on the cross-correlation will be much smaller. In this case, the importance of the Rees-Sciama effect may become significant for high redshift samples.  The effect of the non-linear cross-correlation at scales of arcminutes would contaminate the SZ signal in the CMB and this could confuse its interpretation \\citep[e.g.][]{Fosalba03,Diego03,Myers04, Lieu06, Cao06,Bielby07}.  Our analysis is based on a simulation that assumes a $\\Lambda$CDM cosmology. The non-linear contribution depends on the values of the cosmological parameters. In a flat universe with a cosmological constant, the RS effect will become increasingly dominant relative to the ISW effect as the value of $\\Omega_{\\Lambda}$ decreases. In the most extreme case, if $\\Omega_{\\Lambda}=0.0$, the ISW effect will vanish, leaving only the RS effect \\citep{Seljak96}.  The analysis of this paper could be generalized either using re-normalised perturbation theory \\citep[e.g.][]{Crocce06}, or simulations with different dark energy models. In any case, more general modelling of the non-linear effect will be required for an accurate interpretation of future measurements of the LSS-CMB cross-correlation."
    },
    "0809/0809.3057_arXiv.txt": {
        "abstract": "{Cosmic Rays; High Energy Astrophysics; Acceleration of Particles; Magnetic Turbulence} Ultra High Energy Cosmic Rays (UHECRs) hit the Earth's atmosphere with energies exceeding $10^{18}$ eV. This is the same energy as carried by a tennis ball moving at 100 km/h, but concentrated on a sub-atomic particle. UHECRs are so rare (the flux of particles with $E > 10^{20}$ eV is $0.5$/km$^2$/century) that only a few such particles have been detected over the past 50 years.  Recently, the HiRes and Auger experiments have reported the discovery of a high-energy cut-off in the UHECR spectrum, and Auger has found an apparent clustering of the highest energy events towards nearby active galactic nuclei. Consensus is building that the highest energy particles are accelerated within the radio-bright lobes of these objects, but it remains unclear how this actually happens, and whether the cut-off is due to propagation effects or reflects an intrinsically physical limitation of the acceleration process.  The low event statistics presently allows for many different plausible models; nevertheless observations are beginning to impose strong constraints on them. These observations have also motivated suggestions that new physics may be implicated. We present a review of the key theoretical and observational issues related to the processes of propagation and acceleration of UHECRs and proposed solutions. ",
        "introduction": "Hess (1912) discovered cosmic rays by measuring ionisation at increasing altitudes in a balloon flight. Auger et al. (1938) proved the existence of extended air showers of secondary particles caused by the interaction of primary particles, having energies above $10^{15}$ eV ($1$eV $= 1.602\\times10^{-19}$J), with the Earth's atmosphere. Auger simultaneously observed the arrival of secondary particles in ground detectors set many meters apart. After almost 90 years of research, the origin of cosmic rays and the physical mechanism by which they are accelerated has remained an open question.  The Universe is filled with the remnant radiation emitted by the Big Bang, the Cosmic Microwave Background (CMB) radiation. Since we do not observe degradation of Ultra High Energy Cosmic Rays (UHECRs) due to the interaction with CMB, their sources must be located within 75 Mpc ($1$ Mpc $=10^6$ pc and $1$ pc $\\sim 3.09 \\times 10^{18}$ cm).  The supernova remnant, i.e. the material ejected in the interstellar space by the explosion of a massive star, is believed to be a strong candidate for the source of medium-energy cosmic rays, although direct evidence for this is still lacking. Moreover, the processes of acceleration of particles in the supernova remnant cannot overcome the threshold of $\\sim 10^{15}$eV since higher energy particles cannot be confined in such objects (Lagage \\& Cesarsky 1983). Therefore the sources of UHECRs must be located elsewhere, very likely in nearby extragalactic sources.  This review is intended to cover theoretical and observational aspects of the acceleration mechanisms which lead to Ultra High Energy (UHE), without entering the mathematical details. Inevitably not all works have been properly discussed or mentioned; for a more complete review see for instance Torres  \\& Anchordoqui (2004) and Bhattacharjee \\& Sigl (2000). ",
        "conclusions": ""
    },
    "0809/0809.4677_arXiv.txt": {
        "abstract": "Advances in VLBI instrumentation now allow wideband recording that significantly increases the sensitivity of short wavelength VLBI observations. Observations of the super-massive black hole candidate at the center of the Milky Way, SgrA*, with short wavelength VLBI reduces the scattering effects of the intervening interstellar medium, allowing observations with angular resolution comparable to the apparent size of the event horizon of the putative black hole.  Observations in April 2007 at a wavelength of 1.3mm on a three station VLBI array have now confirmed structure in SgrA* on scales of just a few Schwarzschild radii.  When modeled as a circular Gaussian, the fitted diameter of SgrA* is 37$\\mu$as ($+16$,$-10$; $3\\sigma$), which is smaller than the expected apparent size of the event horizon of the Galactic Center black hole. These observations demonstrate that mm/sub-mm VLBI is poised to open a new window onto the study of black hole physics via high angular resolution observations of the Galactic Center. ",
        "introduction": "At a distance of $\\sim8\\;$kpc (\\cite{reid93}), SgrA*, the compact radio, NIR and X-ray source at the Galactic Center, is thought to mark the position of a super-massive black hole of mass $\\sim4\\times10^6\\;M_\\odot$ (\\cite{schoedel02},\\cite{ghez05}).  Proper motions of SgrA* (\\cite{reid04}) confirm that it traces a significant amount of the mass that is inferred by stellar motions and orbits.  Due to its proximity, SgrA* is the only galactic nucleus that can be studied with VLBI on sub-AU linear scales ($R_{\\mbox{sch}}\\sim10\\mu\\mbox{as}\\sim0.1\\mbox{AU}$).  The ionized ISM, however, scatter-broadens images of SgrA* with a $\\lambda^2$ dependence, and VLBI at the highest frequencies is the only available means to set important limits on intrinsic structures near the event horizon.  VLBI at 7~mm and 3.5~mm has detected evidence for intrinsic structure of SgrA*, but these observations remain dominated by scattering effects, and the intrinsic sizes at these wavelengths (set by the optical depth of the emission) are much larger than the apparent size of the event horizon (\\cite{bower04}, \\cite{shen05}). Evidence from light curves of SgrA* flares from the radio to Xray (\\cite{eckart06}, \\cite{yusef-zadeh06}, \\cite{marrone08}) implicate structures on smaller ($\\sim5-15 R_{\\mbox{sch}}$) scales.  Only at frequencies above 230GHz does the scattering size become smaller than the VLBI array resolution allowing direct measurement of intrinsic structure on these scales corresponding to the innermost accretion region.  Over the past decade, MIT Haystack Observatory has focused on developing next-generation wideband VLBI instrumentation capable of significantly increasing the sensitivity of mm/submm VLBI arrays.  The observations described herein used this instrumentation to show that 1.3mm VLBI of SgrA* can now probe the event horizon of this super-massive black hole candidate. ",
        "conclusions": "The technology to significantly increase the sensitivity of VLBI at wavelengths of 1.3mm and shorter is now enabling observations of SgrA* on Schwarzschild radius scales.  Efforts to extend the capabilities of the current 1.3mm VLBI array include: phasing together radio telescopes on Mauna Kea and at CARMA to increase collecting area (discussed elsewhere in these proceedings), continuing to pursue increases in recording bandwidth, and bringing new mm/submm VLBI sites on-line.  By 2009, the international collaboration (see the author list of \\cite{doeleman08a}) that has carried out these observations, will field a higher sensitivity 1.3mm VLBI array that will be sensitive to time variable structures in SgrA*.  The closure phase is the sum of interferometric phases around a triangle of baselines, is largely immune from calibration errors, and deviates from zero in the presence of asymmetric source structure.  Projected sensitivities will allow monitoring of the closure phase on $\\sim10$ second time scales, and enable tests for periodic structure variation in SgrA* as is predicted by orbiting hot-spot models of the accretion flow (\\cite{avery06}, \\cite{genzel}).  By timing periodicities in the closure phase, one can extract the fundamental black hole spin parameter (\\cite{doeleman08b}).  Thus, by spatially resolving the innermost accretion region surrounding SgrA*, mm/submm VLBI is now positioned to address fundamental issues in black hole physics."
    },
    "0809/0809.1608_arXiv.txt": {
        "abstract": "We present new NICMOS and ACS observations of the quasar jet PKS 0637-752, and we use them, together with existing multiwavelength observations, to produce the most  complete spectral coverage of the source to date. We explore the implications of these observations in the context of models for the jet X-ray emission. By relaxing the assumption of equipartition, we undertake an exhaustive  study of the parameter space for external Compton off the CMB  (EC/CMB) model. We find that  the multiwavelength observations  exclude a magnetic field dominated jet.  Using the method proposed by Georganopoulos et al. (2005) for probing the jet matter content we show that protons are needed for practically all jet configurations, extending a previous application of the method by Uchiyama et al. (2005) that was based on exploring  three particular jet configurations. We also show that equipartition is the only configuration that can reproduce the observations and have one proton per radiating lepton. We finally present a rather model - independent argument that the jet has a spine-sheath flow pattern, with the spine being faster and emitting most of the IR-optical-X-ray emission. ",
        "introduction": "}      When the \\emph{Chandra X-Ray Observatory} used the bright X-ray quasar PKS 0637-752 to  focus  its mirror assembly during its first observation, astronomers expected to detect  a  bright X-ray point source. However, in addition to the bright core of the quasar, \\emph{Chandra}  detected significant X-ray emission from the previously detected  quasar radio jet   \\citep{schwartz00,chartas00}. This bright X-ray emission was found to overproduce the X-ray flux that corresponds to the extension of the radio-optical spectrum to X-ray energies, arguing for a second spectral component, for which the most plausible candidate was initially thought to be  synchrotron - self Compton (SSC) emission. However,  it became immediately apparent  from modeling the  multi-wavelength spectral energy distribution (SED) of the knots of  PKS 0637-752 that  the  expected SSC X-ray flux in equipartition conditions severely underpredicts the observed X-ray flux, with the case for SSC only getting  worse when relativistic beaming effects  are considered \\citep{schwartz00,chartas00}.   Following the discovery of the X-ray jet of PKS 0637-752, there has been extensive observational and theoretical work on powerful quasar jets (for a recent review see Harris \\& Krawczynski 2006). Further X-ray studies using \\emph{Chandra} have confirmed that bright X-ray emitting jets are common among radio-loud quasars and radio galaxies  (e.g. Sambruna et al. 2002, 2004, 2006; Marshall et al. 2005; Siemiginowska et al. 2002; Kataoka et al. 2003; Jorstad et al. 2004; Cheung, Stawarz, \\& Siemiginowska 2006;  Schwartz et al. 2006). They also confirmed that in these jets the X-ray emission is a spectral component separate from the radio optical synchrotron emission, with the X-ray flux being higher that that anticipated by simply extending the radio-optical spectrum to X-ray energies.     A plausible candidate for the X-ray emission mechanism was suggested by \\cite{tavecchio00} and \\cite{celotti01}. These authors argued that the X-ray flux of the knots of PKS 0637-752 is due to external Compton emission  from jet relativistic electrons  in equipartition with the magnetic field of the jet,  up-scattering  cosmic microwave background photons (EC/CMB).   The model, modivated by the superluminal motions observed in  the sub-pc scale jet \\citep{lovell00},  requires that (i) the large scale jet is significantly beamed (bulk Lorentz factor $\\Gamma\\sim 10$) and (ii) the electron distribution has a low energy break or cutoff at $\\gamma \\lesssim 100 $. Both these requirements increase the required jet power uncomfortably close to the Eddington limit \\citep{dermer04}. Because in the EC/CMB model  the X-rays  are produced by very low energy electrons ($\\gamma \\sim 100$), these electrons  have long cooling times that  allow them to reach the terminal points of even megaparsec scale  jets without losing a significant fraction of their energy. Thus,  in the EC/CMB we would expect to see continuous X-ray emission throughout the jet instead of the  knotty X-ray structure we usually observe. A possible solution to this issue is to assume that jets are not continuous and that the X-ray knots are ejected from the central engine during periods of high activity and retain their knot structure as they propagate downstream.  Also, in several X-ray jets, the ratio of the radio to X-ray flux in the knots decreases as we move away from the core. In the context of the EC/CMB model this can be explained if jets gradually decelerate \\citep{georganopoulos04}. In some jets, including PKS 0637-752, there is a low emission bridge of $\\sim 5-10$ arcsec connecting the quasar core to the first jet knot. \\cite{georganopoulos05} argued that, in the context of the EC/CMB model, this can be used to constrain the matter content of the jet through IR observations of the bridge, and \\cite{uchiyama05} used this to show, through {\\sl Spitzer} observations of PKS 0637-752, that an electron-positron jet is not favored  if the X-rays are due to the EC/CMB mechanism.   The fact that  the EC/CMB model requires bulk Lorentz factors much in excess than required by jet-to-counter-jet radio asymmetry arguments (e.g. Arshakian \\& Longair 2004), together with the  issues just discussed, raises the possibility that the X-rays are of synchrotron nature (e.g. review by Harris \\& Krawczynski 2006) by an additional population of very high energy ($\\sim 100$ TeV) electrons. Such high energy electrons may be produced if the jet is characterized by velocity shear \\citep{stawarz02, rieger04}. Alternatively, the decay of a collimated neutron beam produced in the sub-pc scale jet can provide a second high energy electron component \\citep{atoyan04}. The synchrotron  option is viable for sources such as 3C 273, for which the high energy component extends smoothly down to optical energies  \\citep{uchiyama06, jester06, jester07}, but runs into difficulties for sources like PKS 0637-752, in which the high energy component cuts off before the optical regime, because these high energy electrons would cool rapidly to lower energies and would  radiate in the optical,  producing, thereby,  a much higher optical flux than  observed. If, indeed, the X-ray emission is of synchrotron nature, the multi TeV electrons responsible for it will unavoidably up-scatter the CMB photons to TeV energies \\citep{georganopoulos06}. For 3C 273, the upper limit from shallow HESS TeV observations is compatible with the synchrotron interpretation, provided $\\delta \\lesssim 10 $, where $\\delta $ is the usual Doppler factor.    Here we present HST NICMOS and ACS observations of PKS 0637-752, and discuss the constraints they pose, together with available multiwavelength observations,  on our understanding of the jet physics.  In \\S \\ref{obs} we present  the data reduction procedure, in \\S \\ref{results} our observational results, together with existing multiwavelength observations,    in \\S \\ref{discussion}  the  constraints on the jet physics, and in \\S \\ref{conclusions} our conclusions. ",
        "conclusions": "}  We present new NICMOS and ACS observations, which we use, together with existing broadband observations to build  the most complete SED of the jet knots of quasar PKS 0637-752. We relax the equipartition assumption and  show that in the EC/CMB framework for the X-ray emission, the plausible requirement that the $10''$ jet  is at most 1 Mpc long excludes magnetically dominated jets. We also show that limits on the BC emission exclude $e^--e^+$ jets in the EC/CMB framework, regardless of equipartition: the jet needs to carry some, or most of its power with protons. Interestingly, if we require one proton per radiation lepton, only the equipartition jet does not require substantially super-Eddington power. Taking radiative cooling into account, the Lorentz factor $\\Gamma\\sim 18-20$ required to  model the knot SED for the equipartition solution is significantly higher that those used previously, and  is more in agreement with the radio core limit $\\Gamma > 17.4$, derived  from core superluminal motions \\citep{lovell00}. Configurations with significantly more than one protons per radiating lepton are excluded in the EC/CMB model. Finally, we note that the counter jet radio feature and its polarization, together with the limits on the IR-optical and X-ray emission, suggest a spine sheath flow, with most of the IR-optical-X-ray emission coming from the spine."
    },
    "0809/0809.1322_arXiv.txt": {
        "abstract": "$N$-body simulations of Cold Dark Matter (CDM) have shown that, in this hierarchical structure formation model, dark matter halo properties, such as the density profile, the phase-space density profile, the distribution of axial ratio, the distribution of spin parameter, and the distribution of internal specific angular momentum follow `universal' laws or distributions. Here we study the properties of the first generation of haloes in a Hot Dark Matter (HDM) dominated universe, as an example of halo formation through monolithic collapse. We find all these universalities to be present in this case also. Halo density profiles are very well fit by the Navarro et al (1997) profile over two orders of magnitude in mass. The concentration parameter depends on mass as $c \\propto M^{0.2}$, reversing the dependence found in a hierarchical CDM universe. However, the concentration-formation time relation is similar in the two cases: earlier forming haloes tend to be more concentrated than their later forming counterparts. Halo formation histories are also characterized by two phases in the HDM case: an early phase of rapid accretion followed by slower growth. Furthermore, there is no significant difference between the HDM and CDM cases concerning the statistics of other halo properties: the phase-space density profile; the velocity anisotropy profile; the distribution of shape parameters; the distribution of spin parameter, and the distribution of internal specific angular momentum are all similar in the two cases. Only substructure content differs dramatically. These results indicate that mergers do not play a pivotal role in establishing the universalities, thus contradicting models which explain them as consequences of mergers. ",
        "introduction": "\\label{sec:intro}  The mass distribution of the self-gravitating, quasi-equilibrium dark haloes that form in an expanding universe is an issue of fundamental importance. Early work on self-similar spherical collapse predicts virialized structures with power-law density profiles \\citep{ fillmore84,bertschinger85}. As $N$-body techniques improved, it was realised that, in the hierarchical universes, the profile departs significantly from a single power law and is better fitted by a profile with curvature in a log-log plot. \\citep{efstathiou85,dubinski91}. Later, it was established that the density profiles of haloes in CDM and other hierarchical clustering cosmologies, have a universal form which is well represented by the simple fitting formula of \\citet[][hereafter NFW]{nfw96,nfw97} \\begin{equation} \\rho(r)=\\frac{\\rho_s}{(r/r_s)(1+r/r_s)^2} \\end{equation} where $r_s$ is a characteristic radius where the logarithmic density slope is $-2$, and $\\rho_s/4$ is the density at $r_s$. The above equation implies that $\\rho\\propto r^{-1}$ in the inner regions and $\\rho \\propto r^{-3}$ in the outskirts. A useful alternative parameter for describing the shape of the profile is the concentration parameter $c=r_{200}/r_s$ ($r_{200}$ is the virial radius  defined as the radius within which the mean density is $200$ times the critical value). Shortly after the NFW papers, \\citet{huss99b} showed that the NFW model also fit halos which form by monolithic collapse (i.e. without mergers) in a HDM cosmology. Nevertheless much work has focused on the role of mergers in establishing the NFW profile, since mergers contribute substantially to the growth of haloes in the CDM model \\citep{raig98,salvador98,syer98, subramanian00,dekel03}. \\citet{syer98} analysed dynamical processes during repeated mergers, concluding that the universal profile is generated by tidal stripping of small haloes as they merge into larger objects. \\citet{dekel03} extended this model and found the tidal compression from a halo core makes the satellite orbits decay from the radius where $\\rho \\propto r$ to the halo center and causes a rapid steepening of the inner profile to $\\rho \\propto r^{-\\alpha} (\\alpha >1)$. Recent work has confirmed the NFW hypothesis, extending the comparison to a much wider range of densities and radii than the original work, and uncovering small but significant deviation, between the mean density profiles of simulated dark halos and the NFW formula \\citep{power03,diemand04,merritt06,graham06,gao08,hayashi08}. The differences, however, are quite small compared to the scatter between different haloes of the same mass.  In addition to the density profile, other halo properties are found to follow universal profiles or universal distributions  in the hierarchical CDM model. The ``phase-space density'' $\\rho(r)/\\sigma^2(r)$ has a remarkably accurate pure power-law distribution with radius \\citep{taylor01,barnes06}. The distributions of shape parameters (e.g. the axis ratios) have a weak dependence on mass and redshift \\citep{bullock02b,kasun05, allgood06,bett07}. The spin parameter distribution is well described by a log-normal function which varies weakly with halo mass \\citep{barnes87,warren92, cole96, bullock02a,bett07}. \\citet{bullock01b} also found that the cumulative mass distribution of specific angular momentum $j$ is well fitted by a universal function. All these distributions appear to depend little, if at all, on the global cosmological parameters and on the shape of the initial matter power spectrum. However, the origin of these universalities is still not well understood.  In this paper, we will test whether mergers are the dominant physical mechanism to produce these universal profiles or distributions. To this end, we study a range of properties of the haloes which grow by two different paradigms, by  hierarchical aggregation and monolithic collapse.  The concordance cosmology ($\\Lambda$CDM) is  a standard hierarchical universe, and the HDM dominated universe provides an example where formation of the first haloes is monolithic. In a HDM universe, small objects cannot form because free-streaming effects smooth out small-scale structure in the initial condition. The first halos then form by smooth collapse.  We begin in section~\\ref{sec:simulation} with a summary of the $N$-body simulations and halo catalogues used in this paper. In section~\\ref{sec:haloform}, we give a general description of how halos form and evolve in the HDM universe. In section ~\\ref{sec:result}, a range of halo properties are studied in depth in our two model universes. Then, we will summarise and discuss our results in section~\\ref{sec:conclusions}.  \\begin{figure} \\bc \\hspace{-1.4cm} \\resizebox{9cm}{!}{\\includegraphics{figures/halo_evo.ps}}\\\\% \\caption{The evolution of a massive halo in a HDM universe with mass $1.3\\times 10^{15} M_{\\odot}/h$ at redshift zero. All particles within $r_{200}$ at $z=0$ are traced back to seven higher redshifts. In addition the particles in the very inner region $r < 0.01 r_{200}$ are traced back and are shown as the blue points. The size of each box is $\\rm 8 Mpc/h$ in physical (not comoving) units. The redshift and the corresponding $\\sigma_8$ are showed in each panel.} \\label{fig:halo_evo} \\ec \\end{figure} \\begin{figure} \\bc \\hspace{-1.4cm} \\resizebox{9cm}{!}{\\includegraphics{figures/mah.ps}}\\\\% \\caption{Mass assembly histories (MAH) for the 10 most massive haloes at $z=0$ in the HDMRUN sample. The red curve indicates the halo presented in Fig.~\\ref{fig:halo_evo}} \\label{fig:mah} \\ec \\end{figure}  \\begin{figure} \\bc \\hspace{-1.4cm} \\resizebox{9cm}{!}{\\includegraphics{figures/halo_evo2.ps}}\\\\% \\caption{Formation histories of two ``first'' objects in a HDM universe. The right column shows a halo with $N_{200}=2686$ at $z=2.9$ (roughly $M_{200}=3.5\\times 10^{14}M_{\\odot}$); while the left column shows another halo with $N_{200}=1215$ ($M_{200}=1.6\\times 10^{14}M_{\\odot}$) at $z=4.3$. All particles in both halos at the last redshift are traced back to higher redshifts and presented as blue points. Particles in the very inner region $r < 0.1 r_{200}$  are traced back and presented as red points. Other background particles are shown with black points. The coordinate is comoving scale, and the size of each box is $\\rm 6 Mpc/h$. The redshift is labeled in each panel. The simulation glass128 (see Paper I) is used here. Note the artificial discreteness features visible in the filaments in the $z=3.6$ panel.} \\label{fig:halo_evo2} \\ec \\end{figure} ",
        "conclusions": "\\label{sec:conclusions} In this paper , we study the haloes which form by monolithic collapse in HDM cosmological simulations, and we compare a range of properties of these haloes with these of their counterparts in a $\\Lambda$CDM hierarchical universe. From this comparison we explore which physical mechanisms are responsible for the universal profiles or distributions of halo properties. Our simulated HDM universe has an inherent characteristic scale below which the formation of small haloes is suppressed. In this universe, the first generation of halo forms in a ``top-down'' way -- more massive haloes form earlier by smooth accretion. Mergers occur in significant numbers only when building up later halo generations. The HDM cosmology thus provides a good example of monolithic halo growth from Gaussian initial condition. Furthermore, because the HDM characteristic scale is much more sharply imprinted than in the WDM case, we  have a large and well defined monolithic halo sample. As a result, we can explore  monolithic growth in a better controlled way than previous work based on WDM universes. We summarise our conclusions on monolithic halo formation  as follows:  (i) First generation haloes grow by smooth accretion. Such haloes form from roughly spherical collapse of matter around the nodes or ends of filaments. The formation histories of these monolithic haloes are also characterised by two phases: fast initial and slow later accretion, very similar to those found in a CDM universe. In the fast growth phase, smooth infall is the main mechanism rather than the mergers seen in the CDM case.  (ii) The density profiles of our monolithic haloes are well fit by the NFW profile. The concentration parameter $c$ and the characteristic density $\\rho_s$ have strong dependences on halo mass: the more massive a halo, the larger  $c$ and $\\rho_s$. This is the inverse of the CDM dependence, but the concentration-formation time relation is quite similar in the two cases: earlier formed haloes tend to be more concentrated than their later formed counterparts.  (iii) Phase-space density profiles and velocity dispersion anisotropy profiles are very similar in our two cases. This indicates that the kinematic structure of haloes is generated by wide variety of dynamical collapse processes, not just by mergers.  (iv) The distribution of shape parameter is very similar in HDM and CDM universes, as are the distributions of spin parameter $\\lambda'$ and of internal specific angular momentum.  (v) In our HDM simulations, most subhaloes result from the infall of spurious small haloes which form in the filaments as a result of numerical discreteness. The subhalo abundance is much lower than in CDM haloes, and indeed there are very likely no subhaloes in first generation HDM haloes.  The above results show that, except for substructure, the properties of haloes formed by monolithic collapse are very similar to those of haloes formed by hierarchical clustering in a CDM universe. In particular, they have very similar profiles of density, phase-space density and specific angular momentum, and very similar distributions of axial ratios and spin parameter. These results indicate that mergers are not responsible (or are  not required) for the universal properties of CDM haloes. This agrees with results from earlier work based on constrained simulations \\citep{huss99, macmillan06}. Attempts to explain these universal properties with merger-driven models seem unlikely to be on the right track.  The results in this study could also find some applications in realistic dark matter models, especially for the WDM cosmology. In our HDM universe, first generation haloes form by smooth near-spherical accretion with initial mass well below the characteristic free-streaming mass. These haloes form simultaneously with the filaments and sheets. This picture is different from that in the familiar CDM universe. Our results suggest that the structure of the first haloes, in a WDM universe (and thus, presumably, also of their later descendants) will differ little from the structure of CDM haloes, despite the differences in assembly history. This makes it seem unlikely that changing to a WDM model will help resolve the apparent problems of the CDM model on galactic scales."
    },
    "0809/0809.0865_arXiv.txt": {
        "abstract": "We develop a method of constraining the cosmic string tension $G\\mu$ which uses the Canny edge detection algorithm as a means of  searching CMB temperature maps for the signature of the Kaiser-Stebbins effect. We test the potential of this method using high resolution, simulated CMB temperature maps. By modeling the future output from the South Pole Telescope project (including anticipated instrumental noise), we find that cosmic strings with $G\\mu > 5.5\\times10^{-8}$ could be detected. ",
        "introduction": "At very early times it is believed that the universe underwent a series of symmetry breaking phase transitions which led to the formation of different types of topological defects. Among them are linear topological defects known as cosmic strings (for reviews see e.g. \\cite{2000csot.bookV,Hindmarsh:1994re,Brandenberger:1993by}). After creation, the cosmic strings form a random network of infinite strings and closed string loops, the arrangement of which evolves over time through string interactions. Cosmic strings can also have self-interactions that lead to the formation of closed loops via the exchanging of endpoints, or \\emph{intercommutation} \\cite{Shellard:1987bv}. When formed, cosmic string loops break off of the longer segments and continue to oscillate, losing energy via gravitational radiation, until eventually decaying. Infinitely long strings, on the other hand, cannot decay into gravitational radiation and survive indefinitely. The string network eventually approaches a scaling regime in which the number of strings crossing a given Hubble volume is fixed and the strings contribute some fraction of the total energy in the universe. The existence of a scaling solution is supported by independent numerical simulations of the evolution of the cosmic string network \\cite{Albrecht:1984xv,PhysRevLett.60.257,Allen:1990tv,Albrecht:1989mk}. The quantity which characterizes the cosmic strings is their tension, $\\mu$, which is equivalent to the mass per unit length. This tension is directly determined by the energy scale of the symmetry breaking during which the cosmic strings were formed. It is possible that cosmic strings could have been formed at many different epochs, meaning the tension of the cosmic strings can take a wide variety of values. When discussing cosmic strings it is more common to work with the dimensionless parameter $G\\mu$, where $G$ is Newton's constant. Until the late 1990's, cosmic strings were studied as potential seeds for structure formation \\cite{Turok:1985tt,Sato,Stebbins}. The eventual discovery of the acoustic peaks \\cite{Boomerang,WMAP} in the angular power spectrum of the CMB lead to cosmic strings being ruled out as the main origin of structure in favour of the inflationary paradigm, since the angular power spectrum predicted by cosmic strings consists of only a single broad peak \\cite{Periv,Albrecht,Turok}. Despite this, there currently exists a renewed interest in cosmic strings fueled by the study of different cosmological models in which their formation is generically predicted (see \\cite{Brandenberger:1988aj,Jones:2003da,Jeannerot:2003qv} for just a few possibilities). It has also recently been shown that a contribution of less than 10\\% of the observed CMB power on large scales coming from cosmic strings is acceptable \\cite{Pogosian:2003mz,Wyman:2005tu,Fraisse:2006xc,Seljak:2006bg,Bevis:2006mj,Bevis:2007gh,Pogosian:2008am}.  The current bounds on the cosmic string tension come from a variety of measurements. The gravitational waves emanating from many string loops at different times produce a stochastic background which is the focus of current interferometer and pulsar timing experiments. Pulsar timing, specifically, places a bound $G\\mu<10^{-7}-10^{-8}$ on the cosmic string tension \\cite{Damour:2004kw,Jenet:2006sv}. However, we note that in order to place a bound on $G\\mu$ using gravitational wave constraints one must make assumptions about the size of the loops which are formed in the string network, the probability that strings will intercommute when crossing, and even the string model under consideration. Whereas the scaling solution for the long string network is well established, the distribution of loops is uncertain by several orders of magnitude. Therefore, the strength of the bounds obtained by considering gravitational radiation from string loops can be questioned. A more robust bound on the tension comes from the angular power spectrum of the CMB (since the spectrum obtains an important contribution from the long string network). Assuming a scaling solution of the long string network with the parameters from numerical studies of cosmic string evolution, the string contribution to the angular power spectrum of the CMB was determined, and the results of these studies translate directly into a bound $G\\mu<5\\times10^{-7}$ \\cite{Wyman:2005tu,Fraisse:2005hu}.  Along with the above mentioned phenomena, there exists another observational signature unique to cosmic strings which could be directly detected, namely, linear discontinuities in the temperature of the CMB. This signature was first studied by Kaiser and Stebbins \\cite{Kaiser:1984iv} and is usually referred to as the KS-effect. This effect occurs because the space-time around a straight cosmic string is flat, but with a wedge, whose vertex lies along the length of the string, removed. The angle subtended by the missing wedge, $\\phi$, is determined by the tension of the cosmic string as \\cite{Vilenkin:1981zs} \\begin{equation} \\phi=8\\pi G\\mu\\,. \\end{equation} For an observer looking at a source while a cosmic string is moving transversely through the line of sight between the two, the photons passing from the source to the observer along one side of the string will appear to be Doppler shifted relative to those passing along the other side due to this non-trivial geometry (see Figure \\ref{string}). If the source that the observer is viewing happens to be the CMB, this effect will manifest itself as discontinuities in the microwave background temperature along curves in the sky where strings are located. The magnitude of the step in temperature across a cosmic string is \\begin{equation}\\label{KS} \\frac{\\delta T}{T}=8\\pi G\\mu\\gamma_sv_s\\,|\\hat{k}\\cdot(\\hat{v_s}\\times\\hat{e_s})|\\,, \\end{equation} where $v_s$ is the speed with which the cosmic string is moving, $\\gamma_s$ is the relativistic gamma factor corresponding to the speed $v_s$, $\\hat{v}_s$ is the direction of the string movement, $\\hat{e}_s$ is the orientation of the string and $\\hat{k}$ is the direction of observation \\cite{Moessner:1993za}. For cosmic strings formed in a phase transition in the early universe, the ``missing wedge'' produced by a string has, at time $t$, a finite depth given by the Hubble radius $H^{-1}(t)$ at that time \\cite{Joao}. Some work has already been dedicated to searching for the KS-effect in current CMB data \\cite{Jeong:2006pi,Lo:2005xt}, but a cosmic string signal was not found, leading to a constraint on the tension $G\\mu\\lesssim4\\times10^{-6}$.  \\begin{figure} \\begin{center} \\includegraphics[width=0.75\\linewidth]{geometry} \\caption{\\label{string} The geometry of the space-time near a cosmic string. Shown here is a slice of the space-time perpendicular to the orientation of the string. The coloured area represents a missing wedge with deficit angle $\\phi$, while the dashed lines represent the paths of photons travelling from a source to an observer and the arrow shows the direction of motion of the string.} \\end{center} \\end{figure}  In this work we implement a method of detecting the temperature discontinuities in the CMB produced by cosmic strings via the KS-effect using an edge detection algorithm commonly employed in image analysis, the Canny algorithm \\cite{Canny:1986aa}. The motivation behind this choice is clear since the cosmic strings literally appear as edges in the CMB temperature. Depending on the sensitivity of the edge detection algorithm to the temperature edges, a bound on the cosmic string tension could then be imposed. This work is a continuation of a previous study \\cite{Amsel:2007ki} which indicated that an edge detection method may lead to a significant improvement in the sensitivity to the presence of cosmic strings compared to previous direct searches for strings in the CMB. In this paper we improve the method proposed in \\cite{Amsel:2007ki} and we investigate its application to surveys with a different set of specifications than those examined in that initial work.  We are interested in the cosmic strings in the network that survive until later times. The times relevant to the production of an edge signature in the CMB are the time of last scattering until the present day. Based on the evolution of the network, cosmic strings are more numerous around the time of last scattering than later times. On today's sky, those strings correspond to an angular scale of approximately 1$^\\circ$. Therefore, an observation of the CMB with an angular resolution significantly less than 1$^\\circ$ is necessary in order to be able to detect the edges related to these strings. With this in mind, we also focus on the application of the edge detection method to high resolution surveys of the CMB, particularly the future data from the South Pole Telescope.  The South Pole Telescope (SPT) \\cite{Ruhl:2004kv} is a 10m diameter telescope being deployed at the South Pole research station. The telescope is designed to perform large area, high resolution surveys of the CMB to map the anisotropies. The SPT is designed to provide 1$^\\prime$ resolution in the maps of the CMB, making it ideal to search for the KS-effect. Based on previous results \\cite{Amsel:2007ki}, we believe that with such high resolution data our method could provide bounds on the cosmic string tension competitive with those of pulsar timing.  The remainder of this paper is arranged as follows: In section \\ref{secmap}, we discuss the CMB maps used in our analysis with a focus on the anisotropies coming from Gaussian fluctuations and cosmic strings. In Section \\ref{seccanny}, we outline the edge detection algorithm we are using, highlighting the details of our particular implementation. In Section \\ref{seccount}, we discuss how we quantify the edge maps output by the edge detection algorithm and we explain the statistical analysis used to determine if a significant difference has been detected. In Section \\ref{secresults}, we present the results of running the edge detection algorithm on CMB maps and the possible constraints on the cosmic string tension that could be applied. We finish in Section \\ref{secdiscuss} with a discussion of our results. ",
        "conclusions": "\\label{secdiscuss}  We have developed a method of searching for linear discontinuities in the microwave background temperature caused by the presence of cosmic strings along our line of sight to the surface of last scattering. The method which we have developed involves applying an edge detection algorithm to CMB temperature anisotropy maps in order to identify the effect of cosmic strings. We have applied our edge detection method to simulated CMB maps both including cosmic strings, and without cosmic strings, to test its ability to discriminate between the two. This then translates directly into a possible constraint on the cosmic string tension. In particular, we have focused on two different sets of simulations, one which mimics the future output coming from the SPT and one which corresponds to a theoretical survey which covers five times as much sky as the SPT with the same angular resolution. We find that the edge detection method could potentially place a bound on the cosmic string tension of $G\\mu<5\\times10^{-8}$ for a perfect CMB observation from the SPT and that this could be lowered to $G\\mu<3\\times10^{-8}$ for the larger survey\\footnote{Note that the dependence of the limit as a function of angular resolution was studied in \\cite{Amsel:2007ki}}. For more realistic simulations which include instrumental noise, we find that the potential bound corresponding to the SPT weakens by only a small amount to $G\\mu<5.5\\times10^{-8}$ while the possible bound corresponding to the theoretical survey does not change at all, and is still $G\\mu<3\\times10^{-8}$. We consider the constraint corresponding to the SPT specific simulations which include a component of instrumental noise to be the main conclusion of this work. This possible bound is approximately an order of magnitude better than those arising from other methods which use CMB observations and approximately two orders of magnitude better than those arising from other methods which search for the KS-effect. Therefore, we believe that using the output from the SPT along with the edge detection method has the potential to greatly improve the constraint on the cosmic string tension. This bound is not tighter than the constraints arising from current pulsar timing data, although it is competitive, falling directly within the range of values reported by different observations. Nevertheless, as mentioned in the Overview, we consider our method of constraining the tension to be more robust since we make less assumptions about some of the unknown parameters which describe the cosmic string network and its evolution. Therefore, we believe that the possible bound on $G\\mu$ given above would in fact represent a stronger constraint.  We conclude that instrumental noise does not have a very major effect on the ability of the edge detection method to identify the cosmic string signal. We believe that this is an indication that the thresholding with hysteresis performs as it should, since noisy pixels could destroy the edge signal by causing large fluctuations in the gradient magnitude. Furthermore, as one can see from Figure \\ref{comparehist}, the largest difference between histograms occurs at short lengths rather than longer lengths. The instrumental noise leaves this difference in the short edge signal between maps with and without strings relatively unchanged, since the probability of a particularly noisy pixel falling on a short edge, resulting in it being incorrectly detected by the Canny algorithm, is small compared to that for longer edges. While on the topic of instrumental noise, we reiterate that we have included only a simplified white noise component in our simulations. A more complex investigation of instrumental noise would include a low frequency piece which results in stripes appearing in the final map of the CMB. Based on the method described in this paper, it is clear that striping would be crucial, since it would result in maps with more edges than that predicted by the cosmological theory and this could be confused with the effect of cosmic strings. One redeeming feature of this type of low frequency noise is that the stripes which are introduced would lie along the scanning direction, thus, when dealing with actual SPT data, it may be possible to subtract this effect out of the final map or to simply ignore edges lying along the known scanning direction in the edge detection algorithm itself. No other systematic effects due to the instrumental scanning strategy have been included in the current analysis, nor have errors due to foregrounds in the microwave sky. In future work it would be useful to investigate all types types of noise as well as the removal strategies in more detail to determine if they could change the behavior of the edge detection method.  We also found that increasing the simulated survey size increases the statistical significance of the deviations between the histograms for similar values of the cosmic string tension. This behaviour is expected though, since more edge maps were used to compute the mean values in each of the histograms and one can see from Equation \\eqref{t} that the value of $t$ scales as $\\sqrt{n}$. While the p-values are smaller for similar tensions, the final constraint which can be levied by the larger survey is not drastically different from that corresponding to the SPT specific simulations. Increasing the survey size by 5 times only lowered the possible constraint by a factor of roughly $\\sqrt{5}$. Based on this result, we conclude that the survey size does not have a major influence on the ability of the edge detection method.  When generating the simulated CMB maps, we employed a toy model of the cosmic string network which includes only straight strings and no cosmic string loops. More detailed models of the network and its evolution have been developed in other works \\cite{Albrecht:1989mk,PhysRevLett.60.257,Allen:1990tv,Fraisse:2007nu} and can be implemented numerically. Therefore, one obvious way to improve the testing method we have outlined here would be to implement one of these more complex models which would in turn produce a more realistic map of the temperature anisotropies induced via the KS-effect. On the other hand, we stress that a change of this nature would come at a large computational expense. On a similar note, it may also be useful to develop a more robust method of combining the string induced temperature anisotropies with those coming from Gaussian fluctuations, to make sure that the final simulated map agrees with other observations. Furthermore, after applying the Canny edge detection algorithm to the CMB temperature anisotropy maps, we quantify the corresponding edge map by recording the length of every edge appearing in it. As mentioned in Section \\ref{seccount}, this is one of the simplest ways of describing the edge map, and it may be beneficial to investigate an alternative method of image comparison which provides a more powerful way of discriminating between the two edge maps. For now, we leave these improvements as the goal of future work.  While we have chosen to focus on the SPT in this work, the edge detection method is quite versatile and could be used with virtually any high resolution CMB survey. We conclude that this method presents a powerful and unique way of constraining the cosmic string tension which has the potential to perform better than current methods, or, at the very least, to provide a complimentary technique to those already in use."
    },
    "0809/0809.2051_arXiv.txt": {
        "abstract": "Evolved mass-losing stars such as Mira enrich the interstellar medium (ISM) significantly by their dust-rich molecular wind. When these stars move fast enough relative to the ISM, the interaction between the wind and ISM generates the structure known as the astropause (a stellar analog of the heliopause), which is a cometary stellar wind cavity bounded by the contact discontinuity surface between the wind and ISM. Far-infrared observations of Mira spatially resolve the structure of its astropause for the first time, distinguishing the contact surface between Mira's wind and the ISM and the termination shock due to Mira's wind colliding with the ISM. The physical size of the astropause and the estimated speed of the termination shock suggest the age of the astropause to be about $40,000$ yr, confirming a theoretical prediction of the shock re-establishment time after Mira has entered the Local Bubble. ",
        "introduction": "Mira is the archetype of asymptotic giant branch (AGB) stars, which contribute significantly to the enrichment of the interstellar medium (ISM) by a dust-rich molecular stellar wind (e.g.\\ \\citealt{sedlmayr94}). Hence, tracing mass-loss histories of these stars has a profound impact on knowledge of the composition of the ISM and its evolution in host galaxies. Mira is actually a wind-accreting binary system \\citep{rc85}: a mass-losing AGB star is the primary star (Mira A), dominating the system. Mira's wind is, therefore, a cool molecular wind from the primary, in which dust grains convert radiation pressure of the star into the outward momentum of the wind \\citep{kwok80,sss98}.  Previous investigations in CO line emission revealed that Mira's wind at $\\sim 4$ km s$^{-1}$ formed a circumbinary molecular envelope of $0\\farcm3$ in radius, whose otherwise spherical structure was punctuated by bipolar outflows at $\\sim 8$ km s$^{-1}$ \\citep{jos00,fong06}. Mira's space velocity is approximately $150$ km s$^{-1}$ nearly due south as calculated via a spherical trigonometric transformation \\citep{js87} using its proper motion \\citep{hip}, radial velocity \\citep{evans67} and the solar motion \\citep{db98}.  Hence, Mira's wind has been interacting with the ISM flow local to Mira for quite some time. This interaction left a turbulent wake of approximately 2 degrees long, observed in emission in the ultraviolet \\citep{m07} and 21 cm \\ion{H}{1} line \\citep{m08}. The ultraviolet emission is thought to be due to H$_2$ excited collisionally by hot electrons arising from the shocked ISM flow. The 21 cm line emission of atomic hydrogen is, therefore, expected from dissociation of H$_2$ into H.  This is the second case of the wind-ISM interaction around an AGB star after R Hya, in which a bow shock resulting from the interaction was seen in the far-infrared \\citep{ueta06}. While the H$_2$/\\ion{H}{1} wake of Mira has been attributed to the wind-ISM interaction owing to Mira's large space velocity, distinct shock layers of the interaction themselves have not been confirmed observationally, except for the possible bow shock structure seen in the ultraviolet \\citep{m07}. Thus, it has been unclear how the H$_2$/\\ion{H}{1} turbulent wake seen at the large spatial scale is physically related to the CO wind seen at the small spatial scale, except that bright H$_2$ streams \\citep{m07} are correlated spatially with CO bipolar outflows \\citep{jos00}.  In this letter, we present Mira's far-infrared image taken with the {\\sl Spitzer Space Telescope} ({\\sl Spitzer}; \\citealt{sst}), which resolves spatially the structure of Mira's astropause (a stellar analog of the heliopause; \\citealt{nf82}) for the first time. Resolved are (1) Mira's circumbinary dust envelope at the core of the astropause, (2) the termination shock where the stellar wind is slowed down by the ISM and (3) the astrosheath in which a turbulent wake assumes the characteristic cometary shape beyond the termination shock. Thus, the {\\sl Spitzer} $160\\micron$ data, in conjunction with previous multi-wavelength results, provide the most comprehensive view of the wind-ISM interaction for an AGB star to date. ",
        "conclusions": "Figure \\ref{map} displays versions of the {\\sl Spitzer}/MIPS image of Mira at $160\\micron$ in log-scaled false-color. These maps were made at the $12\\farcs0$ pix$^{-1}$ scale ($11\\farcm7 \\times 13\\farcm3$ field of view) with the mean sky coverage of $3.5 \\pm 1.9$ pixel$^{-1}$. The effect of saturation was not removed entirely by adjusting reads in the data ramp, and the $12\\farcs0$ pix$^{-1}$ scale turned out to be the minimum pixel scale for which all pixels register valid data. The one $\\sigma$ sensitivity and the average large-spatial-scale sky emission (the component removed during the median filter) were found to be 0.8 and $8.8 \\pm 1.0$ MJy sr$^{-1}$, respectively. The derived sky emission value was consistent with the estimated value of 6.5 MJy sr $^{-1}$, which is obtained with the {\\sl Spitzer} Planning Observations Tool (SPOT) software.\\footnote{http://ssc.spitzer.caltech.edu/documents/background/} Contours are 80, 60, 40, 20, 10, 5, 2.5 and 0.5$\\%$ of the emission peak, with the lowest contour corresponding to 3 $\\sigma$ ($= 2.5$ MJy sr$^{-1}$).  The light-leak-corrected image is presented in Figure \\ref{map}b, in which the emission core is evidently single-peaked. The leak-corrected emission profile of Mira can be dissected into three distinct emission regions: the emission core, plateau and halo. The core is slightly extended with the full-width at half-maximum (FWHM) of $1\\farcm0 \\times 0\\farcm9$ at the position angle (PA; measured East from North) of 70$^{\\circ}$. Since nearby point sources have the average FWHM of $0\\farcm7$, the extension of the core appears genuine. The plateau is an elliptical region ($3\\farcm4 \\times 2\\farcm8$ at the PA of 70$^{\\circ}$) of relatively flat profile encircling the core (at 5\\% to 30\\% of the peak emission). The halo is a fainter nebulosity surrounding the plateau. If we define the periphery of the halo at 3 $\\sigma$ (the lowest contour in Figure \\ref{map}), then its extent is $8\\farcm5 \\times 6\\farcm5$ at the PA of 15$^{\\circ}$. The scan-map direction is the PA of 160$^{\\circ}$, along which some residual ``streakings'' are still recognizable below 3 $\\sigma$. Hence, the elongation of the $160\\micron$ emission around Mira is not the artifact of scan mapping. Mira has been imaged in the far-IR most recently with the {\\sl AKARI Infrared Astronomy Satellite} \\citep{murakami07}. {\\sl AKARI}'s Mira maps are consistent with the {\\sl Spitzer} map with a single-peaked emission core with extension at the PA of 15$^{\\circ}$ (H.\\ Izumiura et al.\\ 2008, private communication).  The point-spread-function (PSF) in the MIPS $160\\micron$ band has the first Airy pattern out to $\\sim 7\\arcmin$ (at $12\\arcsec$ pixel$^{-1}$) at $\\sim 4\\%$ of the peak surface brightness. A synthesized PSF\\footnote{% Available at http://ssc.spitzer.caltech.edu/archanaly/contributed/stinytim/.} was convolved with an elliptical Gaussian function to conform the observed shape of the core and subtracted from the data (Figure \\ref{map}c). We immediately see that there is still substantial emission ($\\sim 50$ MJy sr$^{-1} = \\sim 20 \\sigma$) in the plateau, suggesting that it is caused by an elliptical ring of $2\\farcm4 \\times 2\\farcm0$ at the PA of 70$^{\\circ}$. The core is oriented in the same way as the circumbinary CO envelope \\citep{jos00}, while the CO envelope is smaller than the core \\citep{fong06}. At the distance of 107 pc to Mira \\citep{knapp03}, the physical size of the CO envelope is $\\sim 3 \\times 10^{16}$ cm in radius. This agrees well with the expected radius of Mira's molecular envelope resulted from mass loss at the rate of $\\sim 2 \\times 10^7$ M$_{\\odot}$ yr$^{-1}$ \\citep{jos00,rs01,fong06}, including the effect of photodissociation by the interstellar radiation field \\citep{mamon88}. Therefore, the core represents Mira's circumbinary dust envelope whose physical size is $\\sim 4.5 \\times 10^{16}$ cm in radius.  Since the Mira system is moving at $\\sim 150$ km s$^{-1}$ the resulting wind-ISM interaction has given rise to the ultraviolet bow shock structure \\citep{m07}, despite Mira's slow wind velocity of 4 km s$^{-1}$ \\citep{jos00,fong06}. Figures \\ref{over}a and \\ref{over}b clearly show that the downstream side of the ultraviolet bow shock coincides with the southern edge of the far-infrared halo. Moreover, the cometary shape of the halo is the most consistent with the bow shock interpretation, in which the far-infrared nebulosity is shaped by a turbulent wake flowing almost due north. According to the theory of the (solar)wind-ISM interaction, the ISM flow approaches supersonically toward Mira, becomes subsonic past the bow shock and streams around the astropause, while Mira's wind expands freely to form a bubble bounded by the termination shock, beyond which the wind becomes compressed and turbulent and flows downstream in the astrosheath \\citep{zank99}. Assuming that the $160\\micron$ emission probes Mira's dusty wind emanating from the circumbinary envelope, we interpret that the leak-corrected, core-subtracted image distinguishes spatially all components of the astropause (Figure \\ref{map}c) with the halo tracing the astrosheath and the emission ring delineating the termination shock (see, also, Figure \\ref{over}d).  The prominent ultraviolet streams (Figures \\ref{over}a and \\ref{over}b) are outflows of H$_{2}$ excited collisionally by hot electrons in the bow-shock-excited ISM flow \\citep{m07}. These H$_{2}$ flows, therefore, must have penetrated the termination shock and been directed downstream while getting excited in the astrosheath. In the astrosheath, H$_2$ can be dissociated by either the interstellar radiation field \\citep{mamon88} or collisional excitation \\citep{m07,m08}. The observed \\ion{H}{1} emission is mostly concentrated in the vicinity of the termination shock, with the peak FWHM of $1\\farcm8$ \\citep{m08}, which is compatible with the size of the molecular radius given the resolution at 21 cm (Figure \\ref{over}c). Similar to H$_2$, atomic H is directed toward downstream beyond the termination shock, lending strong support for the multi-phase characteristics of the flow in the astrosheath.  The stagnation point of the wind-ISM interaction is therefore at the apex of the halo (astropause), which is $3.6 \\times 10^{17}$ cm from Mira. Requiring ram pressure balance at the stagnation point, this implies that the ISM density local to Mira is $n_{\\rm ISM} = 0.02$ cm$^{-3}$. This value is consistent with the result of a two-wind model \\citep{wareing07} that accounts for Mira's entry into the Local Bubble, the region of warm ($\\sim 10^{6}$ K), tenuous ($\\sim 0.01$ cm$^{-3}$) ISM in the solar neighborhood \\citep{l01}. Given $N_{\\rm H I} \\le 1.3 \\times 10^{19}$ cm$^{-2}$ and 25\\% abundance of neutral species \\citep{m08}, we can estimate the optical depth at $160\\micron$ to be $8 \\times 10^{-6}$ assuming a power-index scaling law with the index of 1.2 (i.e. silicate dust). Under the assumption that thermal dust emission dominates at $160\\micron$, the emission map and the optical depth yield the temperature of dust; $\\sim 500 - 1100$ K in the circumbinary dust envelope, $\\sim 200$ K in the termination shock and $\\sim 30 - 200$ K in the astrosheath.  Based on these estimates, the pre-shock temperature of Mira's wind is thought to be $\\sim 50$ K (assuming gas-dust thermalization), implying Mach number of 7.5. With the generalized jump relations (e.g.\\ \\citealt{p81}), the velocity of the termination shock is $-2$ km s$^{-1}$ in Mira's frame and the ratio of post- to pre-shock densities 3.8. These values yield the post-shock gas temperature of $\\sim 900$ K. Thus, in the shock-excited astrosheath, dust is still present (the gas temperature is below the condensation temperature of silicates $\\approx 1200$ K; \\citealt{s77}), but probably de-thermalized from gas. Spectroscopic follow-ups are necessary to investigate the evolution of the dust-gas characteristics upon passage of the termination shock.  At 2 km s$^{-1}$, the termination shock should have taken $\\sim 40,000$ yr to advance to the present position ($1 \\times 10^{17}$ cm) from the astropause ($3.6 \\times 10^{17}$ cm). The total flux of the astrosheath (including the termination shock) is 37 Jy, and this translates to the total mass of $7 \\times 10^{-3}$ M$_{\\odot}$ using the dust emissivity of generic interstellar dust \\citep{dl07}. With Mira's mass loss rate ($\\sim 2 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$; \\citealt{jos00,fong06,rs01})\\footnote{Various authors quote a range of mass-loss rates for Mira, which is a variable star. In the present paper, however, we are concerned with the long-term effects of Mira's mass loss.  Therefore, we adopt the average rate of mass loss in our discussion.}, it should have taken $\\sim 40,000$ yr for the astrosheath to form. These estimates for the age of the termination shock is consistent with the time needed for the termination shock to re-establish itself after entering the Local Bubble ($\\sim 40,000$ yr) as estimated numerically by \\citet{wareing07}, who attributed a kink in Mira's ultraviolet tail \\citep{m07} to Mira's entry into the Local Bubble.  The presence of the bow shock indicates that the outer shock is supersonic. In the warm ($\\sim 10^6$ K), tenuous ($\\sim 0.01$ cm$^{-3}$) Local Bubble \\citep{l01}, however, the shock velocity of $\\sim 150$ km s$^{-1}$ is barely supersonic. This suggests that the pre- and post-shock parameters across the bow shock are nearly continuous. This is consistent with the fact that the ultraviolet emission profile does not indicate any density enhancement (i.e.\\ material pile-up) at the southern edge of the bow. Hence, ultraviolet emission in front of the upstream face of the astropause is due to excitation of the post-shock ISM gas by a shear flow. While Mira's wind and the ISM gas flow are in general separated by the astropause, these flows are inherently turbulent. Thus, these flows are mixed in the downstream wake and H$_2$ is collisionally excited, resulting in the observed far-ultraviolet emission. Excited H$_2$ can be dissociated, leading to the observed \\ion{H}{1} emission. Dust grains help H$_2$ to replenish in the wake, and this cool, turbulent three-phase gas flow results in emission of H$_2$ and \\ion{H}{1} continuing at least for 1.5$^{\\circ}$ long.  Hence, there is an emerging picture of Mira's astropause (Figure \\ref{over}d). The extent of the far-infrared emission outlines the astropause, which is the contact surface between Mira's molecular wind and the ISM flow. The bulk of the halo emission is due to dusty molecular wind material in the astrosheath, shock-excited by the termination shock delineated by the emission ring. The termination shock is the discontinuous interface between the pre- and post-shock regions of Mira's wind, inside which Mira's wind expands undisturbed coming off the circumbinary envelope hosting bipolar outflows.  Since the post-shock regions of both Mira's wind and the ISM flow have a significant physical size, the shock must be non-radiative (i.e., it has not been cooled down and compressed into an unresolvable size). According to \\citet{cloudy}, in a collisionally excited solar wind abundance gas heating dominates cooling below 10$^4$ K (their Fig.\\ 6). The estimated post-shock gas temperature of Mira's wind is $\\sim 900$ K. Thus, the post-shock gas can indeed maintain its temperature, i.e., its physical size. Because previous numerical models (\\citealt{wareing07,raga08}) were concerned mainly with the structure of the turbulent wake of the wind-ISM interaction, follow-up multi-phase hydrodynamical modeling at low temperatures to reproduce the structure of the astropause appears extremely worthwhile.  The total flux of the astropause is 52 Jy, which implies the total mass of $1 \\times 10^{-2}$ M$_{\\odot}$ using the dust emissivity of generic interstellar dust \\citep{dl07}. Since the total atomic mass in the astropause is $1.3 \\times 10^{-3}$ M$_{\\odot}$ \\citep{m08}, the total molecular mass is estimate to be $8.7 \\times 10^{-3}$ M$_{\\odot}$ to be roughly 1 to 9 atomic-to-molecular mass ratio. The observed far-ultraviolet luminosity implies the H$_{2}$ dissociation rate of $\\sim 2.5 \\times 10^{42}$ s$^{-1}$, if collisional excitation accounts for all far-ultraviolet emission \\citep{m07}. This rate, however, would have yielded a factor of four more atomic H than observed \\citep{m08}. Dust grains can potentially provide H$_{2}$ formation sites to reduce excess atomic H. Using parameters of gas in the astrosheath ($n_{\\rm H} \\approx 7$ and $T_{\\rm gas} \\approx 900$ K) and silicate grains ($0.1\\micron$ radius, 3.7 g cc$^{-1}$, 1\\% mass ratio), the H$_2$ formation rate is approximately $4 \\times 10^{40}$ s$^{-1}$, if one assumes that an encounter with two atomic H by a dust grain always results in production of H$_2$. This suggests that H$_2$ far-ultraviolet emission is not entirely due to collision of H$_2$ with hot electrons and by some other means as already implied \\citep{m07,m08} and/or there exist some mechanisms to bolster production of H$_2$ from atomic H.  Clearly, Mira's astropause presents a unique laboratory for investigations of a multi-phase hydrodynamical flow, turbulence and dust-gas processing, which has a profound impact on the chemical makeup of the ISM. Because recent findings suggest that wind-ISM interaction around AGB stars may be common \\citep{ueta06,ueta07}, spatially resolved spectroscopy of multi-phase dusty, low-temperature gas flow will be extremely productive in the coming era of Herschel, SOFIA and ALMA (e.g.\\ \\citealt{herschel,sofia,alma}, respectively)."
    },
    "0809/0809.2790_arXiv.txt": {
        "abstract": "The redshift evolution of the Tully-Fisher Relation probes gravitational dynamics that must be consistent with any modified gravity theory seeking to explain the galactic rotation curves without the need for dark matter. Within the context of non-relativistic Modified Newtonian Dynamics (MOND), the characteristic acceleration scale of the theory appears to be related to the current value of either the Hubble constant, i.e., $\\alpha\\simeq cH_0$, or the dark energy density, i.e., $\\alpha\\simeq (8\\pi G\\rho_\\lambda/3)^{1/2}$.  If these relations are the manifestation of a fundamental coupling of $a_0$ to either of the two cosmological parameters, the cosmological evolution would then dictate a particular dependence of the MOND acceleration scale with redshift that can be tested with Tully-Fisher relations of high-redshift galaxies.  We compare this prediction to two sets of Tully-Fisher data with redshifts up to $z=1.2$. We find that both couplings are excluded within the formal uncertainties. However, when we take into account the potential systematic uncertainties in the data, we find that they marginally favor the coupling of the MOND acceleration scale to the density of dark energy. ",
        "introduction": "The relation between luminosity and spectral line widths in spiral galaxies has been widely established as a probe of the dark matter content of these galaxies (Tully \\& Fisher 1977). This, so called Tully-Fisher Relation (TFR), has subsequently been constructed in various forms with the stellar mass or baryonic mass of a galaxy serving as a proxy for luminosity, and similarly, with the maximum rotational velocity or asymptotic rotation velocity as a substitute for line width. Of the many parameter combinations, the tightest TFR fits occur between baryonic mass and asymptotic rotation velocity (Verheijen 2001).  Galactic velocity profiles that are asymptotically flat appear ubiquitous, yet they were not expected from Newtonian gravitational models of the baryonic matter distribution. Additional sources of the gravitational field, such as dark matter halos (Rubin et al 1985), or modifications to the gravitational theory (e.g., Milgrom 1984) are necessary for their explanation. A successful explanation of the galactic rotation curves must also be compatible with the Tully-Fisher Relation because of the latter's dependence on the asymptotic rotational velocities of galaxies.  The Modified Newtonian Dynamics (MOND) paradigm was first developed as a modification to gravity at low accelerations for spiral galaxies and an alternative to dark matter (Milgrom 1984; for a thorough treatment of MOND see Milgrom 2001, 2002b, 2008 and Sanders \\& McGaugh 2002). One of its early successes has been the fitting of galaxy rotation curves with a parametric function that is determined only by the amount of visible matter in the galaxy and a single parameter $\\alpha_0$ which represents the characteristic acceleration below which gravity becomes non-Newtonian (Bekenstein \\& Milgrom 1984).  Non-relativistic MOND modifies Newtonian gravity such that the net acceleration experienced by a test particle in a gravitational field is $a_{\\rm MOND}=(a_{{\\rm}N}\\alpha_0)^{1/2}$, where $a_{{\\rm}N}$ is the Newtonian gravitational acceleration. This phenomenological model also leads naturally to the Tully-Fisher Relation, since the velocity $V_{\\rm f}$ of an asymptotically flat rotation curve at large distances is related to the baryonic mass $M$ of the galaxy by (Milgrom 1984) \\begin{equation} M=\\frac{V_{\\rm f}^4}{G \\alpha_0}\\;. \\label{TFMond} \\end{equation}  The most unexpected result of fitting rotation curves with MOND is that the acceleration parameter is nearly the same for all galaxies. Moreover, its value is comparable to two combinations of cosmological parameters that have dimensions of acceleration (Milgrom 2001), i.e., \\begin{equation} \\alpha_0\\simeq c H_0 \\simeq c\\sqrt{\\frac{8\\pi G \\rho_\\lambda}{3}}\\;, \\label{coincidence} \\end{equation} where $H_0$ is the value of the Hubble constant and $\\rho_\\lambda$ is the cosmic density of dark energy. In the context of the Cold Dark Matter (CDM) paradigm, this similarity is believed to arise from the cosmological evolution of dark matter halos, which naturally depends on the cosmological expansion. Analytic attempts to account for the MOND acceleration scale within the CDM framework have been indeed made (Kaplinghat \\& Turner 2002) but also refuted (Milgrom 2002a).  MOND predictions have been examined in several other settings, from clusters of galaxies (e.g., Clowe et al.\\ 2006) to larger cosmological scales (e.g., Scott et al.\\ 2001). These studies have established the difficulties of using the MOND phenomenology to account for this wide set of cosmological phenomena. Even without focusing on these issues, if the MOND phenomenology describes a deviation from Newtonian gravity at galactic scales, the numerical similarity of the various physical quantities shown in relation (\\ref{coincidence}) will need to be explained.  MOND is based on a non-relativistic phenomenological function and, therefore, it cannot lead to predictions for the cosmological evolution of the Hubble parameter or of the acceleration scale $\\alpha_0$. Different attempts to devise a relativistic equivalent to MOND lead to varying predictions. For example, in TeVeS, which is a Tensor-Vector-Scalar relativistic theory that accounts for a large part of the MOND phenomenology, the similarity described by equations~(\\ref{coincidence}) is coincidental (Bekenstein \\& Sagi 2008). On the other hand, in a recently proposed theory which attempts to account simultaneously for the dark matter and the dark energy problem (Zhao \\& Li 2008), the acceleration scale couples to the density of dark energy.  If the connection between the acceleration $\\alpha_0$ and either the Hubble parameter $cH$ or the density of dark energy $\\sqrt{\\rho_\\lambda}$ can be proven experimentally, then this will provide important clues towards developing a successful relativistic theory that is consistent with the MOND phenomenology. Indeed, within general relativity, the Einstein equivalence principle guarantees that the cosmological expansion will affect the motion of a test particle only to second order in $(cH_0)^2$ (Peebles 1993). Therefore, the acceleration scale $\\alpha_0$ will couple linearly to the Hubble parameter only if the Einstein equivalence principle is violated at cosmological scales. In this case, MOND will be a modification of inertia rather than of the gravitational field (see Milgrom 2006). On the other hand, the acceleration scale $\\alpha_0$ will couple to the dark energy only if the latter is not a manifestation of a cosmological constant but rather of a dynamical field that couples non-linearly to the metric. This is true because a cosmological constant or a minimally coupled field that describes dark energy introduces a characteristic scale in curvature and not in acceleration. As a result, the relativistic version of MOND in this case will require a modification of the Einstein field equations.  The evolution of the MOND acceleration scale with redshift provides the means with which we can test, in a phenomenological way, the connection represented in equation~(\\ref{coincidence}) (see Milgrom 2008 and Bekenstein \\& Sagi 2008).  This is true because the rate of the cosmological expansion $H\\equiv \\dot{a}/a$ and the density of dark energy have distinct dependences on redshift. In this paper, we use observations of the so-called Tully-Fisher intercept back to a redshift of $z=1.2$ (Weiner et al. 2006; Kassin et al. 2007) as a proxy for the MOND acceleration scale (see Eq.~[\\ref{TFMond}]). This allows us to infer the evolution of $\\alpha_0$ using only the exact phenomena that non-relativistic MOND was designed to explain, without the need for extrapolating the phenomenology to include a treatment of photon trajectories. We then compare the observations to the expected evolution of the MOND acceleration scale when coupled to the Hubble parameter or the density of dark energy, for values of the cosmological parameters that are consistent with recent measurements (Spergel et al. 2007). ",
        "conclusions": "In this paper, we calculated the time-evolution of the MOND acceleration parameter $\\alpha_0$ when coupled to either the Hubble parameter or to the density of dark energy. For a dark energy equation of state consistent with recent data, we expect a positive evolution of $\\alpha_0$ with redshift if it couples to the Hubble parameter and little to no evolution with redshift if it couples to the dark energy density. The latest measurements of the Tully-Fisher Relation support neither coupling scenario within formal uncertainties. Observations to $z=1.2$ indicate a decrease in both the baryonic mass and J-band luminosity Tully-Fisher intercepts, which are coupled to the MOND acceleration parameter as shown in equations~(\\ref{hubble_coupling}) and~(\\ref{de_coupling}). However, the systematic uncertainties introduced by the initial mass function, stellar formation and stellar population evolution affect the data significantly. Sources of systematic uncertainty tend to shift the Tully-Fisher intercept data upward, so that the net trend has a slight downward evolution with redshift. The corrected data would then be consistent with a coupling of $\\alpha_0$ to the dark energy density with an observationally allowed equation of state with $w\\simeq -1$. However, the uncertainties in such a large correction for systematics also makes no evolution in the TF intercept plausible.  Definitive conclusions can only be drawn after the systematic errors are more precisely quantified. Future observations from large telescopes both on the ground and in space of galaxies at higher redshifts will be particularly helpful, since the difference between the MOND acceleration scale evolving coupled to the dark energy density or the Hubble parameter increases with redshift."
    },
    "0809/0809.0890.txt": {
        "abstract": "Based on our intensive spectroscopic campaign with the GoldCam spectrograph on the Kitt Peak National Observatory (KPNO) 2.1-m telescope, %We investigate the probability of existence of the companion galaxies %physically associated with E+A galaxies. we have constructed the first catalogue of E+A galaxies with spectroscopic companion galaxies, and investigated a probability that an E+A galaxy have %physically associated close companion galaxies. We selected 660 E+A galaxies with $4.0{\\rm \\AA} < {\\rm H}\\delta~{\\rm EW}$ at a redshift of $<0.167$ %, %and comparison sample of 11267 normal galaxies at %a redshift of $0.057 \\leq z \\leq 0.062$ %so that the sample satisfies the volume-limited condition %with $-22.5 < M_r < -19.5$, from the Data Release 5 of the Sloan Digital Sky Survey (SDSS). We selected their companion candidates from the SDSS imaging data, and classified them into true companions, fore/background galaxies and companion candidates using the SDSS and our KPNO spectra. We observed %9 %(+ 9  14 2 1) 26 companion candidates of E+A galaxies at the KPNO %using the GoldCam spectrograph %at the Kitt Peak National Observatory (KPNO) 2.1-m telescope, to measure their redshifts. Their spectra showed that %3 %(+ 3  14) 17 targets are true companion galaxies. %As a result, The number of spectroscopically-confirmed E+A's companions are now 34. This becomes the first catalogue of E+A galaxies with spectroscopic companion systems. We found that E+A galaxies have an 54\\% larger probability of having companion galaxies (7.88\\%) compared to the comparison sample of normal galaxies (5.12\\%). A statistical test shows the probabilities are different with 99.7\\% significance. %using statistics based on %the spectroscopic data taken from SDSS, our KPNO observations, %and the SDSS imaging data. %Our strict statistics using both spectroscopic and imaging data %enhances the reliance of discussion about %supports excess %show that E+A galaxies have excess number companions galaxies. Our results based on spectroscopy tightens the connection between the dynamical merger/interaction and the origin of E+A galaxies. ",
        "introduction": "\\label{section:intro}   %Dressler \\& Gunn (1983; 1992) \\citet{dre83,dre92} discovered galaxies with mysterious spectra %while %investigating high-redshift cluster galaxies. in high-redshift clusters of galaxies. These galaxies had strong Balmer absorption lines with no emission in ${\\rm [OII]}$. They were named ``E+A'' galaxies because their spectra resembled a superposition of those of elliptical galaxies (${\\rm Mg}_{5175}$, ${\\rm Fe}_{5270}$ and ${\\rm Ca}_{3934,3468}$ absorption lines) and A-type stars (Strong Balmer absorption)% \\footnote{Because the spectra of elliptical galaxies are characterised by K stars, these galaxies are sometimes called ``k+a'' galaxies \\citep[e.g.,][]{fra93,dre99,bar01}. Following the first discovery, we refer to them as ``E+A'' throughout this paper. }% . The existence of strong Balmer absorption lines shows that these galaxies have experienced starbursts within the last Gyr, since the lifetime of an A-type star is about 1 Gyr. However, they show no sign of ongoing star formation as indicated by non-detection in the ${\\rm [OII]}$ emission line. Therefore, E+A galaxies are interpreted as post-starburst galaxies, i.e., galaxies that have undergone truncated starburst activity %(Dressler \\& Gunn 1983, 1992; Couch \\& Sharples 1987; %MacLaren, Ellis \\& Couch 1988; %Newberry Boroson \\& Kirshner 1990; %Fabricant, McClintock \\& Bautz 1991; %Abraham et al. 1996). \\citep{dre83,dre92,cou87,mac88,new90,fab91,abr96}. Thus, E+A galaxies have attracted a great deal of attention, and some explanations for their origin have been proposed.   In the classical studies, %At first, ``E+A'' galaxies were found in cluster regions in both low-redshift clusters \\citep{fra93,cal93,cal97,cas01,ros01} and high-redshift clusters \\citep{sha85,lav86,cou87,bro88,fab91,bel95,bar96,fis98,mor98,cou98,dre99}. Therefore, a cluster-specific phenomenon %\\citep[e.g.][]{got03b} %(e.g., Goto et al. 2003b) was thought to be responsible for the violent star formation history of E+A galaxies. A ram-pressure stripping model \\citep{spi51,gun72,far80,ken81,aba99,fuj99,qui00,fuj04,fg04} as well as tides from the cluster potential \\citep[e.g.,][]{fuj04} may first accelerate star formation of cluster of galaxies and later turn it off. However, recent large surveys of the nearby Universe found many E+A galaxies in the field regions \\citep{got03,got03b,got05,bla04,qui04,hog06}. %therefore, %At the very least, It is obvious that these E+A galaxies in the field regions cannot be explained by a physical mechanism that works in the cluster region. E+A galaxies have often been thought to be transitional objects during cluster galaxy evolution, involving phenomena such as the Butcher-Oemler effect \\citep[e.g.,][]{but78,rak95,mar01,ell01,kod01,got03a}, %%%%(Butcher, Oemler 1978, 1984; Rakos, Schombert 1995; Couch et al. 1994,1998; %%%%  Margoniner, de Carvalho 2000; Margoniner et al. 2001; Ellingson  et %%%%  al. 2001; Kodama, Bower 2001; Goto et al. 2003a)% the morphology-density relation \\citep[e.g.,][]{dre80,pos84,fas00,got03d,pos05,smi05}, %%%% (Dressler 1980; Dressler et al. 1997; %%%% Postman, Geller 1984;  Fasano et al. 2000; Goto et al. 2003c; %%%% Smith et al. 2005; Postman et al.2005) and the correlation between various properties of the galaxies and the environment \\citep[e.g.,][]{tan04,pop07}. However, explaining cluster galaxy evolution using E+A galaxies may no longer be realistic.  One explanation for E+A phenomena is dust-enshrouded star formation, where E+A galaxies are actually star-forming, but emission lines are invisible in optical wavelengths due to heavy obscuration by dust \\citep[e.g.,][]{pog00}. %As a variant, \\citet{pog00} presented the selective dust extinction %hypothesis, where dust extinction is dependent on stellar age, %since the youngest stars inhabit very dusty star-forming HII regions while %older stars have had time to migrate out of such dusty regions. %If O, B-type stars in E+A galaxies are embedded in dusty regions and %only A-type stars have long enough lifetimes ($\\sim$1 Gyr) to move out %from such regions, this scenario can naturally explain the E+A %phenomena. A straightforward test for this scenario is observation in radio wavelengths in which dust obscuration is negligible. At 20-cm radio wavelengths, synchrotron radiation from electrons accelerated by supernovae can be observed. Therefore, in the absence of a radio-loud active nucleus, the radio flux of a star-forming galaxy can be used to estimate its current massive star formation rate (SFR) \\citep{con92,ken98,hop03}. %\\citet{sma99} %performed such a radio observation and found that among 8 galaxies %detected in radio, 5 %have strong Balmer absorption with no %detection in ${\\rm [OII]}$. %They concluded that massive stars are currently forming in these 5 %galaxies. %\\citet{owe99} investigated the radio properties of galaxies %in a rich cluster at $z \\sim 0.25$ (A2125) and found that optical line %luminosities (e.g., ${\\rm H}\\alpha$+[NII]) were often weaker than one %would expect for the SFRs implied by the radio emission. % \\citet{sma99} observed z=0.4 cluster in radio, and found that 5 out of 10 radio sources show E+A like spectra in optical wavelength. \\citet{cha01} observed 5 nearby field E+A galaxies and detected no radio continuum. % \\citet{mil01} observed radio continua of 15 E+A galaxies and detected moderate levels of star formation in only 2 of them. \\citet{got04a} undertook 20-cm radio continuum observation of 36 E+A galaxies and none of them were detected at 20-cm, suggesting that E+A galaxies are not dusty-starburst galaxies. %(See also \\citet{cha01} but \\citet{buy06}). %(See also \\citealt[][]{cha01}).  %The star formation rates (SFRs) of these two galaxies, however, %are 5.9 and 2.2 M$_{\\odot}$ ${\\rm yr}^{-1}$, %which are an order of magnitude smaller than those in %\\citet{sma99}, consistent with normal to low SFR instead of starburst.  It is well-known that a galaxy--galaxy interaction triggers off explosive star formation %in the pair of galaxies \\citep[e.g.,][]{sch82,lav88,liu95a,liu95b,sch96,nik04}. % \\citet{oeg91} found a nearby E+A galaxy with a tidal tail feature. High-resolution Hubble Space Telescope imaging supports the galaxy--galaxy interaction scenario by showing some post-starburst (E+A) galaxies in high-redshift clusters as having disturbed or interacting signatures \\citep{cou94,cou98,dre94,oem97}. % \\citet{nor01} performed long-slit spectroscopic observations of 21 E+A galaxies, and found that young stellar populations of E+A galaxies are more centrally concentrated than older populations, and old components of E+A galaxies conform to the Faber--Jackson relation. % \\citet{bar01} reported that E+A galaxies, on average, tend to have slightly bluer radial gradients toward the centre compared to normal early-type galaxies. % \\citet{yan04} presented HST observations of the five bluest E+A galaxies with $z \\sim 0.1$ and reported details of disturbed morphologies. Moreover, \\citet{yan04} detected compact sources associated with E+A galaxies consistent with the brightest clusters in nearby starburst galaxies. % \\citet{yam05} not only showed obvious bluer radial colour gradients but also found irregular structures in their two-dimensional (2-D) colour map. % %The spectroscopy also support the merger/interaction origin of E+A galaxies. % %\\citet{liu95a,liu95b} observed 40 merging/interacting systems and found %that some of their spectra resemble E+A galaxies. % \\citet{yag06a} and \\citet{yag06b} investigated the age distribution using long-slit spectroscopy and found a positive gradient in the age of young stars from the centre to the outer regions of the plume. % Recently, some numerical simulations on E+A galaxies have also been presented. % \\citet{bek01} modelled galaxy--galaxy mergers with dust extinction, confirming that such systems can produce spectra that evolve into E+A spectra. % \\citet{bek05} investigated the structural, kinematical and spectrophotometric properties of E+A galaxies, and showed that the 2-D distributions of line-of-sight velocity, velocity dispersion, colour and line index in E+A galaxies formed via the interaction and merging of two gas-rich spirals.  Thus, many studies have been conducted on the large-scale environments and internal properties of E+A galaxies, and recent studies basically support the merger/interaction origin of E+A galaxies. %There is a trend to study detail of internal properties %We can find research for E+A galaxies with a new point of view %\\citep{buy06}, %and the E+A observation using X-ray is also beginning. %Such new studies of E+A galaxies %itself %are important, %However, %we might be missing an important one --- Several studies have focused on the medium-scale environments of E+A galaxies, %the only works that we are aware of them are e.g., \\citet{bla04}, \\citet{got03} and \\citet{got05}. %Actually \\citet{got03,got05} Using SDSS imaging data, \\citet{got03} found that young E+A galaxies have more accompanying galaxies within 100 kpc. %As \\citet{got03,got05} pointed out, there is %an excess in number of accompanying galaxies of E+As using the SDSS %imaging data. %Figure \\ref{fig:nearest_e+a} shows four example images of the E+A %galaxies, it is easily recognized even by eye that these E+A galaxies %have close accompanying galaxy. Results in \\citet{got03} and \\citet{got05} provide strong support for the merger/interaction origin of E+A galaxies, and we are interested in the physical relation between E+A galaxies and their accompanying galaxies. However, their studies are based on the imaging data, and not all accompanying galaxies are spectroscopically observed in the SDSS. In addition, it is unknown which galaxy is a real companion of E+A galaxies in \\citet{got03} and \\citet{got05}.  %\\begin{figure} %\\includegraphics[scale=0.61]{4_ea_systems.eps} %\\caption{Example images of E+A and companion systems %which we found in this study. %An `E+A and E+A' pair is also found in the SDSS data. %Each image size is inlaid. %}\\label{fig:nearest_e+a} %\\end{figure}  Our aim is to extract some clues from medium-scale environments, that is, the physical relations between E+A galaxies and their pair galaxies, to investigate their evolution. %physical properties %might be concealed around medium scale of environments %which belong E+A galaxies. However, we must confirm %Our 1st strategy for this aim is to confirm whether accompanying galaxies in \\citet{got03} and \\citet{got05} are companions, using statistical analyses with spectroscopic data. % before %research of our main %investigating subject. We will begin our studies on the properties of E+A galaxies and pair systems after such verification.  In this paper, we use publicly available {\\it true} E+A galaxies (without ${\\rm H}\\alpha$ nor ${\\rm [OII]}$ emission) selected from the Sloan Digital Sky Survey %(SDSS, York et al. 2000; %Early Data Release, Stoughton et al. 2002; %First Data Release, Abazajian et al. 2003; %Second Data Release, Abazajian et al. 2004; %Third Data Release, Abazajian et al. 2005; %Fourth Data Release, Adelman-McCarthy et al. 2006; %Fifth Data Release, Adelman-McCarthy et al. 2007, hereafter DR5) (SDSS, \\citealt[][]{yor00}; Early Data Release, \\citealt[][]{sto02}; First Data Release, \\citealt[][]{aba03}; Second Data Release, \\citealt[][]{aba04}; Third Data Release, \\citealt[][]{aba05}; Fourth Data Release, \\citealt[][]{ade06}; Fifth Data Release, \\citealt[][]{ade07}, hereafter DR5) by \\citet{got07b}. Both broadband imaging and the spectroscopic survey of 10,000 deg$^2$ of SDSS provide us with the first opportunity to study a very large number of E+A galaxies. \\citet{got07b} analysed $\\sim$670,000 galaxy spectra in the DR5, and the number of %E+A homogeneous E+A galaxies reached 1,062. We created the catalogue of companions/candidates of E+A galaxies using the SDSS SQL service, and we observed the spectra of companion candidates to reveal their redshifts. We then investigated an existing probability of companions of E+A galaxies by a stricter analysis with spectroscopic data. %before studying the relation between the E+A galaxies and companions.  This paper is organised as follows. In section \\ref{section:samples}, the definitions of two samples of %parent galaxies are summarised. %In section \\ref{section:catalogue}, %we explain the construction of companions/candidates catalogue, In section \\ref{section:observation}, we briefly describe our spectroscopic observations on companion candidates of E+A galaxies, and reduction of their data. In section \\ref{section:method}, we explain a somewhat complicated method of statistical analysis for readers to more easily understand. In section \\ref{section:analysis_and_results}, we show the application of the method described in section \\ref{section:method} to both the E+A sample and control sample, and present the results. Lastly, we provide a discussion and summary in section \\ref{section:discussion}.  Unless otherwise stated, we adopt the best-fitting {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} cosmology: $(h,\\Omega_m,\\Omega_L) = (0.71,0.27,0.73)$ \\citep{ben03,kom08}. %Komatsu et al. (2008). %The cosmological %parameters adopted throughout this paper are %$H_0=71 {\\rm km s}^{-1} {\\rm Mpc}^{-1}$, and %$(\\Omega_m,\\Omega_{\\Lambda},\\Omega_k)=(0.27,0.73,0.0)$.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%    \\begin{figure*} \\begin{center} %\\begin{tabular}{cc} % \\includegraphics[scale=0.285,angle=-90]{new_spec/1_target08a_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/1_target06a_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/0_aa08-1_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/0_targetA04_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/4_targetA02_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/8_targetB03_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/3_targetA07_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/5_targetA09_trans_bg.ms.4bin.txt.eps} \\\\ %\\end{tabular} \\includegraphics[scale=0.30]{spectra-a.eps} \\end{center}  \\caption{Spectra of 16 new E+A's companions taken with the KPNO 2.1-m telescope applying a 20\\AA~binning. A sky spectrum in the observed data is shown at the bottom of each panel. This figure includes targets \\#1({\\it top left}), \\#2({\\it top right}), ..., \\#7({\\it bottom left}) and \\#8({\\it bottom right}). %The target No. are 1, 2, and 3 from top to bottom. %The spectrum of Tucson's sky is also plotted, which is provided by \\citet{mas00}. }\\label{fig:spectra} \\end{figure*}  \\begin{figure*} \\begin{center} %\\begin{tabular}{cc} % \\includegraphics[scale=0.285,angle=-90]{new_spec/7_targetA05_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/a_targetB05_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/1_targetB04_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/4_targetA01_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/8_targetB08-2_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/b_targetD09_trans_bg.ms.4bin.txt.eps} \\\\ % \\includegraphics[scale=0.285,angle=-90]{new_spec/1_targetB06_trans_bg.ms.4bin.txt.eps} & % \\includegraphics[scale=0.285,angle=-90]{new_spec/9_targetE06_trans_bg.ms.4bin.txt.eps} \\\\ %\\end{tabular} \\includegraphics[scale=0.30]{spectra-b.eps} \\end{center} \\contcaption{ targets \\#9({\\it top left}), \\#10({\\it top right}), ..., \\#15({\\it bottom left}) and \\#16({\\it bottom right}). } \\end{figure*}  \\begin{figure*} \\begin{center} %\\begin{tabular}{cc} % \\includegraphics[scale=0.285,angle=-90]{new_spec/0_targetE40_trans_bg.ms.4bin.txt.eps} & \\\\ %\\end{tabular} \\includegraphics[scale=0.30,angle=-90]{spectra-c.eps} \\end{center} \\contcaption{ target \\#17. } \\end{figure*}    \\begin{table*} \\begin{center} \\caption{List of E+A and companion pairs. Columns $r$ and $M_r$ are magnitudes and absolute magnitudes with reddening- and {\\it K}-correction by {\\tt kcorrect.v4\\_1\\_4} \\citep{bla03}. Column elements of the ``origin of $z$'' filled by italicised letters show information obtained by the NED search service. ID 25 or ID 26 is an ``E+A and E+A'' pair. }\\label{table:all_list} {\\tabcolsep=0.9mm \\begin{tabular}{cccccccccccccccc} \\hline \\hline \\multicolumn{1}{c}{Pair} & \\multicolumn{6}{c}{E+A galaxies} & \\multicolumn{5}{c}{Companion galaxies} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{~} & \\multicolumn{1}{c}{\\scriptsize projection} \\\\ \\multicolumn{1}{c}{ID} & \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c}{Dec.} & \\multicolumn{1}{c}{$r$} & \\multicolumn{1}{c}{$M_r$} & \\multicolumn{1}{c}{$z$} & \\multicolumn{1}{c}{\\scriptsize ${\\rm H}\\delta~{\\rm EW}$} & \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c}{Dec.} & \\multicolumn{1}{c}{$r$} & \\multicolumn{1}{c}{$M_r$} & \\multicolumn{1}{c}{$z$} & \\multicolumn{1}{c}{origin of $z$} & \\multicolumn{1}{c}{$|\\Delta z|$} & \\multicolumn{1}{c}{\\scriptsize distance(kpc)} \\\\ \\hline % -- 03a 1 & % 587731514216874117 % J011942.23+010751.5 01:19:42.23 & +01:07:51.5 & 16.75 & -21.29 & 0.090 & 6.40 & % J011941.58+010739.9 01:19:41.58 & +01:07:39.9 & 17.63 & -20.41 & 0.089 & SDSS & 0.001 & 25.18 \\\\ % #2 587732156314747278 587732156314747279 2 & %J075323.09+243300.5 07:53:23.09 & +24:33:00.5 & 15.80 & -21.36 & 0.061 & 5.74 & %J075322.80+243302.5 07:53:22.80 & +24:33:02.5 & 15.78 & -21.38 & 0.062 & SDSS & 0.001 & 5.21 \\\\ % B04 587738948273963244 130.421395 26.71089 3 & %SDSSJ084141.13+264239.20 08:41:41.13 & +26:42:39.2 & 16.97 & -21.07 & 0.090 & 4.97 & %SDSSJ084140.75+264241.22 08:41:40.75 & +26:42:41.2 & 17.23 & -20.81 & 0.090 & KPNO Observation \\#11 & 0.000 & 9.11 \\\\ % #1 587727942953992403 587727942953992404 4 & %J090332.77+011236.4 09:03:32.77 & +01:12:36.4 & 16.16 & -20.89 & 0.058 & 4.43 & %J090332.99+011231.7 09:03:32.99 & +01:12:31.7 & 17.39 & -19.66 & 0.058 & SDSS & 0.000 & 6.32 \\\\ % #12 587732153104400523 587732153104400522 5 & % J092503.81+404602.8 09:25:03.81 & +40:46:02.8 & 17.69 & -20.97 & 0.117 & 4.90 & % J092504.65+404555.3 09:25:04.65 & +40:45:55.3 & 16.53 & -22.12 & 0.117 & SDSS & 0.000 & 25.25 \\\\ % --B06 587728949584265418 6 & %SDSSJ095824.62+003203.06 09:58:24.62 & +00:32:03.1  & 17.15 & -20.96 & 0.093 & 5.50 & %SDSSJ095824.08+003206.35 09:58:24.08 & +00:32:06.4  & 17.45 & -20.66 & 0.093 & KPNO Observation \\#15 & 0.000 & 14.89 \\\\ % --A04 587734861608321099 7 & %SDSSJ100522.45+072728.31 10:05:22.45 & +07:27:28.3  & 16.95 & -20.28 & 0.063 & 4.92 & % SDSSJ100521.12+072725.87 10:05:21.12 & +07:27:25.9  & 15.38 & -21.85 & 0.064 & KPNO Observation \\#4 & 0.001 & 23.99 \\\\ % NED 588848899901358270 588848899901358271 8 & % J101629.22-000137.1 10:16:29.22 & -00:01:37.1 & 17.37 & -21.02 & 0.105 & 5.90 & % J101628.55-000134.3 10:16:28.55 & -00:01:34.3 & 17.86 & -20.54 & 0.105 & {\\it MGC 0004420} & 0.000 & 19.43 \\\\ % #13 588017948808446124 588017948808446117 9 & % J105749.08+410035.9 10:57:49.08 & +41:00:35.9 & 17.50 & -21.38 & 0.128 & 4.05 & % J105747.54+410049.2 10:57:47.54 & +41:00:49.2 & 17.16 & -21.72 & 0.128 & SDSS & 0.000 & 49.89 \\\\ % --A07 588017703462043657 10 & %SDSSJ111004.24+115119.66 11:10:04.24 & +11:51:19.7  & 15.45 & -22.25 & 0.078 & 4.13 & %SDSSJ111003.74+115123.09 11:10:03.74 & +11:51:23.1  & 17.10 & -20.61 & 0.077 & KPNO Observation \\#7 & 0.001 & 11.78 \\\\ % #7 587734892751159308 587734892751159405 11 & % J114138.76+094344.6 11:41:38.76 & +09:43:44.6 & 17.64 & -20.07 & 0.078 & 4.64 & % J114136.84+094335.6 11:41:36.84 & +09:43:35.6 & 15.46 & -22.25 & 0.079 & SDSS & 0.001 & 43.52 \\\\ % --A01 588017112367956041 12 & %SDSSJ115208.57+481658.38 11:52:08.57 & +48:16:58.4  & 17.24 & -19.67 & 0.055 & 5.88 & %SDSSJ115207.36+481722.65 11:52:07.36 & +48:17:22.7  & 16.66 & -20.25 & 0.055 & KPNO Observation \\#12 & 0.000 & 28.52 \\\\ % #6 587732482211184777 587732482211184778 13 & % J115628.91+485541.6 11:56:28.91 & +48:55:41.6 & 17.72 & -19.99 & 0.078 & 4.06 & % J115627.66+485539.9 11:56:27.66 & +48:55:39.9 & 16.51 & -21.20 & 0.078 & SDSS & 0.000 & 18.14 \\\\ % #4 587738568702820588 587738568702820587 14 & % J121333.02+142900.1 12:13:33.02 & +14:29:00.1 & 16.26 & -21.01 & 0.064 & 5.94 & % J121334.12+142842.4 12:13:34.12 & +14:28:42.4 & 15.94 & -21.33 & 0.064 & SDSS & 0.000 & 28.90 \\\\ % #5 587735348575469590 587735348575469589 15 & % J122240.46+150205.4 12:22:40.46 & +15:02:05.4 & 17.72 & -19.82 & 0.072 & 8.10 & % J122240.76+150222.3 12:22:40.76 & +15:02:22.3 & 15.67 & -21.88 & 0.072 & SDSS & 0.000 & 23.73 \\\\ % --A09 587733196770705509 16 & %SDSSJ130029.45+545504.02 13:00:29.45 & +54:55:04.0  & 16.60 & -21.41 & 0.089 & 4.81 & %SDSSJ130032.15+545457.89 13:00:32.15 & +54:54:57.9  & 16.65 & -21.36 & 0.088 & KPNO Observation \\#8 & 0.001 & 39.44 \\\\ % #8 587726032795336710 587726032795271362 17 & % J133024.73+022325.8 13:30:24.73 & +02:23:25.8 & 17.33 & -20.41 & 0.079 & 5.66 & % J133023.85+022304.3 13:30:23.85 & +02:23:04.3 & 15.67 & -22.06 & 0.079 & SDSS & 0.000 & 37.22 \\\\ % --B08-2 587729776904372392 18 & %SDSSJ135005.58-024734.68 13:50:05.58 & -02:47:34.7  & 16.71 & -22.16 & 0.128 & 5.81 & %SDSSJ135005.89-024738.82 13:50:05.89 & -02:47:38.8  & 17.94 & -20.93 & 0.129 & KPNO Observation \\#13 & 0.001 & 14.02 \\\\ % NED 587726032797565040 587726032797565039 19 & %J135050.98+021938.4 13:50:50.98 & +02:19:38.4 & 15.85 & -19.90 & 0.033 & 4.59 & %J135053.55+021924.6 13:50:53.55 & +02:19:24.6 & 14.28 & -21.46 & 0.033 & {\\it UGC 08750 NED01} & 0.000 & 26.25 \\\\ % --A05 587739809414250542 20 & %SDSSJ142453.14+230746.79 14:24:53.14 & +23:07:46.8  & 15.70 & -21.89 & 0.074 & 4.05 & %SDSSJ142452.86+230720.89 14:24:52.86 & +23:07:20.9  & 16.82 & -20.77 & 0.074 & KPNO Observation \\#9 & 0.000 & 36.24 \\\\ % #11 587733605330518229 587733605330518227 21 & % J143633.54+555317.8 14:36:33.54 & +55:53:17.8 & 17.60 & -20.78 & 0.104 & 4.08 & % J143631.82+555327.6 14:36:31.82 & +55:53:27.6 & 16.65 & -21.74 & 0.105 & SDSS & 0.001 & 32.96 \\\\ % #3 587726100413677739 587726100413677730 22 & % J145033.38+031820.0 14:50:33.38 & +03:18:20.0 & 17.33 & -19.86 & 0.062 & 4.04 & % J145031.09+031803.5 14:50:31.09 & +03:18:03.5 & 16.89 & -20.30 & 0.062 & SDSS & 0.000 & 44.90 \\\\ % --E40 587736585506586887 23 & %SDSSJ151309.04+334944.41 15:13:09.04 & +33:49:44.4  & 17.47 & -21.84 & 0.155 & 5.05 & %SDSSJ151307.74+334951.39 15:13:07.74 & +33:49:51.4  & 19.48 & -19.83 & 0.155 & KPNO Observation \\#17 & 0.000 & 46.63 \\\\ % --D09 587729160057127148 24 & %SDSSJ153416.00+035934.62 15:34:16.00 & +03:59:34.6  & 17.71 & -21.65 & 0.157 & 4.55 & %SDSSJ153416.19+035933.43 15:34:16.19 & +03:59:33.4  & 18.32 & -21.03 & 0.157 & KPNO Observation \\#14 & 0.000 & 8.27 \\\\ % #9 587736618787602740 587736618787602523 25 & % J155728.87+272545.2 15:57:28.87 & +27:25:45.2 & 16.43 & -21.55 & 0.087 & 4.22 & % J155730.42+272541.1 15:57:30.42 & +27:25:41.1 & 17.30 & -20.68 & 0.088 & SDSS & 0.001 & 34.10 \\\\ % #10 587736618787602523 587736618787602740 26 & % J155730.42+272541.1 15:57:30.42 & +27:25:41.1 & 17.30 & -20.69 & 0.088 & 4.11 & % J155728.87+272545.2 15:57:28.87 & +27:25:45.2 & 16.43 & -21.56 & 0.087 & SDSS & 0.001 & 34.23 \\\\ % NED 587729227690475600 587729227690475598 27 & % J161330.18+510335.5 16:13:30.18 & +51:03:35.5 & 15.48 & -20.34 & 0.034 & 7.57 & % J161332.23+510342.9 16:13:32.23 & +51:03:42.9 & 15.17 & -20.65 & 0.033 & {\\it I Zw 136 NOTES02} & 0.001 & 13.64 \\\\ % --B03 587729231976857930 28 & %SDSSJ162151.96+492859.76 16:21:51.96 & +49:28:59.8  & 17.04 & -20.79 & 0.082 & 4.26 & %SDSSJ162155.07+492908.04 16:21:55.07 & +49:29:08.0  & 18.21 & -19.62 & 0.082 & KPNO Observation \\#6 & 0.000 & 47.84 \\\\ % --A02 588017979453341736 29 & %SDSSJ162301.30+230039.84 16:23:01.30 & +23:00:39.8  & 17.48 & -19.69 & 0.061 & 4.73 & %SDSSJ162301.04+230112.01 16:23:01.04 & +23:01:12.0  & 14.89 & -22.28 & 0.062 & KPNO Observation \\#5 & 0.001 & 37.84 \\\\ % --B05 588018056199012588 30 & %SDSSJ162317.62+324526.75 16:23:17.62 & +32:45:26.8  & 17.73 & -20.32 & 0.090 & 4.65 & %SDSSJ162317.42+324502.30 16:23:17.42 & +32:45:02.3  & 17.24 & -20.81 & 0.091 & KPNO Observation \\#10 & 0.001 & 40.97 \\\\ % --E06 587733431921410283 31 & %SDSSJ163925.01+303709.82 16:39:25.01 & +30:37:09.8  & 17.56 & -20.86 & 0.106 & 8.23 & %SDSSJ163924.87+303715.30 16:39:24.87 & +30:37:15.3  & 18.00 & -20.42 & 0.106 & KPNO Observation \\#16 & 0.000 & 11.09 \\\\ % --aa08 587733609094643902 32 & % J165648.64+314702.3 16:56:48.64 & +31:47:02.3 & 16.70 & -21.58 & 0.100 & 5.75 & % J165648.32+314702.3 16:56:48.32 & +31:47:02.3 & 18.63 & -19.66 & 0.100 & KPNO Observation \\#3 & 0.000 & 13.39 \\\\ % --08a3 587725491597410410 33 & %J170356.71+622848.3 17:03:56.71 & +62:28:48.3 & 17.64 & -21.53 & 0.145 & 7.29 & %J170355.12+622856.0 17:03:55.12 & +62:28:56.0 & 18.03 & -21.14 & 0.144 & KPNO Observation \\#1 & 0.001 & 33.89 \\\\ % --06a 587729781200454076 34 & % J170859.24+322053.1 17:08:59.24 & +32:20:53.1 & 17.58 & -21.14 & 0.121 & 6.03 & % J170859.03+322100.3 17:08:59.03 & +32:21:00.3 & 18.08 & -20.65 & 0.120 & KPNO Observation \\#2 & 0.001 & 16.49 \\\\  %% -- 02 %3 & %%  J011537.59+000534.0 %01:15:37.59 & +00:05:34.0 & 17.78 & -18.80$^{*}$ & 0.044 & ? & %%  J011539.20+000611.0 %01:15:39.20 & +00:06:11.0 & 16.10 & -20.29 & 0.044 & %? & 0.000 & ? \\\\ %% 04a2(ext) %5 & %% J170557.44+630055.7 %17:05:57.44 & +63:00:55.7 & 16.66 & -21.75 & 0.105 & ? & %% J170557.76+630133.2 %17:05:57.76 & +63:01:33.2 & 18.34 & -17.70$^{*}$ & 0.104 & %? & 0.001 & ? \\\\ \\hline \\end{tabular} } \\end{center} \\end{table*}  %1 01:19:42.23 +01:07:51.5 %2 07:53:23.09 +24:33:00.5 %3 09:03:32.77 +01:12:36.4 %4 09:25:03.81 +40:46:02.8 %5 09:58:24.62 +00:32:03.1 %6 10:05:22.45 +07:27:28.3 %7 10:16:29.22 -00:01:37.1 %8 10:57:49.08 +41:00:35.9 %9 11:10:04.24 +11:51:19.7 %10 11:41:38.76 +09:43:44.6 %11 11:52:08.57 +48:16:58.4 %12 11:56:28.91 +48:55:41.6 %13 12:13:33.02 +14:29:00.1 %14 12:22:40.46 +15:02:05.4 %15 13:00:29.45 +54:55:04.0 %16 13:30:24.73 +02:23:25.8 %17 13:50:05.58 -02:47:34.7 %18 13:50:50.98 +02:19:38.4 %19 14:24:53.14 +23:07:46.8 %20 14:36:33.54 +55:53:17.8 %21 14:50:33.38 +03:18:20.0 %22 15:13:09.04 +33:49:44.4 %23 15:34:16.00 +03:59:34.6 %24 15:57:28.87 +27:25:45.2 %25 15:57:30.42 +27:25:41.1 %26 16:13:30.18 +51:03:35.5 %27 16:21:51.96 +49:28:59.8 %28 16:23:01.30 +23:00:39.8 %29 16:23:17.62 +32:45:26.8 %30 16:39:25.01 +30:37:09.8 %31 16:56:48.64 +31:47:02.3 %32 17:03:56.71 +62:28:48.3 %33 17:08:59.24 +32:20:53.1   % ",
        "conclusions": "%\\label{section:discussion}  %The only works about companion galaxies of E+As are %\\citet{got03,got05}, %which investigated accompanying galaxies of E+As %within 100 kpc using SDSS imaging data. %Our analysis adopted more strict statistics %using both spectroscopic and imaging data %to enhance the reliance of discussion about %the existing these companion galaxies. %As a result, we also found that %the probability of existing the companions about E+A galaxies %is larger than that about normal galaxies.  %There are some attractive studies about normal pair galaxies, and %we can consider our result. %\\citet{lam03} examined 1258 galaxy pairs in the 100k public release of %the 2dF galaxy survey and found that star formation in galaxy pairs is %significantly enhanced over that of isolated galaxies for separations %less than 36 kpc and velocity differences less than 100km s$^{-1}$. %In addition, %\\citet{nik04} also investigated %12492 galaxy pairs at %projected separations less than 300 kpc using SDSS DR1, %and reported that the mean specific star formation rate is %significantly enhanced for projected separations less than 30 kpc. %Our result also means that %the fraction of post-starburst galaxies in pairs %is larger than that in solo galaxies, therefore, %we are convinced the results of these studies.  %The studies of E+A galaxies have progressed with increasing speed, %for example, their morphologies, radial photometric/spectroscopic %gradients, %etc. \\citep{bar01,yan04,yam05} %Thus, recent observational studies %and numerical simulations \\citep{bek05} %support the merger/interaction origin of E+A galaxies, %and basically there is no room for doubt as to it. %Our present concern is the evolution of E+A systems, therefore, %more detailed studies \\citep[e.g.][]{yag06a,yag06b,got07a} will be expected. %The next studies for us will be %investigation about properties of companion galaxies %and its dependence of E+A's properties. %We aim to reveal the effect of the local environment for E+A galaxies %and their evolution scenarios via our studies about E+As and companion %galaxies.  %The magnitude range of our sample is limited and the number of E+A %galaxies is too small to obtain volume-limited sample. %In addition, there are a lot of companion candidates without %spectroscopic data. %These factors make systematic study of E+A systems difficult. %A survey of the next generation may solve these problems."
    },
    "0809/0809.1793_arXiv.txt": {
        "abstract": "{ It was found that the advection-dominated accretion flow (ADAF)$+$thin disk model calculations can reproduce the observed spectral energy distributions (SED) of the two low-luminosity active galactic nuclei (AGNs), provided they are accreting at $\\sim 0.01-0.03$ Eddington rates and the thin disks are truncated to ADAFs at $\\sim 100R_{\\rm S}$ ($R_{\\rm S}$ is the Schwarzschild radius) for M81 and NGC4579 \\citep{qdnh99} . However, the black hole masses adopted in their work are about one order of magnitude lower than recent measurements on these two sources. Adopting the well estimated black hole masses, our ADAF$+$thin disk model calculations can reproduce the observed SEDs of these two low-luminosity active galactic nuclei, if the black hole is accreting at $2.5\\times 10^{-4}$ Eddington rates with the thin disk truncated at $R_{\\rm tr}=120R_{\\rm S}$ for M81 ($\\dot{m}=3.3\\times10^{-3}$ and $R_{\\rm tr}=80R_{\\rm S}$ are required for NGC4579). The transition zones with temperature from the thin disk with $\\sim 10^{4-5}$ to $\\sim 10^{9-10}$~K in the ADAF will inevitably emit thermal X-ray lines, which provides a useful diagnosis on their physical properties. The observed widths of the thermal X-ray iron lines at $\\simeq 6.8$~keV are consistent with the Doppler broadening by the Keplerian motion of the gases in the transition zones at $\\sim 100R_{\\rm S}$. We use the structure of the transition zone between the ADAF and the thin disk derived by assuming the turbulent diffusive heat mechanism to calculate their thermal X-ray line emission with standard software package Astrophysical Plasma Emission Code (APEC). Comparing them with the equivalent widths of the observed thermal X-ray iron lines in these two sources, we find that the turbulent diffusive heat mechanism seems to be unable to reproduce the observed thermal X-ray line emission. The test on the evaporation model for the accretion mode transition with the observed thermal X-ray line emission is briefly discussed. ",
        "introduction": "It is well accepted that many astrophysical systems are powered by black hole accretion. The standard thin disks (i.e., geometrically thin and optically thick accretion disks) can successfully explain most observational features of the black hole accretion systems \\citep{ss73}. The ultraviolet/optical continuum emission observed in luminous quasars is usually attributed to the thermal radiation from the standard disks (SD) surrounding the massive black holes in quasars \\citep*[e.g.,][]{sm89}. However, the standard disk model is unable to reproduce the spectral energy distributions (SED) of many sources (e.g., Sgr A$^*$) accreting at very low rates, and the advection dominated accretion flows (ADAF) were suggested to be present in these sources \\citep{ny94,ny95}. In the ADAF model, most released gravitational energy of the gases in the accretion flow is converted to the internal energy of the gas, and the ADAF is hot, geometrically thick, and optically thin. Only a small fraction of the released energy in ADAFs is radiated away, and their radiation efficiency is therefore significantly lower than that for standard thin disks \\citep*[see][for a review and references therein]{nmq98}. The ADAF model can successfully reproduce the observed SEDs in many black hole systems accreting at low rates \\citep*[e.g.,][]{lackny96,nmy96,nmg98}.  There is a critical accretion rate $\\dot{m}_{\\rm crit}$, above which the ADAF is suppressed and a standard thin disk is present. The ADAF can co-exist with the standard thin disk, when it is accreting at rates slightly lower than the critical value $\\dot{m}_{\\rm crit}$. In this case, the ADAF is present in the inner region near the black hole and connects to a standard thin disk at a certain transition radius $R_{\\rm tr}$ \\citep*[see][for a review and references therein]{nmq98}. \\citet{nmy96} proposed the ADAF$+$standard thin disk systems to model the observed spectra of the black hole accretion systems with moderate accretion rates. \\citet{n96} found that many different spectral states observed in black hole X-ray binaries can be understood as a sequence of ADAF$+$thin disk models with varying $\\dot{m}$ and $r_{\\rm tr}$ (where $r_{\\rm tr}$ is the outer radius of ADAF in units of Schwarzschild radius, $R_{\\rm S}=2GM_{\\rm bh}/c^{2}$). This scenario was explored in more detail by several authors \\citep{emn97,encgz98}, and they found that the different spectral states of X-ray binaries (i.e.,  the quiescent state, low state, intermediate state or high state) correspond to different accretion rates. It is believed that standard thin disks (or slim disks) are present in luminous active galactic nuclei (AGN), while ADAFs are present in those low-luminosity AGNs accreting at relatively low rates. The ADAF$+$thin disk systems are required for modeling on a variety of observations of AGNs accreting at the moderate rates \\citep*[e.g.,][]{qdnh99,cao03}.  The physical mechanisms of the transition from a thin disk to an ADAF and its structure are still quite uncertain, though a few different scenarios were suggested \\citep[e.g.,][]{mmh94,h96,lymmx99}. \\citet{mmh94} proposed an \"evaporation\" mechanism, in which the disk is evaporated as heated by electron conduction from a hot corona, and then a quasi-spherical hot accretion flow is formed. An alternative scenario was suggested by \\citet{h96}, in which a \"turbulent diffusive heat\" mechanism is employed to explore the structure of the ADAF$+$thin disk system. This model was further developed, and the global structure of the ADAF$+$thin disk system was derived numerically \\citep{mk2000,mknn2000}. All of these model calculations show that a rapid decrease of temperature occurs within a narrow zone between the inner ADAF and the outer standard thin disk at $\\sim r_{\\rm tr}$.   \\citet{qdnh99} showed that the optical/UV to X-ray emission detected from the nuclei of M81 and NGC 4579 can be well explained by an optically thick, geometrically thin accretion disk extends down to ~100~$R_{\\rm S}$ (inside of which an ADAF is present), provided their accretion rates are $\\dot{m}\\sim 0.01$. The optical/UV bumps observed in these two sources can be attributed to the thermal emission from the outer standard thin disks, while their hard X-ray emissions are dominantly from the inner ADAFs. In their model calculations, the same central black hole mass $M_{\\rm bh}=4\\times10^6$~$M_\\odot$ is adopted for these two sources, which are about an order of magnitude underestimated compared with the recent estimates \\citep{d03,b01}. The iron K$\\alpha$ lines were observed in these two sources \\citep{i96,t98}. In addition to narrow iron K$\\alpha$ lines at 6.4 KeV, the broad line components centered at $E_{{\\rm K}\\alpha}=6.79$~KeV were also observed in these two sources \\citep{dgds04}. The thermal X-ray line emission is a useful diagnosis on black hole accretion systems\\citep*[e.g.,][]{nr99,xu06}. \\citet{dgds04} found that the observed lines from these two sources are too broad to be the thermal line emission from the host galaxy, and they suggested that these broad emission lines may probably be from the transition zones between the ADAFs and the outer standard thin disks. The temperature of the inner edge of the thin disk is around $10^4$~K, and therefore the temperature of the transition zone extends from $\\sim 10^4$ to $\\sim 10^{8-9}$~K while connecting to the ADAF. Thus, the thermal X-ray line emission can naturally originate from such transition zones. The observed widths of the broad lines are roughly consistent with the Doppler broadening by the Keplerian motion of the gases in the transition zones at $\\sim100~R_{\\rm S}$ \\citep{qdnh99}.  In this work, we re-investigate the ADAF$+$thin disk systems in these two low-luminosity active galactic nuclei, NGC 4579 and M81, adopting the black hole masses estimated by \\citet{d03,b01}.  The thermal iron K$\\alpha$ line emissions from the transition zones of the ADAF to the thin disk in these two low-luminosity active galactic nuclei, NGC 4579 and M81, are calculated with the transition model based on the turbulent diffusive heat mechanism developed by \\citet{h96}. ",
        "conclusions": "Our best fits to the observed continuum spectra of these two sources show that the outer thin discs are truncated at $r_{\\rm tr}=120$ and 80, for M81 and NGC4579, respectively. The two important disk parameters, $\\dot{m}$ and $r_{\\rm tr}$ are mainly determined from the comparison of the outer disk spectra with the observed optical/UV continuum emission in these two sources, and the observed X-ray continuum spectra can then be naturally reproduced by the ADAF models. The observed widths of the broad lines are roughly consistent with the Doppler broadening by the Keplerian motion of the gas in the transition zone at ~50-150~$R_{\\rm S}$ \\citep{dgds04}. The ratio of the transition radii of these two sources is $1.5$, which is consistent with the observed line width ratio $\\simeq 0.814$, because the width $\\sigma\\propto r_{\\rm tr}^{-1/2}$.  Comparison of our calculations based on the transition zones given by \\citet{mknn2000} (see Fig. \\ref{fig3}, and \\S 3 and 4 for the detailed description) with the observations show that the accretion rates $\\dot{m}\\sim 7.1\\times10^{-5}$ and $6.5\\times10^{-4}$ for M81 and NGC4579, respectively (see Fig. \\ref{fig3}). These seem to be inconsistent with the optical/UV spectra observed in these two sources, which requires the accretion rates to be $\\dot{m}=2.5\\times10^{-4}$ (M81) and $3.3\\times10^{-3}$ (NGC4579), unless the transition radii deviate significantly from $\\sim 100$ (e.g., an order of magnitude smaller than 100). However, the widths of the thermal X-ray lines provide strict constraints on the transition radii, if the thermal X-ray line emission does originate from the transition zones. The dependence of our results on the metallicity can be understood with Eq.  (\\ref{eq7}). A larger metallicity will increase the line emissivity and line optical depth at the same time, while the line luminosity and EW will slightly increase with a factor ${1-e^{-\\tau_{\\rm line}}}$. Of course, the structure of ADAF$+$thin disk systems based on this model is rather simplified, in which only one temperature is considered (i.e., the electrons have the same temperature as the ions). Thus, we cannot rule out this model only from their thermal X-ray line emission. In the transition zone, the density is very high so that one temperature flow is a good approximation, as the Coulomb interaction between ions and electrons in this region is very efficient. Our calculations of the thermal X-ray line emission from the transition zone have not been affected by this one temperature assumption. The more detailed calculations on the ADAF$+$thin disk systems including energy equilibrium between electrons and ions in the accretion flows may help to test this model, which is beyond the scope of this work.  An alternative model for the transition of a thin disk to an ADAF is the ``evaporation\" mechanism initially suggested by \\citet{mmh94}, in which the thin disk is evaporated as heated by electron conduction from a hot corona, and then a quasi-spherical hot accretion flow is formed. There is a very thin layer between the cold thin disk and the hot corona, of which the electron temperature may vary from $\\sim 10^{6.5}$~K to $\\sim 10^9$~K in the corona \\citep*[e.g.][]{liu95,liu97,liu02}. Such a layer, of course, will emit thermal X-ray line emission. In principle, the evaporation model for the transition of a thin disk to an ADAF can be tested by the observed thermal X-ray line emission. However, the structure of this disk corona system is complicated and is only available by numerically integrating a set of ordinary differential equations \\citep*[e.g.,][]{liu02}, which prevents us from testing this model in this paper."
    },
    "0809/0809.0984_arXiv.txt": {
        "abstract": "{Recent calculations have shown that the UV bump at about 217.5~nm in the extinction curve can be explained by a complex mixture of polycyclic aromatic hydrocarbons (PAHs) in several ionisation states. Other studies proposed that the carriers are a restricted population made of neutral and singly\\textendash ionised dehydrogenated coronene molecules (C$_{24}$H$_{n}$, $n\\leq3$), in line with models of the hydrogenation state of interstellar PAHs predicting that medium\\textendash sized species are highly dehydrogenated.} {To assess the observational consequences of the latter hypothesis we have undertaken a systematic theoretical study of the electronic spectra of dehydrogenated PAHs. We use our first results to see whether such spectra show strong general trends upon dehydrogenation.} {We performed calculations using state\\textendash of\\textendash the\\textendash art techniques in the framework of the density functional theory (DFT) to obtain the electronic ground\\textendash state geometries, and of the time\\textendash dependent DFT to evaluate the electronic excited\\textendash state properties.} {We computed the absorption cross\\textendash section of the species C$_{24}$H$_{n}$ (n=12,10,8,6,4,2,0) in their neutral and cationic charge\\textendash states. Similar calculations were performed for other PAHs and their fully dehydrogenated counterparts.} {$\\pi$\\textendash electron energies are always found to be strongly affected by dehydrogenation. In all cases we examined, progressive dehydrogenation translates into a correspondingly progressive blue shift of the main electronic transitions. In particular, the $\\pi\\to\\pi^*$ collective resonance becomes broader and bluer with dehydrogenation. Its calculated energy position is therefore predicted to fall in the gap between the UV bump and the far\\textendash UV rise of the extinction curve. Since this effect appears to be systematic, it poses a tight observational limit on the column density of strongly dehydrogenated medium\\textendash sized PAHs.} ",
        "introduction": "\\label{introduction}  First discovered by \\citet{ste65a}, the absorption band centered at about 217.5~nm  ($\\sim$5.7~eV) is the most intense interstellar extinction feature. With a relatively constant peak wavelength, the so\\textendash called ``UV bump'' has a Lorentz\\textendash like profile \\citep[e.g.,][]{sea79} whose width varies considerably from one line of sight to another \\citep{fit07}.  Both the strength and the position of the 217.5~nm feature suggest that it originates in some form of carbonaceous material, known to display strong $\\pi\\to\\pi^\\star$ electron transitions in this range. A large number of carrier candidates have been proposed over the years, amongst which size\\textendash restricted graphite particles \\citep{ste65b}, fullerenes, C$_{60}$ in particular, \\citep[e.g.,][]{bra91,igl04}, polycyclic aromatic hydrocarbons \\citep[PAHs,][]{job92a,job92b}, coal\\textendash like material \\citep{pap96}, UV\\textendash processed hydrogenated amorphous carbon particles \\citep{men98}, nano\\textendash sized hydrogenated carbon grains \\citep{sch98}, and carbon onions \\citep{chh03}.  Since each PAH displays strong $\\pi\\to\\pi^\\star$ transition in the 5.7~eV region \\citep{job92a,job92b,dra03,mal04}, mixtures of PAHs must contribute to the 217.5~nm bump. Recently, \\citet{ccp08} showed that such mixtures can account for the entire bump, as well as its strength relative to the non\\textendash linear far\\textendash UV rise, depending on the distribution of charge states in the PAH mixtures. The UV bump was also attributed to a  $\\pi\\to\\pi^\\star$ plasmon resonance from a very limited number of strongly dehydrogenated derivatives of coronene (C$_{24}$H$_{12}$), namely C$_{24}$H$_n$ and C$_{24}$H$_n$$^+$ with $n\\leq3$. This assignment was made on the basis of DDA scattering calculations \\citep{dul98}, supported by laboratory studies of the UV absorption in amorphous carbon \\citep{dul04}. The same molecules were also proposed as carriers of some of the diffuse interstellar bands \\citep{dul06b}. Indeed, models of the hydrogenation state of PAHs predict small and medium\\textendash sized PAHs ($\\sim$20-30 carbon atoms) to be highly dehydrogenated \\citep{lep03}.  Since the effect of dehydrogenation on the electronic spectra of isolated PAHs is poorly characterized, we recently undertook a detailed study of the electronic excitation properties of dehydrogenated PAHs, using state\\textendash of\\textendash the\\textendash art quantum\\textendash chemical tools \\citep{mal08}. As a part of this more detailed work, we present in this paper our first results for the dehydrogenated derivatives of coronene (C$_{24}$H$_{12}$), and completely dehydrogenated perylene (C$_{20}$H$_{10}$) and ovalene (C$_{32}$H$_{14}$). We use them to assess the specific case of the species introduced by \\citet{dul06a}, whose computed spectral properties appear to be incompatible with the observed UV\\textendash bump, and search for general trends in the electron properties of dehydrogenated PAHs.  Technical details of the calculations, as well as the justification for the choice of molecules considered, are given in Sect.~\\ref{methods}. Results are presented and discussed in the astrophysical context in Sect.~\\ref{results}, and our conclusions are reported in Sect.~\\ref{conclusions}.  \\begin{figure}[b!]  \\begin{center} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip, width=2.5cm]{fig1a.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.5cm]{fig1b.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.5cm]{fig1c.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.45cm]{fig1d.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.4cm]{fig1e.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.3cm]{fig1f.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.4cm]{fig1g.eps} \\includegraphics[trim=3cm 3cm 3cm 3cm, clip,width=2.75cm]{fig1h.eps}   \\caption{Geometries of the dehydrogenated coronene molecules considered (C$_{24}$H$_{10}$, C$_{24}$H$_{8}$, C$_{24}$H$_{6}$, C$_{24}$H$_{4}$, C$_{24}$H$_{2}$, and C$_{24}$), and of the completely dehydrogenated perylene (C$_{20}$) and ovalene (C$_{32}$). \\label{mols}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{conclusions}  PAHs are expected to exist in a wide variety of interstellar environments, in a complex statistical equilibrium of different charge and hydrogenation states \\citep[see e.g.,][]{tie05}. Modelling studies suggest that intermediate\\textendash size PAHs in the range of 20\\textendash30 carbon atoms are stripped of most of their hydrogen atoms but still survive under the conditions of the diffuse interstellar medium \\citep{lep03}.   We undertook a systematic theoretical study of the effects of dehydrogenation on the electron properties of neutral and ionised PAHs using state\\textendash of\\textendash the\\textendash art quantum\\textendash chemical techniques in the framework of the real\\textendash time TD\\textendash DFT. For all of the species considered so far, we found that, regardless of charge\\textendash state and molecular size, the $\\pi\\to \\pi^\\star$ collective resonance broadens and shifts to higher energies with increasing dehydrogenation.  When the molecule is half dehydrogenated or more, the band is shifted enough to fall between the bump and the far\\textendash UV rise of the extinction curve. The associated species cannot be the carriers of the bump, contrary to the proposal of \\citet{dul06a}.  While it is overreaching to draw definitive conclusions based on the few cases considered, our upper limits are due to systematic effects, and are thus unlikely to depend strongly on the specific sample of molecules considered. Nonetheless, a larger study, which is underway \\citep{mal08}, will be needed to assess to what extent the effect of dehydrogenation that we observed in UV spectra of PAHs is systematic for the whole class."
    },
    "0809/0809.5045_arXiv.txt": {
        "abstract": "Monte Carlo (MC) simulations and series expansions (SE) data for the energy, specific heat, magnetization, and susceptibility of the three-state and four-state Potts model and Baxter-Wu model  on the square lattice are analyzed in the vicinity of the critical point in order to estimate universal combinations of critical amplitudes. We also form effective ratios of the observables close to the critical point and analyze how they approach the universal critical-amplitude ratios. In particular, using the duality relation, we show analytically that for the Potts model with a number of states $q\\le 4$, the effective ratio of the energy critical amplitudes always approaches unity linearly with respect to the reduced temperature. This fact leads to the prediction of relations among the amplitudes of correction-to-scaling terms of the specific heat in the low- and high-temperature phases. It is a common belief that the four-state Potts and the Baxter-Wu model belong to the same universality class. At the same time, the critical behavior of the four-state Potts model is modified by logarithmic corrections while that of the Baxter-Wu model is not. Numerical analysis shows that critical amplitude ratios are very close for both models and, therefore, gives support to the hypothesis that the critical behavior of both systems is described by the same renormalization group fixed point. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.5103_arXiv.txt": {
        "abstract": "Based on the modeling of the central emission-line width measured over sub-arcsecond apertures with the {\\em Hubble Space Telescope}, we present stringent upper bounds on the mass of the central supermassive black hole, \\mbh, for a sample of 105 nearby galaxies ($D<100\\,\\rm Mpc$) spanning a wide range of Hubble types (E $-$ Sc) and values of the central stellar velocity dispersion, \\sigmac\\ (58 $-$ 419 \\kms). For the vast majority of the objects the derived \\mbh\\ upper limits run parallel and above the well-known \\msc\\ relation independently of the galaxy distance, suggesting that our nebular line-width measurements trace rather well the nuclear gravitational potential. For values of \\sigmac\\ between 90 and 220 \\kms\\ the 68\\% of our upper limits falls immediately above the \\msc\\ relation without exceeding the expected \\mbh\\ values by more than a factor 4.1. No systematic trends or offsets are observed in this \\sigmac\\ range as a function of the galaxy Hubble type or with respect to the presence of a bar. For 6 of our 12 \\mbh\\ upper limits with \\sigmac\\ $< 90$ \\kms\\ our line-width measurements are more sensitive to the stellar contribution to the gravitational potential, either due to the presence of a nuclear stellar cluster or because of a greater distance compared to the other galaxies at the low-\\sigmac\\ end of the \\msc\\ relation. Conversely, our \\mbh\\ upper bounds appear to lie closer to the expected \\mbh\\ in the most massive elliptical galaxies with values of \\sigmac\\ above 220 \\kms. Such a flattening of the \\msc\\ relation at its high-\\sigmac\\ end would appear consistent with a coevolution of supermassive black holes and galaxies driven by dry mergers, although better and more consistent measurements for \\sigmac\\ and $K$-band luminosity are needed for these kind of objects before systematic effects can be ruled out. ",
        "introduction": "\\label{sec:introduction}  Supermassive black holes (SMBHs) have now been discovered in the center of a sufficiently large number of nearby galaxies to probe possible links between the masses of SMBHs (\\mbh) and the global properties of their host galaxies. In fact, it has emerged that \\mbh\\ correlates with the luminosity \\citep{Kormendy1995, Marconi2003a}, mass \\citep{Magorrian1998, Haering2004}, stellar velocity dispersion \\citep{Ferrarese2000, Gebhardt2000, Tremaine2002, Ferrarese2005}, light concentration \\citep{Graham2001}, and gravitational binding energy \\citep{Aller2007} of the host-galaxy spheroidal component, i.e., the entire galaxy in the case of elliptical galaxies or the bulge of disk galaxies. In light of these findings it is now widely accepted that the mass-accretion history of a SMBH is tightly related through feedback to the formation and evolution of the host spheroid \\citep[e.g.,][]{Silk1998, Haehnelt2000, DiMatteo2005}, some studies having suggested a link with the mass of the dark-matter halo \\citep{Ferrarese2002, Pizzella2005}.  The slope and scatter of all these correlations remain quite uncertain \\citep{Novak2006}, however, in particular since they are still based on a limited sample of galaxies with reliable \\mbh\\ that is biased towards early-type systems and that is clustered around a rather limited range of stellar velocity dispersion ($\\sigma$), approximately between $150$ and $250$ \\kms. Given the great theoretical interest spurred by these findings, there is a pressing need to acquire better \\mbh\\ statistics, both in terms of the number of targets and in terms of broadening the range of parent galaxies, in particular towards spiral galaxies.  Secure \\mbh\\ measurements in external galaxies are traditionally obtained through the modeling of the stellar and/or gaseous kinematics, most often as derived using {\\em Hubble Space Telescope (HST)} observations in the optical domain. The advent of adaptive-optics systems working at near-infrared wavelengths has led to more stellar-dynamical measurements of \\mbh\\ from the ground \\citep[][]{Houghton2006, Nowak2007}. Yet, such measurements are still quite expensive, not only because good-quality measurements of the stellar kinematics in the near-infrared require relatively long observations, but also because proper modeling of the stellar kinematics in the immediate vicinity of SMBHs needs robust constraints on the importance of radial orbits and thus additional large-scale observations, possibly with integral-field spectroscopy \\citep{Valluri2004, Cappellari2005}. Water-masers have provided the most accurate extragalactic \\mbh\\ measurements to date, but such gaseous systems are exceedingly rare \\citep{Braatz1994, Greenhill2003}. The modeling of the nuclear ionized-gas kinematics has also led to accurate \\mbh\\ measurements \\citep[e.g.,][]{Barth2001, DallaBonta2008}, and at a relatively cheap cost in terms of observation time compared to stellar-dynamical \\mbh\\ determinations \\citep[e.g.,][]{Verolme2002, Gebhardt2003}. Yet, only an handful of the objects targeted by HST turned out to have sufficiently regular gas velocity fields for the purpose of modeling \\citep[][]{Sarzi2001}. Thus, unless a large number of galaxies pre-selected to have regular nuclear gas kinematics \\citep[for instance following][]{Ho2002} is observed with HST if and when the Space Telescope Imaging Spectrograph (STIS) is successfully refurbished, it is unlikely that the number of galaxies with secure \\mbh\\ measurements will increase dramatically in the near future.  The HST Science Archive already contains an untapped resource that can be used to better constrain the black-hole mass budget across the different morphological types of galaxies, which consists of the vast number of the STIS spectra from which a central emission-line width can be measured. The modeling of this kind of data can indeed lead to tight upper limits on \\mbh, as first shown by \\citet{Sarzi2002}. For this reason we started a program aimed at deriving \\mbh\\ upper limits based on HST spectra for the largest possible number of galaxies and a wide range of morphological types. In this paper we present the results based on a sample of 105 nearby galaxies for which STIS/G750M spectra in the \\ha\\ region and measurements of the stellar velocity dispersion were available from the HST archive and in the literature, respectively. Although we will be able only to set an upper limit on the \\mbh\\ of our galaxies, the lack of exact measurements will be compensated for by the large number of upper limits when studying SMBH mass-host galaxy relationships.  The paper is organized as follows. In \\S~\\ref{sec:analysis} we describe our sample selection and the measurement of central emission-line width, before briefly describing our modeling. We will then present our results and discuss our findings in the context of the \\ms\\ relation between the SBHM mass and central stellar velocity dispersion of the host spheroid in \\S~\\ref{sec:results}. ",
        "conclusions": "\\label{sec:results}  We have determined the $1\\sigma$ upper and lower confidence limits for the \\mbh\\ for randomly orientated disks for 105 galaxies with measurable spectra and stellar velocity dispersions available in the literature. For 19 galaxies of the sample, either \\mbh\\ measurements or upper limits based on resolved kinematics were available (Table~1).  Fig.~\\ref{fig:ULcomparison} shows how such measurements compare with our \\mbh\\ limits, once our values are rescaled accordingly to the distances adopted in these studies. Our \\mbh\\ upper-limits are consistent within 1$\\sigma$ with such estimates, except for NGC~3031 and NGC~4261. Furthermore no systematic offset appears when our upper limits are compared with similar upper bounds in the literature, rather than definite measurements. A particularly complex blend of narrow \\ha+\\nii\\ and broad \\ha\\ lines may have biased our \\mbh\\ estimates in NGC~3031.  To place our \\mbh\\ limits on the various versions of the \\ms\\ relation, we applied to the aperture correction of \\citet{Jorgensen1995} to the literature values of stellar velocity dispersion in order to obtain the values \\sigmac\\ and \\sigmae\\ and that would have been measured within a circular aperture of radius $r_{\\rm e}/8$ and $r_{\\rm e}$, respectively. The effective radii $r_{\\rm e}$ of the spheroidal components of our sample galaxies were taken from various sources in the literature (Table~1) except for few disk galaxies for which $r_{\\rm e}$ was obtained from our own photometric decomposition \\citep[following][]{Mendez2008} of the $K$-band images retrieved from the archive of the Two Micron All Sky Survey \\citep[][hereafter 2MASS]{2MASS}.  In Figs.~\\ref{fig:ULmsigmaFF05} and \\ref{fig:ULmsigmaLauer07} we compare our \\mbh\\ upper limits to the \\ms\\ relation, as given by both \\citet{Ferrarese2005} and \\citet{Lauer2007b}, initially to establish the validity of our method over a wide range of velocity dispersions. Our upper bounds show a well defined trend with both \\sigmac\\ and \\sigmae, running closely above the \\msc\\ and \\mse\\ relations. In the \\msc\\ plane a Spearman's rank coefficient of 0.9 suggests the presence of a correlation at 9$\\sigma$ confidence level whereas a Pearson correlation coefficient of 0.8 supports a linear fit to the logaritmic data, which returns a slope of $3.43\\pm0.21$. At first glance such a value would imply a shallower trend than found by \\citet{Ferrarese2005} and a slope closer to that of the \\citet{Lauer2007b} relation, but we need to keep in mind that the derived slope could be significantly affected by just a few outliers. In particular, for small value of $\\sigma$ our \\mbh\\ upper-limits could be biased owning to a larger stellar contribution to the gravitational potential in small and distant galaxies. On the other hand, we found that our limits appear to parallel particularly well both versions of the \\ms\\ relation for $90 \\leq \\sigma_{\\rm c}, \\sigma_{\\rm e} \\leq 220$ \\kms, whereas at lower and higher $\\sigma$ a substantial fraction of our \\mbh\\ limits lie either considerably above or almost on top of the \\ms\\ relation, respectively.  In the following sections we better quantify and interpret these first considerations.  \\subsection{Main trend in the sample} \\label{subsec:coeff}  In the \\sigmac\\ interval between $90$ and $220$ \\kms\\ our \\mbh\\ upper limits appear to correlate particularly well with \\sigmac, paralleling the \\msc\\ relation. In this \\sigmac\\ region a value of 0.8 for the Spearman's rank correlation coefficient suggests the presence of a correlation at a 7-$\\sigma$ confidence level, whereas a Pearson coefficient of 0.8 indicates that the logarithm values of our \\mbh\\ upper limits and \\sigmac\\ are very likely to be linearly correlated. A linear fit in the $\\log{\\sigma_{\\rm c}}-\\log{M_\\bullet}$ plane delivers a best-fitting slope of $4.52\\pm0.41$ for our \\mbh\\ upper limits, compared to the $4.86\\pm0.43$ slope of the \\cite{Ferrarese2005} relation, with a scatter of 0.39 dex (Figs.~\\ref{fig:histograms}a). In the $90 \\leq \\sigma_{\\rm c} \\leq 220$ \\kms\\ interval we have 66 \\mbh\\ upper limits, which have a median 2.7 times higher than the expected \\mbh\\ value (Fig.~\\ref{fig:histograms}c). These upper-limits can range from falling short of the expected \\mbh\\ values by a factor 3.7 to exceeding them by a factor 17.3, although 68\\% of them actually do not top the expected \\mbh\\ values by more than a factor 4.1 and fall immediately above the \\msc\\ relation. For comparison, by fitting our upper limits in the \\mse\\ plane we obtain a slope of $4.12\\pm0.38$, very close to the value of $4.13\\pm0.32$ found by \\citet{Lauer2007b}. In fact, the parallel trend of our upper limits holds as far as $\\sigma_{\\rm e} \\sim 300$ \\kms, with a Spearman's rank coefficient of 0.8, a Pearson linear correlation coefficient of 0.8 and a linear slope of $3.84\\pm0.28$.  \\placefigthree \\placefigfour \\placefigfive  Fig.~\\ref{fig:ULmsigmaFF05}a shows that such a trend holds independent of galactic distance -- objects as far away as 60 Mpc appear to run parallel to the \\msc\\ relation. In particular, the objects at and below 20 Mpc are well distributed. In fact, if in this range of \\sigmac\\ we perform separate linear regression for the three different populations of upper limits with $D < 30$ Mpc, $30< D <60$ Mpc and $60< D <100$ Mpc, we find slope values that are consistent within the errors, namely of 4.05$\\pm$0.51, 3.51$\\pm$1.01, 4.52$\\pm$1.14, respectively. This finding shows that the observed nuclear line widths do not simply trace an increasingly larger subtended stellar mass, as galaxies with progressively larger stellar velocity dispersions are found preferentially at larger distances. Instead, the fact that our upper limits scale in the same way with $\\sigma_c$ as precisely-measured \\mbh\\ determinations indicates that the nuclear emission we measured arises predominantly in regions of the gravitational potential that are dominated by the influence of the central SMBHs.  This is not completely unexpected given that a number of HST observations revealed that the narrow-line regions of active nuclei appear to be quite concentrated with scales less than 50 pc, much more so than the underlying stellar density profile \\citep[see, e.g.,][]{Ho2008}. Most recently, \\citet{Walsh2008} have mapped the behavior of the narrow-line region for galaxies observed with multiple-slit STIS observations. They found that all galaxies of their sample exhibit a centrally peaked surface-brightness profile, with the majority of them further showing a marked gradient of the emission-line widths within the sphere of influence of the central SMBH. The high degree of concentration of the gaseous tracer of the gravitational potential is what allows us to closely trace the presence of the SMBH even in objects where formally its sphere of influence is not resolved. This is similar to the case of the stellar-dynamical estimates of \\mbh\\ for M32, which have not significantly changed when moving from ground- to space-based observations \\citep[e.g.,][and references therein]{Kormendy2004} due to the exceptional concentration of its stellar light profile. Actually, that fact that our upper limits run so closely to the \\msc\\ relations also suggests that non-gravitational forces do not generally contribute much to the observed line widths (unless for some reason their importance scales with \\sigmac), although their role cannot be firmly excluded on a single-case basis.  Fig.~\\ref{fig:ULmsigmaFF05}b also shows that in the $\\sigma_{\\rm c} = 90 - 220$ \\kms\\ range the upper limits derived in galaxies of different Hubble types lie neither closer nor further away from the \\msc\\ relation, although only a relatively small number of elliptical galaxies falls in this \\sigmac\\ interval. Similarly, even though only $38\\%$ of the spiral and lenticular galaxies in this \\sigmac\\ region are unbarred we do not notice any systematic trend with the presence of a bar, unlike what was found by \\citet{Graham2008}.  Since our upper limits appear to trace quite closely the expected values for \\mbh, we can take advantage of the significant number of galaxies in our sample to understand whether the objects that within the present \\sigmac\\ range appear to show remarkably large or small upper-bounds are in fact exceptional. Assuming that our $1\\sigma$ limits bracket symmetrically the expected values of \\mbh\\ and that our upper bounds lie 2.7 times above the \\msc\\ relation, with the aid of a Monte Carlo simulation we found that $16\\%$ of our \\mbh\\ upper limits should lie above the \\msc\\ relation by more than 3 times its scatter \\citep[adopting 0.34 dex by][]{Ferrarese2005}, while $8\\%$ of them should lie below it by more than its scatter. As Fig.\\ref{fig:ULmsigmaFF05} shows, only four out of 66 objects in the $\\sigma_{\\rm c} = 90 - 220$ \\kms \\ range fall that far above the \\msc\\ relation, with an equal number falling below it by more than its scatter. Both sets of objects correspond to $6\\%$ of the galaxies in the considered \\sigmac\\ range.  The previous considerations strongly argue against the presence of exceedingly large \\mbh\\ (i.e., above the \\msc\\ relation by more than 3 times its scatter) in nearby galactic nuclei, and further suggests that galaxies with considerably smaller \\mbh-budgets (i.e., below the \\msc\\ relation by more than its scatter) may be particularly rare. The presence of undermassive SMBHs in field galaxies have been suggested for instance by \\cite{Vittorini2005}, who argued that in a low galactic-density environment the \\mbh\\ growth may be hampered by the lack of gaseous fuel. A population of undermassive SMBH was discovered also by \\citet{Volonteri2007} in her simulations of the last stages of black-hole mergers, when the binary experiences a recoil due to asymmetric emission of gravitational radiation. According to Fig.~4 of \\citet{Volonteri2007} up to $25\\%$ of the galaxies with $90 \\leq \\sigma_{\\rm c} \\leq 220$ \\kms\\ could contain undermassive SMBHs, with less than $10\\%$ the expectected \\mbh. Unfortunately, only a handful of objects in our sample are massive and close enough (e.g., for $\\sigma_{rm c} > 150$ \\kms\\ and $D< 20$ Mpc) to allow us to probe such a low \\mbh\\ regime, where our simulations indicate that we should expect only $\\sim1$\\% of our upper-limits. Furthermore, although the wide range of Hubble types and of values for \\sigmac\\ spanned by our sample galaxies suggest these are fairly representative of the general properties of the nearby population, our sample is almost certainly incomplete, in particular as we probe the low-end of the luminosity function where most galaxies are found. Our constraints should therefore be regarded with caution.  \\subsection{The lower end of the \\msc\\ relation} \\label{subsec:low}  At small \\sigmac\\ ($< 90$ \\kms) half of our upper limits systematically exceed the expected \\mbh\\ values by 3 times the scatter of the \\msc\\ relation. They are in average larger by more than a factor 40 (Fig.~\\ref{fig:histograms}d), consistent with previous works on much more smaller samples \\citep{Sarzi2002, Sarzi2004, VerdoesKleijn2006}. They are hosted by NGC~3021, NGC~4245, NGC~5347, NGC~5427, NGC~5879, and UGC~1395, which are late-type spirals with different degrees of nuclear activity, as measured by \\citet[][see Table~1]{Ho1997}.  We have considered different possibilities related to the measurement and modeling of the \\niig\\ emission line to explain the high values of \\mbh\\ found in these 6 objects. For instance, the presence of broad or asymmetric components in our spectra could affect the width of the narrow component of the \\nii\\ lines that we measured and consequently the \\mbh\\ upper limits giving larger masses. A similar bias would be introduced if the extent of the flux profile were to be systematically overestimated. Blue asymmetries are observed in the top outliers NGC~5347 and NGC~5427, which are also part of the sample of active galactic nuclei studied by \\citet{Rice2006}, who investigated the resolved kinematics of their narrow-line region with STIS spectra and first reported the presence of blue wings in the \\siipg\\ lines. In our fits, however, the contribution of such additional features was isolated using double Gaussian profiles. As regards the flux profile of the \\nii\\ doublet of the small-\\sigmac\\ outliers, these are not systematically shallower than the other galaxies following the \\msc\\ relation in the same \\sigmac\\ range. Therefore, the \\mbh\\ upper limits that we have calculated are not biased by either of these effects.  To explain the largest \\mbh\\ upper limits found at low \\sigmac\\ values we also considered the impact of the presence of a nuclear star cluster (NC), and in general that of a larger stellar contribution due to a greater distance. NCs are massive stellar clusters coincident with the galaxy photocenter \\citep{Cote2006} that are found in about $75\\%$ of late-type spiral galaxies \\citep{Boker2002}. Their mean effective radius is $\\sim 3.5$ pc \\citep{Boker2004}, small enough for them to be completely enclosed within the central aperture of our spectra. \\citet{Ferrarese2006} found a different \\msc\\ relation for NCs, with similar slope but a normalization that is larger by roughly an order of magnitude than that the one found for SMBHs. The presence of NCs in our low-\\sigmac\\ outliers could therefore explain why they show such high central mass concentrations as indicated by their high \\mbh\\ values. To assess the incidence of NCs in the sample galaxies with $\\sigma_{\\rm c} < 90$ \\kms\\ we analyzed their surface-brightness radial profile obtained with the IRAF task ELLIPSE on the STIS acquisition images. For half of the low-\\sigmac\\ outliers we could recognize the presence of a NC (NGC~3021, NGC~4245, and NGC~5879). On the other hand, we could identify a NC only in one (NGC~4212) of the six galaxies ($17\\%$) which run close to the \\msc\\ relation (IC~342, NGC~2685, NGC~2748, NGC~3982, NGC~4212, and NGC~5194). The presence of a NC in the galaxies at the low-\\sigmac\\ end of our sample is shown in Fig.~\\ref{fig:ULmsigmaFF05}, and in the case of NGC~3021 and and NGC~5879 it was already known \\citep[see][respectively]{Scarlata2004,Seth2008}. If our limits trace indeed the dynamical signature of a NC in these nuclei, better data and more detailed modelling \\citep[e.g.,][]{Barth2008} would be required to disentagle the contribution of the NC and SMBH to the total mass budget. As regards the distance of the low-\\sigmac\\ outliers, although we can only rely on distances inferred from their recessional velocities it is significant that half of them (NGC~5347, NGC~5427, and UGC~1395) are found beyond 30 Mpc, whereas all the other low-\\sigmac\\ galaxies are significantly closer, including those for which our \\mbh\\ upper limits lie well within 3 times the scatter of \\msc\\ relation.  These findings suggest that part, if not all, of the exceedingly large \\mbh\\ values we found at the low-\\sigmac\\ end of the \\msc\\ relation could be ascribed to a more significant stellar contribution to the gravitational potential. This is either because of the presence of a nuclear stellar cluster (in NGC~3021, NGC~4245, and NGC~5879) or due to a larger galactic distance (for NGC~5347, NGC~5427, and UGC~1395) than otherwise required to trace the \\msc\\ relation at these \\sigmac\\ regimes. Therefore, we presently do not need to invoke either non-gravitational forces \\citep[see][]{Sarzi2002} or a population of more massive SMBHs \\citep[see][]{Greene2006} to explain the observed flattening of the \\msc\\ relation at low \\mbh\\ values.   \\subsection{The upper end of the \\msc\\ relation} \\label{subsec:high}  At high \\sigmac\\ ($> 220$ \\kms) our \\mbh\\ upper limits nicely bracket the \\msc\\ relation (Fig.~\\ref{fig:histograms}e) and most of them are consistent with its scatter (Fig.~\\ref{fig:ULmsigmaFF05}). In fact, only four objects ($15\\%$; NGC~2911, NGC~4552, NGC~4594, and NGC~5077) fall above the \\msc\\ relation by more than its scatter, with the same number of galaxies falling as far below the \\msc\\ relation (NGC~3998, NGC~4278, NGC~6861, and UGC~1841). These outliers do not stand out from the rest of the objects with \\sigmac$ > 220$ \\kms\\ for any obvious property such as morphology, nuclear activity or distance. This behavior is suggestive of an actual flattening of the high-mass end of \\msc\\ relation, in particular considering that in the most massive and radio-loud galaxies the ionized-gas velocity dispersion can show a significant excess over a purely gravitational model \\citep[e.g.,][]{VerdoesKleijn2006}\\footnote{In fact, for the two radio-loud galaxies in our sample that were also studied by Verdoes Kleijn et al. (NGC~383 and UGC~7115) the derived \\mbh\\ upper limits lie above the \\msc\\ relation}.   The flattening at high-$\\sigma$ values is less evident when our upper limits are compared to the shallower \\mse\\ relation of \\citet{Lauer2007b}, but is nonetheless present upon closer inspection. In particular, excluding objects with \\sigmae$ < 90$ \\kms\\ where the impact of the stellar potential on our \\mbh\\ estimates could be more important, we found a systematic flattening in the main trend of our upper limits as the high-\\sigmae\\ end of the \\mse\\ plane is progressively populated. Specifically, whereas a linear fit to objects with \\sigmae$=90 - 220$ \\kms\\ yields a slope of $4.12\\pm0.38$ (\\S3.3), extending this range to $280$ \\kms, $340$ \\kms\\ and up to the maximum \\sigmae\\ value in our sample of 386 \\kms\\ results in best-fitting values of $3.86\\pm0.29$, $3.78\\pm0.27$ and $3.56\\pm0.26$, respectively.  This finding would be in agreement with the predictions of semi-analytic models for the coevolution of SMBHs and galaxies at the highest end of the mass spectrum, whereby galaxies and SMBHs grow mainly via gas-poor, dry mergers \\citep{Schawinski2006}. Yet, the behavior of \\msc\\ relation in this regime is still under debate. In particular, the limited number of galaxies with reliable \\mbh\\ measurement in the range \\mbh $ > 10^9$ \\msun\\ are actually consistent with a steepening of the \\msc\\ relation \\citep[e.g.,][]{Wyithe2006,DallaBonta2008}. Furthermore, the cutoff at $\\sigma_{\\rm c} \\sim 400$ \\kms\\ of the local velocity dispersion function \\citep{Seth2008} implies either that SMBHs with \\mbh$>3 \\times 10^9$ \\msun\\ are extremely rare or that if they exists their host galaxies should lay considerably above the present \\msc\\ relation. In fact, at these regimes \\citet{Lauer2007} argue that the stellar luminosity $L$ is better suited than \\sigmac\\ to trace \\mbh.  The \\msc\\ relation should steepen at its high-\\sigmac\\ end if Lauer et al. arguments are correct, since the observed \\sigmac\\ saturates for the most massive of ellipticals while considering increasingly large values of $L$.  Although our results suggest a flattening of the \\msc\\ relation, we need to keep in mind that systematic effects related to the measurement of the bulge properties may be significant at the high-\\sigmac\\ end of the \\msc\\ plane. In particular, the aperture correction for the stellar velocity dispersion may be both more important and more uncertain for the most massive of ellipticals than for smaller elliptical and lenticular galaxies. Indeed, giant ellipticals tend to have shallower central surface brightness profiles than their less massive counterparts, which makes the aperture correction more sensitive to the quality and spatial coverage of the stellar kinematics and to uncertainties on the value of the galaxy effective radius, $r_{\\rm e}$. Incidentally, measurements or $r_{\\rm e}$ are also generally less accurate for giant ellipticals, due the presence of extended stellar halos. Ideally, rather than \\sigmac\\ one would like to have a quantity that is more closely connected to the stellar mass, such as the total $K$-band luminosity, which is also known to relate to \\mbh\\ \\citep{Marconi2003a}, or a {\\it direct} measurement of \\sigmae. Obtaining the $K$-band luminosity of our sample galaxies would require much deeper images than the available 2MASS data, whereas properly measuring \\sigmae\\ requires integral-field observations, such as those derived in the case of the SAURON survey \\citep{Emsellem2007}.  \\subsection{Summary} \\label{sec:concl}  Owing to the exquisite spatial resolution of HST and to the concentrated character of the ionized-gas emission in low-luminosity AGNs, we have been able to set tight upper limits on \\mbh\\ for a sample of 105 nearby galaxies ($D<100\\,\\rm Mpc$) using STIS/G750M spectra. This sample spans a wide range of Hubble types (with 54\\% of spirals) and includes objects with published values for their central stellar velocity dispersion \\sigmac. Our main findings are:  \\begin{itemize} \\item Independent of the galaxy distance, our \\mbh\\ upper limits run parallel and above the \\msc\\ relation, in particular for values of \\sigmac\\ between 90 and 220 \\kms. The median of the 66 \\mbh\\ upper limits in this \\sigmac\\ regime exceeds the expected \\mbh\\ value by a factor 2.7, with 68\\% of our upper limits falling immediately above the \\msc\\ relation and without exceeding the expected \\mbh\\ values by more than a factor 4.1.  \\item That our nebular line-width measurements trace rather well the nuclear gravitational potential, makes large samples of \\mbh\\ upper-limit measurements useful in constraining the frequency of objects with exceedingly low or high values of \\mbh\\ and in probing the black-hole mass budget across the entire Hubble sequence.  \\item No systematic trends or offsets are observed in this \\sigmac\\ range as a function of the galaxy Hubble type, or with respect to the presence of a bar. Furthermore, no evidence was found to suggest that the largest or smallest \\mbh\\ upper limit in the \\sigmac\\ range between 90 and 220 \\kms\\ are actually bracketing exceptionally high or low values of \\mbh. Thus, galaxies with exceedingly high \\mbh\\ budgets must be very rare.  \\item For \\sigmac\\ values below 90 \\kms\\ half of our \\mbh\\ upper limits systematically exceed the expected \\mbh\\ values by more than a factor 40, consistent with previous work on much smaller samples.  The line-width measurements for such low-\\sigmac\\ outliers are most likely affected by the stellar contribution to the gravitational potential, either due to the presence of a nuclear stellar cluster or because of a greater distance compared to the other galaxies at the low-\\sigmac\\ end of the \\msc\\ relation, for which our \\mbh\\ upper limits are closer to the expected \\mbh\\ values.  \\item At the opposite \\sigmac\\ end of the \\msc\\ relation, for values of \\sigmac\\ above 220 \\kms, our \\mbh\\ upper bounds appear to lie much closer the expected \\mbh\\ in the most massive elliptical galaxies, even falling below the \\msc\\ relation. This flattening is less evident when our upper limits are compared with the shallower \\mse\\ relation by \\citet{Lauer2007b}, but is nonetheless present upon closer inspection.  In particular, excluding objects with \\sigmae$ < 90$ \\kms , we found a systematic flattening in the main trend of our upper limits as the high-\\sigmae\\ end of the \\mse\\ plane is progressively populated.  Although such a flattening of the \\msc\\ relations at its high-\\sigmac\\ end would appear consistent with models for the coevolution of supermassive black holes and galaxies driven by dry mergers, we caution that better and more consistent measurements for either the $K$-band luminosity or the integrated value of the stellar velocity dispersion \\sigmae\\ within the bulge effective radius $r_{\\rm e}$ (both better tracers of the bulge mass than \\sigmac) are needed before systematic effects can be ruled out.\\\\  \\end{itemize}  \\bigskip \\noindent {\\bf Acknowledgments.} We acknowledge the anonymous referee for his/her many comments that improved our manuscript. We are indebted with James Binney, St\\'ephane Courteau, Sadegh Khochfar, Lorenzo Morelli, Kevin Schawinski, and Massimo Stiavelli for many useful discussions and suggestions.  We are also indebted with Jonelle Walsh and her collaborators for sharing with us the results of their work prior to publication.  We thank Jairo M\\'endez Abreu for the GASP2D package which we used for measuring the photometric parameters of some of the sample galaxies. We acknowledge the grant CPDA068415/06 by Padua University, which provided support for this research. AB is grateful to the University of Hertfordshire for its hospitality while this paper was in progress. This research has made use of the Lyon-Meudon Extragalactic Database (LEDA), NASA/IPAC Extragalactic Database (NED), and the Two Micron All Sky Survey (2MASS)."
    },
    "0809/0809.0510_arXiv.txt": {
        "abstract": "The progenitors of SN 2008S and the 2008 luminous transient in NGC 300 were deeply dust-enshrouded massive stars, with extremely red mid-infrared colors and relatively low bolometric luminosities ($\\approx5\\times10^4$\\,L$_\\odot$). The transients were optically faint compared to normal core-collapse supernovae, with peak absolute visual magnitudes of $-13\\gtrsim M_V\\gtrsim -15$, and their spectra exhibit narrow Balmer and [Ca II] emission lines. These events are unique among transient-progenitor pairs and hence constitute a new class. Additional members of this class may include the M85 transient, SN 1999bw, 2002bu, and others.  Whether they are true supernovae or bright massive-star eruptions, we argue that their rate is of order $\\sim20$\\% of the core-collapse supernova rate in star-forming galaxies.  This fact is remarkable in light of the observation that a {\\it very} small fraction of all massive stars in any one galaxy, at any moment, have the infrared colors of the progenitors of SN 2008S and the NGC 300 transient.  We show this by extracting mid-infrared and optical luminosity, color, and variability properties of massive stars in M33 using archival imaging.  We find that the fraction of massive stars with colors consistent with the progenitors of SN 2008S and the NGC 300 transient is $\\lesssim10^{-4}$.  In fact, only $\\lesssim10$ similar objects exist in M33 (and perhaps $\\lesssim1$) --- all of which lie at the luminous red extremum of the AGB sequence. That these transients are simultaneously relatively {\\it common} with respect to supernovae, while their progenitors are remarkably {\\it rare} compared to the massive star population, implies that the dust-enshrouded phase is a short-lived phenomenon in the lives of many massive stars.  This shrouded epoch can occur only in the last $\\lesssim10^4$\\,yr before explosion, be it death or merely eruption.  We discuss the implications of this finding for the evolution and census of ``low-mass'' massive stars (i.e., $\\sim$$8-12$\\,M$_\\odot$), and we connect it with theoretical discussions of electron-capture supernovae near this mass range.  Other potential mechanisms, including the explosive birth of massive white dwarfs, and massive star outbursts are also discussed. A systematic census with (warm) {\\it Spitzer} of galaxies in the local universe ($D\\lesssim10$ Mpc) for analogous progenitors would significantly improve our knowledge of this channel to massive stellar explosions, and potentially to others with obscured progenitors. ",
        "introduction": "\\label{section:introduction}  Identifying the progenitors of core-collapse supernovae, the outbursts of Luminous Blue Variables (LBVs), and other massive-star transients is essential for understanding the physics, demographics, variability, evolution, and end-states of massive stars.  The problem of identifying the progenitors of bright transients from massive stars is difficult, and traditionally limited to serendipitous archival imaging of nearby galaxies in the optical and near-infrared (e.g., Van Dyk et al.~2003; Smartt et al.~2004; Li et al.~2007; see the extensive summary in Smartt et al.~2009).  Progenitor searches are complemented by statistical studies of supernova environments within their host galaxies, which provide indirect evidence for associations between certain types of supernovae and broad classes of progenitors (e.g., James \\& Anderson 2006; Kelly et al.~2008; Prieto et al.~2008; Anderson \\& James 2008).  A much more direct method for understanding the relation between types of massive stars and their transients is to catalog all of the massive stars in the local universe ($D\\lesssim10$\\,Mpc) before explosion. While surveys for bright optical transients in the local universe are well-developed (e.g., Li et al.~2001), a fairly complete census of the massive stars in nearby galaxies has only recently been proposed and undertaken (Massey et al.~2006; Kochanek et al.~2008).  Despite the technical challenges required by the depth, area, and cadence of the observations, these surveys are critical for our understanding of the one-to-one correspondence between massive stars and their end-states, whether they are successful or failed explosions (e.g., Kochanek et al.~2008). The long-term promise of these  surveys is to produce a catalog within which the characteristics of progenitors of future supernovae are listed pre-explosion, as in the case of SN 1987A (West et al.~1987; Menzies et al.~1987). They will provide an essential mechanism for understanding the direct causal mapping between individual progenitor types and their transients (see, e.g., Gal-Yam et al.~2007).  Here, we describe a new link in this causal mapping. The discovery by Prieto et al.~(2008a) and Prieto (2008) of the dust-obscured progenitors of the luminous outbursts in NGC 6946 (SN 2008S; Arbour \\& Boles 2008) and in NGC 300 (Monard 2008) with {\\it Spitzer}, opens up qualitatively new possibilities in the study of the connection between massive stars and their explosions. We show that these discoveries allow us to make a strong --- and perhaps unprecedented --- connection between a dust-enshrouded sub-population of massive stars and a new relatively common class of bright transients.  \\\\   The argument presented in this paper can be summarized in four points: \\\\  \\noindent 1.~{\\it The transients SN 2008S and NGC 300\\footnote{Throughout this paper we denote ``the transient in NGC 300'' as ``NGC 300'' (e.g., ``the progenitor of NGC 300'') unless we specifically refer to the host galaxy.} constitute a class.}  Both have peak absolute $V$-band magnitude $M_V\\approx-14\\pm1$ ($\\approx2-3$ magnitudes fainter than normal core-collapse supernovae; e.g., Richardson et al.~2002), strong evidence for internal extinction in their spectra, narrow emission lines (similar to low-luminosity Type IIn supernovae; e.g., Filippenko 1997), and progenitors that are optically obscured and deeply dust-embedded (dust-reprocessed emission giving blackbody temperatures $\\lesssim500$\\,K), and with bolometric luminosities of $\\sim5\\times10^4$\\,$L_{\\odot}$. In addition, they show little infrared variability on a few-year baseline.  The details of this unique class of progenitor-transient pairs and its members, as well as a comparison with other classes of optical transients are presented in \\S\\ref{section:class}.  An in-depth search for analogs to the progenitors of SN 2008S and NGC 300 in M33 is presented in \\S\\ref{section:m33}.\\\\  \\noindent 2.~{\\it Transients of this type are relatively common with respect to core-collapse supernovae.}  A total of $\\approx$\\,22 core-collapse supernovae or supernova-like transients have been discovered within $\\approx10$\\,Mpc since 1999. Sixteen were confirmed supernovae, two were LBV eruptions, one was a Type IIn supernova (SN 2002bu in NGC 4242, $D\\approx8$\\,Mpc; Puckett \\& Gauthier 2002), whose relatively low peak magnitude ($M_V\\approx-15$; Hornoch 2002)  suggests a close similarity with the remaining two bright transients (see \\S\\ref{section:other}), which are of primary interest in this paper: SN 2008S ($D\\approx5.6$\\,Mpc) and NGC 300 ($D\\approx1.9$\\,Mpc), whose physical nature is uncertain.   They may be either true (but optically sub-luminous) supernovae, or a new class of massive star eruptions. Taken at face value, these numbers imply that the rate of SN 2008S-like transients is of order $\\sim10$\\% of the supernova rate.\\footnote{Throughout this paper we use ``supernova rate'' to mean the core-collapse supernova rate unless we specifically mention the contribution from Type Ia supernovae.}  Because of incompleteness, the true rate is likely higher. We discuss the frequency of these events in detail in \\S\\ref{section:rates}.\\\\  \\noindent 3.~{\\it The progenitors of this class are extremely rare among massive stars at any moment, in any star-forming galaxy.} Although the bolometric luminosities of the progenitors of SN 2008S and NGC 300 are unremarkable for massive stars ($\\sim5\\times10^4$\\,L$_\\odot$), their colors put them in a class consisting of less than $10^{-4}$ of all massive stars ($[3.6]-[4.5]\\mu$m color $\\gtrsim2.0$ and $\\approx2.7$, respectively). In a mid-infrared (MIR) color-magnitude diagram (see Figs.~\\ref{fig:cmd} \\& \\ref{fig:cmd2}), these progenitors lie at the extremum of the AGB sequence in both luminosity and MIR color.  We refer to them as ``extreme AGB'' (EAGB) stars throughout this work. Because of their relatively low bolometric luminosities, they are not $\\eta$ Carinae, cool hypergiant, or classical LBV analogs (see, e.g., Smith 2008). They are thus distinct from the ``supernova impostors,'' produced by bright outbursts of optically-luminous LBVs (e.g., SN 1997bs; Van Dyk et al.~2000, 2003). In \\S\\ref{section:m33} we present results from a comprehensive survey of M33 for massive stars with properties  similar (in bolometric luminosity, obscuration, and variability) to the progenitors of SN 2008S and NGC 300. We find remarkably few.  We compare with MIR surveys of the LMC, SMC, and NGC 300.\\\\ ",
        "conclusions": "\\label{section:discussion}  We have shown that in the primary metrics of color, luminosity, and variability, stars analogous to the progenitors of SN 2008S and NGC 300 are exceedingly rare in star-forming galaxies. They have luminosities characteristic of low-mass massive stars, are deeply dust-obscured with extremely red MIR colors, and show little MIR variability (see Figs.~\\ref{fig:cmd}, \\ref{fig:cmd2}, \\ref{fig:sed}, and \\ref{fig:rms}).  In luminosity and color they are distinct from the population of optically-luminous LBV candidates selected from M06.  Although many of the reddest objects selected as EAGB stars in Figure \\ref{fig:cmd2} are highly variable, the few ($\\sim1-5$) least-variable sources most closely resemble the SN 2008S and NGC 300 progenitors.  In this way (but in only this way), they are similar to the LBV candidates.  In this section we discuss the implications of our finding that stars with characteristics analogous to the progenitors of SN 2008S and NGC 300 are rare. In \\S\\ref{section:numbers} we estimate the overall fraction of massive stars that are deeply dust-embedded and the lifetime of stars in that state.  In \\S\\ref{section:theory}, we connect with the evolution of massive stars, including the possibility that SN 2008S-like transients are the result of electron-capture supernovae or massive white dwarf birth.  Section \\ref{section:survey} discusses how many EAGB stars can be found within the local universe ($D\\lesssim10$\\,Mpc) using {\\it Spitzer}.  \\subsection{Numbers \\& Rates} \\label{section:numbers}  The total number of analogs to the progenitors of SN 2008S and NGC 300 in M33 is uncertain.  This uncertainty comes primarily from the fact that we are unable to identify an absolute quantitative criterion for inclusion into this progenitor class based on our three primary metrics: color, luminosity, and variability.  To be conservative, we have identified 18 sources in the region  $M_{4.5}<-10$ and $[3.6]-[4.5]>1.5$ of the CMD that satisfy the minimal criteria of being bright and extremely red (larger open circles in Fig.~\\ref{fig:cmd2}). However, a very strict cut in color and magnitude (i.e., all sources redder than the lower limit to SN 2008S and brighter than NGC 300) yields just two sources.  Among the 18 selected sources, we have shown that roughly half do not have bolometric luminosities indicative of massive stars.  That is, they do not have luminosities as large as one would expect for stars who are traditionally thought to end their lives as supernovae ($L_{\\rm bol}\\gtrsim4\\times10^4$\\,L$_\\odot$). It is important to note, however, that at fixed final luminosity of a massive star its initial ZAMS stellar mass may be multi-valued, implying progenitors with either $\\sim5-7$, $\\sim8-9$, and $\\sim11-14$\\,M$_\\odot$ for $L\\approx5-10\\times10^4$\\,L$_\\odot$. (see Fig.~2 of Smartt et al.~2009; \\S\\ref{section:theory}; see footnote \\ref{foot_p08b}). Of course, the explosions SN 2008S and NGC 300 may not be true supernovae, but rather a new class of bright eruptions from obscured massive stars (see \\S\\ref{section:theory}). Nevertheless, comparing these 18 sources to the progenitors of SN 2008S and NGC 300 in Figure \\ref{fig:sed} we would argue that only roughly half belong to this progenitor class based on SED alone.  Finally, Figure \\ref{fig:rms} shows that only $\\approx1-5$ of the 18 sources vary as {\\it little} as the progenitors of SN 2008S and (potentially) NGC 300.  Importantly, 16 of the 18 sources satisfy the criterion of being highly optically-obscured (see Table \\ref{table:eagb}).  In summary, very few massive stars have the color, luminosity, and variability of the SN 2008S and NGC 300 progenitors.  Our best guess is that the number of true analogs may be as few as zero and as large as $\\sim10-20$.  We denote this number in M33 --- the number of {\\it true} analogs --- as $N_{\\rm EAGB}$. A larger sample of stars, culled from a larger multi-epoch study of local star-forming galaxies (see \\S\\ref{section:survey}) is clearly needed to fill in the parameter space in the extreme red and bright side of the CMD.  This is the most robust way to understand $N_{\\rm EAGB}$ and its uncertainty.  In order to evaluate the {\\it fraction} of stars in M33 that might be analogs to the progenitors of SN 2008S and NGC 300, and so constrain the rate of production of such objects, we must first estimate the total number of massive stars in M33 ($N_\\star$; i.e., with $M_{\\rm ZAMS}\\gtrsim8$\\,M$_\\odot$). This number can be estimated in several ways: (1) extinction-corrected H$\\alpha$ luminosity (e.g., Hoopes et al.~2001; Hoopes \\& Walterbos 2000; Greenawalt 1998), (2) dust-reddening corrected UV continuum luminosity (for {\\it GALEX} observations, see Thilker et al.~2005), (3) total number of main-sequence optical point sources detected with $M_V\\lesssim-2$ (appropriate for stars with ZAMS masses above $\\approx9-10$\\,M$_\\odot$; Lejeune \\& Schaerer 2001), or (4) total number of red supergiants (RSGs) ($M_V\\lesssim-3.5$; M06) times the ratio of the lifetime of a massive star to the time spent as an RSG ($t_\\star/t_{\\rm RSG}\\approx10$; e.g., Schaller et al.~1992). Using the latter method, and selecting RSGs with $V-R>0.5$ and $M_V<-3.5$ from the catalog of M06, we find $\\approx5400$ sources, implying $N_\\star\\approx5.4\\times10^4$.  Taking  a more conservative color cut of $V-R>0.7$ and $M_V<-3.5$, we find that  $N_\\star\\approx3.5\\times10^4$. Similar estimates in the range of $N_\\star\\approx 3-6\\times10^4$ are obtained using method (3) with the M06 catalog, although this estimate suffers significantly from incompleteness. We take $N_\\star=5\\times10^4$ as a fiducial number and include it in our scalings below. Note that estimates of the total star formation rate in M33 range from $\\sim0.3$ to $\\sim0.7$ M$_\\odot$ yr$^{-1}$, consistent with the UV, H$\\alpha$, and FIR luminosities (e.g., Gardan et al.~2007 and references therein), implying a supernova rate for the galaxy of $\\sim0.005$\\,yr$^{-1}$ (e.g., Gordon et al.~1998).  Taking $N_{\\rm EAGB}\\sim5$ and $N_\\star\\approx5\\times10^4$, we find that a fraction \\begin{equation} f_{\\rm EAGB}=\\frac{N_{\\rm EAGB}}{N_\\star}\\sim1\\times10^{-4} \\left[\\frac{N_{\\rm EAGB}}{5}\\right] \\left[\\frac{5\\times10^4}{N_\\star}\\right] \\label{feagb} \\end{equation} of the massive stars in M33 may be analogs to the progenitors of SN 2008S and NGC 300.  As noted in \\S\\ref{section:introduction} (point 2) and \\S\\ref{section:rates}, only a fraction of all massive stars go through this highly dust-enshrouded phase, and produce transients like SN 2008S and NGC 300.   Since, by assumption, roughly all of the massive stars in any galaxy become normal core-collapse supernovae (but, see Kochanek et al.~2008), the rate of SN 2008S-like explosions can be characterized by their fractional rate with respect to the overall supernova rate.  This fraction is determined by dividing the observed rate of SN 2008S-like transients by the total number of supernovae within some volume, times an incompleteness correction that accounts for the fact that SN 2008S-like transients are intrinsically less optically luminous.  Based on the numbers presented in \\S\\ref{section:rates}, we estimate that $f_{\\rm SN}\\approx0.2$, although higher and lower values are not excluded.  For example, it is possible that SN 2008S-like transients are intrinsically rare and the fact that NGC 300 and SN 2008S occurred in the same year was simply chance. Although we cannot exclude this possibility, we note that such an explanation appears improbable in the face of what is known about the rarity of their progenitors ($f_{\\rm EAGB}$; eq.~\\ref{feagb}). Conversely, it is possible that the incompleteness correction exceeds the factor of $\\sim2$ advocated in \\S\\ref{section:rates} within 10\\,Mpc and that such transients are indeed common with respect to supernovae. However, it then becomes increasingly difficult to explain why no more SN 2008S-like transients were observed within 10\\,Mpc in the last 10 years. There is no way to circumvent these uncertainties without a more complete census of progenitors and outbursts.  As stated in \\S\\ref{section:introduction}, the simplest explanation for the fact that SN 2008S- and NGC 300-like transients are simultaneously common with respect to supernovae ($f_{\\rm SN}\\sim0.2$) and that their progenitors are very rare by number at any moment, in any star-forming galaxy ($f_{\\rm EAGB}\\sim10^{-4}$) with respect to massive stars, is that a significant fraction of all massive stars ($\\sim0.2$) go through a brief evolutionary epoch in which they are highly dust-obscured, just before explosion. Taking the average lifetime of massive stars with ZAMS masses in the range of $9-10$\\,M$_\\odot$ to be $t_\\star\\approx3\\times10^7$\\,yr (e.g., Schaller et al.~1992), we find that the duration of this dust-obscured phase is \\begin{equation} t_{\\rm EAGB} \\sim 1\\times10^4 \\left[\\frac{t_\\star}{10^{7.5}\\,\\,\\rm yr}\\right] \\left[\\frac{N_{\\rm EAGB}}{5}\\right] \\left[\\frac{5\\times10^4}{N_\\star}\\right] \\left[\\frac{0.20}{f_{\\rm SN}}\\right]\\,\\,{\\rm yr}. \\label{teagb} \\end{equation} We consider the uncertainty in $f_{\\rm SN}$ to be at the factor of two level and the uncertainty in $N_\\star$ to be at the level of a factor of 1.5.  However, as we have stressed, $N_{\\rm EAGB}$ may be as much as a factor of 5 or more lower ($t_{\\rm EAGB}\\lesssim10^3$\\,yr), or a factor of $\\sim2-4$ higher ($t_{\\rm EAGB}\\sim6\\times10^4$\\,yr). To improve these numbers significantly, a careful monitoring program for optical transients like SN 2008S within the local universe ($D\\lesssim10$\\,Mpc), coupled with a survey of all local galaxies for bright MIR point sources with (warm) {\\it Spitzer} (see \\S\\ref{section:survey}) should be undertaken.  The combination of watching for more transients of this type and associating them with individual progenitors whose luminosities and variability have been cataloged will significantly decrease the uncertainty in both $f_{\\rm SN}$ and $N_{\\rm EAGB}$, and significantly increase our understanding of the causal mapping between progenitors and their outbursts. \\\\  \\subsection{Connection to The Evolution of Massive Stars} \\label{section:theory}  The relation between final luminosity and initial stellar mass may be triply-valued (Smartt et al.~2009) at $\\sim5-7$, $\\sim8-9$, and $\\sim11-14$\\,M$_\\odot$ for $L\\approx5-10\\times10^4$\\,L$_\\odot$. This relation is of course uncertain, particularly in the mass range singled out by the bolometric luminosity of the 2008S and NGC 300 progenitors near $\\sim10$\\,M$_\\odot$.  It is likely further complicated by binarity, and by the mass-loss history, metallicity, and rotation of massive stars.  Because the absolute rate of these outbursts as well as whether or not they should be associated with the death of the progenitor are still uncertain, we consider a number of potential scenarios below.  We list a subset of the possibilties in order of increasing progenitor mass.   \\subsubsection{Massive White Dwarf Birth: $M\\approx6-8$\\,M$_{\\odot}$}  We have referred to the progenitors of SN 2008S and NGC 300 throughout this work as extreme (``E'') AGB stars because they lie at the red extremum of the AGB sequence in the MIR color-magnitude diagram (Figs.~\\ref{fig:cmd} \\& \\ref{fig:cmd2}). Taken literally, these stars may indeed be the progenitors of the most massive O-Ne-Mg white dwarfs, undergoing explosive core-envelope separation as they transition to proto-planetary nebulae (e.g., Riera et al.~1995; Garc{\\'{\\i}}a-Hern{\\'a}ndez et al.~2007). Perhaps the 2008S and NGC 300 progenitors were then akin to the most massive highly-evolved carbon- or oxygen-rich AGB stars (Kwok 1993).  Based on analogy with local proto-planetary nebulae, we would expect bi-polar explosion morphology and eventually the emergence of a hot ionizing continuum source as the newly-born white dwarf begins its cooling phase (perhaps similar to Hen 3-1475/IRAS 17423-1755; Riera et al.~2003).  The initial luminosity of the central source would be of order $\\sim5\\times10^4$\\,L$_\\odot$ for a white dwarf near the Chandrasekhar mass and it should cool on a timescale comparable to $\\sim10^5$\\,yr.  Thus, the bolometric luminosity of the transient should eventually decrease back to approximately pre-outburst levels.  The primary distinguishing characteristic of this particular scenario is the (eventual) emergent hot continuum source and emission lines, bi-polar morphology, and  the fact that the bolometric luminosity should not decrease to pre-outburst levels in the next decades (e.g., Kwok 1993).   \\subsubsection{Electon-Capture Supernova: $M\\approx9$\\,M$_\\odot$} \\label{section:ecsn}  The timescale estimated in equation (\\ref{teagb}) is of the right order of magnitude to be associated with the onset of carbon burning in relatively low-mass massive stars. This is traditionally a very difficult phase to model (see the summary in Woosley et al.~2002; Siess 2006; Poelarends et al.~2008).  One of the most intriguing explanations for the physics of SN 2008S-like transients is that they result from electron-capture SNe (ecSNe) of O-Ne-Mg cores of relatively low-mass massive stars (Miyaji et al.~1980). While speculative, this explanation accounts for many of the observed characteristics of both the transients and their progenitors.  In particular, it accounts for the fact that the progenitors of NGC 300 and SN 2008S were relatively low luminosity and deeply embedded.  Here we follow the scenario detailed by Poelarends et al.~(2008) (see also Nomoto 1984, 1987; Ritossa et al.~1996, Seiss 2006, 2007; as well as Chugai 1997; Wheeler et al.~1998; Woosley et al.~2002; Eldridge et al.~2007; Wanajo et al.~2009).  We know from the properties of the observed progenitors and our analysis of the luminous stars in M33 that the progenitors are extreme AGB stars.  In these systems, the combination of thermal pulses due to He shell burning and dredge-up produces a massive, dusty wind (for lower luminosity and less enshrouded analogs, see the work on carbon stars in the Magellanic clouds by Groenewegen et al.~2007, as well as van Loon et al.~2005, 2006). In the Poelarends et al.~(2008) models, the mass loss peaks at nearly $10^{-4}$\\,M$_\\odot$ yr$^{-1}$ for stars of mass $M\\simeq 9M_\\odot$ and luminosity $L\\simeq 10^5 L_\\odot$, and then drops precipitously for slightly more massive stars which can support core Neon burning, and which eventually become normal iron-core core-collapse supernovae (ccSNe) (see Fig.~13 from Poelarends et al.~2008). The thermal pulses driving the mass loss occur at very high rates for the EAGB stars (on timescales of years), suggesting that mass loss may appear as a steady wind, as seems to be required for the small variability in the lightcurves of the SN 2008S and NGC 300 progenitors (see Fig.~\\ref{fig:rms}; discussion in P08a), rather than as impulsive ejections of optically-thick shells expected for normal carbon stars. The low degree of variability seen in SN 2008S (particularly when contrasted with the EAGB stars in Fig.~\\ref{fig:rms}) might also be explained by the onset of {\\it core} carbon burning as the final phases of stellar evolution commence. Note that the work of Nomoto (1984) and (1987) implies that the EAGB envelope and mass-loss would be carbon-enhanced.  For a narrow mass range near $9$\\,M$_\\odot$, the balance between mass loss and the growth of the core allows the core to become unstable to collapse without igniting core Neon burning, leading to an ecSNe.  This occurs only for a small model-dependent mass range near $9$\\,M$_\\odot$ (see also, e.g., Nomoto 1984; Seiss 2006, 2007). Poelarends et al.~(2008) estimated a mass range of $M_{\\rm min}\\simeq9$\\,M$_\\odot$ to $M_{\\rm max}\\simeq9.25$\\,M$_\\odot$. For a standard Salpeter IMF, and assuming that all stars from $M_{\\rm max}$ to 20\\,M$_\\odot$ form ccSNe, then the ecSN fraction is just $\\approx6$\\% ($9.0\\leq M\\leq9.25$\\,M$_\\odot$). This is below our (albeit, uncertain) fiducial estimate for the rate of SN 2008S-like transients relative to the normal ccSN rate, $f_{\\rm SN}\\sim0.2$. However, other studies have found somewhat broader mass ranges for ecSNe. For example, Seiss (2007) find that $M_{\\rm max}-M_{\\rm min}\\approx1-1.5$\\,M$_\\odot$, which implies a fractional rate for ecSNe more in accord with our nominal estimate for $f_{\\rm SN}$.  While the observed rate of SN 2008S-like transients is uncertain, we have argued that they represent a modest fraction of the normal ccSN rate, consistent with a limited progenitor mass range. The relatively low luminosity of their progenitors implies that they are low-mass massive stars, potentially near the boundary between electron-capture and normal core-collapse supernovae. In addition, Kitaura et al.~(2006) argue that ecSNe should be sub-luminous compared to normal ccSNe, because of their low Ni yields, potentially explaining the low luminosity of SN 2008S-like transients.  Finally, for the fiducial ecSN model of Poelarends et al.~(2008) (see also Nomoto 1984), the  AGB phase lasts for $\\approx4\\times10^4$\\,yr, which although short with respect to the lifetime of the star itself, is of order $t_{\\rm EAGB}$ in equation (\\ref{teagb}) based on the number of analogs to the SN 2008S and NGC 300 progenitors in M33.  Of course, these timescales need not be identical, since the fiducial ecSN progenitors of Poelarends et al.~may evolve significantly in color as a result of mass-loss during their super-(``extreme'') AGB phase, becoming increasingly like the SN 2008S and NGC 300 progenitors as they approach the end of their lives.  Fortunately, this speculative explanation has at least one simple and testable prediction: there should be no surviving progenitor once the transient fades.  This may be testable in the optical, since most of the dust enshrouding the progenitor was likely destroyed by the explosion, but observations in the MIR will be required to be certain that a new shroud has not formed.  A second test is to find strong evidence in the late-time lightcurve for synthesized $^{56}$Ni.  This may be difficult both because ecSN may produce little $^{56}$Ni (Kitaura et al.~2006), and because dust, whether that remaining from the EAGB phase or dust formed in the ejecta (e.g., P08c for M85), may make it difficult to correctly measure the late-time decay rate.  {\\it Spitzer} may again be key in constraining the nature of these events because of this obscuration. A third test is to ensure that {\\it Spitzer} has surveyed all the nearby galaxies so that future examples of these transients can be causally connected to their deeply obscured progenitors.  Finally, as with ccSNe, very nearby ecSNe ($D\\lesssim300$\\,kpc) should produce neutrino signatures characteristic of neutron star formation, which detectors such as SuperKamiokande and its successors would observe (see, e.g., Thompson et al.~2003; Kistler et al.~2008).  The recent paper by Botticella et al.~(2009) in part corroborates the interpretation discussed in Prieto et al.~(2008a) and proposed here that SN 2008S may be an electron-caputre supernova. They present the late-time quasi-bolometric lightcurve of SN 2008S, which shows evidence for a power-law time dependence with a slope indicative of being powered by the radioactive decay $^{56}$Co. Although this need not uniquely signal ecSN as the physical mechanism (see \\S\\ref{section:sn}), it provides some evidence for core-collapse, and something perhaps akin to normal neutron star formation (Kitaura et al.~2006). A lightcurve with similar cadence and photometric coverage has recently been published in Bond et al.~(2009) for NGC 300.  \\subsubsection{Intrinsically Low-Luminosity Iron Core-Collapse Supernova: $M\\sim10-12$\\,M$_{\\odot}$} \\label{section:sn}  Heger et al.~(1997) discuss a mechanism for generating a potentially obscuring ``superwind'' via pulsational mass-loss in red supergiants between 10 and 20\\,M$_\\odot$ during the last $10^4$\\,yr before explosion.  The prediction of enhanced AGB-like obscuration (they compare directly with luminous OH/IR stars), the evolutionary timescale (compare their $10^4$\\,yr with our eq.~\\ref{teagb}), and the secular {\\it increase} in the fundamental mode pulsation period as the star approaches death (see their Fig.~7) are all in good agreement with the requirements on the 2008S and NGC 300 progenitors we discuss in this paper.  The physical mechanism of iron core-collapse supernovae is unknown (Rampp \\& Janka 2000; Liebend\\\"orfer et al.~2001; Thompson et al.~2003; Buras et al.~2003; Burrows et al.~2006). Recent observations hint that low-luminosity Type IIP supernovae may be more common than previously thought (e.g., Chugai \\& Utrobin 2000; Pastorello et al.~2004, 2007), particularly when one accounts for the incompleteness corrections discussed in \\S\\ref{section:rates}.  Because the mechanism of supernovae has yet to be conclusively identified, it is difficult to interpret the diversity in inferred $^{56}{\\rm Ni}$ yield physically.  In fact, that diversity may be larger than previously thought, and we are only now appreciating the existence of a very low-luminosity tail to the Type IIP luminosity function.  If so, it is natural to imagine that these low-luminosity core-collapse events might have analogs that occur in the very dusty circumstellar medium of their massive stellar progenitors, as in Heger et al.~(1997), and thus may give rise to events like SN 2008S and NGC 300.  This scenario yields many of the predictions of the ecSN scenario discussed in \\S\\ref{section:ecsn}.  Indeed, even with a complete sampling of ``ec-'' and ``cc-'' supernovae, it may be difficult to disentangle the two populations since many of the predictions --- radioactive decay powered lightcurves, potentially embedded progenitor, no ``postgenitor''  --- are the same in both.  \\subsubsection{Massive Star Outburst: $M\\approx10-15$\\,M$_{\\odot}$}  On the basis of the relatively low luminosity of their progenitors, we view the ecSN and massive white dwarf birth scenarios discussed above as the most probable explanation SN 2008S and NGC 300. Nevertheless, there is of course the possibility that they are instead a new class of outbursts from relatively low-mass massive stars, potentially analogous to the pulsational instabilities discussed in Poelarends et al.~(2008) or Heger et al.~(1997). The majority of the true ``LBV'' eruptions with documented progenitors (e.g., 1997bs, 2002kg) came from optically bright massive stars significantly more bolometrically luminous than the progenitors of SN 2008S and NGC 300. As we have shown in Figures \\ref{fig:cmd2} and \\ref{fig:sed}, the EAGB population is separate from the sources traditionally classified as LBVs: they are less bolometrically luminous and much more dust-obscured. These facts suggest that if these transients were the outbursts of massive stars then they are distinct from from the classical supernova impostors. If these events are not supernovae, but merely outbursts, then their existence is likely connected to the physics of the transition between stars that become ecSNe and/or ccSNe, and those that do not. The degree of dust-obscuration at outburst is a crucial clue to their evolution.   A simple prediction of this possibility is that the progenitors of SN 2008S and NGC 300 should eventually be re-discovered in the optical and/or infrared after the outburst emission has faded.  For further arguments on the nature of SN 2008S and NGC 300 related to this discussion, see Bond et al.~(2008), Smith et al.~(2009), and Berger et al.~(2008).  \\subsection{A More Complete Census} \\label{section:survey}  Equation (\\ref{feagb}) implies that a fraction $\\sim1\\times10^{-4}$ of the massive star population in any given galaxy appears to be in the evolutionary state that led to the explosions observed as SN 2008S and NGC 300. The simplest explanation, adopted throughout this work, is that the deeply dust-enshrouded phase marks the last $t_{\\rm EAGB}\\lesssim10^4$\\,yr (eq.~\\ref{teagb}) in the life of a fraction $f_{\\rm SN}\\approx0.2$ of the massive star population.  Compilations of star formation and supernova rates in the local universe (e.g., Ando et al.~2005) suggest that the latter is $\\approx2$\\,yr$^{-1}$ within 10\\,Mpc (see \\S\\ref{section:introduction}, point 2), implying that there are $\\sim5\\times10^6$ massive stars and $\\lesssim10^3$ EAGB stars within this volume.  If the lifetime in the pre-explosion, highly dust-obscured phase is $t_{\\rm EAGB}$ (eq.~\\ref{teagb}), we would expect to see one SN 2008S-like transient every few years, in accord with our estimate for $f_{\\rm SN}$.  A multi-epoch survey of all the local star-forming galaxies within 10\\,Mpc with {\\it Spitzer} would allow for a comprehensive census of EAGB stars.  It would significantly increase our knowledge of the variability properties and SED evolution of these objects, and it might allow us to define more strict criteria for inclusion in the class of SN 2008S/NGC 300-like progenitors.  It would therefore decrease the considerable uncertainty in $N_{\\rm EAGB}$ in equations (\\ref{feagb}) and (\\ref{teagb}).  Coupled with the supernova surveys in the local volume, such a study would improve our knowledge of the fraction $f_{\\rm SN}$ of stars that eventually go through the deeply embedded phase just before explosion.  Of course, the most intriguing possibility is that the number of true analogs to the progenitors of SN 2008S and NGC 300 is in fact $N_{\\rm EAGB}\\sim0-1$ in M33 and that $t_{\\rm EAGB}\\lesssim {\\rm few}-10^3$\\,yr.\\footnote{The lower limit here comes from the fact that pre-explosion imaging of the SN 2008S and NGC 300 progenitors establishes a few year baseline.} If so, the final catalog of EAGB stars that would be produced by a {\\it Spitzer} survey would have just $\\sim50-100$ members.  These could be followed up repeatedly, since, given these numbers one would expect to wait just $\\sim10$ years before one of these individual sources exploded. This would give a direct observational link in the causal mapping between a sub-population of massive stellar progenitors and their explosions, connecting them with a short timescale. Indeed, the ability to identify an individual star as marked for imminent death (or eruption) would be an astonishing consequence of this work."
    },
    "0809/0809.3513_arXiv.txt": {
        "abstract": "We study the cosmic velocity--density relation using the spherical collapse model (SCM) as a proxy to non-linear dynamics. Although the dependence of this relation on cosmological parameters is known to be weak, we retain the density parameter $\\Omega_{\\rm m}$ in SCM equations, in order to study the limit $\\Omega_{\\rm m} \\to 0$. We show that in this regime the considered relation is strictly linear, for arbitrary values of the density contrast, on the contrary to some claims in the literature. On the other hand, we confirm that for realistic values of $\\Omega_{\\rm m}$ the exact relation in the SCM is well approximated by the classic formula of Bernardeau (1992), both for voids ($\\delta<0$) and for overdensities up to \\mbox{$\\delta \\sim 2$ -- $3$.} Inspired by this fact, we find further analytic approximations to the relation for the whole range $\\delta \\in [-1, \\infty)$. Our formula for voids accounts for the weak $\\Omega_{\\rm m}$-dependence of their maximal rate of expansion, which for $\\Omega_{\\rm m} < 1$ is slightly smaller that $3/2$. For positive density contrasts, we find a simple relation \\begin{displaymath} \\nabla \\cdot \\mathbf{v} = 3 H_0 \\, \\Omega_{\\rm m}^{0.6} \\left[(1+\\delta)^{1/6}-(1+\\delta)^{1/2}\\right] , \\end{displaymath} that works very well up to the turn-around (i.e. up to $\\delta \\la 13.5$ for $\\Omega_{\\rm m} = 0.25$ and neglected $\\Omega_{\\Lambda}$). Having the same second-order expansion as the formula of Bernardeau, it can be regarded as an extension of the latter for higher density contrasts. Moreover, it gives a better fit to results of cosmological numerical simulations. ",
        "introduction": "\\label{sec:intr} The gravitational instability is commonly accepted as the process of large-scale structure formation in the Universe. According to this scenario, structures formed by the growth of small inhomogeneities present in the early Universe.  Gravitational instability gives rise to a coupling between the density and peculiar velocity fields of matter. On very large, linear scales, the relation between the peculiar velocity $\\mathbf{v}$ and the density contrast $\\delta$ in co-moving coordinates is \\begin{equation} \\nabla \\cdot \\mathbf{v}(\\mathbf{x}) = - H f(\\Omega,\\Lambda) \\, \\delta(\\mathbf{x}) \\,, \\label{eq:linear} \\end{equation} where $H$ is the Hubble constant. [For simplicity of notation, we use the notation ($\\Omega$, $\\Lambda$) instead of ($\\Omega_\\mathrm{m}$, $\\Omega_\\Lambda$).] The coupling constant, $f$, carries information about the underlying cosmological model and is related to the cosmological matter density parameter, $\\Omega$, and cosmological constant, $\\Lambda$, by \\begin{equation} f(\\Omega,\\Lambda) \\simeq \\Omega^{0.6} + \\frac{\\Lambda}{70} \\left(1 + \\frac{\\Omega}{2}\\right) \\label{eq:f_Omega_Lambda} \\end{equation} \\citep{LLPR}. The linear amplitude of peculiar velocities is thus sensitive to $\\Omega$; on the other hand, it is quite insensitive to~$\\Lambda$. Hence, comparing the observed density and velocity fields of galaxies allows one to constrain $\\Omega$, or the degenerate combination $\\beta \\equiv \\Omega^{0.6}/b$ in the presence of so called galaxy biasing (e.g. \\citealt{SW} for a review). This is done by extracting the density field from all-sky redshift surveys -- such as the Point Source Catalogue Redshift survey \\citep[PSCz,][]{Saund}, or the 2MASS Redshift Survey \\citep[2MRS,][]{Huch} -- and comparing it with the observed velocity field from peculiar velocity surveys. The methods for doing this fall into two broad ca\\-te\\-go\\-ries. One can use Equation~(\\ref{eq:linear}), calculating the divergence of the observed velocity field and comparing it directly with the density field from a redshift survey; this is referred to as a \\textit{density--density comparison}. Alternatively, one can use the integral form of Equation~(\\ref{eq:linear}) to calculate the predicted velocity field from a redshift survey, and compare the result with the measured peculiar velocity field; this is called a \\textit{velocity--velocity comparison}. Velocity--velocity comparisons are generally regarded as more reliable, since they involve manipulation of the denser and more homogeneous redshift catalogue data, while density--density comparisons require manipulation of the noisier and sparser velocity data. In both cases, the density and velocity fields need to be smoothed in order to reduce errors and shot noise. Velocity--velocity comparisons require a smaller size of smoothing, of a few \\hmpc. For example, \\citet{Will} used a smoothing scale of $3$ \\hmpc. Such scales are called \\textit{mildly non-linear}: the variance of the density field smoothed over the scale of a few \\hmpc\\ is of order unity.  Mildly non-linear extensions of Equation~(\\ref{eq:linear}) have been developed by a number of workers. These extensions have been based either on various analytical approximations of non-linear dy\\-na\\-mics (\\citealt{ReGe}; Bernardeau 1992, hereafter \\citealt{B92}; \\citealt{Cat}; \\citealt{Ch97}; \\citealt{ChLo}; \\citealt{Ch98}), or numerical (either N-body or hydrodynamic) simulations (\\citealt{Man}; Kudlicki \\etal 2000, hereafter \\citealt{KaCPeR}), or both (\\citealt{NDBB}; \\citealt{Gram}; \\citealt{MY}; Bernardeau \\etal 1999, hereafter \\citealt{B99}). Unlike the linear case~(\\ref{eq:linear}), the non-linear relation between the velocity divergence and the density contrast at a given point is non-deterministic (though in the non-linear regime the two fields remain highly correlated). Therefore, for a full description of the relation, the conditional means (mean $\\nabla \\cdot \\bmath{v}$ given $\\delta$ and vice versa) are not sufficient: one has to describe the full bivariate distribution function for $\\nabla \\cdot \\bmath{v}$ and $\\delta$, or at least the conditional scatter. These aspects of the velocity--density relation were studied by \\citet{Ch98} and more extensively by \\citet{B99}. However, in practical applications the intrinsic scatter in the velocity--density relation is much smaller than the one induced by observational errors, and the conditional means are sufficient.  \\citet{B99} and \\citet{KaCPeR} found that very good fits to the mean relations, obtained for the mildly non-linear fields extracted from numerical simulations, were given by modifications of the formula of \\citet{B92}. This formula describes a non-linear relation between initially Gaussian, random fields of $\\nabla \\cdot \\bmath{v}$ and $\\delta$, under the assumption of a vanishing variance of the density field (so the relation has no scatter). \\citet{B92} claimed his relation to be the same as the one exhibited in the spherical collapse model (hereafter SCM). In practical applications (namely with non-zero variance of the density field), he predicted his formula to work well in voids, but `to become very inaccurate for $\\delta$ larger than 1 or 2'.  In {\\em this\\/} paper we study the velocity--density relation in the SCM. The reason for such an approach is twofold. First, to derive his formula, \\citet{B92} used quite sophisticated methods (summing up first non-vanishing contributions from the reduced part of {\\em all-order\\/} joint moments of $\\nabla \\cdot \\bmath{v}$ and $\\delta$). On the other hand, the dynamics of the SCM is very simple and should allow to re-derive the formula of \\citet{B92} in a straightforward way. More importantly, in the SCM the relation can be easily extended to higher values of $\\delta$, with the hope that this modification will fit better the results of numerical experiments of \\citet{B99} and \\citet{KaCPeR}. The SCM is in principle insensitive to the variance of the density field (and the resulting velocity--density relation is deterministic), but in practice the variance of the smoothed density field dictates how high density contrasts can be reached.  The non-linear relation between $\\delta$ and $f^{-1} \\nabla \\cdot \\mathbf{v}$ (note the scaling $f^{-1}$) depends very weakly on cosmological parameters. \\citet{B92} analysed the $\\Omega$-dependence of the scaled velocity--density relation in the limit $\\lan\\delta^2\\ran \\to 0$ and found it to be very weak. \\citet{Bouch} showed that second and third order expansions for $\\delta$ and $f^{-1} \\nabla \\cdot \\mathbf{v}$ depend extremely weakly on $\\Omega$ and $\\Lambda$. \\citet*{SCF} demonstrated that this is the case for {\\em all\\/} orders. Specifically, they showed that perturbative solutions for the density contrast for arbitrary cosmology are, with a good accuracy, separable: $\\delta_n = D^n(t)\\, \\eps_n(\\bfx)$, where $D(t)$ is the linear growing mode for this cosmology and $\\eps_n$ is the spatial part of the $n$-th order solution for the Einstein--de Sitter model. Using the continuity equation one can then prove, by induction, that the velocity divergence depends on $\\Omega$ and $\\Lambda$ practically only through the factor $f(\\Omega,\\Lambda)$. Most generally, Nusser \\& Colberg (1998) (hereafter \\citet{NuCo}) showed {\\em the equations of motion\\/} of the cosmic pressureless fluid to be `almost independent' of cosmological parameters. The weak dependence of the scaled velocity--density relation on the background cosmological model has been also confirmed by N-body numerical simulations (\\citealt{Man}; \\citealt{B99}).  However, the $\\Omega$-dependence of the equations of motion of the cosmic dust stops to be weak when $\\Omega \\ll 1$ (see eqs.~13--14 of \\citealt{NuCo}). This regime of $\\Omega$ is not physically relevant, since the currently preferred value of $\\Omega$ is much higher. Still, \\citet{B92} derived his formula applying the limit $\\Omega \\to 0$. Therefore, in the present paper we will neglect $\\Lambda$ (setting $\\Lambda = 0$), but will retain the $\\Omega$-dependence of the equations of the spherical collapse and in particular examine the limit of small $\\Omega$.  The paper is organised as follows. Section \\ref{sec:csm} presents general assumptions, terminology and basic formulae of the spherical model. In Section \\ref{sec:pp} we focus on the factor $f$, appearing in Eq.~(\\ref{eq:linear}) and commonly approximated by Formula~(\\ref{eq:f_Omega_Lambda}), or its simplified version $f\\simeq\\Omega^{0.6}$ . Section \\ref{sec:smallOm} contains an analysis of the regime of very small $\\Omega$ and presents the resulting universal velocity--density relation. In Sections \\ref{sec:voids} and \\ref{sec:overd}, basing on analytical considerations, we derive approximations for the relation between the velocity divergence and the density contrast respectively for spherical voids and overdensities, for realistic values of $\\Omega$. These approximations constitute the main results of this paper. Section \\ref{sec:comp} gives a comparison of our fits with results of numerical simulations. We conclude in Section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} The main motivation of this paper was to rederive the formula of \\citet{B92} in a simple way, using the spherical collapse model (SCM), and to extend it to larger density contrasts, where it is no longer valid. The undertaken project abounded in surprises: \\renewcommand{\\labelenumi}{\\roman{enumi}.} \\begin{enumerate} \\item Contrary to the claim of \\citet{B99}, the formula of \\citet{B92} is not exact in the limit of an empty universe. On the contrary, it completely fails in this regime: the exact relation in the SCM is then $f^{-1} \\nabla \\cdot \\mathbf{v} = \\delta$, for an arbitrary $\\delta$. In fact, this is a general result of dynamics in a low-density universe. \\item Although the formula of \\citet{B92} fails for $\\Omega \\to 0$, where it was expected to work best, for realistic values of $\\Omega$ (say, $\\Omega > 0.1$), it describes very well the SCM velocity--density relation in voids. It also works for overdensities up to $\\delta \\sim 2$ -- $3$. \\end{enumerate}  \\ind The velocity--density relation in the SCM is given in a parametric form. Our goal here was to eliminate this parameter (at least approximately) and to provide the relation analytically. We aimed at describing the relation in the whole range $\\rho \\in (0, \\infty)$ (realistically, up to $\\rho_{\\rm vir}$). Therefore, instead of expanding it around $\\rho = \\rho_{\\rm b}$, we adopted an entirely different approach. Namely, we derived asymptotes of the relation in the highly non-linear regime: $\\rho/\\rho_{\\rm b} \\gg 1$ ($\\delta \\gg 1$) for overdensities and $\\rho_{\\rm b}/\\rho \\gg 1$ \\mbox{($0 \\le 1 + \\delta \\ll 1$)} for voids. (Although we also `expanded' around them, in a sense that we also calculated next-to leading-order terms.) These two asymptotes turned out to be qualitatively different. Inspired by their functional forms, we invented semi-phenomenological fits to the exact relation (separately for overdensities and voids), fulfilling the linear theory condition $f^{-1} \\nabla \\cdot \\mathbf{v} = \\delta$.\\\\ \\ind For overdensities, our main result is Formula~(\\ref{form:fit pot}). It describes well the exact relation in the SCM up to the turn-around (for $\\Omega = 0.25$, $\\delta_{\\rm ta} = 13.5$). As already stated, the formula of \\citet{B92} starts to deviate from the exact relation for $\\delta \\sim 3$. We have also fitted the regime $\\delta \\in (\\delta_{\\rm ta},\\delta_{\\rm vir})$, though virialisation and departures from spherical symmetry make practical applicability of the SCM in this regime very limited.\\\\ \\ind In case of voids, the most important results of this paper are Formulae~(\\ref{eq:Theta2_ansatz}) and~(\\ref{eq:Theta3_ansatz}), with $\\Theta_\\mathrm{min}$ given by Equation~(\\ref{eq:Theta_min}). Compared with the SCM, simple Formula~(\\ref{eq:Theta2_ansatz}) has a maximal error of about $2$\\% and is probably sufficient for practical applications. The formula of \\citet{B92} is an even better approximation, except for the limit $\\delta \\to -1$, where for $\\Omega = 0.25$ it has approximately $5$\\% relative error. Our more complicated formula~(\\ref{eq:Theta3_ansatz}) is extremely accurate in the whole range $\\delta \\le 0$: its maximal error is about $0.2$\\%.\\\\ \\ind An ultimate goal of studies such as the present one is to find the relation valid for realistic random cosmic velocity and density fields. Unlike the work of \\citet{B92}, our calculations were greatly simplified by the strong assumption of spherical symmetry. There is therefore no guarantee that better agreement with the SCM implies better agreement with the real relation. In order to check this issue we compared our formulae to fits to results of cosmological numerical simulations, that are present in the literature. We have found that in voids, our formulae, as well as the formula of \\citet{B92}, describe well the real relation. This is partly a consequence of the fact that voids are more spherical than overdensities. In overdensities, both our formula and that of \\citet{B92} require modification, but ours less. This discrepancy is not a failure of the latter of the two, since it has never been intended to work for $\\delta \\ga 2$. Our formula~(\\ref{form:fit pot}), having the same second-order expansion as the formula of \\citet{B92}, can be regarded as its extension into the mildly non-linear regime (for $\\delta$ up to the turn-around). We have also discussed how to (slightly) modify our formula to better fit numerical simulations.\\\\ \\ind In Section~\\ref{sec:intr} we have enlisted arguments for weak dependence of the velocity--density relation on the cosmological parameters. Therefore, in the present analysis we set $\\Lambda = 0$. To study the limit $\\Omega \\to 0$ we have retained $\\Omega$-dependence of the equations of the SCM. Analysing these equations we have confirmed that for realistic values of $\\Omega$, the $\\Omega$-dependence of the relation is indeed very weak. In final formulae it has been therefore neglected, except for Formula~(\\ref{eq:Theta_min}) for $\\Theta_\\mathrm{min}$. The difference between $\\Theta_\\mathrm{min}$ for $\\Omega = 1$ and $\\Omega = 0.25$ is about $5$\\%. In fact, if we want to have better accuracy, there is no guarantee that $\\Lambda$ does not contribute at a comparable level. It is then worth to repeat the analysis with $\\Lambda = 1 - \\Omega$. We plan to undertake such a study in the future."
    },
    "0809/0809.1385_arXiv.txt": {
        "abstract": "We consider multidimensional cosmological model with a higher-dimensional product manifold $M = \\mathbb{R} \\times \\mathbb{R}^{d_0} \\times \\mathbb{H}^{d_1}/\\Gamma$ where $\\mathbb{R}^{d_0}$ is $d_0$-dimensional Ricci-flat external (our) space and $\\mathbb{H}^{d_1}/\\Gamma$ is $d_1$-dimensional compact hyperbolic internal space. M2-brane solution for this model has the stage of accelerating expansion of the external space. We apply this model to explain the late time acceleration of our Universe. Recent observational data (the Hubble parameter at the present time and the redshift when the deceleration parameter changes its sign) fix fully all free parameters of the model. As a result, we find that considered model has too big size of the internal space at the present time and variation of the effective four-dimensional fine structure constant strongly exceeds the observational limits. ",
        "introduction": "Recent astronomical observations abundantly evidence that our Universe underwent stages of accelerating expansion during its evolution. There are at least two of such stages: early inflation and late time acceleration. The latter began approximately at the redshift $z \\sim 0.35 $ (see e.g. \\cite{FTH}) and continues until now. Thus, the construction and investigation of models with stages of acceleration is one of the main challenge of the modern cosmology. Among such models, the models originated from fundamental theories (e.g. string/M-theory) are of the most of interest.   In the present paper we consider a multidimensional cosmological model with a factorizable metric \\ba{1.1} g &=&  -e^{2\\gamma_0}d\\tau \\otimes d\\tau + a_{0BD}^2g^{(0)} + a_1^2g^{(1)} \\\\ &=&  \\Omega^2(-dt\\otimes dt + a_0^2 g^{(0)}) + a_1^2 g^{(1)}\\, , \\nn \\ea which is defined on the manifold with product topology \\begin{equation} M = \\mathbb{R} \\times \\mathbb{R}^{d_0} \\times \\mathbb{H}^{d_1}/\\Gamma\\, , \\end{equation} where $\\mathbb{R}^{d_0}$ is $d_0$-dimensional Ricci-flat external (our) space with metric $g^{(0)}: R[g^{(0)}]=0$ and scale factor $a_0$, and $\\mathbb{H}^{d_1}/\\Gamma$ is $d_1$-dimensional hyperbolic (compact) internal space\\footnote{Negative constant curvature spaces are compact if they have a quotient structure: $ H^{d_i}/\\Gamma_i$, where $H^{d_i}$ and $\\Gamma_i$ are hyperbolic spaces and their discrete isometry group, respectively.} with metric $g^{(1)}: R[g^{(1)}]=-d_1(d_1-1)$ and scale factor $a_1$. Both $a_0$ and $a_1$ depend only on time. First line in \\rf{1.1} is a metric in the Brans-Dicke frame in the harmonic time gauge where $e^{\\gamma_0}=a_{0BD}^{d_0}a_1^{d_1}$ \\cite{IMZ}. Second line in \\rf{1.1} is a metric in the Einstein frame in the synchronous time gauge. The scale factors $a_{0BD}$ of the external space in the Brans-Dicke frame is connected with the scale factor $a_0$ in the Einstein frame as follows: $a_0=\\Omega^{-1}a_{0BD}$, where conformal factor $\\Omega=a_1^{-\\frac{d_1}{d_0-1}}$. Hereafter we consider 3-dimensional external space: $d_0=3$. Harmonic time $\\tau$ is related to synchronous time $t$ as $dt=f(\\tau)d\\tau$, where $f(\\tau)=\\Omega^{-1}a_{0BD}^{3}a_1^{d_1}= a_{0BD}^{3}a_1^{3d_1/2}=a_0^{3}$ \\cite{BaukhZhuk}.  With the standard York-Gibbons-Hawking boundary term $S_{YGH}$, an action for considered model reads \\be{1.2} S = \\frac{1}{2\\kappa^2}\\int\\limits_{M}d^Dx\\sqrt{|g|}R[g]+S_{YGH}\\, , \\ee where $D=1+d_0+d_1 =4+d_1$ is a total number of dimensions and $\\kappa_D$ is a $D$-dimensional gravitational constant.  Substituting the metric \\rf{1.1} into this action and minimizing obtained Lagrangian with respect to the scale factors, we get the the following solutions\\footnote{The general method for this kind of models was elaborated in papers \\cite{Zhuk1,BZnegative}.} of the equations of motion (in the harmonic time gauge): \\be{1.3} a_0(\\tau)=A_1^{\\frac{d_1+2}{6}} \\left(\\sqrt{\\frac{2\\varepsilon}{|R_1|}}\\right)^{\\frac{d_1}{2(d_1-1)}} \\frac{\\exp\\left(-\\sqrt{\\frac{d_1+2}{12(d_1-1)}{2\\varepsilon}}\\; \\tau\\right)} {\\sinh^{d_1/[2(d_1-1]}(-\\sqrt{\\frac{d_1-1}{d_1} 2\\varepsilon}\\; \\tau)}\\, , \\ee and \\be{1.4} a_1(\\tau)=A_1\\left(\\sqrt{\\frac{2\\varepsilon}{|R_1|}}\\right)^{\\frac{1}{d_1-1}} \\frac{\\exp\\left(-\\sqrt{\\frac{3}{(d_1-1)(d_1+2)}2\\varepsilon}{\\; \\tau}\\right)} {\\sinh^{1/(d_1-1)}(-\\sqrt{\\frac{d_1-1}{d_1} 2\\varepsilon}\\; \\tau)}\\, , \\ee where  $A_1$ and $\\varepsilon$ are the constants of integration. The function $f(\\tau)$ can be easily obtained from Eq. \\rf{1.3} via expression $f(\\tau) = a_0^3(\\tau)$.  We should note, that solutions \\rf{1.3}  and \\rf{1.4} for the metric \\rf{1.1} is a particular case of so called S$p$-branes with $(p+1)-$dimensional Ricci-flat external space. In the case $d_0=3$ we obtain $p=2$. Therefore, if underlying model is $(D=11)-$dimensional M-theory, we arrive at M2-branes where the number of internal dimensions is equal to 7. It is well known that such models with hyperbolic internal space undergo the stage of accelerating expansions (see e.g. \\cite{BaukhZhuk} and references therein). However, the parameters of the model were not connected with observational data. So, in the present paper we want to use the modern cosmological data (the present day value for the Hubble parameter and the redshift when our external space transits from deceleration to acceleration) to fix all arbitrary parameters of the considered model and obtain corresponding dynamical behavior for the scale factors, the Hubble parameter, the deceleration parameter and the fine structure \"constant\". ",
        "conclusions": "In the present paper we investigate multidimensional cosmological model with Ricci-flat external space and compact hyperbolic internal space. Such pure gravitational model has the exact solution with the stage of accelerating expansion for our external space (so called Sp-brane solution). This solution depends on a number of free parameters. It is remarkable that observable cosmological parameters (such as the Hubble parameter, the parameter of deceleration) gives a possibility to fix all these free parameters and completely determinate this model. We perform this analysis in the case of M2-brane with $d_1=7$ extra dimensions. For obtained parameters, we describe the dynamical behavior of the considered model. It is shown that our external space really has the finite stage of the accelerating expansion starting at $z=0.35$ and lasting until now. However, this model has two significant drawbacks. On the one hand, the internal space is too big with respect to the standard Kaluza-Klein restrictions $a_{internal}\\leq 10^{-17}$cm and, on the other hand, this space is not sufficiently constant to satisfy the observable limits on the fine structure constant variations. Thus, these drawbacks rule out this model from the scope of viable theories\\footnote{ In papers \\cite{MB}, it is investigated the effects of matter overdensities on the time and space variations of alpha. It is shown that a far slower evolution of alpha will be found in virialised regions (such as the Earth or the Solar System) than in the cosmological background. Thus, this effect gives a possible way to alleviate our conclusions concerning incompatibleness of the considered model with the observations on variations of alpha. However, the problem of too big size of the extra dimensions still remains.}."
    },
    "0809/0809.3380_arXiv.txt": {
        "abstract": "The Lagoon Nebula is an {\\sc Hii} region in the Sagittarius Arm, about 1.3~kpc away, associated with the young (1--3~Myr) open cluster NGC\\,6530, which contains several O stars and several dozen B stars. Lower-mass cluster members, detected by X-ray and H$\\alpha$ emission, and by near-IR excess, number more than a thousand. Myr-old star formation is traced by the optically-visible {\\sc Hii} region and cluster; observations of infrared and submillimetre-wave emission, and of optical emission features, indicate ongoing star formation in several locations across the Lagoon. The most prominent of these are the Hourglass Nebula and M8\\,E. Submillimetre-wave observations also reveal many clumps of dense molecular gas, which may form the next generation of stars. The complex structure of the region has been shaped by the interaction of the underlying molecular gas with multiple massive stars and episodes of star formation. NGC\\,6530 is the oldest component, with the newest stars found embedded in the molecular gas behind the cluster and at its southern rim. A degree to the east of the Lagoon, Simeis~188 is a complex of emission and reflection nebulae, including the bright-rimmed cloud NGC\\,6559; the presence of H$\\alpha$ emission stars suggests ongoing star formation. ",
        "introduction": "\\begin{figure}[thb] \\centering \\includegraphics[draft=False,width=\\textwidth]{M8_M20_GeraldRheman_arXiv.eps} \\caption{A widefield colour image covering M\\,8, M\\,20 and Simeis 188: H$\\alpha$ emission is red and reflection nebulae are blue; north is up and east is to the left; field of view (FOV) $\\sim 2.5^\\circ\\times\\sim 2.2^\\circ$. Courtesy Gerald Rhemann.} \\label{fig-wide} \\end{figure}   The Lagoon Nebula --- M\\,8 --- is the most prominent of a number of star-forming regions and supernova remnants in the section of the Sagittarius-Carina Arm lying near our line of sight towards the Galactic centre (Figs.~\\ref{fig-wide} \\& \\ref{fig-sgr}). Other members include: The Trifid Nebula (M\\,20); the supernova remnant W\\,28, near the Trifid at a kinematic distance of 1.9~kpc \\citep{velazquez}; the nearby optically-invisible {\\sc Hii} region W28\\,A2 (G\\,5.89--0.39), lying between the Lagoon and Trifid Nebulae at a distance of 2.0~kpc \\citep{acord}; and the complex of nebulae to the east of the Lagoon known as Simeis~188. The Eagle Nebula (M\\,16, NGC\\,6611) and M\\,17 lie about $10^\\circ$ further along the Sagittarius-Carina arm.  \\begin{figure}[thb] \\centering \\includegraphics[draft=False,width=\\textwidth]{figure1_arXiv.eps} \\caption{A $2.5^\\circ\\times 2.5^\\circ$ field in Sagittarius at 4 wavelengths: {\\it Upper Left:} Digitised Sky Survey (DSS, blue), dominated by stellar emission; {\\it Upper Right:} DSS2 (red), showing H$\\alpha$ emission as well; {\\it Lower Left:} IRAS 12~$\\mu$m; {\\it Lower Right:} IRAS 100~$\\mu$m. Arrows denote the directions of Galactic longitude and latitude (each arrow is 0.5$^\\circ$ long). The Lagoon Nebula and nearby regions are annotated.} \\label{fig-sgr} \\end{figure}   M\\,8 consists of a rich open cluster with several O-type stars and a prominent {\\sc Hii} region (about half a degree in diameter), the core of the cluster superimposed on the eastern half of the {\\sc Hii} region. The {\\sc Hii} region is surrounded by bright rims and at least one dark `elephant trunk' structure; these are most prominent at the southeastern edge of the {\\sc Hii} region. A dark lane splits the optical nebula from NE to SW (the `Great Rift'); the lack of background stars in the Rift implies that it is an obscuring dust lane rather than a lack of material, but it does not show up clearly in submillimetre- or millimetre-wave maps, either of spectral lines of CO or of the dust continuum; the Rift presumably has a high enough column density to obscure optical wavelengths significantly, but not enough to be obvious in emission.  The open cluster is fairly young (a few Myr). It is centred on 18:04:24, --24:21:12 (J2000.0), with a radius of around 30\\arcmin, but a core radius of only $\\sim$4\\arcmin\\ \\citep{chen}. It contains several O stars and about 60 B stars: One of its probable members, the O4 star \\citep[and probable binary,][]{rauw05} 9\\,Sagittarii, is the chief source of ionising flux for the {\\sc Hii} region. The western half of the {\\sc Hii} region is concentrated into a bright core which contains the Hourglass Nebula, a distinctively-shaped window into a compact {\\sc Hii} region with much denser ionised gas than the main body of the Lagoon Nebula, powered by the young O7 star Herschel~36 \\citep[HD~164740,][]{woodward}.  Southeast of the cluster core, a structure of bright rims and dark lanes stretches west to east. CO and dust maps clearly show this to be a dense molecular cloud; at its eastern end, at least two massive stars are being formed in the optically-invisible M8\\,E cluster which rivals the brightness of the Hourglass at infrared and submillimetre wavelengths (see Fig.~\\ref{fig-irac}).  The nomenclature of nebulae and clusters in the area is quite unclear, since the complex has been observed and catalogued since the 17th century \\citep{burnham}. According to the NGC/IC Project\\footnote{http://ngcic.org}, NGC\\,6533 refers to the whole nebula and NGC\\,6523 is the bright core of the nebula, lying NW of the Great Rift; the SE part of the nebula comprises NGC\\,6526 in the south and NGC\\,6530 in the north; and IC\\,1271 and 4678 are small condensations to the east of the main nebula (IC\\,1271 may refer to the O star HD\\,165052).  Although Messier referred to it as {\\it `amas'}, a cluster \\citep[M\\,8,][]{messier}, we will use M\\,8 to refer to the whole complex of stars, {\\sc Hii} regions and molecular gas, and `Lagoon Nebula' to be synonymous. The NE of the region is dominated by the open cluster, so NGC\\,6530 is now always used to refer to the cluster rather than any surrounding nebulosity. We take NGC\\,6533 to refer to the {\\sc Hii} region only, comprising NGC\\,6523 and 6526. {\\sc Hii} region studies have generally concentrated on the brighter eastern core (i.e.~NGC\\,6523), and so it is this designation that is found in the literature: For most practical purposes, NGC\\,6533 and 6523 are the same. It is not clear that IC\\,1271 and 4678 refer to real structures, and we will not use these designations.   \\begin{figure} \\includegraphics[width=\\textwidth]{m8-irac-4col-0.3-crop_arXiv.eps}\\\\[1ex] \\includegraphics[width=\\textwidth]{M8pic_arXiv.eps}  \\caption{{\\em Upper:} Four-colour infrared mosaic of the Lagoon Nebula from {\\it Spitzer} IRAC data. Red, orange, green and blue correspond to the 4 IRAC wavelengths (8.0, 5.8, 4.5 \\& 3.6~$\\mu$m respectively). North is up and East to the left; FOV is $36^\\prime\\times 21^\\prime$. Based on archive data from {\\it Spitzer} programme 20726, P.I.~J.~Hester. {\\em Lower:} The same field at optical wavelengths, from the Digitised Sky Survey: Red, green and blue correspond to $I$ and $R$-bands (from DSS2) and $B$-band (from DSS) respectively. H$\\alpha$ emission (in the $R$-band) appears greenish. M8\\,E is only visible in the optical $I$-band, but is saturated in the infrared. `The Dragon' (Sect.~\\ref{sec-over-ism}) is the prominent `elephant trunk' to the SE of the core of NGC\\,6530 (see also Fig.~\\ref{fig-hh}).} \\label{fig-irac} \\end{figure}   Much further east of the Lagoon lies an `R association': a scattering of smaller reflection and emission nebulae, often known as Simeis~188. Just to the north of the Lagoon, there is an area of diffuse nebulosity (Fig.~\\ref{fig-sgr}); on purely morphological grounds, this looks like a diffuse northern extension of NGC\\,6533, separated by a dust lane.  This review is structured as follows: An overview of M\\,8 as a whole (Sect.~\\ref{sec-over}), broken down into the main {\\sc Hii} region NGC\\,6523/6533 (Sect.~\\ref{sec-over-hii}), the young stellar cluster NGC\\,6530 (Sect.~\\ref{sec-overview-cluster}), its pre-main-sequence population (PMS, Sect.~\\ref{sec-over-ysos}) and the interstellar medium (Sect.~\\ref{sec-over-ism}), followed by an overview of the distance estimates to M\\,8 (Sect.~\\ref{sec-distance-nebula}). Then, major components of the region are discussed in more detail: NGC\\,6530 (Sect.~\\ref{sec-ngc6530}), including the PMS stars (Sect.~\\ref{sec-ngc6530-pms}) and reviews of age and distance estimates (Secs.~\\ref{sec-age} \\& \\ref{sec-distance-cluster}); the Hourglass Nebula (Sect.~\\ref{sec-hg}); M8\\,E (Sect.~\\ref{sec-m8e}); a few other candidate star-forming regions (Sect.~\\ref{sec-sfrs}); and Simeis\\,188 (Sect.~\\ref{sec-sim188}). Finally, we briefly discuss the structure and evolution of the region as a whole (Sect.~\\ref{sec-discussion}). ",
        "conclusions": "\\begin{figure} \\plotone{scuba-acis_arXiv.eps} \\caption{{\\it Chandra} X-ray sources overlaid on the sub-mm continuum structure of the Lagoon Nebula, from \\citet{gagne}.} \\label{fig-acis} \\end{figure}  \\begin{figure} \\plotone{scuba-acis-correl_arXiv.eps} \\caption{Spatial correlation between 850~\\micron\\ emission and {\\it Chandra} X-ray sources, 2MASS sources, and a random distribution, from \\citet{gagne}.} \\label{fig-acis-correl} \\end{figure}  \\subsubsection{The Structure of Star Formation}  Combining the {\\it Chandra} images towards the Hourglass \\citep{gagne} and NGC\\,6530 \\citep{dam} yields a catalogue of 1482 X-ray sources spanning most of the region observed at 450 and 850~\\micron\\ by \\citet{thesispaper}. Figure~\\ref{fig-acis} suggests that X-ray sources tend to cluster near sub-mm emission cores, which is confirmed by Figure~\\ref{fig-acis-correl}: This shows the spatial correlation of the 850~\\micron\\ flux to the location of each of the 1119 X-ray sources lying within the SCUBA map, and the same statistic for the 9805 2MASS sources in the map and for a set of $10^5$ random sources. The distributions show that the X-ray source locations are non-random and correlated with 850~\\micron\\ flux, as are about 20\\% of the 2MASS sources. In particular, Fig.~\\ref{fig-acis} shows strong clustering around (but slightly offset from) M8\\,E (near 18:04:54, --24:26:30), the central ridge (near 18:04:20, --24:23), and Herschel~36 and the Hourglass (near 18:03:40, --24:23). In the central ridge (M8\\,EC1--5), three lines of X-ray sources are seen separated by $\\sim 1\\arcmin$ from northeast to southwest. A number of smaller X-ray clusters appear close to 850~\\micron\\ cores along the southern rim, specifically: M8\\,SC8, SC1, SE1, SE3, and SE7. \\citet{barba} also found H-H objects associated with M8\\,SC8, SE3 and C3.  There are other bright rims and globules in and around the Lagoon, most of which seem to be due to the action of 9\\,Sgr, such as the bright rims of \\citet{sugitani}, the globule near the Hourglass \\citep{arias05} and the elephant trunk found near M8\\,E \\citep{brand}. These structures, along with the widespread T~Tauri star population \\citep{pris,arias07}, suggest pervasive star formation in the Lagoon, in addition to the young clusters found around the Hourglass and M8\\,E.  \\subsubsection{Star Formation History}  While most authors ascribe an age of a few Myr to NGC 6530, the oldest element of the Lagoon Nebula, \\citet{vda} argue for a much older cluster (by an order of magnitude), based on the probable cluster members found in the giant branch of the H-R diagram. These members seem to be worth investigation: If star formation has really been going on for a few tens of Myr, and is still ongoing at a significant rate, M\\,8 would be a very long-lived star-forming region. \\citeauthor{vda} also suggest, because of the lack of massive stars on the ZAMS, that massive star formation has essentially ceased. This may be true of NGC\\,6530, but massive star formation in M\\,8 as a whole is not over.  \\citet{lightfoot} identified a possible sequence of triggered star formation in M\\,8: NGC\\,6530 is the oldest feature, and some of its members are still ionising the main {\\sc Hii} region, NGC~6523, while younger features are found around the edge of the region, most obviously in the Hourglass and M8\\,E. The O stars in the cluster are also found at the periphery, while the core contains only B stars and later, and it is not clear why this should be so: Do the peripheral O stars represent a later generation of star formation? If so, it is odd that there are no signs of any O stars corresponding to the $\\sim$60 B stars in the core, either as main-sequence O stars or as post-MS objects. However, the basic picture of star formation proceeding outwards from the core of NGC\\,6530 is well-supported by evidence of age gradients \\citep[e.g.~][]{dam,arias07} as well as the prevalence of massive young stars, X-ray sources and H-H objects around the dense molecular cores around the edge of the Lagoon. There are also signs of star formation elsewhere in M\\,8, e.g.~the sample of candidate Class~I sources concentrated to the northeast of the Hourglass \\citep{dam06}.  \\subsubsection{The Structure of the Lagoon Nebula}  \\begin{figure} \\plotone{schemdisc_arXiv.eps} \\caption{Schematic diagram showing a possible structure of the Lagoon Nebula, in the form of two cuts through the Central Ridge and Hourglass regions (shown on a $\\sim 20^\\prime\\times\\sim 10^\\prime$ section of the 850~\\micron\\ emission map). From \\citet{thesis}.} \\label{fig-xsection} \\end{figure}  Compared to the clumps along the southern and southeastern rims of the Lagoon Nebula, the EC clumps seem to fall off fairly shallowly on all sides. This suggests that the central ridge may lie behind the {\\sc Hii} region, and we see the structure face-on, rather than the edge-on view of the southern clumps (see Fig.~\\ref{fig-xsection}). \\citet{woodward} argue convincingly that the Hourglass is embedded in the molecular cloud behind the Lagoon, whereas M8\\,E shows a steep fall-off into the {\\sc Hii} region. So we see ongoing star formation both behind and to the southeast of the ionised gas, suggesting that the ionisation front is moving away from us and to the south and east, compressing and warming the molecular gas (giving rise to submillimetre and CO emission) and presumably triggering star formation, with X-ray emitting YSOs appearing tightly clustered in the wake of the front.  There are indications of a thin screen of material between us and M\\,8, somewhat blueshifted. This might be the last remnants of the front of the molecular cloud, excavated by a blister {\\sc Hii} region on the front side of the cloud, and accelerated towards us by the ionised gas."
    },
    "0809/0809.3290_arXiv.txt": {
        "abstract": "We present an analysis of infrared (IR) echelle spectra of five stars in the TW Hydrae Association (TWA). We model the Zeeman broadening in four magnetic-sensitive \\ion{Ti}{1} lines near $2.2\\ \\mu$m and measure the value of the photospheric magnetic field averaged over the surface of each star. To ensure that other broadening mechanisms are properly taken into account, we also inspect several magnetically insensitive CO lines near $2.3\\ \\mu$m and find no excess broadening above that produced by stellar rotation and instrumental broadening, providing confidence in the magnetic interpretation of the width of the \\ion{Ti}{1} lines. We then utilize our results to test the relationship between stellar magnetic flux and X-ray properties and compare the measured fields with equipartition field values. Finally, we use our results and recent results on a large sample of stars in Taurus to discuss the potential evolution of magnetic field properties between the age of Taurus ($\\sim$2 Myrs) and the age of TWA ($\\sim$10 Myrs). We find that the average stellar field strength increases with age; however, the total unsigned magnetic flux decreases as the stars contract onto the main-sequence. ",
        "introduction": "T Tauri stars (TTSs) are young, low-mass stars, which are newly formed out of their natal molecular clouds and are often associated with OB associations and/or nebulosity. The location of TTSs above the main sequence in the H-R diagram and their lithium abundance also reveal their youth. TTSs display substantial photometric and spectroscopic variability, and their properties have been reviewed by \\citet{appenzeller1989}, \\citet{bertout1989}, \\citet{basribertout1993} and \\citet{menardbertout1999}. Depending on whether they are actively accreting, TTSs are divided into two categories. Classical T Tauri stars (CTTSs) are still surrounded by dusty circumstellar disks with detectable non-stellar emission in the near-infrared (near-IR). Accretion from the disks onto the stars produces many observable features such as strong yet variable H$\\alpha$ emission, often displaying P-Cygni or inverse P-Cygni profile shapes. Naked T Tauri stars (NTTSs), on the other hand, show little or no sign of accretion and do not appear to have close dusty disks around them. It is generally believed that stars evolve through a CTTS phase to become a NTTS on the way to the main sequence.  One essential piece of the puzzle to our understanding of TTSs is the stellar magnetic field. For NTTSs, it is believed that dynamo generated magnetic fields are responsible for most of their peculiarities, e.g., variable X-ray emission, flares, and cool spots. The evolution of CTTSs, especially the interaction between the stars and their circumstellar disks, is believed to be directly controlled by stellar magnetic fields \\citep[see review by][]{bouvier2007}. Various magnetospheric accretion models \\citep{camenzind1990, konigl1991, ccc1993, shu1994, paatz1996} generally agree that the disks of CTTSs are truncated by stellar magnetic fields near the corotation radius and the accreting material flows along the field lines toward the polar regions of the central star. The disks of CTTSs are the birthplace for planets and the radius at which the magnetic field truncates the inner disk may determine where giant planets are halted on their inward migration in the disk \\citep{lin1996, eisner2005}. Magnetic fields probably also play a key role in driving the stellar winds/jets seen in young stars \\citep[e.g.,][]{shu1999}. Therefore, a full understanding of the early evolution of young stars, which will lead to better understanding of both star and planet formation, requires detailed theoretical consideration and empirical measurements of the stellar magnetic field properties.   Direct detection of stellar magnetic fields is difficult, with measurements starting to become plentiful only recently \\citep{cmj2007}. One technique, which led to the first direct detection of magnetic fields on a TTS, Tap 35 \\citep{basri1992} and later on other stars \\citep{guenther1999}, is based on an equivalent width increase in moderately saturated magnetically sensitive lines due to Zeeman splitting. Unfortunately, this method is particularly sensitive to the choice of effective temperature and gravity used when analyzing the stellar line measurements. Another method of detecting stellar magnetic fields is through spectropolarimetry, which looks for net circular polarization in Zeeman-sensitive spectral lines. Observed along the direction that is parallel to the magnetic field, the Zeeman $\\sigma$-components are circularly polarized, with the components of opposite helicity split to either side of the nominal line wavelength. If the magnetic topology on the stars is somewhat organized, there should be a measurable component along the line of sight. Reliable detections of circular polarization in TTS photospheric absoprtion lines are rare, with most recent studies finding only upper limits of $100-200$ G \\citep{cmj1999a, daou2006, smirnov2004}. More recently, significant, but similarly weak, detections have been made \\citep{yang2007, donati2007} and one surprising strong detection on BP Tau has recently been published \\citep{donati2008}. On the other hand, net polarization is detected in the \\ion{He}{1} $5876$ \\AA\\ emission line, first discovered by \\citet{cmj1999a} on BP Tau. This \\ion{He}{1} line is believed to form in the postshock region \\citep{hartmann1994, edwards1994} where material accretes onto the star and the detection on BP Tau yielded a net longitudinal field of $2.46 \\pm 0.12$ kG in the line formation region. Circular polarization in the narrow component of the \\ion{He}{1} line has now been observed in several CTTSs \\citep{valenti2004, symington2005, yang2007, donati2007, donati2008}. These observations suggest that accretion onto CTTSs is indeed controlled by a strong stellar magnetic field.   While spectropolarimetry can provide clear detections of stellar magnetic fields, it is only sensitive to the net field resulting from the predominance of one field polarity over the other. As a result, it can miss the majority of the magnetic flux present on a star. The most successful approach so far for measuring the total magnetic flux has been to measure the Zeeman broadening of spectral lines in unpolarized light \\citep{robinson1980, saar1988, valenti1995, cmj1996, cmj1999b}. For any given Zeeman component, the splitting resulting from the magnetic field is $$\\Delta\\lambda = {e \\over 4\\pi m_ec^2} \\lambda^2 g B = 4.67 \\times 10^{-7} \\lambda^2 g B \\,\\,\\,\\,\\,\\, \\rm m\\mbox{\\AA}\\,\\eqno(1)$$ where $g$ is the Land\\'e $g$-factor of the transition, $B$ is the strength of the magnetic field in kG, and $\\lambda$ is the wavelength of the transition in \\AA. \\citet[][hereafter Paper I]{cmj1999b} detected Zeeman broadening of the \\ion{Ti}{1} line at $2.2233\\ \\mu$m on the CTTS BP Tau and obtained a field strength of $\\bar B = 2.6 \\pm 0.3$ kG, averaged over the entire surface of the star. From analysis of four \\ion{Ti}{1} lines near $2.2\\ \\mu$m, \\citet[][hereafter Paper II]{cmj2004} also detected a strong magnetic field of $2.5 \\pm 0.2$ kG on the NTTS Hubble 4. In addition, these authors examined several CO lines near $2.3\\ \\mu$m, which are magnetically insensitive, and found no excess broadening above that due to rotational and instrumental effects. This technique has now been used to measure the fields of 16 TTSs \\citep{cmj2004,yang2005,cmj2007}.     In the largest study to date of fields on TTSs, \\citet{cmj2007} used a sample of 14 stars to look for correlations between the measured magnetic field properties and stellar properties that might be important for dynamo action.  No significant correlations were found and \\citet{cmj2007} speculated that the fields seen in these young stars may be entrained interstellar fields left over from the star formation process \\citep[e.g.,][]{tayler1987, moss2003}. To supply more observational constraints and further investigate the magnetic properties of TTSs, we present here Zeeman broadening measurements of five NTTSs in the TW Hydrae association (TWA): TWA 5A, TWA 7, TWA 8A, TWA 9A and TWA 9B. At $\\sim 50$ pc \\citep{wichmann1998}, TWA is the nearest region of recent star formation. \\citet{kastner1997} first concluded that a few T Tauri stars in the apparent neighborhood of TW Hya are indeed physically close to each other, based on similar X-ray emission and lithium abundance, and proposed the existence of the TW Hydrae Association. Over 20 stars have been proposed as members of TWA (though only a few of all proposed TWA members have trigonometric parallax confirmation), and a compilation of TWA members can be found in \\cite{zuckerman2004ar}. The age for TWA is estimated at about $10$ Myrs \\citep[see discussion in][]{navascues2006}, and it is intriguing that these pre-main-sequence stars are not associated with any molecular cloud. The origin of the TWA is still under investigation \\citep[e.g., see][]{makarov2007}.  As described earlier, the magnetic field properties of TW Hya (TWA 1) have been investisgated by \\citet[][hereafter, Paper III]{yang2005} and \\citet{yang2007}. Interstingly, \\citet{setiawan2008} report a massive planet very close to TW Hya, and \\citet{huelamo2008} argue that the observed radial velocity variations of TW Hya are best explained by a cool spot instead of a hot jupiter. Here, we study the magnetic fields of 5 additional members of TWA, utilizing infrared (IR) spectroscopy to measure Zeeman broadening in 4 \\ion{Ti}{1} lines in the K band. The observations and data reduction are presented in \\S\\ 2. In \\S\\ 3 we describe our data analysis technique and results, and it is followed by a discussion of the results in \\S\\ 4. ",
        "conclusions": "By fitting the Zeeman broadening in the magnetically sensitive \\ion{Ti}{1} lines, we measure the average surface magnetic field of five TWA stars.  Three different models were used to fit the line profiles and measure the magnetic field strength with generally good agreement in the results between the three.  Model M1 assumes only 1 magnetic component to the stellar atmosphere.  This model provides a good match to TWA 5A in the $2.2228 \\mu$m region (but not the $2.2300 \\mu$m region). This is likely due to the fact that TWA 5A apparently has a relatively large $v$sin$i$ so that rotational broadening, in addition to the magnetic broadening, strongly affects the line profile shapes. Because our magnetic models for TWA 5A do not reproduce line profiles in the $2.2300 \\mu$m region, our derived magnetic properties for TWA 5A may have large systematic errors. Generally, mean field strength increases as more components are added to the model, which suggests that the observed profiles have wings that are too broad to be fit by a single component. For the other stars, the lowest reduced $\\chi^2$ typically occurs for model M2, though M3 is favored in one case.  However, there is generally quite good agreement bewteen the mean field values recovered using these two models with the difference bewteen them averaging 15\\%.  For consistency with previous studies, we use the results from M3 when comparing to results found by \\citet{cmj2007}.  Theoretically, equipartition arguments for the Sun and cool stars suggest that the magnetic field strength should scale with the gas pressure in the surrounding non-magnetic photosphere which sets a limit for the maximum field strength allowed on the star \\citep[e.g.][]{spruit1979,safier1999}. \\citet{rajaguru2002} found that in the regime where convectively stable flux tubes exist, equipartition pressure allows a maximum magnetic field strength of $\\sim 1300$ G for values of \\Teff\\ and \\logg\\ appropriate for most of our TWA stars. We calculate equipartition field strengths for the TWA stars by balancing magnetic pressure, $B_{eq}^2/(8\\pi)$, with the gas pressure in surrounding unmagnetized photosphere, $P_{g}$. We use the gas pressure at the level in the atmosphere where the local temperature is equal to the effective temperature in the model atmospheres used for our stars. This should be an upper limit for gas pressure because this level is approximately where the continuum forms, while the \\ion{Ti}{1} lines form over a wide range of depths above this level in the atmosphere. The equipartition field strengths, $B_{eq}={(8\\pi P_{g})}^{1/2}$, of the TWA stars are listed in the last column of Table \\ref{results}. These values are in good agreement with the estimate from \\citet{rajaguru2002}. Compared with the measured mean field values, they are all smaller by a factor of at least 1.38 and more typically $\\sim 2.0$ (for model M2), suggesting that magnetic pressure dominates over gas pressure in the photospheres of the TWA stars observed here. This in turn implies complete coverage of magnetic fields over surface of these stars. Magnetic regions with excess pressure would expand, while field free regions contract into the lower pressure nonmagnetic regions until equipartition is established at a lower field strength or until the entire star is covered with magnetic fields. As discussed in Papers I and II, our models do not necessarily require field-free regions. The results are essentially the same with or without a 0 kG field component. As an example, we fit the observed spectra of TWA 8A with a model which allows on the stellar surface no field-free region and four magnetic regions with specific field strengths of 1 kG, 3 kG, 5 kG and 7 kG. The measured magnetic field value using this model is 3.4 kG with a reduced $\\chi^2$ of 1.22, while model M3 yields 3.3 kG with a reduced $\\chi^2$ of 1.25 (see Table \\ref{results}). With a field-free region included in our current analysis, all the TWA stars have at least $76\\%$, in most cases $80\\% - 86\\%$, of their surfaces covered with magnetic fields. The surfaces of TTSs are quite possibly fully covered by strong magnetic fields.   \\citet{pevtsov2003} examine the X-ray luminosities and the total unsigned magnetic flux of active stars as well as regions of different activity level on the Sun. They find a nearly linear correlation over many decades of variation in the two quantities. We calculate the magnetic flux of the four TWA stars (excluding TWA 5A since it is an unresolved multiple system), using our magnetic field measurements and the stellar radii estimated from bolometric luminosities given in \\citet{webb1999} and the effective temperatures of our stars. We utilize the X-ray luminosity-magnetic flux relationship of \\citet{pevtsov2003} and calculate the predicted X-ray luminosity. For comparison, we use the X-ray to bolometric luminosity ratio, $L_x/L_{bol}$, reported by \\citet{webb1999} to get the value of $L_x$ for each star. The predicted and measured values of $L_x$ for the three TWA stars (TWA 7, TWA 8A and TWA 9A) presented in this paper and that of TW Hydrae taken from Paper III are plotted in Figure \\ref{xray}. A measured value for TWA 9B is not available in the literature. As shown in Figure. \\ref{xray}, the observed $L_x$ values of the four TWA stars are smaller than the predicted values by a factor of $2.7 \\pm 0.7$ on average. This is similar to the findings of \\citet{cmj2007} in a study of 14 CTTSs in Taurus and Auriga. The stars from \\citet{cmj2007} are also plotted in Figure \\ref{xray}, and the difference between the observed and the predicted $L_x$ values is a bit more obvious. The average ratio of the predicted to observed $L_x$ values is $17.7 \\pm 3.4$ for the 14 Taurus and Auriga stars. At an age of about 10 Myrs, the TWA stars are substantially older than the stars in the Taurus-Auriga region ( $\\sim$2 Myrs). As a result, the TWA stars are substantially closer to the main sequence and this may be the reason why their X-ray luminosities and magnetic flux values are closer to the relationship found by \\citet{pevtsov2003} which was based primarily on main sequence stars. The X-ray emission on active stars such as TTSs is thought to mainly result from energy dissipation in magnetic fields during flares and small-scale flare-like activities \\citep{guedel2003,arzner2007}. Magnetic fields are stressed due to convective gas motions in the lower photosphere and below \\citep{foukal2004}, producing flares of various scales. For PMS stars that are fully covered by magnetic fields, when the magnetic field strength is larger than the equipartition value, it may be more difficult for the gas, at least in the photosphere, to push around the fields and build up magnetic stress. This would be the case for both the Taurus stars and the TWA stars. The ratio of the observed to equipartion magnetic field, $B_{obs}/B_{eq}$, for the Taurus stars on average is greater than that for the TWA stars, so the X-ray production might be more prohibited for the Taurus stars resulting in the observed $L_x$ being further away from the line of equality in Figure \\ref{xray}. It will be interesting to see where TTSs in other age groups fit in this picture relating magnetic flux and X-ray emission.   The origin of magnetic fields on TTSs has sparked various theoretical speculations, but it is still far from clear how the strong fields on these stars are produced. Due to their youth, most TTSs are generally thought to be fully convective, though model calculations have shown that TTSs might have small radiative cores \\citep{granzer2000}. As a result, a solar-like interface dynamo \\citep[e.g.,][]{durney1978,parker1993} is unlikely to be responsible for the fields on TTSs. One potential alternative mechanism is a distributed dynamo amplified by convective motions \\citep[e.g.,][]{durney1993, dobler2006, chabrier2006}.  Such turbulent dynamos are not well understood at this point, but exisiting models seem to only produce field strengths equal to or lower than the pressure equipartition values, and many such models also produce strong fields which occupy only a small filling factor at the stellar surface \\citep[e.g.,][]{cattaneo1999, bercik2005}.  Both of these predicitons appear not to be borne out by the observations.  In addition to a possible dynamo origin, the magnetic fields of TTSs could also result from the primordial flux that survives the star formation process \\citep[e.g.,][]{tayler1987, moss2003}. With the measurements from this paper and those in \\citet{yang2005} and \\citet{cmj2007}, we explore this issue by comparing the average magnetic field strength and magnetic flux of the stars in TWA and in Taurus-Auriga, two regions of different ages. The average magnetic field of the Taurus-Auriga stars is $2.20 \\pm 0.16$ kG, which is smaller than that of the five TWA stars (TW Hya, TWA 7, TWA 8A, TWA 9A and TWA 9B), where $\\bar B =  2.98 \\pm 0.22$ kG. While the average field strength increases slightly from the age of Taurus to that of TWA, we see a sharp decrease of the average magnetic flux, $4\\pi {R_*}^2B$, from $(7.00 \\pm 1.37) \\times 10^{18}$ Wb for the Taurus-Auriga stars compared to $(2.61 \\pm 0.69) \\times 10^{18}$ Wb for the five stars in TWA. This change in the magnetic flux is generally consistent with the decay of primordial magnetic fields on TTSs as suggested by \\citet{moss2003}. With only a limited number of samples on hand, the current evidence is far from sufficient. Further observations and measurements of young stars in other young clusters of different ages, e.g., the Orion Nebula Cluster and the $\\beta$ Pictoris Moving Group, could shed more light on the issue of field generation in young stars."
    },
    "0809/0809.1295_arXiv.txt": {
        "abstract": "Weakly Interacting Massive Particles (WIMPs) are one of the leading candidates for Dark Matter.  We developed a model--independent method for determining the WIMP mass by using data (i.e., measured recoil energies) of direct detection experiments.  Our method is independent of the as yet unknown WIMP density near the Earth, of the form of the WIMP velocity distribution, as well as of the WIMP--nucleus cross section.  It requires however positive signals from at least two detectors with different target nuclei.  At the first phase of this work we found a systematic deviation of the reconstructed WIMP mass from the real one for heavy WIMPs.  Now we improved this method so that this deviation can be strongly reduced for even very high WIMP mass.  The statistical error of the reconstructed mass has also been reduced.  In a background--free environment, a WIMP mass of $\\sim$ 50 GeV could in principle be determined with an error of $\\sim$ 35\\% with only 2 $\\times$ 50 events. ",
        "introduction": "By now there is strong evidence that more than 80\\% of all matter in the Universe is dark (i.e., interacts at most very weakly with electromagnetic radiation and ordinary matter). The dominant component of this cosmological Dark Matter must be due to some yet to be discovered, non--baryonic particles. Weakly Interacting Massive Particles (WIMPs) $\\chi$ with masses roughly between 10 GeV and a few TeV are one of the leading candidates for Dark Matter (for reviews, see Refs.~\\cite{SUSYDM}).  Currently, the most promising method to detect many different WIMP candidates is the direct detection of the recoil energy deposited in a low--background laboratory detector by elastic scattering of ambient WIMPs on the target nuclei \\cite{deta}.  The differential rate for elastic WIMP--nucleus scattering is given by \\cite{SUSYDM}: \\beq \\label{eqn:dRdQ} \\dRdQ = \\calA \\FQ \\int_{v_{\\rm min}}^{\\infty} \\bfrac{f_1(v)}{v} dv\\, . \\eeq Here $R$ is the direct detection event rate, i.e., the number of events per unit time and unit mass of detector material, $Q$ is the energy deposited in the detector, $F(Q)$ is the elastic nuclear form factor, $f_1(v)$ is the one--dimensional velocity distribution function of the WIMPs impinging on the detector, $v$ is the absolute value of the WIMP velocity in the laboratory frame.  The constant coefficient $\\calA$ is defined as \\beq \\label{eqn:calA} \\calA \\equiv \\frac{\\rho_0 \\sigma_0}{2 \\mchi m_{\\rm r,N}^2}\\, , \\eeq where $\\rho_0$ is the WIMP density near the Earth and $\\sigma_0$ is the total cross section ignoring the form factor suppression.  The reduced mass $m_{\\rm r,N}$ is defined by \\beq \\label{eqn:mrN} m_{\\rm r,N} \\equiv \\frac{\\mchi \\mN}{\\mchi+\\mN}\\, , \\eeq where $\\mchi$ is the WIMP mass and $\\mN$ that of the target nucleus.  Finally, $\\vmin = \\alpha \\sqrt{Q}$ is the minimal incoming velocity of incident WIMPs that can deposit the energy $Q$ in the detector and \\beq \\label{eqn:alpha} \\alpha \\equiv \\sfrac{\\mN}{2 m_{\\rm r,N}^2}\\, . \\eeq  In our earlier work \\cite{DMDDf1v}, we developed methods for reconstructing the one--dimensional velocity distribution, $f_1(v)$, and for estimating its moments from the recoil spectrum as well as from measured recoil energies directly in direct detection experiments: \\beqn \\label{eqn:moments} \\expv{v^n} \\eqnequiv \\int_{v_{\\rm min}(\\Qthre)}^\\infty v^n f_1(v) \\~ dv \\non\\\\ \\=        \\alpha^n \\!\\! \\bfrac{2 \\Qthre^{(n+1)/2} \\rthre/\\FQthre+(n+1) I_n} {2 \\Qthre^{1    /2} \\rthre/\\FQthre+      I_0}\\!\\!. \\eeqn Here $\\rthre \\equiv (dR/dQ)_{Q = \\Qthre}$ is an estimated value of the scattering spectrum at the threshold energy, $\\Qthre$, and $I_n$ can be estimated by \\beq \\label{eqn:In_sum} I_n = \\sum_a \\frac{Q_a^{(n-1)/2}}{F^2(Q_a)}\\, , \\eeq where the sum runs over all events in the data set. Note that Eq.(\\ref{eqn:moments}) can be used without prior knowledge of the local WIMP density, $\\rho_0$, of the velocity distribution function of incident WIMPs, $f_1(v)$, as well as of the WIMP--nucleus cross section, $\\sigma_0$. ",
        "conclusions": "In this paper we described the basic ideas of our method for determining the WIMP mass and gave the main formulae.  The algorithmic process for correcting the systematic deviation of the reconstructed WIMP mass has also been discussed. More details and discussions about determining $\\mchi$ can be found in Ref.~\\cite{DMDDmchi}. \\begin{theacknowledgments} This work was partially supported by the Marie Curie Training Research Network ``UniverseNet'' under contract no.~MRTN-CT-2006-035863 as well as by the European Network of Theoretical Astroparticle Physics ENTApP ILIAS/N6 under contract no.~RII3-CT-2004-506222. \\end{theacknowledgments}"
    },
    "0809/0809.4962_arXiv.txt": {
        "abstract": "We report the detection of a point source CXO J172337.5-373442 in a {\\it Chandra} field with a high significance ($26.7\\sigma$), and the discovery ($4\\sigma$) of a $48''$ long X-ray tail emanating from the point source. The X-ray spectra of both the point source and the tail are well described with a single absorbed powerlaw, and the tail is harder (powerlaw index $\\Gamma = 0.14^{+0.59}_{-0.68}$) than the point source ($\\Gamma = 1.78^{+0.13}_{-0.11}$). From this first detailed spatial, spectral and timing X-ray analysis of CXO J172337.5-373442, and from a plausible optical counterpart found from the archives, we conclude that this source is either a Galactic High-Mass X-ray Binary with an X-ray jet or a Galactic pulsar with its ``pulsar wind nebula\" seen as the X-ray tail. Although, the currently available data are not enough to distinguish between these two candidates with certainty, a detailed comparison of their known properties with those of CXO J172337.5-373442 favours the latter type. If this identification is correct, then the pulsar should be middle-aged or old, that has escaped from its supernova remnant, and the X-ray tail should originate from the synchrotron emission from either of the following locations: (1) a shocked region, or (2) a jet emanating from the pulsar's magnetosphere. ",
        "introduction": "The high angular resolution $(0''.5)$ and the low background of the Advanced CCD Imaging Spectrometer (ACIS) of the {\\it Chandra} X-ray observatory is ideal to detect faint sources and to minutely study the structures of faint extended sources. ACIS has observed many serendipitous point sources, some of which are attached to relatively faint X-ray tails/jets/lobes. The detection of such faint extended features requires special care during the data analysis. A point source with X-ray tails/jets/lobes can be a protostar, an active galactic nucleus (AGN), an X-ray binary, or a pulsar wind nebula (PWN), and its properties cannot be studied meaningfully without a proper identification. In this Letter, we report the discovery of an X-ray tail emanating from a point source detected with ACIS. From the first detailed X-ray analysis of this source, we conclude that it is either a high-mass X-ray binary (HMXB) system, or a PWN, with the latter type favoured.  \\section[]{{\\it Chandra} Data Analysis and Results}  The field of the transient low-mass X-ray binary (LMXB) system XTE J1723--376 was observed with the {\\it Chandra} X-ray space mission on 2001 September 4 for 29.7 ks. This observation was made with the ACIS-I mode, ``Timed\" operating mode, and ``Faint\" telemetry format. In order to analyse these data, we have used the Chandra Interactive Analysis of Observations (CIAO) software (ver. 4.0; CALDB ver. 3.2.4) starting from the level 1 event file. We have applied the standard grade filtering, removed the high background times, and used the energy range $0.3-8.0$ keV and a good exposure time of $29.1$ ks for our analysis.  The transient LMXB XTE J1723--376 was in quiescent state, and was not detected with $2\\sigma$ threshold. We have used the CIAO tool ``celldetect\" in order to detect the relatively bright sources in this field. With a threshold of $7\\sigma$, ``celldetect\" has detected only one source. This source has been detected with a very high significance ($\\approx 26.7\\sigma$; see Table 1; also see later for discussion) at the coordinates (J2000): RA = 17 23 37.532 and Dec = -37 34 41.97. Moreover, smoothing of the image clearly shows that an X-ray tail is attached to this source (Fig. 1; see also Fig. 2). We have not found this X-ray source in the archives, except for the very recently published Brera Multi-scale Wavelet (BMW) {\\it Chandra} source catalogue (Romano et al. 2008). The source 1BMC172337.5-373442 in this catalogue is identical with our source. However, the BMW-{\\it Chandra} source catalogue neither gives any detailed properties of this source, nor mentions the X-ray tail of the source. The high significance of the source and the existence of its tail make it very interesting and worth exploring. Therefore, we study the properties of this source in this Letter. Since we have independently detected this source from the {\\it Chandra} field, we will call it CXO J172337.5-373442 using the {\\it Chandra} naming convention. CXO J172337.5-373442 is located on the ACIS-I3 chip. We have applied the exposure map correction, but found that the effects of non-uniform exposure and non-uniform CCD response in the source region is small. We have also found that the bright portion of this source (that is excluding the tail) is consistent with the point spread function (PSF) of a point source at the $\\approx 5'.3$ off-axis location of CXO J172337.5-373442. This strongly suggests that the bright part of CXO J172337.5-373442 is a point source. For our analysis, we have estimated the background from a rectangular region (area $A_{\\rm B} = 41054.1$ arcsec$^2$) in the ACIS-I3 chip (see Fig. 1). This region contains no source when the threshold is set to $3\\sigma$. We have detected 1037 ($= N_{\\rm B}$) counts from this region for the good exposure time, and for the energy range $0.3-8.0$ keV. In order to calculate the net count ($N_{\\rm N}$) from an area $A_{\\rm S}$, we have used the formula $N_{\\rm N} = N_{\\rm T} - (A_{\\rm S}/A_{\\rm B})N_{\\rm B} $, where $N_{\\rm T}$ is the total number of counts detected from that area. Then the error in $N_{\\rm N}$ and the signal-to-noise ratio are given by $\\Delta N_{\\rm N} = \\sqrt{N_{\\rm N} + (1 + A_{\\rm S}/A_{\\rm B})(A_{\\rm S}/A_{\\rm B})N_{\\rm B}}$ and $S/N = N_{\\rm N}/\\Delta N_{\\rm N}$ respectively.  We have used the photons detected in a circle of $\\approx 6''.8$ radius for the spectral analysis of the point source portion of CXO J172337.5-373442. This circle should contain about 98\\% of the source counts, as we have found from the instrument PSF. We have detected $714.4\\pm26.8$ net counts from the $144.3$ arcsec$^2$ area of the circle, which implies a signal-to-noise ratio $26.7$ (Table 1). We note that the pileup is expected to be very small as the counts per frame is low ($\\approx 0.08$). We have used an absorbed powerlaw model ({\\tt wabs*powerlaw} of XSPEC) to fit the corresponding energy spectrum. From our fitting (using $\\chi^2$-statistic), we have found that the best fit values of the neutral hydrogen column density $N_{\\rm H} = 0.37^{+0.10}_{-0.08}\\times10^{22}$ cm$^{-2}$ and the powerlaw photon index $\\Gamma = 1.78^{+0.13}_{-0.11}$ (see Table 1). As the corresponding reduced $\\chi^2 = 26.1/39$, the model of the observed spectrum does not require any other component (see Fig. 3). Here we note that an absorbed blackbody model ({\\tt wabs*bbodyrad} of XSPEC) provides a much worse fit (reduced $\\chi^2 = 56.6/39$).  Fig. 1 shows that an X-ray tail starts from the point source portion of CXO J172337.5-373442, and extends roughly towards the north/northwest. The initial part of this faint tail is probably hidden within the PSF of the point source. We have used the photons detected in a polygon (see Fig. 1) of area $436.8$ arcsec$^2$ for our analysis. Net counts of $24.0\\pm5.9$ have been detected from this area, which means that the tail has been detected with a significance of $4.0\\sigma$. Fig. 1 indicates that the X-ray tail bends towards northeast near its end. The length of the tail is $\\approx 48''$. The tail intensity has a minimum at a location $\\approx 21''$ from the point source, a maximum $\\approx 10''$ away from the minimum, and after that the tail fades. The average width of the tail is $\\approx 12''$. The PSF of the point source portion of CXO J172337.5-373442 cannot explain such size and structure, which strongly suggests that the tail is real. In order to show that the brightness variation and the kink of the tail may be real (and not a result of the smoothing; see Fig. 1), we have displayed the unbinned and unsmoothed {\\it Chandra} image of the CXO J172337.5--373442 region in Fig. 2.  We have fitted the tail spectrum with an absorbed powerlaw model ({\\tt wabs*powerlaw} of XSPEC). As the number of counts is small, we have used C-statistic, and fixed the $N_{\\rm H}$ to a value of $\\approx 0.37\\times10^{22}$ cm$^{-2}$ (which is the best fit value for the point source). The spectral analysis suggests that the X-ray tail is harder than the point source (Table 1). In order to check it in a model independent way, we have compared the hardness ratios ($h/s$) of the point source and the tail. Here $h$ is the number of net counts in the $2-8$ keV energy range and $s$ is that in the $0.3-2$ keV range. We have found that $h/s = 0.749\\pm0.057$ and $1.788\\pm0.929$ respectively for the point source and the tail. This, like the results of the spectral analysis, indicates that the tail is spectrally harder. From the spectral fitting we have also found that the tail is about 14 times fainter than the point source (see Table 1).  We have done the timing analysis of the point source portion of CXO J172337.5-373442. With a frequency bin of 0.034 milli-Hertz, we have found a marginal feature at the frequency $0.0631$ Hz (period $\\approx 15.9$ s) with the peak power of 18.95 and the quality factor of 1874. This period is much smaller than the dither periods (1000 s in Y and 707 s in Z directions), and is not an integer multiple of the lowest time bin size (3.24104 s) of the data. Therefore, this timing feature cannot be a result of dithering or the minimum time bin size. The single trial significance of the feature is $3.95\\sigma$. However, it is not significant if we consider the number of trials. Nevertheless, we mention it in this Letter, because a known frequency may lead to a significant detection from the future observational data.  \\section[]{Multiwavelength view of CXO J172337.5--373442}  An archival search did not yield any radio, ultraviolet or $\\gamma$-ray counterpart of CXO J172337.5--373442. However, we have found a pointlike optical and infrared (IR) source about $0''.15$ away (which we will call ``source A\"). This is the only optical and IR source in the {\\it Chandra} ACIS error circle of the point source portion of CXO J172337.5--373442. The observed $B$, $V$, $R$, $J$, $H$ and $K$ magnitudes of source A are 16.79, 15.63, 15.31, 12.58, 11.86 and 11.60 respectively.  We will now try to understand the nature of source A from its observed optical colours. First we note that the $B-V$ values of almost all the quasars (for a sample of 788 quasars; see the Fig. 1.6 of Peterson 1997) are less than 1, while $B-V = 1.16$ for source A. Moreover, CXO J172337.5--373442, which is likely a Galactic source (see \\S~4), is the only possible X-ray counterpart of source A. Therefore source A is not likely a quasar. Hence it is likely a Galactic source, as it is a point source.  In order to check if source A is a star, we have calculated its $E(R-V)/E(B-V)$ ratio assuming that it is a star (considering various stellar spectral types one by one). In order to calculate this ratio we have used the observed $B-V$ and $R-V$ colours of source A, and the known intrinsic colours $(B-V)_0$ and $(R-V)_0$ of various spectral types of stars. The standard value of $E(R-V)/E(B-V)$ is $-0.78$ (Table 3.21 of Binney \\& Merrifield 1998). We have not found a value for any stellar spectral type that is close to this standard value. The closest value ($\\approx -0.33$) is for O9 supergiants, and the discrepancy is much more for low mass stars. Therefore, source A is plausibly not a star, and at least is very unlikely to be a low mass star.  Since CXO J172337.5--373442 may be a pulsar (see \\S~4), we will now examine if the optical colour of source A is consistent with that of a pulsar. Unlike the stars, the intrinsic optical colours of the pulsars are poorly explored. Therefore we will compare source A with two well known pulsars: Crab and Vela. The intrinsic colours $(B-V)_0$ of Crab and Vela are 0.08 and 0.185 respectively (Sollerman et al. 2000; Golden et al. 2000; Mignani \\& Caraveo 2001; Shibanov et al. 2003). If we assume the same intrinsic colours for source A, then we get $E(B-V) = 1.08$ and $0.975$, implying the neutral atomic hydrogen column density $N_{\\rm H} = 0.52\\times10^{22}$ and $0.47\\times10^{22}$ respectively (inferred from $N_{\\rm H} = 0.48\\times10^{22} E(B-V)$ cm$^{-2}$; Table 3 of Bohlin et al. 1978). These are much less than the Galactic neutral hydrogen column density $1.47\\times10^{22}$ cm$^{-2}$ in this direction (Dickey \\& Lockman 1990; NASA's HEASARC $N_{\\rm H}$ tool), implying that the optical colour of source A is consistent with that of Galactic pulsars.  From the observed optical and IR source density in the CXO J172337.5--373442 region, we find that the probability of the existence of an unrelated optical and IR source within the {\\it Chandra} error circle by chance is about 0.01. This, and the obervational indication that source A is a Galactic source, suggest that this source is the optical and IR counterpart of CXO J172337.5--373442. Assuming this to be true, we have plotted a multiwavelength spectrum of CXO J172337.5--373442 (Fig. 4). However, we note that detailed optical and IR observations will be necessary to confirm this assumption.  \\section[]{Identification of CXO J172337.5--373442}  The newly detected source CXO J172337.5--373442 is not the transient LMXB XTE J1723--376, because {\\it ASCA} measured the position of the latter (during its outburst with $0'.5$ accuracy; Marshall et al. 1999), which is $\\approx 5'$ away from the former. {\\it ASCA} did not detect CXO J172337.5--373442, because (1) the expected $7.762\\times10^{-3}$ count rate of {\\it ASCA} GIS, and $8.065\\times10^{-3}$ count rate of {\\it ASCA} SIS were not enough to significantly detect this source for the exposure of $\\approx 10$ ks; and (2) the PSF of the bright XTE J1723--376 hindered the detection.  The $4\\sigma$ tail is the most striking feature of CXO J172337.5--373442. Therefore, we primarily use this tail in order to identify the source. The tail could be a jet from a protostar. However, we reject this option as no star forming region is found (in optical or IR) at the source location. The tail could also be a jet from an AGN. X-ray jets are common in such systems (Harris \\& Krawczynski 2006). However, we think that CXO J172337.5--373442 is very unlikely to be an AGN for the following reasons. (1) No host galaxy is seen (in optical or IR) at the source location. (2) The Galactic coordinates of the source are L $= 350.250209$, B $= -0.824652$, which shows that the source is very close to the Galactic plane and bulge. It should be difficult to detect an AGN in such a direction because of relatively high absorption. (3) Many AGN spectra show a soft excess below 2 keV (above that expected from a simple powerlaw; Mushotzky et al. 1993). But the spectrum of CXO J172337.5--373442 is adequately fitted with a single (absorbed) powerlaw, and does not require any additional soft X-ray component (see \\S~2 and Fig. 3). (4) More importantly, {\\it Chandra} data analysis gives $N_{\\rm H} = 0.37^{+0.10}_{-0.08}\\times10^{22}$ cm$^{-2}$ (\\S~2), while the total Galactic neutral hydrogen column density in this direction is $1.47\\times10^{22}$ cm$^{-2}$ (\\S~3). This strongly suggests that CXO J172337.5--373442 is a Galactic source. Here we note that to the best of our knowledge the neutral hydrogen data in the source direction is not sufficient to measure the source distance. However, since the source Galactic latitude is low, we assume that the aforesaid $1.47\\times10^{22}$ cm$^{-2}$ value is for a column length of $10-20$ kpc (which is consistent with the data given in Diplas \\& Savage 1994). With this, and with the assumption that the neutral hydrogen density is uniform for a length scale $> 1$ kpc, the observed $N_{\\rm H} = 0.37^{+0.10}_{-0.08}\\times10^{22}$ cm$^{-2}$ gives a crude source distance of $2-6$ kpc.  CXO J172337.5--373442 could also be an X-ray binary system, because X-ray jets have been detected from such sources (Liu et al. 2006; 2007). However, this source may not be an LMXB, because its X-ray spectrum is relatively hard and plausibly entirely nonthermal (\\S~2; Bhattacharya \\& van den Heuvel 1991). Moreover, if source A is the optical counterpart of CXO J172337.5--373442, then it cannot be an LMXB, because (1) source A is very unlikely to be a low mass star (\\S~3); and (2) optical to X-ray luminosity ratio for LMXBs is less than 0.1 (Bhattacharya \\& van den Heuvel 1991), while this ratio is greater than 1 for CXO J172337.5--373442 (see Fig. 4). CXO J172337.5--373442 can be an HMXB, because the observed optical flux (assuming source A is the optical counterpart) is greater than the X-ray flux. However, to the best of our knowledge, so far only one (out of 114) Galactic HMXB is known to have extended X-ray jets/lobes (SS 433; Liu et al. 2006). This suggests that the probability is low for CXO J172337.5--373442 to be an HMXB. The following points strengthen this tentative conclusion to some extent. (1) SS 433 has two X-ray lobes on the opposite sides of the central source (Migliari et al. 2002), while CXO J172337.5--373442 has an X-ray streak (or, tail) on one side (see Fig. 1). (2) The point source component of SS 433 is harder ($\\Gamma = 1.40\\pm0.04$; Namiki et al. 2003) than the point source portion of CXO J172337.5--373442 (see Table 1). (3) The SS 433 lobes are spectrally softer (powerlaw $\\Gamma = 2.1\\pm0.2$; Migliari et al. 2002) than its point source component, unlike CXO J172337.5--373442 (\\S~2). (4) If CXO J172337.5--373442 is a Galactic source (see the previous paragraph), then it is at least about 100 times less X-ray luminous than SS 433. (5) if source A is the optical counterpart of CXO J172337.5--373442, then it  may not be an HMXB, because source A is plausibly not a star (\\S~3).  We think that the point source portion of CXO J172337.5--373442 is likely a pulsar, and the tail is a PWN. This is because many pulsars have PWNe, that are observed as extended X-ray sources (Kargaltsev \\& Pavlov 2008). These PWNe have various components, including long streaks or tails (Pavlov et al. 2006; Kargaltsev \\& Pavlov 2008) like the one observed from CXO J172337.5--373442. Moreover, X-ray streaks are conventionally interpreted as PWNe (Muno et al. 2008). In addition, the following points show that the properties of CXO J172337.5--373442 are consistent with those of pulsars and PWNe. (1) The powerlaw index of the X-ray spectrum of the point source part of CXO J172337.5--373442 is consistent with that of pulsars (see Table 1; and Kargaltsev \\& Pavlov 2008). (2) The powerlaw index of the X-ray spectrum of our source tail is consistent with that of PWNe (e.g., Bhattacharyya S. et al., in preparation). (3) The X-ray point source luminosity to the X-ray tail luminosity ratio ($\\approx 14$) of CXO J172337.5--373442 is consistent with that of PWNe, which can be in the range of $\\sim 0.1 - 45$ (Kargaltsev \\& Pavlov 2008; Bhattacharyya S. et al., in preparation). (4) For a source distance of 5 kpc, the X-ray luminosity of the point source portion of CXO J172337.5--373442 is $\\sim 1.5\\times10^{33}$ ergs s$^{-1}$. Therefore, for a reasonable source distance both the point source luminosity and the tail luminosity are very consistent with those observed from pulsars and PWNe (see the Table 2 of Kargaltsev \\& Pavlov 2008). (5) If the source A is the optical counterpart of CXO J172337.5--373442, then this source can be a pulsar for the following reasons. (a) The optical colour of source A is consistent with that of Galactic pulsars (\\S~3). (b) The $N_{\\rm H}$ values ($0.47\\times10^{22}$ cm$^{-2}$ and $0.52\\times10^{22}$ cm$^{-2}$; \\S~3) inferred from the pulsar identification of source A are well within the 68.3\\% and 90\\% (respectively) confidence limits of the value inferred from the {\\it Chandra} data.  We note that, although CXO J172337.5--373442 can be a pulsar, the identification of source A and CXO J172337.5--373442 as the same pulsar is inconsistent with the fact that, unlike our case, the optical flux of a pulsar is normally less than its X-ray flux. This can be seen from Fig. 4 of Thompson et al. (1999). Therefore, although from the same figure we note that the optical flux may be greater than the X-ray flux for some pulsars (as indicated for PSR B1706-44), CXO J172337.5--373442 cannot be identified as a pulsar with certainty using the currently available data. ",
        "conclusions": "In this Letter, we report the detection of a point X-ray source with high significance, and the discovery of a $4\\sigma$ X-ray tail emanating from the point source. This source can be a Galactic HMXB or a Galactic pulsar. From the discussion in \\S~4, we conclude that the latter identification is favoured. Therefore, we will now briefly discuss some of the properties of this source assuming it to be a pulsar. We notice that the point source (pulsar) X-ray energy spectrum is likely entirely nonthermal (powerlaw; see \\S~2), which implies that the X-ray emission of the pulsar primarily originates in its magnetosphere (Kargaltsev et al. 2005). This means either the thermal emission from the neutron star's surface is largely obscured (possibly by the plasma in the magnetosphere), or the stellar surface temperature is relatively low. Since no supernova remnant (SNR) is seen near CXO J172337.5--373442, the pulsar has probably escaped from its SNR, implying that it is likely middle-aged or old. As the pulsar is moving through the ISM, its speed may be supersonic as the sound speed in ISM is lower than that in SNR. In such a case, the X-ray taillike PWN may originate from synchrotron emission in a shocked region behind the pulsar (Pavlov et al. 2006). Alternatively, such a tail may be a result of synchrotron process in a jet emanating from the pulsar's magnetosphere along the spin axis (Benford 1984). The observed clumpiness and the bending of the tail (Fig. 1) may be caused by the sausage instability and the kink instability respectively. Finally, we note that if source A and CXO J172337.5--373442 are the same pulsar, then it is a unique pulsar with optical luminosity much higher than the X-ray luminosity. On the other hand, if CXO J172337.5--373442 is an HMXB, then, to the best of our knowledge, this will be the second Galactic HMXB known to have a visible X-ray jet. In order to resolve this, further optical and X-ray observations of the source is essential."
    },
    "0809/0809.4015_arXiv.txt": {
        "abstract": "We use the Fifth Data Release of the Sloan Digital Sky Survey to study X-ray-selected galaxy groups and compare their properties to clusters. We search for infall patterns around the groups and use these to measure group mass profiles to large radii.  In previous work, we analyzed infall patterns for an X-ray-selected sample of 72 clusters from the ROSAT All-Sky Survey.  Here, we extend this approach to a sample of systems with smaller X-ray fluxes selected from the 400 Square Degree serendipitous survey of clusters and groups in ROSAT pointed observations.  We identify 16 groups with SDSS DR5 spectroscopy, search for infall patterns, and compute mass profiles out to 2-6 $\\Mpc$ from the group centers with the caustic technique. No other mass estimation methods are currently available at such large radii for these low-mass groups, because the virial estimate requires dynamical equilibrium and the gravitational lensing signal is too weak.  Despite the small masses of these groups, most display recognizable infall patterns.  We use caustic and virial mass estimates to measure the scaling relations between different observables, extending these relations to smaller fluxes and luminosities than many previous surveys.  Close inspection reveals that three of the groups are subclusters in the outskirts of larger clusters.  A fourth group is apparently undergoing a group-group merger.  These four merging groups represent the most extreme outliers in the scaling relations.  Excluding these groups, we find $L_X\\propto\\sigma_p^{3.4\\pm1.6}$, consistent with previous determinations for both clusters and groups.  Understanding cluster and group scaling relations is crucial for measuring cosmological parameters from clusters.  The complex environments of our group sample reinforce the idea that great care must be taken in determining the properties of low-mass clusters and groups. ",
        "introduction": "A large fraction of all galaxies are members of groups.  Groups are a less extreme and much more common type of system than galaxy clusters. Because they are less massive than clusters, groups are currently less well understood than clusters.  Here, we investigate the optical properties of a sample of X-ray-selected groups to determine whether they can be modeled as scaled-down versions of clusters\\footnote{The division between groups and clusters is somewhat arbitrary; here, we define groups to be systems of galaxies with virial masses smaller than $10^{14} h^{-1} M_\\odot$.}.  In particular, we study the outskirts of the groups to determine if their infall regions are readily identifiable as they are in the outskirts of clusters.  With large samples of spectroscopic members, we then estimate the virial masses of the groups and determine if the groups obey the same scaling relations as clusters.  Galaxy clusters and groups are surrounded by infall regions in which galaxies are bound to the system but are not in equilibrium. The Cluster Infall Regions in the Sloan Digital Sky Survey \\citep[][hereafter CIRS]{cirsi} project showed that X-ray-selected clusters display a characteristic trumpet-shaped pattern in radius-redshift phase space diagrams.  These patterns, termed caustics, were first predicted for simple spherical infall onto clusters \\citep{kais87,rg89}, but later work showed that these patterns reflect the dynamics of the infall region \\citep[][hereafter DG]{dg97} and \\citep[][hereafter D99]{diaferio1999}. CIRS and earlier similar studies \\citep[e.g.,][]{gdk99,cairnsi} showed that the amplitude of the caustics yields an estimate of cluster mass profiles consistent with both virial and X-ray mass estimates where the techniques overlap.  More recently, \\citet{diaferio05} and \\citet{lemze09} showed that caustic masses also agree well with mass estimates based on gravitational lensing.  Because neither galaxies nor gas is expected to be in equilibrium outside the virial radius of a cluster, the caustic technique and weak lensing are the only well-studied methods for determining cluster mass profiles at large radii \\citep[see][for a recent review of the caustic technique]{diaferio09}.    CIRS showed that infall patterns are ubiquitous in nearby massive clusters selected by X-ray emission. The CIRS clusters are fairly massive clusters and generally have little surrounding large-scale structure \\citep[but see][]{rines01b,rines02}.  One might suspect that the presence of infall patterns is limited to massive, isolated clusters.  However, other investigators have found infall patterns around the Fornax Cluster \\citep{drink}, the Shapley Supercluster \\citep{rqcm}, an ensemble cluster comprised of poor clusters in the Two Degree Field Galaxy Redshift Survey \\citep{bg03}, and even the galaxy group associated with NGC 5846 \\citep{mahdavi05}. Here, we extend the study of infall patterns by studying a large number of systems with smaller X-ray fluxes than the CIRS clusters. Because these systems typically have smaller masses and fewer member galaxies (more typical of groups), infall patterns might not be identifiable in these systems.  In particular, we use the new 400 Square Degree (400d) survey of clusters and groups in {\\em ROSAT} PSPC pointed observations \\citep{burenin07}.  The 400d survey is the largest area cluster survey extending to flux limits of $\\approx$1.4$\\times 10^{-13}$erg s$^{-1}$cm$^{-2}$.  This flux limit is a factor of 20 smaller than the flux limit of catalogs based on the {\\em ROSAT} All-Sky Survey \\citep[RASS;][]{rass}.  We search for 400d clusters and groups in the spectroscopic footprint of the Sloan Digital Sky Survey Data Release 5 \\citep[][SDSS DR5]{dr5}.  This approach is similar to CIRS and to the RASS-SDSS \\citep{popesso04,popesso05} analysis of clusters in DR2. Compared to both the CIRS and RASS-SDSS samples, the 400d groups are expected to have significantly smaller masses on average.  One motivation for this study is to improve our understanding of cluster scaling relations, in particular to extend these relations to the regime of high-mass groups.  Ambitious cluster surveys like the South Pole Telescope and Atacama Cosmology Telescope require a good understanding of cluster scaling relations to accurately measure cosmological parameters from the abundance and evolution of clusters. Quantifying the scatter in these relations and any Malmquist-like bias is crucial for obtaining robust cosmological constraints from cluster surveys \\citep[e.g.,][]{stanek06}.  Many forecasts for dark energy constraints from cluster surveys assume that scatter in the mass-observable scaling relations has a log-normal distribution \\citep[e.g.,][]{mantz08}.  Deviations from this assumption could significantly impact the constraining power of these surveys.  Recent studies of scaling relations that extend into the group regime have often reached contradictory conclusions.  For instance, a common scaling relation is the $L_X-\\sigma_p$ relation between the X-ray luminosity $L_X$ of a cluster or group and its projected velocity dispersion $\\sigma_p$.  For massive clusters, this relation is typically found to have a steep slope of 4.4$^{+0.7}_{-0.3}$ \\citep{andilxsig} or 3.7$\\pm$0.3 \\citep{popesso05}.  For groups, various studies have found slopes as shallow as 0.37 \\citep[][ the fit is a broken power-law with a slope of 4.02 for clusters]{rasscals} and as steep as 4.7$\\pm$0.9 \\citep{helsdon00b}.  A detailed optical-X-ray study by \\citet{osmond04} recently found a slope of 2.5$\\pm$0.4.  The variety in slopes is produced by many factors, including differing definitions of $L_X$, and in some cases, possible evolution of the galaxies via dynamical friction \\citep{helsdon05}. One problem is that few large, complete group catalogs are available, so many existing studies utilize either heterogeneous samples or include many groups with limited data on each group.  One solution is to utilize X-ray selection \\citep[e.g.,][]{mulchaey98,rasscals}, but existing cluster/group catalogs based on RASS \\citep[][CIRS]{popesso05} have been restricted to systems with relatively high X-ray fluxes (and therefore including few groups). The fainter X-ray flux limits of the 400d survey allow detection of several groups.  Combining this catalog with optical data from the Sloan Digital Sky Survey \\citep{sdss}, we can test whether the scaling relations show similar behavior for systems with smaller fluxes more typical of galaxy groups.  A complementary approach (that we do not apply here) to understanding scaling relations for low-mass systems is to identify the systems optically and use stacking analysis to measure their ensemble properties.  This approach has been applied previously with RASS data \\citep{dai07,rykoff08}.   We describe the data and the group sample in $\\S$ 2.  In $\\S$ 3, we review the caustic technique and use it to estimate the group mass profiles, discuss cluster scaling relations, and compare the caustic mass profiles to simple parametric models.  We discuss some individual groups in $\\S 4$. We compare the caustic mass profiles to X-ray and virial mass estimators in $\\S$ 5.  We discuss our results and conclude in $\\S 6$.  We assume $H_0 = 100~h~\\kms \\Mpc^{-1}, \\Omega _m = 0.3, \\Omega _\\Lambda = 0.7$ throughout. ",
        "conclusions": "We use the Fifth Data Release of the Sloan Digital Sky Survey to study galaxy groups and compare their properties to more massive clusters. In particular,we study the presence of infall patterns around galaxy groups, use these patterns to measure mass profiles, and use these masses to compare the scaling relations of groups and clusters.  This study extends our previous work on infall patterns around X-ray clusters in SDSS \\citep{cirsi} to smaller X-ray fluxes more typical of groups.  These systems typically have smaller X-ray luminosities and masses than the systems studied in previous investigations.  Despite their smaller masses (compared to CAIRNS and CIRS clusters), well-defined infall patterns are present in most of these groups. Four of the 16 groups are contaminated by large-scale structure, preventing an accurate determination of their infall patterns.   We use the infall patterns to compute mass profiles for the groups and compare them to model profiles.  Group infall regions are well fit by NFW and Hernquist profiles and poorly fit by singular isothermal spheres, similar to the results from CAIRNS and CIRS for clusters. Observed clusters and groups resemble those in simulations, and their mass profiles are well described by extrapolations of NFW or Hernquist models out to the turnaround radius.  The shapes of the best-fit NFW group mass profiles agree reasonably well with the predictions of simulations; the average concentration is $c_{200}$$\\approx$6.5, slightly smaller than the value for group size halos extrapolated from \\citet{bullock01} and with similar scatter.  These mass profiles test the shapes of dark matter haloes on a scale difficult to probe with weak lensing or any other mass estimator.  There has been much discussion of whether galaxy groups follow different scaling relations than clusters, particularly with respect to the $L_X-\\sigma_p$ relation \\citep[e.g.,][]{rasscals,helsdon00b,osmond04}.  Using the caustics to determine group membership, velocity dispersions, and mass profiles, the 400d-SDSS groups generally follow the same scaling relations as clusters, but the small number of 400d-SDSS groups does not allow us to exclude results suggesting that groups have a shallower $L_X-\\sigma_p$ relation.   Understanding scaling relations of clusters and groups is critical for future attempts to constrain dark energy with cluster abundance measurements.  Detailed mass estimates are difficult to obtain for groups and low-mass clusters because their X-ray emission is typically faint and because they contain fewer galaxies than massive clusters. This work represents part of an effort to increase the sample of groups and low-mass clusters with detailed mass estimates.  Future work is needed to measure cluster scaling relations robustly across a wide mass range.  In particular, deep high-resolution X-ray observations of groups like those studied here could enable measurements of group masses and X-ray temperatures, while additional optical spectroscopy could improve measurements of velocity dispersions and dynamical masses."
    },
    "0809/0809.2405_arXiv.txt": {
        "abstract": "We present numerical studies of the properties of the stellar velocity distribution in galactic disks which have developed a saturated, two-armed spiral structure.  In previous papers we used the Boltzmann moment equations (BME) up to second order for our studies of the velocity structure in self-gravitating stellar disks. A key assumption of our BME approach is the zero-heat flux approximation, i.e.\\ the neglection of third order velocity terms. We tested this assumption by performing test particle simulations for stars in a disk galaxy subject to a rotating spiral perturbation. As a result we corroborated qualitatively the complex velocity structure found in the BME approach. It turned out that an equilibrium configuration in velocity space is only slowly established on a typical timescale of 5 Gyrs or more. Since many dynamical processes in galaxies (like the growth of spirals or bars) act on shorter timescales, pure equilibrium models might not be fully appropriate for a detailed comparison with observations like the local Galactic velocity distribution.  Third order velocity moments were typically small and uncorrelated over almost all of the disk with the exception of the 1:4 resonance region (UHR). Near the UHR (normalized) fourth and fifth order velocity moments are still of the same order as the second and third order terms. Thus, at the UHR higher order terms are not negligible. ",
        "introduction": "\\label{sectheis:intro}  It is long known that the velocity distribution of stars in galactic disks is non-isotropic \\cite{kobold90}.  Kapteyn and Eddington invoked a superposition of (isotropic) stellar streams with different mean velocities, by this creating an anisotropic velocity distribution (\\cite{kapteyn05}, \\cite{eddington06}).  However, an alternative interpretation by Schwarzschild became the general framework for describing the local velocity distribution of stars of equal age \\cite{schwarzschild07}.  The Schwarzschild distribution is based on a single but anisotropic ellipsoidal distribution function. It is characterized by a Gaussian distribution in all three directions $U$ (radial $r$), $V$ (tangential $\\phi$) and $W$ (vertical $z$) in velocity space, but with different velocity dispersions $\\sigma_{rr}$, $\\sigma_{\\phi\\phi}$, and $\\sigma_{zz}$. In general, the velocity distribution is described by a velocity dispersion tensor \\begin{equation} \\sigma^2_{ij} \\equiv \\overline{(v_i - \\bar{v}_i)(v_j - \\bar{v}_j)}, \\label{mixed_vd} \\end{equation} where $i$ and $j$ denote the different coordinate directions and $v_i$ gives the corresponding velocity. The bar denotes local averaging over velocity space.  The principal axes of this tensor form an imaginary surface that is called the velocity ellipsoid. This ellipsoid is characterized by its anisotropy, measured by the ratio of the velocity dispersions along the principal axes and its orientation.  Because the principal axes need not to be aligned with the coordinate axes, the vertex deviation $l_{\\rm v}$ (which is defined as the angle between the direction from the Sun to the Galactic centre and the direction of the major principal axis of the velocity ellipsoid) needs not to vanish.  In the case of stationary, axisymmetric systems and appropriate distribution functions (DF) the velocity ellipsoid is aligned with the coordinate axes, i.e.\\ $l_{\\rm v} = 0$. However, non-vanishing vertex deviations were found in the solar vicinity in many studies (\\cite{wielen74}, \\cite{mayor72}, \\cite{ratnatunga97}). These vertex deviations could be explained by spiral structures (\\cite{mayor70}, \\cite{yuan71}), by this supporting the importance of non-axisymmetric gravitational potentials.  A major step in analysing the local stellar velocity space was the astrometric satellite mission Hipparcos by ESA \\cite{esa97}. Earlier results concerning the general behaviour of the velocity ellipsoid, e.g.\\ the dependence of the dispersions or the vertex deviations on $B-V$ were corroborated (\\cite{dehnen98b}, \\cite{bienayme99}, \\cite{hogg05}). Moreover, the more numerous known proper motions allowed for a detailed mapping of the velocity space in the solar vicinity.  Studies by Dehnen or by Alcob\\'e \\& Cubarsi turned out that the local velocity distribution of stars exhibits a rich substructure (\\cite{dehnen98a}, \\cite{alcobe05}). For example, two major peaks in the velocity distribution were found where the smaller secondary peak is well detached by at least 30 km\\,s$^{-1}$ from the main peak (the ``u anomaly'').  Dehnen attributed this bimodality to the perturbation exerted by the Galactic bar assuming that its outer Lindblad resonance (OLR) is close to the Sun \\cite{dehnen00}. M\\\"uhlbauer \\& Dehnen showed that a central bar can also explain vertex deviations of about 10$^\\circ$ % \\cite{muehlbauer03}.  In general, a non-axisymmetric gravitational potential might lead to a misalignment of the velocity ellipsoid.  In the case of spiral perturbations, this conclusion was made e.g.\\ by Mayor (\\cite{mayor70}, for a review see Kuijken \\& Tremaine \\cite{kuijken91}) and numerically confirmed recently by Vorobyov \\& Theis (\\cite{vorobyov06}, hereafter VT06). However, it is not clear if the non-axisymmetric gravitational field is entirely responsible for the observed vertex deviation. The existence of moving groups of stars was also shown to produce large vertex deviations \\cite{binney98}. The situation may become even more complicated because moving groups of stars may in turn be caused by the non-axisymmetric gravitational field of spiral arms.  Therefore, a detailed numerical study of the vertex deviation in spiral galaxies is necessary.  In this paper, we present results of our analysis of the velocity structure in disk galaxies caused by spiral perturbations. In earlier papers (VT06, \\cite{vorobyov08}) we presented numerical solutions of the Boltzmann moment equations (BME) up to second order for flat disks.  In Sect.\\ \\ref{sectheis:beads2d} we briefly describe this method and its results with respect to the vertex deviation. A characteristic of our BME approach is the assumption of a zero heat-flux, i.e.\\ we neglect all velocity terms of higher order than 2. In order to test this assumption, we developed a test particle code which measures the velocity moments up to fifth order.  This code and first results about the importance of higher order terms are presented in Sect.\\ \\ref{sectheis:testparticle}. ",
        "conclusions": "\\label{sectheis:summary} We analysed the structure of the velocity space in galactic disks which are subject to a spiral perturbation. In previous papers we used the Boltzmann moment equations up to second order in order to study the growth and saturation of spiral structure in self-gravitating stellar disks. From this analysis we derived the properties of the velocity ellipsoid all over the disk, namely the vertex deviations and the Oort ratios.  A key assumption of the BME approach is the zero-heat flux approximation, i.e.\\ the neglection of third order velocity terms.  We tested this assumption by performing test particle simulations for stars in a disk galaxy subject to a rotating spiral perturbation. As a result we corroborated qualitatively the complex velocity structure found in the BME approach. It turned out that an equilibrium configuration in velocity space is only slowly established on a typical timescale of 5 Gyrs or more. Since many dynamical processes in galaxies (like the growth of spirals or bars) act on shorter timescales, pure equilibrium models might not be fully appropriate for a detailed comparison with observations like the local Galactic velocity distribution.  In our simulations third order velocity moments were typically small and uncorrelated over almost all of the disk with the exception of the 1:4 resonance region (UHR).  Near the UHR (normalized) fourth and fifth order velocity moments are still of the same order as the second and third order terms. Thus, at the UHR higher order terms are not negligible.\\\\  \\noindent {\\bf Acknowledgements}. CT is very grateful to the organizers of the meeting for a very interesting and inspiring conference as well as for the financial support.  Additionally, we would like to thank Christoph Lhotka for his advice in using Mathematica."
    },
    "0809/0809.3686_arXiv.txt": {
        "abstract": "\\noindent In the past year, the HiRes and Auger collaborations have reported the discovery of a high-energy cutoff in the ultra-high energy cosmic-ray (UHECR) spectrum, and an apparent clustering of the highest energy events towards nearby active galactic nuclei (AGNs). Consensus is building that such $\\sim 10^{19}$--$10^{20}$ eV particles are accelerated within the radio-bright lobes of these sources, but it is not yet clear how this actually happens. In this paper, we report (to our knowledge) the first treatment of stochastic particle acceleration in such environments from first principles, showing that energies $\\sim 10^{20}$ eV are reached in $\\sim 10^6$ years for protons. However, our findings reopen the question regarding whether the high-energy cutoff is due solely to propagation effects, or whether it does in fact represent the maximum energy permitted by the acceleration process itself. ",
        "introduction": "Cosmic rays are energetic charged particles traveling throughout the Galaxy and the intergalactic medium under the influence of various physical processes, including deflection by magnetic fields, and collisions with other particles along their trajectory. Their energy spectrum measured at Earth is a steep (roughly power-law) distribution with logarithmic index $\\alpha\\sim 2.6$--3, extending up to a few times $10^{20}$ eV. Other than their spectrum, these particles are characterized by their angular distribution in the sky, and by their mass composition.  A highly significant steepening in the UHECR spectrum was reported by both the HiRes (Abbasi et al. 2008) and \\citet{auger2008c} collaborations. (Distinguished from their lower-energy counterparts, UHECRs have energies in excess of 1 EeV $\\equiv 10^{18}$ eV.) This result may be a strong confirmation of the predicted Greisen-Zatsepin-Kuzmin (GZK) cutoff due to photohadronic interactions between the UHECRs and low-energy photons in the cosmic microwave background (CMB) radiation (Greisen 1966; Zatsepin \\& Kuz'min 1966). Together with the measured low fraction of high-energy photons in the CR distribution, this measurement already rules out so-called top-down models, in which the UHECRs represent the decay products of high-mass dark matter particles created in the early Universe (Semikoz et al. 2007). The measured photon flux is also in conflict with scenarios in which UHECRs are produced by collisions between cosmic strings or topological defects (Bluemer et al. 2008, Auger 2008b). On the other hand, such energetic particles may still be produced via astrophysical acceleration mechanisms (see Torres \\& Anchordoqui 2004 and other references cited therein).  UHECRs are not detected directly, but through the showers they create in Earth's atmosphere (see, e.g., Melia 2009). Depending on the energy and type of primary particle, the ensuing cascade has characteristics that allow the ground-based observatories to determine not only whether the incoming UHECR is a photon, but also its atomic number. It should be pointed out, however, that a determination of the primaries' composition strongly relies on an extrapolation of current phenomenological hadronic interaction models, so it remains rather uncertain. The \\citet{auger2007} data confirmed the dominance of protons in primary cosmic rays, though they also exhibit evidence for a mixed composition extending to energies as high as $\\sim$ 50--60 EeV, with a higher atomic number $Z$, up to $Z \\sim 26$ (Unger et al. 2007).  But the most telling indicator for the possible origin of these UHECRs is the discovery by Auger (Auger 2008a) of their clustering towards nearby ($\\sim 75$ Mpc) AGNs along the supergalactic plane. The significance of this correlation has been further strenghtened by a more recent analysis which weights the AGN spatial distribution by their hard X-ray flux (George et al. 2008). This raises at least two questions: (1) How are the UHECRs accelerated to such high energies? and (2) given these nearby sources, is the sharp suppression of UHECRs in the last decade of their observed energies really due solely to the GZK effect, or does it signal a limitation to the acceleration efficiency?  Previous attempts at understanding how particles are accelerated to EeV energies and beyond have generally been based on first-order Fermi acceleration (see, e.g., Ostrowski 2008, and other references therein) within shocks created by blast waves like those in supernova remnants (Fatuzzo \\& Melia 2003, Crocker et al. 2005). But this process is subject to kinematic restrictions that inhibit the particles from actually reaching ultra-high energies (see, e.g., Nayakshin \\& Melia 1998, Gallant et al. 1999). Recent numerical simulations have shown that an increase in the Lorentz factor $\\gamma$ of ultra-relativistic shock waves steepens the observed spectrum (Niemiec \\& Ostrowski 2006) and reduces its high-energy cutoff.  For these reasons, it is not plausible for UHECRs to emerge from astrophysical environments, such as supernova remnants, where first-order processes are dominant so long as the shock velocity is super-Alfv\\'enic, because they cannot even contain such high-energy particles (Hillas 1984)---the gyration radius of particles with energy $\\sim 10^{20}$ eV for a typical galactic magnetic field is much larger than the size ($<10$ pc) of these structures.  On the other hand, a second-order Fermi process (Fermi 1949) can explain observational features not addressed by the first-order process, as in the case of a supernova remnant itself (see Cowsik \\& Sarkar 1984). Moreover, stochastic particle acceleration through a gyroresonant interaction with MHD turbulence (a second-order Fermi process; see Fermi 1949) can be very efficient if the Alfv\\'en velocity approaches $c$ (Dermer \\& Humi 2001). The stochastic acceleration of particles by turbulent plasma waves has already received some attention in the literature (see Liu et al. 2004, 2006, and references cited therein, and Wolfe \\& Melia 2006). Indeed, the feasibility of second-order Fermi acceleration in radio galaxies has been demonstrated through the steady re-acceleration of electrons in certain hot spots (Almudena Prieto et al. 2002).  Our treatment from first principles, however, avoids many of the previously encountered unknowns and limitations. In this paper, we report (to our knowledge) the first treatment of stochastic acceleration of charged particles in the lobes of radio-bright AGNs by directly computing the trajectory of individual particles. An earlier version of this treatment---for the propagation of charged particles assumed already accelerated at TeV energy through the turbulent magnetic field at the Galactic centre---may be found in Ballantyne et al. 2007, and Wommer, Melia \\& Fatuzzo 2008; by contrast, in the present paper both the propagation and acceleration are taken into account. We show that random scatterings (a second-order Fermi process) between the charges and fluctuations in a turbulent magnetic field can accelerate these particles up to ultra-high energies, provided a broad range of fluctuations is present in the system. ",
        "conclusions": "We have shown that a region containing a Kolmogorov (turbulent) distribution of non-relativistic Alfv\\'en waves can accelerate particles to ultra-high energies. The physical parameters in these regions are compatible with those believed to be operating in the radio lobes of AGNs. We have discussed the predicted differential spectrum within the parameter space of the model, characterized by the size ${\\mathcal R}$ of the acceleration region and the turbulent magnetic energy. Possible tests of this model involve the synchrotron or IC emission by a population of similarly accelerated electrons.  As the Auger observatory continues to gather more data, improving on the statistics, our UHECR source identification will continue to get better. Eventually, we should be able to tell how significant the GZK effect really is, and whether the cutoff in the CR distribution is indeed due to propagation effects, or whether it is primarily the result of limitations in the acceleration itself. Given the fact that energies as high as $\\sim 10^{20}$ eV may be reached within typical radio lobes, it is possible that both of these factors must be considered in future refinements of this work."
    },
    "0809/0809.4259_arXiv.txt": {
        "abstract": "We study model predictions for the abundance and clustering of high-redshift galaxies, and for the properties of their descendants. We focus on three high-redshift populations: Lyman break galaxies at $z\\sim 3$ (LBGs), optically selected star-forming galaxies at $z\\sim 2$ (BXs), and distant red galaxies at $z\\sim 2$ (DRGs).  We select model galaxies from mock catalogues using the observationally defined colour and magnitude criteria.  With plausible dust assumptions, our galaxy formation model can simultaneously reproduce the abundances, the redshift distributions and the clustering of all three observed populations. The star formation rates (SFRs) of model LBGs and BXs are lower than those quoted for the real samples, reflecting different initial mass functions and scatter in model dust properties.  About 85\\% of model galaxies selected as DRGs are star-forming, with SFRs ranging up to $\\sim 10^2M_\\odot$/yr. Model LBGs, BXs and DRGs together account for less than half of all star formation over the range $1.5<z<3.2$; many massive, star-forming galaxies are predicted to be too heavily obscured to appear in these populations.  Model BXs have metallicities which agree roughly with observation, but model LBGs are only slightly more metal-poor, in disagreement with recent observational results.  The model galaxies are predominantly disk-dominated. Stellar masses for LBGs and BXs are typically $\\sim 10^{9.9}M_\\odot$, and for DRGs are $\\sim 10^{10.7}M_\\odot$. Only about 30\\% of model galaxies with $M>10^{11}M_\\odot$ are classified as LBGs or BXs at the relevant redshifts, while 65\\% are classified as DRGs.  Almost all model LBGs and BXs are the central galaxies of their dark halos, but about a quarter of DRGs are satellites.  Half of all LBG descendants at $z=2$ would be identified as BX's, but very few as DRGs.  Clustering increases with decreasing redshift for descendants of all three populations, becoming stronger than that of $L^*$ galaxies by $z=0$, when many have become satellite galaxies and their mean stellar mass has increased by a factor of 10 for LBGs and BXs, and by a factor of 3 for DRGs.  This growth is dominated by star formation until $z\\sim 1$ and thereafter by mergers. Merging is predicted to be more important for LBG and DRG descendants than for BX descendants. Most LBGs and DRGs end up as red ellipticals, while BXs have a more varied fate. 99\\% of local galaxies with $M_*>10^{11.5}M_\\odot$ are predicted to have at least one LBG/BX/DRG progenitor, and over 70\\% above $10^{11}M_\\odot$ ",
        "introduction": "\\label{sec:intro} The redshift interval $1<z<3$ is a very important epoch in the history of galaxy formation. During these several billion years, the star formation rate per unit comoving volume, the abundance of luminous quasars and the specific merger rate of galaxies all reached their peak values. This is when the Hubble sequence of galaxies was established, and most galactic stars were formed.  The last few decades have seen a remarkable development in the observational study of high-redshift galaxies. Using strong the Lyman break feature in the spectrum of star-forming galaxies, \\cite{Steidel96} developed colour-colour criteria to select so-called Lyman Break Galaxies (LBGs) at $z\\sim 3$ using deep $U_nGR$ photometry. Extensions of this technique allowed the selection of large samples of star-forming galaxies at lower redshifts, specifically $z\\sim 2.3$ (BXs) and $z\\sim 1.7$ (BMs) \\citep{Adelberger2004}. Recent observations show that there are many high-redshift galaxies with rather little rest-frame UV luminosity which are missed in such optically selected surveys. \\cite{Franx2003} have successfully developed near-infrared colour criteria to select distant red galaxies (DRGs) at $z\\sim 2$ using deep $JK$ photometry. Follow-up observations at a variety of wavelengths have clarified the physical properties of all these high-redshift galaxy populations by providing star-formation rates \\citep{Erb06a,Reddy06}, stellar masses \\citep{Shapley05,Erb06b,Kriek2006,Papovich2006}, morphologies \\citep{Abraham1996,Papovich2005,Law2007}, dust luminosities \\citep{Webb2003}, kinematics \\citep{Pettini2001,Erb06b} and clustering estimates \\citep{Adelberger2005,Quadri2008}.  Each of these observational samples provides information on a limited subset of the galaxy population at a specific cosmic epoch, and it is obviously of interest to understand their relation to the population as a whole, both at the redshift of observation and at other redshifts, particularly $z=0$ where our knowledge of galaxies is best. Within the current standard cosmological paradigm, structure formation in the gravitationally dominant dark matter distribution can be simulated reliably and is now quite well understood \\citep[e.g.][]{SFW}. For the purposes of modelling galaxy evolution, it can usefully be idealised as a distribution of dark matter halos of ``universal'' structure which grow steadily in mass through accretion and merging. Galaxies form through condensation of gas within this evolving dark halo population, as first set out by \\cite{WR1978}. At any given time their distribution can be well modelled by populating halos with galaxies according to a simple recipe, and then adjusting parameters to fit the observed abundance and clustering. The recipe can be based either on a simplified model of the galaxy formation process or on fitting formulae which make no direct reference to the underlying physics \\citep{PS2000,Seljak2000}.  The flexibility of the latter Halo Occupation Distribution (HOD) approach allows excellent fits to galaxy luminosity functions and clustering properties at any given epoch (see Cooray \\& Sheth 2002 for a review) but provides no natural way to link populations at different epochs (see \\cite{Conroy2008} for a recent attempt to remedy this). In addition, detailed physical models are needed to assess how observational cuts on luminosity and colour affect the properties of high redshift samples.  A more straightforward approach is to follow galaxy formation directly within the evolving dark matter distribution. This extends the semi-analytic technique developed in the early 1990's \\citep{WF1991,KWG1993,Cole1994}, and has made increasingly sophisticated use of information from $N$-body simulations of cosmic structure formation \\citep{KNS1997,Kauffmann1999,Springel2001,Nature2005}. The idea is to design simple parametrised models for the baryonic physics, based either on observation or on more detailed simulations of individual systems, and to implement these recipes in the structure formation framework provided by a dark matter simulation. This provides a powerful tool for studying the formation, evolution and clustering of the galaxy population. It is much less resource-intensive than simulations involving a more direct treatment of baryonic processes, and as a result it allows the treatment of larger volumes and the exploration of a wider range of input parameters and physics recipes.  Recent improvements in such modelling have included the move to higher resolution N-body simulations, enabling use of the substructure information they provide \\citep{Springel2001, Helly2003,Hatton2003,Nature2005,Kang2005} and the inclusion of additional relevant physics, for example, AGN feedback \\citep{Croton2006, Bower2006}, galactic winds \\citep{Bertone2007} and gas stripping in clusters \\citep{Font2008}. These techniques have previously been used to study the properties of LBGs by \\cite{Blaizot2004}.   The semi-analytic models used in this paper are implemented on the \\emph{Millennium Simulation} \\citep[][hereafter MS]{Nature2005} and are a minor modification of those presented in De Lucia \\& Blaizot (2007, hereafter DLB07) and Kitzbichler \\& White (2007, hereafter KW07), which are publicly available at the MS download site\\footnote{http://www.mpa-garching.mpg.de/millennium}. These are refinements of the model originally published as \\cite{Croton2006}. They have successfully matched a wide range of galaxy properties both locally and at high redshift, but there are some notable discrepancies which demonstrate that the description of galaxy formation physics remains incomplete. Our models (and that of KW07) differ from the model of DLB07 only through the introduction of a redshift dependence in the way dust is treated. We construct light-cone surveys of our models as in KW07, and we select LBGs, BXs, and DRGs exactly as in observational studies. We then compare observed and simulated samples in terms of their abundance, their clustering and their distributions of redshift, colour, metallicity and star formation rate, finding good agreement in most cases. We move on to examine model predictions for the relation of these various populations to each other and to the high-redshift galaxy population as a whole, and for the properties of their descendants at lower redshift.   Our paper is organized as follows. In Sec.~\\ref{sec:simulation}, we summarize relevant properties of the \\emph{Millennium Simulation}, of the semi-analytic model and of our light-cone surveys. Sec.~\\ref{sec:mock} deals with the selection of model LBG, BX and DRG samples, and compares them with observed samples.  Sec.~\\ref{sec:result} examines the relation of these model samples to each other and to the full high-redshift population, and the properties of their descendants. We summarize our results in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have used the \\cite{LB2007} model for galaxy formation within the \\emph{Millennium Simulation} to explore the likely physical nature of observed high-redshift galaxies, specifically Lyman Break Galaxies, BX galaxies and Distant Red Galaxies, and their likely descendants at lower redshift. We first built mock catalogues in order to compare observed high-redshift galaxies with similarly selected objects in the simulation. We then used the full galaxy catalogue from the simulation to study the descendants of the high-redshift populations. We found it necessary to modify the original DLB07 dust model in order to match the abundance and colour of the observed high-redshift populations,  Impressively, with a proper dust model it is possible to match the observed abundances, redshift distributions and clustering of all three high-redshift populations in a model which also fits the properties of low-redshift galaxies. The descendants of all three populations become more strongly clustered at lower redshifts, and all are more clustered than $M^*$ galaxies today. A turn-up in clustering strength at small scales is evident at $z\\sim0$, which reflects the high satellite fraction among the descendents. The clustering of DRGs is the least consistent with observation. Angular correlations are well reproduced on scales between 7 and 100 arcsec but our model appears more weakly clustered than real DRGs on both smaller and larger scales. This results in an estimated correlation length for the observed sample which is larger than that of our model DRGs.  Our results show the expected scatter in angular correlation estimates to be quite large for areas as small as that currently observed, so a final judgement on this issue will require significantly larger observational surveys.   Together, model DRGs and BXs account for only 30\\% of the strongly star-forming galaxies ($\\dot{M}_*>20M_\\odot/yr$) at $z\\sim 2$. Most of the rest are gas- and metal-rich systems which are strongly obscured. In contrast, the model suggests that only 15\\% of all galaxies with $M_*>10^{11}M_\\odot$ at $z\\sim 2$ (most of which do indeed have high star formation rates) are missing in current optical and near-infrared surveys. Interestingly, the model predicts most DRGs to be star-forming galaxies. Their average SFR is even higher than those of LBGs and BXs, and they account for more than 65\\% of the galaxies with $M_*>10^{11}M_\\odot$ at $z\\sim2.2$.  There is rather little overlap between these star-forming DRGs and BXs or LBG descendants. On the other hand, half of all LBG descendants are identified as BXs at $z\\sim2.2$. The physical properties and the clustering of these LBG-BXs are quite similar to those of other LBG descendants. These consist of three classes of galaxy: smaller galaxies where the potential is too shallow to retain the gas expelled by supernova feedback, satellite galaxies where the interstellar gas has been exhausted, and gas-rich galaxies where the dust extinction is strong.  Our simulation roughly reproduces the observed relation between stellar mass and gas-phase metallicity for local star-forming galaxies and for BX galaxies, but the gas-phase metallicities predicted for LBGs, although slightly lower than for BXs, are well above the values estimated for real LBGs by \\cite{Maiolino}.  Most model LBGs, BXs and DRGs are disk-dominated systems, residing at the centres of their own halos. LBGs and BXs are selected as blue galaxies and so exclude the significant population of red galaxies which is already present at these redshifts. Nevertheless, by $z\\sim0$ at least half of their descendants are bulge-dominated and red.  The typical stellar masses of LBGs and BXs increase by an order of magnitude by $z\\sim0$, whereas DRG stellar masses increase by a smaller factor, $\\sim 3$. Star formation is the main driver of growth before $z\\sim 1$, then mergers become dominant.  Most low-redshift massive galaxies ($M_*>10^{11}M_\\odot$) descend from at least one of these high redshift populations.  Correspondingly, most descendants are massive galaxies living in massive dark matter halos today. Many of them are satellite galaxies in galaxy clusters. The central galaxies in rich clusters ($M>10^{14.7}M_\\odot$) typically have 8.4 LBG, BX or DRG progenitors, and on average the stars present in the high-redshift galaxies account for around 50\\% of the current stellar mass.  Thus while the observed high-redshift galaxies have contributed significantly to today's massive galaxies, it appears that the relation between the two populations is quite complex."
    },
    "0809/0809.4129_arXiv.txt": {
        "abstract": "{ The eclipsing brown-dwarf binary system \\mbox{2MASS J05352184-0546085} is a case sui generis. For the first time, it allows a detailed analysis of the individual properties of young brown dwarfs, in particular, masses, and radii, and the temperature ratio of the system components can be determined accurately. The system shows a ``temperature reversal\" with the more massive component being the cooler one, and both components are found to be active. We analyze X-ray images obtained by \\chan~ and \\xmm~ containing \\mbox{2MASS J05352184-0546085} in their respective field of view. The \\chan~ observatory data show a clear X-ray source at the position of \\mbox{2MASS J05352184-0546085}, whereas the \\xmm~ data suffer from contamination from other nearby sources, but are consistent with the \\chan~ detection. No indications of flaring activity are found in either of the observations (together $\\approx 70$~ks), and we thus attribute the observed flux to quiescent emission. With an X-ray luminosity of $3\\times10^{28}$~erg/s we find an \\mbox{$L_X/L_{bol}$}-ratio close to the saturation limit of $10^{-3}$ and an \\mbox{$L_{X}/L_{H\\alpha}$}-ratio consistent with values obtained from low-mass stars. The X-ray detection of \\mbox{2MASS J05352184-0546085} reported here provides additional support for the interpretation of the temperature reversal in terms of magnetically suppressed convection, and suggests that the activity phenomena of young brown dwarfs resemble those of their more massive counterparts. } ",
        "introduction": "Magnetic activity is extremely common among late-type stars with outer convection zones \\citep[e.g.,][]{schmitt_1997}. The cause of that activity as diagnosed by chromospheric and coronal emissions is ultimately thought to be a magnetic dynamo operating near the base of the convection zones of these stars. Near the spectral type $\\sim$~M3, stars are thought to become fully convective, thereby preventing the emergence of a solar like dynamo in their interiors, and yet the activity properties seem to remain unchanged across this borderline \\citep[e.g.][]{fleming_1993}. A saturation level of \\mbox{$L_X/L_{bol} \\sim 10^{-3}$} is observed also for fully convective stars, although with substantial dispersion. Ultra cool dwarfs (spectral type M7 and later) show a drop in activity as measured through their H$_{\\alpha}$ emission, and among L-type dwarfs H$_{\\alpha}$ emission is rarely found \\citep{gizis_2000}. For the majority of active M dwarfs the ratio between H$_{\\alpha}$ and bolometric luminosity is \\mbox{$\\log(L_{H\\alpha}/L_{bol}) \\sim -3.8$}, while for most objects of spectral class~L this ratio drops below $-5$. \\citet{mohatny_2002} argue that the high degree of neutrality in the outer atmospheres of very low-mass stars effectively diminishes the activity level with effective temperature.  X-ray detections from ultra cool dwarfs are quite sparse, and only a few objects have been identified as X-ray sources during flaring and quiescent periods. Examples comprise the binary system \\mbox{GJ 569 B} consisting of two evolved (age $\\approx 500$~Myrs) substellar objects of late-M spectral type detected by \\citet{stelzer_2004} and the latest (with respect to spectral type) object being the M9 dwarf LHS~2065 \\citep[e.g.,][]{schmitt_2002, robrade_2008}. Hence, the behavior of X-ray emission is unclear as one approaches the boundary towards the substellar brown dwarfs and also the role played by the different stellar parameters for the activity evolution of these objects remains unclear \\citep{stelzer_2007}. It is particularly surprising that some of the very rapidly rotating substellar objects show little or no signs of activity. During the very early phases of their evolution, however, when the optical characteristics of brown dwarfs resemble those of stars with spectral type M6 or so, they are often found to be quite active. Starting with the first X-ray detection of \\mbox{Cha H$\\alpha$ 1} in the Chamaeleon~I star forming region \\citep{neuhauser_1998}, quite a number of brown dwarfs settled in star forming regions and (young) clusters have been detected in X-rays. The detection rate is higher for young brown dwarf compared to evolved ones, which might be related to a higher intrinsic X-ray luminosity, hotter outer atmospheres, or both.  A crucial advance in the observational endeavor to unravel the nature of low-mass objects was the discovery of the eclipsing brown-dwarf binary \\citep[``2MASS J05352184-0546085\" - in the following \\bd,][]{stassun_2006}, which allowed to directly measure masses, radii, and a temperature ratio for brown dwarfs for the first time. The \\bd~system is thought to be a member of the Orion Nebular Cluster (ONC) and, hence, very young (age $\\approx 10^{6}$~years). It consists of two \\mbox{$M6.5\\pm0.5$} dwarfs in an eccentric ($e = 0.31$) $9.78$~day orbit, such that the two eclipses are separated by a phase difference of $0.67$.  In their analysis of \\bd~ \\citet{stassun_2006} determine masses of \\mbox{M$_1=56\\pm4$~M$_{Jup}$} and \\mbox{M$_2=36\\pm3$~M$_{Jup}$} as well as radii of \\mbox{R$_1=0.67\\pm0.03$~R$_{\\sun}$} and \\mbox{R$_2=0.49\\pm0.02$~R$_{\\sun}$} for the primary and secondary component, respectively. Astonishingly, a ``temperature reversal\" is observed in this system, i.e., the higher-mass primary component was found to be cooler than its lower-mass companion \\citep[T$_1\\approx 2700$~K and T$_2 \\approx 2800$~K,][]{stassun_2007, reiners_2007}.  From the activity point of view \\bd~is particularly interesting since both components appear to be substantially active, and according to \\citet{reiners_2007} the H$_{\\alpha}$-emission of the primary exceeds that of the secondary by a factor of 7. The same authors find a rotational velocity of about $v\\sin(i)=10$~km/s for the primary, while they give only an upper limit of 5~km/s for the secondary. Assuming a magnetic origin of this activity and the same scaling between \\mbox{H$_{\\alpha}$-emission} and magnetic field strength as found for earlier type stars, \\citet{reiners_2007} argue for the presence of magnetic fields of $Bf\\approx 4$~kG for the primary and $Bf\\approx 2$~kG for the secondary.  In this letter we report the discovery of X-ray emission of the eclipsing brown-dwarf binary \\bd, which provides further support for the presence of magnetic activity in this system. ",
        "conclusions": "In summary, we find an unambiguous X-ray detection of the eclipsing brown-dwarf binary \\bd. An X-ray source at the nominal position of \\bd~ is clearly seen with the \\chan~observatory, and the results are also compatible with an \\xmm~image observed about one year earlier. The X-ray luminosity of \\bd~ is of the order of $3\\times 10^{28}$~erg/s, and even though the \\chan~ light curve might indicate some moderate increase in luminosity, no flaring activity is seen in either of the observations. Therefore, we attribute the flux to quiescent X-ray emission.  An inspection of the source photons recorded by \\chan~shows that the observed plasma emits at a temperature of at least $1$~keV and probably more. The relative hardness of the spectrum is compatible with the results obtained by \\citet{stelzer_2007} for a wide and similarly young brown-dwarf binary system. From the bolometric luminosity of the primary component we compute a value of $-3.4$ for the \\mbox{$\\log(L_X/L_{bol})$}-ratio, which is close to the saturation limit of $\\approx -3$ and within the range of values determined by \\citet{preibisch_2005} for other brown dwarf members of the ONC detected in quiescence as well as compatible with values computed by \\citet{stelzer_2006} for other young brown dwarfs. Invoking the findings of \\citet{reiners_2007} we find \\mbox{$\\log(L_{X}/L_{H\\alpha})$} values well consistent with previously published results for low-mass stars. Unfortunately, none of the optical  eclipses is covered by the available X-ray data, so we have no way at the moment to assign the X-ray flux to the individual components of \\bd .  What can our results tell us about the origin of the temperature reversal? The effect may be caused either by a non-coeval formation of the brown dwarfs or by the impact of magnetic fields, which hamper energy transport by convection. The X-ray detection of \\bd~ alone supports the presence of magnetic fields causing the X-ray emission; further evidence is the hardness of the source photons indicating that magnetic processes are the source of the activity. Our results are, therefore, fully in line with those of \\citet{reiners_2007}, pointing towards the presence of strong magnetic activity in \\bd~ near the saturation limit and hence a magnetic origin of the observed temperature reversal.  \\acknowledgement{This work has made use of data obtained from the \\chan~ and \\xmm~ data archives. S.C. acknowledges support from the DLR under grant 50OR0105. P.C.S. acknowledges support from the DLR under grant 50OR703. }"
    },
    "0809/0809.2539_arXiv.txt": {
        "abstract": "We have constrained the extended (delayed and accelerated) models of hydrogen recombination, by investigating associated changes of the position and the width of the last scattering surface. Using the recent CMB and SDSS data, we find that the recent data constraints favor the accelerated recombination model, though the other models (standard, delayed recombination) are not ruled out at 1-$\\sigma$ confidence level. If the accelerated recombination had actually occurred in our early Universe, baryonic clustering on small-scales is likely to be the cause of it. By comparing the ionization history of baryonic cloud models with that of the best-fit accelerated recombination model, we find that some portion of our early Universe has baryonic underdensity. We have made the forecast on the PLANCK data constraint, which shows that we will be able to rule out the standard or delayed recombination models, if the recombination in our early Universe had proceeded with $\\epsilon_\\alpha\\sim-0.01$ or lower, and residual foregrounds and systematic effects are negligible. ",
        "introduction": "The recombination process of cosmic plasma, which have occurred around the redshift $z_{rec}\\simeq 1100$, decouples photons from baryons \\cite{Modern_Cosmology,Foundations_Cosmology}. In the presence of Ly-$\\alpha$ photon sources, the recombination process might proceed with delay \\cite{Delayed_recombination,Constraint_Recombination,Constraint_Recombination2}, or with acceleration in the presence of baryonic clustering on small-scales \\cite{Ionization_history,Baryonic_recombination,Accelerated_recombination}. The Cosmic Microwave Background (CMB) anisotropy, which is sensitive to the ionization history of our Universe, is affected by delay or acceleration of the recombination process \\cite{Delayed_recombination,Constraint_Recombination,Constraint_Recombination2,Ionization_history,Baryonic_recombination}. The major effect on the CMB anisotropy is expected in the CMB polarization and small angular scale temperature anisotropy \\cite{Constraint_Recombination,Constraint_Recombination2,Ionization_history,Accelerated_recombination}. The five year data of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP5:basic_result,WMAP5:powerspectra,WMAP5:parameter} is released and the recent ground-based CMB observations such as the ACBAR \\cite{ACBAR,ACBAR2008} and QUaD \\cite{QUaD1,QUaD2,QUaD:instrument} provide information complementary to the WMAP data. In near future, PLANCK surveyor \\cite{PLANCK:sensitivity} is going to measure CMB temperature and polarization anisotropy with great accuracy over wide range of angular scales. In this paper, we investigate the recombination models and baryonic cloud models, using the recent CMB and SDSS data.  The outline of this paper is as follows. In Section \\ref{recombination_model}, we investigate the extended recombination models and shows that the data constraints favor accelerated recombination models. In Section \\ref{baryonic_cloud}, we discuss and constrain baryonic cloud models. In Section \\ref{forecast}, we make the forecast on the PLANCK data constraint. In Section \\ref{conclusion}, we summarize our investigation with conclusion. ",
        "conclusions": "\\label{conclusion} We have investigated the extended recombination models, using the recent CMB and SDSS data. We find that the data constraints favor accelerated recombination models, though other recombination models (standard, delayed recombination) are not ruled out at 1$\\sigma$ confidence level. By comparing the ionization history of baryonic cloud models with the best-fit accelerated recombination model, we have constrained the baryonic cloud models, from which we find that our early Universe might have slight overdensity of baryonic matter $\\sim1.02\\,\\bar{\\rho}_{\\mathrm {b}}$ for $\\sim 98 \\%$ of total volume and underdensity of baryonic matter $\\sim 0.04\\,\\bar{\\rho}_{\\mathrm {b}}$  for the rest of space. The origin of primordial baryonic clouds, if exists, might be associated with inhomogeneous baryogenesis  \\cite{inhomogeneous_baryogenesis} in our very early Universe.  While we have constrained baryonic cloud models indirectly by fitting the ionization history, more accurate constraint will be obtained only when CMB power spectra of the baryonic cloud models are fitted directly to the data. Once we get enough evidence of accelerated recombination from the upcoming PLANCK data, we plan to constrain baryonic clouds models directly.  \\ack We are grateful to the anonymous referees for thorough reading and comments, which leads to improvements of this paper in many aspects. We acknowledge the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA), ACBAR 2008 and QUaD 2008 data. This work made use of the \\texttt{CosmoMC} package \\cite{CosmoMC} and \\texttt{RECFAST} \\cite{RECFAST1,RECFAST2,RECFAST3}. The numerical analysis was performed on the HPC cluster facility of the Southern Federal University of the Russia. This work is supported by FNU grants 272-06-0417, 272-07-0528 and 21-04-0355."
    },
    "0809/0809.2013_arXiv.txt": {
        "abstract": "Motivated by brane physics, we consider the non-linear Dirac-Born-Infeld (DBI) extension of the Abelian-Higgs model and study the corresponding cosmic string configurations. The model is defined by a potential term, assumed to be of the mexican hat form, and a DBI action for the kinetic terms. We show that it is a continuous deformation of the Abelian-Higgs model, with a single deformation parameter depending on a dimensionless combination of the scalar coupling constant, the vacuum expectation value of the scalar field at infinity, and the brane tension. By means of numerical calculations, we investigate the profiles of the corresponding DBI-cosmic strings and prove that they have a core which is narrower than that of Abelian-Higgs strings. We also show that the corresponding action is smaller than in the standard case suggesting that their formation could be favoured in brane models. Moreover we show that the DBI-cosmic string solutions are non-pathological everywhere in parameter space. Finally, in the limit in which the DBI model reduces to the Bogomolnyi-Prasad-Sommerfield (BPS) Abelian-Higgs model, we find that DBI cosmic strings are no longer BPS: rather they have positive binding energy. We thus argue that, when they meet, two DBI strings will not bind with the corresponding formation of a junction, and hence that a network of DBI strings is likely to behave as a network of standard cosmic strings. ",
        "introduction": "\\label{sec:introduction}  The (Wilkinson Microwave Anisotropy Probe) WMAP5 results~\\cite{2008arXiv0803.0715G,2008arXiv0803.0570H,2008arXiv0803.0732H, 2008arXiv0803.0593N,2008arXiv0803.0586D,2008arXiv0803.0547K} give strong indication in favour of cosmic inflation over other mechanisms for the production of primordial fluctuations~\\cite{inf25}. Since inflation generally takes place at high energy, recently there has been a flurry of activity in constructing models inspired by or derived from string theory (for recent reviews see \\eg Refs.~\\cite{McAllister:2007bg,Burgess:2007pz,Kallosh:2007ig,HenryTye:2006uv}). In a large category of these models, particularly brane-antibrane inflation~\\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad,Lorenz:2007ze,Lorenz:2008je,Lorenz:2008et,Langlois:2008qf,Langlois:2008wt} and D3/D7 inflation~\\cite{Dasgupta:2002ew,Dasgupta:2004dw}, the end result of the inflationary phase is the creation of D-strings (as well as potentially F-strings~\\cite{Copeland:2003bj}), interpreted from the four-dimensional point of view as cosmic strings.  Since cosmic strings are strongly ruled out as the main originator of primordial fluctuations, the D-(and indeed F-) string tension is severely constrained and (under certain assumptions) such that $G_{_{\\rm N}}\\mu\\lesssim 10^{-6}$ in order to preserve the features of the WMAP5 results~\\cite{Pogosian:2003mz}. Nevertheless, the existence and possible detection of the effects of D-strings in the aftermath of an era of brane inflation could be a testable prediction of string theory.  \\par  D-strings themselves have been conjectured to be in correspondence with the $D$-term strings of supergravity~\\cite{Dvali:2003zh}. One of their remarkable properties is that they satisfy a Bogomolnyi-Prasad-Sommerfield (BPS) condition, \\ie they have no binding energy and preserve 1/2 of the original supersymmetries. They also carry fermionic zero modes and are therefore vorton candidates, leading to possible interesting phenomenological consequences~\\cite{Brax:2006zu}.  \\par  The identification between $D$-term strings and D-strings has been made in the low energy limit, when field gradients are small. Inspired by the case of open string modes which can be effectively described by a non-linear action of the Dirac-Born-Infeld (DBI) type, in this paper we construct models of cosmic strings which depart from the low energy approximation and generalise the Abelian-Higgs model to a non-linear one. We will call the resulting topological objects `DBI-cosmic strings', and they are exact solutions of the generalised non-linear DBI action. The action we consider is very different from others which have been discussed in the literature, Refs.~\\cite{Moreno:1998vy,Yang:2000uj,Sarangi:2007mj,Brihaye:2001ag, Babichev:2006cy,Babichev:2007tn}, and in particular does not lead to pathological configurations. In the limit of small field gradients our DBI strings reduce to Abelian-Higgs strings.  We construct the DBI string solutions numerically in a broad range of parameter space, using two numerical methods: a relaxation method and a shooting algorithm. In this way, we show, in particular, that DBI strings with a potential term corresponding to the BPS limit of the Abelian-Higgs model are no-longer BPS.  More specifically, $\\mu_{2n} \\geq 2\\mu_{n}$, where $\\mu _n $ is the action per unit time and length for a string with a winding number $n$: the equality only holds in the low-energy limit. Borrowing language from the standard cosmic string literature~\\cite{Hindmarsh:1994re,Vilenkin:1994,Rubakov}, the strings are therefore in the type II regime (though the deviations from BPS are small, in a sense we will quantify). The network of strings produced will therefore not contain junctions, and all the strings will have the same tension $\\mu_{n=1}$. In the cosmological context we therefore expect the DBI-string network to evolve in the standard way \\eg~\\cite{Bevis:2007gh}, containing infinite strings and loops, radiating energy through gravitational waves, and eventually reaching a scaling solution.  \\par  The paper is organised as follows. In section~\\ref{subsec:abelianmodel} we recall the properties of Abelian-Higgs cosmic strings, while their realisation in the D3/D7 system is discussed in subsection~\\ref{subsec:d3d7system}. In section~\\ref{sec:dbics}, we first briefly review the different non-linear actions which have been put forward so far in the literature. Then in subsection~\\ref{subsec:dbiaction}, we motivate and present our proposed non-linear action for cosmic strings, which we expect to be applicable when field gradients are large. In section~\\ref{sec:dbisol}, we study the DBI-cosmic strings solutions analytically and numerically. In subsection~\\ref{subsec:analytical}, we present simple analytical estimates which allow us to roughly guess the form of the DBI string profiles and, in subsection~\\ref{subsec:numerics}, we compute them numerically by means of two different methods (shooting and over relaxation). In section~\\ref{sec:conclusions} we briefly summarise our main findings and discuss our conclusions. Finally, the appendix gives the full non-linear structure of the DBI cosmic string action. ",
        "conclusions": "\\label{sec:conclusions}  We have considered a natural DBI generalisation of the Abelian-Higgs model whereby the kinetic terms of the Higgs fields do not lead to a linear differential operator in the equations of motion. The particular form of the action is motivated by a specific extra-dimensional model where the Higgs field becomes a complex direction normal to a D3-brane. Although this model leads to nice cosmic string properties, it is not directly related to a string theory model. As such, the closest model of string theory which could have lead to such a DBI action is the D3/D7 system where BPS cosmic strings are formed at the end of an hybrid-like inflation phase. Unfortunately, the charged fields associated to the open string joining the D3- and D7-branes have no obvious geometric meaning and therefore do not lead to our DBI action. It would certainly be very interesting to see if our construction can be embedded within string theory.  \\par  As a four-dimensional model of non-canonical type, the DBI model of cosmic strings does not suffer from any pathology such as divergences or non-single-valuedness of the field profiles (typical of other non-linear actions which have been proposed in the literature). Indeed we find that DBI strings can be continuously deformed to their Abelian-Higgs analogue. In fact, the main difference from the Abelian-Higgs case appears in the BPS case where the DBI strings show a small departure from the BPS property. In particular, we find that the string tension is reduced, a property which may have some phenomenological significance in order to relax the bound on the string tension coming from Cosmic Microwave Background (CMB) data. Moreover we find that the DBI strings have a positive binding energy implying that the string coalescence is energetically disfavoured, leading to the likely formation of networks with singly-wound strings and statistical properties akin to the usual Abelian-Higgs ones.  \\par  In the present article, we have not tackled some important aspects of DBI string dynamics such as string scattering (for which we expect that the higher order terms discussed in the Appendix may play an important r\\^ole), gravitational back reaction, and the coupling to fermions and their zero modes. This is currently under investigation.  \\par  In summary, we have introduced DBI cosmic strings as non-singular solutions derived from a non-linear Lagrangian. We have studied the solutions numerically and found that they differ significantly from their Abelian-Higgs analogues.  However, the network properties of these strings is almost certainly similar to those of type II Abelian-Higgs cosmic strings."
    },
    "0809/0809.2822_arXiv.txt": {
        "abstract": "We have shown that higher dimensional Reissner-Nordstr\\\"om-de Sitter black holes are gravitationally unstable for large values of the electric charge and cosmological constant in $D \\geq 7$ space-time dimensions. We have found the shape of the slightly perturbed black hole at the threshold point of instability. Why only $D=4, 5$ and $6$ dimensional worlds are favorable as to the black stability remains unknown. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0220_arXiv.txt": {
        "abstract": "Using % recent polarimetric observations of the Crab Nebula in the hard X-ray band by INTEGRAL, we show that the absence of vacuum birefringence effects constrains $O(E/M)$ Lorentz violation in QED to the level $|\\xi| < 9\\times10^{-10}$ at $3\\sigma$ CL, tightening by more than three orders of magnitude previous constraints.  We show that planned X-ray polarimeters have the potential to % probe $\\left|\\xi\\right|\\sim 10^{-16}$ by detecting polarization in active galaxies at red-shift $\\sim1$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2964_arXiv.txt": {
        "abstract": "In this note we consider the attitude of astronomers in Argentina in connection with the new problems posed by relativity theory, before and after General Relativity was presented in its final form. We begin considering, very briefly, the sequence of ``technical'' publications related to relativity that appeared in Argentina and use it to attempt to identify who were the relativity leaders and authors in the Argentina scientific community of the 1910-1920s. Among them there are natives of Argentina, permanent resident scientists, and occasional foreign visitors. They are either academic scientists, or high school teachers; we leave aside the {\\it philosophers} and the {\\it aficionados}. For the main characters we discuss, very briefly again, the scientific facts and publications they handled, the modernity of their information and the ``language'' they use to transmit their ideas to their readers.  Finally, we consider astronomers proper; first Charles Dillon Perrine, an astronomer interested in astrophysics, contracted by the government of Argentina in the USA as director of its main observatory. He became interested in testing the possible deflection of light rays by the Sun towards 1912; his Argentine expedition was the first to attempt that test. Perhaps Perrine was not so much interested in Einstein's formulation of relativity theory, which then was perceived as very far away from his own field of research, as in testing the particular astronomical effects it predicted. In any case, he attempted to do so with the acquiescence and financial support of the Argentine state, and as a leading member of its official scientific elite. We briefly contrast his very specific and strictly scientific efforts with those of our second astronomer, Jos\\'e Ubach, SJ, a secondary school teacher of science at a leading Buenos Aires Catholic school who reported in response to Eddington's expedition. Finally, our third astronomer is F\\'elix Aguilar, a leading scientist with a more definite interest in astrometry, who made an effort to contribute to the public understanding of Einstein's theories in Argentina in 1924, when Einstein's visit to Argentina had become a certainty. ",
        "introduction": "Even some years after 1905, it was only a few authors that discussed advanced dynamics, electron theory and radiation in Argentina. In their works they did not necessarily refer to the 1905 work of Einstein. We follow (Ortiz 1995) to present a list of ``technical'' papers connected with relativity. The main responses are those of Lepiney (1906-8), who makes reference to work of Max Abraham on electron dynamics, and again (Lepiney 1907), now with a general discussion of dynamics with a velocity-dependent mass. Also Broggi (1909) discussed Lorentz's electrodynamics and included a mathematical analysis of the works of Minkowski. Following the early experiments of J. J. Thomson, physicists knew that the motion of an electron was modified in the presence of an electromagnetic field, which could be interpreted as an increment of its mass. Lorentz's aether theory and descriptions of the behaviour of the electron, compatible with Maxwell's equations, were also discussed at the time. All these ideas, as well as Abraham's description of the electron as a perfect sphere with surface charge, were descriptions that agreed with ordinary `common sense'.  In 1910 Vito Volterra delivered a lecture at the Sociedad Cient{\\'\\i}fica Argentina (the Argentine Scientific Society; SCA herein) in Buenos Aires; in (Volterra 1910) he discussed the now well-known paper which started the relativity revolution and, with it, Einstein's {\\it annus mirabilis} (see Gangui 2007). In the decade of 1910 Jakob Johann Laub, a physicist of Polish origin, trained in Germany and hired by the Physics Institute at La Plata, offered lectures and gave courses connected with the theory of relativity. Pyenson (1985), who has studied Laub's personality in detail, has indicated that he was the first physicist to co-author a paper with Einstein, and also suggested that a set of lectures on relativity theory given by him at La Plata may have been the first course on that subject given in the Americas. Once in Argentina Laub translated into Spanish some results obtained in Europe. In (Laub 1912) he discussed briefly optical effects in moving bodies in a paper published in the {\\it Anales} of the SCA. The French physicist Camilo Meyer, a former secondary school companion of Henri Poincar\\'e in France, delivered a series of optional courses on mathematical physics at the University of Buenos Aires in 1910-15; even if he did not specialize in relativity, his courses mentioned recent research in physics, including Kaufmann's experiments on the velocity dependence of the electron mass; he also made reference to Einstein's work without entering into details. Again, between 1916 and 1919, Laub (1916; 1919) considered physical and philosophical questions connected with relativity theory from the point of view of a physicist. In the same years there were also translations or adaptations of general articles that reflected an interest on relativity theory; mostly, they were taken from the foreign press or from popular science journals. ",
        "conclusions": ""
    },
    "0809/0809.2101.txt": {
        "abstract": "We present accurate photometric redshifts in the 2-deg$^2$ COSMOS field. The redshifts are computed with 30 broad, intermediate, and narrow bands covering the UV ({\\it GALEX}), Visible-NIR (Subaru, CFHT, UKIRT and NOAO) and mid-IR ({\\it Spitzer}/IRAC). A $\\chi^2$ template-fitting method ({\\it Le Phare}) was used and calibrated with large spectroscopic samples from VLT-VIMOS and Keck-DEIMOS. We develop and implement a new method which accounts for the contributions from emission lines ([\\ion{O}{2}], H$\\beta$, H$\\alpha$ and Ly$\\alpha$) to the spectral energy distributions (SEDs). The treatment of emission lines improves the photo-z accuracy by a factor of 2.5. Comparison of the derived photo-z with 4148 spectroscopic redshifts (i.e. $\\Delta z = z_{\\rm s} - z_{\\rm p}$) indicates a dispersion of $\\sigma_{\\Delta z/(1+z_{\\rm s})}=0.007$ at $i^+_{\\rm AB}<22.5$, a factor of $2-6$ times more accurate than earlier photo-z in the COSMOS, CFHTLS and COMBO-17 survey fields. At fainter magnitudes $i^+_{\\rm AB}<24$ and $z<1.25$, the accuracy is $\\sigma_{\\Delta z/(1+z_{\\rm s})}=0.012$. The deep NIR and IRAC coverage enables the photo-z to be extended to $z\\sim2$ albeit with a lower accuracy ($\\sigma_{\\Delta z/(1+z_{\\rm s})}=0.06$ at $i^+_{\\rm AB}\\sim 24$). The redshift distribution of large magnitude-selected samples is derived and the median redshift is found to range from $z_{\\rm m}=0.66$ at $22 <i^+_{\\rm AB}<22.5$ to $z_{\\rm m}=1.06$ at $24.5 <i^+_{\\rm AB}<25$. At $i^+_{\\rm AB} <26.0$, the multi-wavelength COSMOS catalog includes approximately 607,617 objects. The COSMOS-30 photo-z enable the full exploitation of this survey for studies of galaxy and large scale structure evolution at high redshift. ",
        "introduction": "Photometric redshifts (hereafter photo-z) are an estimate of galaxy distances based on the observed colors (Baum 1962). This method is extremely efficient for assembling large redshift samples for faint galaxies. Despite having a lower accuracy than spectroscopic redshifts (hereafter, spectro-z), photo-z have the advantage of significantly improved completeness down to a flux limit fainter than the spectroscopic limit.  Deep photo-z samples such as COMBO-17 (Wolf et al. 2003), CFHTLS (Ilbert et al. 2006), SWIRE (Rowan-Robinson et al. 2008), and COSMOS (Mobasher et al. 2007) contain more than 1,000,000 galaxies and go as faint as $i\\sim 25$ with a relatively small amount of telescope time.  Typical photo-z with an accuracy of $\\sigma_{\\Delta z/(1+z_{\\rm s})} \\sim0.02-0.04$ ($\\Delta z = z_{\\rm s}-z_{\\rm p}$) are widely used to study the evolution of galaxy stellar masses and luminosities (e.g. Fontana et al. 2000; Wolf et al. 2003; Gabasch et al. 2004; Caputi et al. 2006; Arnouts et al. 2007), for angular clustering analysis (e.g. Heinis et al. 2007; McCracken et al. 2007), to study the relation between galaxy properties and environment (e.g. Capak et al. 2007), to trace large-scale structures (Mazure et al. 2007; Scoville et al. 2007) and to identify clusters at high redshift (Wang \\& Steinhardt 1998). Photo-z are also necessary for dark energy and dark matter weak lensing studies to separate foreground and background galaxies and to control systematic effects such as intrinsic shape alignment, shear shape correlation, and the effects of source clustering (Peacock et al. 2006). All of these applications require strict control of systematic effects in the photo-z estimate.  An efficient way of identifying and removing systematics is to calibrate photo-z on a spectroscopic sample. The most common methods of calibration are neural network methods (e.g. Ball et al. 2004, Colister et al. 2004), a calibration of the color-z relation (Brodwin et al. 2006; Ilbert et al. 2006), or a reconstruction of the SED templates (Budav\\'ari et al. 2000; Feldmann et al. 2006).  As is also true for spectroscopic redshift measurements, photo-z accuracy depends on spectral coverage and resolution. It is also degraded for sources with a low signal-to-noise ratio (Bolzonella et al. 2000).  The Balmer and Lyman breaks contain much of the photo-z information, so accuracy is lower in redshift ranges where these features are not well sampled by the filter set. As the photometric accuracy degrades, it becomes more difficult to constrain the positions of these features, leading to lower accuracy.  This creates a dual dependency of photo-z accuracy on magnitude and redshift which is defined by the survey design (exposure time, filter choice, and calibration accuracy). For this reason, any photo-z sample should be extensively tested and characterized in the same way a spectroscopic sample would be. This analysis is necessary in order to identify the redshift/magnitude ranges over which the photo-z can be trusted and used for a given scientific application.  This paper presents a new version of the photometric redshifts for the Cosmic Evolution Survey (COSMOS: Scoville et al. 2007) and an analysis of their accuracy. COSMOS is the largest {\\it Hubble Space Telescope (HST)} survey ever undertaken - imaging an equatorial 2-deg$^2$ field to a depth of $I_{\\rm F814W} = 27.8$ mag (5 $\\sigma$, AB). The COSMOS field is equatorial to ensure visibility by all ground-based astronomical facilities. This project includes extensive multi-wavelength imaging from X-ray to radio ({\\it XMM}, {\\it Chandra}, {\\it GALEX}, Subaru, CFHT, UKIRT, {\\it Spitzer}, VLA).  New ground-based NIR data (McCracken et al. 2008; Capak et al. 2008), {\\it Spitzer}-IRAC data (Sanders et al. 2007), and medium/narrow band data from the Subaru Telescope (Taniguchi et al. 2007; Sasaki et al. 2008; Taniguchi et al. 2008; Capak et al. 2008) greatly improve the previous photometry catalogue (Capak et al. 2007). These new data are used for the photo-z derived here, yielding a factor of 3 higher accuracy than the first release of COSMOS photo-z (Mobasher et al. 2007).  The COSMOS data are presented in \\S2. The technique used to estimate the photo-z is presented in \\S3. In \\S4, we quantify the photo-z accuracy as a function of the magnitude and redshift. In \\S5, we provide the photo-z distribution of the $i^+$ selected samples. Specialized photo-z for X-ray selected sources, AGN, and variable objects, are discussed in a companion paper (Salvato et al. 2008).  Throughout this paper, we use the standard WMAP cosmology ($\\Omega_{\\rm m}~=~0.3$, $\\Omega_\\Lambda~=~0.7$) with $h=0.7$ and $h~=~H_{\\rm0}/100$~km~s$^{-1}$~Mpc$^{-1}$.  Magnitudes are given in the AB system. ",
        "conclusions": "\\subsection{Comparison with Other Surveys}\\label{other-surveys}  We now compare the accuracy of the COSMOS-30 photo-z (this work) with those obtained previously for COMBO-17 (Wolf et al. 2004) and CFHTLS-DEEP (Ilbert et al. 2006).  The accuracies can be compared as a function of apparent magnitude using the 1$\\sigma$ measured uncertainties in the photo-z probabilities. Once again, we check also the validity of the 1$\\sigma$ uncertainties for other surveys following the method described in \\S4.2.  For the COMBO-17 and CFHTLS photo-z, we use the spectro-z from the VIMOS-VLT Deep Survey (Le F\\`evre et al. 2004, 2005).  Figure~\\ref{cumule68} shows that about 53\\%, 67\\% of the values for $z_{\\rm s}$ are included inside the 1$\\sigma$ uncertainties for COMBO-17 and CFHTLS-DEEP survey, respectively. The 1$\\sigma$ errors for $z_{\\rm p}$ in COMBO-17 are rescaled by $\\sim$1.2 to obtain 68\\% of the $z_s$ within the $1\\sigma$ error (This rescale is small and changes nothing in our conclusions).  Figure~\\ref{error_Ip} shows the redshift dependence of the 1$\\sigma$ uncertainties in photo-z as a function of magnitude for $0.2<z<1.25$. The $z_{\\rm p}$ for the CFHTLS-DEEP were derived from 5 broad bands ($u^*$, $g^\\prime$, $r^\\prime$ , $i^\\prime$ , $z^\\prime$). The accuracy of the COSMOS photo-z is improved by a factor of 3 when compared to CFHTLS-DEEP. This improvement is largely due to the 12 medium bands in this redshift range (NIR data have a real impact only at $z>1.25$). We checked this by deriving photo-z without the 12 medium bands, obtaining an accuracy of $\\sigma_{\\Delta z/(1+z_{\\rm s})}=0.03$ at $i^+ <22.5$ (similar to the COSMOS release of Mobasher et al. 2007 using 8 broad bands $u^*$, $B_{\\rm J}$, $V_{\\rm J}$, $g^+$, $r^+$, $i^+$, $z^+$ and $K_S$). COMBO-17 includes 12 medium bands in addition to the 5 broad bands and Wolf et al. (2004) achieved an accuracy of $\\sigma_{\\Delta z/(1+z)}=0.02$ at $i^+ \\sim 21.5$. The accuracies of COMBO-17 photo-z are intermediate between those of COSMOS-30 and CFHTLS for $i^+_{AB}<22.5$ but become larger that the CFHTLS accuracies at fainter magnitudes, because the COMBO-17 data are about 1.5{\\ts}mag shallower than the CFHTLS data. Since only secure photo-z (one solution) have been kept in the COMBO-17 catalogue, it explains the flattening of the error bars at $i^+_{AB}>24$.   \\begin{figure}[htb] \\centering \\includegraphics[width=8.cm]{f13.ps} \\caption{Redshift distribution from the 2-deg$^2$ COSMOS survey (black solid line), from the 1-deg$^2$ CFHTLS-D2 field (green dashed line), from the 4-deg$^2$ CFHTLS-DEEP fields (blue dotted line, including also D2). CFHTLS-D2 covers 1 deg$^2$ within the COSMOS field. The red long dashed line is the redshift distribution obtained by Fu et al. (2008) who fit the CFHTLS-DEEP photo-z in the magnitude bin $21.5<i^+ <24.5$ and the dashed-dotted line is obtained without the weight applied for the CFHTLS weak lensing selection (J. Coupon, private communication). \\label{distz_convoluee}} \\end{figure}  \\begin{figure}[htb] \\centering \\includegraphics[width=8.cm]{f14.ps} \\caption{Top panel: Evolution of the redshift distribution as a function of $i^+$ magnitude in the COSMOS field. We use the parametrization of Fu et al. (2008) in different redshift bins. Bottom panel: Cumulative redshift distribution. \\label{distz_cosmos}} \\end{figure}    \\begin{table*} \\begin{tabular}{ c c c c c c c c } \\hline \\\\ mag range        &    $a$       &        $b$     &     $c$  & $A$ & average z & median z\\\\ \\\\ \\hline\\\\ $22.0<i^+_{auto}<22.5$  &    0.497$\\pm$0.019  & 12.643$\\pm$0.409  &  0.381$\\pm$0.016  &     4068.19   &   0.66 &  0.66  \\\\ $22.5<i^+_{auto}<23.0$  &    0.448$\\pm$0.016  &  9.251$\\pm$0.218  &  0.742$\\pm$0.030  &     9151.98   &   0.76 &  0.72  \\\\ $23.0<i^+_{auto}<23.5$  &    0.372$\\pm$0.012  &  6.736$\\pm$0.094  &  1.392$\\pm$0.055  &    18232.24   &   0.90 &  0.82  \\\\ $23.5<i^+_{auto}<24.0$  &    0.273$\\pm$0.008  &  5.281$\\pm$0.039  &  2.614$\\pm$0.096  &    35508.58   &   1.05 &  0.92  \\\\ $24.0<i^+_{auto}<24.5$  &    0.201$\\pm$0.005  &  4.494$\\pm$0.024  &  3.932$\\pm$0.134  &    60306.30   &   1.18 &  1.00  \\\\ $24.5<i^+_{auto}<25.0$  &    0.126$\\pm$0.003  &  4.146$\\pm$0.021  &  5.925$\\pm$0.191  &   103340.04   &   1.25 &  1.06  \\\\ \\hline \\end{tabular} \\caption{Galaxy redshift distribution per deg$^2$. The parameters $a$, $b$, $c$ and $A$ of the eqn.\\ref{NZ} function (Fu et al. 2008) are given per apparent bin.} \\label{fu} \\end{table*}   \\subsection{Redshift Distribution of Galaxies}  Fig.~\\ref{distz_convoluee} shows the galaxy redshift distribution per deg$^2$ from the COSMOS and CFHTLS surveys (for $21.5<i^+_{auto} <24.5$).  Although the surveys are largely independent and the photo-z are computed with different codes, the overall agreement in the redshift distributions is excellent. The agreement is particularly good between COSMOS and CFHTLS-D2 which covers 1 deg$^2$ within the COSMOS field.  The density of galaxies at $z>1.5$ in the CFHTLS fields is of a crucial interest for weak lensing analysis (Benjamin et al. 2007). The CFHTLS redshift distribution presents a bump at $z\\sim 3$, but this excess could be due to misidentifications between the $z<0.4$ and $z>1.5$ photo-z (Van Waerbeke et al. 2008) when no NIR data are available (as is the case for CFHTLS). The COSMOS photo-z are computed with NIR data and such catastrophic failures are limited (see \\S4.1). Fig.~\\ref{distz_convoluee} shows that this bump seems created by a deficit of galaxies at $2.5<z<3$, by comparing CFHTLS-D2 and COSMOS.  Fu et al. (2008) used the CFHTLS photo-z distribution to estimate the matter density parameter $\\Omega_{\\rm m}$ and the amplitude of the matter power spectrum $\\sigma_8$.  The redshift distribution from Fu et al. (2008) at $21.5<i^+ <24.5$ and $z<2.5$ is over-plotted in Fig.~\\ref{distz_convoluee} with and without (J. Coupon, private communication) the weight applied for the weak lensing selection (dashed and dashed-dot lines, respectively). Given the 20\\% variation expected from cosmic variance (Ilbert et al. 2006), it is in excellent agreement with the COSMOS-30 redshift distribution. This agreement suggests that the derivation of $\\sigma_8(\\Omega_{\\rm m}/0.25)^{0.64}=0.785\\pm 0.043$ by Fu et al. (2008) is not suffering from biases due to the photo-z or significant cosmic variance.  The COSMOS-30 redshift distribution can be fit with a parametrization similar to that used by Fu et al. (2008): \\begin{equation}\\label{NZ} n(z) = A \\frac{(z^a+z^{ab})}{(z^b+c)} \\end{equation} where A is the normalization factor and $a$, $b$ and $c$ are free parameters. The best fit parameters $a$, $b$ and $c$ are given in Table~\\ref{fu}, as well as the median redshifts. Fig.~\\ref{distz_cosmos} shows the best fit redshift distribution per apparent magnitude bin.  As expected, the median redshift increases at fainter apparent magnitude, ranging from $z_{\\rm m}=0.66$ at $22<i^+ <22.5$ to $z_{\\rm m}=1.06$ at $24.5<i^+ <25$.     \\begin{table*}[htb] \\begin{tabular}{ l l c c c c c } \\hline \\\\ Sample      &                   &   N   &  median $z_s$  &  median $i^+$  &   $\\sigma_{\\Delta z /(1+z_{\\rm s})}$  &  $\\eta$ in \\% \\\\ \\\\ \\hline\\\\ zCOSMOS bright  & $17.5<i^+<22.5$   &  4146 &      0.48      &    21.6        &       0.007                           &   0.7       \\\\ zCOSMOS faint   & $1.5<z_s<3$       &   147 &      2.20      &    24.0        &       0.054                           &   20.4 \\\\ MIPS bright     & $17.5<i^+<22.5$   &   186 &      0.68      &    21.7        &       0.009                           &   0.0  \\\\ MIPS faint      & $22.5<i^+<24.0$   &   116 &      0.90      &    23.1        &       0.011                           &   0.0   \\\\ MIPS very faint & $24.0<i^+<25.0$   &    15 &      1.15      &    24.2        &       0.053                           &   20.0    \\\\ \\hline \\end{tabular} \\caption{Redshift accuracies estimated from the comparison between photometric redshifts and spectroscopic redshifts.} \\label{sigma} \\end{table*}"
    },
    "0809/0809.0234_arXiv.txt": {
        "abstract": "The $^{92}$Mo($\\alpha,n$)$^{95}$Ru, $^{94}$Mo($\\alpha,n$)$^{97}$Ru, and $^{112}$Sn($\\alpha,\\gamma$)$^{116}$Te cross sections were measured at the upper end of the $p$-process Gamow window between 8.2 MeV and 11.1 MeV. Our results are slightly lower than global Hauser-Feshbach calculations from the code NON-SMOKER, but still within the uncertainty of the prediction. The $^{112}$Sn($\\alpha,\\gamma$)$^{116}$Te cross section agrees well with a recently measured thick-target cross section in the same energy range. For the $^{92,94}$Mo($\\alpha,n$) reactions the present data close to the reaction thresholds could eliminate previous uncertainties within a factor of 20, and we can present now useful comparisons to statistical model calculations with different optical potentials. ",
        "introduction": "A ``$p$ process'' was postulated to produce 35 stable but rare isotopes between $^{74}$Se and $^{196}$Hg on the proton-rich side of the valley of stability. Unlike the remaining 99\\% of the heavy elements beyond iron these isotopes cannot be created by (slow or rapid) neutron captures \\cite{bbfh57}, and their solar and isotopic abundances are 1-2 orders of magnitude lower than the respective $s$- and $r$-process nuclei \\cite{AnG89,iupac}. However, so far it seems to be impossible to reproduce the solar abundances of all $p$ isotopes by one single process. In current understanding, several (independently operating) processes seem to contribute.  The largest fraction of $p$ isotopes is created in the ``$\\gamma$ process'' by sequences of photodissociations and $\\beta^+$ decays \\cite{woho78,woho90,ray90,argo03}. This occurs in explosive O/Ne burning during SNII explosions and reproduces the solar abundances for the bulk of $p$ isotopes within a factor of $\\approx$3 \\cite{ray90,raar95}. The SN shock wave induces temperatures of 2-3 GK in the outer (C, Ne, O) layers, sufficient for triggering the required photodisintegrations. More massive stellar models (M$\\geq$20~M$_\\odot$) seem to reach the required temperatures for efficient photodisintegration already at the end of hydrostatic O/Ne burning \\cite{rahe02}. Historically, the $p$ process was thought to proceed via proton captures but today they are found to play no role since the required amount of free protons is not available in the Ne and O layers within the $p$-process timescale of a few seconds and proton captures are too slow for elements with large $Z$. Therefore, the term \"$\\gamma$ process\" is also often used synonymously because seed nuclei pre-existing in the stellar plasma are photodissociated and more proton-rich isotopes are produced, initially by $(\\gamma,~n)$ reactions. When ($\\gamma,~p$) and ($\\gamma,~\\alpha$) reactions become comparable or faster than neutron emission within an isotopic chain, the reaction path branches out and feeds nuclei with lower charge number $Z$. The decrease in temperature after passage of the shock leads to a freeze-out via neutron captures and mainly $\\beta^+$ decays, resulting in the typical $p$-process abundance pattern with maxima at $^{92}$Mo ($N$=50) and $^{144}$Sm ($N$=82).  However, the $\\gamma$-process scenario suffers from a strong underproduction of the most abundant $p$ isotopes, $^{92,94}$Mo and $^{96,98}$Ru, due to lack of seed nuclei with $A$$>$90. For these missing abundances, alternative processes and sites have been proposed, either related to strong neutrino fluxes in the deepest ejected layers of a SNII ($\\nu p$ process \\cite{FML06}), or rapid proton-captures in proton-rich, hot matter accreted on the surface of a neutron star ($rp$ process \\cite{scha98,scha01}).  Modern, self-consistent studies of the $\\gamma$ process have problems to synthesize $p$ nuclei in the regions $A<124$ and $150\\leq A\\leq 165$ \\cite{rahe02}. It is not yet clear whether the observed underproductions are only due to a problem with astrophysical models or also with the nuclear physics input, i.e. the reaction rates used. Thus, the reduction of uncertainties in nuclear data is strictly necessary for a consistent understanding of the $p$ process. Experimental data can improve the situation in two ways, either by directly replacing predictions with measured cross sections in the relevant energy range or by testing the reliability of predictions at other energies when the relevant energy range is not experimentally accessible.  Accordingly, model calculations of the $p$ process require large nuclear reaction networks involving about 1800 isotopes and more than ten thousand reaction rates. Since the largest fraction of such a network contains proton-rich, unstable isotopes, which are not accessible for cross section measurements with present experimental techniques, almost all of these rates have to be determined theoretically by means of the statistical Hauser-Feshbach (HF) formalism \\cite{hafe52}. Hence, there is no alternative to employing statistical model calculations, e.g. those performed with the codes NON-SMOKER \\cite{rath00,rath01} or MOST \\cite{most05}. Comparison with experimental results has shown that theoretical HF predictions with global parameter sets reproduce the neutron- and proton-induced cross sections within a factor of two or even better. This deviation increases for $\\alpha$-induced reactions. It has been noted that the difficulties in calculating these rates are due to the poor knowledge of the $\\alpha$-optical potential at astrophysically relevant energies \\cite{MRO97,DGG02}. The optical potential governs the transmission coefficient in the HF model. Additional input to this model are nuclear level densities and particle separation energies. The latter data are well defined since nuclear masses in the $p$-process network are experimentally known.  It is only one decade ago that charged particle-induced reaction rates at $p$-process energies ($E_p$$\\approx$1-10~MeV, $E_\\alpha$$\\approx$5-15~MeV) became subject of experimental efforts (for a list of reactions, see \\cite{kadonis}).  For example, a series of $\\alpha$-scattering  experiments at energies below the Coulomb barrier helped to improve the optical $\\alpha$-nucleus potential \\cite{DGG02, AOP03, FGM03, Rau03a, Rau03b}. However, firm conclusions are still difficult to draw due to the lack of experimental cross sections.  Nevertheless, extrapolations to the low astrophysical energies remain necessary since the Coulomb barrier hampers scattering experiments in the astrophysically relevant energy range. Scattering data within a confined energy range can be described equally well by different optical potentials with different extrapolation behavior. This ambiguity, which is called the family problem, makes extrapolations very uncertain \\cite{McS66,fuki96,MRO97,GFG05}. These uncertainties can be avoided if experimental cross sections are available for comparison with HF predictions in or close to the astrophysical Gamow window.  Up to now only five $(\\alpha,\\gamma)$ measurements with relevance for the $p$ process were performed within the Gamow window between 5 and 15~MeV: $^{70}$Ge \\cite{fuki96}, $^{96}$Ru \\cite{rapp02}, $^{106}$Cd \\cite{GKE06}, $^{112}$Sn \\cite{OEG07}, and $^{144}$Sm \\cite{somo98a}. With exception of the latter, the first four experiments agree within a factor of three with the NON-SMOKER prediction using a global parameter set with the McFadden-Satchler $\\alpha$ potential \\cite{McS66}. For the $(\\alpha,n)$ particle-exchange reactions a much larger amount of data exists, but almost all of these measurements are restricted to the upper end of the Gamow window. It was noted in \\cite{RGW06,Rau06} that especially in the heavy mass region (beyond $N$=82) the knowledge of $(\\alpha,x)$ cross sections is crucial, whereas at lower masses the $p$-induced reactions have more importance. However, measurements of $\\alpha$-induced reactions are mainly limited -- apart from some exceptions for $(\\alpha,n)$ cross sections -- to the lower mass region below $^{144}$Sm. In an attempt to constrain the $\\alpha$+nucleus optical potential at low energy, it has been suggested to use ($n,\\alpha$) reactions and some reactions around $A\\approx 147-152$ have been studied in this way \\cite{GKA00,KGA01,KGR04}.  Finally, it has to be emphasized that with a few exceptions ($\\alpha,~\\gamma$) and ($p,~\\gamma$) data in the relevant energy range are more useful for a direct application in reaction networks than ($\\gamma,~\\alpha$) or ($\\gamma,~p$) data despite of the fact that the $p$ process proceeds via photodisintegration reactions. This may seem puzzling at first sight but has two reasons. First, reactions in the stellar plasma involve thermally excited targets and not only targets in the ground state as in the laboratory. Measured rates have to be corrected for this effect. It can be shown that the thermal effects are much more pronounced in the photodisintegration than in the capture channel. Thus the measured rate is closer to the actual rate when considering capture reactions. The reverse rate can then be derived applying detailed balance. Second, in order to avoid inconsistencies, forward and reverse rates for a given reaction in a reaction network have to stem from the same source, related to each other by detailed balance. Because of the exponential factor $\\exp (-Q_\\mathrm{forward})$ in the expression for the relation of reverse to forward rate (see, e.g., \\cite{rath00}), it is preferrable to have data for the direction of positive $Q$ value to limit numerical errors in the simulation. In most cases this is the capture direction, exclusively in the case of proton captures close to stability. However, in the intermediate to heavy mass region the $p$ process also includes $\\alpha$ emitters with negative $\\alpha$ capture $Q$ values.  In this work we study ($\\alpha,n$) reactions on light Mo isotopes and $\\alpha$ capture on $^{112}$Sn, motivated by the need to better understand the above described deficiencies in the $p$ process models. Previous investigations of astrophysical relevance were restricted to $\\alpha$ scattering on $^{92}$Mo \\cite{FGM03}, $^{106}$Cd \\cite{KFG06}, and on $^{112}$Sn \\cite{GFG05} attempting to derive an improved optical potential for these cases. The experimental setup for the measurement of $\\alpha$-induced cross sections is described in Sec.~\\ref{exp}. The resulting cross sections and a comparison with previous experimental results for $^{92}$Mo($\\alpha,n$)$^{95}$Ru, $^{94}$Mo($\\alpha,n$)$^{97}$Ru, and $^{112}$Sn($\\alpha,\\gamma$)$^{116}$Te are given in Sec.~\\ref{results}. The impact of using different optical $\\alpha$+nucleus potentials in theoretical calculations for these reactions is investigated in Sec.~\\ref{sec:alphapot}. ",
        "conclusions": "We have measured the $(\\alpha,n)$ cross sections of the $p$ nuclei $^{92,94}$Mo and the $(\\alpha,\\gamma)$ cross section of the lightest tin isotope $^{112}$Sn. For $^{94}$Mo$(\\alpha,n)$ no data for comparison at lower energies was available, but for $^{92}$Mo our energy trend seems to fit well with the results of Denzler et al. \\cite{DRQ95} starting at 12 MeV. For $^{112}$Sn$(\\alpha,\\gamma)$$^{116}$Te we can essentially reproduce the thick-target cross sections from \\\"Ozkan et al. \\cite{OEG07} although we find slightly lower values above 10 MeV. Additionally, we studied the sensitivity of Hauser-Feshbach predictions performed with the code NON-SMOKER for these reactions. We find that the cross sections and $S$ factors are mostly sensitive to the $\\alpha$ widths and thus to the optical $\\alpha$+nucleus potential employed in the calculations. The standard, global predictions -- utilizing the global potential by McFadden-Satchler \\cite{McS66} -- systematically overestimate the data by a factor of two. The potential of Fr\\\"ohlich and Rauscher \\cite{RF03,RFCor03} describes well the reaction data of $^{112}$Sn$(\\alpha,\\gamma)$$^{116}$Te but underestimates the $^{92,94}$Mo data by a factor of about two. The energy dependence of $^{94}$Mo($\\alpha$,n)$^{97}$Ru cannot be reproduced by the statistical model calculations."
    },
    "0809/0809.0002_arXiv.txt": {
        "abstract": "We present a code for solving the coupled Einstein-hydrodynamics equations to evolve relativistic, self-gravitating fluids.  The Einstein field equations are solved in generalized harmonic coordinates on one grid using pseudospectral methods, while the fluids are evolved on another grid using shock-capturing finite difference or finite volume techniques.  We show that the code accurately evolves equilibrium stars and accretion flows.  Then we simulate an equal-mass nonspinning black hole-neutron star binary, evolving through the final four orbits of inspiral, through the merger, to the final stationary black hole.  The gravitational waveform can be reliably extracted from the simulation. ",
        "introduction": "\\label{intro}  Compact object binaries containing neutron stars [i.e. neutron star-neutron star (NSNS) and black hole-neutron star (BHNS) binaries] are perhaps as important in modern astrophysics as binary black holes. Both BHNS and NSNS binaries should be excellent sources of gravitational waves for the ground-based interferometric detectors LIGO, GEO, VIRGO, and TAMA.  It is possible~\\cite{2008ApJ...676.1162S} that the detection rate of these binaries will actually be greater than that of binary black holes. NSNS and BHNS binaries are also interesting because they are leading candidates for explaining the production of short-duration gamma-ray bursts (GRBs)~\\cite{1992ApJ...395L..83N}, especially given observations that locate some short GRBs in elliptical galaxies~\\cite{2005Natur.437..851G,2005Natur.437..855V} or rule out an associated supernova~\\cite{2005Natur.437..845F}.  NSNS and BHNS mergers may also be important for understanding the observed abundances of the heavy elements that are formed by rapid neutron capture in the r-process~\\cite{1974ApJ...192L.145L}.  NSNS and BHNS binary mergers can be accurately modeled only by numerical simulations.  In such systems, the spacetime metric and neutron star fluid are both dynamical.  They are also strongly coupled, and hence must be evolved simultaneously.  In addition, it is clear from the high compactions of the binary objects that only simulations in full general relativity will be adequate.  For NSNS binaries, such simulations are particularly needed to study the possible collapse of the post-merger remnant.  For BHNS binaries, the crucial questions are 1) whether or not the neutron star is tidally disrupted before it plunges into the black hole, and 2) if the neutron star does disrupt, what fraction of the star's matter is promptly swallowed by the hole, what fraction is ejected, and what fraction forms an accretion disk.  To produce a GRB, a substantial accretion disk must remain.  To contribute to the abundance of r-process elements, matter must be expelled from the system.  Studies strongly suggest that the non-Newtonian form of the gravitational potential~\\cite{2005ApJ...634.1202R}, gravitational radiation reaction~\\cite{2005ApJ...626L..41M}, and black hole spin~\\cite{Rantsiou:2007ct} all significantly affect the size of the post-merger accretion disk, underscoring the need for fully relativistic simulations.  Of the two types of systems, NSNS binaries are better studied.  Newtonian simulations have been performed using realistic equations of state together with neutrino radiation effects~\\cite{1996A&A...311..532R, 1997A&A...319..122R,1999A&A...341..499R,2002MNRAS.334..481R,2008arXiv0806.4380D} and magnetic fields~\\cite{2006Sci...312..719P}.  Simulations using complicated equations of state have also been carried out using the conformally flat approximation to general relativity~\\cite{Oechslin:2006uk}.  General relativistic simulations have been performed using $\\Gamma$-law equations of state~\\cite{2000PhRvD..61f4001S,2002PThPh.107..265S,2003PhRvD..68h4020S, 2004PhRvL..92n1101M,2004PhRvD..69f4026M,Baiotti:2008ra} and using more realistic equations of state~\\cite{2005PhRvD..71h4021S,Shibata:2006nm}.  Recently, general relativistic merger evolutions have been performed that include the neutron star magnetic field~\\cite{2008PhRvL.100s1101A,2008PhRvD..78b4012L}.  Numerical modeling of BHNS binary mergers has been carried out using Newtonian or pseudo-Newtonian gravity (in which the black hole is represented by a point mass) using $\\Gamma$-law~\\cite{1999MNRAS.308..780L,1999ApJ...526..178L} and realistic nuclear~\\cite{1999ApJ...527L..39J,2004MNRAS.351.1121R, 2005ApJ...634.1202R} equations of state.  BHNS simulations have also been done in conformal gravity~\\cite{2006PhRvD..73b4012F}, although so far only in the extreme mass ratio limit, in which the black hole is much more massive than the neutron star.  The first fully relativistic BHNS simulation was of a head-on collision~\\cite{Loffler:2006wa}.   Most recently, two groups have independently evolved configurations of BHNS binaries in full general relativity from quasi-circular inspiral through merger.  One group includes Shibata, Taniguchi, Ury\\={u}, and Yamamoto~\\cite{2007CQGra..24..125S,Shibata:2007zm,Yamamoto:2008js}. The members of the other group are Etienne, Faber, Liu, Shapiro, Taniguchi, and Baumgarte~\\cite{2008PhRvD..77h4002E}.  The evolutions produced by these two groups agree qualitatively, with both now finding very small post-merger disks.  However, both groups have so far restricted themselves to $\\Gamma=2$ polytropic equations of state and to initially non-spinning black holes. The case of spinning black holes has been studied by Rantsiou, Kobayashi, and Laguna~\\cite{Rantsiou:2007ct} using a Kerr background metric and a Newtonian approximation for the neutron star self-gravity.  They find that black hole spin has a strong influence on the size of the post-merger accretion disk.  For both types of binary, but especially BHNS binaries, the parameter space remains poorly explored.  Also, there is an unmet need for results to be checked by multiple independent codes.  In each of the above-mentioned calculations in which the metric variables are truly dynamical fields, these fields were evolved numerically using finite differencing (FD).  For a stable evolution, such algorithms should converge to the exact solution as some power of the grid spacing.  Also, shock-capturing techniques have been developed which allow FD codes to evolve fluids with discontinuities stably and accurately.  FD codes usually require very large grids in order to obtain accurate results, although this problem can be mitigated by using mesh refinement or high-order differencing.  Einstein's equations can also be evolved using spectral methods.  In this case, functions are approximated as truncated series expansions in a set of orthogonal basis functions.  Derivatives of the approximated functions are then computed exactly.  The pseudospectral (PS) method is a type of spectral method that uses the values of functions on a spatial grid of collocation points, rather than directly using the spectral coefficients.  This has the advantage that pointwise operations are as straightforward in PS methods as they are in FD methods.  In {\\it multidomain} PS methods, the computational region is divided into domains, each with its own basis functions and corresponding collocation points.  For smooth functions, spectral methods (including PS) converge {\\it exponentially} to the exact solution as the number of basis functions (or, for PS methods, the number of collocation points) is increased.  This allows PS methods to get accurate results with much smaller grids than those used by FD codes. A PS code for solving the Einstein equations has been developed by the Cornell-Caltech relativity group~\\cite{Kidder:2000yq,Scheel:2002yj, Lindblom:2005qh,Boyle:2006ne}.  It has been used to simulate the inspiral of binary black holes for many orbits with very high accuracy for a fairly low computational cost~\\cite{Scheel:2006gg,Pfeiffer:2007yz,Boyle:2007ft}.  There is a difficulty, however, in extending PS methods to evolving non-vacuum spactimes.  Because of the presence of stellar surfaces and, in some cases, hydrodynamic shocks, the evolved variables are not smooth in all derivatives.  In these cases, spectral representations do not converge exponentially to the exact solution.  Rather, they display Gibbs oscillations near the discontinuity that converge away only like some power of the number of collocation points, with the order of convergence given by the order of the discontinuity.  The oscillations can be controlled using special forms of filtering or artificial viscosity (e.g.~\\cite{t89,d94}) but exponential convergence is still lost.  In some cases, the problem can be avoided by placing domain boundaries at discontinuities, e.g. at the stellar surface.  However, this is not practical for complicated shocks or when stars become very deformed (e.g. during tidal disruption).  In this paper, we use the ``mixed'' approach developed by Dimmelmeier {\\it et al}~\\cite{2005PhRvD..71f4023D} for a conformal gravity code which has been used quite successfully to study supernova core collapse~\\cite{Dimmelmeier:2006mh,Ott:2006eu,Dimmelmeier:2007ui, CerdaDuran:2007cr}, and extend it to full general relativity.  In this method, the spacetime metric is evolved on one grid using PS methods, the fluid variables are evolved on a second grid using shock-capturing FD or finite volume techniques, and the two grids communicate by interpolation. Below, these two grids are referred to as the ``PS grid'' and the ``fluid grid''.  This approach would seem to utilize the strongest features of each method.  The spacetime is expected to be smoother than the fluid variables (e.g. at a stellar surface, the discontinuity in metric components appears at a higher derivative than in the density), and so PS techniques should work better on the metric than the fluid.  Also, our code uses many spectral domains, and we expect discontinuities to appear in only a few domains. In the domains without discontinuities, the functions can be represented spectrally with the same accuracy as in smooth problems.  In the domains with discontinuities, convergence will be limited to a power law.  Of course, this error will propagate to the other domains.  However, the domain decomposition can often be chosen so that the slower converging domains take up a small fraction of the overall computational region, and their effect on the overall accuracy is correspondingly small. We can also use higher resolution in these domains if the error is still too large.  Unlike the strategy of fitting domain boundaries to discontinuities, we do not require the exact locations of discontinuities, but only approximate locations.  Even if the evolution of the fields converges rapidly, spectral accuracy is still lost because of the evolution of the fluids, which will at best converge as a fixed power in the grid spacing on the fluid grid.  However, this is not a problem if the resolution on this grid is high.  And, in fact, our mixed technique algorithm allows us to achieve high resolution in the hydrodynamic evolution at a surprisingly low computational cost.  This is because the fluid grid only needs to cover the region containing the matter. This can provide a huge savings for binary inspiral calculations.  For a BHNS inspiral, the fluid grid can be a box centered on the neutron star with outer boundaries slightly outside the star.  For a NSNS inspiral, one would need two such boxes.  For a BHNS system in which the star is shedding mass onto the black hole, the fluid grid would have to cover a region containing both binary objects, but even then, it would not have to extend all the way out into the wave zone, as it would if the metric were being evolved on the same grid.  In this paper, we describe our code and test its ability to evolve BHNS binaries.  In Section~\\ref{code}, we describe our evolution code.  In Section~\\ref{codetests}, we present tests of this code. Next, we apply our code to model the inspiral and merger of a BHNS binary.  We evolve an equal mass binary with an initially non-spinning black hole and irrotational neutron star.  Section~\\ref{bhnsinspiral} describes the inspiral calculation, with particular emphasis on the accuracy and rate of convergence of these simulations.  Section~\\ref{bhnsmerger} describes the merger.  Finally, Section~\\ref{conclusions} gives our conclusions and future directions for our work. ",
        "conclusions": "\\label{conclusions}  We have shown that our code can evolve inspiraling BHNS binaries with high accuracy at fairly low computational cost.  This enables us to begin simulations at relatively large binary separations.  Long inspirals may turn out to be very important for accurately modeling mergers; merger simulations by the Illinois group show some sensitivity to the initial separation~\\cite{2008PhRvD..77h4002E}, with the implication that starting too close to the merger can lead to underestimating the mass of the post-merger accretion disk.  We have also shown that we can simulate the merger of these binaries, although our accuracy of these simulations is not as good so far than that of our inspirals.  Also, there is every reason to believe that the techniques described here would work just as well for binary neutron stars.  An important next step is to demonstrate that our code can evolve more general BHNS binaries.  We have recently begun evolving such systems with different mass ratios and black hole spins, and so far we have had little difficulty in evolving the inspirals.  We hope to report on these simulations in the near future. Also, we plan to study ways to improve the accuracy of our merger simulations.  It has only been two years since the first fully relativistic BHNS merger simulations were reported, and as yet very little of the parameter space has been studied.  In particular, there is a pressing need to study the effects of the black hole spin and the neutron star equation of state. There are preliminary indications (e.g.~\\cite{Rantsiou:2007ct,Shibata:2007zm}) that both of these have important effects on the post-merger accretion disk mass and the gravitational wave signal. Also, both of the other groups currently performing BHNS merger simulations use very similar techniques (BSSN, moving punctures).  It will be useful to compare their results with those obtained using very different techniques, like the ones reported here."
    },
    "0809/0809.3001_arXiv.txt": {
        "abstract": "{Searches for companions of brown dwarfs by direct imaging mainly probe orbital separations greater than 3--10\\,AU. On the other hand, previous radial velocity surveys of brown dwarfs are mainly sensitive to separations smaller than 0.6\\,AU. It has been speculated that the peak of the separation distribution of brown dwarf binaries lies right in the unprobed range. This work extends high-precision radial velocity surveys of brown dwarfs for the first time out to 3\\,AU. Based on more than six years UVES/VLT spectroscopy the binary frequency of brown dwarfs and (very) low-mass stars (M4.25-M8) in Chamaeleon\\,I was determined: 18$^{+20}_{-12}$\\,\\% for the whole sample % and 10$^{+18}_{-8}$\\,\\% for the subsample of ten brown dwarfs and very low-mass stars (M$\\lesssim 0.1\\,\\msun$). Two spectroscopic binaries were confirmed, the brown dwarf candidate \\chaha8 (previously discovered by Joergens \\& M\\\"uller) and the low-mass star CHXR\\,74. % Since their orbital separations appear to be 1\\,AU or greater, the binary frequency at $<$1\\,AU might be less than 10\\%. Now for the first time companion searches of (young) brown dwarfs cover the whole orbital separation range, and the following observational constraints for models of brown dwarf formation can be derived: (i) the frequency of brown dwarf and very low-mass stellar binaries at $<$3\\,AU does not significantly exceed that at $>$3\\,AU; i.e. direct imaging surveys do not miss a significant fraction of brown dwarf binaries; (ii) the overall binary frequency of brown dwarfs and very low-mass stars is 10-30\\,\\%; (iii) the decline in the separation distribution of brown dwarfs towards smaller separations seems to occur between 1 and 3\\,AU; (iv) the observed continuous decrease in the binary frequency from the stellar to the substellar regime is confirmed at $<$3\\,AU providing further evidence of a continuous formation mechanism from low-mass stars to brown dwarfs. } ",
        "introduction": "Searching for companions to brown dwarfs (BDs) is essential for understanding BD formation, for which no widely accepted model exits (e.g. recent review by Luhman et al. 2007). A priori, one might expect that BDs form in the same manner as stars by the direct collapse and fragmentation of molecular clouds, only on a much smaller scale. Standard cloud fragmentation seems to have difficulty in making BDs (Boss 2001; Bate, Bonnell \\& Bromm  2003). One possible explanation is that the simulations lack an important piece of physics; e.g., supersonic magneto-turbulence can create local high densities and facilitate the formation of BDs (Padoan \\& Nordlund 2004). Alternatively, BDs are formed when cloud fragmentation is modified by an additional process that prematurely halts accretion during the protostellar stage, such as dynamical ejection out of the dense gaseous environment (e.g. Reipurth \\& Clarke 2001; Boss 2001; Bate, Bonnell \\& Bromm 2002; Sterzik \\& Durisen 2003; Umbreit et al. 2005) or photoevaporation by ionizing radiation from nearby massive stars (Kroupa \\& Bouvier 2003; Whitworth \\& Zinnecker 2004). A fourth, recently studied scenario foresees BD formation through instabilities in the outer ($\\gtrsim$100\\,AU) disk around a star and release into the field by e.g. secular perturbation (Goodwin \\& Whitworth 2007; Stamatellos et al. 2007).  The frequency and properties of BDs in multiple systems are fundamental parameters in these models, e.g. embryo-ejection scenarios for BD formation predict only a few BD binaries in preferentially close orbits, while isolated fragmentation scenarios are expected to generate BD binaries with similar (scaled-down) properties, as stellar binaries have. However, the binary properties of BDs are poorly constrained for close separations. Most of the current surveys for companions to BDs and very low-mass stars (VLMS, $M \\leq 0.1\\,\\msun$) are done by direct (adaptive optics or HST) imaging (e.g. Reid et al. 2001; Bouy et al. 2003; Burgasser et al. 2003; Gizis et al. 2003; Close et al. 2003; see also Burgasser et al. 2007) and are not sensitive to either close separations ($a\\lesssim2$--3\\,AU and $a\\lesssim10$--15\\,AU for the field and clusters, respectively) nor to low-mass ratio systems (q$\\equiv$M$_2$/M$_1\\lesssim$0.6). These surveys find a lower frequency (10-30\\%) of BD/VLM binaries compared to stars, a separation distribution with a peak around 3--10\\,AU, and a preference for equal-mass components with 77\\% having a mass ratio q$\\geq$0.8. The observed peak of the separation distribution is close to the incompleteness limit; thus, we do not know if there is a significant number of still undetected, very close ($<$3\\,AU) BD/VLM binaries. It has been suggested that BD/VLM binaries with a separation $<$3\\,AU are as frequent or even more frequent than those at $>$3\\,AU (Pinfield et al 2003; Maxted \\& Jeffries 2005; Burgasser et al. 2007). Hence, the peak of the BD/VLM binary separation distribution could lie below 3\\,AU. Furthermore, while it has been argued that the higher frequency of systems with mass ratios close to unity appears to be a real feature of very low-mass binaries, it would be, nevertheless, desirable to actually probe mass ratios lower than 0.6. Given these limits to the current observational data, the question arises as to whether our current picture of BD/VLM binaries is complete, or, whether we miss a significant fraction of very close and/or low-mass ratio systems.  Spectroscopic monitoring of radial velocity (RV) variations provides a means to detect very close binaries. Given sufficient RV precision, the detection of low-mass ratio systems is feasible down to planetary masses (cf. Joergens 2006a; Joergens \\& M\\\"uller 2007). While several spectroscopic surveys for companions of BD/VLMS in young clusters (Joergens \\& Guenther 2001; Kenyon et al. 2005; Joergens 2006a; Kurosawa, Harries \\& Littlefair 2006; Prato 2007a; Maxted et al. 2008) and in the field (Reid et al. 2002; Guenther \\& Wuchterl 2003; Basri \\& Reiners 2006) have been started in recent years, data sampling is sparse in most cases and the number of confirmed close companions to BD/VLMS is still small. To date, there are four spectroscopic BD/VLM binaries known for which an orbital solution has been determined: \\emph{PPl\\,15} (Basri \\& Martin 1999), \\emph{GJ\\,569Bab} (Zapatero Osorio et al 2004; Simon, Bender \\& Prato 2006), \\emph{2MASS~J05352184-0546085} (Stassun, Mathieu \\& Valenti 2006), and \\emph{\\chaha8} (Joergens \\& M\\\"uller 2007). One of them, \\chaha8, was discovered (Joergens \\& M\\\"uller 2007) within an RV survey for (planetary and BD) companions of very young BD/VLMS in the Chamaeleon\\,I star-forming region, which is being carried out with UVES at the VLT (Joergens \\& Guenther 2001; Joergens 2006a). \\chaha8 is presumably the lowest mass RV companion detected so far in a close orbit around a BD/VLMS and only the second very young BD/VLM spectroscopic binary. That survey in Cha\\,I provided the first determination of the binary frequency of very young BD/VLMS at close separations and the indication that the fraction of undetected BD binaries at close separations must be small. These results were based on probing orbital separations below 0.3\\,AU. The present paper reports follow-up spectroscopic observations and RV measurements of the targets of that survey. This significantly enhances the separation range probed. For the first time the binary frequency of (very young) BD/VLMS at separations of 3\\,AU and smaller is determined.  The paper is organized as follows. Section\\,\\ref{sect:rvs} describes the spectroscopic observations and RV measurements, in Sect.\\,\\ref{sect:RVvari} the identification of RV variable objects is presented, in Sect.\\,\\ref{sect:binfrac} the observed binary fraction is derived and the probed separation range assessed, Sect.\\,\\ref{sect:chxr74} gives details for the spectroscopic system CHXR\\,74, and finally, Sect.\\,\\ref{sect:disc} and Sect.\\,\\ref{sect:concl} provide a discussion and conclusions. ",
        "conclusions": "\\label{sect:concl}  For the first time, the binary frequency of BD/VLMS is probed in the orbital separation range 1--3\\,AU. The separation distribution of BD/VLM binaries detected by direct imaging surveys (e.g. Burgasser et al. 2007) has a peak around 3--10\\,AU, which is close to the incompleteness limit of these surveys ($<$2--3\\,AU for the field and $<$10-15\\,AU for star-forming regions, respectively). It seemed, therefore, possible that the largest fraction of BD/VLM binaries has separations in the range of about 1--3\\,AU and still remained undetected. The BD/VLM binary frequency of 10$^{+18}_{-8}$\\,\\% determined here at separations of 3\\,AU and smaller shows that this is not the case because the rate of spectroscopic binaries at $<$3\\,AU does not significantly exceed the rate of resolved binaries at larger separations (10--30\\%, e.g. Burgasser et al. 2007). Thus, the separation range $<$3\\,AU missed by direct imaging surveys is not adding a significant fraction to the BD/VLM binary frequency, and the overall binary frequency of BD/VLMS is in the range 10-30\\%. The finding of no signs (0/10, i.e. $\\leq$10\\%) for companions at $<$1\\,AU is consistent with a decline in the separation distribution of BD/VLM binaries towards smaller separations starting between 1 and 3\\,AU. Furthermore, the previous finding that the observed mass-dependent decrease in the stellar binary frequency (e.g. Delfosse et al. 2004; Shatsky \\& Tokovinin 2002; Kouwenhoven et al. 2005) continuously extends into the BD regime (e.g. Burgasser et al. 2007) is confirmed here for separations $<$3\\,AU. This is consistent with a continuous formation mechanism from low-mass stars to BDs.  These results are based on a relatively small sample. By combining this work in Cha\\,I with surveys of other young regions, which are restricted to separations $\\leq$0.3\\,AU, an overall observed binary fraction of very young BD/VLMS can be determined with better statistics but with a more limited separation range. In this way, a binary frequency of 7$^{+5}_{-3}$\\% (7/97) is found at $\\leq$0.3\\,AU. In the near future, the sample of the survey in Cha\\,I will be substantially enlarged, allowing the re-investigation of the binary frequency at $<$3\\,AU of BD/VLMS in Cha\\,I with significantly improved statistics for a uniformly obtained data set. Current and future observational efforts are, in addition, directed towards determination of orbital parameters of the detected binary systems, including RV follow-up and high-resolution imaging/astrometry.  \\clearpage \\begin{table*} \\begin{minipage}[t]{\\columnwidth} \\begin{center} \\caption{ \\label{tab:rvs} {\\bf New RV measurements for BDs and VLMSs in Cha\\,I.} } \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{lcccl} \\hline \\hline \\myrule Object               &   Date      & HJD           & RV             & ~$\\sigma_{RV}$\\footnote{Error of the \\emph{relative} RVs. In most cases (marked with an asterisk) this is the sample standard deviation derived from two consecutive measurements.} \\\\ &             &               & [km\\,s$^{-1}$] & [km\\,s$^{-1}$] \\\\ \\hline \\myrule  Cha\\,H$\\alpha$\\,1    & 2006 Jan 16 & 2453751.74004 & 16.206 &  0.39\\,* \\\\  \\hline Cha\\,H$\\alpha$\\,2    & 2006 Jan 16 & 2453751.79554 & 16.319 & 0.25\\,* \\\\  \\hline Cha\\,H$\\alpha$\\,3    & 2006 Jan 15 & 2453750.81659 & 16.339 & 0.70\\,* \\\\  \\hline Cha\\,H$\\alpha$\\,5    & 2006 Jan 15 & 2453750.84774 & 15.697 & 0.03\\,* \\\\  \\hline Cha\\,H$\\alpha$\\,6    & 2006 Jan 16 & 2453751.82994 & 16.630 & 0.07\\,* \\\\  \\hline Cha\\,H$\\alpha$\\,7    & 2006 Jan 15 & wrong pointing \\\\  \\hline Cha\\,H$\\alpha$\\,12   % & 2006 Jan 17 & 2453752.70145 & 15.831 & 0.05\\,* \\\\  \\hline B\\,34                & 2005 Nov 25 & 2453699.83976 & 15.687 & 0.10\\,*  \\\\  \\hline  CHXR\\,74             & 2005 Mar 21 & 2453450.58923  & 18.185 & 0.06\\,* \\\\  & 2006 Apr 10 & 2453835.63456  & 17.982 & 0.10 \\\\  & 2006 Jun 15 & 2453901.52267 & 19.063 & 0.10 \\\\  \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  \\begin{table*} \\begin{minipage}[t]{\\columnwidth} \\begin{center} \\caption{ \\label{tab:chi2prob} {\\bf Characteristics of observations and $\\chi^{2}$ probabilities.}} \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{lccccccc} \\hline \\hline \\myrule Object & $\\Delta$ RV\\footnote{half peak-to-peak RV differences} & $\\overline\\mathrm{RV}_{weighted}$\\footnote{ the error of the weighted mean takes an uncertainty of 400\\,m\\,s$^{-1}$ into account for the zero point of the velocity} & no. obs\\footnote{includes data from Table\\,\\ref{tab:rvs}, Joergens (2006a), and Joergens \\& M\\\"uller (2007)} & Min. $\\Delta$t & Max. $\\Delta$t & \\em{p} & \\\\ & [km\\,s$^{-1}$] & [km\\,s$^{-1}$] & & [days] & [days] &\\\\ \\hline \\myrule  Cha\\,H$\\alpha$\\,1    & 0.24 & 16.27 $\\pm$ 0.52 & 3 & 20 & 2113 & 8.3 $\\times 10^{-1}$ \\\\  \\hline Cha\\,H$\\alpha$\\,2    & 0.15 & 16.25 $\\pm$ 0.49 & 3 & 20 & 2113 & 8.7 $\\times 10^{-1}$ \\\\  \\hline Cha\\,H$\\alpha$\\,3    & 0.99 & 14.86 $\\pm$ 0.89 & 3 & 19 & 2111 & 5.6 $\\times 10^{-2}$ \\\\  \\hline Cha\\,H$\\alpha$\\,4    & 0.21 & 14.85 $\\pm$ 0.43 & 11 & 1 &  701 & 1.0  \\\\  \\hline Cha\\,H$\\alpha$\\,5    & 0.13 & 15.59 $\\pm$ 0.48 & 3 & 19 & 2111 & 8.7 $\\times 10^{-1}$ \\\\  \\hline Cha\\,H$\\alpha$\\,6    & 0.28 & 16.52 $\\pm$ 0.56 & 3 & 19 & 2112 & 6.3 $\\times 10^{-1}$ \\\\  \\hline Cha\\,H$\\alpha$\\,7    & 0.58 & 17.09 $\\pm$ 0.98 & 2 & 19 &   19 & 1.5 $\\times 10^{-1}$ \\\\  \\hline Cha\\,H$\\alpha$\\,8    & 1.38 & 16.62 $\\pm$ 0.67 & 11 & 3 & 2542 & 0.0 & vari. \\\\  \\hline Cha\\,H$\\alpha$\\,12   & 0.96 & 15.29 $\\pm$ 0.93 & 3 & 20 & 2113 & 5.6 $\\times 10^{-3}$ \\\\  \\hline B\\,34                & 0.07 & 15.74 $\\pm$ 0.42 & 6 & 6 & 2083 & 1.0  \\\\  \\hline CHXR\\,74             & 2.39 & 16.71 $\\pm$ 0.81 & 13 & 1 & 2285 & 0.0 & vari. \\\\  \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  \\clearpage \\begin{figure} \\centering \\includegraphics[width=0.6\\linewidth,clip]{0413f1a.eps} \\includegraphics[width=0.6\\linewidth,clip]{0413f1b.eps} \\includegraphics[width=0.6\\linewidth,clip]{0413f1c.eps} \\caption{ \\label{fig:rvs1} {\\bf RV measurements of \\chaha1, \\chaha2, \\chaha3} (top to bottom). } \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.6\\linewidth,clip]{0413f2a.eps} \\includegraphics[width=0.6\\linewidth,clip]{0413f2b.eps} \\includegraphics[width=.6\\linewidth,clip]{0413f2c.eps} \\caption{ \\label{fig:rvs2} {\\bf RV measurements of \\chaha4, \\chaha5, \\chaha6} (top to bottom). } \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=.6\\linewidth,clip]{0413f3a.eps} \\includegraphics[width=.6\\linewidth,clip]{0413f3b.eps} \\caption{ \\label{fig:rvs3} {\\bf RV measurements of \\chaha7 and \\chaha12} (top to bottom). } \\end{figure} \\begin{figure} \\centering \\includegraphics[width=.6\\linewidth,clip]{0413f4a.eps}\\hfill \\includegraphics[width=.6\\linewidth,clip]{0413f4b.eps}\\hfill \\caption{ \\label{fig:rvs4} {\\bf RV measurements of B\\,34 and CHXR\\,74} (top to bottom). } \\end{figure}  \\clearpage \\begin{figure} \\centering \\includegraphics[height=\\linewidth,angle=90]{0413f5.eps} \\caption{ \\label{fig:prob} {\\bf Histogram of the $\\chi^2$ probability $p$} for fitting the observed (relative) RV values of the BD/(V)LMS in Cha\\,I studied here with a constant function. Objects with $p < 10^{-5}$ are shown in the right-most bin, \\chaha8 and CHXR74. Based on the threshold $p < 10^{-3}$, they are classified as binaries. Overplotted is the expected distribution for nonvariable objects given the estimated uncertainties. It shows that the error estimates are reliable. } \\end{figure} \\clearpage \\begin{figure*} \\centering \\includegraphics[width=\\linewidth,angle=0]{0413f6.eps} \\caption{ \\label{fig:detectprob} {\\bf Detection probability as function of semi-major axis.} Based on a Monte Carlo simulation of RV measurements for randomly selected binaries using the measurement errors and time separations of the real observations. From top to bottom: \\chaha7; average of \\chaha1 and \\chaha3; average of \\chaha2, \\chaha4, \\chaha5, \\chaha6, \\chaha8, and \\chaha12; B\\,34; and CHXR\\,74. As indicated, some curves are the average detection probability for two or more objects because observations have similar quality for those. } \\end{figure*} \\clearpage  \\begin{figure*} \\centering \\includegraphics[height=\\linewidth,angle=90]{0413f7.eps} \\caption{ \\label{fig:detectprob_comp} {\\bf Dependence of detection probability on assumption about inclination.} Both curves are calculated for $M_1$=0.07\\,$\\msun$, $\\sigma_{RV}$=0.29\\,km\\,s$^{-1}$, $t_{obs}$=[0,20\\,d,2000\\,d], and $q$ selected randomly from [0.2,1]. Assuming a fixed mean inclination value (dotted line), as in Basri \\& Reiners (2006), slightly overestimates the detection efficiency compared to the realistic case of random orientation (solid line). For the fixed inclination case, the mean value of the inclination distribution of 57.3$^{\\circ}$ is used. } \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[height=\\linewidth,angle=90]{0413f8.eps} \\caption{ \\label{fig:detectprob_comp2} {\\bf Dependence of detection probability on number of observations.} Both curves are calculated for $M_1$=0.07\\,$\\msun$, $\\sigma_{RV}$=0.29\\,km\\,s$^{-1}$, a maximum time span of 2000\\,d, $q$ selected randomly from [0.2,1], and random orientation. An observing schedule based on three observations with $t_{obs}$=[0,20\\,d,2000\\,d] (solid line) provides a superior detection efficiency compared to the case of only two observations separated by 2000\\,d ($t_{obs}$=[0,2000\\,d], dotted line). Thus, third-epoch observations between the observations with the largest time difference can have a significant effect on the probability distribution. } \\end{figure*} \\clearpage \\begin{figure*} \\centering \\includegraphics[width=\\linewidth,angle=0]{0413f9.eps} \\caption{ \\label{fig:vgl_surveys1} {\\bf Comparison of detection probability of different BD/VLMS RV surveys.} All curves were calculated for $q$ selected randomly from [0.2,1], random orientation, variability threshold of $p < 10^{-3}$. Solid (red) line: this work (calculated for $M_1$=0.07\\,$\\msun$, $\\sigma_{RV}$=0.29\\,km\\,s$^{-1}$, $t_{obs}$=[0,20\\,d,2000\\,d]); dotted (black) line: Basri \\& Reiners 2006 ($M_1$=0.15\\,$\\msun$, $\\sigma_{RV}$=1.3\\,km\\,s$^{-1}$, $t_{obs}$=[0,1650\\,d]); dashed (blue) line: Guenther \\& Wuchterl 2003 ($M_1$=0.1\\,$\\msun$, $\\sigma_{RV}$=0.64\\,km\\,s$^{-1}$, $t_{obs}$=[0,31\\,d,85\\,d]); dash-dotted (green) line: Kurosawa et al. 2006 ($M_1$=0.07\\,$\\msun$, $\\sigma_{RV}$=0.42\\,km\\,s$^{-1}$, $t_{obs}$=[0,20\\,d]). See online version for a color figure. } \\end{figure*} \\begin{figure*} \\centering \\includegraphics[width=\\linewidth,angle=0]{0413f10.eps} \\caption{ \\label{fig:vgl_surveys2} {\\bf Same as Fig.\\,\\ref{fig:vgl_surveys1} but with logarithmic y-axis to facilitate comparison with other publications. } See online version for a color figure. } \\end{figure*} \\clearpage \\begin{figure*} \\centering \\includegraphics[width=\\linewidth,angle=0]{0413f11.eps} \\caption{ \\label{fig:vgl_surveys3} {\\bf Comparison of detection probability of different BD/VLMS RV surveys for high-mass ratio binaries.} Same as Fig.\\,\\ref{fig:vgl_surveys1} but for $q$ selected randomly from [0.8,1]. See online version for a color figure. } \\end{figure*} \\begin{figure*} \\centering \\includegraphics[width=\\linewidth,angle=0]{0413f12.eps} \\caption{ \\label{fig:vgl_surveys4} {\\bf Same as Fig.\\,\\ref{fig:vgl_surveys3} but with logarithmic y-axis to facilitate comparison with other publications.} See online version for a color figure. } \\end{figure*}   \\clearpage \\begin{figure} \\centering \\includegraphics[height=\\linewidth,angle=90]{0413f13.eps} \\caption{ \\label{fig:deltaRV} {\\bf RV differences vs. object mass.} Plotted are half peak-to-peak differences in the observed RVs. Each data point is labeled with the corresponding object name, the numbers denote the Cha\\,H$\\alpha$ objects. The RV variable objects (CHXR\\,74, \\chaha8) have $\\Delta$RV $\\gtrsim$ 1.4\\,km\\,s$^{-1}$, while the majority of RV constant objects have $\\Delta$RV $\\leq$0.3--0.6\\,km\\,s$^{-1}$, with the exception of \\chaha3 and \\chaha12, which have $\\Delta$RV $\\sim$1\\,km\\,s$^{-1}$. } \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[height=\\linewidth,angle=90]{0413f14.eps} \\caption{ \\label{fig:deltaRV_K} {\\bf RV differences for BD/VLMS in USco and $\\rho$\\,Oph} based on data from Kurosawa et al. (2006). $\\Delta$RV as in Fig.\\,\\ref{fig:deltaRV}. Objects classified as RV variables by these authors are denoted with their names. The two data points with the largest RV difference correspond to GY310 and USco101, which are classified as RV constant. } \\end{figure}"
    },
    "0809/0809.4417_arXiv.txt": {
        "abstract": "The Gaia satellite will be launched at the end of 2011. It will observe at least 1 billion stars, and among them several million emission line stars and hot stars. Gaia will provide parallaxes for each star and spectra for stars till V magnitude equal to 17. After a general description of Gaia, we present the codes and methods, which are currently developed by our team. They will provide automatically the astrophysical parameters and spectral classification for the hot and emission line stars in the Milky Way and other close local group galaxies such as the Magellanic Clouds. ",
        "introduction": "The Gaia space mission will be launched in 2011/2012. It will orbit at the anti-solar L2 point. Its lifetime is expected to be 5 years. Onboard, there are 3 different instruments:  ASTRO, which will provide astrometric measurements (parallaxes, proper-motions) for all stars down to V magnitude 20 with an accuracy 200 times better than Hipparcos. There are 2 spectrophotometers BP/RP (R$\\sim$100, 320--660nm and 650--1000nm) and a low/medium resolution spectrograph called the Radial Velocity Spectrometer (R=5000 to 11500, 847-874nm, designed for GK stars). This last one will provide spectra for stars till V magnitude equal to 17, and will be used to determine the radial velocity of the stars (see Viala et al. 2008). It is expected that 1 billion of stars will be observed by Gaia, on average 70 times in 5 years (but only 40 times for the RVS). Among them, using the IMF from Kroupa (2001),  it is estimated that there are, at maximum, 68 million of hot stars (HS), and 6 million of emission line stars (ELS).  Due to the huge quantity of data (200 teraBytes by year), all the data-reduction and the scientific analysis must  be performed automatically via new software based  on java programming language. ",
        "conclusions": "The Gaia satellite is one of the most ambicious space mission of the next decades. It will provide some astrophysical parameters for 1 billion of stars and among them for several million of hot stars and emission line stars both in the Milky Way and close local group galaxies. They will allow to elaborate 3D static and dynamic maps, statistical studies, giving new clues about the behaviour of the stars (origin, evolution, variability, binarity). Different methods and algorithms are currently developed by our team in order to detect and identify the lines, to determine the fundamental parameters of the stars, and also to provide their spectral classification."
    },
    "0809/0809.4772.txt": {
        "abstract": " ",
        "introduction": "The accurate theoretical prediction of nuclear reaction rates is of  essential importance in the study of big-bang nucleosynthesis (BBN) when particular reaction rates  are very difficult or even impossible to estimate experimentally. The most typical example of such reactions is the production and destruction of light elements catalyzed by a hypothetical long-lived negatively charged massive ($\\gsim 100$ GeV) leptonic particle, here denoted as $X^-$; a candidate for this particle is the supersymmetric (SUSY) counterpart of the tau lepton  $(\\tau)$, i.e., the stau $(\\tilde{\\tau}).$\\footnote{The stau is   %%% foot  %%% one of the particles expected to be discovered at the CERN Large Hadron Collider \\cite{LHC2007}.} Recently,  BBN involving these $X^-$-{\\it catalyzed} nuclear reactions has been extensively studied from the viewpoint of particle-physics catalysis in BBN \\cite{Pospelov2007,Kohri2006,Kaplinghat,Cyburt2006,Hamaguchi, Bird2007,Kawasaki,Jittoh,Jedamzik2007a,Jedamzik2007b, Kusakabe2007a,Kusakabe2007,Pospelov2007-9Be,Pospelov2008-9Be}, which may be briefly reviewed as follows (and precisely reviewed in Refs.~\\citen{Kusakabe2007} and \\citen{Pospelov2008-9Be} and references therein):  \\parindent 10pt i) If  the  $X^-$ particle has a lifetime on the order of $10^3$ s or longer, it can capture a light nucleus, $A$, previously synthesized during BBN, and produce a bound state denoted  as $(AX^-)$, which is a type of exotic atom with the nucleus $A$ in a Coulombic orbit around the massive $X^-$ at the center.\\footnote{For instance,     %%% foot %%% the binding energies of $(^7{\\rm Be} X^-)$ and  $(\\alpha X^-)$ are  $\\sim$1.3  and  $\\sim$0.34 MeV, and  the rms radii are $\\sim$3.6 and $\\sim$7 fm, respectively.}  \\parindent 10pt ii) When the cosmic temperature cools to $T_9 \\sim 0.4$ (in units of $10^9$~K) after the standard BBN $(T_9 \\gsim 0.5)$, $X^-$  captures $^7$Be to form $(^7{\\rm Be} X^-)$, which is subsequently destroyed by the $(^7{\\rm Be} X^-) + p \\to (^8{\\rm B} X^-) + \\gamma$ reaction\\cite{Bird2007} followed by the $\\beta$-decay of $(^8{\\rm B} X^-)$ to $\\alpha + \\alpha +X^-$.  \\parindent 10pt iii) At $T_9 \\sim 0.3$, $X^-$ captures $^7$Li to generate $(^7{\\rm Li} X^-)$, which is subsequently destroyed by the $(^7{\\rm Li} X^-) + p \\to \\alpha + \\alpha +X^-$ reaction.  \\parindent 10pt iv) At $T_9 \\sim 0.1$, $X^-$ particles  capture $\\alpha$ particles to form  $(\\alpha X^-)$ in abundance which triggers the reaction $(\\alpha X^-) + d \\to~^6{\\rm Li} + X^- $ to produce an enormous amount of $^6$Li, as originally proposed by Pospelov\\cite{Pospelov2007} and confirmed by Hamaguchi {\\it et al.}\\cite{Hamaguchi} including two of the present authors (M.~K. and Y.~K.).  \\parindent  10pt v) At $T_9 \\sim 0.01$, although $X^-$ particles start capturing protons to form $(p X^-)$, its abundance remains low due to the charge-exchange reaction $(p X^-) + \\alpha \\to (\\alpha X^-) + p$, which occurs before $(p X^-)$ begins to interact with other nuclei to form heavier elements.\\cite{Pospelov2008-9Be}  \\parindent 20pt The contribution of these reactions is expected to modify standard BBN (SBBN in short) and to solve the calculated underproduction (by a factor of $\\sim$1000) of the primordial abundance of $^6$Li and the overproduction (by a factor of $\\sim$3) of $^7$Li \\cite{6Li7Li}, and, at the same time, to provide information for constraining the lifetime and primordial abundance of the elementary particle $X^-$. However, full quantum mechanical studies of the reactions have not been performed on the basis of nuclear reaction theory except for the work of Ref.~\\citen{Hamaguchi} on the production of $^6$Li.   The purpose of the present paper is to perform a precise quantum three-body calculation of the cross sections of many types of catalyzed BBN (CBBN in short) reactions including those mentioned above in i) -- v) and to provide their reaction rates for use in the BBN network calculation. The adopted theoretical method and numerical technique have been developed by the present authors, reviewed in Ref.~\\citen{Hiyama2003}, and are well established in the field of few-body atomic and nuclear systems.  The CBBN reactions studied in this paper are classified into four types:  \\noindent 1) {\\it Transfer reactions}: \\begin{eqnarray} (AX^-) + a \\to B + X^- , \\end{eqnarray} \\hskip 0.5cm where $A$ is picked up by an incoming nucleus $a$ to produce a nucleus $B (=A+a)$,  \\noindent 2) {\\it Radiative capture reactions}: \\begin{eqnarray} (AX^-) + a \\to (B X^-) + \\gamma \\, , \\end{eqnarray}  \\noindent 3) {\\it Three-body breakup reactions}: \\begin{eqnarray} (AX^-) + a \\to b_1 + b_2  + X^- , \\end{eqnarray}  \\noindent 4) {\\it Charge-exchange reactions}: \\begin{eqnarray} (AX^-) + a \\to (a X^-) + A . \\end{eqnarray} We solve the Schr\\\"{o}dinger equations of CBBN reactions (1$\\cdot$1)$-$(1$\\cdot$4), explicitly taking the three-body degree of freedom into account, and derive their cross sections and reaction rates at incoming energies below 200  keV.  Throughout this paper, it is not necessary to identify $X^-$ with any particular particle such as the SUSY particle stau, although this is assumed, for instance, in Refs.~\\citen{Cyburt2006}, \\citen{Hamaguchi} and \\citen{Pospelov2008-9Be}. The property of $X^-$ we assume here is that it has a charge of $-e$, a mass $m_X \\gsim 100$ GeV and a lifetime $\\tau_X \\gsim 10^3$ s and that it interacts with light elements via the Coulomb force only.   The structure of the present paper is  as follows. We investigate  $X^-$-catalyzed $\\alpha$-transfer reactions that produce $^6$Li, $^7$Li and $^7$Be in \\S 2, resonant and nonresonant radiative capture reactions that destroy $^7$Be in \\S 3, three-body breakup reactions that destroy $^6$Li and $^7$Li in \\S 4, the charge-exchange reactions in \\S 5 and the possibility of the $X^-$-catalyzed primordial production of $^9$Be in \\S 6. A summary is given in \\S 7.    %\\end{document}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "(1) Using a fully quantum three-body method  developed by the authors\\cite{Hiyama2003}, we have calculated the cross sections of  various types of big-bang nucleosynthesis (BBN) reactions that are catalyzed by a hypothetical long-lived negatively charged, massive leptonic particle (called $X^-$) such as the supersymmetric (SUSY) particle {\\it stau}. We also provided their reaction rates for use in the BBN network calculation. The rates are summarized in Table III and in the last column of Table II (the late-time BBN reactions for $T_9 \\lsim 0.05$).    \\vskip 0.2cm (2) In the catalyzed BBN (CBBN) reactions, the strength of the nuclear interaction (on the order of 10 MeV), which causes the transition between the entrance and exit channels, is very much larger than the incident energy ($\\lsim 200$ keV), and therefore the coupling between the two channels is sufficiently strong to induce multistep transitions between the channels. Therefore, any calculation method that does not take into account the above property of the interaction is not suitable for application to the reactions. Also, for the charge-exchange reactions (5$\\cdot$1)$-$(5$\\cdot$3), a fully quantum, nonperturbative treatment is highly desirable since the $\\alpha$-particle is transferred to low-lying states with the principal quantum number $n= 2 - 3$. We have shown that our three-body calculation method \\cite{Hiyama2003} satisfies the above requirements and is useful for the study of all the reactions.   %%%%%%%%%%%%%%%%%%%%  Table III (Summary)  %%%%%%%%%%%% \\begin{table}[th] \\caption{Summary of the calculated reaction rates of  CBBN reactions obtained by the three-body calculation. a) - c)  are for $T_9 \\lsim 0.2$ and d) - g)  are for $T_9 \\lsim 0.5$. %The rate of the   resonant reaction g)  depends on the % mass of $X^-$ particle. } \\begin{center} \\begin{tabular}{ll} \\hline \\hline \\noalign{\\vskip 0.2 true cm} $\\qquad\\qquad$ Reaction  &  $\\qquad\\qquad$$\\qquad$ Reaction rate (${\\rm cm}^3 \\,{\\rm s}^{-1}\\, {\\rm mol}^{-1}$) \\\\ \\hline \\noalign{\\vskip 0.05 true cm} $\\qquad $ {\\it nonresonant} reaction \\\\ \\noalign{\\vskip -0.15 true cm} a) $(\\alpha X^-) + d \\to~^6{\\rm Li} + X^- $ & $\\quad 2.78  \\times 10^8 \\, T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 5.33 \\,T_9^{-\\frac{1}{3}}) ( 1 - 0.62\\, T_9^\\frac{2}{3} - 0.29\\, T_9 ) $\\\\ % -------------- b) $(\\alpha X^-) + t \\to~^7{\\rm Li} + X^- $ & $\\quad 1.4  \\times 10^7 \\, T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 6.08 \\,T_9^{-\\frac{1}{3}}) (  1 + 1.3\\, T_9^\\frac{2}{3} + 0.55\\, T_9  ) $\\\\ % ------------------------- c) $(\\alpha X^-) +~\\!\\!^3{\\rm He} \\to~^7{\\rm Be} + X^- $ & $\\quad 9.4  \\times 10^7 \\, T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 9.66 \\,T_9^{-\\frac{1}{3}}) \\, (  1 + 0.20\\, T_9^\\frac{2}{3} + 0.05\\, T_9   ) $\\\\ % ----------------------- d) $(^6{\\rm Li} X^-) + p \\to~\\alpha +~\\!\\!^3{\\rm He} + X^- $ & $\\quad 2.6  \\times 10^{10} \\,T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 6.74 \\,T_9^{-\\frac{1}{3}}) $\\\\ % ----------------- e) $(^7{\\rm Li} X^-) + p \\to~\\alpha +~\\!\\!\\alpha + X^- $ & $\\quad 3.5  \\times 10^8 \\, \\, T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 6.74 \\,T_9^{-\\frac{1}{3}}) \\,(1 + 0.81 \\, T_9^\\frac{2}{3} + 0.30\\, T_9 )$\\\\ % ------------------ f) $(^7{\\rm Be} X^-) + p \\to~(^8{\\rm B}X^-)+ \\gamma $ & $\\quad 2.3  \\times 10^5 \\,  T_9^{-\\frac{2}{3}} \\, {\\rm exp}\\,(- 8.83 \\,T_9^{-\\frac{1}{3}}) \\,(  1 + 1.9\\, T_9^\\frac{2}{3} + 0.54\\, T_9   ) $\\\\ % ---------------------- \\noalign{\\vskip 0.20 true cm} \\hline % ---------------------------------------- \\noalign{\\vskip 0.05 true cm} $\\qquad $ {\\it resonant} reaction \\\\ \\noalign{\\vskip -0.15 true cm} g) $(^7{\\rm Be} X^-) + p \\to (^8{\\rm B}X^-)_{2p}^{\\rm res}$ & $\\quad$ $1.37  \\times 10^6 \\, T_9^{-\\frac{3}{2}} \\, {\\rm exp}\\,(- 2.28  T_9^{\\,-1}) \\qquad   m_X=\\;50 {\\rm GeV}$ \\\\ $ \\quad \\: \\to (^8{\\rm B}X^-)+\\gamma $ & $\\quad$ $1.44  \\times 10^6 \\, T_9^{-\\frac{3}{2}} \\, {\\rm exp}\\,(- 2.15  T_9^{\\,-1}) \\qquad   m_X=100 {\\rm GeV}$ \\\\ & $\\quad$ $1.48  \\times 10^6 \\, T_9^{-\\frac{3}{2}} \\, {\\rm exp}\\,(- 2.04 T_9^{\\,-1}) \\qquad   m_X=500 {\\rm GeV}$ \\\\ & $\\quad$ $1.51  \\times 10^6 \\, T_9^{-\\frac{3}{2}} \\, {\\rm exp}\\,(- 2.01  T_9^{\\,-1}) \\qquad   m_X \\to \\infty $ \\\\ \\noalign{\\vskip 0.15 true cm} \\hline\\\\ \\end{tabular} \\end{center} \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   \\vskip 0.2cm (3) In the typical photonless CBBN reactions a)$-$c) in Table III, the calculated  $S$-factors have similar orders of magnitude to those of the typical nonresonant photonless reactions in standard BBN (SBBN). As was first pointed out by Pospelov \\cite{Pospelov2007} and then confirmed by Hamaguchi {\\it et al} \\cite{Hamaguchi}, the $S$-factor of CBBN reaction a) that produces $^6$Li is many orders of magnitude larger than that of the radiative capture  $\\alpha + d \\to\\!^6{\\rm Li}+\\gamma$ in SBBN. This imposes strong restrictions on the lifetime and  abundance of  $X^-$. However, this large enhancement is simply because this SBBN reaction is  heavily E1-hindered. On the other hand, the $S$-factors of CBBN reactions b) and c) are found to be only $\\sim 30$ times larger than those of the strong E1 radiative capture SBBN reactions $\\alpha + t (^3{\\rm He}) \\to\\!^7{\\rm Li}(^7{\\rm Be})+\\gamma$. Therefore,  reactions b) and c) do not seem  to change the abundances of $^7$Li-$^7$Be meaningfully.      \\vskip 0.2cm (4) The resonant radiative capture CBBN reaction, g) in Table III (cf. Fig.~5), was proposed by Bird {\\it et al.} \\cite{Bird2007} as a possible solution to the overproduction of $^7$Li-$^7$Be. We have examined their model and result using a $^7{\\rm Be}+p+X^-$ three-body model, which makes it possible to calculate all the steps in the reaction. We have shown that both the very loosely coupled $^7{\\rm Be}+p$ structure of $^8$B and the strongly spin-dependent $^7{\\rm Be}_{\\frac{3}{2}^-}-p$ interaction contribute largely to the energy of the resonance $(^8{\\rm B}X^-)_{2p}^{\\rm res}$. However, the effects work oppositely, almost cancelling each other, and therefore the reaction rate obtained by the simple model \\cite{Bird2007}(assuming a Gaussian charge distribution of $^8$B) is by chance close to ours. The reaction rate is sensitive to $m_X$; the rate with $m_X=100$ GeV is 60\\% of that with $m_X \\to \\infty$ at $T_9=0.3$.  Regarding the recombination process that forms $(^7{\\rm Be}_{\\frac{3}{2}^-}X^-)$ starting from $^7{\\rm Be}_{\\frac{3}{2}^-}$ and a proton, we pointed out in \\S 3.6 that the calculation\\cite{Bird2007} by Bird {\\it et al.} of the energy of the `resonance' state $(^7{\\rm Be}_{\\frac{1}{2}^-}  X^-)_{2s}$ is erroneous  owing to the lack of consideration of the nuclear structure of  $^7{\\rm Be}_{\\frac{3}{2}^-}$ and  $^7{\\rm Be}_{\\frac{1}{2}^-}$; the $\\alpha +~\\!^3{\\rm He}+X^-$ three-body calculation clarifed that the  $(^7{\\rm Be}_{\\frac{1}{2}^-}  X^-)_{2s}$ state cannot be a resonance but must be a bound state below the $^7{\\rm Be}_{\\frac{3}{2}^-} + X^- $  threshold. In Ref.~\\citen{Bird2007} two types of results were reported on the abundance of $^7$Li-$^7$Be calculated {\\it with} and {\\it without} the $(^7{\\rm Be}_{\\frac{1}{2}^-}  X^-)_{2s}$ resonance, but the case {\\it with} the resonance is not acceptable.   \\vskip 0.2cm (5) We have performed fully quantum mechanical study on the late-time BBN reactions (5$\\cdot$1)$-$(5$\\cdot$8), which take place  at $T_9 \\lsim 0.01$ between the neutral bound states, $(p X^-),$ $(d X^-)$ and $(t X^-)$, and light nuclei.  The atomic charge-exchange reactions (5$\\cdot$1)$-$(5$\\cdot$3) are so dominant that they immediately intercept the nuclear reactions (5$\\cdot$4)$-$(5$\\cdot$8).  Therefore, we conclude that change of element abundances due to the late-time BBN reactions is negligible.  \\vskip 0.2cm (6) We have examined the 'resonant' neutron-capture CBBN process (6$\\cdot$1) for producing $^9$Be proposed by Pospelov and coworkers\\cite{Pospelov2007-9Be,Pospelov2008-9Be}. They claimed that the resulting increased output of $^9$Be strongly constrains the abundance and lifetime of $X^-$. However, we found that, as long as the same model as that in Ref.~\\citen{Pospelov2007-9Be} is employed but with an appropriate charge radius of $^8$Be (a too small value of the radius was used in Ref.~\\citen{Pospelov2007-9Be}), the intermediate state $(^9{\\rm Be}_{\\frac{1}{2}^+} X^-)$ appears {\\it below} the $(^8{\\rm Be} X^-) + n $ threshold, thus giving no resonant mechanism. The model is, however, too simple to account for the change in the structure of the {\\it loosely coupled} (resonant) systems $^8$Be and $^9{\\rm Be}_{\\frac{1}{2}^+}$ due to the injection of the $X^-$ particle into them. Therefore, an extended  four-body calculation based on an $\\alpha+\\alpha+n+X^-$ model  is underway to examine (6$\\cdot$1) much more seriously.  \\vskip 0.2cm (7) Finally, we expect that the application of the presently obtained reaction rates to the BBN  network calculation will help  solve the $^6$Li-$^7$Li problem and  simultaneously  impose restrictions on the primordial abundance and  lifetime of the $X^-$ particle.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0809/0809.0552_arXiv.txt": {
        "abstract": "We report basic far-infrared (FIR) properties of eight blue compact dwarf galaxies (BCDs) observed by AKARI. We measure the fluxes at the four FIS bands (wavelengths of 65 $\\mu$m, 90 $\\mu$m, 140 $\\mu$m, and 160 $\\mu$m). Based on these fluxes, we estimate basic quantities about dust: dust temperature, dust mass, and total FIR luminosity. We find that the typical dust temperature of the BCD sample is systematically higher than that of normal spiral galaxies, although there is a large variety. The interstellar radiation field estimated from the dust temperature ranges up to 100 times of the Galactic value. This confirms the concentrated star-forming activity in BCDs. The star formation rate can be evaluated from the FIR luminosity as 0.01--0.5 $M_\\odot$ yr$^{-1}$. Combining this quantity with gas mass taken from the literature, we estimate the gas consumption timescales (gas mass divided by the star formation rate), which prove to span a wide range from 1 Gyr to 100 Gyr. A natural interpretation of this large variety can be provided by intermittent star formation activity. We finally show the relation between dust-to-gas ratio and metallicity (we utilize our estimate of dust mass, and take other necessary quantities from the literature). There is a positive correlation between dust-to-gas ratio and metallicity as expected from chemical evolution models. ",
        "introduction": "\\label{sec:intro}  Dust grains absorb stellar ultraviolet (UV)--optical light and reprocess it into far-infrared (FIR), thereby affecting the energetics of interstellar medium. In particular, \\citet{hirashita_ferrara02} show that the dust produced by supernovae efficiently absorbs UV heating photons and helps gas cooling and H$_2$ production even in primeval galaxies (typical age is $<10^8$ yr). Dust grains also affect the escape fraction of UV photons out of galaxies. Observationally, this means that there is always a correction factor for dust extinction when one estimates star formation rate (SFR) from UV luminosity \\citep{buat96}. Indeed, a large correction factor for dust extinction is required for the cosmic star formation rate traced with UV \\citep{steidel99,takeuchi05,lefloch05,perez05}. Also theoretically, how dust accumulates in the history of the Universe is important in the context of the cosmic reionization, because dust absorption affects the escape fraction of ionizing photons \\citep{ciardi02}.  The origin of the cosmic dust is important in itself, and the formation and destruction of grains are tightly related to the star formation activity of galaxies. The condensation of heavy elements is an important process for the formation of grains. In particular, in young galaxies, dust is predominantly condensed in Type II supernovae (SNe II), whose progenitors have short lifetimes \\citep{dwek80,kozasa89,morgan02,nozawa03}. However, how much dust forms in a SN II is not known. There is a large difference among observationally derived dust mass \\citep{moseley89,dunne03,hines04,sugerman06,meikle07, ercolano07,rho08}. Thus, it is important to constrain the dust mass formed in SNe II by quantifying dust mass in young galaxies.  A large amount of dust has been suggested to exist at high redshift ($z$) (e.g., \\cite{bertoldi03,priddey03}). The estimate of dust mass in high-$z$ galaxies indeed constrains how much dust forms in SNe II \\citep{dwek07}. The rest-frame UV extinction curves also suggest that dust enrichment by SNe II is indeed occurring at high-$z$ \\citep{maiolino04,bianchi07,hirashita08}. However, it is not easy to explore the first dust enrichment in primeval galaxies at high $z$ ($z\\gtsim 5$) with present observational facilities. Therefore, nearby templates of primeval galaxies are useful to test galaxy formation scenarios. The best candidates for such a template are metal-poor ($\\sim 1/10$--1/3 solar) blue compact dwarf galaxies (BCDs), since they are at the early stage of chemical evolution and their typical stellar age is young \\citep{sargent70}. Indeed BCDs harbor appreciable quantities of dust \\citep{thuan99,plante02,hunt06}, and can be used to investigate the dust properties in chemically unevolved galaxies \\citep{takeuchi_nozawa05}.  The dust content in a galaxy is often estimated by measuring the dust emission in FIR. Because of the large all-sky sample, the Infrared Astronomical Satellite (IRAS) 60~$\\mu$m and 100 $\\mu$m bands are most frequently used. However, \\citet{shibai99}, from the study of FIR spectral energy distribution (SED) of the Milky Way, show that the IRAS bands fail to detect a significant fraction of emission from large grains which contribute to wavelengths ($\\lambda$) longer than 100 $\\mu$m (see also \\cite{okumura02}). Moreover, with the size distribution derived for the Milky Way dust grains, the largest contribution to the dust content comes from large grains ($a\\sim 0.1~\\mu$m, where $a$ is the grain radius) \\citep{mathis77}. Large grains also have an importance that their temperature reflects the stellar radiation field of interstellar medium through the equilibrium between the heating from stellar radiation and the radiative cooling in FIR \\citep{draine85}. In summary, FIR observation of emission from large grains at $\\lambda >100~\\mu$m is important to quantify the dust content and the stellar radiation field.  Indeed, \\citet{galliano05} show that some BCDs have a cold dust component contributing to $\\lambda >100~\\mu$m. \\citet{popescu02} also point out that BCDs tend to have the coldest dust temperature by using a Virgo cluster sample observed by the Infrared Space Observatory (ISO), although the sample size is small. This is unexpected from the warm IRAS 60 $\\mu$m/100 $\\mu$m colors of BCDs \\citep{hoffman89}. \\citet{engelbracht08}, by using the 70 $\\mu$m and 160 $\\mu$m bands of Multiband Imaging Photometer for Spitzer (MIPS) on the Spitzer Space Telescope, show that the dust temperature tends to rise as the metallicity decreases from solar metallicity ($Z_\\odot$) to $\\sim 1/10~Z_\\odot$. However, there is a risk of contamination of stochastically heated dust in the 70~$\\mu$m band \\citep{li01}, which leads to overestimation of the large-grain temperature. Thus, a statistical sample at $\\lambda >100~\\mu$m is indeed required to avoid possible contamination from stochastically heated grains, and simultaneously to overcome the small number statistics of long-wavelength ($\\lambda >100~\\mu$m) observations.  In this paper, we report observations of BCDs in FIR wavelengths including $\\lambda > 100~\\mu$m. The data are obtained by AKARI \\citep{murakami07} Far-Infrared Surveyor (FIS) \\citep{kawada07}. The continuous wavelength coverage ($\\lambda\\sim 50$--180 $\\mu$m) and 4 bands in the FIR provided by FIS enable us not only to quantify the long-wavelength ($\\lambda >100~\\mu$m) dust luminosity but also to separate the contribution from stochastically heated grains by using the short-wavelength ($\\lambda\\ltsim 70~\\mu$m) band. We also demonstrate that this advantage of FIS indeed makes it possible to discuss some quantities related to dust properties and dust enrichment.  This paper is organized as follows. Observations and data reduction are described in section \\ref{sec:bcd}. Then, the basic data are shown in section \\ref{sec:result}, and some quantities derived from those data are presented in section \\ref{sec:derived}. In section \\ref{sec:discussion} we discuss the results, and finally in section \\ref{sec:summary} we summarize this paper. ",
        "conclusions": "\\begin{center} \\begin{tabular}{lcccccc} \\hline Name     & SFR\\footnotemark[$*$] & $M_{\\mathrm{H}\\emissiontype{I}}$ & $\\tau_\\mathrm{ SF}$ & $12+\\log\\mathrm{ (O/H)}$ & ${\\cal D}$ \\\\ & [$M_\\odot$ yr$^{-1}$] & [$M_\\odot$] & [Gyr] & & \\\\ \\hline II Zw 40 & 0.22 & $2.0\\times 10^8$ & 0.91 & 8.15 & $2.0\\times 10^{-4}$ \\\\ Mrk 7   & 0.50 & $3.6\\times 10^9$ & 7.2 & 8.54 & $3.9\\times 10^{-4}$ \\\\ Mrk 71  & 0.012 & $1.2\\times 10^9$ & 110 & 7.83 & $4.5\\times 10^{-6}$ \\\\ UM 439  & 0.021 & $1.7\\times 10^8$ & 8.1 & 7.98 & $8.2\\times 10^{-5}$ \\\\ UM 533  & 0.020 & $5.8\\times 10^7$ & 2.9 & 8.10 & $3.8\\times 10^{-4}$ \\\\ II Zw 70 & 0.073 & $3.6\\times 10^8$ & 4.9 & 8.11 & $2.3\\times 10^{-4}$ \\\\ II Zw 71 & 0.080 & $7.3\\times 10^8$ & 9.1 & 8.24 & $1.4\\times 10^{-3}$ \\\\ Mrk 36  & ---\\footnotemark[$\\dagger$] & $1.8\\times 10^7$ & $<2.5$ & 7.81 & $>4.6\\times 10^{-5}$ \\\\ \\hline \\multicolumn{6}{@{}l@{}}{\\hbox to 0pt{\\parbox{120mm}{\\footnotesize Note. Lower and upper limits are shown by ``$>$'' and ``$<$'', respectively. For Mrk 36, SFR cannot be estimated. \\par\\noindent \\footnotemark[$*$] The SFR is derived by using $L_\\mathrm{ FIR}(\\beta =1).$ If $L_\\mathrm{ FIR}(\\beta =2)$ is used instead, the SFR becomes 4--26\\% higher according to the difference in $L_\\mathrm{ FIR}$. \\par\\noindent \\footnotemark[$\\dagger$] The SFR of Mrk 36 cannot be estimated because $L_\\mathrm{ FIR}$ is not available. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  There are two common BCDs between this paper and \\citet{hopkins02}: Mrk 71 and Mrk 36. For the latter galaxy, we do not obtain the SFR because it is not possible to determine $L_\\mathrm{ FIR}$ (section \\ref{subsec:LFIR}). For Mrk 71, our estimate is 2--3 times lower than their estimate (0.065 $M_\\odot$ yr$^{-1}$ from the IRAS 60 $\\mu$m luminosity and 0.041 $M_\\odot$ yr$^{-1}$ from the 1.4 GHz radio luminosity). This confirms the comment of \\citet{hirashita03} that there is a risk of underestimating the SFR for $\\mathrm{ SFR}\\ltsim 1~M_\\odot$ yr$^{-1}$ if we use equation (\\ref{eq:sfr}). This may be because of smaller dust optical depth.  There are three overlapping BCDs between this paper and \\citet{sage92}: II Zw 40, UM 439, and UM 533. Their SFR derived from the FIR luminosity (0.46, 0.085, and 0.049 $M_\\odot$ yr$^{-1}$, respectively) is systematically higher, which is attributed to different $C_\\mathrm{ IR}$ ($6.5\\times 10^{-9}~M_\\odot~\\mathrm{ yr}^{-1}~L_\\odot^{-1}$). This high $C_\\mathrm{ IR}$ is consistent with the SFR derived from H$\\alpha$ luminosity in dwarf galaxies \\citep{thronson86}, which again confirms the necessity of using higher $C_\\mathrm{ IR}$ than proposed by \\citet{hirashita03} for $\\mathrm{ SFR}\\ltsim 1~M_\\odot$ yr$^{-1}$ . Moreover, the SFR derived from H$\\alpha$ in \\citet{sage92} (3.1, 0.20, and 0.19 $M_\\odot$ yr$^{-1}$, respectively) is much higher than SFR derived from the FIR luminosity, which is again explained by small dust optical depth as mentioned in the previous paragraph. However, the SFR derived from H$\\alpha$ luminosity in \\citet{sage92} could be overestimated because they assume that 2/3 of the ionizing photons are absorbed by dust, which might be less in low-metallicity galaxies like BCDs.  \\subsection{Dust mass}\\label{subsec:Mdust}  As stated in section \\ref{sec:intro}, the \\textit{WIDE-S} ($90~\\mu$m) and \\textit{WIDE-L} ($140~\\mu$m) bands are suitable to trace the total mass of large grains which have the dominant contribution to the total dust content. The total dust mass $M_\\mathrm{d}$ is related to the flux as (e.g., \\cite{hildebrand83}) \\begin{eqnarray} M_\\mathrm{ d}= \\frac{F_\\nu (\\lambda )D^2}{\\kappa_\\nu B_\\nu (T_\\mathrm{ d})} \\, , \\end{eqnarray} where $\\kappa_\\nu$ is the mass absorption coefficient of dust grains. For simplicity, we assume spherical grains with a single radius $a$ and with a uniform density $s$. As shown later, the assumption on $a$ is not essential, since $\\kappa_\\nu$ is insensitive to $a$. The cross section of a grain is expressed as $\\pi a^2Q_\\nu$, where $Q_\\nu$ is a dimensionless quantity expressing the emission efficiency. Using $Q_\\nu$, the mass absorption coefficient is written as \\begin{eqnarray} \\kappa_\\nu =\\frac{3Q_\\nu}{4as}\\, . \\end{eqnarray} Since there is a relation of $Q_\\nu\\propto a$ for $a\\ll\\lambda$ \\citep{evans95}, the estimate of $\\kappa_\\nu$ is insensitive to the assumed value of $a$. \\citet{hildebrand83} estimates $Q_\\nu /a$ at $\\lambda =125~\\mu$m to be 3/400 $\\mu$m$^{-1}$. We note that the normalization of $Q_\\nu /a$ at $\\lambda =125~\\mu$m is convenient since it is roughly the middle of the two wavelengths ($\\lambda =90~\\mu$m and 140 $\\mu$m) used for the estimate of dust mass. If we adopt $s=3$ g cm$^{-3}$, we obtain \\begin{eqnarray} \\kappa_\\nu =18.8\\left(\\frac{\\lambda}{125~\\mu\\mathrm{ m}} \\right)^{-\\beta}~\\mathrm{ cm}^{2}~\\mathrm{ g}^{-1}\\, . \\end{eqnarray} In Table \\ref{tab:derived}, we list the dust mass, where we adopt $T_\\mathrm{ d}([140/90])$ for the dust temperature. The dust mass estimated with $\\beta =1$ and 2 are denoted as $M_\\mathrm{ d}(\\beta =1)$ and $M_\\mathrm{ d}(\\beta =2)$, respectively. The dust mass in Mrk 36 is given by using the $T_\\mathrm{ d}([65/90])$ instead of $T_\\mathrm{ d}([140/90])$, for which only a lower limit is given. Since $T_\\mathrm{ d}([65/90])$ tends to overestimate the equilibrium grain temperature (section \\ref{subsec:Tdust}), we regard this dust mass to be a lower limit. {}From Table \\ref{tab:derived}, we observe that $M_\\mathrm{ d}(\\beta =1)$ is systematically lower than $M_\\mathrm{ d}(\\beta =2)$ by a factor of 2.5. This is because higher dust temperature for $\\beta =1$ leads to a lower dust mass with a fixed dust emissivity.  One of our BCDs, II Zw 40, overlaps with the sample in \\citet{galliano05} (section \\ref{subsec:remark}), who model the dust SED. First of all, the mass of the big grain component, which is contributing to the FIR regime, is just the same as our estimate ($1.1\\times 10^5M_\\odot$) if we adopt $\\beta =2$. This means that the AKARI \\textit{WIDE-S} ($90~\\mu$m) and \\textit{WIDE-L} ($140~\\mu$m) bands are indeed useful to trace the dust mass. However, they also show an additional very cold dust component contributing to the emission at submillimeter wavelengths, although the amount of this component is dependent on the FIR emissivity index $\\beta$. This very cold component cannot be investigated by AKARI, and we concentrate on the dust component contributing to the FIR emission in this paper.    \\subsection{Star formation properties} \\label{subsec:consumption}  In section \\ref{subsec:ISRF}, we have shown that the BCD sample tends to show higher UV radiation fields than spiral galaxies. This is consistent with a picture that intense star formation is occurring in a concentrated region in BCDs. In order to examine the properties of such an intense star formation activity (e.g., a burst or a continuous mode of star formation) it is useful to estimate the gas consumption timescale ($\\tau_\\mathrm{ SF}$) as follows.  Here, we adopt the SFR estimated in section \\ref{subsec:sfr} by adopting $L_\\mathrm{ FIR}(\\beta =1)$. (If we adopt $L_\\mathrm{ FIR}(\\beta =2)$, the gas consumption timescale becomes 4--26\\% shorter.) Then the gas consumption timescale is defined as the gas mass divided by the SFR. For the gas mass, we adopt the mass of H\\emissiontype{I} gas ($M_{\\mathrm{H}\\emissiontype{I}}$) in the same way as in \\citet{hirashita02} and \\citet{hopkins02}, since molecular gas is rarely detected for BCDs \\citep{barone00}. We estimate $M_{\\mathrm{H}\\emissiontype{I}}$ of II Zw 71 by using the H\\emissiontype{I} flux measured by \\citet{huchtmeier89} and by using equation (1) in \\citet{lisenfeld98}. The estimated gas consumption time $\\tau_\\mathrm{ SF}$ is listed in Table \\ref{tab:discussion}. We observe that $\\tau_\\mathrm{ SF}$ is distributed with a range of two orders of magnitude, which is consistent with \\citet{hopkins02}. Moreover, the gas consumption time does not correlate with $G_\\mathrm{ UV}$. For example, Mrk 71 has large $G_\\mathrm{ UV}$, which indicates a concentrated intense star formation, but has large $\\tau_\\mathrm{ SF}$ exceeding the cosmic timescale. Such galaxies with long $\\tau_\\mathrm{ SF}$ but high $G_\\mathrm{ UV}$ give us a picture that the star formation is strongly activated only in a limited region. On the contrary, II Zw 40 has large $G_\\mathrm{ UV}$ and small $\\tau_\\mathrm{ SF}$, indicating the whole galaxy may be in an active phase of star formation.  It is possible that the star formation in BCDs is intermittent. In this case, the diversity in $\\tau_\\mathrm{SF}$ is naturally explained: if a BCD is in activated/inactivated phase of star formation, $\\tau_\\mathrm{SF}$ becomes small/large. Indeed, the effect of feedback is considered to be enhanced in small systems such as BCDs, leading to an intermittent nature of star formation \\citep{saito00,carraro01,kobayashi04,kamaya05}. It is also observationally suggested that the star formation in dwarf galaxies may be episodic \\citep{fanelli88,greggio98,grossi07}.  \\subsection{Dust-to-gas ratio and chemical enrichment} \\label{subsec:dust_metal}  The dust production is strongly related to the metal production. Indeed, the relation between dust abundance (usually dust-to-gas ratio is adopted) and metal abundance (metallicity) can be investigated by using chemical evolution models \\citep{dwek98,hirashita02}. Here we consider whether dust-to-gas ratio and metallicity are related or not.  As the metal abundance, we adopt oxygen abundance, which is the most easily observed. The oxygen abundance of each sample BCD is compiled in \\citet{hirashita02} and \\citet{hopkins02}, but for II Zw 71 we take the oxygen abundance from \\citet{kewley05}. The dust-to-gas ratio ${\\cal D}$ is estimated by $M_\\mathrm{d}(\\beta =1)$ divided by $M_{\\mathrm{H}\\emissiontype{I}}$. If we use $M_\\mathrm{d}(\\beta =2)$ instead, we obtain about 2.5 times higher dust-to-gas ratio. Note that the dust-to-gas ratio would shift further upward if there is really a very cold dust component contributing to longer wavelengths than those traced by AKARI (\\cite{galliano05}; section \\ref{subsec:Mdust}). In Table \\ref{tab:discussion}, we list the metallicity (oxygen abundance: $12+\\log (\\mathrm{O/H})$) and the dust-to-gas ratio ${\\cal D}$. Note that the solar abundance is $12+\\log (\\mathrm{O/H})=8.93$ \\citep{anders89}.  In Figure \\ref{fig:dustmetal}, we show the relation between those two quantities. We observe that there is a positive correlation between dust-to-gas ratio and metallicity. This correlation is also shown by \\citet{engelbracht08} with a larger metallicity range. The correlation is consistent with the picture that dust enrichment proceeds as the system is enriched by metals.  \\begin{figure} \\begin{center} \\FigureFile(85mm,85mm){hirashita_fig2.eps} \\end{center} \\caption{Relation between dust-to-gas ratio ${\\cal D}$ and metallicity. The data of the present BCD sample are shown by solid squares. The arrow indicates a lower limit. The solid, dotted, and dashed lines show the model predictions by \\citet{hirashita02} with various dust destruction efficiencies (in their notation $\\beta_\\mathrm{ SN}=1$, 5, and 25, respectively).} \\label{fig:dustmetal} \\end{figure}  It is also interesting to compare the results with theoretical model predictions. Here we adopt a model by \\citet{hirashita02} (originally developed by \\cite{lisenfeld98}), who consider a one-zone chemical evolution model with the instantaneous recycling approximation \\citep{tinsley80}. The model equations consist of the evolutions of gas mass, metal mass, and dust mass. By treating these equations, the relation between dust-to-gas ratio and metallicity is obtained. In particular, there are two important parameters in the model: (i) the fraction of metals condensed into dust grains, $f_\\mathrm{in}$, and (ii) the ``dust destruction efficiency'', $\\beta_\\mathrm{SN}$. The former parameter is fixed to 0.1, and does not contribute to the scatter of the dust-to-gas ratio if the condensation efficiency in the stellar mass loss is independent of age. The ``dust destruction efficiency'' $\\beta_\\mathrm{ SN}$ is defined as $\\beta_\\mathrm{ SN}\\equiv M_\\mathrm{ g}/(\\tau_\\mathrm{ SN}\\psi)$, where $M_\\mathrm{ g}$ is the total gas mass, $\\tau_\\mathrm{ SN}$ is the timescale of dust destruction by supernova shocks, and $\\psi$ is the SFR. This parameter can be expressed as follows \\citep[equation 6]{hirashita02}: \\begin{eqnarray} \\beta_\\mathrm{SN}=\\epsilon M_\\mathrm{s} \\frac{\\gamma}{\\psi}\\frac{1}{X_{\\rm SF}}\\, , \\end{eqnarray} where $\\epsilon$ is the fraction of destroyed dust in a SN blast, $M_{\\rm s}$ is the gas mass accelerated to a velocity large enough for dust destruction (typically 100 km s$^{-1}$), $\\gamma$ is the supernova rate, and $X_{\\rm SF}$ is the gas mass fraction in the star-forming region, where the dust formation and destruction are really occurring (the rest of the gas is considered to be contained in the H \\textsc{i} envelope, which they assume not to be concerned with the dust formation and destruction). Typical values are $\\epsilon\\sim 0.1$ \\citep{mckee89}, $M_{\\rm s}\\sim 6800M_\\odot$ (if we adopt $10^{51}$ erg for the kinetic energy input of a supernova, and 100 km s$^{-1}$ for the threshold velocity for the dust destruction), $\\gamma/\\psi\\sim 1/136M_\\odot$ (for the Galaxy; \\cite{lisenfeld98}), we obtain $\\beta_\\mathrm{SN}=5$ for $X_{\\rm SF}=1$. \\citet{hirashita02} argue that this parameter changes according to the relative contribution between Type Ia supernovae and Type II supernovae. If an intermittent star formation activity is assumed, $\\gamma/\\psi$ can change by a factor of 20 \\citep{bradamante98,hirashita02}. Motivated by this factor, they changed $\\beta_\\mathrm{ SN}$ from 1 to 25.  In Figure \\ref{fig:dustmetal}, we show the theoretical predictions for $\\beta =1$, 5, and 25 with $f_{\\rm in}=0.1$. We observe that those models explains the data points except for the two extremes (II~Zw~71 and Mrk~71 for the upper and lower extremes in Figure \\ref{fig:dustmetal}, respectively). The variation in $\\beta_\\mathrm{ SN}$ increases the scatter of the dust-to-gas ratio at a certain metallicity and makes the correlation less clear. Further change of $\\beta_{\\rm SN}$ may be possible because of the change of $X_{\\rm SF}$. As mentioned in section \\ref{subsec:consumption}, only a small fraction of gas may be concerned with star formation in Mrk 71. Thus, $X_{\\rm SF}\\ll 1$ is expected for Mrk 71, where $\\beta_{\\rm SN}$ could be larger than the above values. Such a large value of $\\beta_{\\rm SN}$ may explain the extremely poor dust-to-gas ratio of Mrk 71.  As commented above, the dust-to-gas ratios are larger if we adopt $\\beta =2$, and/or if the cold dust component contributing to the submillimeter wavelengths, which we fail to trace in FIR, really exists as proposed by \\citet{galliano05}. In these cases, we should adopt a larger $f_{\\rm in}$, but a large variation of $\\beta_{\\rm SN}$ is still required to explain the large scatter of dust-to-gas ratios.    We have reported far-infrared (FIR) properties of eight blue compact dwarf galaxies (BCDs) observed by AKARI. The fluxes at wavelengths of 65 $\\mu$m, 90 $\\mu$m, 140 $\\mu$m, and 160 $\\mu$m are measured and are used to estimate basic quantities about dust grains such as dust temperature, dust mass, and total FIR luminosity. We have found that the typical dust temperature of our BCD sample, except for Mrk 7 and II Zw 71, is significantly higher than that of normal spiral galaxies. This indicates high interstellar radiation field, which proves to range up to 100 times the Galactic value. This confirms the concentrated star-forming activity in BCDs. The star formation rate (SFR) is also evaluated from the FIR luminosity as 0.01--0.5 $M_\\odot$ yr$^{-1}$. Combining this quantity with gas mass taken from the literature, we estimate the gas consumption timescales (gas mass divided by SFR), which span a wide range from 1 Gyr to 100 Gyr. A natural interpretation of this large variety can be provided by intermittent star formation activity. We finally show the relation between dust-to-gas ratio and metallicity (we utilize our estimate of dust mass, and take other necessary quantities from the literature). There is a positive correlation between dust-to-gas ratio and metallicity as expected from chemical evolution models.  \\vspace{3mm}  We are grateful to S. Madden, the referee, for her helpful comments that improved this paper very much. We thank H. Shibai, Y. Y. Tajiri, M. Nagaoka, Y.~Doi, S. Matsuura, and M. Shirahata for useful discussions about observations and data analysis. We thank all members of AKARI project for their continuous help and support. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
    },
    "0809/0809.2557_arXiv.txt": {
        "abstract": "The extraction of a signal from some observational data sets that contain different contaminant emissions, often at a greater level than the signal itself, is a common problem in Astrophysics and Cosmology. The signal can be recovered, for instance, using a simple Wiener filter. However, in certain cases, additional information may also be available, such as a second observation which correlates to a certain level with the sought signal. In order to improve the quality of the reconstruction, it would be useful to include as well this additional information. Under these circumstances, we have constructed a linear filter, the linear covariance-based filter, that extracts the signal from the data but takes also into account the correlation with the second observation. To illustrate the performance of the method, we present a simple application to reconstruct the so-called Integrated Sachs-Wolfe effect from simulated observations of the Cosmic Microwave Background and of catalogues of galaxies. ",
        "introduction": "\\IEEEPARstart{T}{HE} issues of signal reconstruction and component separation have become of major importance in the field of Astrophysics and Cosmology, covering a large range of applications. In particular, Cosmic Microwave Background (CMB) observations require the development of sophisticated image processing techniques, in most cases on the sphere, to extract all the valuable information encoded in this type of data. The CMB is a relic radiation that was emitted shortly after the Big Bang, when the universe was $\\sim$ 380000 years old and has travelled since then -- almost unhindered -- towards us, carrying a wealth of information about the origin and evolution of the universe (see e.g. \\cite{cha05,mar06} and references therein). From the standard theory of inflation, the statistical distribution of the CMB anisotropies is expected to follow an isotropic and Gaussian random field on the sphere. Different physical phenomena leave their imprint in the CMB radiation that we observe today. In particular, a relevant contribution is given by the so-called Integrated Sachs-Wolfe (ISW) effect \\cite{sac67}, which is due to the evolution in time of the gravitational potential, in the linear regime, of the Large Scale Structure (LSS) and its interaction with the CMB. This effect is of great interest since it provides a direct indication of either the presence of dark energy in the case of a flat universe or the existence of spatial curvature \\cite{pee03}.  Using only CMB observations, it is difficult to obtain a direct detection of the ISW signal, since its contribution to the total CMB signal is, in general, small, and is difficult to disentangle it from other physical effects that also constitute the observed CMB picture. However, given the origin of this effect, a spatial correlation between the ISW and the LSS is expected.  Therefore, it was suggested \\cite{cri96} that the signal of the ISW effect could be statistically detected by looking for cross-correlations between CMB and LSS observations. Following these ideas, different groups have obtained a statistical detection of the ISW effect using the CMB data provided by the WMAP NASA satellite and different LSS surveys at levels between $2.5\\sigma$ and $4.5\\sigma$ \\cite{fos03,bou04,fos04,afs04,nol04,vie06,pie06,mce07,mce08,rac08,gia08}.  Nonetheless, it would be desirable to obtain not only a statistical detection of the ISW signal but also the actual map of this effect. A possible procedure would be to apply a reconstruction image technique (e.g. Wiener filter \\cite{wie49}, maximum-entropy method \\cite{ski89}) to the CMB map in order to extract the ISW effect from the rest of contributions of the primordial radiation. However, given the weak level of the signal such a simple procedure is not expected to provide good results. Also, techniques that make use of multifrequency observations (see \\cite{del08} and references therein) are not useful since the signal has the same frequency dependence as the other dominant CMB contributions. Thus, instead, we propose to recover the ISW map applying a filter, that we have called linear covariance-based (LCB) filter, to the CMB data that takes also into account the information coming from the cross-correlation between the CMB and the LSS surveys.  The outline of the paper is as follows. Section~\\ref{sec:method} presents our method to reconstruct a signal embedded in a noisy background and correlated with a second observation. Section~\\ref{sec:isw} shows the performance of the proposed LCB filter in a simple application, the recovery of the ISW signal from simulated data sets. The robustness of the methodology against certain non-idealities is studied in section~\\ref{sec:robustness}. Finally, we present our conclusions in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a linear filter, the LCB filter, which reconstructs a signal from a confusion background taking into account additional information from a second observation correlated with the sought signal. To study the performance of the technique, we have applied it to simulated CMB and LSS observations with the aim of recovering the ISW effect. This is a very challenging problem, due to the weakness of the ISW signal and, as far as we know, this is the first time that a technique is presented that attempts to recover an ISW map instead of searching for a statistical detection of the signal. We have considered different possibilities for the LSS survey, in order to illustrate the dependence of the method on the characteristics of the chosen catalogue of galaxies. From the considered cases, the best reconstruction (mean $\\rho = 0.85$) is obtained for a survey with a wide bin in $z$ and centred in $z=0.8$. For comparison, we have also recovered the ISW signal from CMB simulated maps using a classical WF, finding that the LCB filter provides better reconstructions. This confirms the idea that combining information from the CMB and the LSS survey is useful to recover the ISW map. We have also shown that our methodology is robust when dealing with incomplete sky observations and against the possible uncertainties in our knowledge of the cosmological model. In a future work we will present a more detailed application to the ISW effect that takes into account more realistic physical conditions for the galaxy populations. Finally, we would like to point out that the formalism can be easily generalised to include information from several galaxy surveys at the same time."
    },
    "0809/0809.0308_arXiv.txt": {
        "abstract": "We apply simple optical and SZ cluster finders to mock galaxy catalogues and SZ flux maps created from dark matter halos in a $(1\\,h^{-1}$Gpc${})^3$ dark matter simulation, at redshifts 0.5 and 0.9. At each redshift, the two catalogues are then combined to assess how well they can improve each other, and compared to several variants of catalogues made using SZ flux and galaxy information simultaneously. We use several different criteria to compare the catalogues, and illustrate some of the tradeoffs which arise in tuning the galaxy cluster finders with respect to these criteria. We detail many of the resulting improvements and issues which arise in comparing and combining these two types of data sets. ",
        "introduction": "In order to study galaxy clusters one must first find, and ideally weigh, them.  Catalogues spanning a range of redshifts, with a well understood selection function, are key to learning about many cluster properties -- for a recent review, see e.g.~\\citet{Voi05}. These include counting clusters as a function of redshift to study dark energy and cosmological parameters, characterizing them as environments for galaxy formation, and for understanding galaxy cluster formation and evolution itself. As galaxy clusters correspond to the largest dark matter halos in numerical simulations, many of the statistical properties of the galaxy cluster population are theoretically accessible.  Indeed, predictions for dark matter halo properties in collisionless N-body simulations are converging \\citep{Evr08,CodeCompare,Luk08,Tin08}. The number and positions of cluster sized halos (their counts and clustering) are of interest for dark energy applications, while for cluster and galaxy properties and evolution these halos comprise the backbone for the more complex and less well understood cluster and galaxy gas physics.  In this note we study mock cluster samples produced by finding overdensities of galaxies or distortions in the cosmic microwave background produced by the upscattering of CMB photons by the hot intra-cluster gas \\citep[][hereafter SZ]{SunZel72,SunZel80}. We compare catalogues found from both methods optimized separately, optimized as a pair, and using one catalogue as input for the second catalogue's search algorithm -- a first attempt at joint optical-SZ cluster finding. We identify trends in the resulting samples, as well as new issues and tradeoffs which arise once SZ and optical data are jointly available. The general features we find can be expected to have bearing on surveys using optical and SZ cluster finding together, such as the (optical) Dark Energy Survey (DES\\footnote{http://www.darkenergysurvey.org}) and the (SZ) South Pole Telescope (SPT\\footnote{http://pole.uchicago.edu}) experiment.  The advantage of combining two cluster finding methods is that each method has different strengths and weaknesses.  For example, SZ cluster finders can successfully find high mass halos (e.g.~$\\geq 2 \\times 10^{14} h^{-1} M_\\odot$) but do not give accurate redshifts\\footnote{Though rough estimates are possible using SZ morphology \\citep{SchPfrZar05}.}, while optical cluster finders can miscast a low mass halo as a high mass one, but are much better at obtaining redshifts (albeit still with possible errors). Both optical and SZ cluster finders suffer from projection effects but possibly in different ways as they trace different properties of the baryonic matter. (Projection has been an issue for optical cluster finding since the earliest surveys, e.g.~\\citet{Abe58,Dal92,Lum92,Whi99}, and has been estimated in SZ by \\citet{WhiHerSpr02,HolMcCBab07,Hal07,ShaHolBod07}.) As the largest objects which have had time to virialize in the universe, ``average'' cluster properties evolve with redshift, and so the results of joining SZ and optical may evolve as well.  The outline of our paper is as follows. We start by describing the simulations and (standard) catalogue/map constructions in \\S 2.  In section \\S 3 we describe the (again standard) optical and SZ cluster finders separately, and the resulting catalogues. As our interest is in general properties, we restrict ourselves to relatively simple cluster finders. Along with the finders, criteria for assessing the success of a finder are introduced and applied to a fiducial optical catalogue, for comparison with the catalogues produced with the other finders in the rest of the paper. In \\S 4, we compare the optical and SZ cluster catalogues to the dark matter halo catalogues, and then introduce a matching between the optical and SZ catalogues directly. We measure how each catalogue is improved by the additional information from the other.  Since the matching is not always one-to-one (one optical cluster to one SZ cluster), the augmented SZ catalogue (with optical information) has different information than the improved optical catalogue (with SZ information).  Lastly, a joint finder based on optical galaxies and SZ flux is introduced, and two other steps in complexity are added, and these cases are compared to the previously discussed ones.  We identify the trends in all of these catalogues as a function of finder methods and parameters using our success criteria. More detail about redshift success (a crucial use of the optical catalogue and/or galaxies for SZ) and whether the SZ and optical information can be combined to veto outliers are found in \\S 5. \\S 6 summarizes, discussing the many improvements due to combining the finders, trends which are expected to generalize beyond our simple finders and catalogues, and issues which have been raised. ",
        "conclusions": "Using mock galaxy catalogs and SZ flux maps made from the $z=0.5$ and 0.9 outputs of a $1\\,h^{-1}$Gpc dark matter simulation, we have made a preliminary investigation of improvements in galaxy cluster finding that arise when optical and SZ surveys are combined.  Our errors, beam, etc. were guided by expected properties for the DES and SPT experiments, though they are quite similar to other experiments in progress or soon to begin. We suggested a set of criteria by which to judge cluster finders and showed the corresponding tradeoffs which arose in some simple cases. Catalogues were compared using the number of massive ($\\geq 2\\times 10^{14}\\,h^{-1}M_\\odot$) halos found, the fraction of clusters whose galaxy content was dominated by one underlying halo, the scatter in the lg(richness)- or lg(flux)-lg(mass) relation and the overcounting of massive halos.  As hoped, using two catalogues jointly or using optical cluster finding centered on SZ flux peaks both led to great improvements in the catalogues. When using the two catalogues together to catch more high mass halos, it was preferable to take a more complete optical catalogue, (but more noisy, i.e. including more blends) and let SZ flux cuts weed out some of the outliers. Using SZ peaks as centers for optical searches and using 1 SZ peak per optical cluster worked similarly in terms of blends and finding massive halos, but worked much better in terms of overcounting. Using a relaxed (i.e.~not merged) SZ catalogue for centering the optical search did well for some criteria but led to a large amount of overcounting.  In the combined catalogues we have two estimates of the cluster mass. Since the scatters in the richness- and flux-mass relations are correlated (both arising in part from projection effects) equality of the two estimates did not indicate a correct assignment of mass. However, taking out $>1 \\sigma$ outliers in SZ mass predictions as a function of optical mass predictions did cut out many of the optical outliers (i.e.~optical clusters whose predicted mass is $>2\\sigma$ from the true mass). There was a trend of SZ patches which had more optical clusters within them to be more likely to have their (optical or SZ) predicted and true mass differ.  We also compared the redshift of the found optical cluster to that of the best match halo to the SZ patch.  The redshift was closer if the optical cluster was centered on the SZ peak, but even in the optical case alone there was a scatter due to the photo-$z$ scatter of the central galaxy, and from the best match halo not being the same as the halo of the central galaxy giving the cluster redshift. Except for completeness, the trends were similar at both redshifts. At high redshift, the SZ effect was more effective at finding the massive halos and the optical finder was worse. Similarly, the optical finder worked better in a $\\sigma_8=0.9$ catalogue than in the ones used here, more evidence for trends of FoF vs.~SO finders with cosmology discussed in \\citet{Luk08}.  These promising results are expected, for the most part, to generalize to more sophisticated finders and catalogues.  Our mock catalogs were deliberately designed to be simple, having periodic edges and no light-cone evolution.  The catalogues underestimate the projection effects, as only those within the box are included and do not include catastrophic photo-$z$ failures. In addition it was assumed that the flux map could be perfectly cleaned of foregrounds, and scatter in flux as a function of mass (due to e.g., cluster history or hydrodynamical effects such as shocks) is not included\\footnote{A model to incorporate some of the effects of cluster history, especially merging, in the SZ signal in cosmological maps has been introduced by \\citet{Bod07,ShaHolBod07}.}.  Correspondingly, our simple cluster finders did not take full advantage of the properties of clusters or cluster galaxies, and could certainly be improved. There are many extensions possible to the simple finders used here. A review of optical cluster finders can be found in \\citet{Gal06}, and SZ cluster finders are improving as well \\citep[e.g.][]{SchWhi03,DelMelBar02,GeiKneHob05,PieAnt05,Pie05,Pir05, Sch06a,Sch06b,MelBarDel06}. In particular, using profile information \\citep[such as used in matched filters, e.g.][]{Pos96,WhiKoc02,Koc03,Don07}, changing the way halos are assigned to SZ peaks, using the peak size, etc., can all refine the cluster information. One extreme, if one is only interested in cluster finding as a route to cosmological parameters, would be to bypass the cluster finding entirely and just cross-correlate the SZ flux and galaxy maps, and compare to mock catalog predictions for the cross-correlation.  Several issues we raised for the simple finders will become more pressing with more sophisticated finders -- cluster finders with increased complexity depend more and more upon the expected properties of the targets. In particular, finders can succeed or fail in several different ways and choices must be made about which failures/successes one cares most about. Specifically, as finders will be increasingly tuned and trained on mock catalogues, and no finder will be perfect, the question arises as to which errors are preferable. Ideally the errors should not only be as small as possible, but also well modeled in the mock catalogues. It is preferable to have a finder which relies most heavily upon the most accurate features of the mock catalogues at hand. This is especially important as the finders become more complex and assumed cluster properties become more and more implicit. In addition, there is a question of how the scatter in the finders depends on the target cluster masses, for example higher mass lowers the importance of shot noise errors on the optical richness and is more likely to correspond to a significant SZ signal.  Our focus was on massive ($\\geq 2\\times 10^{14}\\,h^{-1}M_\\odot$) halos, as the SZ methods cannot probe to very low masses, but the particular science goals may argue for a different threshold. The tradeoffs chosen in creating the finder should in part be decided by the intended use of the sample. Usually, the more complex the finder becomes the more tradeoffs are necessary in desirable catalogue features. These decisions are best made with an eye the strengths, features and science goals of the particular observational data set in hand, and the strengths of the mock catalogues available to analyze the data. As such multi-method cluster observations increase, the exciting science within reach will also increase dramatically.   JDC thanks C. Heymans, J. Kollmeier, A. Leathaud, A. Meiksin, R. Nichol, P. Norberg, W. Percival, C. Pfrommer and E. Rozo for discussions.  We would like to thank U.~Portsmouth, the ROE and IUCAA for hospitality and the opportunity to present this work while in progress and completed.  We thank G. Evrard for comments on the draft and the anonymous referee for very helpful comments on the submitted paper. JDC acknowledges NSF-AST-0810820 and DOE for additional support. MW is supported by NASA and the DOE. The simulations used in this paper were performed and analyzed at the National Energy Research Scientific Computing Center.  \\citet{Men08} appeared when we were preparing this paper for publication, which discusses optical clusters found in the Southern Cosmology Survey and their expected SZ signal from ACT."
    },
    "0809/0809.2411_arXiv.txt": {
        "abstract": "We use high-quality Very Large Array (VLA) images of the Fanaroff \\& Riley Class I radio galaxy 3C\\,31 at six frequencies in the range 1365 to 8440\\,MHz to explore the spatial scale and origin of the rotation measure (RM) fluctuations on the line of sight to the radio source.  We analyse the distribution of the degree of polarization to show that the large depolarization asymmetry between the North and South sides of the source seen in earlier work largely disappears as the resolution is increased.  We show that the depolarization seen at low resolution results primarily from unresolved gradients in a Faraday screen in front of the synchrotron-emitting plasma. We establish that the residual degree of polarization in the short-wavelength limit should follow a Burn law and we fit such a law to our data to estimate the residual depolarization at high resolution. We discuss how to interpret the structure function of RM fluctuations in the presence of a finite observing beam and how to address the effects of incomplete sampling of RM distribution using a Monte Carlo approach. We infer that the observed RM variations over selected areas of 3C\\,31, and the small residual depolarization found at high resolution, are consistent with a power spectrum of magnetic fluctuations in front of 3C\\,31 whose power-law slope changes significantly on the scales sampled by our data. The power spectrum $P(f)$ can only have the form expected for Kolmogorov turbulence [$P(f) \\propto f^{-11/3}$] on scales \\la 5\\,kpc. On larger scales we find $P(f)$ \\pa $f^{-2.3}$. We briefly discuss the physical interpretation of these results.  We also compare the global variations of RM across 3C\\,31 with the results of three-dimensional simulations of the magnetic-field fluctuations in the surrounding magnetoionic medium.  We infer that the RM variation across 3C\\,31 is qualitatively as expected from relativistic-jet models of the brightness asymmetry wherein the apparently brighter jet is on the near side of the nucleus and is seen through less magnetoionic material than the fainter jet. We show that our data are inconsistent with observing 3C\\,31 through a spherically-symmetric magnetoionic medium, but that they are consistent with a field distribution that favors the plane perpendicular to the jet axis -- probably because the radio source has evacuated a large cavity in the surrounding medium.  We also apply our analysis techniques to the case of Hydra\\,A, where the shape and the size of the cavities produced by the source in the surrounding medium are known from X-ray data.  We emphasise that it is essential to account for the potential exclusion of magnetoionic material from a large volume containing the radio source when using the RM variations to derive statistical properties of the fluctuations in the foreground magnetic field. ",
        "introduction": "\\label{Introduction}  The hot intracluster medium in rich clusters of galaxies is magnetised, with field strengths which are important for the suppression of heat conduction, even if not dynamically significant \\citep[and references therein]{CT02}. The fields are apparent via synchrotron radiation from cluster halo sources and Faraday rotation of polarized emission from embedded and background sources (the subject of the present paper). At high spatial resolution, the rotation of the {\\bf E}-vector position angle is accurately proportional to the square of the wavelength and depolarization with increasing wavelength is small, indicating that a foreground medium is responsible.  Imaging of these foreground Faraday rotation measure (RM) variations across radio sources in clusters has recently been used to estimate the power spectrum of RM and, by implication, of magnetic-field fluctuations, in turn leading to improved estimates of field strength \\citep{EV03,VE03,VE05,Murgia,Govoni06,Guidetti08}. RM fluctuations are also seen against radio sources in sparser environments such as galaxy groups (e.g.\\ \\citealt*{PBW84,LB87}), but the structure and strength of the magnetic fields in these systems have not been studied in detail.  \\begin{figure} \\begin{center} \\epsfxsize=6cm \\epsffile{i5.5deep2.eps} \\caption{Grey scale at 5.5~arcsec FWHM resolution showing the main regions of 3C\\,31 in total intensity at 1.4 GHz as defined by \\citet[fig.~1b]{lb08a}.  This image shows the radio source in equatorial co-ordinates, while the displays of depolarization and rotation measure below have been rotated anti-clockwise by 19.7 deg to make the inner jet axis vertical. \\label{layout} } \\end{center} \\end{figure}  RM variations in clusters of galaxies have normally been modelled assuming that the magnetic field is tangled either on a single scale \\citep{Burn66,Felten96} or over a range of scales (e.g.\\ \\citealt{EV03,Murgia}) and embedded in a smoothly varying, spherically symmetric density distribution with parameters derived from X-ray observations. This interpretation is complicated by three effects: \\begin{enumerate} \\item It is usually impossible to locate a radio galaxy along the line of sight within a cluster, and the path length over which Faraday rotation occurs is therefore not well determined.  It often has to be assumed that the sources are in the cluster mid-plane. \\item The sources are inclined to the line of sight and may be extended on scales comparable with changes in external density, so the Faraday depth will vary systematically over the radio structure. The observed correlation between depolarization and jet sidedness (in the sense that the side of the source containing the brighter jet depolarizes less rapidly with increasing wavelength) is most easily interpreted in this way: if the nearer jet is brighter as a result of Doppler boosting, then its emission is seen through less magnetoionic material and therefore suffers less Faraday rotation \\citep{L88}.  Although the general association between depolarization asymmetry and jet sidedness is well documented in both weak and powerful sources \\citep{Morg97,GC91}, there are as yet few sources that exhibit asymmetric depolarization for which it has been clearly established that the Faraday rotation is indeed due to a foreground medium (although there are no known counter-examples). The best-observed examples are all in central cluster galaxies with cooling cores: Cygnus\\,A \\citep*{DCP87}; M\\,87 \\citep*{OEK}; Hydra\\,A \\citep{HydraA} and Hercules\\,A \\citep{HerA}. \\item Radio sources must interact with their environments, leading to local deviations from spherical symmetry in the density distribution. There is direct evidence for the presence of evacuated cavities in the hot plasma coincident with the radio lobes of sources in clusters, as predicted by models of radio-source evolution \\citep{Scheuer74,McNN2007}. Magnetic fields in the surrounding medium may be important in stabilizing the cavities against Kelvin-Helmholtz and Raleigh-Taylor instabilities \\citep{DP08}. \\end{enumerate}  In the present paper, we image and model the Faraday rotation distribution across the bright FR\\,I \\citep{FR74} class radio source 3C\\,31. This source is identified with the nearby elliptical galaxy NGC\\,383, which is the brightest member of a rich group of galaxies \\citep{Arp66,Zwi68}. We therefore probe a different environment from the Abell clusters studied by previous authors. Hot gas associated with the galaxy has been detected on both the group and galactic scales by X-ray imaging \\citep{KB99,Hard02}. The nucleus of the source is located at the centre of the hot gas distribution, simplifying the geometry, and the orientation is well constrained, at least for the inner jets \\citep{LB02a}. 3C\\,31 shows a pronounced depolarization asymmetry in the expected sense \\citep{Burch79,Strom83,And92}. Because it is bright, highly polarized, and well resolved in both dimensions \\citep[][and references therein]{lb08a}, the variation of the Faraday rotation measure across it can be studied in unusually rich detail.  We can make a stringent test of  the hypothesis of foreground rotation and evaluate the spatial statistics of RM over different parts of the source.  As the jet orientation is well constrained, we have a unique opportunity to compare the global changes in RM fluctuation amplitude across the structure with the predictions of three-dimensional models, including the effects of departures from spherical symmetry.  3C\\,31 therefore provides an opportunity to examine the kiloparsec-scale structure of the magnetoionic medium in an environment that may be more characteristic of the FR\\,I population as a whole than the clusters which have been the subject of many earlier studies.  We will show that our three-dimensional analysis of 3C\\,31 is most consistent with the hypothesis that the radio source has evacuated large cavities which are essentially devoid of thermal plasma, although existing X-ray observations of the NGC\\,383 group are not deep enough to image such cavities if they are present.  This motivated us to apply our techniques to the RM distribution of a second source, Hydra\\,A, in which the cavities are well characterised \\citep{Wise2007}.  \\begin{figure*} \\epsfxsize=17cm \\epsffile{rmpdplo.eps} \\caption{Depolarization and rotation measure at a resolution of 5.5\\,arcsec FWHM. (a) Burn law $k$ in rad$^2$\\,m$^{-4}$, derived from a weighted least-squares fit to the relation $\\ln [p(\\lambda)] = \\ln [p(0)] - k\\lambda^4$ between 4985 and 1365\\,MHz.  The colour range is from $-400$ to 1000\\,rad$^2$\\,m$^{-4}$. (b) RM from a fit to the four frequencies between 1636 and 1365\\,MHz, primarily to show the RM in the North tail and spur. The range is $-$200 to $-$33\\,rad\\,m$^{-2}$. (c) as (b), but using 5 frequencies between 4985 and 1365\\,MHz. The errors are smaller than in panel (b), but fewer points in the North of the source have sufficient signal to determine the RM.  The rectangles marked `NP2' and `SP3' denote the regions for the RM histograms in Fig.~\\ref{RMhist}(a) and (e) and the structure-function analysis of Section~\\ref{Sfunc_observe}.  The crosses mark the positions for the plots in Figs~\\ref{p-lambdasq-5.5} and \\ref{PA5.5}. \\label{fig:RMDP5.5}} \\end{figure*}   \\begin{figure*} \\epsfxsize=17cm \\epsffile{rmdphi.eps} \\caption{Depolarization and rotation measure at a resolution of 1.5\\,arcsec FWHM.  (a) Burn law $k$ in rad$^2$\\,m$^{-4}$, derived from a weighted least-squares fit to the relation $\\ln [p(\\lambda)] = \\ln [p(0)] - k\\lambda^4$ for six frequencies between 8440 and 1365\\,MHz.  The colour range is from $-400$ to 1000\\,rad$^2$\\,m$^{-4}$, as in Fig.~\\ref{fig:RMDP5.5}(a).  (b) RM determined by a fit to position-angle images at 6 frequencies between 1365 and 8440\\,MHz using a version of the {\\sc aips} task {\\sc rm} modified by G.B. Taylor.  The colour range is $-$180 to $+$20\\,rad\\,m$^{-2}$. The boxes NP1, SP1 and SP2 show the areas used for the RM histograms in Fig.~\\ref{RMhist} and for the structure-function analysis in Section~\\ref{Sfunc_observe}. (c) As (b), but using the {\\sc pacerman} algorithm of \\citet*{Pacerman}.  The crosses mark the positions for the plots in Figs~\\ref{p-lambdasq} and \\ref{PA-lambdasq}. \\label{fig:RMDP1.5}} \\end{figure*}   Section~\\ref{Pol} of the present paper presents the main features of our data on the distributions of the Faraday rotation and depolarization of the polarized emission from 3C\\,31, and Section~\\ref{over} gives an overview of our analysis and general assumptions.  Sections~\\ref{2d} and \\ref{3d} discuss the structure of the magnetic field in the foreground screen that is responsible for the Faraday rotation measure fluctuations.  Section~\\ref{2d} describes our analysis of the two-dimensional spatial statistics of the RM fluctuations. Section~\\ref{3d} explores three-dimensional models of the magnetoionic medium consistent with this two-dimensional analysis. Section~\\ref{Hydra} presents an analysis of the RM distribution in Hydra\\,A. Section~\\ref{Conclusions} summarises our conclusions.  Some mathematical details are given in Appendices~\\ref{App-theory} and \\ref{psformulae} and an analysis of the effects of partially-resolved foreground rotation is described in Appendix~\\ref{SP3depol}.  Throughout this paper, we take NGC\\,383 to have a redshift of 0.0169 (the mean value from \\citealt{Smith2000}, \\citealt*{HVG} and \\citealt{RC3}).  We assume a concordance cosmology with a Hubble Constant $H_0$ = 70\\,km s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_M = 0.3$. At the redshift of NGC\\,383, this gives a linear scale of 0.344\\,kpc\\,arcsec$^{-1}$. For Hydra\\,A, we take z = 0.0549 \\citep{Smith2004} and a linear scale of 1.067\\,kpc\\,arcsec$^{-1}$. ",
        "conclusions": "\\label{Conclusions}  \\subsection{Summary}  We have analysed images of linearly polarized emission in the FR\\,I radio galaxy 3C\\,31 with resolutions of 5.5 and 1.5\\,arcsec FWHM at six frequencies between 1365 and 8440\\,MHz.  Our use of all six frequencies in an analysis of the variation of the degree of polarization with wavelength has enabled us to measure very small depolarizations, which we characterize using the Burn law parameter $k$.  At 1.5-arcsec resolution, we find $\\langle k \\rangle = 3$\\,rad$^2$\\,m$^{-4}$ (DP$^{\\rm 22cm}_{\\rm 6cm}$ = 0.99) within 60\\,arcsec of the nucleus in the North jet and $\\langle k \\rangle = 55 $\\,rad$^2$\\,m$^{-4}$ (DP$^{\\rm 22cm}_{\\rm 6cm}$ = 0.88) in the South.  The low depolarization and absence of deviation from $\\lambda^2$ rotation confirm that almost all of the observed Faraday rotation is due to foreground material.  The RM fluctuation amplitude is significantly larger in the South.  The amplitude and scale of the RM fluctuations in 3C\\,31 appear qualitatively similar to those derived for other well-observed FR\\,I sources in comparable environments, e.g.\\ NGC\\,6251 and 3C\\,449 \\citep{PBW84,Fer99}, but are much smaller in amplitude than those in FR\\,I sources located in cD galaxies with cooling cores such as M\\,87 and Hydra\\,A \\citep{OEK,HydraA} and much larger than for NGC\\,315, which is in a sparse group \\citep{LCCB06}.  We have determined the RM structure function for five regions in 3C\\,31 over which the fluctuation amplitude is roughly constant.  We model the fluctuations on the assumption that the magnetic field responsible for the foreground rotation is an isotropic Gaussian random variable (characterised entirely by its power spectrum). For the range of foreground Faraday rotation seen in 3C\\,31, the observed RM image is very close to the true image convolved with the observing beam. We have used this approximation to incorporate the effects of the beam in the calculation of theoretical structure functions for model power spectra by numerical integration of a Hankel transform relation. We evaluate errors due to imperfect sampling by making multiple realisations of the model power spectrum on the observed grid.  For the four regions with the most reliable RM measurements, our data are consistent with an RM power spectrum which has the same form everywhere but whose amplitude varies with position. The power spectrum cannot be a single power law over the entire range of sampled spatial frequencies. A simple functional form that fits our data is a broken power law $\\hat{C}(f_\\perp) \\propto f_\\perp^{-q}$ with $q = 2.32$ for $f < 0.062$\\,arcsec$^{-1}$ and $\\hat{C}(f_\\perp) \\propto f_\\perp^{-11/3}$, as expected for Kolmogorov turbulence, at higher frequencies. A power spectrum with a similar low-frequency slope and an abrupt high-frequency cut-off gives almost as good a fit to the observed structure functions, however.  We can only determine an approximate lower limit to the outer scale: in terms of spatial frequency, $f_{\\rm min} \\la 0.005$\\,arcsec$^{-1}$ (0.015\\,kpc$^{-1}$).  We conclude that these (and other, similar) models of the power spectrum cannot be distinguished using the available RM data alone. Our depolarization measurements at 1.5-arcsec resolution are also consistent with extrapolations of either power spectrum to higher spatial frequencies. In the South lobe of 3C\\,31, the model power spectra (normalized using the RM structure function) predict more depolarization than is observed at low resolution, however.  In Hydra\\,A, we again fit the RM power spectrum with a power law, but measure a steeper slope, $q = 2.8$, than in 3C\\,31. We find no evidence for a high-frequency cut-off or for a portion of the power spectrum with the Kolmogorov slope at spatial frequencies $\\la$0.8\\,kpc$^{-1}$ (cf.\\ \\citealt{VE05}).  RM analyses for sources in Abell clusters give qualitatively similar results. \\citet{Murgia} found a spectral slope $q \\approx 2$ extending to large spatial scales in Abell 119 and for Abell 2255, \\citet{Govoni06} argued for a spectral slope which steepens from $q \\approx 2$ to $q \\approx 4$ with increasing distance from the cluster centre. In contrast, \\citet{Guidetti08} found that a power-law power spectrum with a Kolmogorov slope and an abrupt long-wavelength cut-off at $f_{\\rm min}$ = 0.014\\,kpc$^{-1}$ gave a very good fit to their RM and depolarization data for A2382, although a shallower slope $q \\approx 2$ extending to longer wavelengths was not ruled out.  All three of these studies assumed high-frequency cut-offs $f_{\\rm max} \\approx 0.1$\\,kpc$^{-1}$.  In all cases except Hydra\\,A (where the range of spatial scales we have investigated is very small), there is evidence for a change in the power-law slope of the power spectrum, but ambiguity about its precise form. Two simple forms of the power spectrum which fit all of the datasets are: \\begin{enumerate} \\item a broken power-law with a high-frequency slope $q_{\\rm high} = 11/3$ (the Kolmogorov value) and no high-frequency cut-off, but flattening to a slope $q_{\\rm low} \\approx$ 2 -- 3 or even cutting off below a frequency $f_{\\rm b} \\sim$ 0.01 -- 1\\,kpc$^{-1}$ and \\item a power-law with index $q \\approx$ 2 -- 3 extending to very low frequencies and a high-frequency cut-off at $f_{\\rm max} \\sim$ 0.1 -- 2\\,kpc$^{-1}$. \\end{enumerate} Other qualitatively similar models would also fit the data.  Numerical simulations predict a wide range of spectral slopes \\citep{Dolag06}, and a detailed comparison between theory and observation seems premature.  Our observation of a change in RM fluctuation amplitude across the nucleus of 3C\\,31 qualitatively supports the relativistic-jet interpretation of the brightness asymmetry between the jets, wherein the fainter (receding) jet is observed through a greater path length of magnetoionic medium than the nearer (approaching) jet. To test this idea quantitatively, we have modelled the global RM variations using three-dimensional simulations, again treating the magnetic field as an isotropic Gaussian variable with a broken power law power spectrum. We assume that the entire structure of 3C\\,31 has an inclination of 52$^\\circ$ as in our kinematic models of the inner jets.  The profile of the RM variance as a function of distance from the nucleus of the galaxy is incompatible with a spherically symmetric distribution for the magnetoionic medium as derived from X-ray observations: this would cause too gradual a variation.  We expect, however, that the radio source will have evacuated cavities in the surrounding hot gas. A model in which the edges of the cavities coincide with the boundaries of the radio emission on large scales is more compatible with the observed RM profile, but current X-ray images are not deep enough to reveal any cavities around 3C\\,31. The derived central rms magnetic-field strength $B_0 \\approx 0.7$\\,nT (7\\,$\\mu$G) and the field varies roughly linearly with density.  The magnetic field is not dynamically dominant over the volumes we model. The ratio of thermal to magnetic pressure, $\\beta_{\\rm P} = p_{\\rm thermal}/(B^2/2\\mu_0) \\approx 10$ at the core radius of the group gas, $r_c = 154$\\,arcsec (52\\,kpc).  We have also modelled the RM distribution in Hydra\\,A, where cavities {\\em are} seen in X-ray images.  A three-dimensional simulation including the cavities gives a good representation of the RM fluctuation profile for an inclination of $45 \\pm 7$\\,deg with $B \\propto n^{0.25 \\pm 0.2}$ and a central magnetic field strength of 1.9\\,nT (19\\,$\\mu$G), so $\\beta_{\\rm P} \\approx 30$ at $r_c = 26$\\,arcsec (27.7\\,kpc).  Close to the nuclei of both sources, our models predict higher RM fluctuation amplitudes than are observed. Although the sampling variances are large, this may be an indication that our assumption of a power-law dependence of magnetic field strength on density is oversimplified and that we have therefore overestimated the field strengths over these volumes.  Few other RM profiles have been published for FR\\,I sources with jets, and it is not yet clear whether the abrupt changes in RM fluctuation amplitude across the nucleus observed in 3C\\,31 and Hydra\\,A are common. In 3C\\,296 \\citep{LCBH06}, there is a smooth variation of rms RM across the nucleus, consistent with the inferred inclination of $58^\\circ$ and Faraday rotation produced by a spherically symmetric hot gas component with a large core radius, as observed using {\\sl XMM-Newton} \\citep{Croston08}. The RM fluctuations observed in NGC\\,315 \\citep{LCCB06} are too small to determine a reliable profile, although clearly larger on the counter-jet side.  \\subsection{Further work}  Our modelling of the RM variations in 3C\\,31 and Hydra\\,A suggests a number of observational and theoretical investigations. \\begin{enumerate} \\item It is clearly necessary to sample a larger range of spatial scales to improve constraints on the RM power spectrum and in particular to decide whether it has the form expected for Kolmogorov turbulence on small scales. Higher {\\em resolution} would probe small scales directly and remove the need to rely on crude estimates from depolarization. Higher {\\em sensitivity} would provide better sampling of large spatial scales by allowing study of low-surface-brightness regions.  Use of analysis techniques such as Bayesian maximum likelihood would give more rigorous error estimates.  We have shown (Fig.~\\ref{fig:RMDP5.5}b and c) that it is possible to derive accurate RM images from a small number of discrete frequencies in a suitably chosen observing band: the new generation of wideband correlators used by EVLA and eMERLIN will enable simultaneous observations over much larger frequency ranges, greatly improving the speed and accuracy of RM determination. \\item Large changes in RM fluctuation amplitude are seen across the nuclei in 3C\\,31 and Hydra\\,A, but not in 3C\\,296. The prevalence of this phenomenon and its dependence on orientation, large-scale radio structure and environment are unknown. It is clear that the RM fluctuation amplitude scales with density, ranging from a few rad\\,m$^{-2}$ in sparse environments (e.g.\\ NGC\\,315) to $\\sim 10^4$\\,rad\\,m$^{-2}$ in cooling cores such as Hydra\\,A. Observations of a larger sample, employing resolutions and frequency ranges matched carefully to the environments, will be needed. \\item We have argued that the presence of cavities in the external medium must be taken into account when modelling RM distributions: failure to do so will lead to biased estimates of the magnetic field strength and its variation with density. Accurate modelling requires knowledge of the source geometry (which we can constrain for 3C\\,31 but not Hydra\\,A) and observations of the cavity structure (which are available for Hydra\\,A but not yet for 3C\\,31).  Deep {\\sl XMM-Newton} observations of the 3C\\,31 group are needed, as are comparable data for further sources. Kinematic models of the radio jets in Hydra\\,A in other sources could be used to constrain their orientations. \\item Deviations in source structure and density from axisymmetry, Faraday rotation from foreground galaxies, X-ray emission from other group members in the field and from relativistic electrons deposited in the cavity by the radio source (via the inverse Compton process) and variations in the Galactic magnetic field across the source structure are potential complications which need to be better understood. \\item The assumption that the magnetic field is an isotropic, Gaussian random variable is questionable. RM images of some radio sources show evidence for preferred directions and there are theoretical reasons to suppose that fields at the edges of cavities and behind radio-source bow shocks should be two-dimensional.  More detailed simulations of radio-source evolution in a magnetised intergalactic medium are needed. \\end{enumerate}"
    },
    "0809/0809.0414_arXiv.txt": {
        "abstract": "{% Toroidal magnetic fields subject to the Tayler instability can transport angular momentum. We show that the Maxwell and Reynolds stress of the nonaxisymmetric field pattern depend linearly on the shear in the cylindrical gap geometry. Resulting angular momentum transport also scales linear with shear. It is directed outwards for astrophysical relevant flows and directed inwards for superrotating flows with ${\\rm d}\\Om/{\\rm d}R>0$. We define an eddy viscosity based on the linear relation between shear and angular momentum transport and show that its maximum for given Prandtl and Hartmann number depends linear on the magnetic Reynolds number $\\Rm$. For $Rm\\simeq1000$ the eddy viscosity is of the size of $30$ in units of the microscopic value. } ",
        "introduction": "Magnetic field configurations in astrophysical objects occur in wide variety. They play an essential role in the formation of flow pattern, are source and result of dynamo processes and enhance transport processes, where especially transport of angular momentum is of interest here. Crucial for magnetic field evolution is the interplay between toroidal and poloidal fields. In a differentially rotating environment for instance a weak poloidal field is wind-up into a strong toroidal one. A strong toroidal field on the other hand can become unstable due to the Tayler instability (\\cite{tayler_1957,tayler_1973,vanda72}) or the Azimuthal magnetorotational instability (AMRI, see \\cite{rued07}). As result also new components of a poloidal field are generated (\\cite{gellert07}). To study in more detail the occurrence of the Tayler instability (TI) and especially the connected transport of angular momentum, we use a model in a cylindrical gap, that starts with an already dominating toroidal field and look for the fields instability in a differentially rotating environment.\\\\ Tayler instability mainly leads to nonaxisymmetric fields. This gives raise to the idea it could be part of a new dynamo mechanism caused only by the magnetic field (\\cite{spruit99}). Even if there is a ongoing discussion if such a dynamo could really work (\\cite{braith06,zahn07,gellert08,rincon08}), the TI might play an essential role for mixing of chemicals (\\cite{rued08}), transport of angular momentum and resulting flattening of the rotation profile of stars. Estimates for angular momentum transport based on the so-called 'Tayler-Spruit' dynamo are in use as input for star evolution models (\\cite{maeder05,yoon06}). Putting such estimates on more robust basis is the aim of this work. It uses nonlinear 3D simulations of the Tayler instability in a cylindrical gap to compute the actual angular momentum transport \\begin{equation} T=\\langle u_R' u_\\phi' - \\frac{1}{\\mu_0\\, \\rho} B_R' B_\\phi'\\rangle, \\label{eq_t} \\end{equation} where field fluctuations are defined as nonaxisymmetric contributions of the whole velocity and magnetic field in cylindrical coordinates $(R, \\phi, z)$. Further on it is shown that $T$ scales linear with the differential rotation, that is angular momentum transport can be expressed in terms of a positive turbulent or eddy viscosity $\\nu_{\\rm T}$. We find that for given shear and magnetic field strength always exists a maximum value of the eddy viscosity, which grows linearly with the magnetic Reynolds number $\\Rm$. ",
        "conclusions": "Strong toroidal fields under the influence of a differentially rotating medium can become unstable to nonaxisymmetric perturbations. This instability, the Tayler instability, appears only if the shear is not to strong. In cylindrical gap geometry for Prandtl numbers $0.1\\leq\\Pm\\leq1$ there is always an upper boundary for the Reynolds number above which the instability is suppressed. The smoothing action of differential rotation dominates. Additionally it needs the crossing of a threshold for the strength of the magnetic field, for a nearly constant toroidal field with $\\mub=1$ this is of the order of $\\Ha\\approx100$. If our astrophysical object in mind provides such requirements, the TI leads to nonaxisymmetric magnetic field pattern.  Using the nonaxisymmetric parts as fluctuations, the resulting angular momentum transport scales linear with the differential rotation, it is positive for astrophysical relevant configurations with $\\muo\\leq0.75$, e.g. angular momentum is transported outwards. This gives raise to the introduction of an eddy viscosity which causes the additional turbulence in the underlying medium.  For a fixed magnetic field there is always an upper and lower critical Reynolds number in between the instability can be observed. And there is also a maximum value of angular momentum transport and eddy viscosity. Taking the maxima of the latter we find its values follow a linear dependence on the magnetic Reynolds number. A direct influence of the field strength is hidden behind the fact, that with the Reynolds number also the Hartmann number is fixed for the point where the eddy viscosity has its maxima.  The linear relation $\\nuT(\\Rm)$ is derived for a simplified model with fixed toroidal field and without any influence from density stratifications for rather high Prandtl numbers $0.1\\leq\\Pm\\leq1$. Although, imagine it hold also for Prandtl numbers one or two orders of magnitude less, this relation could be used to estimate eddy viscosities usable as input for star evolution models. For instance for a young, fast rotating sun it gives values around $\\nut=30$."
    },
    "0809/0809.1762_arXiv.txt": {
        "abstract": "{In the first paper of this series, we presented hydrodynamical simulations with radiative cooling of jet models with parameters representative of the symbiotic system MWC 560. These were jet simulations of a pulsed, initially underdense jet in a high-density ambient medium. They were stopped when the jet reached a length of 50 AU. There, however, a transition of the initially underdense jet towards an overdense jet should occur, which should result in changed kinematics. A few minor differences between the models and the observations were thought to be solved by a model with an increased jet density during the pulses which was calculated only with purely hydrodynamical means in the former paper.} {Therefore, we describe two hydrodynamical simulations with cooling beyond this density balance, one with the same parameters as model i in Paper I (now called model i'), which was presented there with and without cooling, and the second with higher gas densities in the jet pulses (model iv').} {Hydrodynamical simulations, with a further approximated cooling treatment compared to Paper I, were used to be able to enlarge the computational domain.} {The transition causes changes in the expansion of the cocoon and therefore the morphology of the jet, e.g. a larger radial width of the jet knots. We investigate the radiation properties of the jets, the bremsstrahlung and optical emissivities, integrated emission maps, and synthetic absorption line profiles.} {The conclusion that the high observed velocities in CH Cygni, R Aquarii, and MWC 560 favor the models with cooling is unchanged by the transition. The observed parallel features in R Aquarii can be produced by the internal knots or by a variable dense radiative shell of shocked ambient medium. The absorption line profiles show that the real parameters in MWC 560 are closer to model iv' than to model i'.} ",
        "introduction": "The emergence of jets is a very common phenomenon in a variety of astrophysical objects, and it occurs in systems of very different size and mass scales. Jets are observed in active galactic nuclei (AGN) in which they are formed around supermassive black holes, in X-Ray binaries with stellar black holes or neutron stars, in young stellar objects (YSO), in supersoft X-ray sources and in symbiotic stars. Symbiotic systems consist of a red giant undergoing strong mass loss and a white dwarf. More than one hundred symbiotic stars are known, but only about ten systems show jet emission. The most famous systems are R Aquarii, CH Cygni, and MWC 560.  R Aquarii, with a distance of about 200 pc, is one of the nearest symbiotic stars and a well known jet source. The jet has been extensively observed in the optical, at radio wavelengths, and with X-ray observations \\citep[e.g.][]{SoU,PaH,HMKM,HLDF,KPL}. It shows a rich morphology, e.g. a series of parallel features in the jet and the counter-jet, extending a few hundred AU each. In 1984/85 CH Cygni showed a strong radio outburst, during which a double-sided jet with multiple components was ejected \\citep{TSM}. This event allowed an accurate measurement of the jet velocity near 1500 km s$^{-1}$. In HST observations \\citep{EBS}, arcs can be detected that also could be produced by episodic ejection events. While the first two objects are seen at high inclinations, the jet axis in MWC 560 is practically parallel to the line of sight. This special orientation provides an opportunity to observe the outflowing gas as line absorption in the source spectrum. With such observations the radial velocity and the column density of the outflowing jet gas close to the source can be investigated in great detail. In particular we can probe the acceleration and evolution of individual outflow components with spectroscopic monitoring programs, as described in \\citet{SKC}.  In \\citet[][hereafter Paper I]{SCS}, we presented hydrodynamical simulations with and without cooling of jets with parameters that were intended to represent those in MWC 560. In a grid of eight simulations we investigated the influence of different jet pulse parameters. Due to the high computational costs of simulations including cooling, this grid was restricted to adiabatic simulations. Only one model simulation was performed that included a treatment of radiative cooling.  The adiabatic jet models showed some quantitative differences for the different pulse parameters. In comparison with the model including cooling, huge differences were present. Those concern, e.g., a significantly higher bow-shock velocity and a well-defined periodic shock structure with high-density, low-temperature knots separated by hot, low-density beam sections in the radiative model. The expansion velocity of CH Cyg was shown to agree very closely with the jet model with cooling. The situation in R Aqr was less clear. The observationally derived proper motion was more in the range of the gas motion of the adiabatic models, although the origin of the observed emission has not been confirmed yet. The small jet opening angle of 15$^{\\circ}$ favored a jet with cooling more. We showed that the adiabatic jet models are not able to produce low temperature regions, which are, however, necessary to explain strong absorptions from low ionization species in the spectra of MWC 560 as \\ion{Ca}{ii}, but the cooling of the jet gas still did not seem efficient enough. This should be achieved by another model with higher gas densities in the jet outflow. The basic structure of the observed jet absorptions in MWC 560 was reproduced, including the mean velocity, the velocity width and temporal evolution of the highest velocity components, but not their strengths. Again, another model with higher gas densities should solve this.  A drawback of our former simulations was the fact that they were stopped when the jet reached a length of 50 AU. This is unsatisfactory from an observational point of view, as the observed extents of the jet in R Aquarii and CH Cygni are much larger, as well as from the theoretical viewpoint. At this point, the density of the environment has decreased down to the density in the jet nozzle. A transition of the initially underdense jet towards an overdense jet should occur, which should result in changed kinematics. Underdense jets decelerate more than overdense jets, because the position of the jet head can be basically derived from a balance of the ram pressures of jet and external medium and because the relative importance of ram and thermal pressure should also change. As the kinematics of the models was one of the results of the former models, which were compared with the observations, it should be tested whether simulations with larger spatial scales lead to new insights into the physics of jets in symbiotic stars.  Therefore, we performed two hydrodynamical simulations with cooling beyond this density balance, one with the same parameters as model i in Paper I, which was presented both with and without cooling, and the second with higher gas densities in the jet pulses. Due to constraints set by the available computer resources, we were forced to approximate the former treatment of radiative cooling again.  In Sect. \\ref{sec_model}, we briefly describe the changes of our model with respect to Paper I. After a validating of the approximated cooling treatment in section \\ref{sec_valid}, we investigate the global and internal structure of the jet (section \\ref{sec_struct}) and its consequences on the emission (section \\ref{sec_emiss}). Then synthetic absorption line profiles are calculated in section \\ref{sec_abs}. Finally a summary and a discussion are given. ",
        "conclusions": "In this paper, we present results of two new hydrodynamical simulations, including radiative cooling. They were started due to a drawback in our former simulations presented in Paper I \\citep{SCS}, which were stopped close to density balance of the jet and the ambient medium. As underdense and overdense jets behave in different ways and as the kinematics of the jets was a main result in Paper I, we expected new insights.  Another point is that a few minor differences between the models and the observations were thought to be solved by a model with an increased jet density during the pulses -- model iv' -- which was calculated only with purely hydrodynamical means in the former grid of simulations. Therefore, we performed two hydrodynamical simulations with an approximated cooling treatment beyond this density balance, one with the same parameters as model i in Paper I, which was presented there with and without cooling, and the second with higher gas densities in the jet pulses, as in model iv.  Since in the new cooling treatment ionization fractions of collisional ionization equilibrium (CIE) were used and not self-consistent general non-equilibrium (NEQ) rate equations, we discussed the possible effects of underestimating the ionization fraction of hydrogen and validated this cooling treatment by comparing density and temperature values of the gas parcels in the jet, as well as by showing plots of the ionization fraction itself. We saw that the differences are only very small.  In Paper I, the jet in model i with cooling had a constant cross section over the simulated 74 days. This led to an accelerated motion with high velocities of 730 km s$^{-1}$, while the velocity of the adiabatic jet was only about 200 km s$^{-1}$. The extended model i' in this paper now shows that after the first 70 days the transition from an underdense to an overdense jet lets the cross section inflate and therefore compensate for the density profile of the external medium. The acceleration is stopped and the jet then has a constant velocity of about 740 km s$^{-1}$. Thus, the transition results in a completely different motion. The conclusion that the high observed velocities in CH Cygni, R Aquarii, and MWC 560 favor the models with cooling is unchanged by the transition.  The radial inflation mentioned above also changes the internal structure of the jet. The width of the internal knots is also increased, as the basic composition of the jets is not affected. The cocoon of shocked jet matter is again confined to a small shell close to the contact discontinuity and the whole jet lobe is filled with unshocked jet material. In model i', these knots are very pronounced along their width, while their structure is not intact in model iv'.  These knots are the locations of enhanced bremsstrahlung and optical emissivity, which is therefore again more pronounced in model i'. These knots could be identified with the observed parallel features in R Aquarii \\citep{PaH}. After rotating the emissivity plots and integrating them, however, other features become more prominent. The internal knots are blended by the emission of the dense radiative shell of shocked ambient medium. As they are spatially variable, parallel rings of enhanced emission are created by the rotation and also look similar to the observations, if one keeps in mind that axisymmetry was assumed in our model. Without that, the shapes of the features on the shell should be different.  The observed length of the jet in R Aquarii is several hundred AU, thus much larger than in our simulations. At larger distances from the source, however, the density of the ambient medium drops by another order of magnitude, making it likely that the relative contribution of the emission of the knots will increase.  Concerning the absorption line profiles, no large differences between the former model i with cooling and the recent model i' are visible. This means that the influence of the age of the jet on the profiles should only be marginal. In model iv', however, differences arise. The stable absorption trough is shifted toward higher velocities, from -1050 km s$^{-1}$ up to -1300 km s$^{-1}$, and its width is increased from about 200 km s$^{-1}$ to 500 km s$^{-1}$. In the high velocity component, several deep absorptions are present in contrast to the few moderate ones in model i'.  Both results highly increase the ability of the model to reproduce the observed absorption line profiles \\citep[see Fig.1 in Paper I or ][]{SKC}. We can say here that the real parameters in MWC 560 are closer to model iv' than to model i'.  As the pulse structure was not changed, the line widths of the high velocity components are again quite narrow. This could perhaps be improved by using a sinusoidal or Gaussian pulse profile, instead of the rectangular steps in velocity and density, and also a spatial velocity and density profile, instead of the constant values within the jet nozzle."
    },
    "0809/0809.2210_arXiv.txt": {
        "abstract": "The main results from the Auger Observatory are described. A steepening of the spectrum is observed at the highest energies, supporting the expectation that above $4\\times 10^{19}$~eV the cosmic ray energies are significantly degraded by interactions with the CMB photons (the GZK effect). This is further supported by the correlations observed above $6\\times 10^{19}$~eV with the distribution of nearby active galaxies, which also show the potential of Auger to start the era of charged particle astronomy. The lack of observation of photons or neutrinos strongly disfavors top-down models, and these searches may approach in the long term  the sensitivity required to test the fluxes expected from the secondaries of the very same GZK process.  Bounds on the anisotropies at EeV energies contradict hints from previous  experiments that suggested a large excess from regions near the Galactic centre or  the presence of a dipolar type modulation of the cosmic ray flux. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1213_arXiv.txt": {
        "abstract": "I conjecture that the time modulated decay rates reported in single ion measurements of two body electron capture decay of hydrogen like heavy ions at GSI may be related to neutrino spin precession in the static magnetic field of the storage ring. These `GSI Oscillations' arise from interference between amplitudes of decay within and without the magnetic field, a scenario that requires a Dirac neutrino magnetic moment six times lower than the Borexino solar neutrino upper limit of 0.54 x 10E(-10) Bohr magneton. I also show in a way not discussed before that the time modulation associated with interference between massive neutrino amplitudes, if such interference could arise, is of a period at least four orders of magnitude shorter than reported and must average to zero given the time resolution of the GSI measurements. ",
        "introduction": "\\label{sec:intro}  Measurements of weak interaction decay of multiply ionized heavy ions coasting in the ion storage-cooler ring ESR at the GSI laboratory, since the first report in 1992 \\cite{GBB92}, open up new vistas for dedicated studies of weak interactions. In particular, electron capture (EC) decay rates in hydrogen-like and helium-like $^{140}$Pr ions have been recently measured for the first time \\cite{LBG07} by following the motion of the decay ions (D) and the recoil ions (R). The overall decay rates $\\lambda_{\\rm EC}$ of these two-body $^{140}{\\rm Pr} \\to {^{140}{\\rm Ce}} + \\nu$ EC decays, in which no neutrino $\\nu$ is detected, are well understood within standard weak interaction calculations of the underlying $e^-p\\to\\nu_e n$ reaction \\cite{PKB08,IFR07}. However, a time-resolved decay spectroscopy applied subsequently to the two-body EC decay of H-like $^{140}$Pr and $^{142}$Pm single ions revealed an oscillatory behavior, or more specifically a time modulation of the two-body EC decay rate \\cite{LBW08}: \\begin{equation} \\lambda_{\\rm EC}(t)=\\lambda_{\\rm EC}[1+a_{\\rm EC}\\cos(\\omega_{\\rm EC} t + \\phi_{\\rm EC})], \\label{eq:omega} \\end{equation} with amplitude $a_{\\rm EC} \\approx 0.2$, and angular frequency $\\omega^{\\rm lab}_{\\rm EC} \\approx 0.89~{\\rm s}^{-1}$ (period $T^{\\rm lab}_{\\rm EC} \\approx 7.1$~s) in the laboratory system which is equivalent in the rest frame of the decay ion to a minute energy $\\hbar \\omega_{\\rm EC} \\approx 0.84 \\times 10^{-15}$~eV. Subsequent experiments on EC decays of neutral atoms in solid environment have found no evidence for oscillations with periodicities of this order of magnitude \\cite{VCD08,FBH08}. Thus, the oscillations observed in the GSI experiment could have their origin in some characteristics of the H-like ions, produced and isolated in the ESR, and in the electromagnetic fields specific to the ESR which are not operative in normal laboratory experiments. It is suggested here, in Sect.~\\ref{sec:magfield}, that the `GSI Oscillations' could indeed be due to the magnetic field which stabilizes and navigates the motion of the ions in the ESR.  Several works, by Kienle and collaborators, relegated the `GSI Oscillations' to interference between neutrino mass eigenstates that evolve coherently from the electron neutrino $\\nu_e$ \\cite{IRK08,Fab08,KKi08,IKP08,IK09}. This idea apparently also motivated the GSI experiment \\cite{LBW08}. Such interferences, according to these works, lead to oscillatory behavior given by Eq.~(\\ref{eq:omega}) with angular frequency $\\omega_{\\nu_e}$ where, again in the decay-ion rest frame, \\begin{equation} \\hbar\\omega_{\\nu_e}=\\frac{\\Delta (m_{\\nu}c^2)^2}{2M_Dc^2}\\approx 0.29 \\times 10^{-15}~{\\rm eV}. \\label{eq:delE} \\end{equation} Here, $\\Delta (m_{\\nu}c^2)^2 = (0.76 \\pm 0.02) \\times 10^{-4}$~eV$^2$ from accumulated solar $\\nu$ plus KamLAND reactor $\\bar{\\nu}$ data \\cite{SNO08} for the two mass-eigenstate neutrinos that almost exhaust the coupling to $\\nu_e$, and $M_D\\approx 130$~GeV/$c^2$ is the mass of the decay ion $^{140}{\\rm Pr}^{58+}$. Although the value of $\\hbar\\omega_{\\nu_e}$ on the r.h.s. of Eq.~(\\ref{eq:delE}) is about three times smaller than the value of $\\hbar \\omega_{\\rm EC}$ required to resolve the `GSI Oscillation' puzzle, getting down to this order of magnitude nevertheless presents a remarkable achievement if correct.{\\footnote{Eq.~(\\ref{eq:delE}) was also obtained by Lipkin \\cite{Lip08} assuming interference between two unspecified components of the initial state with different momenta and energies that can both decay into the same final state, an electron neutrino and a recoil ion with definite energy and momentum. This scenario was criticized by Peshkin \\cite{Pes08}.}} However, it is shown here in Sect.~\\ref{sec:osc} by following the methodology of Ref.~\\cite{IK09} that the correct energy scale under circumstances allowing oscillatory behavior is given by \\begin{equation} \\hbar\\Omega_{\\nu_e}=\\frac{\\Delta (m_{\\nu}c^2)^2}{2E_{\\nu}}\\approx 0.95 \\times 10^{-11}~{\\rm eV}, \\label{eq:DelE} \\end{equation} where $E_{\\nu}\\approx 4$~MeV is a representative value for neutrino energy in the H-like $^{140}{\\rm Pr}\\to{^{140}{\\rm Ce}}+\\nu_e$ and $^{142}{\\rm Pm}\\to {^{142}{\\rm Nd}}+\\nu_e$ EC decays \\cite{LBW08}. The energy $\\hbar\\Omega_{\\nu_e}$ is larger by over four orders of magnitude than $\\hbar\\omega_{\\rm EC}$ or $\\hbar\\omega_{\\nu_e}$ given by Eq.~(\\ref{eq:delE}), and so it would lead to modulation period shorter by over four orders of magnitude than the 7~s period reported by the GSI experiment. Given a time measurement resolution of order $0.5$~s \\cite{LBW08}, the effect of such oscillatory behavior would average out to zero.  Other authors \\cite{Giu08,BLS08,KKL08,CGL09,Mer09,Fla09} have rejected any link between neutrino mass eigenstates and the EC decay rate oscillatory behavior reported by the GSI experiment \\cite{LBW08}, the underlying argument being that since no neutrino is detected, the EC decay rate sums incoherently over neutrino mass eigenstates, whereas any oscillatory behavior requires interference between amplitudes summed upon coherently. It is instructive, however, to demonstrate this assertion also by adopting the methodology of Ref.~\\cite{IK09}, but with a caveat explained below. To this end the time-dependent EC transition amplitude $A_{\\nu_e}(i \\to f;~t)$, from initial state $i$ (D injected at time $t=0$) to a final state $f$ (R plus a coherent combination of neutrino mass eigenstates at time $t$), is written in terms of transition amplitudes $A_{\\nu_j}(i\\to f;~t)$ that involve propagating mass-eigenstates neutrinos $\\nu_j$ as \\begin{equation} A_{\\nu_e}(i\\to f;~t) = \\sum_jU_{ej}A_{\\nu_j}(i\\to f;~t), \\label{eq:Ae} \\end{equation} where $U_{e j}$ is a neutrino mixing matrix element of the $3 \\times 3$ unitary matrix $U$ \\begin{equation} |\\nu_\\alpha\\rangle = \\sum_{j=1}^3 U_{\\alpha j}^* \\, |\\nu_j\\rangle~~~~ (\\alpha = e, \\mu, \\tau) \\label{eq:U} \\end{equation} between the emitted electron-neutrino $\\nu_e$ and a mass-eigenstate neutrino $\\nu_j$ \\cite{ADA08}. For times of order seconds, appropriate to the `GSI Oscillations', the coherence implied by Eq.~(\\ref{eq:Ae}) is still in effect \\cite{Mer09} and the flavor basis is of physical significance. If the GSI experiment were to detect neutrino $\\nu_{\\beta}$ by a flavor measurement, the corresponding amplitude would have been generated by projecting Eq.~(\\ref{eq:Ae}) onto flavor $\\beta$: \\begin{equation} A_{\\nu_e \\to \\nu_{\\beta}}(i\\to f;~t) = \\sum_jU_{ej}A_{\\nu_j}(i\\to f;~t) U^*_{\\beta j}, \\label{eq:Abeta} \\end{equation} in close analogy with the discussion of neutrino oscillations in dedicated oscillation experiments (Eq.~(13.4) in Ref.~\\cite{ADA08}). The probability associated with the amplitude (\\ref{eq:Abeta}) is then given by \\begin{equation} {\\cal P}_{\\nu_e \\to \\nu_{\\beta}}(i\\to f;~t)=|A_{\\nu_e \\to \\nu_{\\beta}} (i\\to f;~t)|^2. \\label{eq:Pe} \\end{equation} Interference terms $A_{\\nu_j}A^*_{\\nu_{j'}}$ will arise in ${\\cal P}_{\\nu_e \\to \\nu_{\\beta}}(i\\to f;~t)$, leading to oscillations as shown in Sect.~\\ref{sec:osc}. Since the GSI experiment does not detect any neutrino, the overall probability is the sum of probabilities ${\\cal P}_{\\nu_e \\to \\nu_{\\beta}}(i\\to f;~t)$ over {\\it all} three flavors $\\beta$. The probability of observing the transition $i\\to f$, in which D decays to R and a neutrino is emitted but remains undetected, is thus given by \\begin{equation} {\\cal P}_{\\nu_e}(i\\to f;~t) =\\sum_{\\beta}|\\sum_j U_{ej}A_{\\nu_j}(i\\to f;~t) U^*_{\\beta j}|^2. \\label{eq:interf} \\end{equation} Note that the probability ${\\cal P}_{\\nu_e}(i\\to f;~t)$ cannot be written as a square of {\\it one} amplitude, simply because it involves different-flavor final states which require measurement schemes differing from each other and therefore adding up incoherently. Using the unitarity of the mixing matrix $U$, the summation over $\\beta$ in Eq.~(\\ref{eq:interf}) gets rid of the interference terms, leading to the final expression: \\begin{equation} {\\cal P}_{\\nu_e}(i\\to f;~t)=\\sum_{j}|U_{ej}|^2|A_{\\nu_j}(i\\to f;~t)|^2 \\approx |A_{\\nu}(i\\to f;~t)|^2, \\label{eq:incoherent} \\end{equation} where the dependence of the absolute-squared terms $|A_{\\nu_j}(i\\to f;~t)|^2$ on the species $\\nu_j$ was neglected,{\\footnote {This neglect does not hold for interference terms $A_{\\nu_j}A^*_{\\nu_{j'}}$, $j \\neq j'$, which give rise to oscillatory behavior, as discussed in Sect.~\\ref{sec:osc}.}} enabling repeated use of unitarity. The final result, Eq.~(\\ref{eq:incoherent}), is that the probability ${\\cal P}_{\\nu_e}(i\\to f;~t)$ for the two-body EC decay to occur is what standard weak interaction theory yields for a massless electron neutrino, regardless of its coupling to the mass-eigenstate neutrinos. This holds true also for the total EC decay rate which is obtained by time differentiation of ${\\cal P}_{\\nu_e}(i\\to f;~t)$ and which is found identical with the time-independent decay rate $\\lambda_{\\rm EC}$ derived ignoring neutrino mixing. Thus, although the mass-eigenstate components of the emitted neutrino oscillate against each other, the total EC decay rate does not exhibit any oscillatory behavior owing to the unitarity of the matrix $U$, Eq.~(\\ref{eq:U}), which transforms incoherence in one basis into incoherence in the other basis. If $U$ is nonunitary, the above argumentation breaks down, but this does not spoil the more straightforward argumentation that mass-eigenstate neutrinos, as distinct mass particles, have to be summed upon incoherently; one then goes directly from the amplitude $A_{\\nu_e}$, Eq.~(\\ref{eq:Ae}), into the probability ${\\cal P}_{\\nu_e}$, Eq.~(\\ref{eq:incoherent}), which indeed is an incoherent sum over the neutrino mass-eigenstates $\\nu_j$. ",
        "conclusions": "\\label{sec:sum}  In conclusion, it was shown that interference terms between different propagating mass-eigenstate neutrino amplitudes in two-body EC reactions on nuclei cancel out to zero when no neutrinos are detected. Coherence between propagating mass-eigenstate neutrinos is evident within each one of the amplitudes for detecting a given flavor neutrino. It is only when all final flavor rates are summed upon incoherently, as motivated by the different flavor measurements required, that cancelations occur and the overall rate becomes independent of time and is reliably calculable from standard weak interaction theory for massless neutrinos. The underlying logic here is that summing on all possible phase space for flavor neutrinos is equivalent within quantum mechanics to not detecting any specific neutrino.  Interference terms from different propagating mass-eigenstate neutrinos would survive and give rise to oscillatory behavior of the EC decay rate, if and {\\it only} if neutrinos are detected. It was shown that the period of oscillations in such a case is $T\\sim 4\\pi E_{\\nu}/{\\Delta (m_{\\nu}^2)}$ which for $E_{\\nu}\\approx 4$~MeV as in the GSI experiments \\cite{LBW08}, and for $\\Delta (m_{\\nu}^2) \\approx 0.76 \\times 10^{-4}$~eV$^2$ \\cite{SNO08}, assumes the value $T\\sim 4.4\\times 10^{-4}$~s, shorter by over four orders of magnitude than the period reported in these experiments. The oscillation period cited here is in full agreement with the oscillation length tested in dedicated neutrino oscillation experiments, provided the time $t$ is identified with $L/c$ where $L$ is the distance traversed by the neutrino. In particular, besides the $\\Delta (m_{\\nu}^2)$ neutrino input, it depends on the neutrino energy $E_{\\nu}$, not on the mass $M_D$ of the decay ion.  On the positive side, I have proposed a possible explanation of the `GSI Oscillations' puzzle connected with the magnetic field that guides the heavy-ion motion in the ESR, requiring a Dirac neutrino magnetic moment $\\mu_{\\nu}$ about six times smaller than the laboratory upper limit value from the Borexino Collaboration \\cite{Borexino08}. The underlying mechanism is the interference between decay amplitudes not affected by the static magnetic field of the ESR and decay amplitudes affected by this field which induces spin precession of the emitted neutrino. The motion of the recoil ion in the ESR is constrained by the interference long after the neutrino has fled away. This mechanism does not work for Majorana neutrinos that can have no {\\it diagonal} magnetic moments. For nonzero values of {\\it transition} magnetic moments, the resulting spin-flavor precession is suppressed by neutrino mass differences, and it becomes impossible to relate then the GSI Oscillations puzzle to magnetic effects. It is not yet resolved experimentally whether neutrinos are Dirac or Majorana fermions, although the theoretical bias rests with Majorana fermions, in which case the present paper accomplished nothing towards providing a credible explanation of this puzzle.  For experimental verification, note that the time-modulation period $T^{\\rm lab}_{\\rm EC}$ is inversely proportional to $B$, so the effect proposed here may be checked by varying $B$, for example by varying $\\beta=v/c$ for the coasting decay ions. For a fixed value of $\\beta$, $B$ depends on the charge-to-mass ratio of the decay ion which varies only to a few percent with the decay-ion mass $M_D$. Finally, the proposed effect is unique to two-body EC reactions, since three-body weak decays do not constrain the neutrino direction of motion with respect to the fixed direction of $\\vec B$. Indeed, preliminary data on the three-body $\\beta^+$ decay of $^{142}$Pm indicate no time modulation of the $\\beta^+$ decay rate, limiting its modulation amplitude to $a_{\\beta^+} < 0.03(3)$ \\cite{Kienle09}."
    },
    "0809/0809.3020_arXiv.txt": {
        "abstract": "We discuss observing strategy for the Near Infrared Coronagraphic Imager (NICI) on the 8-m Gemini South telescope.  NICI combines a number of techniques to attenuate starlight and suppress superspeckles: 1) coronagraphic imaging, 2) dual channel imaging for Spectral Differential Imaging (SDI) and 3) operation in a fixed Cassegrain rotator mode for Angular Differential Imaging (ADI).  NICI will be used both in service mode and for a dedicated 50 night planet search campaign. While all of these techniques have been used individually in large planet-finding surveys, this is the first time ADI and SDI will be used with a coronagraph in a large survey.  Thus, novel observing strategies are necessary to conduct a viable planet search campaign. ",
        "introduction": "\\label{sec:intro}  %  To date, numerous adaptive optics surveys have attempted to directly image young extrasolar planets in the near infrared\\cite{Biller07, Laf07a, Mas05}. While these surveys have placed strong statistical limits on extrasolar planet distributions, no planet has yet been imaged around a main sequence or pre-main sequence star.  Clearly, directly imaging planets is a nontrivial undertaking. It is important to weigh the survey instrument capability against theoretical predictions as to planet properties and distribution (e.g., from evolutionary models, radial velocity studies, etc.) and design a study which maximizes the chance of detecting a planet while also providing a scientifically interesting null result. In this paper, we will consider how best to maximize scientific output for a 50 night planet finding survey with the Near Infrared Coronagraphic Imager (NICI) at Gemini South.  NICI combines a number of techniques to attenuate starlight and suppress superspeckles: 1) coronagraphic imaging, 2) dual channel imaging for Spectral Differential Imaging (SDI)\\cite{Mar05,Biller07} and 3) operation in a fixed Cassegrain rotator mode for Angular Differential Imaging (ADI) a.k.a. roll subtraction\\cite{Liu04,Mar06}.  While coronagraphic imaging, SDI, and ADI have all been used individually in large planet-finding surveys, this is the first time they will all be used together in a large survey. Thus, it is important to carefully consider how to best carry out an observing campaign with NICI.  Both SDI and ADI techniques seek to separate real objects from speckles. SDI achieves this by exploiting a spectral feature in the desired target (previous implementations utilize the 1.6 $\\mu$m methane absorption feature robustly observed in substellar objects cooler than 1400 K). Images are taken simultaneously both within and outside the chosen absorption feature. Due to the simultaneity of the observations, the star and the coherent speckle pattern are largely identical in both filters, while any faint companion with that absorption feature is bright in one filter and faint in the other. Subtracting the two images thus removes the starlight and speckle patterns while a real companion with the chosen absorption feature remains in the image. In other words, the off-absorption band image acts as an ideal reference point spread function (henceforth PSF) for the absorption band image.  Utilizing a signature spectral feature of substellar objects greatly reduces the number of false positives detected (e.g. a background object, while real, will drop out of the SDI subtraction since it will not have methane absorption).  In previous surveys\\cite{Biller07}, SDI was used at two discrete rotator angles for each object -- a real object on the sky will modulate with a change of rotator angle; however, speckles which are instrumental phenomena will not. Thus, by subtracting data taken at multiple roll angles, speckles are further attenuated and any real object in the frame will display a characteristic jump in roll angle.  ADI employs a similar strategy to build a high quality reference PSF to remove speckles. Instead of observing at discrete roll angles, ADI leaves the rotator off and allows the telescope optics to rotate on the sky (this works best at Cassegrain focus). In a sequence of images taken at different parallactic angles, a real companion will track on the sky with the parallactic angle, while speckles will move randomly. From a series of images, a reference PSF can be constructed for and subtracted from each individual image, attenuating quasi-static speckle structure. Combining both SDI and ADI techniques thus allows an even greater degree of speckle supression.  SDI is the base mode of the NICI instrument. Each exposure produces two frames, one from each of the two channels. Usually, each channel is operated with a different filter, one off the absorption feature (usually on the blue side), and one on the feature (usually redward of the first filter). Whether SDI is combined with ADI or used at discrete roll angles is the choice of the observer, and is not necessarily a straightforward choice. For more details regarding the NICI instrument itself, see Chun et al.\\cite{Chun08}, this volume.  To develop the most appropriate observing strategy for the NICI science campaign, we must first determine a number of quantities (some inherent to the objects observed, some properties of the NICI instrument itself), including:  1) the sky coverage and relative performance of the ADI plus SDI and SDI discrete roll angles observation modes (covered in Section 2)  2) the inherent luminosity of young planets and the required $\\Delta$mag to image them (covered in Section 3.1)  3) the expected methane break strength for cool substellar objects and young planets, which will constrain how well our spectral difference works and how much of a $\\Delta$mag boost we can get from it (also covered in Section 3.2)  4) the achieved contrast of the NICI instrument as a function of separation from the primary star (relative $H$ band vs. $K$ band contrast discussed in Section 3.3, $H$ band contrast discussed in detail in Artigau et al.\\cite{Art08} in this volume)  5) the achieved boost in contrast compared to the standard AO due to the spectral subtraction or due to azimuthal differential imaging (discussed in Artigau et al.\\cite{Art08} in this volume)  6) the throughput of the NICI instrument at various wavelengths  In this document, we will therefore broadly discuss two major trades for the NICI campaign: operating with rotator off and on (ADI plus SDI vs. SDI discrete roll angles, respectively), and $H$ band vs. $K$ band performance. ",
        "conclusions": "Careful scheduling is necessary to use ADI with SDI, since both shear per image and total rotation on the sky need to be taken into account. There will always be parts of the sky where ADI is completely unfeasible. In areas of the sky with enough target rotation but not too high a rate of rotation, ADI plus SDI does seem to provide considerable enhancements in speckle suppression over SDI at discrete rotation angles, although part of the difference for this particular comparison dataset may have resulted from the fact that the SDI dataset at discrete rotation angles was non-optimal.  New, more realistic \"cold start\" models for young planets less than 300 Myr in age\\cite{For08} suggest planets are fainter than originally predicted; thus, somewhat higher contrasts will be necessary to detect them (i.e., 12-16 mag vs. 8-12). Both theoretically necessary contrasts and predicted methane break strengths appear to be optimized in $K$ band relative to $H$ band. However, $H$ band contrast performance appears to be somewhat better than $K$ band inside 1\". Additionally, it is important to consider total NICI instrument throughput in $H$ and $K$ bands when choosing between these, which we have not discussed in this document."
    },
    "0809/0809.3008.txt": {
        "abstract": "We report  proper motion measurements for 425 late-type M, L and T dwarfs, 314 of which have been measured for the first time.  Combining these new proper motions with previously published measurements, we have amassed a sample of 869 M7-T8 dwarfs and combined parallax measurements or calculated spectrophotometric distances to compute tangential velocities for the entire sample.  We find that kinematics for the full and volume-limited 20 pc samples are consistent with those expected for the Galactic thin disk, with no significant differences between late-type M, L, and T dwarfs.  Applying an age-velocity relation we conclude that the average kinematic age of the 20 pc sample of ultracool dwarfs is older than recent kinematic estimates and more consistent with age results calculated with population synthesis models.  There is a statistically distinct population of high tangential velocity sources ($V_{tan}$  $>$ 100 km~s$^{-1}$)  whose kinematics suggest an even older population of ultracool dwarfs belonging to either the Galactic thick disk or halo.  We isolate subsets of the entire sample, including low surface gravity dwarfs, unusually blue L dwarfs, and photometric outliers in $J-K_{s}$ color and investigate their kinematics.  We find that the spectroscopically distinct class of unusually blue L dwarfs have kinematics clearly consistent with old age, implying that high surface gravity and/or low metallicity must be relevant to their spectral properties.  The low surface gravity dwarfs are kinematically younger than the overall population, and the kinematics of the red and blue ultracool dwarfs suggest ages that are younger and older than the full sample, respectively.   We also present a reduced proper motion diagram at 2MASS $K_{s}$ for the entire population and find that a limit of  $H_{K_{s}}$ $>$ 18 excludes M dwarfs from the L and T dwarf population regardless of near-infrared color, potentially enabling the identification of the coldest brown dwarfs in the absence of color information. ",
        "introduction": "Kinematic analyses of stars have played a fundamental role in shaping our picture of the Galaxy and its evolution.  From early investigations (e.g. Schwarzschild  1908 , Lindblad 1925, and Oort 1927) where the large scale structure of the Galactic disk was first explored, through more recent investigations (e.g. Gilmore $\\&$ Reid 1983, Gilmore Wyse $\\&$ Kuiken 1989,  Dehnen 1998, Famaey et al 2005) where the structure of the Galaxy was refined to include a thick disk and prominent features such as streams, moving groups, and superclusters, kinematics have played a vital role in understanding the Galactic origin, evolution, and structure.  Combining kinematics with spectral features, several groups have mapped out ages and metallicities for nearby F,G,K, and M stars (e.g. Nordstrum et al. 2004).  The ages of these stars have become an important constraint on the Galactic star formation history and their kinematics have become a vital probe for investigating membership in the young thin disk, intermediate aged thick disk or older halo portion of the Milky Way. \\\\ \\indent   One population that has yet to have its kinematics exploited is the recently discovered population of brown dwarfs.  These objects, which have core temperatures that do not support stable hydrogen fusion, occupy the late-type M through T dwarf spectral classifications (e.g. \\citealt{2005ARA&A..43..195K}).   Within the late-type M through mid L dwarf spectral types there is a mix of both brown dwarfs and very low mass stars and we collectively refer to these objects as ultra cool dwarfs or UCDs.   They emit the majority of their light in the infrared and thus were only discovered in large numbers with the advent of wide-field near-infrared imaging surveys such as the Two Micron All Sky Survey (hereafter 2MASS; Skrutskie et al 2006), the Deep Infrared Survey of the Southern Sky (hereafter DENIS; Epchtein et al. 1997) and the Sloan Digital Sky Survey (hereafter SDSS; York et al. 2000).  Their very recent discovery has largely precluded astrometric measurements which require several-year baselines to produce useful measurements. Therefore, while brown dwarfs appear to be comparable in number to stars (e.g., Reid et al 1999), their role in the structure of the Milky Way has yet to be explored.  In addition, the thermal evolution of brown dwarfs implies that there is no direct correlation between spectral type and mass, leading to a mass/age degeneracy which makes it difficult to study brown dwarfs as a population.  While some benchmark sources (e.g. cluster members, physical companions to bright stars) have independent age determinations, and spectroscopic analyses are beginning to enable individual mass and age constraints (e.g. Burgasser, Burrows \\& Kirkpatrick 2006; Saumon et al. 2007), the majority of brown dwarfs are not sufficiently characterized to break this degeneracy.  Kinematics can be used as an alternate predictor for the age of the brown dwarf population. \\\\ \\indent Moreover,  kinematics can also be used to characterize subsets of UCDs.  With hundreds of brown dwarfs now known\\footnote{http://www.dwarfarchives.org}, groupings of peculiar objects \\-- sources whose photometric or spectroscopic properties differ consistently from the majority of the population \\-- are becoming distinguishable.  Currently defined subgroups of late-type M, L and T dwarfs include 1) low surface gravity objects (e.g. \\citealt{2004ApJ...600.1020M}, \\citealt{2006ApJ...639.1120K}, \\citealt{2007ApJ...657..511A}, \\citealt{2007AJ....133..439C}), 2) metal poor subdwarfs (e.g. \\citealt{2003ApJ...592.1186B}, \\citealt{2003ApJ...591L..49L}, \\citealt{2006AJ....132.2372G}, \\citealt{2007ApJ...658..557B}), 3) unusually blue L dwarfs (hereafter UBLs; e.g. \\citealt{2003AJ....126.2421C}, \\citealt{2007AJ....133..439C}; \\citealt{2004AJ....127.3553K}, \\citealt{2006AJ....131.2722C}), and 4) unusually red and possibly dusty L dwarfs (e.g. McLean et al 2003).  While observational peculiarities can overlap between these groups (e.g. both young and dusty L dwarfs can be unusually red), they appear to encompass objects with distinct physical traits (e.g., mass, age, composition, and cloud properties) so they are important for drawing a connection between observational characteristics and intrinsic physical properties.  Kinematics can be used to investigate the underlying physical causes for the peculiarities of these groups. \\\\ \\indent \tIn the past decade a number of groups have conducted astrometric surveys of substellar mass objects and have studied subsets of low mass objects  (e.g.  \\citealt{2004AJ....127.2948V}, \\citealt{2002AJ....124.1170D}, \\citealt{2000AJ....120.1085G}, \\citealt{2003AJ....126..975T}, \\citealt{2007AJ....133.2258S}, \\citealt{2007arXiv0710.4786J}, and \\citealt{2007ApJ...666.1205Z}). We have initiated the Brown Dwarf Kinematics  Project (BDKP) which aims to measure the 6D positions and velocities of all known brown dwarfs within 20 pc of the Sun and select sources of scientific interest (e.g. low gravity dwarfs, subdwarfs).  In this article, which will be the first in a series of BDKP papers, we add 314 new proper motion measurements and combine all published proper motion measurements and distance estimates into a uniform sample to examine the brown dwarf population as a whole.  Section 2 of this paper outlines the observed sample and describes how proper motion measurements were made.  Section 3 discusses the expanded sample and how distances and $V_{tan}$ measurements were calculated.   Section 4 examines the full astrometric sample and subsets. Section 5 reviews the high tangential velocity objects in detail.  Finally, section 6 applies an age-velocity relation to the sample and examines resultant ages of the full sample and red/blue outliers. ",
        "conclusions": "We have presented new proper motions for 425 late-type M, L, and T dwarfs and combined all previous proper motion measurements with either parallax measurements or spectrophotometric distances  to compute tangential velocities.  We derive average kinematic and photometric values for individual spectral types as well as for the late-type M,L, and  T populations as a whole.  We conduct a crude U,V,W analysis and find that the full and 20 pc samples examined in this article have space velocities consistent with the Galactic thin disk population.  However, there are 13 objects in the ultracool dwarf population that lie at the tail end of the velocity distribution and are likely to be part of an older Galactic population.  We find a large difference in the kinematics between red and blue photometric outliers and conclude that their velocity dispersions are kinematically distinct from the full or 20 pc samples .  Analysis of the low surface gravity and UBL subgroups also show a distinction from the full and 20 pc samples.  Applying an age-velocity relation we conclude that the red outliers and low surface gravity subgroups are younger than the full and 20 pc samples and the blue outliers and UBLs are older. We conclude that an age-color relation can be derived for the ultracool dwarf population.  We report ages for the 20 pc sample of this kinematic study that are consistent with the 3-6 Gyr values derived in population synthesis models and propose that one reason for prior kinematic reports of mean ages for the L and T dwarf population of $\\sim$1 Gyr is due to small number statistics or a bias in the sample analyzed.  This work is the first paper of the newly initiated Brown Dwarf Kinematics Project (BDKP) where we seek to measure full space motions for the 20 pc sample of ultracool dwarfs with priority given to the lowest temperature L and T dwarfs.  We have two parallax programs underway at CTIO using the ANDICAM IR imager on the 1.3m telescope for brighter targets and the ISPI IR imager on the 4.0m telescope for the fainter targets.  We have plans to obtain radial velocities for a large number of ultracool dwarfs upon completion of FIRE, a high resolution near-IR spectrograph."
    },
    "0809/0809.3350_arXiv.txt": {
        "abstract": "We introduce a new space VLBI project, the Second VLBI Space Observatory Program (VSOP2),  following the success of the VLBI Space Observatory Program (VSOP1). VSOP2 has 10 times higher angular resolution, up to about 40 micro arcseconds, 10 times higher frequency up to 43 GHz, and 10 times higher sensitivity compared to VSOP1. Then VSOP2 should become a most powerful tool to observe innermost regions of AGN and astronomical masers. \\textit{ASTRO-G} is a spacecraft for VSOP2 project constructing in ISAS/JAXA since July 2007. \\textit{ASTRO-G} will be launched by JAXA % H-IIA % rocket in fiscal year 2012. \\textit{ASTRO-G} and ground-based facilities are combined as VSOP2. To achieve the good observation performances, we must realize new technologies. They are large precision antenna, fast-position switching capability, new LNAs, and ultra wide-band down link, etc.. VSOP2 is a huge observation system involving  \\textit{ASTRO-G}, ground radio telescopes, tracking stations, and correlators, one institute can not prepare a whole system of VSOP2. Then we must need close international collaboration to get sufficient quality of resultant maps and  to give a sufficient quantity of observation time  for astronomical community.  We  formed a  new international council  to provide guidance on scientific aspects related of  VSOP2, currently called the VSOP2 International Science Council (VISC2). ",
        "introduction": "The pioneering history of radio telescopes is in some ways the history of the quest for higher angular resolution.  The angular resolution of  a radio interferometer is given by $\\lambda/D$, where $D$ is baseline length or antenna spacing. The upper limit of $D$ was usually much less than 1,000 km for connected interferometers. In order to push this limit, VLBI (Very-Long-Baseline Interferometry) was invented. VLBI is made by connecting existing radio telescopes with magnetic tapes instead of cables. Then VLBI observes radio sources with relatively high brightness temperature like AGN and astronomical masers.  Superluminal motion of AGN jet was shown only by very high angular resolution of VLBI. However, ground-based VLBI has an upper limit of angular resolution because $D$ must not be larger than the diameter of the earth, 13,000 km. In order to push this limit, space VLBI was formulated. Space VLBI is an interferometer connecting radio telescopes in orbit and ground-base ones.  Since the baseline length of space VLBI can extend beyond the diameter of the earth, angular resolution of space VLBI is free from this limit.  However, it is very difficult to set a radio telescope in orbit.  Space VLBI was realized in 1997 as the first space VLBI project, VSOP1 (VLBI Space Observatory Program)[1,2],  after some preceding experiments and vast amounts of effort. The spacecraft for VSOP1 is known as the first dedicated space VLBI satellite, \\textit{HALCA}. This was launched by a M-V rocket of Institute of Space and Astronautical Science (ISAS) from Unchinoura Space Center (USC). The angular resolution of VSOP1 is 3 times better than that of ground-based VLBI at 1.6 and 5 GHz (VSOP1 also had 22-GHz band system, which unfortunately suffered a serious loss of sensitivity at the launch). Seven hundred observations of AGN jets were made in the mission lifetime of 7 years.  Figure 1 shows a schematic display of the comparison between space VLBI and ground-based VLBI.  The Second VLBI Space Observatory Program (VSOP2) is a successor of VSOP1 [3].  \\textit{ASTRO-G} is the satellite for VSOP2, which will be launched by ISAS/JAXA in fiscal year 2012. The construction project of \\textit{ASTRO-G} has been started since July 2007. This project is a joint-project between ISAS, NAOJ (National Astronomical Observatory, Japan) and the universities. Now, basic design of the satellite and developments of several components for \\textit{ASTRO-G} are ongoing.  A review of a basic design of the satellite (PDR: preliminary design review) is scheduled at March of 2009.  After decision of the design (CDR: critical design review) in 2009, we will start the production of the flight model of \\textit{ASTRO-G}. This satellite will be launched with an H-IIA rocket from Tanegashima Space Center (TNSC) of JAXA.  VSOP2 has 10 times higher angular resolution, 10 times higher frequency, and 10 times higher sensitivity compared to VSOP1. VSOP2 will be an astronomical tool with unprecedented angular resolution in all wavelength to explore innermost regions of AGN and astronomical masers. The VSOP2 programme foresees a powerful astronomical instrument involving  \\textit{ASTRO-G}, ground radio telescopes, tracking stations, and correlators.  For this reason, ISAS alone can not all different aspects of this complex project, which will be made in collaboration with international space and ground-base partners.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=28pc]{Fig1.eps} \\hspace{2pc} \\caption{\\label{label}Schematic display of the comparison between space VLBI and ground-based VLBI.} \\end{center} \\end{figure}  \\begin{figure}[h] \\begin{center} \\includegraphics[width=28pc]{Fig2.eps} \\hspace{2pc} \\caption{\\label{label}Artist's impression of the \\textit{ASTRO-G} satellite around perigee.} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0809/0809.3399_arXiv.txt": {
        "abstract": "{% We present binary galaxy merger simulations of gas-rich disks (Sp-Sp), of early-type galaxies and disks (E-Sp, mixed mergers), and mergers of early-type galaxies (E-E, dry mergers) with varying mass ratios and different progenitor morphologies. The simulations include radiative cooling, star formation and black hole (BH) accretion and the associated feedback processes. We find for Sp-Sp mergers, that the peak star formation rate and BH accretion rate decrease and the growth timescales of the central black holes and newly formed stars increase with higher progenitor mass ratios. The termination of star formation by BH feedback in disk mergers is significantly less important for higher progenitor mass ratios (e.g. 3:1 and higher). In addition, the inclusion of BH feedback suppresses efficiently star formation in dry E-E mergers and mixed E-Sp mergers.} ",
        "introduction": "Recent large-scale statistical galaxy surveys (e.g. the Sloan Digital Sky Survey (SDSS) and the two-degree Field (2df) survey) have revealed a robust bimodality in the observed galaxy population. Galaxies above a critical stellar mass of $M_{\\rm crit}\\simeq 3\\times 10^{10} M_{\\odot}$ are typically non-star forming red spheroidal systems with old stellar populations that predominately live in dense environments, whereas galaxies below this critical mass are typically blue, star-forming disk galaxies that lie in the field (e.g. Bell et al. 2003; Kauffmann et al. 2003; Baldry et al. 2004; Balogh et al. 2004).  The observed bimodality can also be seen as a manifestation of cosmic downsizing, in which the most massive galaxies formed a significant proportion of their stars at high redshifts above $z \\gtrsim 2$. The star formation is efficiently quenched in these systems by $z=1$ as manifested in the observed population of extremely red non-starforming massive ellipticals (EROs, e.g. McCarthy 2004; V\\\"ais\\\"anen \\& Johansson 2004). The lower mass systems, on the contrary, have typically been forming stars throughout the whole cosmic epoch (e.g. Glazebrook et al. 2004; Juneau et al. 2005; Cimatti, Daddi, Renzini 2006).   The quenching mechanism responsible for the observed termination of star formation in massive galaxies at redshifts below $z \\lesssim 2-3$ needs to be both energetic enough to trigger the quenching and long-lasting enough to maintain the quenching over a Hubble time. There exists several theoretical explanations for the quenching, including the feedback from AGNs (e.g. Bower et al. 2006; Croton et al. 2006), gaseous major mergers triggering star burst and/or quasar activity (e.g Naab, Jesseit, Burkert, 2006), quenching by shockheated gas above a critical halo mass (Dekel \\& Birnboim 2006; Birnboim, Dekel, Neistein 2007) and gravitational quenching by clumpy accretion (Naab et al. 2007; Dekel \\& Birnboim 2008).  In this paper we study the termination of star formation by black hole (BH) feedback in binary galaxy mergers. The study by Springel, Di Matteo, Hernquist (2005a) showed that BH feedback efficiently terminates star formation in equal-mass disk mergers. Here we confirm this result and expand the analysis to include unequal-mass disk mergers, mixed E-Sp mergers and dry E-E mergers. This paper is structured as follows. In \\S \\ref{Sims} we briefly describe our simulation code and galaxy merger setup. We discuss the effect of varying the initial gas fractions, orbits, merger mass ratios and galaxy progenitors for terminating star formation in \\S \\ref{Results}. Finally, we summarize and discuss our findings \\S \\ref{Conc}. ",
        "conclusions": "\\label{Conc} In this paper we have studied the termination of star formation in merger simulations including BH feedback. We find that the termination of star formation by BH feedback is significantly less important for unequal-mass disk mergers compared to equal-mass disk mergers. The timescale for star formation termination systematically increases with increasing progenitor mass ratios. Similarly, a systematic increase is seen in the half-mass growth timescales of the BHs, with this timescale varying from $\\sim0.1$ Gyr for equal-mass mergers to $\\sim 1$ Gyr for 6:1 mergers. This systematic trend can be used as input in modeling BH accretion more realistically in semi-analytic galaxy formation models (e.g. Croton et al. 2006). For mass-ratios of 3:1 and higher mergers with BH feedback are unable to completely quench the star formation, with the merger remnants showing star formation rates roughly on the pre-merger level even 1 Gyr after completion of the merger. In addition, the star formation is efficiently terminated in mixed E-Sp and dry E-E mergers due to the presence of the fully grown super-massive BHs in the early-type progenitors."
    },
    "0809/0809.2302.txt": {
        "abstract": "Some general properties of the advective-acoustic instability are described and understood using a toy model which is simple enough to allow for analytical estimates of the eigenfrequencies. The essential ingredients of this model, in the unperturbed regime, are a stationary shock and a subsonic region of deceleration. For the sake of analytical simplicity, the 2D unperturbed flow is parallel and the deceleration is produced adiabatically by an external potential. The instability mechanism is determined unambiguously as the consequence of a cycle between advected and acoustic perturbations. The purely acoustic cycle, considered alone, is proven to be stable in this flow. Its contribution to the instability can be either constructive or destructive. A frequency cut-off is associated to the advection time through the region of deceleration. This cut-off frequency explains why the instability favours eigenmodes with a low frequency and a large horizontal wavelength. The relation between the instability occurring in this highly simplified toy model and the properties of SASI observed in the numerical simulations of stellar core-collapse is discussed. This simple set up is proposed as a benchmark test to evaluate the accuracy, in the linear regime, of numerical simulations involving this instability. We illustrate such benchmark simulations in a companion paper. ",
        "introduction": "In the scenario proposed by Bethe \\& Wilson (1985) for core collapse supernovae, the success of the explosion depends on the efficiency of energy deposition by neutrinos below the stalled accretion shock, during the first second after core bounce. Numerical simulations revealed that this mechanism is inefficient in spherical symmetry  (Liebend\\\" orfer et al. 2001). Hydrodynamical instabilities may play an important role by breaking the spherical symmetry and helping the revival of the stalled shock. Indeed, multidimensional simulations allowing for transverse motions, induced by neutrino-driven convection in the gain region, approached the explosion threshold (Burrows et al. 1995, Janka \\& M\\\"uller 1996), although this effect did not seem sufficient (Buras et al. 2003). The discovery of another hydrodynamical instability, named the Standing Accretion Shock Instability (SASI) by Blondin et al. (2003), opened new perspectives which seem very promising for the revival of the shock. The recent simulations of Marek \\& Janka (2007) showed a successful explosion of a $15 M_{\\sun}$ progenitor where neutrino energy deposition is efficient enough owing to the effect of SASI, which lengthens the time spent by the postshock gas in the gain region (Murphy \\& Burrows 2008). SASI is also the starting point of the acoustic mechanism found by Burrows et al. (2006, 2007), where  the excitation of g-mode oscillations inside the proto-neutron star triggers the emission of acoustic waves (see also Weinberg \\& Quataert 2008). Besides the question of the explosion mechanism, the development of SASI seems to have important consequences on the birth conditions of the neutron star, its kick (Scheck et al. 2004, 2006) and also its spin (Blondin \\& Mezzacappa 2007, Yamasaki \\& Foglizzo 2008). \\\\ In view of the spectacular possible consequences of SASI, a fundamental understanding of its mechanism is desired but still a source of debate. It is definitely distinct from convection since it can take place even without any source of heating (Blondin et al. 2003). The ``advective-acoustic cycle'' (Foglizzo \\& Tagger 2000, Foglizzo 2001, Foglizzo 2002) was recognized as the driving mechanism by several authors (Blondin et al. 2003, Burrows et al. 2006, Ohnishi et al. 2006, Foglizzo et al. 2007, Scheck et al. 2008, Yamasaki \\& Foglizzo 2008), but some alternate interpretations were proposed (Blondin \\& Mezzacappa 2006, Blondin \\& Shaw 2007). It should be noted that the analytical arguments of Laming (2007) supporting the existence of a purely acoustic mechanism contained some errors, and the corrected calculation favours the advective-acoustic mechanism (Laming 2008, in press). Part of the difficulty in recognizing the advective-acoustic mechanism in numerical simulations, even in the simplified set up proposed by Blondin et al. (2003), comes from the lack of simple reference models where its properties would be fully understood. The present work aims at providing such a reference, as simple as possible. This reference serves two purposes. It can help us build our physical intuition about the advective-acoustic coupling responsible for an unstable advective-acoustic cycle. It can also be used as a benchmark test to evaluate the accuracy of multidimensional numerical simulations in the linear phase of the instability (Sato et al. 2008, hereafter paper~II).\\\\ Although some of the equations used in the present work have already been described in past publications dealing with more complex flows (e.g. Yamasaki \\& Foglizzo 2008 and references therein), we choose to explain them again in the simpler framework considered here. Analytical calculations are made possible by three main simplifications: (i) the flow is 2D planar, (ii) the region of deceleration is localized spatially at a distance from the shock, and (iii) the deceleration is produced adiabatically. The simplicity of the present set up allows for an exact calculation, even at low frequency. The physical insight in the mechanism at work can be used as a reference, in order to analyze more realistic situations.\\\\ The paper is organized as follows: in Sect.~2, the toy model consisting of a stationary shock in an external step-like potential is described, and its eigenspectrum is computed as an illustrative example. Sect.~3 explains the method used to prove unambiguously the mechanism at work, beyond the determination of the eigenfrequencies. This method requires the determination of the efficiency of the advective-acoustic coupling at the shock, in Sect.~4, and in the region of acoustic feedback in Sect.~5. These local wave coupling efficiencies are combined into a global cycle whose growth rate is computed in Sect.~6 and compared to the direct eigenspectrum of Sect.~2. The importance of the size of the deceleration region is emphasized in both Sect.~5 and 6. The relevance of these results to the problem of core-collapse is discussed in Sect. 7. Results are summarized in Sect.~8. Lengthy analytical derivations are described in Appendices for the sake of clarity. ",
        "conclusions": " \\par -The instability mechanism is based on the amplification of the advective-acoustic cycle. \\par -The purely acoustic cycle is always stable in this flow. \\par -Optimal transverse perturbations are more unstable than longitudinal ones. \\par -The finite size of the coupling region is responsible for a frequency cut-off which favours low frequency perturbations. \\par -This frequency cut-off disfavours high order perturbations, which would lead to an inefficient evanescent pressure feedback: the most unstable modes are expected to be low frequency and low order. \\par -The maximum efficiency $|{\\cal Q}|$ of the advective-acoustic cycle is governed by a compromise between the efficiency of the advective-acoustic coupling, which favours transverse perturbations, and a short cycle time which disfavours too horizontal acoustic propagation. \\par -The eigenspectrum is irregular and sensitive to geometrical factors due to the phase condition which selects a discrete set of frequencies, and the interferences with the acoustic cycle. \\par -A new method has been described to measure the eigenspectrum associated to each cycle considered alone. \\par -This toy model can be used as a benchmark test for numerical simulations, in order to check for undesired numerical artifacts. The important roles played by the shock, the advection of perturbations and the propagation of acoustic waves make this toy model particularly challenging for numerical simulations (paper~II). \\par -The comparison with the properties of SASI suggests that our toy model captures some fundamental features, explaining in particular why SASI is a low frequency, low degree instability.\\\\ A good understanding of the mechanism behind SASI should help perform accurate numerical simulations of core collapse explosions, with particular attention to the mesh size and the boundary conditions."
    },
    "0809/0809.2095_arXiv.txt": {
        "abstract": "In recent years many binary minor planets (BMPs) have been discovered in the Solar system. Many models have been suggested for their formation, but these encounter difficulties explaining their observed characteristics. Here we show that secular perturbations by the Sun (Kozai mechanism) fundamentally change the evolution and the initial distribution of BMPs predicted by such models and lead to unique observational signatures. The Kozai mechanism can lead to a large periodic oscillations in the eccentricity and inclination of highly inclined BMP orbits, where we predict such effects to be observable with current accuracy within a few years (e.g. for the binary asteroid Huenna). In addition, the combined effects of the Kozai mechanism and tidal friction (KCTF) drives BMPs into short period circular orbits. We predict a specific inclination dependent distribution of the separation and eccentricity of BMPs, due to these effects, including a zone of avoidance at the highest inclinations. Specifically the Kozai evolution could explain the recently observed peculiar orbit of the Kuiper belt binary 2001 QW$_{322}$ . Additionally, the KCTF process could lead to BMPs coalescence and serve as an important route for the formation of irregular shaped single minor planets with large axial tilts. ",
        "introduction": "Stable gravitational triple systems require a hierarchical configuration, in which two objects orbit each other in a relatively tight {}``inner binary'', and the third object orbits the binary in a wider {}``outer binary''. Although such triples are stable against disruption, their orbits may change shape and orientation on time scales much longer than their orbital period. In particular, the Kozai-Lidov mechanism\\citep{1962K,1962L} predicts a secular perturbations of the inner binary orbit. So-called {}``Kozai oscillations'' cause the eccentricity and the inclination to fluctuate. The Kozai mechanism is known to be highly important in the evolution of many triple systems\\citep{1962K,1979MS,1968H,1998KEM,car+02,nes+03,2007FT}. It leads to a large (order unity) periodic oscillations (Kozai cycles) in the eccentricity and inclination of inner binaries with high inclinations with respect to the outer binary orbit (hereafter the relative inclination).  In recent years many binary minor planets {[}BMPs; both binary asteroids and binary trans-Neptunian objects (TNOs){]} have been discovered in the Solar system\\citep{ric+06}. BMPs can be regarded as the inner binary members of a triple system in which the sun is the perturbing companion in the outer orbit. Here we study the importance of such perturbations and show that BMPs are susceptible to Kozai oscillations, which play a major role in their evolution. The combined effects of the Kozai mechanism in addition to tidal friction, (Kozai cycles and tidal friction; KCTF\\citep{1979MS,1998KEM}) which becomes important for BMPs with high eccentricities (induced by the Kozai mechanism), change the orbital parameters of the BMPs. These effects could then erase the observational signatures suggested by the various formation scenarios of BMPs discussed in the literature\\citep{wei+89,wei02,mer+02,gol+02,fun+04,ric+06,lee+07}. replacing them with unique and different signatures, that are consistent with current observations. ",
        "conclusions": "\\subsection{KCTF observational signature and BMPs formations scenario}  The Kozai, KCTF and Kozai-Hill processes could appreciably affect the survival of BMPs, their orbits and their orbital phase-space distribution. In fig.~\\ref{f:inclination-separation} we show typical examples of the inclination-SMA separation phase-space in which the different Kozai processes take place. The positions of all BMPs with known inclinations in this phase-space are also shown (detailed discussion of the observed inclinations of BMPs is given elsewhere \\citealp{per+08}). BMPs in the Kozai region are not susceptible to tidal friction during the Kozai cycle. For such BMPs the eccentricity and inclinations continuously and periodically change between the minimal and maximal Kozai eccentricity (see e.g. fig.~\\ref{f:eccentricity-evolution}) and should therefore have a wide distribution of eccentricities which would be very different from the initial conditions set by the formation mechanism of the BMPs. The initial SMA, however, will be conserved in this case, still reflecting the initial conditions.  We note that the current observations (although currently only a small sample exists) are in good agreement with our predictions. All the BMPs in the Kozai region have relatively high eccentricities (none are circular, defined here as $e<0.05$). The zone of avoidance at the highest inclinations in the KCTF region is empty, and BMPs in the KCTF region tend to be circular and with smaller SMA separations. The large difference between binary TNOs and binary asteroids is likely to be related to different formation mechanisms\\citep{ric+06} . Note that BMPs formed only at small separations at which tidal friction is important would never be subjected to the pure Kozai mechanism.  The above predictions should be of major concern when trying to constrain BMPs' formation scenarios by observations. Many formations scenarios for BMPs have been described in the literature (see ref. \\citep{ric+06} for a review) with different predictions for the distribution of BMP SMA separations, eccentricities and inclinations (unfortunately, only few studies explored inclination distribution\\citep{naz+07,sch+08} ; given the importance of the Kozai effect we strongly suggest to explore this in future studies). As we have shown the initial conditions as prescribed by different scenarios will be fundamentally changed by the Kozai and KCTF evolution. In order to find evidence for specific formation scenarios of BMPs, it is therefore essential to take the Kozai evolution of the BMPs into account. For example, typical high eccentricity BMPs produced in some scenarios such as the expected $e>0.8$ in the exchange scenario of Funato et al.\\citep{fun+04} , or $0.2\\lesssim e\\lesssim0.8$ in the chaos assisted capture scenario\\citep{lee+07} , could be observed with very low eccentricities, either due to the random phase in which they are observed during their Kozai cycle or due to their KCTF evolution which would have circularized their orbit (e.g. fig. \\ref{f:inclination-separation}). Recently (and after the initial presentation of this manuscript), \\citet{pet+08} have reported on the peculiar orbit of the Kuiper Belt Binary 2001 QW$_{322}$, with very large SMA but relatively small eccentricity. Such orbit is difficult to explain through previously suggested BMPs formation scenarios, since a binary with such wide SMA is likely to be produced through the exchange scenario scenario, which would typically produce a highly eccentric binary \\citep{fun+04}. Such orbit, however, is naturally produced through the Kozai mechanism presented here. As shown in fig. \\ref{f:eccentricity-evolution}, the orbit of a binary similar to 2001 QW$_{322}$ could evolve due to the Kozai mechanism from initially high eccentricity and small pericenter distance separation (in fact comparable to the separation of the \\textquotedblleft{}regular\\textquotedblright{} Kuiper belt object 2001 QT$_{297}$) to have a low eccentricity and much larger pericenter distance separation, as observed today. Such a binary could therefore have been formed through a binary-single encounter, and then naturally evolve to its current state through Kozai evolution.  Similarly, typical low eccentricity BMPs, such as might be expected to be produced in the dynamical friction capture scenario\\citep{gol+02} , could be observed with very high eccentricities. % Since these effects occur mainly for BMPs in the main Kozai region, we predict an inclination and SMA separation dependence of the eccentricity distribution of BMPs, with a transition between BMPs at the different regimes described in fig.~\\ref{f:inclination-separation}. Similarly the SMA separation distribution would also be changed between the different regimes. In fact, only BMPs outside the Kozai region would preserve the eccentricity and SMA separation signature of their formation scenario, whereas in the other regimes these distributions would be dominated by the Kozai mechanism.  In addition to these predictions, we note the importance of environmental effects when combined with Kozai evolution. In some regions of the solar system BMPs were likely to form and/or to evolve for some time in dense environments where they encountered other minor planets (as suggested or required by most BMP formation and evolutionary scenarios). In such an environment, encounters may change the orbital configuration of the binaries in quite a chaotic way\\citep{heg75} especially for the wider binaries that are more susceptible to encounters (larger cross section). The Kozai timescales (Eq.~\\ref{eq:Pk}) could be shorter than the typical timescale between consequent encounters. In such cases a BMP entering the KCTF regime for the first time following an encounter, could rapidly migrate to smaller SMA separation, thus decreasing its chance for further encounters which could have otherwise potentially export it outside this phase space region. This KCTF region therefore produces a sink in the phase-space distribution and induces a flow from small inclinations to high inclinations in the phase-space distribution of BMPs. We predict that the KCTF mechanism, when at work in a collisional environment, would form a KCTF-collisional signature in which the KCTF and the phase-space regions close to it (i.e. mostly the Kozai region) would be overpopulated relative to the other phase space regions (beside the zone of avoidance which should be depleted of BMPs). The distribution of observed binary TNOs is suggestive of such a signature (fig. \\ref{f:inclination-separation}), with all BMPs at high inclinations, in addition to an empty region at the zone of avoidance, but future observations are required in order to show this at high significance.   \\subsection{Binary spin-orbit correlation and irregularly shaped minor planets }  Since KCTF evolution drives BMPs into close configurations, they may become close enough as to evolve into contact configuration or even mergers (similar suggestion was discussed in the context of merger of inner binary in triple stellar systems due to KCTF; \\citealp{per+09}). These contact/merger products may still show a signature of their KCTF formation scenario producing a relative spin-orbit inclination which is likely to be limited to Kozai inclinations. In addition this Kozai-induced merger scenario could naturally explain the existence of large single minor planets with irregular prolonged shapes (e.g. \\citep{har+80,rom+01} ) and predict their prolonged axis to be aligned perpendicular to their spin axis (where the latter is expected to have the same biased spin-orbit inclination distribution as the Kozai contact binaries). We note that rotational disruption and direct collisions could also produce irregularly shaped minor planets, but are not expected to produce a bias towards high latitude inclinations. The YORP effect could also form a bias toward specific high inclinations, however this process is efficient only for small asteroids (diameter smaller than $50$~km; \\citealp[e.g. ][and references there in]{pol+09}) , and is not expected to affect TNOs at all. If such mergers products are abundant they could effect the relative spin-orbit inclination distribution, and produce tilt axis distribution which is biased towards high latitudes from the orbital plane. Interestingly, such a bias exists in the pole distribution of large binary asteroids \\citep[; not expected to be affected by YORP]{kry+07}."
    },
    "0809/0809.0090_arXiv.txt": {
        "abstract": "{The young $\\sigma$~Orionis cluster is an important location for understanding the formation and evolution of stars, brown dwarfs, and planetary-mass objects. Its metallicity, although being a fundamental parameter, has not been well determined yet.} {We present the first determination of the metallicity of nine young late-type stars in $\\sigma$~Orionis.} {Using the optical and near-infrared broadband photometry available in the literature we derive the effective temperatures for these nine cluster stars, which lie in the interval 4300--6500 K (1--3 \\Msuno). These parameters are employed to compute a grid of synthetic spectra based on the code MOOG and Kurucz model atmospheres. We employ a $\\chi^2$-minimization procedure to derive the stellar surface gravity and atmospheric abundances of Al, Ca, Si, Fe, Ni and Li, using multi-object optical spectroscopy taken with WYFFOS+AF2 at at the William Herschel Telescope ( $\\lambda/\\delta\\lambda\\sim7500$).} {The average metallicity of the $\\sigma$~Orionis cluster is [Fe/H] $ = -0.02\\pm0.09\\pm0.13$ (random and systematic errors). The abundances of the other elements, except lithium, seem to be consistent with solar values. Lithium abundances are in agreement with the ``cosmic'' $^7$Li abundance, except for two stars which show a $\\log \\epsilon(\\mathrm{Li})$ in the range 3.6--3.7 (although almost consistent within the error bars). There are also other two stars with $\\log \\epsilon(\\mathrm{Li})\\sim 2.75$. We derived an average radial velocity of the $\\sigma$~Orionis cluster of $28\\pm4$\\,\\kmso.} {The $\\sigma$~Orionis metallicity is roughly solar.} ",
        "introduction": "\\label{introduction}  The \\object{$\\sigma$~Orionis cluster} is a nearby, very young region that has become a suitable place for searching for and characterizing substellar objects. The age has been estimated at $3 \\pm 2$\\,Ma (Oliveira et~al. 2002; Zapatero Osorio et~al. 2002a; Sherry et~al. 2004) and its heliocentric distance is roughly $360^{+70}_{-60}$\\,pc (Brown et~al. 1994). The cluster is relatively free of extinction ($A_V \\ll$ 1\\,mag -- Lee 1968; B\\'ejar et~al. 2004) and has a moderate spatial member density (Caballero 2008a) and {\\bf a large frequency of intermediate-mass stars with discs}. A compilation of different determinations of the age, distance, and disc frequency at different mass intervals is provided in Caballero (2007). The cluster contains several dozen brown dwarfs with spectroscopic features of youth and with discs (B\\'ejar et~al. 1999; Zapatero Osorio et~al. 2002a; Barrado~y~Navascu\\'es et~al. 2002; Muzerolle et~al. 2003; Kenyon et~al. 2005; Burningham et~al. 2005; Caballero et~al. 2006, 2007). It is also the star forming region with the largest number of candidate isolated planetary-mass objects with follow-up spectroscopy (Zapatero Osorio et~al. 2000, 2002b,c; Mart{\\'\\i}n et~al. 2001; Barrado~y~Navascu\\'es et~al. 2001; Mart{\\'\\i}n \\& Zapatero Osorio~2003). Some of these planetary-mass objects have discs (Zapatero Osorio et al. 2007; Scholz \\& Eisl\\\"offel 2008; Luhman et al. 2008).  \\begin{table*} \\caption[]{Identification number in Caballero (2006), 2MASS coordinates, ASAS-3 $V$, DENIS $i$, and 2MASS $JHK_{\\rm s}$ magnitudes of the 24 pre-selected stars.} \\label{table.coo+phot+teff} $$ \\begin{tabular}{lccccccc} \\hline \\hline \\noalign{\\smallskip} SO\t& $\\alpha$ \t& $\\delta$\t& $V$ \t\t\t& $i$ \t\t\t& $J$  \t\t\t& $H$  \t\t\t& $K_{\\rm s}$  \t          \\\\ & (J2000)\t& (J2000)\t& [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t\t& [mag]\t\t          \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 230838\t& 05 37 36.86\t& --02 08 17.7\t&  $9.898\\pm0.082$\t&  $9.420\\pm1.00$\t&  $8.690\\pm0.029$\t&  $8.144\\pm0.040$\t&  $7.573\\pm0.027$\t   \t  \\\\ % 000041  & 05 40 27.55\t& --02 25 43.1\t& $10.117\\pm0.020$\t&  $9.765\\pm0.04$\t&  $9.062\\pm0.032$\t&  $8.872\\pm0.084$\t&  $8.775\\pm0.024$        \\\\ % 410305  & 05 40 12.09\t& --03 05 28.4\t& $11.035\\pm0.031$\t& $10.565\\pm0.06$\t&  $9.905\\pm0.026$\t&  $9.634\\pm0.023$\t&  $9.607\\pm0.025$        \\\\ % 210734 \t& 05 37 53.04\t& --02 33 34.4\t& $11.649\\pm0.059$\t& $10.798\\pm1.00$\t&  $9.991\\pm0.027$\t&  $9.605\\pm0.024$\t&  $9.474\\pm0.025$\t   \\\\ % 131161 \t& 05 38 56.81\t& --02 04 54.5\t& $12.343\\pm0.112$\t& $11.147\\pm0.04$\t& $10.330\\pm0.027$\t&  $9.763\\pm0.024$\t&  $9.619\\pm0.023$\t   \\\\ % 420039  & 05 38 59.55\t& --02 45 08.1\t& $12.520\\pm0.101$\t& $11.686\\pm0.03$\t& $11.222\\pm0.027$\t& $10.984\\pm0.023$\t& $10.893\\pm0.027$\t   \\\\ % 210434 \t& 05 38 07.85\t& --02 31 31.4\t& $12.715\\pm0.112$\t& $11.634\\pm1.00$\t& $10.566\\pm0.027$\t&  $9.930\\pm0.023$\t&  $9.769\\pm0.025$\t   \\\\ % 000005$^{a}$ \t& 05 38 38.23\t& --02 36 38.4\t& $12.9-14.0$\t        & $12.340\\pm0.03$\t& $11.158\\pm0.026$\t& $10.465\\pm0.023$\t& $10.312\\pm0.022$\t   \t \\\\ % 211394  & 05 37 15.37\t& --02 30 53.4\t& $12.986\\pm0.150$\t& $12.176\\pm1.00$\t& $11.364\\pm0.024$\t& $11.006\\pm0.025$\t& $10.877\\pm0.025$\t   \\\\ % 420316 \t& 05 38 54.11\t& --02 49 29.8  & $13.164\\pm0.204$\t& $11.729\\pm0.03$\t& $10.829\\pm0.026$\t& $10.310\\pm0.024$\t& $10.126\\pm0.019$\t   \\\\ % 111112\t& 05 38 48.04\t& --02 27 14.2\t& $13.340\\pm0.341$\t& $11.359\\pm0.03$\t& $10.156\\pm0.023$\t&  $9.463\\pm0.026$\t&  $9.187\\pm0.019$        \\\\ %  130452\t& 05 39 45.15\t& --02 04 53.9\t& $13.453\\pm0.215$\t& $12.236\\pm0.04$\t& $10.693\\pm0.030$\t&  $9.960\\pm0.035$\t&  $9.556\\pm0.026$        \\\\ %    230062 \t& 05 38 23.88\t& --02 05 42.0\t& $13.675\\pm0.260$\t& $12.378\\pm0.03$\t& $11.451\\pm0.023$\t& $10.866\\pm0.026$\t& $10.715\\pm0.021$        \\\\ % 121137  & 05 38 35.87\t& --02 30 43.3\t& $13.701\\pm0.290$\t& $12.483\\pm0.03$\t& $11.245\\pm0.026$\t& $10.598\\pm0.023$\t& $10.424\\pm0.024$        \\\\ %  000006$^{a}$  & 05 38 49.17\t& --02 38 22.2\t& $13.7-15.0$\t        & $12.891\\pm0.02$\t& $11.389\\pm0.026$\t& $10.663\\pm0.023$\t& $10.511\\pm0.022$\t   \t  \\\\ % 430847  & 05 39 32.57\t& --02 39 44.0\t& $13.756\\pm0.202$\t& $12.385\\pm0.17$\t& $10.820\\pm0.027$\t& $10.104\\pm0.024$\t&  $9.917\\pm0.019$        \\\\ % 430136 \t& 05 40 22.56\t& --02 33 46.9\t& $13.951\\pm0.207$\t& $12.473\\pm0.04$\t& $11.158\\pm0.028$\t& $10.533\\pm0.023$\t& $10.362\\pm0.023$        \\\\ %  420040  & 05 38 27.26 \t& --02 45 09.7\t& $14.192\\pm0.472$\t& $12.847\\pm0.02$\t& $11.955\\pm0.028$\t& $10.792\\pm0.026$\t&  $9.944\\pm0.028$        \\\\ % 440660  & 05 39 54.66\t& --02 46 34.1\t& $14.253\\pm0.354$\t& $12.757\\pm0.03$\t& $11.054\\pm0.028$\t& $10.251\\pm0.024$\t&  $9.832\\pm0.024$        \\\\ %  121112  & 05 38 40.27\t& --02 30 18.5\t& $14.273\\pm0.389$\t& $12.839\\pm0.02$\t& $11.512\\pm0.026$\t& $10.763\\pm0.023$\t& $10.395\\pm0.025$        \\\\ % 420147  & 05 38 52.01\t& --02 46 43.7\t& $14.301\\pm0.305$\t& $12.719\\pm0.02$\t& $11.518\\pm0.026$\t& $10.774\\pm0.024$\t& $10.421\\pm0.021$        \\\\ % 430967  & 05 39 25.20\t& --02 38 22.0 \t& $14.344\\pm0.341$\t& $13.028\\pm0.16$\t& $11.307\\pm0.031$\t& $10.451\\pm0.023$\t& $10.002\\pm0.023$        \\\\ %  140159  & 05 40 05.11\t& --02 19 59.1\t& $14.493\\pm0.324$\t& $12.950\\pm0.03$\t& $11.459\\pm0.026$\t& $10.767\\pm0.024$\t& $10.542\\pm0.021$        \\\\ % 420742  & 05 38 39.82\t& --02 56 46.2\t& $\\gtrsim 14.5$\t& $12.778\\pm0.02$\t& $11.413\\pm0.027$\t& $10.744\\pm0.023$\t& $10.439\\pm0.021$        \\\\ % \\noalign{\\smallskip} \\hline \\end{tabular} $$ \\begin{list}{}{} \\item[$^{a}$] For these stars we provide two values of the $V$ magnitude. The first number corresponds to the ASAS-3 $V$ magnitude and other values from Wolk (1996; SO000005) and Sherry et al. (2004; SO000006). \\end{list} \\end{table*}  A metallicity detemination in the $\\sigma$~Orionis cluster is becoming essential. The derived distance and age of the cluster (from fits to theoretical isochrones) depends on its metallicity, as recently shown by Sherry et~al. (2008). At the age of the cluster ($\\tau \\sim$ 3\\,Ma), the protoplanetary discs are transiting from optically thick to optically thin. Giant planets may be forming at this moment (Lissauer 1993; Pollack et~al. 1996; Boss 1997; Ida \\& Lin 2004). Therefore, the metallicity {\\bf could play an important} r\\^ole in the planetary formation in $\\sigma$~Orionis (e.g. through the planet-metallicity correlation -- Fischer \\& Valenti 2005).  Cunha et~al. (1998) reported that the metallicity of the Orion OB1 association as a whole is [Fe/H] $=-0.16\\pm0.11$\\footnote{[Fe/H]\\,$=\\log [N$(Fe)$/N$(H)]$-\\log [N$(Fe)$/N$(H)$]_\\odot$.}, which is similar to other determinations using different samples of stars (e.g., Cunha \\& Lambert 1994); see a review of abundance ratios and Galactic chemical evolution in McWilliam (1997). However, there have been very few metallicity studies in the $\\sigma$~Orionis cluster. According to Cunha et~al. (2000) and references therein, the iron abundance of the early-G star \\object{HD~294297} (SO000041, Mayrit~1659068) is slightly subsolar ([Fe/H]$_{\\rm HD~294297}$ = --0.18). Nonetheless, although the star is young ($\\log \\epsilon(\\mathrm{Li})= 2.56$\\footnote{$\\log \\epsilon(\\mathrm{Li})=\\log[N$(Li)$/N$(H)]+12.}) and follows the spectrophotometric sequence of $\\sigma$~Orionis, its cluster membership is in doubt because of its abnormally high proper motion (Caballero 2007). Other star whose metallicity has been investigated with optical spectroscopy is the K3 spectroscopic binary OriNTT~429~AB (Mayrit~1415279~AB). Lee, Mart\\'{\\i}n \\& Mathieu (1994) determined [Fe/H] = --0.43 and --0.24$\\pm$0.10 for the primary and the secondary, respectively. A difference of iron abundances like that is questionable if none of the components of the coeval system has not suffered yet from chemical evolution (e.g. departure of main sequence). The remaining metallicity study was presented by Caballero (2006), who measured an average value of [Fe/H] = 0.0$\\pm$0.1 from an ensemble of mid-resolution spectra centred on H$\\alpha$~$\\lambda$6562.8\\,\\AA. Using X-ray spectral energy distributions, coronal abundances of cluster stars have also been found to be roughly solar (Skinner et~al. 2008) or slightly lower (Sanz-Forcada, Franciosini \\& Pallavicini 2004; Franciosini, Pallavicini \\& Sanz-Forcada 2006; L\\'opez-Santiago \\& Caballero 2008).  \\begin{table*} \\caption[]{Names, IRFM effective temperatures, radial and rotational velocities, and remarks from the literature of the 24 pre-selected stars.} \\label{table.names+vel+remarks} $$ \\begin{tabular}{lllcccl} \\hline \\hline \\noalign{\\smallskip} SO$^a$\t& Mayrit$^b$\t& Simbad            \t        & $T_{\\rm eff}$       & $V_{r}^c$ & $v \\sin i$& Remarks$^d$   \\\\ &\t\t& name            \t\t& [K] \t\t      &      [\\kms]  & [\\kms]\t &\t\t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 230838\t& ...\t\t& \\object{HD 290772}           \t& 6000$\\pm$600$^{e}$  & $+23.3\\pm0.4$ & $30\\pm5$ & Herbig Ae/Be, {\\em IRAS}, $\\rho >$ 30\\,arcmin \\\\ 000041  & 1659068\t& \\object{HD 294297}           \t& 6450$\\pm$105        & $+24.6\\pm0.2$ & $<20$\t & [Fe/H], $V_r$ = +25.0\\,\\kms \\\\ 410305  & ...\t\t& \\object{HD 294308}           \t& 6250$\\pm$90\t      & $-37.2\\pm0.2$ & $<20$\t & $\\rho >$ 30\\,arcmin \\\\ 210734  & 789281\t& \\object{2E 1454}     \t\t& 5235$\\pm$100        & $+24.0\\pm0.3$ & $30\\pm5$ & H$\\alpha$ in broad emission \\\\ 131161  & ...\t\t& ...                   \t& 4765$\\pm$115        & $+23.4\\pm0.3$ & $<20$\t   & $\\rho >$ 30\\,arcmin, in NE nebulosity \\\\ 420039  & 591158\t& \\object{[W96] 4771--0026}     & 5935$\\pm$195        & $+33.7\\pm0.3$ & $60\\pm5$ & [S~{\\sc ii}] and [N~{\\sc ii}] in emission \\\\ 210434  & 615296\t& \\object{2E 1459}     \t\t& 4630$\\pm$110        & $+24.5\\pm0.4$ & $<20$\t & H$\\alpha$ in faint emission \\\\ 000005  & 105249\t& \\object{[W96] rJ053838--0236} & 4935--4000$^{f}$    & $+25.1\\pm0.3$ & $<20$\t & H$\\alpha$ in faint emission, $V_r$ = +42$\\pm$10\\,\\kms \\\\ 211394  & 1374283\t& \\object{Mayrit~1374283}       & 5300$\\pm$210        & $+27.8\\pm0.3$ & $<20$\t & high background around H$\\alpha$ \\\\ 420316  & 822170\t& \\object{RX J0538.9--0249}    \t& 4520$\\pm$175        & $+31.8\\pm0.4$ & $<20$\t & H$\\alpha$ in faint emission, $v \\sin i$ = 28\\,\\kms \\\\ 111112\t& 528005 AB\t& \\object{[W96] 4771--899 AB}  \t& $\\lesssim$3500      & $+30.7\\pm0.4$ & $<20$\t & classical T~Tau, resolved binary, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & $V_r$ = +31$\\pm$10\\,\\kms \\\\ 130452\t& ...\t\t& \\object{RX J0539.8--0205 ABC} & $\\lesssim$3800      & $+20.0\\pm0.5$ & $60\\pm5$ & H$\\alpha$ in broad asymmetric emission, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & spectroscopic triple, $\\rho >$ 30\\,arcmin, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & in NE nebulosity, [S~{\\sc ii}] and [N~{\\sc ii}] in emission, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & $V_r$ = +27.8$\\pm$0.9\\,\\kms \\\\ 230062  & ...\t\t& ...                   \t& 4585$\\pm$225        & $+22.6\\pm0.3$ & $<20$\t & H$\\alpha$ in very faint emission, $\\rho >$ 30\\,arcmin \\\\ 121137  & 344337\t& \\object{2E 1468}    \t\t& 4370$\\pm$310        & $+28.9\\pm0.4$ & $70\\pm5$ & H$\\alpha$ in broad asymmetric emission, \\\\ % & \t\t& \t  \t\t\t&\t\t      & \t      & \t & $V_r$ = +35$\\pm$7\\,\\kms, $v \\sin i$ = 80$\\pm$15\\,\\kms \\\\ 000006  & 157155\t& \\object{[W96] rJ053849--0238} & 4450--3500$^{f}$    & $+34.1\\pm0.7$ & $<20$\t & H$\\alpha$ in faint, broad emission \\\\ 430847  & 750107\t& ...                   \t& $\\lesssim$3800      & $+30.7\\pm0.5$ & $<20$\t & H$\\alpha$ in faint emission \\\\ 430136  & 1471085\t& \\object{Kiso A--0904 105}    \t& $\\lesssim$4000      & $+28.6\\pm0.4$ & $60\\pm5$ & H$\\alpha$ in faint asymmetric emission, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & close to Horsehead, [S~{\\sc ii}] and [N~{\\sc ii}] in emission, \\\\ 420040  & 609206\t& \\object{V505 Ori}            \t& $\\lesssim$3800      & $+29.5\\pm0.7$ & $<20$\t & classical T~Tau, $V_r$ = +48$\\pm$10\\,\\kms \\\\ 440660  & 1223121\t& \\object{V606 Ori}            \t& $\\lesssim$3800      & $+29.3\\pm0.4$ & $<20$\t & H$\\alpha$ in strong, broad, asymmetric emission, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & class II, [N~{\\sc ii}] in emission \\\\ 121112  & 348349\t& \\object{Haro 5--13} \t\t& $\\lesssim$3800      & $+34.3\\pm0.4$ & $<20$\t & classical T~Tau \\\\ 420147  & 653170\t& \\object{RU Ori}              \t& $\\lesssim$3800      & $+32.0\\pm0.6$ & $<20$\t & classical T~Tau \\\\ 430967  & 622103\t& \\object{BG Ori}              \t& $\\lesssim$3500      & $+27.0\\pm0.6$ & $<20$\t & H$\\alpha$ in strong, broad, asymmetric emission, \\\\ & \t\t& \t  \t\t\t&\t\t      & \t      & \t & class II, [S~{\\sc ii}] in emission \\\\ 140159  & 1541051\t& \\object{[NYS99] C--05}    \t& $\\lesssim$3500      & $+27.5\\pm0.4$ & $<20$\t & H$\\alpha$ in faint, broad emission \\\\ 420742  & 1248183\t& \\object{[SWW2004] 125} \t& $\\lesssim$3500      & $+29.3\\pm0.5$ & $<20$\t & H$\\alpha$ in faint emission \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} $$ \\begin{list}{}{} \\item[$^{a}$] Identification number in Caballero (2006). \\item[$^{b}$] Mayrit designation in Caballero (2008c). \\item[$^{c}$] Radial velocities extracted from cross-correlation with a K3V template star (see main text). \\item[$^{d}$] Important remarks from the literature (see main text for references). \\item[$^{e}$] Strong X-ray emitter (see Section~\\ref{effective.temperatures}). \\item[$^{f}$] Derived \\teff from the two different $V$ magnitudes given in Table~\\ref{table.coo+phot+teff}. \\end{list} \\end{table*}  In this paper, we determine {\\em photospheric} abundances of iron, lithium and other elements in nine pre-main sequence, late-type stars of $\\sigma$~Orionis, with the aim of providing an average, accurate value of its metallicity. Our study improves the [Fe/H] estimation presented in Caballero (2006), from where we have taken the original optical spectra. ",
        "conclusions": "\\subsection{Stellar parameters and heliocentric distance} \\label{distance}   We display in Fig.~\\ref{fig4} two diagrams involving basic stellar parameters of the selected stars in comparison with several theoretical evolutionary tracks taken from Siess et~al. (2000), for several ages and metallicities. The derived stellar parameters are consistent with almost all the theoretical tracks due to the large error on surface gravity. However, their central values are very well reproduced by the solid line that represents the 3\\,Ma-old, overshooting track. Ages lower than 2\\,Ma or larger than 4\\,Ma are marginally consistent with the stellar parameters of the selected stars. The theoretical tracks are only sensitive to metallicity for effective temperatures higher than \\teff $\\sim$ 4800\\,K.  In the right panel of Fig.~\\ref{fig4}, we plot the luminosity per unit of mass, normalized to this quantity in the Sun:  \\begin{equation} \\mathcal{L} = (L/M)/(L_\\odot / M_\\odot). \\end{equation}  \\noindent Using the well known expressions for the bolometric luminosity and the surface gravity, it is derived~that:  \\begin{equation} L/M = 4 \\pi G \\sigma \\frac{T_{\\rm eff}^4}{g}, \\end{equation}  \\noindent and, therefore,  \\begin{equation} \\mathcal{L} = \\frac{g_\\odot}{g} \\left( \\frac{T_{\\rm eff}}{T_{{\\rm eff,}\\odot}} \\right)^4. \\end{equation}  \\noindent In another way, $\\mathcal{L} \\propto g^{-1} T_{\\rm eff}^4$ at a given metallicity, which means that the normalized luminosity per unit of mass only depends on the derived stellar parameters.  In addition, in Fig.~\\ref{fig6} we display the $\\mathcal{L}$ vs. $J$ diagram and several theoretical evolutionary tracks for different heliocentric distances to the $\\sigma$~Orionis cluster. In this case, all selected stars, except three, are close to the tracks for distances in the range 350\\,pc $< d <$ 500\\,pc. These three stars are the hottest stars in the sample (SO000041 --\\teff = 6450$\\pm$105\\,K--, SO420039 --\\teff = 5935$\\pm$195\\,K-- and SO211394 --\\teff = 5300$\\pm$210\\,K--). SO000041 is marginally consistent with the track at $d=$ 400\\,pc but still too luminous for its expected position in this diagram (see Section~\\ref{lithium}). For the other two stars, one should expect $J$ magnitudes in the range 9\\,mag $\\lesssim J \\lesssim$ 10\\,mag, but surprinsingly they have $J \\sim$ 11.3\\,mag. Remarkably, one of these two stars has a large rotational velocity ($v \\sin{i}$ = 60$\\pm$5\\,\\kmso; SO420039), while the another star falls in the halo of the $\\sigma$~Orionis cluster ($\\rho \\sim$ 23\\,arcmin; SO211394), where the contamination by neighbouring young stellar populations gets higher. In any case, an heliocentric distance of $d \\sim$ 400\\,pc is in agreement with recent results on the distance to the $\\sigma$~Orionis cluster in Sherry et~al. (2008) and Caballero (2008a). These authors claimed that distances of $d=420 \\pm 30$ and $\\sim$ 385\\,pc, respectively, are more plausible than previous estimates at~350\\,pc. However, our error bars are so large that we cannot favour a distance of 450\\,pc rather than~350\\,pc.   \\subsection{Lithium abundance} \\label{lithium}  \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth]{f3.eps} \\caption{Same as Fig.~\\ref{fig4} but for the normalized luminosity per unit of mass, $\\mathcal{L}$, as a function of the apparent magnitude, $J$. Dashed-dotted lines: tracks for fixed age (3\\,Ma), metallicity ($Z$ = 0.02), and overshooting, and variable heliocentric distances $d$ = 500, 450, 350, 300$\\,$pc, from top to bottom; and thick solid line: track for 400$\\,$pc, 3\\,Ma, $Z$ = 0.02, and overshooting. } \\label{fig6} \\end{figure}  Five stars in the selected sample have lithium abundances consistent within the error bars or slightly lower than the solar meteoritic value, $\\log \\epsilon(\\mathrm{Li}) = 3.3$, which is also considered as the current Galactic or ``cosmic'' $^7$Li abundance (Boesgaard \\& Tripicco 1986). Therefore, young, population~I stars are believed to have inherited this lithium abundance at birth. However, two stars (SO131161 and SO210434) show higher lithium abundances, $\\log \\epsilon({\\rm Li}) \\sim 3.6-3.7$ which might be due to wrong stellar parameters. Note that a 100\\,K cooler effective temperature would provide a 0.10--0.15\\,dex lower lithium abundance. We should note that the uncertainties on the lithium abundances are quite high due to the large error bars of the temperature determination. Thus, some stars have lithium abundances that are at least marginally consistent, within the error bars, with the ``cosmic'' $^7$Li abundance.  On the other hand, such high lithium abundances have been found in other T~Tauri stars. For instance, Mart{\\'\\i}n et~al. (1994) showed two stars with similar stellar parameters (\\teff $\\sim 4700$\\,K and \\logg $\\sim 3.6-4.0$) and LTE lithium abundances $\\log \\epsilon({\\rm Li}) \\sim 3.65 - 3.60$. However, their NLTE corrections were between --0.3 and --0.2\\,dex, thus significantly larger than our NLTE corrections, $\\Delta_\\mathrm{NLTE-LTE} \\lesssim -0.10$\\,dex. Those corrections moved their stars to the region of normal lithium abundances. In addition, a high NLTE lithium abundance has also been found in a double-lined spectroscopic binary of the $\\sigma$~Orionis cluster by Lee et~al. (1994): Mayrit~1415279~AB (it was presented in Section~\\ref{introduction}).  Likewise, there are two additional stars with a lithium abundance slightly lower than the others, SO211394 ($\\log \\epsilon({\\rm Li})$ = 2.77$\\pm$0.31) and SO000041 ($\\log \\epsilon({\\rm Li})$ = 2.75$\\pm$0.18). SO211394 is one of the three outliers in the $\\mathcal{L}$ vs. $J$ diagram in Fig.~\\ref{fig6}. The location of this star in the halo of the $\\sigma$~Orionis cluster, its relatively low lithium abundance, and the abnormally faint $J$-band magnitude for its luminosity per unit of mass may suggest classifying the star as a member of a neighbouring, not-so-young population in the \\object{Ori~OB1b} association different from $\\sigma$~Orionis (see, e.g., Jeffries et~al. 2006). Given its position to the west of the cluster, SO211394 may belong to the extended population of the $\\sim$5\\,Ma-old surrounding \\object{Alnilam} (\\object{Collinder~70} or the ``$\\epsilon$~Orionis cluster''; Caballero \\& Solano 2008). Nevertheless, its position in a EW(Li~{\\sc i}) vs. \\teff diagram (e.g. L\\'opez-Santiago et~al. 2006) might also indicate that the star is a member of a distant ($d = 710 \\pm 50$\\,pc) moving group with a Pleiades-like age.  SO000041 has a position only marginally consistent with the theoretical tracks in the $\\mathcal{L}$ vs. $J$ diagram. This fact, together with the rather high proper motion of SO000041 (Section~\\ref{introduction}) and its location even further from the cluster centre than SO211394, bring doubts about its membership to the $\\sigma$~Orionis cluster. SO000041 is the easternmost, bright, cluster member candidate in the Mayrit catalogue, just in the outskirts of the \\object{Horsehead Nebula} (Caballero 2007, 2008c). Membership in a slightly younger population of recently born stars associated to the nebula might explain its relatively low gravity. However, membership in a $\\sim$100\\,Ma-old moving group (just like SO211394) at $d \\sim 630 \\pm 50$\\,pc could explain both its relatively low lithium abundance and large proper motion. Further astrometric and spectrophotometric studies are needed to ascertain the real nature of SO000041 and SO211394.  \\subsection{Metallicity of $\\sigma$~Orionis}  The nine selected stars share the same radial velocity and have high Li abundances consistent with the ``cosmic'' $^7$Li abundance, so we think that they are probably members of the $\\sigma$~Orionis cluster. According to the values shown in Table~\\ref{table.stellarparameters}, the average metallicity of these nine selected stars (i.e. the metallicity of the $\\sigma$~Orionis cluster) is [Fe/H] = --0.02 $\\pm$ 0.09 (random) $\\pm$ 0.13 (systematic). This value, as well as for other element abundances, is consistent with solar metallicity. If we do not take into account SO131161 and SO230062, that are not Mayrit objects (the two no-Mayrit stars are at angular separations $\\rho >$ 30\\,arcmin to the $\\sigma$~Orionis centre, and might not belong to the cluster -- Jeffries et~al. 2006; Caballero 2007, 2008b), the new average metallicity barely varies to [Fe/H] = $-0.05 \\pm 0.08 \\pm 0.13$. Both metallicities would provide a distance to the cluster of $d \\sim$ 440\\,pc from the results presented by Sherry et~al.~(2008). If we also remove from the sample the star SO420039 (that seems to be at a longer heliocentric distance; Section~\\ref{distance}) and the stars SO211394 and SO000041 (that have relatively low lithium abundances; Section~\\ref{lithium}), then the average metallicity is [Fe/H] = $0.00 \\pm 0.07 \\pm 0.13$."
    },
    "0809/0809.2789_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since the original construction of the Randall-Sundrum model \\cite{Randall:1999ee}, many authors aimed to find black hole solutions localized on a brane (see e.g. \\cite{Chamblin:1999by}-\\cite{Gregory:2008br}, and \\cite{Gregory:2008rf} for a review). Although this is still an open problem in the case of a 3-brane, Emparan, Horowitz and Myers found a class of brane-localized black holes in the lower dimensional case of a 2-brane embedded in four dimensional Anti-de Sitter space  $AdS_4$ \\cite{Emparan:1999wa,Emparan:1999fd}. This was found by appropriately cutting the AdS-C metric \\cite{Kinnersley:1970zw,Plebanski:1976gy}, which describes bulk accelerated black holes in $AdS_4$, with a brane. Since the tension of this brane is determined by its acceleration in the $AdS_4$ background, this tension is carefully chosen such that the brane acceleration matches the acceleration of the bulk black hole. Recently, this condition was relaxed by detuning the brane tension from the bulk black hole acceleration to find two new classes of solutions on the 2-brane \\cite{Anber:2008qu}. The first class has time-dependent induced metrics and describes an evolving  lump of dark radiation, while the second is a constant $\\phi$-section of the original four-dimensional C-metric and  describes accelerated black holes by means of strings or struts.  Moreover, it was conjectured in \\cite{Emparan:2002px} that black hole solutions localized on a brane in the $AdS_{D+1}$ braneworld correspond to quantum-corrected black holes in $D$ dimensions. This conjuncture followed naturally from the AdS/CFT correspondence in which classical dynamics in the $AdS_{D+1}$ bulk encodes the quantum dynamics of the dual D-dimensional conformal field theory (CFT). Thus, solving the classical $D+1$ classical equations in the bulk is equivalent to solving Einstein equations $G_{\\mu\\nu}=8\\,\\pi\\,G_D \\left\\langle T_{\\mu\\nu}\\right\\rangle$ on the brane, where $G_D$ is the D-dimensional Newton's constant and $\\left\\langle T_{\\mu\\nu}\\right\\rangle$ is the energy-momentum tensor of a strongly coupled  CFT with a cutoff in the ultraviolet due to the presence of the brane. This conjecture was put to test in \\cite{Emparan:2002px} by comparing the brane-black hole solution found in \\cite{Emparan:1999wa} with the one obtained in the dual $2+1$ CFT coupled to gravity. This can be done by starting with a point particle of mass $M$ in $2+1$ dimension which generates a deficit angle $\\delta=8\\,\\pi\\,G_3 M$, and computing the Casimir energy-momentum tensor \\cite{Souradeep:1992ia}. Then, one can use this energy-momentum tensor to calculate the backreaction on the conical spacetime \\cite{Soleng:1993yh}. The agreement between the two sides is remarkable and gives a strong argument in favor of the AdS/CFT duality in the context of braneworld scenario. Indeed, this conjecture gives us a convenient way to learn about strong quantum effects in curved backgrounds \\cite{Anderson:2004md,Grisa:2007qv,Pujolas:2008rc}.  According to this conjecture, the accelerated black hole solution found in \\cite{Anber:2008qu} was interpreted as quantum corrected black hole(s) accelerating on the brane. In the case of a critical brane, this solution describes a pair of  black holes accelerated by two strings (or a strut) pulling (pushing) them toward infinity. In the CFT picture, the energy per unit length of this string (strut) has both classical and quantum contributions. Indeed, while in pure $2+1$ dimensional gravity point particles do not interact, quantum effects generate a force between the particles. Hence, the tension of the string (strut) takes into consideration both effects.  In the present work, we calculate the quantum-mechanical backreaction on two point particles attached to a strut and accelerating in  $2+1$ D flat background. In our setup, we compute the Green's function of a  conformally coupled scalar. Since we work in an accelerating frame moving with acceleration $A$, it is natural to use a quantum field in equilibrium with a thermal bath at temperature $T=A/2\\,\\pi$, which is the Hawking-Unruh temperature associated with the Rindler (acceleration) horizon. In calculating the thermal Green's function, we impose two different boundary conditions at the location of the strut, namely, transparent and Dirichlet boundary conditions. We find that the numerical coefficient of the induced energy-momentum tensor in the first case is identical to the case of a point particle at rest in a $2+1$ D background. On the other hand, the energy-momentum tensor in the case of Dirichlet boundary conditions is divergent at the position of the strut. This suggests that unless we impose transparent boundary conditions, the location of the strut is susceptible  to the formation of curvature singularity. Further, we use the resulting energy-momentum tensor calculated in the case of transparent boundary conditions to find the gravitational backreaction on the spacetime. Interestingly enough, we find that the resulting quantum corrected metric has the same form as the one obtained before in \\cite{Anber:2008qu} by cutting the $AdS_4$ C-metric with a critical angular dependent brane. Although the former solution describes accelerating black hole dressed with {\\it weakly} coupled CFT (WCFTBH), the latter, according to AdS/CFT+gravity conjecture, describes accelerating black hole dressed with {\\it strongly} coupled CFT (SCFTBH). Moreover, the presence of the CFT at finite temperature gives us a unique opportunity to study finite temperature effects in strongly coupled system in curved background. Contrary to the case of the static black hole constructed previusly in \\cite{Emparan:1999wa}, where it was found that the black hole mass can be as large as $1/4G_3$, the largest mass one can place in $2+1$ D, we find that the maximum mass in the present case is temperature dependent. Studying the behavior of this maximum mass reveals a striking difference between the weakly and strongly coupled theories. Although the mass in the former decreases monotonically with temperature from $1/4G_3$ to zero, we find that in the second case it decreases from  $1/4G_3$ at low temperatures to a minimum value, and then  it increases to $1/4G_3$ again at high temperatures. Further, the black hole horizon circumference diverges in both cases as the mass reaches its maximum value. Beyond this mass, the horizon disappears leaving behind a naked singularity. Actually, it was argued before that quantum effects dresses conical singularities with regular horizon given that these singularities are sufficiently massive. This is known as quantum cosmic censorship \\cite{Emparan:2002px}.  Although this still applies in  the case of accelerating  conical singularities, supermassive singularities (exceeding the maximum allowed black hole mass at a given temperature) violate this censorship.  We start our treatment using a classical background with a vanishing value of the total mass of particles and strut $m_{p}+m_{s}=0$. After computing the quantum-mechanical backreaction on the spacetime, we find that the mass of the strut gets renormalized which in turn violates the above equality. The violation is minimal for small values of the black hole mass, and becomes stronger for larger values.  Our paper is organized as follows. In section 2 we discuss the background geometry used in the setup. Then, in sections 3 and 4 we calculate the thermal Green's function and the induced energy-momentum tensor for both transparent and Dirichlet boundary conditions. The gravitational backreaction due to the weakly coupled CFT is computed in section 5, while in section 6 we show that the same form of the metric found in section 5 can be obtained by cutting the AdS C-metric in 4-D with an angular dependent critical brane \\cite{Anber:2008qu}. The latter describes accelerated black hole dressed with strongly coupled CFT. In section 7 we compute various thermodynamic quantities  of the black hole and comment on the AdS/CFT+gravity interpretation of the strongly coupled solution, and finally we conclude in section 8. ",
        "conclusions": "In this work, we have studied the CFT corrections to two accelerating masses moving with acceleration $A$, by means of  a strut connecting them, in a flat background. To achieve this, we first computed the thermal Green's function for a conformally coupled scalar. This function describes the quantum fluctuations of a quantum field kept in equilibrium with a thermal bath at temperature $T=A/2\\pi$, which is the Hawking-Unruh temperature associated with the acceleration horizon. In calculating the Green's function, we used two different boundary conditions, namely, transparent and Dirichlet boundary conditions. We have found that the induced energy-momentum tensor diverges at the location of the strut in the case of Dirichlet boundary conditions, while the same quantity is regular for transparent boundary conditions. Further, we proceeded using the regular form of the energy-momentum tensor to calculate  the gravitational backreaction on the spacetime. Interestingly enough, we have shown that the  backreaction of the CFT  reproduces the same metric found before in \\cite{Anber:2008qu} which was constructed  by cutting the $AdS_4$ spacetime with angular dependent critical brane. Although the former solution describes accelerated black hole dressed with weakly coupled CFT (WCFTBH), the latter, according to AdS/CFT+gravity, should describe a black hole dressed with strongly coupled CFT (SCFTBH). This is the first use of the duality in a system containing two interacting particles. Moreover, the presence of the CFT at finite temperature gave us an opportunity to study the finite temperature effects in a strongly coupled system.  The existence of the brane implies that high energy states in the dual CFT theory are integrated out, and hence the CFT has a UV cutoff scale $\\mu_{\\mbox{\\scriptsize UV}}\\sim 1/\\ell_4$, or in terms of the  number of CFT degrees of freedom and the $3$ D Newton's constant $\\mu_{\\mbox{\\scriptsize UV}}\\sim 1/g_{*} G_3 $. The solution given by (\\ref{black hole on brane strong CFT}) describes a genuine  black hole if the Compton's wavelength $\\lambda_{C}=1/m_{\\mbox{\\scriptsize BH}}$ satisfies $\\lambda_C < r_0$, where $r_{0}$ is the black hole horizon, otherwise quantum effects would smear it over a volume larger than the horizon radius. As was shown in \\cite{Emparan:2002px}, this requirement sets a new scale $1/\\sqrt{G_4}\\sim 1/\\sqrt{g_*}G_3$ on the brane. In other words, a black hole solution is reliable if its mass ranges in $\\mu_{\\mbox{\\scriptsize UV}}<1/\\sqrt{G_4}<m_{\\mbox{\\scriptsize BH}}<m_{\\mbox{\\scriptsize BH max}}$. The maximum black hole mass $m_{\\mbox{\\scriptsize BH max}}$ is equal to the maximum mass $1/4G_3$ one can place in $2+1$ D in the case of a static black hole. However, we found that $m_{\\mbox{\\scriptsize BH max}}$ is temperature dependent, and can be determined when the strongly or weakly coupled CFT parameters, $G_4MA$ and $g_{*}\\mu$ respectively, reaches the critical value $1/3\\sqrt{3}$. As was shown in figure (\\ref{MBHMAX V LAMBDA}), this mass starts at $1/4G_3$ at low temperatures for both SCFTBH and WCFTBH, and then decreases monotonically until a black hole ceases to exist as $T \\rightarrow\\tilde A/2\\pi$ in the case of WCFTBH. In contrast to this behavior, $m_{\\mbox{\\scriptsize BH max}}$ for SCFTBH develops a minimum  and then increases until it reaches $1/4G_3$ again at $T=\\tilde A/ 2\\pi$. This is a striking deference between strongly and weakly coupled CFT. Nevertheless, in both cases as the black hole mass reaches the maximum allowed value, the proper horizon circumference grows indefinitely. Beyond, $m_{\\mbox{\\scriptsize BH max}}$ the black hole horizon disappears leaving behind undressed singularity in a clear violation of the quantum cosmic censorship. In fact, it was shown in \\cite{Anber:2008qu}  that the time-dependent solutions obtained there do not respect the quantum censorship, as well. Indeed, these results may indicate that the censorship operates only for a static CFT.   The presence of a naked singularity may signal phase transition on the bulk side. However, at this point, it is safe to say that a complete understanding of the nature of this singularity is still an open question. Further study of SCFTBH may reveal rich phenomena of finite temperature CFT coupled to gravity in $2+1$ D."
    },
    "0809/0809.0435_arXiv.txt": {
        "abstract": "{} {Very young stellar clusters and cluster complexes may be embedded in dust lanes along spiral arms in disk galaxies and escape detection in visual bands.  Observations in the near-infrared K-band offer an almost unbiased view of such clusters or complexes due to the small attenuation by dust at this wavelength.  The objective is to determine their population size, absolute K-band magnitude distribution above the limiting magnitude imposed by the data, and location relative to the spiral pattern in disk galaxies.} {All slightly extended sources were identified on deep K-band maps of 46 spiral galaxies reaching at least K=20.3 \\mqas\\ at a signal-to-noise level of 3.  The galaxies had inclination angles $<65\\degr$ and linear resolutions $<100$~pc with seeing better than $1\\arcsec$. The sample includes both barred and normal spirals with a wide spread in types. We also analyzed J- and  H-band colors for 4 galaxies for which such images were available.  An apparent magnitude limit of K = 19~mag was used for the sources analyzed in order to avoid marginal detections.  Furthermore, we derived the source distributions of magnitudes and relative locations with respect to the spiral patterns. } {Almost 70\\% (15/22) of the grand-design spiral galaxies show significant concentration of bright K-band knots in their arm regions corresponding to 30\\% (15/46) of the full sample.  Color-color diagrams for the 4 spirals with JHK photometry suggest that a significant fraction of the diffuse sources found in the arms are complexes of young stellar clusters with ages $<$10~Myr and reddened with several magnitudes of visual extinction. The brightest knots reach an absolute K-band magnitude M$_\\mathrm{K}$ of -15.5~mag corresponding to stellar clusters or complexes with total masses up to at least 10$^5$~M$_{\\odot}$.  Brightest magnitude and number of knots correlate with the total absolute magnitude of the host galaxy.  More knots are seen in galaxies with high far-infrared flux and strong two-armed spiral perturbations.  The bright knots constitute up to a few percent of the total K-band flux from their parent galaxy and account for a star formation rate of $\\sim$1~M$_{\\odot}$ yr$^{-1}$ for the brightest grand-design spiral galaxies.  } {} ",
        "introduction": "The present star formation rate (SFR) in the local Universe is important as a part of the history of galaxy build-up and as a means of comparison for the study of star formation at higher redshifts.  For spiral galaxies, significant variation is observed as a function of Hubble type and luminosity class. Extinction by dust in spiral arms makes it difficult to get a full census from observations in visual bands, whereas near-infrared (NIR) colors offer a much higher transparency.  In a study of the NIR K-band images of spiral galaxies, \\citet{gp98} noticed that several grand-design spirals have bright knots along their arms, suggesting that such knots are related to the spiral structure. Young stellar clusters may expel gas left over after their main star formation phase and may then experience violent relaxation \\citep{bastian06}.  This could lead to their destruction on time scales of a few tens of Myr \\citep{goodwin06}.  The knots are marginally resolved, suggesting sizes of less than 100 pc.  Unbound stellar clusters will exceed this size due to diffusion of stellar orbits within 50~Myr \\citep{wielen77}. These points indicate that the knots are dynamically young.  \\citet{patsis01} studied such knots in two galaxies using narrow-band filters in the K-band and concluded that the major contribution to their K-band fluxes comes from continuum radiation.  In many cases these knots are embedded in dust lanes and invisible in visual bands. Taking advantage of the mutual alignment of eight of these knots along the southern arm of NGC 2997, \\citet{grosbol06} obtained spectroscopy in the J- and K-bands detecting hydrogen Br$_\\gamma$ in emission, undoubtedly demonstrating that the knots are HII regions deeply enshrouded in dust, confirming the results previously suggested by \\citet{patsis01}. They also explored the relative position of the knots with respect to the ridge of the K-band arm to derive kinematic parameters of the density wave to constrain the star formation regime and ages of the knots.  Their absolute luminosity and Br$_\\gamma$ emission suggest they are very young, massive stellar complexes that would host a significant fraction of the heaviest stars formed in a galaxy.  Thus, the statistics of them in spiral galaxies may imply SFR of stars on the upper part of the initial mass function (IMF) and rate of supernovae type II in such systems.  The low attenuation by dust in the K-band makes it possible to obtain almost complete samples while many clusters and complexes may be missed on visual images due to heavy obscuration in dust lanes.  To estimate internal mean physical properties of these NIR starforming knots, we used the theoretical models starburst99 by \\citet[ hereafter SB99]{leitherer99}. The rate of star formation, derived from NIR data, may complement or correct the one derived from the observations of optical HII regions, as done by SINGG \\citep{meurer06, hanish06}. Indeed, due to the dust absorption, few of the NIR knots would contribute significantly to the H$_{\\alpha}$ emission line, used by SINGG to estimate the star formation rate in their sample of galaxies.  Thus, the knots we are studying in this paper may constitute the first phases of the star formation that escape detection in optical bands, or might well be the dusty, densest part of large HII region complexes. We present statistics of such knots detected on K-band images of 46 spiral galaxies.  For the four galaxies observed in the J-, \\mbox{H-,} and K-bands, we also compared their colors with simple SB99 models using different star formation scenarios. ",
        "conclusions": "The K-band images of spiral galaxies provide an almost clear view of their young stellar clusters and starforming complexes that in visual bands are partly obscured by dust absorption (e.g. by dust lanes in spiral arms).  This allows an unbiased study, within the resolution and magnitude limits imposed by the present sample, of the distribution of such complexes in 46 nearby spiral galaxies.  The (H-K)--(J-H) diagrams for the 4 galaxies, for which JHK photometry is available, suggest that the majority of diffuse sources (cs$<$0.2) are starforming complexes that have been reddened with more than 4~mag of visual extinction.  This is supported by their concentration in arm regions in many grand-design galaxies.  Since the clusters complexes are unresolved, the internal attenuation by dust could vary very significantly and may obscure part of it even in the K-band.  Some background galaxies may be included, whereas globular clusters are expected to be significantly fainter. Due to the high extinction and lack of color information for the majority of the sample, detailed conclusions cannot be made on the star formation process and its parameters.  Around 70\\% (15/22) of grand-design spirals in the sample show significant concentration of bright, diffuse knots in their arms corresponding to 30\\% (15/46) of the complete sample.  The 7 grand-design systems that do not display a significant concentration of knots all have weak two-armed spiral perturbations.  This suggests that only disk galaxies with strong spiral perturbation (i.e. 0.1$<$A$_2$) will be able to develop a starforming front associated to its spiral pattern, while galaxies with weaker spirals will have a more random star formation in their disks.  The brightest clusters/complexes as measured by \\MK5, the 5$^\\mathrm{th}$ brightest diffuse knot, reach absolute magnitudes of M$_\\mathrm{K} \\approx -15.5$~mag but depend on absolute magnitude, type, luminosity type, and strength of the spiral perturbation of the host galaxy. Early type spirals have, on average, fainter complexes than later types.  The brightest complexes are found in more luminous late-type spirals in luminosity class I-II with strong two-armed spiral patterns.  The number of clusters and complexes in a galaxy indicated by \\N13, diffuse knots with M$_\\mathrm{K} < -13$~mag, shows the same trends as \\MK5. In addition, the number correlates well with the far-infrared flux, suggesting a dependency on the amount of dust in the galaxy.  Star formation rates were estimated from the total integrated K-band luminosity of the diffuse knots in each galaxy.  The values range up to 1~M$_{\\odot}$ yr$^{-1}$  for grand-design spirals but may only indicate the formation of the most massive complexes."
    },
    "0809/0809.2430_arXiv.txt": {
        "abstract": "We investigate the time average mean square displacement $\\overline{\\delta^2}(x(t))=\\int_0^{t-\\Delta}[x(t^\\prime+\\Delta)-x(t^\\prime)]^2 dt^\\prime/(t-\\Delta)$ for fractional Brownian and Langevin motion. Unlike the previously investigated  continuous time random walk model $\\overline{\\delta^2}$ converges to the ensemble average $\\langle x^2 \\rangle \\sim t^{2 H}$ in the long measurement time limit. The convergence to ergodic behavior is however slow, and surprisingly the Hurst exponent $H=3/4$ marks the critical point of the speed of convergence. When $H<3/4$, the ergodicity breaking parameter $\\mbox{EB} = \\mbox{Var} ( \\overline{\\delta^2} ) / \\langle \\overline{\\delta^2} \\rangle^2\\sim  k(H) \\cdot\\Delta\\cdot t^{-1}$, when $H=3/4$, $\\mbox{EB} \\sim (9/16)(\\ln t) \\cdot\\Delta \\cdot t^{-1}$, and when $3/4<H <1, \\mbox{EB} \\sim k(H)\\Delta^{4-4H} t^{4H-4}$. In the ballistic limit $H \\to 1$ ergodicity is broken and $\\mbox{EB} \\sim 2$. The critical point $H=3/4$ is marked by the divergence of the coefficient $k(H)$. Fractional Brownian motion as a model for recent experiments of sub-diffusion of mRNA in the cell is briefly discussed and comparison with the continuous time random walk model is made. ",
        "introduction": "Fractional calculus, e.g. $d^{1/2}/ dt^{1/2}$, is a powerful mathematical tool for the investigation of physical and biological phenomena with long-range correlations or long memory \\cite{Metzler:00}. For example fractional calculus describes the mechanical memory of viscoelastic materials \\cite{Mainardi:88}. An important application of fractional calculus is in the stochastic modeling of anomalous diffusion. Fractional Fokker-Planck equations describe the long time behavior of the continuous time random walk (CTRW) model, when waiting times  and/or jump lengths have power-law distributions \\cite{Barkai:00,Barkai:01,Metzler:00,Deng:07}. A different stochastic approach to anomalous diffusion is based on fractional Brownian motion (fBM) \\cite{Mandelbrot:68} which is related to recently investigated fractional Langevin equations (see details below) \\cite{Goychuk:07,Min:05,Burov:08}.  Recent single particle tracking  of mRNA molecules \\cite{Golding:06} and lipid granules \\cite{Tolic:04} in living cells revealed that time averaged mean square displacement $\\overline{\\delta^2}$ (defined below more precisely) of individual particles remains a random variable while indicating that the particle motion is sub-diffusive.  This means that the time averages are not identical to ensemble averages. Such breaking of ergodicity was investigated within the sub-diffusive CTRW model \\cite{He:08,Klafter:08}. It was shown that transport and diffusion constants extracted from single particle trajectories remain random variables, even in the long measurement time limit. For a non-technical point of view on this problem see \\cite{Sokolov:08}. Here we consider three stochastic models of anomalous diffusion:  fBM  and the under-damped and the over-damped fractional Langevin equation. Except for limiting cases (i.e. ballistic diffusion) we find that the time average $\\overline{\\delta^2}$, in the long measurement time limit, is identical to the ensemble average $\\langle x^2 \\rangle$, indicating that these models are ergodic. Note however, that experiments  on anomalous dynamics of particles in the cell are always conducted for finite times (due to the life time of the cell).  Here we find the finite time corrections to ergodic behavior, namely we give estimates on how far will a finite time measurement of anomalous diffusion deviate from the ensemble average. Since the convergence to ergodic behavior is slow our results seem particularly important to finite time experiments. The problem of estimating diffusion constants from single particle tracking, for normal diffusion,  is already well investigated \\cite{Saxton:97}.  In recent years there was much interest in non-ergodicity of anomalous diffusion processes. A well investigated system are blinking quantum dots \\cite{Dahan:03,Margolin:05}, which exhibit a L\\'evy walk type of dynamics (a super-diffusive process). Very general formula for the distribution of time averages for weakly non ergodic systems was derived in \\cite{Rebenshtok:07}, and this framework was shown to describe the sub-diffusive CTRW \\cite{Bel:05}.  Bao et al have investigated ergodicity breaking for stochastic dynamics described by the generalized Langevin equation \\cite{Bao:05}. They considered the time averaged velocity variance. The latter converges to $k_B T / m$ in thermal equilibrium if the process is ergodic. It was shown \\cite{Bao:05} that under certain conditions the generalized Langevin equation is non ergodic (see also \\cite{Lee:01,Oliv:03,Costa:06,Dhar:07,Plyukhin:08}). Our work, following the recent experiments \\cite{Tolic:04,Golding:06}, considers the time average of the mean square displacements which yields information on diffusion constants. In contrast with \\cite{Bao:05}, we consider an out of equilibrium situation, because our observable: the coordinate of the particle does not reach an equilibrium, since the system is infinite. ",
        "conclusions": "We showed that the fractional processes $x(t)$, $y(t)$ and $z(t)$ are ergodic. The ergodicity breaking parameter decays as a power-law to zero. In the ballistic limit $\\langle x^2 \\rangle\\sim t^2 $ non ergodicity is found. For the opposite localization limit $\\langle x^2 \\rangle \\sim t^0$ (i.e. $H \\to 0$ for fBM) the asymptotic convergence is reached only after very long times.  Our most surprising result is that the transition between the localization limit and ballistic limit is not smooth. When $H=3/4$ the EB is changed and the amplitude $k(H)$ diverges. Other very different critical exponents of fractional Langevin equations were recently found in \\cite{Burov:08}.  There the critical exponents mark transitions between over-damped and under-damped motion. So stochastic  fractional processes possess a zoo of critical exponents. Let us now compare between our results and those derived based on the CTRW model \\cite{He:08}. The most striking difference is that for the CTRW model we have non ergodic behavior even in the long time limit.  It is attempting to conclude that this indicates that the underlying stochastic motion for the mentioned experiments in the cell is of CTRW nature. However, as mentioned in the introduction experiments are conducted for finite times, and hence what may seem as a deviation from ergodic behavior may actually be a finite time effect. Here we gave analytical predictions for the deviations from ergodicity for finite time measurement, based on three fractional models. The EB parameter depends on measurement time and lag time, and can be used to compare experimental data with predictions of fractional equations (the EB parameter for the CTRW is given in \\cite{He:08}). It should be noted however that for sub-diffusion  in the cell, effects of the boundary of the cell, may be important, and these effects where not considered in this text. Another important difference is that {\\em for an infinite system} we have for the CTRW $\\langle \\overline{\\delta^2} \\rangle \\sim \\Delta/ t^\\alpha$, so the time average procedure yields a linear dependence on $\\Delta$ and an aging effect with respect to the measurement time. Hence for CTRW an anomalous diffusion process may seem normal with respect to $\\Delta$ \\cite{He:08,Barkai:03a,Barkai:03b,Klafter:08}. In contrast for the fractional models we investigated here, we have $\\langle \\overline{\\delta^2} \\rangle \\sim \\Delta^\\alpha$ which is the same as the ensemble average $\\langle x^2 \\rangle \\sim t^\\alpha$. The main difference between the two approaches, is that the CTRW process is non-stationary. It would be interesting to investigate fractional Riemann-Liouville  Brownian motion (Eq. (\\ref{eq:one}) without the integral from $-\\infty$ to $0$) which is a non-stationary process.  {\\bf Acknowledgement} This work was supported by the Israel Science Foundation. EB thanks S. Burov for discussions.       \\newpage %"
    },
    "0809/0809.5026_arXiv.txt": {
        "abstract": "{Mid-infrared spectroscopy of dense illuminated ridges (or photodissociation regions, PDRs) suggests dust evolution. Such evolution must be reflected in the gas physical properties through processes like photo-electric heating or H$_2$ formation. } {With Spitzer Infrared Spectrograph (IRS) and ISOCAM data, we study the mid-IR emission of closeby, well known PDRs. Focusing on the band and continuum dust emissions, we follow their relative contributions and analyze their variations in terms of abundance of dust populations.} {In order to disentangle dust evolution and excitation effects, we use a dust emission model that we couple to radiative transfer. Our dust model reproduces extinction and emission of the standard interstellar medium that we represent with diffuse high galactic latitude clouds called Cirrus. We take the properties of dust in Cirrus as a reference to which we compare the dust emission from more excited regions, namely the Horsehead and the reflection nebula NGC 2023 North.} {We show that in both regions, radiative transfer effects cannot account for the observed spectral variations. We interpret these variations in term of changes of the relative abundance between polycyclic aromatic hydrocarbons (PAHs, mid-IR band carriers) and very small grains (VSGs, mid-IR continuum carriers). } {We conclude that the PAH/VSG abundance ratio is 2.4 times smaller at the peak emission of the Horsehead nebula than in the Cirrus case. For NGC2023 North where spectral evolution is observed across the northern PDR, we conclude that this ratio is $\\sim$\\,5 times lower in the dense, cold zones of the PDR than in its diffuse illuminated part where dust properties seem to be the same as in Cirrus. We conclude that dust in PDRs seems to evolve from \"dense\" to \"diffuse\" properties at the small spatial scale of the dense illuminated ridge. } ",
        "introduction": "Evolution of interstellar dust was first inferred from measurements and modelling of the ultraviolet(UV)-visible extinction \\citep[]{cardelli89, fitzpatrick86, fitzpatrick88, fitzpatrick90, kim94a} and from the infrared (IR) colour variations in IRAS and COBE data \\citep[][]{boulanger90, abergel94, abergel96, laureijs91}. The evolution of interstellar matter is powered by stars which have a profound influence on the chemical and energetic balance of the interstellar medium. Along this lifecycle, dust grains undergo efficient processing by shocks and stellar photons which is primarily reflected in the dust size distribution \\citep[see for instance][]{jones97, jones2004}.  Such evolution is expected to be particularly significant between the dense and diffuse galactic environments.  For instance, the presence of small dust particles in the diffuse galactic medium is necessary to account for the near-infrared (NIR) and mid-infrared (mid-IR) observed emission \\citep[e.g.][]{Sellgren84, draine85} while in dense molecular environment these small particles must coagulate with the larger species \\citep[e.g.][]{stepnik2003}.  Dust has a strong impact on interstellar gas because of efficient coupling such as gas heating by the photoelectric effect or the formation of H$_2$ on grain surfaces. In addition, these gas-grain processes are dominated by the contribution of small dust grains of radius $a<100$ nm \\citep[e.g.][]{habart2001, weingartner2001c, habart2004}.  \\begin{figure*} \\centering \\includegraphics[width=0.62\\textwidth]{09850fig1a.ps} \\includegraphics[width=0.36\\textwidth]{09850fig1b.ps} \\caption{{\\bf Left panel:} The L1630 molecular cloud as seen in the visible (photographic plate obtained using the UK Schmidt Telescope and extracted from the Digitized Sky Survey produced at the Space Telescope Institute). The arrow indicates the direction of the $\\sigma$\\,Orionis binary. {\\bf Right panel:} NGC2023 (top) and the Horsehead nebula (bottom) as seen by ISOCAM-LW2 at 5-8.5\\,$\\mu$m \\citep[][]{abergel2002}. The contours show areas observed with Spitzer-IRS as part of the ``SPECPDR'' program \\citep[][]{joblin2005}. The small and large areas delineated by the white line are those observed with the SL and LL modules, respectively.} \\label{fig:OrionB} \\end{figure*}  Located at the illuminated edges of molecular clouds, photodissociation regions \\citep[hereafter PDRs; for a review see e.g.][]{hollenbach97} are regions of transition between dense and diffuse media.  In addition to the generally large density gradient observed in such regions, there is a gradient of the UV-visible radiation field intensity due to the proximity to a (some) young star(s) and to the absorption of the radiation field by the dust. Consequently, PDRs are the site of significant evolution of both the dust and gas properties which are closely interlinked.  Dust evolution may be a key to the understanding of the unexpectedly large amount of warm H$_2$ in PDRs \\citep[e.g.][]{allers2005, habart2008}.  PDRs are responsible for the reprocessing of most of the UV-visible radiative energy output from stars which is re-emitted in infrared-millimeter wavelengths \\citep[e.g.][]{hollenbach97} and consequently are important IR emitters of normal star forming galaxies. It is thus necessary to understand in detail the emission coming from PDRs in order to interpret IR galaxy spectra.  Infrared Space Observatory (ISO) observations showed a systematic diminution of the emission ratio between aromatic infrared bands (AIBs between 3 and 17\\,$\\mu$m) emitted by polycyclic aromatic hydrocarbons (PAHs) and the mid-IR continuum emission from dense illuminated ridges \\citep[e.g.][]{cesarsky2000,abergel2002}.  \\citet{rapacioli2005} have interpreted this evolution in terms of the chemical evolution of small carbonaceous particle properties by decomposing ISOCAM/CVF spectral cubes (5-16\\,$\\mu$m) with the Single Value Decomposition method. \\citet{berne2007} have extended this study by using the Blind Signal Separation method on Spitzer IRS data.  A limitation of these methods is that they do not take into account the radiative transfer of stellar photons which may affect the shape of the mid-IR spectrum just as dust properties and abundances do.  The aim of this paper is to study the dust evolution in two well known PDRs as seen by Spitzer/IRS and ISOCAM through the modelling of the dust mid-IR emission including radiative transfer. Although other dust properties like the size distribution and/or composition can change, we will interpret observed mid-IR spectral variations in term of PAH/VSG relative abundance.  In section \\ref{sect:presentation} and \\ref{sect:obs}, we present the studied PDRs and the Spitzer data.  The section \\ref{sect:result_obs} is devoted to the observed mid-IR spectral shape in these PDRs.  In section \\ref{sect:dust_PDR_model}, we present our dust model (\\S\\,\\ref{sect:dust_model}) and the radiative transfer (\\S\\,\\ref{sect:transfer_model}) used to model dust emission in PDRs. Sections \\ref{sect:model_HH} and \\ref{sect:model_2023} present results of the modelling of the Horsehead Nebula and NGC2023 respectively. We conclude in section\\,\\ref{sect:conclusion}. ",
        "conclusions": "\\label{sect:conclusion}  We presented mid-IR spectral imaging observations of the Horsehead nebula and NGC2023N obtained with the infrared spectrograph on board the Spitzer Space Telescope. These observations allow us to confirm the AIBs / mid-IR continuum evolution at dense illuminated ridges already observed with the Infrared Space Observatory \\citep[e.g.][]{abergel2002, rapacioli2005}.  We developed a new dust emission model based on the \\citet{desert90} model. This model successfully reproduces the emission and the extinction curve of the diffuse interstellar medium at high galactic latitude (i.e. the Cirrus). The Cirrus dust properties are then used as a reference to be compared to dust properties in the two studied PDRs. This dust emissivity model is coupled to a radiative transfer model in order to properly take into account excitation effects on the mid-IR spectral shape evolution.  For the PDRs studied here, we conclude that excitation effects cannot account for the observed AIBs / mid-IR continuum evolution. We interpret this evolution in term of PAH/VSG relative abundance variation since these two species dominate the 5-35\\,$\\mu$m spectrum of the two studied PDRs except at $\\lambda\\,\\ga25\\,\\mu$m for positions close to the star (in the cavity) in NGC2023 where the BG continuum is significant.  In the Horsehead Nebula case, we do not spatially resolve any spectral variation. We conclude from the study of 7-9\\,$\\mu$m (AIBs) and 22-24\\,$\\mu$m (mid-IR continuum) emission profiles and of the spectrum at the emission peak position that the PAH/VSG relative abundance is 2.4 times lower than in the Cirrus. In NGC2023N, we resolve the spectral variation. We modelled the spectrum of the diffuse medium around the star for a position at the external edge of the dense illuminated ridge traced by H$_2$ $\\nu$\\,=\\,1-0\\,S(1). This spectrum can be reproduced by using Cirrus dust properties. We attributed the decrease of a factor of $\\sim$5 of the 7-9\\,$\\mu$m\\,/\\,22-24\\,$\\mu$m (AIBs / mid-IR continuum) from the diffuse illuminated part to the deep/dense part of the PDR to a decrease of the PAH/VSG relative abundance to $\\sim$1/5 of the Cirrus value. It then seems that dust properties evolve fully \"from dense to diffuse\" properties within the small spatial scale of the dense illuminated ridge. Note that the obtained PAH\\,/\\,VSG relative abundance for the \"unresolved\" Horsehead is almost the median between the diffuse and dense part of NGC2023N, indicating the mean properties of the observed very small particles in the Horsehead.  An increase of the PAH/VSG relative abundance from the dense/deep to the diffuse illuminated part of the PDR is fully consistent with a scenario of photoevaporation PAH clusters (i.e, of the VSGs) \\citep[e.g.][]{cesarsky2000,rapacioli2005, berne2007}. Other physical processes could explain this variation, such as (i) the scenario evoked in \\citet{jones2005} of the evolution of aliphatic hydrocarbons (which do not emit AIBs) which are present in the dense media \\citep[e.g.][]{jones90, jones90b, dartois2007} into aromatic hydrocarbons under the effect of the UV radiation field \\citep[e.g.][]{dartois2005, jones90}, (ii) the release (\"decoagulation\") of the VSGs and PAHs from the dust aggregates (present in dense clouds) at different depths into the PDRs or (iii) a size segregation effect due to the grain dynamics induced by the anisotropic radiation field, as already mentionned for NGC2023 by \\citet{abergel2002}.  All these processes could be at work with different efficiencies depending on the depth into the cloud and/or on the size and nature of dust.  This strong evolution of the PAH relative abundance between the dense and diffuse medium in PDRs could be a clue for the interpretation of the observed (PAH 8\\,$\\mu$m)/24\\,$\\mu$m and (PAH 8\\,$\\mu$m)/160\\,$\\mu$m variations in the SINGS galaxies \\citep[e.g.][]{bendo2008}. Such a PAH relative abundance evolution in the bright IR emitters that are PDRs will have an impact on the estimation of the star formation activity of galaxies using PAH emission."
    },
    "0809/0809.0329_arXiv.txt": {
        "abstract": "Disk galaxies in compact galaxy groups exhibit a remarkable shortfall of neutral hydrogen compared to both isolated spirals and spirals in more loose groups, but the origin of this H{\\sc i} deficiency remains unclear. Based on a sample of highly H{\\sc i} deficient compact galaxy groups, here updated to also include HCG\\,58 and 93, we summarise the first results of a multi-wavelength campaign aimed at understanding the processes responsible for modifying the H{\\sc i} content of galaxy disks in these environments. While tidal stripping, ram pressure stripping by hot intragroup gas, and star-formation induced strangulation could individually be affecting the ISM in some of the group members, these processes each face specific difficulties in explaining the inferred deficiency of H{\\sc i} for the sample as a whole. A complete picture of the mechanisms driving the ISM evolution in the disk galaxies of these groups has thus yet to emerge, but promising avenues for further progress in this field are briefly discussed on the basis of the present sample. ",
        "introduction": "A substantial fraction of all galaxies in the nearby Universe reside in a group or cluster environment. Disk galaxies are not only less common within such environments than in the field, but they also tend to be deficient in neutral hydrogen. For a group or cluster as a whole, this H{\\sc i}~deficiency $\\Delta_{\\rm HI}$ can be quantified as \\begin{equation} \\Delta_{\\rm HI} = {\\rm log} M_{\\rm{HI,pred}} - {\\rm log} M_{\\rm{HI,obs}}, \\label{eq,hi_def} \\end{equation} where $M_{\\rm{HI,obs}}$ is the total observed H{\\sc i} mass of the group or cluster, and $M_{\\rm{HI,pred}}$ is that predicted for a corresponding ensemble of isolated galaxies of similar optical morphology and luminosity.  At the scale of small groups, comprehensive radio studies have shown such H{\\sc i} deficiencies to be particularly pronounced in \\cite{hick82} compact groups (HCGs) (\\cite[Verdes-Montenegro et al.\\ 2001]{verd01}). Tidal interactions can be expected to play a prominent role in affecting the evolution of galaxy disks in these systems, given the compactness and low velocity dispersions ($\\sigma \\sim$ a few hundred km~s$^{-1}$) of the environment.  However, there is also an indication from {\\em ROSAT} data that significant H{\\sc i} deficiency in these groups correlates with the presence of a detectable X-ray emitting intragroup medium (IGM), suggesting that galaxy--IGM interactions such as ram pressure stripping of H{\\sc i} may also be important.  Since a non-negligible fraction of all groups in the Universe (including so-called `fossil' groups) may at some point experience a phase resembling that of HCGs, an improved understanding of the impact of the HCG environment on the group members may well have ramifications extending beyond these somewhat atypical groups. In order to explore the origin of the shortfall of H{\\sc i} in HCGs, and more generally to shed light on disk galaxy evolution in these environments, we have therefore embarked on a multi-wavelength campaign aimed at establishing the detailed properties of the gas and galaxies in the most H{\\sc i} deficient HCGs. The initial focus has been on the properties of any X-ray emitting IGM, in an attempt to constrain the importance of galaxy--IGM interactions in these groups. ",
        "conclusions": "{\\scriptsize \\begin{tabular}{lccclc}\\hline Group & Dist. & $\\Delta_{\\rm HI}$ & $\\sigma$ & X-ray obs. & Expo. time \\\\ & (Mpc) & & (km s$^{-1}$) & & (ks) \\\\ \\hline HCG\\,7     & 54 & 0.60 &  \\hspace{0.4mm} 95  & {\\em XMM}     & 29 \\\\ HCG\\,15   & 92 & 0.46 & 404 & {\\em XMM}     & 26 \\\\ HCG\\,30   & 63 & 1.37 &  \\hspace{0.4mm} 72  & {\\em Chandra} & 29 \\\\ HCG\\,37   & 97 & 0.33 & 446 & {\\em Chandra} & 18 \\\\ HCG\\,40   & 98 & 0.60 & 157 & {\\em Chandra} & 46 \\\\ HCG\\,44   & 23 & 0.69 & 145 & {\\em Chandra} & 20 \\\\ HCG\\,58   & 89 & 0.51 & 178 & {\\em ROSAT}     & 11 \\\\ HCG\\,93   & 64 & 0.99 & 234 &  {\\em ROSAT}/{\\em Swift}   & 14/4 \\\\ HCG\\,97   & 86 & 0.35 & 383 & {\\em Chandra} & 36 \\\\ HCG\\,100  & 69 & 0.27 & 100 & {\\em Chandra} & 42 \\\\ \\hline \\end{tabular} } \\end{center} \\vspace{1mm} \\end{table}   The two groups without {\\em Chandra} or {\\em XMM} data, HCG\\,58 and 93, have been targeted in {\\em ROSAT} pointings. As a by-product of our effort to also obtain UV data for the sample, we have recently acquired a short {\\em Swift} X-ray exposure of HCG\\,93, with similar data underway for HCG\\,58. No obvious diffuse X-ray emission is detected in the {\\em Swift} data of HCG\\,93, consistent with the earlier {\\em ROSAT} result of \\cite{ponm96}, but the sensitivity of the {\\em Swift} data is insufficient to provide strong constraints on the density of any IGM. The IGM in HCG\\,58 also remained undetected in the {\\em ROSAT} analysis of \\cite{osmo04}.      The {\\em Chandra} and {\\em XMM} data, in most cases representing an improvement in sensitivity by 1--2 orders of magnitude over the previous {\\em ROSAT} data, manifestly show that the presence of a substantial hot IGM is not a prerequisite for high H{\\sc i} deficiency in these groups, despite earlier indications of a connection (Verdes-Montenegro et al.\\ 2001).  Although galaxy--IGM interactions can clearly have affected the H{\\sc i} disks of galaxies within X-ray bright groups such as HCG\\,97, such processes cannot generally be dominant in removing H{\\sc i} from the group members in our sample.  Other mechanisms could help account for the observed shortfall of H{\\sc i} within these groups. H{\\sc i}~consumption by star formation, aided by the removal of a continuous supply of H{\\sc i} to the disk, could be playing a prominent role. One problem faced by this scenario, however, is that the time-scale required to exhaust the H{\\sc i} supply to observed levels (assuming $\\Delta_{\\rm HI}=0$ initially), is, on average, at least five~Gyr at the current star formation rates (SFRs) in the groups. Strangulation may therefore not have been important in establishing current H{\\sc i} levels within our sample.  Another possible explanation for the lack of H{\\sc i} are recent tidal interactions, clearly taking place in some of the groups as evidenced by Figure~3.  However, analysis of SDSS data has established that such interactions are commonly accompanied by a clear SFR enhancement (e.g., Li et al.\\ 2008), and yet the SFRs in Hickson groups as determined from {\\em IRAS} fluxes are not elevated relative to values for field spirals (Verdes-Montenegro et al.\\ 1998). Our X-ray analysis also shows no clear evidence of enhanced nuclear activity within galaxies in the X-ray brighter or more H{\\sc i}~deficient groups, suggesting that strong nuclear starbursts or AGN activity triggered by tidally induced gas inflows are not more common or prominent within the more `evolved' groups. Typical indirect signatures expected of tidal interactions are thus generally very modest in these groups, and such interactions may furthermore not themselves destroy H{\\sc i}.  In summary, obvious mechanisms that could be invoked to explain the pronounced deficiency of H{\\sc i} observed for the sample clearly each face some difficulties. The X-ray results indicate that galaxy--IGM interactions may have played a role in destroying H{\\sc i} in the X-ray bright groups, but it remains unclear whether, for example, tidal interactions on their own can explain the reduced H{\\sc i} content in the remaining systems."
    },
    "0809/0809.1928_arXiv.txt": {
        "abstract": "Many optimization techniques have been invented to reduce the noise that is inherent in Monte Carlo radiative transfer simulations. As the typical detectors used in Monte Carlo simulations do not take into account all the information contained in the impacting photon packages, there is still room to optimize this detection process and the corresponding estimate of the surface brightness distributions. We want to investigate how all the information contained in the distribution of impacting photon packages can be optimally used to decrease the noise in the surface brightness distributions and hence to increase the efficiency of Monte Carlo radiative transfer simulations.  We demonstrate that the estimate of the surface brightness distribution in a Monte Carlo radiative transfer simulation is similar to the estimate of the density distribution in an SPH simulation. Based on this similarity, a recipe is constructed for smart detectors that take full advantage of the exact location of the impact of the photon packages. Several types of smart detectors, each corresponding to a different smoothing kernel, are presented. We show that smart detectors, while preserving the same effective resolution, reduce the noise in the surface brightness distributions compared to the classical detectors. The most efficient smart detector realizes a noise reduction of about 10\\%, which corresponds to a reduction of the required number of photon packages (i.e.\\ a reduction of the simulation run time) of 20\\%. As the practical implementation of the smart detectors is straightforward and the additional computational cost is completely negligible, we recommend the use of smart detectors in Monte Carlo radiative transfer simulations. ",
        "introduction": "\\label{introduction.sec}  The Monte Carlo method \\citep[e.g.][]{1959CashEver, 1996ApJ...465..127B, 2001ApJ...551..269G, 2003MNRAS.343.1081B, 2003A&A...399..703N} has become one of the most popular methods to perform radiative transfer simulations. One of the greatest advantages of the Monte Carlo method is its conceptual simplicity. Instead of solving the radiative transfer equations that describe the radiation field, Monte Carlo simulations actually follow the photon packages that make up the radiation field in a very natural way. This ensures that, compared to other approaches to multi-dimensional radiative transfer problems, the practical implementation of Monte Carlo radiative transfer is surprisingly straightforward. This simple, straightforward approach also enables the inclusion of additional ingredients, such as the polarization of scattered light \\citep{1995ApJ...441..400C, 1996ApJ...465..127B}, the kinematics of the sinks and sources \\citep{2001ApJ...548..150M, 2002MNRAS.335..441B} and gas ionization and recombination processes \\citep{2003MNRAS.340.1136E, 2005MNRAS.362.1038E}. On the other hand, probably the most important disadvantage of the Monte Carlo method is the appearance of Poisson noise, which is inherently tied to the probabilistic nature of the method. In pure Monte Carlo radiative transfer simulations, the noise in any observed property goes as $1/\\sqrt{N}$ with $N$ the number of photon packages, which means that one needs to quadruple the number of photon packages to halve the errors. Due to this slow convergence, radiative transfer simulations based on the most simple application of the Monte Carlo method are very inefficient and have difficulties to compete with other methods.  Fortunately, in the many years since the first applications of the Monte Carlo method to radiative transfer simulations, several intelligent optimization techniques have been invented to increase the efficiency of the Monte Carlo techniques. These techniques, equivalent to so-called variance reduction techniques in Monte Carlo integration, aim at reducing the noise by including deterministic elements into the probabilistic simulation. One of the first examples of such techniques is the well-known forced first scattering technique, already included in the first implementations of Monte Carlo radiative transfer \\citep{1959CashEver, 1970A&A.....9...53M}. Other noise-reduction techniques that have strongly increased the efficiency of the Monte Carlo method include the use of weighted photon packages \\citep{1977ApJS...35....1W}, the peel-off technique \\citep{1984ApJ...278..186Y}, the treatment of absorption as a continuous rather than a discrete process \\citep{1999A&A...344..282L}, the frequency distribution adjustment technique \\citep{2001ApJ...554..615B, 2005NewA...10..523B} and the use of polychromatic photon packages \\citep{2005AIPC..761...27B, 2006MNRAS.372....2J}.  One aspect of Monte Carlo radiative transfer where no significant noise reduction techniques have been presented is the last step in the life cycle of the photon packages, namely their detection. The goal of most Monte Carlo radiative transfer simulations is to determine the observed surface brightness distribution at some observer's position. To construct the observed surface brightness distribution, some kind of detector must be simulated, on which the photon packages that leave the system are recorded. In the spirit that Monte Carlo radiative transfer simulations mimic the real physical processes as naturally as possible, the simulated detectors are usually natural (idealized) representations of actual CCD detectors. They basically consist of a two-dimensional array of pixels, which act as a reservoir for the incoming photon packages. When a photon package leaves the system and arrives at the location of the observer, the correct bin is determined and the luminosity recorded in the bin is increased with the luminosity of the photon package. At the end of the simulation, the detector is read out like a CCD detector and the surface brightness distribution is constructed.  While this approach seems the most natural way to simulate the detection of photon packages in a Monte Carlo simulation, it might not be the most efficient. We must be aware that, although we are simulating a real detection as closely as possible, we have {\\em{more}} information at our disposal than real observers. The maximum information that a real observer can obtain (in the academic limit of perfect noise-free observations and instruments) when imaging with a CCD detector, is the number of photon packages that arrive in each of his pixels. As numerical simulators, we have at our disposal the full information on the precise location of the impact of each photon package on the detector. It would be a pity to throw away this useful additional information just in order to mimic the behaviour of a real CCD detector.  The goal of this paper is to investigate how this additional information can be used to decrease the noise in the estimated surface brightness distributions and hence to increase the efficiency of Monte Carlo radiative transfer simulations. In Section~2 we describe the classical way of detecting photon packages as a smoothing process, similar to the averaging processes encountered in the smoothed particle hydrodynamics framework for hydrodynamical simulations. We use this similarity between both processes to construct a new set of {\\em{smart}} detectors, which aim at curing two drawbacks of the classical detectors. In Section~3 we test the accuracy and performance of these smart detectors and compare them with the classical detector. The results are summarized and discussed in Section~4. ",
        "conclusions": "We have focused on an ill-studied area of Monte Carlo radiative transfer simulations where a significant noise reduction can be achieved, namely the detection of photon packages and the corresponding construction of the observed surface brightness distribution. The motivation of this work was the fact that the classical detectors used in Monte Carlo simulations, while closely mimicking a real CCD detector, do not use the full amount of information that is available.  Based on the similarities between the construction of the surface brightness distribution on a detector in Monte Carlo radiative transfer simulations and the calculation of the density in SPH hydrodynamical simulations, we have constructed a set of smart detectors. These smart detectors improve on two aspects of the classical detector: they assign the same weight to all photon packages hitting the detector at the same distance from a grid point and they give more weight to impacts close to a grid point than to impacts at larger distances. We have tested different kinds of smart detectors, based on different smoothing kernels frequently encountered in SPH simulations, namely a gaussian kernel and two spline-based kernels with finite support \\citep{1981csup.book.....H, 1985A&A...149..135M}. We have shown that these new detectors, while preserving the same effective resolution, reduce the noise in the surface brightness distributions compared to the classical detectors. It is demonstrated that the lion's share of this noise reduction is due to a proper weighing of the photon packages with the distance between the impact location and the grid point.  The most efficient smart detector is found to be a detector based on the $M_3$ spline kernel, for which the noise reduction amounts to some 10\\%. While this might seem a modest improvement, one should take into account that the noise reduction in Monte Carlo simulation goes as $1/\\sqrt{N}$. A reduction of the noise with 10\\% is hence equivalent to a reduction of the number of photon packages with 20\\%, which is a significant improvement. Moreover, it should be stressed that this noise reduction basically comes for free, since the practical implementation of the smart detectors is straightforward and the additional computational cost is completely negligible. We hence strongly recommend the use of smart detectors in Monte Carlo radiative transfer simulations.  The link with SPH simulations might stimulate to look for further ways to optimize the estimates of the surface brightness distribution. One major difference between our current problem and the interpolation in SPH simulations is that we use a fixed smoothing length, whereas SPH simulations typically use a spatially (and temporally) varying smoothing length. This way it is possible to take full advantage of the particle distribution to resolve local density structures. Each particle in an SPH simulation typically has an individual smoothing length which is fine-tuned such that each particle interacts with a similar number of neighbors. Unfortunately, it seems hard to introduce variable smoothing lengths here in our current Monte Carlo radiative transfer case. The main difference is that the smoothing procedure in SPH is executed when all particle positions are known, whereas in Monte Carlo radiative transfer the photon packages gradually hit the detector and they must be smoothed out onto the detector before it is known where the other photon packages will arrive. It is impossible to know {\\em{a priori}} how many neighbors a photon package will ultimately have and hence to adapt the smoothing length on an individual basis. One could potentially do a test run with a limited number of photon packages to obtain a first crude estimate of the expected surface brightness distribution and adapt the smoothing lengths accordingly. However, a more elegant option seems to apply an adaptive filtering to the simulated images {\\em{after}} the simulation. Several approaches have been developed for this goal such as wavelet-based algorithms \\citep{1998A&AS..128..397S}, adaptive binning techniques \\citep{1998A&AS..128..397S, 2003MNRAS.342..345C} or adaptive kernel smoothing techniques \\citep{1991AN....312..345R, 1996ApJ...461..622H, 2006MNRAS.368...65E}."
    },
    "0809/0809.4166_arXiv.txt": {
        "abstract": "We present the results of an archival search for Trans-neptunian objects (TNOs) in an ecliptic field observed with Subaru in 2002. The depth of the search allowed us to find 20 new TNOs with magnitudes between $R=24$ and $27$. We fit a double power law model to the data; the most likely values for the bright and faint power law exponents are $\\alpha_1$=$0.73_{-0.09}^{+0.08}$ and $\\alpha_2$=$0.20_{-0.14}^{+0.12}$; the differential number density at $R=23$ is $\\sigma_{23}$=$1.46_{-0.12}^{+0.14}$ and the break magnitude is $R_{eq}$=$25.0_{-0.6}^{+0.8}$.  This is the most precise measurement of the break in the TNO luminosity function to date. The break in the size distribution corresponds to a diameter of $D = 90\\pm30$ km assuming a 4\\% albedo. ",
        "introduction": "\\label{sec:intro} The TNO size distribution is the result of the formation and collisional history of its members. Models of the physical evolution assume the processes at play result in a characteristic power law size distribution. In the model of \\citet{Pan.2005} of strengthless TNOs, the collisional evolution of the objects leads to an evolving size distribution that changes slope at the size at which the collisional lifetime of an object equals the age of the system. This analytical result confirms those obtained from numerical simulations~\\citep{Kenyon.1999,Davis.1999,Kenyon.2004}. The precise location of the break in the size distribution is an important measurement to link TNO formation and evolution models to the current TNO population (see \\citealt{Kenyon.2008} for a review).  The TNO luminosity function is defined by the distance, albedo, and size distributions of TNOs, and it is readily observed.  The distance distribution exhibits a sharp edge at $50\\au$ \\citep{Trujillo.2001} that has been observed in other surveys \\citep{Gladman.2001, Fuentes.2008}. Though it is customary to use a 4\\% albedo, correlations with size have been reported (see \\citealt{Stansberry.2008} for a review). After making sensible assumptions for the albedo and distances of the observed TNOs, the size distribution can be derived from the luminosity function. Several searches have been completed (see \\citealt{Bernstein.2004, Fuentes.2008} for a review), the deepest being the search by \\citet{Bernstein.2004} with the {\\it Hubble Space Telescope} (HST). While previous results sampled the size distribution up to magnitude $R=26$ and were consistent with a power law luminosity function, the \\citet{Bernstein.2004} survey found a significant underabundance of objects at $R\\sim28$, compared to that predicted from extrapolating the luminosity function at brighter magnitudes. They concluded a break in the luminosity function must occur between $R\\sim25$ and $28$.  Since \\citet{Bernstein.2004} announced their results, a number of ground-based surveys for fainter TNOs have been conducted in an effort to bridge the divide between the brighter ($R$$\\sim$$23$) TNOs found through wide-field searches such as the Deep Ecliptic Survey~\\citep{Elliot.2005} and the Canada-France Ecliptic Plane Survey~\\citep{Jones.2006} and those found by \\citet{Bernstein.2004}. \\citet{Fuentes.2008} discovered 82 TNOs in an archival search of Subaru data.  This survey reached a limiting magnitude of $R=25.7$ and successfully detected the break in the luminosity function. In the present work we build upon that earlier work, increasing the limiting magnitude of ground-based surveys to $R\\sim27$ and further narrowing the gap between ground-based and space-based TNO surveys. Constraints on even fainter objects have been placed by stellar occultation surveys \\citep{Bickerton.2008, Zhang.2008}.  We present the details of the data in the next section. \\S~\\ref{sec:mod} outlines the method used to analyze and search for TNOs in the data. The characterization of the search using a synthetic population is presented in \\S~\\ref{sec:effic}. In \\S~\\ref{sec:results} and \\S~\\ref{sec:conc} we present our results and discuss what they imply for the formation and evolution of the TNO population. ",
        "conclusions": "\\label{sec:conc} We have performed a pencil-beam search of a single Subaru ecliptic field. The search covered 0.255$\\sqdeg$, with a limiting magnitude $R=26.76$. We found 20 new TNOs with magnitudes between $R=24.0$ and $26.8$.  As argued by \\citet{Bernstein.2004}, it is not surprising that our faintest detection is near our 50\\% efficiency threshold, given that the luminosity function is much shallower at $R\\sim27$.  Including other surveys in the analysis (see \\S~\\ref{sec:results}) we derive the most likely DPL model for the luminosity function. With only 20 new objects there is an important improvement in the model's PDF from \\citet[Figure 13]{Fuentes.2008} to the current result (See Figure \\ref{fig:mcmc}); due to the new detections constraining the model between $R=26$ and $27$.  The bright end power-law exponent $\\alpha_1$=$0.73_{-0.09}^{+0.08}$ is very close to the results of previous shallower surveys \\citep{Gladman.2001}. The break magnitude $R_{eq}$=$25.0_{-0.6}^{+0.8}$ is more than 1-$\\sigma$ fainter than the initial estimate by \\citet{Bernstein.2004}.  This is also consistent with earlier surveys reporting excellent fits to a single power model for magnitudes brighter than $R\\sim26$~\\citep{Gladman.2001}. Further data between the current ground-based detection limit of $R\\sim27$ and the HST's at $R\\sim28.5$ will determine the break magnitude even more accurately.  Assuming that all objects are located at an heliocentric distance of 42$\\au$, the break in the luminosity function $R_{eq}$ reflects a break in the size distribution $D$. The corresponding diameter at which the size distribution breaks is $D=90\\pm30(p/0.04)^{-0.5}\\km$, where $p$ is the albedo. Theories predict the existence of a break \\citep{Kenyon.2008}, but at smaller diameters ($D\\sim20-40\\km$). The difference reflects the uncertainty in the initial conditions for numerical simulations and the observational assumptions regarding physical properties like the albedo.  In this survey both classical and excited TNO populations are entangled. With better orbital information the size distribution of each can be studied separately. A comparison with other populations in the solar system can shed light on their origin."
    },
    "0809/0809.1096_arXiv.txt": {
        "abstract": "Residual star formation at late times in early-type galaxies and their progenitors must be suppressed in order to explain the population of red, passively evolving systems we see today. Likewise, residual or newly accreted reservoirs of molecular gas that are fuelling star formation must be destroyed. This suppression of star formation in early-type galaxies is now commonly attributed to AGN feedback wherein the reservoir of gas is heated and expelled during a phase of accretion onto the central supermassive black hole. However, direct observational evidence for a link between the destruction of this molecular gas and an AGN phase has been missing so far. We present new mm-wavelength observations from the IRAM 30m telescope of a sample of low redshift SDSS early-type galaxies currently undergoing this process of quenching of late-time star formation. Our observations show that the disappearance of the molecular gas coincides within less than $100$ Myr with the onset of accretion onto the black hole and is too rapid to be due to star formation alone. Since our sample galaxies are not associated to powerful quasar activity or radio jets, we conclude that low-luminosity AGN episodes are sufficient to suppress residual star formation in early-type galaxies. This `suppression mode' of AGN feedback is very different from the `truncation mode' linked to powerful quasar activity during early phases of galaxy formation. ",
        "introduction": "The low redshift galaxy population can be divided into blue starforming and red passively evolving systems \\citep{2004ApJ...600..681B}. The red population is dominated by elliptical and lenticular (early-type) galaxies, and in order to explain the properties of these systems as we observe them, almost all residual star formation in them must be suppressed. This suppression must take place via the heating or expulsion of the molecular gas which is the fuel for star formation \\citep{1986ApJ...303...39D, 2003ApJ...599...38B} and is thought to be driven by the energy input from accreting supermassive black holes at the galaxies' centers in a process called \\textit{AGN feedback} \\citep{1998A&A...331L...1S, 2003ApJ...599...38B, 2005Natur.433..604D, 2005MNRAS.364..407C, 2005MNRAS.358L..16K, 2006Natur.442..888S, 2007MNRAS.382..960K, 2007arXiv0712.3289K}.  The stellar populations of early-type galaxies indicate that the most massive galaxies are the first to undergo this process of suppression, forming their stars earliest and on the shortest time scales \\citep{2005ApJ...621..673T, 2005ApJ...632..137N, 2006AJ....131.1288B, 2007ApJ...669..947J, 2008MNRAS.388...67K, 2008MPLA...23..153K}. These intense episodes of star formation are followed by passive evolution with almost no subsequent star formation, although evolution via processes such as dry mergers is still possible (e.g. \\citealt{2003ApJ...597L.117K, 2005AJ....130.2647V}). It has been suggested that the mass at which galaxies undergo the transition from star-forming to quiescent decreases over cosmic time \\citep{2005ApJ...625...23B, 2006ApJ...651..120B}, and so less massive galaxies have more extended histories of star formation \\citep{2005ApJ...621..673T}.  As new molecular gas is continuously supplied by accretion, mergers and stellar mass loss, most current models of galaxy formation suppress residual star formation by invoking a phase of accretion onto the supermassive black hole at the center of the galaxy \\citep{1998A&A...331L...1S, 2005MNRAS.364..407C, 2006Natur.442..888S, 2006MNRAS.365...11C} as the energy liberated by core-collapse supernovae is insufficient \\citep{1986ApJ...303...39D, 2003ApJ...599...38B}. The resulting active galactic nucleus (AGN) phase includes jets, radiation and outflows liberating sufficient energy to destroy the molecular gas reservoir by heating and expelling it, thus suppressing star formation by depriving the host galaxy of its fuel. Observations indicate that the mass of galaxies and their supermassive black holes are correlated \\citep{2000ApJ...539L...9F, 2000ApJ...539L..13G}, suggesting that the two co-evolve. The destruction of the gas reservoir may then in turn terminate the growth of the black hole as it deprives itself of material for further accretion as these systems transition to a maintenance mode suppressing cooling flows due to stellar mass loss and cold accretion \\citep{1997ApJ...487L.105C, 2007ApJ...665.1038C, 2006MNRAS.370..645B}. Alternative -- or perhaps complementary -- gas heating mechanisms to AGN feedback have been proposed by \\cite{2007MNRAS.380..339B} and \\cite{2008ApJ...680...54K}. Investigating the role of AGN feedback directly is thus an important observational goal.  Early-type galaxies are known to pass through episodes of accretion onto their central black hole, but evidence which links this phase of their evolution to the removal of the molecular gas reservoir has until now been lacking. In this Paper, we present observational evidence for this process occurring in low- to intermediate-mass early-type galaxies at low redshift and identify low-luminosity AGN as the culprit.  Throughout this work, we assume cosmological parameters $(\\Omega_{m} = 0.3, \\Omega_{\\Lambda} = 0.7, H_{0} = 70)$, consistent with the \\textit{Wilkinson Microwave Anisotropy Probe} (WMAP) Third Year results and their combination of results with other data \\citep{2007ApJS..170..377S}. ",
        "conclusions": ""
    },
    "0809/0809.4458_arXiv.txt": {
        "abstract": "We show that the open cluster Trumpler~20, contrary to the earlier findings, is actually an old Galactic open cluster. New CCD photometry and high-resolution spectroscopy are used to derive the main parameters of this cluster. At [Fe/H]=$-$0.11 for a single red giant star, the metallicity is slightly subsolar. The best fit to the color-magnitude diagrams is achieved using a 1.3 Gyr isochrone with convective overshoot. The cluster appears to have a significant reddening at E(B-V)=0.46 (for B0 spectral type), although for red giants this high reddening yields the color temperature exceeding the spectroscopic $T_{\\rm eff}$ by about 200~K. Trumpler 20 is a very rich open cluster, containing at least 700 members brighter than $M_V=+4$. It may extend over the field-of-view available in our study at 20$\\arcmin\\times$20$\\arcmin$. ",
        "introduction": "The number of old open clusters (age $\\agt$1 Gyr) is $\\sim$60 \\citep{fr95} or only $\\sim$4\\% of the all known clusters in the Milky Way \\citep{di04}. Because of this small number of such clusters and the fact that many of them are poorly populated, any addition to this body of old clusters is highly desirable. Here we report a serendipitous discovery of an old open cluster among the known registered clusters in the WEBDA\\footnote{http://www.univie.ac.at/webda/}. It seems to be enlightning and worthy to reflect on what initially led us to this object. In March of 2008 we had an observing run at the 2.2m MPG/ESO telescope with the FEROS spectrograph. For the morning hours, there were no regular targets available and, hence, we started to look for some old open cluster with 12$^{\\rm h}$$<$RA$<$16$^{\\rm h}$. There are 7 such clusters in the WEBDA but none of them appeared to be well-suited for our instrumental setup or offering a compelling science. A search for other potential targets led us to a large photometric survey of 55 young open clusters in the southern sky by \\citet{mc05}. The primary goal of their survey was to examine the age and evolutionary dependence of the Be star phenomenon. Examining the provided color-magnitude diagrams (CMD) our attention was caught by the CMD of Trumpler~20. It was obvious that for this cluster the authors had missed out a characteristic signature of all old open cluster -- that is the red clump, prominently displayed in its CMD at $y=14.5$ and $b-y=0.9$ \\citep[Fig. 59]{mc05}. These authors opted  for a young 160 Myr cluster, instead.  The apparent lack of reliable parameters for Trumpler~20 (Tr~20 hereafter) prompted the present study.  In the discovery paper \\citep{tr30}, Tr~20 (=C1236-603) is described as a \"rich cluster of regular outline composed exclusively of very faint stars nearly evenly scattered\" and classified as a type III~2r cluster. Using the calibrations between the estimated diameter in parsecs and the apparent angular diameter, for this cluster with an angular diameter of 10$\\arcmin$ Trumpler reports a distance equal to 2240~pc. A few decades later, \\citet{ho65} estimated the total number of cluster members from starcounts to be at $\\sim$200 down to $m_{pg}$$\\sim$17.  The Str\\\"{o}mgren $by$ photometry given in \\citet{mc05} has a limited potential in deriving a photometric distance of Tr~20 because it covers as little as $\\sim$1.5 mag of its main sequence. One helpful hint from the paper by \\citet{mc05} is the noted relatively high reddening in this area of the sky at $E(B-V)$$\\sim$0.35.  For the open clusters located within $2\\degr$ from Tr~20, the WEBDA database offers additional reddening estimates, reaching up to $E(B-V)=0.4$.  The access to high-resolution spectroscopy at the 2.2m MPG/ESO gave us an opportunity to measure radial velocities for a few possible Tr~20 members and even attempt to measure its metallicity [Fe/H]. Another fortunate concurrence involved an access to photometric  observations at the 1m telescope of the Cerro Tololo Inter-American Observatory (CTIO). These two facilities allowed us to collect the necessary data and were instrumental in obtaining fundamental properties of Tr~20. ",
        "conclusions": "Trumpler 20 has several features which makes it attractive for the future studies. There is no doubt that Tr~20 is a $\\sim$1.3 Gyr old open cluster. In its age group Tr~20 appears to be the richest Galactic open clusters. Our isochrone fit by no means can be considered final. It also carries a burden of correlations between the distance, age, metallicity, and reddening. Among the parameters derived in this study, the reddening estimate is clearly the least satisfactory. We provide potential clues possibly biasing our reddening estimate, which should help in planning the future photometric observations. The upper part of CMD requires a rigorous clean-up by the means of radial velocities to better delineate the upper main sequence, sub-giant, and giant branches. Although we are confident that the metallicity of Tr~20 is slightly sub-solar, it is desirable to obtain high-resolution spectra for the clump stars and redo the analysis and try to eliminate the apparent discrepancy between the photometric and spectroscopic $T_{\\rm eff}$. It is not clear what causes a considerable width of the main sequence, reaching $\\sim$0.3 mag. Is it differential reddening, large population of binaries, or a combination of both?  The observed spread of the red clump star colors indicates that differential reddening in Tr~20 may not exceed $\\sim$0.1 mag in $B-V$. There is a hint of the blue straggler population which photometrically is impossible to disantangle from numerous field stars. We hope that the future studies will enable refining the cluster parameters and possibly open up new venues of research of this rich open cluster."
    },
    "0809/0809.4172_arXiv.txt": {
        "abstract": "Infrared-Faint Radio Sources (IFRSs) are a class of source which are bright at radio frequencies, but do not appear in deep infrared images.  We report the detection of 14 IFRSs within the {\\it Spitzer} extragalactic First Look Survey field, eight of which are detected near to the limiting magnitude of a deep $R$-band image of the region, at $R\\sim24.5$.  Sensitive {\\it Spitzer Space Telescope} images are stacked in order to place upper limits on their mid-infrared flux densities, and using recent 610-MHz and 1.4-GHz observations we find that they have spectral indices which vary between $\\alpha=0.05$ and 1.38, where we define $\\alpha$ such that $S_{\\nu} = S_{0}\\nu^{-\\alpha}$, and should not be thought of as a single source population.  We place constraints on the luminosity and linear size of these sources, and through comparison with well-studied local objects in the 3CRR catalogue demonstrate that they can be modelled as being compact ($<20$~kpc) Fanaroff-Riley Class II (FRII) radio galaxies located at high redshift ($z\\sim4$). ",
        "introduction": "Infrared-Faint Radio Sources (IFRSs) are a class of radio-bright or infrared-faint objects identified in the Australia Telescope Large Area Survey \\citep*[ATLAS;][]{Norris06}.  The ATLAS survey consists of deep 1.4-GHz radio observations over two of the southern {\\it Spitzer} Wide-area Infrared Extragalactic Survey \\citep[SWIRE;][]{Lonsdale03} fields, which have deep infrared data between 3.6 and 160~$\\mu$m publicly available.  The {\\it Spitzer Space Telescope} \\citep{Werner04} infrared observations of the SWIRE fields are sensitive enough that it was expected that all radio sources detected in the ATLAS survey which were in the local Universe would be visible in the SWIRE images \\citep{Norris06}, whether the radio emission was due to star-formation or an Active Galactic Nuclei (AGN).  However, it was found that some radio sources did not have an infrared counterpart at any of the {\\it Spitzer} wavelengths.  A total of 55 IFRSs of varying flux densities have been detected within the {\\it Chandra} Deep Field South \\citep[CDFS;][]{Norris06} and European Large-Area {\\it ISO} Survey South-1 \\citep[ELAIS-S1;][]{Middelberg08} fields, with several having 1.4-GHz flux densities of greater than 5~mJy, and the brightest having a 1.4-GHz flux density of 26.1~mJy \\citep{Norris06}.  The possibility that these sources are obscured star-forming galaxies with infrared flux densities falling just below the sensitivity limit of {\\it Spitzer} would seem to be ruled out by the production of stacked infrared images centred on the radio source positions of 22 IFRSs \\citep{Norris06} which show no evidence for any infrared counterparts.  \\citet{Norris07} made observations of two of their bright IFRSs using Very Long Baseline Interferometry (VLBI), and found that one of them had a compact core with angular size less than 0.03~arcsec (a linear size below 260~pc at any reasonable redshift), although they were unable to detect the second target.  They concluded that the radio emission is driven by AGN activity rather than star-formation, making IFRSs either high redshift radio-loud quasars, or abnormally highly obscured radio galaxies at more moderate redshifts.  \\citet{Norris07} considered the spectral energy distribution of the radio-loud quasar 3C273, and showed that placing it at low redshift could not reproduce both the detection at 1.4~GHz and non-detection at 3.6~$\\mu$m for the source with the compact core, but moving 3C273 to $z=7$ would fulfil the requirements.  However, the lack of information at other wavelengths makes it difficult to determine exactly what type of source the IFRSs are. These sources are selected through being radio-bright, making the use of radio diagnostics important for determining their type.  The first of the large-scale surveys carried out by the {\\it Spitzer Space Telescope} was the {\\it Spitzer} extragalactic First Look Survey (xFLS).  Four square degrees, centred on $17^{\\rm h}18^{\\rm m}00^{\\rm s}$, $+59\\degr30'00''$ (J2000 coordinates, which are used throughout this work), were observed with two instruments -- the Infrared Array Camera \\citep[IRAC;][]{Fazio04} at 3.6, 4.5, 5.8 and 8~$\\mu$m, and the Multiband Imaging Photometer for {\\it Spitzer} \\citep[MIPS;][]{Rieke04} at 24, 70 and 160~$\\mu$m.  This data is comparable in depth to the SWIRE survey, with catalogue completeness limits ($5\\sigma$, apart from the 3.6-$\\mu$m band which has its completeness limit set at $7\\sigma$) of 20, 25, 100 and 100~$\\mu$Jy in the four IRAC bands.  There are two deep radio surveys of the region, at 1.4~GHz with the Very Large Array \\citep[VLA;][]{Condon03} and at 610~MHz with the Giant Metrewave Radio Telescope \\citep[GMRT;][hereafter Paper I]{Garn07}.  These have comparable resolution (5~arcsec$^{2}$ at 1.4~GHz; $5.8\\times4.7$~arcsec$^{2}$ at 610~MHz) and sensitivity (23~$\\mu$Jy~beam$^{-1}$ at 1.4~GHz; 30~$\\mu$Jy~beam$^{-1}$ at 610~MHz), and cover approximately the same region of sky as the infrared observations.  An $R$-band survey of the xFLS field has been carried out \\citep{Fadda04} with a limiting Vega magnitude of $R$ = 25.5, and estimated 50~per~cent completeness limit of $R$ = 24.5.  The deep and complementary infrared, optical and radio observations make this region a good area to search for other IFRSs, and the availability of multi-frequency radio data will allow the spectral index $\\alpha$\\footnote{where we define the radio spectral index $\\alpha$ such that $S_{\\nu} = S_{0} \\nu^{-\\alpha}$} of any IFRS to be obtained, providing a much greater understanding of the source properties.  In section~\\ref{sec:selection} we describe the selection criteria for forming a sample of Infrared-Faint Radio Sources.  We consider the available radio and infrared properties of these sources in section~\\ref{sec:properties}, create stacked infrared images of the sources in order to test whether the sources lie slightly below the {\\it Spitzer} detection threshold, and place upper limits on their infrared flux between 3.6 and 160~$\\mu$m.  We calculate the radio spectral index of each source, and demonstrate that IFRSs are made up of flat, steep and ultra-steep spectrum sources.  In section~\\ref{sec:modelling} we place limits on the luminosity and linear size of these sources, and compare well-studied sources from the Revised Revised Third Cambridge (3CRR) catalogue of \\citet*{Laing83} to the observed IFRS population.  We demonstrate that it is possible to model the flat, steep and ultra-steep-spectrum IFRSs with the spectral energy distribution (SED) of a local radio galaxy, placed at high redshift, and conclude that it is possible to explain the IFRS population without requiring new, exotic source types.  A flat cosmology with the best-fitting parameters from the Wilkinson Microwave Anisotropy Probe (WMAP) five-year data of $\\Omega_{\\Lambda} = 0.74$ and H$_{0}$ = 72~km~s$^{-1}$~Mpc$^{-1}$ \\citep{Dunkley08} is assumed throughout this work. ",
        "conclusions": "We confirm the existence of a class of bright radio sources which are undetected in deep infrared images.  These sources are rare -- we find a total of 14 sources with $S_{1.4} > 0.5$~mJy and no clear infrared counterparts.  Ten of these have a flux density of greater than 1~mJy, a source density of roughly 2.5~deg$^{-2}$, and comparable to the source density of bright IFRSs found by \\citet{Norris06} and \\citet{Middelberg08}.  There are eight sources with $S_{1.4}>0.5$~mJy that show no evidence for possible infrared counterparts in the deep IRAC images of the {\\it Spitzer} extragalactic First Look Survey field.  We use source stacking to place upper limits on the mid-infrared flux densities of these sources, and use either detections or upper limits in the $R$-band to further constrain the SED of each source.  By using multi-frequency radio data, we demonstrate that the sources vary in type between flat-spectrum quasars with $\\alpha\\sim0.1$, and USS radio sources with $\\alpha\\sim1.4$, and should not be thought of as a single source population.  The possibility that the Infrared-Faint Radio Source population is made up of high-redshift luminous radio galaxies appears increasingly likely -- the lack of detection in infrared images would seem to rule out highly obscured star-forming galaxies, while the existence of flat-spectrum sources implies that at least some of the sources have an AGN origin.  We place upper limits on their overall linear size through the knowledge of their maximum deconvolved angular size from the \\citet{Condon03} 1.4-GHz catalogue, and show that they must be $<20$~kpc at any redshift, making them smaller than the majority of bright radio sources in the local Universe, as given by the 3CRR catalogue.  By looking at the locations of these sources in a P-D diagram, for a range of redshifts, we identify bright local radio sources with similar characteristics, and demonstrate that all three classes of IFRS can be modelled as less-luminous versions of well-studied Fanaroff-Riley Type II radio galaxies placed at high redshift."
    },
    "0809/0809.3752_arXiv.txt": {
        "abstract": "Young massive stars in the central parsec of our Galaxy are best explained by star formation within at least one, and possibly two, massive self-gravitating gaseous discs. With help of numerical simulations, we here consider whether the observed population of young stars could have originated from a large angle collision of two massive gaseous clouds at $R \\simeq 1$ pc from \\sgra. In all the simulations performed, the post-collision gas flow forms an inner, nearly circular gaseous disc and one or two eccentric outer filaments, consistent with the observations. Furthermore, the radial stellar mass distribution is always very steep, $\\Sigma_* \\propto R^{-2}$, again consistent with the observations. All of our simulations produce discs that are warped by between 30$^\\circ$ to 60$^\\circ$, in accordance with the most recent observations. The 3D velocity structure of the stellar distribution is sensitive to initial conditions (e.g., the impact parameter of the clouds) and gas cooling details. For example, the runs in which the inner disc is fed intermittently with material possessing fluctuating angular momentum result in multiple stellar discs with different orbital orientations, contradicting the observed data.  In all the cases the amount of gas accreted by our inner boundary condition is large, enough to allow \\sgra\\ to radiate near its Eddington limit over $\\sim 10^5$ years. This suggests that a refined model would have physically larger clouds (or a cloud and a disc such as the circumnuclear disc) colliding at a distance of a few parsecs rather than 1 parsec as in our simulations. ",
        "introduction": "The power output of the central parsec of our Galaxy is dominated by young (aged $6 \\pm 1$ million years) massive ``He-I'' stars \\citep{KrabbeEtal95}. Many of these stars are located in a well-defined thin disc structure that rotates clockwise as seen on the sky \\citep{Levin03,GenzelEtal03,PaumardEtal06}. The remaining stars can be classified as a second more diffuse disc rotating counter clockwise \\citep{GenzelEtal03,PaumardEtal06} although the statistical significance of this feature is disputed \\citep{LuEtal06}. Stars appear to be on more eccentric orbits in the second disc \\citep{PaumardEtal06,LuEtal06}. The bright He-I stars seem to be excluded from the central arcsecond ($1''\\approx 0.04$ pc when viewed from the $\\approx 8$ kpc distance to the Galactic Centre (GC)), which instead contains at least a dozen less massive but still quite young B-type stars \\citep{GhezEtal03,GhezEtal05}. These stars (commonly referred to as S-stars in the Galactic Centre community) appear to be on rather eccentric orbits distributed without any notable preference to a planar configuration \\citep[e.g.,][]{EisenhauerEtal05,LuEtal06}. At present it is not clear if the S-stars are physically related to the more massive ``disc''-stars, as the initial orbits of the S-stars could have been very different from their present day ones due to resonant relaxation processes \\citep{HopmanAlexander06}. Thus, in this paper we shall concentrate on the puzzle of the young massive stars found outside the inner arcsecond as these present us with a cleaner laboratory environment.  \\cite{Gerhard01} suggested that the stars may have been formed at a distance of tens of parsecs in a massive star cluster. The cluster's orbit would then decay through dynamical friction with the background stars, and eventually the stars would form a disc populated by stars on eccentric orbits \\citep[e.g.,][]{Hansen03,Kim04,Gurkan05,LevinEtal05}. However, infrared observations find no young massive stars outside the central $\\sim 0.5$ pc \\citep{PaumardEtal06} which should be present in this model. In addition, Chandra X-ray limits \\citep{NS05} constrain the number of X-ray emitting Young Stellar Objects to well below what is expected in the cluster disruption scenario.  An alternative model is the in-situ model in which stars form inside a self-gravitating gaseous disc \\citep[e.g.,][]{Levin03,NC05,PaumardEtal06}. In fact, theorists had expected gaseous massive discs around supermassive black holes to form stars or planets \\citep[e.g.,][]{Paczynski78,Kolykhalov80,Shlosman89,Collin99,Gammie01,Goodman03} long before the properties of the young stars in the GC became known. Recently, \\cite{NayakshinEtal07} have numerically simulated the fragmentation process of a geometrically thin gaseous disc of mass $\\simgt 10^4 \\msun$ for conditions appropriate for our GC (albeit with a rather simple cooling prescription) and found a top-heavy mass function for the stars formed there. \\cite{AlexanderEtAl08} extended numerical studies to the fragmentation of eccentric accretion disks.  However, the in-situ model has not so far addressed in detail the origin of the gaseous discs themselves, although similar discs are believed to exist in AGN and quasars. Indeed, sub-parsec massive gaseous discs are invoked as means of feeding supermassive black holes' immense appetite for gaseous fuel \\citep[e.g.][]{Frank02}. Such discs are actually observed as the sites of maser emission in some of the nearby galactic centres \\citep{Miyoshi95}. But \\sgra\\ is currently not a member of the AGN club, let alone the even more powerful and elitist quasar community. Indeed, \\sgra\\ bolometric luminosity is around $\\sim 10^{-9}$ of its Eddington limit \\citep[e.g.,][]{Narayan02,Baganoff03a}. The amount of ionised gas currently present in the inner parsec is estimated at perhaps a mere hundred solar masses \\citep{Paumard04}. There are also tight observational constraints \\citep{Falcke97,Cuadra04} on the presence of an optically thin disc on smaller scales (say 0.01 pc or less).  We believe that the requisite gaseous discs could not have been assembled by viscous transport of angular momentum as is expected to be the case in other systems \\citep{Shakura73,Frank02}.  The viscous time scale of a thin, marginally self-gravitating disc can be estimated as $t_{\\rm visc} = (\\mbh/M_{\\rm disc})^2 (\\alpha \\Omega)^{-1}$, where $\\mbh \\approx 3.5 \\times 10^6\\msun$ and $M_{\\rm disc}$ are masses of the blackhole and the disc, respectively, and $\\Omega$ is the Keplerian angular frequency. For $\\alpha=0.1$ and $M_{\\rm disc} \\approx 10^4 \\msun$ (reasons for this choice of mass are discussed in \\cite{NayakshinEtal06}), $t_{\\rm visc} \\sim 3 \\times 10^7$ years at $r = 0.03$ parsec, the inner edge of the discs \\citep{PaumardEtal06}, and $10^9$ years at 0.3 parsec. These times are very long compared with the age of the stellar systems. In fact, the gaseous discs would have to evolve even faster as the two stellar systems are co-eval within one million years \\citep{PaumardEtal06}.  It is also not clear that one should expect any sort of a viscous quasi-steady state for the gas in our Galactic Centre. There is no strong evidence for other similar star formation events in the central parsec within the last $\\sim 10^8$ years. Therefore, the one-off star formation event appears to be best explained by a one-off deposition of gas within the central parsec. There are several ways in which this could have happened, e.g. a Giant Molecular Cloud (GMC) with a sub-parsec impact parameter (in relation to \\sgra) could have self-collided and become partially bound to the central parsec \\citep[e.g.,][]{NC05}. Alternatively, a GMC could have struck the Circumnuclear Disc (CND) located a few parsecs away from \\sgra, and then created gas streams that settled into the central parsec.  It is very hard to estimate the probability of a cloud-cloud collision in the central part of the Galaxy. \\cite{Hasegawa94} estimated the rate of GMC collisions as one per 20 Myrs in the central 100 pc, but the estimate is uncertain by at least an order of magnitude in either direction due to the lack of knowledge about the size distribution of GMCs. However, there is observational evidence that GMCs can be put on orbits significantly different from the simple circular orbits seen in the inner Galaxy. \\cite{Stolte08} recently measured the proper motions of the Arches star cluster. This star cluster is about 2.5 million years old and has a mass in excess of $10^4 \\msun$ \\citep[e.g.,][]{Figer04}. It is presently about 30 pc away from \\sgra\\ in projection. Its orbit must be strongly non-circular as its 3D velocity is over 200 km/sec in the region where the circular velocity is only $\\sim 110$ km/sec. Therefore, this suggests that massive GMCs can be on highly non-circular orbits. It does not seem implausible that one of these clouds would pass within a few parsec of \\sgra.  In this paper we explore such a one-off collision event in a very simple setup (\\S \\ref{sec:ic}). We allow two massive, uniform and spherical clouds on significantly different orbits to collide with each other one parsec away from \\sgra. The resulting gas dynamics, in particular the way in which gas settles into the inner parsec, is the focus of our effort here. We find that the collision forms streams of gas with varying angular momentum, both in magnitude and direction. Parts of these streams collide and coalesce to form a disc, and remaining parts form one or more orbiting filaments. As the gas cools, it becomes self-gravitating and stars are born, usually in both the disc and the filaments. This overall picture is discussed in \\S \\ref{sec:overall}. A reader mainly interested in a comparison between the simulations and observations may find sections \\ref{sec:radial}, \\ref{sec:accretion} and \\ref{sec:discussion} most relevant\\footnote{For movies of the simulations the reader is directed to \\url{http://www.astro.le.ac.uk/~aph11/movies.html}}. In \\S \\ref{sec:obs} we argue that our simulation results suggest that the location of the original impact was somewhat further away from \\sgra\\ than we assumed here, e.g., perhaps a few parsecs. ",
        "conclusions": "In this paper we presented several simulations of cloud-cloud collisions aimed at reproducing gas flows that could have formed (a) gaseous disc(s) in the central parsec of our Galaxy, as well as the resulting star formation. We found the gas cooling time and the impact parameter of the collision to influence the outcome significantly. Nevertheless, there are several robust results: (a) the inner near-circular and outer eccentric orbital structure of the stars formed there; (b) a sharply peaked mass distribution of stars $\\Sigma_*(R) \\sim 1/R^2$; (c) that the gaseous and stellar discs are warped. These results are in good accordance with the observations. The breakdown of the stellar system into one or more components is sensitive to the initial conditions and also the cooling parameter.  It appears that a GMC with a size of one to a few parsecs, self-colliding on a nearly radial orbit, or striking the CND at the distance of a few parsec from \\sgra\\ could explain the known observational data satisfactorily. Future observations and modelling might put interesting upper limits on the timescale over which the massive stars formed in the GC, and detail the structure and orbit of such a GMC."
    },
    "0809/0809.3939_arXiv.txt": {
        "abstract": "We present first results of a pilot project aimed at exploiting the potentiality of ground based adaptive optics imaging in the near infrared to determine the age of stellar clusters in the Galactic Bulge. We have used a combination of high resolution adaptive optics (ESO-VLT NAOS-CONICA) and wide-field (ESO-NTT-SOFI) photometry of the metal rich  globular cluster NGC 6440 located towards the inner Bulge, to compute a deep color magnitude diagram from the tip of the Red Giant Branch  down to $J\\sim 22$,  two magnitudes below the Main Sequence Turn Off (TO). The magnitude difference between the TO level and the red Horizontal Branch has been used as an age indicator. It is the first time that such a measurement for a bulge globular cluster has been obtained with a ground based telescope. From a direct comparison with  47 Tuc  and with a set of theoretical isochrones, we concluded that NGC 6440 is old and likely coeval to 47 Tuc. This result adds a new evidence that the Galactic Bulge is $\\approx$2 Gyr younger at most than the pristine, metal poor population of the Galactic Halo. ",
        "introduction": "Galactic globular clusters (GCs) are known since a long time to be the oldest stellar population of our Galaxy, hence their accurate age determination is crucial to settle the epoch of the Galaxy formation and a lower limit to the age of the Universe. A major effort has been devoted to obtain reliable ages for the Halo GC system \\citep{ros99,ros00,sw02,dea05}.  With the growing awareness that the Bulge GCs are a distinct sub-system, has come a new urgency to properly characterize their evolutionary sequences and to determine their ages and chemical composition. Moreover, Bulge GCs are {\\it  key templates} of simple stellar populations to study the stellar and chemical  evolution in the high metallicity domain \\citep{mcw97,mat99,wys00,bal07} and  for population synthesis of giant elliptical galaxies. Hence, a proper characterization of the main chemical and evolutionary properties of the Galactic Bulge and its comparison with the other components of the Galaxy is a fundamental step to unveiling the process of Galaxy formation and evolution \\citep{dut02,wys00}.  Despite its importance,  the Galactic Bulge GC system remained mostly unexplored, because of the huge foreground extinction, which severely prevents the usage of optical observations. Also, the limited  photometric and spectroscopic performances of the previous generations of infrared (IR) arrays, prevented  any deep, high resolution imaging and spectroscopy to properly survey the Bulge field and GC population.  Over the last decade a few optical and near IR ground-based photometric studies of bulge  GCs have been performed \\citep{ort96,gua98,fro95,dav00}. Recently, \\citet{val07} presented homogeneous and accurate IR color-magnitude diagrams (CMDs) of the evolved sequences of a sample of 24 GCs in the Bulge direction.  However, only a few attempts exist so far to estimate the age of  Bulge GCs from the direct measurement of the Main Sequence (MS) TO, by using the WFPC2 and NICMOS onboard HST \\citep{ort95,hea00,ort01,zoc01,coh02,fel02}.  The recent advent of adaptive optics (AO) capabilities on 8m-class telescopes allows unprecedent deep ground-based IR photometry down to the TO level in clusters as distant as $\\approx$10 kpc with potential better spatial sampling and coverage than NICMOS onboard HST.    We performed a pilot project aimed at measuring   the age of the Bulge GCs, by taking advantage of the ESO-VLT NAOS-CONICA (NACO)  imaging facility with AO and wavefront sensing in the IR, which is crucial  to find natural guide stars bright enough in such a  reddened environment. In this Letter, we present the first results of such a project, namely the age determination of the  GC NGC 6440.  This cluster has an iron abundance [Fe/H]=-0.56 \\citep{ori08} and a global metallicity [M/H]=-0.4  \\citep{val07} and it is located in the inner Bulge at (l,b)=(7.7,3.8) and 8.2 kpc from the Sun.  The high reddening \\citep[E(B-V)=1.15, ][]{val07} along the line of sight yields an extinction of $\\approx$3.5 mag in the visible, and makes this cluster an ideal target to be investigated in the near IR. ",
        "conclusions": "In this letter we have presented the age determination for NGC 6440: it is the first time that ground based AO imaging in the near infrared has been used to successfully  measure the TO region and derive the age of a Bulge GC. The few previous studies  have been all based on space observations. In fact, since the pioneering work by \\citet{ort95} who first measured the TO region of NGC~6553 and NGC~6528  by using HST-WFC2 and suggested a near-coeval age for the Galactic bulge and Halo,  a number of  subsequent works using HST and the WFPC2 and NICMOS imagers  \\citep{ort01,fel02} refined techniques and procedures to measure the age of these two globular clusters,  by also applying proper motion decontamination to the CMDs \\citep{zoc01,fel02}. Only one other cluster in the inner Bulge, namely Terzan 5 \\citep{ort01,coh02}  has been studied with HST-NICMOS, while \\citet{hea00} analyzed two other GCs observed with HST-WFPC2 in the outer Bulge, namely  NGC~6624 an NGC~6637. All these studies suggest that these clusters are old ($>$10 Gyr) and likely coeval with 47 Tuc.  While our and previous studies agree on the fact that Bulge GCs are likely coeval to 47 Tuc, absolute age estimates range between 11 and 14 Gyr, depending on the adopted theoretical isochrones and/or the absolute age for 47 Tuc itself. In this respect, a significant step forward in setting the absolute age of reference clusters has been performed by \\citet{gra03} who provided accurate distances and absolute ages  for three pillar GCs, namely NGC~6397, NGC~6752 and 47 Tuc, by using the MS Fitting Method and the Hipparcos distance scale.  Their age estimate for 47 Tuc turns out to be 11.2$\\pm$1 Gyr, $\\approx 2.6$ Gyr  younger than the other more metal poor Halo clusters.   A somewhat younger age (by 1-2 Gyr) was previously suggested by \\citet{ros99}, who also found slightly younger ages for other two metal rich clusters in their sample, namely NGC~6352 and NGC~6838, which  possibly belong to the thick disk population, like 47 Tuc \\citep[see also Figure 2 in][]{fer04}.  These first set of age determinations seem to suggest a common  epoch of formation of the tick disk and Bulge   and an overall GC formation process in our Galaxy starting $\\approx 13.5$ Gyr ago \\citep{gra03} with the formation of the  oldest Halo GCs and ending $\\approx 11$ Gyr ago with the formation of the metal rich objects in the thick disk and Bulge. In this scenario, the few younger GCs found in the Halo (like for example Pal~12, Arp~2, Terzan~8 etc.) could not have formed in situ. In fact, a number of recursive accretions of satellites and their GC systems  \\citep[as e.g. the Sagittarius Dwarf Spheroidal, ][]{iba94} could have significantly contributed to form the present Halo stellar populations \\citep{bel03,fer04}.  However, a larger sample of both Bulge and thick disk GCs urgently need to be observed with the purpose of determining their relative and absolute ages, before drawing firm conclusions. This is something that is well within the capabilities of the current generation of instruments, such as the refurbished HST and the improving performances of ground based AO imagers."
    },
    "0809/0809.1427_arXiv.txt": {
        "abstract": "Helioseismic techniques such as ring-diagram analysis have often  been used to determine the subsurface structural differences between solar active and quiet regions. Results obtained by inverting the frequency differences between the regions  are usually interpreted as the sound-speed differences between them. These in turn are used as a measure of temperature and magnetic-field strength  differences between the two regions. In this paper we first show that the ``sound-speed'' difference obtained from inversions is actually a combination of sound-speed difference and a magnetic component. Hence, the inversion result is not directly related to the thermal structure. Next, using solar models that include magnetic fields, we develop a formulation to use the inversion results to infer the differences in the magnetic and thermal structures between active and quiet regions. We then apply our technique to existing structure inversion results for different pairs of active and quiet regions. We find that the effect of magnetic fields is strongest in a shallow region above 0.985$R_\\odot$ and that the strengths of magnetic-field effects at the surface and in the deeper ($r < 0.98R_\\odot$) layers are inversely related, {\\it i.e.}, the stronger the surface magnetic field the smaller the magnetic effects in the deeper layers, and {\\it vice versa}. We also find that the magnetic effects in the deeper layers are the strongest in the quiet regions, consistent with the fact that these are basically regions with weakest magnetic fields at the surface. Because the quiet regions were selected to precede or follow their companion active regions, the results could have implications about the evolution of magnetic fields under active regions. ",
        "introduction": "\\label{sec:intro}  The availability of high-precision helioseismic data and the advancement of local helioseismology techniques have enabled investigations of fine details of solar internal dynamics for small regions of the Sun. Some of the methods, such as ring-diagram and time-distance analyses, have also enabled us examine details of the structure below solar active regions.  Ring-diagram analysis (Hill, 1988; Patron {\\it et al.}, 1997) is used to determine frequencies of  short-wavelength (high degree) modes in a small region of the Sun. These are modes that can be approximated as plane waves over a small area of the Sun. These frequencies can be inverted to determine the structure of the region under study using the same techniques that are used to invert global modes to study the spherically symmetric part of the solar structure. To avoid systematic errors, structure inversions are usually done using frequency differences between two regions to determine the difference in structure between the two regions ({\\it e.g.}, Basu, Antia and Bogart, 2004, 2007). Time-distance helioseismology (Duvall {\\it et al.}, 1993) measures the variations in the travel time of acoustic waves. The travel-time variations can then be used to infer the variations in the wave speed and flow velocity through an inversion procedure ({\\it e.g.}, Kosovichev, Duvall, and Scherrer, 2000; Kosovichev, Duvall, and Birch, 2001).  Basu, Antia, and Bogart (2004) inverted the frequency differences between several active regions and nearby quiet regions to determine the sound-speed difference and the difference in the adiabatic index $\\Gamma_1$ between these pairs of regions. They found that for all the active region -- quiet region pairs, the sound speed below the active regions was smaller than that of the quiet regions for about the first 7 Mm, and then the sound speed became larger in the active regions than in the quiet regions. A similar behavior was seen for the adiabatic index ($\\Gamma_1$). Qualitatively similar results were obtained by Kosovichev, Duvall, and Scherrer (2000) and Kosovichev, Duvall, and Birch (2001) using time-distance analysis. Kosovichev, Duvall, and Scherrer (2000), Kosovichev, Duvall, and Birch (2001), and Basu, Antia, and Bogart (2004) interpreted the results in terms of a temperature difference, and alternatively as a difference in the magnetic field. However, by examining the $\\Gamma_1$ differences, Basu, Antia, and Bogart (2004) also concluded that the results cannot be explained as being caused by temperature changes or magnetic fields alone.   The presence of magnetic fields can affect the frequencies of waves in two ways: Firstly, the magnetic fields change the thermodynamic structure ({\\it i.e.}, pressure and temperature profiles) of the medium that the waves travel through, which, in turn, changes the frequencies of the waves. Secondly, the plasma waves are directly affected by the magnetic fields through the Lorentz force. In other words, the modification to the frequencies results from both structural and non-structural ({\\it i.e.}, Lorentz force) effects of the magnetic fields. Since the two effects are inseparable in the observed frequencies, the ``structures'' revealed by the inversions could partly be the manifestations of the frequency difference caused by Lorentz force on wave propagation. In addition, the inversion kernels to-date have been derived from non-magnetic reference models. Since the structural variables of a magnetic model have both magnetic and non-magnetic (gas) components, it had been uncertain if an inversion using a non-magnetic reference model would reveal the whole or only the gas part of the structural change. By using solar models that include magnetic fields and inversion kernels computed from non-magnetic reference models, Lin, Li, and Basu (2006) showed that the ``sound speed'' revealed by the inversions in the presence of magnetic fields is in fact a combination of both sound speed ($c_g \\equiv \\sqrt{\\Gamma_1 P_{\\rm gas}/\\rho}$) and Alfv{\\'e}n speed ($c_A \\equiv B/\\sqrt{4\\pi\\rho}$). To distinguish this property from the actual sound speed, we call it ``wave speed'', defined as $c_T \\equiv \\sqrt{\\Gamma_1 P_T/\\rho}$, where $P_T$ is the total pressure ($=P_{\\rm gas}+P_{\\rm mag}$, where $P_{\\rm mag} = B^2/8\\pi$, $B$ being the magnetic-field strength). While the speed of sound is directly related to temperature,  there is no simple, direct relationship between the ''wave speed'' $c_T$ and either temperature or magnetic fields. Hence, it has not been possible to determine the magnetic-field and temperature profiles below active regions with any degree of accuracy.   In this paper, we use solar models that include magnetic fields to first confirm the Lin, Li, and Basu (2006) results and then derive a practical way of using the wave speed and $\\Gamma_1$ inversion results to determine the thermal and magnetic structural differences between active and quiet regions. After confirming the reliability of this method to infer the thermal and magnetic structures, we then apply it to the inversion results of Basu, Antia, and Bogart (2004).  The main limitation of our models is that they are one dimensional, and hence the form of magnetic fields allowed is quite restrictive. We cannot, for example, have toroidal magnetic fields or magnetic fields over restricted ranges of latitudes or longitudes. However, we are forced to use such models because the codes to construct self-consistent solar models with arbitrary magnetic fields do not yet exist. A consequence of using 1D models is that we are likely to overestimate the effect of magnetic fields on thermal transport. In the real, three-dimensional case, convection, when obstructed by a magnetic flux tube can bend around the tube, which is not possible in the 1D case.   The rest of the paper is organized as follows: We describe the magnetic models in Section~\\ref{sec:mdl}. A brief description of the inversion technique and the inversion results obtained from pairs of models are given in Section~\\ref{sec:inv}. The strategy to link the inversion results and temperature and magnetic fields is derived in Section~\\ref{sec:calib}. We apply the strategy to solar data in Section~\\ref{sec:sun}, and discuss the results in Section~\\ref{sec:interpretation}. The result of an attempt to model the subsurface magnetic structure is presented in Section~\\ref{sec:model}. We summarize our results in Section~\\ref{sec:summary}. It should be noted that although we apply our method to results from ring-diagram analyses, it can be applied to any helioseismic technique that can determine the adiabatic index ($\\Gamma_1$) below active regions. ",
        "conclusions": "\\label{sec:summary}  The aim of this investigation was to find a practical way to determine the thermal and magnetic structure under active regions from sound speed and $\\Gamma_1$ results obtained by inverting frequency differences between active and quiet regions. We have shown that ``sound speed'' results obtained by inverting frequency differences between active and quiet regions of the Sun represent the ``wave speed'' ($c_T \\equiv \\sqrt{\\Gamma_1 P_{T}/\\rho}$) and not the sound speed ($ c_g \\equiv \\sqrt{\\Gamma_1 P_{\\rm gas}/\\rho}$) or the Alfv\\'{e}n speed. Thus the results cannot be interpreted as being caused by differences in either temperature or magnetic fields alone. A combination of the two is needed to interpret the results correctly.  Using solar models with magnetic fields, we have determined a simple relationship between $\\delta\\Gamma_1/\\Gamma_1$, which is the relative difference of adiabatic index, and $\\delta \\beta$ ($\\beta \\equiv P_{\\rm mag}/P_{\\rm gas}$). $\\delta\\Gamma_1/\\Gamma_1$ was chosen for this work because its feature is well confined in the region of magnetic fields and because it can be obtained unambiguously from inversions. We have had to derive a separate relationship between the two quantities at each depth because the process of ionization, which changes $\\Gamma_1$, varies with depth. The $\\delta \\beta$ - $\\delta\\Gamma_1/\\Gamma_1$ relation can then used to infer the subsurface $\\delta \\beta$. The determined $\\delta \\beta$, along with inversion results of $\\delta c_T^2/c_T^2$ and $\\delta \\Gamma_1/\\Gamma_1$, can  be used to determine $\\delta c_g^2/c_g^2$ and $(\\delta T/T-\\delta\\mu/\\mu)$, provided that $\\delta P_{\\rm gas}/P_{\\rm gas}$ between two structures is negligible ({\\it i.e.}, $\\delta P_{\\rm gas}/P_{\\rm gas} \\ll 1.$). This method has been validated by using models. However, as the assumption $\\delta P_{\\rm gas}/P_{\\rm gas} \\ll 1.$ may not be suitable in the Sun, we only present in this paper $\\delta\\beta$ revealed by applying the method to the solar results obtained by Basu, Antia, and Bogart (2004).  The application of $\\delta \\beta$ -- $\\delta\\Gamma_1/\\Gamma_1$ relation to the solar inversion results shows that, as expected, different active regions behave quite differently. Nevertheless, for all regions studied, $\\delta\\beta$ is always positive immediately below the surface ( $ r > 0.985R_\\odot$) but negative at deeper regions, which implies that the inward decrease of $\\beta$ is faster in an AR with larger MAI than in one with smaller MAI.  The reason for the faster decrease of $\\beta$ under stronger ARs could be because the stronger magnetic fields near the surface suppress convection more, resulting in a larger build-up of gas ({\\it i.e.}, higher density and gas pressure), which leads to lower $\\beta \\equiv P_{\\rm mag}/P_{\\rm gas}$ in the deeper regions. It also appears that the so-called quiet regions have deep-seated magnetic effects.  We find that the near-surface $\\beta$ of AR$\\,9026$ and AR$\\,9393$, which have largest number of flares among all studied regions, is anomalously low. We propose two possible reasons to explain the phenomena. One is that $P_{\\rm gas}$ is larger in the flaring ARs because, as some theories have proposed,  flares might be caused by the larger amount of flux coming from the convection zone. Another possibility is that the relation between $\\delta \\beta$ and $\\delta \\Gamma_1/\\Gamma_1$ derived from our models may not be suitable for the shallow layers of such highly dynamic and eruptive ARs, where the variations in $\\Gamma_1$ can now come from many sources rather than solely from the magnetic fields as in our models. In either case, the inferred $\\delta \\beta$ may not reflect the actual value.  Our attempt to reproduce the active-region profiles of $\\delta \\Gamma_1/\\Gamma_1$ and $\\delta c_T^2/c_T^2$ suggests that the magnetic fields must be both strong and concentrated in order to produce the positive $\\delta c_T^2/c_T^2$ that is seen in the deeper layers. Additionally, the large, positive $\\delta \\Gamma_1/\\Gamma_1$ commonly seen between $0.97R_\\odot$ and $0.985R_\\odot$ can only be caused by a negative $\\delta \\beta$ in that region. The shifting of ionization zones does not produce matching magnitudes and locations of the positive $\\delta \\Gamma_1/\\Gamma_1$."
    },
    "0809/0809.2612_arXiv.txt": {
        "abstract": "The ESSENCE project was a six year supernova search carried out with the CTIO $\\;$ 4-m telescope.  We also obtained spectra with many of the world's largest ground-based telescopes and observed some of our SNe with the Hubble Space Telescope and the Spitzer Space Telescope.  We achieved our goal of discovering over 200 Type Ia SNe in the redshift range 0.2 to 0.8. With these data we determined the cosmic equation of state parameter to $\\pm$ 10 percent.  The data are consistent with a geometrically flat universe whose dark energy is equivalent to Einstein's cosmological constant. ",
        "introduction": "In the early- to mid-1990's astronomers expected that observations of standardizable candles such as Type Ia supernovae (SNe) would reveal to what extent the expansion of the universe was being decelerated by the gravitational attraction of all the matter in it.  Some astronomers predicted that $\\Omega _M \\equiv \\rho _0 / \\rho _{crit}$ = 1, the so-called Einstein-de Sitter universe.  From the dynamics of clusters of galaxies and studies of the large scale structure of the universe, others believed that $\\Omega _M \\approx$ 0.2-0.3.  Then \\cite[Riess \\etal\\ (1998)]{Riess_etal98} and \\cite[Perlmutter \\etal\\ (1999)]{Perlmutter_etal99} independently showed that Type Ia SNe at a redshift of $z \\sim$ 0.5 were systematically ``too faint'' and that their faintness could be attributed to an acceleration of the universe caused by a non-zero vacuum energy density.  We refer to this as dark energy.  \\cite[Einstein (1917)]{Einstein17} introduced the cosmological constant ($\\Lambda$) to characterize a universe that is neither expanding nor contracting.  (This was 12 years before Hubble's discovery of the expansion of the universe.)  Modern cosmologists describe the vacuum energy density by the dimensionless parameter $\\Omega _{\\Lambda} \\equiv \\Lambda c^2 / 3 (H_0)^2$, where H$_0$ is Hubble's constant. 1/$\\sqrt \\Lambda$ has units of length, the ``length scale over which the gravitational effects of a nonzero vacuum energy density would have an obvious and highly visible effect on the geometry of space and time'' (\\cite[Abbott 1988]{Abbott88}).  For H$_0$ = 72 km sec$^{-1}$ Mpc$^{-1}$ and $\\Omega _{\\Lambda}$ = 0.7, the length scale is 2900 Mpc, more than half the size of the observable universe.  Supernova data, combined with the observations of the baryon acoustic oscillations (\\cite[Eisenstein \\etal\\ 2005]{Eisenstein_etal05}) or combined with WMAP satellite observations of the cosmic microwave background radiation (\\cite[Komatsu \\etal\\ 2008]{Komatsu_etal08}), indicate that $\\Omega _M + \\Omega _{\\Lambda}$ = 1.  In other words, the geometry of the universe is flat.  The expansion of the universe was slowing down for the first $\\sim$7 billion years after the Big Bang.  Since then it has been accelerating.  In order to investigate the nature of dark energy, we wish to characterize its equation of state.  Let the equation of state parameter $w \\equiv P/(\\rho c^2)$, where $P$ is the pressure and $\\rho$ is the energy density.  $w$ = +1/3 for radiation.  $w = -1$ for Einstein's cosmological constant.  However, more exotic possibilities are possible.  $w > -1$ could indicate cosmic strings, and $w < -1$ could lead to the eventual instability of all particles in the universe (the so-called ``Big Rip'').  The luminosity distance, measured in Mpc, is given by  \\begin{displaymath} d_{lum} = \\frac{c(1+z)}{H_o}\\int_{0}^{z}\\frac{dz'} {\\sqrt{ \\Omega _M(1 + z')^3 + \\Omega _{\\Lambda}(1 + z')^{3(1+w)} } } \\;\\; . \\end{displaymath}  The distance modulus $\\mu \\equiv m-M$ = 5 log ($d_{lum}$) + 25.  A plot of the distance moduli of supernovae vs. the redshifts (or logarithms of the redshifts) is called the Hubble diagram.  Beyond a redshift of $z \\sim$ 0.2 the loci characterized by various combinations of $\\Omega _M$, $\\Omega _{\\Lambda}$, and $w$ begin to diverge from each other.  If we can obtain accurate distance moduli and spectroscopic redshifts of SNe or their host galaxies, we should in principle be able to determine these parameters.  It is also customary to pick one locus in the Hubble diagram as a reference and to plot the differences of the distance moduli with respect to that locus.  Determining the distance modulus of a given SN involves several steps: 1) imaging it in multiple filters; 2) transforming the photometry to observer rest frame apparent magnitudes; 3) determining the apparent magnitude(s) at maximum light;  4) correcting the apparent magnitudes for dust extinction along the line of sight; and 5) subtracting an estimate of the supernova's absolute magnitude at maximum for each filter.  Since the discovery of the brightness/decline rate relation by \\cite[Phillips (1993)]{Phillips93} and the Cepheid-based calibration, using the Hubble Space Telescope, of the distances of nearby galaxies which have hosted Type Ia SNe (\\cite[Suntzeff \\etal\\ 1999]{Suntzeff_etal99}, \\cite[Freedman \\etal\\ 2001]{Freedman_etal01}), we have come to a greater understanding of the uniformity of the light curves and absolute magnitudes of these objects.  It is important to note that without observations in at least two filters, no host galaxy extinction corrections can be made.  Supernovae observed in only one filter should not be used for the determination of cosmological parameters. Observations of a Type Ia SN in four or more rest frame filters, especially if one is a near-infrared filter, allow an accurate determination of the dust extinction suffered by the SN, even if the dust is very unusual (\\cite[Krisciunas \\etal\\ 2007]{Krisciunas_etal07}). ",
        "conclusions": ""
    },
    "0809/0809.5174_arXiv.txt": {
        "abstract": "We report on the catalog of Gamma--Ray Bursts (GRBs) detected with the Gamma Ray Burst Monitor aboard the \\sax\\ satellite. It includes 1082 GRBs with 40--700 keV fluences in the range from $1.3\\times 10^{-7}$ to $4.5\\times 10^{-4}$~erg~cm$^{-2}$, and 40--700 keV peak fluxes from $3.7\\times 10^{-8}$ to $7.0\\times 10^{-5}$~erg~cm$^{-2}$~s$^{-1}$. We report in the catalog some relevant parameters of each GRB and discuss the derived statistical properties. ",
        "introduction": "\\label{intro} In the last ten years a big step forward has been accomplished in the Gamma Ray Burst (GRB) astronomy. A key role in this advancement has been played by the \\sax\\ satellite \\citep{Boella97a}. Its capability of promptly and accurately localizing GRBs and the possibility of following them up with highly sensitive X--ray telescopes on-board allowed the discovery of the X--ray afterglow emission and, with ground telescopes, of the optical/radio afterglow from GRB sources \\citep{Costa97,Vanparadijs97,Metzger97,Frail97}. In this way the distance scale of long ($>$1~s) GRBs could be eventually established and important information on the GRB sources, their emission mechanisms and their environments could be derived.  The two \\sax\\ instruments that allowed the prompt and accurate GRB localization were the Gamma Ray Burst Monitor (GRBM) and the two Wide Field Cameras (WFC), the first with the role of recognizing the occurrence of a GRB and second with the role of localizing it in the case the event direction occurred in their field of view ($40^\\circ\\times 40^\\circ$ full width at zero response) and was above the instrument sensitivity in their operative energy range (2--28 keV)\\citep{Jager97}. Only a handful of GRBs (50, see Frontera 2004,\\nocite{Frontera04c}) was detected with both instruments, while the GRBM alone detected 1082 GRBs, a number similar to that of GRBs included in the third BATSE (3B) burst catalog~\\citep{Meegan96}. In this paper we present the catalog of these GRBs with their main observational and statistical properties. ",
        "conclusions": "The first complete catalog of the GRBs detected with the \\sax\\ GRBM includes 1082 GRBs, for most of which we have reported complete information (position, duration, peak flux, fluence and spectral hardness).  Concerning the GRB positions, we have reported the most accurate coordinates available for each GRBs included in the catalog. In the next future, for some of them, more accurate (or the first determination of) positions will be available from the Inter--Planetary Network (K. Hurley et al., in preparation), allowing a more accurate determination (or the first determination) of the fluence and peak flux values.  Some relevant results obtained with the BATSE catalogs (e.g., isotropical distribution of GRBs in the sky, deviation of the log$N$--log$S$ and log$N$--log$P$ cumulative distributions from that expected in the case of an homogeneous distribution of GRBs in an Euclidean space, two--peak distribution of the GRBs with time duration, higher hardness of the short GRBs) are confirmed.  For the first time, our catalog includes, in addition to the classical $T_{\\rm 90}$ GRB duration, also the total duration $T_{\\rm det}$ and the integrated time $T_{\\rm a}$ during which the GRB intensity exceeds the 2$\\sigma$ level. We find that the GRB active time has a cutoff at about 200 s, in spite that we find GRBs with $T_{\\rm det}$ up to 600 s. This result can have a relation with the physics of the unknown engine that gives rise to the GRBs. It seems that this engine has a maximum active time of 200 s. The integrated active is found consistent with a log-normal distribution for either short and long GRBs.  The origin of quiescent times is still unknown. The quiescent time properties have been investigated by various authors \\citep{Nakar02,Drago07}. We limit our investigation to the active time $T_{\\rm a}$ and quiescent time $T_q$, both integrated over the entire GRB duration. In the case of $T_{\\rm a}$, we find that its distribution is consistent with a log-normal distribution for short GRBs, while it is inconsistent with the latter in the case of long GRBs. In the case of the integrated $T_q$, we find that the cumulative distribution is  inconsistent with a log-normal distribution in the case of short and long GRBs. The consequences of these results would require further discussion, which is beyond the scope of this paper. We limit us to state that models of GRB engines and of their behavior with time should take into account these statistical results."
    },
    "0809/0809.2562_arXiv.txt": {
        "abstract": "We show that the spectrum of the unusual transient SCP\\,06F6 is consistent with emission from a cool, optically thick and carbon-rich atmosphere if the transient is located at a redshift of $z\\approx0.14$. The implied extragalactic nature of the transient rules out novae, shell flashes, and V838\\,Mon-like events as cause of the observed brightening. The distance to SCP\\,06F6 implies a peak magnitude of $M_I\\simeq-18$, in the regime of supernovae.  While the morphology of the light curve of SCP\\,06F6 around the peak in brightness resembles the slowly evolving Type~IIn supernovae SN\\,1994Y and SN\\,2006\\,gy its spectroscopic appearence differs from all previous observed supernovae. We further report the detection of an X-ray source co-incident with SCP\\,06F6 in a target of opportunity \\textit{XMM-Newton} observation made during the declining phase of the transient. The X-ray luminosity of $L_\\mathrm{X}\\simeq(5\\pm1)\\times10^{42}\\,\\mathrm{ergs\\,s^{-1}}$ is two orders of magnitude higher than observed to date from supernovae. If related to a supernova event, SCP\\,06F6 may define a new class.  An alternative, though less likely, scenario is the tidal disruption of a carbon-rich star. ",
        "introduction": "Studies of local and distant supernovae continue to reveal broad diversity in supernova properties, both spectral and photometric. This includes the recognition of SN with extremely high ejecta velocities (hyper-novae, e.g. \\citealt{galamaetal98-1, mazzalietal02-1, hjorthetal03-1}) and those which reach peak magnitudes markedly brighter than seen previously \\citep{smithetal07-2, quimbyetal07-1}. Indeed, even the ``standard candle\" SN\\,Ia include examples of unusually discrepant SN, e.g.  the extremely faint ($M_B=-15.9$) SN\\,2007ax \\citep{kasliwaletal08-1}. This variety most likely reflects similar variations in progenitor properties such as mass, metallicity and rotation.  Given the dramatically increased rate of SN detection over the past twenty years (there were 20 reported in 1987, 163 in 1997 and 572 in 2007\\footnote{http://cfa-www.harvard.edu/iau/lists/Supernovae.html}), it is perhaps unsurprsing that the classification system has required adaptation. Nonetheless the discovery of supernovae whose properties are broadly different from those seen before remains rare, and it is these examples which could perhaps place the strongest constraints on unusual processes in stellar evolution. In fact, from the theoretical perspective, it appears that the full variety of supernovae has not been discovered, e.g. \\citet{bildstenetal07-1} suggest the existence of faint ($M_V=-15$ to $-18$) thermonuclear supernovae from helium cataclysmic variables, or the tidal disruption and ignition of a white dwarf by intermediate mass black holes \\citep{rosswogetal08-1, rosswogetal09-1}.  It is clear that the growing number of large-area imaging surveys with high temporal cadence (hours to days), and the next generation of wide field space based instruments, will dramatically increase our knowledge about rare transient events. Projects that are currently operating, or coming online in the foreseeable, span a wide range of aperture sizes, e.g. ASAS \\citep{pojmanski97-1} or SuperWASP \\citep{pollaccoetal06-1} with limiting magnitudes $\\sim13-15$, the Catalina Real-Time Transient Survey (CRTS, \\citealt{drakeetal08-1}) or SkyMapper \\citep{kelleretal07-1} with limiting magnitudes of $\\sim19-21$, PanSTARRS \\citep{hodappetal04-1} and LSST \\citep{ivezicetal08-1} with limiting magnitudes of $\\sim24-25$, and ultimately SNAP, which may reach $\\sim 28$th magnitude \\citep{aldering05-1}.  Examples of very unusual events serendipitously discovered by deep supernova surveys are a red transient in M85 [possibly the the result of a stellar merger \\citep{kulkarnietal07-1} but see \\citealt{thompsonetal08-1} for an alternative], and the recent optical transient SCP\\,06F6 reported by \\citet{barbaryetal08-1}, which is so far of unknown nature.  Here, we suggest that SCP\\,06F6 had an extragalactic nature with a redshift of $z\\simeq 0.14$, and may represent a sofar unknown type of supernova or, less likely, a tidal disruption event. ",
        "conclusions": "In the two previous sections, we concluded that the morphology of the optical spectrum of SCP\\,06F6 at maximum brightness indicates an origin in a cool, optically thick and carbon-rich atmosphere at a redshift of $\\simeq0.14$, and that the optical transient was accompanied by a luminous X-ray transient. Here, we discuss the implications for the nature of SCP\\,06F6.  The symmetric light-curve of SCP\\,06F6 prompted \\citet{barbaryetal08-1} to consider micro-lensing as an explanation, although, as they pointed out, the large minimum amplification ($>120$) and the $\\sim 100$ day timescale are hard to understand on such a hypothesis. It can be added that the light-curve does not have the extended wings of microlensing light curves. If our identification of Swan bands is correct, then the microlensing hypothesis seems even more implausible. If microlensing was the cause, then the source would be a carbon star at $z\\simeq0.14$ (dismissing the unlikely combination of a new type of cool transient \\textit{and} microlensing).  The brightest carbon stars have absolute magnitudes around $-4$ \\citep{rebeirotetal93-1}, and so microlensing would have to provide at least 14 magnitudes or a factor of $400,000$ amplification. Roughly the right timescale and amplification result from a $5\\,\\mathrm{M}_\\odot$ object, placed half-way to the carbon star, passing within $4\\,\\mathrm{AU}$ of the line-of-sight to the carbon star at a transverse speed of $200\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$. However, the configuration is highly contrived, and requires the host galaxies of both the carbon star and lensing object to be unusually faint (see below). The problem of two faint galaxies can be avoided if the lensing object is located in either our galaxy or equivalently the host galaxy of the carbon star, but then the lensing object has to be very massive (several hundred solar masses). We therefore reject microlensing as playing any role in explaining SCP\\,06F6.  We suggest as a possible scenario that SCP\\,06F6 is related to a supernova-like event of a carbon-rich star.  The peak magnitude of SCP\\,06F6 is comparable to the peak in the luminosity function of core collapse supernovae. The implied light curve is shown in Fig.\\,\\ref{f-lc}, where it is compared to several other supernovae. As noted by \\citet{barbaryetal08-1} the transient is very slow in comparison to the majority of SN, which reach their peak on time scales of $\\sim 30$ days or less. However, some SN, such as the IIn SN\\,1994Y \\citep{hoetal01-1} and SN\\,2006gy \\citep{smithetal07-2}, do appear to have broad peaks, albeit with poor sampling of the rise. Overall the ensemble of SN light curves seem to broadly incorporate that of SCP\\,06F6.  If SCP\\,06F6 is related to supernovae, it will involve a rapid expansion of the envelope. For a peak magnitude $M_I=-18$, an assumed temperature of 5000\\,K, and a redshift of $z=0.14$, the radius of SCP\\,06F6 would be $\\sim3.5\\times10^{10}$\\,km. Assuming further that the expansion to this dimension occurred over $\\sim100$\\,d, an expansion velocity of $\\sim4000\\,\\mathrm{km\\,s^{-1}}$ is implied. Given the already broad nature of the C$_2$ Swan bands, the velocity gradient across the visible fraction of the envelope would result only in a relatively mild broadening, that would not distort the general appearance of the spectral features, but would just somewhat smooth out the sharp Swan band-heads seen in carbon stars. This is illustrated in the bottom panel of Fig.\\,\\ref{f-spectrum}, where we show the our template spectrum convolved with a  expansion velocity of $4000\\,\\mathrm{km\\,s^{-1}}$.  A puzzling aspect should SCP\\,06F6 lie at moderate redshift is the absence of an apparent host galaxy. The nearest detected object to the location of SCP\\,06F6 is a 6 sigma detection of an object with $z_{850}=25.8$, 1.5 arc-seconds from the transient position. At redshift $z=0.14$ this corresponds to an absolute magnitude of $M_Z \\sim -13.2$, which is extremely faint.  Only a small proportion of stars lie in such low mass galaxies, although examples of star forming galaxies with comparable absolute magnitude (e.g. IC\\,1613; \\citealt{dolphinetal01-1}) can be found. A number of supernovae and GRBs have occurred in faint ($M_R\\ga13$) galaxies \\citep[e.g.][]{levanetal05-1, fruchteretal06-1, drakeetal08-1}, indicating that such an association is not impossible.  The presence of C$_2$ Swan bands in the spectrum of the SCP\\,06F6 requires a carbon to oxygen ratio of C/O$>1$ by number, such that the carbon is not locked up in CO, favouring regions of low metallicity where it is easier for a small amount of carbon production to overwhelm the oxygen abundance. Given the well known relation between mass (or luminosity) and metallicity \\citep{tremonti04etal-1}, such carbon rich events may preferentially occur in faint host galaxies. Deep spectroscopy of the hypothetical nearby galaxy may enable a measure of its redshift, to test the association with SCP\\,06F6 if at $z\\simeq0.14$. The X-ray luminosity of SCP\\,06F6 inferred from the \\textit{XMM-Newton} observations, $L_\\mathrm{X}=(5\\pm1)\\times10^{42}\\,\\mathrm{ergs\\,s^{-1}}$, is much larger than expected from normal core collapse supernovae \\citep{unoetal02-1, kouveliotouetal04-1}, though the SN\\,IIn 1988Z and 2006jd reached $L_\\mathrm{X}\\simeq1-2.5\\times10^{41},\\mathrm{erg\\,s^{-1}}$ \\citep{fabian+terlevich96-1, immleretal07-1}. It is intriguing that both the slowly evolving light curve and large X-ray luminosity of SCP\\,06F6 bear similarities to the behaviours observed in SN\\,IIn, despite the radically different spectral appearance. A speculative scenario is that SCP\\,06F6 is associated with the death of a massive star that underwent mass loss removing the hydrogen envelope, e.g. a WC Wolf-Rayet star prior to the explosion, creating a carbon-rich circumstellar environment, or a Iben \\& Renzini type\\,1.5 supernova inside a metal-poor AGB star \\citep{zijlstra04-1}, which might possess a carbon-rich envelope at the point of core ignition. Interaction of the reverse shock with dense circumstellar matter would then also provide an explanation for the X-ray emission of SCP\\,06F6 \\citep{chevalier+fransson94-1}. We note in passing that \\citet{thompsonetal08-1} related the luminous transients SN\\,2008S and NGC\\,300 to either core-collapse supernovae or bright eruptions of massive dust-enshrouded stars, possibly carbon stars.  An alternative hypothesis of SCP\\,06F6 is that it is not due to stellar collapse but rather to the tidal disruption of a carbon rich star by a black hole. Such tidal disruption events may occur in the core of normal galaxies when stars approach the central black hole \\citep{rees88-1}; or further out in hosts via intermediate mass black holes \\citep[e.g.][]{rosswogetal09-1}. These events can potentially create the light curve shape (long duration flare) and approximate luminosity seen in SCP\\,06F6, for example object D3-3 in \\citet{gezarietal08-1} reaches a peak magnitude of $M_\\mathrm{g}\\sim -19$ and has an transient duration of several hundred days. The X-ray luminosity of SCP\\,06F6 is comparable to that of the candidate tidal disruption events identified in the \\textit{XMM-Newton} Slew Survey \\citep{esquejetal07-1}. This interpretation has the advantage that it more naturally explains the X-ray luminosity in tandem with that in the optical. Finally, tidal disruption events may lead to the ejection of a large fraction of the material of the disrupted star \\citep{ayaletal00-1}, explaining the presence of a large optically thick envelope.  However, this interpretation is not without problems. The lack of any obvious host galaxy to very low luminosities would imply either a very low black hole mass (if black holes do exist at the centres of dwarf irregulars) or that the black hole has somehow been ejected from its host (as has been suggested in a few cases e.g. \\citealt{magainetal05-1, haehneltetal06-1}). These possibilities, combined with the observation that the disrupted object be a carbon rich star, rather than a normal main sequence one appear to make the case for tidal disruption somewhat contrived. Nonetheless, with only one object, and thus an essentially unconstrained rate and space density for such events, it remains a possibility.  Any model of this source will need to explain both its unusual spectral appearance and the implied abundances as well as the high X-ray luminosity."
    },
    "0809/0809.4564_arXiv.txt": {
        "abstract": "We present here the first results of a spectropolarimetric analysis of a small sample ($\\sim20$) of active stars ranging from spectral type M0 to M8, which are either fully-convective or possess a very small radiative core. This study aims at providing new constraints on dynamo processes in fully-convective stars.  Results for stars with spectral types M0-M4 -- i.e. with masses above or just below the full convection threshold ($\\simeq 0.35~\\msun$) -- are presented. Tomographic imaging techniques allow us to reconstruct the surface magnetic topologies from the rotationally modulated time-series of circularly polarised profiles.  We find strong differences between partly and fully convective stars concerning magnetic field topology and characteristic scales, and differential rotation. Our results suggest that magnetic field generation in fully convective stars relies on different dynamo processes than those acting in the Sun and other partly convective stars, in agreement with theoretical expectations. ",
        "introduction": "In partly convective stars such as the Sun, magnetic fields -- the energy source of most activity phenomena -- are induced by plasma motions: the combined action of differential rotation ($\\Omega$ effect) and cyclonic convection ($\\alpha$ effect) manages to generate a self-sustained magnetic field. The so-called $\\alpha\\Omega$ dynamo \\cite{Parker55} is believed to operate mostly through the tachocline, a thin zone of strong shear located at the interface between the radiative inner zone and the convective envelope \\cite{Charbonneau05}.  Stars with masses lower than about $0.35~\\msun$ are fully-convective \\cite{Chabrier97} and thus do not possess a tachocline. However, they manage to trigger magnetic fields \\cite{Saar85,Johns96,Reiners06} and are very active \\cite{Delfosse98,Mohanty03,West04}. Though significant progress was made since first non-solar dynamo mechanisms were proposed \\cite{Durney93}, theoretical and numerical modelling require observational constraints. It is now acknowledged, from observational \\cite{Donati06,Morin08} and theoretical points of view \\cite{Chabrier06,Dobler06,Browning08}, that fully convective stars (FCS) manage to yield large scale magnetic fields. But the properties of such magnetic fields, and their dependency on stellar parameters (in particular mass and rotation rate) are not yet clear.  We present here the first results of a spectropolarimetric analysis of a small sample of active M dwarfs with spectral types ranging from M0 to M8, which are either fully convective or possess a very small radiative core. We aim at exploring the properties of the large-scale magnetic topologies of FCS, and their evolution with main stellar parameters. The stars were selected from the rotation-activity study \\cite{Delfosse98}. We chose only active stars so that the magnetic field is strong enough to produce detectable circularly polarised signatures, allowing us to apply tomographic imaging techniques. More details about the sample are available in Table~\\ref{tab:sample}.  \\begin{table} \\begin{tabular}{cccccccccccc} \\hline Name & ST & \\mstar & log$R_X$ & \\vsini & \\Prot & $\\tau_c$ & $Ro$ & \\rstar & $i$ \\\\ & & (\\msun) & & (\\kms) & (d) & (d) & ($10^{-2}$) & (\\rsun) & (\\degr) \\\\ \\hline GJ 182 & M0.5 & 0.75 & $-3.1$ & 10 & 4.35 & 25 & 17 & 0.82 & 60 \\\\ DT Vir & M0.5 & 0.59 & $-3.4$ & 11 & 2.85 & 31 & 9.2 & 0.53 & 60 \\\\ DS Leo & M0 & 0.58 & $-4.0$ & 2 & 14.0 & 32 & 44 & 0.52 & 60 \\\\ GJ 49 & M1.5 & 0.57 & $< -4.3$ & 1 & 18.6 & 33 & 56 & 0.51 & 45 \\\\ OT Ser & M1.5 & 0.55 & $-3.4$ & 6 & 3.40 & 35 & 9.7 & 0.49 & 45 \\\\ CE Boo & M2.5 & 0.48 & $-3.7$ & 1 & 14.7 & 42 & 35 & 0.43 & 45 \\\\ AD Leo & M3 & $0.42$ & -3.18 & $3.0$ & $2.24$ & 48 & 4.7 & 0.38 & 20\\\\ EQ Peg A & M3.5 & $0.39$ & -3.02 & $17.5$ & $1.06$ & 54 & 2.0 & 0.35 & 60 \\\\ EV Lac & M3.5 & $0.32$ & -3.33 & $4.0$ & $4.37$ & 64 & 6.8 & 0.30 & 60 \\\\ YZ CMi & M4.5 & $0.31$ & -3.09 & $5.0$ & $2.78$ & 66 & 4.2 & 0.29 & 60 \\\\ V374~Peg & M4 & $0.28$ & -3.20 & $36.5$ & $0.446$ & 72 & 0.62 & 0.28 & 70 \\\\ EQ Peg B & M4.5 & $0.25$ & -3.25 & $28.5$ & $0.404$ & 76 & 0.53 & 0.25 & 60 \\\\ \\hline \\end{tabular} \\caption{Fundamental parameters of the stellar sample. Columns 1--4 respectively list the name, the spectral type, the stellar mass, and the logarithmic relative X-ray luminosity log$R_X=$log($L_X/L_{bol}$). The projected rotation velocity and rotation period inferred from Zeeman Doppler Imaging (ZDI) are mentioned in columns 5 and 6. Columns 7--10 respectively list the empirical convective turnover time from \\cite{Kiraga07}, the effective Rossby number $Ro=\\frac{\\Prot}{\\tau_c}$, the theoretical radius suited to the stellar mass from \\cite{Baraffe98}, and the inclination angle used for ZDI. See \\cite{Morin08b,Donati08} for more details.} \\label{tab:sample} \\end{table}  For this study, we used the twin instruments ESPaDOnS on the 3.6-m Canada-France-Hawaii Telescope (CFHT) located in Hawaii and NARVAL on the 2-m T\\'elescope Bernard Lyot (TBL) in southern France. These spectropolarimeters, built on the same design, can produce Stokes I,Q,U and V spectra spanning the entire optical domain (from 370 to 1000~nm) at a resolving power of $\\sim 65000$ \\cite{Donati03}.  We performed monitoring observations of the sample in circularly polarised light (Stokes $V$). Least-squares deconvolution (LSD) \\cite{Donati97} was then applied, resulting, for each spectra, in a synthetic line profile gathering polarimetric information from most photospheric atomic lines. ",
        "conclusions": "Very different large-scale magnetic topologies are observed on each part of the $\\mstar\\simeq0.5~\\msun$ limit. We also note that dynamo processes become suddenly much more efficient at triggering large-scale magnetic fields (see Fig.~\\ref{fig:plotMP} and \\ref{fig:Ro_Bsq_RX}) at approximately the same mass ($\\simeq0.4~\\msun$). This strong step is not visible in the log$R_X$ vs $Ro$ plot. This result suggests that (i) the X-ray emission is sensitive to overall magnetic energy whereas we are only sensitive to the largest scales. (ii) At a given $Ro$ stars with mass above or below $\\sim 0.5~\\msun$ generate comparable magnetic energy, but with very different spatial scales repartition (the less massive stars triggering more magnetic energy in the largest scales).  These strong changes in magnetic field generation occur at masses slightly larger than the theoretical limit to full convection ($\\simeq 0.35~\\msun$). This may be due to strong shrinking of the radiative inner zone predicted by theoretical models, from $\\sim 0.5~\\rsun$ at $\\mstar=0.5~\\msun$ to nothing at $0.35~\\msun$ \\cite{Baraffe98,Siess00}.  We are currently completing this survey to extend it to fast rotators with $\\mstar>0.5~\\msun$, slow rotators with $0.2<\\mstar<0.5~\\msun$, two regimes not explored in the present sample. Present efforts are also directed to ultracool dwarfs ($\\mstar<0.2~\\msun$), to investigate how dynamo processes operate down to the brown dwarf limit, i.e. when stellar atmospheres start to become neutral.  \\begin{figure} \\label{fig:Ro_Bsq_RX} \\begin{tabular}{cc} \\includegraphics[width=.5\\textwidth]{Ro_Bsq} & \\includegraphics[width=.5\\textwidth]{Ro_RX} \\\\ \\end{tabular} \\caption{Left panel: Reconstructed magnetic energy as a function of $Ro$. Right panel: logarithmic relative X-ray luminosity as a function of $Ro$ (see Tab.~\\ref{tab:sample}).} \\end{figure}  \\begin{theacknowledgments} Julien Morin thanks the SOC of Cool Stars 15 and CNRS for providing financial support for attending the conference. \\end{theacknowledgments}"
    },
    "0809/0809.1494_arXiv.txt": {
        "abstract": "We  calculate the rate of photoevaporation of a circumstellar disk by energetic radiation (FUV, 6eV $<h\\nu<$13.6eV; EUV, 13.6eV $<h\\nu<$0.1keV; and Xrays, $h\\nu>0.1$keV)  from its central star. We focus on the effects of FUV and X-ray photons since EUV photoevaporation has been treated previously,  and consider central star masses in the range $0.3-7 {\\rm M}_{\\odot}$.    Contrary to the EUV photoevaporation scenario, which  creates a gap at about $r_g\\sim 7\\  (M_*/1{\\rm M}_{\\odot})$ AU and then erodes the outer disk from inside out, we find that FUV photoevaporation predominantly removes less bound gas from the outer disk.  Heating by FUV photons can cause significant erosion of the outer disk where most of the mass is typically located.  X-rays indirectly increase the mass loss rates (by a factor $\\sim 2$) by ionizing the gas, thereby reducing the positive charge on grains and PAHs and enhancing FUV-induced grain photoelectric heating. FUV and X-ray photons may create a gap in the disk at $\\sim 10$ AU under favourable circumstances.  Photoevaporation timescales for M$_* \\sim 1{\\rm M}_{\\odot}$ stars are estimated to be  $\\sim 10^6$ years, after the onset of disk irradiation by FUV and X-rays.  Disk lifetimes do not vary much for  stellar masses in the range $0.3-3$M$_{\\odot}$. More massive stars ($\\gtrsim 7 {\\rm M}_{\\odot}$)  lose their disks rapidly (in $\\sim 10^5$ years) due to their high EUV and FUV fields. Disk lifetimes are shorter for shallow surface density distributions and when the  dust opacity in the disk is reduced by processes such as grain growth or settling. The latter suggests that the photoevaporation process may accelerate as the dust disk evolves. ",
        "introduction": "Circumstellar disks are closely associated with star formation, with both strong observational evidence for their ubiquitous presence around young stars, and theoretical support as a natural consequence of  angular momentum conservation in  rotating, star-forming cores (e.g., Beckwith et al. 1990, Shu et al. 1994).  Disk  lifetimes are short compared to stellar lifetimes, and  dust disks around solar mass stars are estimated to typically last $\\lesssim 3$ Myrs  (e.g., Haisch et al. 2001).  Gas disk lifetimes are not as well constrained observationally ($\\lesssim 10^7$ yrs, Zuckerman et al. 1995, Pascucci et al. 2006) but are presumably similar to dust disk lifetimes.  These timescales set an upper limit on the time available for the assembly of planetary systems. Furthermore, dust emission studies that are commonly used to probe disk evolution indicate that disks transition from an optically thick stage  characterised by strong accretion and high ratios of total infrared to bolometric luminosities ($L_{IR}/L_{bol} \\sim 10^{-1}$, classical T Tauri star, CTTS) to  a weakly accreting phase with optically thin dust  ($L_{IR}/L_{bol} \\lesssim 10^{-4}$, weak-line T Tauri star, WTTS) (e.g., see recent review on dust disk evolution by Watson et al. 2007).  Based on the  statistics of these two  classes of objects, the transition epoch is estimated to be very brief, $\\sim 0.1 $Myr, and may involve optically thin inner disks with optically thick outer regions (Hillenbrand 2008, Cieza 2008).  Optically thick gas and dust  disks thus originate during  the star formation process,  undergo significant evolution over a few Myrs  to perhaps also form planets, and are rapidly destroyed during the planet formation process.   Various mechanisms have been proposed to date in order to explain the nature of disk dispersal (see Hollenbach et al. 2000 and Dullemond et al. 2007 for recent reviews).  Disk matter builds planetary systems. However, it has long been recognized from considerations of the total gas and solid content of the solar system and by comparing initial disk masses ($\\sim 0.1 {\\rm M}_{*}$) to the observed inventory of planet mass ($\\sim 10^{-3} {\\rm M}_*$) around extrasolar systems that much of the  mass is, in fact, physically removed from the disk either to spiral into the central star or to disperse back into the interstellar medium (e.g., Hayashi et al. 1985).  Viscous accretion is known to occur in the disk and allows transport of  mass onto the central star, but is insufficient as a dispersal mechanism in itself. Viscosity also spreads the disk, as a small amount of matter carries angular momentum outward, and viscous disks do not completely disappear at $\\sim$ 10 Myr, as observed (gas masses $\\lesssim 10^{-2}$ M$_J$, e.g., Zuckerman et al. 1995, Hollenbach et al. 2005, Pascucci et al. 2006; and dust masses $\\lesssim 10^{-4}$ M$_J$, e.g., Beckwith et al. 1990, Haisch et al. 2001). The slow decline in mass of a viscously evolving disk ($M_d \\propto t^{-1}$, e.g. Hartmann et al. 1998) is moreover inconsistent with the observationally inferred rapid transition from the CTTS to WTTS phases. One of the most promising mechanisms heretofore proposed for destroying disks has been photoevaporation, where  the  high energy radiation from a young star heats  gas at the disk surface and results in escaping mass flows  from the outer regions of disks.  In this scenario, viscous evolution causes the inner disk to accrete onto the central star and to spread,  while photoevaporation removes the outer disk mass reservoir.   A number of researchers have worked on the problem of disk photoevaporation in various contexts. Hollenbach et al. (1994, hereafter HJLS94),  Yorke \\& Welz (1996), and Richling \\& Yorke (1997)  examined the photoevaporation caused by the extreme ultraviolet (EUV; $h\\nu > 13.6$ eV) photons from the central star. Johnstone et al. (1998), St\\\"{o}rzer \\& Hollenbach (1999) and Richling \\& Yorke (1998, 2000) modeled photoevaporation of disks around low mass stars caused by the EUV and the far ultraviolet (FUV;  $ 6 < h \\nu < 13.6 $eV) fluxes from nearby massive stars.  Alexander et al. (2004) considered X-rays from the central star and found that they were unlikely to lead to significant photoevaporation rates (also see Ercolano et al. 2008).  Clarke et al. (2001) were the first to combine EUV-induced photoevaporation from a central solar mass star with turbulent viscosity to follow disk evolution. They found that EUV photoevaporation  produces a gap at a characteristic radius $\\sim 3-10$ AU once the viscous accretion through the disk has declined to a rate $\\dot{\\rm M}_{acc} \\lesssim 10^{-9}$ M$_{\\odot}$ yr$^{-1}$,  below the photevaporation rate from that region for typical EUV fields ($\\sim 10^{41-42} {\\rm s}^{-1}$). The timescale for producing the gap was estimated as $\\sim$  a few Myrs for sufficiently high EUV fields and small initial disk masses. At this point the inner disk rapidly accretes onto the central star on a viscous timescale ($\\sim 0.1$ Myr), followed by the  rapid photoevaporation of the outer disk by the unattenuated EUV flux from the central star (Alexander et al. 2005, 2006 a,b).   Matsuyama et al. (2003), Font et al. (2004) and Ruden (2004) have also modeled disk dispersal by a  combination of photoevaporation and viscous spreading and accretion.  Photoevaporation due to the irradiation of a disk by   FUV photons from a central low mass star has not been studied earlier.  FUV and X-ray luminosities of young low mass stars are significant and better determined than their EUV luminosities, which cannot be directly observed. In young low mass stars, FUV, EUV and X-ray fluxes are far in excess of older main sequence low mass stars like the Sun. This is due to both enhanced magnetic activity on stellar surfaces, which leads to much higher chromospheric activity and production of energetic photons, and to the presence of substantial accretion (in accretion columns, Gullbring et al. 1998, Calvet \\& Gullbring 1998) onto the central star from the disk, which leads to accretion shocks producing energetic photons.   Although FUV, EUV, and X-rays are produced near the stellar surface, they may be absorbed before reaching the disk surface. Alexander et al. (2005) have shown that the EUV produced in the accretion shock is not likely to penetrate the accretion column.  Disk winds  accompany accretion with a wind mass loss rate that scales as  $\\sim \\times 0.1$ the accretion rate. In order for the chromospheric EUV  to penetrate the disk wind,  the accretion rate needs to be less than $\\sim 10^{-8}$ M$_{\\odot}$ yr$^{-1}$ (see Hollenbach \\& Gorti 2008). FUV and X-ray photons require much larger column densities than the EUV to be absorbed and therefore begin to penetrate the protostellar wind and  photoevaporate the disk once accretion rates are $\\lesssim 10^{-6}$ M$_{\\odot}$ yr$^{-1}$ (Hollenbach \\& Gorti 2008),  long before the EUV photons from the stellar chromosphere can impact the disk surface.    In this paper, we include the effects of FUV and X-ray radiation from the central star on disk photoevaporation. This is a preliminary study  and our main goal is to assess the importance of FUV radiation on disk  evolution and lifetimes. We therefore make many necessary simplifications, notably the use of steady-state, static models of the chemical abundances, gas and dust temperatures, and density structure followed by a simple ``streamline'' analysis of the photoevaporative flow. We investigate the relative importance of optical, FUV, EUV and X-ray radiation from the central star in inducing photoevaporation by using our recently developed disk models (Gorti \\& Hollenbach 2008, hereafter GH08).  We present models of disks  around central star masses of $0.3-7$ M$_{\\odot}$ that span the range of current observational disk detections (e.g., Luhman et al. 2005, Fuente et al. 2006, Manoj et al. 2007).  We consider disks in their initial stages, with moderate grain growth and power law surface density profiles.  We  calculate photoevaporative mass loss rates from the disk as a function of  relevant model parameters  such as the dust opacity per H nucleus, the FUV luminosity, the X-ray luminosity and the power law exponent of the surface density distribution. Although our results are for a snapshot in time, we  qualitatively estimate the effects of viscosity and discuss disk evolution and dispersal.  The  paper is organized as follows.  In \\S 2 we present a brief overview of disk evolution prior to the onset of photoevaporation and describe our disk models and photoevaporation calculations.   In \\S 3 we discuss our results and  their implications for disk lifetimes, and describe disk evolution as a function of various input parameters.  We summarize and conclude in \\S4. ",
        "conclusions": "We  present  results of a study on the effects of  optical, FUV, EUV and X-ray photons from a star illuminating its own circumstellar disk and their role in dispersing the disk gas. Earlier work on photoevaporation focused on the ionizing EUV photons either from a nearby massive star or from the central star in an early energetic phase in stellar evolution. Here we mainly examine the role of FUV and X-ray photons which are produced even by low-mass stars during and after their accretion phase and their effectiveness in driving mass loss from the disk. We include optical, EUV, FUV and X-ray heating from the central star and find that FUV photoevaporation is most efficient in getting rid of bulk of the disk gas, located in the outer regions of the disk.  We find that FUV photoevaporation induces neutral flows in the disk that are most significant at the outer disk regions ($r\\gtrsim100$AU), where most of the disk mass resides.  Typical mass loss rates for a 1M$_{\\odot}$ star are $\\sim 3\\times 10^{-8} $ M$_{\\odot} {\\rm yr}^{-1}$ at $100-200$ AU, and FUV photons incident on the disk (of initial mass $0.03M_*$) can therefore deplete most of its mass rapidly, on timescales of the order $\\sim 10^6$ yrs.  X-rays do not cause significant flows, but can enhance FUV photoevaporation rates by ionizing the gas and increasing the efficacy of FUV grain photoelectric heating of gas.  EUV photoevaporation is effective in the inner ($r\\lesssim 3$AU for a 1M$_{\\odot}$ star) regions of the disk ($\\dot{M}_{pe} \\sim 10^{-10} $ M$_{\\odot} {\\rm yr}^{-1}$) where the stellar gravitational potential is high, and higher temperatures ($10^4$ K) are needed for escaping flows. EUV photoevaporation may also affect disk evolution at later stages by creating gaps in the disk (at $r \\sim 7$ AU, 1M$_{\\odot}$ case). FUV radiation acting with relatively high X-rays could potentially create gaps at $3-30$ AU at some stages of evolution, but this is not very conclusive from the static model results of this paper.  FUV photoevaporation depletes the disk mass early in its evolution, before EUV photons are able to penetrate the protostellar winds that accompany accretion. We suggest that disk dispersal is initiated by FUV  photons (aided by X-rays) which truncate the disk at the radius where viscous timescales equal photoevaporation timescales (typically at $r\\sim 100$ AU) and remove much of the disk mass.  However, time dependent models that follow the decrease in FUV luminosity as the accretion rate drops in time are needed.  As the disk mass drops,   viscous accretion rates decrease to values where EUV photoevaporation can create a gap. This occurs when the EUV-induced mass loss rate at $r_{g,II} \\sim 7 (M_*/({\\rm 1 M}_{\\odot})$ is higher than the accretion rate here.  The inner disk then drains  viscously and the direct illumination of the inner rim (at $\\sim 7 $AU) by EUV, FUV and X-rays then results in photoevaporation of the remaining outer disk gas. However, it is the FUV photoevaporation of the outer disk mass reservoir that sets the disk lifetime.  We find that FUV photoevaporation rates (and hence disk lifetimes) are sensitive to the FUV luminosity of the star. If the FUV luminosity is mainly generated by accretion activity, the mass loss rates decline with time as accretion diminishes in the disk.  FUV generated by stellar chromospheric activity is not very significant in dispersing the disk, and produces mass loss rates similar to that by stellar optical photons and dust heating of gas at $\\sim 200$ AU.  Disk photoevaporation timescales ($\\tau_{disk}$) are calculated at $\\sim 1$ Myrs after the penetration of the disk wind by FUV and X-rays (at an age $\\sim 1$Myr) and we estimate that disks  may last $\\sim2 $ Myrs.   We also find that the disk dust opacity may  influence the photoevaporation rates by affecting gas density at the flow origin and the degree of flaring.  For a factor of $\\sim 100$ reduction in dust opacity, we find that $\\tau_{disk}$ decreases by a factor of $\\sim 2$.   The surface density distribution of gas in the disk may also affect the photoevaporation rate, as steeper radial distributions decrease the radiation intercepted by the disk and reduce the mass loss rates; changing our surface density from $\\Sigma \\propto r^{-1}$ to $r^{-2}$ increases the disk lifetime by $\\sim$ 50\\%.  Disk lifetimes do not appear to depend on the mass of the central star for stellar masses between $0.3$ and 3M$_{\\odot}$ stars. Here, we derive lifetimes of $\\sim 10^6$ years. High mass stars ($M_* \\gtrsim 7$M$_{\\odot}$)  lose their disks rapidly in $\\sim 10^5$ years because of their very high photospheric FUV and EUV fields.  Our derived  disk lifetimes  of  $\\sim 10^6$ years after the FUV penetration of the protostellar wind for typical solar-mass stars are  consistent with observations, although this study does neglect dynamics and only crudely considers viscosity.  In addition, derived disk lifetimes may be longer if the initial disk mass is greater  than $0.03M_*$.  However,  our qualitative estimates of the ``mean'' disk lifetime by assuming ``mean'' or average stellar properties at a typical age of 1 Myr,  compare very well with observationally inferred disk lifetimes, without requiring very high EUV photon fluxes (Alexander et al. 2006) or special birthplaces near massive OB stars in a cluster environment (e.g., Hollenbach et al. 2000, Adams et al. 2004, Johnstone et al. 2004).  We offer a concluding important caveat.  The present calculations are simplistic and  only qualitatively include the effects of viscosity.  As the FUV field, accretion rates, and disk properties evolve with time, a proper treatment needs to consider all these effects  in a time-dependent model.  A future goal will be to incorporate the present FUV, EUV and X-ray photoevaporation models into a self-consistent and time-dependent solution that includes viscosity (Gorti, Dullemond \\& Hollenbach, in preparation)"
    },
    "0809/0809.1388_arXiv.txt": {
        "abstract": "We report a series of extensive photometric and spectroscopic observations of the luminous M31 nova M31N~2007-11d. Our photometric observations coupled with previous measurements show that the nova took at least four days to reach peak brightness at $R\\simeq14.9$ on 20~Nov~2007~UT. After reaching maximum, the time for the nova to decline 2 and 3 magnitudes from maximum light ($t_2$ and $t_3$) was $\\sim9.5$ and $\\sim13$~days, respectively, establishing that M31N~2007-11d was a moderately fast declining nova. During the nova's evolution a total of three spectra were obtained. The first spectrum was obtained one day after maximum light (5 days post-discovery), followed by two additional spectra taken on the decline at two and three weeks post-maximum. The initial spectrum reveals narrow Balmer and Fe~II emission with P~Cygni profiles superimposed on a blue continuum. These data along with the spectra obtained on the subsequent decline clearly establish that M31N~2007-11d belongs to the Fe~II spectroscopic class. The properties of M31N~2007-11d are discussed within the context of other luminous novae in M31, the Galaxy, and the LMC. Overall, M31N~2007-11d appears to be remarkably similar to Nova LMC~1991, which was another bright, slowly-rising, Fe~II nova. A comparison of the available data for luminous extragalactic novae suggest that the $\\grtsim4$ day rise to maximum light seen in M31N~2007-11d may not be unusual, and that the rise times of luminous Galactic novae, usually assumed to be $\\lessim2$ days, may have been underestimated. ",
        "introduction": "Classical Novae are a sub-class of cataclysmic variable systems where a Roche-lobe-filling star (typically a cool, near-main-sequence star) transfers mass to a white dwarf companion (Warner 1995, 2008). Eventually, a thermonuclear runaway (TNR) in the accreted material ensues, which drives substantial mass loss from the system, and leads to the nova eruption (e.g. Starrfield et al. 2008). Absolute magnitudes as bright as M$_V \\simeq -10$ have been observed at the peak of eruption. Their high luminosities and high rates of occurrence ($\\sim30$~yr$^{-1}$ in a galaxy like the Milky Way [Shafter 2002]), make novae powerful probes of the evolution of binary systems in different (extragalactic) stellar populations. The most throughly studied extragalactic system is M31, where more than 700 novae have been discovered since Hubble (1929) began his pioneering work in the early 20th century (e.g. see Darnley et al. 2006; Pietsch et al. 2007; Shafter 2008, and references therein).  In November of 2007, Nakano (2007) reported the discovery of a slowly-rising, and particularly luminous nova in M31. The nova, M31~2007-11d, was initially detected at $m=17.7$ (unfiltered) on 2007~Nov~16.51 before reaching $m=14.9$ on Nov~20.385. It was discovered independently by Quimby et al. (2007) as part of the ROTSE~IIIb patrol, and a finding chart based on these data showing the position of M31~2007-11d within M31 is shown in Figure~1. According to the compilation of M31 novae by Pietsch et al. (2007), only six (including M31~2007-11d) of the more than 700 novae recorded in M31 have been observed with $m<15$. Shortly after it became clear that M31~2007-11d was an unusually luminous nova, we initiated a series of photometric and spectroscopic observations to follow its subsequent evolution. Here, we report the results of this campaign. ",
        "conclusions": "Our principal conclusions can be summarized as follows:  1) M31N~2007-11d was one of the more luminous novae to be discovered in M31 over the past century. After a relatively slow rise to maximum light, the nova reached $R=14.9$ on 20~Nov~2007~UT.  2) Extensive photometry obtained on the subsequent decline from maximum light showed that the nova could be classified as a relatively ``fast\" nova, that took $\\sim9.5$ and $\\sim13$ days to by decline 2 and 3 magnitudes from maximum light, respectively.  3) Spectroscopic observations obtained the day following maximum light, and then again $\\sim2$ and $\\sim3$ weeks post maximum reveal relatively narrow (FWHM $\\sim2300$~km~s$^{-1}$) Balmer and Fe~II emission lines that clearly establish M31N~2007-11d as a member of the Fe~II class in William's (1992) system.  4) Among the most luminous Galactic novae with $M_V\\leq-9$, only one in four is of the Fe~II type despite the fact that $\\sim80$\\% of all Galactic novae are members of this class. Thus, M31N~2007-11d appears to be relatively rare in this respect.  5) The photometric and spectroscopic properties of M31N~2007-11d appear to be remarkably similar to that of another extragalactic nova, Nova LMC~1991 (Della Valle 1991, Schwarz et al. 2001). In particular, both were Fe~II novae that reached absolute magnitudes brighter than $M_V=-9$, and were characterized by relatively slow rises to maximum light.  6) Although there is considerable scatter at the upper end of the MMRD relation, we find no compelling evidence for a {\\it distinct} class of super-bright novae. The large scatter in luminosity among novae with $M_V\\lessim-9.0$ in Figure~6, coupled with the smooth drop-off in the luminosity distribution in Figure~5, is consistent with the uniform variation expected from a combination of observational uncertainties and intrinsic variability in fundamental nova parameters.  7) The rise time to maximum light for Galactic novae is poorly constrained, but has been generally regarded to be of order one to two days for the most luminous novae. A comparison with available data for luminous M31 and LMC novae such as M31N~2007-11d and Nova LMC~1991 shows rise times up to 4 days may not be unusual, and suggests that Galactic nova rise times may have been underestimated."
    },
    "0809/0809.4939_arXiv.txt": {
        "abstract": "We extend the results of a previous paper where a model of interacting dark energy, with a cosmological term decaying linearly with the Hubble parameter, is tested against the observed mass power spectrum. In spite of the agreement with observations of type Ia supernovas, baryonic acoustic oscillations and the cosmic microwave background, we had shown previously that no good concordance is achieved if we include the mass power spectrum. However, our analysis was based on the {\\it ad hoc} assumption that the interacting cosmological term is strictly homogeneous. Now we perform a more complete analysis, by perturbing such a term. Although our conclusions are still based on a particular, scale invariant choice of the primordial spectrum of dark energy perturbations, we show that a cosmological term decaying linearly with the Hubble parameter is indeed disfavored as compared to the standard model. ",
        "introduction": "The elucidation of the nature of dark energy is one of the most important challenges of modern cosmology, requiring at present the attention of theoreticians and observational teams. From the theoretical viewpoint, a crucial problem is to understand the role of vacuum in the cosmological scenario, its possible relation with dark energy and, in this case, why its observed density is so small as compared to the value theoretically expected by quantum field theories \\cite{Pad}.  Among the different approaches to this problem, one can find the suggestion that dark energy is the manifestation of quantum vacuum in the curved space-time, and that its density depends on the curvature, decaying from a huge initial value as space expands. As the total energy must be conserved, the vacuum decay is concomitant with matter production, a general feature of this kind of models and, more generally, of interacting dark energy models \\cite{Ozer,Horvat,SC,Grande,Holografia,Zimdahl,Ma,Fabris}.  However, it is difficult to derive the vacuum contribution in the expanding background, and some phenomenological approaches are needed in order to implement the above idea. A thermodynamical analysis in de Sitter space-time, in line with the holographic conjecture, has suggested a particular dependence of the vacuum density $\\Lambda$ on the Hubble parameter $H$, which is a good approximation in expanding, quasi-de Sitter space-times \\cite{SC2}.  Such a dependence was obtained by noting that a free particle in the de Sitter background presents, superposed to its normal modes, a thermal motion with a characteristic temperature equals to $H$, which is an expression of the known association of a temperature $H$ to the de Sitter horizon. On the other hand, if we regularize the vacuum energy density in the flat space-time by postulating a thermal distribution of the vacuum fluctuation modes at a temperature $m$ (which is equivalent to impose a superior cutoff $m$ on the modes frequencies), we obtain $\\Lambda_0 \\approx m^4$. Now, if in de Sitter space-time we shift this vacuum temperature from $m$ to $m+H$ we obtain, after subtracting the flat space-time contribution, the ansatz $\\Lambda \\approx (m+H)^4 - m^4$.  This is a phenomenological ansatz, in the sense that there is still no rigorous derivation of it in the realm of quantum field theories in curved space-time. Other interesting ansatzen have been proposed as well, also on phenomenological or semi-phenomenological basis (see, for example, \\cite{Grande,Holografia,Fabris}). In this context the comparison with observations is relevant, since it can rule out or constrain some models while we do not have a rigorous theoretical answer to the problem.  If we use for the cutoff $m$ the energy scale of the QCD phase transition (the latest known cosmological vacuum transition), in the limit of very early times we have $H>>m$, and the cosmological term is proportional to $H^4$. This leads to a non-singular inflationary solution, with an initial quasi-de Sitter phase giving origin to a radiation-dominated universe, with a matter content generated at the expenses of vacuum energy \\cite{SC2}.  For late times all depends on the masses of the produced particles. The time-energy uncertainty relation suggests that particles of mass $M$ can only be produced if $H > M$. Hence, massive particles as the baryonic matter, axions and supersymmetric candidates for dark matter should stop to be produced at the time of electro-weak phase transition (whereas the production of photons and massless neutrinos must be forbidden by some selection rule, otherwise our universe would be completely different at present). Therefore, if no other particle is taken into account, for late times we have a genuine cosmological constant, like in the standard $\\Lambda$CDM model. However, if we consider the possibility of very light dark particles (as massive gravitons, for example), the vacuum decay could still happen. Since in this limit we have $H<<m$, now the vacuum density varies linearly with $H$, that is, $\\Lambda \\approx m^3H$ \\cite{SC2}. On the other hand, as in the present universe the cosmological term is dominant, the Friedmann equation gives $\\Lambda \\approx H^2$. Hence, we have $H \\approx m^3$ and $\\Lambda \\approx m^6$. With $m$ of the order of the energy scale of the QCD chiral phase transition (approximately the pion mass), these expressions lead to numerical coincidences (in orders of magnitude) with the observed current values of $H$ and $\\Lambda$.  In this case, the resulting scenario is similar to the standard one, with the radiation phase followed by a long matter era, with the cosmological term dominating for large times \\cite{Borges}. A detailed analysis of the redshift-distance relation for type Ia supernovas, the baryonic acoustic oscillations (BAO) and the position of the first peak in the spectrum of anisotropies of the cosmic microwave background (CMB) has shown a very good concordance \\cite{Jailson2}, with present values for the Hubble parameter, the relative matter density and the universe age inside the limits imposed by other, non-cosmological observations.  Nevertheless, a late-time matter production may dilute the density contrast during the process of structure formation \\cite{Reuven}. Therefore, the study of the evolution of perturbations and the comparison of the model predictions with the observed mass power spectrum is also necessary. In a previous paper \\cite{I} we have shown that, in the case of a strictly homogeneous cosmological term (which means that matter production is homogeneous as well), the matter contrast is indeed suppressed at late times, after achieving a maximum near the present epoch. This would be a potential explanation for the cosmic coincidence, but it also leads to a suppression in the power spectrum. As a consequence, a good accordance with observations is only possible with a relative matter density above the concordance value obtained from the joint analysis of supernovas, BAO and CMB.  This would rule out the late-time vacuum decay in the context of the present model, but the homogeneity of the interacting cosmological term is an {\\it ad hoc} hypothesis, which should be relaxed before taking a definite conclusion. If matter is perturbed, the interacting dark energy must also be, and we have to verify whether such a perturbation is negligible or not. That is what we do in the present paper. By writing our ansatz for the variation of vacuum density in a covariant form, we derive a natural expression for the linear perturbations in the cosmological term, which can be integrated together with the relativistic equations for perturbations in matter and radiation. We then construct the predicted mass power spectrum, comparing with the observational data. Our conclusion is that, assuming a scale invariant primordial spectrum for the vacuum perturbations, a late-time vacuum decay linear with $H$ is clearly ruled out also in this case.  The paper is organized as follows. In the next section we briefly review the main features of the present model and its accordance with other kind of observations. In Section III we obtain and integrate the set of coupled perturbations equations for the vacuum, matter, radiation and metric, obtaining the corresponding power spectrum. In Section IV we outline our conclusions. ",
        "conclusions": "From a theoretical point of view, the hypothesis of vacuum decay is interesting in several aspects. First of all, it may naturally conciliate the observed cosmological term with a huge initial value for the vacuum density through a relaxing mechanism in the expanding space-time. Secondly, in the realm of our thermodynamic ansatz, we have a non-singular and inflationary early phase in the universe evolution, with the presently observed matter content generated by a primordial vacuum transition \\cite{SC2}.  A late-time vacuum decay depends on the mass of the produced particles. This possibility would be interesting as an explanation for the cosmic coincidence, since the suppression of the matter contrast owing to matter production begins to have importance when the cosmological term starts to dominate the cosmic expansion \\cite{I}. Until now it has also been survived to a precise joint analysis of supernovas, BAO and CMB observations, with a good concordance for the two free model parameters, i.e., the present values of the Hubble constant and of the relative matter density \\cite{Jailson2}.  In this paper we have extended a previous study of structure formation in this context \\cite{I}, investigating the consequences of matter production for the mass power spectrum. We had already shown that, in the case of a strictly homogeneous cosmological term, the suppression in the matter contrast leads to a disagreement with the observed spectrum, unless the matter density parameter is as high as $\\Omega_{m0} = 0.48$ (for which the accordance is excellent). This value is outside the current dynamical limits on $\\Omega_{m0}$ and above the concordance value obtained from the joint analysis of supernovas, BAO and CMB, $\\Omega_{m0} = 0.36$ \\cite{Jailson2}.  Now we have considered the possibility of perturbations in the cosmological term, a more natural assumption. Assuming that such a term presents the same scale-invariant primordial spectrum as dark matter, we conclude that the agreement with the observed spectrum, for any acceptable value of $\\Omega_{m0}$, is even worse. We have also considered the case in which the initial vacuum perturbations are zero, but the fitting with observations is still bad, unless for a very high matter density. This seems to rule out a late-time running of the vacuum term, at least in the recipe of the present model.  Of course, one could test other possible amplitudes and forms for the vacuum primordial spectrum (despite its unnaturalness). In fact, we have tested a large range of amplitudes, but the obtained power spectrums are not significantly better. We have also tested values for the present mass density of decoupled particles above the baryonic mass density, supposing the presence of a massive dark particle aside the light one produced by the vacuum decay. Also in this case the fitting between the observed and predicted power spectrums is poor.  Let us remind, however, that discarding a late-time variation of the cosmological term in the present model does not mean to discard the general idea of vacuum decay. If dark energy is a manifestation of quantum vacuum in the curved, expanding background, the inclusion of a decaying cosmological term in Einstein equations is as natural as the inclusion of a genuine cosmological constant. In the realm of the present model, the disagreement with large structure observations indicates that the cosmological term does not vary at late times. Nevertheless, a different ansatz for the vacuum variation may lead to a better accordance.\\footnote{Another possible scape for this situation would be supposing that our ansatz is not a good approximation for the present universe, because it is still not a quasi-de Sitter space-time. In this case, however, it would be surprising that a genuine cosmological constant fits so well the observations.}  Even if the vacuum decay is restricted to very early times - due to the masses of the produced particles - we have interesting consequences, as already discussed. A precise comparison of such a scenario with observations, in the context of our thermodynamic ansatz \\cite{SC2}, is still in order. Particularly, its necessary to find the primordial spectrum of matter perturbations generated in the inflationary phase of the model, to be tested against the observed spectrum of anisotropies in the CMB. This research is already in progress."
    },
    "0809/0809.1800_arXiv.txt": {
        "abstract": "{ This article summarizes the basic facts and ideas concerning the formation and evolution of cataclysmic variables (CVs). It is shown why the formation of CVs must involve huge losses of mass and orbital angular momentum, very likely via a common envelope evolution. A brief discussion of the principles of the long-term evolution of semi-detached binaries follows. Finally a brief sketch of CV evolution is given. ",
        "introduction": "Cataclysmic variables (CVs) are short-period semi-detached binary systems in which a white dwarf (WD) primary accretes matter from a low-mass companion star \\citep{Warner1995}. CVs are intrinsically variable and that on a wide range of time scales (from seconds to $\\gtrsim 10^6\\,{\\rm yr}$) and with a huge range of amplitudes (of up to $10^6$ and possibly even more). The rich phenomenology of CV variability which includes, among other things, phenomena like flickering, dwarf nova and classical nova outbursts, can to a large extent  be understood as either immediate or long-term consequences of the mass transfer process. Interesting as all these phenomena are, they are of no particular interest here. Rather, in the following I shall concentrate on evolutionary aspects, i.e. on the formation and evolution of CVs. Readers who are mainly interested in CVs as variable stars should instead turn to the monographs by \\citet{Warner1995} or \\citet{Hellier2002}. ",
        "conclusions": ""
    },
    "0809/0809.2803_arXiv.txt": {
        "abstract": "{} {We present a survey of the  ortho--\\HTDP (1$_{1,0}$--1$_{1,1}$) line toward a sample of 10 starless cores and 6 protostellar cores, carried out at the Caltech Submillimeter Observatory. The high diagnostic power of this line is revealed for the study of the chemistry, and the evolutionary and dynamical status of low-mass dense cores.} {The derived ortho--\\HTDP \\ column densities ($N(ortho-\\HTDP)$) are compared with predictions from simple chemical models of centrally concentrated cloud cores.} {The line is detected in 7 starless cores and in 4 protostellar cores.  $N(ortho-\\HTDP)$ ranges between 2 and 40$\\times$10$^{12}$~cm$^{-2}$ in starless cores and between 2 and 9$\\times$10$^{12}$~cm$^{-2}$ in protostellar cores. The brightest lines are detected toward the densest and most centrally concentrated starless cores, where the CO depletion factor and the deuterium fractionation are also largest. The large scatter observed in plots of $N(ortho-\\HTDP)$ vs. the observed deuterium fractionation and vs. the CO depletion factor is likely to be due to variations in the ortho--to--para (o/p) ratio of \\HTDP \\ from $>$ 0.5 for $T_{\\rm kin}$ $<$ 10~K gas in pre--stellar cores to $\\simeq$ 0.03 (consistent with $T_{\\rm kin}$ $\\simeq$ 15~K for protostellar cores).  The two Ophiuchus cores in our sample also require a relatively low o/p ratio ($\\simeq$ 0.3). Other parameters, including the cosmic-ray ionization rate, the CO depletion factor (or, more in general, the depletion factor of neutral species), the volume density, the fraction of dust grains and PAHs also largely affect the ortho--\\HTDP \\ abundance. In particular, gas temperatures above 15~K, low CO depletion factors and large abundance of negatively charged small dust grains or PAHs drastically reduce the deuterium fractionations to values inconsistent with those observed toward pre--stellar and protostellar cores. The most deuterated and \\HTDP--rich objects (L~429, L~1544, L~694-2 and L~183) are reproduced by chemical models of centrally concentrated (central densties $\\simeq$10$^{6}$ cm$^{-3}$) cores with chemical ages between 10$^4$ and 10$^6$~yr. Upper limits of the para-\\HTHOP (1$_1^-$--2$_1^+$) and para--\\DTHP (1$_{1,0}$--1$_{0,1}$) lines are also given. The upper limit to the para-\\HTHOP \\ fractional abundance is $\\simeq$10$^{-8}$ and we find an upper limit to the  para--\\DTHP /ortho--\\HTDP \\ column density ratio equal to 1, consistent with chemical model predictions of high density (2$\\times$10$^6$ cm$^{-3}$) and low temperature ($T_{\\rm kin}$ $<$ 10~K) clouds. } {Our results point out the need for better determinations of temperature and density profiles in dense cores as well as for observations of para-\\HTDP .} ",
        "introduction": "In the past decade, astrochemistry has become more and more crucial in understanding the structure and evolution of star forming regions.  There are no doubts that stars like our Sun form in gas and dust condensations within molecular clouds and that the process of star formation can only be understood by means of detailed observations of the dust (good probe of the most abundant and elusive molecule, \\MOLH) and molecular lines (unique tools to study kinematics and the chemical composition).  Millimeter and submillimeter continuum dust emission observations \\citep[see][for a detailed review of this topic]{wac07} have provided a very good probe of the density structure of dense cores, although uncertainties are still present regarding the dust opacity and temperature, both likely to change within centrally concentrated objects \\citep[but such variations have so far been hard to quantify observationally, see e.g.][]{bga03,plb03,pbm04}. Stellar counts in the near-infrared provide an alternative way of measuring the dust (and \\MOLH ) column and cloud structure \\citep{llc94}, independent of any variation in dust properties, but they cannot probe regions with extinctions above about 40-50 mag \\citep{all98}, i.e. the central zone of very dense cores, such as L~1544, where $A_{\\rm V} \\simeq$ 100 mag within 11$^{\\arcsec}$ \\citep{wma99}. It appears that many starless cores can be approximated as Bonnor--Ebert spheres \\citep{e55,b56}, with values of the central densities ranging from about 10$^5$ \\percc \\ \\citep[as in the case of B~68;][]{all01} to 10$^6$ \\percc \\ (for e.g., L~1544, L~183, and L~694-2; \\citealt{wma99}, \\citealt{plb03}, \\citealt{hwl03}). At the lower end of the central density range, dense cores appear to be isothermal, with gas temperatures close to 10~K \\citep{gwg02,tmc04}, whereas higher density cores have clear evidences of temperature drops in the central few thousand AU, with dust temperatures approaching about 7~K \\citep[][see also Bergin \\& Tafalla 2007 for a comprehensive review on starless cores]{ers01,plb03,pbm04,sg05,pbc07,ccw07}  \\citet{kf05} have proposed that the \"shallower\" cores (such as B~68) are in approximate equilibrium and will not evolve to form protostars, whereas the centrally concentrated ones (such as L~1544) are unstable cores that are proceeding toward gravitational collapse and the formation of protostars. Indeed, this is in agreement with the findings of \\cite{lba02}, who claim that B~68 is oscillating around an equilibrium state, and those of \\citet{cwz02a} and \\citet{vcc05}, who studied the kinematic structure of L~1544 and found that it is consistent with contraction in the core nucleus (or central contraction).  To trace the gas properties, \\AMM \\ and \\NTHP \\ have been extensively used for several years \\citep{bm89} and they seem to trace quite similar conditions, having comparable morphologies and line widths \\citep{bcm98,cbm02,tmc02,tmc04}, despite of the (two orders of magnitude) difference in the critical densities of the most frequently observed transitions (\\AMM (1,1) and \\NTHP (1-0), see \\citealt{pbc07}). This is quite lucky for astronomers considering that only relatively recently it has been realized that these two species are among the few that are left in the gas phase at volume densities above $\\sim$10$^5$ \\percc . In fact, CO, CS and, in general, all the carbon bearing species so far observed \\citep[with the exception of CN;][]{hwp08} are heavily affected by freeze-out in the central parts of dense starless cores \\citep{klv96,wlv98,cwt99,bcl01,blc02,bah02,cwz02b,tmc02,blc03,ccw04,tmc04,ppa05,ccw05,tsm06,pbc07}. The freeze-out of neutral species \\citep[in particular of CO,][]{dl84,rm00a,rm00b}, boosts the deuterium fractionation in species such as \\NTHP , \\AMM , H$_2$CO , and \\HCOP \\citep[see e.g.][]{bll95,trf00,cwz02b,blc03,ccw05,lgr06,glp06}. Also in star forming cores, in particular in the direction of Class 0 sources, the first protostellar stage \\citep[e.g.][]{awb00}, right after the pre-stellar phase, the deuterium fractionation is found to be very large \\citep{ccl98,lcc02,lgp02,lrg02,vsm02,pct02,vpc03,pch04,ccw04,mcr05,pct06,cch07}. This is thought to be the signature of a (recent) past in which the star forming cloud core experienced the low temperature and high density conditions typical of the most centrally concentrated starless cores, where some deuterated molecules (e.g. deuterated ammonia) are formed in the gas phase (and stored on dust surfaces) whereas others (such as deuterated methanol and formaldehyde) are likely formed onto dust surfaces and then released to the gas phase partially upon formation \\citep{gpc06,gwh07} as well as via the interaction with the newly born protostar, which can (i) heat dust grains, leading to mantle evaporation \\citep[as in Hot Cores and Corinos;][]{t90,ctc03,bcl04,bcn04,bcw07}, and/or (ii) \"erode\" dust mantles via sputtering in shocks produced by the associated energetic outflows \\citep[e.g.][]{lgp02}.  Questions that are still open are: (1) are \\NTHP , \\AMM \\ and their deuterated forms really tracing the inner portions of centrally concentrated cores on the verge of star formation?  Although Bergin et al. (2002), Pagani et al. (2005, 2007) found evidences of depletion in the center of B~68 and L~183, respectively, there are no signs of freeze--out for NH$_3$ in L~1544 (Crapsi et al. 2007) and for both nitrogen bearing species in L~1517B and L~1498 (Tafalla et al. 2004). At densities above $\\sim$10$^6$ \\percc , the freeze--out time scale is quite short ($\\sim$1,000 yr) and all heavy species are expected to condense onto grain mantles. Moreover, recent laboratory measurements clearly show that \\MOLN , the parent species of both \\AMM \\ and \\NTHP , should freeze--out at the same rate as CO \\citep[having similar binding energies and sticking probabilities;][]{ovf05,bfo06}. (2) For how long the high degree of deuterium fractionation observed in starless cores is maintained after the formation of a protostellar object?  The detection of strong ortho-\\HTDP \\trans \\ emission in the direction of L~1544 \\citep{cvc03}, and the conclusion that \\HTHP , with its deuterated counterparts, is one of the most abundant molecular ions in core centers, have opened a new way to study the chemical evolution \\citep{rhm03,rhm04,wfp04,fpw04,fpw05,fpw06a,fpw06b,ahr05} and the kinematics \\citep{vcc05} of the central few thousand AU of starless cores. Thus, \\HTDP \\  is an important tool to understand the chemical and physical properties of the material out of which protoplanetary disks and ultimately planetary systems form.  In the gas phase at temperatures below $\\sim$20~K, the deuterium fractionation is mostly regulated by the proton--deuteron exchange reaction: \\begin{eqnarray} \\HTHP + {\\rm HD} \\rightarrow \\HTDP + \\MOLH + \\Delta {\\rm E} , \\label{e_htdp} \\end{eqnarray} where $\\Delta$E \\citep[= 230~K,] []{mbh89} prevents the reverse reaction to be fast in cold regions, unless a significant fraction of \\MOLH \\ is in ortho form \\citep{ghr02,wfp04}. \\citet{fpw06a} showed that values of the ortho--to--para (o/p) \\MOLH \\ ratio much higher than 0.03 are inconsistent with the observed high levels of deuteration of the gas (see their Fig.~6). The above reaction, together with the freeze--out of neutral species \\citep[which boosts up the production rate of \\HTDP \\ compared to \\HTHP ; e.g.][]{aoi01}, allows the \\HTDP /\\HTHP \\ ratio to overcome the cosmic D/H ratio by several orders of magnitude. In fact, not only ortho--\\HTDP \\trans \\ has been found to be $\\simeq$1~K strong in L~1544, but also \\DTHP \\ has been detected toward another starless condensation \\citep[16293E, in Ophiuchus;][]{vpy04}.  The ortho--\\HTDP \\trans \\ line has also recently been mapped in L~1544 \\citep{vcc06}, finding that the \\HTDP \\ emitting region has a radius of about 5000 AU, comparable to the size of the \\NTDP (2-1) map made in the same region by \\citet{cwz02a}.  In this paper we present new ortho--\\HTDP \\trans \\ observations, carried out at the Caltech Submillimeter Observatory (CSO) antenna, in the direction of 10 starless cores and 6 cores associated with very young protostellar objects.  As it will be shown, the line has been detected in 7 of the 10 starless cores and in 4 out of 6 star forming cores. In Sect.~\\ref{sobs} the observational details are given. Results and ortho--\\HTDP \\ spectra are shown in Sect.~\\ref{sres}, together with a brief discussion on the upper limits of the para-\\DTHP (1$_{1,0}$-1$_{0,1}$) lines (para-\\HTHOP (1$_1^-$--2$_1^+$) upper limits can be found in the on-line appendix). \\HTDP \\ column densities are derived in Sect.~\\ref{sana}.  A chemical discussion, aimed at interpreting the observations, is described in Sect.~\\ref{schem} and conclusions are in Sect.~\\ref{scon}. ",
        "conclusions": "\\label{scon}  Low--mass dense cloud cores have been observed with the CSO antenna at the frequency of the ortho--\\HTDP (1$_{1,0}$--1$_{1,1}$) line.  The main conclusions of this work are:  \\noindent 1)  In starless cores, the line has been detected in 7 (out of 10) objects. The brightest lines ($T_{\\rm mb}$ = 0.7-0.9~K) are observed toward the densest and more dynamically evolved starless cores (L~1544, L~183, Oph~D, L~429, and L~694-2), where the deuterium fractionation and the CO depletion factor are largest. Non-detections are found in cores less centrally concentrated and dense than the rest of the sample (L~1498, TMC--2, and B~68). In B~68, recent APEX observations have detected a faint ortho--\\HTDP \\ line with intensities consistent with our upper limit.  \\noindent 2) Significant differences in line widths are observed also among starless cores, with the narrowest lines ($\\Delta {\\rm v}$ $\\simeq$ 0.4~\\kms ) found in TMC--1C, L~1517B, Oph~D, and L~183, whereas L~1544, L~429, and L~694-2 show $\\Delta {\\rm v}$ = 0.5--0.7~\\kms.  \\noindent 3) The ortho--\\HTDP \\trans \\ line has been detected in 4 out of 6 protostellar cores, where we find $T_{\\rm mb}$ $\\leq$ 0.5~K (the largest value observed is toward L~1521F, which hosts a very low luminosity object recently detected by Spitzer). The broadest lines are observed toward the two cores in Perseus NGC~1333--\\DCOP \\ and B~1, where $\\Delta {\\rm v}$ $\\simeq$1~\\kms. The ortho--\\HTDP \\trans \\ lines have broader non--thermal widths than \\NTDP (2--1) lines, even in L~1544, where the two transitions have similar extension and morphology.  \\noindent 4) The ortho--\\HTDP \\ column densities range between 2 and 40$\\times$10$^{12}$~cm$^{-2}$ in starless cores and between 2 and 9 $\\times$10$^{12}$~cm$^{-2}$ in protostellar cores. Thus, protostars in the earliest stages of their evolution appear to have already changed the chemical structure of their envelopes, by lowering the ortho--\\HTDP \\ fractional abundance but not their deuterium fractionation. This is probably due to (a few degree) variation of the gas temperature, which strongly affects the o/p--\\HTDP \\ ratio. A similar effect is also found in the two Ophiuchus cores, suggesting a (slightly) larger gas temperature than in the other (mainly Taurus and Perseus) cores.  \\noindent 5) Simple models suggest that variations in the gas kinetic temperature, CO depletion factor, grain size distribution, cosmic--ray ionization rate and volume densities can largely affect the \\HTDP \\ abundance relative to \\HTHP . In particular, gas temperatures above 15~K, low CO depletion factors and large abundance of negatively charged small dust grains or PAHs drastically reduce the deuterium fractionation to values inconsistent with those observed toward pre--stellar and protostellar cores.  \\noindent 6) Plots of the ortho--\\HTDP \\ column density vs. (i) the deuterium fractionation observed in species such as \\NTHP \\ and \\AMM , and (ii) the observed CO depletion factor, show a large scatter. The data can be reproduced by chemical models of centrally concentrated cores assuming variations in the o/p ratio of \\HTDP \\ (ultimately linked to kinetic temperature variations). Changes in the cosmic--ray ionization rate between 6 and 30$\\times$10$^{-18}$s$^{-1}$ can also explain part of the scatter, but not those objects with large deuterium fractionations {\\it and} low ortho--\\HTDP \\ column densities (such as the Ophiuchus pre--stellar cores and the protostellar ones), for which a decrease of the o/p--\\HTDP \\ ratio is required. The most deuterated and \\HTDP--rich objects are better reproduced by chemical models of centrally concentrated cores with $n_c$ = a few $\\times$ 10$^6$ \\percc \\ and chemical ages between 10$^4$ and 10$^6$~yr. Protostellar cores, plus the two Ophiuchus cores, are better matched by  lower o/p-\\HTDP \\ ratios, more dynamically evolved (central densities $\\ge$10$^7$ cm$^{-3}$) models and chemical ages of $\\simeq$10$^3$-10$^4$ years. To better constrain dynamical and chemical ages, the rate coefficient for the proton-deuteron exchange reaction needs to be well defined.  \\noindent 7) The para-\\HTHOP (1$_1^-$--2$_1^+$) upper limits are consistent with radiative transfer calculations if the fractional abundance of \\HTHOP \\ is $\\la$10$^{-8}$.  The para--\\DTHP (1$_{1,0}$--1$_{0,1}$) upper limits provide an upper limit of the para--\\DTHP \\ to ortho--\\HTDP \\ column density ratio ($\\equiv p/o$).  We find $p/o$$<$1 in all the sources (except for L~183 and NGC~1333--\\DCOP , where $p/o < 3$), consistent with  chemical model predictions of high density (2$\\times$10$^6$~\\percc) and low temperature ($T_{\\rm kin} < 10$~K) clouds (\\citealt{fpw04}).  More accurate determinations of temperature and density profiles, as well as observations of the para--\\HTDP (1$_{0,1}$-1$_{0,0}$) line at 1.37~THz, are sorely needed to put more stringent constraints on gas--grain chemical processes in dense cloud cores.  \\appendix"
    },
    "0809/0809.2518_arXiv.txt": {
        "abstract": "{ The local stellar mass density is observed to be significantly lower than the value obtained from integrating the cosmic star formation history (SFH), assuming that all the stars formed with a Salpeter initial mass function (IMF). Even other favoured IMFs, more successful in reconciling the observed $z=0$ stellar mass density with that inferred from the SFH, have difficulties in reproducing the stellar mass density observed at higher redshift. In this study we investigate to what extent this discrepancy can be alleviated for any universal power-law IMF. We find that an IMF with a high-mass slope shallower ($2.15$) than the Salpeter slope ($2.35$) reconciles the observed stellar mass density with the cosmic star formation history, but only at low redshifts. At higher redshifts $z>0.5$ we find that observed stellar mass densities are systematically lower than predicted from the cosmic star formation history, for any universal power-law IMF. } ",
        "introduction": "A number of studies have noted that the observed local stellar mass density, assuming a universal initial mass function (IMF) equivalent to the Salpeter (1955) initial mass function (IMF), is significantly smaller than that inferred from the cosmic star formation history (SFH). Integration of the SFH of Hopkins \\& Beacom (2006; hereafter HB06) assuming a Salpeter IMF suggests a local stellar mass-density of $\\Omega_{*;SFH}\\sim 0.0066 \\pm 0.0015$ (in units of the critical density) whereas analysis of large galaxy surveys suggest a value $\\Omega_{*;obs}\\sim 0.0041 \\pm 0.0010$ (Wilkins, Trentham \\& Hopkins 2008 hereafter WTH08). This behaviour continues at higher redshifts where the observed stellar mass-density history (SMH) remains systematically lower than that inferred from the SFH (HB06, WTH08) as shown in Figure 1.  Both the observed SMH and that predicted from the SFH are dependent upon the assumed IMF, although the scaling for each is different (see WTH08, also \\S\\,3 and \\S\\,4 below). The commonly assumed IMFs of Kroupa (2001) and Chabrier (2003) result in a marginally better correspondence between the observed SMH and that predicted from the SFH (as shown in the lower panel of Figure 1 for the Kroupa 2001 IMF). Note the lower normalisation of the data points in this panel compared to those assuming the Salpeter IMF, and the fact that while the Kroupa (2001) IMF scales down the SMH inferred from the SFH, it also scales down the observed SMH data, although not by as much. While this IMF choice reduces the $z=0$ discrepancy slightly, there remains a significant inconsistency at higher redshift. The IMF of Baldry \\& Glazebrook (2003, hereafter BG03), goes further toward resolving this inconsistency (Hopkins \\& Beacom 2008) although even then the higher redshift discrepancy remains. It is possible that there remains some systematic errors in the SFH or SMH that may explain the observed discrepany. These include issues such as uncertainties in the extent of dust obscuration on the SFH, or redshift-dependent effects associated with observable rest-frame wavelengths for the SMH. These are discussed by WTH08 who conclude that significant systematic errors, at the level required to resolve the discrepancy, seem unlikely.  There have recently been several indications that the IMF may indeed not be universal, but instead evolves with redshift. These have arisen from a variety of independent considerations. Van Dokkum (2008), for example, invokes evolution of the IMF to reconcile the rate of evolution in both luminosity and colour for early-type cluster galaxies at moderate redshifts ($z<0.8$). Dav{\\'e} (2008) finds that an evolving IMF better explains the observed evolution in the relationship between individual galaxy star formation rates and stellar masses. Baugh et al.\\ (2005) requires an IMF with a flatter high-mass slope at higher redshifts to explain the observed numbers of sub-mm galaxies, and several other studies in the past few years also favour a flatter IMF at higher redshift (see references in WTH08).  We put aside the issue of an evolving IMF for the purposes of the current study. Here we are motivated by the different sensitivity in the scaling of the SMH and the SFH to the choice of IMF, and we extend the analysis of WTH08 by considering the extent to which any universal power-law IMF can reconcile the observed SFH and SMH data. In \\S\\,2 we introduce the stellar, integrated galaxial and cosmic initial mass functions, and discuss the possible distinctions. In \\S\\,3 and \\S\\,4 we quantify the effect of the IMF upon both the observed stellar mass density and that inferred from the SFH. In \\S\\,5 we use a $\\chi^{2}$ minimisation to establish the universal IMF that best reconciles the observed SMH with the SFH and in \\S\\,6 we present our conclusions. Throughout this work, we assume a flat $\\Lambda$ CDM cosmology with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\rm matter}=0.3$ and $H_{0} = 70\\,\\, {\\rm kms}^{-1}\\, \\rm{Mpc^{-1}.}$   \\begin{figure} \\includegraphics[width=20pc]{salpeter.eps} \\caption{The evolution of stellar mass density assuming a Salpeter IMF (top panel) and a Kroupa (2001) IMF (bottom panel). The shaded region is the stellar mass density inferred from the uncertainty envelope of the SFH of HB06. The points are observed values from compilation of WTH08.} \\end{figure} ",
        "conclusions": "We have investigated the effect of a universal cosmic initial mass function on the correspondence between the observed build up of stellar mass density (WTH08) and that predicted from the star formation history of HB06. We find that a cosmic IMF with a high-mass slope of $\\alpha_{2}=2.15\\pm0.15$ can produce a statistical reconciliation between the star formation history and observed stellar mass density at low redshifts ($z<0.5$). At higher redshifts, however, the observed stellar mass density lies systematically below that predicted from the star formation history. The remaining discrepancy, of order a factor of 3 at redshifts $2\\lesssim z \\lesssim 4$, may be possible to reconcile through systematic errors in the SMH or SFH measurements, although this seems unlikely (see discussion in WTH08). An evolving IMF remains an attractive solution to this problem.  Any possible evolving cosmic IMF must resolve not only the discrepancy between the SMH and SFH, but many other aspects of galaxy evolution (e.g., Dav{\\'e} 2008, van Dokkum 2008) together with maintaining consistency with the extragalactic background light (Fardal 2007), and the chemical evolution of galaxies, together with matching the observed local IGIMF.  \\vskip 20pt  \\noindent We would like to thank Karl Glazebrook for discussions leading to this paper. We would also like to thank Pavel Kroupa, Carsten Weidner and Thomas Maschberger for useful discussion about the IGIMF. Further, we also thank the anonymous referee for helpful comments and suggestions. SMW acknowledges support of an STFC studentship and of King's College, NT acknowledges support provided by STFC, AMH acknowledges support provided by the Australian Research Council in the form of a QEII Fellowship (DP0557850) and RT is funded by the Funda\\c{c}\\~{a}o para a Ci\\^{e}ncia e a Tecnologia under the reference PRAXIS SFRH/BD/16973/04."
    },
    "0809/0809.5081_arXiv.txt": {
        "abstract": "The Carina Nebula (NGC~3372) is our richest nearby laboratory in which to study feedback through UV radiation and stellar winds from very massive stars during the formation of an OB association, at an early phase before supernova explosions have disrupted the environment. This feedback is triggering new generations of star formation around the periphery of the nebula, while simultaneously evaporating the gas and dust reservoirs out of which young stars are trying to accrete. Carina is currently powered by UV radiation from 65 O-type stars and 3 WNH stars, but for most of its lifetime when its most massive star ($\\eta$ Carinae) was on the main-sequence, the Carina Nebula was powered by 70 O-type stars that produced a hydrogen ionizing luminosity 150 times stronger than in the Orion Nebula.  At a distance of 2.3 kpc, Carina has the most extreme stellar population within a few kpc of the Sun, and suffers little interstellar extinction.  It is our best bridge between the detailed star-formation processes that can be studied in nearby regions like Orion, and much more extreme but also more distant regions like 30 Doradus.  Existing observations have only begun to tap the tremendous potential of this region for understanding the importance of feedback in star formation --- it will provide a reservoir of new discoveries for the next generation of large ground-based telescopes, space telescopes, and large submillimeter and radio arrays. ",
        "introduction": "Feedback from young massive stars may play an integral role in star and planet formation.  Most stars are born in OB associations (Blaauw 1964; Lada \\& Lada 2003), in the vicinity of hot massive stars spawned only from giant molecular clouds.  The effects of feedback from these massive stars cannot be studied in nearby quiescent regions of star formation.  \\begin{figure}[!ht] \\centering \\includegraphics[draft=False,width=0.95\\textwidth]{smithbrooks.f1.eps} \\caption{A large field-of-view photograph of the Carina Nebula made from UK Schmidt plates by David Malin. Anglo-Australian Observatory/Royal Obs.\\ Edinburgh.} \\label{fig:big} \\end{figure}  The Carina Nebula (Fig.~\\ref{fig:big}; NGC~3372) is the southern hemisphere's largest and highest surface brightness nebula --- some central parts of the nebula are even brighter than the Orion Nebula. It provides an ideal laboratory in which to study ongoing star formation in the vicinity of some of the most massive stars known, including our Galaxy's most luminous known star, $\\eta$ Carinae. Carina's star clusters are not as densely packed as those of 30~Dor in the LMC (Massey \\& Hunter 1998), or objects in our own Galaxy like the Arches cluster near the Galactic center (Najarro et al.\\ 2004; Figer et al.\\ 1999, 2002), NGC~3603 (Moffat et al.\\ 2002), or W49 (Welch et al.\\ 1987).  However, these other regions are too distant for detailed studies of small-scale phenomena like irradiated protoplanetary disks and jets, and their study is hampered by considerably more extinction. The low extinction toward Carina combined with its proximity and rich nebular content provide a worthwhile trade-off.  For the most massive stars, lifetimes of $\\sim$3 Myr before they explode as supernovae (SNe) are comparable to the time it takes to clear away their surrounding large-scale distribution of molecular material.  Consequently, we already have a situation in Carina where the most massive members like $\\eta$~Car are approaching their imminent demise while new stars are being born from dense molecular gas only 5--20 pc away.  In the next 1--2 Myr, there will be several energetic SNe in Carina.  These will carve out an even larger cavity in the ISM and form a giant superbubble in the Galactic plane, and may pollute protoplanetary disks with nuclear-processed ejecta.  In the mean time, studying this rich region provides a snapshot of a young proto-superbubble (Smith et al.\\ 2000) energized only by UV radiation and stellar winds, just before the disruptive SN-dominated phase.    \\begin{figure}[!ht] \\centering \\includegraphics[draft=False,width=0.95\\textwidth]{smithbrooks.f2.eps} \\caption{An {\\it HST}/WFPC2 image of $\\eta$ Carinae and its surrounding nebulosity (N.~Smith/NASA).} \\label{fig:eta1} \\end{figure} \\begin{figure}[!ht] \\centering \\includegraphics[draft=False,width=0.8\\textwidth]{smithbrooks.f3.eps} \\caption{The historical visual light curve of $\\eta$ Carinae (from Frew 2004).} \\label{fig:lightcurve} \\end{figure}  \\subsection{The Star:  Eta Carinae}  From before the time of Herschel up to the present, most observational effort in the Carina Nebula had focussed on the variable star $\\eta$ Carinae (Fig.~\\ref{fig:eta1}) and its immediate surroundings.  The unusual variability of $\\eta$~Car was noted as early as the 17th century by Halley, when it was reputed to vary between 4th and 2nd magnitude.  Later, in the early and mid 19th century, it was observed by John Herschel (1847) as it brightened significantly to become the second brightest star in the sky.  It then faded from view by the 1870s.  Frew (2004) gives a careful and detailed account of the historical light curve (Fig.~\\ref{fig:lightcurve}).  During that event, the star ejected more than 10~$M_{\\odot}$ with almost 10$^{50}$ erg of kinetic energy (Smith et al.\\ 2003b), which has since expanded to form the so-called ``Homunculus'' nebula seen today in {\\it HST} images like the one in Figure~\\ref{fig:eta1}.  After more than 160 years and intensive study at all wavelengths, the underlying cause of this so-called ``Great Eruption'' is still the most enduring mystery associated with $\\eta$~Car.  Additionally, much current work centers around the 5.5-yr periodic variability that is likely the product of $\\eta$ Carinae being an eccentric binary system (e.g., Damineli et al.\\ 2000).  Our focus here is on star formation in the surrounding giant H~{\\sc ii} region, so we do not aim to review or summarize the vast and complicated literature concerning $\\eta$ Carinae and its ejecta (see Davidson \\& Humphreys 1997; although much has been added to our understanding in the subsequent decade making that review outdated in some respects).  However, the wild variability of $\\eta$~Car, its extreme nature, and its imminent demise are relevant for nebular structures in the region.  From the point of view of understanding star formation in the region, two important facts are most worth remembering:  (1) Although $\\eta$ Carinae is probably the most massive and luminous star known in the Milky Way, {\\it at the current time it contributes essentially nothing to the radiative energy budget of the surrounding region}.  This strange twist arises because $\\eta$ Carinae is currently surrounded by a nearly opaque dust shell ejected 160 years ago, which absorbs nearly all of the star's UV and visual luminosity.  That luminosity is then re-radiated in the thermal IR, and it escapes from the nebula, exerting little influence on the surrounding gas in the star-forming complex.  (2) However, before 1843 (for about 3 Myr prior), $\\eta$~Carinae had a tremendous influence on the UV energy budget of the entire region. Its UV luminosity dominated the ionization of the main cavity and it sculpted many of the most prominent elephant trunks in the region.  When the giant eruption ended about 150 years ago and $\\eta$ Car faded due to dust obscuration, the UV luminosity of the region dropped by 20\\% (Smith 2006a) and the remaining O stars that are scattered throughout the region now dominate the ionization.  This change, which arises because of $\\eta$ Car's precarious instability and its evolved state, is unique among Galactic H~{\\sc ii} regions.    \\begin{figure}[!h] \\centering \\includegraphics[draft=False,width=0.95\\textwidth]{smithbrooks.f4.eps} \\caption{Left: Herschel's drawing of the Keyhole Nebula, made during $\\eta$ Carinae's 19th century outburst.  Right: A modern color photograph of the Keyhole made by David Malin.} \\label{fig:herschel} \\end{figure}  \\subsection{Historical Perspective: The Keyhole}  The famous ``Keyhole'' nebula at the heart of the greater Carina Nebula also has a rich history, aided by its location immediately adjacent to $\\eta$ Carinae, as well as its extremely high surface brightness.  The physical origin of the Keyhole shape is still not understood (although see Smith 2002b).  While the old-fashioned keyhole shape is easily recognized in Herschel's drawing, more creativity is required to identify this structure in modern photographs (see comparison in Fig.~\\ref{fig:herschel}).  The most dramatic difference is in the lower part of the Keyhole, where much of the diffuse emission is brighter in Herschel's drawing.  This has caused some consternation to observers (e.g., Bok 1932), making some suspect that Herschel may have had an over-active imagination.  In fact, it is likely that Herschel's drawing was largely correct, and that the appearance of the nebula has actually changed with time.  Herschel's drawing in Figure~\\ref{fig:herschel} was made around 1840, during the Great Eruption of $\\eta$~Car, before the star was obscured by its dusty nebula.  Thus, at that time the Keyhole was illuminated by a much brighter star than is seen today, and the reflection nebulosity around the Keyhole could have differed in appearance.  We now know from both historical and modern data that the reflected spectrum of $\\eta$~Car is seen in the Keyhole (Walborn \\& Liller 1977; Lopez \\& Meaburn 1986), confirming that it is at the same distance as $\\eta$~Car. Given that the larger Carina Nebula is more than 150 light years across, one might wonder if light echoes from the Great Eruption can still be seen in archival images.  As for the diffuse emission from the larger Carina Nebula, Herschel did have a few choice words to describe it: ``It would be manifestly impossible by verbal description to give any just idea of the capricious forms and irregular gradations of light affected by the different branches and appendages of this nebula.''  The rich, complex structure that Herschel was referring to is associated with the many dust pillars and dark globules, the likes of which are not seen in the Orion Nebula.  Perhaps this is an early testament to feedback from the extreme stellar population in Carina, providing us with a snapshot of a molecular cloud getting shredded by its progeny.  In the 20th century, observers began to take a broader view of Carina, realizing that it spanned several degrees of the sky.  In this review, we are concerned with star formation across the full extent of the Carina Nebula.  The nomenclature for this region is interchanged loosely, being referred to variously as NGC~3372, the ``Carina Nebula'', the ``Great Carina Nebula'', the ``Keyhole Nebula'', the ``Eta Carinae Nebula'', or (incorrectly) as simply ``Eta Carinae''. We reserve ``Eta Carinae'' for discussions of the massive star itself, the ``Homunculus'' nebula (Gaviola 1950), plus surrounding ejecta nebulosity within about 1\\arcmin\\ (e.g.\\ Smith et al.\\ 2005a).  We use the name ``Keyhole Nebula'' to refer to the structures in Figure~\\ref{fig:herschel}, within about 10\\arcmin\\ of $\\eta$ Car, while we use the term ``Carina Nebula'' to refer to the larger giant H~{\\sc ii} region spanning several square degrees (Fig.~\\ref{fig:big}), of which $\\eta$ Car and the Keyhole are only a small component at the center.   \\subsection{Pioneering Observations of the Nebulosity}  The first modern interference-filter photographs of the region were made by Walborn (1975) and Deharveng \\& Maucherat (1975).  These images showed the complex structures of the Keyhole and the rest of NGC~3372 in ionized gas, and revealed many silhouette structures, such as dust pillars and globules, and the large, dark V-shaped dust lane that bisects the nebula.  It was established early on that the V-shaped dark lanes are the result of obscuration by dust intermixed with molecular gas (\\citealt{Dickel741}; \\citealt{Feinstein73}; \\citealt{Smith87}).  Far-infrared (far-IR) emission was detected from the dark lanes by \\citet{Harvey79} at 80~$\\mu$m and \\citet{Ghosh88} at 120--300 $\\mu$m, implying that the dust was warm, with T$\\simeq$30--40 K.  The first single-dish molecular line studies of Carina by Gardner et al.\\ (1973) and \\citet{Dickel742} revealed that these dark dust lanes coincided with the bulk of the molecular gas.  Pioneering radio observations of the Carina nebula focused on the central region with angular resolutions (FWHM) that were typically 3\\arcmin\\ (e.g., \\citealt{Gardner68}; \\citealt{Shaver70}; \\citealt{Gardner70}; \\citealt{Jones73}; \\citealt{Dickel741}; \\citealt{Huchtmeier75}; \\citealt{Retallack83}; \\citealt{Tateyama91}). \\citet{Gardner70} first identified the two bright concentrations known as Car~I and Car~II (see Fig.~\\ref{fig:co}).  Subsequent observations established that these two sources were thermal and represented two separate H~{\\sc ii} regions. Car I was found to be located toward the north-western part of the inner nebula, and Car II was identified with the central Keyhole Nebula. \\citet{Retallack83} was the first to separate $\\eta$~Car from the bright radio continuum emission of Car~II, and to highlight the similar ring-shaped structures of the radio source Car~II and the Keyhole Nebula seen at visual wavelengths.  Detections of non-thermal emission toward the nebula were reported by \\citet{Jones73} and \\citet{Tateyama91}. However, no subsequent work has found convincing evidence to confirm these non-thermal sources. The non-thermal source G287.7--1.3 first reported by \\citet{Shaver70} has since been shown to be extragalactic \\citep{Caswell75}.  The magnetar studied by Gaensler et al.\\ (2005) is located at a larger distance and well outside the Carina Nebula 1$^{\\circ}$ to the east.  It is amusing to note that the first large-scale CO study of the Carina Nebula at 115 GHz used the {\\it optical} ESO 3.6-m telescope at La Silla in 1978 \\citep{deGraauw81}. These data were soon followed by the CO survey undertaken by \\citet{ Whiteoak84} using the 4-m radio telescope in Sydney.  The resulting maps with a $\\sim$2\\arcmin\\ (FWHM) beam offered the most detailed view of the Carina molecular cloud for more than a decade, and showed that the bulk of the molecular material was concentrated into a northern and a southern cloud.  The total extent of the molecular cloud was revealed later in the Columbia CO survey of the Galactic plane \\citep{Grabelsky88}. The survey showed that the Carina molecular cloud is part of a GMC complex that is situated on the near side of the Carina arm and extends over 150 pc between $l$=284.7 and 289, with a total estimated mass of 6.7$\\times$10$^5$~M$_{\\odot}$.  Results from the first studies of the velocity field within the Carina Nebula showed complex structure that is still not understood. Interpretations of the large-scale dynamics measured across the nebula involved merging spiral arms \\citep{Tateyama91}, rotating neutral clouds \\citep{Meaburn84} or old \\HII\\ regions \\citep{Cersosimo84}. On a smaller scale, the complicated dynamics measured in the vicinity of Car~II alluded to the tremendous influence of the stellar wind from $\\eta$~Car. Motions of ionized gas were studied via radio hydrogen recombination lines (e.g. \\citealt{Gardner70}; \\citealt{Huchtmeier75}; \\citealt{Cersosimo84}) and optical-line emission profiles (e.g. \\citealt{Deharveng75}). All studies detected double-peaked profiles in the vicinity of Car II. \\citet{Deharveng75} proposed that the ionized gas lies in an expanding shell whose center of expansion is $\\eta$~Car.  Additional evidence for peculiar, high-speed gas motions in the vicinity of $\\eta$ Car has been seen in resonance-line absorption profiles (\\citealt{Walborn75}).  These unusual multi-component profiles have received continued study in the optical and UV (Walborn 1982; Walborn \\& Hesser 1982; \\citealt{Laurent82}; Garcia \\& Walborn 2000; Walborn et al.\\ 2002a, 2007), but are still not understood. Velocity components range over 500 km s$^{-1}$ and have sometimes been interepreted as evidence for a supernova remnant along the line of sight, although proof of this hypothesis remains elusive.  Elliot (1979) interpreted broad emission components as evidence for a supernova remnant in the Keyhole, but these were actually broad lines in the wind spectrum of $\\eta$ Car itself that were reflected by dust in the Keyhole (see Walborn \\& Liller 1977; Lopez \\& Meaburn 1986).  Nevertheless, there are a few observational clues that hint at a possible young supernova remnant (SNR) inside Carina.  In addition to the complex blueshifted absorption lines already noted, IR spectroscopy with ISO revealed an unusual and strong broad 22~$\\mu$m emission feature, attributed to silicates by Chan \\& Onaka (2000), who note that the only other source where a similar emission feature is seen is the young supernova remnant Cas~A.  Carina also displays spatially extended diffuse soft X-ray emission across much of its central $\\sim$1$^{\\circ}$ region.  The diffuse X-ray emission was first detected in {\\it Einstein} data by Seward et al.\\ (1979), although Seward \\& Chlebowski (1982) concluded that the diffuse emission could be accounted for simply by mechanical energy input from stellar winds.  We discuss this further in Sect.\\ 8.2.   \\subsection{Prevailing View of Star Formation in Carina}  Early far-IR and molecular surveys described above found no luminous, embedded sites of active star formation in the bright central region of the nebula.  Additionally, no bright maser sources were identified in the region (e.g., \\citealt{Caswell98}).  Hence, the prevailing view until the mid-1990s was that Carina was an evolved H~{\\sc ii} region devoid of active star formation.  This view has changed dramatically in recent years, due in part to larger spatial coverage and higher-quality data at IR and radio wavelengths.  In just the past decade, the Carina Nebula has been recognized as a hotbed of active, ongoing star formation where the star formation is occuring primarily at the periphery of the nebula (Smith et al.\\ 2000; Megeath et al.\\ 1996).  It provides a laboratory to study several star formation phenomena in great detail, all of which have recently been identified here: (1) evaporating protoplanetary disks (the so-called ``proplyds''), small cometary clouds, or globules (Smith et al.\\ 2003a, 2004b; Smith 2002$b$; Brooks et al.\\ 2000, 2005; Cox \\& Bronfman 1995), (2) the erosion of large dust pillars and the triggering of a second generation of star formation within them (Smith et al.\\ 2000, 2005b; Rathborne et al.\\ 2004; Megeath et al.\\ 1996), (3) irradiated Herbig-Haro (HH) jets that signify active accretion (Smith et al.\\ 2004$a$), (4) photodissociation regions (PDRs) on the surfaces of molecular clouds across the region (Brooks et al.\\ 2003, 2005, 1998; Rathborne et al.\\ 2002; Smith et al.\\ 2000; Smith \\& Brooks 2007; Mizutani et al.\\ 2004), and on the largest scales, (5) the early stages of a fledgeling superbubble (Smith et al.\\ 2000; Smith \\& Brooks 2007).  Most of the current star-formation activity seems to be occurring near the edges of the nebula, although this is not strictly true as some of the most recent data show.  In general, the central clusters Tr~14 and Tr~16 tend to be devoid of star formation.  DeGioia-Eastwood et al.\\ (2001) and Tapia et al.\\ (2003) noted the presence of young stars with IR excess in these clusters, but these were consistent with age spreads of several million years among the low/intermediate-mass stellar population.  Similarly, Sanchawala et al.\\ (2007) noted that $\\sim$300 X-ray sources in these clusters have properties consistent with low/intermediate-mass pre-MS stars.  In any case, the most active sites of ongoing star formation are not the central clusters, but the outlying areas of the nebula.  Because of its proximity, large size on the sky, and its extreme nature, the level of detail we can observe in the Carina Nebula makes it appear bewilderingly complex, and it is just beginning to reveal its secrets to us.  The following sections focus on some of the recent work that has shed new light on Carina as an active star forming region.   \\begin{figure}[!ht] \\centering \\includegraphics[draft=False,width=0.95\\textwidth]{smithbrooks.f5.eps} \\caption{The Milky Way in Carina observed by IRAS at 25 $\\mu$m (blue), 60~$\\mu$m (green), and 100~$\\mu$m (red).  Several other well-known regions are labeled.} \\label{fig:iras} \\end{figure} ",
        "conclusions": ""
    },
    "0809/0809.5048_arXiv.txt": {
        "abstract": "{}{The earliest stages of intermediate- and high-mass star formation remain poorly understood. To gain deeper insights, we study a previously discovered protostellar source that is deeply embedded and drives an energetic molecular outflow.}{The source, UYSO\\,1, located close to IRAS\\,07029--1215 at a distance of about 1\\,kpc, was observed in the (sub)millimeter and centimeter wavelength ranges, as well as at near-, mid-, and far-infrared wavelengths.}{The multi-wavelength observations resulted in the detection of a double intermediate-mass protostar at the location of UYSO\\,1. In addition to the associated molecular outflow, with a projected size of 0.25~pc, two intersecting near-infrared jets with projected sizes of 0.4~pc and 0.2~pc were found. However, no infrared counterparts to the driving sources could be detected in sensitive near- to far-infrared observations (including \\textsl{Spitzer}). In interferometric millimeter observations, UYSO\\,1 was resolved into two continuum sources with high column densities ($>10^{24}$~cm$^{-2}$) and gas masses of 3.5~$M_\\odot$ and 1.2~$M_\\odot$, with a linear separation of 4200~AU. We report the discovery of a H$_2$O maser towards one of the two sources. Within an appropriate multi-wavelength coverage, the total luminosity is roughly estimated to be $\\approx$50~$L_\\odot$, shared by the two components, one of which is driving the molecular outflow that has a dynamical timescale of less than a few thousand years. The jets of the two individual components are not aligned.  Submillimeter observations show that the region lacks the typical hot-core chemistry.}{We find two protostellar objects, whose associated circumstellar and parent core masses are high enough to suggest that they may evolve into intermediate-mass stars. This is corroborated by their association with a very massive and energetic CO outflow, suggesting high protostellar accretion rates. The short dynamical timescale of the outflow, the pristine chemical composition of the cloud core and absence of hot core tracers, the absence of detectable radio continuum emission, and the very low protostellar luminosity argue for an extremely early evolutionary stage.} ",
        "introduction": "In the investigation of the earliest stages of intermediate- and high-mass star formation, the relative importance of formation processes as different as disk accretion and coalescence is still an open question (for a recent review, see e.g.~\\citealp{beu07}). Detailed multi-wavelength studies are needed to illuminate the often confusing observational picture.  Since high-mass protostars are less ubiquitous than their low-mass counterparts, they are on average more distant. The resulting low linear resolution often leads to an oversimplified picture. This is well illustrated by the study of \\citet{bsg02}, which reported a case where a single bipolar molecular outflow discovered previously with single-dish radio telescopes in a massive-star--forming (MSF) region was resolved into a multiple outflow system in interferometric data (see also \\citealp{bss03}). \\citet{dav04} analyzed near-infrared data of two collimated jets in a similar MSF region. They conclude that these jets are very similar to their low-mass counterparts.  We report the results of a comprehensive multi-wavelength follow-up study of a candidate massive protostellar source that was discovered close to IRAS~07029--1215 during a millimeter continuum survey of the surroundings of luminous IRAS sources in the outer galaxy \\citep{kle05}. UYSO\\,1 is a deeply embedded, very young source at a distance of only 1~kpc, powering a high-velocity bipolar CO outflow; \\citet{jan04}\\defcitealias{jan04}{Paper~I}\\citepalias{jan04} derived the mass of the UYSO\\,1 cloud core from the CO(3-2) line emission to be 44~$M_{\\odot}$, inferring the hydrogen column density from the integrated CO emission, while the mass derived from optically thin dust continuum emission ($\\lambda=850\\um$), tracing high column densities, is 15~$M_{\\odot}$. The mass of the outflow was estimated to be 5.4~$M_{\\odot}$. For details on these earlier mass estimates as well as for a discussion of the distance, we refer to \\citetalias{jan04}. Compared to other massive molecular outflows, this mass is at the lower end of the range reported by \\citet{bss02}, but most of their sources are much more distant, at an average distance of $4.7\\pm3.4$~kpc. No plausible driving source could be identified in IRAS and MSX mid- and far-infrared data. In single-dish observations, \\citet{beu08} detected C$_2$H submillimeter emission towards UYSO\\,1, a transition that may be a tracer of the earliest stages of (massive) star formation.  The aim of our follow-up observations of UYSO\\,1 was to study the region at higher angular resolution and to search for the driving source of the enormous CO outflow. For these purposes, we conducted observations in the infrared regime as well as in the millimetric to centimetric wavelength ranges. We note that for very early evolutionary stages, outflow mass entrainment rates and column densities are more conclusive in determining whether a massive star forms than a mass determination since material is still rapidly accreted. In Section~\\ref{sec_obse}, we describe the variety of observations carried out before presenting the results in Section~\\ref{sec_resu}. We discuss our findings in Section~\\ref{sec_disc} and conclude in Section~\\ref{sec_conc}.  \\begin{figure} \\centering \\includegraphics[width=\\linewidth, angle=-90]{0598_fig1.eps} \\caption{H$_2$ S(1) emission, as observed with OmegaPrime at the Calar Alto 3.5m telescope (filter NB2122), in logarithmic scaling. The box denotes the close-up VLT view in Fig.~\\ref{jcmtisaac}. The B star HD\\,53623 is marked by an arrow and the location of IRAS 07029-1215 is marked by its uncertainty ellipse.} \\label{calar_vlt} \\end{figure}  \\begin{figure*} \\centering \\includegraphics[width=17cm, angle=-90]{0598_fig2.eps} \\caption{H$_2$ S(1) emission (NB2.13) in logarithmic scaling, overlaid with the JCMT CO(3-2) observations -- total line 10\\% to 90\\% in 10\\% steps, line wings 30\\%, 50\\%, 70\\% and 90\\% -- as well as the two continuum sources seen with the PdBI (green ellipses indicating the FWHM synthesized beam). Four arrows indicate the locations of the jets, and an additional small arrow indicates the eastern bow shock of the smaller jet.} \\label{jcmtisaac} \\end{figure*}   \\begin{figure*} \\centering  \\includegraphics*[width=\\linewidth]{0598_fig3.eps}  \\caption{Left: continuum-subtracted H$_2$ S(1) emission (NB2.13--NB2.09) in logarithmic scaling. The subtraction of the brightest stars left a few artefacts. The box indicates the field of view of the right panel; ellipses indicate the positions of the two millimeter continuum sources as in Fig.~\\ref{jcmtisaac}. Right: NIR-continuum--only (NB2.09) close-up view of the central area in linear scaling with the 3\\,mm PdBI continuum data shown as contours in steps of 1\\,mJy.} \\label{nircont} \\end{figure*} ",
        "conclusions": "\\label{sec_conc} We report the detection of two highly collimated jets intersecting close to the position of the previously identified candidate massive protostar UYSO\\,1. Follow-up high-resolution millimeter radio observations show two continuum sources with masses of 3.5~$M_\\odot$ (UYSO\\,1a) and 1.2~$M_\\odot$ (UYSO\\,1b) and high column densities ($>10^{24}$~cm$^{-2}$) close to the geometric centers of the jets which probably are their powering sources. The projected distance between the two sources is 4200~AU. UYSO\\,1a appears to power the energetic molecular outflow that was previously discovered towards UYSO\\,1. UYSO\\,1b contains a previously unknown H$_2$O maser source while a search for CH$_3$OH maser emission remained unsuccessful. The two millimeter protostars are not detected in deep near-infrared imaging and remain undetected also at mid- to far-infrared wavelengths. Also, no developing H\\,{\\sc II} regions were found. Submillimeter observations show that the region is rather cold and does not show typical hot-core chemistry. Curiously, not even emission in NH$_3$ or N$_2$H$^+$ was detected. An attempt to constrain the combined SED of the two sources yields an estimated luminosity of roughly $L\\approx50\\,L_\\odot$, which is about an order of magnitude lower than a tentative estimate of the accretion luminosity. The key to combining these new results into a coherent picture probably lies in the unknown outflow inclination angle with respect to the line of sight. A large inclination angle would lead to configurations in which the circumstellar disk is seen nearly edge-on, obscuring the central object and helping to explain their infrared non-detections. For moderately large inclination angles of about $>25\\,^\\circ$, where a thick circumstellar disk and a dense envelope would obscure the central object, the molecular outflow powered by UYSO\\,1a has a dynamical timescale of $<10000$ years with enormous outflow mass entrainment rates of $>\\dot{M}=3\\times10^{-4}\\,M_\\odot$\\,yr$^{-1}$. The luminosity of the corresponding near-infrared jet is comparable to the most luminous jets found in a survey of Orion A. In addition to the peculiar chemistry and the still relatively low luminosity, this indicates a very early evolutionary stage."
    },
    "0809/0809.0451_arXiv.txt": {
        "abstract": "{}{We report a serendipitous detection of rapid, large amplitude flux density variations in the highly core-dominated, flat-spectrum radio quasar 1156+295 during an observing session at the Very Long Baseline Array (VLBA).}{The source was observed as a part of the MOJAVE survey programme with the VLBA at 15\\,GHz on February 5, 2007. Large amplitude variability in the correlated flux density, unexplainable in terms of the source structure, was first discovered while processing the data, and later confirmed by calibrating the antenna gains using 24 other sources observed in the experiment.}{The source shows variations in the correlated flux density as high as 40\\% on a timescale of only 2.7 hours. This places 1156+295 between the classical IDV sources and the so-called intra-hour variables. The observed variability timescale and the modulation index of 13\\% are consistent with interstellar scintillation by a nearby, highly turbulent scattering screen. The large modulation index at 15\\,GHz implies a scattering measure that is atypically high for a high galactic latitude source such as 1156+295.}{} ",
        "introduction": "Intraday variability (IDV) of compact, radio-loud AGN at cm-wavelengths was first discovered in the mid-1980s \\citep{wit86,hee87} and both source-intrinsic and extrinsic mechanisms for IDV have been vigorously studied \\citep[for a review, see e.g.][]{wag95}. Detection of time delays in the variability pattern arrival times between widely separated telescopes, as well as observed annual modulation of the variability timescale, have shown conclusively that interstellar scintillation (ISS) is the cause of intra-hour flux density variations observed in the three most extreme IDV sources \\object{PKS\\,0405-385} \\citep{jau00}, \\object{J1819+3845} \\citep{den02,den03}, and \\object{PKS\\,1257-326} \\citep{big03,big06}. Evidence for the ISS origin of IDV was found also in the cases of \\object{0917+624} \\citep{ric01,jau01,fur02}, \\object{PKS\\,1519-273} \\citep{jau03}, and \\object{J1128+5925} \\citep{gab07}. This provided the possibility to study both small-scale spatial fluctuations in the interstellar medium and the structure of compact radio sources on the microarcsecond scale. Unfortunately, the extreme IDV sources, which are the most suitable for these studies, are rare. By ``extreme'', we mean sources showing variability on timescales of a few hours or less and with an rms amplitude of modulation of over 10\\%. In this Letter, we report the serendipitous discovery of very large amplitude IDV in quasar \\object{1156+295} with a timescale of variations shorter than 3\\,h. This discovery is also unusual because it was achieved by a VLBI experiment.  \\begin{figure} \\centering \\resizebox{0.75\\hsize}{!}{\\includegraphics{0423fig1.eps}} \\caption{Naturally weighted 15\\,GHz VLBA image of \\object{1156+295} observed on February 5, 2007. The map peak brightness is 1.46\\,Jy\\,beam$^{-1}$. The contours begin at 0.7\\,mJy\\,beam$^{-1}$ and increase in steps of 2. Since the source flux density varied significantly during the observation, an amplitude self-calibration had to be used early in the imaging process. This may affect the image quality.} \\label{map} \\end{figure}  \\begin{figure*} \\centering \\resizebox{0.9\\hsize}{!}{\\includegraphics{0423fig2.eps}} \\caption{{\\it Top:} Amplitude self-calibration solutions of three example antennas (Brewster, Mauna Kea, and Pie Town) for the VLBA experiment on February 5, 2007. There is one solution per scan, and IFs are kept separate. Low elevation scans ($<15$\\degr) are excluded from the figure. Red symbols correspond to scans of \\object{1156+295} and blue symbols correspond to all other sources. The obvious discrepancy between the solutions for \\object{1156+295} and the other sources is present at all ten antennas.  {\\it Bottom:} Examples of the correlated flux density curves of \\object{1156+295} at three baselines after calibration of antenna gains by amplitude correction factors interpolated from the nearby scans of other sources. Time is in UT.} \\label{calib} \\end{figure*}  \\object{1156+295} (\\object{4C\\,+29.45}) is an optically violently variable quasar at $z=0.729$ \\citep{bur68}, which has Galactic coordinates $l=199.4$\\degr, $b=+78.4$\\degr. The source is strongly variable throughout the electromagnetic spectrum from radio to $\\gamma$-rays.  At radio frequencies, \\object{1156+295} shows significant long-term variability on a timescale of months \\citep[e.g.][]{kov02}. \\citet{wil83} reported a range of at least 5 magnitudes in the optical brightness, and in gamma-rays \\citet{sre96} detected an order of magnitude variability using the EGRET instrument onboard the Compton Gamma Ray Observatory.  During high radio brightness states, \\object{1156+295} appears to be highly core-dominated in VLBI images (95\\% at some epochs \\citep{kov05}; see also Fig.~\\ref{map}). Superluminal motion at speeds ranging from $3.5 h^{-1} c$ to $14.1 h^{-1} c$ were reported by e.g. \\citet{pin97}, \\citet{jor01}, \\citet{hon04}, and \\citet{kel04}.  In the optical bands, \\object{1156+295} shows rapid variations on timescales of between less than 30 minutes and 3 days \\citep{wil83,rai98}. The first IDV detection of \\object{1156+295} at radio frequencies was reported by \\citet{lov03}, who observed the source with the VLA at 5\\,GHz in January 2002. Their study measured rms flux density variations of 167\\,mJy during 3 nights of observations and a variability timescale of $\\sim24$\\,h. The relative amplitude of variations observed by \\citet{lov03} was far smaller than reported here.  Throughout this paper, we use the following cosmological parameters: $H_0$ = 73 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.27, and $\\Omega _\\Lambda$ = 0.73. The angular scaling conversion for $z=0.729$ is 7.07\\,pc mas$^{-1}$. ",
        "conclusions": "If we assumed a source-intrinsic origin of the observed variability in \\object{1156+295}, it would imply, by light travel time arguments, a brightness temperature of $\\gtrsim 2 \\times 10^{19}$\\,K, which is far in excess of the inverse Compton (IC) limit of $10^{12}$\\,K \\citep{kel69}. A Doppler factor higher than $\\sim 270$ would be then required to avoid the IC catastrophe. This would cause severe problems with the energy requirements in the source \\citep{beg94}. Extremely fast jets are also unlikely in the light of VLBI monitoring surveys, which indicate that the maximum jet Lorentz factor $\\Gamma$ is $30-40$ \\citep{coh07}. The standard model of incoherent synchrotron radiation from relativistic electrons in the jet would therefore have significant difficulties in explaining the observed variability, if it were intrinsic. For this reason, we concentrate on two possible extrinsic causes: interstellar scintillation and an extreme scattering event.  Propagation effects in the ionised medium of the Milky Way will cause a sufficiently compact source to scintillate \\citep[for review of ISS, see e.g.][]{ric90,goo97}. Since the observed $m$ in \\object{1156+295} is clearly below 100\\%, it is reasonable to assume that diffractive scintillation is fully quenched and to consider only refractive ISS. The observed $t_\\mathrm{var}$ and $m$ constrain the properties of a possible scattering screen and the size of the scintillating source. Because of the high $m$ in our case, it appears likely that our observing frequency is close to the critical frequency $\\nu_\\mathrm{s}$ between strong and weak scattering regimes. \\citet{goo97} provided a formula for $m$ that can be used to interpolate solutions close to $\\nu_\\mathrm{s}$, where the asymptotic solutions for strong and weak regimes break down. Combining his Eqs.  10, 12, 18, 19, and 20, and assuming a Gaussian brightness profile for the source, we calculated the screen distance, $D_\\mathrm{scr}$, and the FWHM size of the scintillating source $\\theta_\\mathrm{s}^{FWHM}$ as a function of scattering measure \\mbox{{\\it SM} $=\\int_0^d C_N^2 dx$}, where $C_N^2$ is the strength of the electron-density fluctuations\\footnote{Note that $\\theta_\\mathrm{s}$ in \\citet{goo97} is approximately $\\theta_\\mathrm{s}^{FWHM} / 2.35$.} (Fig.~\\ref{model}). The results were parameterised by the screen velocity. We estimated the uncertainties in $D_\\mathrm{scr}$ and $\\theta_\\mathrm{s}^{FWHM}$ by Monte Carlo methods and the hatched regions in Fig.~\\ref{model} show the 1$\\sigma$ error for these quantities. In addition, the results given by the interpolation formula of \\citet{goo97} were confirmed by numerically integrating the intensity covariance function for refractive scintillation \\citep[see Appendix A in][]{ric06}. In the bottom panel of Fig.~\\ref{model}, we indicate an approximate lower limit to the source size, $\\theta_\\mathrm{s}^{FWHM} \\ge 17$\\,$\\mu$as, given by the IC catastrophe limit and by assuming an upper limit of 50 for the Doppler factor \\citep{lah99,coh07}. As can be seen from Fig.~\\ref{model}, if the variability is due to ISS, we have either a rather nearby screen ($D_\\mathrm{scr} \\lesssim 300$\\,pc) with {\\it SM} $\\gtrsim 0.5$\\,m$^{-20/3}$\\,pc \\emph{or} $\\theta_\\mathrm{s}^{FWHM}$ so small that it would require Doppler factor exceeding 50.  We can estimate $\\theta_\\mathrm{s}^{FWHM}$ independently by assuming that there is approximate equipartition between the energy densities of the magnetic field and the radiating electrons \\citep{sco77}. We assume the synchrotron peak frequency to be 15\\,GHz and the corresponding peak flux density to be 1.5\\,Jy. This is based on the assumption that the scintillating component is the core of a \\citet{bla79} type jet and corresponds to the $\\tau=1$ surface at our observing frequency. The resulting equipartition size is $\\theta_s^{FWHM} \\sim 270 \\cdot \\delta^{-7/17}$\\,$\\mu$as, where $\\delta$ is the Doppler factor. If we again take $\\delta \\le 50$, this results in $\\theta_\\mathrm{s}^{FWHM} \\gtrsim 55$\\,$\\mu$as, {\\it SM}\\,$\\gtrsim 4$\\,m$^{-20/3}$\\,pc, and $D_\\mathrm{scr} \\lesssim 100$\\,pc for screen velocities between 10 and 50\\,km\\,s$^{-1}$. Since the Galactic electron distribution model by \\citet{cor02} predicts a scattering measure of only 0.1\\,m$^{-20/3}$\\,pc for the line-of-sight towards \\object{1156+295}, our results indicate a nearby, localised region of highly turbulent ionised gas in that direction.  The frequency-dependence of $m$ is slightly surprising: one would not expect $m$ to increase from 5.8\\% at 5\\,GHz \\citep{lov03} to 13\\% at 15\\,GHz, unless the scattering is strong \\emph{and} the source is smaller than the scattering angle. In that case the scintillating component would contain only about $30-50$\\% of the total flux density of the core, because otherwise the observed $m$ would be significantly higher. Another possibility is that \\object{1156+295} was in a more compact stage during our observation than it was 5 years earlier.  Simultaneous multi-frequency measurements of $m$ are required to clarify this.  \\begin{figure} \\centering \\resizebox{1.00\\hsize}{!}{\\includegraphics{0423fig5.eps}} \\caption{Parameters of the possible scattering screen. {\\it Top:} The distance to the screen as a function of the scattering measure {\\it SM}. Different screen velocities are shown in different colours. {\\it Bottom:} The scintillating source size as a function of {\\it SM}. The right edge of the panel gives the Doppler factor if $\\theta_\\mathrm{s}^{FWHM}$ is equal to equipartition size. The grey shaded area corresponds to $\\theta_\\mathrm{s}^{FWHM} < 17$\\,$\\mu$as. In both panels, the hatched region around the solid lines gives 1$\\sigma$ error for the plotted quantities. The +1$\\sigma$ lines in the $\\theta_\\mathrm{s}^{FWHM}$ plot roughly correspond to -1$\\sigma$ lines in the $D_\\mathrm{scr}$ plot and vice versa.} \\label{model} \\end{figure}  It is also possible that the observed variations are not due to standard ISS, but instead correspond to an isolated extreme scattering event \\citep[ESE;][]{fie87}. Unfortunately, our short, single-frequency observation does not provide us with sufficient information to distinguish between these two cases. However, the shape of the flux density curve resembles an ESE with two maxima symmetrically surrounding a deep minimum. Although ``classical'' ESEs have durations of between weeks and months, \\citet{cim02} reported an event in \\object{0954+658} with a duration of less than 2 days. The extremely short timescale of our event would, in the case of an ESE, imply a cloud size $<0.01$\\,AU."
    },
    "0809/0809.5104_arXiv.txt": {
        "abstract": "We present the SpineWeb framework for the topological analysis of the Cosmic Web and the identification of its walls, filaments and cluster nodes. Based on the {\\it watershed segmentation} of the cosmic density field, the \\textit{SpineWeb }method invokes the local adjacency properties of the boundaries between the watershed basins to trace the critical points in the density field and the separatrices defined by them. The separatrices are classified into \\textit{walls} and the {\\it spine}, the network of filaments and nodes in the matter distribution. Testing the method with a heuristic Voronoi model yields outstanding results. Following the discussion of the test results, we apply the {\\it SpineWeb} method to a set of cosmological N-body simulations. The latter illustrates the potential for studying the structure and dynamics of the Cosmic Web. ",
        "introduction": "The large scale distribution of matter revealed by galaxy surveys features a complex network of interconnected filamentary galaxy associations. This network, which has become known as the {\\it Cosmic Web} \\citep{Bond96}, contains structures from a few megaparsecs up to tens and even hundreds of Megaparsecs of size. The weblike spatial arrangement of galaxies and mass into elongated filaments, sheetlike walls and dense compact clusters, the existence of large near-empty void regions and the hierarchical nature of this mass distribution -- marked by substructure over a wide range of scales and densities -- are its three major characteristics. Its appearance has been most dramatically illustrated by the recently produced maps of the nearby cosmos, the 2dFGRS, the SDSS and the 2MASS redshift surveys \\citep[e.g.][]{Colless03,Huchra05}.  In this paper we introduce the SpineWeb formalism for analyzing the structure and topology of the Cosmic Web. It identifies the sheets and filaments in the Cosmic Web, along with the large underdense void regions, and their mutual connection into the {\\it Spine} of the cosmic matter distribution. The method is based on the {\\it Watershed Transform} \\citep[WST,][]{Beucher82}, and is largely free of user-specific parameters and artificial smoothing scale(s). Its output will enable the study of the physical properties and dynamics of the individual morphological components, along with their topology and hierarchical characteristics.  \\subsection{the Cosmic Web} The Cosmic Web is the most salient manifestation of the anisotropic nature of gravitational collapse, the motor behind the formation of structure in the cosmos \\citep{Peebles80}. N-body computer simulations have profusely illustrated how a primordial field of tiny Gaussian density perturbations transforms into a pronounced and intricate filigree of filamentary features, dented by dense compact clumps at the nodes of the network \\citep[][]{Colberg05,Springel05}. The filaments connect into the cluster nodes and act as the transport channels along which matter flows into the clusters.  Fundamental understanding of anisotropic collapse on cosmological scales came with the seminal study by \\citet{Zeldovich70}, who recognized the key role of the large scale tidal force field in shaping the Cosmic Web \\citep[also see][]{Icke73}. The collapse of a primordial cloud (dark) matter passes through successive stages, first assuming a flattened \\textit{sheetlike} configuration as it collapses along its shortest axis. This is followed by a rapid evolution towards an elongated \\textit{filament} as the intermediate axis collapses and, if collapse continues along the longest axis, may ultimately produce a dense, compact and virialized \\textit{cluster} or \\textit{halo}. The hierarchical setting of these processes, occurring simultaneously over a wide range of scales and modulated by the expansion of the Universe, complicates the picture considerably. Recent state-of-the-art computer experiments like the Millennium simulation \\citep[][]{Springel05} clearly show the hierarchical nature in which not only the clusters build up but also the filamentary network itself \\citep[see][]{Aragon07b}.  The Cosmic Web theory of \\cite{Bond96} succeeded in synthesizing all relevant aspects into a coherent dynamical and evolutionary framework. It is based on the realization that the outline of the cosmic web may already be recognized in the primordial density field. The statistics of the primordial tidal field explains why the large scale universe looks predominantly filamentary and why in overdense regions sheetlike membranes are only marginal features \\citep{pogosyan98}. Of key importance is the observation that the rare high peaks, which will eventually emerge as clusters, are the dominant agents for generating the large scale tidal force field: it is the clusters which weave the cosmic tapestry of filaments \\citep[][]{Bond96,WeyEdb96,WeyBond08a}. They cement the structural relations between the components of the Cosmic Web and themselves form the junctions at which filaments tie up. This relates the strength and prominence of the filamentary bridges to the proximity, mass, shape and mutual orientation of the generating cluster peaks: the strongest bridges are those between the richest clusters that stand closely together and point into each other's direction.  The emerging picture is one of a primordially and hierarchically defined network whose weblike topology is imprinted over a wide spectrum of scales. Weblike patterns on ever larger scales get to dominate the density field as cosmic evolution proceeds, and as small scale structures merge into larger ones. Within the gradually emptying void regions, however, the topological outline of the early weblike patterns remains largely visible.  \\subsection{Closing in on the Cosmic Web} Despite a large variety of attempts, as yet no generally accepted descriptive framework has emerged for the objective and quantitative analysis of the geometry and topology of the Cosmic Web. The great complexity of both the individual structures as well as their connectivity, the lack of structural symmetries, its intrinsic multiscale nature and the wide range of densities that one finds in the cosmic matter distribution has prevented the use of simple and straightforward techniques.  Historically, the quantitative analysis of the Cosmic Web has been dominated by a description in terms of statistical measures of clustering of galaxies and matter. While correlation functions have been the mainstay of the cosmological analysis of large scale structure, a direct interpretation in terms of the patterns and texture of the Cosmic web has largely remained elusive. Over the years a variety of heuristic measures have been forwarded to analyze specific aspects of the spatial patterns in the large scale Universe, but only in recent years there have been attempts towards developing complete descriptors of the intricate spatial patterns that define the Cosmic Web. Nearly without exception these methods borrow extensively from other branches of science such as image processing, mathematical morphology, computational geometry and medical imaging.  \\begin{figure*}[!ht] \\centering \\includegraphics[width=0.49\\textwidth,angle=0.0]{fig_01.eps} \\includegraphics[width=0.49\\textwidth,angle=0.0]{fig_02.eps} \\includegraphics[width=0.49\\textwidth,angle=0.0]{fig_03.eps} \\includegraphics[width=0.49\\textwidth,angle=0.0]{fig_04.eps} \\caption{Top left: a slice of 100 $h^{-1}$ Mpc of side showing the density field computed from a N-body simulation. Top right: the same slice as the left panel but here showing the density field as a shaded landscape where high density regions correspond to ridges while underdense regions correspond to valleys. The light source used to shade the surface is located at the north-east. Bottom left: the 2D watershed transform computed from the 2D density field using a discrete-intensity  bucket algorithm (black contours). Each individual valley is randomly colored for visualization purposes. Bottom right: the same landscape as the top right panel but here we also show the watershed transform superimposed as thick black lines. The watershed contour has been slightly smoothed for visualization purposes.} \\label{fig:density_landscape} \\vskip 0.35truecm \\end{figure*}  Noteworthy examples include filament detection with the help of the Candy model \\citep{Stoica05} and wavelet analysis of the Cosmic Web \\citep{Martinez05}. Several methods seek to relate morphological features to singularities in the density field, usually invoking information on the gradient and Hessian of the density field, or of the tidal field \\citep[see e.g.][]{Sousbie08,Aragon07a,Aragon07b,Hahn07a,Hahn07b,Bond09}. A classification scheme on the basis of the manifolds in the tidal field -- involving all morphological features in the cosmic matter distribution -- has been presented by \\cite{Hahn07a,Hahn07b,Forero08}. However, its success may depend strongly on the correct choice of the smoothing scale. Another concept addressing the gradient and Hessian of the density field is that of the {\\it skeleton analysis}, a direct application of Morse theory to cosmological density fields \\citep[see][]{Colombi2000,Pogosyan09}. The skeleton formalism has been developed for the morphological analysis of the Megaparsec Cosmic Web, in redshift surveys like SDSS as well as in N-body simulations \\citep{Novikov06,Sousbie08,Sousbie08b,Sousbie09}. Its present implementation refers to features identified at one single specific scale and suffers from Gaussian smoothing. The multiscale nature of the cosmic matter distribution is explicitly addressed by the Multiscale Morphology Filter, which is based on a scale-space analysis of the Hessian of the density field \\cite{Aragon07a,Aragon07b} to identify cluster, filaments and sheets on the scale where they are locally most prominent.  \\subsection{Watershed and Cosmic Spine} One technique that implicitly addresses the topology of the Cosmic Web is the Watershed Void Finder (WVF) developed by \\cite{Platen07}. The WVF is an application of the {\\it Watershed Transform} for the identification of underdense basins in the megaparsec-scale matter distribution. The Watershed transforms segments the density field into isolated basins and delineates the boundaries of cosmological voids.  The method presented in this paper is the natural extension of the Watershed Void Finder. It includes the watershed transform into a wider context as a framework for studying both the morphology and topology of the cosmic web and its various constituents. The result is the \\textit{SpineWeb} method, a complete framework for the identification of voids, walls and filaments. Via the practical role of the watershed transform in computing the Morse complex it is intimately related to Morse theory, in which it finds its mathematical foundation. It is important to note that although Morse theory can be used to describe the same topology traced by the SpineWeb method there is not a univocal relation between the two as the SpineWeb is based on the observed properties of the Cosmic Web.  An important aspect of our method is that it is an intrinsically scale-free method, starting from a scale-free reconstruction of the density field. We use the DTFE method of \\cite{Schaap00}, which guarantees an optimal and unbiased representation of the hierarchical nature and anisotropic morphology of cosmic structure \\citep[see][for an extensive description]{WeySchaap09}. Having guaranteed the capability of invoking a full scale-free {\\it Scale-space} representation of cosmic structure, our watershed procedure not only traces the outline of filaments and sheets, but may also be extended towards doing so over a range of scales in order to address their hierarchical structure.  \\subsection{Outline} The principal rationale behind the SpineWeb analysis of the cosmic matter distribution is the interest in relating the geometry of the matter and galaxy distribution in a more meaningful fashion to the underlying dynamical evolution. One particular aspect of this dynamically motivated disentanglement of structure is the attempt to identify various evolutionary stages of the tidally induced anisotropic collapse of structure in the Universe.  In this paper we will focus specifically on the description of the basic SpineWeb formalism, confined to a density field sampled on a regular grid. We start by discussing the topological background of our study in section~\\ref{sec:watershed}, focusing on the watershed transform, its connection to the general context of Morse theory and the related issues of practical interest to the SpineWeb formalism. The overall cosmological background of the structure, formation and dynamics of the Cosmic Web follows in section~\\ref{sec:cosmicweb}, amongst others to establish the link between the structures identified by the SpineWeb technique and the filamentary identity of tidal bridges in the theory of the Cosmic Web \\citep{Zeldovich70,Bond96}. The technical aspects of the SpineWeb formalism are outlined in some detail in section~\\ref{sec:spineweb}. Subsequently, the formalism is tested by applying it to two different classes of spatial particle distributions. The first testbed concerns two simple heuristic Voronoi clustering models which model aspects of cellular and/or weblike spatial distributions. Visual and quantitative tests are described in section~\\ref{sec:voronoi}. The operation of SpineWeb in a more realistic setting of a $\\Lambda$CDM simulation is the subject of section~\\ref{sec:lcdm}. In this section we also stress the fundamental differences between a structural selection based on density thresholds or one based on topological criteria. To illustrate the potential for analyzing cosmological structures, in section~\\ref{sec:analysis} we shortly describe three quantitative measures for the matter distribution in the $\\Lambda$CDM simulation we used for testing. Finally, section~\\ref{sec:conclusions} summarizes our results, and discusses prospects and further developments of the SpineWeb formalism.    \\bigskip ",
        "conclusions": "\\label{sec:conclusions} The Spine of the Cosmic Web is the cosmic web's framework, consisting of the network of filaments and walls and their connections at the clusters nodes. In this study we present a topological technique based on the discrete watershed transform of the cosmic density field for the identification and characterization of voids, walls and filaments. Our method is closely related to a variety of concepts from computational topology, and has a strong mathematical foundation in Morse theory of singularities.  The SpineWeb method is ideally suited for morphological and dynamical studies of the Large Scale Structure. Amongst others, it will allow a better insight into the formation and dynamics of the anisotropic filamentary and wall-like structures in the Large Scale Universe. Another immediate application is in addressing the question whether and which influences the large scale environment has on the halos and galaxies that are forming within their realm.  As a first test of its viability, we applied our method on a set of heuristic Voronoi clustering models. The SpineWeb procedure succeeds in reconstructing the original properties of the cellular galaxy distribution. In the implementation presented in this work, we effectively restrict ourselves to a single spatial scale determined by the voxel scale of the regular grid on which the density field is sampled. In a forthcoming paper we will discuss the effect of the multiscale nature of the matter distribution. The scale-space formulation of the SpineWeb method will enable us to identify fainter features in the density field and establish their connections with other objects into a truly hierarchical weblike pattern. In other words, it provides an effective way towards characterizing the hierarchy of structures in the Cosmic Web.  A crucial aspect of the watershed transform and of our method is the definition of local neighborhood. In the case of regular grids the immediate neighborhood of 26 pixels is arguably the best option. However, for unmeshed data such as galaxy surveys and N-body simulations, other neighborhood definitions offer a better choice. Among these, the Voronoi contiguous cell defined by the Delaunay tessellation of the point distribution represents a promising option. In the third paper of this series we will present the result of a Delaunay implementation of the Spine method.\\\\"
    },
    "0809/0809.2038_arXiv.txt": {
        "abstract": "A recent high angular resolution extinction map toward the most opaque molecular globule, Globule 2, in the Coalsack Nebula revealed that it contains a strong central ring of dust column density. This ring represents a region of high density and pressure that is likely a transient and possibly turbulent structure. Dynamical models suggest that the ring has formed as a result of a sudden increase in external pressure which is driving a compression wave into the Globule. Here we combine the extinction measurements with a detailed study of the \\ceol\\, molecular line profiles toward Globule 2 in order to investigate the overall kinematics and, in doing so, test this dynamical model.  We find that the ring corresponds to an enhancement in the \\ceo\\, non-thermal velocity dispersion and non-thermal pressure. We observe a velocity gradient across the Globule that appears to trace two distinct systematic subsonic velocity flows that happen to converge within the ring. We suggest, therefore, that the ring has formed as two subsonic flows of turbulent gas merge within the Globule. The fact that the outer layers of the Globule appear stable against collapse yet there is no centrally condensed core, suggests that the Globule may be evolving from the outside in and has yet to stabilize, confirming its youth. ",
        "introduction": "Understanding the initial conditions that lead to the formation of dense cores is important as it has direct implications for theories of star formation.  Recent studies of starless cores reveal that the kinematics in the early stages of core collapse can be quite complex: in some cases the molecular tracers indicate simultaneously opposing motions (expansion {\\it {and}} infall; \\citealp{Lada03a}). In the case of the starless Bok globule B68, the presence of both inward and outward motions, traced by molecular line emission, has been interpreted to arise due to the oscillation of the outer layers of the cloud \\citep{Lada03a,Redman06,Maret07}. This idea is strengthened by modeling that suggests that the complex motions within starless cores can occur due to oscillating perturbations about a thermally-supported, pressure bounded structure \\citep{Keto06,Broderick07}. Indeed, recent numerical simulations reveal that small amounts of rotation can cause initially unstable and collapsing cores to stabilize and oscillate \\citep{Matsumoto03}.  The combination of high angular resolution extinction and molecular line observations toward cold, dense starless cores can trace the properties and kinematics of the dense gas and allow us to gain insight into the physical nature of core and, thus, star formation. Extinction measurements are particularly useful for tracing the internal structure of cores because they probe the material along the line-of-sight and are not susceptible to optical depth effects (as in the case of \\co\\, molecular line emission) or assumptions about the dust temperature or emissivity (as in the case of millimeter continuum emission). Thus, the combination of dust extinction observations with optically thin molecular line observations can accurately reveal the masses and densities of the dust as well as the kinematics and dynamics of the gas.  In this paper we combine extinction observations with a detailed kinematic study of the molecular line profiles observed toward Tapia's Globule 2 in the Coalsack Nebula. Because Globule 2 does not contain a central protostar, is one of the densest cores within the Coalsack complex, and is thought to be in a very early stage of condensing, a detailed study of its kinematic structure could provide important details on the processes leading to core formation.  The Coalsack Nebula is a nearby dark molecular cloud complex that extends almost 6\\degree\\, on the sky ($\\sim$ 15 pc at its distance of 150 pc; \\citealp{Cambresy99}) and appears completely devoid of star formation activity \\citep{Nyman89,Bourke95,Kato99}. It has a total mass of $\\sim$ 3550\\,\\Msun\\, and is characterized by complex, filamentary molecular structures which contain many dark cores with kinetic temperatures, \\tk, of 8--11 K \\citep{Nyman89}. The darkest and densest of these cores are Tapia's Globules 1, 2, and 3 \\citep{Tapia73,Bok77}.  Using deep near-infrared imaging observations obtained with the European Southern Observatory (ESO) 3.6m New Technology Telescope (NTT), \\cite{Lada04} were able to examine the structure of Globule 2 in significantly greater detail than previously possible. The NICE method \\citep{Lada94,Alves98} was used to generate a visual extinction (\\av) map that was constructed by measuring the individual infrared extinctions to approximately 24,000 stars located behind Globule 2.  This resulted in an extinction map across the core with significantly improved angular resolution ($\\sim$ 15\\arcsec) compared to previous maps ($\\sim$ 60\\,\\arcsec; \\citealp{Jones80,Jones84,Racca02}).  The extinction map revealed Globule 2 to have a mass of $\\sim$6.1\\,\\Msun, a mean density of $\\sim$3$\\times$10$^{3}$\\,\\cmc\\, and Jeans mass, \\Mj, of $\\sim$10.4\\,\\Msun. From near-infrared polarimetry observations, \\cite{Jones84} suggested that Globule 2 is supported either by turbulence or a magnetic field and appears stable. This is also consistent with the fact that the derived mass is less than \\Mj\\, suggesting it is unlikely to fragment further. Moreover, a Bonnor-Ebert (BE) fit to the outer regions of the Globule's radial extinction profile suggests that it is globally stable against collapse \\citep{Lada04}. The Globule's derived BE mass is $\\sim$ 6\\,\\Msun. Even though this is comparable to the derived mass calculated from the dust, the high angular resolution infrared extinction map toward the Globule does not reveal a centrally condensed core as would be expected if it was a true BE sphere. Instead, the most prominent feature within the extinction map is a strong central ring of dust column density. This ring likely represents a transient structure that has formed from a region of high density and pressure \\citep{Lada04}. Because there is no source of internal pressure to support the ring against gravitational collapse, it is thought that the ring is not in dynamical equilibrium with its surroundings and, given the short timescales for inward collapse (2$\\times$10$^5$ years), most likely represents an extremely early stage of evolution.  From a single \\ceoh\\, spectrum toward the center of the globule, \\cite{Lada04} noted a double-peaked profile comprising a bright narrow line and a fainter broad line. Based on the absence of this double-lined profile in lower angular resolution spectra \\citep[e.g.,][]{Brooks76,Nyman89,Kato99}, it was suggested that the narrow line component most likely arises within a small region and that the broad line component traces a more spatially extended feature \\citep{Lada04}. The line width of each blended component was found to be super-thermal indicating that significant non-thermal, or turbulent, motions most likely dominate.  Dynamical models for Globule 2, however, suggest that the ring structure has formed as a result of a sudden increase in external pressure which is driving a compression wave into the Globule \\citep{Hennebelle06}.  Using the Mopra 22-m telescope, we have obtained the first complete high angular and spectral resolution molecular line map of Globule 2. Because \\co\\, is optically thick toward most dense cores, and is thus only tracing the outer layers, we choose to conduct this study using the optically thin \\ceol\\, transition which is far superior in tracing the dense gas associated with the central core.  In combination with the extinction map, these data should allow us to test the various formation mechanisms for the ring feature; is it a high-density transient structure formed via the dissipation of turbulent gas \\citep{Lada04} or does it trace the propagation of a compression wave through the Globule \\citep{Hennebelle06}?  Moreover, these data will enable the relative spatial distributions of the two molecular line components to be definitively traced.  In particular, does the narrow-line component originate exclusively in the dense ring?  Does the broad-line component primarily trace the turbulent envelope of the Globule?  Answers to these questions will directly test the idea that Globule 2 is in perhaps the earliest stage of evolution: a dense core in the process of formation. ",
        "conclusions": "Globule 2 is one of the densest cores within the Coalsack Nebula, but contains no evidence for star-formation. Despite the fact that overall it appears stable against collapse, a recent high angular resolution extinction map reveals Globule 2 contains a central ring of enhanced column density and, thus, may in fact be a dynamically unstable, transient core on the verge of condensing to form a protostar.  It has been speculated that this central ring of column density has formed either as a result of turbulence or from a shock induced via an external compression wave \\citep{Lada04,Hennebelle06}.   Here we have combined dust extinction measurements with a high angular resolution \\ceo\\, molecular line map in order to study the detailed kinematics of Globule 2.  The \\ceo\\, map reveals that the Globule is marginally gravitationally bound and not virialized.  Moreover, the \\ceo\\, map reveals a complex velocity profile that appears to arise from two distinct velocity features that are present both toward the ring and external to it. In addition, these velocity features appear to trace material on opposing sides of the extinction ring.  Each of these velocity flows have non-thermal subsonic velocity dispersions.  Outside the extinction ring the Globule is thermally dominated, however, both the measured velocity dispersion and non-thermal pressure increase toward the extinction ring. Of all the measured properties, we find the best spatial correspondence between the one-dimensional velocity dispersion and the visual extinction.  We suggest, therefore, that the extinction ring is an enhancement in column density due to the merging of the two subsonic flows. Thus, it appears as though we are seeing the convergence of two subsonic flows within a bound object. Moreover, because the outer regions of the Globule appear stable against collapse and there is no centrally condensed core, it may be that the Globule is evolving from the outside in and has yet to stabilize.  Finally, although we cannot entirely rule out recent dynamical models for the formation of the extinction ring via an external compression wave \\citep{Hennebelle06}, these models are unable to completely characterize the observed \\ceo\\, kinematics. The discrepancies between the model and the data may arise because of the simplified assumption of a spherical compression wave."
    },
    "0809/0809.3722_arXiv.txt": {
        "abstract": "Soft phase lags, in which X-ray pulses in lower energy bands arrive later than pulses in higher energy bands, have been observed in nearly all accretion-powered millisecond pulsars, but their origin remains an open question.  In a study of the 2.5~ms accretion-powered pulsar \\saxj, we report that the magnitude of these lags is strongly dependent on the accretion rate. During the brightest stage of the outbursts from this source, the lags increase in magnitude as the accretion rate drops; when the outbursts enter their dimmer flaring-tail stage, the relationship reverses.  We evaluate this complex dependence in the context of two theoretical models for the lags, one relying on the scattering of photons by the accretion disk and the other invoking a two-component model for the photon emission.  In both cases, the turnover suggests that we are observing the source transitioning into the ``propeller'' accretion regime. ",
        "introduction": "Accretion-powered millisecond pulsars (AMPs) provide a unique laboratory for studying the process of disk accretion onto magnetic stars (\\citealt{Psaltis99} and refrences therein).  The $\\sim$$10^8$~G stellar magnetic field truncates the accretion disk only a few neutron star radii above the surface, channeling the infalling matter toward the magnetic poles. The resulting X-ray--emitting hot spots and accretion shocks are modulated by the stellar rotation to produce the observed pulsations.  Such compact accretion geometry poses a particularly interesting challenge.  We want to understand the details of the interaction between the magnetic field and the disk, and the effects that result from the proximity of different emitting regions.  The accretion-powered 2.5~ms pulsar \\saxj\\ is a particularly good system for this type of study: in addition to having multiple outbursts, it spans orders of magnitude in accretion rate.  In this paper we report our study of the energy-dependence of pulsations from this source and show that it gives new insight into these problems.  \\saxj\\ was the first detected AMP \\citep{Wijnands98}, and it remains the best studied.  Persistent pulsations have been detected throughout the four $\\sim$1 month outbursts from this source that have been observed by the {\\em Rossi X-ray Timing Explorer} (\\RXTE) during 1998--2005.  The timing of these pulsations establishes a 2.01~hr binary orbital period \\citep{Chakrabarty98} that is increasing on a timescale of $P_{\\rm{orb}}/\\dot{P}_{\\rm{orb}} = (66\\pm4)\\times10^6$~yr, possibly due to the ablation of its low-mass companion during X-ray quiescence (\\citealt{Hartman08}, hereafter \\SpinPaper; \\citealt{DiSalvo08}).  The presence of X-ray pulsations during the peak accretion rate and the long-term spin-down of this source provide a tight constraint on its magnetic field: $B = (0.4$--$1.5)\\times10^8$~G (\\SpinPaper; \\citealt{Psaltis99}).  A distance of 3.4--3.6~kpc is estimated from thermonuclear X-ray burst properties \\citep{Galloway06}.  The outbursts from \\saxj\\ had similar light curves, which we divide for convenience into four stages (see Fig.~2 of \\SpinPaper\\ for a schematic): a rise from quiescence lasting $\\approx$5~days; a short-lived peak at a 2--25~keV flux of $(1.9$--$2.6)\\times10^{-9}$\\fluxunits; a 10--15~day slow decay in luminosity that lasts until the source reaches $\\approx$$0.7\\times10^{-9}$\\fluxunits; and finally a rapid fall in luminosity followed by weeks of low-luminosity flares with a $\\sim$5~d periodicity, during which the observed flux can vary by a factor of 1000 \\citep{Wijnands01}.  This sudden drop has been interpreted as a transition into a ``propeller'' accretion regime \\citep[e.g.,][]{Gilfanov98} in which accretion is only partially inhibited \\citep{Spruit93, Rappaport04}, although it may also be caused by a change in the viscosity of the accretion disk as it transitions into quiescence \\citep{Gilfanov98, Gierlinski02}.  Shortly after the discovery of pulsations in \\saxj, \\citet{Cui98} noted the presence of soft phase lags, the tendency of the pulses at lower (``softer'') energies to arrive later than the corresponding pulses at higher (``harder'') energies.  The magnitude of the lags increased approximately linearly between 2--10~keV and saturated at 0.08 rotational cycles (200~\\us) above 10~keV. Similar soft lags have been detected in six other AMPs:  IGR~J00291$+$5934 \\citep{Galloway05}, XTE~J0929$-$314 \\citep{Galloway02}, XTE~J1751$-$305 \\citep{Gierlinski05}, XTE~J1807$-$294 \\citep{Kirsch04}, XTE~J1814$-$338 \\citep{Watts06}, and HETE~J1900.1$-$2455 \\citep{Galloway07}.  A number of explanations for the AMP soft lags have been proposed.  Models that include only Doppler boosting, due to the stellar rotation \\citep{Ford00,Weinberg01} or the orbiting material in the accretion disk \\citep{Sazonov01}, have shortcomings that are summarized by \\citet{Gierlinski02}.  The two most plausible explanations invoke a two-component emission model with differing angular dependence and spectra (\\citealt{Poutanen03}; hereafter PG03) or the scattering delay of hot photons from the accretion shock off the cooler disk or stellar surface (\\citealt{Falanga07}; hereafter FT07).  In this paper we present an energy-resolved analysis of the persistent pulsations for all the outbursts of \\saxj\\ detected by \\RXTE.  We will discuss the implications of our results for these models of the soft lags. ",
        "conclusions": "The energy dependence of the pulsations from \\saxj\\ exhibit strong variability both between and during the four outbursts observed by \\RXTE.  The overall fractional amplitudes and the shape of the amplitude vs. energy curves differ substantially between outbursts, with no obvious correlation with flux; the only clear trend is that the pulsations generally become less sinusoidal with energy in the observed 2--30~keV range.  In contrast, the soft lags of the pulses have a consistent energy dependence, increasing in magnitude in the 2--10~keV band and saturating at 10--30~keV, with a strong dependence on flux. At the highest fluxes the lags are near zero, but they increase steadily as the flux falls.  Below a 2--25~keV flux of $\\approx$$0.7\\times 10^{-9}$\\fluxunits, this trend reverses and lags get smaller again.  This particular flux also marks the point in the outburst where the luminosity decays rapidly and the source enters the flaring tail.  Soft lags are a common feature of AMP pulsations, and two plausible models have emerged to explain their existence.  PG03 suggested a two-component model involving soft emission from a surface hot spot and a hard Comptonized component from the shocked layer in the accretion funnel.  Each component has a different beaming pattern, which is affected in a different way by the Doppler effect.  This results in the soft pulse lagging the hard pulse.  While this model works well for \\saxj\\ and XTE~J1751$-$305 \\citep{Gierlinski05}, it does not predict the reduction in soft lags at very high energies seen in IGR~J00291$+$5934 \\citep{Falanga05}.  As a result FT07 proposed an alternative model in which the disk plays a prominent role.  In this model, the main pulsed element is the hard component from the accretion column: this radiation is then downscattered by cold plasma in the inner disk (or on the neutron star surface) to generate a soft lagging component.\\footnote{\\citet{Ibragimov09} recently argued that photoelectric absorption in the accretion disk would prevent the large number of scattering events required by the FT07 model. However, standard disk models give a temperature of $\\sim$300~eV in the inner disk \\citep{Shakura73}; the actual temperature will be higher due to the irradiation.  At these temperatures and the expected electron densities ($n_e \\sim 3\\times10^{19}$~cm$^{-3}$; see FT07 and Fig.~\\ref{fig:FalangaFitVsMDot} here), all elements through oxygen will easily be fully ionized by the hard tail of the thermal distribution.  The only significant source of photoelectric absorption is Fe, but for this $n_e$ only Fe~{\\sc xxv} and higher remain at $T \\ge 300$~eV, with the transition to full ionization occuring at 400--600~eV \\citep{Rubiano04}.  Its He-like and H-like forms have K edges at 9~keV, so they will not absorb the softer photons responsible for the lags.  Since the FT07 model is not ruled out {\\em a priori}, this discussion attempts to address both models without preference.}  These soft lags get larger as the density of the scattering plasma falls.  Let us now explore the factors that might cause the flux dependence of the soft lags, assuming that flux directly traces accretion rate.  We then need to ask how the properties of the system are expected to change as the accretion rate falls.  For a star with a magnetic field that is strong enough to channel the accretion flow, the reduction in accretion rate should cause an increase in inner disk radius (e.g., \\citealt{Ghosh77}).  This will have two relevant effects.  First, the density of scattering plasma near the star will fall: within the model of FT07 this would cause the soft lags to increase.  In addition the relative angular velocity of the inner disk and star will reduce. This will affect the angle at which the accretion funnel impacts the surface \\citep{Romanova04}, changing the column density above any surface hot spot. This type of change might work within the model of PG03 to affect the soft lags.  If these changes continue steadily, however, it is hard to see what could cause the turnover seen in Figure~\\ref{fig:PhaseLagSlopeVsFlux}.  One intriguing possibility that would give a very natural explanation for such a turnover is that the system reaches the point where the magnetospheric radius $r_m$ exceeds the corotation radius $r_c$.  This is often referred to as the ``propeller'' regime, since it was originally thought that matter would then be ejected from the system, with the star going from spin-up to spin-down \\citep{Illarionov75}.  In fact, the transition is probably not sharp, and $r_m$ needs to be substantially larger than $r_c$ for ejection to occur \\citep{Spruit93}.  At intermediate accretion rates, continued accretion is possible even while the neutron star starts to spin down \\citep{Spruit93, Rappaport04}.  Simulations by \\citet{Long05} indicate that the transition between spin up and spin down should occur when $r_m / r_c \\approx 0.7$ for AMPs.  The observed turnover happens at an accretion rate of $\\dot{M} = 1.8\\times10^{-10}\\ M_\\sun$~yr$^{-1}$, giving $r_m / r_c = 0.6$ for the magnetic field $B = 1.5\\times10^8$~G from \\SpinPaper.\\footnote{Accretion rates were estimated in the usual manner, by equating the bolometric and infall luminosities: $4\\pi d^2 f_x c_{\\rm{bol}} = GM\\dot{M}/R$, using $d = 3.5$~kpc and the bolometric correction $c_{\\rm{bol}} = 2.12$ from \\citet{Galloway06}. The Alfv\\'en radius was estimated as $r_m = k_A (2GM)^{-1/7} \\dot{M}^{-2/7} (BR^3)^{4/7}$, with $k_A = 0.5$ as per \\citet{Long05}.  $M = 1.4\\ M_\\sun$ and $R = 10$~km were used in all cases.}  This close agreement between the simulations and the $\\dot{M}$ at which the flux rapidly drops and the soft lags turnover strongly suggests that \\saxj\\ passes through this transition as the accretion rate falls.  In this eventuality there are two effects that could cause the turnover in soft lags.  \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.47\\textwidth]{FalangaFitVsMDot.eps} \\end{center} \\caption{ Electron density $n_e$ of the inner accretion disk, as fit using the model of FT07, versus the mass accretion rate for the 2002 and 2005 outbursts of \\saxj.  The vertical dotted line marks the $\\dot{M}$ at which the magnitude of the soft lags is greatest and below which the outburst is in its flaring-tail stage; a transition from normal accretion to the ``propeller'' regime can explain this change in the observed behavior. The dashed lines give the best fits of $n_e \\propto \\dot{M}^\\gamma$ in each region. \\label{fig:FalangaFitVsMDot}} \\end{figure}  The first is due to the fact that the relative angular velocity of the inner disk and the star changes sign.  The accretion funnel will go from being dragged forwards to being dragged backwards (assuming that it settles rapidly; \\citealt{Romanova04}).  The column depth above any surface hot spot should therefore go through a maximum near the zero torque point, causing a turnover in the scattering properties of the accretion column.  If the accretion column plays a major role as discussed in PG03, this could cause the observed turnover.  Another explanation within the PG03 model is the suppression of the lagged soft component at high fluxes due to the increased optical depth of the accretion shock when $\\dot{M}$ is large.  However, this mechanism would require the normalization of the soft component to decrease during the outburst peaks, contrary to what is seen (cf. Fig.~2 of \\citealt{Gierlinski02}).  A second possibility stems from the fact that the properties of the disk must also change at this point.  For accretion to continue after zero-torque point, the inner edge stops retreating, and the density at the inner edge must rise \\citep{Spruit93}.  In the scenario of FT07, this rise in inner disk density would cause the magnitude of the soft lags to decrease again.  In the absence of Compton upscattering, which would introduce hard lags not observed in \\saxj, their model predicts that the scattering from the disk will modify the arrival times as $\\Delta t(E) \\approx (m_e c / \\sigma_T n_e) (E_0^{-1} - E^{-1})$, where $\\sigma_T$ is the Thompson cross-section and $n_e$ is the electron density of the scattering region.  This model provides reasonable fits to the observed phase lags; Figure~\\ref{fig:FalangaFitVsMDot} shows the resulting $n_e$.  The fitted electron densities probe the relationship between the inner disk and the accretion rate.  $n_e \\sim 10^{19.5}$~cm$^{-3}$ gives a mean free path of $\\ell = (\\sigma_T n_e)^{-1} \\sim 0.5$~km and time $\\ell/c \\sim 2$~\\us, requiring $\\sim$100 scattering events to produce the observed soft lags.  This $\\ell$ rules out the thin neutron star atmosphere as a scattering medium.  On the other hand, simulations by \\citet{Romanova04} suggest an inner disk thickness of order the neutron star radius, and $(R / \\ell)^2 \\sim 400$ is more than adequate to accomodate the number of needed scattering events.  From the fits in Figure~\\ref{fig:FalangaFitVsMDot}, we find that $n_e \\propto \\dot{M}^{1.0\\pm0.1}$ during the main outburst; during the presumed propeller stage, $n_e \\propto \\dot{M}^{-0.4\\pm0.1}$.  An interesting consequence of the former relationship is that at higher accretion rates before the turnover, when $r_m / r_c \\lesssim 0.7$, the flow rate (the volume of material per unit time crossing the inner edge of the disk) must be constant: since $\\dot{M} = \\Omega m_p n_e r_m^2 v_r$, where $\\Omega$ is the solid angle of the inner edge and $v_r$ is the radial velocity of material crossing it, the flow rate $\\Omega r_m^2 v_r$ is fixed.  If the phase lags in \\saxj\\ are entirely due to Compton downscattering, as suggested by FT07, then this model must explain why the magnitudes of the soft lags for the second harmonic are roughly half the magnitudes for the fundamental over a wide range of $\\dot{M}$, as can be seen in Figure~\\ref{fig:PhaseLagSlopeVsFlux}.  Following the formalism of \\citet{Sunyaev80}, we model the timing effects of Comptonization with a distribution $P(\\Delta t)$ of the photon escape times from the scattering medium.  For the $k$th harmonic of the spin frequency $\\nu_s$, we convolve the harmonic component $x_k(t) = A_k \\exp(2\\pi i k\\nu_s t)$ with $P(\\Delta t)$ to find the new amplitude $A'_k$ and the time lag $\\Delta t$ of the scattered pulses: \\begin{equation} A_k'e^{2\\pi i k\\nu_s (t - \\Delta t)} = \\int_{-\\infty}^t A_k e^{2\\pi i k\\nu_s t'} P(t' - t) dt' \\label{eq:ST80Formalism} \\end{equation} In general, the choice of $P(\\Delta t)$ accounts for the geometry of the emitting and scattering regions.  However, for the large number ($\\sim$100) of scatterings previously noted, the geometry becomes unimportant and a photon diffusion distribution of $P(\\Delta t) = \\tau^{-1} \\exp(\\Delta t/\\tau)$ is appropriate \\citep{Sunyaev80}.  Applying this form to equation~(\\ref{eq:ST80Formalism}), the time lag of the $k^{\\rm{th}}$ harmonic is \\begin{equation} \\Delta t_k = (2\\pi k\\nu_s\\tau)^{-1} \\tan^{-1} 2\\pi k\\nu_s\\tau \\, . \\end{equation} The diffusion timescale cannot be longer than the lags themselves:  $\\tau \\lesssim \\Delta t_{\\rm{max}} = 200$~\\us.  This sets a lower limit of $\\Delta t_2 / \\Delta t_1 \\gtrsim 0.85$ due to scattering alone, contrary to the observation that $\\Delta t_2 / \\Delta t_1 \\approx 0.5$.  Thus Compton scattering cannot by itself fully explain the observed lags: a change in the energy dependence of the initial pulses is still necessary.  Clearly, detailed calculations are required to check whether the two models really would predict the outcomes that we suggest, but they are at least in principle plausible.  One would also need to test whether the models would predict any reversal in other pulsation properties at propeller onset, given that no such changes are observed.  The interesting behavior of the soft lags in this source suggests that the lags of the other accretion-powered millisecond pulsars should be revisited to check for flux dependence.  The only source for which this has been done is XTE~J1814$-$338: an increase in soft lags was observed in the tail of the outburst \\citep{Watts06}, but the rapid drop off and shorter tail of this source made it very difficult to study the pulsations at the lowest fluxes. The other AMPs may be more promising as candidates for further study."
    },
    "0809/0809.3208_arXiv.txt": {
        "abstract": "We have found a photoevaporated disk in the Orion Nebula that includes a wide binary. HST/ACS observations of the proplyd 124-132 show two point-like sources separated by $0\\farcs15$, or about 60 AU at the distance of Orion.  The two sources have nearly identical $I$ and $z$ magnitudes. We analyze the brightest component, Source~N, comparing the observed magnitudes with those predicted using a 1\\,Myr Baraffe/NEXTGEN isochrone with different accretion luminosities and extinctions. We find that a low mass ($\\simeq\\!0.04~M_\\odot$) brown dwarf $\\sim$1~Myr old with mass accretion rate $\\log\\dot{ M}\\simeq-10.3$,  typical for objects of this mass, and about 2~magnitudes of visual extinction provides the best fit to the data.  This is the first observation of a circumbinary disk undergoing photoevaporation and, if confirmed by spectroscopic observations, the first direct detection of a wide substellar pair still accreting and enshrouded in its circumbinary disk. ",
        "introduction": "We report on the discovery of a circumbinary disk seen in silhouette against the bright nebular background of the Orion Nebula. Using multicolor observations taken with the Hubble Space Telescope as part of the HST Treasury Program on the Orion Nebula Cluster, we show that source 124-132, a photoevaporated disk $\\approx\\!1\\farcm5$ north of the Trapezium previously imaged in H$\\alpha$ \\citep{Smith+05}, harbors a binary.  Binaries represent a typical product of the gravitational collapse of cores with high angular momentum. In fact, the  majority of stars in star-forming regions are in binary or multiple systems (see Duchene 1999; Monin et~al.\\ 2007; and for a review Mathieu 2000 and references therein). Also the ONC contains a large number of binary stars \\citep{Prosser+94, Padgett+97, 1998ApJ...500..825P, Koehler+06,Reipurth+07}, although the current counts indicate that the binary frequency, especially at the low-mass end, is lower by a factor 2 to 5 than in star-forming T~Associations like Taurus-Auriga \\citep{Koehler+06, Reipurth+07}.  Whether the relative paucity of binary systems in the ONC is due to the initial conditions of the cloud or to a ``feedback\" effect, like the ejection of binary companions in the cluster core caused by close dynamic encounters, is still a subject of debate.  There can be three disks in a young binary system: two circumstellar disks and a circumbinary one \\citep{1993prpl.conf..749L, 1994ApJ...421..651A, 1997MNRAS.285...33B}, with a complex geometry due to the potential variety of alignments between the disks and orbital planes \\citep{Monin+07}, and interacting through transfer of energy and angular momentum \\citep{Gunther+02}. Disks have been detected around many spectroscopic binaries. Direct imaging of disks around wide, well separated binary systems, either circumstellar or circumbinary, is much rarer, the two most notable examples being the circumbinary disks around GG~Tau \\citep{1994A&A...286..149D} and UY~Aur \\citep{1998A&A...332..867D}.  The data presented here are part of the large dataset (520 images) of ACS exposures obtained for the \\textit{HST Treasury Program on the Orion Nebula Cluster} (GO-10246). We mapped with ACS/WFC an area of about 450 square arcmin nearly centered on the Trapezium Cluster in five filters: F435W (Johnson~$B$, 420s); F555W (Johnson~$V$, 385s); F685N (H$\\alpha$+[\\ion{N}{2}]  $\\lambda 6583$, 340s); F775W (Cousins~$I_\\mathrm{C}$, 385s); and F850LP ($z$-band, 385s). Due to the adopted dithering strategy, most of the field has been exposed two times so the total integration time is typically twice that listed.   The drizzled ACS images have been visually inspected and a master catalog of 219 circumstellar disks has been compiled \\citep{Ricci+08}. Figure~1 is extracted from this last paper, which also provides  the fits files in electronic form. Further details on the \\textit{HST Treasury Program} observing strategy and data products will be given elsewhere  (Robberto et~al.\\ 2008, in preparation). ",
        "conclusions": "Our images reveal the presence of a dark silhouette disk harboring two faint, similar  objects.  In the F775W and F850LP filters the sources appear clearly  point-like, well separated, and without any evidence of bipolar extended  emission from either sides of the disk, as normally observed when edge-on disks obscure their central stars \\citep[][and references therein]{Luhman+07}. We will thus assume that we are seeing the reddened central stars rather than their diffuse  light scattered by the disk surface.  To constrain the nature of the two sources, we plot their position in a color-magnitude diagram (Fig.~\\ref{fig:isochrone}) and (for Source~N only) in a color-color diagram (Fig.~\\ref{fig:color_color}) with respect to the 1~Myr PMS isochrone of \\cite{1998A&A...337..403B}, calculated in our passbands using the \\textsc{NextGenII} synthetic atmosphere models \\citep{nextgenII}. We show the stellar population of the ONC, we also show the displacement from the isochrone caused by extinction and mass accretion, both combined with the synthetic stellar spectra before convolution with the filter passbands. For the extinction we use both the standard $R_V=3.1$ reddening law and the $R_V=5.5$ reddening law, considered more appropriate for the Orion Nebula \\citep{BaadeMinkowsky37, Johnson67}. The accretion spectrum is reproduced following the \\cite{CalvetGullbring98} recipe,  combining $3/4$ of a $T_\\mathrm{eff}=7000$~K black body, accounting for the optically thick emission from the photospheric region heated by the shock of accreting material, with $1/4$ of optically thin emission from a dense slab ($n=10^8$~cm$^{-3}$) of photoionized gas also at $T_\\mathrm{eff}\\simeq7000$~K, accounting for the pre-shock region. The  accretion luminosity, $L_\\mathrm{accr}$, and the stellar bolometric luminosity, $L_\\mathrm{phot}$, are both calculated integrating the model spectra, with $L_\\mathrm{accr}$ finally multiplied by a normalizing factor. This approximate treatment is adequate for our purposes, as we are interested in the intensity in broad-band filters rather than in the details of the true accretion spectrum.  Figures~\\ref{fig:isochrone} and \\ref{fig:color_color} show that the observed  magnitudes of Source~N can be reproduced by different combinations of $A_V$, $\\log(L_\\mathrm{accr}/L_\\mathrm{phot})$ and stellar mass. To find the best combinations of parameters, we calculate a large 3-d grid of models covering the  full range of viable $A_V$, $\\log(L_\\mathrm{accr}/L_\\mathrm{phot})$ and stellar mass, derive the synthetic photometry and search for the best match with the data.  The results are summarized in Table~\\ref{table:intersections} and  marked as stars in Figures~\\ref{fig:isochrone} and \\ref{fig:color_color}. It turns out that for each reddening law, there are typically two intersection points. For each reddening law, the two intersections provide combinations of lower and higher mass, extinction and accretion, (the \\textit{low} and \\textit{high} solutions). All four solutions place Source~N in the brown dwarf regime. To test the robustness of this result and to find the optimal solution, we have assumed that the uncertainties are dominated by photometric errors and performed a Monte Carlo simulation. We generated a set of 10,000 trial stars populating a gaussian error distribution for the three magnitudes; we then iterated our procedure to find the corresponding  mass, accretion luminosity and extinction. Figure~\\ref{fig:mass_accr_relation} shows the results for the two choices of $R_V$. The two areas at the top and bottom of each figure  correspond to the \\textit{high} and \\textit{low} solutions, respectively. The  height of the peaks shows  that the \\textit{high} solution is more compatible with $R_V=3.1$, vice versa if $R_V=5.5$. Only if  $R_V=5.5$ the strongly degenerate \\textit{high} solution may lead to masses above the hydrogen burning limit ($M\\simeq0.08~M_\\odot$).  Another clue to the nature of the sources may come from a comparison of our estimated mass accretion rate with the values normally found for similar stars. Using Eq.~(2) of \\cite{Muzerolle+03} with the stellar parameters from our Baraffe/NEXTGEN models, we derive for $\\dot{M}$ the values reported in Table~\\ref{table:intersections}. A comparison with Figure~8 of \\cite{Muzerolle+03}, or Figure~2 of \\cite{Mohanty+05}, shows that only our \\textit{low} solutions provide mass accretion rates similar to those found in low-mass BDs, whereas the \\textit{high} solutions give values in excess by 2~orders of magnitudes. If we exclude the \\textit{high} values on this basis,  the most consistent solution  for Source~N is the \\textit{low} case, corresponding to a $M\\simeq 0.04~M_\\odot$ brown-dwarf accreting at a rather standard $\\dot{M}\\simeq 5\\times10^{-11}~M_\\odot$~yr$^{-1}$, with about $A_V=2$ and $R_V=5.5$, indicative of grains larger than those typically found in the interstellar medium. Similar conclusions should hold for Source~S, given the nearly identical $I$ and $z$~magnitudes. Even if IR spectroscopy will be needed to firmly establish the nature of these objects, the possibility that the two sources have sub-stellar mass appears quite robust. If confirmed, this would be the first direct observation of a young brown-dwarf binary still accreting in its circumbinary disk.  Whereas our images shows the presence of a large-scale circumbinary disk, the evidence for mass accretion points to to the presence of at least one circumstellar disk. Theory provides a consistent scenario, predicting that circumstellar and circumbinary disks emerge as the mature outcome of the evolution of the original circumstellar disk in which the binary system formed \\citep{Gunther+02}. Due to the exchange of angular momentum from the binary to the disk, an inner gap develops between the inner and outer disk regions. Material can flow through the gap along spiral arms, feeding the circumstellar disks and therefore sustaining mass accretion into the central stars. A detailed comparison of our system with the extant theoretical models is hindered, however, by the fact that our circumbinary disk is being photoevaporated, the first time such a phenomenon is observed. The influence of UV flux from an external source on a circumstellar disk has been studied only for single stars by  \\cite{2002ApJ...578..897R}, who have shown that the extra heating of a disk face produces an increase of the disk flaring angle, leading to photoevaporation and possibly a warping of the system. This may explain why Source~S remains  undetected in our bluest filters:  the southern side of the disk is the one more directly exposed to the UV flux and the stronger photoevaporation may locally increase the optical depth on this side. Assuming the sources are identical, a line-of-sight difference $\\Delta A_V\\simeq 0.3^{m}$  would bring Source~S below our detection limit at the shortest wavelengths while being compatible with the magnitudes observed in the red filters. \\vspace{-12pt}"
    },
    "0809/0809.3664_arXiv.txt": {
        "abstract": "First results from a high-resolution three-dimensional nonlinear numerical study of the kink oscillation are presented. We show in detail the development of a shear instability in an untwisted  line-tied magnetic flux tube. The instability produces significant deformations of the tube boundary. An extended transition layer may naturally evolve as a result of the shear instability at a sharp transition between the flux tube and the external medium. We also discuss the possible effects of the instability on the process of resonant absorption when an inhomogeneous layer is included in the model. One of the implications of these results is that the azimuthal component of the magnetic field of a stable flux tube in the solar corona, needed to prevent the shear instability, is probably constrained to be in a very specific range. ",
        "introduction": "Coronal loops and filament threads are magnetic flux tubes whose field lines are anchored to the dense photosphere. Usually these structures are modeled as circular tubes whose properties change mainly in the radial direction with a sharp or a smooth transition. Using the ideal magnetohydrodynamic (MHD) equations the eigenmodes of cylindrical magnetic tubes can be calculated. A fundamental eigenmode of oscillation is the kink mode, with an azimuthal number $m=1$, which produces a transverse displacement of the tube.  Examples of the calculations of the kink mode for sharp interfaces can be found in \\cite{spruit81,edrob83,cally86} and for smooth transition layers in \\cite{gooss92, rudrob02, vand04, terr06}. In the last case, the coupling of compressional fast magnetohydrodynamic (MHD) waves and shear Alfv\\'en waves leads to the formation of resonances in the inhomogeneous layers and this mechanism is a possible candidate to explain the damping of transverse coronal loop oscillations \\citep{hollyang88, rudrob02, gooss02} and filament threads \\citep{arr08}.  So far, most of the theoretical studies about the kink mode are in the linear regime and second order perturbations are neglected in the MHD equations, however, there are some observational indications suggesting that this assumption might be not fully justified in all the oscillating loops \\citep[][]{terrof04}. Nonlinearity adds new and interesting effects. For example, the ponderomotive force creates a flow along the field lines \\citep{rank94,tik95}, and for the standing kink oscillation it tends to accumulate mass at the loop apex \\citep[][]{terrof04}. When gas pressure is taken into account, the otherwise unlimited accumulation of mass is prevented by the pressure forces which limit the secular growth and produce saturation in the density.  Another possible consequence of the nonlinearity is the generation of instabilities. It is well known that even in the linear regime, a sharp interface between two media in relative motion is liable to the Kelvin-Helmholtz instability (KHI). In the presence of a magnetic field, KHI of a flow parallel to the field lines in a homogeneous medium sets in if the velocity jump exceeds the maximum Alfv\\'en speed. The criteria for the KHI for an axial flow is modified in the presence of a boundary layer of finite width, which also induces the possibility of resonant flow instabilities \\citep[see][]{holl90,yangh91,and01a,and01b}. The effect of an azimuthal flow is less explored in the literature. \\citet{heyvpri83,browpri} showed that azimuthal shear motions in the presence of a smooth transition layer can be KH unstable. The most likely place for the KHI to occur is where the velocity is largest and the magnetic field is perpendicular to the velocity \\citep[see for example][]{rank93}. For the fundamental standing kink mode this is precisely the antinode of the velocity, located at half the loop length from the photosphere.  In this Letter we report on the nonlinear evolution of the kink oscillation of a tube without twist. The full nonlinear three-dimensional ideal MHD equations are solved numerically. From the simulations we investigate the motions in the magnetic tube and show the first results of the development of the nonlinear shear instability at the tube boundary. ",
        "conclusions": "The numerical results shown in this Letter indicate that the shear motions involved in the kink oscillations of a magnetic tube might be unstable. The instabilities are found to create small length scales in the azimuthal direction and grow rapidly in time. For a cylindrical discontinuous interface between two homogeneous stationary rotating fluids the following growth rates (imaginary part of the frequency) of the Kelvin Helmholtz instability can be readily derived: \\begin{equation}\\label{kh} \\omega_i^2=\\frac{\\rho_{\\rm in}\\rho_{\\rm ex}}{(\\rho_{\\rm in}+\\rho_{\\rm ex})^2}\\frac{m^2}{R^2} 4   v_0^2-2 k_z^2 \\frac{B_0^2}{\\mu (\\rho_{\\rm in}+\\rho_{\\rm ex})}. \\end{equation} Here $2v_0$ is the amplitude of the velocity shear at the boundary. The derivation involves the assumption that the azimuthal length scales are much smaller than the longitudinal ones (which is the case for the instabilities that formed in our study). Clearly the background equilibrium for which  equation~(\\ref{kh}) is obtained is very different from the shear motions associated to the kink mode oscillations, which depend on time, $\\phi$ and $z$. Nevertheless, the small scales and the localization of the instabilities in the azimuthal direction, and the fast growth rates, suggest that equation~(\\ref{kh}) may be considered as a local analysis and we need not to worry about the azimuthal and time dependence of the shear motions. For the parameters considered in \\S\\ref{tube}, the first unstable mode corresponds to $m=5$ (which satisfies the assumption of azimuthal localization), and its growth rate is $1/\\omega_i\\approx 5.0\\tau_{\\rm A}$. The important point here is that it is smaller than the kink period, around $16.3\\tau_{\\rm A}$ (satisfying the assumption of localization in time). Moreover, this calculation also clarifies that the most unstable modes will have small azimuthal length scales but large longitudinal length scales, since this minimizes the stabilizing force induced by the bending of the field lines. However, the large longitudinal length scales prevent to interpret equation~(\\ref{kh}) as a local approximation in $z$ since the velocity shear is less than $2v_0$ throughout most of the loop. Growth rates are thus expected to be somewhat smaller, and instability will set in only for azimuthal length scales smaller than those predicted by equation~(\\ref{kh}). An appropriate analysis thus needs to take into account the longitudinal variation of the shear profile and hence necessarily involves a 2D model.  We have found that the evolution of the tube is very sensitive to the amplitude of the initial perturbation. As expected, the larger the amplitude of the initial perturbation the faster the development of the instability. By changing the initial amplitude of the perturbation in a broad range, an extended transition layer (even much larger than the developed in Fig.~\\ref{density}) may naturally evolve as the result of the shear instability of a sharp transition between the flux tube and the external medium. In addition, for very large amplitudes the tube might show severe deformations, a wake can even form behind the tube and it can interact with the main body in each oscillation, the final shape of the tube being rather irregular.  When an inhomogeneous layer between the tube and the environment is included then the motions are characterized by phase-mixed scales and the instability is also present. Nevertheless, the instability develops more slowly than in the sharp transition case. This is probably due to the fact that the thicker the layer the later the generation of small length scales due to the phase mixing process, and thus the onset of the instability. This situation seems to be related to the development of the KHI for torsional Alfv\\'en waves ($m=0$) described by \\citet{browpri} \\citep[see also][]{walker}. Interestingly, for the regime studied here (basically thick layers) the attenuation of the central part of the tube due to resonant absorption is not significantly altered by the changes at the boundary due to the shear instability.  In the context of coronal loops an immediate question that arises from the results presented here is why, up to now, there is no clear evidence of such instability from the observation of oscillating loops. There are several partial answers. Magnetic twist, not included in our model, might decrease or even suppress the instability since the presence of a magnetic field component along the flow stabilizes the KHI. Several examples of the stabilizing effect of a helical magnetic field component can be found in rising tubes in the convection zone \\citep[see for example][]{fer96} or in stellar jets. On the other hand, a tube with very large azimuthal magnetic field  is subject to the instabilities of the linear pinch (typically for a twist larger than $2.5\\pi$, although it depends on the details of the equilibrium and boundary conditions). Therefore, the azimuthal component of the magnetic field of a stable flux tube in the solar corona is probably constrained to be in a specific range (being $2.5\\pi$ an upper bound for the twist).  Other factors that could explain the absence of observational evidences of the instability are that it is not spatially resolved with the current spatial resolution of the telescopes or simply that the amplitude of the oscillating loops is not strong enough to develop the instability. However, the tube displacement produced by the initial perturbation in our simulations is of the order of and even smaller than the observed amplitudes of oscillation in loops, so this last possibility seems to be ruled out.  We have given a qualitative description of the shear instability, but a quantitative analysis (with improved numerical resolution) about the growth rates of the modes, and a detailed study of the effect of the instability on the damping rates under different regimes is still needed. In addition, the equilibrium configuration should be improved in several aspects. For example, a more accurate model should incorporate a realistic variation of the density and temperature from the photosphere to the corona. However, the inclusion of a twisted magnetic field for the reasons mentioned before seems to be the most relevant aspect that needs to be addressed."
    },
    "0809/0809.1661_arXiv.txt": {
        "abstract": "We measure the rest-frame B-band luminosity function of red-sequence galaxies (RSLF) of five intermediate-redshift ($0.5 < z < 0.9$), high-mass ($\\sigma > 950$ km s$^{-1}$) clusters. Cluster galaxies are identified through photometric redshifts based on imaging in seven bands (five broad, and two narrow) using the WIYN 3.5m telescope.  The luminosity functions are well-fit down to $M_B^*+3$ for all of the clusters out to a radius of $R_{200}$.  For comparison, the luminosity functions for a sample of 59 low redshift clusters selected from the SDSS are measured as well.  There is a brightening trend ($M_B^*$ increases by 0.7 mags by z=0.75) with redshift comparable to what is seen in the field for similarly defined galaxies, although there is a hint that the cluster red-sequence brightening is more rapid in the past ($z>0.5$), and relatively shallow at more recent times.  Contrary to other claims, we find little evidence for evolution of the faint end slope.  Previous indications of evolution may be due to limitations in measurement technique, bias in the sample selection, and cluster to cluster variation.  As seen in both the low and high redshift sample, a significant amount of variation in luminosity functions parameters $\\alpha$ and $M^*$ exists between individual clusters. ",
        "introduction": "In the nearby Universe, galaxy clusters are dominated by early type galaxies (Dressler 1980) along the so-called ``red sequence,'' the name for which derives from the tight color-magnitude relation observed for these galaxies (Visvanathan \\& Sandage 1977).  The most luminous cluster galaxies, located at the tip of the red sequence, appear to have relatively little star formation since $z \\sim 3$ (Bower et al. 1992; Ellis et al. 1997; Stanford et al.  1998; Kelson et al. 2001).  Observations of high redshift ($z>1$) clusters are consistent with this inference, indicating at least the bright end of the color-magnitude relationship has only passively evolved since $z\\sim1.5$ (Stanford et al. 1997, Mullis et al. 2005, Stanford et al. 2005). However, evidence for recent ($\\leq 5$ Gyr) bursts of star formation in today's lower-mass cluster galaxies (Poggianti et al. 2001, Conselice et al. 2003) indicates not all cluster galaxies have the same homogeneous star formation histories. While the luminous red sequence in clusters may have long been in place, the extension of the red sequence to lower luminosity may be a relatively recent phenomenon.  Recent, deep imaging programs have led to the claim of a deficit of red sequence at low luminosities in high redshift clusters (Nakata et al. 2001, de Lucia et al. 2004, Goto et al. 2005, Tanaka et al. 2005). This would suggest many of the galaxies seen along today's red sequence may be late additions, perhaps culled from a quenched and faded blue population. This fits neatly with a notion that the red sequence is not monolithic, as evinced by the seminal studies of the Virgo cluster luminosity function (Sandage, Binggeli \\& Tammann 1985). However, the deficit may not be ubiquitous (e.g., Andreon 2006).  Differences in the analysis methods, clusters samples, and the inherent limitations in different data sets (i.e., the depth of available data) have left the observational evidence for a deficit inconclusive.  Differences between studies are also complicated by the possible existence of evolutionary phenomenon that is a function of cluster mass.  Substantial differences exist between the cluster and field red sequences (de Propris et al. 2003, Croton et al. 2005), and there is evidence for differences in the luminosity function with cluster mass, both locally (Hansen et al. 2005, Hilton et al. 2005) and at intermediate redshift (Koyama et al. 2007, Gilbank et al. 2007). However, evidence to the contrary has also appeared in the literature (Barkhouse et al. 2007, Andreon 2007). Given these disparate claims, it is difficult to assemble a clear picture of the red-sequence evolution -- a challenge again exacerbated by differences in cluster samples and measurement techniques between studies.  A common theme repeated in many of the studies is the large cluster to cluster variation seen in any of the parameterizations of the cluster populations, particularly the luminosity function.  Estimates of the dwarf-to-giant ratio (DGR, de Lucia et al. 2006) and Schechter-function (Schechter 1976) fits to the luminosity function (Barkhouse et al. 2007, Andreon 2007) both show significant cluster-to-cluster variation. Unique features also have long been identified in the shape of luminosity functions of local clusters (Biviano et al. 1995, Yagi et al. 2002), indicative of a multi-component population even among the red-galaxy population.  Traditionally, deep studies of the so-called red-sequence luminosity function (RSLF) in local clusters have fit two functions (either double Schechter functions or a Gaussian and Schechter function) to the distribution of galaxies (Sandage et al. 1985, Jerjen \\& Tammann 1997, Popesso et al. 2005). The bright end of the luminosity function is well described by a functional form with a sharp turnover at lower luminosities, while the fainter luminosity-function component is generally found to be rapidly rising at lower luminosities. Variations in the relative amplitude of these components may well drive overall variations in observed cluster luminosity functions. Studies of intermediate redshift clusters are dominated by this brighter component, due to depth limitations. While such studies rarely reach deep enough to well-characterize the rapidly-rising faint end, its modulation may significantly impact the observed count-parametrization.  In this paper, we explore the red-sequence population in five high-density regions containing massive, bona fide, and well-studied clusters between $0.5<z < 0.9$. Our study is based on deep, multi-band imaging data from the WIYN 3.5m telescope and extant spectroscopic data. The depth of our WIYN data is comparable to other cluster studies at similar redshifts, and allows an independent assessment of the cluster RSLF significantly below the knee in the bright end of the luminosity function. The extent to which this, or any other extant study at intermediate redshift, can address the faint-end of the cluster RSLF is a point we address. For comparison, we include measurements of the luminosity function for a large sample of low redshift clusters from the Sloan Digital Sky Survey measured in the same manner as our intermediate redshift sample.  In \\S 2 we summarize the observation data and highlight our measurement techniques.  In \\S \\ref{sec:lf}, we present the luminosity functions measured for each of the clusters as well as the tests to confirm and assess the reliability of the measured luminosity functions. In \\S \\ref{sec:lowz}, the low-redshift cluster sample is described. Finally, we discuss the observed evolution in our luminosity functions (\\S 4), compare to results in the literature, and discuss the implications of our results in \\S 5.  Throughout this work, we adopt H$_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{Matter} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$; all absolute magnitudes are in the rest-frame (Johnson) B-band in the Vega system. ",
        "conclusions": "  \\begin{itemize} \\item Selection effects can be critical to the determination of any signal of evolution.  Clusters do show variations with the luminosity function in terms of mass and radius, which can lead to erroneous conclusions in terms of evolution, if not carefully accounted. Although small, this survey predominately measures the RSLF in massive system which seems to display different behaviors than low mass systems.  \\item Significant cluster to cluster variations exists, even at a given mass. The dispersion in cluster luminosity-function parameters, measured for individual clusters, is typically an order of magnitude greater than the error estimate on those parameters from fitting to an ensemble-averaged luminosity function.  Significantly more work needs to be done to better understand these cluster to cluster variations.  \\end{itemize}  A clear picture of the evolution of the cluster RSLF remains elusive, yet such a picture is critical for understanding the processes that drive the growth and transformations of cluster-galaxies over cosmic time. It is conceivable that a sufficiently delineated map of the RSLF evolution with time, cluster-mass, and location within the local cluster environment can help confirm or refute predictions of $\\Lambda CDM$ models on time-scales relevant to the assembly of the most massive, virialized systems in the universe. While the observational challenge has yet to be met, the theoretical models also fall short on definitive predictions due to the complexity of the gas-physics (including star-formation) that channels the transformation of galaxies from the blue cloud onto the red sequence. For example, clean predictions of how and when the different sub-populations along the red sequence (e.g., E, S0, dE) are formed are yet to be had. As such, the red sequence, as interpreted simply as a mass sequence, will likely continue as a critical observational foil thrown up as a test of hierarchical models. Tracking the fate of the blue population of galaxies that cause the Butchler-Oemler effect is a key to understanding the physics behind the transformative processes in clusters; at least some of these systems are likely to be the progenitors of the red-sequence galaxies. The complexity of the astrophysics will likely stymie a definitive observational picture of the transformation process, however headway can be made by juxtaposing the blue and red populations in the context of environment.  In future work we will investigate the cluster blue-galaxy population in this context.  We would like to thank Vy Tran and Chris Moran for providing access to their spectroscopic redshifts for MS1054 and MS0451, the anonymous referee for comments that dramatically expanded and improved this work, SAAO for support (SMC), and U. Toronto for hospitality while pursuing this work (MAB).  Research was supported by STScI/AR-9917, NSF/AST-0307417, NSF/AST-0607516 and a Wisconsin Space Grant. We acknowledge use of the Sloan Digital Sky Survey (SDSS and SDSS-II; see //www.sdss.org/ for funding, management and participating institutions).   \\appendix"
    },
    "0809/0809.3178_arXiv.txt": {
        "abstract": "In a former paper (Garcia-Rissmann \\etal 2005; hereafter Paper I), we have presented spectra of 64 active, 9 normal and 5 Starburst galaxies in the region around the near-IR Calcium triplet absorption lines (CaT) and the [SIII]$\\lambda$9069 line. In the present paper we analyze the CaT strength ($W_{\\rm CaT}$), and kinematical products derived in that study, namely stellar ($\\sigma_{\\star}$) and ionized gas ($\\sigma_{\\rm gas}$) velocity dispersions. Our main results may be summarized as follows: (1) Seyfert 2s show no sign of dilution in $W_{\\rm CaT}$ with respect to the values spanned by normal galaxies, even when optical absorption lines such as the CaII K band at 3933 \\AA\\ are much weaker than in old, bulge-like stellar populations. (2) The location of Seyfert 2s in the $W_{\\rm CaT}$-$W_{\\rm CaK}$ plane is consistent with evolutionary synthesis models. The implication is that the source responsible for the dilution of optical lines in these AGN is a young stellar population, rather than an AGN featureless continuum, confirming the conclusion of the pioneer study of Terlevich, D\\'{\\i}az \\& Terlevich. (3) In Seyfert 1s, both $W_{\\rm [SIII]}$ and $W_{\\rm CaT}$ tend to be diluted due to the presence of a non-stellar component, in agreement with the unification paradigm. (4) A comparison of $\\sigma_\\star$ with $\\sigma_{\\rm gas}$ (obtained from the {\\it core} of the [SIII] emitting line) confirms the existence of a correlation between the typical velocities of stars and clouds of the Narrow Line Region. The strength and scatter around this correlation are similar to those previously obtained from the [OIII]$\\lambda$5007 line width. ",
        "introduction": "\\label{sec:Introduction}   The stellar absorption lines from the near-IR Calcium triplet (CaT) made their debut in the field of active galactic nuclei (AGN) with the work of Terlevich, D\\'{\\i}az \\& Terlevich (1990, hereinafter TDT).  At that time, it was thought that the optical spectrum of Seyfert 2s contained a featureless continuum (FC) from the AGN, which, in light of unification scenarios (Antonucci 1993), should be associated with scattered light from the hidden Seyfert 1 nucleus (Cid Fernandes \\& Terlevich 1995).  TDT realized that this idea could not be fully correct, since their data indicated that, unlike for absorption lines in the optical, the strength of the CaT in these sources shows no signs of dilution with respect to the values found in normal galaxies. Their interpretation of this fact was that what was called a ``non-stellar FC'' was in fact a stellar component associated to young stellar populations. This conclusion was thoroughly confirmed by numerous studies in the past decade (see Cid Fernandes \\etal 2004 for a review).  Since then, interest in the CaT shifted towards its use as a tracer of stellar kinematics, particularly in AGN.  The location of the CaT in a relatively clean spectral region makes it an ideal feature to measure stellar velocity dispersions ($\\sigma_\\star$), which trace the gravitational potential of the host galaxy's bulge.  In an influential work, Nelson \\& Whittle (1996; NW) used the CaT in a comparison between stellar and gaseous kinematics of AGN, finding that the typical velocity Narrow Line Region (NLR) clouds ($\\sigma_{\\rm gas}$) correlates with $\\sigma_\\star$, albeit with significant scatter.  More recently, a strong correlation between $\\sigma_\\star$ and black-hole mass (Tremaine \\etal 2002 and references therein) indirectly boosted interest in observations of the CaT (Nelson 2000; Botte \\etal 2004).   For completely different reasons, recent work on normal galaxies has also raised the interest in the CaT. Empirical investigations have revealed rather puzzling behavior of the CaT in bulges and elliptical galaxies. Firstly, detailed population synthesis models tend to over-predict the CaT strength (Saglia \\etal 2002; Cenarro \\etal 2003, 2004), particularly for giant ellipticals, though for dwarf ellipticals the match between data and models is satisfactory (Michielsen \\etal 2007).  Secondly, while classic metallicity tracers like the Mg$_2$ index are known to correlate with $\\sigma_\\star$ (eg, Terlevich \\etal 1981), tracing the well-known mass-metallicity relation, the strength of the CaT appears to be {\\it anti-correlated} with $\\sigma_\\star$ (Cenarro \\etal 2003, 2004; Falc\\'on-Barroso \\etal 2003; Michielsen \\etal 2003). The interpretation of these results is still not clear, and the situation is likely to be even more complex in systems with varied star formation histories like AGN (Cid Fernandes \\etal 2004; Wild \\etal 2007).   One thus sees that, more than 15 years after TDT introduced the CaT in AGN research, there are still plenty of reasons to keep studying it. With this general motivation, in Paper I we have carried out a spectroscopic survey of Seyfert galaxies in the region including the CaT and the [SIII]$\\lambda$9069 emission line. Paper I concentrated in the presentation of the data and the derivation of four main data products: stellar velocity dispersions ($\\sigma_\\star$), [SIII] emission line widths ($\\sigma_{\\rm [SIII]}$, which is representative of the highly ionized component of the gas of the Narrow Line Region in AGNs), [SIII] equivalent widths ($W_{\\rm [SIII]}$), and the CaT strength ($W_{\\rm CaT}$).  Here we use these data to address the following issues: (1) examine the contribution of a non-stellar component in the near-IR spectra of Seyfert galaxies and test the consistency of $W_{\\rm [SIII]}$ and $W_{\\rm CaT}$ data with the unified model, (2) study the sensitivity of $W_{\\rm CaT}$ to stellar population properties, (3) investigate the connection between NLR and stellar kinematics and CaT strength, and (4) report results of spatially resolved spectroscopy, not reported in Paper I.  To study these questions, we complement the data in Paper I with values of the equivalent width of the CaII K line ($W_{\\rm CaK}$) obtained from previous studies. As shown by Cid Fernandes \\etal (2001), the CaII K line is a powerful tracer of stellar populations even in type 2 Seyferts. By comparing $W_{\\rm CaT}$ with $W_{\\rm CaK}$ we can assess whether the CaT is a useful stellar population tracer, and determine whether the combination of CaK and CaT strengths is explained as due to stars alone or if an FC component is necessary. We also add in literature information on the width ($\\sigma_{\\rm [OIII]}$) of the [OIII]$\\lambda$5007 line, which is useful to test whether our results are somehow affected by the choice of [SIII]$\\lambda$9069 as a tracer of NLR motions.   This paper is organized as follows. Section \\ref{sec:CaT_NonStellarLight} presents a series of studies related to the CaT strength. After showing CaT observational properties in \\ref{subsec:CaT in Sy}, we compare $W_{\\rm CaT}$ with $W_{\\rm CaK}$ in \\ref{subsec:Dilution} and investigate whether the location of the points in the $W_{\\rm CaT}$ versus $W_{\\rm CaK}$ diagram is consistent with the existence of an FC at near-IR wavelengths. In \\ref{Stellar Models} we use evolutionary synthesis models to track the behavior of both $W_{\\rm CaT}$ and $W_{\\rm CaK}$ as a function of age and metallicity, and overlay the models onto the data in the $W_{\\rm CaT}$ versus $W_{\\rm CaK}$ diagram. In \\ref{subsec:Spatial CaT Profiles} we study the spatial behavior of $\\sigma_{\\star}$ and $W_{\\rm CaT}$ for about half of our sample. In Section \\ref{sec:Kinematics} we analyze stellar and ionized gas kinematics: stellar velocity dispersions and its relation to CaT are studied in \\ref{subsec:Stellar velocity dispersions}, while in \\ref{subsec:Stellar vs. NLR Kinematics} we discuss the link between stellar and gas kinematics. Finally, section \\ref{sec:Summary} summarizes our main results. ",
        "conclusions": "\\label{sec:Summary}  In Paper I we presented a spectroscopic atlas of 78 galaxies in the region of CaT lines, and measured the equivalent widths of CaT and [SIII]$\\lambda$9069 lines, and stellar and gaseous velocity dispersions. In this Paper we used these data, as other measurements, to investigate the CaT strength and the kinematical properties of Seyfert nuclei.  Our conclusions can be summarized as follows:  \\begin{enumerate}   \\item The equivalent widths of CaT and [SIII] are diluted in Seyfert 1s, due to the presence of a non-stellar component, while in Seyfert 2s there is no sign of dilution. We show that the nuclear dilution of CaT lines in Seyfert 1s and the non observed dilution in Seyfert 2s holds spatially, i.e., within approximately the central kpc. At this distance the effects disappears, thus implying that the stellar lines are diluted due the nuclear non-stellar continuum directly seen in Seyfert 1s.  \\item The CaT strength turns out not to be a good tracer for the mean ages and metallicities of metal rich stellar populations. Nonetheless, its combination with CaK line yields a useful constraint on the nature of the continuum emission from optical to near-IR wavelengths.  \\item We show that the location of Seyfert 2s in the CaT-CaK plane could be explained satisfactorily by stellar population synthesis models, by considering mixtures of an old plus very young stellar populations with a extinction of just two or three magnitudes. The hypothesis of Seyfert 2 nuclei composition as a mixing of an old stellar population plus power-law central source cannot account for the majority of the data.  \\item There is a correlation between nuclear $\\sigma_\\star$ and $\\sigma_{\\rm [SIII]}$, as well as an anticorrelation between $\\sigma_{\\rm [SIII]}$/$\\sigma_{\\star}$ and $\\sigma_{\\star}$, both with substantial scatter. This means that care must be exercised when using $\\sigma_{\\rm [SIII]}$ as a proxy for $\\sigma_\\star$. These results are compatible with previous results of NW and GH.  \\end{enumerate}"
    },
    "0809/0809.3949_arXiv.txt": {
        "abstract": "We present high resolution $4.7\\micron$ CO fundamental spectroscopy of V836 Tau, a young star with properties that are between those of classical and weak T Tauri stars and which may be dissipating its circumstellar disk. We find that the CO line profiles of V836 Tau are unusual in that they are markedly double-peaked, even after correcting for stellar photospheric absorption in the spectrum. This suggests that the CO emission arises from a restricted range of disk radii ($< 0.5$\\,AU), in contrast to the situation for most classical T Tauri stars where the CO emission extends out to much larger radii ($\\sim 1-2$\\,AU).  We discuss whether the outer radius of the emission in V836 Tau results from the physical truncation of the disk or an excitation effect.  We also explore how either of these hypotheses may bear on our understanding of disk dissipation in this system. ",
        "introduction": "V836 Tau is an interesting young star with properties that are between those of classical and weak T Tauri stars. Its optical spectrum displays relatively weak, variable H$\\alpha$ emission, with measured equivalent widths in the range 1--25\\AA\\ (White \\& Hillenbrand 2004; Beristain, Edwards, \\& Kwan 2001; Kenyon et al.\\ 1998; Wolk \\& Walter 1996; Mundt et al.\\ 1983). Similarly, V836 Tau also shows little veiling and has one of the smallest measured stellar accretion rates among T Tauri stars (Hartigan, Edwards, \\& Ghandour 1995). Nevertheless, the detection of H$\\alpha$ emission with an inverse P Cygni profile (Wolk \\& Walter 1996) demonstrates the continued, possibly intermittent, accretion of material onto the star. Millimeter wavelength CO emission has also been detected from V836 Tau (Duvert et al.\\ 2000) despite the relatively low disk mass $\\sim 0.01\\Msun$ inferred from the submillimeter dust continuum properties (e.g., Andrews \\& Williams 2005).  At infrared wavelengths, V836 shows little or no excess in the $K$- and $L$-bands, but a much stronger excess indicative of an optically thick disk is apparent beyond 10$\\micron$ (Strom et al. 1989; Skrutskie et al. 1990; Padgett et al.\\ 2006). Interpreting the unusual spectral energy distribution (SED) and low stellar accretion rate of V836 Tau as indications that the system is on the verge of dissipating its disk, Strom et al.\\ (1989) consequently referred to V836 Tau as a ``transitional T Tauri star'', a term reflecting the view that classical T Tauri stars evolve into weak T Tauri stars with the dissipation of the disk playing a central role in that evolutionary process. More generally, transition objects, systems that can be modeled as an optically thick disk that has an optically thin region (a hole or a gap) at smaller radii, have been suggested to be in the process of dissipating their disks from the inside-out (Strom et al.\\ 1989; Skrutskie et al.\\ 1990), possibly as a result of the formation of planetesimals (Strom et al.\\ 1989), a giant planetary companion (e.g., Skrutskie et al.\\ 1990), disk photoevaporation (Clarke et al.\\ 2001; Alexander et al.\\ 2006), or potentially other processes.   We report here on the properties of the 4.7$\\micron$ CO fundamental ($\\Delta v$=1) emission from V836 Tau to see whether it can provide any insights into the extent and nature of disk dissipation in the system. CO fundamental emission has been previously characterized as a probe of circumstellar disk gas (Najita et al.\\ 2003, 2007a; Blake \\& Boogert 2004; Brittain et al.\\ 2007). In the case of disks surrounding low mass stars, CO fundamental emission is estimated to arise from within $1-2$\\,AU of the star (Najita et al.\\ 2003). As a result, the emission is well-suited to probing the radial distribution of gas in the terrestrial planet region of the disk. Recent thermal-chemical models of the inner regions of disks surrounding low mass stars suggest that the CO emission arises from the warm surface region of the disk that is heated by stellar (X-ray) irradiation and possibly mechanical heating processes (e.g., Glassgold et al.\\ 2004). The CO fundamental emission from V836 Tau was discussed briefly in an earlier study (Najita et al.\\ 2003). Here we report a new, higher signal-to-noise observation of the CO emission, which we use to characterize the radial distribution of the gas in inner $\\sim$1--2\\,AU of the disk surrounding V836 Tau. ",
        "conclusions": "\\subsection{Truncated Excitation or Truncated Disk?}  The CO emission from V836 Tau shows similarly double-peaked line profiles for lines spanning a large range in excitation temperature.  The line profile shapes indicate that the CO emission is truncated beyond $\\sim 0.4$\\,AU. The similarity in the relative strengths of the 1--0 transitions show that the 1--0 transitions are optically thick over the emitting region. As described in the previous section, these properties could be produced by a {\\it physical} truncation of the gaseous disk beyond a radius of $\\sim 0.4$\\,AU or a truncation of the CO emission (but not the disk gas column density) at the same radius.  Although it is difficult to determine which of these interpretations is correct based on the available information, some possibilities can be ruled out.  For example, the truncation of the CO emission is unlikely to result from the thermal dissociation of CO since the excitation temperature of the gas is much less than the thermal dissocation temperature of CO ($\\sim 4000$\\,K at the densities of inner disks).  We might also consider the possible explanations that have been put forward for the origin of the double-peaked line profiles observed in the $v$=2--0 CO overtone bandhead emission from T Tauri stars and Herbig Ae stars at $2.3\\micron$ (e.g., Carr et al.\\ 1993; Najita et al.\\ 1996, 2000). Such emission is detected only from sources with high accretion rates, probably a consequence of the temperatures and high column densities that are needed to produce the overtone emission. Since these systems have disks that are believed to be radially continuous, the double-peaked lines are unlikely to arise from a radially truncated disk. We have previously described two possible explanations for the double-peaked lines that make up the CO $v$=2--0 bandhead: (1) the outer radius is a dust sublimation front, or (2) the transition from atomic H to H$_2$ at decreasing temperature depopulates the higher CO vibrational levels.  In the first scenario, a dust sublimation front renders the line emitting layer optically thick in the continuum at temperatures below $\\sim 1500$\\,K, eliminating the contrast of the CO $v$=2--0 emission above the continuum at these radii (e.g., Carr 1989). This explanation is difficult to apply to the CO fundamental lines, because typical dust temperatures at the inferred outer radius of $\\sim 0.4$\\,AU for the CO fundamental emission are much below the dust sublimation temperature (e.g., D'Alessio et al.\\ 1999) and measured dust sublimation radii are much smaller than the inferred outer radius of the CO fundamental emission (Eisner et al.\\ 2005; Muzerolle et al.\\ 2003). Therefore, the outer radius observed for the CO fundamental emission is unlikely to result from dust sublimation.  As an alternative scenario, we previously speculated that the transition from atomic to molecular hydrogen at decreasing disk temperature leads to the depopulation of the higher vibrational levels of CO due to the lower collisional cross-section of molecular hydrogen with CO compared to that of atomic hydrogen with CO (Najita et al.\\ 1996). In chemical equilibrium, disks would transition from a mixture of CO and atomic hydrogen at small disk radii ($T > 2000$\\,K) to a mixture of CO and molecular hydrogen at large disk radii ($T < 2000$\\,K). In this situation, we estimated that radiative trapping would be able to maintain the $v \\ge 2$ vibrational populations down to a temperature of $\\sim 1500$\\,K, below which significant depopulation would occur. Using a non-LTE model of this kind, we were able to reproduce the line intensities and shapes of CO overtone emission lines covering a wide range of excitation conditions ($v$=2--0 to $v$=5--3) in the spectra of two Herbig Ae stars.  The same effect is unlikely to apply in detail to the truncation of the CO fundamental emission seen in V836 Tau ($\\Rout \\sim 0.4$\\,AU). While chemical equilibrium may be relevant at the large column densities needed to produce the overtone emission, disks are expected to have significant vertical structure and to depart significantly from chemical equilibrium at the smaller column densities needed to produce the fundamental emission. Thermal-chemical models of T Tauri disks irradiated by stellar X-rays (Glassgold et al.\\ 2004; Meijerink et al.\\ 2008) indicate that overlying the large column densities where the overtone lines form is a warm surface layer that is expected to be conducive to the production of CO fundamental emission over a large range in radii (Glassgold et al.\\ 2004).  The X-ray irradiated disk models, which have currently studied the structure of disks over the region 0.25--2\\,AU, find that throughout this range of radii disks possess a warm ($\\sim 1000$\\,K) surface layer ($\\gtrsim 10^{21}\\persqcm$) of CO mixed with atomic hydrogen (Glassgold et al.\\ 2004; Meijerink et al.\\ 2008). Thus, the CO fundamental transitions could plausibly be excited over radii within and beyond 1\\,AU. This expectation is in agreement with observations of CO fundamental emission from T Tauri stars. Empirically, we find that almost all accreting T Tauri stars show CO fundamental emission and that the majority of CO fundamental emission profiles are centrally peaked.  This indicates that the emission arises from a wide range of disk radii ($\\Rout/\\Rin > 20$, $\\Rout \\simeq 1-2$\\,AU), without a sharp truncation in excitation as a function of velocity (Najita et al.\\ 2003).    However, it is possible that these general trends, obtained for typical T Tauri stars, do not apply to V836 Tau. Whereas typical T Tauri stars have strong near-infrared excesses and stellar accretion rates $\\sim 10^{-8}\\Msunperyr$, V836 Tau has a weak near-infrared excess and a low stellar accretion rate $\\sim 10^{-9}\\Msunperyr$ (Hartigan et al.\\ 1995 scaled to Gullbring et al.\\ 1998; Herczeg et al.\\ 2006). The weak near-infrared excess might indicate either a significant settling of grains out of the disk atmosphere or a lack of dust at small disk radii.  In this situation, one might imagine that a lack of small grains could reduce the heating of the gaseous atmosphere, if gas-grain collisions dominate the heating of the gaseous disk.  In a cooler atmosphere, vibrational CO emission would be more difficult to produce, perhaps contributing, thereby, to the truncation of the CO emission. This does not seem likely in the context of recent X-ray irradiated disk atmosphere models, in which the gas is heated directly by X-rays or accretion-related processes and the grains function primarily as a coolant for the gas through collisions. In such a model, we might expect grain growth and settling to produce warmer gaseous atmospheres, rather than a radial truncation of the CO emission.  Alternatively, one might imagine that CO might be less abundant in a grain-poor disk atmosphere if the CO is synthesized from H$_2$ that forms on grains.  However, in the X-ray irradiated disk atmosphere models, the grains are typically warmer than 100\\,K within 1\\,AU, and have therefore been assumed to play no significant role in the synthesis of CO and other molecules (Glassgold et al.\\ 2004). As a result, grain growth and settling is not expected to significantly reduce the strength of the CO emission from the disk atmosphere.  Another possibility is that the low accretion rate of V836 Tau compared to the average T Tauri accretion rate ($\\sim 10^{-8}\\Msunperyr$; Hartmann et al.\\ 1998) might play a role in truncating the CO emission. Since accretion-related processes may heat the gaseous atmosphere (Glassgold et al.\\ 2004), the gaseous atmosphere may experience reduced heating at the lower accretion rate of V836 Tau. While the available heating is clearly able to produce detectable $v$=1--0 CO fundamental emission from V836 Tau, could reduced heating radially truncate the emission, perhaps via a non-LTE effect?  A simple estimate indicates that the vibrational populations could be in non-LTE. The critical density $n_{\\rm cr}$ for the $v$=1 level of CO, obtained from the Einstein A-values for the $v$=1--0 transitions (e.g., Goorvitch \\& Chackerian 1994) and the $v$=1--0 collision rates for CO with atomic hydrogen (Glass \\& Kironde 1983), varies as $T^{-1/2}$ and equals $\\sim 5 \\times 10^{12}\\percc$ at a temperature of 1000\\,K (see discussion in Najita et al.\\ 1996). Our modeling of the CO emission from V836 Tau shows that the $v$=1--0 CO emission arises from a gas column density $\\sim 0.01 \\gpersqcm$ at disk radii $\\lesssim 0.4$ AU.  In the D'Alessio et al.\\ (1999) disk model, the density in the disk atmosphere over this range of column density is $\\nH \\sim 10^{11}-10^{12} \\percc$ at radii 0.1--0.4 AU. Therefore, a rough estimate is $\\nH/n_{\\rm cr} \\sim 0.1$ in the CO emitting region.  Our modeling of the CO emission further shows that the $v$=1--0 CO lines are optically thick, with $\\tau \\sim 10$. Since the escape probability for Gaussian lines depends asymptotically on the line optical depth as $\\sim \\tau^{-1}$ (Mihalas 1978), the line optical depth approximately compensates for the low value of $\\nH/n_{\\rm cr}$ so that the $v=1$ level could be in LTE. The situation for the higher vibrational levels is possibly similar, with the critical densities for these levels being comparable to the critical density for $v$=1. This because the A-values of the higher vibrational levels are larger (proportional to $v$), but the collision rates may also be larger because of the contribution from $\\Delta v > 1$ collisions (A. Glassgold, personal communication). Thus, the CO vibrational level populations could plausibly be in LTE, consistent with the assumption made in \\S 3.3.  However, both lower temperatures and lower densities in the CO emitting region will tend to drive departures from LTE.  Reduced accretion heating can produce both lower temperatures and lower densities.  As the accretion heating is reduced, the vertical temperature inversion that puts the CO into emission will become increasing limited to the lower column density surface region that can be heated by external irradiation (e.g., by stellar X-rays).  This surface region will be characterized by lower densities. Detailed calculations of the thermal, chemical, and density structure of disk atmospheres are needed to address this issue quantitatively.  Because such calculations are currently lacking, we might take instead a more empirical approach and compare the CO fundamental line profiles of V836 Tau with those of other T Tauri stars with low stellar accretion rates and/or low near-infrared excesses. As an example of such a comparison, Figure 7 (top panel) shows the CO fundamental emission from LkCa\\,15, a T Tauri star with a stellar accretion rate similar to that of V836 Tau (Hartmann et al.\\ 1998; see also Najita, Strom, \\& Muzerolle 2007b). The spectrum shows the data reported in Najita et al.\\ (2003), but re-reduced using the approach described in \\S 2. The $v$=1--0 CO emission from LkCa\\,15 is centered at the stellar velocity (short vertical lines, bottom panel), but compared to V836 Tau, it shows a centrally peaked profile.  Although the signal-to-noise ratio of the spectrum is limited, weak stellar photospheric features appear to be present. In the bottom panel of Figure 7, we show the CO emission from LkCa\\,15 after correction for a stellar photospheric contribution, following the approach described in \\S 3.1. The stellar photospheric model assumes an Allard stellar atmosphere model with a gravity of $\\log g = 4.0$, an effective temperature of 4400\\,K to match the K5 spectral type of LkCa\\,15 (Herbig \\& Bell 1988), a stellar rotational velocity $v \\sin i =12.5\\kms$ (Hartmann, Soderblom, \\& Stauffer 1987), and an observed (topocentric) radial velocity of $-39.8 \\kms$, which is appropriate for the measured stellar radial velocity (Hartmann et al.\\ 1987) and the observation date. A veiling of 2.5 times the stellar continuum roughly reproduces the strength of the stellar photospheric features near the 1--0 P30 line. This level of veiling is also consistent with the veiling at $5\\micron$ implied by the SED (e.g., Furlan et al. 2006). Subtracting the veiled stellar photospheric component produces a disk CO emission profile that is even more centrally peaked (Fig.\\ 7, bottom panel).  In addition to its low stellar accretion rate, LkCa\\,15 is also similar to V836 Tau in that it has the characteristics of a transition object with a weak near-infrared continuum and an optically thick outer disk (e.g., Bergin et al.\\ 2004; Espaillat et al.\\ 2007). Its properties therefore probe empirically how both the reduction in small grains and reduced accretion-related heating might affect the CO fundamental emission from the disk. The LkCa\\,15 spectrum shows that in at least some cases, these effects do not radially truncate the CO fundamental emission within 1 AU.   To summarize, the double-peaked line profile of the CO fundamental emission from V836 Tau may indicate that the gaseous disk extends from close to the star ($\\sim 0.05$\\,AU) out to a physical truncation radius ($\\sim 0.4$\\,AU). If the gaseous disk in V836 Tau is instead continuous beyond this radius, the double-peaked CO profile would indicate a sudden truncation of the CO emission beyond a radius of 0.4\\,AU. As discussed in \\S 4.2, an abrupt decrement in excitation is not expected empirically, since the majority of CO emission profiles observed to date (both typical classical T Tauri stars and transition objects like LkCa\\,15) show centrally peaked CO profiles with no comparable decrement in excitation with radius.  However, a truncated emission profile may result from either an (anomalously) steep temperature gradient or departures from LTE that become important in low accretion rate systems. These two possible explanations can be explored and potentially distinguished with a higher signal-to-noise spectrum that measures the strengths of the $v$=2--1 lines. A theoretical study of the possibility of non-LTE CO level populations would also be welcome in sorting out whether an excitation effect is a possible explanation for the outer radius of the emission. Additional observations of low accretion rate sources would be useful in exploring this issue empirically.  Most definitive of all would be observations of spectral line transitions that robustly probe the disk atmosphere of V836 Tau at excitation temperatures $< 400$\\,K.  These observations would complement the insensitivity of the CO fundamental transitions to low temperature gas. If little emission is detected with these diagnostics at disk radii $>0.4$\\,AU, that would strongly suggest that the gaseous inner disk in V836 Tau is physically truncated at $\\sim 0.4$\\,AU. Some of the mid-infrared molecular emission diagnostics recently reported in the spectrum of AA Tau (C$_2$H$_2$, HCN, H$_2$O, OH; Carr \\& Najita 2008) may prove useful in this regard.   \\subsection{Nature of V836 Tau}  The possibility that the disk of V836 Tau is physically truncated beyond 0.4\\,AU may bear on our understanding of the nature of the system. As noted in \\S 1, V836 Tau has been previously classified as a transition object (Strom et al.\\ 1989), a system with an optically thin inner disk (within $\\Rhole$) and an optically thick outer disk (beyond $\\Rhole$). The SEDs of these systems have been variously explained as a consequence of grain growth and planetesimal formation (e.g., Strom et al.\\ 1989; Dullemond \\& Dominik 2005), giant planet formation (e.g., Skrutskie et al.\\ 1990; Marsh \\& Mahoney 1992), or photoevaporation (e.g., Clarke et al.\\ 2001; Alexander et al.\\ 2006).  While all of these processes can produce optically thin regions in the disk (inner holes or gaps), they make different predictions for stellar accretion rates and disk masses, as well as the radial distribution of the gaseous disk. As a result, measuring the radial distribution of disk gas and comparing the stellar accretion rates and disk masses of transition objects with those of accreting T Tauri stars of comparable age can potentially sort among the possible explanations for a transitional SED.  For example, a recent study using the latter approach showed that transition objects in Taurus (including V836 Tau) have stellar accretion rates that are on average $\\sim 10$ times lower than those of non-transitional T Tauri stars with comparable disk masses (Najita et al.\\ 2007b). Such a reduced stellar accretion rate is predicted for disks that have formed Jovian mass planets (e.g., Lubow \\& D'Angelo 2006), suggesting that giant planet formation may play a role in explaining the origin of at least some transition objects.  Studies of the radial distribution of the gaseous disk in the system provide an additional way to distinguish the nature of individual transition objects. For example, although grain growth can render the inner disk optically thin (within $\\Rhole$), the gas in the same region of the disk is not expected  to be altered significantly; gas would therefore fill the region within $\\Rhole$. If a giant planet has formed with a mass sufficient to open a gap, both the gas and dust would be expected to be cleared dynamically from the vicinity of the orbit of the planet, creating an inner disk (within $\\Rinner < \\Rhole$) that is fed by accretion streams from an outer disk (beyond $\\Rhole$; Lubow, Seibert, \\& Artymowicz 1999; Bryden et al.\\ 1999; Kley 1999; D'Angelo et al.\\ 2003; Lubow \\& D'Angelo 2006). Accretion streams are not expected in the case of a massive giant planet ($\\sim 10 M_J$; e.g., Lubow et al.\\ 1999), and no significant gas or dust is expected anywhere within $\\Rhole$ in this case. A lack of gas or dust within $\\Rhole$ is also expected in the photoevaporation case; such systems are further expected to have a very low disk mass ($\\sim 0.001 \\Msun$; e.g., Alexander \\& Armitage 2007).  Our results for V836 Tau are intriguing in this context. Compared to the SEDs of well-studied transition objects such as GM Aur and DM Tau, where a strong infrared excess appears only beyond $\\sim 10\\micron$ indicating an optically thin region 3--20\\,AU in size (Calvet et al.\\ 2005), the SED of V836 Tau shows a significant infrared excess at a shorter wavelength $\\sim 5-10\\micron$. Therefore, if the V836 Tau system has an optically thin inner region, it is comparatively small and plausibly within the range of disk radii probed by CO fundamental emission (within $\\lesssim 2$\\,AU; Najita et al.\\ 2003). The SED of V836 Tau can be fit with a simple model of a flared optically thick disk with an inner hole $\\sim 1$\\,AU in radius (see Appendix). In comparison, the SED of the transition object LkCa\\,15 can be interpreted as indicating an optically thin region within $\\sim 3$\\,AU (Bergin et al.\\ 2004), or a radial gap extending from 5--46\\,AU (Espaillat et al.\\ 2007). A planet orbiting at such large distances ($\\sim 3$\\,AU or $\\sim 40$\\,AU) may create a gap in the disk, but not radially truncate the gaseous disk significantly within 1\\,AU. In contrast, systems like V836 Tau, in which the SED may indicate an optically thin region at much smaller radii ($< 1$\\,AU), are the ones for which the CO fundamental emission would in principle be capable of diagnosing a physically truncated inner disk if an orbiting companion is present. The radially truncated CO emission that we observe may support this picture.  As an alternative interpretation of the available data, the short wavelength SED of V836 Tau ($\\lambda < 10\\micron$) can also be fit with a simple model of an inclined ($i=60$), geometrically flat disk (see Appendix) that possibly results from significant grain growth and settling. In such a situation, the gaseous disk would be radially continuous and we might expect to observe a CO profile like those of other non-transition T Tauri stars.  Such a profile, typically centrally peaked, is not observed.  In the photoevaporation and massive giant planet scenarios for the origin of a transitional SED, little gas is expected to be present within $\\Rhole$, in contrast to the observed situation where a gaseous disk is present at 0.05--0.4\\,AU and ongoing (possibly intermittent) stellar accretion is observed.  The relatively high disk mass of V836 Tau ($\\sim 0.01\\Msun$; Andrews \\& Williams 2005), compared to the much smaller masses at which photoevaporation is expected to be able to create an inner hole ($\\sim 0.001\\Msun$), further argues against the photoevaporation scenario.  The possibility of a massive planet is also restricted by current limits on the stellar radial velocity of V836 Tau, which constrain the mass of a companion within $0.4 - 1$\\,AU to $<5-10 M_J$ (L. Prato, personal communication).  These arguments are schematic in that they rely on theoretical predictions that have not been verified observationally.  For example, the sizes of gaps that will be induced by an orbiting companion of a given mass, and the extent to which stellar accretion will be reduced, are poorly known from an observational point of view.  Stellar companions have been found to establish large inner holes and to terminate stellar accretion in some transition objects (CoKu Tau/4---Ireland \\& Kraus 2008; D'Alessio et al.\\ 2005) and not in others (CS Cha---Guenther et al.\\ 2007; Espaillat et al.\\ 2007).  The situation for lower mass companions is essentially unexplored.  As theoretical predictions are tested, it will be useful to examine in detail the range of companion masses that are consistent with the properties of V836 Tau."
    },
    "0809/0809.1457_arXiv.txt": {
        "abstract": "Dense, small molecular cloud cores have been identified as the direct progenitors of stars. One of the best studied examples is Barnard 68 which is considered a prototype stable, spherical gas core, confined by a diffuse high-pressure environment. Observations of its radial density structure however indicate that Barnard 68 should be gravitationally unstable and collapsing which appears to be inconsistent with its inferred long lifetime and stability. We argue that Barnard 68 is currently experiencing a fatal collision with another small core which will lead to gravitational collapse. Despite the fact that this system is still in an early phase of interaction, our numerical simulations imply that the future gravitational collapse is already detectable in the outer surface density structure of the globule which mimicks the profile of a gravitationally unstable Bonnor-Ebert sphere. Within the next $ 2 \\times 10^5$ years Barnard 68 will condense into a low-mass solar-type star(s), formed in isolation, and surrounded by diffuse, hot interstellar gas. As witnessed in situ for Barnard 68, core mergers might in general play an important role in triggering star formation and shaping the molecular core mass distribution and by that also the stellar initial mass function. ",
        "introduction": "Barnard 68 is considered an excellent test case and the prototype of a dense molecular cloud core (Alves et al. 2001b). Because of its small distance ($\\sim$ 125 pc), this so called Bok globule /Bok \\& Reilly 1947) with a mass of M=2.1 $M_{\\odot}$, contained within a region of R = 12,500 AU has been observed with unprecedented accuracy (Alves et al. 2001b; Lada et al. 2003; Bergin et al. 2006; Redman et al. 2006; Maret et al. 2007).  Deep near-infrared dust extinction measurements of starlight toward individual background stars observed through the cloud provided a detailed 2-dimensional map of its projected surface density distribution (Fig. 1) from which a high-resolution density profile was derived over the entire extent of the system. The striking agreement of the inferred density structure with the theoretical solution of an isothermal, pressure confined, hydrostatic gas sphere (so called Bonnor-Ebert sphere) was interpreted as a signature that the globule is old, thermally supported and stable, with the pressure gradient balancing the gravitational force.  This conclusion has received additional support from molecular line observations (Lada et al. 2003) that show complex profile shapes which can be interpreted as signatures of stable oscillations (Redman et al. 2006; Broderick et al. 2007) with subsonic velocities of order $V \\approx$ 0.04 km/s which is 20\\% of the isothermal sound speed $c_s \\approx$ 0.2 km/s. The age of Barnard 68 should therefore be larger than one dynamical oscillation timescale $\\tau_{dyn}$ = 2R/V = 3 $\\times 10^6$ yrs which is long compared to its gravitational collapse timescale $\\tau_{coll}=(3 \\pi/32 G \\rho)^{1/2} = 0.17 \\times 10^6$ years, where $\\rho=1.5 \\times 10^{-19}$ g cm$^{-3}$ is the average mass density. The observed hydrostatic profile and the inferred large age are strong arguments for stability.  Other observations however are in conflict with this conclusion. The best fitting hydrostatic model leads to the conclusion that Barnard 68 should be gravitationally unstable (Alves et al. 2001b).  Pressure confined, self-gravitating gas spheres with radii $R$ and isothermal sound speeds $c_s$ have self-similar density distributions (see Appendix) that are characterized by the dimensionless parameter  \\begin{equation} \\xi_{max} = \\frac{R}{c_s} \\sqrt{4 \\pi G \\rho_c} \\end{equation}  \\noindent where $\\rho_c$ is the central density.  Although all values of $\\xi_{max} > 0$ represent equilibrium solutions where the gravitational force is exactly balanced by pressure forces, a stability analyses (Bonnor 1956) shows that small perturbations should lead to gravitational collapse in cores with $\\xi_{max} > 6.5$.  Barnard 68's surface density distribution is characterised by $\\xi_{max} = 6.9 \\pm 0.2$. It should collapse on a free-fall timescale which is much shorter than its oscillation timescale. In addition, the question arises how this globule could ever achieve an unstable equilibrium state in the first place.  Dense molecular cores like Barnard 68 are the last well characterized configuration of interstellar gas before gravitational collapse towards star formation. They have masses from a tenth to tenths of solar masses and typical sizes from a tenth to a few tenths of a parsec (Alves et al. 2007).  They are found in isolation, known then as Bok globules (Bok 1948), or embedded in lower density molecular cloud complexes. Because of their relatively simple shapes they have long been recognized as important laboratories to the process of star formation. Barnard 68 is one such dense core.  Despite being isolated and surrounded by warm and diffuse gas it is part of the Pipe Nebula complex (Alves et al. 2007; Lombardi et al. 2006; Lada et al. 2008) which consists of an ensemble of about 200 cores within a region of order 10 pc.  Many cores appear distorted and asymmetric. This could be a result of their interaction with the highly turbulent diffuse gas environment that leads to non-linear and non-radial oscillations (Broderick et al., 2007, 2008). Another possibility which we will investigate here are collisions between cores. We propose that Barnard 68 is currently experiencing such a fatal collision that triggers gravitational collapse. Its peculiar and seemingly contradictory properties are early signatures of this process that cannot be understood if the globule is treated as an isolated and stable Bonnor-Ebert sphere. Observations show that the distribution of molecular core masses is surprisingly similar to the stellar initial mass function (Alves et al. 2007). This indicates that the masses of stars originate directly from the processes that shape the molecular core mass function. Barnard 68 demonstrates that core-core collisions could play an important role in this process. ",
        "conclusions": "We are then in the fortunate situation of witnessing the collapse of a Bok globule and the formation of a star like our Sun, or a low-mass multiple stellar system, in the relative nearby solar neighborhood. Given the distance to this cloud this would make Barnard one of the nearest star forming clouds to Earth, a perfect laboratory to investigate the early phases of gravitational cloud collapse. Barnard 68 has probably been stable for several million years. However its current fatal impact with an object of 10\\% its mass has made this globule gravitationally unstable which is revealed by a dimensionless structural parameter $\\xi_{max} = 6.9 \\pm 0.2$ that already exceeds the critical value expected for a stable cloud.  The impact is also mixing gas from the chemically less evolved outer parts of the main cloud as well as the CO-rich gas of the bullet with the CO depleted gas in the center of Barnard 68.  Within the next $10^5$ years the bullet will be completely engulfed by Barnard 68 which at the same time will develop a centrally peaked surface density profile (dashed curve in Fig. 3) and a velocity field consistent with large-scale gravitational collapse. $10^5$ years later a new infrared source will appear in its central region.  An isolated star like our Sun will be born in our immediate Galactic neighborhood, probably surrounded by a residual dusty accretion disk where planets might start forming."
    },
    "0809/0809.0876_arXiv.txt": {
        "abstract": "{Low and intermediate frequency quasi-periodic oscillations (QPOs) are thought to be due to oscillations of Comptonizing regions or hot blobs embedded in Keplerian disks. Any movement of these perturbations is expected  systematically to change the QPO frequency.} {Our goal is to find systems where such a systematic drifts have been observed. We also try to find the real cause of such drifts and whether they shed some light on the accretion disk dynamics.} {Using archival data of the recent outburst of GRO J1655-40, we report the presence of such systematic drifts not only during the rising phase from the $25^{th}$ of February 2005 to the $12^{th}$ March 2005, when the QPO frequency monotonically increased from $82$mHz to $17.78$Hz but also in the decline phase from the $15^{th}$ September 2005 to the $5^{th}$ of October 2005, when the QPO frequency decreased from $13.14$Hz to $34$mHz.} {We fitted the frequency drifts with the propagatory oscillating shock solution. In the shock-oscillation solution, the frequency is inversely proportional to the infall time scale from the shock location. We obtained the shock location and strength through such a fit. } {The astonishing smoothness of the variation of the QPO frequency over a period of weeks directly supports the view that it may due to the drift of an oscillating shock rather than the movements of a blob inside a differentially rotating disk.} } ",
        "introduction": "The galactic nano-quasar GRO J1655-40 is an interesting low mass X-ray binary (LMXB) with a primary mass $M = 7.02\\pm0.22~M_\\odot$ (Orosz \\& Bailyn 1997) and the companion star mass = $2.3~M_\\odot$ located at a distance of $D = 3.2 \\pm 0.2$~kpc (Hjellming \\& Rupen 1995). The disk has an approximate inclination angle of $\\theta = 69.5^\\circ\\pm0.1^\\circ$ (Orosz \\& Bailyn 1997) to the line of sight. In the last week of February 2005 it became X-ray active (see Shaposhnikov et al. 2007 and references therein) and remained so for the next 260 days before returning to the hard state. In this Letter, we thoroughly analyse the data of the first two weeks of the very initial stage (rising phase) and the last three weeks of the final stage (decline phase) of the 2005 outburst. We clearly observe very smooth day to day variation of the QPO frequency in these two phases. We propose that a satisfactory explanation of this behavior can be obtained if we assume that an oscillating shock which is sweeping inward through the disk in the rising phase and outward in declining phase is responsible for the QPO. In the next Section, we briefly present an overview of QPOs observed in the black hole candidates. In \\S 3,  we present the shock oscillation solution for the generation of QPOs. In \\S 4, we present the observational results in detail and show our model fit of the QPO frequencies from day to day. We interpret the results  and extract the shock parameters. In \\S 5, we give a coherent description of the rising and the declining phases of the outburst. ",
        "conclusions": "In this Letter, we show that during the rising phase of the outburst of GRO J1655-40, the QPO frequency increases very slowly in the first few days and then rapidly increases to about $18$Hz before it disappears altogether. Our slowly drifting shock oscillation solution explains this variation very accurately. Our estimation suggests that the shock was at $1270$ Schwarzschild radii when the first observation of QPO was made and it went to $r=59$ Schwarzschild radii when it was last observed. Within $15$ hours of this, the shock front went inside the horizon and thus when the observation was made on the next day, the QPO was absent. The inward drift velocity of the shock was slow, only about $20$ m s$^{-1}$. As the shock proceeds close to the black hole, the Keplerian disk follows the shock and the energy spectrum gradually became softer with the black body component becoming stronger each day. Thus the whole scenario is consistent with our theoretical understanding (MSC96) that this drift is due to the reduction of the post-shock pressure by the increased cooling effects in addition to the higher upstream ram pressure. Both MSC96 and CAM04 computed the cooling timescale using bremsstrahlung and Comptonization respectively and showed that oscillations occur when the infall time scale is comparable to these coolings in super-massive and staller mass black holes respectively. As a consistency check, one could also use the fitted spectrum to compute the electron temperature and calculate the cooling time scale from $E_{th}/{\\dot E}$ where, $E_{th}$ is the thermal energy content of the electrons and ${\\dot E}$ is the rate of cooling due to Comptonization. However, $E_{th}$ depends on the accretion rate and electron temperature and ${\\dot E}$ depends on the enhancement factor (Dermar et al. 1991). However, the PCA data from 3-25keV does not allow one to compute the these unknowns unambiguously.  In the decline phase, the QPO frequency was found to be decreasing monotonically. This means that an oscillating shock is propagating outwards. The nature of the reduction of QPO frequency is very curious. For the first $\\sim 3.5$ days, the shock location did not change very much, from $\\sim 40$ to $\\sim 59$ or so, as if it was stalling. In this phase an additional cooler Comptonized component was required. After this, the shock suddenly moved away with almost constant acceleration for another $16$ days before the QPO disappeared completely. The spectrum remained hard and the intensity decreased monotonically. Towards the end, the spectrum is dominated only by the power-law component. We conjecture that in the first $\\sim 3.5$ days the shock remained inside the disk drifting slowly outward, perhaps due to interaction of the receding shock with the still incoming Keplerian flow. After that, the receding Keplerian disk created a vacuum in the system and accelerated the shock wave outward. This is indicated by almost square-law behavior of the shock location.  In the two component advective flow model that we are using, the shock forms in the sub-Keplerian component since a Keplerian flow is subsonic and cannot have a shock. At the onset of the rising phase of the outburst, only the sub-Keplerian flow rate increases rapidly since it is almost freely falling. The shocks we consider may have actually formed farther out (than $1270r_g$, for example) we could detect them only when the rms value of the QPO is high enough. Similarly, when the shock comes closer to the black hole horizon, the noise also rises and QPOs may not be detectable as well.  The observational result we described here is unique in the sense that we are able to connect the QPO frequency of one observation with that of the next by a simple analytical formalism. To our knowledge no competing model exists which uses a true solution of the flow, such as shocks in our case, to explain such a behavior. In any case, even if the QPOs were generated by certain non-axisymmetric feature (such as a blob), it is impossible that it will survive beyond a few orbital time-scales. Furthermore, the smooth decrease of QPO frequency requires that the blob moves outwards through a differentially rotating disk for some weeks, odds of which is insignificant. If our explanation of the propagation of shocks is correct, then this outburst shows convincingly how a shock wave smoothly disappears behind the horizon after $\\sim 16$ days of its initial detection. Thus, such an observation would be unlikely in neutron star candidates. Our interpretation is generic and other outburst sources should also show such systematic drifts. A discussion of other similar sources will be dealt with in future."
    },
    "0809/0809.2991.txt": {
        "abstract": "Surface photometry detections of red and exceedingly faint halos around galaxies have resurrected the old question of whether some non-negligible fraction of the missing baryons of the Universe could be hiding in the form of faint, hydrogen-burning stars. The optical/near-infrared colours of these red halos have proved very difficult to reconcile with any normal type of stellar population, but can in principle be explained by advocating a bottom-heavy stellar initial mass function. This implies a high stellar mass-to-light ratio and hence a substantial baryonic mass locked up in such halos. Here, we explore the constraints imposed by current observations of ordinary stellar halo subdwarfs on a putative red halo of low-mass stars around the Milky Way. Assuming structural parameters similar to those of the red halo recently detected in stacked images of external disk galaxies, we find that a smooth halo component with a bottom-heavy initial mass function is completely ruled out by current star count data for the Milky Way. All viable smooth red halo models with a density slope even remotely similar to that of the stacked halo moreover contain far too little mass to have any bearing on the missing-baryon problem. However, we note that these constraints can be sidestepped if the red halo stars are locked up in star clusters, and discuss potential observations of other nearby galaxies that may be able to put such scenarios to the test. ",
        "introduction": "The quest to unravel the nature of dark matter, estimated to make up around 90\\% of the total matter content \\citep[e.g.][]{Komatsu et al.}, remains one of the most important tasks of modern cosmology. Dark matter appears to exist in at least two separate forms: one baryonic, and one non-baryonic. While the non-baryonic component is the dominant one, a substantial fraction of the baryons in the low-redshift Universe \\citep[$\\approx 1/3$--2/3;][]{Fukugita,Fukugita & Peebles a,Nicastro et al.,Prochaska & Tumlinson} are also at large.  These missing baryons could in principle be hiding in a variety of different forms: as faint/failed stars or stellar remnants \\citep[so-called Massive Astrophysical Compact Halo Objects or MACHOs;][]{Griest}, as cold gas clouds \\citep{Pfenniger et al.,Pfenniger & Combes}, as a warm/hot intergalactic medium \\citep{Cen & Ostriker,Dav\u00e9 et al.} or as hot gaseous halos around galaxies \\citep{Maller & Bullock,Fukugita & Peebles b,Sommer-Larsen}. While current simulations seem to favour the latter two alternatives as the main reservoirs, observations are still unable to confirm this hypothesis \\citep[see][for a review]{Bregman}.  The old idea of baryonic dark matter in the form of faint, low-mass stars has recently gained new momentum through surface photometry detections of very red and exceedingly faint structures -- ``red halos'' -- around galaxies of different types. The history of this topic goes back to the mid-90s, when deep optical and near-IR images indicated the presence of a faint halo around the edge-on disk galaxy NGC 5907 \\citep[e.g.][]{Sackett et al.,Lequeux et al. a,Rudy et al.,James & Casali}. The colours of this structure were much too red to be reconciled with any normal type of stellar population, and indicative of a halo population with an abnormally high fraction of low-mass stars. At around the same time, \\citet{Molinari et al.} also announced the detection of a red halo around the cD galaxy at the centre of the galaxy cluster Abell 3284. Skepticism grew with the discovery of what appeared to be the remnants of a disrupted dwarf galaxy close to NGC 5907 \\citep{Shang et al.}, leading to suggestions that this feature, in combination with other effects, could have resulted in a spurious halo detection \\citep{Zheng et al.}. While follow-up observations with the Hubble Space Telescope (HST) took some of the edge out of this criticism \\citep{Zepf et al.}, the field fell into disrepute. Deeper images have later revealed a wealth of tidal streams in the halo of NGC 5907 \\citep{Martinez-Delgado et al.}, but this does not by itself explain the red excess originally detected.  The red halos would not die quietly though, and new reports started to surface a few years later. First, \\citet{Bergvall & \u00d6stlin} and \\citet{Bergvall et al.} presented deep optical/near-IR images of faint and abnormally red structures around blue compact galaxies. \\citet*{Zibetti et al.} then stacked images of 1047 edge-on disk galaxies from the Sloan Digital Sky Survey (SDSS) and detected a halo population with a strong red excess and optical colours curiously similar to those previously derived for NGC 5907 -- again very difficult to reconcile with standard halo populations. The halo detected around an edge-on disk galaxy at redshift $z=0.322$ in the Hubble Ultra Deep Field shows similarly red colours \\citep{Zibetti & Ferguson}, and \\citet{Tamm et al.} argue that even the halo of Andromeda displays a pronounced red excess.  \\citet{Zackrisson et al. a} analyzed the colours of some of these new detections and found that the halos of both blue compact galaxies and stacked edge-on disks could be explained by a stellar population with a very bottom-heavy initial mass function ($dN/dM\\propto M^{-\\alpha}$ with $\\alpha\\approx 4.50$). The high mass-to-light ratio of such a population makes it a potential reservoir for at least part of the baryons missing from current inventories. While the stellar initial mass function (IMF) is often assumed to be universal, recent observational studies suggest that it may vary both as a function of environment \\citep{Hoversten & Glazebrook}, and as a function of cosmic time \\citep{van Dokkum}. Corroborating evidence for an IMF as extreme as that advocated by \\citet{Zackrisson et al. a} also comes from star counts in the field population of the LMC, where a slope of $\\alpha\\approx 5$--6 was derived for masses $\\geq 1 \\ M_\\odot$ \\citep{Massey,Gouliermis et al.}.  Taking the red halo detections at face value, one may ask whether the Milky Way itself could be surrounded by a hitherto undetected red halo of low-mass, hydrogen-burning stars, with photometric properties similar to the halo detected around stacked edge-on disks. The known baryonic components (thin disc, thick disk, bulge and standard stellar halo) of the Milky Way contribute around 5--$6\\times 10^{10}\\ M_\\odot$ \\citep[e.g.][]{Sommer-Larsen & Dolgov,Klypin et al.,Flynn et al.} to the Milky Way's virial mass of $\\approx 1\\times 10^{12} \\ M_\\odot$ \\citep[][]{Klypin et al.}. A cosmic baryon fraction of $\\Omega_\\mathrm{baryons}/\\Omega_\\mathrm{M}\\approx 0.17$ \\citep{Komatsu et al.}, combined with the theoretical prediction that the baryon fraction should be $\\approx 90\\%$ of the cosmic average for a Milky Way-sized halo \\citep{Crain et al.}, on the other hand suggests the presence of some $\\approx 1.5\\times 10^{11}\\ M_\\odot$ of baryonic material within its virial radius, leaving at least $60\\%$ of its baryons to be found.  Low-mass stars would in principle be detectable through microlensing effects, and such putative MACHOs may already have been discovered in the halos of both the Milky Way and M31 \\citep[e.g.][]{Alcock et al.,Calchi Novati et al. a,Riffeser et al.}. Their masses (0.1--1 $M_\\odot$) and inferred contribution to the mass of the dark matter of galaxies ($\\approx 20 \\%$) are, however, very difficult to reconcile with any kind of stellar MACHO candidate \\citep[e.g.][]{Freese}. This result, coupled to the fact that competing teams have failed to confirm these detections \\citep[][]{Tisserand et al.,de Jong et al.} have led to the suspicion that the observations must have been misinterpreted \\citep[e.g.][]{Belokurov et al.} or that the compact objects detected through this technique may be of non-baryonic origin (e.g. primordial black holes, mirror matter objects, preon stars or scalar dark matter miniclusters -- see \\citealt{Zackrisson & Riehm} for a more thorough discussion).  If some of the missing baryons of the Milky Way are locked up in the form of hydrogen-burning stars in a red halo, such a structure must also have evaded the faint star counts aimed to constrain the luminosity function of halo subdwarfs \\citep[e.g.][]{Gould et al., Gould,Digby et al.,Brandner}, since no significant excess of low-mass stars has yet been detected with this method. In fact, direct observations of this kind are expected to impose much stronger constraints on main sequence stars in the halo than what current microlensing surveys can achieve.  In this paper, we explore to what extent a smooth red halo of low-mass stars similar to that detected by \\cite{Zibetti et al.} might already be ruled out by these observations of halo subdwarfs in the Milky Way. In section 2, we describe the observational data used. Section 3 outlines the method used to test red halo models against these observations. The resulting constraints on red halo models are presented in section 4. Section 5 discusses the robustness of these constraints and section 6 summarizes our findings. ",
        "conclusions": "Our results indicate that a smooth halo with a bottom-heavy IMF and structural parameters similar to those of the stacked halo is completely ruled out in the Milky Way by current star count data. A halo component with a bottom-heavy IMF would have to have an overall scaling or density profile that differs substantially from that of the stacked halo to remain viable. Moreover, all permitted smooth red halo models with a density slope {\\it even remotely similar} to that of the stacked halo contain far too little mass to have any bearing on the missing-baryon problem in the Milky Way.  These conclusions are admittedly based on a large number of assumptions regarding the properties of the red halo. Since one can in principle consider many formation scenarios for a halo-like structure dominated by low-mass stars -- for example in situ formation, population III stars, dynamical mass segregation \\citep[see][for a more details discussion]{Zackrisson et al. b} -- the properties of the red halo population may be very different from those thus far adopted. Here, we therefore investigate the robustness of our conclusions in light of the most important assumptions made.  \\subsection{Age, star formation history and metallicity of the halo} \\begin{figure*} \\plottwo{f5a.eps}{f5b.eps} \\caption{$N_\\mathrm{sub}/L_i$ as a function of age for stellar populations with different IMF slopes $\\alpha$ (indicated by labels).  {\\bf a)} Populations with $Z=0.001$ and $\\tau=1$ Gyr. {\\bf b)} Populations with $\\tau=1$ Gyr at metallicities of $Z=0.0001$ (thin dashed), $Z=0.001$ (thick solid) and $Z=0.008$ (thin solid). To avoid cluttering, only a subset of IMF slopes are included. \\label{N/L}} \\end{figure*} In deriving the constraints presented in section 4, fixed $N_\\mathrm{sub}/L_i$ ratios have been adopted for every IMF slope considered, based on the assumed age, prior star formation history and metallicity of the red halo. In reality, of course, we have but very poor constraints on these properties. For a halo population, one naively expects a relatively high age and a low metallicity, and we therefore restrict the discussion to ages $t\\geq 1$ Gyr and metallicities $Z\\leq 0.008$. In Fig.~\\ref{N/L}a and b, we demonstrate the dependence of $N_\\mathrm{sub}/L_i$ on age and $Z$ in the case of a stellar population with $\\tau=1$ Gyr. The $N_\\mathrm{sub}/L_i$ ratio generally increases as a function of age.  However, for the more bottom-heavy IMFs (i.e. high IMF slope $\\alpha$), the age dependence is weak, and an age lower than the 10 Gyr assumed in section 4 would allow $N_\\mathrm{sub}/L_i$ to be lowered by no more than $\\approx 25\\%$ (for $\\alpha\\geq 4$). Allowing a metallicity higher than the assumed value of $Z=0.001$ would just increase $N_\\mathrm{sub}/L_i$ and hence strengthen the constraints. A metallicity as low as $Z=0.0001$ would decrease $N_\\mathrm{sub}/L_i$, but by no more than $\\approx 30\\%$. The maximum decrease is moreover attained at the highest ages, which means that this drop in $N_\\mathrm{sub}/L_i$ does not add to the age effect -- you cannot have both at the same time. Hence, the combined effect of age and metallicity is no more than $\\approx 30\\%$ for IMFs with $\\alpha\\geq 4$.  Alterations in the star formation history affect the $N_\\mathrm{sub}/L_i$ ratio in a way that is very similar to age variations. The maximised $N_\\mathrm{sub}/L_i$ are attained in the case of an instantaneous burst of star formation (i.e. a single-age population). While the perfectly coeval onset and quenching of star formation associated with this scenario seems unrealistic given the huge spatial scales involved in the halo, an instantaneous burst would just make the constraints on the red halo stronger. Allowing a more prolonged star formation history than the $\\tau=1$ Gyr adopted in section 4 on the other hand lowers $N_\\mathrm{sub}/L_i$. In the extreme case of having a star formation that increases as a function of time, the $N_\\mathrm{sub}/L_i$ ratio stays almost constant at the value attained at the lowest ages. Hence, such scenarios do not provide the means for lowering $N_\\mathrm{sub}/L_i$ more than the age effects already discussed.  In summary, $N_\\mathrm{sub}/L_i$ for these alternative scenarios may be somewhat lower ($\\approx 30\\%$) than the values listed in Table \\ref{NL_table} and used in the constraints presented in section 4. A detailed inspection of the entries in Table \\ref{NL_table} reveals that this may shift the constraint levels by at most one level, in the sense that the constraints given for an IMF with slope $\\alpha$ may relaxed to mimic those for an IMF with a slope one step lower among those tested. As an example, changing the age from 10 to 1 Gyr for an $\\alpha=4.00$ model would lower $N_\\mathrm{sub}/L_i$ from $N_\\mathrm{sub}/L_i\\approx 1800$ to $\\approx 1300$ (see Fig.~\\ref{N/L}a), i.e. a decrease by $\\approx 30\\%$. This $N_\\mathrm{sub}/L_i$ ratio is intermediate between those listed for $\\alpha=3.50$ and 4.00. Hence, the relaxed constraints on the allowed structural parameters for a $\\alpha=4.00$ halo would fall between those of the $\\alpha=3.50$ and 4.00 contours in Fig.~\\ref{paramspace}. This shift is too small to allow a halo population with $\\alpha=4.0$--5.0 into regions of the ($q_\\mathrm{red}$, $\\beta_\\mathrm{red}$) parameter space where it would contribute significantly to the missing-baryon problem.  Other spectral synthesis codes may produce slightly different values for $N_\\mathrm{sub}/L_i$, but the qualitative impact of such variations on the final constraints is easy to assess from Table~\\ref{NL_table} and Fig.~\\ref{N/L}.  Once the constraints on $q_\\mathrm{red}$ and $\\beta_\\mathrm{red}$ have been determined, the mass of the viable red halo models are computed using equation (\\ref{Meq}), based on the $M/N_\\mathrm{sub}$ ratio. This ratio also has a slight dependence on age and star formation history, but varies by no more than $\\approx 30\\%$ for IMFs with slopes $\\alpha\\geq 3$. The inferred red halo masses change insignificantly because of such effects.  Hence, we argue that our conclusions are robust with respect to the assumptions made about the age, metallicity and star formation history of the red halo.  \\subsection{Disk scale length} The constraints presented in section 4 have been derived using a scale length for the Milky Way of $r_\\mathrm{exp}=2.5$ kpc. The effect of changing $r_\\mathrm{exp}$ is to spatially shift the projected distance $r_\\mathrm{proj}=2\\times r_\\mathrm{exp}$ between the plane of the disk and the region where the density scale of the red halo models are set by the assumed $\\mu_{\\mathrm{red},\\; i}$. Adopting a smaller value of $r_\\mathrm{exp}$ weakens the constraints imposed on red halo models by star counts in volumes A and B, whereas a larger $r_\\mathrm{exp}$ would make the constraints stronger. Allowing $r_\\mathrm{exp}=2.0$ kpc \\citep[at the lower limit of what is allowed by the constraints set by][]{Juric et al.} only allows an increase of the overall mass of red halos with IMFs with $\\alpha=4.0$--5.0 by a factor of $\\approx 2$, which still places the red halo mass substantially below the range relevant for the missing-baryon problem in the Milky Way. Therefore, our conclusions seem very robust with respect to uncertainties in the scale length of the Milky Way disk.  \\subsection{Halo IMF} So far, we have only considered power-law IMFs with single-valued slopes $\\alpha$ for the red halo population. In principle, far more complicated IMFs (e.g. broken power laws, lognormal IMFs, IMFs with mass spikes) could be considered. While this is not warranted given the large observational uncertainties associated with current red halo data, improved measurements may indeed require modifications of the assumed IMF. The revised constraints can be assessed by comparing $N_\\mathrm{sub}/L_i$ and $M/N_\\mathrm{sub}$ of populations with more complicated IMFs to the values listed in Table~\\ref{NL_table}. However, as long as the IMF is not tweaked to move a very large fraction of the overall population mass outside the stellar mass range to which the current subdwarf star counts are sensitive ($\\approx 0.15-0.7 \\ M_\\odot$), the conclusions presented here will not be qualitatively altered.  \\subsection{Halo density profile}  The constraints presented are based on the reasonable assumption that the volume density of red halo stars decreases monotonically with distance from the Galactic centre. Shells and tidal tails are sometimes seen in the outskirts of galaxies, and represent overdensities of matter that formally violate this assumption. If one entertained the notion that the abnormal colours of the red halo come from low-mass stars in a shell-like structure, then the constraints presented here can in principle be completely sidestepped by pushing the radius of this shell to distances outside the volume probed by the \\citet{Gould et al.} star counts (volume B in Fig.~\\ref{schematic}). However, while the density of such a structure can be chosen freely to correspond to the measured surface brightness level at any single projected radius, the decrease in surface brightness as a function of projected distance from the centre would be far slower than that reported by Z04.  So far, we have assumed the red halo density profile to be a single-valued power law, whereas more complicated density profiles can of course be considered. A broken power-law profile with a shallow slope (or a constant density core) close to the centre ($R_\\mathrm{c}< R_\\mathrm{proj}$) would result in constraints on the outer slope and flattening identical to those presented here, except that the overall red halo mass would become even lower than implied by current constraints for the allowed parameter combinations. The situation becomes more complicated if the power-law density profile changes slope inside (or even outside) the volumes A and B. However, a break of this type would be imprinted in the red halo surface brightness profiles at some level, and there is no compelling evidence for such features at the current time.  \\subsection{Subdwarf detectability as a function of distance}  The luminosity functions of \\citet{Digby et al.} and \\citet{Gould et al.} have for simplicity been adopted throughout volumes A and B for all subdwarfs in the luminosity range $7.5\\lesssim M_V\\lesssim 12.4$. In reality, the luminosity function is not sampled equally well at all distances, since the faintest (i.e. least massive) subdwarfs typically cannot be detected out to as large distances as the brightest ones. In fact, $M_V\\approx 12.4$ stars are detected only out to a distances of order 5 kpc from the Sun in volume B, i.e. substantially smaller than the 40 kpc adopted for its outer boundary. The constraints imposed by these volumes may therefore be too strict for red halo populations composed almost entirely of the subdwarfs at the lowest masses. However, even if we completely disregard all constraints based on volume B, thereby restricting ourselves to very small distances from the Sun ($\\leq 2.8$ kpc), we are still left with reasonably strong constraints imposed by volume A (as depicted in Fig.~\\ref{paramspace2}a). All $\\alpha=4.50$ populations with $\\beta_\\mathrm{red}\\geq 2.75$ are constrained to contribute less than $10^{10}\\ M_\\odot$ to the baryon budget for all $\\mu_i$ normalizations considered. Shallower density profiles $\\beta_\\mathrm{red}< 2.75$ do allow red halo masses in excess of $10^{10}\\ M_\\odot$, but these would be in poor agreement with the values derived by Z04.   \\subsection{Halo smoothness}  One way of sidestepping the Milky Way star counts would be to assume that the red halo stars are not smoothly distributed, but clustered. Scenarios of this type have been carefully investigated by \\citet{Kerins a,Kerins b}, and remain a viable way of hiding the red halo population from detection in the Milky Way. The volumes A and B in which the star counts have assumed to be valid (Fig.~\\ref{schematic}), are crude representations of the actual volumes probed. In reality, the subdwarf samples used in the studies by \\citet{Gould et al.} and \\citet{Digby et al.} have been obtained in a limited number of fields, which would more accurately be represented by a large number of narrow cones probing the volume. If red halo stars were clustered, we may have missed them simply because they have so far happened to fall between the cones.  Such clusters of low-mass stars in the Milky Way halo might already be detectable as faint, extended objects in a wide-angle survey such as the SDSS (provided that a sufficient number of such clusters is located at sufficiently small distances from us), or in external galaxies, using a combination of surface photometry and star counts.  As an example, consider a star cluster of mass $4\\times 10^4 \\ M_\\odot$ \\citep[the mass favoured by][]{Kerins a,Kerins b} with $Z=0.004$ and IMF slope $\\alpha=4.50$ \\citep[i.e. the values favoured by][]{Zackrisson et al. a}. The cluster would have an absolute magnitude of $M_I\\approx -3.3$ in the Cousins $I$-band, which is $\\approx 0.5$ mag below the tip of the red giant branch (RGB) in a 10 Gyr old population. Objects this bright are readily detectable with the HST at distances out to at least $\\approx 10$ Mpc. However, objects of this type could possibly -- depending on their compactness and exact distance -- be mistaken for single stars when observed in the commonly used $V$ and $I$ filters alone, since the $V-I$ colour of such a population is predicted to be $V-I\\approx 1.7$ (assuming $\\tau=1$ Gyr and an age of 10 Gyr), similar to that of RGB stars. Multicolour data should nonetheless be able to reveal their exotic nature, since the spectra of integrated stellar populations typically differ substantially from those of individual stars.  The idea of red halo star clusters at these luminosities is very interesting in the light of recent work by \\citet{Yan et al.}, who in the halo of the nearby galaxy M60 (distance $\\approx 16$ Mpc) may have detected a population of halo objects with luminosities of giants, but $izJHK$ colours that are difficult to reconcile with current models of such stars. Yan et al. argue that these objects, whatever their nature, may be responsible for the anomalous colours of red halos. It remains to be investigated, however, whether the colours of these objects may be consistent with the expectations for star clusters with bottom-heavy IMFs.  \\subsection{The typicality of the Milky Way}  Since the analysis by Z04 does not directly reveal what the variance of halo properties within the stacked sample is, we do not know what fraction of the stacked galaxies have a red halo. Perhaps some do not, and the Milky Way is just one such case. If so, we are confronted with red halos possibly contributing significantly to the baryonic mass in some disk galaxies -- but not in others. In disks at least, the baryonic Tully-Fisher relation (i.e. the relation between the total inferred stellar and gas masses vs. the rotational velocity) can be tuned via the stellar population mass-to-light ratio to show very small scatter \\citep[][]{McGaugh et al.,McGaugh}, which has been used as an argument that most of the baryons of disk galaxies have already been identified. Having considerable baryonic reservoirs in red halos for some disk galaxies, but not for others, would supposedly increase the scatter, but the magnitude of this effect is unclear. Moreover, there are observational indications that the Milky Way may be offset from the baryonic Tully-Fisher relation \\citep{Flynn et al.}. If this is the case, constraints based on the Milky Way may not be able to say anything definitive on the red halo masses of disk galaxies in general."
    },
    "0809/0809.2042_arXiv.txt": {
        "abstract": "We explore the amount of obscured star-formation as a function of environment in the A901/902 supercluster at $z=0.165$ in conjunction with a field sample drawn from the A901 and CDFS fields, imaged with {\\it HST} as part of the STAGES and GEMS surveys. We combine the \\combo\\ near-UV/optical SED with \\spitzer\\ 24\\micron\\ photometry to estimate both the unobscured and obscured star formation in galaxies with $\\rm M_\\ast > 10^{10}M_\\odot$. We find that the star formation activity in massive galaxies is suppressed in dense environments, in agreement with previous studies. Yet, nearly 40\\% of the star-forming galaxies have red optical colors at intermediate and high densities. These red systems are not starbursting; they have star formation rates per unit stellar mass similar to or lower than blue star-forming galaxies. More than half of the red star-forming galaxies have low IR-to-UV luminosity ratios, relatively high Sersic indices and they are equally abundant at all densities. They might be gradually quenching their star-formation, possibly but not necessarily under the influence of gas-removing environmental processes. The other $\\ga$40\\% of the red star-forming galaxies have high IR-to-UV luminosity ratios, indicative of high dust obscuration. They have relatively high specific star formation rates and are more abundant at intermediate densities. Our results indicate that while there is an overall suppression in the star-forming galaxy fraction with density, the small amount of star formation surviving the cluster environment is to a large extent obscured, suggesting that environmental interactions trigger a phase of obscured star formation, before complete quenching. ",
        "introduction": "\\label{sec:intro} Much observational evidence gathered so far has established that the environment in which galaxies live plays an important role in shaping their properties, such as their star formation activity, gas content, and morphology, in the sense that galaxies in regions of high galaxy density tend to have less ongoing star formation, less cold gas and more bulge-dominated morphology \\citep{oemler74,dressler80,lewis02,gavazzi02,gomez03,balogh04a,kauffmann04,mcintosh04,baldry06}. Yet, a real concern is that most star formation indicators used to date are based on optical properties and are susceptible to the effects of dust attenuation. Indeed a number of studies using mid-infrared or radio-derived star formation rates (SFRs) have found evidence for some unexpectedly intense bursts of star formation in intermediate-density regions \\citep[e.g.][]{MO02,coia05,fadda08}. The object of this paper is to use wide-field photometric redshift data, deep \\spitzer\\ data and wide-field HST imaging of the $z=0.165$ Abell 901/902 supercluster to explore the incidence of dust-obscured star formation for low-SFR galaxies: is dust-obscured star formation important even at low SFRs, and how does it vary with environment?  \\subsection{Environment, star formation and morphology}\\label{sec:intro1} Historically, the first clear evidence that environment influences galaxy properties is the observed predominance of early-type galaxies in low redshift clusters with respect to the field, alongside with a paucity of late-type, emission-line galaxies \\citep[e.g.][]{morgan61,dressler80}. The so-called morphology-density relation appears to be in place already at $z\\sim1$, but varies quantitatively with redshift: between $z\\sim0.5$ and the present, the fraction of late-type spirals in intermediate-density regions decreases in favour of the population of S0 galaxies \\citep{dressler97,smith05,postman05}. This has suggested that spiral galaxies evolve into smooth and passive systems such as S0s as they enter the dense environment of galaxy clusters.  Connected to the morphology-density relation is the decrease of the average SFR, as derived from optical colors or emission lines, with increasing environmental density \\citep[e.g.][]{balogh98,gavazzi02,pimbblet02,gomez03}. Among the two, the relation between color (or stellar age) and environment appears to be the most fundamental one: at fixed color, morphology shows only a weak residual dependence on environment \\citep{blanton05,wolf07}. Moreover, the link between the morphology-density and the SFR-density relations has significant scatter: not all spirals in clusters appear to be star-forming, at least on the basis of their optical spectra \\citep{poggianti99,goto03}. The question remains whether these spirals are really passive or they have star formation activity that escapes detection in the optical. Indeed, selection of passive spirals on the basis of their emission lines can be contaminated by dusty early-type spirals with low level of star formation activity, that could instead be detected using e.g. mid-infrared colors \\citep{wilman08}.  The SFR-density relation extends to very low local galaxy number densities \\citep[e.g.][]{lewis02,gomez03} and dark matter densities \\citep{gray04}, corresponding to the outskirts of clusters and the densities of groups. This suggests that not only the cores of clusters impact galaxy properties but galaxies may experience significant pre-processing in systems with lower density and lower velocity dispersion such as groups, before entering the denser and hotter environment of the cluster \\citep[e.g.][]{zabludoff02,fujita04}.  \\subsection{Environmental physical processes}\\label{sec:intro2} Several processes can act on galaxies as they interact with their surrounding environment. The intensity and timescale of individual processes may also vary with galaxy mass and during the galaxy lifetime as it moves through different density environments \\citep[for a review see][]{boselli06}. The gas content and hence star formation activity of galaxies can be affected by interaction with the intra-cluster medium (ICM).  The cold gas reservoir can be stripped due to the ram-pressure experienced by galaxies falling at high velocities in the dense ICM of the cluster \\citep{gunngott72,quilis00}. Ram-pressure stripping can lead to fast truncation of star formation and its action can be recognized from truncated H$\\alpha$ profiles \\citep{koopmann04}, asymmetric gas distribution and deficiency in the cold HI gas \\citep{giovanelli85,cayatte90,solanes01} of many spiral galaxies in local clusters. It is possible that on the front of compression of the cold gas due to ram-pressure a burst of star formation is induced \\citep[e.g.][]{gavazzi85,gavazzi03}. Another gas-stripping process, that affects star formation on longer timescales (of few Gyrs) than ram-pressure, is the so-called `strangulation' or `starvation': assuming that galaxies are surrounded by a halo of hot diffuse gas, this can be removed when galaxies become satellites of larger dark matter halos \\citep{larson80,balogh00}. Star formation can continue consuming the cold disk gas, but will eventually die out for the lack of supply of new fresh gas.  The gas distribution, star formation activity and morphology of galaxies can be altered also via interaction with other galaxies. Mergers between two equally-massive gas-rich galaxies can lead to the formation of a spheroidal system \\citep[e.g.][]{toomre72,barnes88,kauffmann93}. The merger can trigger an intense burst of star formation \\citep[e.g.][]{kennicutt87}, rapidly consuming the cold gas and then exhausting due to feedback processes \\citep{springel05}. Merging and slow galaxy-galaxy encounters are favoured in groups and in the infall region of clusters \\citep[e.g.][]{moss06}. At higher densities, galaxies can be affected by the cumulative effect of several rapid encounters with other cluster members, a mechanism known as `galaxy harassment' \\citep{moore98}. After a transient burst of star formation, galaxy harassment leads to substantial change in morphology. This mechanism can start to operate at intermediate densities, inducing density fluctuations in the gas \\citep{porter08}.  \\subsection{Dust-obscured star formation in dense environments?}\\label{sec:intro3} The net effect of the various mechanisms of interaction of galaxies with environment is an accelerated depletion or exhaustion of the gas reservoir and hence a suppression of the star formation activity. Many of these mechanisms, however, can lead to a temporary enhancement of star formation, either due to gas compression (e.g. ram-pressure) or density fluctuations that funnel the gas toward the center triggering nuclear activity (e.g. tidal interactions). The gas and dust column density is likely to increase during such processes and star formation can be to a large extent obscured and escape optical detection. Star formation indicators that are not affected by dust attenuation need to be adopted in order to quantify the occurrence of these obscured star formation episodes.  Already several studies based on observations in the thermal infrared (IR) or in the radio have identified significant populations of IR-bright or radio-bright galaxies in the outer regions of nearby and intermediate-redshift galaxy clusters \\citep[e.g.][]{smail99,MO02,MO03,best04,coia05}. \\cite{MO02} find that up to 20\\% of the galaxies in 20 nearby Abell clusters have centrally-concentrated dust-obscured star formation. These galaxies have different spatial distribution with respect to normal star-forming galaxies or active galactic nuclei (AGN): they are preferentially found in intermediate-density regions. In the A901/902 cluster at $z=0.165$, \\cite{WGM05} have identified an excess of dusty red galaxies with young stellar populations in the intermediate-density, infalling region of the cluster. Other studies have identified a population of red star-forming galaxies both in the field \\citep{hammer97} and in clusters \\citep{verdugo07}. These galaxies could be mistakenly classified as post-starburst on the basis of their weak emission lines \\citep{poggianti99,bekki01}. It is interesting to notice that populations of red, IR-bright star-forming galaxies are often found in filaments \\citep[e.g.][]{fadda00,fadda08,porter08} and in unvirialized or merging clusters \\citep[e.g.][]{MO03,geach06,moran07}. Significant populations of starburst, IR-bright galaxies have been also found in a dynamically young cluster at $z=0.83$ by \\cite{marcillac07}. These systems could in fact be more abundant at higher redshift \\citep{saintonge08} as expected from the increase in cosmic star formation activity. Recently, \\cite{elbaz07} have shown that the detection of these galaxies with the use of dust-independent SFR indicators can even lead to a reversal of the star-formation--density relation at $z\\sim1$.  In this work we want to explore as a function of local galaxy density the importance in the local Universe of the star formation `hidden' among red galaxies, that would be missed by optical, dust-sensitive SFR indicators. Uniquely, we wish to push to modest SFRs ($\\rm \\sim 0.2M_\\odot~yr^{-1}$), in order to constrain the star formation mode of typical (not rare starbursting) systems. There are two key requirements for such a study: 1) obscuration-free SFR indicators, ideally given by the combination of deep thermal IR and UV, in order to obtain a complete census of the total (obscured and unobscured) SFR; 2) a long baseline in environmental density covering from the cluster cores to the field in order to quantitatively characterize the SFR-density relation. We analyse the \\combo\\ CDFS and A901 fields at $z<0.3$, complementing the UV/optical photometry from \\combo\\ with {\\it Spitzer} 24\\micron\\ data and with {\\it HST} V-band imaging from the Galaxy Evolution from Morphology and SEDs (GEMS) survey and the Space Telescope A901/902 Galaxy Evolution Survey (STAGES). The A901 field is particularly interesting in that it contains the supercluster A901/902 at $z=0.165$, a complex system with four main substructures probably in the process of accreting or merging, where mechanisms altering the star formation and morphological properties of galaxies might be favoured \\citep[e.g.][]{gray02,gray04,gray08,WGM05,wolf07,heymans08}.  We present the sample and the data in Section~\\ref{sec:sample} and describe the derivation of SFR and environmental density in Sections~\\ref{sec:parameters} and~\\ref{sec:dens}. After discussing the classification into star-forming and quiescent galaxies in Section~\\ref{sec:fractions}, we explore the dependence on local galaxy density of the fraction of (obscured and unobscured) star-forming galaxies and their contribution to the total star-formation activity as a function of environment (Section~\\ref{sec:fractions2}). The properties of red star-forming galaxies, such as their SFR, mass, morphology and dust attenuation, are compared to those of unobscured star-forming galaxies in Section~\\ref{sec:redSF}. We summarize and discuss our results in Section~\\ref{sec:conclusion}. Throughout the paper we assume a cosmology with $\\rm \\Omega_m=0.3$, $\\rm \\Omega_\\Lambda=0.7$ and $\\rm H_0=70km~s^{-1}~Mpc^{-1}$. ",
        "conclusions": "\\label{sec:conclusion}  \\subsection{Star formation among red galaxies} We have combined \\combo\\ optical data with MIPS 24\\micron\\ data for a sample of low-redshift ($0.05<z<0.3$) galaxies in the CDFS and A901 fields with the aim of studying the occurrence of obscured star formation as a function of environment. The 24\\micron\\ information allows us to recover directly the flux from young stellar populations absorbed and re-emitted by dust, and thus to trace, in combination with the UV/optical information, the total (unobscured and obscured) star formation activity in galaxies. The A901 field is particularly suited from this kind of analysis, not only for the exceptional multiwavelength coverage, but also because it includes the complex A901/902 supercluster at $z=0.165$ extending over an area of 5$\\times$5~Mpc$^2$h$_{70}^{-2}$. The supercluster is composed of four main substructures, probably in the process of merging. The complex dynamical state of the A901/902 supercluster potentially makes it an ideal case for identifying galaxies in their process of evolution under the influence of environment. The CDFS, with the same multiwavelength coverage, offers instead a control sample of field galaxies at similar redshift as the cluster.  In this work we have focused on galaxies with stellar masses larger than $10^{10}M_\\odot$, above which the red sequence is complete out to $z=0.3$ (our limiting redshift). This mass limit roughly corresponds to $0.1\\times M^\\ast$ over the redshift range $z<0.3$ \\citep[]{bell03b,borch06}. We define as star-forming those galaxies with a specific SFR (derived from UV and IR luminosities) above $\\rm 2\\times 10^{-11} yr^{-1}$. Our focus is on star-forming galaxies populating the red sequence, either because they show low levels of star formation insufficient to alter the color of the underlying older population or because their star formation activity is highly obscured by dust. Studies based on the UV and optical emission of galaxies have identified a significant amount of low-level star-formation in low-mass ellipticals \\citep[]{yi05,kaviraj07}, with a hint of a peak in `frosting' activity at group densities \\citep[]{rogers07}. Star formation indicators that are less sensitive to dust attenuation, such as the 24\\micron\\ emission that we exploit in this work, are instead required to detect dust-obscured star formation.  We have studied the abundance of blue and red star-forming galaxies as a function of environment, as expressed by the galaxy number overdensity in a radius of 0.25~Mpc, focusing on the contribution of star formation on the red sequence compared to optically-detectable star formation. Our results can be summarized as follows. \\begin{itemize} \\item[-] The overall fraction of star-forming galaxies decreases from $\\sim$60\\% in underdense regions to $\\sim$20\\% in high-density regions. The stellar mass fraction contributed by star-forming galaxies also decreases going from the field to the cluster cores. The decline is steeper than for the number fraction because, while no significant environmental evolution in stellar mass occurs for star-forming galaxies, the mass function of quiescent galaxies reaches higher stellar masses at higher densities. \\item[-] The fraction of blue star-forming galaxies decreases monotonically from $\\sim$40\\% at low densities to less than 20\\% at higher densities. On the contrary, red SF galaxies do not show a monotonic behaviour as a function of environment. After an initial decline of the red SF fraction from the field to higher \\dens, we identify an overabundance of obscured star formation at intermediate densities, those typical of the outskirts of the A901/902 supercluster cores. At both intermediate and high densities, red SF galaxies represent 40\\% of all star-forming galaxies and contribute $20-30$\\% of the total star formation activity at these densities. \\end{itemize}  To a first order, our results confirm the well-known SFR-density relation \\citep[e.g.][]{gavazzi02,lewis02,gomez03,balogh04b,kauffmann04} and morphology-density relation \\citep[e.g.]{dressler80,dressler97,treu03,vdW08}. In addition to this, we find a significant contribution by red-sequence galaxies, identified as star-forming through their IR emission, to the total star formation activity up to the highest densities of the cluster. This would be at least partly missed by optical studies. This result is consistent with \\cite{WGM05} who already found an enhancement of optically-classified dusty red galaxies in the medium-density outskirts of the A901/902 supercluster. In this work, supported by deep 24\\micron\\ data, we directly measures the amount of star formation going on in these galaxies.  Our results are also in line with recent studies of clusters at similar redshifts as A901/902 or higher, which have identified a population of IR-bright galaxies in filaments and the infalling regions of the clusters. \\cite{fadda00} found a population of 15\\micron-detected galaxies with high 15\\micron-to-optical flux ratio suggesting star formation activity in the cluster A1689 at $z=0.18$, in excess with respect to the Virgo and Coma cluster \\citep[see also][]{duc02}. In the cluster A2667 at $z=0.23$ \\cite{cortese07} have identified a IR-bright $L^\\ast$ spiral galaxy in the process of being transformed by the cluster environment which triggers an intense burst of star formation. At similar redshift, \\cite{fadda08} find two filamentary structures in the outskirts of the A1763 cluster at $z=0.23$ (probably undergoing accretion events), which are rich in actively star-forming galaxies. \\cite{geach06} find an excess of mid-infrared sources in an unvirialized cluster at $z\\sim0.4$, where star formation might be triggered via mergers or interactions between gas-rich spirals. However, they also note that significant cluster-to-cluster variations are possible: they do not find any significant excess in another cluster at similar redshift, of similar mass but with a hotter and smoother ICM. Moving to higher redshift, \\cite{marcillac07} studied 24\\micron\\ sources in a massive, dynamically young, unvirialized cluster at $z=0.83$. They find that IR-detected galaxies tend to lie in the outskirts of the cluster, while they avoid the merging region. Finally, \\cite{elbaz07}, utilizing 24\\micron\\ imaging in the GOODS fields at redshift $0.8<z<1.2$, have identified for the first time a reversal of the SFR-density relation observed at lower redshifts. This result has been recently confirmed by \\cite{cooper08} with a spectroscopic analysis using DEEP2 data.   \\subsection{Dusty or old?} The relative abundance of red SF galaxies at intermediate and high densities suggests that they are transforming under the influence of some environment-related process. What are the star formation activity, morphology and dust attenuation of these red star-forming galaxies? \\begin{itemize} \\item[-] The red SF galaxies in our sample are not in a starburst phase. The few starburst galaxies (with $\\rm \\log (SFR/M_\\ast)>-9.7$, corresponding to a birthrate parameter $b>1$, assuming a formation redshift of 4) in our sample all populate the blue cloud. We find that red SF galaxies have similar SFR as blue SF galaxies, and slightly lower specific SFR. While the overall fraction of star-forming galaxies decreases with density, we do not identify any significant evolution in their level of activity, either obscured or not. \\item[-] The morphology of star-forming galaxies is not very sensitive to their color. Red SF have similar distribution in Sersic index as blue SF: they are predominantly disc-dominated. Moreover, the morphology of star-forming galaxies depends little on the environment in which they live. This suggests that, on average, changes in stellar populations and changes in morphology happen on different timescales, as hinted at by the fact that color seems to be more sensitive to environment than morphology \\citep{blanton05}. The rise of red massive spirals in the infalling regions of the A901/902 cluster has also been interpreted by \\cite{wolf08} as due to SFR decline not accompanied by morphological change. A two-step scenario in which star-formation is quenched first and morphological transformation follows on longer timescale is also supported by the analysis of \\cite{sanchez07} of the A2218 cluster at $z=0.17$. \\item[-] Red SF galaxies have IR-to-UV luminosity ratios ($\\rm L_{IR}/L_{UV}$), a proxy for the level of UV attenuation by dust, on average higher than blue SF galaxies. The distribution in their IR-to-UV luminosity ratios suggests however the presence of two different populations, hence possibly two distinct mechanisms affecting star formation activity of red galaxies. Roughly half of the red SF galaxies in our sample have relatively low $\\rm L_{IR}/L_{UV}$, similar to the average value of the bulk of blue SF galaxies, without evolution with environment. The other half of the red SF galaxies have instead systematically higher $\\rm L_{IR}/L_{UV}$. The range in dust attenuation of this second population becomes narrower at higher densities, suggesting a trend of decreasing attenuation with density. \\end{itemize}  On the basis of the IR properties of red SF galaxies we tentatively distinguish them into two subpopulations. Low-attenuation red SF galaxies have low specific SFR ($\\la10^{-10.3}$) independent of environment. Among star-forming galaxies they tend to have higher Sersic indices ($\\langle n\\rangle\\sim2.5$). These properties suggest that these galaxies are dominated by rather old stellar populations but have some residual star formation. They {\\it could} be anemic/gas-deficient spirals gradually suppressing their star formation as a consequence of the removal of their gas reservoir as they move into higher-density environments \\citep{fumagalli_gavazzi08}. Their star formation could be suppressed on relatively long timescales (of few Gyrs) if strangulation of the hot, diffuse gas occurs while the galaxies enter a more massive halo \\citep[e.g.][]{balogh00,vdB08}. The gradual fading of the disc would make the morphology of these galaxies appear of earlier type. In addition to strangulation, when the density of the surrounding medium becomes sufficiently high, ram-pressure can act on smaller mass galaxies to remove the remaining gas on the disc and lead to fast quenching \\citep[e.g.][]{gunngott72,quilis00,boselli06b}. However, there is no significant evidence that the relative abundance of low-attenuation red SF galaxies varies with environment. Therefore, we cannot exclude that these galaxies are suppressing their SF due to internal feedback processes only, without any additional environmental action required.  The other subpopulation of red SF galaxies have systematically higher $\\rm L_{IR}/L_{UV}$, indicative of higher levels of dust attenuation.  They represent $\\ga$40\\% of all red SF galaxies in the sample, even after accounting for purely edge-on spirals. By visual inspection of their {\\it HST} V-band images, we can say that the majority of them are spiral galaxies with a bright nucleus or inner bar/disk, suggesting intense star formation activity in the galaxy core (we cannot exclude AGN contribution in some cases). We also find few cases of interacting galaxies and merger remnants. In comparison to the low-attenuation red SF class discussed above, they have systematically higher specific SFR and lower values of Sersic index ($\\langle n\\rangle\\sim1.5$). As opposed to low-attenuation red SF galaxies, they tend to be more abundant at intermediate densities where their stellar mass is $\\sim$50\\% and $\\sim$70\\% higher than at low and high densities, respectively. This suggests that environmental interactions are particularly efficient in triggering episodes of obscured, often centrally concentrated, star formation in these massive late-type spirals. While we find few cases of interacting galaxies, violent processes, such as mergers, leading to intense starbursts cannot be the dominant phenomenon. These galaxies are likely more sensitive to more gentle mechanisms that perturb the distribution of gas inducing star formation (but not a starburst) and at the same time increase the gas/dust column density. This process should not alter morphology as long as star formation is still detectable. The fact that galaxies undergoing this phase are preferentially found at intermediate densities and with relatively high stellar masses might indicate longer duration of the dust-obscured episode of SF for more massive galaxies, which are thus more likely to be caught in this phase than low-mass galaxies. Harassment can act on massive spirals, funnelling the gas toward the center and leading to a temporary enhancement of star formation \\citep{moore98,lake98}. The timescales of this process could be relatively long if it occurs at group-like densities, rather than in the cluster. Tidal interactions between galaxies at low and intermediate densities can also produce gas funnelling toward the center \\citep{mihos04}. \\\\  In summary, we have identified a significant amount of star formation `hidden' among red-sequence galaxies, contributing at least 30\\% to the total star formation activity at intermediate and high densities. The red SF population is composed partly of disc galaxies  dominated by old stellar populations and with low-level residual star formation, and partly of spirals or irregular galaxies undergoing modest (non-starburst) episodes of dust-obscured star formation. This means that, while we confirm the general suppression of star formation with increasing environmental density, the small amount of star formation surviving the cluster happens to a large extent in galaxies either obscured or dominated by old stellar populations. Low-attenuation red SF galaxies seem to be a ubiquitous population at all densities. Therefore an environmental action is not necessarily required to explain their ongoing low-level star formation. On the contrary, dusty SF galaxies are relatively more abundant at intermediate densities. They might be experiencing harassment or tidal interactions with other galaxies, which funnel gas toward the center inducing a (partly or totally obscured) episode of star formation. Ram-pressure can also be partly responsible for the population of relatively more massive dusty SF galaxies in the cluster: while it is not effective in removing the gas from the disc in massive galaxies, it could perturb it inducing obscured star formation. The complex dynamical state of the A901/902 supercluster could favour a combination of different processes producing a temporary enhancement of obscured star formation."
    },
    "0809/0809.0792_arXiv.txt": {
        "abstract": "We study a scenario that a hidden gauge boson constitutes the dominant component of dark matter and decays into the standard model particles through a gauge kinetic mixing. Interestingly, gamma rays and positrons produced from the decay of hidden gauge boson can explain both the EGRET excess of diffuse gamma rays and the HEAT anomaly in the positron fraction. The spectra of the gamma rays and the positrons have distinctive features; the absence of line emission of the gamma ray and a sharp peak in the positron fraction. Such features may be observed by the FGST and PAMELA satellites. ",
        "introduction": "\\label{sec:1} % The concept of symmetry has been the guiding principle in modern physics. The structure of the standard model (SM) is dictated by $SU(3)_{C}\\times SU(2)_{L}\\times U(1)_{Y}$ gauge symmetries. The electroweak symmetry, $SU(2)_{L}\\times U(1)_{Y}$, is spontaneously broken by a non-vanishing vacuum expectation value (vev) of the Higgs boson, and massive $W$ and $Z$ bosons are generated. It is quite natural in the string landscape that there are many other gauge symmetries as well as discrete ones realized in nature, and some of the gauge symmetries may be spontaneously broken, leading to massive gauge bosons, as in the SM.  Suppose that some of the hidden gauge bosons are `odd' under a $Z_{2}$ parity, whereas all the SM particles are `even'. If the parity is exact, and if the hidden gauge symmetries are spontaneously broken, the lightest parity-odd gauge boson is stable and can be a good candidate for cold dark matter. Some recent examples are the T-odd $U(1)$ gauge boson in the little Higgs model with T-parity~\\cite{ArkaniHamed:2002qy,Cheng:2003ju,Cheng:2004yc,Csaki:2008se} and the KK photon in the minimal universal extra dimension model~\\cite{Appelquist:2000nn}. However, it is not known yet which of unbroken or (explicitly or spontaneously) broken discrete symmetries are more common in the string landscape. If a discrete symmetry breaking is a general phenomenon, we may expect that the dark matter is not absolutely stable and ultimately decays into the SM particles.  In this paper we consider a simplest case that there exist a hidden $U(1)^{\\prime}$ gauge symmetry which is spontaneously broken and a parity under which only the hidden gauge boson changes its sign. We assume that the parity is broken by a kinetic mixing between the $U(1)^{\\prime}$ and the SM $U(1)_{Y}$ gauge symmetries\\ \\cite{Holdom:1985ag,Foot:1991kb}. As a result, the hidden gauge boson decays into the SM particles through the parity violating interactions induced by the kinetic mixing. If such a violation is so tiny that the lifetime of the hidden gauge boson is much longer than the age of our universe, the hidden gauge boson can be the dominant component of dark matter. Furthermore, the subsequent decays of the SM particles will form continuous spectra of the gamma rays and the positrons in the high-energy cosmic ray. With an appropriate amount of the parity violation, we see that those gamma rays and positrons could be the source of the excesses of the gamma ray observed by Energetic Gamma Ray Experiment Telescope (EGRET)~\\cite{Sreekumar:1997un,Strong:2004ry} and the positron flux observed by High Energy Antimatter Telescope (HEAT)\\ \\cite{Barwick:1997ig}, MASS\\ \\cite{Grimani:2002yz} and AMS\\ \\cite{Aguilar:2007yf} experiments.  Recently, the gravitino dark matter scenario was extensively studied in the framework of supersymmetry with $R$-parity violation~\\cite{Takayama:2000uz,Buchmuller:2007ui}. The decay of the gravitino into the gamma rays was discussed in Refs.~\\cite{Takayama:2000uz,Buchmuller:2007ui}. Furthermore, if the gravitino is heavier than $W$ boson, it decays predominantly into a $W$ or $Z$ boson and a lepton. With the gravitino mass of ${\\cal O}(100)$ GeV and the lifetime of ${\\cal O}(10^{26})$ sec, the gravitino decay can simultaneously explain the EGRET anomaly in the extragalactic diffuse gamma ray background and the HEAT excess in the positron fraction\\ \\cite{Bertone:2007aw,Ibarra:2007wg,Ibarra:2008qg,Ishiwata:2008cu}. Our decaying hidden gauge boson has therefore many parallels with the gravitino with $R$-parity violation. However, in our model, the decay branching ratio of the hidden gauge boson to $W$ boson is highly suppressed, which is significantly different from the gravitino case.  This paper is organized as follows. In Sec.\\ \\ref{sec:model}, we present the effective Lagrangian and Feynman rules used in our calculations. In Sec.\\ \\ref{sec:numerical}, we calculate the spectra of gamma ray and positron flux from the decay of the hidden gauge boson and compare them with the observed data. Sec.\\ \\ref{sec:Conclusion} is devoted to discussions and conclusions. ",
        "conclusions": "\\label{sec:Conclusion} % Let us here discuss briefly how such a small kinetic mixing in Eq.~(\\ref{eq:kinetic}) can arise. If there is an unbroken $Z_{2}$ parity symmetry under which the hidden gauge boson flips its sign, the kinetic mixing would be forbidden. This parity may arise from a more fundamental symmetry or it may be an accidental one. In order to have a non-vanishing kinetic mixing, we need to break the parity by a small amount and transmit the breaking to the kinetic mixing~% \\footnote{ With the presence of the parity symmetry, a small breaking is natural in the sense of `t Hooft. % }. To this end we introduce messenger fields $\\xi_{1}$ and $\\xi_{2}$ that have both $U(1)_{Y}$ and $U(1)^{\\prime}$ charges as $(1,1)$ and $(1,-1)$, respectively. Under the parity, they transform as: \\begin{equation} \\xi_{1}\\leftrightarrow\\xi_{2}.\\,\\label{eq:parity_messenger}\\end{equation} Note that the charge assignment is consistent with the parity transformation (see the discussion below Eq.\\ (\\ref{eq:gauge_scalar})). We assume that the messenger mass scale $M_{m}$ is very large, e.g. the grand unified theory (GUT) scale. Suppose that the parity is broken in the messenger sector in such a way that $\\xi_{1}$ and $\\xi_{2}$ obtain different masses. Let us denote the mass difference by $\\Delta M_{m}$, which parametrizes the amount of the parity violation. Integrating out the messengers, we are then left with a small kinetic mixing, $\\epsilon$, given by % \\begin{equation} \\epsilon\\;\\sim\\;\\frac{g_{h}g^{\\prime}}{16\\pi^{2}}\\frac{\\Delta M_{m}^{2}}{M_{m}^{2}},\\end{equation} where $g_{h}$ is the coupling constant of the hidden $U(1)$ gauge symmetry. For $\\Delta M_{m}\\sim{\\cal O}({\\rm TeV})$ and $M_{m}\\sim10^{15}{\\rm \\, GeV}$, we obtain the tiny kinetic mixing of the right magnitude of ${\\cal O}(10^{-26})$.  It is tempting to identify the origin of the parity violation with the spontaneous breaking of the $U(1)^{\\prime}$. As an illustration, we will present a toy model below. As we will see later, there are some dangerous couplings which may spoil the stability of the $\\ap$ in this model. However, these problems could be solved by embedding the model into a theory with supersymmetry or an extra dimension(s).  Suppose that all the matter fields in the hidden sector are neutral under the parity. Then, the gauge field $B_{\\mu}^{\\prime}$ cannot have any interactions with those hidden matter fields, since they are forbidden by the parity. Let us introduce two scalar fields, $\\phi_{1}$ and $\\phi_{2}$, which transform under the parity as \\begin{equation} \\phi_{1}\\leftrightarrow\\phi_{2},\\label{eq:parity_scalar}\\end{equation} and we assume that $\\phi_{1}$ and $\\phi_{2}$ are neutral under the SM gauge symmetries. Then the following interactions are allowed:\\begin{equation} {\\cal L}\\supset ig_{h}B_{\\mu}^{\\prime}(\\phi_{1}^{*}\\partial^{\\mu}\\phi_{1}-h.c.)-ig_{h}B_{\\mu}^{\\prime}(\\phi_{2}^{*}\\partial^{\\mu}\\phi_{2}-h.c.).\\,\\label{eq:gauge_scalar}\\end{equation} The parity is spontaneously broken if one of the two scalars develops a non-vanishing vev. To this end, we consider the following potential:\\begin{equation} \\lambda(|\\phi_{1}|^{2}-v_{h}^{2})^{2}+\\lambda(|\\phi_{2}|^{2}-v_{h}^{2})^{2}+2\\kappa|\\phi_{1}|^{2}|\\phi_{2}|^{2},\\,\\label{eq:potential}\\end{equation} where $v_{h}$ , $\\lambda$ and $\\kappa$ are real and positive. For $\\kappa>\\lambda$, there are four distinct vacua, $(\\phi_{1},\\phi_{2})=(\\pm v_{h},0)$ and $(0,\\pm v_{h})$. We take one possibility of them, $(\\phi_{1},\\phi_{2})=(v_{h},0)$, as an example. Therefore, the hidden $U(1)$ gauge symmetry is spontaneously broken by the vev of $\\phi_{1}$, and the associated gauge boson acquires a mass, $m=\\sqrt{2}g_{h}v_{h}$, by eating the imaginary component of $\\phi_{1}$.  In order to transmit the parity violation, we introduce couplings between $\\phi_{1,2}$ and the messenger fields, % \\begin{equation} -{\\cal L}\\supset M_{m}^{2}(|\\xi_{1}|^{2}+|\\xi_{2}|^{2})+\\kappa^{\\prime}(|\\phi_{1}|^{2}|\\xi_{1}|^{2}+|\\phi_{2}|^{2}|\\xi_{2}|^{2}),\\,\\label{eq:Scalar_messengers}\\end{equation} where $M_{m}$ is the messenger mass, and $\\kappa^{\\prime}$ is a real and positive constant. After the $U(1)^{\\prime}$ is spontaneously broken, masses of $\\xi_{1}$ and $\\xi_{2}$ are slightly different: $m_{\\xi_{1}}^{2}=M_{m}^{2}+\\kappa^{\\prime}v_{h}^{2}$ and $m_{\\xi_{2}}^{2}=M_{m}^{2}$. After integrating out these heavy messengers, we obtain the kinetic mixing in Eq. (\\ref{eq:kinetic}) with $\\epsilon$ given by\\[ \\epsilon\\sim\\frac{g_{h}g^{\\prime}}{16\\pi^{2}}\\frac{\\kappa^{\\prime}v_{h}^{2}}{M_{m}^{2}}.\\] For $g_{h}\\sim\\kappa^{\\prime}\\sim{\\cal O}(1)$, $v_{h}\\sim1{\\rm \\, TeV}$, and $M_{m}\\sim10^{15}{\\rm \\, GeV}$, we obtain $\\epsilon\\sim10^{-26}$. We can therefore realize more or less the correct magnitude of $\\epsilon$ needed to account for the excesses of the gamma rays and the positrons in this toy model.  In the above toy model, we have assumed that the SM particles couple to the hidden sector only through the messenger fields. It is also possible to introduce direct interactions between $\\phi_{1,2}$ and the visible sector. For instance, we can couple them to the SM fermions as \\begin{equation} {\\cal L}\\supset\\frac{\\alpha}{M_{p}^{2}}(|\\phi_{1}|^{2}+|\\phi_{2}|^{2})\\bar{f}_{L}f_{R}H+{\\rm h.c.}\\,,\\label{eq:scalar_SM}\\end{equation} where $\\alpha\\sim{\\cal O}(1)$, $M_{p}\\simeq2.4\\times10^{18}{\\rm \\, GeV}$ is the reduced Planck scale, and $f_{R(L)}$ is the right-(left-)handed fermion and $H$ is the Higgs field. Although this interaction does not lead to the gauge kinetic mixing, it induces the decay of the hidden gauge boson into a fermion pair after breaking the parity. However, the decay branching ratio through such interaction is negligible, compared to that through the kinetic mixing with $\\epsilon\\sim10^{-26}$.  More dangerous direct couplings are those between $\\phi_{1,2}$ and the Higgs field, $(|\\phi_{1}|^{2}+|\\phi_{2}|^{2})|H|^{2}$. If there exist such interactions, the $\\ap$ could decay into the Higgs bosons immediately. If we extend the model into a supersymmetric one, the direct couplings to the Higgs bosons can be suppressed with the aid of the $R$-symmetry. In a theory with an extra dimension, it is also possible to suppress the direct couplings with appropriate configurations of the branes (e.g. the visible particles on one brane, while the hidden particles on the other).  It is also known that many $U(1)$ symmetries appear in the string theory, and the kinetic mixing can similarly arise in the low energy theory by integrating out heavy string states that have charges of two $U(1)$ gauge symmetries. It has been extensively studied how large the kinetic mixing can be in e.g. Ref.\\ \\cite{Dienes:1996zr} (see also Ref.~\\cite{Abel:2008ai} and references therein). For instance, in a warped background geometry, we can have an exponentially small kinetic mixing~\\cite{Abel:2008ai}.  Let us also have some comments on the the GUT. So far we have assumed the existence of the kinetic mixing between $U(1)_{Y}$ and the hidden $U(1)^{\\prime}$. If one of the $U(1)$ gauge symmetries actually sits within an unbroken non-Abelian gauge symmetry, such kinetic mixing is not allowed. This does not necessarily mean that kinetic mixings are incompatible with GUT, because the GUT gauge group is spontaneously broken. Indeed, kinetic mixings between the $U(1)^{\\prime}$ and the GUT gauge fields can arise below the GUT scale by picking up non-vanishing vevs of the Higgs bosons responsible for the GUT breaking.  Let us also discuss how the hidden gauge boson could be generated in the early universe to account for the observed abundance of dark matter. For simplicity, we neglect a numerical coefficient of order unity in the following discussion. In the presence of the messenger fields, there appears at one-loop level a following interaction between the $U(1)^{\\prime}$ and $U(1)_{Y}$ gauge fields: % \\begin{equation} {\\cal L}\\;\\supset\\;\\frac{g_{h}^{2}g^{\\prime2}}{16\\pi^{2}M_{m}^{4}}(B_{\\alpha\\beta}B^{\\alpha\\beta})(B_{\\gamma\\delta}^{\\prime}B^{\\prime\\gamma\\delta}).\\end{equation} The hidden gauge boson $A^{\\prime}$ will be produced through the above interaction most efficiently at the reheating. The $\\ap$ abundance is roughly estimated to be % \\begin{equation} \\frac{n_{\\ap}}{s}\\;\\sim\\;10^{-12}\\, g_{h}^{4}\\lrfp{T_{R}}{3\\cdot10^{13}{\\rm \\, GeV}}{7}\\lrfp{M_{m}}{10^{15}{\\rm \\, GeV}}{-8},\\end{equation} where $s$ and $T_{R}$ are the entropy density and the reheating temperature, respectively. For the mass $m_{\\ap}\\sim{\\cal O}(100){\\rm \\, GeV}$, a right abundance of $\\ap$ can be generated from the above interaction for $T_{R}\\sim10^{13}$ GeV. Also the $\\ap$ can be non-thermally produced by the inflaton decay~\\cite{Endo:2006qk,Endo:2007ih,Endo:2007sz}.  In the above model, the parity is spontaneously broken by the vev of $\\phi_{1}$. If the breaking occurs after inflation, domain walls connecting two of the four vacua in Eq.\\ (\\ref{eq:potential}) will be formed~\\cite{Zeldovich:1974uw,Kibble:1976sj,Vilenkin:1981zs}, which can be the cosmological disaster. There are several means to get around this problem, and one of which is to introduce a tiny explicit breaking of the parity symmetry. Therefore the domain walls are not stable, and eventually annihilate after collisions~\\cite{Vilenkin:1981zs,Coulson:1995nv,Larsson:1996sp}. Another solution is to assume that the breaking occurs before inflation. The last one is to assume that the initial positions of the scalars $\\phi_{1}$ and $\\phi_{2}$ are deviated from the origin. In the last case, the domain walls, if formed, will be annihilated eventually~% \\footnote{ The significant amount of the gravitational waves may be produced in the collisions of domain walls~\\cite{Takahashi:2008mu}. The hidden gauge bosons are also produced by the annihilation processes of the domain walls.% }.  In this paper we have considered a possibility that a hidden gauge boson $A'$, which constitutes the dominant component of dark matter, decays into the SM particles through the kinetic mixing term that breaks the $Z_{2}$ parity symmetry. As a result, the branching ratios are solely determined by the mass of the hidden gauge boson. Continuum spectra of photons and positrons are generated from $A'\\to\\ell^{+}\\ell^{-}$ ($\\ell=e,\\,\\mu,\\,\\tau$), and from the decays of hadrons, mainly $\\pi^{0,\\,\\pm}$, produced in the subsequent QCD hadronization process~% \\footnote{There also exist antiprotons produced from the $A'$ decay, whose flux can also be calculated and compared with the present data. However, we did not pursue the detailed comparison in this paper due to the large experimental uncertainties and our poor understanding of diffusion models.% }. If the mass of $A'$ is about ${\\cal O}(100)$ GeV and its lifetime is of order ${\\cal O}(10^{26})$ seconds, those gamma rays and positrons from $A'$ decay may account for the observed excesses in the extra galactic diffuse gamma ray flux and the positron fraction. Interestingly, in our model, the spectra of the gamma rays and the positrons have distinctive features: the absence of line emission of the gamma ray and a sharp peak in the positron fraction. Such features may be observed by the FGST and PAMELA satellites."
    },
    "0809/0809.2797_arXiv.txt": {
        "abstract": "Minimizing the scatter between cluster mass and accessible observables is an important goal for cluster cosmology. In this work, we introduce a new matched filter richness estimator, and test its performance using the maxBCG cluster catalog.  Our new estimator significantly reduces the variance in the $\\Lx-$richness relation, from $\\sigma_{\\lnl}^2=(0.86\\pm0.02)^2$ to $\\sigma_{\\lnl}^2=(0.69\\pm0.02)^2$.  Relative to the maxBCG richness estimate, it also removes the strong redshift dependence of the richness scaling relations, and is significantly more robust to photometric and redshift errors.  These improvements are largely due to our more sophisticated treatment of galaxy color data.  We also demonstrate the scatter in the $\\Lx-$richness relation depends on the aperture used to estimate cluster richness, and introduce a novel approach for optimizing said aperture which can be easily generalized to other mass tracers. ",
        "introduction": "The dependence of the halo mass function on cosmology is a problem that is well understood both analytically \\citep{pressschechter74,bondetal91,shethtormen02} and numerically \\citep{jenkinsetal01,warrenetal06,tinkeretal08}.  In principle, this detailed understanding allows one to place tight constraints on the amplitude of the primordial power spectrum and on dark energy parameters \\citep[e.g.][]{holderetal01,haimanetal01}.  In practice, life is not so simple. Cluster mass is not an observable, and so we must rely on other quantities that trace mass to estimate the halo mass function.  In this context, observables that are tightly correlated with mass and whose scatter is well understood are highly desirable, as they permit a more accurate measurement of the mass function.  One such mass tracer, and the subject of interest for this work, is the so called cluster richness, a measure of the galaxy content of a cluster.  Relative to other popular mass tracers such as X-ray properties, SZ-decrements, and galaxy velocity dispersion, optical richness has unique advantages and disadvantages. Its unique advantages are: \\begin{enumerate} \\item cluster richness can be easily estimated with inexpensive, photometric optical data. \\item cluster richness can be estimated for both massive clusters and low mass groups. \\end{enumerate} The first of these two properties is significant because it implies that cluster richness estimates are readily available given any large, photometric optical survey such as the SDSS \\citep{yorketal00} , DES\\footnote{http://www.darkenergysurvey.org/}, or LSST\\footnote{http://www.lsst.org/lsst\\_home.shtml}. The latter property, on the other hand, is an important advantage for a much more interesting reason.  Beginning with \\citet{whiteetal93}, cosmological constraints from galaxy clusters have been presented as a degeneracy relation $\\sigma_8\\Omega_m^{\\gamma}=constant$ where $\\gamma\\approx 0.5$, $\\sigma_8$ is a parameter specifying the amplitude of the primordial power spectrum, and $\\Omega_m$ is the matter density of the universe in units of the critical density. The existence of this degeneracy is easy to explain~\\citep[][]{rozoetal04}: suppose that we only measured the abundance of galaxy clusters at a single mass scale.  Since the halo mass function depends on both $\\sigma_8$ and $\\Omega_m$, it is evident that with just one observable there must be a degeneracy between these two parameters.  But what if we measure the halo mass function over a range of scales?  This is roughly equivalent to measuring the amplitude and slope of the halo mass function at the statistical pivot point.  If the mass range probed is small, then the slope of the mass function is not well constrained, and the degeneracy between $\\sigma_8$ and $\\Omega_m$ will remain.  In order to break this degeneracy, a measurement of the halo mass function over a large range of masses is necessary. Currently, only spectroscopic velocity measurements and optical richness estimates can probe a mass range wide enough to successfully break this degeneracy, but the former requires considerably more observing resources.  There are, however, important disadvantages to using cluster richness as a mass tracer. For instance, historically, the fact that the relation between cluster richness and mass cannot be predicted a priori based on simple physical arguments was viewed as a significant drawback.  Nowadays, however, this argument holds little sway, since the level of accuracy required for precision comsology in our a priori knowledge of cluster scaling relations is pushing current research towards a self-calibrating approach, in which both cosmology and cluster scaling relations are simultaneously constrained from the data \\citep[][]{limahu04,majumdaretal04, limahu05,hucohn06,wuetal08}.  Thus, in so far as self-calibration is necessary to insure one-self against possible biases in cosmological estimates, the lack of a simple physical model for predicting cluster richness is no longer a serious drawback.  Another reason why optical richness estimates fell out of favor relative to other mass tracers is that, in the past, richness estimates were known to suffer from significant projection effects, which resulted in impure cluster samples as well as large scatter in the mass-richness relation. Abell made one of the first systematic attempts at measuring richness \\citep{abell58,abell89} in defining his richness classes. He tried to minimize projection by only counting galaxies dimmer than $m_3$, the magnitude of the third brightest cluster galaxy, but brighter than $m_3+2$. The bright cut is aimed at foreground interlopers, while the dim cut reduces the contribution of the galaxy background. Later methods used similar counting techniques but included a proper account of the background \\citep[e.g.][]{bahcall81} . Since then, more sophisticated algorithms have been developed and applied to CCD-based imaging \\citep[e.g. matched-filter methods][]{postmanetal96,bramel00,yee99,kochaneketal03,dongetal07}.  Projection effects are now a much more benign problem thanks to these more sophisticated richness measurement techniques, the advent of accurate photometric data enabled by modern CCDs, and most recently, the well-known observations that ellipticals and cluster E/S0 galaxies in particular tend to form a tight ridgeline in color-magnitude space \\citep{visvanathan77,boweretal92,gladdersyee00,kmawe07a}. This color clustering has been integral to richness measurements in the SDSS \\citep{gotoetal02,milleretal05,kmawe07a} and the Red Sequence Cluster Survey  \\citep[RCS:][]{gladders05}, and such color-based measures have been shown to be effective mass tracers \\citep{yee03,muzzin07,sheldonetal07,johnstonetal07,rmbej08,beckeretal07}.  While richness estimates show a strong correlation with other mass proxies \\citep[e.g.][]{yee03,daietal07,sheldonetal07,johnstonetal07,beckeretal07,rembj08}, considerable scatter in the mass--richness relation still remains. For instance, the richness measure used in the RCS cluster catalog has a logarithmic scatter of $\\sigma_{\\ln M}\\approx 0.8$ \\citep{gladdersetal07}, while for maxBCG clusters the number is closer to $\\sigma_{\\ln M}\\approx 0.5$~\\citep{rozoetal08a}. This is to be compared to the scatter for X-ray mass tracers, which is expected to be as low as $\\approx 8\\%$ for $Y_X$ based on simulations \\citep{kravtsovetal06}, or as high as $\\approx 25\\%$ for non-core extracted soft X-ray band luminosities~\\citep[e.g.][]{sebsn06,vbefh08}. Clearly, much improvement is needed to bring the scatter of richness measures to the level of X-ray mass tracers.  This work is aimed at reducing the variance in the richness-mass relation.  We do this by explicitly constructing a new richness estimator that significantly reduces the scatter in mass at fixed richness for maxBCG clusters. Relative to $\\Nobs$ of maxBCG, we introduce two significant differences.   The first of these involves using a matched filter algorithm to estimate cluster richness.  Matched filters have been used in the literature before \\citep{postmanetal96,kochaneketal03}. Unlike those works, however, our matched filter includes a color component, which is of critical importance for reducing projection effects over the redshift range spanned by our cluster sample.\\footnote{In \\citet{kochaneketal03}, the low redshift of the clusters make single band magnitudes better proxies for distance than colors, so the lack of a color filter in the richness estimator is less important for their work.} In that sense, our filter is closer in spirit to that of \\citet{dongetal07}, who include a photometric redshift filter into their richness estimate.  We also note here that group-scale studies suggest that some measure of the average color in the cluster is indicative of mass, particularly below $ \\sim 10^{14} M_{\\odot}$ \\citep[][]{martinez02,martinez06,weinmann06,hansenetal07} .  The second difference we introduce is the way in which the aperture used to estimate cluster richness is determined. Generically, cluster richness estimators involve counting the number of galaxies within some specified aperture, which can thus be interpreted as defining the ``size'' of the cluster.  This begs the question, then, of how is one to select the correct size of a cluster a priori?  Theoretically, halo sizes are usually defined in terms of $R_\\Delta$, a radius which encompasses a mean density that is $\\Delta$ times either the mean or the critical density of the universe (conventions vary from author to author).  Unfortunately, not only is such a definition not applicable observationally, authors vary both on the reference background density (critical versus mean mass density), and on the specific overdensity value.  Thus, even though significant progress has been made \\citep{cuestaetal08}, a definitive definition of halo size remains elusive.  In this work, we approach this question with observations in mind.  That is, rather than coming with a preconceived notion of what the radius of a cluster is, we let the data tell us what the optimal radii for our clusters is by demanding that optical richness be as tightly correlated as possible with X-ray luminosity. The idea is as follows: first, one posits a scaling relation between cluster richness and cluster radius.  When estimating cluster richness, one then demands that the richness-radius scaling relation be satisfied.  For instance, given a cluster, one can simply make an initial guess for its richness.  Using the richness-radius scaling relation, one can then draw a circle of the appropriate radius, and count the number of galaxies within it.  If the richness was underestimated, one will find too many galaxies, signaling that the richness estimate must be increased.  Proceeding in this way, one can quickly zero in on the appropriate richness for the object.  This does, however, leave open the question of what the correct richness-radius relation is. Since we are interested in finding a new richness estimator that is tightly correlated with halo mass, we can use the scatter in the mass--richness relation as our figure of merit to determine the ``correct'' richness-radius relation.  In practice, we use the $\\Lx-$richness scatter rather than the mass--richness scatter because the scatter in mass is not directly observable. We emphasize that since the mass scatter at fixed X-ray luminosity \\citep[see e.g.][]{vikhlininetal08} is considerably tighter than the corresponding scatter at fixed richness~\\citep{rozoetal08a}, the use of X-ray luminosity as a mass tracer for our purposes is well justified.  The layout of the paper is as follows. We describe the data sets used in this work in \\S~\\ref{sec:data}.  Our matched filter estimator is introduced in \\S~\\ref{sec:matchfilter}, followed by our method for determining the optimal radius-richness relation in \\S~\\ref{sec:methods}.  We present our results in section \\S~\\ref{sec:results}.  In investigating the properties of our new richness measure, we have discovered that the redshift evolution of the richness-mass relation of our new estimator is much more mild than that measured for $N_{200}$. These results and the corresponding discussion are presented in \\S~\\ref{sec:zdep}.  We summarize our results and present our conclusions in \\S~\\ref{sec:conclusions}.  Throughout, whenever needed a flat $\\Lambda$CDM cosmology with $\\Omega_M=0.3$ and $h=1.0$ was assumed. ",
        "conclusions": "\\label{sec:conclusions}  We have introduced a new matched filter richness estimator $\\lambda$ whose correlation with mass is significantly tighter than that of $\\Nobs$, the original maxBCG richness estimator. Relative to other matched filter estimates, our estimator has two significant differences: \\begin{enumerate} \\item The richness is measured on a scale that is optimized in the sense that it minimizes the scatter in $\\Lx$ at fixed richness. \\item In addition to a radial and magnitude filters, we include a color filter.  This is of crucial importance for differentiating between member and non-member (projected) galaxies.  \\end{enumerate} The first these points is important since we have demonstrated that a poor choice of aperture increases the scatter in mass at fixed richness, while the latter minimizes the impact of projection effects in richness estimates.  Of the two, however, the improved treatment of galaxy color is the principal reason for the marked reduction of the scatter in the $\\Lx-$richness relation.  Our procedure for aperture optimization can be easily generalized to any mass tracer for which one can construct a calibrating data set.  In our particular case, we minimize the scatter in the $\\Lx-\\lambda$ relation by measuring both $\\Lx$ and $\\lambda$ within an aperture $R_c(\\lambda)=R_0(\\lambda/100)^\\alpha$, and varying the model parameters $R_0$ and $\\lambda$.  Given the small richness range probed by our sample, we have not been able to isolate unique values for $R_0$ and $\\alpha$, finding instead a degeneracy region corresponding to a fixed mean cluster radius for the clusters in the sample.   Based on \\emph{a priori} assumptions about the radius--richness scaling, we have fixed $\\alpha=1/3$, which yields a normalization of $R_0 = 1.27\\pm0.03$.  We note, however, that the degeneracy region intersects $\\alpha=0$ at $R_0\\approx 850\\ \\kpc$. Although we expect that this fixed scaling will not be ideal at the rich group/poor cluster scale, it does work as a ``first guess'' richness and may be applicable to future cluster finding techniques. At this point, it is unclear whether the cluster radii selected by our technique reflects a true physical property of the maxBCG clusters, or whether it is driven primarily by a compromise between the the increase signal one expects at larger aperture, and the smaller noise one expects for smaller apertures. Regardless of the source, it is likely that similar aperture dependences exist for other mass tracers.  We have also found our new richness estimator has scaling relations whose redshift evolution is much more mild than those exhibited by $\\Nobs$.  This difference arises due to two effects: first, $\\Nobs$ uses a top-hat filter to select cluster galaxies, where as our matched filter estimator $\\lambda$ uses gaussian color filters.  Second, $\\Nobs$ centered its color filter at the color of the BCG, whereas $\\lambda$ centered its color using an observationally calibrated color--redshift relation.  The fact that the color of the BCG does not always agree with the observationally calibrated redshift-color relation leads to a systematic difference between the two richness measures.  Moreover, we also found that the sharp edges of the top-hat filter result in a richness estimator that is very sensitive to the details of the color model, whereas our gaussian filter is much more robust to moderate changes in the model parameters.  Restricting ourselves to the clean cluster sample, which excludes cooling flow clusters and clusters with obvious foreground contamination in their X-ray luminosities, we have found that the scatter in the $\\Lx-$richness relation of the 2000 richest clusters is is $\\sigma_{\\lnl|\\lambda}=0.69$ for $\\lambda$, compared to $\\sigma_{\\lnl|\\Nobs}=0.86$ for $\\Nobs$. Assuming a slope of $\\approx 1.6$ for the $\\Lx-M$ relation~\\citep{sebsn06,rembj08,vbefh08}, these amount to a logarithmic scatter in the mass--richness relation of $\\approx 0.43$ and $0.54$ respectively. While this is a very significant improvement, we expect that further tightening of the scatter in mass at fixed richness must be possible. For instance, assuming the intrinsic scatter in the richness-mass relation is Poisson, the logarithmic scatter possible for clusters with 20 galaxies or so should be roughly $\\approx 0.2$.  Fortunately, there are still many options left for us to explore in our quest to define optical mass proxies that can be competitive with other mass tracers in terms of the tightness of the correlation with mass. As we have defined it here, our richness estimates only makes use of the number of galaxies in the cluster. One could, for instance, weigh our cluster richness by other optical mass tracers such as the luminosity of the brightest cluster galaxy \\citep{reyesetal08}, the abundance of baryons contained in the intracluster light \\citep[e.g.][]{gonzales07}, or other aspects of the cluster galaxy morphology \\citep[e.g. Bautz-Morgan Type, ][]{bautz70}.  In addition, one could weigh each galaxy's contribution to the richness by physical observables such as galaxy luminosity.  Such a luminosity weighted richness estimate would be a measure of the optical luminosity of the cluster as a whole, and might be better correlated with mass than richness itself \\citep[see also][]{lin03,milleretal05, pbbrv05}. It is also likely that further improvements in richness estimates can arise with more accurate filters, a possibility we intend to explore in future work.  Finally, we know that even with today's filters, part of the scatter we observe must be due to systematics effects such as failures of the cluster finding algorithm in identifying the correct center of a cluster, a problem which we have not addressed in this work. For the time being, the fact that naive theoretical expectations result in a scatter much lower than previously observed, and the fact that on our first attempt at defining a better richness estimator resulted in a highly significant ($\\approx 11\\sigma$) improvement over $\\Nobs$, suggest that the future is rife with opportunities for this kind of work."
    },
    "0809/0809.0337_arXiv.txt": {
        "abstract": "{ \\ddscatv\\ is a freely available open-source Fortran-90 software package applying the ``discrete dipole approximation'' (DDA) to calculate scattering and absorption of electromagnetic waves by targets with arbitrary geometries and complex refractive index.  The targets may be isolated entities (e.g., dust particles), but may also be 1-d or 2-d periodic arrays of ``target unit cells'', which can be used to study absorption, scattering, and electric fields around arrays of nanostructures.  The DDA approximates the target by an array of polarizable points. The theory of the DDA and its implementation in \\ddscat\\ is presented in Draine (1988) and Draine \\& Flatau (1994), and its extension to periodic structures (and near-field calculations) in Draine \\& Flatau (2008). \\ddscatv\\ allows accurate calculations of electromagnetic scattering from targets with ``size parameters'' $2\\pi \\aeff/\\lambda \\ltsim 25$ provided the refractive index $m$ is not large compared to unity ($|m-1| \\ltsim 2$). \\ddscatv\\ includes support for {\\tt MPI}, {\\tt OpenMP}, and the \\Intel\\ Math Kernel Library (MKL).  \\ddscat\\ supports calculations for a variety of target geometries (e.g., ellipsoids, regular tetrahedra, rectangular solids, finite cylinders, hexagonal prisms, etc.). Target materials may be both inhomogeneous and anisotropic. It is straightforward for the user to ``import'' arbitrary target geometries into the code. \\ddscat\\ automatically calculates total cross sections for absorption and scattering and selected elements of the Mueller scattering intensity matrix for specified orientation of the target relative to the incident wave, and for specified scattering directions. \\ddscatv\\ can calculate scattering and absorption by targets that are periodic in one or two dimensions. \\ddscatv\\ also supports postprocessing to calculate $\\bE$ and $\\bB$ at user-selected locations near the target; a Fortran-90 code {\\bf DDfield} for this purpose is included in the distribution.  This User Guide explains how to use \\ddscatv\\ to carry out electromagnetic scattering calculations. }  \\newpage \\tableofcontents \\newpage ",
        "introduction": "}  DDSCAT is a software package to calculate scattering and absorption of electromagnetic waves by targets with arbitrary geometries using the ``discrete dipole approximation'' (DDA). In this approximation the target is replaced by an array of point dipoles (or, more precisely, polarizable points); the electromagnetic scattering problem for an incident periodic wave interacting with this array of point dipoles is then solved essentially exactly. The DDA (sometimes referred to as the ``coupled dipole approximation'') was apparently first proposed by Purcell \\& Pennypacker (1973). DDA theory was reviewed and developed further by Draine (1988), Draine \\& Goodman (1993), reviewed by Draine \\& Flatau (1994) and Draine (2000), and recently extended to periodic structures by Draine \\& Flatau (2008).  \\ddscatv, the current release of DDSDCAT, is an open-source Fortran 90 implementation of the DDA developed by the authors.\\footnote{ The release history of {\\bf DDSCAT} is as follows: \\begin{itemize} \\vspace*{-0.4em} \\item{\\bf DDSCAT 4b}: Released 1993 March 12 \\vspace*{-0.4em} \\item{\\bf DDSCAT 4b1}: Released 1993 July 9 \\vspace*{-0.4em} \\item{\\bf DDSCAT 4c}: Although never announced, {\\tt DDSCAT.4c} was made available to a number of interested users beginning 1994 December 18 \\vspace*{-0.4em} \\item{\\bf DDSCAT 5a7}: Released 1996 \\vspace*{-0.4em} \\item{\\bf DDSCAT 5a8}: Released 1997 April 24 \\vspace*{-0.4em} \\item{\\bf DDSCAT 5a9}: Released 1998 December 15 \\vspace*{-0.4em} \\item{\\bf DDSCAT 5a10}: Released 2000 June 15 \\vspace*{-0.4em} \\item{\\bf DDSCAT 6.0}: Released 2003 September 2 \\vspace*{-0.4em} \\item{\\bf DDSCAT 6.1}: Released 2004 September 10 \\vspace*{-0.4em} \\item{\\bf DDSCAT 7.0}: Released 2008 September 1 \\end{itemize} } It extends DDSCAT to treat structures that are periodic in one or two dimensions, using methods described by Draine \\& Flatau (2008).  DDSCAT is intended to be a versatile tool, suitable for a wide variety of applications including studies of interstellar dust, atmospheric aerosols, blood cells, marine microorganisms, and nanostructure arrays. As provided, \\ddscatv\\ should be usable for many applications without modification, but the program is written in a modular form, so that modifications, if required, should be fairly straightforward.  The authors make this code openly available to others, in the hope that it will prove a useful tool.  We ask only that: \\begin{itemize} \\item If you publish results obtained using {{\\bf DDSCAT}}, please acknowledge the source of the code, and cite relevant papers, such as Draine \\& Flatau (1994), Draine \\& Flatau (2008), and this UserGuide (Draine \\& Flatau 2008).  \\item If you discover any errors in the code or documentation, please promptly communicate them to the authors.  \\item You comply with the ``copyleft\" agreement (more formally, the GNU General Public License) of the Free Software Foundation: you may copy, distribute, and/or modify the software identified as coming under this agreement. If you distribute copies of this software, you must give the recipients all the rights which you have. See the file {\\tt doc/copyleft.txt} distributed with the DDSCAT software. \\end{itemize} We also strongly encourage you to send email to {\\tt draine@astro.princeton.edu} identifying yourself as a user of DDSCAT;  this will enable the authors to notify you of any bugs, corrections, or improvements in DDSCAT. Up-to-date information on DDSCAT can be found at \\\\ \\hspace*{2em}{\\tt http://www.astro.princeton.edu/$\\sim$draine/DDSCAT.html}\\\\ Information is also available at {\\tt http://ddscat.wikidot.com/start}  The current version, \\ddscatv, uses the DDA formulae from Draine (1988), with dipole polarizabilities determined from the Lattice Dispersion Relation (Draine \\& Goodman 1993, Gutkowicz-Krusin \\& Draine 2004). The code incorporates Fast Fourier Transform (FFT) methods (Goodman, Draine, \\& Flatau 1991). \\ddscatv\\ includes capability to calculate scattering and absorption by targets that are periodic in one or two dimensions -- arrays of nanostructures, for example. The theoretical basis for application of the DDA to periodic structures is developed in Draine \\& Flatau (2008).  We refer you to the list of references at the end of this document for discussions of the theory and accuracy of the DDA [in particular, reviews by Draine and Flatau (1994) and Draine (2000), and recent extension to 1-d and 2-d arrays by Draine \\& Flatau (2008).]. In \\S\\ref{sec:whats_new} we describe the principal changes between \\ddscatv\\ and the previous releases.  The succeeding sections contain instructions for: \\begin{itemize} \\item compiling and linking the code; \\item running a sample calculation; \\item understanding the output from the sample calculation; \\item modifying the parameter file to do your desired calculations; \\item specifying target orientation; \\item changing the {\\tt DIMENSION}ing of the source code to accommodate your desired calculations. \\end{itemize} The instructions for compiling, linking, and running will be appropriate for a Linux system; slight changes will be necessary for non-Linux sites, but they are quite minor and should present no difficulty.  Finally, the current version of this User Guide can be obtained from\\newline {\\tt http://arxiv.org/abs/0809.0337} -- you will be offered the options of downloading either Postscript or PDF versions of the User Guide. ",
        "conclusions": ""
    },
    "0809/0809.3791_arXiv.txt": {
        "abstract": "In alternative theories of gravity, designed to produce cosmic acceleration at the current epoch, the growth of large scale structure can be modified. We study the potential of upcoming and future tomographic surveys such as DES and LSST, with the aid of CMB and supernovae data, to detect departures from the growth of cosmic structure expected within General Relativity. We employ parametric forms to quantify the potential time- and scale-dependent variation of the effective gravitational constant, and the differences between the two Newtonian potentials. We then apply the Fisher matrix technique to forecast the errors on the modified growth parameters from galaxy clustering, weak lensing, CMB, and their cross-correlations across multiple photometric redshift bins. We find that even with conservative assumptions about the data, DES will produce non-trivial constraints on modified growth, and that LSST will do significantly better. ",
        "introduction": "\\label{Intro} Observations strongly favor a universe that has recently entered a phase of accelerated expansion~\\cite{Riess:1998cb,perlmutter}. This poses a puzzle for modern cosmology as standard \\emph{General Relativity} (GR), applied to a universe which contains only radiation and dust, has difficulties fitting the data. One can view this as evidence for the existence of \\emph{Dark Energy} (DE) -- a yet unknown component with a negative equation of state, such as a cosmological constant, $\\Lambda$. An alternative explanation could involve modifying GR in a manner that leads to accelerating solutions.  Popular examples include the so-called $f(R)$ class of models~\\cite{Starobinsky:1980te,Capozziello:2003tk,Carroll:2003wy,Starobinsky:2007hu,Nojiri:2008nk,Boisseau:2000pr}, Chameleon type scalar-tensor theories~\\cite{Khoury:2003aq}, the \\emph{Dvali-Gabadadze-Porrati} (DGP) model~\\cite{Dvali:2000hr}, and models motivated by DGP, such as the recently introduced Degravitation scenario~\\cite{Dvali:2007kt}.  Although the $\\Lambda$CDM model, consisting of GR  with $\\Lambda$ and \\emph{Cold Dark Matter} (CDM), is currently the best fit to the data, it faces some challenges on the theoretical side, such as the coincidence and the fine-tuning problems. On the other side, alternative theories of gravity have to satisfy the multitude of existing experimental tests passed by GR~\\cite{Will,Will:2005va}. This typically requires a degree of fine-tuning which, at best, does not improve on that involved in setting $\\Lambda$ to the value required by current data. Moreover, once these modified theories are tuned to avoid conflicts with existing constraints (when it is possible), their predictions for the expansion history of the universe are often identical to that of the $\\Lambda$CDM model\\footnote{The modifications of gravity discussed in this work do not attempt to replace dark matter. We assume existence of CDM.}\\cite{Dvali:2007kt,Hu:2007nk,Appleby:2007vb,Pogosian:2007sw,Brax:2008hh,Capozziello:2008fn}. However, this degeneracy is typically broken at the level of cosmological structure formation; indeed, models of modified gravity that closely mimic the cosmological constant at the background level can still give significantly different predictions of the growth of structure.  The large scale structure of the universe therefore offers a promising testing ground for GR and it is important to explore to what extent one can detect departures from GR in the growth of structure with present and upcoming cosmological data.  By definition, the term `modified gravity' implies that the form of the Einstein-Hilbert action is different from that of GR, and, as a consequence, the Einstein equations are changed. At the background level, the modifications allow for a late-time acceleration, which is typically degenerate with $\\Lambda$CDM after the required tuning. However, since the equations describing the evolution of cosmological perturbations are modified as well, models with the same expansion history as in $\\Lambda$CDM can lead to different dynamics for the growth of cosmic structure.  Structure formation has been studied for f(R) models in~\\cite{Song:2006ej, Bean:2006up, Pogosian:2007sw,Tsujikawa:2007xu}, for Chameleon models in~\\cite{Brax:2005ew}, and for the DGP model in~\\cite{Koyama:2005kd,Song:2006jk,Song:2007wd,Cardoso:2007xc}. One way to test the consistency of the $\\Lambda$CDM model is to compare values of the cosmological parameters extracted from distance measures, such as supernovae magnitudes and baryon acoustic oscillations, to the values found from growth measures, such as galaxy counts and weak lensing \\cite{Ishak:2005zs}. One can also introduce general parametrizations of the modified evolution of gravitational potentials and matter perturbations~ \\cite{Zhang:2007nk,Amin:2007wi,Hu:2007pj,Hu:2008zd,Bertschinger:2008zb} for the purpose of detecting/constraining departures from GR.  Scalar metric perturbations in the Newtonian gauge are described by two potentials, $\\Psi(\\vec{x},t)$ and $\\Phi(\\vec{x},t)$, which correspond to perturbations in the time-time and space-space components of the metric tensor, respectively. In the $\\Lambda$CDM model, these two potentials are equal during the epoch of structure formation, and their time dependence is set by the same scale-independent linear growth function $g(a)$ that describes the growth of matter density perturbations. This, generally, is no longer true in theories of modified gravity, where one can have scale-dependent growth patterns. The two Newtonian potentials need not be the same, and their dependence on matter perturbations can be different. Working in Fourier space, we parametrize the ratio between the Newtonian potentials and the dependence of $\\Psi$ on the matter density perturbation with two time- and scale-dependent functions $\\gamma(a,k)$ and $\\mu(a,k)$. We then study the potential of upcoming and future tomographic surveys to constrain departures of these functions from their GR values. Dealing with unknown functions implies working with an infinite number of degrees of freedom. Similarly to the more widely studied problem of determining the dark energy equation of state, $w(z)$, one can proceed in several ways: \\begin{enumerate} \\item One can assume a functional form motivated by a certain class of theories, with a few parameters, and forecast the constraints on the parameters. This would tell us the extent to which one can constrain these theories and also allow us to reconstruct the shape of the functions based on the chosen form. The results, of course, would depend on the choice of the parametrization. However, they would be good indicators of the power of current and upcoming surveys to constrain departures from GR. Moreover, as we will discuss shortly, it is possible to employ a parametrization that accurately represents a broad class of modified theories. This is the approach we take in this paper. \\item Another, non-parametric approach, consists of performing a \\emph{Principal Component Analysis} (PCA) to determine the eigenmodes of $\\mu$ and $\\gamma$ that can be constrained by data~\\cite{PCA1,PCA2}. This method allows one to compare different experiments and their combinations according to the relative gain in information about the functions. PCA can also point to the ``sweet spots'' in redshift and scale where data is most sensitive to variations in $\\mu$ and $\\gamma$, which can be a useful guide for designing future observing strategies.  The PCA method does not allow one to reconstruct the shape of the functions from data. However, one can still reproduce the errors on parameters of any parametrization from the eigenvectors and eigenvalues found using PCA~\\cite{PCA2}. Hence, in terms of forecasting the errors, the PCA method can do everything that the first method can, plus the benefits mentioned above. It is, however, more demanding computationally, and one needs a criterion for deciding which modes are well-constrained. We consider the PCA method in a separate publication \\cite{modgrav_pca}. \\item A third approach, which can work for certain estimators of $\\gamma$, is a direct reconstruction from data. In~\\cite{Zhang:2007nk,Zhang:2008ba}, it was proposed to consider the ratio of the peculiar velocity-galaxy correlation with the weak lensing - galaxy correlation. In such a ratio, the dependence on the galaxy bias cancels out. Then, since the peculiar velocities are determined by the potential $\\Psi$, while the weak lensing is controlled by $\\Phi+\\Psi$, such ratio, if appropriately constructed, would directly probe any difference between $\\Phi$ and $\\Psi$. This is a more direct and model-independent way of testing GR with the growth of structure than the first two methods. Its power, however, will depend on how well future experiments will be able to measure  peculiar velocities. Also, while it may allow the extraction of $\\gamma$, it does not directly probe $\\mu$. Still, this is a novel and promising method that should be pursued in parallel with the first two. \\end{enumerate}  Our parametric forms for $\\gamma(a,k)$ and $\\mu(a,k)$ are analogous to those introduced in~\\cite{Bertschinger:2008zb}, and contain a total of five parameters. These forms are highly accurate in describing the linear growth in a wide class of scalar-tensor theories and also flexible enough to capture features of modified dynamics in other theories. We use the Fisher matrix technique to forecast the errors on the modified growth parameters, along with the standard set of cosmological parameters. We represent the data in terms of a set of all possible two-point correlation functions (both auto- and cross-correlations) between the galaxy counts, weak lensing shear, and \\emph{Cosmic Microwave Background} (CMB) temperature anisotropy, across multiple redshift bins, in addition to the CMB E-mode polarization autocorrelation, and the CMB E-mode and temperature cross-correlation. We find that even with a conservative treatment of data, such as using only the modes that are well within the linear regime, the upcoming and future tomographic surveys, like the \\emph{Dark Energy Survey} (DES) and \\emph{Large Synoptic Survey Telescope} (LSST), in combination with CMB and future \\emph{SuperNovae} (SNe) luminosity-distance data, will be able to produce non-trivial constraints on all five parameters. We also show that LSST will significantly improve our ability to test GR, compared to DES.  The rest of the paper is organized as follows. In Sec.~\\ref{theory}, we motivate our choice of parametric forms for $\\gamma(a,k)$ and $\\mu(a,k)$. In Sec.~\\ref{observables}, we describe the set of observables expected from tomographic weak lensing surveys, their dependence on the underlying gravitational potentials, and the Fisher matrix technique we used to forecast the parameter errors. We also describe the experiments considered -- DES~\\cite{DES}, LSST~\\cite{LSST}, the \\emph{SuperNova/Acceleration Probe} (SNAP)~\\cite{SNAP}, and Planck~\\cite{Planck}. We present our results in Sec.~\\ref{results} along with a discussion of their dependence on the assumptions made in the analysis. We conclude with a summary in Sec.~\\ref{summary}. ",
        "conclusions": "\\label{summary}  In this paper, we have investigated the power of future tomographic surveys to constrain modifications in the growth of structure w.r.t. to that predicted in $\\Lambda$CDM. Models of modified gravity, as well as models of coupled dark energy and dark matter, in general introduce a scale-dependence in the growth of structure and a time- and scale-dependent slip between the gravitational potentials. These modifications are expected to leave characteristic imprints on the observables, which in principle could be used to break the background degeneracy among different models of cosmic acceleration. It is therefore useful and important to explore to which extent the upcoming experiments will be able to detect and constrain modified growth patterns.  We have used a five-parameter description for the rescaling of the Newton constant, $\\mu(a,k)$, and for the ratio of the metric potentials, $\\gamma(a,k)$, equivalent to that introduced in~\\cite{Bertschinger:2008zb}. From the point of view of scalar-tensor theories (e.g.$f(R)$ and Chameleon models) these parameters are  related to the coupling in the dark sector and to the characteristic mass scale of the model. In these cases, the five parameters are not all independent; they need to satisfy the consistency conditions given by Eq.~(\\ref{st_consistency}).  This can be used to constrain the scalar-tensor models and potentially rule them out. In Sec.~\\ref{theory} we have described the $f(R)$ and Chameleon theories which we have used as fiducial models for our error forecasting. We have then studied in detail the constraints on the five parameters based on four fiducial models (two $f(R)$ and two Chameleon) expected from Weak Lensing (WL), Galaxy Counts (GC), CMB and their cross-correlation spectra as seen by Planck, DES and LSST (additionally using Planck and SNAP to constrain the standard set of cosmological parameters). We have found that for scalar-tensor type models, DES can provide $20-40$\\% level constraints on individual parameters, and that LSST can improve on that, constraining the MG parameters to better than $10$\\% level.  We have also found a strong degeneracy between the two functions $\\mu$ and $\\gamma$, which was consistent with our expectations. Despite this degeneracy, however, the error ellipses are sufficiently tight for us to still find non-trivial bounds on the parameters of both functions. Overall, with DES, and especially with LSST, one will be able to significantly reduce the volume of the allowed parameter space in scalar-tensor type models.  We have also found that the dilution of the constraints due to the degeneracy with linear bias is minimal when the full set of auto- and cross-correlation tomographic datasets is considered.    The method we have employed  relies on the choice of fiducial values for the MG parameters, and therefore is model-dependent. Nevertheless, the results we have obtained are encouraging as they show that upcoming and future surveys can place non-trivial bounds on modifications of the growth of structure  even in the most conservative case, i.e. considering only linear scales. This motivates us to apply the Principal Component Analysis (PCA) to the functions $\\mu$ and $\\gamma$. This method is more demanding computationally but, on the other hand, allows for a model-independent comparison of  different experiments according to the relative gain in information about the functions. Furthermore, this method can point to the ``sweet spots'' in redshift and scale where data is most sensitive to variations in the functions $\\mu$ and $\\gamma$. This would be important information for designing future observing strategies. We present the results of the PCA analysis in a separate publication~\\cite{modgrav_pca}.  Another direction for future work is to study the effects of allowing for a small fraction of hot dark matter, such as neutrinos, and dynamical dark energy. One can, in principle, model the effects of dark energy perturbations and massive neutrinos by appropriately choosing the time-scale dependence of $\\mu$ and $\\gamma$. It is generally difficult to design a test that will definitively distinguish modified gravity from an exotic form of dark energy. However, it is clear that future data will have the potential to detect or, at the very least, significantly reduce the range of, departures from the $\\Lambda$CDM model."
    },
    "0809/0809.4664_arXiv.txt": {
        "abstract": "The characteristics of irradiated solar system planetary atmospheres have been studied for decades, consequently modern planetary science benefits from an exhaustive body of ground- and space-based data.  The study of extrasolar planetary atmospheres, by contrast, is still in its infancy and currently rests on a few score of datapoints, mostly of the transiting planets.  This short survey aims not to review this dynamic field but rather stresses the importance of a few theoretical concepts and processes for our understanding of exoplanet atmospheres.  Topics covered include atmospheric structure and dynamics, cloud processes and photochemistry of planetary atmospheres. Influences on the albedos, spectra, and colors of extrasolar planets are reviewed and caution is urged in the interpretation of exoplanet colors. ",
        "introduction": "The atmosphere controls a planet's evolution through time and provides a window into the chemical and physical conditions under which the planet formed.  Except for airless worlds, like Mercury, the atmosphere also mediates the flow of information we receive about the nature of the planet.  In this contribution I aim to discuss a sampling of a few of the more important concepts which relate to our understanding of extrasolar planetary atmospheres.   I will focus on topics that are important for understanding spectral and photometric data from extrasolar planets, with an eye towards providing examples from the solar system.  In some sense the trajectory of extrasolar planet science is recapitulating the history of solar system planetary science.   Three decades ago in the University of Arizona Space Science series book ``Jupiter'' Wallace (1976) reviewed the sparse data then available on Jupiter's thermal emission spectrum and presented models of the planet's atmospheric structure.  His effort to piece together a consistent picture of the atmosphere from a handful of thermal infrared photometric measurements, reflection spectra, and transmission data from stellar occultations bears a striking resemblance to current efforts to understand the transiting extrasolar planets.  Today, with our eyes now fixed on exoplanets, we are fortunate to have at hand the insights gained from such work over the past half century of planetary exploration. For a few transiting planets we again find ourselves in the situation where we have a few broad band brightness temperature measurements from which we are attempting to infer global atmospheric processes.  As we carry out our science we would do well to remember that this ground has been crossed before: irradiated planetary atmospheres have been encountered before the $21^{\\rm st}$ century (a fact sometimes overlooked in the modern literature).  The excitement of our time, of course, is that the range of parameter space, from planetary masses to the dizzying range of incident radiation is far larger than previously encountered.  Nevertheless we should not be surprised that well understood processes familiar from solar system planetary science, including photochemistry, hot stratospheres, and cloud processes also play important roles in extrasolar planet atmospheres (Marley 1997).  To aid those attempting to traverse this new landscape of theory and data of exoplanet atmospheres for the first time,  I discuss in this chapter a few elementary themes that are helpful for understanding the conceptual scenery. While I will illustrate concepts by drawing from work in the field, I am not in any sense attempting to provide a comprehensive review of exoplanet science, as many such recent reviews are available (e.g., Charbonneau (2008), Deming (2008), and Seager et al.\\ (2008) review recent progress; Burrows et al. (2001, 2006) and Barman et al. (2005) present useful surveys of much of the theory of these objects; and Marley et al.\\ (2007) review the pre-2007 field).  Instead I present a brief, illustrated (if somewhat idiosyncratic) guidebook to a few of the important processes encountered at these new worlds.  Hence I will touch on atmospheric structure, cloud formation, the interpretation of brightness temperature and the importance of photochemistry. I will mention some of the processes contributing to stratospheric heating and conclude with some brief comments about albedos and colors of planets. ",
        "conclusions": "As we enter the age of the direct detection and characterization of extrasolar planets it is important not to overlook the lessons learned from the past half century of planetary exploration.  Solar system planets give ample demonstration of the importance of atmospheric dynamics, cloud and condensation processes, and photochemistry in controlling the face of planets.  The great diversity of atmospheres seen in our own solar system, from the hazy skies of Titan to the turbulent atmosphere of Jupiter to the blue vistas of Earth, are emblematic of the diversity of processes that can affect properties of planets.  Exoplanets will exhibit even larger ranges of properties and we should not be surprised by hazy yellow Earths, red, white, or blue Jupiters, or other unexpected worlds. Certainly extrasolar planet atmospheres will be influenced by some of the processes mentioned here.  (Other yet-to-be-discovered influences will undoubtedly be important as well.) Ultimately, however, such discoveries will extend the journey of planetary exploration begun by the {\\it Mariners} and {\\it Voyagers} of the last century out into the galaxy.  {\\bf Acknowledgements:} The author thanks Jonathan Fortney and Kevin Zahnle for helpful comments on the manuscript and Thierry Montmerle and the entire Les Houches staff for organizing an exceptionally productive and informative program."
    },
    "0809/0809.4949_arXiv.txt": {
        "abstract": "We present a purely-radiative hydrodynamic model of the $\\kappa$-mechanism that sustains radial oscillations in Cepheid variables. We determine the physical conditions favourable for the $\\kappa$-mechanism to occur by the means of a configurable hollow in the radiative conductivity profile. By starting from these most favourable conditions, we complete nonlinear direct numerical simulations (DNS) and compare them with the results given by a linear-stability analysis of radial modes. We find that well-defined instability strips are generated by changing the location and shape of the conductivity hollow.  For a given position in the layer, the hollow amplitude and width stand out as the key parameters governing the appearance of unstable modes driven by the $\\kappa$-mechanism. The DNS confirm both the growth rates and structures of the linearly-unstable modes. The nonlinear saturation that arises is produced by intricate couplings between the excited fundamental mode and higher damped overtones. These couplings are measured by projecting the DNS fields onto an acoustic subspace built from regular and adjoint eigenvectors and a 2:1 resonance is found to be responsible for the saturation of the $\\kappa$-mechanism instability. ",
        "introduction": "Eddington (1917) discovered an excitation mechanism of stellar oscillations that is related to the opacity in ionisation regions: the $\\kappa$-mechanism. This mechanism can only occur in regions of a star where the opacity varies so as to block the radiative flux during compression phases (Zhevakin 1953; Cox 1958). Ionisation regions correspond to a strong increase in opacity, leading to the ``opacity bumps'' that are responsible for the local driving of modes. These ionisation regions have nevertheless to be located in a very precise region of a star, neither too close to the surface nor to deep into the stellar core, in order to balance the damping that occurs in other regions. It defines the so-called \\textit{transition region} which is the limit between the quasi-adiabatic interior and the strongly non-adiabatic surface. For classical Cepheids that pulsate on the fundamental acoustic mode, this transition region is located at a temperature $T\\simeq 40\\ 000$ K corresponding to the second helium ionisation (Baker \\& Kippenhahn 1965).  However, the bump location is not solely responsible for the acoustic instability. A careful treatment of the $\\kappa$-mechanism would involve dynamical couplings with convection, metallicity effects and realistic equations of state and opacity tables (Bono et al. 1999). The purpose of our model is to simplify the hydrodynamic approach while retaining the \\textit{leading order} phenomenon -the opacity bump location- such that feasible DNS of the $\\kappa$-mechanism can be achieved. ",
        "conclusions": "Direct numerical simulations (DNS) of the $\\kappa$-mechanism that excite stellar oscillations are performed. We first compute the most favourable setups using a linear-stability analysis of radial modes propagating in a 1-D layer of gas. In our model, a configurable hollow in the radiative conductivity profile mimics the opacity bump responsible for the layer destabilisation. The instability strips found in the linear study are outstandingly confirmed by the DNS and we show that the nonlinear saturation that arises involves a 2:1 resonance between the linearly-unstable fundamental mode and the linearly-stable second overtone."
    },
    "0809/0809.1870_arXiv.txt": {
        "abstract": "{The current stellar atmosphere programs still cannot match some fundamental observations of the brightest stars, and with new techniques, such as optical interferometry, providing new data for these stars, additional development of stellar atmosphere codes is required.} {To modify the open-source model atmosphere program \\textsc{Atlas} to treat spherical geometry, creating a test-bed stellar atmosphere code for stars with extended atmospheres.} {The plane-parallel \\textsc{Atlas} has been changed by introducing the necessary spherical modifications in the pressure structure, in the radiative transfer and in the temperature correction.} {Several test models show that the spherical program matches the plane-parallel models in the high surface gravity regime, and matches spherical models computed by \\textsc{Phoenix} and by \\textsc{MARCS} in the low gravity case.} {} ",
        "introduction": "\\textsc{Atlas} \\citep{1970SAOSR.309.....K, 1979ApJS...40....1K} is a well-documented, well-tested, robust, open-source computer program for computing static, plane-parallel stellar atmospheres in local thermodynamic equilibrium (LTE).  Since \\textsc{Atlas} came to maturity in the 1970s, stellar atmosphere codes have progressed in a number of directions to include important additional physics.  For example, the \\textsc{Phoenix} program \\citep{1999ApJ...512..377H} includes advances such as the massive use of statistical equilibrium (NLTE) in place of LTE, spherically symmetric extension in place of plane-parallel geometry, the inclusion of the dust opacities needed for brown dwarf temperatures, and the ability to include blast wave velocity fields.  These advances, while obviously moving toward greater reality, have not eliminated some persistent problems.  For example, a detailed study of Arcturus using \\textsc{Phoenix} models \\citep{2003ApJ...596..501S} found that their spherical NLTE models \\textit{increased} the discrepancy between the observed and computed spectral irradiance.  A similar analysis of models with solar parameters \\citep{2005ApJ...618..926S} also concluded that important opacity and/or other physics is still missing.  This evidence that the significant improvements contained in the state-of-the-art programs have not achieved closer agreement with the observations of these fundamental, bright stars indicates a need to continue exploring additional physics.  This is a valuable role that \\textsc{Atlas} can play because it is open source and freely available from Kurucz's web page (\\textit{http://kurucz.harvard.edu}), where anyone can download the code and use it as the starting point for studying other possible improvements.  An example of the advantage provided by this starting point is the work of \\citet{1996ASPC..108...73K} who explored convection as represented by the full spectrum of turbulence \\citep{1991ApJ...370..295C, 1996ApJ...473..550C} as an alternative to the standard mixing length theory.  More recently, \\citet{2007IAUS..239...71S} have developed a GNU-Linux port of \\textsc{Atlas} for use in a range of applications.  In that spirit, we have developed versions of the \\textsc{Atlas} program that replace the assumption of plane-parallel geometry by spherical symmetry.  These codes, comparable to the LTE, spherical version of \\textsc{Phoenix} or of the spherical version of \\textsc{MARCS} \\citep{2008A&A...486..951G}, can serve as the basis for the study of additional physics needed to understand luminous stars that are cool enough that NLTE effects are not dominant.  Such continued studies are certainly necessary given the revolutionary achievements of optical interferometry in imaging the surfaces of stars, thus providing powerful new observational tests of models of stellar atmospheres. ",
        "conclusions": "We have modified the robust, open-source, plane-parallel model atmosphere program \\textsc{Atlas} to treat spherically extended geometry.  The resulting spherical code, \\textsc{SAtlas}, which is available in both opacity distribution function and opacity sampling versions, was used to compute several test models.  At high surface gravity the spherical model structure is essentially identical to the plane-parallel model structure.  At low surface gravity, the \\textsc{SAtlas} models agree very well with the spherical model structures computed by \\textsc{Phoenix} and by \\textsc{MARCS}.   The \\textsc{SAtlas} program, which runs easily on a desktop workstation, offers a viable alternative for modeling the atmospheres of low surface gravity stars.  As an example of the utility of \\textsc{SAtlas}, we have used it to compute more than 2500 models to create model cubes with fine parameter spacing covering the specific $L_{\\ast}$, $M_{\\ast}$ and $R_{\\ast}$ values needed for an analysis of the optical interferometry of just three stars (Neilson \\& Lester submitted).  We are in the process of computing more models for our own application.  These codes and models are available at \\texttt{http://www.astro.utoronto.ca/$\\sim$lester/Programs/}.  \\textit"
    },
    "0809/0809.3099_arXiv.txt": {
        "abstract": "We present the properties of a large sample (12,282) of nearly face-on low surface brightness (LSB) disk galaxies selected from the main galaxy sample of SDSS-DR4. These properties include $B$-band central surface brightness $\\mu_0(B)$, scale lengths $h$, integrated magnitudes, colors, and distances $D$. This sample has $\\mu_0(B)$ values from 22 to 24.5 mag arcsec$^{-2}$ with a median value of 22.42 mag arcsec$^{-2}$, and disk scale lengths ranging from 2 to 19 kpc. They are quite bright with $M_B$ taking values from -18 to -23 mag with a median value of -20.08 mag. There exist clear correlations between log$h$ and $M_B$, log$h$ and log$D$, log$D$ and $M_B$.  However, no obvious correlations are found between $\\mu_0(B)$ and log$h$, colors etc. The correlation between colors and log$h$ is weak even though it exists. Both the optical-optical and optical-NIR color-color diagrams indicate that most of them have a mixture of young and old stellar populations. They also satisfy color-magnitude relations, which indicate that brighter galaxies tend generally to be redder. The comparison between the LSBGs and a control sample of nearly face-on disk galaxies with higher surface brightness (HSB) with $\\mu_0(B)$ from 18.5 to 22 mag arcsec$^{-2}$ show that, at a given luminosity or distance, the observed LSB galaxies tend to have larger scale lengths. These trends could be seen gradually by dividing both the LSBGs and HSBGs into two sub-groups according to surface brightness. A volume-limited sub-sample was extracted to check the incompleteness of surface brightness. The only one of the property relations having an obvious change is the relation of log$h$ versus $\\mu_0(B)$, which shows a correlation in this sub-sample. ",
        "introduction": "Low surface brightness galaxies (LSBGs) are galaxies that emit much less light per area than normal galaxies. Yet, owing to their faintness compared with the night sky, they are hard to find. Hence their contribution to the local galaxy population  has been underestimated for a long time. It has been suggested that LSBGs may comprise up to half of the local galaxy population (McGaugh et al. 1995).  The discovery and searching history of the LSBGs and the understanding on them have been summarized by Impey $\\&$ Bothun (1997) and Bothun et al.(1997). The first quantitative suggestion about LSBGs is the so-called Freeman's law. It was noticed by Freeman (1970) that the central surface brightness of 28 out of 36 disk galaxies fell in the range of $\\mu_0(B)$=21.65$\\pm$0.3 mag~arcsec$^{-2}$. This could be caused by selection effects (Disney 1976; Zwicky 1957).  The successful hunting for very diffuse galaxies (Romanishin et al. 1983 (from UGC); Sandage et al. 1984; Ellis et al. 1984; Bothun et al. 1987) prompted some new surveys. The first one was based on the second Palomar Sky Survey (POSS-II, Schombert and Bothun 1988; Schombert et al. 1992). A second one was initiated in the Fornax cluster using the Malinization technique in order to compare the results to Virgo (Bothun et al. 1992; Caldwell $\\&$ Bothun 1987). The third one used the Automated Plate Measuring (APM) to scan UK Schmidt plates using an algorithm optimized to find LSB galaxies. This formed the most extensive catalog of LSB galaxies to that date (Impey et al. 1996). Morshidi-Esslinger et al. (1999a,b) extended such sample. The ``Texas survey'' (O'Neil et al. 1997a,b) increases the surface density of LSBGs in the general field, in which they firstly found the red LSB galaxy populations.  Following these catalogs, various aspects of LSBGs have been investigated. Especially, in a series of publications, Impey and his colleagues analyzed their APM survey sample (693 field galaxies in various morphologies, see Impey et al. 1996; Sprayberry et al. 1996; Sprayberry et al. 1997; Impey et al. 2001; Burkholder et al. 2001). Numerous other studies about LSBGs have been done based on smaller samples, relating to surface photometry and color distributions (McGaugh $\\&$ Bothun 1994; de Blok et al. 1995; O'Neil et al. 1997a); number, luminosity and mass density (McGaugh 1996); AGN properties (Mei et al. 2008); and in the optical (Beijersbergen et al. 1999; Bergvall et al. 1999; Galaz et al. 2006, Zackrisson et al. 2005, Pizzella et al. 2008), near-IR (NIR, Bergvall et al. 1999; Galaz et al. 2002,2006; Monnier Ragaigne et al. 2003, Zackrisson et al. 2005), infrared (Hinz et al. 2007) and UV (Boissier et al. 2008) wavelengths.  In addition, Kniazev et al. (2004) developed a method to search for LSBGs from modern survey and used the sample of Impey et. al. (1996) to test their method.  These studies show that LSBGs may be unevolved systems with low metallicity (McGaugh 1994; de Blok $\\&$ van der Hulst 1998a; O'Neil et al. 1998; Bell et al. 2000), low stellar formation rate (van der Hulst et al. 1993; van Zee et al. 1997), small stellar density, a relatively high gas fraction (de Blok et al. 1996; McGaugh $\\&$ de Blok 1997) and large amounts of dark matter (de Blok $\\&$ McGaugh 1997). Despite these impressive progresses, there are still several challenges about LSBGs, such as many aspects of their formation and evolution, in particular, considerable uncertainty regarding their star formation history and so on. Moreover, the previous studies were based on small samples, up to about several hundreds, while the present modern digital sky surveys, such as the Sloan Digital Sky Survey (SDSS), greatly improved the observed numbers of astronomical objects, and could provide a much larger sample of LSBGs, hence giving much more information. Therefore, we propose a project to search for a large sample of LSBGs from the SDSS database, and then to study their properties in detail.  In this first paper of our series about the large sample of LSBGs selected from the SDSS, we will present how to select the sample and then show their basic properties and stellar populations from photometric colors. Then, we will present their spectroscopic properties (Liang et al. 2008, Chen et al. 2008, in preparation) and their detailed surface photometry and color gradients (Zhong et al. 2008, in preparation) etc. This paper is organized as follows. In Sect.~\\ref{sec.2}, we describe how to select the sample. In Sect.3, we present the relations between some property parameters. In Sect.4, we compare the LSBGs with the HSB disk galaxies sample selected simultaneously. In Sect.5, we discuss the incompleteness of the sample at the low surface brightness end and present the property relations of a volume-limited sub-sample. In Sect.6, we do cross correlation for the sample with the 2MASS database to obtain their optical and NIR color-color diagrams and the color-magnitude relations. The discussions and conclusions are given in Sect.7. Throughout the paper, a cosmological model with $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3  and $\\Omega_\\lambda$=0.7 is adopted. The units of surface brightness is mag~arcsec$^{-2}$.  \\section[]{The Sample} \\label{sec.2}  The SDSS is the most ambitious astronomical survey ever undertaken in imaging and spectroscopy (York et al. 2000; Stoughton et al. 2002; Abazajian et al. 2003, 2004). The imaging data are done in drift scan mode and are 95\\% complete for point sources at 22.0, 22.2, 22.2, 21.3, and 20.5 in five bands ($u$, $g$, $r$, $i$ and $z$) respectively. The spectra are flux- and wavelength-calibrated with 4096 pixels from 3800 to 9200 \\AA~at $R\\approx1800$.  The sample used in this work is selected from the main galaxy sample (MGS) of SDSS-DR4 (Strauss et al. 2002), which comprises galaxies with  $r$-band Petrosian magnitude $r\\leq$17.77 (corrected for foreground Galactic extinction using the reddening maps of Schlegel et al. 1998) and $r$-band Petrosian half-light surface brightness $\\mu_{50}\\leq$24.5 mag arcsec$^{-2}$. This sample has a median redshift of about 0.10.  The MGS is a spectroscopic sample selected from the SDSS photometric data. The completeness for objects with spectroscopy observations is high, exceeding 99\\%, and the fraction of galaxies eliminated by surface brightness cut is very small ($\\sim0.1\\%$). The only significant source of incompleteness identified is bending with saturated stars, which is higher for higher brightness because they subtend more sky (Strauss et al. 2002). Relative to all the SDSS targets, the SDSS spectroscopic survey is 90\\% complete (Blanton et al. 2003; Hogg et al. 2004; Strauss et al. 2002; McIntosh et al. 2006). The 7\\% incompleteness coming primarily from galaxies missed due to fiber collisions since the minimum separation of fiber centers is 55$\\arcsec$ (Blanton et al. 2003). It does lead to a slight underrepresentation of high-density regions (Hogg et al. 2004), however, where bright early-type galaxies usually dominate. Additionally, there are a small faction of redshift failures ($\\sim1\\%$) and some bright star contamination ($\\sim2\\%$) that becomes more significant for brighter galaxies (Strauss et al. 2002; McIntosh et al. 2006). However, the completeness of the sample at the low surface brightness end needs to be discussed carefully because of the known problem of surface brightness limit in a redshift survey and the focus of our work on the low surface brightness ones. We will specially discuss this incompleteness in Sect.~\\ref{sec.55} and will extract a volume-limited sub-sample, which is complete, from our whole sample galaxies to check what changes appearing in the related properties of the galaxies.   We prefer to use the DR4 rather than the latest DR6 here. That is because DR4 has been large enough for statistical analysis of LSBGs; other property parameters for galaxies in DR4 have also been obtained by the MPA/JHU group\\footnote{http://www.mpa-garching.mpg.de/SDSS/}, such as the emission-line measurements etc.; and the NYU-VAGC\\footnote{http://sdss.physics.nyu.edu/vagc/} catalog (Blanton et al. 2005) has made some useful cross correlations with other survey databases. These can provide us more information on both photometry and spectroscopy. Our sample galaxies are selected as follows.  \\subsection[]{The whole sample: the nearly face-on disk galaxies} \\label{sec.2.1}  To obtain reliable surface brightness values of the disk galaxies and minimize the effect of dust extinction inside the galaxies, we select the nearly face-on disk galaxies as a working sample by following the criteria below.  \\begin{enumerate}  \\item  $fracDev_r$ $<$ 0.25  The parameter $fracDev_r$ indicates the fraction of luminosity contributed by the de Vaucouleurs profile relative to exponential profile in the $r$-band. The topic in this work focuses on disk galaxies, whose surface brightness profiles usually can be well described by an exponential formula (e.g., Bernardi et al. 2005; Chang et al. 2006b; Shao et al. 2007). Therefore, we select the sample galaxies almost having an exponential light profile by requiring their $fracDev_r$ $<$ 0.25. This can minimize the effect of bulge light on the disk galaxies. Using this selection criterion, we can obtain 111,479 ($\\sim$28\\%) objects from the total 401,007 galaxies in the SDSS-DR4 main galaxy sample.  \\item $b/a$ $>$ 0.75  To select nearly face-on galaxies, we also require objects with the axis ratio $b/a$$>$0.75 (corresponding to the inclination $i<$41.41 degree), where $a$ and $b$ are the semi-major and semi-minor axes of the fitted exponential disk respectively. After this step is applied, 32,226 ($\\sim$29\\% of 111,479) galaxies are left.  \\item $M_B$ $<$ -18  Keeping the $B$-band absolute magnitude $M_B$ $<$ -18 in mind excludes the few dwarf galaxies contained in the sample. The $B$-band absolute magnitude $M_B$ is calculated using \\begin{equation} \\label{eq.Mb} M = m - 5\\log (D_L) - K(z), \\end{equation} where $m$ is the apparent magnitude that corrected for Galactic extinction using the reddening maps of Schlegel et al. (1998), $D_L$ is the luminosity distance calculated by using the Local Group relative redshift values from the NYU-VAGC catalog (Blanton et al. 2005, see Sect.~\\ref{sec.3.1}). The $K$-correction $K$($z$) are calculated by using the $K$-corrections program given on the NYU-VAGC website (Blanton et al. 2005). Finally, we obtain 30,333 galaxies as our whole sample after excluding $\\sim$6\\% of 32,226 nearly face-on disk galaxies by this criterion.  \\end{enumerate}  \\subsection[]{The subsamples: the LSBGs and HSBGs} \\label{sec.2.2}  A commonly used parameter to classify a galaxy to be one with low (or high) surface brightness is its $B-$band central surface brightness $\\mu_0(B)$. In this work, we adopted $\\mu_0(B)$$\\geq$22 mag arcsec$^{-2}$, a commonly applied criterion (Boissier et al. 2003), as LSBGs and $\\mu_0(B)$$<$22 mag arcsec$^{-2}$ as high surface brightness galaxies (HSBGs) following Impey, Burkholder \\& Sprayberry (2001) and O'Neil et al. (1997a,b).  The surface brightness profiles of disk galaxies are well approximated by an exponential disk profile with the form: \\begin{equation} \\Sigma(r) = \\Sigma_0 exp(-{{r}\\over{a}}), \\end{equation} where $\\Sigma_0$ is the surface brightness of the disk in units of $M_\\odot$ pc$^{-2}$ and $a$ is the disk scale-length measured in units of arcsec. In the logarithmic units, this equation becomes: \\begin{equation} \\mu_r = \\mu_0 + 1.086({{r}\\over{a}}), \\end{equation} where $\\mu_0$ is the central surface brightness in mag arcsec$^{-2}$.  In the analysis, the disk is assumed to be infinitely thin. The total flux is given by: \\begin{equation} F_{tot} = 2 {\\pi} a^2 \\Sigma_0. \\end{equation} Then, the total apparent magnitude can be determined from this equation. Converting this to a logarithmic scale is: \\begin{equation} \\mu_0 = m + 2.5\\log {(2 \\pi a^2)}, \\end{equation} where $m$ refers to the total apparent magnitude and $\\mu_0$ refers to the central surface brightness.  The central surface brightness is corrected by inclination and cosmological dimming effects by following: \\begin{equation} \\mu_0 = m + 2.5\\log {(2 \\pi a^2)} + 2.5\\log{(b/a)} - 10\\log{(1+z)}, \\end{equation} where $z$ is the redshift of the target.  \\begin{figure} \\centering \\includegraphics [width=7.5cm, height=5.2cm] {./fig/histgram.LSB.eps} \\caption{Histogram distribution of the selected nearly face-on disk galaxies (30,333) in $B$-band central surface brightness $\\mu_0(B)$. The whole sample is divided into two parts: the white region with $\\mu_0(B)$ $\\geq$22 mag arcsec$^{-2}$ as the LSBGs (12,282), and the shadowed region with $\\mu_0(B)$ $<$22 mag arcsec$^{-2}$ as the HSBGs (18,051). } \\label{fig.1} \\end{figure}   The $B$-band (also $U$, $V$, $R$ ) central surface brightnesses and magnitudes of the SDSS galaxies can be calculated from the $g$ and $r$-band quantities by using the conversions provided by Smith et al. (2002). The histogram distribution of $\\mu_0(B)$ for our selected whole sample (30,333 galaxies) is shown in Fig.~\\ref{fig.1}. The blank histogram is for the LSBGs and the shadowed one is for the HSBGs with the separating value of 22 mag arcsec$^{-2}$. We call the LSBGs as {\\bf Sample-L} and the HSBGs as {\\bf Sample-H} in this work. Sample-L includes 12,282 ($\\sim$40\\%) galaxies with 22$\\leq$$\\mu_0(B)$$\\leq$24.5 mag arcsec$^{-2}$ with median and mean values of 22.42 and 22.50 mag arcsec$^{-2}$ respectively. Sample-H includes 18,051 ($\\sim$60\\%) galaxies with 18.5$\\leq$$\\mu_0(B)$$<$22 mag arcsec$^{-2}$ with median and mean values of 21.41 and 21.29 mag arcsec$^{-2}$ respectively. Since the LSBGs are the major topic of this paper, we will carefully present the properties of Sample-L in Sect.~\\ref{sec.3} and roughly compare the HSBGs with the LSBGs in Sect.~\\ref{sec.4}. However, when looking for the relations of some property parameters with surface brightness, it is useful to show all of the sample, including LSBGs and HSBGs, therefore, in Sect.3.2 we present both the LSBGs and HSBGs in the relations of $m_B$, distance, $M_B$ with $\\mu_0$(B). The cross-correlations with near infrared 2MASS data will be emphasized in Sect.~\\ref{sec.5}.  \\section[]{Property parameters and relations of the large sample of LSBGs} \\label{sec.3}  Some property parameters of LSBGs have been obtained and discussed in previous work, such as Impey et al. (1996, 2001), de Blok et al. (1995), Bothun et al. (1997), Galaz et al. (2002, 2006) etc. In this section, we present some property parameters and relations for our Sample-L (the LSBGs) and compare them with the previous results. The Sample-H (HSBGs) are also given in the relations of $m_B$, distance, $M_B$ with $\\mu_0$(B) (see Sect.3.2 and Fig.~\\ref{fig.3}), which is useful to show the entire sample.  \\subsection[]{The histogram distributions of some property parameters} \\label{sec.3.1}  We show the distributions of some property parameters of our LSBGs in Fig.~\\ref{fig.2}. Fig.~\\ref{fig.2}a shows the histogram distribution of their redshift $z$ obtained from the NYU-VAGC catalog (Blanton et al. 2005), which are the peculiar velocity corrected Local Group relative redshifts from SDSS. It shows that the redshifts are from 0 to 0.3 with median and mean values of 0.080 and 0.089, respectively. Fig.~\\ref{fig.2}b shows the histogram distributions of the disk scale-length $a$ in arcsec, and Fig.~\\ref{fig.2}c shows the histogram distributions of the disk scale-length $h$ in kpc. The range of $a$ is from 2 to 13 arcsec with median and mean values of 3.78 and 4.19 arcsec, and $h$ is from 2 kpc to 19 kpc with median and mean values of 5.98 and 6.30 kpc, respectively. The scale lengths of LSBGs investigated here are in the same ranges as in McGaugh (1992) and de Blok et al. (1995). Fig.~\\ref{fig.2}d shows the histogram distribution of the $B$-band absolute magnitudes, which is calculated using Eq.~(\\ref{eq.Mb}). It shows that the LSBGs are not necessary faint galaxies, but rather they are quite bright, $M_B$ is from -18 to -23 mag with the median and mean values of -20.08 and -20.06 mag, respectively.  \\begin{figure} \\centering \\includegraphics [width=7.8cm, height=7.5cm] {./fig/histgram.about_LSB.eps} \\caption{Histogram distributions of some property parameters of Sample-L (the LSBGs): (a). the redshift, obtained from NYU-VAGC catalog and is the peculiar velocity-corrected Local Group relative redshift from SDSS; (b). the disk scale length (the disk semi-major axis) in arcsec as $a$; (c). the disk scale length (the radius) in kpc as $h$; (d). the $B$-band absolute magnitude $M_B$, which obviously shows that LSBGs does not necessarily have low luminosity, and the LSBGs can be very luminous.} \\label{fig.2} \\end{figure}  \\subsection{Distributions of some property parameters with $B$-band central surface brightness}  We present the relations between the total $B$-band apparent magnitudes, distance, absolute magnitude $M_B$ and the $B$-band central surface brightness $\\mu_0(B)$ for Sample-L and Sample-H.  \\subsubsection{central surface brightness and apparent magnitude}   Figure~\\ref{fig.3}a presents the relations between the $B$-band apparent magnitudes and the $B$-band central surface brightness $\\mu_0(B)$. The vertical lines separate the LSBGs and HSBGs with $\\mu_0(B)$=22 mag arcsec$^{-2}$. It shows that these galaxies have the $m_B$ from 14 to 19 mag. At a given magnitude, most galaxies are detectable over a wide range of surface brightnesses, for example, 18$<$$\\mu_0(B)$$<$24.5 mag arcsec$^{-2}$ at $m_{B}$$=$18 mag. There exists almost no correlation within the ranges. Some objects missed probably at faint $\\mu_0$ could be due to the poor contrast with the sky brightness. These results are similar to those shown by the samples of APM survey (693 LSB field galaxies) given by Impey et al. (1996).  \\subsubsection{central surface brightness and distance}  Figure~\\ref{fig.3}b presents the distribution of the distance versus $\\mu_0(B)$ of the sample galaxies. For the LSBGs, the left part of the vertical line, it shows that the range of surface brightness for the nearest galaxies is substantially larger than that for the more distant galaxies. We expect this to influence surface brightness selection, where galaxies with higher surface brightness could be seen at the largest distances (Disney $\\&$ Phillips 1983, Impey et al. 1996). But for the HSBGs, the right part of the vertical line, it does not obviously show such trend. This may mean that the selection bias affects much the LSBGs. This distribution is consistent with that of the sample in the APM survey presented by Impey et al. (1996). However, the volume-limited sub-sample of LSBGs do not show such a trend (see Sect.5).  \\subsubsection{central surface brightness and absolute magnitude}  Figure~\\ref{fig.3}c presents the distribution of the absolute $B$-band magnitudes versus $\\mu_0(B)$ of the sample galaxies. The left part from the vertical line is for the LSBGs. There exists a correlation among them showing that the fainter galaxies tend to have lower surface brightnesses though there is large scatter. It also shows over the whole range of 22$<$$\\mu_0(B)$$<$24.5 mag arcsec$^{-2}$, galaxies are being discovered over the entire range of -23$<$$M_B$$<$-18 mag. The right part from the vertical line is for the HSBGs, which show that there is no obvious correlation between their $M_B$ and $\\mu_0(B)$. This could mean that the selection bias affects the LSBGs greatly. These are also similar to that presented by the APM sample given by Impey et al. (1996).   \\begin{figure} \\centering \\includegraphics[bb= 56 128 252 728, width=6.0cm, height=11.2cm, clip] {./fig/all_about_uB.eps} \\caption{Distributions of some property parameters with $\\mu_0(B)$ for our LSBGs (the left part of the vertical line) and HSBGs (the right part of the vertical line): (a). the total $B$-band apparent magnitudes; (b). the distance of the galaxies; (c). the $B$-band absolute magnitudes. The solid lines are the median-value points within the bins of 0.10 of $\\mu_0(B)$ less than 23.6 mag arcsec$^{-2}$, and we take $\\mu_0(B)$ greater than 23.6 mag arcsec$^{-2}$ as one bin because there are very few sources. The lines refer to the median-value points within the bins. } \\label{fig.3} \\end{figure}  \\subsection{Distributions of some property parameters with disk scale length and distance}  Some of the previous studies have shown there are obvious correlations between the luminosity and disk scale length log$h$ (Impey et al. 2001; Bergvall et al. 1999), distance log$D$ and disk scale length log$h$ (Impey et al. 2001) for LSBGs. These relations become more clear from our much larger and homogeneous sample.  \\subsubsection{disk scale length and luminosity} \\label{sec.3.3.1} We show the relation between absolute magnitudes $M_B$ and disk scale length log$h$ for our Sample-L in Fig.~\\ref{fig.4}. It can be seen that there is a very clear correlation between them, which can be fitted by a least-square fit as \\begin{eqnarray} \\log h = -0.150(\\pm 0.0007) M_B - 2.245(\\pm 0.014), \\label{eqhMb} \\end{eqnarray} with a small standard deviation of 0.070 dex. The Spearman-Rank Order correlation coefficient is -0.893.  This correlation has been shown in Impey et al. (2001, for their 238 galaxies from APM) and Bergvall et al. (1999, for the gathered samples from literature), but is more obvious for our larger and homogeneous sample. It shows brighter galaxies tend to have larger scale lengths.  It is worth comparing such relations of our LSBGs with the corresponding relations of HSBGs. They show some discrepancies, which will be specially discussed in Sect.~\\ref{sec.4}.  \\begin{figure} \\centering \\includegraphics [width=8.0cm, height=5.8cm] {./fig/MB_h.LSB.eps} \\caption{Relations of the LSB sample galaxies about their disk scale length in logarithm versus $B$-band absolute magnitudes. The solid line refers to the least-square fit for the sample galaxies given by Eq.(\\ref{eqhMb}). The Spearman-Rank Order correlation coefficient is -0.893.} \\label{fig.4} \\end{figure}  \\subsubsection{disk scale length and distance}  Our LSBGs sample also show a very clear correlation between their distance and disk scale length, which is given in Fig.~\\ref{fig.5}. This correlation can be given by  a least-square fit as: \\begin{eqnarray} \\log h = 0.511(\\pm 0.004) \\log D - 0.536(\\pm 0.010), \\label{eqhD} \\end{eqnarray} with the standard derivation of 0.10 dex. The Spearman-Rank Order correlation coefficient is 0.767.  This relation obviously show the selection effect, i.e., in the more distant Universe, the galaxies with larger scale lengths have the benefit that they can be detected and observed.   \\begin{figure} \\centering \\includegraphics [width=8.0cm, height=5.8cm] {./fig/Dl_h.LSB.eps} \\caption{Relations of the LSB sample galaxies about their distance versus disk scale length. The solid line refers to the least-square fit for the sample galaxies given by Eq.(\\ref{eqhD}). The Spearman-Rank Order correlation coefficient is 0.767.} \\label{fig.5} \\end{figure}  \\subsubsection{distance and luminosity}  Figure~\\ref{fig.13} shows the distribution between the distance log$D$ and the $B$-band absolute magnitude $M_B$. It could be seen that there is an obvious correlation between them, though there is some scatter. This correlation can be shown by a least-square fit as: \\begin{eqnarray} \\log D = -0.211(\\pm 0.001) M_B - 1.675(\\pm 0.025), \\label{eq.DMb} \\end{eqnarray} with the standard derivation of 0.13 dex. The Spearman-Rank Order correlation coefficient is -0.847.  \\begin{figure} \\centering \\includegraphics [width=8.0cm, height=5.8cm] {./fig/MB_D.LSB.eps} \\caption{Distribution of the distance log$D$ and the $B$-band absolute magnitude $M_B$ for our LSBGs. The solid line refers to the least-square fit for the sample galaxies given by Eq.(\\ref{eq.DMb}). The Spearman-Rank Order correlation coefficient is -0.847.} \\label{fig.13} \\end{figure}  This correlation show the selection effects of the survey observation. That means the brighter galaxies can be observed at larger distances, and the relatively fainter ones can only be observed at smaller distances.   \\subsubsection{disk scale lengths and central surface brightness}  Figure~\\ref{fig.6} shows the distribution between disk scale length log$h$ and the central surface brightness $\\mu_0(B)$, which does not show obvious correlation. This result is consistent with those of Beijerbergen et al. (1999) and O'Neil et al. (1997b) for smaller samples.  \\begin{figure} \\centering \\includegraphics [width=8.0cm, height=5.8cm] {./fig/uB_lo_h.LSB.eps} \\caption{Distribution of the disk scale length versus the $B$-band surface brightness for our LSBGs. The solid line is the median-value points within the bins that is the same as Fig.~\\ref{fig.3}. } \\label{fig.6} \\end{figure}  \\subsection{Relations between colors and disk scale length, colors and central surface brightness }  In this section we show the relations between colors and disk scale length, as well as between colors and central surface brightness for the LSBGs.  \\subsubsection{colors and disk scale lengths }  Figure~\\ref{fig.7}a presents the distribution between $(B-V)$ color and disk scale length of our large sample of LSBGs. The correlation between them is weak even though it exists, and the Spearman-Rank Order correlation coefficient between them is only 0.302. Our result is not much different from the previous studies on small samples (e.g., de Blok et al. 1995; Beijersbergen et al. 1999), which showed the colors do not depend on sizes of LSBGs. But Bergvall et al. (1999) have showed an obvious correlation for a small sample of spirals, large LSBGs and their blue LSBGs.  \\begin{figure} \\centering \\includegraphics [width=7.8cm, height=8.6cm] {./fig/about_BV.LSB.eps} \\caption{(a). Distribution of $B - V$ color as a function of disk scale length for the Sample-L. The solid line is the median-value points within the bins of 0.10 of $\\log h$ from 0.5 to 1.0 ($h$ in kpc). But $\\log h$ less than 0.5 or greater than 1.0 ($h$ in kpc) as one bin, respectively, since the sources are very few in those two bins. (b). Distribution of $B - V$ color as a function of $B$-band central surface brightness for the Sample-L. The solid line is the median-value points within the bins that is same as Fig.~\\ref{fig.3}. } \\label{fig.7} \\end{figure}  \\subsubsection{colors and central surface brightness}  Figure~\\ref{fig.7}b presents the distribution between the $(B-V)$ colors and central surface brightness $\\mu_0(B)$ for our Sample-L, which does not show obvious correlation in the whole ranges of the two parameters. The results presented from smaller samples of LSBGs studied by Beijersbergen et al. (1999), de Blok et al. (1995) and Bothun et al. (1997) are consistent with our result here from a large sample of LSBGs.   \\subsection{The optical color-color diagram and stellar populations}  \\subsubsection{The negligible affect of dust extinction on colors} \\label{sec.3.5.1}  In this work, we have selected the nearly face-on disk galaxies to minimize the affection of dust extinction. Thus their properties of stellar population won't be affected much by dust extinction, which could be confirmed by the independence of the colors with inclination. Fig.~\\ref{fig.12} shows such independence of the LSB sample galaxies (12,282). Indeed the HSB sample galaxies present similar independence, which were not presented as plots here to save space. Additionally, we also do cross-correlations for these nearly face-on disk galaxies with the IRAS infrared PSCz catalog within a range of 5 arcsec, but only find very few IR-detected objects, which also confirm the negligible affect of dust on their properties. This negligible affect of dust on colors is consistent with other studies (e.g., Bell et al. 2000; Galaz et al. 2006).  \\begin{figure} \\centering \\includegraphics [width=8.0cm, height=5.8cm] {./fig/about_i.eps} \\caption{ Relations between disk inclination and colors of the Sample-L: (a). $U-B$, (b). $B-V$. The solid lines are the least-square fits for the data, which show almost no slopes. } \\label{fig.12} \\end{figure}  \\subsubsection{$UBV$ diagram and stellar populations}  O'Neil et al. (1997a,b) have performed a digital survey for LSBGs in the spiral-rich Cancer and Pegasus clusters as well as the low density regime defined by the Great Wall (the ``Texas survey\"). They found a total of 127 galaxies of angular diameter larger than 15 arcseconds with $\\mu_0(B)$$\\geq$22 mag arcsec$^{-2}$. They discussed the stellar populations of this sample of LSBGs by optical color-color relations and firstly found the very red population of LSBGs. They divided the galaxies into three categories: the very blue (with $(U-B) < -0.2, (B-V) < 0.6$), the very red (with $(U-B) > 0.3, (B-V) > 0.8$), and the rest that have other color values.  In our Fig.~\\ref{fig.8}, we shows the $(U-B)$ vs. $(B-V)$ diagram for our Sample-L. The average values of our sample galaxies are $<(U-B)> = -0.0012 \\pm 0.08$, $<(B-V)> = 0.64 \\pm 0.06$ in the ranges of $-0.6 \\leq (U-B) \\leq 0.6$ and $0 \\leq (B-V) \\leq 1.2$. These are not very different from O'Neil's sample, and is similar to those of Romanishin et al. (1983) and McGaugh $\\&$ Bothun (1994) as well.  \\begin{figure} \\centering \\includegraphics [width=8cm, height=5.8cm] {./fig/UB_BV.LSB.eps} \\caption{$U-B$ versus $B-V$ diagram of our LSBGs. The solid and dashed lines are taken from O'Neil et al. (1997b) to define the ``very blue\" and ``very red\" populations of galaxies (see text). } \\label{fig.8} \\end{figure}  We also plot the separation lines suggested by O'Neil et al. (1997a,b) for the ``very blue\" and ``very red\" LSBGs in Fig.~\\ref{fig.8} (the solid and dashed lines, see above). Our Sample-L contains 535 ($\\sim 4.4\\%$) ``very blue\" LSBGs and 69 ($\\sim 0.48\\%$) ``very red\" LSBGs. Most of the sample galaxies are the ``rest case\" with mixed stellar populations.  Heller $\\&$ Brosch (2001) also present such $UBV$ colors for their 29 LSB dwarf irregular galaxies from the Virgo cluster, and found their ``very blue\" is about 30\\%  and no ``very red\" ones. These two different fractions could be easily understood since their samples are dwarf irregular galaxies, which should be blue. de Blok et al. (1995) also present such a $UBV$ diagram for a small sample of 21 LSB galaxies gathered from literature, and they did not find any sample belonging to the ``very red\" range that O'Neil et al. suggested.  The colors of these LSBGs are obviously different from those of the E0/S0 galaxies but not much different from the Sc or later spiral galaxies with active star formation. As O'Neil et al. (1997b) commented, the average colors of E0/S0 galaxies (good examples of old stellar populations) have $(U-B)$=0.54, $(B-V)$=0.96 (Tinsley 1978); Sc or later spiral galaxies with active star formation typically have colors in the range of 0.35 $\\leq$ $(B-V)$ $\\leq$ 0.65, -0.2 $\\leq$ $(U-B)$ $\\leq$ 0.4 (Huchra 1977; Bothun 1982), with mean values $<(B-V)> \\sim$0.50, $<(U-B)> \\sim$-0.20 (Bothun 1982), and $<(V-I)> \\sim$1.0 (Han 1992).  The colors of our galaxies range continuously from very blue to very red and include a group of old galaxies, which show evidence of recent star formation. These LSBGs at the present epoch define a wide range of evolutionary states. This wide range is similar to that of the high surface brightness galaxies (see Sect.~\\ref{sec.4}). ",
        "conclusions": "\\label{sec.6}  Low surface brightness galaxies (LSBGs) are important populations in the Universe. They were studied little in the past due to the observational limits of their low surface brightness. The fantastic modern SDSS survey of large sky coverage area is unique for such studies because it detects as many low surface brightness galaxies. We select a large sample of 12,282 LSBGs from the main galaxy sample of SDSS-DR4 database, and present their basic properties in this paper. The results are summarized as follows.  \\begin{enumerate}  \\item Firstly, we select 30,333 nearly face-on disk galaxies with $M_B<-18$ from the main galaxy sample of SDSS-DR4. Their surface brightness profiles can be fitted well by an exponential disk. The $B-$band central surface brightness of these disk galaxies distribute in a range of 18.5 $<$$\\mu_0(B)$$<$ 24.5 mag arcsec$^{-2}$. This sample was then divided into two sub-samples: Sample-L (12,282, the LSBGs) and Sample-H (18,051, the HSBGs) by the surface brightness value of 22 mag arcsec$^{-2}$. The median $\\mu_0(B)$ of Sample-L is 22.42 mag arcsec$^{-2}$ and the mean value is 22.50 mag arcsec$^{-2}$. The median and mean values of $\\mu_0(B)$ of Sample-H are 21.41 and 21.29 mag arcsec$^{-2}$, respectively.  \\item The scale length $h$ of Sample-L are from 2 kpc to 19 kpc with the median scale length of 5.98 kpc. These values are similar to those scale lengths of Sample-H, which means that the LSBGs and HSBGs are comparable in size.  \\item The absolute $B$-band magnitude $M_B$ of Sample-L are from -18 to -23 mag with the median value of -20.08 mag and the mean value of -20.06 mag. These values are similar to those $M_B$ of Sample-H (-20.47 and -20.41 mag respectively), which means that the LSBGs are not necessarily faint, and their luminosities are comparable to those of the HSBGs.  \\item Some relations between the property parameters show the selection effects of the survey observations, such as the apparent magnitude vs. $\\mu_0(B)$, distance versus $\\mu_0(B)$,  absolute magnitude versus $\\mu_0(B)$, distance versus disk scale length, and the distance versus absolute magnitude. These relations show that the galaxies with higher surface brightness tend to be detected at farther distance; and the galaxies observed at farther distance are also brighter and have larger disk scale lengths.  \\item A fundamental correlation to the sample galaxies is the disk scale length log$h$ versus absolute magnitude $M_B$. For Sample-L, the scale length log$h$ and absolute magnitude $M_B$ show very tight correlation, meaning that the brighter galaxies tend to be larger. The tight correlation between log$h$ and $M_B$ also exists in the HSGBs, but with a slightly steeper slope, which means that, at a given $M_B$, the observed LSBGs tend to be larger than the HSBGs, and for a given size of galaxies, the observed HSBGs tend to be brighter than the LSBGs.  \\item  There is no obvious correlation between colors and central surface brightness $\\mu_0(B)$ for our large sample of LSBGs. The correlation between colors and disk scale length (log$h$) is weak even though it exists.  \\begin{figure} \\centering \\includegraphics [width=7.8cm, height=9.0cm] {./fig/MBhD.eps} \\caption{(a). Correlation of disk scale length and luminosity for the whole sample galaxies: LSBGs (the solid line); ISBGs (the dot-long dash line); HSBGs (the long-dash line); VHSBGs (the short-long dash line). (b). Correlation of disk scale length and distance of the whole sample galaxies, and the lines represent the same means as (a). } \\label{fig.14} \\end{figure}  \\item The color-color diagrams of $(U-B)$ vs. $(B-V)$ show that most of our LSBGs have a mix of young and old stellar populations. Following the definitions of O'Neil et al. (1997a,b), our sample may contain about 4.4\\% (535 out of 12282) ``very blue\" LSB galaxies with $(U-B)$$<$-0.2 and $(B-V)$$<$0.6, and $\\sim$0.48\\% (69 out of 12282) ``very red\" LSB galaxies with $(U-B)$$>$0.3 and $(B-V)$$>$0.8. These fractions are 4.0\\% and 2.1\\% for our HSBGs, respectively.  \\item A volume-limited sub-sample was extracted by considering $z<0.1$ and brighter than the corresponding $M_r$, which is complete and can be used to check the incompleteness of surface brightness (especially at the low end), and to check how the property relations change. The only one with obvious change is the relation of log$h$ versus $\\mu_0$(B), which shows a correlation in this sub-sample.  \\item The optical-NIR color-color diagrams can break the degeneracy of age and metallicity. By doing cross-correlations between the SDSS and 2MASS datasets, we obtain the NIR color sample of Sample-L2 (15.29\\% of the Sample-L), and such sample of Sample-H2 (29.47\\% of the Sample-H). The smaller detected fraction of LSBGs is not unexpected.  The color-color diagram of $(R-K)$ vs. $(B-R)$ clearly shows the stellar populations of the sample galaxies by comparing with the predictions of the stellar population synthesis models. This shows that our LSBG samples have a wide range of ages and metallicities, from Z=0.0004 to Z=Z$_{\\odot}$, even Z=0.05, and the exponential infall time-scale $\\tau$ is from 0 to 16 Gyrs. It confirms that most LSBGs have a mix of young and old stellar populations.  \\begin{figure} \\centering \\includegraphics [width=7.5cm, height=7.2cm] {./fig/UBV_counter.eps} \\caption{ $U-B$ versus $B-V$ diagram: (a). LSBGs, (b). ISBGs, (c). HSBGs, (d). VHSBGs. The horizontal and vertical lines are the same as in Fig.~\\ref{fig.8}. The contours and the squares are the same as in Fig.~\\ref{fig.UBV2mass} (please see the online color version for more details). } \\label{fig.15} \\end{figure}  \\item The color-magnitude relations ($(r-K)$ vs. $M_r$) of our LSBGs confirm that there exist the CMR for the spiral galaxies with low surface brightness, showing that the brighter LSBGs tend to be redder. And, the HSBGs do not show obvious differences from the LSBGs on their CMRs.  \\item We have attempted to divide our total nearly face-on disk galaxies (30,333) into four sub-groups according to their surface brightness by following McGaugh (1996): the VHSBGs with $\\mu_0(B)$$<$21.25 mag arcsec$^{-2}$, the HSBGs with 21.25$<$$\\mu_0(B)$$<$22 mag arcsec$^{-2}$, the ISBGs with 22$<$$\\mu_0(B)$$<$22.75 mag arcsec$^{-2}$) and the LSBGs with 22.75$<$$\\mu_0(B)$$<$24.5 mag arcsec$^{-2}$. Similar to the comparisons between LSBGs and HSBGs in Sect.~\\ref{sec.4}, these four sub-groups do show obvious difference in the relations of log$h$ vs. $M_B$, log$h$ vs. log$D$, and some difference in the color-color diagram of $(U-B)$ vs. $(B-V)$ as given in Fig.~\\ref{fig.14} and Fig.~\\ref{fig.15}. Fig.~\\ref{fig.14}a shows that, at a given $M_B$ or scale length log$h$, the observed galaxies in the four sub-groups show the gradual increase in their scale length or gradual decrease in their luminosity following the decreasing surface brightness. The four lines refer to the linear least-square fits for the data in the four sub-groups. To show the scatter of the data as well, we obtained their standard derivations as 0.057, 0.049, 0.050, 0.093 dex for the LSBGs, ISBGs, HSBGs and VHSBGs, respectively. Fig.~\\ref{fig.14}b also shows such gradual changes in their relations of scale length and distance. Similarly, the four lines refer to the linear least-square fits for the data in the four sub-groups, and the standard derivations of the data to the fittings are 0.092, 0.096, 0.098, 0.13 dex for the LSBGs, ISBGs, HSBGs and VHSBGs, respectively. Fig.~\\ref{fig.15}a-d presents the discrepancies between their stellar populations among the four sub-groups. Even with large scatter, the LSBGs do generally show slightly bluer ($B-V$) color than others, and there are almost no the ``red populations\" in the LSBGs. It also shows the VHSBGs consist of more ``red populations\"  than other groups. The counters refer to the 68.3\\% confidence level of the data points, which show the gradually redder $(B-V)$ colors of the sub-groups following their increasing surface brightness.  In summary, this large sample (12,282) of LSBGs greatly extends the known sample of LSBGs up to date. This is the first paper of our series work about the large sample of LSBGs to present the distributions of some property parameters of them, and to provide more information about some of their fundamental properties, e.g. the tight correlations between log$h$ and $M_B$, log$h$ and log$D$, moreover, it presents their stellar populations by using the color-color diagrams of them.  In the following work, we will discuss more about the major issues of these ``populations\", such as why they are LSBGs with small star formation rates, what is their total mass content and M/L ratios, and what is their total contribution to the baryonic/total mass etc. We will also obtain their metallicities in interstellar gas from emission lines. The detailed surface photometry and color gradients will be also studied in detail. Moreover, some more efforts on other aspects are also needed, such as to extend the studies of edge-on galaxies, but the dust effects must be considered carefully then.  \\end{enumerate}"
    },
    "0809/0809.3785_arXiv.txt": {
        "abstract": "{ We present the detection of deuterated molecular hydrogen (HD) in the remote Universe in a damped Lyman-$\\alpha$ cloud at $z_{\\rm abs}=2.418$ toward the quasar SDSS\\,J143912.04$+$111740.5. This is a unique system in which H$_2$ and CO molecules are also detected. The chemical enrichment of this gas derived from Zn\\,{\\sc ii} and S\\,{\\sc ii} is as high as in the Sun. We measure $N({\\rm HD})/2 N({\\rm H_2})=1.5^{+0.6}_{-0.4}\\times10^{-5}$, which is significantly higher than the same ratio measured in the Galaxy and close to the primordial D/H ratio estimated from the WMAP constraint on the baryonic matter density ($\\Omega_{\\rm b}$). This indicates a low astration factor of deuterium that contrasts with the unusually high chemical enrichment of the gas. This can be interpreted as the consequence of an intense infall of primordial gas onto the associated galaxy. Detection of HD molecules at high-$z$ also opens the possibility to obtain an independent constraint on the cosmological-time variability of $\\mu$, the proton-to-electron mass ratio. } ",
        "introduction": "Deuterium is produced by primordial nucleosynthesis and is subsequently destroyed in stars. Therefore measurements of D/H from primordial gas provide important constraints on the baryonic matter density $\\Omega_{\\rm b}$ in the framework of Big-Bang cosmology \\citep{Wagoner73}. Measurements at different redshifts in turn provide important clues on the star formation history \\citep{Daigne04,Steigman07}. All the available D/H measurements at high-$z$ are based on the determination of the $N($D$^0)/N($H$^0)$  ratio in low-metallicity QSO absorption line systems. These measurements are difficult mainly because the velocity separation between D\\,{\\sc i} and H\\,{\\sc i} absorption lines is small ($\\Delta v_{\\ion{D}{i}/\\ion{H}{i}}\\sim80$~km\\,s$^{-1}$) implying the lines are easily blended. An additional difficulty is the presence of the Lyman-$\\alpha$ forest, making hard to find the true continuum position and to discern between D\\,{\\sc i} absorption lines and intervening Ly-$\\alpha$ forest lines. This explains why, despite more than a decade of efforts, only seven robust measurements of D/H at high-redshift have been performed \\citep{OMeara06,Pettini08a}. It happens that these measurements are consistent with the value [D/H]$_{\\rm p}$ = 2.55$\\pm 0.10\\times10^{-5}$ derived from the Wilkinson Microwave Anisotropy Probe (WMAP) five year data \\citep{Komatsu08} together with the $\\eta\\leftrightarrow$ [D/H]$_{\\rm p}$ conversion from \\citet{Burles01}, where $\\eta$ is the baryon-to-photon ratio. We consider this [D/H]$_{\\rm p}$ value as the primordial D abundance in this work. \\par Damped Lyman-$\\alpha$ systems (DLAs) are absorbers with the highest neutral hydrogen column densities among H\\,{\\sc i} absorption line systems ($N($H$^0)\\ga10^{20}$~cm$^{-2}$) and are thought to arise from the neutral interstellar medium (ISM) of distant galaxies \\citep{Wolfe05}. Only a small fraction ($\\sim$10-15\\%) of them show detectable amounts of H$_2$ \\citep{Ledoux03,Noterdaeme08}. Among the 14 high-$z$ H$_2$-bearing DLAs known to date, only two show HD absorption \\citep{Varshalovich01,Srianand08}. The $z_{\\rm abs}=2.418$ system toward SDSS\\,J143912$+$111740 we present here is the only one where, in addition to H$_2$ and HD, CO is detected. From the excitation of CO it is possible to derive in a straightforward way an accurate estimate of the Cosmic Microwave Background Radiation temperature at the redshift of the absorber \\citep{Srianand08}. Here we focus on the HD/H$_2$ ratio and its implications for the star-formation history in the DLA galaxy. ",
        "conclusions": "We reported the detection of HD molecules at $z_{\\rm abs}=2.418$ toward SDSS\\,J143912$+$111740, following a careful selection of quasars in the SDSS database and intensive observations with UVES at the Very Large Telescope. The system presents characteristics very similar to what is observed in the solar neighbourhood. We find HD/2H$_2=1.5\\times10^{-5}$, consistent with an astration factor of deuterium less than 1.7 which is contrasting with the high chemical enrichment. This is best explained by a scenario in which the gas that goes through star formation is replenished by continuous infall of ambient primordial gas. Similar results arise from recent numerical simulations and semi-analytical models \\citep[e.g.][]{Keres05,Erb08}. Interestingly, dynamical studies of nearby galaxies \\citep{Fraternali08} as well as interpretation of Ly-$\\alpha$ blobs at high redshift \\citep{Weidinger05,Nilsson06,Dijkstra06} provide independent observational evidences for accretion of gas onto massive galaxies.  We finally stress the importance of detecting similar systems to probe the time-variation of the proton-to-electron mass ratio from HD lines. Although such a independent test would be welcome to characterize possible unknown systematics, significantly higher signal-to-noise ratio is required to obtain limits comparable to those obtained with other techniques.  \\acknowledgement{We thank an anonymous referee for comments and suggestions that improved the quality of the paper. We warmly thank the ESO Director Discretionary Time allocation committee and the ESO Director General, Catherine Cesarsky, for allowing us to carry out these observations. PPJ and RS gratefully acknowledge support from the Indo-French Centre for the Promotion of Advanced Research (Centre Franco-Indien pour la Promotion de la Recherche Avanc\\'ee).}"
    },
    "0809/0809.1439_arXiv.txt": {
        "abstract": "We use a set of high-resolution simulations of scale-free Einstein-de Sitter cosmologies to investigate the logarithmic slope of the phase-space density profile $Q(r) = \\rho(r)/\\sigma^3(r)$ of dark matter (DM) haloes. The initial conditions for the simulations are determined by a power law power spectrum  of the form $P(k) \\propto k^n$.  We compute the $Q(r)$ profiles using the radial, tangential and full velocity dispersion, and the velocity anisotropy parameter, $\\beta(r)$. We express $Q(r)$ as a single power-law $Q(r) \\propto r^\\alpha$ and derive a median slope $\\alpha$ in each simulation and for each definition of $Q$.  Our main findings are: 1. The various $Q(r)$ profiles follow a power law to a good approximation.  2. The slopes depend on the concentration parameter $c$ of the DM haloes, where for $c \\gtrsim 10$ the slopes steepen with rising concentration and for $c \\lesssim 10$ the trend flattens and even turns around.  3. The asymptotic value of $\\beta$ as $r\\rightarrow R_{\\mathrm{vir}}$ increases with the value of $c$.  4. In accordance with \\citet{Zait2007} $\\alpha_{\\mathrm{rad}}$ becomes more negative as the asymptotic value of $\\beta$ at the virial radius increases.  5. This introduces a weak dependence of the $Q(r)$ slopes on the slope of the power spectrum. ",
        "introduction": "For the density profile $\\rho(r)$ of dark matter (DM) haloes a variety of empirical fitting formulae exist.  The one applied mostly is the so-called NFW \\citep{Navarro1996,Navarro1997} profile, which gives a fair description for haloes found in a wide range of $N$-body simulations.  Recent studies though are hinting to use a different functional form for the density profile that differs from the usual NFW form mainly in the central regions \\citep[see, e.g.][]{Navarro2004,Merritt2006}, but common to all suggested profiles is the fact that they are not described by a simple power-law but rather require a smoothly changing logarithmic slope.  It comes therefore as a surprise that the coarse-grained phase-space density\\footnote{See, e.g., \\citet{Dehnen2005} for clarifications about the applicability of the term 'phase-space density', as the term is commonly used, however, we will continue to refer to $Q(r)$ as the phase-space density profile.} profile, calculated from the density- and the velocity dispersion profile $\\sigma(r)$ as \\begin{equation} \\label{eq:qr} Q(r) = \\frac{\\rho(r)}{\\sigma^3(r)}, \\end{equation} very closely follows a {\\em single} power-law $Q(r) \\propto r^\\alpha$ with $\\alpha \\approx -1.9$.  This was first noted by \\citet{Taylor2001} and has since been confirmed by many other works \\citep[e.g.][]{Boylan-Kolchin2004, Rasia2004, Ascasibar2004, Austin2005, Dehnen2005, Hoffman2007, Ascasibar2008}.  Different velocity dispersions have been used with equation~\\ref{eq:qr}, generally the radial velocity dispersion $\\sigma_{\\mathrm{rad}}(r)$ is used, but also the total velocity dispersion $\\sigma_{\\mathrm{tot}}(r)$ has been employed.  Both lead to separate definitions for $Q$ that nevertheless can be well fitted with a single power-law, albeit marginal differences in the slopes \\citep{Dehnen2005, Faltenbacher2007, Ascasibar2008}, see also the discussion in~\\citet{Zait2007}.  The power law phase-space density profile has been found mostly in simulations of structure formation in the standard model of cosmology, namely the flat $\\Lambda$CDM model.  It is not known how general its behaviour in other cosmologies is and the dynamical origin of the power-law is not yet understood.  However, the structure of (spherical) haloes in equilibrium obeys the (spherical) Jeans equation which relates the density profile, the phase-space density profile and the velocity anisotropy parameter $\\beta$ (to be defined below).  As mentioned above, the universality of the density profile is well-established across a large variety of cosmological models.  The description of the phase-space density profiles as a single power-law, however, is not that well checked.   Hence we aim here at studying the (phase-space) structure of DM haloes in different cosmogonies, in particular scale-free Einstein-de Sitter cosmologies with varying spectral index $n$, to explicitly probe universality of the phase-space profile.  We use the numerical simulations of \\mbox{\\cite{Knollmann2008}} who studied the structure of haloes in pure DM flat cosmologies with a primordial power spectrum of the form $P(k) \\propto k^n$, finding the equilibrium haloes to be well described by a universal density profile.  Scale-free models provide a clean probe of the possible dependence of the structure of DM halos on the primordial initial conditions.  This provides insight  to the concordance $\\Lambda$CDM model in the sense that the power spectrum for this cosmology exhibits a varying slope depending on the scale.  We shall further investigate the behaviour of the phase-space density profile with different components of the total velocity dispersion $\\sigma_\\mathrm{tot}$, namely the radial and tangential dispersions.  Given the close connection between the $Q(r)$ power law and the velocity anisotropy \\citep{Zait2007}, expressed by the $\\beta$ parameter, the $\\beta$ profile of the DM haloes is to be explicitly studied. The paper is organized as follows. The numerical experiments are briefly discussed in \\S 2. The fitting of the phase-space density profile is described  in \\S 3 and the results are presented in \\S 4. The paper concludes with a summary and discussion in \\S 5. ",
        "conclusions": "We have analysed five high-resolution simulations of scale-free cosmologies and investigated the logarithmic slope $\\alpha$ of the radial, angular and total phase-space density profiles $Q(r)=\\rho(r)/\\sigma^3(r)$ when fitted by a single power-law $Q(r) \\propto r^\\alpha$.  Additionally, we used a halo sample drawn from a $\\Lambda$CDM simulation as a control sample to verify our analysis procedure and reproduced previous results.  Focussing on a well-defined sample of haloes in each scale-free simulation, chosen to be in a comparable dynamical state, we derived the slopes $\\alpha_{\\mathrm{rad}}$, $\\alpha_{\\mathrm{tan}}$ and $\\alpha_\\mathrm{tot}$ for each halo and computed the median slopes in each simulation.  We found that the slopes for the three definitions of the phase-space density profile $Q$ vary between the models, namely $Q$ becomes shallower with steeper spectral index $n$ of the initial power spectrum $P(k) \\propto k^n$, with the exception of the 512-2.75 model, which shows flatter $Q$ profiles than found in the 512-2.50 model.  We also find a trend for $\\alpha$ to first become flatter with decreasing concentration; this trend flattens at $c\\approx 10$ and even turns around for the tangential and total $Q$ profiles.  This is connected to previous findings showing a correlation between the power spectrum index $n$ and the concentration $c$.  But how can the $n$ - $c$ relation affect the $Q(r)$ profiles?  The scaled density profiles appear to follow a universal form, independent of the cosmological model~\\citep[see e.g.][]{Navarro1996, Navarro1997, Navarro2004}; in particular this has been shown explicitly for the models studied here \\citep{Knollmann2008}.  The dependence on the cosmological model (i.e. the slope of the power spectrum of initial density perturbations) is introduced via the concentration parameter, and thereby by the dependence of the $\\beta$ profile on the model.  The Jeans equation, which dictates the structure of spherical DM haloes in equilibrium, relates the density, phase space density and the velocity anisotropy profiles \\citep[e.g.][and references therein]{Zait2007}. Zait et al. showed that for a given density profile a more radial velocity dispersion at the outer parts of a halo implies a more negative $\\alpha_\\mathrm{rad}$, which is precisely what we find and have shown in figure~\\ref{fig:beta_all}.  It is hence interesting to note that, even though the density profile follows a universal form, the $\\beta$-profile does not, but rather depends on the concentration \\citep{Knollmann2008}.  Summarizing past and present results the  following conclusions can be drawn about DM haloes in scale-free cosmologies. The density profile shows a universal profile, upon an appropriate scaling, but the concentration parameter depends on the power spectrum index $n$ \\citep{Knollmann2008}.  Here it has been found the  phase-space density profile obeys a power law.  The slope of its power law depends on whether the phase-space is defined in terms of the radial, tangential or the full velocity dispersion.  Furthermore, the slopes depends on the shape of the power spectrum.  Also, the velocity anisotropy at $R_\\mathrm{vir}$ depends strongly on the power spectrum index.  The main problem addressed here is the universality of the phase-space density profile, and a mixed answer has been found.  On the one hand it has been very robustly shown that indeed a power law profile provides a good fit in all models considered here.  Yet, the power law slope $\\alpha$ has been found to vary with power spectrum index $n$.  Given our lack of understanding of the origin of the phase-space density power law it is interesting to trace the reason for the dependence of the slopes $\\alpha$ on $n$.  It has been shown here that at least part of that dependence is attributed to the concentration parameter variation with $n$.  However, it is unclear whether there is a direct dependence of the slope $\\alpha$ on $n$.  This can be tested by comparing the slope and $n$ relation for a sub-sample of haloes of different models with the same value of the concentration parameter. Regrettably, the sample of haloes used in the present work is too small to provide a clear answer."
    },
    "0809/0809.2783_arXiv.txt": {
        "abstract": "Ultra-compact dwarf galaxies (UCDs) are stellar systems with masses of around $10^7$ to $10^8$ M$_\\odot$ and half mass radii of 10-100 pc. They have some properties in common with massive globular clusters, however dynamical mass estimates have shown that UCDs have mass-to-light ratios which are on average about twice as large than those of globular clusters at comparable metallicity, and tend to be larger than what one would expect for old stellar systems composed out of stars with standard mass functions.  One possible explanation for elevated high mass-to-light ratios in UCDs is the existence of a substantial amount of dark matter, which could have ended up in UCDs if they are the remnant nuclei of tidally stripped dwarf galaxies, and dark matter was dragged into these nuclei prior to tidal stripping through e.g. adiabatic gas infall.  Tidal stripping of dwarf galaxies has also been suggested as the origin of several massive globular clusters like Omega Cen, in which case one should expect that globular clusters also form with substantial amounts of dark matter in them.  In this paper, we present collisional N-body simulations which study the co-evolution of a system composed out of stars and dark matter. We find that the dark matter gets removed from the central regions of such systems due to dynamical friction and mass segregation of stars. The friction timescale is significantly shorter than a Hubble time for typical globular clusters, while most UCDs have friction times much longer than a Hubble time. Therefore, a significant dark matter fraction remains within the half-mass radius of present-day UCDs, making dark matter a viable explanation for the elevated M/L ratios of UCDs. If at least some globular clusters formed in a way similar to UCDs, we predict a substantial amount of dark matter in their outer parts. ",
        "introduction": "\\label{sec:intro}  Ultra-compact dwarf galaxies (UCDs) were discovered in the late 1990s in spectroscopic surveys of the Fornax galaxy cluster \\citep{hetal99, detal00} and have since then been found in other nearby galaxy clusters as well \\citep{hetal05,metal05,metal07,jetal06,fetal07,retal07}. They are bright ($-11 < M_V < -13.5$) and compact ($7 < r_h < 100$ pc) stellar systems which have ages of at least several Gyr and possibly up to 10 Gyr \\citep{metal06,eetal07}.  The masses and sizes of UCDs are larger than those of Galactic globular clusters, but similar to those of nuclei in dwarf elliptical galaxies (Drinkwater et al. 2003, Bekki et al. 2003).  One of the most remarkable properties of UCDs is that their dynamical mass-to-light ratios are on average about twice as large than those of globular clusters of comparable metallicity, and also tend to be larger than what one would expect based on simple stellar evolution models that assume a standard stellar initial mass function, like e.g. \\citet{k01} \\citep{hetal05,dhk08,metal08}. If not due to a failure of stellar evolution models, this points either to unusual stellar mass functions \\citep{mk08, dbk08} or possibly to the presence of a significant amount of dark matter in UCDs. We note that most methods used to determine M/L ratios for UCDs rely on the assumptions that mass follows light and isotropic velocity dispersions. How well these assumptions are fulfilled is currently not known. Also due to the large distances of UCDs, only integrated velocity dispersions can be obtained, which in most cases are intermediate between the central and the global velocity dispersions. In order to determine the mass-to-light ratio, the mass modeling has to take the density profiles of the UCDs as well as the effects of seeing and a finite slit size into account, as done for example in \\citet{hetal07}.  Several formation scenarios have been discussed for UCDs, like e.g.\\ that UCDs are simply massive globular clusters and form in the same way \\citep{hetal99,eetal07,fetal08}, that they are the nuclei of tidally stripped, originally much more extended galaxies \\citep{bcd01,bcd03,tde08,getal08}, or that they are merged globular clusters \\citep{ol00,fk02}. \\citet{getal08} have shown that funneling of dark matter to the central region of a disk galaxy, due to gas-infall, can significantly increase the M/L ratios in the nuclear region, and hence may explain the elevated M/L ratios of UCDs, provided that UCDs formed by tidal stripping. Indeed, it has been suggested that also GCs may have originated as centers of individual primordial dark matter halos (e.g. Carraro \\& Lia 2000, Lee et al. 2007, Bekki et al. 2007). If dark matter funneling is an efficient mechanism \\citep{getal08}, one may therefore expect both UCDs and GCs to be formed with a significant fraction of dark matter. It is important to note that such an increase of dark matter density by some kind of funneling mechanism is necessary to explain a significant amount of dark matter in UCDs or GCs, since their present-day stellar (and hence implied dark matter) densities are up to 2-3 orders of magnitude higher than expected for cuspy dark matter halos of dwarf galaxy mass \\citep{giletal07}. This is shown in Fig.~\\ref{fig:density}. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{UCD_trelax_density.eps} \\end{center} \\caption{The mean mass density within the half-mass radius of the joint sample of GCs and UCDs from Fig.~\\ref{fig:mtol2} is plotted vs. their relaxation time \\citep{metal08}. The dotted line indicates the approximate central ($r\\lesssim10$pc) dark matter densities expected for cuspy dwarf galaxy CDM halos \\citep{giletal07}. } \\label{fig:density} \\end{figure}  In this paper, we start from the working hypothesis that both GCs and UCDs are formed with the same non-zero dark-to-stellar-mass-fraction. We then investigate how the dynamical co-evolution of dark matter and stars changes the observed dark matter fraction as a function of time. We assess whether the observed rise of M/L ratios from the regime of GCs to that of UCDs can be explained by our working hypothesis and the subsequent dynamical evolution. ",
        "conclusions": "\\label{sec:concl}  We have performed collisional $N$-body simulations of the evolution of compact systems composed out of a mix of stars and dark matter particles. Our simulations show that dark matter is depleted from the centers of these systems due to dynamical friction and energy equipartition between stars and dark matter particles. The inspiral time of stars is short enough that only 20\\% of the original dark matter would remain within the half-mass radius in typical globular clusters. If mainly stars from the inner cluster parts are used to determine mass-to-light ratios, the resulting increase in the mass-to-light ratio is within the errors with which mass-to-light ratios are typically determined for globular clusters and would therefore be difficult to detect.  If not tidally stripped, dark matter should also reside in the outer parts of globular clusters. For a number of globular clusters, \\citet{setal07} have indeed reported a flattening of the velocity disperion in the outer cluster parts. This could however be due to a number of reasons like contamination of the sample by background stars or the tidal interaction of a star cluster with the gravitational field of the Milky Way \\citep{detal98, cetal05}. Detailed simulations would be necessary to exclude these possibilities and confirm that the observed flattening is due to a dark matter halo.  UCDs on the other hand have inspiral times significantly longer than a Hubble time and therefore still contain most of the dark matter in their centers. Dark matter therefore seems a viable explanation for the elevated M/L ratios of UCDs, provided that UCDs originate from the centers of dark matter halos and have seen their dark matter content being increased by dark matter funneling, through e.g. adiabatic gas infall \\citep{getal08}.  A prediction of our simulations, which can in principle be tested by observations, is that globular clusters which have expelled the dark matter from their centers should also be mass segregated. Non-mass segregated clusters with velocity dispersions and mass-to-light ratios in agreement with simple stellar population models, would therefore have formed without significant amounts of dark matter in their centers.  Also, if dark matter existed in a globular cluster at the time of its formation, it should still reside in its outer parts, especially if tidal stripping due to external tidal forces from the host galaxy \\citep{ms05} and relaxation driven internal mass loss was not important for the cluster evolution.  In this case, the measured mass-to-light ratio should increase towards the outer cluster parts, which can in principle be detected with dedicated radial velocity or proper motion surveys. The future astrometric satellite {\\it GAIA} would be an excellent tool for such a search since it will provide accurate proper motions for thousands of stars in the halos of nearby globular clusters."
    },
    "0809/0809.0265_arXiv.txt": {
        "abstract": "This paper is the 4th in a series describing the latest additions to the BaSTI stellar evolution database, which consists of a large set of homogeneous models and tools for population synthesis studies, covering ages between 30~Myr and $\\sim$20~Gyr and 11 values of $Z$ (total metallicity). Here we present a new set of low and high resolution synthetic spectra based on the BaSTI stellar models, covering a large range of simple stellar populations (SSPs) for both scaled solar and $\\alpha$-enhanced metal mixtures.  This enables a completely consistent study of the photometric and spectroscopic properties of both resolved and unresolved stellar populations, and allows us to make detailed tests on specific factors which can affect their integrated properties. Our low resolution spectra are suitable for deriving broadband magnitudes and colors in any photometric system.  These spectra cover the full wavelength range (9--160000nm) and include all evolutionary stages up to the end of AGB evolution. Our high resolution spectra are suitable for studying the behaviour of line indices and we have tested them against a large sample of Galactic globular clusters.  We find that the range of ages, iron abundances [Fe/H], and degree of $\\alpha$-enhancement predicted by the models matches the observed values very well. We have also tested the global consistency of the BaSTI models by making detailed comparisons between ages and metallicities derived from isochrone fitting to observed colour-magnitude diagrams, and from line index strengths, for the Galactic globular cluster 47~Tuc and the open cluster M67. For 47~Tuc we find reasonable agreement between the 2 methods, within the estimated errors.  From the comparison with M67 we find non-negligible effects on derived line indices caused by statistical fluctuations, which are a result of the specific method used to populate an isochrone and assign appropriate spectra to individual stars. ",
        "introduction": "\\label{sec:intro}  Large grids of stellar models and isochrones are necessary tools to interpret photometric and spectroscopic observations of resolved and unresolved stellar populations. This, in turn, allows one to both test the theory of stellar evolution, and investigate the formation and evolution of galaxies, one of the major open problems of modern astrophysics.  To this purpose, we started in 2004 a major project aimed at creating a large and homogeneous database of stellar evolution models and isochrones (BaSTI -- a Bag of Stellar Tracks and Isochrones) that cover a large chemical composition range relevant to stellar populations in galaxies of various morphological types. In addition, BaSTI includes the option of choosing between different treatments of core convection and stellar mass loss employed in the model calculations. In its present form, BaSTI contains a grid of models that cover stellar populations with ages ranging between 30~Myr and $\\sim$20~Gyr, which include all evolutionary phases from the Main Sequence (MS) to the end of the Asymptotic Giant Branch (AGB) evolution or carbon ignition, depending on the value of stellar mass.  They are calculated for two metal mixtures (scaled solar and $\\alpha$-enhanced) and 11 values of $Z$ (total metallicity), each with two choices for the Reimers mass loss parameter $\\eta$ \\citep{reimers}, and two mixing prescriptions during the Main Sequence (without and with overshooting from the Schwarzschild boundary). Broadband magnitudes and colors for the BaSTI models and isochrones are provided in various photometric systems (Johnson-Cousins, Sloan, Str\\\"omgren, Walraven, ACS/$HST$), making use of bolometric corrections derived from an updated set of model atmosphere calculations, for element mixtures consistent with those employed in the stellar evolution computations\\footnote{All the results presented in this work are based on the 2008 release of the BaSTI archive. We refer to the official BaSTI website for more details on this new release.}.  In addition to stellar models and isochrones, BaSTI also provides a series of Web tools that enable an interactive access to the database and make it possible to compute user-specified evolutionary tracks, isochrones, stellar luminosity functions (star counts in a stellar population as a function of magnitude) plus synthetic color-magnitude diagrams (CMDs) for arbitrary star formation histories (SFHs). The results of this major effort have been published in \\citet[][ hereinafter Papers~I and II]{basti1,basti2} and \\citet[][hereinafter Paper~III]{basti3}. All Web tools, models and isochrones can be found at the BaSTI official website \\url{http://193.204.1.62/index.html}.  The database has been extensively tested against observations of local stellar populations and eclipsing binary systems (see, for some examples, the tests presented in Papers~I, ~II and ~III; \\citealt{ms07,tomasella}). In its present form BaSTI can be used to investigate resolved stellar populations, and indeed it has been widely employed in studies of Galactic and extragalactic resolved star clusters \\citep[see, e.g.][for just a few examples]{deangeli,villanova,mackey} and in the determination of SFH and chemical enrichment history of resolved Local Group galaxies \\citep[see, e.g.][]{gallart,barker,gulli,carrera}.  Integrated colors and magnitudes for simple (single-age, single metallicity) stellar populations (SSPs) can be easily determined from the isochrones. The only additional pieces of information needed are an Initial Mass Function (IMF) and the relationship between the initial and final (remnant) stellar masses.  Indeed, some integrated colors from BaSTI models have already been tested against empirical constraints in Papers~II and III. Moreover, \\citet{paj} and \\citet{salcas} have already employed integrated colors from BaSTI isochrones to address issues related to the SFH of unresolved dwarf and elliptical galaxies, and extragalactic globular cluster ages.  A further development of BaSTI, which we present in this paper, is the inclusion of integrated spectra (including spectral line indices) for SSPs spanning the same parameter space -- in terms of ages, metallicities and heavy element mixtures -- covered by the isochrones of Papers~I and II. Corresponding quantities for composite stellar populations with an arbitrary SFH can then be easily determined by integration over a range of discrete SSPs.  We present both low and high resolution versions suited to different purposes.  Our low resolution spectra cover all evolutionary stages and are based on libraries of (predominantly) theoretical stellar spectra, consistent with those used to produce the bolometric corrections employed in the isochrone calculations. This ensures that our SSP integrated broadband magnitudes and colors obtained from either adding up the broadband fluxes of the individual stars populating the appropriate isochrone, or from their composite integrated spectrum, will be exactly the same (see test in section~\\ref{sec:lrs}). Our high resolution spectra are suitable for deriving line strengths and making detailed tests of their behaviour as a function of age and chemical composition, and also horizontal branch morphology and sampling in the case of resolved stellar clusters (see tests in sections~\\ref{sec:47tuc} and \\ref{sec:m67}).  With the inclusion of predictions of both photometric and spectroscopic integrated properties of SSPs, BaSTI will provide a well tested, fully homogeneous and up-to-date set of theoretical tools that enable us to investigate self-consistently both resolved and unresolved stellar populations. We present a global test of this self-consistency and reliability of the database, by considering both the observed CMD and integrated spectrum of the well-studied clusters 47~Tuc and M67, to intercompare ages and metallicities obtained from our high-resolution spectra with the results from fitting theoretical isochrones to the observed CMD, and with independent chemical composition estimates from spectroscopy of individual stars.  The structure of the paper is as follows. A description of the method used to create our theoretical integrated spectra is given in Section~\\ref{sec:intspec}, including separate detailed descriptions of our low and high resolution spectra in Sections~\\ref{sec:lrs} and \\ref{sec:hrs} respectively. This is followed in Section~\\ref{sec:comparison} by a comparison of our derived line indices with the Galactic globular cluster sample of \\citet{schi05}.  A global consistency test of our isochrones, integrated colors and spectra using photometric and spectroscopic data for 47~Tuc and M67 is presented in Sections~\\ref{sec:47tuc} and \\ref{sec:m67}. A brief summary in Section~\\ref{sec:sum} closes the paper. ",
        "conclusions": "\\label{sec:sum}  In this 4th paper in the BaSTI series we have presented a new set of low and high resolution synthetic spectra based on the BaSTI stellar models, covering a large range of SSPs, for both scaled solar and $\\alpha$ enhanced metal mixtures.  This enables a completely consistent study of the photometric and spectroscopic properties of both resolved and unresolved stellar populations, and allows us to make detailed tests on specific factors which can affect their integrated properties.  \\begin{itemize} \\item We have produced low resolution spectra suitable for deriving broadband magnitudes and colors in any photometric system and ensured that the derived colors are fully consistent with those predicted from the BaSTI isochrones.  \\item We have produced high resolution spectra suitable for studying the behaviour of line indices and tested them against a large sample of Galactic globular clusters, making particular note of the effect of horizontal branch morphology.  \\item We have tested the global consistency of the BaSTI models by making detailed comparisons between ages and metallicities derived from isochrone fitting and from line index strengths for the Galactic globular cluster 47~Tuc and the open cluster M67.  \\item We find non-negligible effects on derived line indices caused by statistical fluctuations, which are a result of the specific method used to populate an isochrone and assign appropriate spectra to individual stars. \\end{itemize}"
    },
    "0809/0809.0862_arXiv.txt": {
        "abstract": "{In spite of all the advances in multi-dimensional hydrodynamics, investigations of stellar evolution and stellar pulsations still depend on one-dimensional computations. This paper devises an alternative to the mixing-length theory or turbulence models usually adopted in modelling convective transport in such studies.} {The present work attempts to develop a time-dependent description of convection, which reflects the essential physics of convection and that is only moderately dependent on numerical parameters and far less time consuming than existing multi-dimensional hydrodynamics computations.} {Assuming that the most extensive convective patterns generate the majority of convective transport, the convective velocity field is described using two parallel, radial columns to represent up- and downstream flows. Horizontal exchange, in the form of fluid flow and radiation, over their connecting interface couples the two columns and allows a simple circulating motion. The main parameters of this convective description have straightforward geometrical meanings, namely the diameter of the columns (corresponding to the size of the convective cells) and the ratio of the cross-section between up- and downdrafts. For this geometrical setup, the time-dependent solution of the equations of radiation hydrodynamics is computed from an implicit scheme that has the advantage of being unaffected by the Courant-Friedrichs-Lewy time-step limit. This implementation is part of the TAPIR-Code (short for {\\bf T}he {\\bf a}da{\\bf p}tive, {\\bf i}mplicit {\\bf R}HD-Code).} {To demonstrate the approach, results for convection zones in Cepheids are presented. The convective energy transport and convective velocities agree with expectations for Cepheids and the scheme reproduces both the kinetic energy flux and convective overshoot. A study of the parameter influence shows that the type of solution derived for these stars is in fact fairly robust with respect to the constitutive numerical parameters. } {} ",
        "introduction": "Convection is one of the persistent problems in stellar astrophysics. Almost all stars contain regions where convective transport is important; in the photosphere, in the envelope, or in the interior where nuclear burning occurs. A description of convection is therefore an essential ingredient to all types of investigations of stellar structure and evolution. Unfortunately, due to its nonlinear, nonlocal, and multi-length-scale nature, modelling convection turns out to be an intricate problem.  Conceptually, there are several different approaches to the numerical simulation of convective transport in stars. The most straightforward approach is the time-dependent solution of the equations of radiation hydrodynamics in a 3D (or at least 2D) domain to compute the convective flow patterns directly. However, this process is very expensive in computing time, prohibitively so for some applications. This particularly applies to problems with a large difference in relevant timescales, for instance between the thermal and acoustic timescale in the stellar interior, or the hydrodynamics and radiative timescale in the outer layers of luminous stars. An additional limit is imposed by the restricted spatial resolution. Even the most elaborate, high-resolution simulations of stellar convection are only capable of resolving the largest scales in the convective velocity field. The effect of turbulence on smaller length scales is effectively ignored, even though it is a possible interpretation to attribute the intrinsic numerical dissipation of the scheme to some (unknown) unresolved turbulence. In particular, hydrodynamics codes often include artificial viscosity for numerical stability, and `unresolved turbulence' is the only physical mechanism that could be used to justify the inclusion of this additional dissipation.  Despite the poor description of turbulence in hydrodynamics computations, multi-dimensional simulations of solar granulation \\citep[e.g.][]{SN98,Asplund2000,Wedemeyer2004} achieve remarkable quantitative agreement with observations.   A completely different and more subtle approach than trying to resolve turbulence on large numerical grids are convection models that use an equation, or a set of equations, to describe convective transport either with a heuristic parametrization or based on turbulence theory.  The most widely used of these 1D-descriptions is the well known mixing-length theory (MLT) \\citep{MLT,CoxandGiuli}, which originates in ideas of \\citet{Prandtl1925}. The MLT has been remarkably successful in stellar astrophysics in application to stars from white dwarfs to super giants, probably because of its simple yet flexible parametrization and its robust reference to the adiabatic temperature gradient.   Modern alternatives to the MLT are convection models in which the original single-eddy assumption of the MLT has been replaced with a full spectrum of turbulence \\citep[e.g.][]{CM91,C96,CGM96} by either assuming or computing a turbulent energy spectrum. For some types of stars, these models can also avoid the mixing-length scale as a free parameter. Convection models of this type are included in many stellar evolution codes as an alternative to MLT.  However, all of these models are, in a similar way to the original MLT, local theories that do not provide any information about overshoot. In stellar-evolution codes, this deficit is overcome by adding overshoot by means of a separate (typically diffusive) parametrization.   Another type of one-equation models consists of a time-dependent equation for the turbulent kinetic energy \\citep[e.g.][]{stell,kuhfuss,GM92} using heuristic approximations for individual terms. Containing a diffusion term for the turbulent kinetic energy, they also allow a simple form of non-locality. These convection models are mainly geared towards computations of stellar pulsations \\citep[e.g.][]{BS1994, FeuchtingerptII, KBSC2002} where one is interested in the time-dependence of the convective transport.   Finally, the most complete way to model turbulence is the `Reynolds stress approach' by solving a set of moment equations \\citep{canuto1992, canuto1993, canuto1997, Xiong1989, Xiongetal1997} that are terminated by a closure model at the third or fourth order. These closures are based on either the quasi-normal approximation for the fourth-order moments or use a parametrization with reference to measured data, hydrodynamics (large-eddy) simulations, or concepts such as the `plume model' (see below). Turbulence models of this type are able to describe convective transport in a time-dependent and non-local way. For applications of the Reynolds stress model to stellar surface convection zones, see \\citet{Kupka99,KuMo2002,MoKu2004}.   The two-column scheme, presented in this paper, is in-between the above categories, combining a hydrodynamics simulation of the convective fluid flow with a parametrized, predefined geometry of the flow patterns. This setup is almost as simple as a 1D description; it describes up- and downstream with two parallel radial columns, and fluid flow over an interface between those two columns allows a basic circulating convective motion. The two-column model could therefore be regarded as a simplistic 2D hydrodynamics scheme, limiting the horizontal range of the grid to just two cells. The very coarse description of the convective velocity field effectively implies that the most extensive convective flow patterns generate the majority of the convective transport. However, this assumption does not differ significantly from what is assumed in multi-dimensional hydrodynamics computations that are also unable to resolve the full spectrum of convective turbulence. Since multi-dimensional hydrodynamics achieve, in spite of this limitation, a good agreement with observations, the scenario of macroscopic convective patterns with distinct up- and downstream regions and little sub-structure (as also observed in the solar granulation) appears to be a sufficient description of the actual physics \\citep{nordlundetal1997}. The two-column approach may therefore also give reasonable results.   In the two-column scheme, the convective flux is computed directly from hydrodynamics and not from a heuristic model. Although it is not without numerical parameters, there is no `mixing-length parameter', nor anything equivalent. The method is also intrinsically non-local and the thickness of convective regions and the amount of overshoot are obtained consistently.   The basic idea of modelling convection using separate radial stratifications for up- and downstream regions is in fact an old one. In the 1960's to 1970's, predating advances in computing power that enabled 2D and 3D hydrodynamics computations to become possible, several similar two- or multi-stream models were devised. From observations, the solar granulation pattern appeared to be separable to almost distinct hot and cool areas; it was therefore a logical first step to place two stratifications next to each other in order to construct more realistic models of the solar photosphere \\citep{MG1969, nordlund76}. More recently a two-stream model was applied by \\cite{lesaffre2005} to investigate the convective Urca process in supernova-progenitor white dwarfs. In geophysics, a similar concept, known as `plume model', was introduced by \\citet{MTT1956}. In its basic form, this model considers only plumes that are immersed in a static, surrounding medium, although there are also models that consider both up- and downdrafts and their interaction \\citep[e.g.][]{Telford1970, WA1986, CB1987, RSM1992}. The idea of separated up- and downwards streams was also used to construct closures for turbulence models \\citep{AMF1997, ZGLM1999, LR2001, GH2002, GHRS2005, CCH2007}   However, all existing multi-stream models differ from the present attempt in that the `two-column-scheme' introduced below is based on fully implicit time-dependent radiation hydrodynamics in \\emph{both} radial and horizontal directions without any ad-hoc assumptions or parametrizations for the physical coupling of the two columns. The term `two-column' (in contrast to `two-stream') was chosen intentionally because the discretization scheme resembles that of two 1D discretizations placed beside each other in two parallel columns. In analogy to `2D' for `two-dimensional', we will occasionally refer to `two-column' as `2C' in the following sections.  The remaining paper is structured as follows. The next section, Sect.~\\ref{sect.2C-scheme}, introduces the two-column discretization scheme and its geometric derivation.  % The equations of radiation hydrodynamics are given in Sect.~\\ref{sect.RHDeq-anal} in analytical form,                        % while Sect.~\\ref{sect.RHDeq-discr} describes their discretization and considers details                                      % such as artificial viscosity and radiative transport. Section~\\ref{sect.method} presents the deployed solution algorithm,                                                          % followed by a demonstrating example and some parameter studies in Sect.~\\ref{sect.results}.                                  % Finally, Sect.~\\ref{sect.concl} draws the conclusions and summarizes the paper.                                              % A verification of the method by comparison with detailed 2D hydrodynamics computations as well as applications in time-dependent calculations of Cepheids' pulsations will be given in the forthcoming part II paper of this series. ",
        "conclusions": "\\label{sect.concl}  The scheme proposed in this paper has a number of advantages: \\begin{itemize} \\item It is a non-local description of convection and therefore provides a consistent computation of the depth of the convective region including convective overshoot. \\item The convective flux is computed directly from hydrodynamics and not from a heuristic, parametrized model. \\item The two-column convection has stationary solutions and in principle allows arbitrarily large time steps. This implies that it is suitable for application to problems involving long time series, such as stellar pulsations or stellar evolution. \\item The 2C-model is much faster than multi-dimensional hydrodynamics computations; stationary solutions can be obtained within minutes. \\item Radiative transport is an intrinsic part of the scheme, i.e.\\ no hydrodynamical model with plugged-in radiation effects. \\item The main parameters of the scheme have a straightforward geometrical meaning that also provides indications of reasonable values for these parameters. \\end{itemize} However, there are also shortcomings to be considered: \\begin{itemize} \\item The 2C-scheme is basically still a parameter-dependent model. These parameters require proper adjustment. \\item Horizontal advection and radiative transport are poorly represented because of the very coarse `two-cell' spatial resolution in the horizontal direction. \\item The 2C-model uses a simplistic description of the full spectrum of vertical and horizontal convective motion, which ignores turbulence effects and limits the investigation of more subtle features of convection. \\end{itemize}  In its present form, the two-column scheme provides a simple, yet physically sound and consistent, non-local, radiation-hydrodynamics description of the convective circulation. The model still contains free parameters, but their geometrical interpretation provides at least reasonable indications of proper values, and  they do not change the results by magnitudes. For applications in which more detail and higher certainty is required, more elaborate methods, such as multi-dimensional hydrodynamics or turbulence models, remain the most appropriate alternative."
    },
    "0809/0809.0053_arXiv.txt": {
        "abstract": "The Australia Telescope Compact Array (ATCA) has been used at 1.38 and 2.38~GHz to survey  seven southern Abell clusters of galaxies with high X-ray luminosities: A2746, A2837, A3126, A3216, A3230, A3827 and A3836. The clusters have also been surveyed at 0.843~GHz with the Molonglo Observatory Synthesis Telescope (MOST). We have listed a complete 1.38-GHz sample of 149 radio sources within the Abell circles centred on their X-ray centroids. We compare their identification fractions, emitted 1.38-GHz and optical powers, radio spectral indices and radial variation in projected source density with those of the radio-selected samples of Slee et al.\\ (1998). We compare our fractional radio luminosity function with that of the radio-selected samples of Ledlow \\& Owen (1996) and Slee et al.\\ (1998). Three significant differences are noted between X-ray and radio-selected samples of clusters; (1) the X-ray sample has an excess of flat-spectrum radio sources; (2) the fractional radio luminosity function for the FR\\,I sources in the X-ray selected sample is much steeper, implying that  fewer of their cluster galaxies  become hosts for the stronger FR\\,I radio galaxies; (3) a complete absence of FR\\,{\\sc II} radio galaxies in the X-ray selected sample. The average excess projected density of radio sources near our cluster centres is $\\sim$5 times the background source density. ",
        "introduction": "\\label{intro}  Clusters of galaxies provide an interesting environment in which to test theories of how radio galaxies form and evolve. In particular, the presence of a relatively dense, X-ray emitting ionised gas permeating clusters provides an ideal laboratory to study jet propagation and the expansion of radio lobes.  The examination of these processes requires carefully-defined  samples incorporating data on the individual cluster members from the radio, optical and X-ray regimes. Further, selection of the cluster samples may have an impact upon the conclusions that can be drawn. Clusters selected primarily on the radio power of the objects within them may represent a distinct environment, with a radio source population different from that in clusters selected on the basis of integrated optical luminosities or X-ray luminosity.  A number of synthesis radio surveys of large samples of clusters have been published in the last 17 years. In addition, numerous publications on individual clusters and individual cluster radio galaxies exist in the literature. Comprehensive collections of 1.4-GHz maps of cluster radio galaxies taken with the Very Large Array (VLA) of the National Radio Astronomy Observatory (NRAO, USA) were published by Zhao, Burns \\& Owen (1989), Owen et al.\\ (1992, 1993), Ledlow \\& Owen (1995a), and a detailed analysis of these data were made by Ledlow \\& Owen (1995b, 1996). A second set of VLA maps at 1.5 and 4.9~GHz with complete source lists were published by Slee, Perley \\& Siegman (1989), Slee, Roy \\& Savage (1994) and Slee, Roy \\& Andernach (1996), who later drew general conclusions about the relationships  between radio, optical and X-ray parameters (Slee, Roy \\& Andernach 1998). A comprehensive survey at 0.843~GHz of 39 nearby southern clusters was made by Unewisse (1993), but none of the more distant clusters in the present paper was included.  The present paper presents complete source lists and analysis for seven rich clusters from Abell et al's (1989) catalogue south of Declination $-50^{\\circ}$, selected on account of their high ROSAT X-ray luminosities in the ``X-ray brightest Abell-type clusters of galaxies'' (XBACs) catalogue of Ebeling et al.\\ (1996). The radio survey was made at 1.38 and 2.38~GHz with the Australia Telescope Compact Array (ATCA) with the initial object of identifying a suitable candidate cluster in a search for radio gravitational lensing (Tsarevsky et al.\\ 2008, in preparation). In addition, as an X-ray-selected sample, the source list and its general analysis are of interest in their own right.  Our selection criteria included the following:  \\noindent 1.~Clusters south of Declination $-50^{\\circ}$, enabling observations in the ``cut'' mode (cf.\\ Sect.\\ 2.1) with good coverage of the uv plane.  \\noindent 2.~Cluster X-ray luminosities (0.1--2.4~keV) above the median of 1.56$\\times10^{37}$\\,W of the XBACs sample; the exception is A3126, whose $L_X$ falls only 13\\% below the median, and was included because it has a well-determined redshift and a high velocity dispersion.  \\noindent 3.~Contamination of the ROSAT image by active galactic nuclei (AGN) is minimal as assessed by Ebeling et al.\\ (1996).  The clusters were also surveyed with the Molonglo Observatory Synthesis Telescope (MOST) at 0.843~GHz with the primary object of detecting as many as possible of the 1.38 and 2.38-GHz sources in order to obtain more accurate radio spectral indices and, if possible, identify steep-spectrum radio relics, which are known to favour the cluster environment.  Section~2 briefly describes the observations, while Section~3 presents the radio source lists and optical identifications. Section~4 discusses the derived parameters of the sources, including spectral indices, relationships between emitted radio, optical and X-ray luminosities and the clusters' radio luminosity function, and Section~5 presents a discussion of our results.  In the following text, frequent reference is made to cluster samples selected by X-ray emission, radio emission or using neither of these criteria. To avoid confusion, we call samples selected by their X-ray emission ``X-ray samples'' and those selected by their radio emission are termed ``radio samples''; naturally, both samples imply that optical clusters are present.  For consistency with previously published radio and X-ray data, to which we refer in this paper, we use the Einstein-de Sitter cosmological model with H$_0$=75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and q$_0$=0.5 throughout this paper. ",
        "conclusions": "\\label{discussion}   %  Our results do not encompass a sufficiently large range of redshift, nor are there a sufficient number of clusters, to search for evolution in strong X-ray emitting clusters, but we can draw a few conclusions about the effects of the cluster environment on the evolution of galaxies and their associated radio sources.  First, 25\\% of the identified cluster radio sources possess spectral indices more positive than $\\alpha=-$0.50, which are distributed over the full range of angular distances out to the Abell radius. Such flat or even rising spectra are also associated with the active nuclei of isolated E/S0 galaxies (cf.\\ Slee et al.\\ 1994); therefore it seems likely that a significant proportion of radio sources in strong X-ray clusters may be undergoing an active phase, perhaps because their black holes are fuelled by the copious intracluster gas, although there is no clear tendency for them to favour the inner regions of these clusters where the gas density is highest.  Both Slee et al.\\ (1998) and Ledlow \\& Owen (1996) found a positive relationship between radio and optical powers, and the effect is present, but with lower significance, in our small X-ray sample. The radio and optical powers are not strongly correlated for individual galaxies; cuts through Figures~7a and 7b at constant $R$-band power shows a large variation of at least an order of magnitude in radio power. The interpretation of these data is very model-dependent. One can reasonably assume that the ages of FR\\,I radio sources are at least an order of magnitude lower than that of their host galaxies, but the effect of the host galaxy's gas content, and of the intracluster gas density, on the evolution of an FR\\,{\\sc I}'s jets and lobes are largely unknown. If one also assumes that the optical luminosity of the host galaxy is fairly constant over the lifetime of the radio source, and that its radio power evolves with time as the source expands into its local environment, then the density of points along a line of constant optical power represents the amount of time radio sources spend in that state. Unfortunately, we do not have enough data from this small X-ray-selected sample of clusters to construct a bivariate luminosity function, which would show the rate at which radio power evolves at various values of optical luminosity.  The fraction of E/S0 galaxies in the optical and radio samples of Ledlow \\& Owen (1996) and Slee et al.\\ (1998), respectively, that host FR\\,{\\sc I} radio galaxies at a particular epoch is generally 0.013--0.025. Ledlow \\& Owen (1996) show that this fraction seems to be independent of the environment, i.e., whether the E/S0 galaxies are in clusters or in the general field. In addition, the integrated bivariate radio luminosity functions of Ledlow \\& Owen (1996) and Slee et al.\\ (1998) indicate that the detectable fraction increases with the optical luminosity of their host galaxies over the complete range of FR\\,{\\sc I} power. If we add to these facts that Ledlow \\& Owen (1996) find from surface photometry that the hosts of cluster FR\\,I galaxies are optically indistinguishable from radio-quiet cluster galaxies down to their sensitivity limit at 1.40 GHz of 10~mJy, then clearly the optical power output is a sufficient and necessary condition for the production of an FR\\,I radio galaxy. It then follows that fainter  E/S0 galaxies with $M_R>-$20.5 (the limit of the present samples) will be detectable with similar frequency at 1.38-GHz radio powers below 10$^{22}$\\,W\\,Hz$^{-1}$. None of these parameters is inconsistent with the suggestion that all E/S0 galaxies, whether in or out of clusters, will host an FR\\,{\\sc I} radio galaxy at least once during their lifetimes.  A notable difference between our X-ray sample and the radio/optical samples is our failure to detect a single FR\\,{\\sc II} cluster radio galaxy.  FR\\,{\\sc II} sources are found with equal frequency both in clusters and in the general field (Ledlow \\& Owen, 1996) and are characterized by two parameters: (1)~log\\,(P$_{1.4}$/W\\,Hz$^{-1})\\ge$24.0, (2)~edge-brightened lobes, caused by the interaction of the jets with the surrounding gas. The FR\\,{\\sc II} phenomenon is not confined to the brightest galaxies, but is found also in ellipticals down to a luminosity of $M_R=-$21.0, which encompasses practically all the identified cluster galaxies in our X-ray sample and the radio samples of Slee et al.\\ (1998) and Ledlow \\& Owen (1996). If FR\\,{\\sc II} sources had been contained within our X-ray selected sample, we had the necessary angular resolution, especially at 2.38~GHz, to resolve their outer hot spots from the general lobe structure to confirm the evidence provided by the steepening of the RLF in the samples of Ledlow \\& Owen (1996) and Slee et al.\\ (1998). Ledlow \\& Owen (1996) listed 18 FR\\,{\\sc II} radio galaxies in their complete sample, which surveyed a projected sky area of 155\\,Mpc$^2$. Slee et al.\\ (1998), in their sample out to r/R$_A$=0.5, surveyed a projected area of 71\\,Mpc$^2$. Therefore, if the radio sources in these two samples had been drawn from the same populations, we would have expected to detect (71/155)$\\times$18=8.2 FR\\,{\\sc II}'s and we detected nine FR\\,{\\sc II}'s. The present X-ray sample surveyed a projected area of 88\\,Mpc$^2$, so we should have expected to detect (88/155)$\\times$18=10.2 FR\\,{\\sc II}'s, but we found none.  It seems probable that the higher luminosity X-ray clusters are deficient in FR\\,{\\sc II} radio galaxies, although no apparent reason is evident, as there is more than sufficient intracluster gas with which the jets can interact to cause edge-brightening. Perhaps the higher gas density inhibits the formation of long jets. If this is a valid phenomenon among clusters, then the clusters containing FR\\,{\\sc II} sources should possess significantly lower values of $L_X$ than those containing only FR\\,{\\sc I}s. We have searched for this effect in the radio samples of Slee et al.\\ (1998). If we consider the median values of $L_X$ in their radio sample out to 0.50~r/$R_A$, for the seven clusters containing FR\\,{\\sc II} sources and the 21~clusters containing only FR\\,Is, we find (including upper limits of 0.2$\\times10^{37}$\\,W for the 12~clusters that were well surveyed with ROSAT in the 1RXS catalogue, but not detected) that the medians for the FR\\,{\\sc II} and FR\\,{\\sc I} clusters are $<$0.2 and 0.81$\\times10^{37}$\\,W respectively. This result suggests that low values of $L_X$ are associated with FR\\,{\\sc II} sources and intermediate values of $L_X$ with FR\\,{\\sc I}s. This conclusion is reinforced when we compare the above results with the median value of $L_X=2.31\\times10^{37}$\\,W in the present X-ray selected sample, in which no FR\\,{\\sc II}s are found. There is, however, an exception to this trend in Slee et al's radio sample, in which an FR\\,{\\sc II} is present in one cluster with $L_X=6.48\\times10^{37}$\\,W.  Clearly, larger samples of X-ray selected clusters and their associated sources need to be examined before a firm decision can be made.  The lack of FR\\,{\\sc II} radio galaxies in our present sample is consistent with the scenario advanced by Hardcastle et al.\\ (2007) for the powering of jets in low and high-powered radio galaxies. The FR\\,{\\sc I}'s are seen as low-excitation objects, which lack the narrow-line optical emission expected from conventional AGN. In a diagram of K-band vs.\\ 151-MHz luminosity (Fig.~1 of Hardcastle et al.), these galaxies congregate in the area occupied by low to medium values of K-band and radio luminosity. The FR\\,{\\sc II}'s, on the other hand, show the more typical parameters of AGN with high-excitation, narrow-band emission and occupation of the medium-to-high area of the K-band -- radio luminosity diagram.  Hardcastle et al.\\ (2007) present a convincing argument that the jets of FR\\,{\\sc I}'s are powered by accretion on to a black hole of hot ionized gas from its surrounding medium. These so-called hot-mode systems require a plentiful supply of hot gas and a massive central black hole. By contrast, the jets of high-excitation galaxies (FR\\,{\\sc II}'s) are maintained directly by cold-mode accretion, and neither need a rich, hot-gas environment nor to be at the bottom of a deep potential well, allowing galaxy mergers (ellipticals with spirals) to take place.  It therefore seems likely that galaxy clusters with strong X-ray emission, implying the presence of a dense, hot ionized gas, will possess the required features to fuel the jets of FR\\,{\\sc I}'s. If this model is true, then there should be an X-ray luminosity below which the jets of FR\\,{\\sc I}'s are not so readily fueled. This means that the radio-selected samples of galaxy clusters will contain a mixture of FR\\,{\\sc I}'s and FR\\,{\\sc II}'s, as is reported in the present paper.  In a survey based on ROSAT-PSPC data of a large optically selected sample of clusters with ACO richness class R$\\ge$2, David et al.\\ (1999) showed that the dominant correlation was between $L_X$ and ACO galaxy count (i.e., richness); for a given optical richness there is also a correlation between $L_X$ and Bautz-Morgan type, with BM\\,I clusters being the most luminous. No obvious relationships of this kind are seen in our present sample of clusters, which encompasses Bautz-Morgan types from I to III, ACO richness classes from 0 to 2, and with its X-ray luminosities restricted to a small range of 2.65:1. It can be said, however, that the median X-ray luminosity of 2.32$\\times10^{37}$\\,W is comparatively high for a sample that has such diverse values of Bautz-Morgan type and optical richness."
    },
    "0809/0809.1443_arXiv.txt": {
        "abstract": "We have assembled a sample of 64 late-type spiral galaxies ($T$ types 6.0--9.0, corresponding to Hubble types Scd--Sm) with archival {\\em Chandra} data. At a signal-to-noise (S/N) threshold of 3, we find 12 objects with X-ray point-source detections in close proximity with the optical or near-infrared position of the nucleus (median offset $\\delta$ = 1\\farcs 6), suggestive of possible low-luminosity active galactic nuclei (AGNs). Including measurements with 3 $>$ S/N $>$ 1.5, our detections increase to 18. These X-ray sources range in luminosity from $L_{\\rm X}$(2--10 keV) = $10^{37.1}$ to $10^{39.6}$ ergs~s$^{-1}$. Considering possible contamination from low-mass X-ray binaries (LMXBs), we estimate that $\\sim$5 detections are possible LMXBs instead of true AGNs, based on the probability of observing a LMXB in a nuclear star cluster typically found in these late-type spiral galaxies. Given the typical ages of nuclear star clusters, contamination by high-mass X-ray binaries is unlikely.  This AGN fraction is higher than that observed in optical surveys, indicating that active nuclei, and hence central black holes, are more common than previously suggested. The incidence of AGN activity in such late-type spiral galaxies also suggests that nuclear massive black holes can form and grow in galaxies with little or no evidence for bulges.  Follow-up multiwavelength observations will be necessary to confirm the true nature of these sources. ",
        "introduction": "There is now overwhelming evidence that the vast majority of massive early-type galaxies harbor a supermassive black hole (BH) at their centers, and that BHs play an important role during the assembly and evolution of their host galaxies (see reviews in \\citealt{Ho04}). The existence of the $M_{\\rm BH}-\\sigma_\\star$ correlation between BH mass and central stellar velocity dispersion links the growth of the BH with the formation of the host bulge \\citep{Gebhardt00,FM00}, perhaps through feedback mechanisms associated with an active galactic nucleus (AGN). The detection of BHs can be achieved by several methods, including direct kinematics of surrounding gas and stars (e.g., \\citealt{Kormendy04,Ghez05}), and optical spectroscopic identification of nuclear emission lines originating from an AGN (e.g., \\citealt{Ho_2_95}).  The first method is limited to nearby, inactive or mildly active galaxies, however, while optical spectroscopic identification is generally limited to moderately bright nuclei that are not strongly confused by circumnuclear star formation or heavily obscured by dust.  An extensive spectroscopic survey of nearby galaxies has revealed that approximately $\\sim 60$\\% of galaxies of Hubble type E -- Sb contain an AGN \\citep{Ho_5_97}, broadly consistent with the idea that BHs are ubiquitous in these systems \\citep{Ho08}. For types later than Sc, where classical bulges disappear \\citep{KK04}, only $\\sim 15$\\% of galaxies show optical AGN activity. This leads to the notion that classical bulges or ellipticals are necessary for BH formation. There are, however, a few striking counter examples. NGC\\,4395, an essentially bulgeless Sm galaxy, exhibits all the hallmarks of a Seyfert 1 nucleus \\citep{FS89,FHS93}, and it has a BH mass of $\\sim 10^5 M_\\odot$ \\citep{FH03,Peterson05}. POX\\,52 shares an almost identical optical spectrum to NGC\\,4395, but is hosted in a spheroidal galaxy \\citep{Barth04}\\footnote{\\cite{Barth04} called POX 52 a dwarf elliptical galaxy, but here we follow the nomenclature advocated by \\cite{Kormendy08} and adopted in \\cite{GHB08}.}, a visually similar but distinct morphological type compared to elliptical galaxies or classical bulges. More recently, an AGN was discovered in NGC\\,3621, a late-type Sd galaxy, using infrared spectroscopic observations \\citep{Satyapal07}. On the other hand, the nearby Scd galaxy M\\,33 has a firm upper limit on the mass of any central dark object of $<1500 M_\\odot$ \\citep{Gebhardt01}, far below the typical $10^5 M_\\odot$ BHs found in galaxies such as NGC\\,4395. Thus BHs can indeed exist in bulgeless systems, but exactly how common such systems are remains unknown.  From a spectral analysis of SDSS data, \\cite{GH07c} have found a population of low-mass BHs ($M_{\\rm BH} \\approx 10^5$--$10^6 M_\\odot$) accreting at appreciable Eddington ratios. The host galaxies from the initial sample of 19 sources \\citep{GH04}, using {\\em HST} imaging, are found to be mid- to late-type spirals or spheroidals similar to POX\\,52 \\citep{GHB08}. None are as late-type as NGC\\,4395, however, and so the incidence of BHs in very late-type systems remains uncertain.  For such systems, there is also the possibility that BHs are replaced by nuclear star clusters (NCs). Both \\cite{Ferrarese06} and \\cite{WH06} suggest that BHs in the $M_{\\rm BH}-\\sigma_\\star$ relation are replaced by compact massive objects at low mass, with a generalized correlation $M_{\\rm CMO} - \\sigma_\\star$. Such compact objects could include NCs instead of BHs, commonly found in late-type spirals \\citep{Boker02,Boker04}. It is therefore not immediately obvious that such late-type spirals {\\em should} contain BHs. Establishing a BH census for late-type spirals is important to determine if BHs are indeed formed in earlier-type galaxies only (ellipticals or early-type spirals), or if they also exist in very late-type galaxies.  One complication in late-type spirals is the abundance of gas and dust, with strong levels of star formation.  If these systems contain BHs, they are likely to be of relatively low mass, and hence the AGN, even when fully active, will be of quite low luminosity.  While the mere presence of an AGN indicates a BH, energetically weak or heavily embedded AGNs are extraordinarily difficult to observe, particularly at optical wavelengths where most of these observations have historically occurred.  Dust obscuration severely attenuates optical and ultraviolet flux, and central star formation can dominate the optical/infrared emission.  X-ray observations allow us to peer through the obscuring gas and dust column, since high column densities are required to suppress X-ray flux (particularly energetic X-rays). X-ray detections are therefore a useful indicator of AGN activity, and since the launch of the {\\em Chandra X-ray Observatory}, the sensitivity with which to carry out such studies has been greatly improved. In particular, the Advanced CCD Imaging Spectrometer (ACIS; \\citealt{Garmire03}) provides better than 0\\farcs5 resolution with very low background contamination ($\\sim 10^{-7}$ counts pix$^{-1}$ sec$^{-1}$ between 0.5 and 8 keV). This allows low flux limits to be achieved with very short (few ks) exposures, enabling surveys of non-X-ray selected samples \\citep{Ho01}. \\cite{Gallo07} have used X-ray detections successfully to identify BHs in low-mass bulges, and have begun to establish statistics on AGN X-ray activity as a function of host bulge mass. X-ray observations have also been useful in searching for nuclear activity in quiescent (i.e. optically faint) BHs residing in massive elliptical galaxies \\citep{DiMatteo00,Loewenstein01,Pellegrini05,Soria06,WTH08}, and in star-forming early-type spiral galaxies \\citep{TG07}.  In this paper we investigate the nuclear activity in a sample of late-type spiral galaxies using archival {\\em Chandra} data, which provide a clearer method of detecting AGN activity in such star formation-dominated, potentially heavily obscured environments.  Systems with nuclear X-ray detections coincident with the optical nucleus are identified as AGN candidates, and after considering the possible contamination of our results from X-ray binaries, we estimate the prevalence of AGNs in late-type spiral galaxies. We discuss our sample selection in \\S~\\ref{sec:sample} and our results in \\S~\\ref{sec:results}. We conclude with a discussion of possible non-AGN sources and directions for future work in \\S~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  The evidence for the majority of these candidates being true AGNs is circumstantial. For most of these systems, the physical scale of an X-ray point source in {\\em Chandra} data is of order tens of parsecs, making it extremely unlikely that the vast majority of detections are serendipitous field X-ray binaries along the line of sight to or near the nucleus. Furthermore, the probability of detecting a background/foreground X-ray source within the {\\em Chandra} PSF is extremely low, of order $10^{-7}$ \\citep{Gallo07}. With a few exceptions, the positional agreement between the X-ray source and the 2MASS position of the nucleus is excellent (less than 2\\arcsec, compared to the {\\em Chandra} and 2MASS positional uncertainties of $\\sim$1\\arcsec). NGC\\,4945 has a larger offset ($\\delta \\approx$ 3\\farcs2), but the presence of bright diffuse X-ray emission makes point-source detection more challenging, and once again this object is a known Seyfert 2.  Many of these galaxies have \\hii \\ nuclei as identified in the Palomar survey, but most of the optical observations have slit apertures of $2 \\arcsec \\times 4 \\arcsec$ \\citep{Ho_2_95}, much larger than the {\\em Chandra} PSF. The Palomar data are therefore probing a much larger region than the X-ray point sources, and we must be careful when making comparisons. As these are gas- and dust-rich systems, there is clearly some level of star formation in the circumnuclear region. If we consider, for example, typical \\ha \\ luminosities emitted by such low-luminosity AGNs, they are significantly lower than the measured \\ha \\ luminosities in the Palomar data. Assuming the bolometric, X-ray, and \\ha \\ luminosities of AGNs are related by $L_{\\rm bol} = 16 L_{\\rm X} = 220 L_{\\rm H\\alpha}$ \\citep{Ho08}, a range $L_{\\rm X} = 10^{37}$ -- $10^{40}$ ergs s$^{-1}$ corresponds to $L_{\\rm H\\alpha} = 7 \\times 10^{35}$ -- $7 \\times 10^{38}$ ergs s$^{-1}$. Palomar objects in our sample with \\hii \\ nuclei have $L_{\\rm H\\alpha} = 10^{37.5}$ -- $10^{38.8}$ ergs s$^{-1}$, which can explain why such low-luminosity AGNs, if present, were missed in the Palomar data.  The galaxies in our sample are low-mass, late-type spirals, and thus likely harbor a NC at the center of the gravitational potential well. \\cite{Boker02,Boker04} found that $\\sim$75\\% of late-type spirals ($T$ = 6--9) host a compact, photometrically distinct source located near the photometric center of the galaxy, with the surface brightness profile deviating from a pure disk profile. These NCs have mean effective radii of $r_{\\rm eff} = 3.5$ pc and $I$-band luminosities of $L_I=10^{6.2} L_\\odot$. \\cite{Rossa06} found that for NCs in galaxies within a similar range in $T$ type, the average cluster mass is $M=10^{6.25} M_\\odot$. The increased stellar density at the center therefore implies an increased chance of detecting, in particular, a low-mass X-ray binary (LMXB) as a nuclear X-ray point source, confusing our results. M\\,33, for example, likely has no central massive BH \\citep{Gebhardt01}, but does possess an X-ray source coincident with the optical nucleus, with properties suggestive of an X-ray binary \\citep{Dubus04}. This is why we do not included M\\,33 as part of our AGN candidates. NGC\\,2403 also hosts an X-ray binary in its known nuclear star cluster \\citep{Yukita07}, although the X-ray luminosity is not high enough to be confused with an AGN (in this work, no detection was found above the limiting luminosity of $10^{36.1}$ ergs s$^{-1}$). X-ray binaries must therefore be carefully considered when attempting to detect low-luminosity AGNs. Although NCs also occur frequently in more massive early-type spiral galaxies \\citep{Carollo97,Carollo98} and intermediate-luminosity elliptical galaxies \\citep{Cote06}, LMXB contamination is not important in these cases because at these larger BH masses $L_{\\rm X} > 10^{40}$ ergs s$^{-1}$ due to nuclear activity, much greater than typical LMXB X-ray luminosities.  An X-ray detection in a NC is much more likely to be a LMXB rather than a high-mass X-ray binary (HMXB) because the mean age of NCs in late-type hosts is $\\sim 10^{8 {\\rm -} 9}$ yrs \\citep{Rossa06}, older than typical high-mass stars in HMXBs. Although many late-type galaxies have \\hii \\ nuclei \\citep{Ho_3_97}, the slit aperture used, in most cases, is significantly larger than the {\\em Chandra} PSF. Given that nearly all galaxies in the survey exhibit some emission lines in their nuclear spectra, and that late-type galaxies in particular are gas and dust rich with strong levels of star formation, the prevalence of \\hii \\ nuclei is not surprising. This does not necessarily imply, however, that the NCs have ongoing star formation (and therefore HMXBs are more likely), as evidenced by the inferred NC ages from \\cite{Rossa06}.  For our purposes, we will approximate the NCs as globular clusters (GCs) in order to estimate the incidence of LMXBs. Although the environments of both types of clusters are widely different, NCs and GCs share remarkably similar structural and dynamical properties. Viewed in the fundamental plane (relating effective radius, surface density, and velocity dispersion), NCs occupy a region very distinct from classical bulges and elliptical galaxies, and fall on the high-mass extension of the GC sequence \\citep{Walcher05}. With some important exceptions (see below), approximating NCs as GCs is therefore well justified.  As in \\cite{Gallo07}, we can estimate the expected number of LMXBs brighter than $3.2 \\times 10^{38}$ ergs s$^{-1}$ for a GC (and thus NC) of given mass, size, and $g-z$ color, using an expression derived by \\cite{Sivakoff07}:  \\begin{equation} n_{\\rm X} = 8 \\times 10^{-2} \\left( \\frac{M}{10^6 M_\\odot} \\right)^{1.237} 10^{0.9(g-z)} \\left( \\frac{r_{\\rm h,corr}}{1 {\\rm \\ pc}} \\right)^{-2.22}. \\end{equation}  \\noindent The quantity $r_{\\rm h,corr}$ specifies the corrected half-mass radius, where $r_{\\rm h,corr} = r_{\\rm h} \\times 10^{0.17[(g-z)-1.2]}$. Since this depends rather weakly on color, we can approximate $r_{\\rm h,corr} \\approx r_e = 3.5$ pc as a mean value for NCs in late-type spirals. \\cite{Rossa06} find a mean $B-V$ color for such NCs of order $\\sim 0.5$--$0.6$ mag. In order to get $g-z$ colors, we use approximate color transformations from \\cite{Fukugita95} by approximating the stellar population in a NC as that between an Scd and an Sbc galaxy, which have $B-V$ colors similar to those of NCs. This gives $g-z$ colors ranging from 1.05 to 1.17 mag; we adopt $g-z = 1.1$ mag. Therefore the mean number of LMXBs we expect is $n_{\\rm X} \\approx 0.1$. This estimate is specifically for sources brighter than $\\sim 10^{38}$ ergs s$^{-1}$, which corresponds to the completeness limit of \\cite{Sivakoff07}. The limiting luminosity for our X-ray detections varies from $L_{\\rm X} = 10^{36.2}$ to $10^{38.1}$ ergs s$^{-1}$. For objects in the \\citeauthor{Sivakoff07} sample with a lower limiting luminosity ($\\sim 10^{37}$ ergs s$^{-1}$ or lower), however, they detected only a factor of a few more sources.  An important caveat to this estimate is the question of stellar age. GCs generally have old stellar populations, formed via a single burst of star formation, whereas NCs have repeated star formation episodes since the reservoir of gas can be replenished. This leads to, for a given cluster mass, a younger mean stellar age for NCs compared to GCs, and thus bluer colors \\citep{Rossa06,Walcher06}. As a result, the expected number of LMXBs estimated above is an upper limit, since it was derived for GCs, and the true number of LMXBs in NCs is likely lower. This is due to the $\\sim 1$ Gyr timescale necessary for the stellar companion to come into Roche lobe contact with its BH companion (and thus turn on the LMXB), either by stellar evolution or via orbital decay \\citep{WG98}.  Using $n_{\\rm X} \\approx 0.1$ as our estimate, this would imply only $\\sim 5$ X-ray detections from LMXBs in a sample of 64 late-type spirals (assuming 75\\% have NCs with mean properties), of which M\\,33 is a typical example. This is a factor of 3--4 lower than our observed detection rate. The observed X-ray luminosities in our sample are similar to the known AGNs, making our other detections reasonable AGN candidates. This implies that NCs and BHs need not be necessarily mutually exclusive, if the majority of our AGN candidates are true AGNs. Indeed, \\cite{Seth08} find that over all Hubble types, roughly 10\\% of all galaxies hosting a NC also host an AGN, based on optical spectroscopy, with a fraction that increases strongly with galaxy and NC mass.  It is important to note that these detections are AGN {\\em candidates}, and are worthy of further detailed observations to properly identify these sources. \\cite{Ghosh08} provide some evidence from a multiwavelength analysis that NGC\\,3184, NGC\\,4713, and NGC\\,5457 are true AGNs. Since most of our candidates have very few counts, long-integration X-ray spectroscopy will be useful in differentiating true AGNs from LMXBs, particularly at very soft and hard energies. As described above, AGNs have X-ray spectra composed of an absorbed power law (with $\\Gamma \\approx 1.8$), whereas X-ray binaries have spectra that can be separated into three states: a high/soft state with a $\\sim$1 keV thermal component and a weak nonthermal tail at high ($>$ 10 keV) energies; a very high state with a steep power law ($\\Gamma \\approx 2.5$) combined with a $\\sim$1 keV thermal blackbody, which can constitute a significant amount of the total flux; and a low/hard state with a single power law ($1.4 < \\Gamma < 2.1$) \\citep{RM06}. Of these three states, only the low/hard state could be confused with an AGN, but X-ray binaries in these states typically have extraordinarily low luminosities ($L_{\\rm X} = 10^{30.5}$ -- $10^{33.5}$ ergs s$^{-1}$), which would distinguish them from true AGNs.  Our spectral fits (Table~\\ref{tab:spectra} and Figure~\\ref{fig:spectra}) are insufficient to definitively rule out one model over another, though they remain very suggestive. For NGC\\,4039, a pure thermal blackbody is likely ruled out, whereas an absorbed power law is plausible, and an absorbed power law + blackbody model (indicative of a very high state LMXB) is consistent with the data. Given that the reduced $\\chi^2$ of the latter model is $<$1, however, we may be overfitting the data, and so we must be cautious. Although it may appear as though this is more likely an X-ray binary (of extremely high luminosity $>10^{39}$ ergs s$^{-1}$), longer-integration data are needed to distinguish between the two models.  NGC\\,4490 and NGC\\,6946 are both well fit by simple absorbed power law models ($\\chi^2 \\approx 1$), consistent with the AGN interpretation. In the case of NGC\\,4490, the unabsorbed blackbody model results in a poorer fit, particularly at low energies, while for NGC\\,6946, both models are consistent with the data. Unfortunately, the data are not of high enough quality to differentiate between models with more parameters, such as an absorbed power law + blackbody, since $\\chi^2 \\lesssim 1$ in all cases. Once again, these results underscore the need for higher S/N data if we are to use a $\\chi^2$ statistic to differentiate between various models.. As a check, we also analyzed NGC\\,4395, a known Seyfert 1, which is unsurprisingly well fit by an absorbed power law, whose parameters are consistent with those given by \\cite{Moran05}.  In Figure~\\ref{fig:colors} we show the X-ray colors of those AGN candidates with sufficient counts. We also include various absorbed power-law models and a 1 keV blackbody (representative of the high/soft state from \\citealt{RM06}), computed for observations with ACIS-S using {\\tt XSPEC}. All AGN candidates are consistent with absorbed power-law models, and inconsistent with the 1 keV blackbody. Although the uncertainties are large, it appears as though most of these candidates exhibit softer X-ray spectra than is typically assumed for low-luminosity AGNs. Low-mass AGNs can exhibit a range of X-ray spectra, however, with $\\Gamma \\approx$ 1.0 to 3.0 \\citep{GH07a}. Clearly, detailed X-ray spectroscopy is needed to further characterize these systems.  Once AGNs are positively identified, determining a BH mass estimate will be challenging. Dynamical measurements are currently not feasible in these objects. Reverberation mapping \\citep{Peterson07} or virial estimates \\citep{GH05} based on the strength of optical emission lines (if detected) will likely require excellent seeing conditions with a narrow slit or adaptive optics. As these are all nearly bulgeless systems, using the bulge velocity dispersion $\\sigma_\\star$ or the bulge luminosity $L_B$ \\citep{MH03} to determine $M_{\\rm BH}$ is challenging and highly uncertain. NCs with BHs, however, fall on the $M_{\\rm BH}-\\sigma_\\star$ relation---for example NGC\\,4395 \\citep{FH03,Peterson05} and G1 \\citep{GRH02,GRH05}---consistent with the extrapolation of $M_{\\rm BH}-\\sigma_\\star$ to lower masses \\citep{BGH05}. The velocity dispersions of NCs \\citep{Walcher05} could therefore be used to estimate the BH mass. BH masses can also be estimated from X-ray variability timescales derived from complete power density spectra (\\citealt{McHardy06}), obtained over long monitoring campaigns. Alternatively, the BH mass can be estimated from the excess variance measured in shorter X-ray data sets \\citep{NPC04,GNC08}.  Our detection rate suggests an AGN fraction in late-type spiral galaxies of $\\sim$20\\%--25\\% (depending on the LMXB contamination), higher than that observed by \\cite{Ho_5_97}. (We should caution, however, that several detections are at low S/N. Using only S/N$>$3 detections, our AGN fraction becomes $\\sim$17\\%.) This would imply that we are indeed missing many low-mass AGNs in optical surveys, and that BHs in such systems are more prevalent than previously suggested. Optical surveys of low-mass AGNs can only find systems with high accretion rates, such that their bolometric luminosity $L_{\\rm bol}$ approaches a substantial fraction of the Eddington luminosity $L_{\\rm Edd} \\equiv 1.26 \\times 10^{38} (M_{\\rm BH} / M_\\odot)$ ergs s$^{-1}$ \\citep{GH07b}. \\cite{GH07c} found a median $L_{\\rm bol}/L_{\\rm Edd} = 0.4$ and a minimum $L_{\\rm bol} / L_{\\rm Edd} \\approx 10^{-2}$ in a sample of 174 optically selected low-mass AGNs. In a pilot study drawn from this sample, \\cite{GH07a} measured $L_{\\rm X}$(0.5--2 keV) $\\approx 10^{41}$--$10^{43}$ ergs s$^{-1}$, much higher than any of our measurements. Most of our candidates with available $M_{\\rm BH}$ estimates exhibit $L_{\\rm bol}/L_{\\rm Edd}$ ratios orders of magnitude lower than those of \\citeauthor{GH07c}, assuming a correction factor of 16 to convert from $L_{\\rm X}$ to $L_{\\rm bol}$ \\citep{Ho08}. The one exception is NGC\\,4395, a known Seyfert nucleus, with $L_{\\rm bol} / L_{\\rm Edd} \\approx 10^{-2.1}$, consistent with the value of $2 \\times 10^{-3}$ to $2 \\times 10^{-2}$ found by \\cite{FH03}. (In comparison, POX\\,52 has $L_{\\rm bol}/L_{\\rm Edd} \\approx 1$; \\citealt{Barth04}.) Our estimated Eddington ratios are similar, however, to those seen in massive BHs in earlier-type hosts with low-luminosity AGNs \\citep{Ho08}, suggesting that we are detecting the low-mass analogs of weakly accreting BHs. X-ray surveys therefore allow us to probe much fainter down the luminosity function of low-mass AGNs, detecting weakly accreting systems that would be completely swamped by galactic emission at optical wavelengths, and obtain a much higher detection rate. The \\citeauthor{GH07c} sample, for instance, identified only 174 low-mass AGNs out of 8435 optically identified AGNs (from a parent sample of $\\sim$600,000 galaxies).  Given the low estimated accretion rates of our AGN candidates, radio observations can help to clarify the nature of these X-ray sources. If these are indeed low-mass and low-accretion-rate systems, we should expect them to have a radio signature, much like the more massive low-luminosity BHs \\citep{Ho08}. Extrapolating from higher masses, we would expect 2 cm radio fluxes of order 10 to 100 $\\mu$Jy \\citep{Nagar02}. If these X-ray sources are X-ray binaries instead, however, they would be accreting at Eddington or super-Eddington rates to achieve such high luminosities. X-ray binaries in such states do no have a radio counterpart \\citep{RM06}.  \\cite{Satyapal08} have conducted a similar survey of AGNs in the infrared, searching for emission-line signatures of AGN activity (in particular, \\nev \\ emission; see \\citealt{AS08}) in late-type galaxies that are optically quiescent. This should identify weakly accreting systems similar to our candidates. In a sample of 32 objects, nine of which overlap with our {\\em Chandra} sample, they find seven AGN candidates, a similar detection fraction. Interestingly, of the nine overlap objects, six (IC\\,342, NGC\\,925, NGC\\,3184, NGC\\,4490, NGC\\,4559, and NGC\\,6946) are identified in this paper as AGN candidates, but not by \\citeauthor{Satyapal08} It is possible that the weakly active nuclei picked out by our X-ray search method are too faint to emit much \\nev \\ emission. We can estimate the total \\nev \\ luminosity $L_{\\rm NeV}$ for our sample assuming the bolometric luminosity $L_{\\rm bol} = 16 L_{\\rm X}$ and $L_{\\rm bol} = 9 L_{5100}$ \\citep{Ho08}, where $L_{5100}$ is the continuum luminosity at 5100 \\AA, coupled with the correlation between $L_{5100}$ and $L_{\\rm NeV}$ found by \\cite{Dasyra08}. For our range in X-ray luminosities ($L_{\\rm X} = 10^{37}$ -- $10^{40}$ ergs s$^{-1}$), this corresponds to \\nev \\ luminosities of $L_{\\rm NeV} = 10^{35}$ -- $7 \\times 10^{37}$ ergs s$^{-1}$. These luminosities are orders of magnitude fainter than those measured by \\cite{Satyapal08}. Even NGC\\,3184, which appears to be a {\\em bona fide} AGN \\citep{Ghosh08}, was not detected. The \\citeauthor{Satyapal08} detections of \\nev \\ are in excess of the \\citeauthor{Dasyra08} relation, however, thus it is possible these relations do not hold for very late-type spiral galaxies.  Another factor that might be relevant is the ionization parameter; since \\nev\\ is a high-ionization line, its strength decreases dramatically when the ionization parameter decreases \\citep{AS08}.  By analogy with the local population of LINERs and Seyferts \\citep{Ho08}, the very low Eddington ratios of our low-mass systems should correspond to low ionization parameters, and thus plausibly weak \\nev\\ emission.   Finally, we note that NGC\\,3556 is identified by \\citeauthor{Satyapal08} as an AGN candidate but has no X-ray detection in this work. \\cite{WCI03} identify an X-ray source as a candidate for the galactic nucleus in NGC\\,3556 (which they number source 35); we chose not to include this source in our candidate list as the offset from the optical nucleus is 10\\arcsec. More detailed study is clearly warranted, and it will likely take the combined efforts of large multiwavelength surveys to find such low-mass and low-luminosity AGNs, and to properly assess the global AGN fraction in these late-type systems."
    },
    "0809/0809.1169_arXiv.txt": {
        "abstract": " ",
        "introduction": "Basic parameters of the close binary system AE~Aqr are listed in Tab.\\,\\ref{ikhsanov-t1}. The red dwarf overflows its Roche lobe and transfers material through the L1 point towards the white dwarf. This material, however, is neither accreted onto the surface of the white dwarf nor stored in a disk around its magnetosphere. Instead, it is streaming out from the system with an average velocity of $300\\,{\\rm km\\,s^{-1}}$. The spin-down power of the white dwarf exceeds the bolometric luminosity of the system by a factor of a few (see Tab.\\,\\ref{ikhsanov-t2}). Finally, the system shows flaring activity whose properties are absolutely unique among all classes of flaring astrophysical objects \\citep[for a review see][and references therein]{Beskrovnaya-etal-1996,Ikhsanov-etal-2004}.  Theoretical studies of the system during the last decade have led to a conclusion that AE~Aqr does not fit in any of the accretion-based models developed for Cataclysmic Variables \\citep[see, e.g.][]{Wynn-etal-1997}. In other words, the white dwarf in the system is not an accretion-powered pulsar. This situation can be explained in terms of the centrifugal inhibition (propeller) model provided that the magnetic field of the white dwarf is strong enough for its magnetospheric radius to exceed the corotational radius (see Tab.\\,\\ref{ikhsanov-t3}). However, how strong is the magnetic field\\,? A lack of the white dwarf's photospheric lines in the system spectra makes impossible a direct measurement of the field strength through Zeeman effect. Among other methods, which can be used to answer this question, are the simulation of Doppler H$\\alpha$ tomogram of the system, the modeling of the rapid braking of the white dwarf, the analysis of the circularly polarized optical emission, and a reconstruction of power source of the pulsing hard X-ray emission recently discovered by SUZAKU space telescope. Following these methods one comes to a conclusion that the dipole magnetic moment of the white dwarf in AE~Aqr is close to $1.5 \\times 10^{34}\\,{\\rm G\\,cm^3}$ and its rapid braking is governed by the pulsar-like spin-down mechanism.  \\begin{table}[t] \\caption[]{Parameters of AE~Aquarii$^{*}$} \\label{ikhsanov-t1} \\begin{tabular}{l|ccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\it System parameters} &  Distance  & Binary period & Inclination & Eccentricity & Mass ratio  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Value &  $(100\\pm30)$\\,pc  &  $9.88$\\,hr &  $50^{\\degr}<i<70^{\\degr}$ & $0.02$ & $0.6-0.8$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\it Stellar parameters}   &  Type   & Mass ($M_{\\sun}$) & Spin period &  $\\dot{P}$ (${\\rm s\\,s^{-1}}$) & $\\left[\\vec{\\bf \\Omega}\\,\\wedge\\,\\vec{\\bf m}\\right]^{**}$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Secondary &  K3V--K5V   & $0.41\\sin^{-3}{i}$ & $\\sim 9.88$\\,hr & -- & --\\\\ Primary   & White Dwarf &  $0.54\\sin^{-3}{i}$ & $33.08$\\,s & $5.64\\times 10^{-14}$ & $74^{\\degr} - 76^{\\degr}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\newline $^*$~~For references see e.g. \\citet{Ikhsanov-etal-2004} \\newline $^{**}$~The angle between the rotational and magnetic axes \\end{table} ",
        "conclusions": "Simulations of Doppler H$\\alpha$ tomogram, analysis of the rapid braking of the white dwarf, and the modeling of particle acceleration in AE~Aqr suggest that the dipole magnetic moment of the white dwarf in this system is $\\mu \\simeq 1.5 \\times 10^{34}\\,{\\rm G\\,cm^3}$. Under these conditions the spin-down of the white dwarf is governed by the pulsar-like mechanism, i.e. its spin-down power is released predominantly in a form of the magneto-dipole waves and accelerated particles and significantly exceeds the bolometric luminosity of the system. This makes the degenerate companion of AE~Aqr the first white dwarf in the family of spin-powered pulsars.  The above evaluation of the magnetic field of the white dwarf does not contradict observations of the circularly polarized optical emission as well as all presently established properties of the system. Moreover, the pulsar-like approach appears to be an effective tool in explanation of the origin of hot polar caps at the white dwarf surface \\citep[in terms of the dissipation of the backflowing current in the magnetosphere, see, e.g.][]{Ikhsanov-etal-2004, Ikhsanov-Biermann-2006} and properties of the system derived from observations in high-energy part of the spectrum.  At the same time, the above conclusion rises a problem about the history of the system. The age of the white dwarf evaluated from its surface temperature  $(1-1.6) \\times 10^4$\\,K \\citep{Eracleous-etal-1994} is limited to $\\ga 10^8$\\,yr \\citep[see, e.g.][]{Schoenberner-etal-2000}, while the spin-down time scale is only $P/\\dot{P} \\sim 10^7$\\,yr. This indicates that the system history contains an accretion-driven spin-up epoch. However, the spin period to which a white dwarf with $\\mu \\sim 10^{34}\\,{\\rm G\\,s^{-1}}$ could be spun-up by a disk accretion is substantially larger than the currently observed 33\\,s period \\citep[for discussion see][]{Ikhsanov-1999}. It, therefore, appears, that the magnetic field of the white dwarf has been amplified during a previous epoch. As shown by \\citet{Ikhsanov-1999}, this could occur at the end of the accretion-driven spin-up phase due to gravitational waves emission instability of the degenerate core of the star. In this light, AE~Aqr can be considered as a precursor of a polar. The detailed study of the corresponding evolutionary track isnow work in progress.  \\vspace{1cm}  Nazar Ikhsanov acknowledge the support of the European Commission under the Marie Curie Incoming Fellowship Program. This work was partly supported by Russian Foundation of Basic Research under the grant number 07-02-00535a."
    },
    "0809/0809.3446_arXiv.txt": {
        "abstract": "In this Letter we explore the hypothesis that the quasar SDSSJ092712.65+294344.0 is hosting a massive black hole binary embedded in a circumbinary disc. The lightest, secondary black hole is active, and gas orbiting around it is responsible for the blue--shifted broad emission lines with velocity off-set of 2650 km s$^{-1}$, relative to the galaxy rest frame. As the tidal interaction of the binary with the outer disc is expected to excavate a gap, the blue-shifted narrow emission lines are consistent with being emitted from the low-density inhomogeneous gas of the hollow region. From the observations we infer a binary mass ratio $q\\approx 0.3$, a mass for the primary of $M_1\\approx 2 \\times 10^9\\,M_{\\odot},$ and a semi-major axis of 0.34 pc, corresponding to an orbital period of 370 years. We use the results of cosmological merger trees to estimate the likely-hood of observing SDSSJ092712.65+294344.0 as recoiling black hole or as a binary.  We find that the binary hypothesis is preferred being one hundred times more probable than the ejection hypothesis. If SDSSJ092712.65+294344.0 hosts a binary, it would be the one closest massive black hole binary system ever discovered. ",
        "introduction": "If massive black holes (MBHs) are ubiquitous in galaxy spheroids and galaxies experience multiple mergers during their cosmic assembly, MBH {\\it binaries} (MBHBs hereon) should be common, albeit transient features of galactic bulges.  Observationally, the paucity of dual active nuclei points toward rapid inspiral of the MBHs down to parsec scales where they form a gravitationally bound pair (Mayer et al. 2007).  After birth, the MBHB further hardens under the action of gas/dynamical and/or stellar torques (Begelman, Blandford \\& Rees 1980) and depending on their efficiency (Milosavljevic \\& Merritt 2001; Armitage \\& Natarajan 2002; Berczik et al. 2006; Sesana et al. 2007; Perets et al. 2007; Colpi et al. 2007) the binary may remain in its hard state for a relatively long time, before coalescing under the action of gravitational waves.  MBHBs can coalesce shortly, and a clear prediction of general relativity is that the relic MBH receives a kick in response to the asymmetric pattern of gravitational waves emitted just prior coalescence.  For slowly spinning MBHs recoil velocities of a few $\\times 100$ km s$^{-1}$ were predicted, but recent numerical advances in general relativity have indicated that velocities as extreme as $\\sim 4000$ km s$^{-1}$ can be attained for maximally rotating Kerr MBHs with spin vectors lying in the orbital plane (Schnittman \\& Buonanno 2007; Lousto \\& Zlochower and references therein).  These velocities far exceed the escape speed even from the brightest galaxies and for this reason MBHs receding from their parent host have been considered as electromagnetic targets of a coalescence event (Devecchi et al. 2008; Loeb 2007; Volonteri \\& Madau 2008).   Recently, the quasar SDSS J092712.65+294344.0 (hereafter SDSSJ0927; Adelman--McCarthy et al. 2007) has been considered as first candidate of a recoiling MBH by Komossa, Zhou \\& Lu (2008, KZL).  SDSSJ0927 exhibits, in its optical spectrum, two distinct sets of lines offset by 2650 km s$^{-1}$ relative to each other. The first set of narrow emission lines (NELs) has redshift $z=0.713$ (identified by KZL as the redshift of the galaxy that hosted the MBH at its birth). We will refer to this set that is typical of an AGN as rest-frame NELs (rf--NELs hereon; the red NELS in KZL).  The rf--NELs do not show any evidence of ionization stratification (as in few other quasars), possibly because these lines are not resolved in the SDSS spectrum. The second set of lines comprises two blue--shifted systems featuring different FWHM: the broad Mg II and Balmer emission lines (b--BELs hereon) with FWHM $\\approx 4000$ km s$^{-1}$, and the high-ionization narrow lines (b--NELs) with FWHM $\\approx 460-2000$ km s$^{-1}$, both consistent with a redshift $z=0.698$. The b--BELs and b--NELs refer to a line system in coherent motion relative to the rest-frame of the host, with a light-of-sight velocity of 2650 km s$^{-1}$.  In KZL, the permitted b--BELs result from a broad line emission system gravitationally bound to the recoiling MBH, while b--NELs remains problematic for the reasons explained in Section~\\ref{sec:prob}.  The aim of this Letter is to suggest an alternative interpretation to SDSSJ0927: that SDSSJ0927 is an unequal--mass MBHB embedded in a circumbinary gaseous disc at the centre of its host galaxy.  In Section ~\\ref{sec:model} and ~\\ref{sec:orb} we describe the main properties of the model that are requested to reproduce the observed properties of b--BELs.  In Section ~\\ref{sec:gap} we discuss the possible origin of b--NELs and of their Doppler shift. In Section~\\ref{sec:concl} we present our conclusions. ",
        "conclusions": "\\label{sec:concl}  The binary hypothesis for SDSSJ0927 relies on a simple set of assumptions.  We can relax the assumption of circular orbits, and assume a non zero eccentricity ($e<0.5$), expected to be present after gap opening (Armitage \\& Natarajan 2005; Dotti et al. 2006a, 2007; Cuadra et al. 2008).  Eccentricities of this magnitude are still compatible with the binary model.  Larger values are instead unlikely in SDSSJ0927: if the binary has a large eccentricity and $M_2$ is near apocentre, an extremely massive $M_1$ is needed to justify such a large orbital velocity.  On the other hand, Shields et al. (2009) set an upper limit on the change of the Doppler shift of the b-BEL system ($<10$ km/s/year). In the binary scenario, this prevents $M_2$ to be caught too close to the pericentre. We note that assuming circular orbit and the set of MBHB parameters discussed above, the b-BEL system moves only of 4.2 km/s over three years, corresponding to an average shift of $\\approx 1.4$ km/s/year.   SDSSJ0927 is clearly exceptional. Either it is the first example of a MBH ejected from its nucleus due to the gravitational recoil, or it harbours one of the smallest separation MBHBs ever detected. Only the possible binary in OJ287 could have closer MBHs ($\\sim 4.5\\times 10^{-2}$ pc, Valtonen 2007). SDSSJ0927 would provide clear evidence that MBHBs can indeed reach coalescence in less than a Hubble time, and that Nature is able to solve the last-parsec problem. If the model discussed in this letter is confirmed, SDSSJ0927 will represent the first detection ever of emission from a circumbinary gap around a MBHB.   The proof of existence of binaries with such a small orbital separation is of extreme importance for gravitational wave experiments such as LISA (Bender et al. 1994), and the Pulsar Timing Arrays (PTA, Sesana et al. 2008).  Since our theoretical expectations for a binary with the characteristics of SDSSJ0927 are of order $10^{-4}$ (binary hypothesis) down to $10^{-6}$ (ejection hypothesis) per year, either we have been extraordinarily lucky, or MBHBs are {\\it more common} than predicted, opening exciting possibilities for LISA and PTA."
    },
    "0809/0809.1396_arXiv.txt": {
        "abstract": "The Monoceros R2 region was first recognized as a chain of reflection nebulae illuminated by A- and B-type stars. These nebulae are associated with a giant molecular cloud that is one of the closest massive star forming regions to the Sun. This chapter reviews the properties of the Mon~R2 region, including the namesake reflection nebulae, the large scale molecular cloud, global star formation activity, and properties of prominent star forming regions in the cloud. ",
        "introduction": "The Monoceros R2 region is distinguished by a chain of reflection nebulae that extend over 2\\deg\\ on the sky. The brightest of these nebulae were studied as far back as \\citet{Seares20}, and were included in a larger study of nebulous objects by \\citet{Hubble22}, who demonstrated that the extended emission can be attributed to the associated stars. The nebulae are included in the Catalog of Bright Diffuse Galactic Nebulae constructed by \\citet{Cederblad46}. Various lists of nebulae identified from inspection of the Palomar Observatory Sky Survey appear in \\citet{DG63,DG66}, \\citet{VDB66}, and \\citet{Herbst76}. \\citet{GGD78} identified seven additional nebulous objects as candidate Herbig-Haro objects, although they were later shown to be reflection nebulae \\citep{Cohen80}. The nomenclature ``Mon~R2'' originated with \\citet{VDB66} to indicate the second association of reflection nebulae in the constellation Monoceros\\footnote{The Mon~R1 reflection nebulae are part of the Mon~OB1 association, which includes the Galactic cluster NGC~2264.}.  \\begin{table} \\caption{Optical Nebulae in the Mon~R2 Region\\label{tbl:nebula}} \\smallskip \\begin{center} {\\small \\begin{tabular}{lcclrrrr} \\tableline \\noalign{\\smallskip} Source & $\\alpha$ & $\\delta$ & SpT$^a$  & vdB      & DG       & Ced      & NGC      \\\\ \\cline{2-3} \\\\ & \\multicolumn{2}{c}{(J2000)}\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} HR  1     & 6:06:58 & $-5$:55:10 &         &    &    &    &      \\\\ HR  2     & 6:07:26 & $-6$:38:13 & A0:     &    &    &    &      \\\\ HR  3     & 6:07:30 & $-6$:29:16 &         &    &    &    &      \\\\ HR  4     & 6:07:32 & $-6$:24:03 & B2 V    & 67 & 88 & 63 & 2170 \\\\ HR  5a    & 6:07:46 & $-6$:21:47 & A5:     &    &    &    &      \\\\ HR  5b    & 6:07:46 & $-6$:23:02 &         &    &    &    &      \\\\ HR  6     & 6:07:49 & $-6$:16:35 &         &    &    &    &      \\\\ HR  7     & 6:07:48 & $-6$:26:46 &         &    &    &    &      \\\\ HR  8$^b$ & 6:08:05 & $-6$:21:34 & B1 V    & 69 & 90 & 66 &      \\\\ HR  9$^b$ & 6:08:04 & $-6$:13:38 & B2 V    & 68 & 89 & 65 &      \\\\ HR 10     & 6:08:26 & $-5$:20:19 & B1 V    & 70 & 91 &    &      \\\\ HR 11     & 6:09:15 & $-6$:30:29 &         &    &    &    &      \\\\ HR 12     & 6:09:25 & $-6$:18:35 &         &    &    &    &      \\\\ HR 13     & 6:09:31 & $-6$:19:36 & B3 V    & 72 & 93 & 68 & 2182 \\\\ HR 14     & 6:09:45 & $-6$:18:35 & B5 V    &    &    &    &      \\\\ HR 15     & 6:10:01 & $-6$:18:42 & B8 V    &    &    &    &      \\\\ HR 16     & 6:10:15 & $-6$:27:32 &         &    &    &    &      \\\\ HR 17     & 6:10:37 & $-6$:09:58 &         &    &    &    &      \\\\ HR 18     & 6:10:47 & $-6$:12:41 &         &    & 94 & 69 & 2183 \\\\ HR 19     & 6:10:56 & $-6$:14:22 &         &    &    &    &      \\\\ HR 20     & 6:10:57 & $-6$:16:21 &         &    &    &    &      \\\\ HR 21     & 6:11:00 & $-6$:14:36 &         &    &    &    &      \\\\ HR 22     & 6:11:06 & $-6$:12:33 & B8-A0   &    & 95 & 70 & 2185 \\\\ HR 23     & 6:11:49 & $-6$:09:22 & B4 V    & 74 & 96 & 71 &      \\\\ HR 24     & 6:12:45 & $-6$:12:31 &         &    &    &    &      \\\\ HR 25     & 6:12:46 & $-6$:10:49 &         &    &    &    &      \\\\ HR 26     & 6:14:46 & $-6$:20:36 &         &    &    &    &      \\\\ HR 27     & 6:14:53 & $-6$:22:44 & B7 V    &    & 97 &    &      \\\\ HR 28     & 6:15:21 & $-6$:25:50 & B5 V    &    & 98 &    &      \\\\ HR 29$^c$ & 6:12:50 & $-6$:13:11 & K4      &    &    &    &      \\\\ HR 30     & 6:10:58 & $-6$:14:39 & Be      &    &    &    &      \\\\ DG 92     & 6:08:30 & $-6$:30:30 &         &    & 92 &    &      \\\\ GGD 11    & 6:08:33 & $-6$:18:16 &         &    &    &    &      \\\\ GGD 12    & 6:10:45 & $-6$:12:45 &         &    &    &    &      \\\\ GGD 13    & 6:10:46 & $-6$:13:04 &         &    &    &    &      \\\\ GGD 14    & 6:10:50 & $-6$:12:01 &         &    &    &    &      \\\\ GGD 15    & 6:10:50 & $-6$:11:15 &         &    &    &    &      \\\\ GGD 16    & 6:12:44 & $-6$:15:24 &         &    &    &    &      \\\\ GGD 17    & 6:12:50 & $-6$:12:57 &         &    &    &    &      \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip}  \\multicolumn{8}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^a$ Spectral types from \\citet{Herbst76} and \\citet{Carballo92}. }}\\\\ \\multicolumn{8}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^b$ \\citet{Herbst76} associated HR~8 with vdB~69 and HR~9 with vdB~68. These designations are inconsistent with the coordinates of these sources and the finding chart in that paper (cf. Figure~8 in Downes \\etal~1975). We have maintained the same designations in \\citet{Herbst76}, and have updated the finding chart in Fig.~\\ref{fig:poss} accordingly. }}\\\\ \\multicolumn{8}{l}{\\parbox{0.8\\textwidth}{\\footnotesize $^c$ T~Tauri star also known as Bretz~4 \\citep{her72}. }}\\\\ \\multicolumn{8}{l}{\\parbox{0.8\\textwidth}{\\footnotesize Ced : \\citet{Cederblad46}\\\\ DG  : \\citet{DG63}\\\\ GGD : \\citet{GGD78}\\\\ HR  : \\citet{Herbst76}\\\\ vdB : \\citet{VDB66}\\\\ }}\\\\  \\end{tabular} } \\end{center} \\end{table}  \\begin{figure} \\begin{center} \\includegraphics[angle=0,height=0.9\\textheight]{poss.SR.eps} \\end{center} \\vspace{-5mm} \\caption{ Location of the optical nebulae in Mon~R2 (see Table~\\ref{tbl:nebula}) marked on the blue print from the Palomar Observatory Sky Survey. The regions numbered 1-30 refer to the sources cataloged by \\citet{Herbst76}. The ``GGD'' and ``DG'' sources are the nebulous objects identified by \\citet{GGD78} and \\citet{DG63} respectively. The main Mon~R2 star forming cloud core is also indicated. \\label{fig:poss} } \\end{figure}  Table~\\ref{tbl:nebula} lists the nebulous sources in the Mon~R2 region identified in the studies mentioned above. Coordinates were determined from the digitized Palomar Sky Survey using previously published finding charts \\citep{Downes75,Herbst76,Carballo88} and coordinate lists \\citep{Rodriguez80,Carballo92}. Figure~\\ref{fig:poss} marks the location of the nebulae on the blue print from the Palomar Observatory Sky Survey. The sources GGD 12 through 15 are located within a small area of the sky and are often collectively referred to as GGD~12-15. The most extensively studied region is the active star forming site embedded in a dense molecular core (identified as the ``core'' in Fig.~\\ref{fig:poss}) that lies midway between the reflection nebulae vdB~67 and vdB~69 (see Fig.~\\ref{fig:croman} for a more detailed optical color image of this region). In many contexts, the name ``Mon~R2'' refers to this particular star forming region.  \\begin{figure}[tb] \\begin{center} \\includegraphics[angle=-90,scale=0.54]{ngc2170_croman-cor.eps} \\end{center} \\caption{ Optical photograph of the Mon~R2 region that highlights the presence of reflection nebulae, embedded star forming regions, and dark nebulae. The Mon~R2 core is the red region on the right between the NGC~2170 (vdB~67; lower right) and vdB~69 (center right) blue reflection nebulae. North is to the upper right, and east is to the upper left. Photograph courtesy of R. Croman. \\label{fig:croman} } \\end{figure}  \\citet[see also Herbst \\& Racine~1976]{Racine68} conducted the first detailed spectroscopic and photometric study of the Mon~R2 nebulae and found that the illuminating stars have mainly B spectral types, where B1~V is the earliest type star in the association. \\citet{Herbst76} estimated an age of 6 to 10~Myr for the association. The lower limit is ascertained since the color-magnitude diagram for the Mon~R2 stellar population does not deviate from the zero-age main-sequence for stars with $(B-V)_0 \\le 0.05$~mag, while a pre-main-sequence population is observed for such stars in other clusters such as the Orion Nebula Cluster and NGC~2264. The upper age limit was inferred since the B1 stars in the association still lie on the main-sequence. As noted by \\citet{Herbst76}, the young age for the association suggests that the B-stars likely formed within the cloud, and the reflection nebulae are not the result of a chance encounter of B-stars with a gas-cloud (c.f. the Pleiades; White~2003).  \\citet{Racine68} estimated a distance of 830$\\pm$50~pc to Mon~R2 using spectroscopy and/or $UBV$ photometry for 10 stars in the reflection nebulae. Using similar techniques, \\citet[][as reported in Downes \\etal~1975]{RK68} derived a distance of 700~pc. \\citet{Racine70} revised the distance from \\citet{Racine68} to 950~pc, but did not provide further details on the factors leading to the new estimate. \\citet{Herbst76} fitted the zero-age-main-sequence locus from \\citet{Johnson63} to dereddened $UBV$ photometry and also obtained a distance of 830$\\pm$50~pc, which is the distance typically adopted for the Mon~R2 association. ",
        "conclusions": ""
    },
    "0809/0809.4320_arXiv.txt": {
        "abstract": "The masses of star clusters range over seven decades, from ten up to one hundred million solar masses. Remarkably, clusters with masses in the range $10^4M_\\odot$ to $10^6M_\\odot$ show no systematic variation of radius with mass. However, recent observations have shown that clusters with $\\mcl\\gtrsim 3\\times10^6M_\\odot$ do show an increase in size with increasing mass. We point out that clusters with $\\mcl\\gtrsim 10^6M_\\odot$ were optically thick to far infrared radiation when they formed, and explore the hypothesis that the size of clusters with $\\mcl\\gtrsim3\\times10^6M_\\odot$ is set by a balance between accretion powered radiation pressure and gravity when the clusters formed, yielding a mass-radius relation $\\rcl\\sim 0.3(\\mcl/10^6M_\\odot)^{3/5}\\pc$. We show that the Jeans mass in optically thick objects increases systematically with cluster mass. We argue, by assuming that the break in the stellar initial mass function is set by the Jeans mass, that optically thick clusters are born with top heavy initial mass functions; it follows that they are over-luminous compared to optically thin clusters when young, and have a higher mass to light ratio $\\Upsilon_V=\\mcl/L_V$ when older than $\\sim 1$ Gyr. Old, optically thick clusters have $\\Upsilon_V\\sim \\mcl^{0.1-0.3}$. It follows that $L_V\\sim\\sigma^{\\beta}$, where $\\sigma$ is the cluster velocity dispersion, and $3.6\\lesssim\\beta\\lesssim4.5$.  It appears that $\\Upsilon_V$ is an increasing function of cluster mass for compact clusters and ultra-compact dwarf galaxies. We show that this is unlikely to be due to the presence of non-baryonic dark matter, by comparing clusters to Milky Way satellite galaxies, which are dark matter dominated. The satellite galaxies appear to have a fixed mass inside a fiducial radius, $M(r=r_0)=const.$. ",
        "introduction": "Most stars in the Milky Way are believed to have formed in star clusters, e.g., \\citet{L295}.  In their review \\citet{Clarke} note that $96\\%$ of stars in the Orion B cloud are in clusters, while $50-80\\%$ of stars in Orion A are. The latter fraction is consistent with that found by \\citet{1993AJ....105.1927G} in Tarus. \\citet{2007prpl.conf..361A} suggest that only $\\sim20-25\\%$ of stars form outside clusters.  There is independent evidence, based on simulations of field binary star synthesis rather than on star counts in star forming regions, that the majority of stars formed in clusters \\citep{1995MNRAS.277.1491K}. This result applies to stars formed over the lifetime of the Milky Way disk. The similarity between the stellar initial mass function (IMF) inferred from field stars and that measured in young clusters implies that most low mass stars form in clusters, and hence along side massive stars, rather than in an isolated mode, at least in the Milky Way.  Observations of nearby galaxies suggest that a significant fraction of stars form in clusters in these galaxies as well. For example, \\citet{meurer95} find that $20\\%$ of the UV light in seven starburst galaxies they surveyed comes from young star clusters. Similarly, the cluster fraction of the B band light in the galaxy merger NGC 3256 found by \\citet{zepf99} was $19\\%$, while \\citet{2005ApJ...631L.133F} found that $20\\%$ of the total H$\\alpha$ emission in the Antennae galaxies comes from clusters.  In the case of these external galaxies, the estimates of the cluster fraction ($\\sim20\\%$) are likely to be lower limits, a point noted by \\citet{2005ApJ...631L.133F}. For example, \\citet{meurer95} estimated that the cluster fraction of young stars in NGC 5253 was $\\sim14\\%$ since that fraction of the 2200 \\AA\\ UV light came from clusters. However, more recent radio and infrared observations of NGC 5253 have revealed the presence of a deeply embedded star cluster, with a luminosity $\\sim 4\\times10^{42}\\ergs$ \\citep{2003Natur.423..621T}, and a correspondingly large ionizing flux $Q\\approx  7\\times10^{52}\\s^{-1}$ \\citep{2004ApJ...602L..85T} indicating an age of order 1 Myrs or less. This is $43\\%$ of the bolometric luminosity of the galaxy, given by \\citet{2001ApJ...554L..29G} as $9.2\\times10^{42}\\ergs$, showing that {\\em at least} $43\\%$ of the star formation in this galaxy occurs in a single cluster. This extremely luminous cluster was not detected in the UV by \\citet{meurer95}. It seems likely that much of the clustered star formation in starburst galaxies is similarly obscured, so that the clustered star formation estimate of \\citet{meurer95} is significantly low.  Observations of young star clusters find a power law mass distribution of the form \\be  \\label{eqn:mass distribution} % {dN_{cl}(m)\\over dm}=N_0m^{-\\alpha} \\ee   % with $1.5<\\alpha<2$ in both the Milky Way and in other galaxies \\citep{1989ApJ...337..761K,1997ApJ...476..144M}. This implies that most massive stars are found in the few most massive clusters. If one believes that massive stars are relevant to galaxy formation, understanding star formation in massive clusters is crucial.  In this paper we show that clusters with initial $\\mcl\\gtrsim3\\times10^6M_\\odot$ were likely supported by radiation pressure just before and during the time their stars were forming; it follows that such clusters will exhibit a mass radius relation $\\rcl\\sim \\mcl^{3/5}$, unlike globular clusters, which have radii independent of their masses (by cluster radius we mean the projected radius that encloses half of the cluster light).  In addition, we argue that star formation in massive ($\\mcl\\gtrsim10^6M_\\odot$) clusters will produce an IMF with a larger characteristic mass than star formation in less massive clusters. The IMF in the Milky Way is described by a broken power law \\citep{2001MNRAS.322..231K,2002ApJ...573..366M}, or by a log-normal distribution \\citep{1979ApJS...41..513M}. Similarly, the IMF in young Milky Way clusters such as Orion is described by a broken power, e.g., \\citep{2002ApJ...573..366M}, similar to that of the Milky Way as a whole.  While the classic \\citet{1955ApJ...121..161S} IMF has only an upper and lower mass cutoff, these more recent power-law based estimates of the IMF involve at least one additional characteristic mass. In either the log-normal or powerlaw models, the characteristic mass in the Milky Way is of order $0.5M_\\odot$.  The origin of the characteristic mass is disputed in the literature. For example, \\citet{adams} state that ``\\ldots the Jeans mass has virtually nothing to do with the masses of forming stars''. In contrast, \\citet{McKee07} state ``A recurrent theme in star-formation theory is that the characteristic mass--defined by the peak of the IMF--is the Jeans mass at some preferred density.'' We argue that the second point of view is correct.  The gas in the interstellar medium is not smoothly distributed; a substantial fraction is found in the form of giant (tens of parsec size and $10^6M_\\odot$ mass, in the Milky Way) molecular clouds, which contain parsec scale clumps, which in turn contain sub-parsec scale cores; see, e.g., \\citet{McKee07,bergin07}. The density increases as one moves down the size and mass scale. The masses typically follow a power law distribution similar in form to that of massive star clusters (eqn. \\ref{eqn:mass distribution} above). This behavior is consistent with the notion that the density and mass distributions are controlled by supersonic turbulence. Both analytic theory and simulations predict log-normal density distributions in the gas density and gas surface density of supersonic turbulent flows; observations find surface densities that are consistent with log-normal distributions.  Measured velocity differences are scale dependent, again consistent with turbulent motions; they are large on large scales, but decrease with decreasing scale. The length scale on which the velocity is equal to the thermal sound speed is called the sonic length. Below this scale turbulent motions cannot compress the gas, so that the density is no longer controlled by turbulence; rather, thermal or magnetic pressure, along with the self-gravity of the gas, will control the density. In the Milky way the sonic length is $\\sim0.03\\pc$. Star forming cores are typically about this size or smaller.  We assume that on scales below the sonic length the fragmentation properties are controlled by the thermal properties of the gas.  For an ideal gas equation of state $P=\\rho kT/\\mu\\sim \\rho^\\gamma$, where $\\gamma$ is the effective adiabatic index of the gas. The Jeans mass scales as $T^{3/2}/\\rho^{1/2}\\sim \\rho^{(3\\gamma-4)/2}$. If $\\gamma<4/3$ an increase in the gas density (as, for example, when the clump is self-gravitating) leads to a decrease in the Jeans mass and the likelihood of further fragmentation.  When $\\gamma>4/3$ the Jeans mass is an increasing function of density, and so in the absence of other physics the gas will not fragment. For an isothermal equation of state $\\gamma=1$, while for an adiabatic equation of state $\\gamma=5/3$.  A hard lower bound to the characteristic mass is given by the opacity limit \\citep{LLB,Rees}, expressed by the condition that the accretion luminosity of a contracting sphere of gas not exceed the black body luminosity of a sphere of the same radius. The critical mass in this case is just the Jeans mass, at the radius and density when the sphere becomes optically thick. This bound applies whether the protostellar core is formed by turbulence and transient, or if it self-gravitating. In either case the luminosity is given roughly by $L\\sim v^5/G$; the velocity is $v\\sim\\sqrt{GM/r}$ if the clump is contracting, and larger if it is transient.  Other theories associate the characteristic mass with the Jeans mass at some other density, e.g., with the density at which the gas first becomes thermally coupled to dust grains \\citep{larson73,larson05,jappsen05}. The argument is that the gas is nearly but not quite isothermal; at low density the temperature decreases with increasing density, but above the critical density temperature begins to increase with increasing density. This is modeled as a change in the effective adiabatic index $\\gamma$ of the gas, from less than unity to slightly more than unity; numerical experiments \\citep{li03} indicate that gas fragments readily when $\\gamma<1$, but less readily for $\\gamma>1$.  \\citet{2006MNRAS.368.1296B} perform numerical experiments, using an isothermal equation of state, that show that the mass of the break in the IMF is equal to the initial Jeans mass, and increases with increasing initial Jeans mass. When they employ a Larson style equation of state, they find a fixed characteristic mass in the resulting IMF, corresponding to the density where the adiabatic index suddenly increases from $\\gamma<1$ to $\\gamma>1$.  Competitive accretion \\citep{zinnecker82,bonnell01} schemes also predict that the characteristic mass is the mean Jeans mass, i.e., the Jeans mass using the mean density of the star forming region, and $T\\approx10$ K \\citep{1998ApJ...501L.205K,2006MNRAS.368.1296B}.  Yet other theories for the origin of the IMF start from the notion described above, that the observed fragmentation in giant molecular clouds arises from turbulence, e.g., \\cite{PN}. The turbulence establishes a powerlaw in clump mass at the high mass end, with an index similar to the Salpeter values. However, even in this work the {\\em characteristic mass} is related to the Jeans mass (calculated using the mean density) divided by the Alfvenic or fast magnetosonic Mach number ${\\cal M}_F$ to the $2/3$ power \\citet{PN}, their equation 30.  \\citet{elmegreen08} note that numerical simulations consistently show a proportionality between the characteristic mass and the thermal Jeans mass, and point out that the observed constancy of the characteristic mass then requires a constant Jeans mass, despite variations in the environments where stars from. We argue below that under conditions found in ultraluminous infrared galaxies (ULIRGs) and other massive galaxies neither the Jeans mass nor the characteristic mass is like that found in the Milky Way.  This paper is organized as follows. In \\S\\ref{sec: Cluster} we show how the Jeans mass $\\mj$ varies with cluster mass. When the star forming clumps have $\\rcl\\sim1\\pc$ and $\\mcl\\gtrsim10^6M_\\odot$, the finite optical depth to the far infrared radiation released by the contraction of the protocluster gas or by the dissipation of turbulent motion leads to an increase in temperature, and, we show, in the Jeans mass. The role of radiation pressure in setting the radius of very massive clusters is described in \\S\\ref{sec:radsupport}. In \\S\\ref{sec:compare} we compare our results to observations of massive young clusters, of ultra-compact dwarfs (UCDs) and of central massive objects such as those in \\citet{2002AJ....124.3073G} and \\citet{walcher05}. We give a short discussion of our results in the context of previous models for very massive star clusters in \\S\\ref{sec:discussion}, and offer our conclusions in \\S\\ref{sec:conclusions}.  In the appendices we discuss the log-normal distribution, the initial mass function (or IMF) that we use and the calculation of the associated light to mass ratio, we outline the calculation of the mass-radius relation for a radiation pressure supported cluster, and finally we discuss the dark matter density on $10\\pc$ scales predicted by recent numerical simulations. ",
        "conclusions": "\\label{sec:conclusions} We have shown that the Jeans mass in a cluster is roughly independent of the cluster mass as long as the cluster is optically thin to the FIR, and assuming that the mass radius relation seen in embedded clusters reflects the primordial mass radius relation. Clusters forming today in the Milky War are optically thin, so this finding is consistent with the observed constancy of the IMF in our galaxy.  Many Milky Way globular clusters were also optically thin, or had low enough velocities that their accretion luminosities would not have raised the gas temperature significantly. The Jeans mass in such clusters would then depend only on their radii at the time of formation.  However, we have pointed out that many clusters in other galaxies, and in the Milky Way in the past, are or were optically thick to the FIR; such clusters often (though not always) have high enough accretion luminosities that the gas temperature would have been higher than $10$K. We went on to argue that the Jeans mass is larger than a solar mass in such clusters. Using the assumption that the break in the IMF is associated with the Jeans mass, we concluded that massive clusters should have high $L/M$ ratios when young, and high $M/L$ ratios when more than a few Gyrs old.  The prediction of high $L/M$ is consistent with observations of apparently enhanced light to mass ratios in the superstar clusters M82-F and M82-11, although the optically thick clusters M82-9 and NGC 4038:W99-15 appear to have normal $M/L$ ratios. Other superstar clusters in M82 and other nearby galaxies are optically thin, and so they should have normal IMFs. Clearly, more data on superstar clusters would be helpful.  A top heavy IMF, which we showed should occur in the most massive and compact globular clusters, will tend to produce an excess of LMXBs and pulsars in metal rich globular clusters in both the Milky Way and in nearby galaxies, and at the same time, a higher $M/L$ ratio.  Compelling support for the notion of a high mass to light ratio in massive clusters is supplied by Figure \\ref{Fig:M_over_L vs M_cl}, which compares the predicted and observed mass to light ratio for objects with masses ranging up to $10^8M_\\odot$. As the solid lines in the Figure indicate, this result is consistent with the prediction of a top-heavy IMF resulting from the elevated temperatures associated with accretion in an optically thick environment.  We also argued that in the most massive clusters, with $\\mcl\\gtrsim 3\\times10^6M_\\odot$, the substantial luminosity associated with the contraction of the cluster must have been dynamically important. The predicted mass-radius relation, $\\rcl\\sim \\mcl^{3/5}$, leads to a number of relations between the properties of massive star clusters, involving the cluster velocity dispersion, the cluster luminosity, and the surface density; for example, we found $L_V\\sim \\sigma^4$. These relations appear to be consistent with the observed properties of massive star clusters.  The combination of strong evidence for a cluster-mass dependent mass to light ratio and for a mass-radius relation $\\rcl\\sim\\mcl^{3/5}$ provides solid support for the idea that a contraction powered radiation field has left its mark on the most massive star clusters we see around us.  It is fitting to close by emphasizing again the striking fact that globular clusters are observed to have radii of $2-3$ parsecs, with only weak dependence on cluster mass over a range $10^4M_\\odot$ to $3\\times10^6M_\\odot$. The origin of this (lack of a) mass-radius relation remains a major puzzle."
    },
    "0809/0809.0503_arXiv.txt": {
        "abstract": "We present two series of MOST (Microvariability \\& Oscillations of STars) space-based photometry, covering nearly continuously 10 days in 2004 and 30 days in 2007, of selected variable stars in the upper Main Sequence of the old open cluster M67. New high-precision light curves were obtained for the blue-straggler binary/triple systems AH~Cnc, ES~Cnc and EV~Cnc. The precision and phase coverage of ES~Cnc and EV~Cnc is by far superior to any previous observations. The light curve of ES~Cnc is modelled in detail, assuming two dark photospheric spots and Roche geometry. An analysis of the light curve of AH~Cnc indicates a low mass ratio ($q \\sim 0.13$) and a high inclination angle for this system. Two new long-period eclipsing binaries, GSC~814--323 and HD~75638 (non-members of M67) were discovered. We also present ground-based DDO spectroscopy of ES~Cnc and of the newly found eclipsing binaries. Especially interesting is HD~75638, a member of a visual binary, which must itself be a triple or a higher-multiplicity system. New light curves of two $\\delta$~Scuti pulsators, EX~Cnc and EW~Cnc, have been analyzed leading to detection of 26 and 8 pulsation frequencies of high temporal stability. ",
        "introduction": "\\label{intro}  M67 (NGC 2682) is one of the best studied open clusters. It is also one of the oldest known galactic clusters, with an estimated age of about 4 Gyr. Other parameters of the cluster are: $E(B-V) = 0.05$ and distance modulus $(m-M_V) = 9.59$ \\citep{mmj1993}. M67 contains a high percentage of blue stragglers, stars which are bluer and more luminous than the main sequence turnoff point of the cluster. For M67, the turnoff occurs near $B-V = 0.55$, and $V = 13.0$. The formation mechanism of blue straggles is still debated, but the most probable scenario is via mass transfer or mergers in close binaries \\citep{bail1995}. Several blue stragglers in M67 are known to be eclipsing binaries, which supports this hypothesis. The large number of close binaries in M67 is suggested by the relatively large number (25) of X-ray sources \\citep{bell1998}. In such an old cluster, the X-ray sources cannot be single rapidly-rotating stars, but -- more likely --  are tidally locked components of close binaries. In addition, long-term spectroscopic surveys of M67 \\citep{milo1992,lath1996} revealed among the blue stragglers, 5 spectroscopic binaries with orbital periods between 850 -- 5000 days.  In this paper, we present two time series of rapid, quasi-continuous photometry for selected targets in M67 obtained with the MOST (Microvariability \\& Oscillations of STars) satellite in 2004 and 2007. The targets were selected primarily to study two domains of cluster objects accessible to a small telescope like MOST: the upper Main Sequence (mostly blue stragglers) and red giants. Here we describe the results for the former group of objects, including two $\\delta$~Scuti, while a future paper (Kallinger et al., in preparation) will describe analyses and asteroseismology of the p-mode pulsations of the red giants in the cluster.  The main goals of the observations in the context of this paper were twofold: (i)~a search for new variables with very small amplitudes, and (ii)~obtaining high-quality light curves, particularly for blue stragglers with periods close to one day which are difficult to sample properly with ground-based observations. As we describe in Section~\\ref{observe}, the 2004 observations were experimental and short (compared to other MOST runs), while the scope of the 2007 observations was much expanded to include time series photometry of a large number of guide stars in the M67 field. The photometric observations of MOST have been augmented by medium-dispersion spectroscopy performed at the David Dunlap Observatory (DDO).  The paper is organized as follows: Section~\\ref{observe} presents MOST time-series photometry and spectroscopic observations performed at the DDO. The light variations are described in Section~\\ref{lightcurves} while a detailed discussion of individual binary systems known or found to be variable is given in Section~\\ref{results}. The $\\delta$~Scuti pulsating stars are discussed separately in Section~\\ref{deltascuti}. The paper is concluded with the summary of the results in Section~\\ref{sum}. ",
        "conclusions": "\\label{sum}  Two long, almost continuous photometric runs (10 days in 2004 and 30 days in 2007) of stars in the M67 field were obtained by the MOST space mission, followed by ground based spectroscopic observations at DDO. The targets include eclipsing blue stragglers in M67, $\\delta$~Scuti-type pulsating star, and variable non-cluster members in the cluster field.  Among the main results, our analysis of the light curve of the contact binary AH~Cnc supports its low mass ratio and a high inclination angle, as found by \\citet{sand2003a}. Contrary to a recent suggestion, there is no compelling evidence for a third component in the system. The MOST light curves of the blue straggler triple system ES~Cnc show a clear asymmetry which we interpret as the presence of two dark photospheric spots, about 1200~K cooler than the surrounding photosphere. A comparison of the 2004 and 2007 photometry reveals almost no changes in the fitted sizes and positions of the spots, suggesting a surprising stability in the photospheric features of the components over three years. However, our analysis of ES~Cnc is only preliminary. The system deserves a dedicated campaign of high-resolution spectroscopy and simultaneous high-precision photometry, but the proximity of its period to one day is a somewhat complicating factor.  Two new field eclipsing systems were discovered: HD~75638 and GSC~814-323. The former is a known spectroscopic triple. The eclipses are very shallow ($A \\sim 0.013$ mag) and would be barely detectable from the ground. The photometric period (the mean interval between the eclipses) of 5.81 days is compatible with the spectroscopic period of 5.8167 days found from observations spanning 3.68 years by \\citet{nord1997}. High-resolution spectroscopy of the system combined with our light curve could lead to reliable determination of masses of all three components. Spectra of GSC~814-323 show that it is a spectroscopic binary. Since the ratio of intensities of the components in the broadening functions is about 0.29, we expect the secondary eclipses to be 0.09 mag deep.  Period analysis of two known $\\delta$~Scuti variables, EX~Cnc and EW~Cnc yielded a very reliable frequency spectrum without any obvious instrumental periodicities. For the brighter of the two variables, EX~Cnc, we could reliably determine as many as 26 pulsational frequencies. In spite of the larger photometric scatter of about 0.02 mag for the fainter variable, EW~Cnc, we see 8 pulsation frequencies in the amplitude spectrum. Pulsational frequencies in both stars are temporary highly stable. The guide star photometry of other blue stragglers (Sanders 752, 968, 977, 1066, and 1466) was affected by the strong background variations which limited amplitudes of any detectable pulsations to $<0.001$ mag."
    },
    "0809/0809.0673_arXiv.txt": {
        "abstract": "Quasi-equilibrium models of rapidly rotating triaxially deformed stars are computed in general relativistic gravity, assuming a conformally flat spatial geometry (Isenberg-Wilson-Mathews formulation) and a polytropic equation of state. Highly deformed solutions are calculated on the initial slice covered by spherical coordinate grids, centered at the source, in all angular directions up to a large truncation radius. Constant rest mass sequences are calculated from nearly axisymmetric to maximally deformed triaxial configurations. Selected parameters are to model (proto-) neutron stars; the compactness is $\\compa = 0.001, 0.1, 0.14, 0.2$ for polytropic index $n=0.3$ and $\\compa=0.001, 0.1, 0.12, 0.14$ for $n=0.5$. We confirmed that the triaxial solutions exist for these parameters as in the case of Newtonian polytropes. However, it is also found that the triaxial sequences become shorter for higher compactness, and those may disappear at a certain large compactness for the $n=0.5$ case. In the scenario of the contraction of proto-neutron stars being subject to strong viscosity and rapid cooling, it is plausible that, once the viscosity driven secular instability sets in during the contraction, the proto-neutron stars are always maximally deformed triaxial configurations, as long as the compactness and the equation of state parameters allow such triaxial sequences. Detection of gravitational waves from such sources may be used as another probe for the nuclear equation of state. ",
        "introduction": "Rapidly rotating compact objects are expected to be formed as new born neutron stars after stellar core collapses, or as differentially rotating hypermassive neutron stars after binary neutron star mergers.  Accretion onto neutron stars in X-ray binaries can also lead to rapid rotation. All of these have been extensively studied as strong sources of gravitational waves for the ground based laser interferometers LIGO/GEO600/VIRGO/TAMA (See e.g. \\cite{reviews} and references therein).  Classical models of rotating stars are a class of ellipsoidal figures of equilibrium; self-gravitating rotating incompressible fluids in Newtonian gravity. Such solutions include sequences of axisymmetric Maclaurin ellipsoids, or non-axisymmetic Jacobi, Dedekind, or Riemann S-type ellipsoids \\cite{Ch69}. These models are used to study the secular evolutions of rapidly rotating stars due to the viscosity and the radiation back-reaction of gravitational waves \\cite{ellipsoidal}. Lai and Shapiro \\cite{LS95} have developed an ellipsoidal approximation to the rotating polytropes, and applied the model to clarify the secular evolution of rapidly rotating neutron stars in detail, and more recently focused on the viscosity driven secular instability \\cite{STU04}.  As discussed in \\cite{reviews,ellipsoidal,LS95}, and shown by a number of numerical simulations of rapidly rotating compact stars, core collapse, and binary neutron star mergers \\cite{SimDI,SimCC,SimBNS}, a transient triaxially deformed compact object may survive within a secular time scale. In this paper, we consider uniformly rotating models of such triaxially deformed compact objects, an extension of the Jacobi ellipsoid in general relativity. In Newtonian gravity, such solutions exist for rotating polytropes with polytropic index $n < 0.808$ \\cite{James64}.  Here, the polytropic equation of state (EOS) $p=\\kappa\\rho^{1+1/n}$ relates the pressure $p$ with the baryon rest mass density $\\rho$.  In general relativity, such configurations are not in equilibrium due to the back-reaction of gravitational radiation. However as in the case of quasi-equilibrium initial data of binary neutron stars, triaxially deformed uniformly rotating stars are in quasi-equilibrium, as long as the gravitational luminosity is small enough that the energy radiated away within a rotational period is small compared to the binding energy of the star, which is always the case, and if the viscosity is strong enough for the flow field to become uniformly rotating during the evolution. Therefore, as an important application, a sequence of uniformly rotating quasi-equilibrium solutions may model a secular evolution from the proto-neutron star to the neutron star in the strong viscosity limit, and each solution may serve as the initial data for the general relativistic hydrodynamic simulations of such objects.  Models of rapidly rotating neutron stars have been extensively studied as stationary, axisymmetric, perfect-fluid spacetimes \\cite{Nick03}, while less attention has been paid to uniformly rotating triaxial solutions, not only because they are not exact equilibria due to gravitational radiation reaction% , but also because, in early calculations \\cite{FIP86}, models of the EOS above nuclear density did not allow large enough values of $T/|W|\\sim 0.14$ where a triaxial sequence is expected to bifurcate from an axisymmetric sequence. Here, $T/|W|$ is the ratio of kinetic energy $T$ to gravitational potential energy $W$.  However, as seen, for example, in \\cite{Haensel}, a value $n\\sim 0.5$ (an effective adiabatic index $\\Gamma\\sim 3$) may be possible for recent models of EOS for high density nuclear matter, above $\\rho > \\rho_{\\rm nuc} \\sim 2\\times 10^{14} {\\rm g/cm}^3$, and, as we will see below, triaxial quasi-equilibrium solutions do exist even in strong gravity for relatively small polytropic indexes, such as $n=0.5$ or $n=0.3$, as in the Newtonian case.  About a decade ago, Nozawa succeeded in computing uniformly rotating triaxial polytropes in general relativistic gravity in his thesis \\cite{Nozawa}, although his calculations were limited by the computational resources to low resolutions. A few studies approximating the fluid as an ellipsoidal configuration in general relativistic gravity have been made \\cite{GRellpisoid}, and perturbative analyses have located the bifurcation point suggesting the existence of solutions having triaxial bar-mode deformations \\cite{BFGall,Dorota,Saijo06,SL96,YE97,SF98}. Our computations of triaxially deformed stars can be used to locate the instability points on the axisymmetric sequence as well as to estimate the gravitational wave amplitude and luminosity from such objects.  In this paper, we present our first results on triaxial configurations of rapidly rotating general relativistic stars as models of neutron stars in extreme rotation. We assume a conformally flat spatial slice, and solve the constraints and spatial trace of the Einstein equation (Isenberg-Wilson-Mathews (IWM) formulation) \\cite{ISEN78,WM89} \\footnote{The validity of the IWM formulation for axisymmetric configurations is discussed in \\cite{CST96,WM_TEST}}. This is different from the formulation used by Nozawa \\cite{Nozawa} in which the line element is chosen to be the same form as that of stationary axisymmetric spacetime, but an azimuthal dependence is allowed (see \\cite{usui99} for the same formulation). The formulations and the code are described in the next section. For testing the code, selected axisymmetric solutions are compared with the results in the literature, and the bifurcation points of axisymmetric and triaxial sequences in weak gravity are examined.  Then, we present the results of deformation sequences of constant rest mass systematically in the range of two parameters, the compactness $\\compa$ and the polytropic index $n$, appropriate for realistic neutron stars.  Applications of such triaxial solutions in the contraction of a newly born proto-neutron star are briefly discussed in the final section. Throughout the paper, we use units such that $G=c=1$.  For our tensor notation, we adopt the use of Greek letters for spacetime indices, and Latin letters for spatial indices. ",
        "conclusions": "As a result of massive stellar core collapses, proto-neutron stars are formed and contract to more compact neutron stars within the time scale of cooling of a few tens of seconds \\cite{protoNS}. Even for the small rotation rate of the collapsing stellar core, the  ratio $T/|W|$ of the proto-neutron star becomes much higher than the value where the axisymmetric solution becomes secularly unstable against the viscosity driven $\\ell=m=2$ bar mode instability \\cite{LS95,STU04}. Therefore, uniformly rotating triaxial solutions discussed in this paper may describe a quasi-stationary model of proto-neutron star contraction in the range of $\\compa \\sim 0.1-0.2$, assuming the following: (1) a certain mechanism of strong viscosity operates during the contraction, (2) the time scale rapid cooling is shorter than that of gravitational radiation reaction, (3) the effective polytropic (adiabatic) index $n$ $(\\Gamma=1+1/n)$ of the EOS for the realistic neutron star matter is small (large) enough to allow uniformly rotating triaxial solutions, and (4) those triaxial solutions are dynamically stable.  Such an evolutionary track of a proto-neutron star contraction has been considered using a compressible ellipsoidal model\\footnote{In their work, changes in the entropy during the evolution is modeled by the changes in the adiabatic constant $\\kappa$ of the polytropic EOS.} \\cite{STU04}. Our results add two further important features to this. First, the sequences of triaxial solutions terminate at the maximally deformed models, at the mass-shedding limits, and the changes in ADM mass or total angular momentum are small along the triaxial sequences from the bifurcation points to the termination points, even for the stiffer EOS such as $n=0.3$, as seen in Table \\ref{tab:n03n05all}; the triaxial sequences are not very long at all. Secondly, the triaxial sequences may become shorter and disappear as the compactness becomes larger for a relatively less stiff EOS such as $n=0.5$.  The angular velocity near the braking limit is estimated as $\\Omega M \\sim (M/R)^{3/2} \\sim (M/R)^{2}J/M^2$.  Therefore once the secular bar mode instability sets in, conserving $M$ and $J$, the proto-neutron star evolves towards a maximally deformed triaxial configuration as it contracts, say, from $M/R\\sim 0.1$ to $0.2$. And then, it is likely that {\\it it always evolves along the sequence of maximally deformed configurations during the contraction} as long as such triaxial solutions exist and are dynamically stable in the parameter region of the effective $\\Gamma$ and $\\compa$. Excess angular momentum arising from contraction may be transported outward by mass ejected from the Lagrange point at the cusp of the longest semi-major axis.  (Note that the time scale of the mass ejection may be that of cooling, which is much longer than the dynamical time scale.) Furthermore, if the effective $\\Gamma$ satisfies $\\Gamma \\alt 3$, the triaxial solution may disappear when the solution reaches a certain value of the compactness $\\compa$ and higher. Dynamically stability of such uniformly rotating solutions are not known, but it is unlikely that the dynamical instability appears within such short triaxial sequences, along which the ratio $T/|W|$ is nearly constant.  It is estimated that the amplitude of gravitational wave (GW) signals from such objects may be detectable using the ground based laser interferometric detectors, if the source is within a few tens of Mpc \\cite{LS95,Saijo06}.  Detection of the persistent GW signals even after the proto-neutron star contraction phase suggests a large effective $\\Gamma \\agt 3$, while the shutdown of the signal during the contraction implies the relatively smaller $\\Gamma \\alt 3$. Detection of such GW signal may set another constraint on the EOS parameter of high density matter.  Source modeling for constructing the wave templates may be straightforward because one can concentrate on calculating the maximally deformed configurations. Our next plan is to include more realistic nuclear EOS in the code, then to estimate the gravitational wave amplitude for those EOSs that allow the triaxial solutions."
    },
    "0809/0809.0390_arXiv.txt": {
        "abstract": "{The application of the virial theorem provides a tool to estimate supermassive black hole (BH) masses in large samples of active galactic nuclei (AGN) with broad emission lines at all redshifts and luminosities, if the broad line region (BLR) is gravitationally bound. In this paper we discuss the importance of radiation forces on BLR clouds arising from the deposition of momentum by ionizing photons. Such radiation forces counteract gravitational ones and, if not taken into account, BH masses can be severely underestimated. We provide virial relations corrected for the effect of radiation pressure and we discuss their physical meaning and application. If these corrections to virial masses, calibrated with low luminosity objects, are extrapolated to high luminosities then the BLRs of most quasars might be gravitationally unbound. The importance of radiation forces in high luminosity objects must be thoroughly investigated to assess the reliability of quasar BH masses. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0359_arXiv.txt": {
        "abstract": "{The origin of magnetism around asymptotic giant branch (AGB) stars is still uncertain. These stars may drive an important dynamo, but if the magnetic energy dissipates all in X-rays, the observed X-ray luminosities are too low for a strong, dynamically important global field. Other explanations for the circular polarization of SiO masers in the AGB atmospheres may thus be needed.} {The interaction of the AGB wind with previously ejected matter and with planets is expected to bear complex magneto-hydrodynamic phenomena on a short time scale, such that strong magnetic fields can be maintained locally. Here we provide observational evidence for the corresponding magnetic fluctuations.} {We use the circular polarization of the $v=1,J=2-1$ SiO masers as tracer for magnetic activity. A correlation polarimeter allows us to simultaneously record all Stokes parameters. An SiO maser survey of 80 AGB stars was performed, out of which eight sources with the strongest circular polarization were selected for further monitoring.} {In two AGB stars, V~Cam and R~Leo, we find evidence of pseudo-periodic fluctuations of the fractional circular polarization on a timescale of a few hours, from which we infer magnetic fluctuations of $\\sim 1$\\,G. The phenomenon is rare and, if detected in an SiO star, restricted to a narrow range of velocities. It seems to be associated with planetary wake flows suggested by VLBI maps.} {While scenarios involving magnetic activity in the extended stellar atmosphere have problems to explain all observed features, precessing Jovian magnetospheres predict all of them without difficulty. For the case of R\\,Leo, we constrain the orbit of the planet (estimated period 5.2\\,years), derive a stellar mass estimate of $0.7\\,M_{\\sun}$ from it, and discuss the impact of planetary magnetism on the survival of planets. Smooth velocity variations of the fluctuating circular polarization feature are predicted as the planet moves along its orbit.} ",
        "introduction": "Towards the end of their lifetime, low- to intermediate mass stars undergo a phase in which they burn Helium in a shell on top of a Carbon-Oxygen core, and Hydrogen in another shell above the Helium shell. During this phase, when the stars appear in the Hertzsprung-Russell diagram in the upper asymptotic giant branch (AGB), they lose, due to pulsations, at least half of their mass, which forms a circumstellar envelope (see Herwig, 2005, for a review). The inner part of this envelope, also called extended atmosphere, is expected to bear complex magneto-hydrodynamic phenomena, due to the interaction of the wind with the previously ejected matter and with planets (Struck et al., 2004, 2002, Struck-Marcell, 1988). As in the solar system, where space weather changes on timescales of hours (e.g. Prang\\'e et al., 2004), fluctuations of the magnetic field about a mean value can be expected due to this interaction, but the observational evidence is still lacking. Here we will show that for a narrow range of velocities the circular polarization of SiO masers, generally accepted as a tracer of the magnetic field in the extended atmosphere of AGB stars (Nedoluha \\& Watson, 1994, Elitzur, 1996), varies in two AGB stars with a period of a few hours. Previous multi-epoch observations (Diamond et al., 2003, Pardo et al., 2004, Kang et al., 2006) of SiO masers were not polarimetric, while the sampling of the polarization variability study of Glenn et al. (2003) was not dense enough to detect such intraday magnetic fluctuations.  Our knowledge about magnetism within the extended atmosphere of AGB stars relies on circular polarization measurements of SiO masers at 1.5 to 7 AU distance from a star with a typical radius of 0.7 to 1\\,AU. Maser theory implies molecular excitation in dense pockets of gas and amplification of light in narrow tubes. Assuming that the circular polarization is caused by the Zeeman effect in the non-paramagnetic SiO molecules in the tubes, the reported fractional circular polarizations of up to 9\\,\\% for the $v=1, J=1-0$ transition ($v$ and $J$ are the vibrational and rotational quantum numbers, respectively) yield line-of-sight averaged magnetic flux densities of up to 100\\,G (Barvainis et al., 1987). More recent observations (Herpin et al., 2006) of the $v=1, J=2-1$ masers (excited in gas layers very close to those where $v=1, J=1-0$ is excited) also show the presence of circular and linear polarizations in many stars. These observations result in lower magnetic flux density estimates in the range 1 to 15\\,G, again assuming the circular polarization is due to the Zeeman effect. In the outer envelope, H$_2$O and OH masers reside at $\\sim 100$ to 400\\,AU, respectively $\\sim 1000$ to $\\sim 10000$ AU, from the star. The Zeeman effect of H$_2$O masers has been used (Vlemmings et al., 2002) to estimate, by extrapolation assuming a solar-type field topology, a magnetic flux density at the stellar surface of 100\\,G. In summary, it is now observationally evident that AGB stars maintain a magnetic flux throughout their envelope. Blackman et al. (2001) showed that AGB stars can generate a magnetic field via dynamo action at the interface between the rapidly rotating core and the extended convection zone. The importance of the magnetic field for global envelope dynamics is still a matter of debate, though (e.g. Soker, 2006).  Once the SiO maser spots are formed (often distributed along arcs around the star, hereafter referred to as maser shell, see e.g. Cotton et al., 2008), they are subject to magnetospheric events which are known to be rapidly variable if we refer to the solar system. Here we provide first evidence for such fluctuations in the atmospheres of two AGB stars by frequent polarization sampling. The polarization monitoring by Glenn et al. (2003) was only sensitive to slow variations and therefore to a long-term readjustment of the magnetic field (on a timescale of several months). ",
        "conclusions": "In summary, in our opinion none of the non-planetary alternatives satisfactorily explains the present observations of R\\,Leo and V\\,Cam, while a single scenario, namely precessing Jovian-type magnetospheres, are consistent with all of them. Since the discovery of a planet around a star in the post-red giant phase (Silvotti et al., 2007), we know that planets can survive in such an environment. Whether a given planet survives the AGB phase of its parent star depends on several factors (Villaver \\& Livio, 2007). Evolution of the orbit is the main one - depending on its distance from the star, the planet is either dragged inwards by tidal friction and evaporated, or repelled due to stellar mass loss, while for Jupiter-sized planets the gain of mass due to Roche flow from the star or accretion from the stellar wind barely affect their orbits. If the suggested planet around R~Leo, one of the best studied Mira stars, of solar mass and at 113.5\\,pc distance (Fedele et al., 2005), does not spiral in and evaporate during the remainder of the upper AGB phase, it will survive the subsequent planetary nebula stage against evaporation (Villaver \\& Livio, 2007) only if it has at least two Jupiter masses and orbits at least at $\\sim 3$\\,AU. We note that the fluctuating SiO maser feature in R\\,Leo has an orbital radius of 2.7\\,AU, close to the critical distance. We have demonstrated that the orbital elements of its host planet can be determined by measuring the velocity of the fluctuation and the position of the corresponding maser spot. The orbital period of the planet is 5.2\\,years, which provides an estimate of the stellar mass of $0.7\\,M_{\\sun}$. This agrees with mass estimates found in the literature, of about $1\\,M_{\\sun}$ for R\\,Leo. On the long term, our method has the potential to reveal, together with an independent estimate of the stellar mass, sub-Keplerian orbital motion and thus migration of the planet through the atmosphere towards the star. This requires correspondingly accurate observations. The situation may be complicated by the presence of several planets, whose orbits could be constrained by our method provided that the polarization fluctuations are not affected by a spectral blend of independent features at the same radial velocity and that the VLBI observations are free from a spatial blend of independent features along the same line of sight. The estimated mass loss of R\\,Leo, $\\sim 10^{-7}\\,M_{\\sun}/{\\rm y}$, is too weak to measure orbit evolution within a reasonable amount of time ($\\sim 3\\,\\mu$as in 1000 years). The situation may change if the opposite effect due to tidal friction, neglected in this estimate, is by several orders of magnitude more important than that of stellar mass loss.  Without further observational evidence at hand, one can only speculate on whether the suggested Jovian magnetic fields are decisive for evaporation or survival of the planets. They suppress turbulent viscosity and therefore the tidal friction induced by the latter (Zahn, 1977), and also lead to spin angular momentum loss due to magnetic braking. This contrasts with scenarios where the coupling between the stellar magnetic field and the planetary one lead to an increased drag area and therefore infall of the planet as compared to the case of an unmagnetized planet. The finding that our observations most likely imply a planetary magnetopause may turn out to be important here and motivate studies of planetary evolution in late-type stars that include magnetic star-planet interaction (omitted by Villaver \\& Livio, 2007, but included by Nordhaus \\& Blackman, 2006 who predict stellar outflows). The case for exoplanetary magnetism proposed here can only be strengthened by a combination of extensive simultaneous polarization monitoring and repeated spatially resolved observations. The experimentum crucis for an unambiguous evidence of planets will be the detection of orbital motion in the proper motion of maser spots showing the magnetic fluctuations, by contemporaneous direct imaging of the suggested wake flows at the same velocity as the fluctuations. Such a double evidence would rule out ambiguities: there may be planets without wake flows, if the wind density is not sufficient to produce observable features, and there are radial alignments of maser spots not related to wake flows. The coincidence of a persisting pseudo-periodic magnetic fluctuation with a radial alignment of maser spots strongly suggests that these phenomena are associated and naturally explained by a precessing Jovian magnetosphere and a planetary wake flow, respectively."
    },
    "0809/0809.5260_arXiv.txt": {
        "abstract": "We report high spatial resolution 11.2 $\\mu$m and 18.1 $\\mu$m imaging of the eruptive variable V838 Monocerotis, obtained with Gemini Observatory's Michelle in 2007 March.  The 2007 flux density of the unresolved stellar core is roughly 2 times brighter than that observed in 2004.  We interpret these data as evidence that V838 Mon has experienced a new circumstellar dust creation event.  We also report a gap of spatially extended thermal emission over radial distances of 1860-93000 AU from the central source,  which suggests that no prior significant circumstellar dust production events have occurred within the past $\\sim$900-1500 years. ",
        "introduction": "V838 Monocerotis experienced three dramatic photometric outbursts in 2002 January-March in which it brightened from v $\\sim$15 to v $\\sim$6.6.  The star displayed a F-type spectrum characterized by strong P-Cygni profiled H$\\alpha$ emission shortly after outburst (see e.g. Wisniewski et al 2003a), but rapidly cooled to a L-type supergiant by 2002 October (Evans et al 2003).  A spectacular light emerged around the system in 2002 February, and the time evolution of scattered light from this echo has been traced in exquisite detail by the Hubble Space Telescope Advanced Camera for Surveys (see e.g. Bond et al. 2003).  Spatially resolved thermal emission from the light echo material was also detected by the Spitzer Space Telescope's MIPS (Banerjee et al. 2006).  The origin of the light echo material (circumstellar versus interstellar) is not well known, but the 2004 epoch detection of the echo in thermal emission enabled Banerjee et al (2006) to place initial constraints on the dust mass of the material.  The large mass inferred led Banerjee et al (2006) to suggest that at least some of the dust might be interstellar in nature.  \\subsection{Key Questions}  While a multitude of research groups have observed and monitored V838 Mon over a wide-range of wavelengths and using a wide-variety of diagnostic techniques, several fundamental questions regarding the object remain unanswered.  These include:  \\begin{itemize} \\item \\textit{What caused V838 Mon's outbursts?}  A diverse array of mechanisms have been proposed to explain V838 Mon's behavior, as summarized in the proceedings of the 2006 conference ``The Nature of V838 Mon and Its Light Echo'' (ASP Conf Proc Vol 324).  These mechanisms can be classified into two groups: events which are singular in nature (i.e. a stellar merger) and events which are recurrent (i.e. a nova-like explosion).  \\item \\textit{How much of V838 Mon's light echo material is remnant circumstellar material and how much is interstellar material?} \\end{itemize} ",
        "conclusions": ""
    },
    "0809/0809.2676_arXiv.txt": {
        "abstract": "It is commonly assumed that reaction measurements for astrophysics should be preferably performed in the direction of positive $Q$ value to minimize the impact of the stellar enhancement factor, i.e.\\ the difference between the laboratory rate and the actual stellar rate. We show that the stellar effects can be minimized in the charged particle channel, even when the reaction $Q$ value is negative. As a demonstration, the cross section of the astrophysically relevant $^{85}$Rb(p,n)$^{85}$Sr reaction has been measured by activation between $2.16\\leq E_\\mathrm{c.m.}\\leq 3.96$ MeV and the astrophysical reaction rate for (p,n) as well as (n,p) is directly inferred from the data. The presented arguments are also relevant for other $\\alpha$- and proton-induced reactions in the $p$ and $rp$ processes. Additionally, our results confirm a previously derived modification of a global optical proton potential. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.5056_arXiv.txt": {
        "abstract": "The recently discovered population of ultra-faint extended line emitters, with fluxes of a few times $ 10^{-18} \\ergs \\cm$ at $z\\sim 3$, can account for the majority of the incidence rate of Damped \\lya systems (DLAs) at this redshift if the line emission is interpreted as \\lya. We show here that a model similar to that proposed by \\citet{2000ApJ...534..594H}, which explains the incidence rate and kinematics of DLAs in the context of $\\Lambda$CDM models for structure formation, also reproduces the size distribution of the new population of faint \\lya emitters for plausible parameters. This lends further support to the interpretation of the emission as \\lya, as well as the identification of the emitters with the hitherto elusive population of DLA host galaxies. The observed incidence rate of DLAs together with the observed space density and size distribution of the emitters suggest a duty cycle of $\\sim 0.2-0.4$ for the \\lya emission from DLA host galaxies. We further show that \\lya cooling is expected to contribute little to the \\lya emission for the majority of emitters. This leaves centrally concentrated star formation at a rate of a few tenths \\Msol yr$^{-1}$, surrounded by extended \\lya halos with radii up to 30-50 kpc, as the most plausible explanation for the origin of the emission. Both the luminosity function of \\lya emission and the velocity width distribution of low ionization absorption require that galaxies inside Dark Matter (DM) halos with virial velocities $\\la 50-70 \\kmsec$ contribute little to the incidence rate of DLAs at $z\\sim 3$, suggesting that energy and momentum input due to star formation efficiently removes gas from these halos. Galaxies with DM halos with virial velocities of $100-150 \\kmsec$ appear to account for the majority of DLA host galaxies. DLA host galaxies at $z\\sim 3$ should thus become the building blocks of typical present-day galaxies like our Milky Way. ",
        "introduction": "Quasar absorption spectra provide excellent probes of the distribution of baryons in the high-redshift Universe. Damped Lyman Alpha systems (DLAs, historically defined as having a neutral hydrogen column density $N_{\\ro{HI}} > 2 \\times 10^{20}$cm$^{-2}$) are particularly useful as they are likely to play an important role as a reservoir of gas for the formation of stars and galaxies at high redshift. They dominate the neutral gas content of the Universe between $z \\sim 0-5$, and at $ z \\sim 3.0-4.5$ their neutral gas content is comparable to visible stellar mass in present-day galaxies \\citep{1986RSPTA.320..503W,1996MNRAS.283L..79S,2000ApJ...543..552S}. DLAs thus form an important link between primordial plasma and the stellar structures that form from it.  In spite of observations of over 1000 DLAs, \\citet{2005ARA&A..43..861W} still conclude that the question ``what is a damped \\lya system?'' has not yet been answered conclusively. One of the reasons for this is that absorption spectra provide only indirect information in velocity space, and only probe the gas along one line-of-sight through the galaxy. The \\lya absorption feature itself provides no information about the velocity structure of the DLA, because of the large optical depth even in the damping wings. This has led observers to look at the associated low ionization metal absorption features. Low ionization species like SiII, CII and FeI are believed to be good tracers of the neutral gas in DLAs. \\citet{1997ApJ...487...73P} developed the velocity width distribution of metal absorption features into an important diagnostic tool for DLAs. The much lower absorption optical depth allows us to extract detailed kinematical information for the gas in DLAs. The absorption profiles are normally clumpy and asymmetric, with the strongest absorption feature often occurring at one edge of the profile. Velocity widths range from 30$\\kmsec$ up to several hundred $\\kmsec$. Note, however, that there is a distinctive lack of very narrow absorption profiles.  Kinematical models aiming to reproduce the velocity width data fall into two categories. \\citet{1986RSPTA.320..503W} suggested a close connection between DLAs and disks of present-day spiral galaxies. \\citet{1997ApJ...487...73P} modelled DLAs as thick rotating disks with a rotation speed ($\\sim 200\\kmsec$) typical of present-day galaxies. \\citet{1998ApJ...495..647H} challenged this interpretation and demonstrated that it is not unique. The merging of proto-galactic clumps expected in Cold Dark Matter (CDM)-like models for structure formation can explain the shape of the profiles equally well. Galactic winds are also likely to play an important role \\citep[e.g.][]{1999MNRAS.308L..39F}. They further showed that for the same virial velocity width, merging clumps produce significantly larger velocity width and argued that the latter interpretation is favoured by the observed velocity width distribution in the context of the CDM paradigm for structure formation \\citep{2000ApJ...534..594H}.  The level of enrichment with metals provides another important clue as to the nature of DLAs --- see \\citet{2004cmpe.conf..257P, 2006fdg..conf..319P,2005ARA&A..43..861W} for reviews. Significant metal absorption is found in all DLAs, though as a population they are metal poor. DLAs are also relatively dust free \\citep{2001A&A...379..393E,2006MNRAS.367..211W}. Initially, the low metallicity was used to argue that DLA host galaxies are the chemically unevolved but otherwise very similar counter-parts of typical present-day spiral galaxies \\citep{1986RSPTA.320..503W}. In the model of \\citet{1998ApJ...495..647H}, DLAs instead preferentially probe the outer parts of much less massive galaxies, many of which end up as building blocks of typical present-day galaxies that form by hierarchical merging in CDM-like models for structure formation. Recently, \\citet{2008arXiv0804.4474P} demonstrated that such a model fits the observed metallicity of DLAs at $z\\sim 3$ very well.   Models of DLAs invariably predict more than quasar absorption features can in principle test. As already mentioned, absorption is measured along a single line of sight and thus is a one-dimensional probe of the properties of the DLA. To explore the spatial extent and structure of the DLA, we need to observe DLAs in emission. Attempts to do this have focussed on both line and continuum emission. \\lya emission holds great potential in this respect. Star-formation, cooling radiation and fluorescent re-emission of the meta-galactic UV background are all expected to contribute at different levels. In addition, stellar continuum emission from the newly formed stars should also be bright enough to be detectable.  Many observers have attempted to find the galaxy counterparts of DLAs at high redshift in emission by searching adjacent to quasar sightlines with known absorption systems \\citep{1999MNRAS.305..849F,1999MNRAS.309..875B,2000ApJ...536...36K, 2001ApJ...551...37K,2001MNRAS.326..759W,2007A&A...468..587C}. This is a difficult task, as the light of the extremely bright quasar must be accurately subtracted to study the light of the galaxy, which is very faint in comparison. \\citet{2000ApJ...536...36K} and others caution of the possibility that a given emission feature is a Point Source Function (PSF) artefact rather than a real source. \\citet{2007A&A...468..587C} report that, for $z > 2$, six DLA galaxies have been confirmed through spectroscopic observation of \\lya emission, with other techniques producing a few additional candidates. \\citeauthor{2007A&A...468..587C} added another six \\lya emission candidates to this group. A quantitative statistical interpretation of the many (largely unsuccessful) searches is difficult if not impossible, but the rather low success rate appears to be consistent with their interpretation as galaxies of rather low mass and star formation rate.  \\citet{2008ApJ...681..856R} reported the results of a long-slit search for low surface brightness \\lya emitters at redshift $z \\sim 3$, which reached flux levels that are about a 10 times lower than previous \\lya surveys at this redshift. They found 27 faint line emitters, many of which are extended in wavelength and real space. They argue that the majority of the emitters are likely to be \\lya (rather than low-$z$ interlopers), powered by a central region of star formation and processed by radiative transfer through surrounding neutral hydrogen. \\citet{2008ApJ...681..856R} note that the incidence rate inferred from the space density and the size distribution of the emitters is similar to that of DLAs, and suggest that they are the host population of DLAs and high column density Lyman Limit Systems (LLS). If this is indeed the case then the observations of \\citet{2008ApJ...681..856R} give us --- for the first time --- the size distribution and space density of DLA host galaxies \\citep[Figure 19 of ][]{2008ApJ...681..856R}.  Models predicting the velocity width and size distribution of DLAs have been constructed based on the observed luminosity function of galaxies \\citep{1999MNRAS.305..849F,2008ApJ...683..321F} and on the Press-Schechter formalism \\citep{1974ApJ...187..425P} in conjunction with numerical simulations \\citep{1997ApJ...484...31G,1997ApJ...486...42G,1998ApJ...495..647H, 2000ApJ...534..594H,2004MNRAS.348..421N,2007ApJ...660..945N}. More recently numerical simulations have attempted to model the entire DLA population self-consistently \\citep{2006ApJ...645...55R,2007arXiv0710.4137R,2008arXiv0804.4474P}.  In this paper, we revisit models for the absorption properties of DLAs in the light of the size distribution data of \\citet{2008ApJ...681..856R}, improved data on the velocity width distribution from metal absorption lines (which has presented a challenge to purely numerical simulations --- see \\citet{2007arXiv0710.4137R}) and numerical simulations with increased resolution, box size and sophistication incorporating the additional physics of radiative transfer, gas chemistry and star formation.  Throughout this work we use cosmological parameters of the 5 year WMAP data \\citep{2008arXiv0803.0732H}: $(h, \\Omega_M, \\Omega_b, \\Omega_{\\Lambda}, \\sigma_8, n) = (0.701, 0.279, 0.046, 0.721, 0.817, 0.96)$. ",
        "conclusions": "We have considered an updated version of the Haehnelt et al. model for the kinematics of DLAs, in light of the discovery of a new population of extended low surface brightness \\lya emitters with a total inferred incidence rate similar to that of DLAs. The main differences with the modelling of \\citet{2000ApJ...534..594H} are the use of the \\st modification to the \\ps formalism, an update of cosmological parameters, and the use of an exponential suppression of the cross section for damped absorption for low virial velocities instead of an sharp cut-off.  Our main results are the following.  \\begin{itemize}  \\item{The observed velocity width distribution of the associated metal absorption can be fit with a model where the cross section for damped absorption scales with the virial velocity of the halo as $\\sigma \\propto v_\\ro{c}^{2.5}$, the absorption cross section is suppressed in halos with $ v_\\ro{c}\\le 70\\kmsec$, and the conditional velocity width is given by that in the simulations of Haehnelt et al. (or the very similar distribution of the simulations of \\citealt{2008arXiv0804.4474P}).}  \\item{The same model can fit the cumulative incidence rate for absorption inferred from the size distribution of the \\lya emitters if the \\lya emission has a duty cycle of $f_\\ro{d} = 0.2-0.4$, and the emission extends over an area which is larger than the cross section for damped \\lya absorption by a factor $f_\\ro{d}^{-1} = 2.5-5$.}  \\item{The maximum expected \\lya cooling luminosity due to collapsing gas in DM halos falls short of the observed \\lya luminosities by a factor of three to five, even for optimistic assumptions regarding the expected emission due to \\lya cooling. Furthermore, the expected dependence of the \\lya cooling luminosity on the virial velocity of DM halos thereby maps into a luminosity function with a slope shallower than observed. \\lya cooling radiation should thus not contribute significantly to the \\lya emission for the majority of the objects, especially at the faint end of the \\lya luminosity function.}  \\item{The \\lya luminosity function is well fit by a simple model where the \\lya luminosity scales linearly with the mass of the DM halo and the emission is suppressed for low mass DM halos. The \\lya luminosity function can be fit for a wide range of duty cycles including the duty cycle required to simultaneously explain the kinematic properties of DLAs and the cumulative incidence rate inferred from the observed size distribution of the \\lya emitters.}  \\item{Our model predicts that the bulk of the contribution to the incidence rate of DLAs comes from absorption systems hosted in DM halos with virial velocities in the range from 50-200 $\\kmsec$ and masses in the range $10^{10}$ to $10^{12} \\Msol$. The cut-off for damped absorption occurs at somewhat higher virial velocities than suggested by numerical simulations, which attempt to simulate the gas distribution and kinematics of DLAs self-consistently. These simulations, however, fall short of reproducing the observed velocity widths of the associate low-ionization metal absorption of DLAs. If the suppression of the cross section for damped \\lya absorption in halos with virial velocities up to $70 \\kmsec$ is indeed real, then feedback due to star formation at high redshift has to be more efficient in removing gas --- even from rather deep potential wells --- than is assumed in most models of galaxies formation and numerical simulations. Alternatively the simulation (and our model) may underestimate the effect of stellar feedback on the velocity width of absorbers hosted by DM halos with small virial velocities. This may not be implausible given the fact that the simulations still fail to reproduce realistic galactic winds.} \\end{itemize}   Our modelling here adds to the mounting evidence that DLAs and LLs predominantly probe the outer regions of dwarfish\\footnote{``In English the only correct plural of dwarf is dwarfs, and the adjective is dwarfish \\ldots dwarves and dwarvish are used \\ldots only when speaking of the ancient people to whom Thorin Oakenshield and his companions belonged.'' \\citet{Tolkien}.} galaxies. This picture had been suggested for a while by the kinematic and metallicity data. With the discovery of a faint population of \\lya emitters, most plausibly identified with the population of DLA/LL host galaxies, we finally have a handle on their space density, sizes and (in conjunction with models like the one presented here) masses and virial velocities. In the picture that emerges, DLAs are hosted by the galaxies that, in the context of the now well established $\\Lambda$CDM paradigm for structure formation, are expected to become the building blocks of typical galaxies like our Milky Way."
    },
    "0809/0809.3247_arXiv.txt": {
        "abstract": "{} {We use VLT/UVES high-resolution optical spectroscopy of seven GRB afterglows (at \\zgrb\\ between 2.20 and 3.97) to investigate circumburst and interstellar plasma in the host galaxies of the bursts.} {GRBs 050730, 050820, 050922C, 060607, 071031, and 080310 were each detected by the \\emph{Swift} satellite. With the minimum possible time delay (as short as eight minutes), follow-up optical spectroscopy at 6.0~\\kms\\ resolution began with UVES. We present Voigt-profile component fits and analysis of the high-ion absorption detected in these spectra and in the pre-\\emph{Swift} UVES spectrum of GRB~021004, at velocities between \\zgrb--5\\,000~\\kms\\ and \\zgrb.} { We identify several distinct categories of high-ion absorption at velocities close to \\zgrb: (i) Strong high-ion absorption at \\zgrb\\ itself is always seen in \\os, \\cf, and \\sif, usually (in six of seven cases) in \\nf, and occasionally in \\sfo\\ and \\ssix. Three of the cases show log\\,$N$(\\nf)$>$14 in a single-component, suggesting a circumburst origin, but we cannot rule out an interstellar origin. Indeed, using the non-detection of \\sfo$^*$ at \\zgrb\\ toward GRB~050730 together with a UV photo-excitation model, we place a lower limit of 400~pc on the distance of the \\sfo-bearing gas from the GRB. (ii) Complex, multi-component \\cf\\ and \\sif\\ profiles extending over 100--400~\\kms\\ around \\zgrb\\ are observed in each spectrum; these velocity fields are similar to those measured in \\cf\\ in damped \\lya\\ systems at similar redshifts, suggesting a galactic origin. (iii) Asymmetric, blueshifted, absorption-line wings covering 65--140~\\kms\\ are seen in the \\cf, \\sif, and \\os\\ profiles in four of the seven GRB afterglow spectra. The wing kinematics together with the ``Galactic'' \\cf/\\sif\\ ratios measured in two cases suggest that the wings trace outflowing interstellar gas in the GRB host galaxies. (iv) High-velocity (HV; 500--5\\,000~\\kms\\ relative to \\zgrb) components are detected in six of the seven spectra; these components are not necessarily a single homogeneous population, and many of them may arise in unrelated foreground galaxies. However, in the cases of GRBs~071031 and 080310, the ionization properties of the HV components (very high \\cf/\\sif\\ ratios and absence of neutral-phase absorption in \\siw\\ or \\cw) are suggestive of a circumburst origin; models of Wolf-Rayet winds from the GRB progenitors can explain both the kinematics and ionization level of these HV components. }{} ",
        "introduction": "For many years the only backlights available for studying the distant Universe through absorption-line spectroscopy were AGN. Since the discovery that gamma-ray bursts (GRBs) are extragalactic \\citep{vP97, Me97}, GRBs have also been employed as backlights, with the advantages of extremely high luminosity and power-law continua, but the disadvantage of rapid temporal fading, implying that rapid-response observations are necessary for their study. Not only do analyses of GRB afterglow spectra allow intervening systems to be studied, they also enable the properties of the interstellar gas in the host galaxy of the GRB to be investigated, particularly its kinematics, chemical abundances, dust content, molecular gas content, and ionized gas content \\citep{Ca03, Kl04, Fi05, Sv03, Sv04, Pe06, Fy06, Pr06, Pr07a}. Many GRB spectra show damped \\lya\\ (DLA) absorption at \\zgrb, i.e. they show a neutral gas column density log\\,$N$(\\hi)$>$20.3 \\citep{Je01, Hj03, Ja04, Ja06, Vr04, St05, Be06, Wa06, Pr07b}, typically with higher \\hi\\ column densities than are seen in DLAs toward QSOs \\citep[QSO-DLAs, reviewed by][]{Wo05}. The higher $N$(\\hi) in GRB-DLAs may be related to their intense star-formation activity \\citep{Bl02}, or to the GRB sight-lines passing preferentially through the inner regions of the host galaxies \\citep{Fy08, Pr08a}. Photometric studies of GRB host galaxies have shown them to be star-forming, dwarf galaxies \\citep{Ch04, Wi07, Th07, Th08a}, with no evidence that the hosts are peculiar \\citep{Sv08}.  In this paper we use rapid-response mode spectra of seven high-redshift long-duration GRB afterglows to investigate high-ion absorption (in \\os, \\nf, \\cf, \\sif, \\sfo, and \\ssix) in the ISM of the host galaxy, as well as in the circumburst medium immediately surrounding the GRB. Observations of ionized interstellar gas in galactic environments provide important constraints on several physical processes, including star formation and subsequent feedback, galactic winds, interactions with satellite galaxies, and accretion. In the pre-\\emph{Swift} era, studying high ions in galaxy environments at high redshifts was complicated by the difficulty of associating individual absorbers with individual galaxies. There are several \\emph{indirect} ways to explore high-ion absorption in high-redshift galaxies. Firstly, outflowing \\cf\\ and \\sif\\ is seen in the composite spectrum of high-redshift Lyman-break galaxies \\citep[LBGs;][]{Sh03}. Secondly, there is a statistical correlation between the redshifts of \\cf\\ absorbers seen in QSO spectra and the redshifts of galaxies near the QSO lines-of-sight \\citep{Ad05}. And thirdly, high-ion absorption is detected in all DLAs \\citep{WP00, Fo07a, Fo07b, Le08}, which trace galactic structures. However, because of the faintness of high-redshift galaxies, only in a single published case \\citep[the lensed LBG cB58;][]{Pe00, Pe02} are interstellar high-ion absorption lines seen at moderate-to-high resolution in the spectra of \\emph{individual} galaxies at $z\\!>\\!2$.  Now, with the ability to be on-target within a few minutes of a \\emph{Swift} GRB trigger, a new means to directly probe the interstellar (as well as circumburst) medium in individual high-redshift galaxies is available \\citep[see][]{Fi05, DE07, Ch07}. High-resolution optical spectroscopy of GRB afterglows now routinely produces spectra with high-enough signal-to-noise to uncover the velocity sub-structure in the absorbing gas. Evidence has even been found for time-variation in the strength of absorption in fine-structure lines at \\zgrb\\ \\citep{DZ06, Vr07}. In this paper, we present a survey of high-ion absorption in the highest-quality GRB optical afterglow spectra taken to date. We pay particular attention to the high-ion kinematics, since velocity measurements can constrain the presence of gaseous outflows, either from the GRB progenitor or from the host galaxy. We structure this paper as follows. In \\S2 we discuss the GRB observations, data processing, and our line profile fitting procedures. In \\S3 we describe the observational properties of the high-ion absorption at velocities near \\zgrb\\ in each afterglow spectrum. We discuss four distinct observational categories of high-ion absorption in the GRB afterglow spectra in \\S4, and we present a summary in \\S5. ",
        "conclusions": "The six high-ion absorption lines reported in this paper at velocities near \\zgrb\\ are \\sif, \\sfo, \\cf, \\ssix, \\nf, and \\os, which trace the Si$^{+3}$, S$^{+3}$, C$^{+3}$, S$^{+5}$, N$^{+4}$, and O$^{+5}$ ions, requiring energies of 33.5, 34.8, 47.9, 72.6, 77.5, and 113.9~eV for their creation, respectively \\citep{Mo03}. A galactic ionizing spectrum, dominated by the integrated radiation from O and B stars, drops strongly above 54~eV, the \\ion{He}{ii} ionization edge \\citep{BH86}. So whereas the ions Si$^{+3}$, S$^{+3}$, and C$^{+3}$ can be photoionized in the ISM by starlight, the ions S$^{+5}$, N$^{+4}$, and O$^{+5}$ cannot. We expect this to remain true in GRB host galaxies, even though they have been shown to harbor significant numbers of Wolf-Rayet stars \\citep{Ha06}, since Wolf-Rayet spectra also show a strong break at 54\\,eV \\citep{Cr07}. Instead, the detection of the three ions S$^{+5}$, N$^{+4}$, and O$^{+5}$ implies the presence of either a hard, non-stellar radiation source or hot, collisionally ionized gas.  In the seven GRB afterglow spectra in our sample, the median and standard deviation of log\\,$N$(\\os) at \\zgrb\\ is 15.12$\\pm$0.54. For comparison, the median log\\,$N$(\\os) measured in 100 sight lines passing through the Milky Way halo is 14.38 \\citep{Wa03, Sa03}, rising to 14.80 when including the contribution from high-velocity clouds \\citep{Se03}. Turning to \\nf, we report a median log\\,$N$(\\nf)=14.03$\\pm$0.49 in the seven GRB spectra, whereas the median log\\,$N$(\\nf) in 32 Galactic halo sight lines is 13.45 \\citep{IS04}. High-ion absorption in \\nf\\ and \\os\\ has also been observed in $z$=2--3 DLA galaxies, with a median log\\,$N$(\\os)=14.77 \\citep{Fo07a}, and 76 per cent of DLAs showing log\\,$N$(\\nf)$<$13.50 (Fox et al. 2008, in preparation). % Thus the column densities of \\os\\ and \\nf\\ seen at \\zgrb\\ are, in the mean, higher than those seen in the halo of the Milky Way and those in DLAs. The highly ionized Galactic ISM is a complex environment, with many physical processes contributing to the production of the high ions, including conductive interfaces, turbulent mixing layers, shocks, expanding supernova remnants, and galactic fountain flows \\citep{Sa03, Zs03, IS04, Bo08}, and there is no reason why the ISM in the host galaxies of GRBs should be any less complex.  With this in mind, we now discuss the various categories of high-ion absorption seen near \\zgrb. Throughout this section, we refer the reader to Table 9, which summarizes the high-ion to high-ion column density ratios measured in the GRB afterglow spectra, and compares them to the ratios measured in various other astrophysical environments.  \\begin{table} \\begin{minipage}[t]{\\columnwidth} % \\caption{High-ion column density ratios} \\begin{tabular}{lccc} \\hline\\hline GRB\\footnote{We report the observed column density ratios in different velocity regions: in the strong components seen at \\zgrb, in the blueshifted wings, and in the HV absorbers. All limits are 3$\\sigma$. In cases marked $^*$, saturation, blending, or low S/N prevents the ratio from being measured reliably.} & $\\frac{N({\\rm \\cf})}{N({\\rm \\os)}}$ & $\\frac{N({\\rm \\cf})}{N({\\rm \\sif)}}$ & $\\frac{N({\\rm \\nf})}{N({\\rm \\os)}}$\\\\ \\hline \\bf{At $z$(GRB)}\\\\ 050922C   & $^*$         & $^*$         & 0.21$\\pm$0.04\\\\ % 060607    & $^*$         & 7.6$\\pm$1.2     &  $^*$     \\\\ % \\hline \\bf{Blueshifted Wings}\\\\ 050730    & 0.40$\\pm$0.25 & 3.7$\\pm$0.9 & $<$0.09  \\\\ % 060607    & $^*$          & $>$13       & $^*$     \\\\ % 071031    & 1.6$\\pm$0.4   & 3.2$\\pm$0.9 & $<$0.11  \\\\ % 080310    & 0.55$\\pm$0.11 & 25$\\pm$4    & $<$0.05  \\\\ % \\hline \\bf{HV components}\\\\ 021004 $-$2900~\\kms\\ & $^*$          & 5.1$\\pm$3.1 & $^*$\\\\ % 050730 $-$1550~\\kms\\ & $^*$          & 3.7$\\pm$0.4 & $^*$\\\\ % 050820 $-$3660~\\kms\\ & $^*$          & $>$24       & $^*$\\\\ % 060607 $-$1850~\\kms\\ & $^*$          & 3.2$\\pm$0.5 & $^*$\\\\ % 071031 $-$560~\\kms\\  & 0.19$\\pm$0.06 & $>$110   & $<$0.06\\\\ % 071031 $-$510~\\kms\\  & 0.35$\\pm$0.10 & $>$140   & $<$0.07\\\\ % 071031 $-$370~\\kms\\  & $^*$          & 40$\\pm$9    & $^*$\\\\ % 080310 $-$1400~\\kms\\ & $^*$          & $>$36       & $^*$\\\\ % \\hline Galactic Halo\\footnote{Mean$\\pm$1$\\sigma$ in 16 Milky Way halo sight lines \\citep{Zs03}.} & 0.60$\\pm$0.47 &  3.46$\\pm$1.09 & 0.12$\\pm$0.07\\\\ LMC\\footnote{Range observed in LMC gas (not including \\hw\\ regions) in four sight lines \\citep{LH07}.} & $<$0.09--0.59 & 1.5--2.5 & $<$0.33\\\\ % DLAs, $z$=2--3\\footnote{Mean$\\pm$1$\\sigma$ in 10 DLAs \\citep{Fo07a}.} & 2.1$\\pm$0.7 & 5.4$\\pm$2.2 & ... \\\\ IGM, $z$=2--4 & 0.35\\footnote{Median value among four IGM absorbers measured by \\citet{Si06}.} & 16$\\pm$4\\footnote{Median$\\pm$1$\\sigma$ in 188 IGM absorbers \\citep{Bo03}.} & $<$0.23\\footnote{Measured in composite QSO spectrum by \\citet{LS93}.}\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  \\subsection{Strong absorption at \\zgrb} We always detect strong high-ion absorption exactly at \\zgrb\\footnote{In two cases the velocity centroid of the strongest high-ion absorption component differs from the nominal GRB redshift (measured from the low-ion lines) by more than 10~\\kms: for GRB~050922C the offset is 65~\\kms, and for GRB~071031 the offset is 25~\\kms. However, since the GRB redshifts are not known to better than a precision of several tens of \\kms, these offsets are not significant, and we treat the strong high-ion absorption as arising exactly at \\zgrb.}, in the form of saturated \\os, \\cf, and \\sif\\ components, and (where the data are unblended) also in \\sfo\\ and \\ssix. \\nf\\ is present in 6/7 cases, with the advantage of being less saturated than the other high-ion lines. Because of this, \\nf\\ is the best line to investigate the line width and optical depth in the strong absorption components. In the six cases where \\nf\\ is detected, we fit eight \\nf\\ components, with $b$-values of 7$\\pm$3, 9$\\pm$3, 11$\\pm$2, 14$\\pm$8, 16$\\pm$3, 18$\\pm$3, 19$\\pm$3 and 28$\\pm$3~\\kms. The median of these eight values is 16~\\kms. % A \\nf\\ component formed in gas at 200\\,000~K, the temperature at which the production of \\nf\\ by collisions is maximized \\citep{GS07}, has a $b$-value of 15.4~\\kms. Therefore, we are unable to rule out collisional ionization for the gas traced by \\nf.  \\subsubsection{Photoionized circumburst gas?} Noticing that the \\nf\\ absorption at \\zgrb\\ tends to be stronger and narrower (in total velocity width) than the \\nf\\ absorption observed in the Milky Way ISM and in DLAs, \\citet{Pr08b} have recently argued that the \\nf\\ absorption components at \\zgrb\\ do not trace ISM material in the halo of the host galaxy, but rather arise in circumburst gas in the immediate vicinity of the GRB. Our high-resolution UVES dataset, covering a wider range of ionization states than has been observed before, allows us to investigate this intriguing idea.  Our data are qualitatively consistent with the circumburst hypothesis in the cases of GRBs~021004, 050922C, and 071031, where we see strong, single components of \\nf\\ at \\zgrb, each with log\\,$N$(\\nf)$>$14.0, aligned with other high ions including \\os\\ with log\\,$N$(\\os)$\\ga$14.5. For these cases, our new measurements of the \\os, \\sfo, and \\ssix\\ column densities at \\zgrb\\ (listed in Tables 2, 4, and 7) could be used to constrain photoionization models of the circumburst region. However, in the remaining four spectra in our sample (GRBs~050730, 050820, 060607, and 080310) the situation is more complex. In each of these four cases, log\\,$N$(\\nf)$<$14.1, and in two cases log\\,$N$(\\nf)$<$13.2, i.e the observed \\nf\\ is \\emph{not} in excess of what is expected from the host galaxy's ISM. In two of these cases (GRBs 050730 and 080310) the \\nf\\ absorption profiles show two components separated by a few tens of \\kms\\ -- this is not the signature of a single burst. In another case (GRB~050820) there is an offset of 110~\\kms\\ between a single \\nf\\ component and the strongest absorption in the other high ions. And finally, in the case of GRB~060607, no \\nf\\ is detected at all. In addition to these case-by-case details, we repeat that the $b$-values of the detected \\nf\\ absorption components at \\zgrb\\ are \\emph{not} particularly narrow, so photoionization is not required by the line widths (though it is not ruled out either).  \\subsubsection{Total ionized column density} We can further investigate the nature of the strong components seen at \\zgrb\\ by calculating the total column density of ionized hydrogen $N$(hot~\\hw)\\footnote{Here we use the word ``hot'' simply to label the \\nf-bearing gas; the upper limit on the temperature is 215,000\\,K for the median \\nf\\ $b$-value of 16~\\kms.} contained in these components. To do this we have to correct for ionization and metallicity, so the calculation is only possible for cases where the nitrogen abundance has been measured in the DLA at \\zgrb. This is true for three GRBs in our sample: [N/H]=$-$3.16$\\pm$0.10 for GRB~050730, $>\\!-$1.35 for GRB~050820, and $<\\!-$4.09 for GRB~050922C \\citep{Pr07b}. The relationship between the \\nf\\ and hot \\hw\\ column densities can be written in the following way: \\begin{equation} N({\\rm hot~\\hw})=\\frac{N{\\rm (\\nf)}}{{\\rm (\\nf/N)(N/H)_n}}\\frac{{\\rm (N/H)_{n}}}{({\\rm N/H)_{\\rm i}}}, \\end{equation} where (N/H)$_{\\rm i}$ denotes the nitrogen abundance in the ionized gas and (N/H)$_{\\rm n}$ denotes the nitrogen abundance in the neutral gas (which is measured). We take the solar nitrogen abundance of (N/H)$_\\odot=10^{-4.22}$ from \\citet{Gr07}, and then (N/H)$_{\\rm n}=10^{\\rm [N/H]}$(N/H)$_\\odot$. If the gas is collisionally ionized, as would be appropriate for \\nf\\ in the host galaxy ISM, \\nf/N$<$0.25 at all temperatures \\citep{GS07}\\footnote{This is true for either equilibrium or non-equilibrium models. However, in the \\citet{Pr08b} GRB photoionization models, a \\nf\\ ionization fraction as high as 0.6 is predicted.}. Therefore we adopt 0.25 as the maximum allowed ionization fraction, corresponding to a minimum ionization correction of a factor of 4. If the ionized gas has the same metallicity as the neutral gas, so that the term (N/H)$_{\\rm n}$/(N/H)$_{\\rm i}$ in Equation 1 is equal to one, then $N$(hot~\\hw) is:\\\\ $>$21.8 for GRB~050730 [which has log\\,$N$(\\hi)=22.15],\\\\ $>$19.3 for GRB~050820 [which has log\\,$N$(\\hi)=21.00], and \\\\ $>$22.7 for GRB~050922C [which has log\\,$N$(\\hi)=21.55],\\\\ where the results are lower limits since (a) we have used the smallest allowed ionization correction, and (b) they only apply to the \\nf-bearing gas, and \\hw\\ may exist at other temperatures as well. Comparing the values of $N$(hot~\\hw) and $N$(\\hi), we find the hot-ionized-to-neutral ratio $N$(hot~\\hw)/$N$(\\hi) takes values of $>$0.4 and $>$0.02 for GRBs~050730 and 050820, and $>$13 for GRB~050922C. Whereas the first two values are modest, the GRB~050922C value is extremely (and implausibly) high -- it requires over ten times as much mass in hot ionized gas in a single component as there is in neutral gas in the entire host galaxy. However, if the nitrogen abundance in the ionized gas was higher than in the neutral gas, so that (N/H)$_{\\rm n}$/(N/H)$_{\\rm i}<1$, then solutions with lower $N$(hot~\\hw) would be possible. Thus our calculation of the total ionized column density at \\zgrb\\ suggests that the metallicity in the highly ionized gas is higher than the metallicity in the neutral gas (particularly for GRB~050922C); however, this calculation is inconclusive in determining whether the gas is circumburst or interstellar.  \\subsubsection{Photo-excitation modeling of $N$(\\sfo) and $N$(\\sfo$^*$)} \\begin{figure} \\includegraphics[width=9cm]{f7.eps} \\caption{Analysis of the \\sfo\\ level populations at \\zgrb\\ toward GRB~050730. The solid and dashed lines show the predictions from UV photo-excitation models for the time-variation of the \\sfo\\ and \\sfo$^*$ column densities (see text for details). The data points show the observed $N$(\\sfo) and upper limit to $N$(\\sfo$^*$) in the UVES spectrum, plotted at the $x$-position corresponding to the mid-observation epoch (five hours after the trigger). The solid lines show the model with the best-fit distance of 418$\\pm$9~pc. The dashed lines show a model with the distance $d$ fixed at 200~pc, for illustration. In this case $N$(\\sfo$^*$) would be 0.5 dex stronger than observed, so this model is clearly ruled out. The result that the \\sfo-absorbing gas must lie at $d>400$~pc implies an interstellar rather than circumburst origin.} \\end{figure}  Another method for distinguishing between circumburst gas and more distant interstellar gas is to look for absorption from fine-structure levels. In circumburst gas, UV pumping by GRB photons will populate the excited electronic levels, which decay and cascade to populate the fine-structure levels of the ground state \\citep{Pr06, Ch07, Vr07}. Because of the rapid fading of the afterglow, the level populations evolve in a non-equilibrium, time-dependent manner. Our data allow us to investigate this process in the \\sfo\\ ion, since in four GRB spectra the combination of a detection of \\sfo\\ $\\lambda$1062.662 and a non-detection of \\sfo$^*$~$\\lambda$1072.973 allows us to place an upper limit on the \\sfo$^*$/\\sfo\\ ratio at \\zgrb. The Einstein $A$-value of the forbidden transition between the ground state of \\sfo\\ and its fine-structure level is $7.7\\times10^{-3}$~s$^{-1}$ \\citep[taken from NIST;][]{Ra08}\\footnote{See http://physics.nist.gov/PhysRefData/ASD/index.html}. In other words, in the absence of stimulated photo- and collisional de-excitation, the fine-structure level probed by the \\sfo$^*$ 1072.973 transition will decay in 1/$A$=130\\,s. In fully ionized gas (as expected near GRBs) there is no \\lya\\ opacity and so the \\sfo\\ lines (which lie shortward of \\lya) will not be shielded from UV radiation. For GRBs~021004, 050730, and 080310 we measure 3$\\sigma$ limits to the $N$(\\sfo$^*$)/$N$(\\sfo) ratio of $<$0.08, $<$0.12, and $<$0.41 respectively (for the other four afterglows we are unable to constrain this ratio due to blending). %  Following the technique described in \\citet{Vr07}, we can model the \\sfo\\ excitation level in gas subject to UV radiation from the GRB, assuming that photo-excitation is the dominant excitation process. We chose to model the \\sfo\\ level population measured toward GRB~050730, because this case has the strongest constraint on $N$(\\sfo$^*$). The model has five free parameters: the distance $d$ between the GRB and the absorbing gas (assuming all the \\sfo\\ arises at the same distance), the total \\sfo\\ column density $N_{\\rm tot}$ (equivalent to the pre-burst ground-state \\sfo\\ column density), the afterglow spectral slope $\\beta$ (where $F_\\nu\\propto\\nu^{\\beta}$), the Doppler $b$-parameter of the absorbing gas, and the rest-frame time when the calculation begins $t_0$ (i.e. the time when the afterglow photons begin to be tracked in the code). For a given set of these parameters, the model predicts the \\sfo\\ and \\sfo$^*$ column densities as a function of time (both in the observer frame and in the rest frame). Since the time of mid-observation of GRB~050730 is known, five hours (observer frame) after the trigger, predicted values for $N$(\\sfo) and $N$(\\sfo$^*$) can be extracted from each model.  The parameters in our GRB~050730 model were determined as follows. $N_{\\rm tot}$ was chosen to equal $N$(\\sfo)+$N$(\\sfo$^*$), assuming the actual $N$(\\sfo$^*$) is equal to the measured upper limit (if the actual value is lower, our conclusions are strengthened; see below). The value for $\\beta$ ($-$0.56) was determined using the observed light curve for this afterglow. The parameter $t_0$ is poorly constrained, but our runs showed that the \\sfo\\ level populations are very insensitive to $t_0$ for values between 0 and 300~s; we used a value $t_0$=10\\,s \\citep[see][]{Vr07}. Finally a $b$-value of 10~\\kms\\ was chosen as a typical value for interstellar gas, and indeed matches (within the 1-$\\sigma$ error) the measured $b$-value for the \\sfo\\ component at 0~\\kms\\ toward GRB~050730. This leaves the distance $d$ as the one genuinely floating parameter in the model; a grid of models at different distance $d$ was generated, and the predicted and observed $N$(\\sfo) and $N$(\\sfo$^*$) were compared using a chi-squared minimization routine.  The model results are shown in Figure 7. The best-fit distance $d$ is 418$\\pm$9~pc. Because we assumed that the actual \\sfo$^*$ column density was equal to the measured upper limit, this represents a strong \\emph{lower limit} of 400~pc (2$\\sigma$) on the distance of the absorbing gas from the GRB. In other words, if the gas was at $d\\!<\\!400$~pc, we would have seen \\sfo$^*$ absorption, whereas we clearly do not. This represents an important result, since it implies that the \\sfo\\ absorption has an interstellar rather than circumburst origin. The distance limit on the \\sfo\\ absorption does not necessarily have to apply to the \\nf\\ absorption (or the other high ions). However, the \\nf\\ and \\sfo\\ profiles at \\zgrb\\ toward GRB~050730 are similar, both showing the component near 0~\\kms\\ and a weaker component at 25~\\kms, suggesting that the two ions are co-spatial and hence that the \\nf\\ absorption also arises at $>$400~pc from the GRB.  \\subsubsection{Origin of strong absorption} In summary, we find that while the circumburst hypothesis explains the observation that in three afterglow spectra, strong single-component \\nf\\ is aligned with saturated absorption in the other high ions, it is challenged by the \\nf\\ profiles in the other four datasets, it is not strictly required by the \\nf\\ line widths, and the expected high-ion fine-structure lines and time-variable column densities are yet to be observed. Furthermore, in one case (GRB~050730) we are able to place a lower limit on the distance to the high-ion absorbing gas of 400~pc, using the simultaneous detection of \\sfo\\ and non-detection of \\sfo$^*$. Therefore, based on our current dataset, we cannot rule out an interstellar explanation for the strong high-ion components at \\zgrb, noting that since the higher $N$(\\hi) values in GRB-DLAs than in QSO-DLAs suggest that GRB sightlines pass through the inner star-forming regions of the host galaxies \\citep{Pr08a, Fy08}, one would \\emph{expect} high column densities of interstellar plasma at \\zgrb, so the strength of the \\nf\\ absorption does not, in our view, argue against an interstellar origin. We also note that a strong, apparently narrow (though unresolved) \\nf\\ component is observed in the spectrum of LBG cB58 at the galaxy redshift \\citep[$z$=2.73;][]{Pe02}, showing that such components can be formed even in the absence of a GRB. Time-series observations of \\nf\\ in GRB afterglow spectra, which unfortunately are not available in the current dataset, would help to resolve this issue. A significant detection of time-variation in the \\nf\\ column density at \\zgrb\\ \\citep[see models by][]{Pr08b} would represent the discovery of gas that is unambiguously close to the GRB. %  \\subsection{Multi-component absorption in \\cf\\ and \\sif} The multi-component structure seen in \\cf\\ and \\sif\\ spreading over several hundred~\\kms\\ around \\zgrb\\ is not seen in the other high ions. However, the \\cf\\ and \\sif\\ profiles lie out of the \\lya\\ forest, so are free from blending and show higher S/N than \\os\\ or \\nf, and it is unclear how the \\os\\ and \\nf\\ profiles would appear if observed at the same S/N. These complex, multi-component \\sif\\ and \\cf\\ profiles are reminiscent of the profiles of these ions in damped \\lya\\ (DLA) absorbers \\citep{Lu96, Le98, WP00, Fo07b, Le08}, which represent galaxy halos seen in absorption toward a background source. The median (mean) value of the total \\cf\\ velocity width in 74 DLAs and sub-DLAs at $z$=2--3 reported by \\citet{Fo07b} is 255~\\kms\\ (342~\\kms), where the width is the total observed range of absorption, regardless of optical depth. In the seven GRB spectra in our current sample, the median (mean) total \\cf\\ velocity width measured in an identical manner is 280~\\kms\\ (320~\\kms), % similar to the DLA value (note we have excluded the HV components, which we treat separately). A precise comparison of the median \\cf\\ column density between the GRB sample and the DLA sample is not possible, because these measurements are affected by saturation, but we can state that log\\,$N$(\\cf) at \\zgrb\\ is $\\ga$15 in all seven GRB spectra, which is significantly stronger than the median DLA value log\\,$N$(DLA \\cf)=14.27$\\pm$0.57 measured by \\citet{Fo07b}. Nonetheless, the similarities in total \\cf\\ line width between GRB and DLAs support a galactic origin for the multiple components seen in \\cf\\ and \\sif\\ in a velocity range of several hundred \\kms\\ around \\zgrb.  \\subsection{Negative-velocity (blueshifted) absorption-line wings} One notable feature present in four of the seven GRB afterglow spectra in our sample is a negative-velocity absorption-line wing seen in \\cf, \\sif, and (at lower S/N) \\os. These wings are characterized by relatively smooth, asymmetric absorption extending for 65--140~\\kms\\ % to negative velocities relative to the strong absorption at \\zgrb. To explore the high-ion velocity structure in the wings, we present in Figure 8 apparent column density profiles of the high ions over the velocity of the wings, calculated using $N_a(v)=3.768\\times10^{14}\\tau_a(v)/f\\lambda$, where the apparent optical depth $\\tau_a(v)={\\rm ln}\\,[F_c(v)/F(v)]$, $F(v)$ and $F_c(v)$ are the observed flux level and estimated continuum level as a function of velocity, and where $\\lambda$ is in \\AA\\ \\citep{SS91}. We follow the error treatment of \\citet{SS92}. This form of display allows the velocity structure seen in each ion to be closely compared. In two cases we add the profile of \\cw\\ $\\lambda$1334.532, which also shows wing absorption. The high-ion column densities in the wings are summarized in Table 10.   \\begin{table*} \\begin{minipage}[t]{18cm} \\caption{Properties of absorption in the blueshifted wings} \\begin{tabular}{lcccc ccc} \\hline\\hline GRB & \\zgrb\\ & $v_{\\rm min}, v_{\\rm max}$\\footnote{On velocity scale relative to \\zgrb.} & $\\delta v$\\footnote{Velocity extent of wing $\\delta v\\!\\equiv\\!v_{\\rm max}$--$v_{\\rm min}$, shaded light grey on Figures 2, 3, and 4.} & log\\,$N$(\\os) & log\\,$N$(\\cf) & log\\,$N$(\\sif) & log\\,$N$(\\nf)\\\\ & &       (\\kms) & (\\kms) & ($N$ in \\sqcm) & ($N$ in \\sqcm) & ($N$ in \\sqcm) & ($N$ in \\sqcm)\\\\ \\hline 050730 &  3.9686 & $-$145,$-$ 80 &  65 &                 14.47$\\pm$0.35 & 13.79$\\pm$0.03  &                13.21$\\pm$0.02 &             $<$13.56 \\\\ 060607 &  3.0749 & $-$140,$-$ 60 &  80 &                            ... & 13.00$\\pm$0.02  &                      $<$12.27 &             $<$12.91 \\\\ 071031 &  2.6922 & $-$150,$-$ 40 & 110 &                 14.25$\\pm$0.06 & 14.45$\\pm$0.09  &                13.95$\\pm$0.06 &             $<$13.59 \\\\ 080310 &  2.4274 & $-$260,$-$120 & 140 &                 14.31$\\pm$0.08 & 14.05$\\pm$0.02  &                12.65$\\pm$0.06 &             $<$13.27 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\begin{figure*} \\includegraphics[width=18cm]{f8.eps} \\caption{High-ion apparent column density profiles in the four negative-velocity wings in our sample. Errors reflect both statistical noise and continuum placement uncertainties. The \\cf\\ and \\sif\\ profiles have each been scaled by the factors annotated on the plot to allow for comparison of their shape with \\os. Only lower limits (arrows) can be presented for saturated pixels. The \\os\\ data have been rebinned by four pixels in the case of GRB~071031, eight pixels in the case of GRB~050922C, and five pixels in the other cases. \\nf, \\sfo, and \\ssix\\ are not detected in the wings.} \\end{figure*}  We find that the \\cf\\ and \\sif\\ profiles track one another closely in three of the four wings, indicating that the two ions form in the same regions of gas. The exception is the very weak wing seen toward GRB~060607, in which \\sif\\ is not detected, but \\sit\\ is seen. For the wings seen toward GRBs~050730 and 080310, the \\os\\ and \\cf\\ profiles share a common gradient d$N_a$/d$v$. However, for GRB~071031, the \\os\\ absorption shows a shallower slope. We measure a $N$(\\cf)/$N$(\\sif) ratio of $\\approx$3 in two of the four wings (GRBs~050730 and 071031), similar to the Galactic halo average of 3.46$\\pm$1.09 measured by \\citet{Zs03}. This supports an interstellar origin for the absorption-line wings. In the wings toward GRBs~060607 and 080310 the $N$(\\cf)/$N$(\\sif) ratio is much higher, taking values of $>$13 (3$\\sigma$) and 25$\\pm$4 respectively. %  \\os\\ absorption-line wings with similar velocity extent have been observed at zero redshift in the ultraviolet spectra of quasars taken with the \\emph{Far-Ultraviolet Spectroscopic Explorer (FUSE)} satellite \\citep{Se01, Se03, Fo05, Fo06, Sa05a, Ke06}\\footnote{Analogous \\cf\\ and \\sif\\ wings are yet to be observed at zero redshift, but these two ions fall in the near-UV, requiring \\emph{HST} observations rather than \\emph{FUSE}, and the number of high-resolution QSO spectra taken with \\emph{HST} is small compared to the number of extragalactic sight-lines observed with \\emph{FUSE}.}. \\citet{Fo06} measured eleven Galactic \\os\\ wings, and report a mean and standard deviation \\os\\ column density of log\\,$N$(\\os)=13.85$\\pm$0.45, with each wing extending over $\\sim$100~\\kms; these authors successfully modeled the \\os\\ absorption-line wings as tracing Galactic outflows moving under ballistic trajectories. Furthermore, a strong \\os\\ wing covering a similar velocity range is seen in a sub-DLA (an absorber with 19.0$<$log\\,$N$(\\hi)$<$20.3) at $z$=2.67 \\citep{Fo07c}. The \\cf\\ profile observed in the LBG cB58 at $z$=2.7 shows a clear absorption wing extending from $-$250 to $-$750~\\kms\\ ($\\approx$ nine resolution elements) in the galaxy rest-frame \\citep{Pe02}, and an asymmetric blueshifted absorption feature in \\ion{Mg}{ii} has been reported in the composite spectrum of 1400 $z\\!\\approx\\!1.4$ galaxies by \\citet{We08}, who also interpret it as a galactic outflow signature.  \\begin{figure*} \\includegraphics[width=18.5cm]{f9.eps} \\caption{Montage of apparent column density profiles of \\os\\ absorption-line wings measured in various environments: near \\zgrb\\ in three GRB afterglow spectra (this paper, three left-most lines), at zero redshift in the Milky Way halo in the UV spectra of three AGN \\citep[][ three right-most lines]{Fo06}, and in a sub-DLA at $z$=2.6562 \\citep{Fo07c}. The Milky Way and sub-DLA cases have been inverted and velocity offsets have been added for ease of comparison, but the gradient of each profile is preserved. Linear fits are shown with color-coded dashed lines, and the gradient of each fit in units of $10^9$~ions~\\sqcm~(\\kms)$^{-2}$ is annotated with its 1$\\sigma$ error on the panel. In these diverse galactic environments sampling over twelve billion years, the \\os\\ absorption-line profiles can be used to trace and quantify galactic outflows.} \\end{figure*}  Whereas the $z$=0 \\os\\ wings are always detected at \\emph{positive} velocities (redshifted) relative to the Local Standard of Rest (LSR), the GRB wings are seen at \\emph{negative} velocities (blueshifted) relative to \\zgrb. This behavior is strongly suggestive of galactic outflows: in the Milky Way case (at least for high-latitude and anti-center sight lines, where Galactic rotation effects are insignificant), outflowing gas appears redshifted relative to the LSR, whereas in the GRB case, outflowing gas moves toward us along the line-of-sight and hence appears blueshifted relative to the host galaxy. The possibility remains that some wing absorbers represent chance alignments of unrelated components separated by small velocities, which are unresolved by the spectrograph. However, we find this explanation unlikely for the strongest wings in our sample, seen toward GRBs~071031 and 080310, since in these cases the column density in the wing increases monotonically with velocity\\footnote{In the case of GRB~071031, the \\cf\\ profile shows a smooth gradient with velocity once the $-$115~\\kms\\ component has been subtracted off.} in an interval of over 100~\\kms, and there is no reason why a series of random components should be aligned in this way.  To illustrate the similarity in the wing absorption features observed at both low and high redshift, we show in Figure 9 a comparison of the \\os\\ apparent column density profiles in the zero redshift and GRB sight line wings. We have inverted the positive-velocity wings to compare their gradient with the wings observed in the GRB spectra. While the similarity in the overall shapes of the wing profiles is striking, the high-$z$ wings in the GRB host galaxies show higher \\os\\ column densities (by a factor of $\\approx$5) % and steeper gradients d$N_{\\rm a}$/d$v$ (by a factor of 5--10) % than the Milky Way wings, implying higher mass flow rates.  \\subsection{High-velocity components at 500--5\\,000~\\kms} In six of our seven GRB spectra (all GRBs except 050922C), high-ion components at 500--5\\,000~\\kms\\ relative to \\zgrb\\ are observed (see Figure 5). Such HV components have been observed before in GRB~020813 \\citep{Ba03}, GRB~021004 \\citep{Mo02, Mi03, Sc03, Fi05, La06}, GRB~030226 \\citep{Kl04}, GRB~050505 \\citep{Be06}, and GRB~050730 \\citep{DE07, Ch07}. \\citet{Ch07} report that \\cf\\ components with $W_{\\rm r}\\!>\\!200$~m\\AA\\ occur at 1\\,000--5\\,000~\\kms\\ in 20 per cent of GRB afterglow spectra. In our data, the incidence of HV \\cf\\ absorbers (6/7, or 86\\%) is much higher because (a) we have a higher sensitivity: the weakest HV \\cf\\ absorber in our sample shows $W_{\\rm r}$=48~m\\AA, and (b) we classify absorbers in the range 500--1\\,000~\\kms\\ as HV.  The HV absorbers seen in GRB afterglow spectra do not represent a homogeneous population, as do the blueshifted absorption-line wings and (potentially) the strong, narrow absorbers at \\zgrb. On the contrary, we observe a broad range of properties among the HV absorbers. In three sight lines (toward GRBs~021004, 050730, and 060607), % the HV components exhibit low-ionization absorption (clearest in \\cw\\ $\\lambda$1334.532 and \\siw\\ $\\lambda$1260.422). However, the HV components toward GRBs~050820, 071031 and 080310 show no (or very weak) neutral-phase absorption, in either \\cw\\ or \\siw. So whereas the presence of neutral gas in the first three cases implies the absorbing material is not formed in the circumburst medium \\citep[because this medium is expected to be fully ionized by the burst; e.g.][]{Vr07, Ch07, Pr08b}, we find that the properties of the HV components in the latter three cases \\emph{are} consistent with an origin in the close circumburst medium. This is particularly true for the HV gas toward GRB~071031, which is detected in the form of strong, \\os\\ and \\cf\\ components that are not only absent in \\cw\\ and \\siw, but also show non-detections of \\sif. We have placed limits of $N$(\\cf)/$N$(\\sif) $>$110 at $-$560~\\kms, and $>$140 at $-$510~\\kms\\ in the GRB~071031 spectrum. These extreme ratios are far higher than typical interstellar values \\citep[Galactic ISM values for this ratio are $\\approx$3;][]{Zs03} implying these HV components have a different, non-interstellar origin.  Given the connection between GRBs and massive stars \\citep{MW99, WB06}, a leading explanation for the HV components in GRB afterglow spectra is that they trace outburst episodes from the massive-star progenitors of the GRBs \\citep[e.g.][]{Sc03}, and thus arise in the circumburst region. As pointed out by \\citet{Be06}, a Wolf-Rayet outflow model would give rise to absorbers that are enriched in carbon and deficient in silicon, naturally explaining the high \\cf/\\sif\\ ratios measured in the HV absorbers.  If the wind-blown bubbles around GRB progenitors are similar in size to those observed around Wolf-Rayet stars in the Milky Way, their typical radius will be $\\approx$2--10~pc \\citep{Gr00}. Several groups have recently modeled Wolf-Rayet winds and their relationship to the HV components in GRB afterglow spectra \\citep{Ra01, Ra05, Cv04, El06}, and in particular, predictions for the high-ion signature of Wolf-Rayet winds are given by \\citet{vM05,vM07,vM08}. Depending on observing epoch and viewing angle, components at 150-700~\\kms\\ and 1\\,000-2\\,000~\\kms\\ are predicted at various stages of Wolf-Rayet evolution \\citep{vM05}; these velocities match the observed range of HV components in our sample. Recent models of the absorption-line signatures of galactic winds \\citep[as opposed to single-star winds;][]{Fa07, KR07, Sa08} generally predict high-ion components at velocities of only 100--200~\\kms\\ relative to the galaxy, though observations support the idea that galactic winds can reach speeds of $\\approx$1\\,000~\\kms\\ \\citep{Fr02, Tr07, We08}. Nonetheless, the fairly narrow $b$-values measured in the HV components toward GRBs~071031 and 080310 (not to mention the lack of \\sif) are at odds with the broader, collisionally ionized components expected in galactic winds \\citep{OD06}. Thus we conclude that among the available models, Wolf-Rayet winds are the best explanation for the highly ionized HV components toward GRBs~071031 and 080310."
    },
    "0809/0809.3906_arXiv.txt": {
        "abstract": "Axisymmetric, orbit-based dynamical models are used to derive dark matter scaling relations for Coma early-type galaxies. From faint to bright galaxies halo core-radii and asymptotic circular velocities increase. Compared to spirals of the same brightness, the majority of Coma early-types -- those with old stellar populations -- have similar halo core-radii but more than 2 times larger asymptotic halo velocities. The average dark matter density inside $2 \\, \\reff$ decreases with increasing luminosity and is $6.8$ times larger than in disk galaxies of the same $B$-band luminosity. Compared at the same stellar mass, dark matter densities in ellipticals are $13.5$ times higher than in spirals. Different baryon concentrations in ellipticals and spirals cannot explain the higher dark matter density in ellipticals. Instead, the assembly redshift (1+z) of Coma early-type halos is likely about two times larger than of comparably bright spirals. Assuming that local spirals typically assemble at a redshift of one, the majority of bright Coma early-type galaxy halos must have formed around $z \\approx 2-3$. For about half of our Coma galaxies the assembly redshifts match with constraints derived from stellar populations. We find dark matter densities and estimated assembly redshifts of our observed Coma galaxies in reasonable agreement with recent semi-analytic galaxy formation models. ",
        "introduction": "\\label{sec:intro} Present-day elliptical galaxies are known to host mostly old stellar populations (\\citealt{Tra00}, \\citealt{Ter02}, \\citealt{DTho05}). Whether their stars have formed in situ or whether ellipticals assembled their present-day morphology only over time (for example by mergers) is less clear. An important clue on the assembly redshift of a galaxy is provided by its dark matter density. For example, in the simple spherical collapse model \\citep{Gun72} the average density of virialized halos is proportional to the mean density of the universe at the formation epoch: halos which form earlier become denser. Similarly, in cosmological $N$-body simulations, the concentration (and, thus, the inner density) is found to be higher in halos that have assembled earlier (e.g. \\citealt{Nav96}, \\citealt{Wec02}). In addition to this connection between formation epoch and halo density, the final halo mass distribution also depends on the interplay between dark matter and baryons during the actual galaxy formation process (e.g. \\citealt{Blu86}, \\citealt{Bin01}). Then, the properties of galaxy halos provide valuable information about when and how a galaxy has assembled its baryons.  Despite its cosmological relevance, the radial distribution of dark (and luminous) mass in early-type galaxies is not well known: because of the lack of cold gas as a dynamical tracer, masses are difficult to determine. Stellar dynamical models require the exploration of a galaxy's orbital structure and have only recently become available for axisymmetric or more general systems (\\citealt{Cre99}, \\citealt{Geb00}, \\citealt{Tho04}, \\citealt{Val04}, \\citealt{Cap05}, \\citealt{deLor07}, \\citealt{Cha08}, \\citealt{vdB08}). Scaling relations for the inner dark matter distribution in early-types have by now only been reported for round and non-rotating galaxies (\\citealt{Kr00}, \\citealt{G01}) and spirals (\\citealt{PSS1,PSS2} and \\citealt{Kor04}). The aim of the present paper is to provide empirical scaling relations for generic cluster early-types (flattened, with different degrees of rotation). In particular, this paper is focussed on the inner dark matter density and its implications on the assembly redshift of elliptical galaxy halos.  The paper is organized as follows. In Sec.~\\ref{sec:data}, we review the galaxy sample and its modelling. Dark matter scaling relations are presented in Sec.~\\ref{sec:scaling}. Sec.~\\ref{sec:density} is dedicated to the dark matter density. The effect of baryons on the dark matter density is discussed in Sec.~\\ref{sec:contraction}, while Sec.~\\ref{sec:assembly} deals with the halo assembly redshift. In Sec.~\\ref{sec:semi-analytic} our results are compared to semi-analytic galaxy formation models. A summary is given in Sec.~\\ref{sec:summary}. In the following, we assume that the Coma cluster is at a distance of $d = 100 \\, \\Mpc$. ",
        "conclusions": "\\label{sec:summary} We have presented dark matter scaling relations derived from axisymmetric, orbit-based dynamical models of flattened and rotating as well as non-rotating Coma early-type galaxies. Dark matter halos in these galaxies follow similar trends with luminosity as in spirals. Thereby, the majority of Coma early-types -- those with old stellar populations -- have halo core-radii $\\rh$ similar as in spirals with the same $B$-band luminosity, but their asymptotic halo velocities are about $2.4$ times higher. In contrast, four Coma early-types -- with young central stellar populations -- have halo velocities of the same order as in comparably bright spirals, but their core-radii are smaller by a factor of 4. Differences between spirals and ellipticals increase, when the comparison is made at the same stellar mass. The average halo density inside $2 \\, \\reff$ exceeds that of comparably bright spirals by about a factor of $6.8$. If the higher baryon concentration in ellipticals is taken into account, the excess density reduces to about a factor of 3, but if ellipticals and spirals are compared at the same stellar mass, then it is again of the order of $6.5$.  Our measured dark matter densities match with a comparison sample of simulated cluster ellipticals constructed from the semi-analytic galaxy formation models of \\citet{deL07}. These synthetic ellipticals have $\\zform \\approx 0.5-4$ and higher dark matter densities than simulated field spirals, which are appear on average around $\\zform\\spir \\approx 1$.  Assuming for local spirals $\\zform\\spir = 1$ as well, and assuming further that the inner dark matter density scales with the formation redshift like $(1+\\zform)^3$, our results imply that ellipticals have formed $\\Delta \\zform \\approx 1-2$ earlier than spirals. Without baryon correction, we find an average formation redshift around $\\zform \\approx 3$, which is slightly larger than in semi-analytic galaxy formation models. Accounting for the more concentrated baryons in ellipticals, the average formation redshift drops to $\\zform \\approx 2$.  \\begin{figure*}\\centering \\begin{minipage}{140mm} \\centering \\includegraphics[width=135mm,angle=0]{f9.eps} \\caption[] {Average halo density $\\rhoavdm$ versus assembly redshift. Large symbols: Coma galaxies without baryon correction (a) and with baryon correction (b). Small symbols: simulated cluster ellipticals and simulated field spirals from the semi-analytic models of \\citet{deL07}.} \\label{fig:rho_zf_sam} \\end{minipage} \\end{figure*}   For about half of our sample, dark halo formation redshifts match with constraints derived from stellar populations \\citep{Meh03}: the assembly epoch of these (old) early-types coincides with the epoch of formation of their stellar components."
    },
    "0809/0809.2846_arXiv.txt": {
        "abstract": "A general framework for tests of Lorentz invariance with electromagnetic waves is presented, allowing for operators of arbitrary mass dimension. Signatures of Lorentz violations include vacuum birefringence, vacuum dispersion, and anisotropies. Sensitive searches for violations using sources such as active galaxies, gamma-ray bursts, and the cosmic microwave background are discussed. Direction-dependent dispersion constraints are obtained on operators of dimension 6 and 8 using gamma-ray bursts and the blazar Markarian 501. Stringent constraints on operators of dimension 3 are found using 5-year data from the Wilkinson Microwave Anisotropy Probe. No evidence appears for isotropic Lorentz violation, while some support at 1$\\si$ is found for anisotropic violation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.0558_arXiv.txt": {
        "abstract": "{We present results on the modelling of the ejection of a superluminal component in the jet of \\object{3C~111}. We propose that the component is generated by an injection of dense material followed by a decrease in the injection rate of bulk particles in the jet. Our model is supported by 1D relativistic hydrodynamics and emission simulations, and is capable of reproducing the brightness evolution of two features, as revealed by 15 GHz VLBA observations. We show that other scenarios, such as an increase of the Lorentz factor in the material of the perturbation, fails to reproduce the observed evolution of this flare. ",
        "introduction": "Flaring events at radio frequencies are known to take place in Active Galactic Nuclei (AGN), usually followed by the observation of new radio features in the parsec-scale jets \\citep[e.g.,][]{sa02}. It has been shown that the ejection of those features, or components, is related to dips in the X-ray emission from the active nucleus in the case of 3C~120 \\citep{Mar02}, and perhaps also in 3C~111 \\citep{Mar06}. The dips in X-rays precede the observations of new radio-components. The decrease in X-ray emission may be caused by the loss of the inner regions of the disc. In this scenario, a fraction of the accreted material is injected in the jet and a new component is later observed in VLBI images, after the material becomes detectable at the observing frequencies, as it evolves downstream. The components are interpreted as the shocks produced by the ejection of denser and/or faster plasma in the flaring event from the accretion disc \\citep{mar85}. The conditions for triggering the ejection of the material in those radio features are still unknown. Hydrodynamical simulations \\citep[][ A03 hereafter]{Alo03} have shown that such jet perturbations produce a forward and a reverse structure, which would be expected to be observed as a fast front and a slower back component.   In the jet of 3C~111 ($z=0.049$, $1\\,\\rm{mas}\\simeq 1\\,\\rm{pc}$), a very strong flaring event in early 1996 gave rise to the ejection of two jet features observed at 15 GHz with the Very Long Baseline Array \\citep[labeled as components E and F -see Fig.~\\ref{fig:0} and][ K08 hereafter]{Kad08}. Both component trajectories can be back-extrapolated to similar ejection epochs within 3 months (around 1996.10). However, they show different speeds and the time evolution of their brightness is different (see Fig.~\\ref{fig:0}): whereas the inner component F is initially brighter (1996.82 and 1997.19) and fades out very rapidly (1997.66 and 1998.18), the leading component E shows a slower decrease in flux density. After 1999, F has disappeared and E evolves accelerating. In \\cite{pe08} we propose the possibility that these components are the front and rear region of a single perturbation. Here we review this result and discuss on the problems that other possible scenarios could face. \\begin{figure}[t] \\includegraphics[clip,angle=0,width=0.23\\textwidth]{fig1a} \\includegraphics[clip,angle=0,width=0.24\\textwidth]{fig1b} \\caption{Core distance and flux density evolution with time of components E and F in 3C~111, based on the results from K08.} \\label{fig:0} \\end{figure} ",
        "conclusions": "In any of the scenarios mentioned above, the perturbed regions have enhanced emission with respect to the underlying jet. However, only 1) an overpressured perturbation with the same Lorentz factor as the underlying flow avoids the front region to be brighter from the beginning, and 2) a rarefaction behind the perturbation avoids the formation of a strong reverse shock behind the perturbation. With these restrictions, the second component fades out rapidly and then the first one dominates the emission, as observed for components E and F in 3C111.  We note that a denser medium injected after the perturbation would lead to the formation of a brighter feature behind component F, produced by the reverse shock formed between the end of the perturbation and the medium injected afterwards. As stated above, this shock appears also in our simulations but it has a negligible effect in terms of emission. Observational support for the inclusion of this tenuous material in the simulations can be found in the prominent emission gap following behind the E/F complex in 3C~111 (see K08). New ejection of emitting material is detected on the time scale of more than 2 years, corresponding to a gap width of up to 2\\,mas in 1999 (cf. K08). We have exagerated the effect of the dilute medium in order to focus on the evolution of the perturbation alone. However, such a dilute medium (a factor 10 less dense than the steady jet) is not needed for the conclusions of this work to be valid: a denser medium, but underdense with respect to the steady jet, would produce a similar effect, with little emission contributed from the reverse shock formed.  Thus, our result explains the evolution of these radio-components on the basis of the scenario explained above. From this work, we can derive a recipe for distinguishing between ejection events composed by denser material only or denser and/or faster material. On top of this, we want to remark that these results are based on the hypothesis that the injection of bulk particles in the jet is reduced after the injection of the perturbation, which could be related to the processes taking place in the accretion-disk/black-hole system, in the line of the results in \\cite{Mar02}, who showed the existence of a relation between the processes taking place in the accretion-disk and the jet \\citep[see][for discussion on this issue]{pe08}.  At present we are following a new radio-flare in this source with denser sampling. The aim is to perform a deeper analysis of the early stages of evolution of the expected radio-components and to compare this evolution with the model presented here."
    },
    "0809/0809.1654_arXiv.txt": {
        "abstract": "We present a highly reliable flux-limited census of 18,949 point sources in the Galactic mid-plane that have intrinsically red mid-infrared colors. These sources were selected from the \\textit{Spitzer Space Telescope} GLIMPSE~I and II surveys of 274\\sqdeg of the Galactic mid-plane, and consist mostly of high- and intermediate-mass young stellar objects (YSOs) and asymptotic giant branch (AGB) stars. The selection criteria were carefully chosen to minimize the effects of position-dependent sensitivity, saturation, and confusion. The distribution of sources on the sky and their location in IRAC and MIPS 24\\microns color-magnitude and color-color space are presented. Using this large sample, we find that YSOs and AGB stars can be mostly separated by simple color-magnitude selection criteria into approximately $50-70$\\percent of YSOs and $30-50$\\percent of AGB stars. Planetary nebulae and background galaxies together represent at most $2-3$\\percent of all the red sources. 1,004 red sources in the GLIMPSE~II region, mostly AGB stars with high mass-loss rates, show significant ($\\ge$0.3\\mag) variability at 4.5 and/or 8.0\\microns. With over 11,000 likely YSOs and over 7,000 likely AGB stars, this is to date the largest uniform census of AGB stars and high- and intermediate mass YSOs in the Milky-Way Galaxy. ",
        "introduction": "The \\textit{Spitzer Space Telescope} has recently completed a number of surveys of the Galactic mid-plane using the InfraRed Array Camera (IRAC; \\citealt{Fazio:04:10}) at 3.6, 4.5, 5.8, and 8.0\\microns, and the Multiband Imaging Photometer for \\textit{Spitzer} (MIPS; \\citealt{Rieke:04:25}) at 24 and 70\\microns. In the context of star formation, the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire \\citep[GLIMPSE;][]{Benjamin:03:953} surveys - which to date includes GLIMPSE~I, GLIMPSE~II, and GLIMPSE 3D - and the MIPS Galaxy (MIPSGAL) surveys - which include the MIPSGAL I and II surveys - have so far been used for studies of individual star formation regions \\citep[e.g.][]{Whitney:04:315,Indebetouw:07:321,Mercer:07:242,Shepherd:07:464,Povich:08}. However, the full potential of these surveys is that they provide a uniform view of the Galactic mid-plane - not only do they cover large and well-studied star formation regions, but they also show the distributed star formation between these regions. Therefore, in addition to focusing on specific regions, a whole continuum of star formation environments can now be studied. When seen in this light, these observations have the potential to revolutionize our view of Galactic star formation.  These surveys are not the first of their kind at mid- and far-infrared wavelengths. Previous major surveys covering the Galactic plane include the InfraRed Astronomical Satellite (IRAS; \\citealt{Neugebauer:84:L1}) all-sky survey in 1983 at 12, 25, 60, and 100\\microns, the Infrared Space Observatory (ISO) Galactic plane survey \\citep{Omont:03:975} at 7 and 15\\microns, and the Midcourse Space Experiment (MSX) survey of the Galactic plane \\citep{Price:01:2819} at 8.28, 12.13, 14.65, and 21.3\\microns. However, the combination of coverage, sensitivity and resolution of the \\textit{Spitzer} observations is unprecedented: the full-width half-maximum (FWHM) of the point spread function (PSF) is 2\\arcsec at 8\\microns, and 6\\arcsec at 24\\microns, compared to the detector resolution of 18.3\\arcsec for MSX 8.28 and 21.3\\microns and a FWHM of approximately 3-5\\arcmin for IRAS 12 and 25\\microns. The point source sensitivity at 8\\microns is 100 and 1,000 times better than MSX 8.28\\microns and IRAS 12\\microns respectively, and the sensitivity at 24\\microns is also approximately 100 and 1,000 times more sensitive than MSX 21.3\\microns and IRAS 25\\microns respectively. At 7\\microns, the sensitivity and resolution of the ISOGAL observations (9\\mJy and 6\\arcsec respectively) approach those of the \\textit{Spitzer} GLIMPSE survey, but the coverage of the ISOGAL survey is only 6\\percent of that of the GLIMPSE and MIPSGAL surveys (16\\sqdeg for ISOGAL versus 274\\sqdeg for GLIMPSE).  These three previous surveys have been used to search for young stellar objects (YSOs), which are brighter and redder at infrared wavelengths than field stars due to thermal emission from circumstellar dust. \\citet{Wood:89:265} selected 1,717 candidate embedded massive stars with UCHII regions from the IRAS data, 1,646 of which are associated with the Galactic plane; \\citet{Felli:02:20971} used the ISOGAL survey in conjunction with radio observations to compile a list of 715 YSO candidates; and the MSX survey was used to compile a Galaxy-wide sample of candidate massive YSOs which were followed up to eliminate contaminants via the Red MSX Source (RMS) survey \\citep{Hoare:04:156}. However, the increased sensitivity of the recent \\textit{Spitzer} observations of the Galactic mid-plane will allow many intermediate mass YSOs and more distant massive YSOs to be seen.  In this and future papers, our aim is to construct a photometrically reliable catalog of sources that is likely to contain many YSOs, and to use it to study the distribution of star formation regions in the Galaxy, the environments in which stars form, and to estimate the present rate of star formation in the Galaxy. The current paper describes the initial compilation of a red source catalog which is photometrically very reliable, and is affected as little as possible by any biases due for example to position-dependent sensitivity or saturation limits. In addition to YSOs, a number of asymptotic giant branch (AGB) stars, which are also red at mid-infrared wavelengths due to the dust that surrounds them, are present in the red source catalog presented here. Because interstellar extinction is low at mid-infrared wavelengths, and because of the selection criteria used for this catalog, these are amongst the reddest and most distant AGB stars in the Galaxy.  In \\S\\ref{sec:observations}, we describe the GLIMPSE observations that are used to compile the red source catalog, the various issues that affect the completeness of the GLIMPSE Catalogs, and the complementary surveys that are used to construct SEDs from 1.25 to 24\\microns. In \\S\\ref{sec:selection}, we describe the selection of the red sources from the GLIMPSE Catalog, including the procedure used to increase the reliability and improve the uniformity of the selection across the Galaxy. In \\S\\ref{sec:analysis} we show the angular distribution and colors of the red sources, we identify over 1,000 red variable sources, and we study the composition of the red source catalog. Finally, in \\S\\ref{sec:summary}, we summarize our findings. ",
        "conclusions": "\\label{sec:summary}  We have compiled a flux-limited catalog of nearly 19,000 sources in the Galactic plane that are intrinsically red at mid-infrared wavelengths, using data from the \\textit{Spitzer} GLIMPSE~I and II surveys and IRAC observations of the Galactic center. The sources were required to satisfy the brightness and quality selection criteria from equation (\\ref{eq:brightness}) and (\\ref{eq:quality}) to improve the reliability of the red source catalog (\\S\\ref{sec:initial_selection}), and were required to have $[4.5]-[8.0]\\ge1$ to be considered `red' (\\S\\ref{sec:red_selection}). The latter criterion was determined to be the most straightforward and reproducible way of selecting red sources (\\S\\ref{sec:qualitative_selection}). In particular, the interstellar extinction law is mostly flat between 4.5 and 8.0\\microns, meaning that any red color between the two bands is likely to be intrinsic rather than due to interstellar extinction.  Independent mosaic photometry was performed on all sources selected by these criteria to validate the GLIMPSE 4.5 and 8.0\\microns photometry, and each source was visually examined to ensure that the independent flux could be trusted. Sources for which reliable independent photometry could not be performed were rejected from the red source catalog. In addition, sources for which the independent flux was more reliable than the original GLIMPSE flux and for which the new fluxes no longer satisfied the brightness and color selection criteria from equations (\\ref{eq:brightness}) and (\\ref{eq:color}) were rejected from the red source catalog. In total, 18,949 red sources satisfied all the selection criteria and were determined to have reliable photometry. The final catalog is given in Table~\\ref{tab:red}, and includes JH\\ks, IRAC, and MIPS 24\\microns magnitudes. It is at least 65\\percent complete (for point sources) and close to 100\\percent reliable, meaning that every source in the catalog is a genuine red source.  The near- and mid-infrared color-magnitude and color-color distribution of the red sources was presented (\\S\\ref{sec:spatial}). One particular feature of the IRAC color-magnitude distribution of the red sources was the presence of two distinct populations, one peaking at bright [4.5] and [8.0] magnitudes ($[8.0]<6.5$ and $[4.5]<8.0$), and one peaking at faint [4.5] and [8.0] magnitudes ($[8.0]>8$ and $[4.5]>8.5$). The bright peak consists mostly of bright AGB stars, while the remaining sources consist mostly of fainter AGB stars and YSOs. Using simple color and magnitude selection criteria, the red sources were separated into three distinct populations (\\S\\ref{sec:decontamination}): `extreme' or `obscured' C- and O-rich AGB stars (\\xagb stars), `standard' C- and O- rich AGB stars (\\sagb stars), and YSOs.  The angular distribution, near- and mid-infrared colors, and magnitudes of these three populations were found to differ. The \\xagb stars appear to be seen at least as far as the Galactic center, show a rapid drop-off with Galactic longitude, and as expected for distant sources, their near-infrared colors therefore suggest higher values of interstellar extinction than the other populations. The \\sagb stars show a somewhat shallower dropoff with Galactic longitude, and show less reddening than the \\xagb stars, consistent with their closer distance. While they are also expected to be seen at least as far as the Galactic center (\\S\\ref{sec:distances}), no concentration of sources is seen at the Galactic center itself, but this could be an artifact due to the requirement for a valid MIPS 24\\microns flux for the classification (\\S\\ref{sec:decontamination}). Finally, the YSOs show a shallow dropoff with Galactic longitude, and their distribution is highly clustered, unlike the two populations of AGB stars. The approximate separation of the three populations suggests that approximately $30-50$\\percent of sources in the red source catalog are likely to be AGB stars, and approximately $50-70$\\percent are likely to be YSOs. The fraction of red sources that are galaxies and PNe was found to be very small, on the order of a few \\percent at most (\\S\\ref{sec:pne} and \\S\\ref{sec:galaxies}).  The AGB stars in the red source catalog are likely to form one of the largest samples of mid-infrared selected AGB stars in the Galaxy to date, with over 7,000 AGB star candidates. In particular, the coverage of the GLIMPSE~II region at two epochs has allowed us to uncover over a thousand sources with significant variability ($>0.3$\\mag) at 4.5 and/or 8.0\\microns, which we identify as \\xagb stars with Mira variability. These represent one fifth of all red sources in the GLIMPSE~II region. Of all the AGB stars in the Galaxy that fall in the GLIMPSE survey area (but are not necessarily detected by \\textit{Spitzer}), the red source catalog is likely to only contain a small fraction ($\\sim1$\\percent) of all sAGB stars, but may contain up to a quarter of all xAGB stars.  In parallel, over 11,000 YSO candidates have been uncovered. These do not provide a \\textit{complete} picture of Galactic star formation as seen by \\textit{Spitzer}. For example, the red source catalog does not include blended sources, extended sources, point source YSOs for which excess emission at 4.5\\microns due to H$_2$ and CO bandhead emission from outflows makes the $[4.5]-[8.0]$ too blue for the selection criterion used in this paper, and sources that are so embedded that they are not detected at IRAC wavelengths. Nevertheless, it is a first step towards a study of star formation - as seen by \\textit{Spitzer} - on a Galactic scale, and is to date the largest consistently selected sample of YSOs in the Milky-Way. These thousands of YSOs trace many previously known and some previously unknown sites of star formation, including large star formation complexes, smaller star formation regions, and dark clouds. From this perspective, the red source catalog can be thought of not only as a large sample of AGB stars and YSOs, but as the most detailed map to date of the birth sites of intermediate and massive stars in the Galactic plane."
    },
    "0809/0809.4592_arXiv.txt": {
        "abstract": "We report {\\it Hubble Space Telescope} imaging obtained 155 days and 449 days after the 2006 outburst of RS Ophiuchi.  Both epochs show evidence of extended emission, consistent with that seen in earlier radio observations, and a maximum expansion rate of $3200\\pm300$ km s$^{-1}$ (in the plane of the sky).  The extended structure is consistent with the remnant having a bipolar morphology with an inclination similar to that determined for the binary. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1979_arXiv.txt": {
        "abstract": "{In this paper we describe our convective hydrocodes for radial stellar pulsation. We adopt the Kuhfu\\ss{} (1986) model of convection, reformulated for the use in stellar pulsation hydrocodes. Physical as well as numerical assumptions of the code are described in detail. Described tests show, that our models are numerically robust and reproduce basic observational constraints.  We discuss the effects of different treatment of some quantities in other pulsation hydrocodes. Our most important finding concerns the treatment of the turbulent source function in convectively stable regions. In our code we allow for negative values of source function in convectively stable zones, which reflects negative buoyancy. However, some authors restrict the source term to non-negative values. We show that this assumption leads to very high turbulent energies in convectively stable regions. The effect looks like overshooting, but it is not, because turbulence is generated by pulsations. Also, turbulent elements do not carry kinetic nor thermal energy, into convectively stable layers. The range of this artificial overshooting (as we shall call it) is as large as 6 local pressure scale heights, leading to unphysical internal damping through the eddy-viscous forces, in deep, convectively stable parts of the star.}{hydrodynamics -- convection -- overshooting -- stars: oscillations  -- methods: numerical} ",
        "introduction": "%   Pulsation hydrocodes play a key role in understanding the variability of classical pulsators: $\\delta$~Cephei and RR~Lyrae stars. They enable us to model the light and radial velocity curves and to study the modal selection in these stars.  The first hydrocodes were purely radiative, as it was believed, that convection should not alter the pulsation properties of the stars, specially close to the blue edge. Indeed, radiative hydrocodes were very successful in reproducing many of the observed features (see \\eg Buchler 1998 and references therein). However several problems remained (Buchler 1998, Kov\\'acs \\& Kanbur 1998), with modeling of the double-mode pulsations being the most severe among them. Success of convective hydrocodes in solving this longstanding problem (Koll\\'ath \\etal 1998, Feuchtinger 1998) focused attention on modeling classical pulsators with convective hydrocodes.  Turbulent convection is an important physical process acting in many types of stars. It is essentially three dimensional process transporting energy through many length-scales: from macroscopic eddy cells to microscopic molecular scales, were energy is dissipated. Stellar convection acts in hard-turbulence regimes (extremely high Rayleigh numbers ($\\gtrsim 10^{12}$) and extremely small Prandtl numbers ($\\sim 10^{-9}$, Gehmeyr \\& Winkler 1992a) which furthermore complicates its modeling. However, many essential features of convection may be described with simple one-dimensional models. Of these, the famous mixing-length theory (MLT) of B\\\"ohm-Vitense (1958) is most commonly used and underlies many other more complicated models. MLT however, is not suitable for stellar pulsation problems, since it is a local and time-independent theory. In pulsating variable stars, large-scale motions of gas interact with smaller-scale turbulent motions and time-dependent models are necessary to describe the coupling between them. Such one-dimensional models of turbulent convection were developed by many authors. Appropriate equations may be obtained through the Reynolds averaging technique (see \\eg Stani\\v{s}i\\'c 1985). All quantities are decomposed into mean and fluctuating parts, for example for velocity and temperature we have $\\boldsymbol{U}=\\bar{\\boldsymbol{U}}+\\boldsymbol{U}'$, $T=\\bar{T}+T'$, respectively. Hydrodynamic equations decouple into mean and fluctuating (turbulent) equations, which are coupled by second order correlations, like $\\boldsymbol{U}'\\boldsymbol{U}'$ or $\\overline{\\boldsymbol{U}'T'}$, that need to be modeled in order to close the system. Such procedure introduces several dimensionless, order of unity, free parameters, usually denoted by $\\alpha$-s. In one-equation models, turbulent equations are reduced to one equation for turbulent energy, $e_t=\\overline{\\boldsymbol{U}'^2/2}$. An excellent review showing in detail the derivations and approximations made in different models, is given by Baker (1987). We will focus on two models that are commonly  adopted in some recent linear as well as non-linear calculations, namely the Stellingwerf model (1982) and the models based on the work of Kuhfu\\ss{} (1986).  Stellingwerf (1982) truncated the set of three turbulent equations derived by Castor (1968, unpublished) to a one-equation model for the turbulent energy, applying MLT motivated closure relations. To model the small-scale turbulent dissipation, Stellingwerf introduced the eddy-viscous pressure term, in an ad hoc way. Kuhfu\\ss{} (1986) model is self-consistent, with all necessary modeling based on physical arguments. Eddy-viscous terms result from first-order modeling of the Reynolds tensor. All the correlations are modeled in a consistent way, using diffusion approximation. This leads to fully differentiable formulation, contrary to the Stellingwerf model, in which effects of buoyant deceleration of the turbulent eddies must be neglected in convectively stable regions. This leads to extreme overshooting in Stellingwerf model as was shown by Gehmeyr \\& Winkler (1992a,b). These authors performed a detailed comparison of both models. They favour the Kuhfu\\ss{} model pointing other shortcomings of the Stellingwerf treatment. Some drawbacks of the original Stellingwerf theory are overcome by Bono \\& Stellingwerf (1992, 1994), who propose a better treatment of convectively stable regions (\\cf Section~5). In our hydrocodes we adopt the Kuhfu\\ss{} model as it is self-consistent and based on firm physical grounds.  In this paper we present detailed description of our convective hydrocodes. In the first part of the paper (Sections 2 and 3) we discuss the turbulent convection model we adopt in our code: physical formulation in Section 2 and numerical implementation in Section~3 and Appendices. In the second part of the paper (Sections~4~and~5) we present results of some tests of our hydrocode (Section~4), and we discuss the effects of different formulations used in other hydrocodes (Section~5). Conclusions are collected in Section~6. ",
        "conclusions": "%  In this paper we describe, our convective hydrocodes for radial stellar pulsation. Convection model we use, is based on the Kuhfu\\ss{} (1986) model. In the first part of the paper we briefly describe the model and list all model equations and necessary quantities. Technical details, concerning numerical schemes and methods we use, are fully described in Section~3 and in Appendices. Many tests we have done, part of which are briefly described in Section~4, prove that our code works properly. Computed models are numerically robust and reproduce well basic features observed in classical Cepheids, like Hertzsprung bump progression.  There are several other hydrocodes, that adopt similar convection model as our code. However, in different hydrocodes different treatments of some quantities, such as the turbulent source function, $S$, or eddy-viscous terms are used. Consequences of these differences were not fully studied up to date. In Section~5 we compare some of these treatments. Our most important finding, concerns the treatment of source function in convectively stable zones. In our hydrocode, we allow for negative source function in convectively stable zones, which reflects negative buoyancy, and is physically well motivated. However, in some other codes (\\eg Florida-Budapest code), source function is restricted to non-negative values. This corresponds to the neglect of buoyant forces in convectively stable layers of the model. Similarly, convective flux is also restricted to non-negative values. We find that such approach has several serious drawbacks, that we list below:  ({\\it i}) Due to neglect of negative buoyancy effects (by assuming non-negative source function), significant turbulent energies are present in convectively stable layers of the model. When negative source function is absent in convectively stable layers, balance between eddy-viscous driving and turbulent dissipation sets the turbulent energies at relatively high level, $10^9-10^{10}$erg/g. We call this phenomenon artificial overshooting, as turbulent energies are generated by pulsations. Also, turbulent eddies do not transport heat into convectively stable layers, as convective flux is equal to 0.  ({\\it ii}) The range of this artificial overshooting is very large. Significant turbulent energies extend to more than 6 local pressure scale heights below the envelope convection zone.  ({\\it iii}) Significant turbulent energies in the deep, convectively stable parts of the model, lead to strong eddy-viscous damping, and consequently to lower pulsation amplitudes of the models in comparison to models computed with negative source function.  ({\\it iv}) Physical overshooting, due to turbulent flux, plays a minor role, and described, high turbulent energies are present in convectively stable layers, regardless if turbulent flux is included in the model or not.   In the next paper (Smolec \\& Moskalik 2008) we will show further consequences of assuming non-negative source function. We show crucial role of this assumption in double-mode Cepheid models computed with the Florida-Budapest code (Koll\\'ath \\etal 2002). We will explain in detail the mechanism that leads to double-mode behaviour and we will show, that this mechanism does not work if source term is treated properly, that is, if we allow for negative values of the source function.   \\Acknow{We are grateful to prof. Wojciech Dziembowski for fruitful discussions, specially concerning Sections~2~and~5. Alosza Pamyatnykh is acknowledged for the permission to use the opacity interpolating subroutines. Differences  between our formulation, and formulation used in Florida-Budapest code were clarified during the RS stay in Budapest. We wish to thank Robert Szab\\'o for computation of several models with the Florida-Budapest hydrocode, used for comparison with our models. We are grateful to Robert Szab\\'o and Zoltan Koll\\'ath for fruitful discussions concerning the Florida-Budapest hydrocode. Hospitality of Konkoly Observatory staff is acknowledged. This work has been supported by the Polish MNiSW Grant No. 1 P03D 011 30.}"
    },
    "0809/0809.4137_arXiv.txt": {
        "abstract": "This article summarizes a Splinter Session at the Cool Stars XV conference in St.~Andrews with 3 review and 4 contributed talks. The speakers have discussed various approaches to understand the structure and evolution of the gas component in protoplanetary disks. These ranged from observational spectroscopy in the UV, infrared and millimeter, through to chemical and hydrodynamical models. The focus was on disks around low-mass stars, ranging from classical T\\,Tauri stars to transitional disks and debris disks. Emphasis was put on water and organic molecules, the relation to planet formation, and the formation of holes and gaps in the inner regions. ",
        "introduction": "Most of the current knowledge about the structure and evolution of protoplanetary disks (photometry, SED class, images in scattered light, etc.) originate from the dust component, whereas the gas in protoplanetary disks is, in many cases, more difficult to observe and more difficult to model (see Fig.~1). Yet, it is the gas component that contains 99\\% of the disk mass and sets the initial conditions for planet formation. The gas is responsible for the delivery of water and organic material to terrestrial planets and, by its dispersal, drives the evolution to debris disks.  With existing and new telescopes ({\\sc Spitzer, Herschel, Sofia, Jwst, Alma}), tracers of the gas in form of emission lines of atoms, ions and molecules can be observed. These observations allow for a direct determination of the gas mass, the temperature distribution and the chemical composition of the gas independent of the dust, provided that we can trust the models. By comparing objects of different ages, the evolution of the gas in protoplanetary disks can be revealed.  Since UV, infrared, and millimeter lines probe very different regions in the disks, observations generally need to be combined to obtain a complete and consistent picture of the physical and chemical state of the gas in protoplanetary disks. Here, models play an essential role to convert observations into understanding. Given these new challenges, astronomers and theoreticians around the world have started to ``hunt'' the gas in protoplanetary disks.  Key questions are: \\begin{itemize} \\item How does the gas-to-dust ratio vary with age, starting from the initial interstellar value of 100:1 to virtually 0:1 in debris disks? \\item What is the density structure, temperature, and dynamics of the gas in protoplanetary disks at different ages? \\item What is the chemical composition of the gas in the regions where terrestrial and where gas giant planets form? \\item What is the driving mechanism for the gas mass loss?  How does the dust component react on the gas removal? \\item Does the gas removal truncate the process of planet formation? \\end{itemize} ",
        "conclusions": "The gas component in protoplanetary disks can be observed by emission lines from the ultraviolet to millimeter. The structure, dynamics, temperature structure and chemical composition of the gas in the disks provide the initial for planet formation. The evolution of the gas component from massive gas-rich disks to almost gas-free debris disk, including the formation of central holes, limits the timescale for giant planet formation, and the delivery of water and organic material to terrestrial planets.  Spectroscopic measurements of gas in disks are generally more challenging than broadband observations of dust both in terms of observations and theory, but they provide complementary information on the disk structure, in addition to unique constraints on kinematics and chemistry. Future facilities, operating at near-infrared to millimeter wavelengths such as {\\sc Jwst, Herschel, Sofia}, and {\\sc Alma}, will open up this field and lead to dramatic advances in our understanding of disk evolution and planet formation."
    },
    "0809/0809.5077_arXiv.txt": {
        "abstract": "We investigate the utility of a new, self-similar pressure profile for fitting Sunyaev-Zel'dovich (SZ) effect observations of galaxy clusters.  Current SZ imaging instruments---such as the Sunyaev-Zel'dovich Array (SZA)---are capable of probing clusters over a large range in physical scale.  A model is therefore required that can accurately describe a cluster's pressure profile over a broad range of radii, from the core of the cluster out to a significant fraction of the virial radius. In the analysis presented here, we fit a radial pressure profile derived from simulations and detailed X-ray analysis of relaxed clusters to SZA observations of three clusters with exceptionally high quality X-ray data: A1835,  A1914, and CL~J1226.9+3332.  From the joint analysis of the SZ and X-ray data, we derive physical properties such as gas mass, total mass, gas fraction and the intrinsic, integrated Compton $y$-parameter. We find that parameters derived from the joint fit to the SZ and X-ray data agree well with a detailed, independent X-ray-only analysis of the same clusters. In particular, we find that, when combined with X-ray imaging data, this new pressure profile yields an independent electron radial temperature profile that is in good agreement with spectroscopic X-ray measurements. ",
        "introduction": "The expansion history of the universe and the growth of large-scale structure are two of the most compelling topics in cosmology.  As galaxy clusters are the largest collapsed objects in the universe, taking a Hubble time to form, their abundance and evolution are critically sensitive to the details of that expansion history. Cluster surveys can therefore provide fundamental clues to the nature and abundance of dark matter and dark energy \\citep[see, e.g.,][]{white1993b,frenk1999,haiman2001,weller2002}.  While clusters have been extensively studied using X-ray observations of the hot gas that comprises the intracluster medium (ICM), radio measurements of the same gas via the Sunyaev-Zel'dovich (SZ) effect \\citep{sunyaev1972} provide an independent and complementary probe of the ICM \\cite[e.g.,][]{carlstrom2002}.  The SZ effect arises from inverse Compton scattering of cosmic microwave background (CMB) photons by the electrons in the ICM, imparting a detectable spectral signature to the CMB that is independent of the redshift of the cluster. As the SZ effect is a measure of the integrated line-of-sight electron density, weighted by temperature, i.e., the integrated pressure (see \\S~\\ref{xray_sze}), it probes different properties than the X-ray emission, which is proportional to the square of the electron density. Cluster surveys exploiting the redshift independence of the SZ effect are now being conducted by a variety of instruments, including the SZA \\citep{muchovej2007}, the South Pole Telescope \\citep{ruhl2004}, the Atacama Cosmology Telescope \\citep{kosowsky2003} and the Atacama Pathfinder Experiment (APEX) SZ instrument \\citep{dobbs2006}.  To maximize the utility of clusters as cosmological probes we must understand how to accurately relate their observable properties to their total masses.  The integrated SZ signal from a cluster is proportional to the total thermal energy of the cluster, and is therefore a measure of the underlying gravitational potential, and ultimately the dark matter content, within a cluster. The SZ flux of a cluster thereby should provide a robust, low scatter proxy for the total cluster mass, \\Mtot\\ \\citep[see, for example,][]{dasilva2004,motl2005,nagai2006,reid2006}.  To date, cosmological studies combining SZA and X-ray data have relied almost exclusively on the isothermal $\\beta$-model \\citep[first used in][]{cavaliere1976,cavaliere1978} to fit the SZ signal in the region interior to $r_{2500}$ \\citep[see][for applications of the isothermal $\\beta$-model to SZ+X-ray data]{grego2000, reese2002,laroque2006,bonamente2006,bonamente2008}. Here $r_{2500}$ is the radius within which the mean cluster density is a factor of 2500 over the critical density of the universe at the cluster redshift.  While the isothermal $\\beta$-model recovers global properties of clusters quite accurately in this regime \\citep{laroque2006}, deep X-ray observations of nearby clusters show that isothermality is a poor description of the cluster outskirts ($r\\sim r_{500}$) \\citep[see e.g.][and references therein]{piffaretti2005,vikhlinin2005a,pratt2007}.  An improved model for the cluster gas, accurate to large radii, is therefore critical for the analysis and cosmological interpretation of SZ data obtained with new instruments that are capable of probing the outer regions ($r \\sim r_{500}$) of clusters. Such a model must be simple enough that it can be constrained by SZ data with limited angular resolution and sensitivity typical of data sets acquired by SZ survey instruments optimized for detection, rather than imaging.  While the $\\beta$-model has the virtue of simplicity, previous attempts to relax the assumption of isothermality typically required high-significance, spatially-resolved X-ray spectroscopy; such data are seldom obtained in short X-ray exposures of high-redshift clusters. This is particularly true for the cluster outskirts \\citep[see, e.g.,][]{laroque2006}. Attempts to move beyond the $\\beta$-model have typically improved the modeling of the gas density only within the core ($r\\lesssim0.15\\,r_{500}$).  In this work, we investigate a new model for the cluster gas pressure by using it to fit SZ data from the SZA and X-ray data from {\\em Chandra} to three well-studied clusters: Abell 1835,  Abell 1914, and CL~J1226.9+3332. The model was derived by \\citet{nagai2007b} from simulations and from detailed analysis of deep {\\it Chandra} measurements of nearby relaxed clusters.  The simplicity of this model---and the fact that SZ data are inherently sensitive to the integrated electron pressure---allow it to be used either in conjunction with X-ray imaging data, or fit to SZ data alone.  The outline of the paper is as follows: in \\S~\\ref{modeling}, we present the details of the model, and couple it with an ICM density model that allows the inclusion of X-ray imaging data.  In \\S~\\ref{realobs}, we apply this method to three clusters, combining new data from the SZA with {\\em Chandra} X-ray imaging data.  We demonstrate the utility of this model by applying it to SZ+X-ray data without relying on X-ray spectroscopic information. Results from the joint SZ+X-ray analysis are then compared to results from an X-ray-only analysis, including spectroscopic data, in \\S~\\ref{results}. We offer our conclusions in \\S~\\ref{conclusions}. ",
        "conclusions": "In this work, we demonstrated the application of a new pressure profile, proposed by \\citet{nagai2007b}, in fitting SZ data taken by the Sunyaev-Zel'dovich Array. By combining the pressure profile constraints from the SZ data with constraints on density from \\emph{Chandra} X-ray imaging, using a simplified form of the density model proposed by \\citet{vikhlinin2006}, we were able to determine the properties of three clusters (A1835,  A1914, and CL~J1226.9+3332) spanning a wide range in redshift and dynamical states.  This technique was compared with two others: a joint analysis of the same SZ + X-ray data relying on the commonly used isothermal $\\beta$-model, which assumes the pressure and density profiles have the same form, and an X-ray only analysis that independently models temperature and density using both X-ray imaging and spectroscopic measurements (but excluding SZ data). We find that the global cluster properties at both $r_{2500}$ and $r_{500}$ determined from ICM profiles fit in the joint pressure analysis are in excellent agreement with properties determined in the X-ray analysis. By contrast, the isothermal $\\beta$-model tends to overestimate with respect to the other models the gas pressure at $r > r_{2500}$, where isothermality is an increasingly poor assumption. The $\\beta$-model thus leads to an overestimate of the cylindrically-integrated Compton $y$-parameter at $r_{500}$. Since the isothermal $\\beta$-model does not provide a good description of the ICM profile in the cluster outskirts, we caution against its use in deriving global properties of clusters even at $r \\sim r_{500}$, which is a large fraction of the virial radius.  We tested the ability to recover the ICM electron pressure profiles from SZ data by analyzing the SZ data together with the X-ray imaging data alone, ignoring the X-ray spectroscopic information. Assuming the ideal gas law, we derive electron temperature profiles by coupling the pressure fits to the SZ data with the density fits to the X-ray imaging data. We find that these derived temperature profiles show broad agreement with those determined spectroscopically from deep X-ray observations, even for the highest redshift cluster in our sample, at $z=0.89$.  This method therefore provides an independent technique for determining the radial electron temperature distribution in high-redshift clusters, for which deep spectroscopic X-ray data may be unavailable and are impractical to obtain."
    },
    "0809/0809.0464_arXiv.txt": {
        "abstract": "We propose a regular classical field theory realisation of the Dvali--Gabadadze--Porrati mechanism by considering our universe to be the four-dimensional core of a seven dimensional 't~Hooft--Polyakov hypermonopole. We show the existence of metastable gravitons trapped in the core. Their mass spectrum is discrete, positive definite, and computed for various values of the field coupling constants: the resulting Newton gravity law is seven-dimensional at small and large distances but can be made four-dimensional on intermediate length scales. There is no need of a cosmological constant in the bulk, the spacetime is asymptotically flat and of infinite volume in the extra-dimensions. Confinement is achieved through the local positive curvature of the extra-dimensions induced by the monopole-forming fields and for natural values of the coupling constants of order unity. ",
        "introduction": "\\label{sec:intro}  Gravity occupies a central role in high energy physics and cosmology. On one hand, the unification of the fundamental interactions in the context of String Theory suggests that we may live in a more than four-dimensional world~\\cite{Nordstrom:1914, Polchinski:1998rq}. On the other hand, the recent acceleration of our universe has been confirmed by different experiments and it is now a widely accepted important result of modern observational cosmology~\\cite{Komatsu:2008hk}. The idea that such an unexplained acceleration may be the signature of extra-dimensions has been intensively explored in the recent years~\\cite{Arkani-Hamed:1998rs, Randall:1999vf, Deffayet:2001pu}. In the Dvali--Gabadadze--Porrati (DGP) model, the extra-dimensions (bulk) may actually be non-compact and of infinite volume~\\cite{Dvali:2000hr, Dvali:2000xg}. Gravitons are reflected back onto our universe (brane) due to a different gravity coupling constant on the brane and in the bulk. The original DGP action in $\\nextra + 4$ dimensions reads \\begin{equation} \\label{eq:action_dgp} S=\\dfrac{\\Mp^2}{2} \\int{|\\sqrt{\\gbrane}|} \\Rbrane \\ud^4 x + \\dfrac{\\Mpbulk^{2+\\nextra}}{2} \\int{ \\sqrt{|g|} R \\ud^{\\nextra+4} X}, \\end{equation} where $\\gbrane$ and $\\Rbrane$ are respectively the determinant and scalar curvature of the induced metric along our brane, while $g$ and $R$ are the corresponding quantities in the bulk. It has been shown that this model could actually explain the observed acceleration of the universe, although some works suggest that it may be spoiled by instabilities~\\cite{Deffayet:2000uy, Deffayet:2006wp, Gregory:2007xy}. Although the form of Eq.~(\\ref{eq:action_dgp}) has been originally explained by quantum effects, ``regularised models'' have been proposed to justify it from a more classical and tractable point of view, free of instabilities. Such an approach has been explored in Refs~\\cite{Kolanovic:2003am, Kolanovic:2003da, Shaposhnikov:2004ds} and shown to confine gravitons by explicitly choosing some profile for $\\Mpbulk(X)$ or $g(X)$. In a complete physical framework, both of these functions are however not free and it is not clear that the DGP mechanism could indeed appear in any classical system. This question is of crucial importance in order to assess the viability of both infinite volume extra-dimensions and instability-free DGP-like mechanism.  In this letter, we answer this question in the context of canonical classical field theory. Our approach is motivated by condensed matter physics: topological defects are a direct consequence of the symmetry breaking mechanism and can model smooth branes~\\cite{Kibble:1976, Akama:1982jy, Rubakov:1983bb}. Assuming the space-time to be seven-dimensional, an $SO(3)$ spontaneous symmetry breaking in $\\nextra=3$ codimensions generically forms 't~Hooft--Polyakov hypermonopoles~\\cite{'tHooft:1974qc,Polyakov:74}. In the following, we prove the existence of a DGP-like mechanism in the core (assumed to be our universe) of such a monopole.  Compared to lower dimensional defects~\\cite{Ringeval:2004ju}, the existence of positively curved $\\nextra-1$ dimensional regions in the bulk is crucial to allow metastable gravitons to be trapped inside the core. Six is indeed the minimal number of spatial dimensions for which there exists a foliation of the extra-dimensions by two-dimensional positively curved surfaces. In order to allow for a varying Planck mass, we have for completeness included a dilaton $\\dilaton$ having a mass $\\md$ in the Einstein frame. In the Jordan frame, the action associated with this system is \\begin{equation} \\label{eq:action_pole} \\begin{aligned} S & = \\dfrac{1}{2 \\kappa^2} \\int \\ue^{\\dilaton}\\sqrt{-g}  \\left[R -g^{AB} \\partial_A \\dilaton \\partial_B \\dilaton - \\dilapot(\\dilaton) \\right] \\ud^7 x \\\\ & + \\int \\sqrt{-g} \\left[-\\dfrac{1}{2} g^{AB} \\calD_A \\higgsvect \\cdot \\calD_{B} \\higgsvect -\\dfrac{1}{4} \\vect{H}_{AB} \\cdot \\vect{H}^{AB} \\right. \\\\ & - \\left. \\dfrac{\\lambda}{8} \\left(\\higgsvect \\cdot \\higgsvect - \\vev^2 \\right)^2 \\right]\\ud^7 x , \\end{aligned} \\end{equation} where the dilaton potential reads $\\dilapot=\\md^2 \\dilaton^2 \\exp(2\\dilaton/5)$. The $SO(3)$ Higgs field $\\higgsvect=\\{\\higgs^a\\}$ is in the triplet representation ($a \\in \\{1,2,3 \\}$). Its vacuum expectation value $\\vev$ breaks $SO(3)$ into $U(1)$. The covariant derivatives $\\calD_A$ enforce gauge invariance and incorporate the gauge fields $\\vect{C}_A=\\{C_A^a\\}$, \\begin{equation} \\calD_A \\higgsvect  = \\partial_A \\higgsvect - q \\vect{\\gauge}_A \\wedge \\higgsvect, \\end{equation} $q$ being the charge, while the field strength tensor $\\vect{H}_{AB}$ is \\begin{equation} \\vect{H}_{AB}  = \\partial_A \\vect{\\gauge}_B - \\partial_B \\vect{\\gauge}_A - q \\vect{\\gauge}_A \\wedge \\vect{\\gauge}_B. \\end{equation} As for the dimensional analysis, we have $[\\kappa^2]=M^{-5}$, $[q]=M^{-3/2}$, $[\\lambda]=M^{-3}$ and $[C_A]=[\\Phi]=[v] = M^{5/2}$. ",
        "conclusions": "We proposed here a canonical classical field theory model which describes a seven-dimensional monopole, at the core of which gravitons get trapped. Their mass spectrum being positive definite, there are no instabilities for the tensor modes. This phenomenon turns out to be a natural way, in the context of field theory, to implement the DGP idea. The required field configurations can be obtained without fine-tuning from a dense set of coupling constant values of order unity. However, to obtain a four-dimensional gravity behaviour over a wide range of length scales, some amount of fine-tuning is certainly required to confine an almost massless mode. Notice that the smaller the graviton mass, the larger the effective four-dimensional Planck mass, possibly addressing the mass hierarchy problem. It would be interesting, if possible, to find a condensed matter system for which this mechanism could be experimentally explored."
    },
    "0809/0809.0187_arXiv.txt": {
        "abstract": "We present a preliminary analysis of the photospheric activity of CoRoT-Exo-2a, a young G7V star accompanied by a transiting hot Jupiter recently discovered by CoRoT.  We apply spot modelling techniques developed for the analysis of the Sun as a star and  capable to extract from CoRoT high precision light curves information on the variation of  the total spotted area and the longitude of active regions along the 142 days of the observations.  This preliminary analysis shows that the active regions form within two active longitudes separated  by about $180^{\\circ}$ and rotating with periods of 4.5221 and 4.5543 days, respectively,  and that the total spotted area oscillates with a period  of about 28.9 days. ",
        "introduction": "CoRoT is a space experiment devoted to asteroseismology and extrasolar planet search through the observations of planetary transits. It has recently discovered a hot Jupiter, CoRoT-Exo-2b,  orbiting with a period of 1.743 days around a main-sequence G7 star which displays a remarkable photospheric activity \\citep{Alonsoetal08,Bouchyetal08}. Given the late spectral type and short rotation period (about 4.5 days) of the star, its activity is regarded as the manifestation of magnetic fields in the atmosphere, amplified and modulated by a hydromagnetic dynamo. In this paper, we present some preliminary results about the spot modelling of such a star, indicated as CoRoT-Exo-2a, which is a good proxy for the young Sun, probably at an age of approximately 0.5 Gyr \\citep{Bouchyetal08} . A detailed account of the results obtained from the spot modelling of the light curve of CoRoT-Exo-2a will be provided in \\citep{Lanzaetal08}. ",
        "conclusions": ""
    },
    "0809/0809.2657_arXiv.txt": {
        "abstract": "We examine the problem of particle acceleration at a relativistic shocks assuming pitch-angle scattering and using a Hartree-Fock method to approximate the associated eigenfunctions. This leads to a simple transcendental equation determining the power-law index, $s$, given the up and downstream velocities. We compare our results with accurate numerical solutions obtained using the eigenfunction method. In addition to the power-law index this method yields the angular and spatial distributions upstream of the shock. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2369_arXiv.txt": {
        "abstract": "The intrinsic geometry of the Kerr ergosurface on constant Boyer-Lindquist (BL), Kerr, and Doran time slices is characterized. Unlike the BL slice, which had been previously studied, the other slices (i) do not have conical singularities at the poles (except the Doran slice in the extremal limit), (ii) have finite polar circumference in the extremal limit, and (iii)  for sufficiently large spin parameter fail to be isometrically embeddable as a surface of revolution above some latitude. The Doran slice develops an embeddable polar cap for spin parameters greater than about 0.96. ",
        "introduction": "The ergosurface of the Kerr spacetime is the locus at which the asymptotic time translation symmetry becomes lightlike. This symmetry is spacelike inside the ergosurface, so negative energy states of matter are allowed there~\\cite{Penrose:1969pc}. The existence of these negative energy states is probably central to the mechanism by which rotational energy is extracted from astrophysical black holes~\\cite{Begelman:1984mw, Meier01}. The Penrose process~\\cite{Penrose:1969pc,Penrose:1971uk} is an idealized example of such a mechanism, whose discovery led directly to the realization that black holes behave as thermodynamic systems in equilibrium, with an entropy proportional to the surface area of the event horizon and a temperature proportional to the surface gravity~\\cite{Bekenstein:1973ur}.  Given the importance of the ergosurface, it is natural to examine its intrinsic two-dimensional (2d) geometry. While this geometry plays no direct role in dynamical processes, it could be useful for physical intuition and visualization, and is in any case mathematically interesting. Unlike the event horizon, however, the ergosurface is not a null surface, and therefore different spacelike slices of it determine different 2d geometries. In this paper we examine and compare the geometry of three different slices, those defined by the  Boyer-Lindquist (BL), Kerr, and Doran~\\cite{Doran:1999gb} time coordinates.  The geometry of the BL slice has already been extensively studied~\\cite{Sharp81,Kokkotas:1988fx,Pelavas:2000za}. Peculiar features of that geometry are that the poles have conical singularities, and that in the limit of maximal rotation (extreme Kerr geometry) the surface becomes infinitely long in the polar direction. This happens because the ergosurface coincides with the horizon at the poles, and the horizon recedes to infinite distance on a surface of constant BL time coordinate in the extremal limit. By contrast, the horizon remains at a finite separation on constant Kerr or Doran time surfaces, since these surfaces are defined by geodesics freely falling across the horizon.  The remainder of this paper is organized as follows. In section 2 we present the Kerr metric in the three different coordinate systems and obtain the 2d metrics for the corresponding ergosurface slices. Section 3 discusses general properties of axisymmetric 2d geometries, and section 4 presents the results for the slices of the Kerr ergosurface. Mathematica was used for most of the computations. ",
        "conclusions": "We have analyzed and compared the geometries of three different spatial slices of the ergosurface of the Kerr spacetime, corresponding to constant values of Boyer-Lindquist, Kerr and Doran time coordinates. We found that, unlike the BL slice which has a conical singularity, both the Kerr and Doran slices are smooth at the pole, except for the extremal case of the Doran slice. Also unlike the BL slice, the distance from equator to pole remains finite as the spin parameter $a$ approaches the extremal value $M$. All the slices develop negative curvature above some latitude for sufficiently high spin, and the Doran slice develops also a positive curvature polar cap above $a\\approx 0.95 M$. In the extremal limit $a=M$ this becomes a conical singularity at the pole, with a deficit angle $2\\pi(1-\\sqrt{2/3})$.  The slices differ in where they can be isometrically embedded as surfaces of revolution in a 3d Euclidean space. The BL slice is fully embeddable for all spin parameters, and the Kerr and Doran slices fail to be embeddable above some latitude when the spin parameter is sufficiently large. For spin parameter $a\\gtrsim 0.96M$ the Doran slice is also embeddable at high latitudes, in a polar cap. We showed some embedding diagrams, and indicated the embeddable regions on plots of circumference vs. radius. These plots allow easy comparison of the geometries of the three slices."
    },
    "0809/0809.4912_arXiv.txt": {
        "abstract": "We present preliminary results and observables from a model of microquasar based on a theoretical framework where stationary, powerful, compact jets are launched and then accelerated from an inner magnetized disk. This model aim at providing a consistent picture of microquasars in all their spectral states. It is composed of an outer standard accretion disk down to a variable transition radius where it changes to a magnetized disk, called the Jet Emitting Disk (JED). The theoretical framework providing the heating, we solve the radiative equilibrium and obtain the JED structure. Our JED solutions are rich, and reproduce the already known scheme where a cold optically-thick and a hot optically-thin solutions bracket a thermally unstable one. We present the model and preliminary results, whith a first attempt at reproducing the observed SED of XTE~J1118+480. ",
        "introduction": "Microquasars are stellar binaries with a short period in which one of the two components is a stellar-mass black-hole. The secondary component is a normal star filling its Roche lobe, and loosing mass through the first Lagrangian point. The ejected matter organizes itself in an accretion disk, from which jets are launched. Microquasar are one piece of a chain of objects with Active Galactic Nuclei and Gamma Ray Bursts, sharing the same ingredients and physics (Mirabel 2004).  Various theoretical works have been made to model the accretion disk, with the goal of reproducing the observations. The first kind of observations that is attempted to be reproduced is the spectral energy distribution (SED), which results from the combination of model {\\it components} and the consistency with which they are combined, and the physical radiative processes occuring therein. Other observables, such as for example time-lag observed in X-rays or Quasi-Periodic Oscillations, are usually being modeled with more or less dedicated work, and reflects more the behavior of a given component (the disk for QPOs, e.g. Tagger \\& Varni\\`ere 2006, and the jet for time-lags, see for instance Kylafis et al. 2008).  Our own approach is the following: based on a {\\it consistent} dynamical framework describing how jets are launched from a disk (see Ferreira 1997 and references therein), we build a microquasar model comprising an outer standard accretion disk down to a transition radius, and an inner disk from which jets are lauched, called the Jet Emitting Disk (JED). This model is a global attempt at explaining the general behavior of microquasars in all their spectral states (see Ferreira et al. 2006, Petrucci et al. 2008). Our goal here is to explore the richness of the solutions given by the radiative equilibrium of the JED, and produce theoretical SEDs that can be tested against observations. ",
        "conclusions": ""
    },
    "0809/0809.1706_arXiv.txt": {
        "abstract": "With the goal of deriving the physical and chemical conditions of star forming regions in the Large Magellanic Cloud (LMC), a spectral line survey of the prominent star forming region N113 is presented. The observations cover parts of the frequency range from 85\\,GHz to 357\\,GHz and include 63 molecular transitions from a total of 16 species, among them spectra of rare isotopologues. Maps of selected molecular lines as well as the 1.2\\,mm continuum distribution are also presented. Molecular abundances in the core of the complex are consistent with a photon dominated region (PDR) in a nitrogen deficient environment. While CO shows optical depths of order $\\tau$$\\sim$10, $^{13}$CO is optically thin. The most prominent lines of CS, HCN, and HCO$^+$ show signs of weak saturation ($\\tau$$\\sim$0.5). Densities range from 5$\\times$10$^3$\\,cm$^{-3}$ for CO to almost 10$^6$ for CS, HCN, and a few other species, indicating that only the densest regions provide sufficient shielding even for some of the most common species. An ortho- to para-H$_2$CO ratio of $\\sim$3 hints at H$_2$CO formation in a warm ($\\ga$40\\,K) environment. Isotope ratios are $^{12}$C/$^{13}$C $\\sim$ 49$\\pm$5, $^{16}$O/$^{18}$O $\\sim$ 2000$\\pm$250, $^{18}$O/$^{17}$O $\\sim$ 1.7$\\pm$0.2 and $^{32}$S/$^{34}$S $\\sim$ 15. Agreement with data from other star forming clouds shows that the gas is well mixed in the LMC. The isotope ratios do not only differ from those seen in the Galaxy. They also do not form a continuation of the trends observed with decreasing metallicity from the inner to the outer Galaxy. This implies that the outer Galaxy, even though showing an intermediate metallicity, is not providing a transition zone between the inner Galaxy and the metal poor environment of the Magellanic Clouds. A part of this discrepancy is likely caused by differences in the age of the stellar populations in the outer Galaxy and the LMC. While, however, this scenario readily explains measured carbon and oxygen isotope ratios, nitrogen and sulfur still lack a self-consistent interpretation. ",
        "introduction": "The Magellanic Clouds are two southern irregular galaxies that provide unique opportunities to study astrophysical processes (e.g., Westerlund 1990). Noteworthy are their extremely small distances ($\\sim$50 and 60\\,kpc), low heavy element contents, low dust-to-gas mass ratios, high [O/C] and [O/N] elemental abundance ratios, high atomic-to-molecular hydrogen (H/H$_2$) ratios, and an intense ultraviolet (UV) and far-ultraviolet (FUV) radiation field. The Magellanic Clouds, which are much smaller than the Milky Way, are characterized by a lower degree of nuclear processing than the Galaxy. Right now, however, they are undergoing an episode of vigorous star formation.  In the past couple of decades there have been an enormous number of studies of the Magellanic Clouds (see, e.g., the IAU Symposia No. 108, 148 190, and 256). The first systematic molecular surveys, covering the central parts of the Magellanic Clouds in the CO $J$=1--0 line were those of Cohen et al. (1988) and Rubio et al. (1991) with angular resolutions of 8\\ffam8. More detailed CO surveys were carried out with the NANTEN telescope at 2\\ffam6 resolution (e.g., Yamaguchi et al. 2001; Mizuno et al. 2006). CO data with even higher resolution, obtained in the $J$=1--0 and 2--1 lines with 50$''$ and 25$''$ beamwidth, were taken with the SEST (Swedish-ESO Submillimeter Telescope). Several articles were published as part of the SEST Key-Program on CO in the Magellanic Clouds (e.g., Israel et al. 2003).  While these studies provide an overall view of the well shielded molecular medium in the Magellanic Clouds, multiline studies investigating the physical and chemical properties of the gas only exist for a very small number of targets in the vicinity of prominent HII regions. Following the pioneering study of Johansson et al. (1994) on N\\,159 in the LMC, Chin et al. (1997, 1998) and Heikkil{\\\"a}, et al. (1998, 1999) published molecular multi-line studies on star forming regions in the Magellanic Clouds. First detections of deuterated molecules outside the Galaxy were reported by Chin et al. (1996b) and Heikkil{\\\"a}, et al. (1997).  N\\,113, located in the central part of the LMC more than 2$^{\\circ}$ west of 30\\,Dor, is hosting the most intense H$_2$O maser of the Magellanic Clouds (Whiteoak \\& Gardner 1986; Lazendic et al. 2002; Oliveira et al. 2006). OH maser emission was also observed (Brooks \\& Whiteoak 1997). An IRAS (InfraRed Astronomy Satellite) point source, IRAS~0513--694, is associated with N\\,113. While Chin et al. (1996b, 1997) and Wong et al. (2006) collected single-dish and interferometric data from some $\\lambda$3\\,mm transitions, a dedicated molecular multiline study of this star forming region covering the $\\lambda$1--3\\,mm band is still missing. To improve our understanding of the interstellar medium (ISM) associated with massive star formation in an environment with low metallicity and to obtain a comprehensive view onto one of the most remarkable star forming regions of the LMC, we mapped the 1.2\\,mm continuum, obtained a spectral survey of the central region, and also mapped the cloud in several molecular lines. ",
        "conclusions": "Toward the prominent star-forming region N\\,113 in the Large Magellanic Cloud, we obtained with the SEST a map of the the $\\lambda$1.2\\,mm dust continuum as well as spectral line data covering 63 transitions from a total of 16 molecular species. These include 50 detections, 7 tentative detections, and 6 undetected transitions. The APEX and SEST telescopes also contribute molecular line maps.  (1) The total mass of the N\\,113 molecular complex is estimated to be a few 10$^5$\\,M$_\\odot$ with column densities of $N$(H$_2)$ $\\sim$ 2--10$\\times$10$^{22}$\\,cm$^{-2}$ for the central 45\\arcsec\\ (11\\,pc) and 24\\arcsec (5.5\\,pc), respectively. Assuming virial equilibrium, we obtain for a 45\\arcsec\\ beam a $N$(CO)/$N$(H$_2$) abundance ratio of about 4$\\times$10$^{-5}$ and a conversion factor of $X$ = $N$(H$_2$)/$I_{\\rm CO 1-0}$ = 3.4$\\times$10$^{20}$\\,cm$^{-2}$ (K\\,km\\,s$^{-1}$)$^{-1}$.  (2) The virial theorem suggests an average density of $n$(H$_2$) $\\sim$ 500\\,cm$^{-3}$ for the central 45\\arcsec\\ of the cloud. From CO we obtain a density of $n$(H$_2$) $\\sim$ 5000, for C$_3$H$_2$ and CH$_3$OH the density becomes several 10$^4$\\,cm$^{-3}$, while other molecular species (CS, HCN, HCO$^+$, and H$_2$CO) mainly trace a gas density of several 10$^5$\\,cm$^{-3}$. This indicates a high degree of clumping. Efficient shielding for most of the molecular species seems to occur in the densest regions only.  (3) Among the molecular species observed, only two remain undetected, NO and HNCO. HC$_3$N is only tentatively detected. All these molecules are nitrogen bearing. Chemically, the N\\,113 molecular complex can be described by a photon dominated region in an environment that lacks nitrogen. An ortho- to para-H$_2$CO column density ratio of $\\sim$3 indicates that at least formaldehyde was formed in a warm ($T_{\\rm kin}$ $\\ga$ 40\\,K) medium.  (4) Comparing isotope ratios with line intensity ratios we find that CO is optically thick ($\\tau$$\\sim$10), while $^{13}$CO is optically thin. The main lines of CS, HCN, and HCO$^+$ are only moderately opaque ($\\tau$$\\sim$0.5).  (5) The interstellar medium of the LMC appears to be well mixed. Carbon, nitrogen, and oxygen isotope ratios demonstrate that the outer Galaxy does not provide a ``bridge'' between the interstellar medium of the solar neighborhood and that of the LMC. This is likely caused by the high age of the stellar population of the LMC relative to that of the outer Galaxy. Adopting this scenario, observed carbon and oxygen isotope ratios are qualitatively understood, while nitrogen and sulfur isotope ratios remain an enigma."
    },
    "0809/0809.4065_arXiv.txt": {
        "abstract": "A non-universal scalar mass supergravity type of model is explored where the first two generation of scalars and the third generation of sleptons may be very massive.  The lighter or vanishing third generation  of  squarks as well as Higgs scalars at the unification scale cause  the radiative electroweak  symmetry breaking constraint to be less  prohibitive. Thus, both FCNC/CP-violation problems as well as the naturalness problem are within control.  We identify a large slepton mass effect in the RGE  of $m_{H_D}^2$ (for the down type of Higgs) that may turn the later negative at the electroweak scale even for a small $\\tan\\beta$.  A hyperbolic branch/focus point like effect is found for $m_A^2$ that may result in very light Higgs spectra.  The  lightest stable  particle is dominantly a bino that pair annihilates via Higgs exchange, giving rise to  a  WMAP  satisfied  relic  density region for all $\\tan\\beta$. Detection  prospects  of  such  LSPs in  the upcoming  dark matter experiments  both of  direct and  indirect types (photon flux) are interesting.  The  Higgs bosons and  the third generation of  squarks are light in  this scenario and these may be  easily probed besides charginos and neutralinos in the early runs of LHC. \\\\ PACS No: 04.65.+e, 13.40Em, 14.60Ef, 13.85.-t, 14.80.Ly ",
        "introduction": "\\label{introduction} Low energy supersymmetry (SUSY)\\cite{SUSY} is one of the most active fields of research for physics beyond the standard model (SM)\\cite{standardmodel}. A minimal extension of the Standard Model when supersymmetry is incorporated is the Minimal Supersymmetric Standard Model (MSSM)\\cite{KaneKingRev,SUSY} that includes two Higgs doublets. The model however has a large number of SUSY breaking parameters and this motivates one into studying models with specific mechanisms for breaking SUSY. The later involves high scale physics input and renormalization group analyses. This in general leads to a large reduction of the number of unknown parameters. The minimal supergravity (mSUGRA)\\cite{msugra} model is one of the well studied SUSY models. It requires a very few input parameters at the gauge coupling unification scale or grand unified theory (GUT) scale, $M_G \\simeq 2\\times 10^{16}$~GeV. The parameters indeed quantify our ignorance of the exact nature of SUSY breaking. The model incorporates radiative breaking of electroweak symmetry (REWSB). The unification scale universal input parameters are: i) the gaugino mass parameter $\\mhalf$, ii) the scalar mass parameter $m_0$ and iii) the tri-linear SUSY breaking parameter $A_0$. Additionally, one has to provide with $\\tanb$, the ratio of Higgs vacuum expectation values and the sign of the Higgsino mixing parameter $\\mu$. Renormalization group evolutions are used to obtain the electroweak scale parameters of MSSM.  There is however no {\\em a priori} necessity of considering such universalities of parameters. Indeed, it is worthwhile to explore scenarios with non-universalities in the scalar or in the gaugino masses at the unification scale\\cite{berez,Nath:1997qm,Cerdeno:2004zj,ellis-all,baer-all-non,baer-higgs,biswarup-sourov-my-uc} . Non-universal scalar masses may appear because of non-flat K\\\"ahler potential\\cite{soni}. However, one must be careful to accommodate the stringent constraints from phenomena involving flavor changing neutral currents (FCNC). Satisfying FCNC constraints demands near-degeneracy of the first two generations of scalar masses but the requirements on the third generation of scalars as well as the Higgs scalars are not so stringent \\cite{FCNC,ArkaniHamed:flavor}. Thus it is seen that FCNC constraints as well as constraints from CP-violating phases (for example, those arising from the electric dipole moments of the electron and neutron) may be managed by introducing multi-TeV scalar masses for the first two generations of scalars\\cite{Gabbiani:1996hi}. The third generations of scalar masses and the Higgs scalar masses however should be adequately light in order to satisfy naturalness\\cite{BG}. There have been several efforts for obtaining the desired features as outlined above. Analyses with Radiatively Generated Inverted Mass Hierarchy Models (RIMH) were made in Ref.\\cite{RIMH,ArkaniHamed:flavor,Baer:2000xa}. In Ref.\\cite{Baer:2000xa} the authors achieved the above-mentioned requirements at the electroweak scale by using $t-b-\\tau$ Yukawa unification with special non-universal relationship among the scalar masses at $M_G$. Here, the Higgs and the third generation of scalar masses are rapidly diminished at the electroweak scale via RG evolutions. However, the Yukawa unification and the consideration of REWSB constrain such models heavily. A second realization of the above idea is partially possible via the hyperbolic branch (HB)/ focus point (FP)\\cite{focus,hyper} scenarios where the scalars may become considerably massive (multi-TeV) in a subset of the typical mSUGRA parameter space satisfying universal boundary conditions while fine-tuning\\cite{BG} still remains small. A third possibility was considered in Ref.\\cite{Barger:1999iv}, where plain decoupling arguments motivated the authors in using explicit splitting at $M_G$ between the scalars belonging to the first two generations and the same of the third generation (along with the Higgs scalars). Here the first two generations of scalars were chosen in the multi-TeV domain whereas the third generation of scalars as well as the Higgs scalars were considered to be in the sub-TeV zone. Additionally, a large value of the tri-linear coupling parameter was chosen in Ref.\\cite{Barger:1999iv} for the first two-generations. Universality of scalar masses in the first two-generations along with a choice of a different scalar mass parameter for the third generation as well as the Higgs scalars, or even splitting of squarks and sleptons within the third generation itself have also been considered in Refs.\\cite{Baer:2004xx,King:2006tf,Cerdeno:2001up}. In Ref.\\cite{King:2006tf} the authors additionally considered non-universality in the gaugino masses and analyzed the fine-tuning aspect of computing the relic density of dark matter in addition to obtaining a parameter zone of the MSSM corresponding to a well tempered neutralino\\cite{{Arkani-Hamed:2006mb}}. Similarly, analyses with only non-universalities in the  Higgs scalar may be seen in Refs.\\cite{ellis-all,baer-higgs,Nath:1997qm,McMullen:2006mj}. A comprehensive set of characteristic possibilities for varieties of non-universal SUGRA scenarios may be seen in Ref.\\cite{pncompLHC}.  In this analysis we explore a SUGRA scenario with non-universal scalar masses that i) allows to have very large first two-generation of scalar masses so as to obey the FCNC and the CP-violation limits easily ({\\em ie.} without requiring any ultra-small phases) and that would not impose any additional price on fine-tuning, ii) spans a large amount of MSSM parameter space satisfying the neutralino relic density constraint from WMAP data by not requiring any delicate mixing of bino and Higgsinos, so that we would find a bino-dominated lightest neutralino for most of the parameter space and, iii) satisfies the Higgs mass lower bound from LEP2 data as well as other low energy constraints. Certainly, a model with REWSB and universal scalar mass like mSUGRA is not friendly to achieve these objectives if we consider a common gaugino mass parameter below a TeV or so.  The above phenomenologically inspired requirements in combination motivate us to introduce non-universality between the third and the first two-generation of scalars as well as the Higgs scalars. This has to be such that the REWSB conditions would not become prohibitive to have a multi-TeV first two-generation of scalars. Here we would like to point out that we would not be able to satisfy the mentioned objectives by considering non-universalities only in the Higgs scalar masses.  A simpler and purely phenomenological attempt in this direction in a universal gaugino mass framework could be to generate all the third generation of scalar masses and the Higgs scalar masses radiatively, starting from zero values at the unification scale, while keeping the first two-generations of scalar masses universal (and this may be in the multi-TeV zone). We point out that starting with a similar range of values for the Higgs scalars as with the third generation of scalars at $M_G$ would be desirable since this would lead to similar range of mass values for all the soft-SUSY breaking terms contributing to REWSB after RG evolutions. However, if we consider all the third generation of scalars at $M_G$ to be light, we would find stau ($\\tilde \\tau$) becoming the lightest stable particle (LSP) or even tachyonic at the electroweak scale. On a similar note, we remind that such an appearance of tachyonic sleptons also arise in the Anomaly Mediated Supersymmetry Breaking (AMSB) analyses, and it is avoided purely by a phenomenologically inspired way of adding an appropriate non-zero mass value to all the scalars at the unification scale in the minimal AMSB scenario\\cite{mAMSB}. Thus with a simple motivation of managing with FCNC and CP-violation, naturalness as well as the dark matter constraints concurrently we consider a non-zero mass value for the third generation of sleptons, while having the masses of squarks of the third generation and the Higgs scalars vanishing at $M_G$. Additionally, for convenience we set the third generation of slepton mass parameters the same as that of the first two-generation of scalars at $M_G$.  Thus, the parameters of our Non-universal Scalar Mass model which will henceforth be called the NUSM model is given by \\beq \\tan\\beta, \\mhalf, A_0, sign(\\mu),~{\\rm and}~m_0 \\eeq where the scalar mass input assignments at $M_G$ are as follows.  \\\\ i) The unification scale mass parameter for the first two-generations of squarks,sleptons and the third generation of sleptons is $m_0$, where $m_0$ is allowed to span up to a very large value.\\\\ ii) The mass parameters for the third generation of squarks and Higgs scalars are set to zero. \\\\ We could have also chosen a non-vanishing value for ii) as long as it is sufficiently small. We may note here that different values of mass parameters for squarks and sleptons at $M_G$ may appear in orbifold models with large threshold corrections. This is related to having different modular weights associated with the squarks and the sleptons in a given generation\\cite{Brignole:1993dj,Cerdeno:2007hn}.  As we will see below, the NUSM model provides with a highly bino-dominated neutralino dark matter over almost its full range of parameter space. It has an interesting feature of having a large funnel region, a region of parameter space where the associated annihilation channels are characterized by the direct channel pole $2m_{\\tilde \\chi_1^0} \\simeq m_A, m_H$. We note that unlike mSUGRA where one finds the funnel region only for a large value of $\\tan\\beta$, here in the NUSM one finds it for almost all possible values of $\\tan\\beta (\\gsim 5)$. We will also see that the NUSM  is further characterized by a lighter $m_A$ or $m_H$ particularly when $m_0$ is large and this is found even for small values of $\\tan\\beta (\\sim 10)$. A small part of the parameter space may be far from the decoupling\\cite{decouplingHiggs} region of Higgs boson mixing that causes the lighter CP-even Higgs boson ($h$-boson) to be non-Standard Model like\\cite{nondecouplingHiggs}. This in turn reduces the lower bound of $m_h$ much below the LEP2 Higgs boson mass limit.  In this work we include a semi-analytic calculation that first points out a large mass effect in the solution of the renormalization group equations (RGE) of $m_{H_D}^2$.  The effect causes a hyperbolic branch/focus point like behavior\\footnote{We remind that the well known HB/FP effect that occurs in $m_{H_U}^2$ is associated with the first minimization condition of Radiative electroweak symmetry breaking. In this analysis we do not have such an effect. On the contrary we have a similar HB/FP effect for small values of $\\tan\\beta$ in connection with the second minimization condition that is associated with $m_A^2$.} in $m_A^2 (\\simeq m_{H_D}^2 - m_{H_U}^2)$ that may cause $m_A$ to become smaller even for a small $\\tan\\beta$ as mentioned above. Here we remind that $m_A$ typically becomes smaller in mSUGRA only for large $\\tan\\beta$.  The paper is organized as follows.  In Sec.\\ref{nusmhbfp} we will primarily discuss the large mass RGE effects in the NUSM  on $m_A^2$ and its consequent reduction for large $m_0$ via an HB/FP-like effect. We will obtain the semi-analytic results which will also be verified by numerical computation. In Sec.\\ref{cdmandfunnel} we will study the cold dark matter constraint from WMAP data including also the constraints from the LEP2 Higgs bound, $b \\rightarrow s+\\gamma$, and $ B_s \\to \\mu^+ \\mu^- $. We will also discuss a few sample points in the context of LHC reach. In Sec.\\ref{detectionrates} we will discuss the direct and indirect detection (via continuous gamma ray) rates of the LSP. Finally we will conclude in Sec.\\ref{concl}. ",
        "conclusions": "\\label{concl} In this analysis we have worked with a non-universal scalar mass scenario in a supergravity framework. We started with purely phenomenological motivations namely, i) to manage the FCNC and CP-violation type of constraints by decoupling, ii) to obtain WMAP satisfied values for neutralino relic density for a broad region of parameter space without depending on any delicate mixing of gauginos and Higgsinos, iii) to have radiative electroweak symmetry breaking and iv) to keep naturalness within control. Keeping the above in mind and considering a unified gaugino mass scenario, we used a common scalar mass parameter $m_0$ at the gauge coupling unification scale for the first two-generation of scalars as well as the third generation of sleptons. The item (i) mentioned above would require $m_0$ to be large, and the item (iv) would prefer light third generation of squarks and light Higgs scalars. For simplicity we used vanishing third generation of squarks and Higgs scalar masses at the unification scale. In such a scenario with a possibly multi-TeV $m_0$, we first found a large mass effect or in particular large slepton mass effect in the RGE of $m_{H_D}^2$ ({\\em ie.} for the down type of Higgs scalar) that turns the later negative at the electroweak scale almost irrespective of a value of $\\tan\\beta$. Extending the semi-analytic solution of $m_A^2$ by considering the $h_b$ and the $h_\\tau$ terms (this is required here even for a small $\\tan\\beta$) relevant for the large slepton mass effect we found a hyperbolic branch/focus point like effect in $m_A^2~(\\simeq m_{H_D}^2 - m_{H_U}^2)$ for small $\\tan\\beta$. This causes $m_A$ to be almost independent of $m_0$ for a large domain of the later. But, with a very large $m_0$ this causes $m_A$ to become very light or this may even turn $m_A^2$ negative giving rise to no radiative electroweak symmetry breaking. We further found that because of such large slepton mass effect in the RGE, the Higgs sector may reach an intense coupling region with all the Higgs bosons becoming very light and this may evade the LEP2 limit of $m_h$ for a limited region of parameter space. However, constraint from $Br(B_s\\rightarrow \\mu^+ \\mu^-)$ becomes stringent in this region because of lighter $m_A$.  In general, we find relatively lighter $m_A$ or $m_H$ and this fact leads to a large $A$-pole annihilation region or funnel region of dark matter even for a small $\\tan\\beta$. This is in contrast to minimal supergravity type of scenarios where funnel region may occur only for large values of $\\tan\\beta$. The nature of the LSP is bino-dominated, thus there is no need of any delicate mixing of binos and Higgsinos in order to satisfy the neutralino relic density constraint. We have also computed the direct detection rates of LSP-nucleon scattering. The upcoming detectors like XENON-1T would be able to probe almost the entire region of parameter space. We have further estimated the indirect detection prospect via computing continuous photon fluxes. The ongoing GLAST experiment will be succesfully able to probe the parameter space even for a less cuspy halo profile. We have also briefly discussed the detection prospect of sparticles in the LHC. The  Higgs bosons and  the third generation of squarks are light in this scenario. In addition to charginos and neutralinos the above may be  easily probed in the early runs of LHC."
    },
    "0809/0809.3473_arXiv.txt": {
        "abstract": "{} {We study the molecular content and chemistry of a circumstellar disk surrounding the Herbig Ae star AB Aur at (sub-)millimeter wavelengths. Our aim is to reconstruct the chemical history and composition of the AB Aur disk and to compare it with disks around low-mass, cooler T Tauri stars. } {We observe the AB Aur disk with the IRAM Plateau de Bure Interferometer in the C- and D- configurations in rotational lines of CS, HCN, C$_2$H, CH$_3$OH, HCO$^+$, and CO isotopes. Using an iterative minimization technique, observed columns densities and abundances are derived. These values are further compared with results of an advanced chemical model that is based on a steady-state flared disk structure with a vertical temperature gradient, and gas-grain chemical network with surface reactions. } {We firmly detect HCO$^+$ in the 1--0 transition, tentatively detect HCN, and do not detect CS, C$_2$H, and CH$_3$OH. The observed HCO$^+$ and $^{13}$CO column densities as well as the upper limits to the column densities of HCN, CS, C$_2$H, and CH$_3$OH are in good agreement with modeling results and those from previous studies. } {The AB Aur disk possesses more CO, but is less abundant in other molecular species compared to the DM Tau disk. This is primarily caused by intense UV irradiation from the central Herbig A0 star, which results in a hotter disk where CO freeze out does not occur and thus surface formation of complex CO-bearing molecules might be inhibited.} ",
        "introduction": "The rich variety of the detected exoplanetary systems cannot be fully understood without knowledge about the physical and chemical evolution of their precursors -- protoplanetary disks. The question  what role chemistry  plays during planet formation and how it is coupled to disk dynamics remains loosly constrained.  Nowadays it is widely believed that disk evolution is controlled by redistribution of angular momentum due to turbulent viscosity. A promising source of turbulence is the magnetorotational instability that works if the disk matter is sufficiently ionized \\citep{BH_91}. The chemistry of ionization fraction in protoplanetary disks was studied in detail, both theoretically and observationally \\citep[see, e.g.,][]{Gammie_96,GNI_97,IG_99,Red2,Ilg06a,Pascucci_ea07}. To a large extent it is dominated by stellar radiation in the upper disk region and high-energy cosmic ray particles in their midplanes. For the inner, planet-forming disk regions, X-ray radiation from a young star may play a crucial role. Recently, \\citet{Tur07} % investigated the evolution of the ionization degree in the inner region of a protoplanetary disk, using a coupled 3D magneto-hydrodynamical and chemical code and found that dynamic transport increases electron concentration within the dead zone. It has been widely discussed that the interplay between chemical and transport processes must have been important during the formation of our own solar system \\citep[e.g.,][]{Cyr98,BoMo02,Woo05}.  \\begin{table*}[th]% \\caption{List of observations using the Plateau de Bure Interferometer. In the column of the configuration, the abbreviation of  e.g.\\ ``5D'' refers to the number of antennas and the used configuration. } \\label{obspara} \\centering \\begin{tabular}{c c c c c c c c c} \\hline\\hline  % Line  & transition &  frequency &configu- & baseline  & integration   &  synthezied beam  & 1 $\\sigma$ rms\\\\ &                    &     [MHz]       &ration     &  range [m] &  time ($T_{\\rm ON}$)  [h]       &   size, position angle &  [mJy/beam] \\\\ \\hline HCO$^+$     &1--0                                                     & 89188.523 & 5D,C2   &  24 -- 175  & 5.9+10   &    5.2$''$ \\x\\ 4.8$''$, 6$^\\circ$ &  18 \\\\          % HCN       &  $J$=1$-$0, $F$=2$-$1                        &  88631.847&6CD &  24 -- 175 & 5.9           & 6.0$''$ \\x\\ 4.2$''$, 105$^\\circ$  &  24 \\\\      % CS                 & 2--1                                                    & 97980.950 & 5D    &  24 -- 160   & 1.8          & 9.0$''$ \\x\\ 4.2$''$, 120$^\\circ$   & 170\\\\        % C$_2$H & $N$ = 1--0,  $J$ = 3/2--1/2, $F$=2--1 & 87316.925 & 6D    & 24 -- 113 & 1        &  9.8$''$ \\x\\ 3.9$''$, 110$^\\circ$  & 37 \\\\          % CH$_3$OH &15$_{(3,13)}-$14$_{(4,10)}$ A$^+$    & 88594.960 & 6CD &   24 -- 175 & 3.9         &  6.1$''$ \\x\\ 4.2$''$, 105$^\\circ$  &  24\\\\       % CH$_3$OH &15$_{(3,12)}-$14$_{(4,11)}$ A$^-$     & 88940.090 & 5D    &   24 -- 175 & 3.9+2.2  &  6.3$''$ \\x\\ 4.3$''$, 105$^\\circ$  &  24\\\\         % CH$_3$OH &   2$_{(1,1)}-1_{(1,0)} A^-$                    & 97582.830 & 5D    &   32 -- 73   & 1.8          &  9.4$''$ \\x\\ 4.3$''$, 148$^\\circ$   & 170  \\\\        % \\hline \\end{tabular} \\end{table*}   %    Due to high computational demands, the studies of disk chemistry coupled to disk dynamics are in their early stages \\citep{Ilgner_etal2004,Willacy_etal2006,Semenov_etal2006,TG07}. Chemical models were successfully applied to study disk chemistry for steady-state accretion disks \\citep[e.g.,][]{Will98, vanDishoeck_Blake1998,Aikawa_Herbst1999,Markwick_ea02,Sem05,Dut07}. The most important result is the chemical stratification in such disks, where many molecules are abundant at an intermediate, slightly UV irradiated layer, while being frozen out in the cold and dark midplane and broken apart by unshielded high-energy radiation in disk atmospheres. \\citet{Bergin_ea03} have shown that the non-thermal UV radiation from T Tauri stars may partly come in Ly$_\\alpha$ photons, which significantly affects abundances of molecules dissociated by these photons. In contrast, Herbig Be/Ae stars emit much stronger thermal UV radiation, with circumstellar disks being hotter and more ionized than the disks around T Tauri stars.  A detailed understanding of physical and chemical conditions of the planet-forming environment requires not only sophisticated models, but also high-resolution observations of protoplanetary disks, both in dust continuum and various molecular lines. Unfortunately, such studies are challenging due to the limited sensitivity and spatial resolution of available observational facilities, and small disk sizes ($\\sim 100-1\\,000$~AU). \\citet{Dut97} and \\citet{Kas97} have first detected several molecular species towards protoplanetary disks by using the IRAM 30m single-dish telescope, followed by observations of \\citet{vZad01,Thi04,Sem05}. Interferometric observations by \\citet{Qi01,Aik03,Pie06,Pie07} allowed to study disk gas at outer radii $\\ge 50$~AU in lines of several abundant species (CO, $^{13}$CO, C$^{18}$O, \\hco, CS, CN, H$_2$CO). These studies have confirmed that photochemistry plays an important role even at large disk radii and that abundances of many gas-phase molecules are depleted.  In 2005 a joint Heidelberg-Bordeaux ``Chemistry In Disks'' (CID) project was established to investigate the spatial distribution of various molecular tracers across well-studied T~Tauri and Herbig~Ae disks of various age, followed by comprehensive modeling. In the first paper of this series, we have presented the results of a deep search for N$_2$H$^+$ and HCO$^+$ towards two T Tauri stars (DM~Tau, LkCa~15) and one Herbig Ae star (MWC~480), see \\citet{Dut07}. The N$_2$H$^+$ emission has been detected in LkCa 15 and DM Tau, with the N$_2$H$^+$ to HCO$^+$ ratio similar to that of cold dense cores and the disk ionization degree as predicted by chemical models. In this second paper, we investigate the molecular content of the well-studied disk around the Herbig Ae star AB Aurigae \\citep[][distance = 145 pc]{vdAnc98,Gra99,Rob01,Fuk04}, which is an intermediate-mass analog of the young T~Tauri stars and evolutionary precursor of the main-sequence debris disk sources like $\\beta$\\,Pic, $\\alpha$\\,Lyr, and $\\alpha$\\,PsA. This disk is peculiar in comparison with other known circumstellar disks since it is still embedded in a rather extended envelope and has clumpy density sub-structures \\citep{Fuk04,Cor05,Pie05,Lin06}. This points toward a low-mass companion or a giant planet \\citep{Rod07}, a recent stellar encounter, or imprint of earlier non-steady evolution (see discussion in \\citet{Pie05}, hereafter Paper~II). Using the IRAM 30-m telescope, \\citet[][hereafter Paper I]{Sem05} have studied the molecular content towards this source at low resolution and derived basic parameters of the envelope. They have derived a low disk inclination of about $17^{\\rm o}$ (face-on orientation) and low disk mass of about 0.013\\,$M_{\\sun}$. \\citet[][]{Pie05} have resolved the disk structure in different CO isotopes using the Plateau de Bure Interferometer (PdBI). They have found that the disk is inclined by about $30^{\\rm o}$ and its rotation departs from the Keplerian law, with an exponent for the rotation velocity as low as 0.41. The outer radius of the gaseous disk is close to $\\sim 1000$~AU, its inner hole is at a radius of  $\\approx$\\,100~AU in the dust emission,  and at $\\approx$\\,70~AU in CO lines, and the disk mass is $\\approx$ 0.02\\,$M_{\\sun}$.   The main aim of the present study is to search for C$_2$H, CS, HCN, HCO$^+$, CH$_3$OH, and CO isotopes in the disk of AB Aur using compact interferometer configurations of the PdBI, to derive their column densities, and to compare these values with those from a robust chemical model and the disk of a cooler T~Tauri star DM~Tau. ",
        "conclusions": "The AB Aur environment is clearly dominated by the envelope, as shown by the detections made with the 30-m telescope and the lack of detections with the IRAM array in short (1-3 hours) integrations. The disk has an inner hole in the gas distribution with a radius of $\\approx$ 70 AU. Using a gas-grain chemical model with surface reactions coupled to a flared passive disk model, we reproduce all observed column densities and upper limits. Note that such a comparison between the advanced chemical model and the interferometric observations are hampered by the lack of knowledge on the H$_2$ surface density. The surface density of the AB Aur disk is derived from the millimeter dust emission, a procedure prone to uncertainties. Part of the discrepancies between the model and the observations may also be due to inhomogeneous structure of the AB Aur disk and/or grain evolution in central disk regions that is not considered in the calculations. Modeled and observed column densities relative to $^{13}$CO are both lower for the AB Aur disk than the values measured in DM Tau. The absolute amounts of CO gas are similar in  both disks. The poor molecular content of the AB Aur system can be explained by much more intense UV irradiation by the central A0 star and less effective dust shielding in the low-mass disk compared to the more massive disk around the less luminous M1 star DM Tau. The indirectly inferred cosmic ray ionization rate that is needed to fit the HCO$^+$ data is a factor of 2--5 lower than the standard value, which can be explained by scattering of CRPs in a magnetized AB Aur disk or a clumpy disk structure."
    },
    "0809/0809.2580_arXiv.txt": {
        "abstract": "We present the longest, by a factor of two, near-infrared lightcurve from Sgr~A* -- the supermassive black hole in the Galactic center. Achieved by combining Keck and VLT data from one common night, which fortuitously had simultaneous Chandra and SMA data, this lightcurve is used to address  two outstanding problems.  First, a putative quasi-periodicity of $\\sim$20\\,min reported by groups using ESO's VLT is not confirmed by Keck observations.   Second, while the infrared and  mm-regimes  are thought to be related based on reported time lags between lightcurves from the two wavelength domains, the reported time lag of 20\\,min inferred using the Keck data of this common VLT/Keck night only is at odds with the lag of $\\sim100$\\,min reported earlier.  With our long lightcurve, we find that (i) the simultaneous 1.3 millimeter observations are in fact consistent with a $\\sim100$\\,min time lag, (ii) the different methods of NIR photometry used by the VLT and Keck groups lead to consistent results, (iii) the Lomb-Scargle periodogram of the whole NIR lightcurve is featureless and follows a power-law with slope -1.6, and (iv) scanning the lightcurve with a sliding window to look for a transient QPO phenomenon reveals for a certain part of the lightcurve a 25\\,min peak in the periodogram. Using Monte Carlo simulations and taking the number of trials into account, we find it to be insignificant. ",
        "introduction": "The near-infrared (NIR) regime at high angular resolution has proven to be of great value in Galactic center research. The proper motion of stars detected in this waveband demonstrates the existence of a supermassive black hole (BH) at the center of our Galaxy: Sagittarius~A* \\citep[Sgr~A*; see, e.g.,][]{eckigenzel, ghez98, gheznature,ghez,genzel00,rainer1}. In 2003, NIR emission associated with Sgr~A* was detected \\citep{genzel,ghez04}, which is important since it is the most underluminous BH accretion system observed thus far (with a bolometric luminosity nine orders of magnitude lower than its Eddington luminosity). The NIR emission is highly variable: intensity changes by factors $\\leq 10$, lasting between about 10 and 100 min, occur at least 4 times a day, and are highly polarized \\citep[e.g.][]{ecki1,ecki2,ich2,ich,tuan,trippe}. Simultaneous NIR and X-ray observations have revealed that each X-ray flare is accompanied by a NIR flare with zero time lag, but not vice versa \\citep{ecki04,ecki1,ecki08,belanger05,hornstein07}. Recent campaigns that included mm observations, reported a characteristic time lag between the NIR/X-ray and longer wavelengths that has been interpreted in terms of an expanding synchrotron plasmon. The limited overlap between data sets, however, has led to debates over the nature of this time lag \\citep[e.g.][]{ecki1,yusef06,yusef07,danmultwav}.  The characteristics of the NIR emission from Sgr~A* is of particular interest given that there have been reports of periodic modulations with a period of $\\sim$20\\,min present in NIR flares detected during VLT observations \\citep{genzel,ecki2,ich2,ich,trippe}. In a recent study, however, \\citet{tuan} used robust statistical estimators and found no significant peaks in the periodograms of the Sgr~A* lightcurves observed with the Keck II telescope.  Moreover, quasi-periodicities in the X-ray regime claimed to be present by \\citet{aschenbach1}, are not statistically significant \\citep{belanger}.  In this Letter, we address the questions of a periodic component in the NIR flux, and the relation of the NIR to the mm-regime, by combining for the first time contiguous Keck and VLT data. During that 10 hour session, there was also simultaneous coverage by Chandra and SMA. These observations were published by \\citet{danmultwav}, but interpreted taking only the Keck data into account. ",
        "conclusions": "In this Letter we report on a 600 minute NIR lightcurve of Sgr~A* that combines Keck II and VLT data. We showed that the Lomb-Scargle periodogram of the overall lightcurve is featureless, and is statistically consistent with a single stochastic process that has a power-law spectrum of index $-1.6$. A certain subset of the data has a prominent peak in its corresponding periodogram which, however, is not significant when analyzed in the context of the whole lightcurve. This points to the following dilemma in Sgr~A* research: if a periodic component exists, it is clearly a weak and transient phenomenon that persists over very few cycles, and probably has an evolving period. As we have shown, certain parts of longer pure red noise lightcurves can easily mimic such a behavior. To be able to distinguish between both scenarios, many lightcurves are needed and the range of frequencies under consideration must be narrowed down, e.g. to $0.04 - 0.07\\;\\mbox{min}^{-1}$ if a noticeable peak continues to occur in this range only.  We furthermore showed the difficulty with establishing the simple expanding plasmon model frequently proposed to relate NIR and mm-flares. Sgr~A* is such an active source in the NIR (and maybe not every NIR flare has a millimeter counterpart) that very long simultaneous observations are needed for a meaningful cross-correlation analysis."
    },
    "0809/0809.0585_arXiv.txt": {
        "abstract": "% In this chapter we review the properties of the Orion outlying clouds at $b < -21^\\circ$. These clouds are located far off the Orion giant molecular cloud complex and are in most cases small cometary-shaped clouds, with their head pointing back towards the main Orion clouds. A wealth of data indicate that star formation is ongoing in many of these clouds. The star formation in these regions might have been triggered due to the strong impact of the massive stars in the Orion OB association. Some of the clouds discussed here may be part of the Orion-Eridanus bubble. An overview on each individual cloud is given. A synthesis of the Pre-Main Sequence stars discovered in these clouds is presented. We also discuss the millimeter and centimeter data and present a review of the outflows and Herbig-Haro objects so far discovered in these clouds. ",
        "introduction": "\\label{intro}  The Orion Outlying Clouds  are typically small clouds located in the outskirts of the Orion molecular cloud complex. Most of the known clouds of this kind are located to the west of the Orion OB association. They often appear as small cometary clouds, which point back towards the Orion OB association and in which star formation might have been triggered due to the strong impact of the massive stars in the Orion OB association. The illumination of dense clumps in molecular clouds by OB stars could be responsible for their collapse and subsequent star formation. The UV radiation from the OB stars may sweep the molecular cloud material into a cometary shape with a dense core located at the head of the cometary clouds. Bright Rimmed Clouds (BRCs) associated with HII regions are examples where the UV flux of a nearby OB star ionizes the external layers of the cloud and causes the BRC to collapse. The shape of a BRC (curved or cometary-like) is governed by the ionization fronts from the OB stars that, compressing the head of the BRC, enhance the density of the outer layers of the cloud. This process may lead to sequential star formation. \\citet{Sugi91} and \\citet{Sugi94} catalogue 89 BRCs associated with IRAS point sources in both the northern and southern hemispheres, some of which are also associated with Herbig-Haro objects and molecular outflows. These BRCs are sites in which triggered star formation might have taken place.   \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=\\textwidth]{FIG/fig1_13_34cm.jpg} \\caption{Map of the 100~$\\mu$m IRAS dust emission of the Orion OB association. The main Orion molecular clouds A and B are indicated. The Orion outlying clouds discussed in this chapter are labelled. The big circles represent the  ROSAT X-ray clumps that satisfy the \\citet{Sterz95} criteria.} \\label{clouds} \\end{center} \\end{figure}  \\begin{table}[!ht] \\vspace{8mm} \\caption{\\label{tab:coordinates} Orion outlying clouds discussed in this chapter} \\label{coordinates} \\begin{center} {\\small \\begin{tabular}{lccccll} \\tableline \\noalign{\\smallskip} Cloud & $\\alpha$(2000)   & $\\delta$(2000)  & $l$      & $b$      &  other IDs  & refs.  \\\\ & h~~~m~~~s   & $^{\\circ}$~~~~'~~~~''& ($^\\circ$) & ($^\\circ$) &         &        \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} CB~29               & 05:22:12 & --03:41:34  & 205.8 & --21.5 &                  & 1     \\\\ L~1634              & 05:19:49 & --05:52:05  & 207.6 & --23.0 & MBM~110          & 2, 3  \\\\ L\\,1616$^\\dagger$    & 05:07:00 & --03:21:06  & 203.5 & --24.7 &                  & 2     \\\\ L\\,1615$^\\dagger$    & 05:05:30 & --03:26:00  & 203.4 & --25.1 &                  & 2     \\\\ CB~28               & 05:06:20 & --03:56:22  & 204.0 & --25.1 & LBN~923          & 1, 2  \\\\ IC~2118             & 05:03:55 & --08:23:59  & 208.1 & --27.7 &                  & 4     \\\\ LBN~991$^\\dagger$   & 05:11:00 & --12:24:00  & 213.1 & --27.8 &                  & 2     \\\\ LBN~917             & 04:47:42 & --05:55:00  & 203.5 & --30.1 & DIR~203$-$32     & 2, 5  \\\\ LBN~906             & 04:41:00 & --05:24:00  & 202.1 & --31.4 &                  & 2     \\\\ L~1642$^\\dagger$    & 04:35:03 & --14:13:57  & 210.9 & --36.6 & LBN~981, MBM~20  & 2, 3  \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{7}{l}{\\parbox{0.9\\textwidth}{\\footnotesize $^\\dagger$ coinciding with an X-ray clump (see Sect.~1).}}\\\\[2ex] \\multicolumn{7}{l}{\\parbox{0.9\\textwidth}{\\footnotesize References: 1:~\\citet{Clem88}; 2:~\\citet{Lyn62}; 3:~\\citet{MagBliMun85}; 4:~\\citet{Drey1908}: {\\it Index Catalog}; 5:~\\citet{reach98}   }}\\\\ \\end{tabular} } \\begin{flushleft} \\end{flushleft}  \\end{center} \\end{table}  In this chapter we discuss the outlying clouds in Orion at $b < -21^\\circ$, focusing mainly on the four best studied cases, namely L\\,1615/L\\,1616, L\\,1634, the IC~2118 region and L\\,1642. However, few aspects of other high-galactic latitude clouds are also mentioned. Some of these clouds may be inside the so called Orion-Eridanus bubble, a cavity of the interstellar medium towards Orion which is filled with hot ionized gas surrounded by an expanding shell of neutral hydrogen. For a detailed description of the Orion-Eridanus bubble we refer the reader to \\citet{Brown95}. We use the 100~$\\mu$m IRAS map in order to identify the Orion outlying clouds discussed here and refer to Figure~\\ref{clouds} for their spatial location and to Table~\\ref{tab:coordinates} for their approximate coordinates, in order of decreasing Galactic latitude.  In addition to the Orion outlying clouds we also discuss a few topics of the four small bright-rimmed clouds and cometary globules No. 27, 35, 40 and 41 by  \\citet{Ogu98}, which are located in the vicinity of the O7~V star $\\sigma$~Ori.  The PMS objects in each of the outlying clouds and globules are also discussed. A number of H$\\alpha$-emission stars are found on the clouds, but other are scattered around them. Their coordinates, designations and other informations are synthesized and reported in a table (Table~\\ref{tab:tts}). For many of these stars a proper motion determination is provided in the catalogue by \\citet{Ducour05}. The objects in the Orion OB1a association are not discussed here, but they are treated in other chapters of this book. While some H$\\alpha$-emission stars in the outer regions of Orion were discovered in the objective-prism survey by \\citet{Step86}, many others were revealed in the Kiso H$\\alpha$ survey  \\citep{Wir91, Nak95}. Other weak H$\\alpha$-emission stars in these regions were identified as optical counterparts of the ROSAT All-Sky Survey sources in spectroscopic follow-ups \\citep[e.g.][]{Alc96, Alc00}.  The complete sky coverage of the ROSAT All-Sky Survey  permits an unbiased analysis of the spatial distribution of X-ray active stars, though it is flux limited. \\citet{Sterz95} established a criterion for selecting young star candidates from ROSAT All-Sky Survey sources based on X-ray hardness ratios and the ratio of X-ray to optical flux, and used this to trace their spatial distribution in a $\\sim$ 700~deg$^2$ field around the Orion molecular clouds. The surface density distribution reproduces the major clusters associated with the OB sub-group associations (OB1a, OB1b, OB1c, and $\\lambda$-Ori), and the spatial extent of the clusters is consistent with dispersal times between 2 and 10~Myrs, e.g. the respective ages of the stellar components in those regions. The same analysis revealed several overdensities of X-ray sources with high probability of being low-mass PMS stars \\citep{Sterz95, Walt00}. Several of these X-ray clumps are located far from the main Orion molecular clouds and some of them are coincident with or very close in position to the Orion outlying clouds, in particular L\\,1615/L\\,1616, LBN~991 and L\\,1642. Some of these clumps are indicated with circles in the 100~$\\mu$m IRAS map in Figure~\\ref{clouds}. It is interesting to note that L\\,1634 and IC~2118 were not revealed as X-ray enhancements in the analysis by \\citet{Sterz95} (see Figure~\\ref{clouds}). This may be an indication that the X-ray emitting young stellar population in these regions is less conspicuous, or that the objects are much younger, and/or more embedded, than in other clouds revealed as X-ray clumps like L\\,1616.  The chapter is structured as follows: an overview of the L\\,1634, L\\,1616/L\\,1615, IC~2118 and L\\,1642 clouds and their young populations are presented in Sections~2, 3, 4 and 5 respectively, while other small clouds are discussed in Section~6. The four small clouds around $\\sigma$~Ori are then discussed in Section~7, while a synthesis of the distances of the outlying clouds is presented in Section~8.  Finally, a summary is presented in Section~9.\\\\ ",
        "conclusions": "\\label{sum}  As a summary we present in Table~\\ref{tab:summary} the numbers of YSOs and outflows in each region discusses above and a bibliographic guide to the observational studies performed in different wavelength ranges.   The apparent small number of YSOs found so far in the region of  L\\,1634 compared to L\\,1616 may result from the fact that the studies conducted in this region have been mainly devoted to the investigation of the outflow, rather than to the search of PMS stars. Many more low-mass PMS stars in L~1634 are hence expected to be identified in the coming years. It is not surprising that the clouds with most active star formation are closer in angular distance to the main Orion molecular clouds, in particular L\\,1634, L\\,1615 and L\\,1616. Star formation in the IC~2118 region is also ongoing, although not so vigorously as in L\\,1616. IC~2118 may be connected to the star-forming complex in Orion. However, a true link has not been clarified yet. Given the distance of 210~pc for the reflection nebula and its associated molecular clouds, the stellar association in IC~2118 may be located within the Orion--Eridanus Bubble \\citep{Kun04}. The other clouds at higher galactic latitude, like L\\,1642, are probably related to the Orion-Eridanus bubble rather than to the Orion star forming region. In fact, those clouds are closer to the Sun, consistently with the idea that they may be associated with the nearest side of the bubble outskirt. Further investigations of the PMS populations of these clouds in the coming years will help to shed light on their relationship with the Orion-Eridanus bubble.  \\clearpage   \\begin{table}[!ht] \\setlength{\\tabcolsep}{0.5\\tabcolsep}  \\caption{\\label{tab:summary} Summary of observations in the Orion outlying clouds and other globules} \\label{coordinates} \\begin{flushleft} {\\footnotesize \\begin{tabular}{l@{\\hskip3pt}r@{\\hskip5pt}c@{\\hskip3pt}c@{\\hskip8pt}r@{\\hskip8pt}c@{\\hskip8pt}c@{\\hskip8pt}l} \\tableline \\noalign{\\smallskip} &   \\multicolumn{2}{c}{Number of} & ~ &\t \\multicolumn{4}{c}{References~~} \\\\ Cloud   & YSOs  & outflows & ~  &   X-rays    &\toptical     &\tIR  &  radio   \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} L~1634, CB~29              & 9      &  2     & ~   & 15, 16     & 1, 5, 10     &  2, 10, 13,     &  3, 22, 23,   \\\\ &        &        & ~   &\t        & 12, 16, 43   & 17, 24, 29, 44  &  4, 7, 17     \\\\ \\noalign{\\smallskip} L\\,1616 and 1615, CB~28      & 57     & 1      & ~   & 15, 16     & 16, 37, 43  & 16, 25, 30, 37  &  25, 7, 14, 32 \\\\  \\noalign{\\smallskip} IC~2118                    & 10     &        & ~  & 15, 16     & 16, 26, 43    & 21, 30          &  31, 4, 9, 21 \\\\  \\noalign{\\smallskip} LBN~991                    & 1      &        & ~  & 15, 16     & 16\t       &   --            &   --\t   \\\\  \\noalign{\\smallskip} LBN~917, LBN~906           & 1      &        & ~  &  15\t       & 11\t       &   30            &   9           \\\\  \\noalign{\\smallskip} L~1642                     & 6      & 1      & ~  &  15\t       & 6, 18, 19     &  20, 28, 30     &  3, 4, 27     \\\\  \\noalign{\\smallskip} anonymous X-ray clump      & 4      &        & ~  &  15, 16    & 16\t       &     --          &  --           \\\\  \\noalign{\\smallskip} OS98-27                    & 2      &        & ~  &            &  33, 34       &     --          &    --         \\\\  \\noalign{\\smallskip} OS98-35                    & 1      &        & ~  &            &  33           &    --           &    --        \\\\  \\noalign{\\smallskip} OS98-40 = Ori~I-2          & 1      &  1     & ~  &   38       &  35, 40        &    35, 42       &    36, 39, 40, 41 \\\\   \\noalign{\\smallskip} \\tableline \\end{tabular} } \\begin{flushleft} {\\footnotesize 1:~\\citet{Cohen80}; 2:~\\citet{Coh85}; 3:~\\citet{MagBliMun85}; 4:~\\citet{Mad86}; 5:~\\citet{Step86}; 6:~\\citet{Sand87}; 7:~\\citet{Clem88}; 8:~\\citet{ReipHeat90}; 9:~\\citet{Bally91}; 10:~\\citet{Sugi91}; 11:~\\citet{GrHe92}; 12:~\\citet{Bohi93}; 13:~\\citet{HoLa95}; 14:~\\citet{Ram95}; 15:~\\citet{Sterz95}; 16:~\\citet{Alc96, Alc00, Alc04}; 17:~\\citet{Davis97}; 18:~\\citet{Hearty00}; 19:~\\citet{LiCh00}; 20:~\\citet{Vert00}; 21:~\\citet{Kun01}; 22:~\\citet{Beltr02}; 23:~\\citet{DeVr02}; 24:~\\citet{Nisi02}; 25:~\\citet{Stank02}; 26:~\\citet{Kun04}; 27:~\\citet{Russeil03}; 28:~\\citet{Leht04}; 29:~\\citet{OConn04}; 30:~\\citet{reach98}; 31:~\\citet{Yon99}; 32:~\\citet{Park04}; 33:~\\citet{Yun97}; 34:~\\citet{Bri05}; 35:~\\citet{Mad99}; 36:~\\citet{Sug89}; 37:~\\citet{Gan08}; 38:~\\citet{Car98}; 39:~\\citet{Lar99}; 40:~\\citet{Cer92}; 41:~\\citet{Bot96}; 42:~\\citet{Hod94}; 43:~\\citet{Lee07}; 44:~\\citet{Sea08} } \\end{flushleft}  \\end{flushleft} \\end{table}  \\vspace{0.5cm}     {\\bf Acknowledgements.} We thank the referee, K. Sugitani, for his careful reading and helpful comments and suggestions. We are also grateful to Bo Reipurth for his manifold inputs and comments on an earlier version of the manuscript. We thank T. Stanke and C. Davis for providing figures from their work. We are grateful to Steve Rodney for his help with LaTex and the scaling of the format of some of the figures. We also thank A. Frasca, L. Spezzi, D. Gandolfi, E. Marilli, L. Testi and A. Natta for fruitful discussions, as well as for their collaboration in some of the works mentioned in this chapter. We acknowledge the use of the color image of IC~2118 by Noel Carboni. This work was partially financed by the Italian Ministery of University and Research (MIUR). Financial support from INAF (PRIN-INAF-2005 project ''Stellar clusters: a benchmark for star formation and stellar evolution'') and from {\\em Regione Campania} is also acknowledged. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.   \\clearpage"
    },
    "0809/0809.0316_arXiv.txt": {
        "abstract": "I model profiles of the [Ne\\,{\\sc ii}] forbidden emission line at 12.81$\\mu$m, emitted by photoevaporative winds from discs around young, solar-mass stars.  The predicted line luminosities ($\\sim10^{-6}$L$_{\\odot}$) are consistent with recent data, and the line profiles vary significantly with disc inclination.  Edge-on discs show broad (30--40\\,km\\,s$^{-1}$) double-peaked profiles, due to the rotation of the disc, while in face-on discs the structure of the wind results in a narrower line ($\\simeq10$\\,km\\,s$^{-1}$) and a significant blue-shift (5--10\\,km\\,s$^{-1}$).  These results suggest that observations of  [Ne\\,{\\sc ii}] line profiles can provide a direct test of models of protoplanetary disc photoevaporation. ",
        "introduction": "\\label{sec:intro} The manner in which gas is removed from discs around young stars has critical consequences for theories of planet formation, as the formation of gas-giant planets must precede the dissipation of gas discs.  Currently, models of protoplanetary disc evolution suggest that the most of the gas is removed by a combination of viscous accretion through the disc and photoevaporation by irradiation from the central star [e.g.,~\\citet*{cc01,acp06b}; see also the recent reviews by \\citet{dull_ppiv} and \\citet{rda08}].  Such models agree reasonably well with current data, but the ionizing luminosities of T Tauri stars (required to drive the photoevaporation) are very poorly constrained.  \\citet*{acp05} used archival data to estimate the chromospheric ionizing fluxes from a small sample of bright T Tauri stars, and found values in the range $\\Phi \\sim 10^{42}$--$10^{43}$ ionizing photons per second.  However, recent data suggests that this may over-estimate more typical values of $\\Phi$ by 1--2 orders of magnitude \\citep{herczeg07b}.  Photoevaporative flows have been observed clearly in ultracompact H\\,{\\sc ii} regions around massive stars \\citep{holl94,lugo04}, and in the externally illuminated ``proplyds'' in the Orion nebula \\citep*[e.g.][]{jhb98}.  However, the evidence for or against the existence of central star-driven photoevaporative winds from discs around solar-mass stars is rather tenuous.  \\citet{font04} modelled the profiles of optical forbidden emission lines and found reasonable agreement with both the line fluxes and profiles obtained in the spectroscopic survey of \\citet*{heg95}.  However, some discrepancies were also present (notably between the predicted and observed fluxes from [O\\,{\\sc i}]), and no further studies of such emission lines exist.  For several reasons, the [Ne\\,{\\sc ii}] fine-structure line at 12.81$\\mu$m offers perhaps the best opportunity to confirm or deny the existence of a slow ($\\sim10$\\,km\\,s$^{-1}$) ionized wind from a T Tauri disc.  The high ionization potential of neon (21.56eV) means that Ne$^+$ only exists close to T Tauri stars in photoionized gas, so forbidden emission lines from neon ions are unlikely to arise elsewhere in the T Tauri system.  The 12.81$\\mu$m line also falls in an atmospheric window (unlike the [Ne\\,{\\sc iii}] line at 15.55$\\mu$m), and can therefore be observed from the ground at echelle resolution.  Moreover, the continuum emission from the star-disc system is orders of magnitude weaker in the mid-infrared than in the optical, so the line-to-continuum ratio for [Ne\\,{\\sc ii}] can often be higher than for optical forbidden lines such as [N\\,{\\sc ii}] or [S\\,{\\sc ii}].  In addition, as will be shown, the critical density of the [Ne\\,{\\sc ii}] line is such that most of the emission arises in the ``launching region'' of the photoevaporative wind, making it an excellent tracer of disc photoevaporation.  Recently the {\\it Spitzer Space Telescope} has detected the [Ne\\,{\\sc ii}] 12.81$\\mu$m line towards more than 20 young, solar-mass stars \\citep{pasc07,lahuis07,esp07}, and in most cases it is thought to arise in the star-disc system.  \\citet{herczeg07a} observed the [Ne\\,{\\sc ii}] line from the nearby source TW Hya at high resolution ($\\lambda/\\Delta\\lambda \\approx 30,000$), and found that the line (with width $21\\pm4$\\,km\\,s$^{-1}$) was narrower than lines that originate in the accretion funnel, but significantly broader than lines emitted from the disc at larger radii: this strongly suggests that the line originates in a hot disc atmosphere.  \\citet*{gni07} and \\citet{gh08} showed that X-ray heating of the disc can reproduce the observed line fluxes, but a photoevaporative disc wind should give rise to a similar line flux \\citep{pasc07,gh08}.  While both ionization mechanisms (X-ray and UV) can likely explain the observed line fluxes, the velocity structure of a photoevaporative wind should result in a line profile that is distinct from that produced by X-ray irradiation of a bound disc atmosphere.  In this {\\it Letter} I model the profiles of the [Ne\\,{\\sc ii}] 12.81$\\mu$m line emitted by photoevaporative winds, and find that high-resolution spectroscopy can provide an unambiguous means of detecting such winds in T Tauri systems. ",
        "conclusions": "In general, these results compare favourably with the previous modelling of \\citet{font04}.  These authors modelled the emission from optical forbidden lines ([S\\,{\\sc ii}] and [N\\,{\\sc ii}]), and for face-on inclinations found that the line profiles were blue-shifted by $\\simeq10$\\,km\\,s$^{-1}$.  The blue-shift decreased with increasing inclination angle, and in the edge-on discs the lines were centred on zero velocity.  They did not, however, find any double-peaked lines, even in the edge-on case.  The reasons for this are not entirely clear, but it is likely attributable to the different critical densities of the lines considered.  The critical densities of the [S\\,{\\sc ii}] and [N\\,{\\sc ii}] lines are lower that that of the [Ne\\,{\\sc ii}] 12.81$\\mu$m line, so these lines are dominated by emission from larger radii, where the rotation speed of the disc is smaller.  Consequently, as in the case of a disc with a large inner hole, the lines do not appear double-peaked.  So far I have neglected the effects of X-ray ionization in producing [Ne\\,{\\sc ii}] emission, but when discussing observations this cannot be ignored.  \\citet{gni07} modelled the effects of X-ray irradiation in detail \\citep[see also][]{gh08}.  Unlike in the UV case, the X-ray luminosities of T Tauri stars are well-known, and the predicted [Ne\\,{\\sc ii}] emission from X-ray ionized gas is comparable in strength to the emission from the photoevaporative wind ($\\sim 10^{-6}$L$_{\\odot}$).  X-ray heating likely creates a hot, but static, disc atmosphere, and in real systems the emission from this atmosphere will be added to the line profiles modelled here.  The line profiles from such a bound atmosphere have not been modelled in detail, but their general shape was discussed by \\citet{gni07}.  They find that emission is dominated by gas at small radii $(\\lesssim15$AU), with most of the emission arising at radii where the rotational speed of the disc gas is $\\sim10$\\,km\\,s$^{-1}$.  The gas temperature in this region is 1000--5000K, resulting in an intrinsic (Doppler) line-width of $\\sim 1$\\,km\\,s$^{-1}$.  The line should therefore be double-peaked when viewed edge-on, with a profile similar to that predicted here.  In the face-on case, however, the profile should be different: narrow ($\\sim 1$\\,km\\,s$^{-1}$), and centred on zero velocity.  The broader ($\\sim 10$\\,km\\,s$^{-1}$) line-width predicted here is unlikely to be matched by a static disc atmosphere unless any turbulence in the disc is highly supersonic \\citep[which is generally not the case in protoplanetary discs, e.g.,][]{bh98}, or unless the emission originates very close to the star ($\\lesssim 1$AU).  In addition, a static atmosphere results in line emission centred on zero velocity, distinct from the blue-shifted profile of the photoevaporative wind.  Critical to observing this blue-shift are the relative line fluxes from the static atmosphere and the wind.  If these are comparable, as predicted for $L_X\\sim10^{30}$erg s$^{-1}$ and $\\Phi \\sim 10^{41}$s$^{-1}$, then the diagnostic blue-shift will likely be 2--5\\,km\\,s$^{-1}$.   Recently \\citet{herczeg07a} observed the [Ne\\,{\\sc ii}] 12.81$\\mu$m line from the nearby face-on ($i\\simeq7^{\\circ}$) disc TW Hya, at a resolution of $\\lambda/\\Delta\\lambda \\approx 30,000$.  The emission was unresolved at an effective spatial resolution of $\\sim40$AU, suggesting that it arises within this distance of the central star.  \\citet{herczeg07a} measured a line-width (FWHM) of $21\\pm4$\\,km\\,s$^{-1}$, centred on $-2\\pm3$\\,km\\,s$^{-1}$.  The blue-shift is not statistically significant, in part because of the limitations of the wavelength calibration, but the line width is rather larger than expected from a static disc atmosphere alone.  In addition there was some evidence for asymmetry in the profile, with an enhanced flux on the blue side of the line, but not at a statistically significant level.  The observed profile is consistent with emission from a photoevaporative wind (centred on $\\simeq -5$--10\\,km\\,s$^{-1}$) combined with the emission from a heated disc atmosphere (centred on zero velocity).  Further such observations are expected in the near future, and in this {\\it Letter} I have shown that high-resolution observations of the [Ne\\,{\\sc ii}] 12.81$\\mu$m line can provide a critical test of models of disc photoevaporation.  Detection of blue-shifted [Ne\\,{\\sc ii}] emission would provide unambiguous evidence of a photoevaporative wind, and observations of these line profiles may represent the most readily observable diagnostic of central star-driven disc photoevaporation."
    },
    "0809/0809.2914_arXiv.txt": {
        "abstract": "{One of the main obstacles in extracting the Cosmic Microwave Background (CMB) signal from observations in the mm-submm range is the foreground contamination by emission from galactic components: mainly synchrotron, free-free and thermal dust emission. Due to the statistical nature of the intrinsic CMB signal it is essential to minimize the systematic errors in the CMB temperature determinations.   Following the available knowledge of the spectral behavior of the galactic foregrounds simple, power law-like spectra have been assumed. The feasibility of using a simple neural network for extracting the CMB temperature signal from the combined CMB and  foreground signals has been investigated. As a specific example, we have analysed simulated data, like that expected from the ESA Planck Surveyor mission. A simple multilayer perceptron neural network with 2 hidden layers can provide temperature estimates, over more than 80 percent of the sky, that are to a high degree uncorrelated with the foreground signals. A single network will be able to cover the dynamic range of the Planck noise level over the entire sky.} ",
        "introduction": "The Cosmic Microwave Background (CMB) is by far the most important data set available for cosmological investigations. The angular power spectrum of its temperature and polarization anisotropies contain unique information about the basic cosmological parameters. Since the discovery of the CMB by Penzias \\& Wilson (1965), tremendous efforts on the instrumental side have been made to improve sensitivity and angular resolution. The latest major experiment is the successful, ongoing, \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP Bennett et al.\\,2003a) and the next  step forward in CMB measurements will no doubt be the ESA Planck mission.  Due to the statistical nature of the CMB signal it is essential to subtract carefully all signals from non-cosmological sources, before the data can be used in a cosmological context. A clear demonstration of this problem has been given by Tegmark et al.\\ (2000). Their model included frequency dependence, angular scale dependence for each physical component, and variations in the frequency dependence across the sky.  From these simulations they calculated angular power spectra and extracted the basic cosmological parameters. For the foregrounds they assumed power laws, and made 3 different sets of assumptions about the power law parameters: Optimistic (O), Middle-of-The-Road (MID), Pessimistic (PESS). With observational errors like those for Planck, they found that the accuracy of most cosmological parameters is degraded by a factor of about 2 for the MID model, and by a factor of about 5 for the PESS model. Therefore, in order to exploit fully the scientific capability of an experiment like Planck a procedure for reliable removal of all non-cosmological signals is mandatory.  On angular scales larger than about 30' the non-CMB signal is dominated by diffuse emission from our own galaxy (De Zotti et al.\\ 1999). Synchrotron (Haslam et al.\\ 1982) and free-free (Haffner, Reynolds \\& Tufte 1999; Finkbeiner 2003)  emissions dominate below about 60--80 GHz, while at higher frequencies thermal dust emission takes over. On smaller angular scales the foreground fluctuations are dominated by several populations of extragalactic sources, each with different spectral behaviour: radio sources, dusty galaxies and the Sunyaev-Zeldovich effect from clusters of galaxies.  The CMB radiation has a virtually perfect black body spectrum (Mather 1999), while all known non-cosmological signals have a clear non-thermal frequency dependence. Therefore, it should be possible to split the observed microwave signal into its different components. This depends critically, of course, on the observational accuracy and the number of frequencies observed.  A lot of different methods for component separation of CMB signals have been investigated, including: \\begin{itemize} \\item Maximum Entropy Method (MEM): Hobson et al.\\ \\cite{hobs98}, Stolyarov et al.\\ \\cite{stol02}, Barreiro et al.\\ \\cite{barr04}, Stolyarov et al.\\ \\cite{stol05} \\item Internal Linear Combination method (ILC): e.g.\\ Bennett et al.\\ \\cite{benn03b} \\item Wiener Filtering: e.g.\\ Tegmark \\& Efstathiou \\cite{tegm96} \\item Independent Component Analysis (ICA) method: e.g.\\ Maino et al.\\ \\cite{main02} \\end{itemize} It is beyond the scope of this paper to give a review of all relevant component separation methods, but an excellent review can be found in  Delabrouille \\& Cardoso \\cite{dela07}.  The statistical properties of the CMB temperature fluctuations on a sphere are normally expressed as a sum of spherical harmonics: \\begin{equation} T(\\theta,\\varphi) = \\sum_{l=0}^{\\infty}\\sum_{m=-l}^{l} a_{lm} Y_{lm}(\\theta,\\varphi), \\end{equation} where $Y_{lm}(\\theta,\\phi)$ are the spherical harmonic functions. The simplest versions of the inflation paradigm predict that the CMB temperature fluctuations will be isotropically, randomly, Gaussianly distributed on a sphere. With this assumption, the statistical properties are then completely specified by the second order statistics, the angular power spectrum $C_{l}$, \\begin{equation} C_{l}~=~\\frac{1}{2l + 1}~\\sum_{m=-l}|a_{lm}|^{2}. \\end{equation} Therefore, a lot of effort has been put into constructing algorithms for foreground removal from CMB data, which assure minimum systematic residuals in the derived angular power spectrum. But, of course, an optimal method will also assure minimal systematic residuals in the CMB map itself.  As demonstrated by Chiang \\& Naselsky (2007), important information can be derived from the phases of the spherical harmonic coefficients. They find clear evidence for systematics in the distribution of phases at high latitudes in the WMAP 3y ILC map, and interpret them as due to residual foreground signals. For the Planck mission, one of the key scientific issues will be the search for primordial non-gaussianity in the CMB map. If clear evidence is not found, then tight upper limits on non-gaussianity will be important, emphasizing the need not only power for spectra with no systematic errors, but also maps.  The main purpose of this paper is to investigate the feasibility of removing galactic foreground emission from CMB data by means of a neural network. As an example we will specifically examine simulated temperature data as expected from the ESA Planck CMB mission. See ESA-SCI(2005)1 for a detailed technical description of the payload and a discussion of the expected scientific capability.  Planck will provide maps of the microwave sky with unprecedented sensitivity and angular resolution. The 2 detector systems have been designed to cover a very broad frequency range in order to facilitate the removal of foregrounds. Therefore, it is important to analyse how well the different foreground signals can be removed by using only Planck data, and so-called `blind' extraction methods.  For Planck, there is still uncertainty about how this component separation shall be performed in practice. Ideally, for each pixel on the sky the component separation procedure should provide an estimate of the temperature of the CMB, with an accidental error derived from the observational accuracy and with  negligible systematic errors (i.e.\\ an error uncorrelated with any of the parameters of the different foreground signals).  Several methods have been developed to estimate the CMB temperature without any assumptions about priors. Totally 'blind' methods like the 'Independent Component Analysis' method used by  FastICA (Maino et al.\\ 2002) and SMICA (Delabrouille et al.\\ 2003; Moudden et al.\\ 2005) can, in principle, estimate both the CMB map and  the spatial and frequency dependence of the foreground sources, only relying on the CMB data itself. Similarly, in the 'Internal Linear Combination' (ILC) method used by the WMAP team (Bennett et al.\\  2003b; Hinshaw et al.\\ 2007), no assumptions are made \\emph{a priori} about the foregrounds. Another approach was investigated by Brandt et al.\\ (1994). They parameterized the spectra of the different foregrounds and then fitted them, sky pixel by sky pixel, with a non-linear least squares Levenberg-Marquardt algorithm (Levenberg 1944; Marquardt 1963). This approach has been further developed by Linden-V{\\o}rnle \\& N{\\o}rgaard-Nielsen (1998) and Eriksen et al.\\ (2006).  In this paper we will elaborate on the ILC method by incorporating the non-linear features of neural networks. To set up these neural networks we parameterize the spectra of the different foregrounds in line with the Brandt et al.\\ (1994) approach.  This paper is organized in the following way: In Section 2, we outline the available information about spectral behaviour of the galactic foregrounds. We present the details of the modeling used in Sections 3 and 4. In Section 5, we briefly discuss the results found previously with the ILC method. Since neural networks are not commonly used in astrophysical literature, a brief introduction to the neural network concept is given in Section 6. We present the neural network used in Section 7 . In Section 8, we discuss our results in relation to similar previous investigations. ",
        "conclusions": "In this paper we investigate mainly the concept of using simple neural networks (MLPs) to remove the galactic foregrounds from the CMB signal, making some simplifying assumptions: only 3 galactic components are taken into account (synchrotron, free-free and thermal dust emission); the spectra of the foregrounds follow simple modifications to power laws; the observational noise is white (Gaussian); Gaussian beams are assumed. Under these assumption it has been demonstrated that for more than 80 per cent of the sky a neural network can provide reliable estimates of the CMB temperature with errors, which to a high degree are uncorrelated to the basic parameters of the sky model used.  The spectra of the foregrounds in the Planck frequency range, close to the galactic plane, are at present known to be complex, but little is known about the complexity at high latitudes. Therefore, in the case of Planck the details of the modeling of the foregrounds will first be established when the data has been carefully reduced and analysed. This will, no doubt, result in the need for including more galactic components (e.g.\\ spinning dust) and, most likely, a revision of the spectral behavior of the components will be required. From the available knowledge of the Planck detectors, it is clear that the noise will not be white, and it will be a challenging task to investigate how to minimize the influence of non-white noise features. It could well be that more sophisticated neural networks than MLP's will be needed in the Planck pipeline.  To further test the capabilities of neural networks to clean CMB temperature maps, we will, in a forthcoming paper, analyse the Planck Sky Model simulations of both Planck and WMAP data, as well as some real data, the recently released WMAP 5yr data."
    },
    "0809/0809.1126_arXiv.txt": {
        "abstract": "Growing evidence indicate supermassive black holes (SMBHs) in a mass range of $M_{\\rm BH}$$\\sim 10^6-10^{10}M_{\\odot}$ lurking in central stellar bulges of galaxies. Extensive observations reveal fairly tight power laws of $M_{\\rm BH}$ versus the mean stellar velocity dispersion $\\sigma$ of the host stellar bulge. Together with evidence for correlations between $M_{\\rm BH}$ and other properties of host bulges, the dynamic evolution of a bulge and the formation of a central SMBH should be linked. In this Letter, we reproduce the empirical $M_{\\rm BH}-\\sigma$ power laws based on our recent theoretical analyses (Lou \\& Wang; Wang \\& Lou; Lou, Jiang \\& Jin) for a self-similar general polytropic quasi-static dynamic evolution of bulges with self-gravity and spherical symmetry and present a sensible criterion of forming a central SMBH. The key result is $M_{\\rm BH}={\\cal L}\\ \\sigma^{1/(1-n)}$ where $2/3<n<1$ and ${\\cal L}$ is a proportional coefficient characteristic of different classes of host bulges. By fitting and comparing several empirical $M_{\\rm BH}-\\sigma$ power laws, we conclude that SMBHs and galactic bulges grow and evolve in a coeval manner and most likely there exist several classes of galactic bulge systems in quasi-static self-similar evolution and that to mix them together can lead to an unrealistic fitting. Based on our bulge-SMBH model, we provide explanations for intrinsic scatter in the relation and a unified scenario for the formation and evolution of SMBHs in different classes of host bulges. ",
        "introduction": "It is now widely accepted that supermassive black holes (SMBHs) within a mass range of $10^6\\sim10^{10} M_{\\odot}$ ($M_{\\odot}=2\\times 10^{33}$g is the solar mass) form at the centres of spiral and elliptical galaxies (e.g., Lynden-Bell 1969; Kormendy \\& Richstone 1995; Kormendy 2004). Observationally, SMBH masses $M_{\\rm BH}$ correlate with various properties of spiral galaxy bulges or elliptical galaxies, including bulge luminosities (e.g., Kormendy \\& Richstone 1995; Magorrian et al. 1998; Marconi \\& Hunt 2003), stellar bulge masses $M_{\\rm bulge}$ (e.g., Magorrian et al. 1998; Marconi \\& Hunt 2003; H$\\ddot{\\rm a}$ring \\& Rix 2004), $M_{\\rm BH}$ versus $M_{\\rm bulge}$ relations in active (AGN) and inactive galaxies (e.g., Wandel 1999, 2002; McLure \\& Dunlop 2002), galaxy light concentrations (e.g., Graham et al. 2001), the S$\\acute{\\rm e}$rsic index (S$\\acute{\\rm e}$rsic 1968) of surface brightness profile (e.g., Graham \\& Driver 2007), inner core radii (e.g., Lauer et al. 2007), spiral arm pitch angles (e.g., Seigar et al. 2008), bulge gravitational binding energies (e.g., Aller \\& Richstone 2007) and mean stellar velocity dispersions $\\sigma$ (e.g., Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Tremaine et al. 2002; Ferrarese \\& Ford 2005; Hu 2008). These empirical correlations strongly suggest a physical link between SMBHs and their host bulges (e.g., Springel, Matteo \\& Hernquist 2005; Li, Haiman \\& MacLow 2007).  Among these empirical relations, $M_{\\rm BH}$ and $\\sigma$ correlate tightly in a power law with an intrinsic scatter of $\\lsim 0.3$ dex (e.g., Novak et al. 2006). This relation was explored theoretically (e.g., Silk \\& Rees 1998; Fabian 1999; Blandford 1999) before observations (e.g., Ferrarese \\& Merritt 2000; Gebhardt et al. 2000) and the emphasis was on outflow effects for galaxies. The idea was further elaborated by King (2003). A model of singular isothermal sphere with rotation (Adams, Graff \\& Richstone 2001) was proposed for the $M_{\\rm BH}-\\sigma$ relation. This relation was also studied in a semi-analytic model (Kauffmann \\& Haehnelt 2000) with starbursts while SMBHs being formed and fueled during major mergers. Accretion of collisional dark matter onto SMBHs may also give the $M_{\\rm BH}-\\sigma$ relation (e.g., Ostriker 2000; see Haehnelt 2004 for a review). There are also numerical simulations to model feedbacks from SMBHs and stars on host galaxies.  Observationally, there are two empirical types of bulges: classical bulges (spiral galaxies with classical bulges or elliptical galaxies) and pseudobulges (e.g., Kormendy et al. 2004; Drory \\& Fisher 2007). While SMBHs in classical bulges are formed after major mergers, pseudobulges do not show obvious merger signatures. Interestingly, pseudobulges also manifest a $M_{\\rm BH}-\\sigma$ power law yet with a different exponent (e.g., Kormendy \\& Gebhardt 2001; Hu 2008).  Self-similar dynamics of a conventional polytropic gas sphere has been studied earlier (e.g., Yahil 1983; Suto \\& Silk 1988; Lou \\& Wang 2006 (LW06) and references therein). LW06 obtained novel self-similar quasi-static dynamic solutions which approach singular polytropic spheres (SPS) after a long time. Such polytropic dynamic solutions have been further generalized and applied to study protostar formation, ``champagne flows\" in H II regions, stellar core collapse, rebound shocks and the formation of compact stellar objects in a single fluid model (LW06; Lou \\& Wang 2007; Wang \\& Lou 2007, 2008; Hu \\& Lou 2008) as well as galaxy clusters in a two-fluid model (Lou, Jiang \\& Jin 2008; LJJ hereafter). Here, we take such quasi-static solutions in a general polytropic fluid to describe long-time evolution of host stellar bulges and the formation of central SMBHs and to establish $M_{\\rm BH}-\\sigma$ power laws. ",
        "conclusions": "On the basis of a self-similar quasi-static dynamic evolution of a general polytropic sphere (LW06, LJJ), we establish $M_{\\rm BH}=\\mathcal{L}\\ \\sigma^{\\ 1/(1-n)}$ with $2/3<n<1$ by equation (\\ref{relation}). Our first conclusion is $1/(1-n)>3$ which appears consistent with observations so far. Secondly, uncertainties in the formation criterion of a SMBH [i.e., factor $f$ in $M_{\\rm BH}=f r_sc^2/(2G)$] and in the choice of $r_c$ and thus $\\rho_c$ will not change the form of equation (\\ref{relation}) and the $n$ value but will only affect the value of ${\\cal L}$. Thirdly, the tight $M_{\\rm BH}-\\sigma$ power laws and other relations among the SMBH mass $M_{\\rm BH}$ and known properties of host stellar bulges strongly suggest coeval growths of SMBHs and galactic bulges (e.g., Page et al. 2001; Haehnelt 2004; Kauffmann \\& Haehnelt 2000; Hu et al. 2006). Fourthly in our model, while forming a SMBH at the bulge centre (e.g., by core collapse of gas and stars or by merging), the spherical general polytropic bulge evolves in a self-similar quasi-static phase for a long time. We then reproduce well-established empirical $M_{\\rm BH}-\\sigma$ power laws. Different energetic processes appear to give rise to different scaling index $n$ values, which finally determines the slope of the $M_{\\rm BH}-\\sigma$ relations in a logarithmic presentation.  Besides classical bulges and pseudobulges, there are also `core' elliptical galaxies (i.e., those with apparent `cores' of relatively flat brightness; Lauer et al. 1995; Hu 2008), thought to have formed by `dry' mergers (i.e., almost without gas). A steeper $M_{\\rm BH}-\\sigma$ relation exists in these galaxies as compared to that for classical bulges (e.g., Lauer et al. 1995; Laine et al. 2003; Lauer et al. 2007).  This can also be accommodated in our unified scenario that all hosts of SMBHs may finally evolve into self-similar quasi-static phase with different scaling parameters (i.e., different index $n$ for the slope and different $\\rho_c$ for the normalization of the $M_{\\rm BH}-\\sigma$ relation). While these bulges can be of quite different kinds in galaxies, they have these similar tight relations and we thus provide a unified self-similar dynamic framework to model the relatively quiescent evolution phase of SMBH host bulges and the growth of SMBH masses. As the observed $M_{\\rm BH}-\\sigma$ relation for classical bulges is tight, the elliptical galaxies and spiral galaxies appear to take on close $n$ values for merging processes.  In our model, $n$ is a key scaling index to determine the exponent of the $M_{\\rm BH}-\\sigma$ power law. The smaller the value of $n$ is, the steeper the density profile is and the smaller the index of the $M_{\\rm BH}-\\sigma$ relation is. If SMBHs are formed by collapse of stars and gas and a less steeper density distribution may provide a more effective mechanism to form SMBHs, then we conclude that for a certain value of velocity dispersions, \\footnote{The mean velocity dispersion $\\sigma={\\cal Q}K^{1/2}$ is not sensitive to $n$.} the smaller the mass of an initially formed SMBH is, the smaller the value of $n$ is.  After a long lapse, our quasi-static solutions approach the static SPS solution as the leading term that is independent of the timescale. So all the described relations here are nearly independent of time, as long as the systems have evolved for a long enough time. It is not obvious to decide when the host bulges began to take the described self-similar evolution. But even if we take different times in our model, our results would remain largely the same. The only difference is that for an early time, we may have a chance to observe accretion shocks within the region of stellar bulges. Such shocks are characterized by a rapid inner density rise of several times and a rapid inner rise of stellar velocity dispersions depending on the strength of accretion shocks."
    },
    "0809/0809.3912_arXiv.txt": {
        "abstract": "{The analysis of the spectral energy distribution variability at frequencies from radio to TeV is a powerful tool in the investigation of the dynamics, the physics and the structure evolution occurring in the most exotic flavour of active galaxies, the blazars. In particular, the presence of {\\sl Fermi-GST} is providing a unique opportunity for such studies delivering $\\gamma$-ray data of unprecedented quality. Here we introduce a monitoring program that has been running at the Effelsberg 100~m telescope since January 2007, underpinning a broad multi-frequency collaboration of facilities that cover the band from radio to infrared. Sixty one selected blazars are observed monthly between 2.64 GHz and 43 GHz. The calibration accuracy is better than a few percent as it is demonstrated with some preliminary examples. ",
        "introduction": "\\begin{table}[h] \\caption{The monitored sources.} \\label{angelakis_gamma_tab1} \\begin{center} \\begin{tabular}{lll} \\hline \\\\ \\multicolumn{3}{c}{Target sources}\\\\ \\hline PKS\\,0003$-$066  &TXS\\,0814$+$425    &OE\\,355  \\\\ PKS\\,0215$+$015  &1E\\,0317.0$+$18    &OJ\\,248  \\\\ PKS\\,0235$+$164  &H\\,1426$+$428      &OJ\\,287  \\\\ PKS\\,0238$-$084  &1ES\\,1959$+$650    &OP\\,313  \\\\ PKS\\,0336$-$019  &1ES\\,2344$+$514    &3C\\,84   \\\\ PKS\\,0420$-$014  &1ES\\,1544$+$820    &3C\\,111  \\\\ PKS\\,0528$+$134  &1ES\\,0502$+$675    &3C\\,120  \\\\ PKS\\,0735$+$178  &0059$+$581         &3C\\,273  \\\\ PKS\\,0748$+$126  &0219$+$428         &3C\\,279  \\\\ PKS\\,1038$+$064  &1128$+$592         &3C\\,345  \\\\ PKS\\,1127$-$145  &Mkn\\,180           &3C\\,371  \\\\ PKS\\,1406$-$076  &Mkn\\,421           &3C\\,446  \\\\ PKS\\,1510$-$089  &Mkn\\,501           &3C\\,454.3\\\\ PKS\\,1730$-$130  &4C\\,21.35   &CTA\\,102\\\\ PKS\\,2155$-$152  &4C\\,28.07   &Cyg\\,A  \\\\ PKS\\,2155$-$304  &4C\\,31.63   &BL\\,LAC \\\\ PKS\\,2345$-$16   &4C\\,38.41   &WCom    \\\\ S5\\,0836$+$71    &4C\\,56.27   &M87     \\\\ S5\\,1803$+$78    &4C\\,47.08   &OS\\,319 \\\\ S5\\,0716$+$71    &TON\\,0599   &\\\\ S4\\,0954$+$65    &NRAO\\,150   &\\\\ \\hline \\end{tabular} \\end{center} \\end{table} Being among the most dramatic manifestations of the activity induced in the nuclei of active galaxies, blazars comprise a unique probe of the exotic physics at play in such systems. Their phenomenology is dominated by extreme characteristics such as high degree of linear polarisation, intense variability -- both in total power and polarisation -- at all wavebands, apparent motions of features in their radio structure, which are often highly superluminal and brightness temperatures exceeding the Compton limit \\cite[see e.g.][]{angelakis_gamma_Urry1999APh}. This violent behaviour is attributed to relativistic jets oriented very close ($\\le 20^{\\circ}$to $30^{\\circ}$) to the line-of-sight \\cite[see e.g.][]{angelakis_gamma_Urry1995PASP}. Their spectral energy distribution (SED) is characterised by the presence of two peaks. The lower energy one, spreading from radio to far ultraviolet and soft X-rays, is believed to be due to relativistic electrons emitting via the synchrotron mechanism. The high energy part, reaching up to TeV energies is believed to be produced by synchrotron self-Compton or external Compton scattering.  Despite the overall understanding, many details remain unclear and are subject to models attempting their clarification. For instance, several ideas have been put forth to explain the origin of their variability. Two of them are like the shock-in-jet model \\citep[e.g.][]{angelakis_gamma_Marscher1985ApJ,angelakis_gamma_Aller1985ApJ,angelakis_gamma_Marscher1996ASPC} and the relativistic plasma shells \\citep[e.g.][]{angelakis_gamma_Spada2001MNRAS,angelakis_gamma_Guetta2004AnA}. Alternatively, it has been suggested that light-house effect could be causing variability in cases of rotating helical jets or the helical trajectories of plasma elements \\citep[e.g.][]{angelakis_gamma_Begelman1980Natur,angelakis_gamma_Camenzind1992AnA}.  A powerful tool in the investigation of these details and for understanding the dynamics, the physics and the structure of the radiating regions is the analysis of the variability of their SEDs at frequency ranges as broad as possible (from radio to TeV). Such an approach can shed light on linear scales inaccessible even to interferometric techniques and most importantly discriminate between different variability scenarios. Particularly, the presence of {\\sl Fermi-GST} is providing a unique opportunity for these studies providing the $\\gamma$-ray data of unprecedented quality.  To benefit fully from the newly flying {\\sl Fermi-GST} telescope, a broad collaboration of facilities that aim at understanding the blazar phenomenon via the broad-band SED variability monitoring has been initiated by \\cite{angelakis_gamma_Fuhrmann2007AIPC}. A sample of 61 selected sources are monitored monthly and in tight coordination with other facilities covering the cm, mm, optical and infrared bands (Fuhrmann et al. in prep.). Here we present a very brief introduction of the activities in the centimetric band with the 100~m Effelsberg telescope and show some preliminary examples. ",
        "conclusions": "Provided that the calibrators are presumably not intrinsically variable, figure~\\ref{angelakis_gamma_3C286} serves as an excellent means of estimating the system repeatability. In table~\\ref{angelakis_gamma_t2} the calibration precision is of the order of a few percent. Factors that induce such variations are primarily tropospheric variations, opacity effects, confusion and system instabilities. At higher frequencies (23.05 GHz and above), the atmospheric opacity becomes particularly important.  The plots in figure~\\ref{angelakis_gamma_f2} show how intensive and fast the variability of the spectrum may be. This very fact urges for dense frequency coverage especially of the millimetre regime from where the evolution of a flaring event in the radio band is expected to start. The evolution of the spectrum over time is among the first-priority studies. That requires tight coordination of the different participating stations as has been managed so far. The activity seen in the centimetre band is noteworthy. Whether this activity is correlated with that at other bands and especially that at high energies, or with changes in the structure is one of the first questions to be addressed. The behaviour of the polarisation when such an event occurs is also among the most important issues to be investigated."
    },
    "0809/0809.4190_arXiv.txt": {
        "abstract": "{ In June--September 2006 the Be/X-ray binary EXO~2030+375 experienced the second giant outburst since its discovery. The source was shown to have a complicated pulse-averaged X-ray spectral continuum with possible evidence of cyclotron absorption features. In this paper we present the first pulse-phase resolved analysis of the broad band X-ray spectra of EXO~2030+375 obtained with the \\textsl{INTEGRAL} observatory close to the maximum and during the decay phase of the giant outburst. We report a strong variability of the spectrum with pulse phase. Alternative spectral continuum models are discussed. The dependence of the spectral parameters on pulse phase during the maximum of the outburst and the evolution of the pulse profiles with time are qualitatively consistent with the pulsar's emission diagram changing from the fan-beam geometry close to the maximum of the outburst to a combination of pencil and fan beams (of comparable intesities) at the end of the decay phase. Evidence of a cyclotron absorption line around 63\\,keV at the pulse phase interval preceeding the main peak of the pulse profile is present in the spectrum obtained close to the maximum of the outburst. } ",
        "introduction": "}  The transient accreting pulsar EXO 2030+375 belongs to the most common type of X-ray pulsar systems -- the Be/X-ray binaries. Such systems form a subclass of high mass X-ray binaries. They consist of a pulsar and a Be (or Oe) companion, a main-sequence star of spectral type B (or O) that shows Balmer emission lines \\citep[see e.g.][for a review]{Slettebak88}. The line emission is believed to be associated with an equatorial outflow of material expelled from the rapidly rotating Be star that probably forms a quasi-Keplerian disk around its equator \\citep{Hanuschik96,Quirrenbach_etal97}. If the disk reaches a radius comparable to the periastron separation, then disk gas accreted by the neutron star can power a significant (and usually transient) X-ray source.  Be/X-ray binaries typically show two types of outburst behavior: \\begin{enumerate} \\item Normal (or type I) outbursts. They are characterized by relatively low X-ray luminosities $L_{\\rm X} \\sim 10^{36}{\\rm-}10^{37}\\,{\\rm erg~s}^{-1}$, low spin-up rates (if any), and recurrence at the orbital period (or its multiples). Such outbursts last from days to weeks and are associated with the periastron passages of the neutron star. \\item Giant (or type II) outbursts. They are characterized by higher X-ray luminosities $L_{\\rm X} \\gtrsim 10^{37}\\,{\\rm erg~s}^{-1}$ and high spin-up rates. Such outbursts occur irregularly. They last several weeks and are not correlated with any particular orbital phase. The typical time between outbursts is around several years. Giant outbursts are thought to stem from a dramatic expansion of the disk surrounding the Be star, leading to the formation of an accretion disk around the compact object. \\end{enumerate}  EXO 2030+375 is one of the best-studied Be/X-ray binaries. It was discovered with the \\textsl{EXOSAT} satellite during a giant outburst in 1985 \\citep{Parmar_etal89a}. The optical companion of the pulsar is a B0 Ve star identified by optical and infrared observations of the \\textsl{EXOSAT} error circle \\citep{MotchJanot87,Janot_etal88,Coe_etal88}. The orbital period and eccentricity of the system are $\\sim$46\\,d and $\\sim$0.42, respectively \\citep{Wilson_etal02}. The period of X-ray pulsations is $\\sim$42\\,s. There were two giant outbursts in the history of observations of EXO~2030+375. During the first one in 1985 (when the pulsar was discovered), the X-ray luminosity of the source reached a value of $L_{\\rm 1-20~keV}\\sim 2 \\times 10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ (assuming a distance of 7.1 kpc, \\citealt{Wilson_etal02}). The spin frequency of the pulsar changed dramatically, with a spin-up timescale $-P/\\dot P \\sim 30$\\,yr indicating the formation of an accretion disk around the neutron star. The second giant outburst took place in June--September 2006 \\citep{CorbetLevine06,Klochkov_etal07} and was again accompanied by a strong spin-up of the neutron star. The X-ray luminosity at the maximum of the outburst was slightly lower than during the 1985 giant outburst: $L_{\\rm 1-20~keV}\\sim 1.2 \\times 10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ \\citep{WilsonFinger06}.  During the 2006 giant outburst the source was observed several times with the \\textsl{INTEGRAL} satellite. A preliminary pulse-averaged spectral analysis of some of these observations are presented in \\citet{Klochkov_etal07}. A detailed analysis of pulse-averaged \\textsl{RXTE} spectra obtained during the outburst was performed by \\citet{Wilson_etal08}. It has been shown that the pulse averaged X-ray continuum of the source has a complicated shape and cannot be modeled by a simple power law/cutoff model. \\citet{Wilson_etal08} include an absorption line at $\\sim$10\\,keV (which they interpreted as a cyclotron line) in their spectral model, whereas \\citet{Klochkov_etal07} have shown that the spectrum can be fitted equally well without the absorption line, but including a broad emission ``bump'' at $\\sim$15\\,keV.  Here we present for the first time pulse-phase resolved broad band (3--150\\,keV) spectra of EXO~2030+375 during a giant outburst. For our analysis we used all the available \\textsl{INTEGRAL} data taken during the 2006 outburst. The X-ray continuum of the source shows strong variability with pulse phase, with some features present only at particular pulse-phase intervals. The description of observations that we used is provided in Sect.~\\ref{obs}. Details of data processing are described in Sect.~\\ref{data}. Sections~\\ref{max} and \\ref{decay} are devoted to the analysis of the data close to the maximum and during the decay of the outburst, respectively. The results are discussed in Sect.~\\ref{discuss} and briefly summarized in Sect.~\\ref{concl}. ",
        "conclusions": "}  We used \\textsl{INTEGRAL} observations to study the pulse-phase dependence of the broad band X-ray spectrum of EXO~2030+375 close to the maximum and during the decay of its 2006 giant outburst. This is the first pulse-phase resolved spectral study of the source. In all observations, significant pulse phase variability of the X-ray continuum was observed.  Alternative spectral continuum models are discussed. We argue that the interpretation of the feature at $\\sim$10\\,keV as a cyclotron absorption line proposed previously is questionable.  Pulse-phase dependencies of the continuum parameters close to the maximum of the outburst, as well as the evolution of the pulse profiles from the maximum to the end of the outburst, is qualitatively consistent with the picture where the pulsar's emission diagram changes from the fan-beam configuration close to the maximum of the outburst to a combination of pencil and fan beams (whose amplitudes are comparable) at the end of the decay phase.  Evidence of an absorption line at $\\sim$63\\,keV is found during the maximum of the outburst at a narrow phase interval preceeding the main peak of the pulse profile. This feature can be interpreted as the first harmonic of the previously reported cyclotron line at $\\sim$36\\,keV."
    },
    "0809/0809.4645_arXiv.txt": {
        "abstract": "Modern radio telescopes are extremely sensitive to plasma on the line of sight from a radio source to the antenna.  Plasmas in the corona and solar wind produce measurable changes in the radio wave amplitude and phase, and the phase difference between wave fields of opposite circular polarization.  Such measurements can be made of radio waves from spacecraft transmitters and extragalactic radio sources, using radio telescopes and spacecraft tracking antennas.  Data have been taken at frequencies from about 80 MHz to 8000 MHz.  Lower frequencies probe plasma at greater heliocentric distances.  Analysis of these data yields information on the plasma density, density fluctuations, and plasma flow speeds in the corona and solar wind, and on the magnetic field in the solar corona. This paper will concentrate on the information that can be obtained from measurements of Faraday rotation through the corona and inner solar wind.  The magnitude of Faraday rotation is proportional to the line of sight integral of the plasma density and the line-of-sight component of the magnetic field.  Faraday rotation provides an almost unique means of estimating the magnetic field in this part of space.  This technique has contributed to measurement of the large scale coronal magnetic field, the properties of electromagnetic turbulence in the corona, possible detection of electrical currents in the corona, and probing of the internal structure of coronal mass ejections (CMEs).  This paper concentrates on the search for small-scale coronal turbulence and remote sensing of the structure of CMEs.  Future investigations with the Expanded Very Large Array (EVLA) or Murchison Widefield Array (MWA) could provide unique observational input on the astrophysics of CMEs. ",
        "introduction": "Understanding the physics of the solar corona and its transition to the solar wind requires specifying the plasma physics properties in the region from the coronal base to heliocentric distances of tens of solar radii. We require knowledge of parameters such as the plasma density, vector magnetic field, plasma flow velocity, and ion and electron temperatures. Properties of turbulence are also important, so that we may test various theories which invoke turbulence as an important agent.  It is in this part of space that radioastronomical propagation measurements can play an important role.  The basic idea is to observe a radio source which is viewed through the corona.  This radio source can be an extragalactic radio source like a radio galaxy or a quasar, or the transmitter of a spacecraft.  A cartoon illustrating the geometry is shown in Figure 1 of \\cite{Spangler02}. The radio waves propagate through the corona or the inner solar wind, and are modified in their amplitude and phase.  A radio telescope on Earth can measure these amplitude or phase changes, and infer properties of the intervening plasma. Further details are discussed in the literature; illustrative papers are those of \\cite{Bird90} and \\cite{Spangler02}.  This paper will emphasize observation of Faraday rotation in the plasma of the outer corona and inner solar wind.  A linearly polarized radio wave propagating through a magnetized plasma will undergo a rotation in the plane of linear polarization, with the change in polarization position angle $\\Delta \\chi$ given by \\begin{equation} \\Delta \\chi = \\left[ \\left( \\frac{e^3}{2 \\pi m_e^2 c^4}\\right) \\int_L n_e \\vec{B} \\cdot \\vec{dz} \\right] \\lambda^2 \\end{equation} Equation (1) is in cgs units.  The term in square brackets on the right hand side is referred to as the {\\em rotation measure} ($RM$).   It is conventionally reported in SI units of radians/m$^2$; the cgs value of the rotation measure is converted to SI by multiplying by a factor of $10^4$.  The geometry of coronal Faraday rotation is illustrated in Figure 1 of \\cite{Spangler05}.  Radio waves from a linearly polarized radio source, in this case a radio galaxy or quasar, propagate towards the Earth and pass close to the Sun.  Because of the higher plasma density and magnetic field strength close to the Sun, the measured $RM$ is mainly determined by plasma close to the ``proximate point'', i.e. the point on the line of sight which is closest to the Sun.  The heliocentric distance of the proximate point is called the ``impact parameter''.  Our investigations (e.g. \\cite{Sakurai94,Mancuso99,Mancuso00,Spangler00,Spangler05,Ingleby07,Spangler07}) have used the Very Large Array (VLA) of the National Radio Astronomy Observatory to make such measurements.  With the VLA, there are many radio galaxies and quasars which are sufficiently strong to make these measurements. An observation of one or more such sources juxtaposed to the Sun can be made almost every day.  Radio galaxies and quasars radiate by synchrotron radiation, and are linearly polarized at the level of several percent or more. ",
        "conclusions": "Faraday rotation measurements on radio waves which have propagated through the solar corona and inner solar wind provide a unique plasma diagnostic on that region of space. Measurements to date have provided information on the coronal magnetic field, the amplitude and spatial power spectrum of MHD turbulence in the corona, estimates of electrical currents, and the plasma structure of coronal mass ejections.  Two new instruments which will become available in the next few years, the Extended Very Large Array (EVLA) and the Murchison Widefield Array (MWA), will permit much better measurements, as well as measurements of a different type.  \\noindent {\\bf Acknowledgements.}  We thank the US National Science Foundation, Division of Atmospheric Sciences, for supporting this research via grant ATM03-54782. We also thank the National Radio Astronomy Observatory for allocating significant amounts of VLA observing time to this research program."
    },
    "0809/0809.1988_arXiv.txt": {
        "abstract": "We review the expected properties of Pop III and very metal-poor starburst and the behaviour the \\lya\\ and \\Heiiuv\\ emission lines, which are most likely the best/easiest signatures to single out such objects. Existing claims of Pop III signatures in distant galaxies are critically examined, and the searches for \\Heiiuv\\  emission at high redshift are summarised. Finally, we briefly summarise ongoing and future deep observations at $z > 6$ aiming in particular at detecting the sources of cosmic reionisation as well as primeval/Pop III galaxies. ",
        "introduction": "Finding the first stars and galaxies remains a major observational challenge in astrophysics. Searches for metal-free (Pop III) or extremely metal-poor individual stars or for ensembles of stars (clusters, proto-galaxies, populations of galaxies, etc.) are ongoing both nearby (Pop III stars in the halo of our Galaxy), and in galaxies out to the highest redshifts currently known.  Astronomers have been quite inventive in searching for Pop III or metal-poor stars, exploiting many possible direct and indirect signatures; some of these methods will be discussed below. A more detailed account on primeval galaxies, some of the physics related to these objects, an overview on searches etc.\\ is presented in the lectures of Schaerer (2007).  Although some studies may have found signatures of Pop III stars, none of those claims is very strong (see discussion below), and my personal opinion -- taking a somewhat conservative attitude -- is that Pop III still remain to be found. However,  there is good hope that this goal should be reached in the quite near future with current or forthcoming instrumentation, as I will sketch below. ",
        "conclusions": ""
    },
    "0809/0809.3669_arXiv.txt": {
        "abstract": "I review the status of the Radio Cerenkov detection technique in searches for ultra-high energy (UHE) neutrinos of cosmic origin. After outlining the physics motivations for UHE neutrino searches, I give an overview of the status of current and proposed experiments in the field. ",
        "introduction": "\\label{intro} Over the past century, cosmic radiation from sources outside the earth's atmosphere has been measured over 14 orders of magnitude in energy. At $10^{18}$~eV, cosmic rays are incident on the earth's atmosphere at a rate of one particle/km$^{2}$/year. Cosmic ray experiments HiRes and Auger~\\cite{hiresauger} both see a break in the spectrum near $10^{19.5}$~eV and this was expected; at this energy, protons and heavy nuclei reach their threshold for interacting with cosmic microwave background (CMB) photons through what is known as the Greisen-Zatsepin-Kuzmin (GZK) process~\\cite{gzk}. This process can, for example, occur through a $\\Delta$ resonance: \\begin{center} \\begin{equation} \\label{eq:gzk} p + \\gamma_{\\mathrm{CMB}} \\rightarrow \\Delta^{*} \\rightarrow n ~\\pi^{+} \\rightarrow n ~e^{+} ~\\nu_{\\mu} ~\\nu_{e} ~\\bar{\\nu_{\\mu}} \\end{equation} \\end{center} The GZK process slows a cosmic ray above threshold within approximately 50~Mpc of its source. Above a few TeV, gamma rays too are absorbed before reaching us, but in the infrared background radiation through pair production~\\cite{gammaraycutoff}.  While the GZK process prevents us from observing distant cosmic rays above the cutoff, the same process is a source of neutrinos, as seen in Equation~\\ref{eq:gzk}, which can travel cosmic distances unattenuated. In addition, any photohadronic process that produces cosmic rays within an astronomical source would be expected to produce neutrinos from the decay of charged pions. Figure~\\ref{fig:neutrinos} includes a range of models for the expected neutrino flux from the GZK process in the UHE region and current limits on diffuse neutrino fluxes.   To date, the only observed sources of neutrinos outside of the earth's atmosphere have been the sun~\\cite{sun} and Supernova~1987a~\\cite{supernova1987a}. No diffuse cosmic neutrino flux has been measured, though such a measurement is crucial to a complete picture of the UHE universe. Neutrinos are the only messengers above $10^{19.5}$~eV that can probe distances greater than hundreds of Mpc. Undeflected by magnetic fields, neutrinos point back to their place of origin, and above $10^{17}$~GeV would produce interactions with center-of-mass energies exceeding that of typical interactions at the LHC.  Several experiments are searching for a cosmic flux of neutrinos by detecting the visible Cerenkov radiation produced by particle tracks~\\cite{otherexperiments}, the largest being the AMANDA-II experiment and its km$^3$ successor IceCube~\\cite{amandaiiicecube}.  Their observations are consistent with the expectation from atmospheric neutrino background.  They constrain the neutrino flux below $10^{18}$~eV. ",
        "conclusions": "The field of UHE neutrino detection using the radio Cerenkov technique has become a mature field, with existing experiments beginning to probe the expected neutrino flux from the GZK process. The technique holds the capability of moving beyond the discovery stage and into an era of making particle- and astrophysics measurements with 10-100~UHE neutrinos per year.  I have described several proposed projects that are being designed to measure neutrinos rates at that level once operating at full scale."
    },
    "0809/0809.3802_arXiv.txt": {
        "abstract": "Yan and Lazarian have proposed a new technique through which the magnetic field geometry in the diffuse interstellar medium, or in circumstellar matter, could be determined from the linear polarization of interstellar absorption or fluorescence emission lines from ions pumped by an anisotropic illuminating flux.  New long-slit spectroscopic observations of the reflection nebula NGC2023, obtained with the Southern African Large Telescope Robert Stobie Spectrograph (RSS), have detected a number of atomic fluorescence  lines of OI, NI, SiII, and FeII for the first time in a neutral medium.  A model which predicts these lines and others illustrates which lines would be appropriate targets for an RSS spectropolarimetric investigation of this new diagnostic. ",
        "introduction": "Atomic resonance/ fluorescence lines may be linearly polarized when an\\-iso\\-tro\\-pic optical pumping produces an an\\-iso\\-tro\\-pic angular momentum distribution (``alignment'') within the ground state \\citep{yl2006,yl2007,yl2008} (hereafter YL06,YL07,YL08).  This occurs if the pumping photon rate is greater than the collision rate, for instance within 1-10 pc of an OB star.  In the presence of a magnetic field, ``magnetic realignment'' then occurs when the Larmor frequency is greater than the photon rate. For the ``saturated'' case (valid in the diffuse ISM, where any field $> 0.1$  $\\mu$Gauss causes realignment), the polarization depends on the 3D geometry of the magnetic field, the ion ground state configuration and sometimes on the pumping spectrum.  In circumstellar matter, $10-10^4$ $\\mu$Gauss is required for realignment, and the polarization can depend also on the field strength. Related effects have been seen in the sun \\citep{s1994}, but have not yet been demonstrated elsewhere.  As a diagnostic of the magnetic field, magnetic realignment is potentially more powerful than the 21 cm Zeeman Effect, since it is sensitive to weaker fields and works in hot gas, and more powerful than dust alignment, since it is sensitive to 3D geometry, gas properties and velocity.  This paper explores observational aspects of this effect.  \\begin{figure}[!ht] \\plotfiddle{NordsieckFig1.eps}{0.75in}{0}{45}{45}{-140}{-250} \\caption{Geometry for Magnetic Realignment} \\end{figure}  Figure 1 shows the observational geometry.  A gas cloud with a magnetic field at an angle  $\\theta_B$ from the line of sight is illuminated by a nearby star (the ``pumping star'') at an angle  $\\theta_r$ from the field ($\\theta$ from the line of sight).  In the absorption case, flux from a star behind the cloud (the ``target star'', which may be the same as the pumping star) is absorbed by ions in the cloud.  If the ions are aligned, the resultant absorption line is linearly polarized, with a position angle either parallel or perpendicular to the projected magnetic field, and a degree that depends on  $\\theta_B$ and $\\theta_r$.  For absorption, YL06 show that a non-zero polarization requires ions with at least three fine states in the ground level $(J \\geq 1)$.  In neutral gas, the most common ions are NI, OI, SII, and FeII.  The spectral resolution required is $R = \\lambda / \\Delta\\lambda \\ga 20,000$, to resolve the interstellar lines.  Since the absorption lines are almost entirely in the FUV, this requires a high resolution FUV spectropolarimeter \\citep{n2003}.  In the emission case, the observed light is light from the pumping star which is scattered into the line of sight.  The unique polarization signature of magnetic realignment is a distortion of position angles of the scattered emission lines from the centrosymmetric pattern expected for pure reflection polarization. YL07 and YL08 show that ions need at least three fine \\emph{or hyperfine} states in the ground level ($F \\geq 1$) for a nonzero magnetic field effect.  This adds the resonance lines of NaI, KI, and AlII to the roster.  Also, transitions that do not return to the ground state (fluorescence) exhibit the effect, thus opening up numerous targets in the visible and NIR.  The spectral resolution required is more modest: for sensitivity against scattered dust continuum $R \\ga 5,000$ is adequate.   Thus in principle, this effect can be observed with a ground-based moderate-resolution spectropolarimeter.  This paper describes the beginnings of such an investigation, targeted at fluorescence lines in the visible. ",
        "conclusions": ""
    },
    "0809/0809.3205_arXiv.txt": {
        "abstract": "The origin of ultra-high energy cosmic rays is discussed in light of the latest observational results from the Pierre Auger Observatory, highlighting potential astrophysical sources such as active galactic nuclei, gamma-ray bursts, and clusters of galaxies. Key issues include their energy budget, the acceleration and escape of protons and nuclei, and their propagation in extragalactic radiation and magnetic fields. We briefly address the prospects for Telescope Array and future facilities such as JEM-EUSO, and also emphasize the importance of multi-messenger X-ray and gamma-ray signatures in addition to neutrinos as diagnostic tools for source identification. ",
        "introduction": "\\label{sec:intro} Several decades after their discovery, the origin of ultra-high energy cosmic rays (UHECRs), cosmic particles with energies $10^{18}$-$10^{20}$ eV and above, continue to be one of the biggest mysteries in astrophysics \\cite{nw00}. During the last few years, the field has witnessed revolutionary advances, thanks to the advent of new generation, hybrid detector facilities spearheaded by the Pierre Auger Observatory in the southern hemisphere \\cite{wat07}. The Telescope Array \\cite{fuk07} has begun operation in the northern hemisphere and we eagerly await their data. Expectations are already high for future projects such as Auger North \\cite{nit07}, or the Extreme Universe Space Observatory on the Japanese Experiment Module (JEM-EUSO), to be deployed on the International Space Station \\cite{tes07}. This article, by no means a thorough review, discusses selected topics in the rapidly developing field of UHECRs, focusing on the theory of potential astrophysical sources confronted with the latest observational results. Some of the issues below are covered in more detail in earlier articles \\cite{ino07}. ",
        "conclusions": "\\label{sec:conc} Besides AGNs, GRBs and clusters, other proposals for the astrophysical origin of UHECRs not covered in this brief review include starburst galaxies \\cite{anc99}, extragalactic magnetars \\cite{aro03}, dormant black holes \\cite{bol99}, etc. (see \\cite{hil84}). The true source of UHECRs could be one or more of these; alternatively, we should not discount the possibility that it is actually none of them, and the truth something totally unexpected.  In this regard, we may recall the history of GRB research. Before the launch of BATSE in 1991, the sample of GRBs numbered $\\lesssim 200$, with no clear anisotropy observed in the sky distribution. However, most believed, based on theoretical plausibility, that they were generated by Galactic neutron stars, while a very small minority advocated neutron star mergers at cosmological distances. After several years and a sample of more than 3000 bursts, the currently most favored scenario for the progenitors of long GRBs is related to the collapse of massive stars \\cite{woo93}, a concept that didn't exist before BATSE. The lesson is that we should be cautious about jumping to conclusions on the origin of UHECRs based on just 20 observed events and our limited theoretical understanding of cosmic objects.  Regardless of theoretical prejudice, the quest for the solution of the UHECR mystery is progressing in earnest with ongoing measurements by Auger and Telescope Array. An order of magnitude greater statistics for the highest energy events expected from JEM-EUSO or Auger North+South, and the combined effort of upcoming neutrino, X-ray and gamma-ray observations is expected to lead us toward the true answer in the near future.   \\ack The author is thankful for past and ongoing collaborations with F. Aharonian, E. Armengaud, K. Asano, N. Kawakatu, P. Meszaros, F. Miniati, K. Murase, K. Sato, G. Sigl, N. Sugiyama, H. Takami and T. Yamamoto."
    },
    "0809/0809.4688.txt": {
        "abstract": "The Swift mission has discovered an intriguing feature of Gamma-Ray Burst (GRBs) afterglows, a phase of shallow decline of the flux in the X-ray and optical lightcurves. This behaviour is typically attributed to energy injection into the burst ejecta. At some point this phase ends, resulting in a break in the lightcurve, which is commonly interpreted as the cessation of the energy injection. In a few cases, however, while breaks in the X-ray lightcurve are observed, optical emission continues its slow flux decline. This behaviour suggests a more complex scenario. In this paper, we present a model that invokes a double component outflow, in which narrowly collimated ejecta are responsible for the X-ray emission while a broad outflow is responsible for the optical emission. The narrow component can produce a jet break in the X-ray lightcurve at relatively early times, while the optical emission does not break due to its lower degree of collimation. In our model both components are subject to energy injection for the whole duration of the follow-up observations. We apply this model to GRBs with chromatic breaks, and we show how it might change the interpretation of  the GRBs canonical lightcurve. We also study our model from a theoretical point of view, investigating the possible configurations of frequencies and the values of GRB physical parameters allowed in our model. ",
        "introduction": "\\label{intro}  Since its launch, the \\textit{Swift} mission (Gehrels et al. 2004) has allowed us to observe the emission from Gamma-Ray Burst (GRB) afterglows in the X-ray and UV/Optical from as early as $\\sim$1 minute after the burst trigger by means of the X-ray Telescope (XRT, Burrows et al. 2004) and UV/Optical Telescope (UVOT, Roming et al. 2005). This unprecedented response time has allowed us to unveil the early behaviour of GRB afterglow lightcurves, which turn out to be more complex than expected. Typically, at the end of the prompt emission the X-ray flux $F$ exhibits a rapid decay. This can be modeled with a powerlaw $F\\sim t^{-\\alpha{_1}}$ with slope $\\alpha_{1} \\sim 3 - 5$ (Tagliaferri et al. 2005). This phase, which usually lasts hundreds of seconds, is widely interpreted as the tail of the prompt emission (Kumar \\& Panaitescu 2000; for a review, see Zhang et al. 2006). After that, the X-ray flux decays in a much shallower way, forming a ``plateau'' with a slope $\\alpha_{2} \\sim 0.1 - 0.8$. The spectrum in this phase can be different from that observed during the fast decay, which indicates a different physical origin. The duration of the slow decline is a few thousands of seconds (O'Brien et al. 2006, Willingale et al. 2007). After this time, a break occurs and the lightcurve becomes steeper, with a powerlaw slope of $\\alpha _{3}\\sim1.3$. Indeed, this latter phase was studied well prior to the launch of \\textit{Swift} (e.g. De Pasquale et al. 2006; Gendre et al. 2006) and it is understood to be emission from synchrotron radiation, resulting from a shock produced by the expansion of the burst ejecta into the circumburst medium (Meszaros \\& Rees 1997). Occasionally, a further break may occur a few days after the trigger, leading to a segment with decay slope of $\\alpha_{4}\\sim2$ . This steep decay can be interpreted as the signature of collimated outflow (Sari et al. 1999). Overall, this evolution of the X-ray flux  is now referred as the ``canonical'' X-ray lightcurve (Nousek et al. 2006). In the optical band, the flux decays with a similar range of slopes to those of the X-ray, with the exception of the initial fast decay phase, which is usually absent (Oates et al. in preparation).\\\\  The slow decay is probably the most perplexing among the novel aspects discovered by \\textit{Swift}, and several models have been proposed to explain it (see e.g. Zhang 2007 for a complete review). These models in general fall into three main classes: i) energy injection into the burst ejecta, either in the form of Poynting flux or late time shells of jecta; ii) a non uniform angular energy distribution in the jet or a jet seen off-axis, so that a fraction of the early afterglow emission is not fully beamed towards the observer; iii) a change of the microphysical parameters that leads to an increase in the conversion efficiency of the ejecta energy to radiation. % In the majority of GRBs, achromatic breaks are observed, %both in the X-ray and in the optical bands, between the slow decay %phase (with index $\\alpha_2$) and the normal decay phase (with index %$\\alpha_3$).  Puzzlingly, in a few \\textit{Swift} GRBs the slow decline phase ends with a ``chromatic break'' (Panaitescu et al. 2006a; see also Melandri et al. 2008): i.e. a transition from the shallow to the normal decay appears in the X-ray band but is absent in the optical band, where the flux continues to decline at a slow rate. This feature is very hard to explain with any model that predicts a single origin for the X-ray and optical emission. In the attempt to solve this problem, Ghisellini et al. (2007) suggested a model in which the optical emission is caused by the interaction between the ejecta and the circumburst medium, while the X-ray radiation is produced by internal shocks occurring in collimated shells emitted by the GRB central engine at relatively late times. If the Lorentz factor $\\Gamma$ of these shells decreases with time, a ``jet-like'' break will be detected (in the X-ray band only) at the time in which $\\Gamma^{-1} = \\theta$, where $\\theta$ is the opening angle of ejecta. %This model, however, is not easily reconciled with the %observed spectral changes between the fast and slow decay phases. An alternative scenario, proposed by Genet et al. (2007) and Uhm \\& Beloborodov (2007), assumes that both the X-ray and optical emission is due to reverse shocks crossing the shells. % In this scenario, if %the Lorentz factor of late shells decreases down to $\\Gamma \\sim$ a %few and only a small fraction ($<5\\%$) of the kinetic energy is %transferred to the electron population, the model predicts %lightcurves similar to those observed. However, this model requires that the external shock emission is basically turned off. This may need conditions difficult to meet. Other authors argue that the jet breaks are actually hidden in the optical lightcurves (Curran et al. 2007) and/or less evident than expected (Panaitescu et al., 2007a, Liang et al. 2007). In Panaitescu (2007b), the author proposes a complex scenario, in which the plateau, the flares and the chromatic breaks seen in the X-ray lightcurve are caused by scattering of the forward-shock synchrotron emission by a relativistic outflow, located behind the leading blast-wave. Efforts have also been made to reconcile the chromatic breaks with the scenario of an unique outflow (Panaitescu et al. 2006a), hypothesizing an evolution of the microphysical parameters, including the fractions of blast wave energy given to electrons and to the magnetic field. However, as the authors themselves pointed out, the required evolution is assumed ``ad hoc'', and still lacks a self-consistent physical explanation.   Recently, Oates et al. (2007) have investigated the case of \\textit{Swift} GRB050802, one of the bursts in the dataset of Panaitescu (2006a), which shows a very evident chromatic break. They found that the observed late SED cannot be reproduced by models based on single component outflow, and proposed a model based on two outflows: a narrow one responsible for the X-ray emission, and a wider one that powers the optical emission. Both outflows receive continuous energy by means of shells emitted at late times or in the form of Poynting flux. The break in the X-ray lightcurve, in this scenario, is interpreted as a jet break, and there is no discontinuation of energy injection. The ``normal'' decay phase is then a post jet break phase with a slope less steep than usual because of the energy injection. The fact that the optical lightcurve does not show a break within the time of the follow-up observations is naturally explained by the lower degree of collimation of the outflow responsible for it. In this paper, we conduct a detailed analysis of a sample of other GRBs that are reported to have chromatic breaks, showing that this model can potentially interpret the observed behaviour. We also discuss how this scenario may change our interpretation of the canonical lightcurve of GRBs and the deep implications that this change of perspective may have on our understanding of GRB physics. This paper is organized as follows. In \\S~\\ref{data_analysis} we introduce the dataset and the data analysis,  while in \\S~\\ref{res} we present the application of the model to the GRB sample. Discussion and conclusions follow in \\S~\\ref{disc} and \\S~\\ref{conc}, respectively. ",
        "conclusions": "\\label{conc}  In this paper, we have reanalysed the full sample of \\textit{Swift} GRBs whith chromatic breaks, originally discussed by Panaitescu et al. (2006a). In addition, we have also studied GRB 060605, another \\textit{Swift} burst with good quality XRT and UVOT data and a chromatic break in the XRT lightcurve. We have shown how our model, based on a prolonged energy injection into a double component outflow and a jet break, is physically plausible and can well explain the behaviour of the optical and X-ray emission of GRB050319, 060505 and GRB050802 (see also Oates et al. 2007). %While GRB060505 shows that the optical and X-ray emission should be %emitted by physically different mechanisms, this burst cannot be %interpreted in our model, at least with the simplest assumptions. GRB050922c has been shown not to require a chromatic break. We note that our model can also be applied to the other GRBs with claim of chromatic breaks published in Panaitescu et al.~(2006a) and might, in principle, be extended to all GRBs featuring the canonical lightcurve (Nousek et al. 2006), even those without chromatic breaks. % we remind, though, %that some degree of tuning of physical parameters values are needed. We emphasize that it would have not been possible to derive our conclusions if we had considered the X-ray data only, since GRB050319, GRB060605 and GRB050802 exhibit a canonical X-ray lightcurve. Instead, the combined optical and X-ray analysis has shown that the component responsible for the optical is uncoupled from the outflow that produces the X-ray emission. In our model, the ejecta responsible for the X-ray emission are narrowly beamed, and undergo an early jet break that explains the chromatic break seen in the X-ray only. Our model of combined jet expansion and energy injection may have deep consequences on our understanding of the GRB, since it calls for a revision of the physics processes that take place in these objects. % Our model could be applied to all GRBs, including those displaying %the ``usual'' achromatic breaks, and may thus have deep consequences %on our understanding of the GRB physics."
    },
    "0809/0809.0206_arXiv.txt": {
        "abstract": "We report the identification of a source coincident with the position of the nearby type II-P supernova (SN) 2008bk in high quality optical and near-infrared pre-explosion images from the ESO Very Large Telescope (VLT). The SN position in the optical and near-infrared pre-explosion images is identified to within about $\\pm$70 and $\\pm$40 mas, respectively, using post-explosion $K_{\\rm s}$-band images obtained with the NAOS CONICA adaptive optics system on the VLT. The pre-explosion source detected in four different bands is precisely coincident with SN 2008bk and is consistent with being dominated by a single point source. We determine the nature of the point source using the STARS stellar evolutionary models and find that its colours and luminosity are consistent with the source being a red supergiant progenitor of SN 2008bk with an initial mass of 8.5$\\pm$1.0 M$_{\\odot}$. ",
        "introduction": "The red supergiant progenitors of several type II-P supernovae (SNe) have now been directly identified in pre-explosion observations. All these are moderate mass red supergiants in the range $\\sim$7-16 M$_{\\odot}$ (Smartt et al. 2004; Van Dyk, Li \\& Filippenko 2003; Maund \\& Smartt 2005; Maund, Smartt \\& Danziger 2005; Hendry et al. 2006; Li et al. 2006, Li et al. 2007). Recently, Smartt et al. (2008) have presented a volume limited systematic study of the progenitors of 20 type II-P SNe with high-quality pre-explosion images available. They find a minimum initial mass of 8.5$^{+1}_{-1.5}$ M$_{\\odot}$ for the progenitors and suggest that there is a 'red supergiant problem' with the red supergiant progenitors more massive than $\\sim$17 M$_{\\odot}$ remaining undetected.  In this letter we report the identification of the progenitor of the type II-P event, SN 2008bk, making use of adaptive optics (AO) assisted target of opportunity (ToO) observations of the SN using the ESO Very Large Telescope (VLT). SN 2008bk was discovered by Monard (2008) on 2008 Mar. 25.14 UT in a nearby (3.9 Mpc, Karachentsev et al. 2003) Scd-type galaxy NGC 7793. Li et al. (2008) obtained a more precise position for the SN which is 9\".2 east and 126\".4 north of the host galaxy nucleus. It was spectroscopically classified by Morrell and Stritzinger (2008) as a type II-P similar to SN 1999em at 36 days after explosion on 2008 Apr. 12.4 UT. This classification is also supported by amateur photometry$\\footnote{http://www.astrosurf.com/snweb2/2008/08bk/08bkMeas.htm}$ showing a very flat plateau, consistent with a normal type II-P SN. NGC 7793 has a wealth of pre-discovery images available and based on their astrometry, Li et al. (2008) identified a possible progenitor star in an archival $I$-band image from VLT/FORS. Subsequently Maoz and Mannucci (2008) estimated the $J$ and $K_{\\rm s}$ magnitudes for the progenitor from archival near infra-red (NIR) pre-explosion images from VLT/ISAAC. Using accurate relative astrometry between our high-resolution post-explosion AO images and the pre-explosion images from the VLT we show that the possible progenitor identified by Li et al. and Mannucci \\& Maoz is precisely coincident with the position of SN 2008bk and characterise its properties using stellar evolutionary models. ",
        "conclusions": "We have identified a source coincident with SN 2008bk in pre-explosion VLT images in four different optical and near-IR bands making use of adaptive optics $K_{\\rm s}$-band images of the SN from VLT. The colours and luminosity of the pre-explosion source are consistent with it being a red supergiant having an initial mass of 8.5$\\pm$1.0 M$_{\\odot}$. The coincidence of the pre-explosion source with SN 2008bk makes it the fourth intermediate mass ($\\sim$8 M$_{\\odot}$) red supergiant progenitor for a type II-P SN directly detected in pre-explosion images. Our observations also demonstrate the potential of 8m-class telescopes equipped with adaptive optics (see also Gal-Yam et al. 2005, Crockett et al. 2008) in precisely identifying SN progenitors in pre-explosion images."
    },
    "0809/0809.0030_arXiv.txt": {
        "abstract": "We report deep Sub-Millimeter Array observations of 26 pre-main-sequence (PMS) stars  with evolved  inner disks. These observations measure the mass of the outer disk (r  $\\sim$20-100 AU) across every stage of the dissipation of the inner disk (r $<$ 10 AU) as determined by the IR  spectral energy distributions (SEDs). We find that only  targets with  high mid-IR excesses are detected and have disk masses in the 1-5 M$_{Jup}$ range, while  most of our objects remain undetected to sensitivity levels  of  M$_{DISK}$ $\\sim$0.2-1.5 M$_{Jup}$.   To put these results in a more general context, we collected publicly available data to construct the optical to millimeter  wavelength SEDs of over 120 additional PMS stars.  We find that the near-IR and mid-IR emission remain optically thick  in objects whose disk masses span  2 orders of magnitude ($\\sim$0.5-50 $M_{Jup})$.  Taken together,  these results  imply that, \\emph{in general,  inner disks start to dissipate only after the outer disk has been significantly depleted of mass.} This provides strong support for photoevaporation being one of the dominant processes driving disk evolution. ",
        "introduction": "The vast majority of PMS stars in any given population are either accreting Classical T Tauri Stars (CTTSs)  with excess emission extending from the  near-IR to the millimeter or weak-line T Tauri Stars (WTTSs) with bare stellar photospheres. The fact that very few objects lacking near-IR excess show mid-IR or (sub)millimeter excess emission implies that, once the inner disk dissipates,  the entire disk goes away very rapidly (Skrutskie et al. 1990; Wolk $\\&$ Walter 1996; Cieza et al. 2007).  To first order, the evolution of primordial disks is driven by viscous accretion. Viscous evolution models (Hartmann et al. 1998;  Hueso $\\&$ Guillot 2005) are broadly consistent with the observational constraints for disk masses, disk sizes,  and accretion rates as a function of time;  however, they also predict smooth,  power-law  evolution of the disk properties. This smooth evolution is inconsistent  with the very rapid disk dissipation that usually occurs after a much longer disk lifetime. Pure viscous evolution models also fail to explain the variety of SEDs of  the so called  ``transition  objects\"  (the few objects that are caught in the short transition between typical CTTSs and bare stellar photospheres).  Recent disk evolution models, known as ``UV-switch'' models, combine viscous evolution with photoevaporation by the central star (Clarke et al. 2001; Alexander et al. 2006) and are able to account for both the disk lifetimes of several million years and the short disk dissipation timescales  ($\\tau$ $<$ 0.5 Myr). According to these models, extreme ultraviolet  (EUV) photons originating in the stellar chromosphere, ionize and heat the circumstellar hydrogen. Beyond some critical radius, the thermal velocity of the ionized hydrogen exceeds its escape velocity and the material is lost in the form of a wind. At early stages, the accretion rate dominates over the evaporation rate and the disk undergoes standard viscous evolution.  Later on, as the accretion rate drops, the outer disk is no longer able to resupply the inner disk with material. At this point, the inner  disk drains on a viscous timescale and an inner hole is formed. Once this inner hole has formed, EUV radiation very efficiently photoevaporates the inner edge of the disk and the disk rapidly dissipates from the inside out. Thus, the UV-switch model naturally accounts for disk lifetimes and dissipation timescales as well as for the SEDs of  some disks suggesting the presence of large inner holes.  Photoevaporation, however, is not the only mechanism that has been proposed to explain the opacity holes of  transition disks, some of which have now been confirmed by direct submillimeter imaging  (e.g., Brown  et al. 2008). Theoretical models of the dynamical interaction between forming planets and the disk (Lin $\\&$ Papaloizous 1979, Artymowicz $\\&$ Lubow 1994) also predict the formation of inner holes and gaps, and thus planet formation quickly became one of the most exciting explanations proposed for the inner holes of transition disks (Calvet et al. 2002; D'Alessio et al. 2005). However, the planet formation process is not required to be far along in order to affect the SED of a circumstellar disk.  Once primordial sub-micron dust grains grow into somewhat larger bodies (r $\\gg$ $\\lambda$), most of the solid mass ceases to interact with the radiation, and the opacity function decreases dramatically. Dullemond $\\&$ Dominik (2005) find that grain growth is a strong function of radius it is more efficient in the inner regions where the surface density is higher and the dynamical timescales are shorter,  and hence can also produce opacity holes.  Understanding  the processes operating in  transition disks is crucial for understanding disk evolution and planet formation. However, since multiple processes can result in similar  IR SEDs, additional observational constraints are necessary to establish their relative importance. Here, we  report deep SMA observations of a sample of 26 PMS stars whose IR SEDs trace the dissipation of the inner disk (r$<$ 10 AU).  These observations provide information on  the  mass of  their  cold outer disks (r  $\\sim$20-100 AU)  and help us to distinguish between different evolutionary scenarios. ",
        "conclusions": "{\\label{DIS}}  \\subsection{The SED  evolution of PMS Stars}{\\label{SEDevo}}  In order to put the results of our SMA survey in a more general context of disk evolution,  we combined publicly available data to construct the optical, near-, mid-IR, and  (sub)millimeter wavelength SEDs of over 120 additional PMS stars. We started from the (sub)millimeter targets in Taurus and Ophichus studied by Andrews $\\&$ Williams (2005, 2007) and then searched for their  \\emph{Spitzer} fluxes in the catalogs produced  by the ``Cores to Disks\" (Evans et al. 2003) and Taurus (Padget et al. 2006) Legacy Projects \\footnote{\\tt http://ssc.spitzer.caltech.edu/legacy/all.html}. For objects with both \\emph{Spitzer} and (sub)millimeter data, we collected the near-IR photometry from the 2MASS database and, when available, the optical data from the literature. The complete SEDs of all these targets will be presented in a followup paper (Cieza et al., in prep). In this Letter, we focus on the analysis of the [K$_S$]-[24] colors of the extended sample (our SMA sample plus the objects discussed above) as a function of  disk mass. For consistency, we calculate all disk masses from Eq. 1 (or Eq. 2 if  850 $\\mu$m data are not available), assuming the following distances:  50 pc for the TW Hydra association, 125 pc for Ophuichus, and 140 pc for the Taurus and Upper Sco  regions.  In Fig.  2 we plot disk mass as a function of [K$_S$]-[24] color for the extended sample. This figure shows that our SMA study is significantly more sensitive than previous surveys.  One striking result from Figure 2 is that all the objects with disk masses larger  than $\\sim$2 M$_{Jup}$ have  [K$_S$]-[24] colors $>$ 3.5,  consistent with optically thick 24 $\\mu$m  emission (i.e. no object  in the sample has a massive disk and  optically thin 24 $\\mu$m emission). This result excludes any scenario in which most disks evolve exclusively from the inside out (e.g., an evolution dominated by grain growth).   If that were the case, we would expect to see some massive disks with optically thin 24 $\\mu$m emission.  Since the 24 $\\mu$m data probes the inner $\\sim$10 AU of the disk, while the vast majority of the disk mass is outside this radius, any disk evolving strictly from the inside out would still retain most of its mass at the point where its 24 $\\mu$m emission is transitioning from optically thick to optically thin.  Fig. 2 suggests  a scenario in which the levels of mid-IR excess emission remain constant while the mass of the disk is being depleted by 2 orders of magnitude through accretion onto the star.  Such a  scenario is illustrated by the SEDs in Fig 3. We note that the properties of the SEDs on the bottom of Fig. 3 and  the top of Fig. 1 overlap. Together, they form a sequence in which a disk loses mass,  maintaining a constant near- and mid-IR SED (Fig 3) until the mass of the disk reaches a critical level around 1 M$_{Jup}$, at which point the inner disk starts to dissipate from  the inside out  through photoevaporation (Fig 1). The SEDs in Figs. 1 and 3 are ordered,  left to right and top to bottom, to \\emph{illustrate} this evolutionary sequence.  However,  their order should not be taken literately. Other factors,  in addition to the evolutionary status of the disk, play a role in the morphology of an SED (e.g.,  disk inclination and stellar luminosity).  \\subsection{Implications for Disk Evolution}  As discussed above,  the data in Fig.  2 are inconsistent with disk evolution occurring strictly from the inside out as expected from disk evolution models dominated by grain growth (Dullemond $\\&$ Dominik, 2005). Planet formation is a considerably more complex process than grain growth. As such, its effect on the observable properties of disks are  more uncertain. However, planet-formation models still make some testable predictions. In particular, forming planets with masses larger than $\\sim$0.5-1 M$_{Jup}$ should be able to open a gap in the disk independently of the disk mass   (Edgar et al.  2007).  The low masses of  \\emph{all} the disks  in our SMA sample  suggests that  their inner holes are not driven mainly by the formation of Jovian planets.  However, the fact that the vast majority of our targets are non-accreting objects could introduce a strong bias against disks with planet-induced inner holes because, unless the planet is very massive ($>$10 M$_{Jup}$), some disk material is expected to flow  across the planet's orbit and reach the star (Lubow et al. 1999). The SEDs of our sample are most consistent with those of photoevaporating disks. Photoevaporation models, such as those presented by Alexander et al. (2006), can simultaneously explain  not only  the properties of our SMA sample, but also many of the observational results discussed in this Letter, most notably:  { \\it{a) The duration of the CTTS stage, the fact that most of their IR SEDs look alike,  and their wide range of  disk masses.}}  According to photoevaporation models,  during the first few Mys of  evolution viscous accretions dominate over photoevaporation. The  mass of the  disk is depleted by accretion onto the star,  but   the accretion rate through the disk is large enough to replenish the inner disk.  Thus,  the IR SED remains unaffected,  as seen in Fig 3.  {\\it{b) The fraction of WTTS with disks, their SEDs, and their low disk masses.}}  Photoevaporation  predicts that all disks pass through a short ($\\tau$ $<$ 0.5 Mys) inside-out clearing phase \\ once the accretion rate matches the photoevaporation rate  ($\\sim$10$^{-10}$ M$_{\\odot}$yr$^{-1}$).  During this clearing stage, the mass of the disk is predicted to be 0.05-0.5 M$_{Jup}$, depending on the exact ionizing  flux and the viscosity law.   This phase is in excellent agreement with the properties and incidence of WTTS disks (See Fig. 1).  Also, we note  that FW Tau, the transition object highlighted in Fig. 2, has a disk with a mass of $\\sim$0.4 M$_{Jup}$, within the  predicted range.  {\\it{c) The low incidence of  holes in disks that are massive or strongly accreting.}} Photoevaporation predicts that the inner disk will drain \\emph{only} after the outer disk has been significantly depleted of mass  and the accretion rate becomes very small.  Counter examples to this prediction exist, such as GM Aur and DM Tau (M$_{DISK}$ $\\sim$25 M$_{Jup}$, M$_{ACC}$= $\\sim$10$^{-8}$ M$_{\\odot}$   yr$^{-1}$, Najita et al. 2007), but their incidence seems to be of the order of a few percent. Other processes, such as grain growth or planet formation, must be responsible for their inner holes.  In practice, all the processes discussed in this Letter (grain growth, planet formation, and photoevaporation) are expected to operate simultaneously and affect one another.  Other processes such as dynamical interactions in binary stellar systems  are also likely to play a role (Ireland $\\&$ Kraus, 2008).  However, the remarkable success of the photoevaporation models accounting for the many observational results listed above strongly suggests that,  together with viscous accretion, photoevaporation is one of the dominant processes driving disk evolution. This conclusion seems to contradict the recent results by Najita et al. (2007), who find that only 2 out of the 12 transition objects they consider are consistent with photoevaporating disks. However, the different sample selection biases likely account for the discrepancy (Alexander 2008). Quantifying the relative importance of disk evolution mechanisms requires establishing specific observational metrics to distinguish among them and taking into consideration the details of  the sample selection.  Such a task is beyond the scope of this Letter, but will be attempted in a followup paper."
    },
    "0809/0809.2942_arXiv.txt": {
        "abstract": "If dark matter is made of mirror baryons, they are present in all gravitationally bound structures. Here we investigate some effects of mirror dark matter on neutron stars and discuss possible observational consequences. The general-relativistic hydrostatic equations are generalized to spherical objects with multiple fluids that interact by gravity. We use the minimal parity-symmetric extension of the standard model, which implies that the microphysics is the same in the two sectors. We find that the mass-radius relation is significantly modified in the presence of a few percent mirror baryons. This effect mimics that of other exotica\\eg quark matter. In contrast to the common view that the neutron-star equilibrium sequence is unique, we show that it depends on the relative number of mirror baryons to ordinary baryons. It is therefore history dependent. The critical mass for core collapse\\ie the process by which neutron stars are created, is modified in the presence of mirror baryons. We calculate the modified Chandrasekhar mass and fit it with a polynomial. A few percent mirror baryons is sufficient to lower the critical mass for core collapse by $\\sim 0.1$~M$_\\odot$. This could allow for the formation of extraordinary compact neutron stars with low mass. ",
        "introduction": "Observations indicate that non-luminous matter contributes a significant part of the mass in galaxies. This ``dark matter'' (DM) cannot be made of known particles, but is of unknown type. In this paper we describe some potential effects of DM on neutron stars (NS), which can be tested with future observations of NS. These effects could be important also for the interpretation of the observational results. If DM exists, it should be present in all astrophysical objects, unless there is an efficient and unexpected segregation mechanism. DM can be present already in the formation processes and it can be accreted subsequently from the environment. Both of these processes should naturally occur, but the amount of DM that is trapped inside different objects should depend on the nature of the DM particles, the type of objects and their history. In general we expect that critical phases of stellar evolution could be modified in the presence of relatively small amounts of DM.  The most studied hypothesis is that DM is made of weakly interacting massive particles (WIMPs), which naturally appears in supersymmetric extensions of the standard model (SM) of particle physics. Should they exist, WIMPs accumulate in NS due to elastic scattering on nucleons. If the mass of accreted WIMPs reaches a critical value they become self-gravitating and form a dense core supported by degeneracy pressure. The dark core collapses into a black hole if it reaches the Chandrasekhar mass, $M \\sim M_{Pl}^3/m_{DM}^2$. This limits the possible density, cross section and mass of WIMPs \\cite{Goldman:1989nd,Bertone:2007ae}. See also \\cite{Gould:1989gw}. The expected high mass of WIMPs gives a low Chandrasekhar mass for the DM core, they therefore do not affect the mass and radius of NS significantly. The annihilation of WIMPs in the DM core produces heat and old NS should therefore have a steady-state temperature of about $10^4$~K \\cite{Kouvaris:2007ay}. The coldest observed NS have temperatures in the range $10^5-10^6$~K. It is observationally difficult to detect NS with significantly lower temperatures, because the luminosity scales as $T^4$. The effects of WIMPs on NS therefore seems to be of little practical concern. WIMPs are, however, only one class of DM particles and some observations suggest that the model is oversimplified, see\\eg \\cite{Perivolaropoulos:2008ud} and references therein.  Mirror matter is another viable DM candidate that emerges if one, instead of (or in addition to) assuming a symmetry between bosons and fermions -- supersymmetry, assumes that nature is parity symmetric. It is a matter of fact that the weak nuclear force is not parity symmetric. The main theoretical motivation for the mirror-matter hypothesis is that it constitutes the simplest way to restore parity symmetry in the physical laws of nature. When Lee and Yang proposed the non-parity of weak interactions in 1956, they mentioned also the possibility to restore parity by doubling the number of particles in the SM \\cite{Lee:1956qn}. Thereby the Universe is divided into two sectors with opposite handedness that interact mainly by gravity. In the minimal parity-symmetric extension of the SM \\cite{Foot:1991bp,Pavsic:1974rq}, the group structure is $G\\otimes G$, where $G$ is the gauge group of the SM. In this model the two sectors are described by the same lagrangians, but where ordinary particles have left-handed interactions, mirror particles have right-handed interactions. Except for gravity, mirror matter could interact with ordinary matter via so-called kinetic mixing of gauge bosons, or via unknown fields that carry both ordinary and mirror charges. If such interactions exist they must be weak and we therefore neglect them here. Since photons do not interact with mirror baryons, or interact only via the weak kinetic mixing, mirror matter constitute a natural candidate for the DM in the Universe. The study of the cosmological implications of mirror matter is well-defined, because the microphysics of mirror matter is the same as that of ordinary matter. Many consequences of mirror matter for particle physics and astrophysics have been studied during the last decades. The reader can refer to \\cite{Okun:2006eb} for a review of the history of mirror matter and a list of most relevant papers published before 2006.  Like their ordinary counterparts, mirror baryons can form atoms, molecules and astrophysical objects such as planets, stars and globular clusters. However, even though the microphysics is the same in both sectors, the chemical content of the mirror sector should be different, because the cosmology must be different. In particular, Big Bang nucleosynthesis (BBN) requires that the mirror sector has lower temperature than the ordinary sector \\cite{Berezhiani:1995am, Berezhiani:2000gw}. This has implications for the thermodynamics of the early Universe \\cite{Berezhiani:1995am,Ciarcelluti:2008vs} and for the key cosmological epochs. Analytical results and numerical calculations of BBN \\cite{Berezhiani:1995am,Ciarcelluti:2008vm} show that the mirror sector should be helium dominated, and the abundance of heavy elements is expected to be higher than in the ordinary sector. Invisible stars made of mirror baryons are candidates for Massive Astrophysical Compact Halo Objects (MACHOs), which have been observed via microlensing events \\cite{Blinnikov:1996fm,Foot:1999hm,Mohapatra:1999ih}. Mirror stars contain more helium and less hydrogen, and therefore have different properties than ordinary stars \\cite{Berezhiani:2005vv}. The consequences of mirror matter on primordial structure formation, the cosmic microwave background and the large-scale-structure distribution of matter have been studied \\cite{Ignatiev:2003js,Berezhiani:2003wj, Ciarcelluti:2003wm,Ciarcelluti:2004ij,Ciarcelluti:2004ik,Ciarcelluti:2004ip}. These studies provide stringent bounds on the mirror sector and prove that it is a viable candidate for DM. In addition, mirror DM is one of the few potential explanations for the recent DAMA/LIBRA annual modulation signal \\cite{Bernabei:2008yi,Foot:2008nw,Ciarcelluti:2008qk}.  The properties of NS are intimately connected to the equation of state (EoS) of matter at densities well beyond nuclear saturation density, $n_0\\sim 0.16$~fm$^{-3}$. NS therefore are natural laboratories for the exploration of baryonic matter under extreme conditions, complementary to those created in terrestrial experiments with atomic nuclei and heavy-ion collisions. A stiff EoS at high density is needed to explain NS with high mass \\cite{Champion06062008} and large $R_\\infty$ \\cite{Trumper:2003we}. Heavy-ion collision data for kaon production and elliptic flow provide an upper limit on the stiffness of the EoS \\cite{Fuchs:2005zg, Danielewicz:2002pu}. By combining these constraints it is possible to test and exclude certain models of high-density EoS \\cite{Klahn:2006ir,Klahn:2006iw}. In this context it is important to be aware of the potential effects of DM on the properties of NS. A long-debated question in compact-star physics is whether quark matter exist in the core of NS and whether there are unambiguous observables associated with that. Other exotic states of matter\\eg meson condensates and hyperons could also exist in the core. For recent reviews, see \\cite{Weber:2007ch,Page:2006ud,Weber:2004kj}. For an example model of hybrid stars (stars with a quark-matter core enclosed in a nuclear-matter shell) that are consistent with present observational and experimental constraints, see \\cite{Klahn:2006iw}. A typical consequence of exotica is that the maximum NS mass and the radii of NS with typical masses, $M \\sim 1.35$~M$_\\odot$, becomes lower. Constraints on the mass and mass-radius relation of compact stars are therefore used as indicators for the presence/absence of exotica in compact stars. As an example, see the claim in \\cite{Ozel:2006bv} and the counter-examples provided in \\cite{Alford:2006vz}. In this paper we show that if mirror matter (or, in principle, some other form of stable self-interacting DM) accumulates in NS, the equilibrium sequence could be  significantly modified and the effect is similar to that of traditional exotic phases of matter. The NS equilibrium sequence is directly related to the ground-state equation of state and it is therefore commonly assumed to be unique. We show that if DM accumulates in NS this is not necessarily the case. See also \\cite{Khlopov:1989fj}, where potential effects of mirror matter on NS are discussed qualitatively. ",
        "conclusions": "If DM is made of mirror baryons, they are present in all gravitational structures. In particular, they should exist in stars at the final stages of stellar evolution. In this paper we show that this could have significant effects on the formation and structure of NS. Our knowledge of NS is limited and highly dependent on observations. The situation is complicated by the fact that the conditions inside NS cannot be reproduced in laboratories. Ideally, we should learn how to calculate the ground-state EoS from first principles using the SM and then test it with observations of NS, but this task is presently too complicated. An intermediate goal is to unify ``low-density'' results from nuclear-physics experiments and lattice-QCD calculations with observed properties of NS. In that context it is important to understand the potential effects of DM on NS, so that the observations are not misunderstood and false conclusions about the microphysics are made. After all, DM apparently is dominant over ordinary matter and we do not know what it is. In this paper we illustrate that the NS sequence is significantly modified in the presence of a few percent mirror DM and the effect mimics that of other exotic forms of matter\\eg quark matter. An observation of an extraordinary compact NS is therefore not sufficient to conclude that there must be some form of exotic phase of ordinary matter in the star, since it can be explained also by the presence of mirror matter.  We find that the NS equilibrium sequence is a function of the mirror DM content\\ie in contrast to the equilibrium sequence of ordinary NS it is not a one-parameter sequence. The maximum mass of the sequence and the radii of stars with typically observed masses, $M \\sim 1.35$~M$_\\odot$, decrease with increasing mirror baryon number. The {\\it non-uniqueness} of the equilibrium sequence is the main result of this work, because it is not clear that it can be explained with traditional physics, nor with self-annihilating DM candidates such as supersymmetric particles. New measurements of NS masses and mass-radius relations are frequently made, thanks to the high resolution and sensitivity of modern observatories. The uniqueness property of the NS equilibrium sequence could therefore be put to the test in the near future. If it turns out that NS masses and radii cannot be explained with a one-parameter sequence (after effects of rotation are accounted for), one could search for a correlation between the orbits of the observed NS and the expected density of DM along those orbits. The accumulation of mirror matter in stars could have a significant effect also on the critical mass for core collapse. This is a potential formation mechanism for NS with low masses and extraordinary small size. These results can be used in future studies to test models of the mirror Universe. \\newline \\vspace{0.2cm}"
    },
    "0809/0809.0495_arXiv.txt": {
        "abstract": "We present high-quality J, H and $K_s$ photometry of four Small Magellanic Cloud stellar clusters with intermediate ages in the 1--7 Gyr range (namely NGC~339, 361, 416 and 419)  . We obtained deep Color-Magnitude Diagrams to study the evolved sequences and providing a detailed census of the Red Giant Branch (RGB), Asymptotic Giant Branch (AGB) and Carbon star populations in each cluster and their contribution to the total cluster light. We find that in the $\\sim$5-7 Gyr old clusters AGB stars account for $\\sim$6\\% of the total light in $K_s$-band, Carbon stars are lacking and RGB stars account for $\\sim$45\\% of the total bolometric luminosity. These empirical findings are in good agreement with the theoretical predictions. Finally, we derived photometric metallicities computed by using the properties of the RGB and finding an iron content of [Fe/H]=~-1.18, -1.08, -0.99 and -0.96 dex for  NGC~339, 361, 416 and 419 respectively. ",
        "introduction": "\\label{intro}   The history and the evolution of the Large and Small Magellanic Clouds (LMC and SMC, respectively) are intimately related to the gravitational interactions between the two Clouds and the Milky Way. In particular, the main episodes of star formation enhancement in the SMC are triggered by the near, perigalactic passages of the LMC and of the Milky Way \\citep{hz04, zh04}. The epoch and rate of the main star formation episodes in the Magellanic Clouds (MCs) are still matter of debate and several scenarios have been proposed. For the SMC, while \\citet{hz04} suggest 3 episodes occurred 400 Myr, 3 Gyr and 9 Gyr ago, \\citet{dolphin} favor a more continuous star formation in the halo with a dominant episode 5--8 Gyr ago. \\citet{raf} argue that the cluster age distribution shows a few peaks, but no significant {\\sl gaps} as in the LMC \\citep[see also][]{chiosi06}.  MCs host a globular cluster system which includes objects with different age and metallicity, thus representing a formidable probe of the various stellar populations in the MCs as well as ideal templates for the study of stellar evolution and population synthesis. In this respect, the near-infrared (IR) spectral range is particularly suitable to sample the evolved stellar sequences, whose giant stars are characterized by low surface gravities and effective temperatures.   In our previous papers \\citep[][hereafter Paper I and Paper II, respectively]{f04,m06} the near-IR Color-Magnitude Diagrams (CMDs) of 19 young-intermediate age LMC clusters (from $\\sim$80 Myr to $\\sim$3 Gyr) have been analysed, providing a quantitative estimate of the AGB and RGB contributions to the total light, as a function of the cluster age. The AGB contribution to the total luminosity starts to be significant at $\\sim$200 Myr, with a maximum ($\\sim$80\\%) at $\\sim$500-600 Gyr. At this same epoch the RGB phase transition occurs and for ages older than 1 Gyr the RGB itself becomes fully developed, while the contribution of AGB is progressively reduced.  The present paper reports the results for four SMC clusters belonging to the intermediate-age population of the SMC. The principal aims of this work are the study of the main features of their near IR CMDs and the contribution to the total cluster luminosity of the AGB and RGB stars. This sample of SMC clusters allows to check the contribution of the AGB and RGB stars for clusters in an age range ($\\sim$5-7 Gyr) not covered by the LMC cluster system (because it corresponds to the so-called {\\sl Age-Gap}), thus providing complementary information.   The paper is organized as follows: Sect.~\\ref{obs} describes the observations and photometric analysis. Sect.~\\ref{cmd} presents the near-IR CMDs and their main features, the inferred metallicities and the integrated magnitudes. Sect.~\\ref{count} describes the procedure adopted to estimate the completeness correction and the field decontamination. Sect.~\\ref{agb} describes the detailed census of the AGB and Carbon stars in each cluster, and their contribution to the total luminosity, while Sect.~\\ref{rgb} analyzes the RGB stellar population. ",
        "conclusions": "\\label{conc}    By using high-quality near IR photometry of four SMC stellar clusters with intermediate ages we derived new photometric metallicities. All the observed clusters show similar metallicities ([Fe/H]$\\sim$-1) with only a weak dependence with the age. Actually, the Age-Metallicity Relation for the SMC is not well known and the different metallicity indicators, namely stellar clusters, planetary nebulae, field stars, still exhibit strong discrepancies. The most recent survey of SMC field giants based on Ca II triplet by \\citet{carrera} indicates a value of [Fe/H]$\\sim$--1 dex in the age range between $\\sim$3 and $\\sim$10 Gyr, and [Fe/H]$\\sim$--0.7 dex at $\\sim$1 Gyr.  Furthermore, we investigated the contribution of the AGB and RGB evolutionary stages to the total cluster luminosity. The cluster NGC~419, with an age of $\\sim$1 Gyr, exhibits population ratios for AGB and RGB that follow the behaviour already observed in the LMC clusters for objects of similar age. The other 3 clusters, with older ages in the  $\\sim$5-7 Gyr range show a negligible ($\\sim$6\\%) luminosity contribution by the AGB, lacking  bright C-stars, and an increasing contribution by the RGB population with respect to clusters of younger ages like NGC~419 and those in the LMC. We find a general agreement between the empirical population ratios and those predicted by theoretical models at [M/H]=~--1.35 dex."
    },
    "0809/0809.4944_arXiv.txt": {
        "abstract": "We review the study of inhomogeneous perturbations about a homogeneous and isotropic background cosmology. We adopt a coordinate based approach, but give geometrical interpretations of metric perturbations in terms of the expansion, shear and curvature of constant-time hypersurfaces and the orthogonal timelike vector field. We give the gauge transformation rules for metric and matter variables at first and second order. We show how gauge invariant variables are constructed by identifying geometric or matter variables in physically-defined coordinate systems, and give the relations between many commonly used gauge-invariant variables. In particular we show how the Einstein equations or energy-momentum conservation can be used to obtain simple evolution equations at linear order, and discuss extensions to non-linear order. We present evolution equations for systems with multiple interacting fluids and scalar fields, identifying adiabatic and entropy perturbations. As an application we consider the origin of primordial curvature and isocurvature perturbations from field perturbations during inflation in the very early universe. ",
        "introduction": "\\setcounter{equation}{0}   The standard model of hot big bang cosmology is based upon the spatially homogeneous and isotropic Friedmann-Robertson-Walker (FRW) model. This successfully describes the average expansion of the universe on large scales according to Einstein's theory of general relativity, and the evolution from a hot, dense initial state dominated by radiation to the cool, low density state dominated by non-relativistic matter and, apparently, vacuum energy at the present day. This standard model is described by just a handful of numbers specifying the expansion rate, the temperature of the present microwave background radiation, the density of visible matter, dark matter and dark vacuum energy.   But an homogeneous model cannot describe the complexity of the actual distribution of matter and energy in our observed universe where stars and galaxies form clusters and superclusters of galaxies across a wide range of scales. For this we need to be able to describe spatial inhomogeneity and anisotropy. But there are few exact solutions of general relativity that incorporate spatially inhomogeneous and anisotropic matter and hence geometry. Therefore we use a perturbative approach starting from the spatially homogeneous and isotropic FRW model as a background solution with simple properties within which we can study the increasing complexity of inhomogeneous perturbations order by order.   The introduction of a homogeneous background model to describe the inhomogeneous universe leads to an ambiguity in the choice of coordinates. In the FRW model the homogeneous three-dimensional hypersurfaces provide a natural time slicing of four-dimensional spacetime. For instance, hypersurfaces of a uniform density coincide with hypersurfaces with uniform spatial curvature. However in the real, inhomogeneous universe spatial hypersurfaces of uniform density do not in general have uniform spatial curvature, and hypersurfaces of uniform curvature do not have uniform density. In general relativity there is, a priori, no preferred choice of coordinates. Choosing a set of coordinates in the inhomogeneous universe which will then be described by an FRW model plus perturbations amounts to assigning a mapping between spacetime points in the inhomogeneous universe and the homogeneous background model. The freedom in this choice is the gauge freedom, or gauge problem, in general relativistic perturbation theory. It can lead to apparently different descriptions of the same physical solution simply due to the choice of coordinates. This freedom can be a powerful tool as it allows one to work in terms of variables best suited to the problem in hand, however it is important to work with gauge-invariant definitions of the perturbations variables if the results are to be easily assimilated and cross-referenced to the work of others.   In this review we will focus on how one can construct a variety of gauge-invariant variables to deal with perturbations in different cosmological models at first order and beyond. We will emphasise the geometrical meaning of metric and matter perturbations and their gauge-invariant definitions.   Most work to date has been done only to linear order where the perturbations obey linear field equations. Even then we must consistently solve the linear evolution equations, subject to the constraint equations of general relativity which relate the matter variables to the geometry. Beyond first order, the non-linearity of Einstein's equations becomes evident, making progress much more difficult. But in limiting cases, notably the large scale limit, it is possible to extend some of the simple results of linear theory to higher orders.   Previous reviews on the topic of linear perturbation theory in cosmology include, for example, Refs.~\\cite{KS,MFB,Durrer1993,Bertschinger1995,Riotto:2002yw,Hu_review}, and the relevant chapters in Refs.~\\cite{weinberg,Peebles80,Landau:1987gn,LLBook}. For second-order perturbation theory see Refs.~\\cite{Noh2004,Bartolo:2004if,Nakamura:2004rm}. For reviews on inflation in particular see, e.g., Refs.~\\cite{Lindebook,KolbTurner,Lidsey:1995np,LRreview,LLBook,Langlois:2004de,Bassett:2005xm,Wands:2007bd} and for cosmic microwave background anisotropies see, e.g., Refs.~\\cite{Kodama:1986fg,Lewis:1999bs,cmbeasy2,Giovannini:2004rj} Finally, for reviews on perturbation theory in the context of the higher-dimensional brane-world scenario see, e.g., Refs. ~\\cite{Kodama:2000fa,Bridgman:2001mc,Riazuelo:2002mi,Maartens:2003tw,Brax:2004xh} and for perturbation theory in the context of the low-energy string effective action see, e.g., Refs.~\\cite{Lidsey:1999mc,Cartier01}.     For simplicity we work with a flat background spatial metric which is compatible with current observations. For generalisation to spatially hyperbolic or spherical FRW models see other reviews, e.g.,~Ref.\\cite{KS}. We use a prime to denote derivatives with respect to conformal time, and we use a comma to denote partial derivatives with respect to comoving spatial coordinates, i.e., \\be T_{,i} \\equiv \\frac{\\partial T}{\\partial x^i} \\,. \\ee For further definitions and notation see Appendix \\ref{defandnot_sect}. ",
        "conclusions": "Linear perturbations have become part of the standard toolbox of modern cosmology. Earlier confusion surrounding apparently different behaviour found in different coordinate bases has largely been resolved through the use of variables which have gauge-invariant definitions. In Section \\ref{gi_var_sec} we have emphasised the physical interpretation of these gauge-invariant variables.  The power spectrum of primordial perturbations revealed by the cosmic microwave background and large scale galaxy surveys is a powerful probe of inflationary models of the very early universe, and a challenge for alternative theories. Linear theory enables us to relate the primordial spectra to quantum fluctuations in the metric and matter fields at much higher energies. In the simplest single field inflation models, it is possible to equate the primordial density perturbation with the curvature perturbation, $\\zeta$ defined in \\eq{defzeta}, during inflation which remains constant on large scales for adiabatic density perturbations. More generally, if one allows for non-adiabatic perturbations then it becomes necessary to allow for variation in the curvature perturbation, even on super-Hubble scales. Nonetheless it is still possible in many cases to identify the primordial density perturbation with a perturbed expansion in the $\\delta N$ approach described in Section~\\ref{non_linear_sec} where the integrated expansion, $N$, is a function of the local field values on spatially flat slices during inflation. As one goes beyond the textbook examples, it becomes necessary to have a clear physical definition of the perturbation variables to consistently extend the background FRW equations to the inhomogeneous perturbations.  The new frontier in the study of cosmological perturbations is the study of non-linear primordial perturbations, at second-order and beyond. Many of the familiar certainties of linear perturbation theory no longer apply. We have shown in Section \\ref{gaugesec} that quantities that were automatically gauge-independent at first order (including the non-adiabatic pressure perturbation, anisotropic stress, and the tensor metric perturbation) become gauge dependent at second order. We have shown in Section \\ref{gi_var_sec} that it is possible to use the same methodology to construct gauge-invariant variables at second (and higher) order. A variable with an unambiguous physical meaning will have a gauge invariant definition.  The resulting gauge invariant definitions inevitably become more complicated than those at first order and we have only presented explicit definitions at second order for a few cases. Likewise we have not attempted to present the second order dynamical equations in a comprehensive manner as was done at first order in Section~\\ref{sect:dynamics}. Our aim has been to provide an introduction to some of the issues that arise at higher order. Early works on non-linear and second order perturbation theory include Refs.~\\cite{Tomita1967,Peebles80}, more recently see also Refs.~\\cite{Pyne:1995bs,Tomita:2005nu,Noh2004,Nakamura:2004wr,Nakamura:2004rm,Nakamura:2006rk}.    Non-linearities allow additional information to be gleaned from the primordial perturbations. Much effort is currently being devoted to the study of higher-order correlations. Non-Gaussianity in the distribution of primordial density perturbations would reveal interactions beyond the linear theory. Such interactions are minimal (suppressed by slow-roll parameters) in the simplest single field inflation model, so any detection of primordial non-Gaussianity would cause a major upheaval in our thinking about the very early universe. In principle the $\\delta N$ approach can be easily extended to higher-orders, simply by extending the Taylor series for the integrated expansion as a function of the field perturbations beyond linear order. This enables one to compute higher-order correlation functions for $\\zeta$ as shown in Section \\ref{non_linear_sec}. The challenge is then to develop transfer functions to relate the primordial $\\zeta$ to observables beyond linear order~\\cite{Bartolo:2004if,Bartolo:2005kv,Bartolo:2006cu,Bartolo:2006fj}, although large primordial non-Gaussianity (e.g., $\\fNL\\gg1$) is expected to dominate over non-linearity in the transfer functions.     The Klein-Gordon equation in closed form at second order shows that at second order scalar perturbations will also be sourced by terms quadratic in first order field perturbations \\cite{M2006,Malik:2007nd}. This and second order perturbation theory in general, provides an alternative to using the $\\delta N$ formalism in calculating the primordial non-gaussianity $\\fNL$ \\cite{Seery:2008qj}. The main advantage of perturbation theory is that it can also be extended to sub-horizon scales, whereas the $\\delta N$ formalism is only valid on super-horizon scales, and in some cases it has been shown to be numerically more efficient \\cite{Malik:2006pm}.  Typically non-linear effects are going to be small, given that scalar metric perturbations are only of order $10^{-5}$. Primordial tensor modes are smaller, and vector modes are effectively zero during scalar field driven inflation. But the additional information available from higher-order correlations and the use of optimised filters for specific forms of non-Gaussianity means that there are already impressive constraints on the degree of primordial Gaussianity.  Qualitatively new effects appear beyond linear order. The non-linearity of the field equations inevitably leads to mixing between scalar, vector and tensor modes and the existence of primordial density perturbations then inevitably generates vector and tensor modes \\cite{Mollerach:2003nq,Nakamura:2004rm,Ananda:2006af,Baumann:2007zm,Lu:2007cj}. As shown in Section \\ref{gaugesec}, if we continue studies of scalar perturbations to higher order then the distinction between the different types of perturbations becomes gauge dependent and consistent computation will require careful (gauge invariant) definition of the variables being used."
    },
    "0809/0809.3755_arXiv.txt": {
        "abstract": "Magnetic reconnection plays a crucial role in violent energy conversion occurring in the environments of high electrical conductivity, such as the solar atmosphere, magnetosphere, and fusion devices. We focus on the morphological features of the process in two different environments, the solar atmosphere and the geomagnetic tail. In addition to indirect evidence that indicates reconnection in progress or having just taken place, such as auroral manifestations in the magnetosphere and the flare loop system in the solar atmosphere,  more direct evidence of reconnection in the solar and terrestrial environments is being collected. Such evidence includes the reconnection inflow near the reconnecting current sheet, and the outflow along the sheet characterized by a sequence of plasmoids. Both turbulent and unsteady Petschek-type reconnection processes could account for the observations. We also discuss other relevant observational consequences of both mechanisms in these two settings. While on face value, these are two completely different physical environments, there emerge many commonalities, for example, an Alfv\\'{e}n speed of the same order of magnitude, a key parameter determining the reconnection rate. This comparative study is meant as a contribution to current efforts aimed at isolating similarities in processes occurring in very different contexts in the heliosphere, and even in the universe. ",
        "introduction": "Magnetic reconnection is one of the fundamental physical processes that occur in magnetized plasmas, and manifestations range from laboratory to magnetospheric plasmas, and to solar and astrophysical plasmas. Reconnection has as its main characteristic feature the conversion of magnetic field energy into other forms of energy, and transfer of plasma and magnetic flux between separate flux systems: a re-arrangement of magnetic topology. This process is associated with phenomena such as heating of the solar corona, solar eruptions, coupling between the solar wind and the terrestrial magnetosphere, substorms in the geomagnetic tail, and on a more speculative level, high-energy phenomena observed in astrophysical systems [e.g., see {\\it Priest and Forbes} 2000].  Because these phenomena occur in environments of high electrical conductivity, the energy conversion process is usually confined to a small, local region, such as the X-type neutral point, the current sheet, the quasi-separatrix layer, and so on. Among these types of magnetic neutral regions, the X-point is probably the most basic configuration from which the others could develop. For example, a current sheet could develop either through the collapse of an X-point [{\\it Dungey} 1953] or through the severe stretching of the magnetic field around an X-point [{\\it Forbes and Isenberg} 1991]; and a quasi-separatrix layer [{\\it D\\'{e}moulin et al.} 1996] might be produced by skewing the magnetic fields on both sides of an X-point.  At the  current sheet separating the shocked solar wind and the terrestrial magnetosphere (the magnetopause), magnetic reconnection can transfer solar wind momentum and energy into the magnetosphere on the dayside. Figure \\ref{fig:magnetosphere} displays a schematic of the topological changes that are initiated by magnetic reconnection between solar wind field and terrestrial field on the dayside (stage 1). Reconnected (or ``open\") field lines are then dragged by the shocked solar wind flow into the nightside magnetosphere, where they form the tail lobes (stage 2). Addition of magnetic flux from the dayside to the nightside increases the energy stored in the tail. A thin current sheet forms in the tail and at substorm onset reconnection takes places there (stage 3), heating plasma and jetting it in bursty fashion both earthward [{\\it Baumjohann et al.}, 1990; {\\it Angelopoulos et al.}, 1992] (stage 4) and tailward (stage 5) at high speeds. On the earthward side, reconnected field lines rapidly snap back toward the Earth due to magnetic tension, constituting a process known as the shrinkage of field lines [e.g., see {\\it \\v{S}vestka et al.} 1987; {\\it Lin et al.} 1995a; {\\it Forbes and Acton} 1996; {\\it Lin} 2004; {\\it Reeves et al.} 2008]. These field lines eventually return to the dayside and replace the eroded magnetic flux.  Magnetic reconnection and a general convective cycle, including frontside and nightside reconnections, in the process described above were proposed for the first time by {\\it Dungey} [1961]. Later, the substorm model was developed, particularly by {\\it McPherron et al.} [1973], {\\it Hones} [1977], and {\\it McPherron} [1991]. (See also {\\it Baker et al.} [1996] for a more recent review.) As emerges from the discussions above, the substorm is an explosive energy release and is accompanied by a reconfiguration of the near-tail magnetic field (dipolarization) and other phenomena to be discussed later. If the magnetosphere is immersed in a southward-pointing interplanetary magnetic field of long (many hours) duration, such as what typically happens during passage of interplanetary coronal mass ejections (ICMEs), then repetitive substorms [{\\it Farrugia et al.} 1993a] occur at intervals of about 2--3 hours [{\\it Huang et al.} 2003].  The formation and development of a current sheet and/or a quasi-separatrix in the solar magnetic field is usually associated with solar eruptive processes, including solar flares, eruptive prominences, and coronal mass ejections (CMEs). In fact, substorms have been called the flares in the magnetosphere; i.e., loosely speaking, due to the similarity and explosive character in the energy conversion processes, substorms may be considered as the magnetospheric counterparts of solar eruptions. As we show below, a comparison of reconnection in the solar and terrestrial contexts reveals many common features.  Table \\ref{tbl:VA} lists some important parameters for the solar atmosphere and for the magnetosphere. Those for the solar atmosphere can be found in the book by {\\it Priest} [1982], those for the low latitude dayside magnetosphere were taken from the {\\it in situ} measurements reported by {\\it Paschmann et al.} [1986], those for the middle distant tail lobe and the plasma sheet are from {\\it Slavin et al.} [1985], and those for the near Earth sheet are from {\\it Baumjohann et al.} [1989]. We notice the same orders of magnitude of large $V_{A}$ and small $\\beta$ in the solar corona and in the tail lobes of the magnetosphere, but small $V_{A}$ and larger $\\beta$ in the solar photosphere, in the dayside magnetopause, and in the central plasma sheet in the geomagnetic tail. Those in the central sheet (e.g., see {\\it Baumjohann et al.} [1989]) suggest a rapid increase of the reconnection rate when low density, high Alfv\\'{e}n speed, lobe field becomes reconnected (e.g., {\\it Hesse et al.} [1996]). However, the corresponding values for the solar counterpart of the geomagnetic tail sheet are not available due to the lack of {\\it in situ} measurements in the solar environment. On the other hand, the low $V_{A}$ and high $\\beta$ in the solar photosphere and in the dayside magnetopause have different consequences that will be discussed shortly.  So even though Table \\ref{tbl:VA} indicates some similarities of the current sheet in the magnetosphere to that in the solar corona, the levels of our understanding of the details of magnetic reconnection in the magnetosphere and in the solar corona are very disparate. In the magnetosphere, extensive {\\it in situ} measurements reveal tremendously rich information of the internal features of the current sheets and the magnetic reconnection processes at the magnetopause (e.g., {\\it Rijnbeek et al.} [1989]; {\\it Sonnerup et al.} [1981]; {\\it Phan et al.} [2000]; {\\it Paschmann} [2005]), and in the magnetotail [e.g., {\\it Murata et al.} 1995; {\\it Baker et al.} 1985, {\\it Nakamura et al.} 2006]. In the solar corona, on the other hand, directly observing the current sheet turns out to be very hard, if not impossible. Difficulties stem from two aspects: the low level of emission from the sheet plasma compared to that of other nearby bright features (see the introduction of {\\it Ko et al.} [2003] for a brief review), and the possible thinness of the sheet.  In theory, it has often been stated that the sheet is too thin to be observable since its thickness is believed to be limited by the Larmor radius of particles that is only a few tens of meters in the coronal environment [{\\it Litvinenko} 1996; {\\it Wood and Neukirch} 2005; and references therein]. Another reason that might have discouraged direct observations of the current sheet could be ascribed to the ``circuit model\" of eruptions developed by {\\it Martens and Kuin} [1989], in which the theoretical current sheet length is short. Not until the theoretical works of {\\it Forbes and Lin} [2000], {\\it Lin and Forbes} [2000], and {\\it Lin} [2002], as well as a series subsequent observational ones [{\\it Ciaravella et al.} 2002; {\\it Ko et al.} 2003; {\\it Raymond et al.} 2003; {\\it Sui and Holman} 2004; {\\it Sui et al.} 2004; {\\it Webb et al.} 2003; {\\it Lin et al.} 2005; {\\it Bemporad et al.} 2006; {\\it Lin et al.} 2007], have we started realizing that a fairly long current sheet could develop in the major eruptive process.  In this work we undertake a comparison of the reconnection processes in these two environments, putting emphasis on geomagnetic substorms, on the one hand, and solar flares, on the other. Such a comparison consists of two aspects. First, we look into the detailed internal structures of the reconnecting sheets observed in the geomagnetic environment and in the solar corona; second, we compare two theoretical mechanisms that may account for the observed features of the reconnecting current sheet. This brings the known results for magnetic reconnection obtained from the investigations of various communities together, and helps people of different research areas to better understand the work and the progress made by others. ",
        "conclusions": "In this work we have engaged in a comparative study of the reconnection process in the solar and terrestrial contexts. Properties of the magnetized plasma determines that the rate of reconnection in both environments is controlled by the local Alfv\\'{e}n speed. This further yields that the reconnection process in the low $V_{A}$ and high $\\beta$ region transports magnetic flux and energy to the region of high $V_{A}$ and low $\\beta$ gradually, and that reconnection in the high $V_{A}$ and low $\\beta$ region releases the stored magnetic energy violently. This may explain why the consequences occurring in the photosphere and in the dayside magnetopause (low $V_{A}$ and high $\\beta$) are much less significant than those occurring in the corona (flares and CMEs) and in the magnetotail (substorms) (high $V_{A}$ and low $\\beta$).  Of course, because of the huge difference in scale length and magnetic field strength, the absolute amount of energy conversion and the absolute rate of reconnection (usually measured by the strength of reconnection electric field along the X-line) in the solar and terrestrial contexts could greatly differ from one another. For example, a typical solar eruption involves energy release of up to 10$^{32}$ ergs (e.g., see {\\it Priest} [1982]) and the absolute reconnection rate typical varies from 1 V/cm to 10 V/cm (e.g., see {\\it Poletto and Kopp} [1986] and {\\it Qiu et al.} [2004]), and a typical substorm releases energy of only $10^{21}\\sim 10^{22}$ ergs [{\\it Baker et al.} 1997] and the absolute reconnection rate of 1 $\\sim$ 10 mV/m (e.g., see {\\it Blanchard et al.} [1997] and {\\it Vaivads et al.} [2006]).  On the other hand, the rate of reconnection is also often measured by the Alfv\\'{e}n Mach number $M_{A}$, which is the reconnection inflow speed $v_{i}$ compared to the local Alfv\\'{e}n speed $V_{A}$ near the reconnection region. So $M_{A}$ is known as the relative reconnection rate as well (e.g., see discussions of {\\it Lin and Forbes} [2000] and {\\it Priest and Forbes} [2000]). In this sense, $M_{A}$ is within the same order of magnitudes between 0.01 and 0.1 for the reconnection processes in both solar and magnetospheric environments (cf. {\\it Ko et al.} [2003], {\\it Lin et al.} [2005] and {\\it Vaivads et al.} [2006]) although the absolute rates of reconnection in the two environments are vastly different. This result should help us study and well understand the similar energy conversion processes taking place on other objects in the universe.  Another emphasis related to reconnection has been the formation of plasmoids/plasma blobs in both environments. The magnetic reconnection process manifests almost the same features in various different magnetized plasma environments, including the solar corona and the magnetosphere. Usually flows are observed in two opposite directions along the sheet in both the magnetotail and the CME/flare sheet, and only the antisunward flow is observed in the helmet streamer. This is probably because the reconnection in the helmet streamer is weaker and the sunward reconnection outflow is easily stopped by the closed loops below. Plasmoids or plasma blobs flowing along the reconnecting current sheet are the ubiquitous features observed in these environments although mechanisms for the formation the current sheet vary from case to case.  These plasmoids are usually ascribed to multiple X-line reconnection as a result of the tearing mode instability developing in a long thin current sheet (e.g., see {\\it Furth et al.} [1963] and {\\it Priest} [1985] for the review; and {\\it Drake et al.} [2006] and {\\it Loureiro et al.} [2007] for the most recent works on this topic). Reconnection occurs at each of these X-lines, but one of them eventually develops to become the dominant one at which the fastest reconnection takes place, and its high speed flow expels both plasmoids and other X-lines to move in two opposite directions. The mechanism determining the location of this dominant X-line is an open question, but it should depend on the parameters of the plasma and magnetic field around.  An alternative process that may account for the plasmoid feature is time-dependent Petschek-type reconnection extensively studied by Semenov and his co-authors. In this framework, diffusion initially commences somewhere in a pre-existing current sheet where a finite electrical resistivity appears. The impact caused by the diffusion is not confined to the diffusion region but propagates as a disturbance at the local Alfv\\'{e}n speed to the medium at large. Determined by the time profile of the diffusion, the disturbance propagates along the current sheet with two pairs of the slow mode shocks surrounding the reconnected plasma and magnetic field in the outflow region. Departing from the diffusion region, the two pairs of the slow shock and the associated plasma flows propagate in opposite direction. In this case, multiple X-lines and blobs may appear alternatively along the sheet, but the dissipation just occurs only at the slow shocks and the X-line where the initial diffusion commences. Such a treatment applied by Semenov and his co-workers is based on considerations of mathematical simplicity in order to find analytic solutions to the problem.  This scenario is quite like that of the multiple X-line reconnection due to the tearing mode in the stage when one X-line eventually becomes dominant and two reconnection outflows in opposite direction emanate from it. One difference is that the dissipation occurs everywhere in the sheet undergoing tearing, and the dominant X-line develops in a self-consistent fashion from many X-lines during the process (e.g., see {\\it Fu and Lee} [1985]). But the dominant X-line in the Petschek-type reconnection is fixed artificially from the very beginning. Another difference comes from our understanding of the dissipation mechanisms in the process. In the tearing mode version, the dissipation occurs in the whole current sheet as a result of the plasma turbulence caused by the instability. In the Petschek version, on the other hand, in addition to the diffusion region, the slow shocks play the main role in the energy conversion. These two approaches should be physically equivalent. Comparing the arguments and discussions of {\\it Forbes and Malherbe} [1991] and {\\it Priest and Forbes} [2000, pp. 391--393] on the same results from the numerical experiment of reconnection in the two-ribbon flare process seems to be suggestive of such an equivalence. But rigorous studies of this equivalence are not yet available due to mathematical difficulties.  In addition to the interior structure of the reconnecting current sheets, it is also necessary to pay attention to the environment outside the sheet and its possible impacts on the energy conversion process. In discussions of magnetotail versus corona reconnection, one topic that has come up repeatedly is whether there is any evidence for the strong increases in Alfv\\'{e}n speed at the Sun as one moves normal to the reconnecting sheet. The gradients are large in the Earth's tail with Alfv\\'{e}n speeds of $< 100$ km s$^{-1}$ in the central plasma sheet where reconnection commences, but increasing to $> 1000$ km s$^{-1}$ over 100 to 1000 km distant in the outer plasma sheet and lobes. This strong gradient is believed to be the source of the ``explosive\" appearance of magnetotail reconnection as first the low Alfv\\'{e}n speed flux tubes in the central plasma sheet reconnect followed later by the high Alfv\\'{e}n speed lobe flux tubes (e.g., see {\\it Hesse et al.} [1996]).  In principle, fast reconnection in the solar corona also occurs during the explosive event, like flares, in the current sheet where a large gradient of Alfv\\'{e}n speed exists in the direction normal to the sheet. This is represented indirectly by the dramatic change in plasma $\\beta$ from well below unity outside the sheet to around unity inside the sheet (e.g., see {\\it Ko et al.} [2003]), and by the significant difference between the inflow and the outflow speeds of the plasma measured around the sheet developed in an eruption (e.g., see {\\it Lin et al.} [2005]). Unlike the geomagnetic tail sheet in the ``explosive\" stage, however, the coronal sheet in the ``explosive\" phase does not exist beforehand, but forms and develops during the eruptive process. So the strong gradient of the Alfv\\'{e}n speed across the CME/flare sheet is not the very source of the ``explosive\" appearance of reconnection. Instead the CME/flare current sheet forms as the closed magnetic field is severely stretched by the eruption, separating magnetic field lines of opposite polarity, and the temporal vacuum around the sheet quickly brings these field lines and plasma into the sheet driving reconnection to occur (e.g., see {\\it Lin et al.} [2003]). In this process it is not necessarily the magnetic field of the low Alfv\\'{e}n speed that reconnects before that of high Alfv\\'{e}n speed. But it is true that the Alfv\\'{e}n speed inside the sheet is low compared to that outside (e.g., see {\\it Forbes and Malherbe} [1991])."
    },
    "0809/0809.1750_arXiv.txt": {
        "abstract": "{We present the results from a study of the X--ray variability and the near-IR to X--ray spectral energy distribution of four low-luminosity, Seyfert 1 galaxies.} {We compared their variability amplitude and broad band spectrum with those of more luminous AGN in order to investigate whether accretion in low-luminosity AGN operates as in their luminous counterparts.} {We used archival \\xmm\\ and, in two cases, \\asca\\ data to estimate their X--ray variability amplitude and determine their X--ray spectral shape and luminosity. We also used archival \\hst\\ data to measure their optical nuclear luminosity, and near-IR measurements from the literature, in order to construct their near-IR to X--ray spectra.} {The X--ray variability amplitude of the four Seyferts is what one would expect, given their black hole masses. Their near-IR to X--ray spectrum has the same shape as the spectrum of quasars which are $10^{2}-10^5$ times more luminous.} {The objects in our sample are optically classified as Seyfert 1-1.5. This implies that they  host a relatively unobscured AGN-like nucleus. They are also of low luminosity and accrete at a low rate. They are therefore good candidates to detect radiation from an inefficient accretion process. However, our results suggest that they are similar to AGN which are $10^2 - 10^5$ times more luminous. The combination of  a ``radiative efficient accretion disc plus an X--ray producing hot corona\" may persist at low accretion rates as well.} ",
        "introduction": "\\smallskip  The current paradigm for active galactic nuclei (AGN) comprises a supermassive black hole (BH) which is accreting material through a disc. Accretion in  a geometrically thin, optically thick disc (Shakura \\& Sunyaev 1973) is the favorite mechanism to achieve the high accretion efficiency (of the order of $\\sim 0.1$ or so) which is necessary to explain the large power emitted by AGN. Theoretical modeling of low accretion rate systems suggests that, when the accretion rate is below a few percent of the Eddington limit, L$_{\\rm Edd}$, AGN may switch to a different accretion mode characterized by low radiative efficiency. Most ``radiatively inefficient accretion flow\" models (RIAFs; see, e.g., Yuan (2007) and references therein) suggest that the kinetic  energy associated with the gas is either advected with the matter into the BH or redirected into an outflow.  Many of the nearby AGN are intrinsically faint (exhibiting bolometric luminosities less than $10^{43-44}$ ergs/s or so) while the central BH in a large fraction of them is as large as $10^7-10^{8}$ M$_{\\odot}$ (see, e.g., Panessa \\etal\\ 2006). Consequently, most of these nuclei accrete at a rate lower than a few percent of L$_{\\rm Edd}$. It is then possible that  accretion in these objects operates in a radiative  inefficient mode.  For example, according to Merloni et. al.  (2003), the accretion mode should change below  L$_{\\rm X}$/L$_{\\rm Edd} \\sim 10^{-3}$ (where L$_{\\rm X}$ and L$_{\\rm Edd}$ are the $2-10$ keV and  Eddington luminosity, respectively), while Chiaberge, Capetti \\& Macchetto (2005) suggest that RIAFs may operate in active nuclei with  L$_{\\rm O}$/L$_{\\rm Edd} \\le 10^{-4}$ (where L$_{\\rm O}$ is the optical luminosity in the $R-$band). However, conclusive evidence for the presence of RIAFs in low luminosity AGNs (LLAGNs) is still lacking.  One of the most frequently applied tests to investigate whether a RIAF operates in LLAGNs is to compare their nuclear spectral energy distribution (SED) in the near-infrared (IR) to ultraviolet (UV) part of the spectrum with the average SED of powerful quasars.  At these frequencies  different accretion disc models are expected to show large differences in spectral shape. RIAFs should lack both the ``big blue bump\" and the IR (reprocessed) bump, which instead characterize the emission from  an optically thick, geometrically thin accretion disc and the surrounding heated dust. Application of this test to various samples of LLAGNs has yielded contradictory results. While Maoz (2007) finds that the broad band SEDs of AGNs with luminosities as low as $\\sim 10^{40}$ erg/s are quite similar to the SED of more luminous  AGN, Ho (1999)  found that the low-luminosity AGN SEDs have a weak or absent blue bump and are ``radio loud\". Quataert, di Matteo and Narayan (1999) find that the optical/UV spectrum of M81 and NGC~4579 decrease with increasing frequency, in contrast to the ``canonical\" quasar spectrum. Ptak \\etal\\ (2004) find that NGC~3998 (a ``type 1\" LINER galaxy) is lacking the ``big blue bump\", while RIAF models can fit its UV to X--ray SED reasonably well. Chiaberge \\etal\\ (2006) also find that the SED of a low luminosity Seyfert 2 galaxy, namely NGC~4565, is different from that of luminous quasars.  Another powerful test to identify the presence of RIAFs in low luminosity AGNs is to compare their  variability properties with the properties of more luminous AGNs. Ptak \\etal\\ (1998) found that LLAGNs tend to show little or no significant short-term variability, despite their small luminosity. This is opposite to what is observed in powerful AGN, where the variability amplitude increases with decreasing luminosity. This difference has been interpreted as an indication that the size of the  X-ray producing region in LLAGNs is significantly  larger than the size of the X-ray source in ``normal\" Seyferts. This is in support of the hypothesis that advection-dominated accretion operates in LLAGNs, as in this case one can explain the larger size of the nuclear source.  Our aim is to compare the X--ray variability amplitude and the optical/X--rays SED of low luminosity AGN and powerful quasars. We used the Palomar optical spectroscopic survey of nearby galaxies (Ho, Filippenko \\& Sargent 1995) to find suitable objects for our study. This sample has the advantage of having uniform and high-quality data that allowed classification of the galactic nuclei to be determined with well defined and objective criteria (Ho, Filippenko \\& Sargent 1997). We only considered  objects which are classified as Seyfert 1 - Seyfert 1.5, because in this case we can be certain of the AGN-like nature of their nuclear activity, and that the central source is unobscured. There are 11 such objects in the Palomar sample (Ho \\etal\\ 1997).  We chose four from these sources, namely  NGC~4235, NGC~4639, NGC~5033 and NGC~5273, because there exist estimates of their central black hole mass (hence we can estimate their L$_{\\rm Edd}$) and  their nuclear luminosity is small: their L$_{\\rm X}$/L$_{\\rm Edd}$ ratio is smaller than $10^{-3}$. Therefore, it is possible that accretion in these objects operates in a radiatively inefficient mode.  First, we used archival \\xmm\\ and \\asca\\ data to measure their X-ray variability amplitude and we compared it with the variability amplitude of more luminous Seyferts. Then, by i) fitting  models to the \\xmm\\ spectra of the four sources, ii) measuring their nuclear optical luminosity using archival \\hst\\ data, and iii) using near-IR  luminosity measurements from the literature, we constructed their near-IR to X--ray SEDs and we compared them with the mean, broad band SED of luminous PG quasars. Our final result that the properties of the four objects in our sample are similar to those of more luminous AGNs, should have interesting implications regarding the nature of the accretion mode in low luminosity active galaxies. ",
        "conclusions": "We have used archival \\xmm and \\asca data in order to measure the X--ray variability amplitude of four low luminosity Seyfert 1-1.5 galaxies, and compare it with the amplitude of brighter AGN. We have used the same  \\xmm data to determine accurately their nuclear  X--ray luminosity and spectral shape, and HST data to measure the optical/UV flux of their active nucleus. We then combined the \\xmm spectra, the HST measurements, and near-IR data from literature to construct  their  broad band, near-IR to X--ray SED and compare it  with the mean SED of radio-quiet quasars. Our main aim is to  investigate whether accretion in these objects operates in a mode which is different from the accretion mode in their more luminous counterparts.  Such comparisons are frequently used to investigate the nature of the nuclear source in low luminosity AGNs.  Most of the similar work in the past has focused mainly  on the comparison between  the properties of galaxies with low-ionization nuclear emission-line regions (LINERs; Heckman 1980) and luminous AGN. However, the nature of LINERS and their relation to AGN has been debated for several years. For this reason, we decided to study objects with a firm optical classification as a Seyfert 1 - 1.5 galaxy. In this way, we are certain that they host an accretion driven, AGN-like, UV to X--ray nuclear source. In addition, the clear presence of broad lines in their optical spectra implies that we can view their central source directly, and hence determine their nuclear SED reliably.  Furthermore, using the X--ray luminosities listed in Table~3, and the BH mass estimates listed in Table~2, we find that the (L$_{\\rm X}$/L$_{\\rm Edd}$) ratio of NGC~5273 is $6\\times 10^{-4}$. The value of the same ratio in the case of NGC~4235, NGC~4639 and NGC~5033 is even smaller: $7.7, 1.8$ and $3.3 \\times 10^{-5}$, respectively. This is much smaller than the same ratio of powerful quasars, and even  20 to 80 times smaller than the ratio of other Seyfert 1 galaxies in the Palomar survey (like for example NGC~4051, NGC~4151 and NGC~5548). It is also below the limit of $10^{-3}$ where Merloni \\etal\\ (2003) suggest that  the accretion mode may change from a thin, optically thick disc to that of a RIAF. We also converted the measured F547 fluxes of NGC~4639, NGC~5033 and NGC~5273, and the measured F606W flux of NGC~4235, to flux at 7000 \\AA, assuming an optical spectral index of 1, and we estimated the monochromatic luminosity at this wavelength. Chiaberge \\etal\\ (2005) suggest that if L$_{\\rm O}$/L$_{\\rm Edd}\\le 10^{-4}$, then a RIAF may operate in an AGN. In the case of NGC~4235, NGC~4639 and NGC~5033, this ratio is $6-60 \\times 10^{-6}$, while in the case of NGC~5273,  L$_{\\rm O}$/L$_{\\rm Edd}\\sim1.5\\times 10^{-4}$.  In summary, the four objects of our study  i) should host a relatively unobscured AGN-like nucleus, ii)  are of low luminosity, and iii) accrete on a low rate as well. Hence they are good candidates to detect radiation from RIAF-like processes.  We detect significant X--ray variations in the \\xmm light curves of NGC~5033 and NGC~5273, and in the \\asca light curve of NGC~4235. The  variability amplitude is 5--15\\% of the light curve's mean,  and is  exactly of the ``right\" order, given their BH mass. In agreement with Ptak \\etal\\ (1998), we find that these objects do not follow the ``variability amplitude -- luminosity\" relation of other more luminous AGN. However, this result does  not necessarily imply a larger characteristic size for the X--ray producing region than is the case in ``normal\" AGN. On the contrary, if the size is determined by BH mass, as expected, then our results suggest that, when scaled to BH mass,  the size of their X--ray sources is similar to the size of the nuclei in more luminous AGN. Their  ``misplacement\" in the ``amplitude -- luminosity\" plane is not due to their variability amplitude being too low but due  to their luminosity being low,  most probably because of their low accretion rate. Their X--ray spectral shape is also consistent with the shape of more luminous AGN. Even the best fitting spectral slope in the case of NGC~5273  is not exceptionally unusual; NGC~4151, a well studied, luminous Seyfert galaxy, also shows a very flat X--ray spectrum of $\\Gamma\\sim 1.5$ (see e.g. de Rosa \\etal\\ 2007). We therefore conclude that, from the X--ray  point of view, the central engine in these low luminosity Seyferts operates in the same way as in the more luminous AGN.  The near-IR to X--ray SEDs of the objects in our sample are rather sparse, and their qualitative comparison with either the average SED of radio-quiet quasars or the SED of other low-luminosity objects which appear to host a RIAF, is not conclusive. If we take into account the fact that we used measurements taken several years apart, and the possibility of large amplitude variations, at least in the X--ray band, then it is reasonable to conclude that the SEDs of the AGN in our sample are similar to the average SED of quasars. On the other hand, the SED of at least NGC~4235 and NGC~4639 are almost identical to the SED of NGC~3998, which is consistent with a model of a truncated disc and a RIAF at smaller radii  (see also the case of NGC~4579 and M81 whose SEDs are consistent with a similar model; Quataert \\etal\\ 1999).  The situation becomes clearer when we compare the optical luminosity, L$_{\\rm V}$, with the X--ray luminosity, L$_{\\rm X}$, in the 4 Seyferts of our sample (and a few other LLAGNs) and in quasars (i.e. AGN which are 10$^2$ to $10^5$ times more luminous). We find that the quasar optical luminosity increases almost proportionally with  L$_{\\rm X}$. The extrapolation of the quasar ``L$_{\\rm V}$ vs L$_{\\rm X}$\" best fit line to lower luminosities agrees very well with the LLAGN data. It is then reasonable to assume that, the observed ``continuity\" in the  (L$_{\\rm V}$, L$_{\\rm X}$) distribution of quasars and LLAGNs, also suggests that the optical/X--ray SEDs in AGN are similar, hence, there is no accretion mode  ``phase transition\",  over almost seven orders of magnitude.  Our results are in agreement with the results of Panessa \\etal\\ (2006) who study the X--ray and optical emission line correlation in a sample of AGN with a wide range of Eddington ratios, and conclude that low-luminosity Seyferts are simply scaled-down versions of luminous quasars. Similar results were obtained by  Maoz (2007) who studied the SEDs of 13 low luminosity LINERs, and found that radiative efficient discs may operate even in these objects, which are characterized by very low accretion rates.  Although we have tried to use the best available data to determine the SED for the four Seyferts, the fact that the measurements we use  in the various bands are separated by years, introduces some uncertainty. Significant improvement can be achieved if we manage to determine the average nuclear flux of each object in various bands. This can be achieved only if we observe individual objects many times over a time scale of a few years at least, which is not an easy task. However, the ``variability amplitude test\" results  are not affected by these limitations. Even if we consider only them, the main result of our work is that accretion operates in the ``normal\" way even in objects where the ratio  L$_{\\rm X}$/L$_{\\rm Edd}$ is as low as a few$\\times 10^{-5}$.  Perhaps, it is not just the accretion rate that determines the accretion mode in AGN. Black hole spin and/or mass may also play a role as to how accretion operates in active galaxies. We plan to investigate this issue by studying the variability properties of a large sample of low luminosity AGNs, with a wide range of BH masses. A dichotomy in the ``variability amplitude -- BH mass\" plane, could provide us with important clues as which are the main parameters which dictate the ``switch\" in the AGN accretion mode.  \\vskip 0.4cm"
    },
    "0809/0809.3339_arXiv.txt": {
        "abstract": "A report is presented on Suzaku  observations of the ultra-luminous X-ray source X-1 in the starburst galaxy M82, made  three time in 2005 October for an exposure of $\\sim 30$ ks each. The XIS signals from a  region of radius $3'$ around the nucleus defined a 2--10 keV flux of $2.1 \\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ attributable to point sources. The 3.2--10 keV spectrum was slightly more convex than a power-law with a photon index of 1.7. In all  observations, the HXD also detected signals from M82 up to $\\sim 20$ keV, at a 12--20 keV flux of $4.4 \\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$. The HXD spectrum was  steeper than that of the XIS. The XIS and HXD spectra can be jointly reproduced by a cutoff power-law model, or  similar curved models. Of the detected wide-band signals, 1/3 to 2/3 are attributable to X-1, while the remainder to other discrete sources in  M82. Regardless of the modeling of these contaminants, the spectrum attributable to X-1 is more curved than a  power-law, with  a bolometric luminosity of $ (1.5-3) \\times 10^{40}$ erg s$^{-1}$. These results are interpreted as Comptonized emission from a black hole of $100-200$ solar masses, radiating roughly at the Eddington luminosity. ",
        "introduction": "\\label{sec:intro}  Ultra-luminous compact X-ray sources (ULXs; e.g., \\cite{makishima2000}), found in many nearby galaxies, are often considered as candidate  intermediate-mass black holes (BHs), because their typical luminosity reaching $3 \\times 10^{39}$ to $10^{41}$  erg s$^{-1}$, calculated assuming isotropic radiation, largely exceeds the Eddington limit for known stellar-mass BHs. Since the first discovery by the {\\rm Einstein} satellite  \\citep{fabbiano1989}, ULXs have been studied extensively in soft X-rays with {\\rm ROSAT} \\citep{colbert1999,colbert2002,liu2005, liu2006}, and then in energies up to  $\\sim 10$ keV with {\\rm ASCA} \\citep{makishima2000,mizuno2001}, {\\rm Chandra} (e.g. \\cite{swartz2004}), {\\rm XMM-Newton} (e.g. \\cite{feng+kaaret2005, stobbart2006}), and {\\rm Suzaku} \\citep{mizuno2007,isobe2008}. Searches for their counterparts have also been carried out in the optical (e.g. \\cite{pakull2003, ptak2006}) and radio (e.g. \\cite{sanchez-sutil2006}) wavelengths. In contrast, we have little information on ULXs in hard X-rays above 10 keV, where the established mass-accreting BHs, including both stellar-mass ones and active galactic nuclei (AGNs), are known to emit a considerable fraction of their radiative luminosity.  Among a number of ULXs known so far, so-called source X-1 in the starburst galaxy M82 is the most extreme object, as the  luminosity  reached $10^{41}$ erg s$^{-1}$. Although the host galaxy M82 has been observed repeatedly with the past X-ray observatories due to its nearby location ($\\sim 3.6$ Mpc; \\cite{freedman2004}) and its active star formation, the bright diffuse thermal X-rays from this galaxy hampered the detection of X-1 till the {\\rm ASCA} era. The presence of this object, X-1,  was first suspected by \\citet{tsuru1997}, based on differences of fluxes measured with EXSOAT, Ginga, and ASCA. Subsequently, the object was confirmed as a variable hard source through  multiple {\\rm ASCA} observations, and interpreted at first as an active galactic nucleus veiled by the thermal plasma emission \\citep{matsumoto1999,ptak1999}. However, thanks to the superb angular resolution of  {\\rm Chandra}, it was later revealed to be a non-AGN source located $\\sim 170$ pc off  the dynamical center of the galaxy \\citep{matsumoto2001,kaaret2001,tsuru2004}.  To explain  the highest luminosity ($1\\times 10^{41}$ erg s$^{-1}$ in 3--20 keV; \\cite{rephaeli2002}) recorded from M82 X-1, a mass of at least 700 $M_{\\odot}$ is necessary assuming isotropic radiation with a sub-Eddington luminosity. Moreover, the 54 mHz X-ray quasi-periodic oscillation (QPO), detected with {\\rm XMM-Newton} and {\\rm RXTE} \\citep{strohmayer2003,fiorito2004,mucciarelli2006,dewangan2006,kaaret2006}, suggests that the source, presumably a BH, has a mass of 100--300 $M_{\\odot}$, assuming the QPO frequency to scale inversely with the BH mass. These  results, together with a recent  study of variability pattern \\citep{casella2008}, make M82 X-1 the most promising candidate for an intermediate-mass BH. It is also suggested to have a 62 day periodicity \\citep{kaaret2006}.  Spectral studies of M82 X-1 have been conducted extensively. With {\\rm ASCA} and {\\rm XMM-Newton}, however, such attempts were limited to relatively narrow bands, e.g., 2--10 keV, because of the surrounding bright thermal X-rays. As a result, the observed spectra were reproduced by a variety of spectral models, including a power-law (PL) of photon index $\\Gamma = 1.7-2.5$ \\citep{matsumoto1999,dewangan2006}, a thermal plasma emission model with a temperature 4--11 keV \\citep{matsumoto1999}, an unsaturated Comptonization model \\citep{fiorito2004,agrawal2006}, and  slim-disk emission \\citep{okajima2006}. Since these modelings imply vastly different physical conditions of the emission region, the nature of M82 X-1 has remained ambiguous. For example, the slim-disk modeling by \\citet{okajima2006} leads to a BH mass of $\\sim 30~M_\\odot$, which is significantly lower than is implied by the other arguments. Another complication with  X-1 is that it is heavily confused with other point sources in the core region of M82.  These confusion problems have been  resolved by {\\rm Chandra} with its high angular resolution. In fact, \\citet{kaaret2006} found that the 0.3--8 keV band  spectrum of X-1 can be represented by a PL model of a photon index $\\Gamma \\sim 1.7$. Neverthelss, the information is still insufficient to arrive at a clear scenario as to the nature of X-1.   Evidently, the detection of X-1 in harder energies is of crucial importance. The Large Area Counter onboard {\\rm Ginga}, sensitive over the 2--30 keV energy band, detected signals up to $\\sim$20 keV from the M82 galaxy, and yielded a spectrum which is better described by a thermal bremsstrahlung model of temperature $\\sim 5.7$ keV  than by a single PL model \\citep{tsuru1992}. The emission,  interpreted at that time as arising from a thin hot plasma, is now more likely to be attributable primarily to X-1. \\citet{cappi1999}  successfully described a  3--30 keV BeppoSAX spectrum of  M82 by a thermal plasma model of temperature $\\sim 8.2$ keV. Furthermore, {\\rm RXTE} observations  of M82 on 30 occasions yielded 3--50 keV spectra \\citep{rephaeli2002}; the spectra acquired in brighter and fainter phases were both reproduced by thermal plasma models, with a temperature  of $\\sim 6.6$ and $\\sim 7.4$ keV, respectively. These existing results in harder energies commonly indicate mildly curving broad-band spectra, rather than a PL-like one. However, the fields-of-view (FOVs) of these non-imaging instruments also contained the spiral galaxy M81, hosting a low-luminosity AGN \\citep{ishisaki1996,pellegrini2000,laparola2004}, located $37'$ away from M82. Therefore, possible contamination to the observed signal was not excluded.  The Hard X-ray Detector (HXD; \\cite{takahashi2007,kokubun2007}) onboard {\\rm Suzaku} \\citep{mitsuda2007}, with its unprecedented sensitivity in a broad (10--600 keV) energy band and a tightly collimated FOV of $34' \\times 34'$ (FWHM), is expected to provide improved hard X-ray information on M82 X-1. Here, we report on the {\\rm Suzaku} detection of M82 X-1 up to $\\sim 20$ keV, and interpret the results in terms of an intermediate-mass BH. ",
        "conclusions": "Based on our discussion, we conclude that M82 X-1 was in the PL state during the present Suzaku observation, and was radiating roughly close to, or slightly below, the Eddington luminosity. This leads to a mass estimate of  at least $100-200~M_\\odot$, which is in a good agreement with an independent estimate based on the QPO frequency (section~\\ref{sec:intro}). We therefore consider M82 X-1 as a genuine intermediate-mass BH, in the sense that it is significantly more massive than those BHs that can be explained by the standard stellar evolutionary scenario."
    },
    "0809/0809.2423_arXiv.txt": {
        "abstract": "A method to compute the full hierarchy of the critical subsets of  a density field  is presented. It is based on a watershed technique and uses a probability propagation scheme to improve the quality of the segmentation by circumventing  the discreteness  of the sampling. It can be applied within spaces of arbitrary dimensions and geometry. This recursive segmentation of space yields, for  a $d$-dimensional space, a  $d-1$ succession of  $n$-dimensional subspaces that fully characterize the topology of the density field.  The final  1D manifold  of the hierarchy is the fully connected network of  the primary critical lines of the field : the skeleton. It corresponds to the subset of lines  linking maxima to saddle points, and provides  a definition of the filaments that compose the cosmic web as a precise physical object, which makes it possible to compute any of its properties such as its length, curvature, connectivity etc...  When the skeleton extraction is applied to initial conditions of cosmological N-body simulations and their present day non linear counterparts, it is shown that the time evolution of the cosmic web, as traced by the skeleton, is well accounted for by the  Zel'dovich approximation. Comparing this skeleton to the initial skeleton undergoing the Zel'dovich mapping shows that  two effects are competing during the formation of  the cosmic web: a general dilation of the larger filaments that is captured by a simple deformation of the skeleton of the initial conditions on the one hand, and the shrinking, fusion and disappearance of the more numerous smaller filaments on the other hand. Other applications of the N dimensional skeleton and its peak patch hierarchy are discussed.\\\\ ",
        "introduction": "The web-like pattern certainly is the most striking feature of matter distribution on megaparsecs scale in the Universe. The existence of the ``cosmic web'' \\citep{zeldo70} \\citep{bbks} has been confirmed more than twenty years ago by the first CfA catalog \\citep{delap86} and the more recent catalogs such as SDSS \\citep{sdss} or 2dFGRS \\citep{2df}. These observations, together with the dramatic improvement of computer simulations (e.g. \\cite{teyssier} \\cite{ocvirk}) have largely improved the picture  of a Universe formed by an intricate network of voids ({\\it i.e.} globular under-dense regions) embedded  in a complex filamentary web which nodes are the location of denser halos. The traditional way of understanding large scale structures (LSS) formation and evolution relies on Friedman equations and assumes  that LSS are the outcome of the growth of very small primordial quantum fluctuations by gravitational  instability (see {\\it e.g.} \\cite{peebles80} or \\cite{peebles93} and references therein). In this theory, the solution for structure formation is described in terms of a mass distribution that one needs to grasp ({\\it i.e.} by  following the evolution of its most important features) and compare these to observations. Comprehending the mass distribution  as a whole, especially at non-linear stages, is a very difficult task. A possible solution therefore  consists in extracting and studying simple characteristic features of matter distribution such as voids, halos and filaments as individual physical objects. So far, mainly because of  the relatively higher complexity of the filaments, most theoretical and computational researches have focused on the voids and halos.  The dark matter halos have arguably been the most studied component of the cosmic web. Their density profiles for instance are very well described by so-called NFW profiles \\citep{NFW97} and non-parametric models are still under investigation \\citep{merritt06}. The dependence of these density profiles on the halos mass ({\\it e.g.} \\cite{bbks}, \\cite{lacey93}) has also been investigated thoroughly and its relationship with redshift and environmental properties  are a very active topics ({\\it e.g.}  \\cite{harker06}, \\cite{aubert06}, \\cite{wang07}, \\cite{aragon-calvo07}, \\cite{3dskel} or \\cite{hahn07}). From a computational point of view,  much effort has been put into the development of various algorithms to identify halos in simulations and galaxies  in spectroscopic redshift galaxy surveys. The friend-of-friend algorithm \\citep{FOF} is now widely spread, as well as more complex hierarchical sub-structures identifiers such as HFOF \\citep{gottlober98}, SUBFIND \\citep{springel01}, VOBOZ \\citep{neyrinck05} or  ADAPTAHOP \\citep{ADHOP}.  Voids are another feature of cosmological matter distribution that also have a long history of theoretical and computational modeling. The first voids were observed by \\cite{kirshner81} and are in some sense the counterpart of halos: the initial quantum perturbations collapsing into halos at non-linear stages leave room to voids in the under-dense regions. The first theoretical voids models where developed by \\cite{hoffman82}, \\cite{icke84} or  \\cite{bertschinger85} among others, while numerical void finders exist, such as the one described in \\cite{el-ad97}, ZOBOV \\citep{neyrinck08}, based on Voronoi tessellation, or the recent Watershed Void Finder, based on the Watershed transform ({\\it e.g.} \\cite{beucher79}, \\cite{beucher93}), by \\cite{platen07} (see the introduction and references therein for a more complete review of the subject). The improvements in our understanding of voids and halos properties led to the formulation  of powerful theories such as the patches theory \\citep{bbks} the extended Press-Schechter theory ({\\it e.g.}  \\cite{bond91} and \\cite{sheth98}) or the skeleton-tree formalism \\citep{hanami01}.  But our investigation of the filaments as individual objects is not yet  as thorough as for the halos and voids:  the definition of a well established mathematical framework  for their study could therefore lead to significant improvements in our understanding of matter distribution in the Universe. The first attempts date from \\cite{barrow85}, who used a graph-theory construction: the minimal spanning tree (MST). This method defines the cosmic web as the network linking galaxies (or particles from a numerical simulation), having the property of being loop-free and of minimal total length. This technique was later developed in order to try quantifying in an objective way the properties of the cosmic web (see {\\it e.g.}  \\cite{graham95}, \\cite{colberg07} and a review on the subject can be found in \\cite{martinez02}). The so-called shape finders (\\cite{sathyaprakash96}, \\cite{sahni98} or \\cite{bharadwaj00}) allow a statistical study of the filaments and another method, based on the CANDY model, commonly used to detect road networks, uses a marked point process and a simulated annealing algorithm to trace the filaments \\citep{stoica05}. More recently, the skeleton formalism and its local approximation, that describe the filaments as particular field lines of the density field, was introduced by \\cite{2dskel} and \\cite{3dskel} with the advantage of framing a well-defined mathematical ground for theoretical predictions of the filaments properties as well as an efficient numerical identification algorithm. Finally, an interesting first attempt to unify  halos, voids and filaments identification using the Multiscale Morphology Filter (MMF) technique was also proposed by \\cite{aragon-calvo07a}.  In this paper, we introduce a framework and algorithm to identify the full hierarchy of critical lines, surfaces, volumes... of density distribution in the general case of $d$-dimensional spaces. For 3D space, these critical subspaces can be identified to the void and peak patches, as well as filaments and other primary critical lines of the distribution. The algorithm extracts the filaments as a differentiable and, by definition, fully connected networks that traces the backbone of the cosmic web. This method is closely related to the skeleton formalism presented in \\cite{2dskel} and \\cite{3dskel} and is also based on both Morse theory (see {\\it e.g.} \\cite{colombi00}, \\cite{milnor63} or \\cite{jost}) and an improved Watershed segmentation  algorithm that uses a probability propagation scheme.\\\\  This paper is organized as follows. In section \\ref{sec:algo}, we present a general definition of the critical sub-spaces that we use as well as a method to extract them from sampled density  field with a sub-pixel precision (focusing more specifically on the filaments in the 2D and 3D case). In section \\ref{sec:applications}, we use  this formalism to study the time  evolution of the cosmic web,  and understand the change of its properties as a specific object via the truncated Zel'dovich approximation \\citep{zeldo70}. Finally, in section \\ref{sec:conclusion}, we summarize our findings and discuss a few possible applications to N-body simulations and observational spectroscopic galaxy surveys. The details of  a general simplex minimization algorithm used in section \\ref{sec:algo} are presented in Appendix A while the general behavior of the inter-skeleton pseudo-distance as defined in section \\ref{sec:applications} is given in appendix B. ",
        "conclusions": "\\label{sec:conclusion} We have presented a method, based on an improved watershed technique, to efficicently compute the full hierarchy of critical subsets from a density field within spaces of arbitrary dimensions. Our algorithm uses a fast one pass probability propagation scheme that is able to improve significantly the quality of the segmentation by circumventing  the discreteness  of the sampling. We showed that, following Morse theory, a recursive segmentation of space yields, for  a $d$-dimensional space, a succession of  $d-1$ $n$-dimensional subspaces that characterize the topology of the density field. In 3D for cosmological matter density distribution, we particularly focused on the 3D subspaces which are the peak and void patches of the field ({\\rm i.e. the attraction/repulsion pools}) and the 1D critical lines which trace the filaments as well as the whole primary cosmic web structure ({\\it i.e.} a  {\\sl fully-connected}, non-local skeleton as defined in \\cite{2dskel}). For  the primary critical lines, we also demonstrated that it is possible to use the probabilities distribution from our algorithm to derive a smooth and differentiable skeleton with a sub-pixel level resolution. Thus this method  allows us to consider the cosmic web as a precise physical object  and makes it possible to compute any of its properties such as length, curvature, halos connectivity etc...\\\\  As an application, we used our algorithm to study the evolution of the cosmic web, while comparing the time evolution of the skeleton (a proxy to the cosmic web) of a simulation, to those corresponding to different versions of its Zel'dovich approximation.  We first compared the evolution of the respective lengths of the different skeletons and then introduced a method to compute pseudo-distances between different skeletons. This pseudo distance makes it possible to compare different features of the skeleton such as the {size} of their filaments and the similarity of their locations. Using these measurements, we showed that two effects were competing, with net result a decrease of the cosmic length with time: a general dilation of the larger scales filaments that could be captured by a simple deformation of the skeleton of the initial conditions on the one hand,  and the shrinking, fusion and disappearance of the more numerous smaller scales filaments on the other hand. We also showed that a simple Zel'dovich approximation could accurately capture most features of the evolution of the cosmic web  for all scales larger than a few megaparsecs  (provided an effective smoothing introduced by the approximation  is taken into account). Hence in this context, the skeleton has proven to be a useful tool both for insight and as a quantitative probe and diagnostic. Conversely, the match between the simulated and the mapped skeleton confirms and extends geometrically the (point process elliptical)  peak patch theory \\citep{bbks} since both the peaks and their frontiers (the skeleton in 2D and the peak patch volumes in 3D) are well mapped by the  Zel'dovich approximation. \\\\   The domain of interest of the skeleton is quite vast and offer the prospect of  a number of applications.  From a theoretical point of view, using the points presented in this paper and in \\cite{3dskel}, we are presently developing a general theory of the skeleton and its statistical properties \\citep{pogo} that aims to understand the properties of the critical lines of scale invariant Gaussian random fields as mathematical objects.  In particular, this companion paper provides quantitative analytic predictions for the length per unit volume (resp. curvature) of the critical lines and its scaling with the shape parameter of the field, and checks successfully the current algorithm against these. In this paper we focussed on the skeleton. One could clearly investigate on the rest of the peak patch hierarchy and measure, say, the surface or volume of the (hyper)-surfaces of the recursion (whose last intersection is given by the primary critical lines). Another interesting issue would be to estimate the fraction of special (degenerate) points which do not satisfy the Morse condition, where the fields behaves pathologically  \\citep{pogo}.  For instance, one of the shortcoming of the present algorithm concerns special fields where critical lines disappear, a situation which occurs, say, in the context of tracing dendrites in a neural network, or blood vessels within a liver. Note also that there exist other sets of (geometrical) critical lines that are not topological invariants such as the lines of steepest ascent connecting directly minima to maxima which are not accounted for by the present formalism. In contrast, the algorithm is well suited to identify bifurcation points, and the connectivity of the network. In particular, in an astrophysical context, it would be worthwhile to make use of this feature and study statistically how the skeleton connects onto dark matter halos as a function of, say, they mass or spin, and investigate the details of local spin accretion in the context of the cosmic web superhighways, hence completing the spin alignment measurements of  \\cite{3dskel}  on smaller scales. More generally, the algorithm provides a neat bridge, via the provided connectivity, between the theory of continuous fields on the one hand, and graph theory for discrete networks on the other. This could prove to be of importance in the context of percolation theory. For instance, the percolation threshold can be explained in terms of the properties of the connectivity of the relevant nodes (see {\\it e.g.} \\cite{colombi00}).\\\\ Here, as argued in section~2.4 we deliberately chose not to consider the issue of shot noise and its consequences on segmentation, for which no definitive  solution yet exists, though many improvements have been proposed in the literature (see {\\it e.g.} \\cite{jos01}). Instead, we  followed  the approach of \\cite{3dskel}, that simply  involves convolving the sampled density field with a large enough (in terms of sampling scale) Gaussian kernel so that the field can be considered smooth and differentiable; the probabilistic algorithm allows for the removal of sampling effects and small intensity residual shot noise. In appendix~C we show that the corresponding fully connected skeleton is nonetheless quite robust (the core of the network remains quasi unchanged), so long as the SNR is above one. A possible drawback of this method is that it introduces a smoothing scale attached to the skeleton. This is not necessarily a problem in cosmology as the scaling of the skeleton properties with scale yields information on the distribution over these scales. Moreover, one is usually interested  by the properties of the skeleton on a given scale (typically larger than the halo scale, a few megaparsecs).  Nonetheless, there exist more complex multi-scale sampling and smoothing techniques such as the one presented in \\cite{platen07} or \\cite{colberg07} that could  straightforwardly  be adapted to our implementation. All the algorithm requires is  a structured sampling grid  where one can recover a one to one pixel neighbourhood ({\\it i.e.} one needs to be able to find the neighbouring pixels of any pixel and these pixels must have the former as neighbour as well).  For instance, we already implemented the algorithm for an {\\tt Healpix} \\citep{hivon} pixelisation of the sphere (see Figure~\\ref{healpix}),  while a direct implementation on a delaunay tesselation network is clearly an option\\footnote{for instance to segment regions on the surface of skull}.\\\\ A natural extension of the theoretical component of this work would be to investigate numerically the properties of the bifurcation points in  abstract space or anisotropic settings (see \\cite{pogo} for a theoretical discussion for isotropic Gaussian random fields). For instance, in  the context of cosmic structure formation, \\cite{hanami01}, relied  on the parallel between the skeleton of the density field in position-time 4D space and in position-scale 4D space to relate the two. In the former, the skeleton is a natural way of computing what is known as a halos merger tree, commonly used in semi-analytical galaxy formation models (see \\cite{hatton} for instance): the skeleton traces the evolution of the critical points of the density field in time. The peak theory \\citep{bbks} tells us that the smoothing scale can be linked to time evolution on scales where gravitational effects remain weakly non-linear. A worthwhile goal is to establish the parallel between the properties of 4D skeleton in this position-smoothing scale space (which can be computed from the Gaussian initial conditions only) and the halo merger tree.% Finally, note that classical bifurcation theory is concerned  with the evolution of a critical point as a function of a control parameter. In the language of the skeleton, this evolution may correspond to the skeleton in the extended ``phase space''.  From the physical and observational point of view, an other interesting venue would be to apply the skeleton to actual galaxy catalogs such as the SDSS \\citep{sdss}  to characterize the (universal) statistics of filaments as physical objects, like halos or voids, and describe them in terms of their thickness, length, curvature and environmental properties (galaxies types, halo proximity, color and morphology gradient...), both in virtual and observed catalogs. It could also be used as a diagnosis tool for inverse methods which aim at reconstructing the three dimensional distribution of the IGM from say QSO bundles \\citep{caucci} or upcoming radio surveys (LOFAR, SKA etc...) Clearly the peak patch segmentation developed in this paper will also be useful in the context of the upcoming surveys such as the LSST, or the the  SDSS-3 BAO surveys, for instance to identify rare events such as large walls or voids and study their shape. Its application to CMB related full sky data, such as WMAP or Planck should provide insight into, e.g. the level of non Gaussianity in these maps. Similarly, upcoming large scale weak lensing surveys (Dune, SNAP...) could be analysed in terms of these tools (see \\cite{pichon} for the validation of a reconstruction method in this context). Using the skeleton, the geometry of cold gas accretion that fuel stellar formation in the core of galaxies could be probed. The properties of the distribution of metals on smaller scales could be  also investigated using peak patches, to see how they influence galactic properties; one could  compare these statistical results to those obtained through WHIM detection by Oxygen emission lines (\\cite{aracil} \\cite{caucci}). Indeed it has been claimed (see  e.g. \\cite{ocvirk} \\cite{dekel}) that the geometry of the cosmic inflow on a galaxy (its mass, temperature and entropy distribution,  the connectivity of the local filaments network etc. ) is strongly correlated to its history and nature.\\\\  In closing, let us emphasize  again that the scope of application of the algorithm presented in this paper extends well beyond the context of the large scale structure of the universe: it could be used in any scientific of engineering context (medical tomography, geophysics, drilling ...) where the geometrical structure of a given field needs to be characterized.    \\begin{figure*} \\centering   \\includegraphics[width=5.5cm]{img/CMB}  \\includegraphics[width=5.5cm]{img/CMB-pp}  \\includegraphics[width=5.5cm]{img/CMB-pp-smooth} \\caption{{\\sl from left to right}: the 5 year WMAP release of the CMB temperature map, the corresponding peak patches  and the peak patches of the same field smoothed over a FHWM of 420 arcmin. Different colours represent different patches. The algorithm described in section~\\ref{sec:algo}  is implemented here on the {\\tt healpix} pixelisation. \\label{healpix}} \\end{figure*}     \\subsection*{Acknowledgements} { \\sl We thank  D. Pogosyan, D.~Aubert, J.~Devriendt, J.~Blaizot, S.~Peirani and S.~Prunet for fruitful comments during the course of this work, and D.~Munro for freely distributing his Yorick programming language and opengl interface (available at {\\em\\tt http://yorick.sourceforge.net/}).  This work was carried within the framework of the Horizon project, \\texttt{www.projet-horizon.fr}.   }"
    },
    "0809/0809.5035_arXiv.txt": {
        "abstract": "{% Weak-lensing surveys need accurate theoretical predictions for interpretation of their results and cosmological-parameter estimation. }{% We study the accuracy of various approximations to cosmic shear and weak galaxy-galaxy lensing and investigate effects of Born corrections and lens-lens coupling. }{% We use ray-tracing through the Millennium Simulation, a large $N$-body simulation of cosmic structure formation, to calculate various cosmic-shear and galaxy-galaxy-lensing statistics. We compare the results from ray-tracing to semi-analytic predictions. }{% (i) We confirm that the first-order approximation (i.e. neglecting lensing effects beyond first order in density fluctuations) provides an excellent fit to cosmic-shear power spectra as long as the actual matter power spectrum is used as input. Common fitting formulae, however, strongly underestimate the cosmic-shear power spectra (by $>30\\%$ on scales $\\ell>10000$). Halo models provide a better fit to cosmic shear-power spectra, but there are still noticeable deviations ($\\sim10\\%$). (ii) Cosmic-shear B-modes, which are induced by Born corrections and lens-lens coupling, are at least three orders of magnitude smaller than cosmic-shear E-modes. Semi-analytic extensions to the first-order approximation predict the right order of magnitude for the B-mode. Compared to the ray-tracing results, however, the semi-analytic predictions may differ by a factor two on small scales and also show a different scale dependence. (iii) The first-order approximation may under- or overestimate the galaxy-galaxy-lensing shear signal by several percent due to the neglect of magnification bias, which may lead to a correlation between the shear and the observed number density of lenses. }{% (i) Current semi-analytic models need to be improved in order to match the degree of statistical accuracy expected for future weak-lensing surveys. (ii) Shear B-modes induced by corrections to the first-order approximation are not important for future cosmic-shear surveys. (iii) Magnification bias can be important for galaxy-galaxy-lensing surveys. } ",
        "introduction": "During the past few years, weak gravitational lensing has developed rapidly from mere detection to an important cosmological tool \\citep{MunshiEtal2008_wl_review}. Measurements of cosmic shear help us to constrain the properties of the cosmic matter distribution \\citep[e.g.][]{SemboloniEtal2006,HoekstraEtal2006,SimonEtal2007,BenjaminEtal2007,MasseyEtal2007_cosmic_scaffolding,FuEtal2008}, the growth of structure \\citep[e.g.][]{BaconEtal2005,MasseyEtal2007_3D_WL}, and the nature of the dark energy \\citep[e.g.][]{TaylorEtal2007,SchimdEtal2007,AmendolaKunzSapone2008}. Weak galaxy-galaxy lensing can be used to study the properties of galactic dark-matter halos and the relation between luminous and dark matter \\citep[e.g.][]{MandelbaumEtal2006_Galaxies,SimonEtal2007,GavazziEtal2007}.  The accuracy that can be reached in weak-lensing surveys is determined by several factors. On the observational side, high accuracy requires large field sizes and deep observations with a high number density of galaxies with measurable shapes. Moreover, it is crucial to obtain an accurate and unbiased measurement of galaxy ellipticities. Finally, for the interpretation of the resulting data and the inference of cosmological parameters, an accurate theoretical model is needed. A thorough understanding of systematic effects in weak lensing will become particularly important with the advent of very large weak-lensing surveys such as CFHTLS\\footnote{\\texttt{http://www.cfht.hawaii.edu/Science/CFHLS}}, KIDS\\footnote{\\texttt{http://http://www.astro-wise.org/projects/KIDS}}, Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu}}, and LSST\\footnote{\\texttt{http://www.lsst.org}}, or the planned Dark Energy Survey\\footnote{\\texttt{http://www.darkenergysurvey.org}}, DUNE \\citep{RefregierEtal2006_DUNE}, and SNAP\\footnote{\\texttt{http://snap.lbl.gov}}. For these surveys, the statistical uncertainties will be very small, so the accuracy will be limited by the remaining systematics in the data reduction and theoretical modeling.  While significant improvement on image-ellipticity measurements are expected in the near future \\citep{MasseyEtal2007_STEP2}, one still needs to investigate, how uncertain current theoretical predictions are, and how much improvement can be expected for these. Presently, the most accurate way to obtain predictions for weak-lensing surveys is to perform ray-tracing through large high-resolution $N$-body simulations of cosmic structure formation \\citep[see, e.g.,][]{WambsganssCenOstriker1998,JainSeljakWhite2000,WhiteHu2000,VanWaerbekeEtal2001,HamanaMellier2001,ValeWhite2003,White2005}. The drawback of this approach is that large $N$-body simulations are computationally demanding, so using them to explore the whole parameter space of cosmological models is currently unrealistic. On the other hand, ray-tracing simulations enable one to check the approximations and assumptions made in computationally less demanding (semi-)analytic models, and adjust and extend these models where necessary.  Numerous ray-tracing methods have been developed to study the many aspects of gravitational lensing. Tree-based ray-tracing methods \\citep[][]{AubertAmaraMetcalf2007} that adapt to the varying spatial resolution of $N$-body simulations have been used to study the impact of substructure on strong lensing by dark matter halos \\citep[][]{PeiraniEtal2008}. Cluster strong lensing simulations, which require good mass modelling of galaxy clusters, usually ignore the matter distribution outside clusters and use the thin-lens approximation in the ray-tracing \\citep[e.g.][]{BartelmannWeiss1994,MeneghettiEtal2007,RozoEtal2008}.  Many simulations of weak lensing by clusters and large-scale structure \\citep[e.g.][]{WambsganssCenOstriker1998,JainSeljakWhite2000,ValeWhite2003,PaceEtal2007} employ algorithms that are based on the multiple-lens-plane approximation \\citep{BlandfordNarayan1986} to trace light rays through cosmological $N$-body simulations. Others \\citep[e.g.][]{CouchamnBarberThomas1999,CarboneEtal2008_I} perform ray-tracing though the three-dimensional gravitational potential. In a simpler approach \\citep[e.g.][]{WhiteVale2004,HeymansEtal2006,HilbertMetcalfWhite2007}, the matter in the $N$-body simulation is projected along unperturbed light paths onto a single lens plane, which is then used to calculate lensing observables. Recent simulations of CMB lensing use generalisations of the single- or multiple-plane approximation that take the curvature of the sky into account \\citep[e.g.][]{DasBode2008,TeyssierEtal2008_arXiv,FosalbaEtal2008}.  In this work, we employ multiple-lens-plane ray-tracing through the Millennium Simulation \\citep{SpringelEtal2005_Millennium} to study weak lensing. One of the largest $N$-body simulations available today, the Millennium Simulation provides not only a much larger volume, but also a higher spatial and mass resolution than simulations used for earlier weak-lensing studies. In order to take full advantage of the large simulation volume and high resolution, the ray-tracing algorithm used here differs in several aspects from algorithms used in previous works \\citep[e.g.][]{JainSeljakWhite2000}. Here, we give a detailed description of our ray-tracing algorithm.  Semi-analytic weak-lensing predictions are usually based on the first-order approximation, in which light deflections are only considered to first order in the peculiar gravitational potential and hence, to first order in the matter fluctuations. The ray-tracing approach allows us to look at effects neglected in the first-order approximation such Born corrections and lens-lens coupling. Here, we investigate the cosmic-shear B-modes induced by these effects and compare the ray-tracing results to semi-analytic estimates \\citep{CoorayHu2002,HirataSeljak2003,ShapiroCooray2006}, whose accuracy has not been confirmed by numerical simulations yet. Moreover, we investigate how well fitting formulae \\citep{PeacockDodds1996,EisensteinHu1999,SmithEtal2003} and halo models \\citep{Seljak2000,CooraySheth2002} reproduce cosmic-shear power spectra. Finally, we investigate the accuracy of the first-order approximation for weak galaxy-galaxy lensing.  The paper is organised as follows. In the next section, we introduce the theoretical background and notation used in our lensing analysis. In Sec.~\\ref{sec:algorithm}, we discuss our ray-tracing algorithm. The results from our ray-tracing analysis are presented in Sec.~\\ref{sec:results}. We conclude our paper with a summary in Sec.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this work, we have described a new variant of the multiple-lens-plane  algorithm, which is particularly suited for ray-tracing through very large cosmological $N$-body simulations. The algorithm differs in some important details from previous works. This allows us to take full advantage of the unprecedented statistical power offered by the large volume and high spatial and mass resolution of the Millennium Simulation. The features discussed include: a tilted line-of-sight (to avoid periodic repetition of structures along the line-of-sight), adaptive slice boundaries (to avoid the slicing and duplication of bound structures), adaptive smoothing of the projected matter distribution on the lens planes (to reduce shot noise from the particles), a mutliple-mesh method for calculating the light deflections and distortions at the lens planes (which takes into account the small-scale and large-scale structure simultaneously), and a  method to include galaxies (as lenses and sources) from semi-analytic galaxy-formation models in the ray-tracing process.  We have used the ray-tracing code and the Millennium Simulation to investigate the impact of lens-lens coupling and multiple ray deflections on various cosmic shear two-point statistics. We have computed convergence power spectra from a set of ray-tracing realisations. For testing and comparison, we have also computed a first-order prediction of the convergence power spectrum using the measured three-dimensional power spectra of the mass distribution in the Millennium Simulation. We find that this first-order prediction agrees very well with the ray-tracing results except for very small scales (the difference is $>5\\%$ only for $\\ell>20000$), where smoothing on the lens planes becomes important.  Comparing the convergence power spectrum from the ray-tracing to the predictions based on the fitting formulae for the matter power spectrum by \\citet{PeacockDodds1996} and \\citet{SmithEtal2003}, we find significant discrepancies ($>30\\%$ for $\\ell>10000$), casting the usefulness of these fitting formulae for cosmological parameter estimation for future surveys into doubt. A prediction based on the popular halo model and the halo concentration-mass relation of \\citet{NetoEtal2007} fits better, but there are still noticeable deviations, in particular for higher source redshifts ($\\sim10\\%$ for sources at $\\zS=2$). This indicates a need for more accurate descriptions of matter power spectra.  Furthermore, we have computed the E- and B-mode aperture mass dispersion using our ray-tracing algorithm. We find the B-mode to be finite, but at least three orders of magnitude smaller than the E-mode. The amplitude of the B-mode is slightly larger and shows a different scale dependence than the predictions of \\citet{CoorayHu2002} and \\citet{HirataSeljak2003}. We have performed various tests to exclude numerical artifacts as the origin of the deviations. Despite these discrepancies, we can confirm the finding of \\citet{ShapiroCooray2006} that the lensing-induced B-mode can be safely neglected even in an all-sky survey.  Corrections to the first-order approximation can have a considerable impact on galaxy-galaxy lensing. In the simple case of a volume-limited sample of unbiased lens galaxies and all sources at redshifts $z=1$, the first-order approximation overestimates the mean tangential shear around lenses by $\\approx 10-20\\%$ at an angular separation of $1\\,\\arcmint$ due to its failure to incorporate the magnification bias. The impact of the magnification bias on the galaxy-galaxy lensing signal depends on the survey selection criteria and the luminosity and redshift distribution of the sources and the lenses. A detailed investigation of this effect should be carried out in future work."
    },
    "0809/0809.2747_arXiv.txt": {
        "abstract": "The central regions of galaxies harbor some of the most extreme physical phenomena, including dense stellar clusters, non-circular motions of molecular clouds and strong and pervasive magnetic field structures.  In particular, radio observations have shown that the central few hundred parsecs of our Galaxy has a striking magnetic field configuration. It is not yet clear whether these magnetic structures are unique to our Milky Way or a common feature of all similar galaxies.  Therefore, we report on (a) a new radio polarimetric survey of the central 200 pc of the Galaxy to better characterize the magnetic field structure and (b) a search for large-scale and organized magnetized structure in the nuclear regions of nearby galaxies using data from the Very Large Array (VLA) archive. The high angular resolution (1--5$^{\\prime\\prime}$) of the VLA allows us to study the central 1\\,kpc of the nearest galaxies to search for magnetized nuclear features similar to what is detected in our own Galactic center. Such magnetic features play a important role in the nuclear regions of galaxies in terms of gas transport and the physical conditions of the interstellar medium in this unusual region of galaxies. ",
        "introduction": "\\vskip 0.1in Magnetic fields play crucial roles in the interstellar environments of galaxies and are observed on on a number of scales in galaxies, up to scale lengths of 8\\,kpc for the regular, organized fields in the disks (e.g., Beck 2004, Beck and Gaensler 2006). Radio synchrotron observations are one of the best ways to probe magnetic fields across galaxies as radio waves are not subject to interstellar extinction. In addition to revealing maps of the magnetic field structure, radio observations give estimates for the mean equipartition magnetic field strengths in galaxies from $\\sim$ 10-20 $\\mu$G in galactic spiral arms, and $\\sim$40 $\\mu$G in galactic nuclear regions. The central regions of galaxies harbor some of the most extreme physical phenomena, including dense stellar clusters, non-circular motions of molecular clouds and strong and pervasive magnetic field structures. In addition, studies across the electromagnetic spectrum reveal that galactic nuclei are major sites of episodic and energetic phenomena related both to active galactic nuclei (AGN) and also to intense bursts of star formation.  Recently, studies of the circumnuclear regions of several galaxies have revealed that magnetic fields in this part of the galaxy may help to transport materials into the nuclear region (which then may fuel a starburst or AGN). In the barred galaxies, NGC\\,1097 and NGC\\,1365, radio emission is enhanced along the bar regions, indicating the presence of a shock front. Further, magnetic stresses in the circumnuclear ring can then drive mass inwards at rates high enough to feed nuclear activity (Beck \\textit{et al} 1999; 2005). In the spiral ringed galaxy M\\,94 (NGC\\,4736) the polarized radio emission reveals a pattern of ordered magnetic field arising from the central regions of the galaxy that may be where magnetic amplification occurs (Chyzy and Buta 2008). In addition, the expulsion of materials from the nucleus may also be regulated by magnetic fields. Energetic events arising from star-burst or AGN activity in the nucleus often result in explosive vertical ``fountains'' of million degree gas, rising 1--2\\,kiloparsec (kpc) above galactic nuclear regions (e.g., NGC\\,1569 and NGC\\,4631; Heckman \\textit{et al} 1995). These fountains may well be related to vertical magnetic field structures and outflow from the galactic disk (e.g., NGC\\,4631; Golla and Hummel 1994).  Detailed knowledge of the violent interplay between stellar and magnetized interstellar components is limited for most galactic nuclei because of their great distances, even for the most sophisticated observational instrumentation. At a distance of only 8.0\\,kpc, the center of our Milky Way galaxy provides an excellent laboratory for detailed studies of the interplay of massive stars, gas and, importantly, magnetic fields. Because of the nearly 30 magnitudes of visual extinction toward this region of our Galaxy, studies at radio wavelengths have proved to be fruitful in understanding the details of the magnetized interstellar medium and for understanding the nuclei of nearby, normal galaxies. ",
        "conclusions": ""
    },
    "0809/0809.0432_arXiv.txt": {
        "abstract": "{Radiation from accretion discs in cataclysmic variable stars (CVs) provides fundamental information about the properties of these close binary systems and about the physics of accretion in general. Of particular interest are dwarf-nova outburst cycles during which variations of the disc properties allow a detailed study of the physical processes in accretion flows.} {The detailed diagnostics of accretion disc structure can be achieved by including in its description all the relevant heating and cooling physical mechanism, in particular the convective energy transport that, although dominant at temperatures $\\lta 10^4\\,\\rm K$, is usually not taken into account when calculating spectra of accretion discs. The disc's self-consistently calculated structure and emission allow testing models of dwarf-nova outbursts and accretion-disc models in general.} {We constructed a radiative transfer code coupled with a code determining the disc's hydrostatic vertical structure.} {We have obtained for the first time model spectra of cold, convective accretion discs. As expected, these spectra are mostly flat in the optical wavelengths with no contribution from the UV, which in quiescence must be emitted by the white dwarf. The disc structures obtained with our radiative-transfer code compare well with the solutions of equations used to describe the dwarf-nova outburst cycle according to the thermal-viscous disc instability model thus allowing the two to be combined. For high-temperature radiative discs our spectra are compatible with models obtained with Hubeny's code \\textsl{TLUSTY}.} {Our code allows calculating the spectral evolution of dwarf nova stars through their whole outburst cycle, providing a new tool for testing models of accretion discs in cataclysmic variables. We show that convection plays an important role in determining the vertical disc structure and substantially affects emitted spectra when, as often the case, it is effective at optical depths $\\tau\\sim 1$. The emergent spectrum is independent of the parameters of the convection model. We confirm that, as required by the disc instability model, quiescent discs in dwarf novae must be optically thick in their outer regions. In general, no emission lines are present in the absence of external irradiation.} \\keywords {accretion,accretion discs -- Stars: dwarf novae} ",
        "introduction": "\\label{sec:intro}  In most cataclysmic variable stars (CVs) matter lost by the Roche-lobe filling low-mass secondary star forms an accretion disc around the white-dwarf primary. In bright systems such as nova-like stars and dwarf novae in outburst, the disc is the dominant source of luminosity. In quiescent dwarf novae, although rather dim, the disc emission provides crucial information about the physics of accretion discs, in particular about the mechanisms driving the instabilities that are responsible for the dwarf-nova outbursts. It is generally accepted \\citep[see, however,][]{sl07} that properties of the dwarf-nova outburst cycles are reproduced well by the thermal-viscous disc instability model \\citep[DIM; see][for a review]{l01}, but in reality the description of some phases of theses cycles is far from satisfactory. This is especially true of the quiescence state. Until now models of disc emission spectra have been calculated consistently only for hot, bright accretion discs \\citep[see][and references therein]{wh98}. In such a scheme the vertical structure of an accretion ring is calculated using the program TLUSDISK \\citet{hubeny90,hubeny91}, whereas the spectra are synthesized by the code SYNSPEC \\citep{synspec}. A different approach to radiative transfer in accretion discs was presented by \\citet[][hereafter SW]{sw91}, but also in this case it was applied  only to hot discs. An attempt to apply the SW code to cold quiescent discs in \\citet{Idan1} was not conclusive.  The main obstacle in solving the vertical structure of cold accretion discs is the presence of convection, which was difficult to include in radiative transfer codes (Hubeny, private communication); models of \\citet{wh98} are restricted to temperatures for which convection is unimportant. In their pioneering work \\citet{pringletal86} used vertically integrated disc structures and Kurucz's stellar spectra to test the DIM versus the model attributing dwarf-nova outbursts to an increase of mass-transfer rate from the secondary. More recently \\citet[][see also \\citet{nagel2}]{nagel3} have calculated non-LTE accretion-disc spectra throughout the outburst cycle of the dwarf nova SS~Cyg, taking into account white-dwarf irradiation. They found that in order to obtain agreement with observations of this system in quiescence one has to assume a radial disc that is optically thin in its outer regions.  Such a structure, however, is inconsistent with the DIM. The AcDc NLTE code by  \\citet{nagel04} used to calculate these spectra does not take into account the (often dominant)\\,vertical convective energy transport.  Of related interest are radiative-transfer codes describing emission from cold proto-stellar accretion discs such as \\citet{Hugel09} and \\citet{min09}. Neither of these codes includes convection.  There is therefore a need for a radiative transfer code consistent with the DIM. This is the main motivation behind the present work. In particular we have obtained a scheme that is totally consistent with the code of \\citet[][hereafter HMDLH]{hmdl98}, which allows following the dwarf-nova spectra during various phases of the outburst with special emphasis put on quiescence. One should stress that the principal aim of our project is testing the DIM, not trying to reproduce the observed spectra of dwarf-nova stars at all costs. In particular our code should allow calculating various delays in rising to (and decaying from) outburst between light at different wavelengths (commonly known a the ``UV-delay\") and determining what properties of the quiescent dwarf-nova disc can be accounted for by the DIM).  The HMDLH version of the DIM uses a 1+1\\,D scheme in which the time-dependent radial evolution equations use as input a pre-calculated grid of hydrostatic vertical structures. These structures are calculated using the standard equations of stellar structure (or the grey-atmosphere approximation in the optically thin case). The angular-momentum viscous transport mechanism is described by the $\\alpha$-ansatz of \\citet{ss73}. For a given $\\alpha$, $M/R^3$ and effective temperature $\\Teff$ there exists a unique solution describing the disc vertical structure. Such solutions are very handy for solving the time-dependent equations of the disc evolution but do not produce realistic emission spectra. These can be calculated, however, by using the same input parameters ($\\alpha$, $M/R^3$ and $\\Teff$) in a radiative-transfer code on the condition that the vertical structures calculated by the two methods are the same up to the photosphere. In this way one can reproduce spectra of the whole cycle of a dwarf-nova outburst.  The outline of the article is as follows. In Section \\ref{sec:model} we present the model used, with special stress put on marking the differences with the original SW code on which it is based. We discuss there also in details the opacities used. Section \\ref{sec:onering} is dealing with the solutions obtained for various types of physical set-ups. In this section we compare our vertical structure solutions with those obtained with the HMDLH code for a hot stationary and cold non-stationary discs. These two models represent the main phases of the dwarf-nova outburst cycle respectively: the outburst and quiescence. We compare the corresponding thermal-equilibria (``S-curves\") in \\ref{subsec:sc}. The importance of convection in determining the vertical disc structure is discussed in detail in section \\ref{subsec:Conv}. In section \\ref{sec:full_disc} we show our hot whole-disc solutions as examples of spectra emitted during the decay from maximum of a dwarf-nova outburst. We also show that, despite differences in treatment of vertical viscosity stratification our solutions compare very well with the UV spectra obtained by \\citet{wh98}. Then in section \\ref{sec:non-stat} we present and discuss the structure and spectrum of cold quiescent accretion disc. Finally we conclude the paper in Section \\ref{sec:concl} by shortly discussing future developments and necessary improvements. A report relating the progress of our work on the radiative-transfer code was published in the conference proceedings \\citet{Idan2}. Several solutions presented in this report were only preliminary and have been substantially improved, some results we found to be incorrect. We address these points in the present article. ",
        "conclusions": "\\label{sec:concl}  We have constructed a numerical code solving radiative transfer in accretion discs for a range of effective temperatures reaching values as low as 3000\\,K. As far as we know, our code is the only one taking into account convective energy transport, including the cases when convection is dominating the energy transport. At high temperatures, in the case of fully radiative accretion discs, our spectra are in a very satisfactory agreement with the calculations of \\citet{wh98} despite the different viscosity prescriptions respectively used. The $\\alpha$-viscosity prescription of our code is used in the Disc Instability Model and in general to characterize the structure of accretion discs.  Although our code includes convective energy transport it fails to account for important effects that have been successfully taken into account by other authors. For example we assume LTE, whereas the AcDC code \\citep{nagel04} describes the NLTE structures. These different approaches are therefore only milestones on the road towards a complete description of accretion disc emission. Unfortunately the main obstacle one will have to overcome is the persisting lack of a reliable description of the disc's ``viscosity\" stratification.  Our code has been designed to work in synergy with the DIM code of HMDLH. This goal has been achieved for most of the range of the relevant parameters. We find, in agreement with the predictions of the DIM, that quiescent dwarf-nova discs must be optically thick. In absence of spontaneously formed disc corona emission lines observed in quiescent dwarf-novae have to be the result of irradiation by the white dwarf (and the boundary layer) of an optically thick disc as suggested by \\citet{Smak91}.  In a subsequent article we will include the effects of limb darkening and produce a grid of disc spectra to be compared with observations of various types of cataclysmic variables and various phases of the dwarf-nova outburst cycle. In particular we will focus on the quiescence discs and the effect of the irradiation by the white dwarf and the formation of emission lines. Combining ILHS with the time dependent HMDLH code will provide a tool for the study of dwarf nova outburst cycles.  In particular it will be possible to compare observed quiescent dwarf nova spectra with the model spectra of cold non-equilibrium accretion discs and hot white dwarfs, instead of assuming \\citep[as in e.g.][and references therein]{urbsion} that the disc in such a state is hot and in equilibrium.  The structure of our code allows easy inclusion of various tabulated opacity data. In a future work we plan calculating spectra of helium discs and very cold discs in which molecular opacities are important."
    },
    "0809/0809.2601_arXiv.txt": {
        "abstract": "We present the dark matter (DM) extension (by $N$) of the minimal supersymmetric standard model to give the recent trend of the high energy positron spectrum of the PAMELA/HEAT experiments. If the trend survives by future experiments, the MSSM needs to be extended. Here, we minimally extend the MSSM with one more DM component  $N$ together with a heavy lepton $E$, and introduce the coupling $e_R E_R^c N_R$. This coupling naturally appears in the flipped SU(5) GUT models. The model contains the discrete symmetry $Z_6$, and for some parameter ranges there result two DM components. For the MSSM fields, the conventional $R$-parity, which is a subgroup of $Z_6$, is preserved. We also present the needed parameter ranges of these additional particles. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1434_arXiv.txt": {
        "abstract": "We have used the Parkes Multibeam system and the Sloan Digital Sky Survey (SDSS) to assemble a sample of 195 galaxies selected originally from their \\hi signature to avoid biases against unevolved or low surface brightness objects. For each source 9 intrinsic properties are measured homogeneously, as well as inclination and an optical spectrum. The sample, which should be almost entirely free of either misidentification or confusion, includes a wide diversity of galaxies ranging from inchoate, low surface brightness dwarfs to giant spirals. Despite this diversity there are 5 clear correlations among their properties. They include a common dynamical mass-to-light ratio within their optical radii, a correlation between surface-brightness and Luminosity and a common \\hi surface-density. Such correlation should provide strong constrains on models of galaxy formation and evolution. ",
        "introduction": "Searching for systematic correlations among the global properties of galaxies, analogous to the H-R diagram for stars, may offer the best hope of having a better understanding about their formation and evolution. The recent availability of large and systematic data-sets at several wavelengths makes this a good time to search. Such searches go back to the optical pioneers of galaxy exploration such as Hubble (1937), Holmberg (1965), de Vaucouleurs (1991) and Zwicky (1942). The trends and correlations they discovered, re-emerge very clearly in the Sloan Digital Sky Survey, as described by Blanton et al. (2003) in their analysis of 200.000 galaxies at redshifts of $\\sim$ 0.1.  This paper deals with a much smaller sample of galaxies, but ones found by an entirely different technique, i.e. in a blind search for neutral Hydrogen gas at 21-cm. The principal motivations for such a blind search are two. First of all, by definition, optically selected galaxies already contain many stars, which may not be the case in the younger, or less evolved galaxies, which may have the most to tell us about their formation and evolution. Apart from X-ray searches, which detect the very hot gas in the potential wells of giant ellipticals, most galaxy searches rely, either directly or indirectly, on the presence of stars and so are biased against young or unevolved objects. Secondly, a blind \\hi search offers a unique way round the strong optical surface brightness selection effects, which could disguise the very correlations one is looking for. By definition a galaxy must be separately both luminous enough \\emph{and} large enough to distinguish it above a sky background which, by galaxy standards, is quite bright. These two \\emph{separate} requirements squeeze optically selected galaxies into an extremely narrow range of surface brightnesses (SB) centred, for a given catalogue, on $\\Sigma_{cat}$ where: \\begin{equation} \\Sigma_{cat} \\equiv \\frac{l_{ap}}{\\pi\\theta^{2}_{ap}} \\end{equation} where $l_{ap}$ and $\\theta_{ap}$ are the minimum apparent luminosity, and the minimum apparent angular radius which the galaxy catalogue will accept (Disney 1999). The Full Width Half Maximum (FWHM) of the detectable range is typically only 3 magnitudes wide. Such a narrow stricture certainly impoverished the old photographic surveys (Disney 1976, Disney \\& Phillipps 1983, Dalcanton 1997). It might be naively thought that CCDs would have cured this stricture. Not so, because they detect galaxies that are both fainter \\emph{and} smaller than emulsions, their $\\Sigma_{cats}$'s are (see equation (1)) scarcely any dimmer. On the other hand, the limitation with an \\hie-selected sample is that it will miss the early type galaxies - which contain very little neutral gas. According to the SDSS survey late types make up the large majority of all galaxies, although they may emit less than their fair share of light (Blanton et al. 2001).  Blind \\hi surveys face four main problems. To reach sources of a given column density of \\hi (N$_{H\\!I}$ in atoms cm$^{-2}$) any survey must be at least sensitive enough to find them in the most favourable case, i.e. when they fill the beam. In that case (Minchin 2003, Disney 2008) t$_{obs}$ (per beam) $\\ge$ $k$/[N$_{H\\!I}$ (min)]$^{2}$ where $k$ is independent of dish-diameter D (because a larger dish projects the same system noise onto a smaller area of sky). Since there is a loose correspondence between N$_{H\\!I}$ and SB (e.g. Swaters et al. 2003). Mathematically: \\begin{equation} N_{H\\!I} \\cong 10^{20.1} \\left(\\frac{M_{H\\!I}}{L_{B}}\\right)_{\\odot} 10^{[0.4(27-{\\overline{\\mu}}_{21})]}~atoms~cm^{-2} \\end{equation} (Disney \\& Banks 1997) where $\\left(\\frac{M_{H\\!I}}{L_{B}}\\right)_{\\odot}$ is in solar units and $\\overline{\\mu}_{21}$ is the average SB in B mag arcsec$^{-2}$, taken over the same area as N$_{H\\!I}$. Since the effective SB ($\\mu_{ef\\!f}$) of an exponential disc is 2.1 magnitudes brighter than $\\overline{\\mu}_{21}$ (Salpeter \\& Hoffman 1996) column density sensitivity at the $N_{H\\!I} < 10^{19.5} cm^{-2}$ level is needed to detect LSBGs (i.e. $\\mu_{ef\\!f}$ $\\>$ 25.0 B mag arcsec$^{-2}$) with $\\left(\\frac{M_{H\\!I}}{L_{B}}\\right)_{\\odot}$ $\\approx$ 0.3 (see Section 3), which in turn requires integration times per beam of several hundred seconds - which rendered earlier blind \\hi surveys either too insensitive or too impractical to detect LSBGs. Since this difficulty was not recognised until recently, some earlier claims, based on small blind \\hi surveys, to set rigid upper limits to the amount of cosmic \\hi or the numbers of LSBGs in the Universe, must be set aside [e.g. Shostak 1977, Fisher and Tully 1981, Zwaan et al. 1997].  The coming of the Multibeam \\hi detector (Staveley-Smith et al. 1996) was essential to carry out such blind \\hi surveys. HIPASS, the \\hi Parkes All Sky Survey (Meyer et al. 2004) used the first multibeam system to survey the entire Southern Sky, and the north up to +25$^{\\circ}$. More than 4000 sources were identified in the Southern Sky alone. The Equatorial Survey, described in this paper, was initially part of HIPASS, but with an accelerated search so as to exploit the SDSS-DR2 optical data when it emerged. HIPASS sources are typically more than a degree apart which makes for a real challenge in obtaining complementary optical, and other data, simply because, before SDSS, each source required a separate observing campaign.  The strong clustering of galaxies, and of \\hi galaxies in particular, makes the identification of the source with an optical candidate quite challenging. This is true even when both 21-cm and optical velocities are known, for galaxy velocities are strongly clustered too. To be certain that every Parkes-detected source was correctly identified with its optical counterpart would require interferometric follow-up in every case - which is infeasible when one is dealing with hundreds of sources, as here. We have used interferometry, coupled with simulations, to be certain that only a handful of the remaining sources (less than 10) are still misidentified. This handful is too small to affect the main correlations.  The Equatorial Strip Survey (ES henceforth), reported on here, was a search through HIPASS cubes between declinations $-$6\\arcdeg and +10\\arcdeg. Thus 5780 deg$^{2}$ of \\hi data, in the velocity range between $-$1280 and +12700 \\kmse, were searched largely by eye to come up with 1107 sources (Garcia-Appadoo PhD Thesis 2005). The Equatorial Strip was chosen because: (i) it was approximately perpendicular to the Galactic Plane and so it is mostly dark enough to make LSBGs detectable; (ii) it is accessible for follow up with a large range of instruments including the VLA, and for some of its area the SDSS-DR2; (iii) it includes the area searched for LSBGs by Impey et al. (1996) using the APM machine (Cawson et al. 1987) to scan UK-Schmidt plates. This should eventually lead to an estimate of the number of LSBGs still missing from wide-scale optical surveys. So long as LSBGs remain a putative reservoir for missing baryons (e.g. Fukugita et al. 1998) it is important to know this. Of the ES area, 50 per cent will eventually be scanned by SDSS, and 35 per cent of it is already publicly released. We concentrated our search in the 1700 deg$^{2}$ of the ES released earlier in DR2 (Abazajian et al. 2004). 370 cross identifications were made, and then refined down to only 195 to reduce possible misidentifications to a minimum.  These 195 sources have SDSS-DR2 optical diameters (50 per cent light diameter and 90 per cent light diameter) between 0.16 and 6 arcmins (West et al. 2008). The SDSS pipeline photometry on sources of this size was known to be extremely unreliable so we had to devise new techniques to reduce the data (West PhD Thesis 2005), which delayed the project by nearly two years. Eventually we will analyse many more ES sources already observed with SDSS. But the results found for the first 195 are clear and interesting enough to merit being published now.  We have measured, at 21-cm, the peak flux, the integrated flux, the line width, and the spectral shape (e.g. two horned or Gaussian mainly), and in the optical the luminosities at $u, g, r, i, z,$ two radii R$_{50}$($g$) and R$_{90}$($g$) containing respectively 50 and 90 per cent of the $g$-band light, the inclination, the morphology and a fibre nuclear spectrum set either on the nucleus or the brightest HII region. Given the strong correlations between some of the colours this amounts to about a dozen independent measurements, not far short of the 14 desired. Chiefly missing is a rotation curve to spell out the distribution of Dark Matter. Nevertheless we have sufficient information to hope that some important aspects of galaxy systematics will emerge. With 13 properties there will be $\\sim$ (13$\\times$12)/2 $\\sim$ 80 possible correlations to look for.  The previous largest survey of discs was reported in a remarkable, but largely unremarked paper by Gavazzi et al. in 1996 entitled \\emph{The phenomenology of disc galaxies}. They found: (a) the mass-to-light ratio in the NIR is virtually constant and equal to 4.6 in H-band solar units; (b) Population I indicators are all anti-correlated with mass; (c) Conversely the luminosity and surface brightness of old stellar populations all increase with mass; (d) Bulge components all increase non-linearly with mass; (d) All the above properties are independent of either morphological type or environment.  Previous blind \\hi surveys, with follow-up optical or NIR data, have been carried out by Zwaan et al. (1997), Spitzak and Schneider (1998), Rosenberg and Schneider (2000), Minchin et al. (2003),  Davies et al. (2004), Giovanelli et al. (2005). Some of the results are discussed in Rosenberg et al. (2005), Rosenberg and Schneider (2003) and Minchin et al. (2004). The main results can be summarised briefly as follows: \\begin{enumerate} \\item all but one galaxy (Minchin et al. 2005, 2007) seem to have optical counterparts; \\item although spirals predominate, the counterparts cover a wide variety of morphological types and luminosities; \\item there are significant, but not overwhelming, numbers of low surface brightness galaxies; \\item to within the measurement errors all \\hi galaxies have the same \\hi column densities ($\\sim 10^{20.65 \\pm 0.38} cm^{-2}$) and this is not a selection effect. \\end{enumerate} The various surveys differ from one another partly in radio sensitivity, but mainly in the quality of their optical or NIR follow-up data. Our Equatorial Survey is distinct in two main respects: a) we have been fastidious in rejecting all sources where there might be more than one galaxy in or near the beam; b) we have available high quality multi-band SDSS-DR2 data for every source, data which is both sensitive and uniform, across a wide dynamic range. Rosenberg at al. (2005) have attempted a similar analysis using, instead of optical, 2MASS J-band data.  Our most interesting results can be summarised as follows: a) All galaxies have the same mass-to-light ratio in H-band solar units. In other words Gavazzi's Law holds over a dynamic range of 8 mag in absolute luminosity even in \\hie-selected galaxies free of optical selection effects. b) There is a clear correlation between SB and Luminosity. Combined with (a) it reveals a correlation between dynamical mass and optical radius cubed, again over a large dynamic range, but now with more scatter. In other words there is roughly constant global-density for galaxies (\\mbox{$\\sim$ 1 $\\times$ 10$^{-24}~gm~cm^{-3}$}). The Virial Theorem then implies the same angular velocity \\mbox{$\\sim$ 4 $\\times$ 10$^{-16} ~radians~s^{-1}$}, implying they could have rotated no more than 40 times in a Hubble epoch. (c) All the galaxies have the same global surface density in \\hie, that is to say the same \\hie-mass divided by optical-radius squared (R$^{2}$); (d) All the galaxies have exponential profiles throughout most of their extents so that $R_{90}$ is correlated very tightly with $R_{50}$; (e) There is a clear colour-luminosity correlation in the sense that more luminous galaxies are redder, which is well known in optically selected samples.  The remainder of this paper is arranged by section as follows:  (2) \\textit{The Radio Observations} describe the \\hi survey, the detection process and the contents of the resulting catalogue.  (3) \\textit{The Optical Observations} looks into the sensitivity of the SDSS for picking up large angular-size low surface brightness galaxies, and galaxies with low $\\left(\\frac{M_{H\\!I}}{L_{B}}\\right)$ ratios. It then discusses the general problem of identifying reliable optical counterparts to \\hi sources in blind \\hi surveys.  (4) \\textit{The Diversity of Sources} displays complete data sets for a handful of detected objects, ranging from giant spirals to inchoate low surface brightness galaxies, which represent important \\hie-selected class types. A montage of the optical counterparts illustrates their very wide variety. Finally the complete data set is laid out in the form of tables.  (5) \\textit{The Correlations in Galaxy Properties} reveals the 5 separate correlations among their properties.  (6) \\textit{The Discussion} argues that the correlations, despite the small size of the sample, are significant, not unduly affected by distance uncertainties, and compatible with the Tully-Fisher relation: Line-width $\\sim$ L$^{\\alpha}$, provided $\\alpha$ $\\simeq$ 1/3. If no more than 6 physical invariants control galaxies, as we argue, then 5 correlations suggest a degree of organisation among gas-rich galaxies which is surprising, and which must strongly constrain theories as to their formation and evolution. ",
        "conclusions": "We set out to look for correlations within a sample of galaxies free of optical selection effects. Apart from inclination and velocity we measured 9 parameters for each galaxy. Since the 4 colours appear degenerate we are left with 6 $\\textit{a priori}$ independent observables, between which we find the following 5 correlations, 4 old and 1 new (A) which are shown in Table \\ref{correlations}, where the columns are as follows.\\\\ \\indent \\textit{Column (1)}.-- Correlations observed.\\\\ \\indent \\textit{Columns (2) and (3)}.-- Denote the distance dependance, d$^{x}$, of the left and right-handed parameters in the correlation.\\\\ \\indent \\textit{Column (4)}.-- Indicates the difference between the two exponents in Cols. (2) and (3) (see text for explanation).\\\\ \\indent \\textit{Column (5)}.-- Shows the $\\textit{observables}$ introduced into the correlation, that have not been used before.  \\begin{table} \\caption{Correlations in the ES sample.} \\label{correlations} \\begin{center} \\begin{tabular}{lccccl} \\hline & \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} \\\\ \\hline \\multicolumn{1}{c}{A} & \\multicolumn{1}{l}{SB/Lum} & \\multicolumn{1}{c}{d$^{0}$} & \\multicolumn{1}{c}{d$^{2}$} & \\multicolumn{1}{c}{2} & \\multicolumn{1}{l}{SB, $\\theta_{90}$, d}\\\\ \\multicolumn{1}{c}{B} & \\multicolumn{1}{l}{Col/Lum} & \\multicolumn{1}{c}{d$^{0}$} & \\multicolumn{1}{c}{d$^{2}$} & \\multicolumn{1}{c}{2} & \\multicolumn{1}{l}{$(g - r)$}\\\\ \\multicolumn{1}{c}{C} & \\multicolumn{1}{l}{M$_{dyn}$/L$_{H}$} & \\multicolumn{1}{c}{d$^{1}$} & \\multicolumn{1}{c}{d$^{2}$} & \\multicolumn{1}{c}{1} & \\multicolumn{1}{l}{$\\Delta V$}\\\\ \\multicolumn{1}{c}{D} & \\multicolumn{1}{l}{R$_{90}$/R$_{50}$} & \\multicolumn{1}{c}{d$^{1}$} & \\multicolumn{1}{c}{d$^{1}$} & \\multicolumn{1}{c}{0} & \\multicolumn{1}{l}{$\\theta_{50}$} \\\\ \\multicolumn{1}{c}{E} & \\multicolumn{1}{l}{M$_{H\\!I}$/R$_{90}^{2}$} & \\multicolumn{1}{c}{d$^{2}$} & \\multicolumn{1}{c}{d$^{2}$} & \\multicolumn{1}{c}{0} & \\multicolumn{1}{l}{S~(21cm)}\\\\ \\hline \\end{tabular} \\end{center} \\end{table}   Column (5) in Table \\ref{correlations} is pertinent to the $\\textit{independence}$ of the 5 correlations. Unless $\\textit{a priori}$ reasons can be advanced for a dependance of the new observables on previous ones, then the 5 correlations must be independent. It might, for instance, be argued that $\\theta_{50}$ should depend on $\\theta_{90}$, but that ignores the real differences in `concentration' q $\\equiv$ $\\theta_{90}$/$\\theta_{50}$ between galaxies of different morphological type (e.g. q = 2.3 for pure exponential but q = 4.2 for de Vaucouleur profiles). Concentration is a crude but quantitative proxy for morphological type, and there is no $\\textit{a priori}$ case for this to be the same among HI-sources. Thus we can argue that the 5 correlations are truly independent.  Although the number of galaxies is modest the dynamic range of most of their properties is large (i.e. more than 7 magnitudes in luminosity, area, M$_{H\\!I}$ and M$_{dyn}$) whilst the correlations are, in a statistical sense, highly significant. So far as the $\\textit{existence}$ of the correlations (our main interest here) is concerned the sample-size is not an embarrassment, though larger samples will be needed to pin down the $\\textit{slopes}$ of the correlations more precisely. There is moreover a great diversity of galaxy types in the sample, ranging between early-type giant spirals and tiny `Inchoates' of such low surface brightness that they would usually be, and indeed were, missed in optical surveys. The difficulty of making a reliable census of galaxies in light alone is illustrated by Fig. \\ref{SBlum} where it can be seen that at all luminosities galaxies can vary in SB over 3 or 4 magnitudes, making the dimmest of them, even the giants, hard to pick out from the background by optical means alone. The narrow SB range of earlier catalogues (Freeman 1970, Disney 1976) was almost certainly a selection artefact.  Two correlations in particular need explaining. Correlation (E), i.e. M$_{H\\!I}$ $\\sim$ R$^{2}_{g}$ could be related to the stability of a gas layer to gravitational collapse, and hence to a trigger for star formation (e.g. Kennicutt 1989, Schaye 2008). The other, the SB/Luminosity correlation (A) has considerable scatter (Fig. \\ref{SBlum}) making it difficult to pin down an accurate value for $\\beta$ in $\\Sigma \\sim L^{\\beta}$, where $\\beta$ is not far from 1/2. Combining our data with the GOLDmine sample, in order to increase sample-size and hopefully reduce scatter, yields $\\Sigma \\sim L^{1/3}$, or $\\Sigma \\sim R$, or $L_{H}/R_{g}^{3}$ = constant or, considering Gavazzi's Law, M$_{dyn}$/R$_{g}^{3}$ = constant ($\\sim$ 1 H atom cm$^{-3}$ within the optical radius) which, see Fig. \\ref{21b}, is a pretty good fit. So the mass - and luminosity - density of galaxies is more or less independent of their Luminosity. This is at least a better mnemonic for recalling the rather scattered SB/Lum law. If true the Virial Theorem implies that all HI galaxies spin at the same angular velocity and rotate $\\sim$ 40 times in a Hubble epoch.  What about the Tully-Fisher correlation $\\Delta V \\sim L^{\\alpha}$? Gavazzi's Law (C): $L_{H} \\sim M_{dyn} \\equiv R_{g} \\Delta V^{2}$ while above: $R_{g} \\sim L_{H}^{1/3}$, hence $\\Delta V \\sim L_{H}^{1/3}$. In other words our data are consistent with the TF law provided $\\alpha$ = 1/3 [see also Fig. \\ref{TFg}]. The problem with using TF directly is demonstrated in Table \\ref{ranges} where it can be seen that the low dynamic range in $\\Delta V$ will always make it difficult to find the exponent $\\alpha$ accurately, and it makes more sense to incorporate $\\Delta V$ into the Dynamical Mass ($\\equiv$ R $\\Delta V^{2}$) instead, as it is in Gavazzi's Law (C) where it leads to both a clear correlation with Luminosity (Fig. \\ref{gavHGM}) and perhaps a natural physical explanation [`all galaxies are made of the same stuff'].  The distance to most galaxies in our radial velocity range can be problematic so it is important to ask if and how distance uncertainties will affect the correlations. Where there is a difference (last column) between the distance-dependant exponents $x$ in our list one expects distance uncertainties to introduce extra scattering to correlations, as for instance in (A) but not in (E). It is interesting to note that the tightest correlations observed i.e. (D) and (E) are the ones unaffected by distance uncertainties, whilst the loosest, i.e. (A) [Fig. \\ref{colourmag}] and (B) [Fig. \\ref{SBlum}] should indeed be the strongest affected. This hints that distance uncertainties probably are increasing the scatter. On the other hand look at Gavazzi's relation (B) and in particular Figures \\ref{dynlum} and \\ref{gavHGM}. There is no increase in scatter when we come to lower radial velocities where one might expect the relative distance-uncertainties to have the most pernicious effect on the correlations. This suggests that although distance uncertainty can increase scatter, that increase must be rather limited here, and exculpates the uniform but rather crude distances we have adopted. Note also that only 1 observable is size, and hence distance dependant (R $\\equiv d \\theta_{90}$) while 5 are not, i.e. $\\Sigma_{opt}, \\nhi, (g - r), \\Delta V$ and Concentration $R_{90}/R_{50}$. This implies strong correlations among $\\textit{intrinsic}$ galaxy properties (e.g. colour and SB, $\\Delta V$ and \\nhi) i.e. ones that have nothing to do with either scale or distance.  How constraining could the five discussed correlations (A) to (E) be on theories of galaxy formation and evolution? That depends on how many fundamental, independent invariant physical properties galaxies might posses. At present it is hard to imagine more than 7: total mass, baryon fraction, age, angular momentum, random energy, radius and central condensation. Once a galaxy has virialised, the Virial Theorem will provide one relation between them, leaving only 6 independent properties. Ignoring interactions one might expect, to first order, that these properties might be conserved by each galaxy though, on a secular time-scale some random energy might be dissipated, while some baryons might be ejected. Thus five correlations (if they are truly independent) within a six-parameter set hint at a high degree of organisation among gas-rich galaxies, something partially foreshadowed by the early \\hi pioneers [Brosche 1973, Balkowski 1973]. We are presently carrying out a Principal Component Analysis to investigate the degree of organisation in this data.  On the observational front work in progress by us includes enlarging the sample size as more SDSS and \\hi data [e.g. from the Bonn 100 metre Multibeam Survey (Kerp et al. in preparation) becomes available. Our automated SDSS photometric pipeline for large galaxies should greatly accelerate progress (West et al. 2008). \\hi interferometry is being obtained, first to sort out identifications in tight groups and second to measure more actual \\hi radii to get a better understanding of the Minchin-Rosenberg effect. So far we have exploited neither the colours nor the fibre spectra but work is in hand to include them too (West et al. in preparation). Statistical work on mean cosmic properties, such as the mean \\hi cosmic density,  depends on a clear understanding of ones selection effects and this is underway too (Garcia-Appadoo et al. in preparation). We are also working on the reverse List, that is to say a list of objects in the SDSS survey which we would have expected to contain measurable \\hie, and thus show up in the Equatorial Survey, but which do not (West \\& Garcia-Appadoo in preparation)."
    },
    "0809/0809.3788_arXiv.txt": {
        "abstract": "Ultra high-energy cosmic rays (UHECRs) are believed to be protons accelerated in magnetized plasma outflows of extra-Galactic sources. The acceleration of protons to $\\sim10^{20}$~eV requires a source power $L>10^{47}~{\\rm erg~s^{-1}}$.  The absence of steady sources of sufficient power within the GZK horizon of 100~Mpc, implies that UHECR sources are transient.  We show that UHECR \"flares\" should be accompanied by strong X-ray and $\\gamma$-ray emission, and that X-ray and $\\gamma$-ray surveys constrain flares which last less than a decade to satisfy at least one of the following conditions: {\\it (i)} $L>10^{50}{\\rm erg~s^{-1}}$; {\\it (ii)} the power carried by accelerated electrons is lower by a factor $>10^2$ than the power carried by magnetic fields or by $>10^3$ than the power in accelerated protons; or {\\it (iii)} the sources exist only at low redshifts, $z\\ll1$. The implausibility of requirements {\\it (ii)} and {\\it (iii)} argue in favor of transient sources with $L>10^{50}{\\rm erg~s^{-1}}$. ",
        "introduction": "\\label{sec:intro}  The origin of ultra high-energy ($>10^{19}$~eV) cosmic-rays (UHECRs) remains a mystery \\citep{Bhattacharjee:1998qc,Nagano:2000ve}.  The sources have not been robustly identified, and the models of particle acceleration are challenged by the fact that the energy spectrum extends to $>10^{20}$~eV. Several observational clues suggest that the UHECR flux is dominated by extra-Galactic light nuclei: the spectrum flattens at $\\sim10^{19}$~eV \\citep{Nagano:2000ve}, there is evidence for a composition change from heavy to light nuclei at $\\sim10^{19}$~eV \\citep{Bird93,Abbasi05}, and the UHECR arrival direction distribution is nearly isotropic \\citep{Finley04,HiResIso}. The recent detection of a weak anisotropy in the arrival distribution of $>6\\times10^{19}$~eV cosmic-rays \\citep{auger}, is consistent with that predicted by assuming that the spatial distribution of UHECR sources correlates with the large-scale distribution of galaxies \\citep{Fisher97,Kashti08}.  Although the identity of the UHECR particles is uncertain, we will assume here that they are protons. This assumption is motivated by two arguments. First, the observed spectrum of cosmic-rays with energies $>10^{19}$~eV is consistent with a cosmological distribution of proton accelerators producing (intrinsically) a power-law spectrum of high energy protons with $d\\log N/d\\log E\\approx -2$, for the number $N$ as a function of energy $E$ \\citep{Waxman95,Bahcall03,Kashti08}.  This intrinsic power-law spectrum is consistent with that expected in models of particle acceleration in collisionless shocks, for both non-relativistic~\\citep{Blandford87} and relativistic shocks (Waxman 2006; see however Keshet 2006).  Second, the leading candidates for extra-Galactic sources, namely gamma-ray bursts and active galactic nuclei (see below), are expected to accelerate primarily protons.  Robust model-independent considerations imply that UHECR protons can only be produced by sources with an exceedingly high power output \\citep{Waxman_CR_rev}, $L>\\Gamma^2\\beta^{-1}10^{46}~{\\rm erg~s^{-1}}$, where $\\Gamma$ and $\\beta c$ are the Lorentz factor and characteristic velocity associated with plasma motions within the source\\footnote{Somewhat more stringent limits may be obtained by specifying the acceleration process; see \\citet{Norman95}, \\citet{Waxman95prl}, and \\S~\\ref{sec:rates} below.}. Since no steady source above this power threshold is known to exist within the 100~Mpc {\\it GZK horizon}, the distance to which the propagation of $\\sim10^{20}$~eV protons is limited by their interaction with the cosmic microwave background \\citep{G66,ZK}, the UHECR sources are most likely transient. A possible alternative is, of course, an unknown class of \"dark sources\", which produce little radiation and therefore remain undetectable by telescopes.  Only two types of sources are known to satisfy the above minimum power requirement: active galactic nuclei (AGN) -- the brightest known steady sources, and gamma-ray bursts (GRBs) -- the brightest known transient sources\\footnote{It was recognized early on (\\cite{Hillas} and references therein) that while highly magnetized neutron stars may also satisfy the minimum power requirement, it is difficult to utilize the potential drop in their electro-magnetic winds for proton acceleration to ultra-high energy (see, however, \\cite{Arons}).}. The absence of AGN with $L>10^{46}~{\\rm erg~s^{-1}}$ within the GZK horizon had motivated \\citet{Gruzinov} to suggest that UHECRs may be produced by a new, yet undetected, class of short duration AGN flares resulting from the tidal disruption of stars or accretion disk instabilities.  We show in \\S~\\ref{sec:flares} that if electrons are accelerated together with the protons in UHECR-producing flares, then their radiative losses will produce a bright flare of X-ray and $\\gamma$-ray photons. We then show in \\S~\\ref{sec:LFs} that existing X-ray and $\\gamma$-ray surveys already put stringent constraints on the properties of UHECR flares. In \\S~\\ref{sec:hidden} we discuss the possibility of \"hiding\" the X-ray emission. Our conclusions are summarized in \\S~\\ref{sec:conclusions}.  Throughout our discussion, we consider a scenario in which the flare is associated with ejection of magnetized plasma from the source, and where the charged particles are accelerated within the magnetized outflow. The non thermal emission from a wide range of sources is described within the framework of such a scenario. This includes AGN jets and GRBs, as well as the transient AGN flares proposed by \\citet{Gruzinov}. We parametrize the UHECR flares by their power, $L$, duration, $\\Delta t$, characteristic ejection speed $\\beta c$, and rate per unit volume in the local Universe, $\\dot{n}$. The fractions of the total energy output carried by protons, electrons and magnetic fields are denoted by $\\epsilon_p$, $\\epsilon_e$ and $\\epsilon_B$, respectively, and we assume that the energy spectrum of accelerated electrons is similar to that of accelerated protons with a power-law index, $d\\log N/d\\log E\\approx -2$.  This index is expected for astrophysical sources which accelerate particles in strong collisionless shocks \\citep{Blandford87,Waxman_CLS_rev}, and is inferred from the radiation observed from a variety of high energy sources, such as supernova remnants \\citep{SN} and GRBs \\citep{WAG97a}.  It should be pointed out that the composition of the jets of high energy sources is unknown. Two classes of models are generally discussed: jets where the energy flux is dominated by the plasma kinetic energy, and jets where most of the energy is carried by electromagnetic flux. For the \"kinetic\" jets, a plausible mechanism exists for energy dissipation, particle acceleration and radiation emission, namely internal collisionless shocks within the outflow. Within this mechanism, a particle distribution following $d\\log N/d\\log E\\approx -2$ is naturally expected. For the \"electromagnetic\" jets, the mechanism of energy dissipation and particle acceleration is not well understood. We will assume that a $d\\log N/d\\log E\\approx -2$ particle distribution is generated in this model too, as suggested by observations.  As explained at the end of \\S~\\ref{sec:rates}, our conclusions are valid for both spherical and jetted flows.  Throughout the paper, $L$ stands for the \"isotropic-equivalent power\" (i.e. for a flow which is conical rather than spherical, $L$ stands for the power that would have been carried by the flow had it been spherically symmetric), and $\\dot{n}$ stands for the \"isotropic-equivalent rate density\" (i.e. the rate inferred under the assumption of spherically symmetric emission). ",
        "conclusions": "\\label{sec:conclusions}  The absence of steady sources of sufficient power to accelerate UHECRs within the GZK horizon of 100~Mpc, implies that UHECR sources are transient.  We have shown that UHECR \"flares\" should be accompanied by strong X-ray and $\\gamma$-ray emission. Figure 1 demonstrates that X-ray and $\\gamma$-ray surveys constrain flares which last longer than $\\sim5$~min and less than a decade to satisfy at least one of the following conditions: {\\it (i)} $L>10^{50}{\\rm erg~s^{-1}}$; {\\it (ii)} the power carried by accelerated electrons is lower by a factor $>10^2$ than the magnetic field power or by $>10^3$ than that carried by accelerated protons; or {\\it (iii)} the sources exist only at low redshifts, $z\\ll1$. The implausibility of requirements {\\it (ii)} and {\\it (iii)} argue in favor of transient sources with $L>10^{50}{\\rm erg~s^{-1}}$.  The required luminosity is well above the brightest luminosity ever recorded in an AGN flare, and exceeds by two orders of magnitude the Eddington limit for a black hole of $10^{10}M_\\odot$, the highest mass expected to exist within a distance of 100Mpc \\citep{Lauer,Natarajan}. The results shown in Fig. 1 exclude the regime of low flare luminosities considered by Farrar \\& Gruzinov (2008).  The lower bound of $\\sim 300$s on the window of flare durations over which our constraints apply, originates from the integration time of the X-ray data used in Fig. 1.  For flare durations $\\Delta t\\lesssim 300$s, Eq. (\\ref{eq:L_Bmin1}) requires $\\epsilon_BL>0.3\\times 10^{50}~{\\rm erg~s^{-1}}$, not significantly different from the minimum luminosity inferred for longer flare durations.  We have also explored potential caveats to the above conclusions. Long-duration ($\\gtrsim 100$ years) flares which are electromagnetically dominated could evade the constraints illustrated in Fig. 1 if their gamma-ray emission peaks outside the EGRET energy band. In addition, an unknown population of ``electromagnetically-dark'' flares is in principle possible (see \\S \\ref{sec:hidden} for details), although there is no physical motivation to make its existence natural.  Future gamma-ray observations with {\\it GLAST} and X-ray observations with {\\it EXIST} would improve the statistical constraints on the source population of UHECRs and potentially shed more light on their nature."
    },
    "0809/0809.1058_arXiv.txt": {
        "abstract": "We study the spectral energy distributions and evolution of a large sample of optically selected quasars from the Sloan Digital Sky Survey (SDSS) that were observed in 323 \\Chandra\\, images analyzed by the \\Chandra\\, Multiwavelength Project. Our highest-confidence matched sample includes 1135 X-ray detected quasars in the redshift range $0.2<z<5.4$, representing some 36\\,Msec of effective exposure.  We provide catalogs of QSO properties, and describe our novel method of calculating X-ray flux upper limits and effective sky coverage. Spectroscopic redshifts are available for about 1/3 of the detected sample; elsewhere, redshifts are estimated photometrically. We detect 56 QSOs with redshift $z>3$, substantially expanding the known sample. We find no evidence for evolution out to $z\\sim$5 for either the X-ray photon index $\\Gamma$ or for the ratio of optical/UV to X-ray flux \\aox.  About 10\\% of detected QSOs show best-fit intrinsic absorbing columns $>10^{22}\\cmsq$, but the fraction might reach $\\sim$1/3 if most non-detections are absorbed. We confirm a significant correlation between \\aox\\, and optical luminosity, but it flattens or disappears for fainter ($M_B\\gax -23$) AGN alone.  We report significant hardening of $\\Gamma$ both towards higher X-ray luminosity, and for relatively X-ray loud quasars. These trends may represent a relative increase in non-thermal X-ray emission, and our findings thereby strengthen analogies between Galactic black hole binaries and AGN. For uniformly-selected subsamples of narrow line Seyfert~1s and narrow absorption line QSOs, we find no evidence for unusual distributions of either \\aox\\, or $\\Gamma$. ",
        "introduction": "\\label{intro}  Interest in the properties of active galaxies and their evolution has recently intensified because of deep connections being revealed between supermassive black holes (SMBH) and galaxy evolution, such as the relationship between the mass of galaxy spheroids and the SMBH they host (the $M_{BH}-\\sigma$ connection; \\citealt{Ferrarese00, Gebhardt00}). A feedback paradigm could account for this correlation, whereby winds from active galactic nuclei (AGN) moderate the SMBH growth by truncating that of their host galaxies (e.g., \\citealt{Granato04}).  Feedback models may explain the correspondence between the local mass density of SMBH and the luminosity density produced by high redshift quasars \\citep{Yu02,Hopkins06b} as well as the `cosmic downsizing' (decrease in the space density of luminous AGN) seen in AGN luminosity functions \\citep{Barger05,Hasinger05,Scannapieco05}. If quasar activity is induced by massive mergers (e.g., \\citealt{Wyithe02, Wyithe05}), then the jigsaw puzzle now assembling may merge smoothly with cosmological models of hierarchical structure formation.  Many, if not most, of the accreting SMBH in the Universe may be obscured by gas and dust in the circumnuclear region, or in the extended host galaxy.  The obscured fraction may depend on both luminosity and redshift \\citep{Ueda03, Brandt05, LaFranca05}, and is indeed likely to evolve on grounds both theoretical (e.g., \\citealt{Hopkins06a}) and observational \\citep{Treister06,Ballantyne06}.  Such evolution seems to be required for AGN populations to compose the observed spectrum of the cosmic X-ray background (CXRB; \\citealt{Gilli07}).  However, a full census of all SMBH remains observationally challenging, since some are heavily obscured, or accreting at very low rates, below the sensitivity limits of current telescopes even at low-redshifts.  AGN unification models explain many of the observed differences in the spectral energy distributions (SEDs) of AGN as being due to the line-of-sight effects of anisotropic distributions of obscuring material near the SMBH \\citep{Antonucci93}. The intrinsic number ratio of obscured to unobscured AGN may evolve, and is almost certainly a function of luminosity. Indeed, the ratio in the Seyfert (low-) luminosity regime is currently estimated to be $\\sim$4, whereas for the QSO (high-) luminosity regime, it may be closer to unity \\citep{Gilli07}.  Astronomers, like most people, usually look where they can see. Type~1 quasars are the easiest AGN to find in large numbers via either spectroscopic or color selection because of their broad emission lines (FWHM$\\gax$1000\\,\\kms) and generally blue continuum slopes and brighter magnitudes.  Most of these show few other signs of obscuration such as infrared excess or weakened X-ray emission.  Large samples of Type~1 QSOs - as the brightest high redshift objects - have served to probe intervening galaxies, clusters, and the intergalactic material (IGM) along the line of sight, right to the epoch of reionization.  Because observed SEDs are thought to be less affected by obscuration and therefore more representative of intrinsic accretion physics, the evolution of this Type~1 sample is of interest as well.  The SEDs and clustering properties of Type~1 QSOs have been studied in increasing detail, probing farther into the Universe and wider across the sky and the electromagnetic spectrum.  SEDs including mid-infrared photometry from {\\em Spitzer} were compiled and characterized recently for 259 quasars from the Sloan Digital Sky Survey (SDSS) by \\citet{Richards06b}. These studies are most useful for calculating bolometric luminosities and $K$-corrections towards understanding the energetics of the accretion process, and its evolution across cosmic time. Studies that target SDSS QSO pairs with small separations find a significant clustering excess on small scales ($\\lax 40$\\,kpc$h$) of varying strengths (e.g., \\citealt{Hennawi06,Myers07,Myers08}), which could be due to mutual triggering, or might simply result from the locally overdense environments in which quasars form \\citep{Hopkins08}.  Optical luminosity function studies of optically selected quasars (e.g., \\citealt{Richards06a,Richards05,Croom04}) date back decades  (e.g., \\citealt{Boyle88}). Accurate luminosity functions are needed to trace the accretion history of SMBH and to contrast the buildup of SMBH with the growth of galaxy spheroids.  An increasing number of optical surveys not only select AGN photometrically, but also determine fairly reliable photometric  redshifts for them. These samples stand to vastly improve the available statistical reliability and the resolution available in the  luminosity/redshift plane.  X-ray observations have been found to efficiently select AGN of many varieties, and at higher surface densities than ever \\citep{Hasinger05}.  Independent of the classical AGN optical emission line criteria, X-rays are a primary signature of accretion onto a massive compact object, and the observed X-ray bandpass corresponds at higher redshifts to restframe energies capable of penetrating larger intrinsic columns of gas and dust.\\footnote{The observed-frame, effective absorbing column is $N_{\\mathrm H}^{\\mathrm eff}\\sim N_{\\mathrm  H}/(1+z)^{2.6}$ \\citep{Wilman99}.}  Even for Type~1 QSOs, X-ray observations have revealed new connections (e.g., between black holes from 10 to 10$^9$ $M_{\\odot}$; \\citealt{Maccarone03, Falcke04}) and new physical insights such as the possibility of ubiquitous powerful relativistic outflows \\citep{Middleton07} or of relativistically broadened fluorescent Fe\\,K$\\alpha$ emission (see, e.g., the review articles by \\citealt{Fabian00, Reynolds03}).  Quasars with broad absorption  lines (BALQSOs) blueward of their UV emission lines turn out to be highly absorbed in X-rays \\citep{Green95, Green96,Green01, Gallagher02}. Quasars that are radio loud are 2 -- 3 times brighter in X-rays for the same optical magnitude \\citep{Zamorani81,Worrall87,Shen06}, and also may have harder X-ray spectra (e.g., \\citealt{Shastri93}, although see \\citealt{Galbiati05}).  The high sensitivity and spatial resolution of \\Chandra\\, and \\XMM\\, open other avenues for exploration of quasars and their environments. Clusters of galaxies have been discovered in the vicinity of, or along the sightlines to quasars \\citep{Green05,Siemiginowska05}.  Lensed quasars have now been spatially resolved in X-rays, unexpectedly showing significantly different flux ratios than at other wavebands \\citep{Green02,Blackburne06,Lamer06}.  The expected evolution in the environment, accretion rates, and masses of SMBH in AGN should correspond to observable evolution in their SEDs. The two most common X-ray measurements used are the X-ray power-law photon index $\\Gamma$\\footnote{$\\Gamma$ is the photon number index  of an assumed power-law continuum such that $N_E(E)=N_{E_0}\\,E^{\\Gamma}$. In terms of a spectral index $\\alpha$ from $f_{\\nu}=f_{\\nu_0}\\nu^{\\alpha}$, we define $\\Gamma= (1-\\alpha)$.} and the X-ray-to-optical spectral slope, \\aox.\\footnote{$\\aox$\\, is the slope of a hypothetical power-law from 2500\\,\\AA\\, to 2~keV; $\\aox\\, = 0.3838~{\\mathrm log} (\\opteml/\\xeml)$}. Many of the apparent correlations have been challenged as being artifacts of selection or the by-products of small, heterogeneous samples, which impedes progress in our understanding of quasar physics and evolution.  The archives of current X-ray imaging observatories like \\Chandra\\, and XMM-{\\em Newton} are growing rapidly. and several large efforts for pipeline processing, source characterization \\citep{Ptak03,Aldcroft06}, and catalog generation are underway.  The 2XMM catalogue \\citep{Page06} is available, and the \\Chandra\\, source catalog\\footnote{The \\Chandra\\, source catalog webpage is {\\tt http://cxc.harvard.edu/csc}}. is due out in 2008 \\citep{Fabbiano07}.  Serendipitous wide-area surveys with \\Chandra\\, were pioneered by the \\Chandra\\, Multiwavelength Project (ChaMP; \\citealt{Green04,DKim04a}) for high Galactic latitude, while ChaMPlane \\citep{Grindlay05} has studied the stellar content of Galactic plane fields. The ChaMP, described in more detail below, also performs multiwavelength source matching and spectroscopic characterization. These efforts are greatly augmented by large surveys such as the Sloan Digital Sky Survey  (SDSS; \\citealt{York00}), which will obtain spectroscopy of $\\sim100,000$ QSOs (e.g., \\citealt{Schneider07}), and can additionally facilitate the efficient extrapolation of photometric quasar selection with photometric redshift estimation almost a magnitude fainter, towards a million QSOs (e.g., \\citealt{Richards08}).  The current paper studies the X-ray and optical properties of the subset of these QSOs imaged in X-rays by \\Chandra\\, as part of the \\Chandra\\, Multiwavelength Project \\citep{Green04}. ",
        "conclusions": "In this study, we have presented the largest homogeneous study to date of optically selected broad line quasars (from the SDSS) with sensitive X-ray flux limits (from \\Chandra; mode 2$\\times 10^{-14}$\\fcgs; 0.5-8\\,keV). Our large sample highlights the large dispersion in quasar properties that is unveiled whenever sensitive limits and wide sky areas are combined.  We confirm and extend several well-known multiwavelength relations. The relationship between \\aox\\, and 2500\\AA\\, luminosity is confirmed for high luminosities ($M_B\\lax -23$, or ${\\mathrm log}\\,(l_{\\mathrm 2500\\AA})\\gax$29.8 (log\\,($\\nu_{\\mathrm 2500}l_{\\mathrm 2500\\AA})\\gax $44.9) with slopes similar to those found previously (e.g., S06).  Including a wider luminosity range inevitably produces a flatter relation across a range of subsamples which test the effects of excluding \\Chandra\\, PI targets, RL QSOs, and QSOs with BALs, or NALs as well as NLS1s. No significant $\\aox$($\\opteml$) correlation exists for objects (68\\% detected) at lower luminosities. Another possibility is that the relation simply flattens at low luminosities, or has a higher order luminosity dependence (e.g., \\citealt{Kelly07a}).  We find for the first time a significant, robust and rather steep dependence of X-ray continuum slope $\\Gamma$ on X-ray luminosity \\xeml\\, in the sense that the spectrum hardens with increasing \\xeml. A trend in the opposite sense has been reported recently for AGN in the \\Chandra\\, Deep Fields \\citep{Saez08}, but may be dominated by differences between Type~1 and Type~2 AGN.  We also report a shallow but significant trend that $\\Gamma$ is harder for relatively X-ray bright (low \\aox) QSOs. We note that X-ray bright QSOs include the RL QSOs, and that RL QSOs also have predominantly flatter (harder) $\\Gamma$. We thus speculate that the overall trends of $\\Gamma$(\\aox) and $\\Gamma$(\\xeml) both reflect an increase in the non-thermal emission fraction toward higher X-ray luminosities.  Not all QSOs with a strong non-thermal X-ray emission component are necessarily radio loud - detectable radio loudness may pertain only to a fraction of these objects.  Radio bright phases may be short compared to the overall QSO lifetime and/or episodic.  The black hole masses in AGN are 5--8 orders of magnitude larger than those in Galactic black hole X-ray binaries (GBH).  Since for a given $L/L_{Edd}$ the disc temperature scales with mass as $M^{-1/4}$, AGN disks are cooler than in GBH.  The thermal accretion disk emission component that dominates the soft X-ray emission ($\\sim$2\\,keV) in GBH corresponds to optical/UV ($\\sim$2500\\AA) emission in AGN.   Similarly, the non-thermal emission (probably from Comptonized emission from the disk's corona) sampled as X-ray emission in AGN comes from harder ($\\sim$20\\,keV) X-rays in GBH.  Type~1 QSOs may be analogous to Galactic black hole binaries (GBH) in the high/soft state \\citep{Sobolewska08} where a similar trend is seen in $\\Gamma$ vs. a disc/Comptonization index $\\alpha_{GBH}$ (analogous to \\aox\\, for QSOs) as we report here. \\cite{Jester06} also compare the disk vs. non-thermal emission fraction of GBH and AGN, and find that AGN segregate by radio loudness similarly to GBH in regions where luminosity and/or non-thermal fraction are high.  While the correlations we find lend significant support to these interpretations, the scatter is large, and could be significantly reduced if extrinsic effects can be identified and corrected for.  To better understand the intrinsic physics of accretion requires identifying and quantifying extrinsic effects such as absorption and non-thermal processes. Absorption may occur close to the SMBH gravitational radius, in the BLR, a molecular torus, surrounding starforming regions, the outer host galaxy or at intervening redshifts.  The intrinsic absorption may be orientation-dependent, may evolve with redshift, and may be a function of luminosity as well.  Contributions from non-thermal processes certainly play a role, whether or not it is reflected in detectable radio emission, and that role may change with SMBH mass, spin, and environment, so consequently with lookback time as well.  Let's face it - quasars are complicated, but on average they don't change much.  If quasar SED changes with luminosity were large compared to the dispersion in the population and relatively immune to selection effects, we would have been using them for standard candles a long time ago.  Large samples, uniformly observed and analyzed, offer the greatest hope to disentangle the intertwining mysteries."
    },
    "0809/0809.2484_arXiv.txt": {
        "abstract": "We present a detailed multi-wavelength analysis and interpretation of the evolution of an M7.6 flare that occurred near the south-east limb on October 24, 2003. Preflare images at TRACE 195~{\\AA} show that the bright and complex system of coronal loops already existed at the flaring site. The X-ray observations of the flare taken from the RHESSI spacecraft reveal two phases of the flare evolution. The first phase is characterized by the altitude decrease of the X-ray looptop (LT) source for  $\\sim$11 minutes. Such a long duration of the descending LT source motion is reported for the first time. The EUV loops, located below the X-ray LT source, also undergo contraction with similar speed ($\\sim$15~km~s$^{-1}$) in this interval. During the second phase the two distinct hard X-ray footpoints (FP) sources are observed which correlate well with UV and H$\\alpha$ flare ribbons. The X-ray LT source now exhibits upward motion as anticipated from the standard flare model. The RHESSI spectra during the first phase are soft and indicative of hot thermal emission from flaring loops with temperatures ~$T>25$~MK at the early stage. On the other hand, the spectra at high energies ($\\varepsilon \\gtrsim$25~keV) follow hard power laws during the second phase ($\\gamma = 2.6-2.8$). We show that the observed motion of the LT and FP sources can be understood as a consequence of three-dimensional magnetic reconnection at a separator in the corona. During the first phase of the flare, the reconnection releases an excess of magnetic energy related to the magnetic tensions generated before a flare by the shear flows in the photosphere. The relaxation of the associated magnetic shear in the corona by the reconnection process explains the descending motion of the LT source. During the second phase, the ordinary reconnection process dominates describing the energy release in terms of the standard model of large eruptive flares with increasing FP separation and upward motion of the LT source. ",
        "introduction": "The active Sun displays a stunning variety of transient energetic phenomena ranging from the smallest microjets and microflares to the largest flares and coronal mass ejections (CMEs). It is generally accepted that the energy released during flares and CMEs is stored in the corona prior to the event in the form of stressed or non-potential magnetic fields. We believe that this stored free energy in the coronal magnetic fields is released explosively through the process of magnetic reconnection. Multi-wavelength measurements of flares and CMEs made over the years have provided many valuable pieces of information about the response of coronal energy release in the different atmospheric layers of the Sun. During an eruptive flare, the whole magnetic configuration of coronal loops crossing an inversion line is disrupted and is followed by a newly rebuilt loop system from pre-flare to post-flare stages. The expansion of flare ribbons and growth of flare loop system are the clearest observational findings which support the flare models involving magnetic reconnection \\citep[for a review see][]{Priest02,Lin03}.  Yohkoh spacecraft significantly improved our overall understanding about the flare physics \\citep{Hudson04}. Observations made by Yohkoh confirmed the presence of hard X-ray sources at the footpoints of the loop system \\citep{Sakao94} and detected a new hard X-ray source above the apex of the hot flaring loop observed in soft X-rays \\citep{Masuda94,Somov05}. The above-the-looptop source is believed to be intimately connected to the primary energy release site and location of the electron acceleration in the corona, again supporting the reconnection models. RHESSI spacecraft with superior observing capabilities in X-ray energy levels, further refined our knowledge about the X-ray LT coronal sources \\citep[see review by][]{Krucker08}. It detected a downward motion of LT source during the initial phase of impulsive rise of hard X-ray flux \\citep{Sui03,Sui04,Liu04,Ji06,Veronig06,Joshi07}. We still do not have a plausible explanation of this phenomenon. RHESSI observations have also revealed the existence of double coronal sources \\citep{Sui03,Sui04,Sui05,Veronig06,Li07,Liu08}. The double coronal source has been interpreted as an indirect evidence for the formation and expansion of a large scale current sheet in the corona \\citep{Sui03,Sui04,Sui05}.  In this paper we present a comprehensive multi-wavelength analysis of an M7.6 flare on 2003 October 24. A detailed X-ray imaging and spectroscopic analysis of RHESSI observations exhibit two phases of flare evolution in view of the X-ray energy release process. A significant and unusually long downward motion of the LT source characterizes the first phase. During the whole second phase we observe footpoints sources and upward expansion of the LT source, as it is commonly observed in solar flares. The event was associated with a fast CME. We compare flare observations at different X-ray wavelengths, and also with other chromospheric and coronal images with the aim to study the consequences of the magnetic reconnection process at various atmospheric layers. In section 2, we describe the observational characteristics of the event. We interpret our observational findings in the final section of the paper. ",
        "conclusions": "\\subsection{The first phase} Preflare images of the active region taken by TRACE at 195~{\\AA} show that a bright and complex system of coronal loops already existed in the flaring region. The X-ray and EUV observations during the first phase of flare evolution reveal many important observational findings. The X-ray LT source appeared right at the flare onset and starts moving to lower altitudes. The LT source is characterized by unusually long duration of altitude decrease (from 02:24 to 02:35 UT). About 5 minutes after the appearance of X-ray LT source, an intense and diffuse brightening is observed at the top of TRACE~195~{\\AA} loops, located below RHESSI LT source. \\cite{Warren01} and \\cite{Gallagher02} reported very similar bright and diffuse emission from TRACE~195~{\\AA} loops associated with X-class flares, which was interpreted as being due to the Fe XXIV contribution of plasma at a temperature of 15-20 MK. \\cite{Phillips05} further confirmed that the Fe XXIV line emission is the dominant contributor to emission in the TRACE~195~{\\AA} channel for flare associated high temperature features (such as hot LT emission). These bright EUV loops showed a large scale contraction. The descending motion of X-ray LT source ceased at 02:35 UT while the EUV coronal loops continued to contract several minutes later, which we interpret as an effect of plasma cooling (from X-rays to EUV). The speeds of downward motion for EUV loops and X-ray LT source are comparable during the descending phase of X-ray LT source. Another noticeable feature in 25--50 keV X-ray images is the appearance of two bright FP sources in this phase for a short interval (02:30:00-02:31:00).  The downward motion of LT source in the early phase of flare evolution has been established by many RHESSI observations {\\citep{Krucker03,Sui03,Sui04,Liu04,Ji06,Veronig06,Joshi07,Liu08}. Some of the previously studied events also show energy dependent structure of the LT source during the phase of altitude decrease, i.e. the higher energy sources were located higher in the corona and showed a faster speed of downward motion than lower energy sources (Sui et al. 2004; Veronig et al. 2006; Liu et al. 2008). In our event, however, the difference of LT altitude between the sources observed at low and high energy bands is not significant during this stage. Further the LT altitude decrease is observed in RHESSI observations for $\\sim$11 minutes, which is significantly longer than reported previously. The contraction of flaring EUV coronal loops has been reported by \\cite{Li06}, \\cite{Ji07}, and \\cite{LiuR09}. \\cite {Li05} reported the shrinkage of flare loops from radio observations at 34 GHz.  Finally we note that the duration of downward motion of the LT source is marked by high plasma temperature. RHESSI observations indicate that the plasma temperature reached to high values ($\\sim$27 MK) at the early stage. The temperature determined from GOES channels is lower (maximum temperature $\\approx$ 20 MK) and shows gradual rise with the LT shrinkage. Large difference between temperature computed from RHESSI and GOES measurements could be due to different sensitivity and response of these two instruments. This is also consistent with the fact that the observed plasma is multi-temperature. These results suggest a possible connection between LT heating and loop contraction.  \\subsection{The second phase}  The flare evolution during the second phase shows impulsive hard X-ray emission at high energies. During this interval various observational signatures in H$\\alpha$, (E)UV and hard X-ray wavelengths can be well understood by the standard CSHKP \\citep{Carmichael64,Sturrock66,Hirayama74,Kopp76} model of eruptive flares. These observational features include: appearance of two parallel flare ribbons in TRACE 1600~{\\AA} and H$\\alpha$ images; two hard X-ray FPs with increasing separation, one lying on each flare ribbons; growth of X-ray LT height; formation of relatively cool loop system below X-ray LT source at UV and H$\\alpha$ wavelengths connecting the flare ribbons as well as post flare loop configuration observed in the later stages of gradual decline of X-ray intensity. RHESSI and GOES measurements suggest that there is no further increase in temperature during the second phase. The flat X-ray spectra at high energies ($\\gamma$ $<$ 3) obtained during the second phase indicate strong non-thermal emission, % dominated by energy released into acceleration of electrons.  The beginning of the second phase is marked by the appearance of a HXR FP source which lie on the northern UV flare ribbon. At the flare maximum, second FP source developed on the southern flare ribbon. However, the intensity of the southern FP increases continuously after the flare maximum and eventually both the FPs became of comparable brightness (Fig. \\ref{rhessi_second_ph}). The consistency between the timings of peaks in light curves and morphological evolution of FP sources can be understood in terms of thick-target model \\citep{Brown71,Hudson72,Syrovatskii72} in which the X-ray production at the FPs of the loop system takes place when high energy electrons, accelerated in the reconnection region, come along the guiding magnetic field lines and slam the denser transition region and chromospheric layers producing hard X-ray emission via nonthermal bremsstrahlung produced by fast electrons scattered off ions. The evolution of FP sources from asymmetric to symmetric brightness can be seen as an evidence for different injection conditions of high energy electrons along the two legs of the loop system \\citep{Siarkowski04} which could be mainly because of the highly complicated systems of loops. We find that the upward motion of 12--25~keV~LT source starts at the flare maximum. This upward growth indicates progression of magnetic reconnection in the higher coronal loops rooted successively further apart from the magnetic inversion line. The beginning of the ascending motion of X-ray LT source at the overall maximum of HXR time profile has been reported in several flares observed by RHESSI \\citep{Sui03,Sui04,Liu04,Veronig06,Liu08}. On the other hand the lower loops eventually cool down (and fade away) and begin to be visible in low energies in UV and H$\\alpha$ lines. The structure of flaring region seen in H$\\alpha$ filtergrams at this stage also reveal a loop system connecting two H$\\alpha$ flare ribbons and is very similar to the 1600~{\\AA} images. We find that intense emission is produced from the top of the H$\\alpha$ loops. The appearance of bright looptop source in H$\\alpha$ emission against the solar disk was also observed in the LT event studied in \\cite{Veronig06} and \\cite{Joshi07}. It is rather an unusual feature and considered as an evidence for high pressure, i.e. high density ($n \\gtrsim 10^{12}$~$cm^{-3}$) in the observed postflare loops \\citep{Heinzel87}.  \\subsection{Overall picture of the flare evolution}  In order to interpret our observational results, we have to recall some general properties of energy accumulation before a solar flare and energy release during a flare. On the basis of data obtained by the Hard X-ray Telescope {\\em HXT\\/} on the satellite {\\em Yohkoh\\/}, \\cite{Somov02,Somov03} suggested that the large-scale structure and dynamics of two-ribbon flares, as seen in hard X-rays (HXR), can be explained in terms of the three-dimensional reconnection at a separator in the corona. More specifically, they suggested that, before a large two ribbon flare the bases of magnetic field lines are moved by the large-scale photospheric flows of two types. First, the converging flows, i.e., the flows directed to the photospheric neutral line ($ PNL $), create the pre-flare current layers in the corona and provided an excess of magnetic energy sufficient to produce a flare. Second, the shear flows, which are parallel to the $ PNL $, increase the length of field lines in the corona and, therefore, produce an excess of energy too.  During a flare, both excesses of magnetic energy are released completely or partially. In this way, the model describes two kinds of apparent motions of the HXR kernels. One involves the increasing distance between the chromospheric flare ribbons (and associated HXR kernels) which results from reconnection in a coronal current layer. The other involves the approaching kernels and is associated with the relaxation of magnetic tensions generated by the photospheric shear flows before a flare. We call the latter process {\\em shear relaxation\\/} \\citep[for more details][] {Somov07}.  Let us clarify the general situation in terms of a specific model called ``rainbow reconnection'' \\citep{Somov86}. Figure~\\ref{mag_photo}a shows the simplest presentation of a bipolar distribution of the vertical component~$ B_z $ of magnetic field in the photosphere. The neutral line~$ PNL $ divides the region into two zones with field polarities~$ N $ and $ S $. This region may be deformed by horizontal photospheric flows with the velocity field~$ {\\bf V} $ in such a way that the $ PNL $ gradually acquires the shape of the letter $ S $ as shown in Figure~\\ref{mag_photo}b. Beginning at a critical value of the curvature of the $ PNL $ \\citep[see][]{Gorbachev88} in the magnetic field above the photosphere, calculated in the potential approximation, there appears a topologically featured field line, a separator, as illustrated by Figure~\\ref{rainbow}. The separator~$ X $ is located above the $ PNL $ like a rainbow above a river which makes a bend. The separator is the place where a reconnecting current layer can be created, and therefore the pre\\-flare energy accumulation can begin.  On the other hand, Figure~\\ref{mag_photo}c shows that a large-scale vortex-type flow generates two components of the velocity field in the photosphere: the velocity components~$ {\\bf V}_{\\|} $ and $ {\\bf V}_{\\bot} $ are parallel and perpendicular to the $ PNL $. The first component provides a shear of field lines (magnetic shear) above the $ PNL $. The shear flow before a flare creates the longer magnetic loops which, being reconnected mainly during the first phase of a flare, provide the bright HXR kernels with a large foot\\-point separation (Figure~\\ref{flare_cartoon}). The second component $ {\\bf V}_{\\bot} $ of the velocity field in the photosphere tends to drive reconnection in the corona. To demonstrate the basic physics in the simplest way, we consider only a central region~$ C $ in the vicinity of the $ S $-shaped neutral line~$ PNL $ in Figure~\\ref{mag_photo}b. Here we put the $ y $-direction along the $ PNL $; the separator is nearly parallel to $ PNL $ as was shown in Figure~\\ref{rainbow}.  Being reconnected at the separator, each magnetic field line as well as each tube of magnetic flux~$ f $ is initially accelerated to high velocity ($ \\stackrel{ > }{ _\\sim } 1000 $ km/s) inside a super-hot turbulent-current layer \\citep[SHTCL;][Chapter~6]{Somov07}. Each tube of reconnected field lines, being frozen into super-hot ($ T_{e} \\stackrel{ > }{ _\\sim } 10^8 $~K) plasma, moves out of the SHTCL; in the down\\-flow it forms a magnetic loop with properties of a collapsing magnetic trap \\citep[][Chapter~7]{Somov97,Somov07}. A collapsing magnetic trap model has been first applied to one of the RHESSI events showing a descending LT source by \\cite {Veronig06}. The longitudinal and transverse sizes of the trap decrease, causing the trapped particles to acquire an additional energy under Fermi and betatron acceleration respectively \\citep{Somov03a}. The energy distribution of trapped electrons and their HXR emission can be calculated as a function of the trap length and its thickness \\citep{Bogachev07}.  If the electrons injected into the trap from the SHTCL have a power-law energy distribution, then their spectrum remains a power-law one throughout the acceleration process for both the Fermi and betatron mechanisms. For electrons with a thermal injection spectrum, the model predicts two types of HXR coronal sources. (a) Thermal sources are formed in traps dominated by the betatron mechanism. (b) Non\\-thermal sources with a power-law spectrum appear when electrons are accelerated by the Fermi mechanism. With account of rare Coulomb collisions inside the trap, a double-power-low spectrum is formed from a power-law spectrum of electron injection from SHTCL \\citep{Bogachev09}. Fermi acceleration has significant advantages in collapsing magnetic traps as compared with the betatron mechanism which mainly heats the low-energy electrons \\citep[][Chapter~7]{Somov07}.  It is important to note that the emission measure of emitting fast electrons, trapped and accelerated inside the collapsing loop, initially grows slowly with decrease of the loop length $ L (t) $. As calculated by \\cite{Bogachev07}, see their Fig. 3a, the emission measure reaches its maximal value only when the loop length becomes as small as of about 0.2--0.1 of its initial length $ L (0) $. With further shrinkage of the loop, the emission measure decreases quickly to zero. As a consequence, the flux of hard X-ray emission from the collapsing trap attains a maximum at approximately the same values of remaining length of the loop \\citep[see l (t) = L(t)/L(0) $\\sim$ 0.2 in Fig. 4a in][]{Bogachev07}. On the other hand, when the loop is just created by reconnection process at the X-point inside a thin reconnecting current layer \\citep[e.g,][]{Ugai08}, it is strongly stressed by magnetic tensions along magnetic field lines in direction from the X-point to the edge of the current layer. That is why the loop shrinks and becomes less stressed. Moreover, the loop goes down less quickly with the progress of time because the magnetic stress is decreasing continuously. In other words, the loop becomes less non-equilibrium, more close to a state with minimal energy, i.e. the potential state \\citep{Somov06}. In such a state, the height of the loop is naturally proportional to the distance between its feet. Therefore, it is reasonable to assume that, at the time when the collapsing loop has the maximal brightness in hard X-rays, the height of the LT coronal HXR sources is proportional to the distance between the conjugate FP sources in the chromosphere.    Now we shall discuss the consequences of the last assumption coming back to the rainbow reconnection model. Figure~\\ref{flare_cartoon}a shows different flux tubes~$ f_{1} $, $ f_{2} $ etc reconnected at different times~$ t_{1} $, $ t_{2} $ etc; here $ t_{2} > t_{1} $ etc. The first flux tube, the loop~$ f_{1} $ manifests two FPs~$ P_{a} $ and $ P_{b} $ and a loop\\-top HXR source~$ LT $. Figure~\\ref{flare_cartoon}a illustrates two effects. The first one is the well-known classical effect, an increase of the distance between the flare ribbons because of reconnection in the coronal SHTCL. The displacements~$ \\delta x $ of FPs are antiparallel to the converging components~$ {\\bf V}_{\\bot} $ of the velocity field in the photosphere. This means that two flare ribbons move out from the $ PNL $ as predicted by the standard model of two-ribbon flare.  The second effect is less trivial. The displacements~$ \\delta y $ of the FPs are parallel to the PNL because of relaxation of the non\\-potential component of the field created by the photospheric shear flows before a flare. So this is the magnetic shear relaxation. The displacements~$ \\delta y $ of FPs are antiparallel to the photospheric shear velocity~$ V_{\\parallel} $. Since the photospheric shear mainly dominates in the vicinity of the $ PNL $, during the first phase of a flare $ \\delta y \\gg \\delta x $.  In order to interpret our observational results related to the first phase, let us consider a magnetic loop~$ f_{1} $ which after the magnetic shear relaxation process evolves into loop~$f_{2}$. The magnetic loop~$ f_{1} $ has a larger altitude of the loop\\-top HXR source (LT) because the distance between its FPs, $ P_{a} $ and~$ P_{b} $, is larger than the distance between the FPs of loop $f_{2}$, $ P_{a}^{\\, \\prime} $ and $ P_{b}^{\\, \\prime} $. As a result, an apparent motion of the coronal HXR source is directed downward ($ \\delta z < 0 $) and the two conjugate FPs sources converge. The descending motion of the coronal source is typical for the first phase in the flare of October 24, 2003, as well as presumably in many other flares \\citep[see][]{Sui03,Sui04,Li05,Ji06,Veronig06,Ji07,Joshi07,Liu09}. The observations of larger separation between the two conjugate X-ray FPs observed in the first phase than that in the second phase (cf. Figure~\\ref{rhessi_25-50}) also seems to be consistent with our model. However, the detailed analysis FPs motion during the first phase was not possible for this event as FPs were visible for a very short time. %         The second phase of a flare involves reconnection of the magnetic field lines (and flux tubes) whose FPs are located at larger distance from the $ PNL $ and, therefore, have a smaller shear or practically none as schematically illustrated by Figure~\\ref{flare_cartoon}b. Here the apparent displacements~$ \\delta x $ of FPs are directed away from each other and from the $ PNL $. The distance between FPs becomes larger with time, and the LT~HXR source moves upward in agreement with the standard model. This seems to be fairly true for eruptive flares in the later stage when the primary energy source (i.e. reconnection site) is located high enough in the corona. It looks like that the standard model of the two-ribbon flares seems to be ``asymptotically'' true at later stages. However, this is not necessarily a charitable trust. Here it is noteworthy that same physical concepts concerning the photospheric motions may apply to other events such as the Bastille Day Flare of 2000 July 14, although more complicated models are required to explain the observational characteristics of such a large event \\citep[]{Somov02,Somov05b}.  In 72 flare events analyzed by So\\-mov et al. (2005a) on the basis of the {\\em Yohkoh\\/} {\\em HXT\\/} data, it was not simple to distinguish a flare with a decreasing FP separation parallel to the $ PNL $, as discussed above, from a flare with increasing separation also parallel to the $ PNL $, because both kinds of separation were present in the same flare. In the onset of such a flare, the FP sources move toward each other and the distance between them decreases. Then they pass through a ``critical point''. At this moment, the line connecting the sources is nearly perpendicular to the straight $ PNL $; in general, we should characterize the photospheric magnetic field by a {\\em smoothed, simplified\\/} neutral line \\citep[$SNL$, see][]{Gorbachev89} which is not a straight line. After that moment, the FP sources move away from each other with increasing separation between them \\cite [for example, the M4.4 flare on 2000 October~29 at 01:46~UT;][see Figure~2] {Somov05}. Such a motion pattern seems to be similar to that one predicted by the rainbow reconnection model after the critical moment shown in Figure~\\ref{flare_cartoon}b. Starting that moment, the shear relaxation must continue ($ \\delta y > 0 $) if the excess of coronal magnetic tensions have not been completely released during the first phase. Otherwise the ordinary reconnection process ($ \\delta y = 0 $) dominates during the second phase describing the ordinary energy release in terms of the classical standard model. Note that in both cases $ \\delta x > 0 $ and, therefore, during the second phase the FP separation increases and, as a consequence, the height of the loop\\-top HXR sources increases too.   In summary, recent multi-wavelength studies of solar flares have revealed a kind of two evolutionary phases of solar flares in terms of the motion of LT source. In the initial phase (or the first phase) the LT source exhibits a downward motion for a few to several minutes and the explanation of this phase is beyond the scope of classical CHSKP scenario of eruptive flares. The analysis of 2003 October 29 M7.6 flare, presented in this paper, provides one of the clearest observations of the descending LT source with X-ray and EUV images. Recent studies have revealed another important aspect of the first phase in which the descending LT source motion in temporally associated with the converging motion of FPs sources \\citep{Ji06,Ji07,Liu09} which is consistent with the rainbow reconnection model discussed in this paper. \\cite{Ji07} suggested another explanation of the first phase which is based on magnetic implosion conjecture \\citep{Hudson00} that predicts contraction of coronal field lines simultaneously with the energy release. In the framework of sheared linear force-free arcade, \\cite{Ji07} showed that the release of magnetic energy will reduce magnetic shear of the arcades and that the less sheared arcades will have smaller height and span. \\cite{Liu09} further found that the double FPs first do move toward and then away from each other, mainly parallel and perpendicular to the $ PNL $, respectively, and that the transition from the first to the second type of these apparent FP motions coincides with the direction reversal of the motion of the loop\\-top source. Therefore an important aspect of future investigations would be to examine the relationship between the converging motion of FPs sources and the descending LT source motion for more solar flares using existing data and new observations."
    },
    "0809/0809.0162_arXiv.txt": {
        "abstract": "Assuming that the positron excess in PAMELA satellite data is a consequence of annihilations of cold dark matter, we consider from a model-independent perspective if the data show a preference for the spin of dark matter. We then perform a general analysis of annihilations into two-body states to determine what weighted combination of channels best describes the data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.2810_arXiv.txt": {
        "abstract": "We present radio observations of the optically bright Type IIn supernova SN 1995N. We observed the SN at  radio wavelengths with the Very Large Array (VLA) for 11 years. We also observed it at low radio frequencies with the Giant Metrewave Radio Telescope (GMRT) at various epochs within $6.5-10$ years since explosion. Although there are indications of an early optically thick phase, most of the data are in the optically thin regime so it is difficult to distinguish between synchrotron self absorption (SSA) and free-free absorption (FFA) mechanisms. However, the information from other wavelengths indicates that the FFA is the dominant absorption process. Model fits of radio emission with the FFA give reasonable physical parameters. Making use of  X-ray and optical observations, we derive the physical conditions of the shocked ejecta and the shocked CSM. ",
        "introduction": "\\label{sec:introduction}   Type IIn supernovae (hereafter SNe IIn) show narrow H$\\alpha$ emission line atop a broad emission feature, with no broad absorption. This classification was  suggested by \\citet{sch90} who noticed the above properties in 8 supernovae (SNe). The H$\\alpha$ and bolometric luminosities of most SNe IIn are large, which can be explained by the shock interaction of supernova (SN) ejecta with a very dense circumstellar (CS) gas \\citep{chu90}. Many SNe IIn also show an excess late time infrared emission, which is  a signature of dense CS dust \\citep{ger02}. CS interaction in SNe is typically observed by radio and X-ray emission \\citep{che03}. However, in many cases, radio emission from SNe IIn has not been detected \\citep{van96}.  One of the important issues for SNe IIn is their high presupernova mass loss rates. For example, for SN 1997ab the derived mass loss rate for a presupernova wind  velocity of $90\\,{\\rm km\\, s^{-1}}$ is $\\sim 10^{-2}\\, M_\\odot\\,{\\rm yr^{-1}}$ \\citep{sal98}. For SN 1994W the mass loss rate was $\\sim 0.2\\, M_\\odot\\,{\\rm yr^{-1}}$ \\citep{chu04}, and for SN 1995G $\\sim 0.1\\, M_\\odot\\,{\\rm yr^{-1}}$ \\citep{chu03}. One possibility for these high mass loss rates is that they are related to an explosive event a few years before the SN outburst \\citep{chu03,pas07}. Radio observations may put better constraints on the mass loss of the progenitor stars.  SN 1995N was discovered in MCG $-$02$-$38$-$017 (Arp 261) on 1995 May 5 \\citep{pol95} at a distance of 24 Mpc  \\citep[see][]{fra02}. \\citet{pol95} estimated that the supernova was at least 10 months old upon discovery by comparing it with spectroscopic chronometers, and classified it as a Type IIn supernova. We assume the date of explosion to be 1994 July 4 throughout the paper \\citep{fra02}. SN 1995N turned out to be relatively bright at late times in most wavebands. It was detected in the radio with the Very Large Array \\citep[VLA,][]{van96} and the Giant Metrewave Radio Telescope \\citep[GMRT,][]{cha05}. In X-rays it was detected first by  ROSAT in Aug 1996 and Aug 1997,  and later by  ASCA in Jan 1998 \\citep{fox00}. In July 2003 XMM-Newton \\citep{zam05} and in March 2004  Chandra X-ray Observatory \\citep{cha05} detected the late time X-ray emission from the SN. SN 1995N has declined very slowly in the optical band \\citep{li02,pas05}, with only a 2.5 mag change in the V-band over $\\sim2500$ days after explosion. This  is consistent with the slow spectral evolution reported in \\citet{fra02}. Ground based optical and HST observations of the late-time spectral evolution of SN 1995N were used by \\citet{fra02} to argue that the late-time evolution is most likely powered by the X-rays from the interaction of the ejecta and the circumstellar medium of the progenitor. They in turn proposed that the progenitors of Type IIn supernovae are similar to red supergiants in their superwind phases, when most of their hydrogen-rich gas is expelled in the last $10^4\\, \\rm yr$ before explosion.   In this paper, we report extensive radio observations of SN 1995N taken with the VLA and GMRT. We discuss the observations and data analysis in \\S \\ref{sec:observations}. We fit various models to the data and discuss the results in \\S \\ref{sec:absorption}. In \\S \\ref{sec:interpretation}, we utilize the results from other wavelength observations along with the radio bands to derive constraints on the physical parameters. We summarize our results briefly in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  A useful way to characterize the radio emission from supernovae is by considering their peak radio luminosity and time of the peak; the more luminous Type II supernovae peak at a later time, indicating that they have a denser circumstellar medium \\citep{che06}. In such a plot, SN 1995N is close to the Type IIL SN 1979C and Type IIn SN 1978K, but is less luminous that the Type IIn SNe 1986J and 1988Z \\citep{wei90,wil02}. However, it is more radio luminous than many of the SNe IIn observed by \\citet{van96}.  The optically thin radio evolution of SN 1995N is fairly well defined and yields a radio spectral index $\\alpha\\approx 0.6$. This value is in the range found for Type II SNe and is smaller than the values found for Type Ib/c SNe \\citep{wei02}. The rate of decline is also similar to that found for other SNe II and, in a simple model for the emission, indicates an expansion factor for the emitting region $m=0.83$. Assuming interaction with a steady wind, the implied value of the power law index of the outer supernova density profile is $n\\approx8$. The same value of $n$ is consistent with the X-ray properties of SN 1995N.  Although there are clear indications of early low frequency absorption in the radio light curves of SN 1995N, there is little information on the evolution during the optically thick phase so that the nature of the absorption cannot be deduced from the radio light curves. Based on the emission measure of the gas deduced from optical observations and the late  peak radio luminosity, we infer that free-free absorption is the likely absorption mechanism. The implied mass loss rate for the supernova progenitor is $6 \\times 10^{-5}\\, {\\rm M_\\odot \\, yr^{-1}}$, but this value has a number of uncertainties. The assumed velocity of the radio emitting region is $5000$ km s$^{-1}$; however, there are some indications from the H$\\alpha$ line that there are higher velocities \\citep{fra02}, which would lead to a higher value of $\\dot M$. The assumed wind velocity of 10 km s$^{-1}$ is lower than has been inferred from the narrow H$\\alpha$ lines observed in some SNe IIn, e.g., 90 km s$^{-1}$ in SN 1997ab \\citep{sal98}, 160 km s$^{-1}$ in SN 1997eg \\citep{sal02}, and 100 km s$^{-1}$ in SN 2002ic \\citep{kot04}. An order of magnitude increase in $u_w$ would increase $\\dot M$ by an order of magnitude. Finally, the optical line emission suggests clumping of the ionized wind gas \\citep{fra02}, which would reduce the estimate of $\\dot M$ found here.  Regardless of the uncertainty in $\\dot M$, the values deduced for SNe IIn from optical observations listed in \\S~1 are significantly larger than the value deduced for SN 1995N. In some of those cases, the mass loss may be a short lived event before the supernova. The long time coverage of SN 1995N, especially at radio wavelengths, shows that the circumstellar gas is extended in this case, out to radii $\\ga 2\\times 10^{17}$ cm. The Type IIn supernovae are a heterogeneous group of objects."
    },
    "0809/0809.0738_arXiv.txt": {
        "abstract": "We analyse the results of recent measurements of nonthermal emission from individual supernova remnants (SNRs) and their correspondence to the nonlinear kinetic theory of cosmic ray (CR) acceleration in SNRs. It is shown that the theory fits these data in a satisfactory way and provides the strong evidences for the efficient CR production in SNRs accompanied by significant magnetic field amplification. Magnetic field amplification leads to considerable increase of CR maximum energy so that the spectrum of CRs accelerated in SNRs is consistent with the requirements for the formation of Galactic CR spectrum up to the energy $\\sim 10^{17}$~eV. ",
        "introduction": "The main reason why supernova remnants (SNRs) are considered as a cosmic ray (CR) source is a simple argument about the energy required to sustain the Galactic cosmic ray (GCR) population against loss by escape, nuclear interactions and ionization energy loss. The mechanical energy input to the Galaxy from each supernova (SN) is about $10^{51}$ erg so that with a rate of about one every 30 years the total mechanical power input from supernovae is of the order $10^{42}$ erg/s \\citep[e.g.][]{berezinetal90}. Thus supernovae have enough power to drive the GCR acceleration if there exists a mechanism for channeling about 10\\% of the mechanical energy into relativistic particles. The high velocity ejecta produced in the supernova explosion interacts with the ambient medium to produce a system of strong shocks. The shocks in turn can pick up a few particles from the plasma flowing into the shock fronts and accelerate them to high energies.  The only theory of particle acceleration which at present is sufficiently well developed and specific to allow quantitative model calculations is diffusive acceleration applied to the strong shocks associated with SNRs \\citep[e.g. see][for review]{drury83, blei87, bk88,jel91,md01}. Considerable efforts have been made during the last years to empirically confirm the theoretical expectation that the main part of GCRs indeed originates in SNRs.Theoretically progress in the solution of this problem has been due to the development of the kinetic nonlinear theory of diffusive shock acceleration \\citep{byk96, bv97, bv00a}. The theory includes all the most relevant physical factors, essential for SNR evolution and CR acceleration, and it is able to make quantitative predictions of the expected properties of CRs produced in SNRs and their nonthermal radiation.  Detail information about high-energy CR population in SNRs can be obtained from the observations of the nonthermal emission produced by accelerated CRs in SNRs. The electron CR component is evident in a wide wave length range by the radiation that they produce in SNRs,  from radio to $\\gamma$-ray emission, whereas in the case of the nuclear CR component a $\\gamma$-ray detection is the only possibility to find it. If this nuclear component is strongly enhanced inside SNRs then through inelastic nuclear collisions, leading to pion production and subsequent decay, $\\gamma$-rays will be produced at the detectable level.  High-energy $\\gamma$-rays can also be produced by CR electrons due to the inverse Compton (IC) scattering of the background photons. Therefore it is not so obvious which CR component (nuclear or electron) produces TeV emission, detected from Cassiopeia~A (Cas~A) \\citep{aha01}, RXJ1713.7-3946 \\citep{mur00,enomoto02,aha05,aha06a} and Vela~Jr \\citep{aha06b}. For this determination the SNR magnetic field value plays the crucial role. The relative role of electrons in $\\gamma$-ray production is low if SNR magnetic field is strongly amplified. In addition, if the magnetic fields in SNRs are considerably larger than the typical interstellar value, then SNRs are able to generate a power-law spectrum of accelerated CRs at least up to the knee energy \\citep{bk99,bv04}. In fact, there is strong evidence that the actual magnetic field in SNRs is indeed intrinsically enhanced: the analysis of the nonthermal emission of young SNRs show that all of them have amplified fields \\citep{vbk05, pariz06}. Required strength of the magnetic field, that is significantly higher than the interstellar medium (ISM) field, has to be attributed to nonlinear field amplification at the SN shock by CR acceleration itself, in fact by the dominant nuclear CR component \\citep{lb00,bell04}.  The application of kinetic theory to individual SNRs \\citep[see][for a review]{ber05} has demonstrated its power in explaining the observed SNR properties and in predicting new effects like the extent of magnetic field amplification, leading to the concentration of the highest-energy electrons in a very thin shell just behind the shock. The theory also predicted the enhancement of the high-energy galactic \\gr background radiation due to the contribution of CRs confined in parent SNRs, that was partly confirmed by the measured background TeV-emission.  In order to perform the detail comparison of theoretical expectation with the experiment one needs sufficient number of individual SNRs with known values of relevant physical parameters such as the age, distance, ISM density. Unfortunately the number of such SNRs are very limited, especially thous which are seen in all wavelength from radio to \\gr. It is why every new experimentally established property of nonthermal emission of SNR represents considerable interest for theoretical analysis. Therefore we start this review with a brief consideration of three SNRs for which essentially new information about the efficiency of CR production was recently obtained. Then we shall consider the correspondence of the interior magnetic field values extracted from the analysis of the experimental data with theoretical expectation. Finally, we discuss the maximum energy of CRs which can be achieved during their acceleration in SNRs, taking into account realistic values of magnetic field in SNRs. ",
        "conclusions": "Detailed consideration performed within a frame of nonlinear kinetic theory demonstrates that the CR production efficiency in young SNRs is consistent with requirements for the Galactic CR sources. Recent experimental and theoretical studies of Tycho's SNR, \\rxj and SN~1987A provides new very strong evidence for that.  Synchrotron radiation from young SNRs provides evidence that efficient CR acceleration followed by the significant shock modification and magnetic field amplification takes place there. Magnetic field amplification leads to the considerable increase of maximum energy of CRs, accelerated in SNRs. Calculations performed within kinetic nonlinear theory demonstrate that the expected GCR spectrum, produced in SNRs, in satisfactory way fits the existing GCR data up to the energy $10^{17}$~eV. The first knee in the observed all particles GCR spectrum is attributed to the maximum energy of protons, produced in SNRs. The steepening of the all particle GCR spectrum above the knee energy $3\\times 10^{15}$~eV is a result of progressive depression of contribution of light CR nuclei with the energy increase. Such a scenario is confirmed by KASCADE experiment which shows relatively sharp cutoff spectra of GCR species at energies $\\epsilon_{max}\\approx 3Z\\times10^{15}$~eV \\citep{kascade}.  The nature of the whole GCR spectrum, which does not require any significant contribution of galactic CR sources other then SNRs, can be the following. GCR spectrum up to the energy $\\epsilon=10^{17}$~eV is dominated by the contribution of galactic sources, whereas at $\\epsilon>10^{18}$~eV GCRs are predominantly extragalactic. Such a picture is consistent with the analysis, presented here, and with the analysis of ultra high energy CR properties \\citep{aloiso06} as well. Detail experimental study of GCR spectrum and composition could find whether indeed the transition from galactic to extragalactic component takes place in energy range $10^{17}-10^{18}$~eV.  Due to high interior magnetic field in young SNRs the $\\pi^0$-decay $\\gamma$-rays generated by the nuclear CR component as a rule dominates over $\\gamma$-rays, generated by electron CR component, and calculated $\\gamma$-ray flux fits existing data.  New generation of ground-based $\\gamma$-ray detectors CANGAROO~III, H.E.S.S., MAGIC, VERITAS will presumably measure significantly more accurate $\\gamma$-ray spectra from already detected SNRs, like Cas~A, RXJ1713.7 and Tycho's SNR, as well as from a number of new sources. In this regard considerable interest as a potential \\gr sources represent Tycho's SNR, SN~1987A and Kepler's SNR \\citep{bkv06} eventhough the expected \\gr fluxes from these objects only slightly above the sensitivity of modern TeV \\gr instruments. Together with data, which soon will be obtained by space $\\gamma$-ray observatory GLAST, they will reproduce the full $\\gamma$-ray spectra of some number of individual SNRs, as well as the Galactic background spectrum. Although the existing data together with their theoretical interpretation presents a consistent picture of GCR origin, these new measurements will allow for detailed study of SNRs and the CR production inside them."
    },
    "0809/0809.1955_arXiv.txt": {
        "abstract": "The first high-resolution ($\\sim$5~mas) VLBI observations of 6.7-GHz methanol masers in DR21(OH)N, a candidate circumstellar disc around a very young massive star, are presented. Previous observations of these masers at 50~mas angular resolution revealed a rotating structure at the position of a candidate massive protostar, with a well-sampled position-velocity diagram suggesting Keplerian rotation. Observations presented here using the European VLBI Network (EVN) have provided the first high angular resolution maps of the masers, providing a test for the disc hypothesis and the Gaussian centroiding technique. The EVN maps have confirmed the shape of the disc and its rotation curve. Weaker maser emission seen previously with MERLIN between the two main spectral peaks is seen in the EVN total power spectrum, but is absent in the cross-power spectrum. This suggests that the spatially extended emission is resolved out by the EVN. The rotating disc is coincident with a Class I massive (proto)star and at the implied centre of an outflow traced by two bow shocks. We discuss the impact of this result on the massive stellar accretion disc hypothesis and on the validity of the centroiding technique to determine the structures of unresolved masers using compact radio interferometric arrays. ",
        "introduction": "Massive stellar accretion is an important astrophysical process, the mechanics of which are still poorly understood (for a thorough review of the topic, see Zinnecker \\& Yorke, 2007). Observing masers in massive star-forming regions allows us to trace the bulk motions of molecular gas in the vicinity of nascent massive stars. This can give us important information about the infall, outflow and rotational motions associated with massive star formation, in particular identifying possible circumstellar discs around massive protostars (Cesaroni et al. 2007).  Recent observations of 6.7-GHz methanol masers using the Multi-Element Radio Linked Interferometer Network (MERLIN) in the massive star-forming region DR21(OH)N led to the discovery of a highly unusual maser with a double peaked spectrum (Harvey-Smith et al. 2008). The maser region is elongated in an east-west direction and has an estimated diameter of 60~au. The resulting position-velocity diagram showed a clear Keplerian rotation profile, suggesting that the methanol masers may be tracing a rotating disc of material falling onto a massive star. The masers coincide with the very young Class I candidate massive star in the infra-red and dust core ERO3 (Marston et al. 2004). This extremely reddened object is at the centre of (and perpendicular to) a possible molecular outflow described by Davis et al. (2007). The magnetic field direction implied from the linear polarization of the methanol masers is roughly perpendicular to the implied plane of the disc. The indirect evidence for a young, massive circumstellar disc in DR21(OH)N is considerable, but to remove any uncertainty we must be able to confirm the rotation curve observed by Harvey-Smith et al. (2008). This is the single piece of observational evidence that can (i) directly demonstrate the motion of molecular gas around the young massive star and (ii) produce an estimate of the current mass of the star. In order to test the reliability of the position-velocity diagram and therefore the disc interpretation, the 6.7-GHz methanol masers in DR21(OH)N were re-observed using the European VLBI Network at the same spectral resolution but with an angular resolution ten times greater than that of MERLIN (5~mas compared with 50~mas). In doing this, it would become possible to detect the effect of spatial blending in the MERLIN centroid maps and to verify (or otherwise) the velocity curve of the maser disc. ",
        "conclusions": "Observations of the DR21(OH)N region using the European VLBI Network have confirmed the existence of an unusual methanol maser feature at 6.7~GHz with a double-peaked spectral profile and a velocity gradient of 32 m~s$^{-1}~au^{-1}$. Analysis of the single dish spectrum reveals weaker maser emission linking the two peaks. Gaussian centroid analysis of the maser emission in each spectral channel revealed a velocity profile resembling a Keplerian rotation curve. The masers coincide with a Class I (very young) B-type protostar, which is located at the centre of an outflow seen by the bow-shocks in H$_2$. The ouflow lies parallel to the magnetic field direction in the maser disc, implied by observations of the linear polarization of the methanol masers. If this object is not a disc around a massive protostar then it would be very difficult to explain the origin of the large velocity gradient and clear Keplerian rotation profile. A model is currently under development to constrain the central mass of the protostar in DR21(OH)N, which will prove very important in verifying the origin of this disc, resulting in wider implications in testing current models of massive star-formation ."
    },
    "0809/0809.0249_arXiv.txt": {
        "abstract": "Studying the internal structure of exoplanet-host stars compared to that of similar stars without detected planets is particularly important for the understanding of planetary formation. The observed average overmetallicity of stars with planets is an interesting point in that respect. In this framework, asteroseismic studies represent an excellent tool to determine the structural differences between stars with and without detected planets. It also leads to more precise values of the stellar parameters like mass, gravity, effective temperature, than those obtained from spectroscopy alone. Interestingly enough, the detection of stellar oscillations is obtained with the same instruments as used for the discovery of exoplanets, both from the ground and from space. The time scales however are very different, as the oscillations of solar type stars have periods around five to ten minutes, while the exoplanets orbits may go from a few days up to many years.  Here I discuss the asteroseismology of exoplanet-host stars, with a few examples. ",
        "introduction": "I can see at least four different reasons for which astrophysicists interested in exoplanets should also bother about the asteroseismology of the central stars of planetary systems.  First reason: the observations for stellar oscillations and exoplanet searches are done with the same instruments. In some cases, the same observations, analysed on different time scales, can lead to both planet detection and seismic studies. This was the case for the star $\\mu$ Arae, observed with HARPS during eight nights in June 2004: these observations, aimed for asteroseismology, lead to the discovery of the exoplanet $\\mu$ Arae d (Santos et al. \\cite{santos04}).  Second reason : \"Some people's noise is other people's signal\". Indeed, when searching for exoplanets, the signal to noise ratio is limited by the stellar oscillations, which appear as a noise for the radial velocity variations induced by the planetary motions, while they represent in fact the stellar oscillations signal. A better knowledge of seismology and a better treatment of this low time scale signal could help determining the planetary parameters.  Third reason: obtain precise values of the parameters of exoplanet-hosts stars. Asteroseismic studies, combined with spectroscopic observations, can lead to values of the stellar parameters which are much more precise than from spectroscopy alone. In the case of $\\iota$ Hor, for example, the mass determined from spectroscopy and position in the HR diagram was $10\\%$ wrong (Vauclair et al. \\cite{vauclair08}).  Fourth reason: links between asteroseismology and planets discoveries. In two cases at least, asteroseismic studies lead to exoplanet discoveries. One is the case, already mentioned, of $\\mu$ Arae. Studies of seismic period variations (the so-called ``time delay method\") also lead recently to the spectacular discovery of a planet around the extreme horizontal branch star V391 Peg (Silvotti et al. \\cite{silvotti07}).  Generally speaking, the goals of the asteroseismology of exoplanets-host stars are to derive their masses, ages, evolutionary stages, outside and inside metallicities, to compare them with the Sun and with stars without detected planets, to obtain hints about the theories of planetary formation and migration, and to try and model the oscillations well enough to decrease the \"noise\" for planet detection. ",
        "conclusions": "From the few examples already available, asteroseismology has proved to be a powerful tool for determining stellar parameters, in particular for exoplanet-host stars. The scientific community is well prepared for future observations with on-going or planned projects. Space missions like COROT, and later on KEPLER, are expected to give a large amount of new data for seismology, besides planet searches. Meanwhile ground based instruments devoted to exoplanets like HARPS or SOPHIE can be used for seismology. All these observations will give the possibility of using the seismic tests with a precision never obtained before. They will lead to better parameters for exoplanet-host stars and help for a better understanding of planetary formation and evolution."
    },
    "0809/0809.2599_arXiv.txt": {
        "abstract": "We report on the first detection of hard X-ray photons ($E>$2.5 keV)  in the X-ray transient Narrow-Line Seyfert 1 galaxy WPVS 007 which was the AGN with the softest X-ray spectrum during the ROSAT All-Sky Survey. The AGN is clearly detected at a level of about 2$\\times 10^{-17}$ W m$^{-2}$ in the observed 0.3-10.0 keV band by \\swift\\ in a 50 ks observation in 2007 September. For the first time since the ROSAT All-Sky Survey observation in 1990 it was possible to derive an X-ray photon distribution by adding together all \\swift\\ observations that have been performed so far (85.5 ks in total). This photon distribution is consistent with an X-ray spectrum  of an AGN with a partial covering absorber with a column density in the order of $\\sim 1\\times10^{23}$ cm$^{-2}$ and a covering fraction of about 90\\%.  A comparison with the 2002 \\chandra\\ data suggests that WPVS 007 has become brighter by a factor of about 4. The \\swift\\ data also suggest that the absorber which is causing the current low-state may have started to disappear. This disappearance is indicated by a significant change in the hardness ratio from a very soft X-ray state during the 2005 October to 2007 January observations to a rather hard X-ray state in the 2007 September observations. In the UV, WPVS 007 seems to become fainter by up to 0.5 mag over the last two years. The optical to X-ray spectral slope derived from the spectral energy distribution is \\aox=2.5 which classifies WPVS 007 as an X-ray weak AGN. After correcting for reddening and X-ray absorption, \\aox\\ becomes 1.9 and the luminosity in the Big-Blue-Bump is log $L_{\\rm BBB}$=37.7 [W], which translates into an Eddington ratio $L/L_{\\rm Edd}\\approx 1$. ",
        "introduction": "The Narrow-Line Seyfert 1 galaxy (NLS1) \\wpvs\\ \\citep[1RXS J003916.6$-511701$, RBS 0088; $\\alpha_{2000}$ = $00^{\\rm h} 39^{\\rm m} 15.^{\\rm s}8$, $\\delta_{2000}$ = $-51^{\\circ} 17' 03 \\farcs 0$, z=0.029; ][]{wamsteker85,grupe95} is a unique X-ray transient AGN. During the ROSAT All-Sky Survey \\citep[RASS, ][]{voges99}, when WPVS 007 was in its X-ray high-state, the X-ray-to-optical flux ratio was log $f_{\\rm x}/f_{\\rm o}$\\footnote{$f_{\\rm x}/f_{\\rm o}$ is defined by \\citet{beu99} as log($f_{\\rm x}/f_{\\rm o}$)= log CR + 0.4 V - 5.86, with CR = the ROSAT PSPC count rate in the 0.1-2.0 keV range.}=0.22 which is typical for a Seyfert 1 galaxy \\citep[e.g. ][]{beu99, maccacaro88}. Nevertheless, in all X-ray follow-up observations between 1993 and 2007 using ROSAT, \\chandra\\, and \\swift\\ WPVS 007 has been found in an extreme X-ray low-state that made WPVS 007 almost vanish from the X-ray sky \\citep{gru01, grupe07}. It had been last detected in 2002 August by Chandra with $1\\times 10^{-3}$ ACIS-S counts s$^{-1}$ \\citep{vaughan04}. While this strange transient behavior had been a mystery for a decade, \\citet{leighly05} and \\citet{leighly08}  discovered that it is probably absorption that causes the extreme low-state in this AGN. They found in a 2003 November {\\it FUSE} observation that the UV spectrum of WPVS 007 had changed dramatically compared with HST observations obtained in 1996 July    \\citep{goodrich00, constantin03}:   the {\\it FUSE} observation revealed the emergence of a  BAL flow.   \\citet{leighly08} concluded that WPVS 007 is a low-luminosity, low black hole   mass cousin of Broad Absorption Line Quasars (BALQSOs). The relatively short timescale for the development of the BAL inferred in WPVS 007 suggests that it could re-brighten at any time on relatively short timescales. We started a monitoring campaign using \\swift\\ in 2005 October \\citep{grupe07}. The purpose of this campaign is to detect a possible rebrightening and therefore the disappearance of the absorber. By the end of the second year of monitoring WPVS 007 with \\swift\\ in 2007 January, only a 3$\\sigma$ upper limit of the observed X-ray flux could be derived at a level of $2.3\\times 10^{-17}$ W m$^{-2}$ in the rest-frame 0.2-2.0 keV band.  \\citet{brandt00} found that about 10\\% of optically selected quasars appear to be X-ray weak with an optical to X-ray spectral slope \\aox\\footnote{The X-ray loudness is defined by \\citet{tananbaum79} as \\aox=--0.384 log($f_{\\rm 2keV}/f_{2500\\AA}$).}$>$2.0. X-ray weakness can either be intrinsic, such as seen in e.g. PHL 1811 \\citep[][]{leighly07}, or caused by intrinsic absorption. Low X-ray flux states have been found in several NLS1s, such as Mrk 1239, 1H 0707--495, and most recently in Mrk 335 \\citep{grupe04b, boller02, gallo04, grupe07b}. In Mrk 335 \\citet{grupe07b} showed that this dramatic drop in the X-ray flux goes along with a strong change in the X-ray spectrum. One of the possible scenarios to explain the X-ray spectra of low-state NLS1s is a partial covering absorber. An alternative solution to the strange looking low-state spectra of these NLS1s is reflection as suggested for example by \\citet{fabian89, fabian04}. The X-ray spectra by themselves do not provide a definite solution to distinguish between both models because with the quality of the spectra typically obtained and the number of free model parameters, both models fit the X-ray spectra well and they are statistically indistinguishable. In addition, the cases of several well-studied AGN have shown that low X-ray flux states are temporary events and not necessarily persistent \\citep[e.g.][]{komossa97, guainazzi98, costantini00, komossa01, bianchi05, grupe07b, grupe08, ballo08a}.    The \\swift\\ Gamma-Ray Burst  explorer mission \\citep{gehrels04} was launched on 2004 November 20.  With  its  simultaneous multi-wavelength capacity and its scheduling flexibility, \\swift\\ has shown itself to be the ideal observatory to study highly variable objects like AGN.  Even though \\swift's main purpose is to detect and observe Gamma-Ray bursts, \\swift\\ has also done an excellent job in observing AGN as fill-in targets or Target-of-Opportunity (ToO) observations \\citep[e.g. ][]{markwardt05, grupe06, grupe07, grupe07b, kataoka07, sambruna07, giommi07}. \\swift's UV/Optical Telescope  \\citep[UVOT, ][]{roming04} and X-ray Telescope \\citep[XRT, ][]{burrows04} cover the electromagnetic spectrum between 6500\\AA\\ on the low energy side  to 10 keV at the high energy end simultaneously.  We started the WPVS 007 monitoring campaign again in 2007 July. On 2007 August 10 a failure of one of the \\swift\\ gyros forced the spacecraft into a safe-hold \\citep{gehrels07}. During the spacecraft recovery process, \\swift\\ was slewing only around a small circle in the sky that included the position of WPVS 007.  We took advantage of this limited slewing capacity and obtained about 50ks of X-ray observations of the field of WPVS 007 in 2007 September.  During these observations we were not only able to clearly detect WPVS 007 in X-rays, but even more importantly, we discovered a significant number of photons at hard energies above 2 keV, and derived a photon distribution that gives us information for the first time since the RASS about the X-ray spectrum.  In this paper we present and discuss these new observations.  The outline of this paper is as follows: in \\S\\,\\ref{observe} we describe the \\swift\\  observations and the data reduction, in \\S\\,\\ref{results} we present the results of the \\swift\\ XRT data, and in \\S\\,\\ref{discuss} we discuss the results. Throughout the paper spectral indexes are denoted as energy spectral indexes with $F_{\\nu} \\propto \\nu^{-\\alpha}$. Luminosities are calculated assuming a $\\Lambda$CDM cosmology with $\\Omega_{\\rm M}$=0.27, $\\Omega_{\\Lambda}$=0.73 and a Hubble constant of $H_0$=75 km s$^{-1}$ Mpc$^{-1}$ using a luminosity distances D=118 Mpc given by \\citet{hogg99}. All errors are 1$\\sigma$ unless stated otherwise. ",
        "conclusions": "Discussion}  We presented new \\swift\\ observations on the X-ray transient NLS1 \\wpvs. We were able to detect the AGN in X-rays at a level of 5$\\times 10^{-4}$ XRT counts s$^{-1}$ in the 0.3-10.0 keV band which is equivalent to an X-ray flux of 2$\\times 10^{-17}$ W m$^{-2}$. The most interesting aspect of this detection is that the majority of source photons are detected at energies $E>$2 keV which is extremely surprising considering that WPVS 007 used to be the AGN with the softest X-ray spectrum in the sky, having no photon above 0.5 keV detected during the RASS.  Compared with the ACIS-S count rate measured during the 2002 \\chandra\\ observation \\citep[$1\\times 10^{-3}$ ACIS-S counts s$^{-1}$;][]{vaughan04} this indicates that WPVS 007 has become brighter by a factor of a few. This variability, however, only occurs in the hard band. The 0.2-2.0 keV flux between the \\chandra\\ 2002 and \\swift\\ 2007 observations remains constant within the errors. The non-variability of the optical spectra seen by \\citet{winkler92} and \\citet{grupe95} suggests that the intrinsic continuum spectrum may not have changed. This non variability in the optical spectrum is somewhat similar to Mrk 335 where the optical spectra taken in 1999 during the X-ray high state and 2007 September during the low-state are identical \\citep{grupe08}. The non variability of the optical spectrum together with the strong variability seen in X-rays and the UV may suggest that the inclination angle in which we observe the nuclear region is rather high assuming that the variability is a consequence of variable absorption. This is expected for BAL QSOs according to the AGN geometry model presented by \\citet{elvis00}. However, this is not exclusive. Recently there have been reports by e.g. \\citet{berrington07} and \\citet{ghosh08} of BAL QSOs with polar outflows.   Although the best-fit model is the partial covering absorber model, in principle a reflection model \\citep[e.g.][]{fabian89,fabian04} results in a similar shape of the X-ray spectrum. While it is impossible with the current data set to distinguish betwen the partial covering absorber and refelction models, the presence of a strong absorber in the UV spectrum \\citep{leighly08} makes the partial covering absorber model more plausible as the explanation for the observed \\swift\\ X-ray spectrum.  One question is if the X-ray spectrum has changed intrinsically, or whether the absorber has become more transparent sometime between 2002 and 2007. The strong UV absorption in the FUSE observations also suggests that WPVS 007 was highly absorbed in 2003 November. There was only one photon found above 2 keV in the \\chandra\\ 2002 observation, but 24 in the combined \\swift\\ data. In order to answer the question we simulated a 10 ks ACIS-S spectrum using the spectral parameters of a partial covering absorber fit derived from the \\swift\\ data. Using these parameters we expect to see 13 photons with energies above 2 keV in the \\chandra\\ data. However, only one single photon with $E>$2keV was found. This is a statistically significant lower value than the expected number of photons. We performed another simulation assuming a column density of the partial covering absorber of $N_{\\rm H,pc} = 1\\times 10^{24}$ cm$^{-2}$ and leaving all other spectral parameters at the same values. This simulation showed that the hard X-ray photons could be suppressed significantly down to one photon above 2 keV in a 10 ks Chandra ACIS-S observations. Even though we are dealing with low number photon statistics, the hardness ratios found from the \\swift\\ data prior 2007 January observations the Chandra 2002 data both indicate a rather soft X-ray spectrum, while the 2007 September data suggest a rather hard spectrum. Because the flux during the 2007 September \\swift\\ observation is a factor of about 4 higher than during the 2002 \\chandra observation, it suggests that the column density of the absorber of the high energy spectrum had decreased significantly sometime in the time between 2007 January and September.  This result suggests that the absorber has started to disappear again. Note, that a change in the absorber column density of a fully-covered absorber has exactly the opposite effect on the hardness ratio than what a partial covered absorber does: while in increase of the absorption column density of a fully-covered absorber causes the hardness ratio to become harder, because primarily soft X-ray photons are absorbed, an increase of the absorption column density of a partial covering absorber causes the source to appear softer, because the soft X-ray photons are dominated by the unabsorbed fraction of the X-ray continuum while the hard X-ray photons become suppressed. Considering that we observed the flux to decrease significantly between the RASS observation in 1990 and the pointed ROSAT observation in 1993, we may suggest that we will be able to see WPVS 007 as a bright X-ray NLS1 again within the next few years.  The optical to X-ray spectral slope \\aox=2.5 clearly classifies WPVS 007 as an X-ray weak AGN following the definition by \\citet{brandt00} who define 'X-ray weak' AGN by \\aox$\\ge$2.0. The reddening-corrected \\aox=1.9 is at the borderline. Note that this value has large uncertainties. Nevertheless, this result is consistent with the findings by \\citet{leighly04} that NLS1s with outflows seen in the blueshifted emission lines have steeper \\aox\\ than NLS1s without outflows.  One question remains: why was the X-ray spectral slope during the RASS observation so steep? An energy index of \\ax=8.7 as measured from a single absorbed power law model to the RASS spectrum \\citep{grupe95} can not be explained by standard accretion disk models. We can also exclude a tidal disruption event by a star orbiting closely to the central black hole. First, the optical to X-ray flux ratio during the RASS was log $f_{\\rm X}/f_{\\rm o}$=0.22\\footnote{The typical optical to X-ray spectral slope \\aox\\ cannot be applied to these data because no X-ray photons with $E>$0.5 keV were detected in the RASS data making any estimate of \\aox very uncertain.} which is typical for a Seyfert 1 galaxy \\citep{beu99,maccacaro88}, and second we do not detect any significant changes in the optical spectrum. This is different than what has been found in the X-ray transient Seyfert 2 galaxy IC 3599 \\citep{grupe95a, brandt95} or as recently reported for SDSS J095209.56+214313.3  by \\citet{komossa08} where strong changes in the optical emission lines were found as a result of an X-ray outburst caused by a dramatic accretion event. Note also that the X-ray to optical flux ratio of IC 3599 during the RASS was log $f_{\\rm X}/f_{\\rm o}$=1.45. One possible explanation for the absence of any photon above 0.5 keV during the RASS observation is the presence of an ionized, 'warm', absorber in the line of sight. This type of absorber is transparent at soft energies but absorbs at intermediate energies, and therefore provides an efficient mechanism to produce very steep soft X-ray spectra (e.g., Komossa \\& Meerschweinchen 2000, Done \\& Nayakshin 2007). Using {\\it absori} within XSPEC we simulated a 300s spectrum of WPVS 007 during the RASS with an absorber column density $N_{\\rm H}=5\\times 10^{22}$ cm$^{-2}$ of the warm absorber with an ionization parameter $\\xi$=1000. Fitting this simulated spectrum with an absorbed  single power law model results in \\ax=6.2$^{+2.0}_{-1.2}$, which is similar to what has been found during the RASS \\citep{grupe95}.  Our detection of WPVS 007 with \\swift\\ and the extraction of a photon distribution from the 2005 to 2007 data opens a new window in our understanding of WPVS 007 and X-ray weak AGN in general. While in the UV we were able to follow the increase of the absorber column density \\citep{leighly05, leighly08} we are now able to follow the disappearance of the absorber in X-rays. This will need future follow-up observations by \\xmm\\ or \\chandra\\ and \\swift. While \\xmm\\ or \\chandra\\ will be able to obtain a detailed low-state spectrum, \\swift\\ will continue monitoring \\wpvs\\ in order to determine when it will become X-ray bright again. In addition, optical polarimetry or spectropolarimetry is needed to measure how much of the optical emission is seen directly and how much is scattered (polarized) emission. Spectropolarimetry will put strong constraints on the geometry of the AGN and will confirm (or not) the assumption that WPVS 007 is viewed at a rather high inclination angle. We also have an approved HST COS observation along with a Chandra observation in order to study the developments in the UV absorption lines and in the X-ray continuum."
    },
    "0809/0809.0887_arXiv.txt": {
        "abstract": "If Type-{\\sc ii} supernovae -- the evolutionary end points of short-lived, massive stars -- produce a significant quantity of dust ($>$0.1\\,\\Msun) then they can explain the rest-frame far-infrared emission seen in galaxies and quasars in the first Gyr of the Universe. Submillimetre (submm) observations of the Galactic supernova remnant, Cas~A, provided the first observational evidence for the formation of significant quantities of dust in Type-{\\sc ii} supernovae. In this paper we present new data which show that the submm emission from Cas~A is polarised at a level significantly higher than that of its synchrotron emission. The orientation is consistent with that of the magnetic field in Cas~A, implying that the polarised submm emission is associated with the remnant. No known mechanism would vary the synchrotron polarisation in this way and so we attribute the excess polarised submm flux to cold dust within the remnant, providing fresh evidence that cosmic dust can form rapidly. This is supported by the presence of both polarised and unpolarised dust emission in the north of the remnant, where there is no contamination from foreground molecular clouds. The inferred dust polarisation fraction is unprecedented ($f_{\\rm pol}\\sim 30$ per cent) which, coupled with the brief timescale available for grain alignment ($<$300\\,yr), suggests that supernova dust differs from that seen in other Galactic sources (where $f_{\\rm pol}=2-7$ per cent), or that a highly efficient grain alignment process must operate in the environment of a supernova remnant. ",
        "introduction": "\\label{sec:intro}  The large quantities of dust seen in high-redshift quasars and galaxies \\citep{priddey03, priddey08, wang07, wang08}, at a time when the Universe was only $\\sim$1\\,Gyr old, suggests that a rapid mechanism for dust production must operate. Type-{\\sc ii} supernovae (SNe) are good candidates for these dust factories as they evolve to a dust-producing phase in only a few hundred Myr and contain a high abundance of heavy elements. Theory predicts that each Type-{\\sc ii} SNe should produce $\\sim$0.1--1\\,\\Msun\\ of dust \\citep{tf01, nozawa03, bs07} and, if true, this can account for the dust observed at high redshift \\citep{me03, dwek07}.  Evidence for this quantity of dust forming in young SNe in the local Universe has been scant and controversial. Observations from mid-/far-infrared (IR) satellites ({\\em IRAS\\/}, {\\em ISO\\/} and {\\em Spitzer\\/}) detect only warm dust and find orders of magnitude less than predicted by theory \\citep{dwek87, lagage96, douvion01, blair07, borkowski06, williams06, sugerman06, meikle07}, e.g.\\ the maximum warm dust mass inferred for Cas~A from {\\em Spitzer\\/} data is $0.03-0.05$\\,\\Msun\\ \\citep{rho08}. Submillimetre (submm) observations of Cas~A seemed to provide the first evidence for large ($\\sim$2\\,\\Msun) quantities of colder dust manufactured in the supernova explosion \\citep[][hereafter D03]{dunne03}\\defcitealias{dunne03}{D03}. However, subsequent observations of CO and OH towards Cas~A suggested that some or all of the submm emission may originate from dust in a foreground molecular cloud complex rather than the remnant \\citep{krause04, wb05}. Although the level of foreground contamination is highly uncertain -- because of the difficulties in converting a molecular line intensity into a submm flux density -- the molecular data have cast doubt on the idea that significant quantities of dust can form rapidly in SNe.  Another explanation for the submm excess, this time without a significant mass of dusty material, was proposed by \\citet{dwek04}. Elongated iron needles with a much higher submm emissivity than canonical dust can also account for the far-IR and submm spectral energy distribution (SED) of Cas~A. If present in the general ISM, they should contaminate the polarisation signal from the CMB which could have importance consequences \\citep[e.g.][]{bowden04}.  Polarimetry at 850\\,\\mic\\ is extremely challenging but provides one way to test the competing hypotheses for the submm emission in Cas~A: cold dust in Cas~A, foreground contamination, or iron needles. We report such observations here, with the paper organised as follows: \\S2 describes the submm and radio polarimetry; \\S3 presents the results and investigates the robustness of the submm polarimetry; in \\S4, we use the polarimetry to place constraints on the fraction of submm emission arising from the remnant and comment on the implications for grain alignment theories. ",
        "conclusions": "\\label{dust}  Since we cannot explain the submm polarisation as arising from synchrotron emission we will now consider the most likely alternative -- that it is due to dust aligned with the magnetic field in the remnant. We will estimate $p_{\\rm d}$, the polarisation fraction of the dust in the remnant. We know that the submm flux we measure is a combination of radio synchrotron and thermal emission from dust, such that $I_{\\rm s} = I_{\\rm r} + I_{\\rm d}$. Similarly we can write that $Q_{\\rm s} = Q_{\\rm r} + Q_{\\rm d}$ and $U_{\\rm s} = U_{\\rm r} + U_{\\rm d}$. We can also express $Q$ and $U$ in terms of $I$, $p$ and $\\theta$, as $Q=Ip\\cos(2\\theta)$ and $U=Ip\\sin(2\\theta)$. Combining these equations and solving for dust polarisation, $p_{\\rm d}$, gives:  \\begin{equation} \\label{RM} p_{\\rm d} = \\frac{\\left[(I_{\\rm s}p_{\\rm s} - I_{\\rm r}p_{\\rm r})^2 + 2\\, I_{\\rm s}p_{\\rm s}I_{\\rm r}p_{\\rm r}(1-\\cos 2(\\theta_{\\rm s} - \\theta_{\\rm r}))\\right]^{1/2}}{I_{\\rm d}} \\end{equation}  \\noindent which for $\\theta_{\\rm s} = \\theta_{\\rm r}$ simplifies to  \\begin{equation} \\label{dp0} p_{\\rm d} = (p_{\\rm s} - p_{\\rm r})\\frac{I_{\\rm r}}{I_{\\rm d}} + p_{\\rm s} \\end{equation}  The fractional dust polarisation calculated using Eqn.~\\ref{dp0} is unprecedented, with an average $p_{\\rm d}= 30\\pm 2$ per cent. A histogram of $p_{\\rm d}$ is shown in Fig.~\\ref{phistdF}. Eqn.~\\ref{dp0} assumes that the polarised synchrotron and dust radiation is emitted with the same intrinsic PA (i.e.\\ aligns with the $B$ field in the same way) and that any rotation between 5-GHz and 850-\\mic\\ vectors is due to Faraday rotation at 5\\,GHz. This is a reasonable assumption given Fig.~\\ref{PAF}. We obtain a slightly higher value for $p_{\\rm d}$ (33 per cent) if we use the difference between the RM corrected 5-GHz PA and the submm PA in Eqn.~\\ref{RM} instead.  \\begin{figure} \\includegraphics[width=8.5cm]{phist_dust.ps} \\caption{\\label{phistdF} Histogram of dust polarisation fraction, $p_{\\rm d}$, calculated using Eqn.~\\ref{dp0}.} \\end{figure}  We can create a map of 850-\\mic\\ dust flux density, $I_{\\rm d}$, by scaling the 5-GHz map to 850\\,\\mic\\ as described in \\S~\\ref{submmradio}. Subtracting the scaled synchrotron map from the 850-\\mic\\ image leaves the 850-\\mic\\ emission due to dust. This dust map\\footnote{This map differs slightly from that in \\citetalias{dunne03} as we have used a more accurate, spatially varying spectral index to extrapolate the radio synchrotron flux to the submm waveband.} is shown in Fig.~\\ref{dustpolF}, with the dust polarisation calculated from Eqn.~\\ref{dp0}. Over-plotted are CO(2--1) contours from \\citeauthor{eales} (in preparation). This shows three things: first, there is dust emission in the northern part of the remnant, but no corresponding CO emission; second, the level of submm polarisation anti-correlates with the CO contours; the dust polarisation being lowest at the western peak, consistent with the suggestion by \\citet{krause04} and \\citet{wb05} that its submm emission is contaminated by foreground clouds; third, there is polarised dust emission in both the north and the west.  The polarisation signal is correlated with the magnetic field in the remnant, which means that some of the emission in the western region must be associated with the remnant. None of dust intensity or polarised flux in the north is associated with molecular material. Thus the finding of \\citet{krause04} that there is `no cold dust in Cas~A' is demonstrably incorrect. The average dust polarisation around the western peak is $13.7\\pm 1.4$ per cent; in the north, away from the CO, it is $39.4\\pm 3.2$ per cent. If we assume that the intrinsic value of $p_{\\rm d}$ is similar throughout the remnant, we can use the uncontaminated values in the north to estimate the contribution of the foreground material to the western peak. Since $p \\propto 1/I$, we simply correct the dust flux in the western peak by the ratio $13.7/39.4 = 0.34$, which gives a flux density of $\\sim 0.8\\,$Jy intrinsic to Cas A at the Western peak. In order to make a comprehensive correction for the intervening foreground material a full set of molecular gas tracers must be observed over the whole remnant. This data set does not yet exist and therefore we cannot make a definitive measurement of the total submm flux associated with the remnant at this time.  We can, however, estimate a conservative lower limit by assuming that only the polarised flux density at 850\\,\\mic ($\\sim 30$ percent) is intrinsic to the remnant. The integrated flux density from the dust map in Fig.~\\ref{dustpolF} is $S_{850}= 20.1$\\,Jy\\footnote{This is higher than the value of 15.9 Jy quoted in \\citetalias{dunne03} due to the lower synchrotron contribution at 850\\,\\mic\\ when using a varying radio spectral index}, and taking the average 30 per cent polarised fraction as the minimum gives us a total lower limit on the dust flux for Cas~A at 850\\,\\mic\\ of 6.0\\,Jy, based on our submm polarimetry.\\/  To convert this flux density into a dust mass we must assume a value for the dust mass absorption coefficient $\\kappa_{850}$. In our previous work \\citepalias{dunne03} we used a value which was derived from laboratory studies of cosmic dust analogues which were amorphous, non-spherical or aggregates. This gave a high value of $\\kappa_{850}=0.76\\,\\rm m^2\\,kg^{-1}$. If we apply this value to our flux measurement and use a temperature of $T_d=20$ K (which comes from fitting the SED from the IR-submm) we get a dust mass estimate for Cas~A of $\\sim 1.0$\\,\\Msun. This value is consistent with theoretical predictions for supernova dust yields \\citep[e.g.][]{tf01,nozawa03,bs07} and sufficient to explain the dust at high redshift \\citep{me03, dwek07}. This is also an upper limit to the possible mass of dust in Cas~A given the masses of condensible elements formed in supernovae with progenitor masses in the range applicable to Cas~A (13-20 \\Msun -- as suggested by the recent determination of the type of Cas~A as a IIb by \\citet{krause08}). According to nucleosynthetic models by \\citet{lc03}, supernovae with progenitors in this mass range could produce 0.5-1.0 \\Msun of dust if the condensation efficiency were 100 per cent.  We can also look at new values for $\\kappa_{850}$ which have been derived from models of dust formation in supernovae similar to Cas~A. Values range from $\\kappa_{850}=0.2\\,\\rm m^2\\,kg^{-1}$ from \\citet{bs07} to $\\kappa_{850} = 0.049 \\,\\rm m^2\\, kg^{-1}$ for a Cas~A-like case from \\citep[][in preparation]{nozawa03, kozasa}. These $\\kappa$ values produce dust masses in the range 3.8--15\\, \\Msun. Given the above discussion on the available mass of condensible elements, these values are not physically plausible. This suggests that the dust which is emitting the polarised radiation at submm wavelengths is not currently predicted by the dust formation models and is not necessarily the same as the dust which is emitting at the shorter FIR wavelengths seen by {\\em Spitzer\\/} and {\\em IRAS\\/}. The supernovae formation models currently only deal with spherical grains, but in order to produce polarised emission the grains must be elongated to some degree. Further development of the dust formation models may be required in order to produce grains which can reproduce the quantity and polarised nature of the submm emission in Cas~A. In order to have a physically plausible mass of dust the emissivity of the dust must be high - at least as high as that found in the laboratory studies of amorphous and aggregate grains.\\/  \\begin{figure} \\includegraphics[width=8.5cm]{dustpol_COcont.ps} \\caption{\\label{dustpolF} Synchrotron-subtracted 850-\\mic\\ image, showing the location of the dust. Overlaid are the dust polarisation vectors calculated using Eqn.~\\ref{dp0}. Also shown in blue are CO(2--1) contours, at the same resolution \\citep[][(in preparation)]{eales}. CO contours coincide with the western dust peak, but the dust polarisation at these positions indicates that some of the 850-\\mic\\ emission must be due to dust in the remnant. The decrease in dust polarisation at the CO peak in the west is consistent with some fraction of the submm emission at this location being due to foreground material, associated with the CO. Note that there is no CO emission coincident with the dust emission in the north.} \\end{figure}  The highly elongated iron needle grains proposed by \\citeauthor{dwek04} alleviate the emissivity issue, as for long enough axial ratios they have emissivities many orders of magnitude higher than spheroidal non-conducting grains. Graphite whiskers may also provide a similar solution, especially in light of their recent discovery in meteorites \\citep{fries08}. Iron or graphite needles could reduce the dust mass required although a full analysis of the properties of such grains which match the new observations is beyond the scope of this paper. We therefore do not rule out such exotic grains and will consider them in more detail in a future paper \\citep[][in preparation]{gomez}.\\/  \\subsection{Grain alignment}  What does this polarisation fraction tell us about the properties of the dust in Cas~A? Average polarisation fractions in typical interstellar and molecular clouds are of the order 2--7 per cent \\citep{hildebrand00, matthews02, cc07} though some clouds have values as high as 15--20 per cent \\citep{benoit04}. We are unaware of any measurements as high as those quoted here, which suggests that either the dust or the alignment mechanisms in Cas~A differ - perhaps unsurprisingly -- from those in the general ISM.  Observations of polarised starlight and thermal emission from dust mean that there must be a population of dust grains which are non-spheroidal and a mechanism capable of aligning the grains so that an appreciable polarisation signal can be detected. The details of how interstellar dust aligns so efficiently, despite impacts by gas atoms acting to randomise their orientations, has kept theorists busy for half a century. The original proposal by \\citet{dg50, dg51} suggested paramagnetic alignment, where the axis of rotation of the slightly elongated grains becomes aligned with the $B$ field though magnetic dissipation. This theory was unable to explain the polarisation levels observed in the ISM and star-forming regions because the timescales for alignment are longer than the timescales for randomisation by collisions with the gas.  There have been many significant improvements to the theory of grain alignment in recent years -- see \\citet{lazarian07} for an excellent review. Several mechanisms are now thought to operate in different astrophysical environments. Briefly these comprise radiative torques \\citep{dolginov72, dw96, lh07b, hl08a}, mechanical alignment in a supersonic or subsonic gas flow \\citep{gold52, rhm95, lazarian97, yl03, lh07a} and super-paramagnetic alignment \\citep{lh08}.  The strong and turbulent magnetic field in Cas~A \\citep[$B\\sim 0.5$\\,mG --][]{wright99}, together with the large abundances of heavy elements and the presence of a hot X-ray plasma, means that certain mechanisms may be particularly applicable to the environment of Cas~A. Pinwheel torques arising from the interaction of grains and electrons in a hot plasma \\citep{hl08b}, supersonic mechanical alignment, sub-sonic mechanical alignment through MHD turbulence \\citep{yl03} or grains with super-paramagnetic (e.g.\\ Fe) inclusions \\citep{lh08} can all lead to highly efficient alignment. The timescales they have to operate is small in astrophysical terms, as the explosion of Cas~A occurred only $\\sim$300 years ago. We note that the low synchrotron polarised fraction indicates a turbulent and disordered magnetic field on sub-arcsecond angular scales. This acts to reduce the observed polarisation as random orientations are averaged within the synthesised beam. If the dust grains were aligned with this tangled small-scale magnetic field then we might expect a similar beam depolarisation effect to operate at submm wavelengths. However, this would give rise to unphysical intrinsic dust polarisation fractions and so we believe that the alignment mechanism must be operating on larger scales. The radial $B$ field in Cas~A is ordered on large scales and thought to arise from Raleigh-Taylor instabilities at the reverse shock boundary \\citep{gull75}. These large-scale radial motions may also be responsible for aligning the dust in Cas~A.  Even if full alignment is theoretically possible, the line of sight angle to the B-field in Cas~A together with the random alignment of orientations along the line of sight and within the beam should all act to reduce the measured polarisation. That we still observe such a high $f_{pol}$ indicates that the grains themselves are emitting strongly polarised radiation. For ideal alignment and line-of-sight viewing conditions $f_{pol} \\sim 30$ percent can be acheived with only moderately non-spherical grains of astronomical silicate, with axial ratios of between 1.4 -- 1.7 for oblate and prolate grains respectively \\citep{padoan01}. Thus the level of polarisation {\\em per se\\/} is not, on its own, strong evidence for needle like grains. An estimate of the aligned fraction would be required in order to constrain the shape of the grains.\\/  A full application of the various grain alignment models to Cas~A is beyond the scope of this paper but these data should provide an interesting test of aligment theories and in time should also shed light on the nature of the grains responsible for the polarised submm emission."
    },
    "0809/0809.2620_arXiv.txt": {
        "abstract": "We consider the emission processes in the large-scale jets of powerful quasars based on the results obtained with the VLA, {\\it Spitzer}, {\\it Hubble}, and {\\it Chandra}. We show that two well-known jets, 3C~273 and PKS~1136$-$135, have two distinct spectral components on large-scales: (1) the low-energy (LE) synchrotron spectrum extending from radio to infrared, and (2) the high-energy (HE) component arising from optical and extending to X-rays. The X-ray emission in quasar jets is often attributed to inverse-Compton scattering of cosmic microwave background (CMB) photons by radio-emitting electrons in a highly relativistic jet. However, recent data prefer synchrotron radiation by a second distinct population as the origin of the HE component. We anticipate that optical polarimetry with {\\it Hubble} will establish the synchrotron nature of the HE component. Gamma-ray observations with \\emph{GLAST} (renamed as the \\emph{Fermi} Gamma-ray Space Telescope), as well as future TeV observations, are expected to place important constraints on the jet models. ",
        "introduction": "The emission processes responsible for the spectral energy distributions (SEDs) of large-scale quasar jets, constructed using \\emph{Spitzer}, \\emph{HST} and \\emph{Chandra} data, are the subject of active debate \\cite{HK06}. The X-ray intensity relative to the radio synchrotron flux is generally too high to be explained by synchrotron-self-Compton \\cite{Cha00}. The SED at radio, optical, and X-ray traces an inflected shape, which disfavors the interpretation of X-rays as due to synchrotron radiation from a single population of  electrons.  Inverse-Compton (IC) scattering of CMB photons by high-energy electrons (with electron Lorentz factor of $\\gamma_e \\sim 30$) in a highly relativistic jet with bulk Lorentz factor $\\Gamma \\sim 10$ initially seemed an attractive way to explain the observed X-ray emission \\cite{Tav00,CGC01}, but this process is also not free of problems \\cite{AD04}. Finally, the X-rays may arise from synchrotron radiation from extremely energetic protons \\cite{Aha02}. Determining which of these emission mechanisms produces the observed X-ray jets in powerful quasars is a strong motivation for more observations of radio-loud quasars, and has resulted in a rapid increase in the number of known X-ray jets.  In this work, we highlight new results based on the multiwavelength data from the \\emph{VLA}, \\emph{Hubble}, \\emph{Spitzer} and \\emph{Chandra}, aiming at identifying the radiation mechanisms operating in powerful quasar jets. The jets in the nearest quasar 3C~273 ($z=0.158$) and  in the lobe-dominated quasar PKS~1136$-$135 ($z=0.554$) are selected for our detailed analysis. ",
        "conclusions": "\\subsection{LE component}  Let us deduce some physical parameters of the LE spectral component, which is synchrotron radiation by relativistic electrons in the jet. The observed jet radiation is enhanced by  Doppler beaming. (The Doppler factor is defined as $\\delta \\equiv [\\Gamma(1-\\beta\\cos\\theta)]^{-1}$ with $\\beta c$ the velocity of the jet, $\\Gamma = (1-\\beta^2)^{-1/2}$ the bulk Lorentz factor of the jet, and $\\theta$ the observing angle with respect to the jet direction.)  The LE synchrotron emission has a high-energy cutoff at $\\nu_c \\simeq 5\\times 10^{13}$ Hz, presumably associated with the maximum energy of the electron population. The maximum energy can be deduced as $E_{\\rm max} \\sim 0.2\\ (B_{-4}\\delta)^{-0.5}\\ \\rm  TeV$, where  $B_{-4} = B/(0.1\\ \\rm mG)$ is the comoving magnetic field strength. The equipartition magnetic field is $B \\delta \\sim 0.1$ mG for 3C~273 and $B \\delta \\sim 0.7$ mG for PKS~1136$-$135.  \\begin{figure}[H] % \\centerline{\\psfig{file=f3.eps,width=8cm}} \\caption{SEDs of the large-scale jets in 3C~273 (taken from Fig.~\\ref{fig:3c273SED}) and PKS~1136$-$135 (black filled circles: Ref.~\\refcite{Uch07}). The jet in PKS~1136$-$135 (fluxes extracted from knots A to E) has a two-component SED similar to 3C~273. The curves represent synchrotron radiation models.} \\label{fig:pks1136SED} \\end{figure}   By extrapolating VLA radio flux of knot H1/H2 in 3C~273 to lower frequencies (10--100 MHz), we found that a low energy   ``cutoff\"  or ``break\" should be present as well. Otherwise, the jet radio flux around $\\sim 100$ MHz  exceeds  even the total radio flux of 3C~273. It would be reasonable to assume that such a cutoff/break is common to all the knots. We adopt $E_{\\rm min} \\sim 1000 m_ec^2$ in Fig.~\\ref{fig:3c273SED} to introduce the low energy cutoff. Also, the observed synchrotron spectra from the radio to optical frequencies are likely to be formed in a  ``cooling regime\", and therefore a cooling break may appear in low-frequency radio bands.  Using the relation $s_{\\rm LE} = 2\\alpha_{\\rm LE}$, where $s_{\\rm LE}$ is acceleration index, the observed index of $\\alpha_{\\rm LE} \\simeq 0.7$ is translated into the electron acceleration index of $s_{\\rm LE} \\simeq 1.4$, close to the hardest possible index  ($s=1.5$) in the nonlinear shock acceleration theory. Here we assumed particle acceleration taking place in each knot, where relativistic electrons are accumulated over a certain timescale of $10^5$ to $10^7$ yr. Assuming that the number of radio-emitting electrons is balanced by the number of cold protons in the jet, we estimated a jet kinetic power of $L_{\\rm kin} \\sim 5\\times 10^{44}\\ \\rm erg\\ s^{-1}$ with $\\delta = \\Gamma = 4$ and the equipartition magnetic field.  \\subsection{HE component: IC/CMB or synchrotron?}  The beamed IC scenario \\cite{Tav00,CGC01} was initially considered to be an attractive explanation of the X-ray-dominated component of the SED for jet knots in radio-loud quasars, like 3C~273 (Ref.~\\refcite{Sam01}) and PKS~1136$-$135 (Refs.~\\refcite{Sam06a,Tav06}). We argued for a direct connection between the optical and X-ray fluxes \\cite{Uch06}. The same mechanism should be responsible for the optical and X-ray emissions. If one adopts the beamed IC model for the X-rays, the energy distribution of electrons has to continue down to very low energies of $E_{\\rm min} \\sim m_e c^2$ to account for the optical emission.  In the framework of the beamed IC radiation, it is interesting to note that the jet exhibits the evolution of multiwavelength emission along the jet as shown in Fig.~\\ref{fig:3c273image}. The increase of radio brightness is accompanied by the decrease of X-ray brightness towards downstream, which can be understood in terms of deceleration of the jet \\cite{GK04,Sam06a,Tav06}. If the beamed IC scenario is correct, we are able to study the flow structure of the large-scale jets.  New results indicate, however, that the beamed IC scenario faces several difficulties. First, in the outer knots of the 3C~273 jet, the radio spectral index of $\\alpha_{\\rm LE}\\simeq 0.7$ disagrees with  the HE spectral index of $\\alpha_{\\rm HE}\\simeq 1.0$ (see Refs.~\\refcite{Uch06,Jes06}). Second, the parameters required in the beamed IC/CMB model are often too extreme. In the case of 3C~273, we obtained a Doppler factor of $\\delta = 30$ for knot A and a jet kinetic power of $L_{\\rm kin} \\simeq 1\\times 10^{48}\\ \\rm erg\\ s^{-1}$. The obtained Doppler factor seems too large for 3C~273, since relatively large contributions from the un-beamed components are observed in the core emission (see Fig.~\\ref{fig:3c273SED}) and accretion-disk emission appears to be comparable to the jet emission in the X-ray regime \\cite{GP04}. In the case of  PKS~1136$-$135, we obtained a Doppler factor of $\\delta = 19$, which is also difficult to reconcile with the lobe-dominated nature of the quasar.  Instead of the beamed IC, synchrotron radiation by a second population of high-energy electrons may account for the HE component. Then the spectral index of the HE component, $\\alpha_{\\rm HE} \\simeq 0.5\\mbox{--}1.0$ in 3C~273 and PKS~1136$-$135, corresponds to $s_{\\rm HE} \\simeq 1.0\\mbox{--}2.0$. However, $s_{\\rm HE} \\simeq 1.0$ would be incompatible with diffusive shock acceleration theory. This brings into question the idea that the second electron population is accelerated through diffusive shock acceleration. The second synchrotron component may instead be due to turbulent acceleration operating in the shear layers \\cite{SO02}. Unlike the shock acceleration, a hard spectrum, $s < 1.5$, can be expected to form in the case of turbulent acceleration (i.e., second-order Fermi acceleration). This scenario should be tested by future TeV observations (see below).  \\subsection{GLAST (Fermi-LAT) observations}  As detailed in Ref.~\\refcite{Geo06}, GLAST measurements of the quasar jets will provide us with new  information (or constraints) about the jet emission. The infrared-emitting electrons in a relativistic  jet inevitably emit GeV $\\gamma$-rays through an IC/CMB process (see Fig.~\\ref{fig:3c273SED}). The total radiative output is dominated by the multi-GeV $\\gamma$-rays and their predicted flux with $\\delta \\sim 7$ exceeds  the 1-yr sensitivity offered by the GLAST LAT in the case of 3C 273. Also, the observations of  PKS~1136$-$135 with GLAST may be able to test the IC/CMB hypothesis, since the beamed IC model with $\\delta \\sim 20$ predicts the observable GeV flux \\cite{Uch07}. Note that steadiness and hard spectrum can distinguish the large-scale emission from the  core emission.  Moreover, if the HE fluxes are due to synchrotron radiation by electrons, the UV/X-ray-emitting electrons in a relativistic  jet produce IC-upscattered TeV $\\gamma$-rays. This offers a potential diagnostic tool to distinguish the origin of the HE component.  \\subsection{HST polarimetry}  Finally we emphasize  that optical polarimetry can be an effective way of discriminating the radiation models responsible for the optical-to-X-ray emission of  the jets. In the beamed IC interpretation, the optical and X-ray emission are due to Compton up-scattering off the amplified CMB by high-energy electrons of $\\gamma_{e} \\sim 3$ (optical) and $\\gamma_{e} \\sim 100$ (X-ray). Unlike in the case of synchrotron models, the X-rays are expected to be {\\it unpolarized} and the optical light is nearly unpolarized at most a few percent of polarization (see Fig.~\\ref{fig:pol}). Precise polarization measurements in the optical can in principle verify (or discard) the beamed IC model. Unfortunately, there have been no useful polarization observations of quasar jets with {\\it HST} so far. Only for 3C 273, early {\\it HST} polarimetry of the jet was done but with the pre-COSTAR FOC and low significances \\cite{Tho93}. Recently, \\emph{HST} polarimetry of PKS~1136$-$135 has been approved (PI: Eric Perlman). New polarimetry of quasar jets on large-scales will provide  important clues as to the origin of the HE component.  \\begin{figure}[H] % \\centerline{\\psfig{file=f4.eps,width=8cm}} \\caption{The polarization degree in the synchrotron$+$IC model of knot A in the 3C 273 jet. The Doppler and Lorentz factor of a jet are $\\delta = \\Gamma = 15$. The minimum electron energy: $E_{\\rm min} = 2m_ec^2$.  } \\label{fig:pol} \\end{figure}"
    },
    "0809/0809.2416_arXiv.txt": {
        "abstract": "We investigate the impact of the star formation efficiency in cluster-forming-cores (i.e. local SFE) on the evolution of the mass in star clusters over the age range 1-100\\,Myr, when star clusters undergo their infant weight-loss/mortality phase. Our model builds on the $N$-body simulation grid of \\citet{bau07}. Assuming a constant formation rate of gas-embedded clusters and a weak tidal field, we show that the ratio between the total mass in stars bound to the clusters over that age range and the total mass in stars initially formed in gas-embedded clusters is a strongly increasing function of the averaged local SFE, with little influence from any assumed core mass-radius relation.  Our results suggest that, for young starbursts with estimated tidal field strength and known recent star formation history, observed cluster-to-star mass ratios, once corrected for the undetected clusters, constitute promising probes of the local SFE, without the need of resorting to gas mass estimates.  Similarly, the mass ratio of stars which remain in bound clusters at the end of the infant mortality/weight-loss phase (i.e. age $\\gtrsim 50$\\,Myr) depends sensitively on the mean local SFE, although the impacts of the width of the SFE distribution function and of the core mass-radius relation require more careful assessment in this case.  Following the recent finding by \\citet{bas08} that galaxies form, on the average, 8\\,\\% of their stars in bound clusters regardless of their star formation rate, we raise the hypothesis that star formation in the present-day Universe is characterized by a near-universal distribution for the local SFE.  A related potential application of our model consists in tracing the evolution of the local SFE over cosmological lookback times by comparing the age distribution of the total mass in star clusters to that in field stars in galaxies where field stars can be resolved and age-dated.  We describe model aspects which are still to be worked out before achieving this goal. ",
        "introduction": "Star formation efficiencies (SFE) are not easy to come by observationally since they require reliable determinations of both the gaseous and stellar masses of a star forming region.  Yet, the SFE is a fundamental parameter of star and star cluster formation processes and of the early evolution of star clusters.  {\\it Local} SFEs, i.e. the mass fraction of a giant molecular cloud core turned into stars, have been estimated for a handful of Galactic nearby gas-embedded clusters~\\citep{ll03}.  They range from 10 to 30 per cents, which is an order of magnitude higher than what is inferred when the SFE is averaged over an entire giant molecular cloud (the so-called {\\it global} SFE).  Attempts to compare SFE estimates in the Milky Way and its neighbours, where they are best studied, to those in major merger-induced starbursts, involves the additional difficulty of scale differences.  While for the former, SFEs are defined down to individual giant molecular clouds (GMC) and their dense star forming cores, only large-scale -- often galaxy-wide -- SFEs are measured for the latter owing to their large distances. Estimates of SFEs over up to kpc size nuclear regions in Ultra Luminous IR Galaxies (=ULIRGs), all of which are late stages of massive gas-rich mergers, are reported to be $1-2$ orders of magnitude higher than  global SFEs measured in spiral galaxies~\\citep[][their fig.~2b]{gao04}.  Because the formation of compact and massive star clusters is observed to be closely associated to active star forming environments, prominent star formation episodes in the history of galaxies are best traced by the age distribution of their star clusters~\\citep{West+04,Fritze04,BrodieStrader06,Kotulla+08}. Although the vast majority of stars actually form in gas-embedded star clusters, galaxies are predominantly made of field stars.  In this respect, a peak in a star cluster age distribution only tells us {\\em when} an epoch of enhanced star cluster formation took place, not yet {\\em how much total stellar mass was involved}.  To derive the total mass in stars formed in a starburst older than $\\sim 3-4$ Gyr from the integrated spectrum of a galaxy is, in the present state-of-the-art, hopeless~\\citep{FritzeLilly07}. Knowing the mass fraction of all star formation still in bound clusters as a function of star cluster age, by properly accounting for cluster destruction and mass loss, would allow us to carry the analysis of galaxies'evolution one step further and to estimate the history of their overall star formation rate. We now have a fairly mature view of how to decipher the rates of {\\it bound} cluster formation and evolution of a given star cluster system, in essence based on the cluster age and mass distributions~\\citep{par08a}.  That is, we are able to assess the amount of field stars arising from the secular evolution (i.e. the preferential depletion of low-mass stars due to tidal stripping and internal two-body relaxation) of star clusters which survived the expulsion of their residual star forming gas.  As pointed out by \\citet{mcl99}, however, secular evolution does not necessarily affect the total mass in star clusters significantly.  As for the star clusters of the Galactic Old Halo and of the Large Magellanic Cloud for instance, the observed total mass in clusters is less than an order of magnitude smaller than their total mass at the onset of secular evolution~\\citep[see table 3 and table 2 in][respectively]{par05, par08a}.  If the vast majority of stars form in gas-embedded star clusters, this implies that the bulk of field stars is given off at an earlier stage, very likely as a result of cluster dynamical response to the expulsion of the residual star forming gas out of their parent cores due to massive star activity.  Because of its resulting shallower potential, an exposed cluster attempts to reach a new equilibrium, which results into the loss of stars (``infant weight-loss``) and, when the local SFE is too low, the eventual disruption of the cluster (``infant mortality``)~\\citep[see Fig.~1 in ][]{goo06, goo08}.  This violent relaxation phase, which lasts for at most 50\\,Myr, appears therefore as the prime driver of the ratio of the total mass in bound clusters to the total mass in stars.  In what follows, we assume that all stars form in gas-embedded clusters, i.e. we assume that there is no mode of field star formation or loose stellar association formation.  We will further comment about this hypthesis in our discussion (section 3).  Young starbursts are the best places to infer the impact of cluster violent relaxation on the cluster-to-star mass ratio.  Based on {\\sl HST} $UV$ imaging of a sample of starburst galaxies, the morphology of which ranges from blue compact dwarfs to ultra-luminous merging far-infrared galaxies, \\citet{meu95} find that, on average, about 20 per cents of the $UV$ luminosity comes from star clusters.  \\citet{deg03} study the Mice and Tadpole interacting systems by means of a pixel-by-pixel analysis of their color-magnitude and color-color diagrams. In the Mice interacting galaxies, they find that star clusters are predominantly associated with regions of active star formation, where they contribute $\\geq 40$\\,per cents of the total flux, irrespective of wavelength (but see our discussion section for caveats regarding these values). Similar flux ratios are inferred for the Tadpole system, an example of a galaxy encounter between two unequal-mass galaxies. These high luminosity fractions, presumably representative of large mass fractions, demonstrate that star clusters contribute significantly to the total stellar light given off by young starbursts.  While being higher than the initial Galactic halo globular cluster mass fraction as compared to halo field stars, i.e. $\\simeq 0.1$, these ratios are also conspicuously smaller than unity, thus confirming that the dissolution of star clusters takes place on a short time-scale.  Analyses of star cluster systems on the basis of multi-band imaging are feasible out to Virgo cluster distances and have the potential to provide a comprehensive view of the individual star formation, chemical enrichment and mass assembly histories of their parent galaxies over the past Hubble-Time, provided that the last missing issue, establishing the ratio of star formation ending up in bound star clusters, is solved. This is what we attempt here, that is, we investigate what is/are the driving parameter(s) behind the mass fraction of clusters as compared to field stars.  Specifically, we build on the $N$-body simulation grids recently obtained by \\citet{bau07} to compute the evolution over the age range 1-100\\,Myr of the mass in clusters and to compare it with the overall star formation rate arising from the same star formation episode.  We make predictions about the cluster formation efficiency (i.e. the ratio between the total mass in clusters over the age range of relevance and the total mass in star forming gas cores) and the ratio between the total mass in clusters and the total mass in stars.  And we highlight their relation to the local SFE. ",
        "conclusions": "Building on the $N$-body simulation grid recently obtained by \\citet{bau07}, we have performed detailed simulations of the early (i.e. up to an age of 100\\,Myr) evolution of star cluster systems.  In this first contribution, we have adopted the following assumptions: star forming cores with a power-law mass spectrum of slope $-2$, constant gas-embedded cluster formation rate, Gaussian distribution function of the local SFE and weak external tidal field.  Besides, star formation is assumed to mostly take place in gas-embedded clusters.  Under these assumptions, we have highlighted the strong impact of the local SFE on the temporal evolution of the mass in clusters and on the ratio between the total mass in clusters and the total mass in stars.  Within less than 10\\,Myr after gas expulsion, the age distribution of the mass in clusters decreases as a result of cluster violent relaxation and the amplitude of the decline depends strongly on the adopted SFE distribution function.  This is because the fraction of stars remaining bound to a cluster after its violent relaxation is sharply increasing with the SFE $\\epsilon$ once $\\epsilon > 0.33$~\\citep[Fig.~1 in][]{par07}.  The ratio $M_{cl}/M_{st}$ between the total mass in stars still bound to the clusters and the total mass in stars formed in gas-embedded clusters over the age range of $1-100$\\,Myr constitutes a sensitive probe of the local SFE which characterized their progenitor cores.  Besides, this ratio is only weakly affected by the choice of a core mass-radius relation. We caution however that {\\it (i)} observationally derived cluster-to-star mass ratios account for detected clusters only; {\\it (ii)} ratios displayed in Table \\ref{tab:clusterfrac} will be different for other tidal field strengths and non-constant gas-embedded cluster formation history. Measuring the flux ratio given off by star clusters in a young starburst and converting it in a cluster-to-star mass ratio therefore constitutes a promising way of inferring the local SFE, provided that the observed $M_{cl}/M_{st}$ is turned into the actual one and that time-varying gas-embedded cluster formation rate and stronger tidal fields are -- if necessary -- accounted for via generalized versions of the present model.  Another potential application of our model is to trace the evolution of the local SFE over cosmological lookback times.  This can be done, for galaxies close enough so that their field stars can be resolved and age-dated, by comparing the history of the bound cluster formation rate (i.e. the age distribution of the cluster mass at the end of violent relaxation, obtained by correcting the observed age distribution for secular evolutionary mass losses and fading effects, see e.g. \\citet{par08a}) on the one hand and the history of the SFR on the other hand.  This bound cluster-to-star mass ratio is our bound fraction $F_{bound}$ (see equation \\ref{eq:fbound}), equivalent to model outputs over the age range 50-100\\,Myr.  This application, however, requires tighter constraints on the slope of the core mass-radius relation, the width of the local SFE distribution and an estimate of the external tidal field strength.  Combined with the recent finding of \\citet{bas08} that, on the average, galaxies form 8\\,\\% of their stars in bound star clusters, the relation between the bound fraction $F_{bound}$ and the local SFE $\\epsilon$ constitutes a hint towards an almost universal local SFE in the present-day Universe.  In a forthcoming paper, we will apply our model to the observed relation between the present-day SFR of galaxies and the absolute visual magnitude of their brightest young clusters so as to shed better constraints on the present-day local SFE distribution function."
    },
    "0809/0809.4006_arXiv.txt": {
        "abstract": "The Orion Molecular Cloud 2/3 region (hereafter, OMC-2/3) and the reflection nebula NGC~1977 encompass a section of the Orion A molecular cloud undergoing vigorous star forming activity.  One of the richest assemblages of protostars in the nearest 500 pc is seen in OMC-2/3, while NGC 1977 contains a cluster of over 100 young stars. In this review, we present a census of the protostars, pre-main sequence stars, and young brown dwarfs in these regions.  These are identified through sub-millimeter surveys, far-red to near-infrared imaging and spectroscopy with ground-based telescopes, mid-infrared photometry from the {\\it Spitzer Space Telescope}, and X-ray observations made with the Chandra X-ray Observatory.  We present an overview of the distribution of molecular gas associated with these regions and the rich complex of shock heated nebulae created by the young stars interacting with the molecular gas.  Finally, we discuss the relationship of OMC-2/3 and NGC~1977 to the neighboring Orion Nebula Cluster and the Orion OB1 association. ",
        "introduction": "The section of the Orion A molecular cloud directly north of the Orion Nebula is a region of intense star formation activity.  Although the number and density of young stars in this area does not rival those found in the Orion Nebula, it is still one of the most active regions of ongoing star formation within 500 pc of the Sun.  It is typically divided into two regions: the Orion Molecular Cloud 2/3 region, which contains a filamentary molecular cloud rapidly forming new stars, and NGC~1977, a composite reflection and emission line nebula.  In this chapter, we will review the observations which delineate the stellar, sub-stellar, protostellar and gas content in these two regions. We will begin by discussing the distribution of stars, gas and dust in these two regions and their relationship to the Orion Nebula Cluster. ",
        "conclusions": "The northern end of the Orion A molecular cloud is the most active region of low and high mass star formation within 500~pc of our Sun and is one of the most intensively studied star forming regions in the sky. However, most of these observations exclusively target the Orion Nebula and the region immediately surrounding the nebula.  This chapter reviews the rich concentration of pre-main sequence stars and protostars north of the Orion Nebula in OMC-2/3 and NGC~1977.  These regions provide examples of star formation in environments much different than that in the Orion Nebula; in particular, in environments without the extreme-UV radiation found in the Orion Nebula, and showing lower stellar densities than the densities found in the center of the Orion Nebula.  Future studies of the ONC should encompass OMC-2/3 and NGC~1977 in order to fully characterize the wide variety of environments within this single contiguous star forming region and how these environments affect star and planet formation."
    },
    "0809/0809.4230_arXiv.txt": {
        "abstract": "We study the collective kinklike normal modes of a system of several cylindrical loops using the T-matrix theory. Loops that have similar kink frequencies oscillate collectively with a frequency which is slightly different from that of the individual kink mode. On the other hand, if the kink frequency of a loop is different from that of the others, it oscillates individually with its own frequency. Since the individual kink frequency depends on the loop density but not on its radius for typical $1$ MK coronal loops, a coupling between kink oscillations of neighboring loops take place when they have similar densities. The relevance of these results in the interpretation of the oscillations studied by \\citet{schrijver2000} and \\citet{verwichte2004}, in which transverse collective loop oscillations seem to be detected, is discussed. In the first case, two loops oscillating in antiphase are observed; interpreting this motion as a collective kink mode suggests that their densities are roughly equal. In the second case, there are almost three groups of tubes that oscillate with similar periods and therefore their dynamics can be collective, which again seems to indicate that the loops of each group share a similar density. All the other loops seem to oscillate individually and their densities can be different from the rest. ",
        "introduction": "Transverse coronal loop oscillations were discovered by the Transition Region and Coronal Explorer (TRACE) in 1998 \\citep[see, e.g.][]{aschwanden1999,aschwanden2002, nakariakov1999}. These oscillations were initiated shortly after a solar flare that disturbed the loops. Since their first observation, transverse oscillations have been routinely observed and studied. Much before TRACE observations, the theory of loop oscillations was developed \\citep{spruit1981,edwin&roberts1983,cally1986} and the different kinds of oscillations were studied. The observed transverse motions have been interpreted in terms of the fundamental kink mode of the fast magnetohydrodynamic (MHD) oscillation \\citep{nakariakov1999}, which is the only mode that can produce the observed transverse loop displacement.  In many cases the observed coronal loops belong to complex active regions and are not isolated but forming bundles or arcades of loops. For example, in \\citet{schrijver2000} antiphase transverse oscillations of adjacent loops were reported. In addition, in \\citet{verwichte2004} phase and antiphase motions were observed in a post-flare arcade. On the other hand, it is currently debated whether active region coronal loops are monolithic or multistranded \\citep[see, e.g.][]{aschwanden2005, klimchuk2006,deforest2007}. In the multistranded model, it is suggested that loops are formed by several tens or hundreds of strands considered as miniloops for which the heating plasma properties are approximately uniform in the transverse direction \\citep{klimchuk2006}. Most analytical studies about transverse loop oscillations have only considered the properties of individual loops. However, from the information provided by the observations, it is necessary to study not only individual loops but also how several tubes can oscillate as a whole, since their joint dynamics can be different from that of a single loop. Only a few works have considered composite structures. \\citet{berton1987} studied the MHD normal modes of a periodic magnetic medium. \\citet{kris93} and \\citet{kris94} studied numerically the propagation of fast waves in two slabs unbounded in the longitudinal direction. In \\citet{diaz2005} the oscillations of the prominence thread structure were investigated. These authors found that in a system of equal fibrils the only non-leaky mode is the symmetric one, which means that all the fibrils oscillate in spatial phase with the same frequency.  \\citet{luna2006} studied a system of two coronal slabs and found that the symmetric and antisymmetric modes can be trapped. A more complex system of two coronal cylinders was studied in \\citet{luna2008}. Four trapped normal modes were found and the interchange of energy between loops was shown by solving the time-dependent problem. In \\citet{terradas2008} a multistranded loop formed by ten strands was considered. The composite loop oscillates transversely as a whole with a global motion of the strands after an external disturbance. This work shows that the bundle of strands oscillates with a combination of collective modes. On the other hand, an analytical approximation to the normal modes of a loop pair has been carried out by \\citet{doorsselaere2008}. The authors assume the long wavelength approximation and obtain an analytical dispersion relation for two different tubes together with the four kink mode polarizations described in \\citet{luna2008}.  In this work we aim to study the normal modes of a loop set with different physical and geometrical properties by using the scattering theory. The scattering theory, or its matricial formulation called T-matrix theory \\citep[see, e.g.][]{waterman&truell1961,waterman1969,ramm}, was first applied to magnetic tubes by \\citet{bogdan&zweibel1985}. These authors studied the interaction of acoustic plane waves with an ensemble of parallel magnetic fibrils distributed uniformly in the so-called spaghetti sunspot model. The authors derived and solved the dispersion relation in the long wavelength limit. In \\cite{bogdan&cattaneo1989} the frequency shifts and velocity eigenfunctions were calculated for the case of random fibril distributions of up to 100 flux tubes. Many other papers were published studying the cross section of a fibril spot insonified by external acoustic waves \\cite[see][]{bogdan&fox1991,keppens1994}. In all these papers a non-magnetized external medium was considered and the eigenfrequencies and eigenmodes of the acoustic oscillations were obtained.  In this paper we generalize the method to a system with an external magnetized medium, in order to extend previous works to coronal loop conditions. Our model consists of an ensemble of parallel cylinders, without gravity and curvature. We consider uniform magnetic field in the internal loop medium and in the external or coronal medium. This assumption allows the interaction of the tubes through fast MHD waves. In addition, we explicitly calculate the eigenvalues and eigenfunctions of the normal modes of the model.  This paper is organized as follows. In \\S \\ref{model} the loop ensemble model and the equations for its dynamics are presented. In \\S \\ref{normal_modes} we briefly describe the T-matrix theory and apply it to our model. With this theory the exact eigenfrequencies and eigenmodes of two non-identical loops are investigated in \\S \\ref{2loops}. We study the dependence of the interaction with the relative density and radii of the loops. The study of three identical aligned loops is presented in \\S \\ref{3loops}. In the same section the interaction between three non-identical loops is considered. Finally in \\S \\ref{disc_conc} the results are summarized and the main conclusions are drawn. ",
        "conclusions": "\\label{disc_conc}   In this work we have investigated the kinklike normal modes of a system of several loops with the help of the T-matrix theory. The results of this work can be summarized as follows:  \\begin{enumerate}  \\item In the system of two non-identical loops, we have found four kinklike normal modes $P_x$, $AP_y$, $P_y$, and $AP_x$. The frequencies of the $P_x$ and $AP_y$ solutions are very similar as well as the frequencies of the $P_y$ and $AP_x$ modes. This result agrees with \\citet{doorsselaere2008}, who considered thin tubes (i.e. long wavelength approximation). For fat loops the $P_x$ and $AP_y$ modes, as well the $P_y$ and $AP_x$, have different frequencies, as was shown in \\cite{luna2008}.  \\item For a system of two loops we have investigated the dependence of the interaction between kink oscillations as a function on the relative density of the loop pair. For $\\rho_\\mathrm{1}=3 \\rho_\\mathrm{0}$ we have found that the oscillations of the loops are coupled in the range of $\\rho_\\mathrm{2}$ between $2 \\rho_\\mathrm{0}$ to $4 \\rho_\\mathrm{0}$ and that the coupling is maximum at $\\rho_\\mathrm{2} = \\rho_\\mathrm{1}=3\\rho_\\mathrm{0}$. Outside this density range the loops are essentially decoupled and oscillate independently. This is qualitatively similar to the behavior of the anomalous modes described by \\citet{doorsselaere2008}.  \\item We have also studied the dependence of the interaction with the relative radii of the loops. We have seen that in the range of radii for which transverse loop oscillations have been observed the interaction depends very little on this parameter and the loops strongly interact for all the radii considered. The explanation of this behavior is that in our loops the thin loop approximation can be applied and in this situation the kink frequency depends on the tube density and not on the radii.  \\item In the case of a system of three equal, aligned loops there are eight kinklike normal modes. The lower frequency mode corresponds to the three loops oscillating in phase in the $x$-direction, i.e. along the direction in which their axes are aligned, in agreement with the results of two identical loops. On the other hand, the upper frequency mode corresponds to the three loops oscillating in phase in the $y$-direction. This does not agree with the two identical loops situation, in which the upper mode corresponds to the two loops oscillating in antiphase in the $x$-direction. In fact, this property of the three-loop system is also true for ensembles of four or more aligned loops.  \\item We have made a parametric study of the kinklike modes in a system of three loops with equal radii and different densities by changing the density of loop $3$, $\\rho_\\mathrm{3}$. We have chosen $\\rho_\\mathrm{1}=3 \\rho_\\mathrm{0}$ and $\\rho_\\mathrm{2}=2 \\rho_\\mathrm{0}$ so that the interaction between loops $1$ and $2$ is negligible. We have found that the oscillations of loop $3$ are coupled with loop $2$ when $\\rho_\\mathrm{3} \\approx \\rho_\\mathrm{2}$, whereas loop $1$ oscillates independently. Furthermore, loop $3$ couples with loop $1$ when $\\rho_\\mathrm{3} \\approx \\rho_\\mathrm{1}$ with loop $2$ oscillating independently. If $\\rho_\\mathrm{3}$ takes different values, the system is decoupled and the three loops oscillate independently.   \\end{enumerate}  In this work, we have found that the interaction between loops regarding kinklike motions depends strongly on the their individual kink frequencies. If these frequencies are similar, loop motions are coupled and the normal modes are collective. On the other hand, if the loop kink frequencies are quite different their motions are not coupled. Since the individual frequencies depend on the loop density and radius, we have studied separately the influence of the two parameters. We have found that if the densities are quite similar, loops are coupled and the oscillations are collective. On the other hand, if the densities are quite different, the tubes oscillate independently. The range of densities for which the loops are coupled depends on the system properties and in the configuration of three loops this range is narrower than in the two tubes configuration.  From the results shown in this paper we suggest that the antiphase motions reported in \\citet{schrijver2000} and \\citet{schrijver2002} are collective motions and, therefore, that the individual kink frequencies are similar but different from the collective observed frequency. If the loop model presented here is valid, both loop densities are also similar. In addition, in \\citet{verwichte2004} a loop arcade is studied and three groups of tubes oscillating with similar frequencies can be appreciated. The dynamics of each group of tubes can be interpreted as collective, although a detailed study of such configuration is needed to relate the loop characteristics and the frequency of oscillation of the group. On the other hand, loops not belonging to these three groups do not share their frequencies with other loops, and so oscillate independently. This has to be interpreted as a sign that these loops have different densities from those of the rest of the loops. It must be mentioned that in \\citet{verwichte2004} all the oscillations are assumed as individual, but this is only true in the case of loops that do not share their frequency. If this assumption is applied to loop with a collective behavior it produces wrong results for the loop parameters. For example, if the loops actually oscillate with the lowest frequency collective mode, the assumption of these authors produces an underestimation of the magnetic field or an overestimation of the loop density.  The T-matrix method shown in this paper can be easily applied to more complex configurations with gas pressure and tubes with flows or more complex systems of loops, i.e. arcades with myriads of loops or multistranded loops. It is expected that in such systems loops oscillate essentially independently except for loops with similar individual oscillation frequencies."
    },
    "0809/0809.1881_arXiv.txt": {
        "abstract": "We analyze two pre-supernova (SN) and three post-SN high-resolution images of the site of the Type II-Plateau supernova SN~2006my in an effort to either detect the progenitor star or to constrain its properties.  Following image registration, we find that an isolated stellar object is not detected at the location of SN~2006my in either of the two pre-SN images.  In the first, an $I$-band image obtained with the Wide-Field and Planetary Camera 2 on board the {\\it Hubble Space Telescope}, the offset between the SN~2006my location and a detected source (``Source 1'') is too large: $\\geq 0.08^{\\prime\\prime}$, which corresponds to a confidence level of non-association of $96\\%$ from our most liberal estimates of the transformation and measurement uncertainties.  In the second, a similarly obtained $V$-band image, a source is detected (``Source 2'') that has overlap with the SN~2006my location but is definitively an extended object.  Through artificial star tests carried out on the precise location of SN~2006my in the images, we derive a $3\\ \\sigma$ upper bound on the luminosity of a red supergiant that could have remained undetected in our pre-SN images of $\\log L/L_\\odot = 5.10$, which translates to an upper bound on such a star's initial mass of $15 {\\rm\\ M}_\\odot$ from the STARS stellar evolutionary models.  Although considered unlikely, we can not rule out the possibility that {\\it part} of the light comprising Source 1, which exhibits a slight extension relative to other point sources in the image, or {\\it part} of the light contributing to the extended Source 2, may be due to the progenitor of SN~2006my.  Only additional, high-resolution observations of the site taken after SN~2006my has faded beyond detection can confirm or reject these possibilities. ",
        "introduction": "\\label{sec:1}  The most common class of core-collapse supernovae (SNe) displays a distinct plateau in its optical light curve, and is therefore dubbed Type II-Plateau (II-P; see \\citealt{Filippenko97} for a review of SN classifications).  This type of stellar explosion has long been thought to result from the core collapse and subsequent envelope ejection of isolated red supergiant (RSG) stars, but it is only in recent years that direct observational evidence of the progenitor-SN~II-P connection has begun to accumulate \\citep{Vandyk03,Smartt04,Maund05,Li06a,Eldridge07,Schawinski08}.  By registering pre-SN and post-SN images, usually taken at high resolution using either space-based optical detectors, or ground-based infrared detectors equipped with laser guide star adaptive optics systems, progenitor star identifications have now been proposed for seven SNe~II-P (for a contemporary review, see \\citealt{Smartt08}, and references therein).  Although different in detail, all seven proposed progenitors have properties consistent with those of supergiant stars.  Because the field is still in its infancy --- at this point none of the proposed progenitor objects has been definitively confirmed by its absence in high-resolution images of the SN site taken after the SN has faded beyond detection --- it is imperative to carefully examine every new progenitor claim.  Here we investigate the status of the progenitor of SN~2006my, for which \\citet[][hereafter L07]{Li07} recently proposed the identification of an RSG progenitor in pre-SN images and derived a zero-age main-sequence mass for it of $M_{\\rm ZAMS} = 10^{+5}_{-3} {\\rm\\ M}_\\odot$.  Although discovered several months after explosion \\citep{Stanishev06}, the photometry and spectroscopy of SN~2006my presented by L07 clearly establish it as an SN~II-P.  To identify the progenitor star, L07 register post-SN ground-based optical images taken with the Canada-France-Hawaii Telescope in the Sloan $r^\\prime$ band under excellent seeing conditions (typical stellar full width at half-maximum [FWHM] $= 0.6^{\\prime\\prime}$) with pre-SN {\\it Hubble Space Telescope} ({\\it HST}) Wide-Field and Planetary Camera 2 (WFPC2) images (typical stellar FWHM $= 0.15^{\\prime\\prime}$ on the WF2 chip, on which the SN site is located), and identify a source in the pre-SN images within the $1\\ \\sigma$ error circle estimated from the transformation uncertainty.  Here, we reexamine this identification with the benefit of two new sets of high-resolution, post-SN images: one taken with the wide-field channel of the Near Infrared Camera 2 operated behind the laser guide star assisted adaptive optics system (\\citealt{Wizinowich06}) on the Keck II 10-m telescope (stellar FWHM $= 0.10^{\\prime\\prime}$; note that SN~2006my served as its own tip-tilt star for the observations), and the other taken with the {\\it HST} WFPC2 camera (with the SN centered on the PC chip; typical stellar FWHM $= 0.08^{\\prime\\prime}$) as part of a study of the progenitors of core-collapse supernovae (GO 10803; PI: Smartt).  This paper is organized as follows.  In \\S~\\ref{sec:2.1} we present the pre-SN and post-SN images and details of the photometric measurements performed on them; in \\S~\\ref{sec:2.2} we describe the image registration process along with our conclusion that no isolated stellar source is detected at the location of SN~2006my in either of the pre-SN images; in \\S~\\ref{sec:2.3} we estimate detection limits in the pre-SN images; and in \\S~\\ref{sec:2.4} we derive an upper mass limit on the progenitor of SN~2006my under the assumption that it was a RSG.  We summarize our findings in \\S~\\ref{sec:3}. ",
        "conclusions": "\\label{sec:3}  We analyze two pre-SN and three new post-SN high-resolution images of the site of the Type II-Plateau supernova SN~2006my in an effort to either detect the progenitor star or to constrain its properties.  Our primary result is that we do not detect an isolated stellar object at the location of SN~2006my in either of the two pre-SN images.  From our image registration, we therefore do not confirm the association found by L07 between a stellar source (Source 1) in pre-SN $I$-band images (I1) and the location of SN~2006my.  Using new high-resolution post-SN images, we derive an offset between SN~2006my and Source 1 of $\\geq 0.08^{\\prime\\prime}$ from the SN location, which represents a confidence level of non-association of more than $96\\%$ from our most liberal estimates of the image transformation and measurement uncertainties.  Through artificial star tests carried out at the precise location of SN~2006my in image I1, we derive a $3\\ \\sigma$ upper bound on the mass of the progenitor of SN~2006my of $M_{\\rm ZAMS} = 15 {\\rm\\ M}_\\odot$ from the STARS stellar evolutionary models.  Although considered unlikely, we can not rule out the possibility that {\\it part} of the light comprising Source 1, which exhibits some extension relative to other point sources in image I1, or {\\it part} of the light contributing to Source 2, a definitively extended source detected in pre-SN $V$-band images (V1) that has overlap with the SN~2006my location, may be due to the progenitor of SN~2006my.  Only additional, high-resolution observations of the site taken after SN~2006my has faded beyond detection can confirm or reject these possibilities."
    },
    "0809/0809.0696_arXiv.txt": {
        "abstract": "The radio galaxy Fornax\\,A (NGC\\,1316) is a prominent merger remnant in the outskirts of the Fornax cluster. Its giant radio lobes suggest the presence of a powerful AGN, and thus a central supermassive black hole (SMBH). Fornax\\,A now seems to be in a transition state between active black hole growth and quiescence, as indicated by the strongly declined activity of the nucleus. Studying objects in this evolutionary phase is particularly important in order to understand the link between bulge formation and black hole growth, which is manifested in the $M_\\bullet$-$\\sigma$ relation between black hole mass and bulge velocity dispersion. So far a measurement of the SMBH mass has not been possible in Fornax\\,A, as it is enshrouded in dust which makes optical measurements impossible. We present high-resolution adaptive optics assisted integral-field data of Fornax\\,A, taken with SINFONI at the Very Large Telescope in the $K$ band, where the influence of dust is negligible. The achieved spatial resolution is $0.085$~arcsec, which is about a fifth of the diameter of the expected sphere of influence of the black hole.  The stellar kinematics was measured using the region around the CO bandheads at $2.3$~$\\umu$m. Fornax\\,A does not rotate inside the inner $\\sim3$~arcsec. The velocity dispersion increases towards the centre. The weak AGN emission affects the stellar kinematics in the inner $\\sim0.06$~arcsec only. Beyond this radius, the stellar kinematics appears relaxed in the central regions. We use axisymmetric orbit models to determine the mass of the SMBH in the centre of Fornax\\,A. The three-dimensional nature of our data provides the possibility to directly test the consistency of the data with axisymmetry by modelling each of the four quadrants separately. According to our dynamical models, consistent SMBH masses $M_\\bullet$ and dynamical \\emph{Ks} band mass-to-light ratios $\\Upsilon$ are obtained for all quadrants, with $\\langle M_{\\bullet}\\rangle=1.3\\times10^{8}$~M$_{\\odot}$ ($\\mathrm{rms}(M_{\\bullet})=0.4\\times10^{8}$~M$_{\\odot}$) and $\\langle\\Upsilon\\rangle=0.68$ ($\\mathrm{rms}(\\Upsilon)=0.03$), confirming the assumption of axisymmetry. For the folded and averaged data we find $M_{\\bullet}=1.5_{-0.8}^{+0.75}\\times10^8$~M$_{\\odot}$ and $\\Upsilon=0.65^{+0.075}_{-0.05}$ ($3\\sigma$ errors). Thus the black-hole mass of Fornax\\,A is consistent within the error with the \\citet{Tremaine-02} $M_{\\bullet}$-$\\sigma$ relation, but is a factor $\\sim4$ smaller than expected from its bulge mass and the \\citet{MarconiHunt-03} relation. ",
        "introduction": "Studies of the dynamics of stars and gas in the centres of nearby galaxies with a massive bulge component have established the presence of supermassive black holes (SMBHs) in the $10^6$--$10^9$~M$_\\odot$ range. The mass of the central SMBH is tightly correlated with the bulge mass or luminosity (\\textrm{e.g.} \\citealt{MarconiHunt-03}) and with the bulge velocity dispersion $\\sigma$ \\citep{Gebhardt-00a,Ferrarese-00}. These correlations suggest that bulge evolution and black hole growth are closely linked.  Indeed there is increasing theoretical evidence that galaxy merging (or other processes that lead to gas inflow, like secular evolution), nuclear activity and feedback are all somehow linked to bulge and SMBH evolution (\\textrm{e.g.} \\citealt{diMatteo-05,Johansson-08,Younger-08,Hopkins-08a}).  To observationally constrain present theories of bulge and black hole evolution, detailed studies of AGN and merger remnants in different evolutionary stages are essential. NGC\\,5128 (Cen\\,A) is presently the only galaxy with a powerful AGN that underwent a recent major merger which has a measured black-hole mass \\citep{Silge-05,Marconi-06,Neumayer-07}.  A galaxy very similar to Cen\\,A is NGC\\,1316 (Fornax\\,A).  Fornax\\,A is a giant elliptical galaxy located in the outskirts of the Fornax galaxy cluster. It is one of the brightest radio galaxies in the sky with giant double radio lobes and $\\mathcal{S}$-shaped nuclear radio jets \\citep{Geldzahler-84}. The peculiar morphology with numerous tidal tails, shells and loops \\citep{Schweizer-80,Schweizer-81} suggests that a major merger happened about $3$~Gyr ago \\citep{Goudfrooij-01}, followed by some minor mergers \\citep{Mackie-98}. Its nucleus however is surprisingly faint in X-rays, which can only be explained if the nucleus became dormant during the last $0.1$~Gyr \\citep{Iyomoto-98}. This makes it an ideal target for a dynamical black-hole mass measurement, as its stellar absorption lines are probably not or just marginally diluted by non-stellar emission from the AGN.  Fornax\\,A is classified as an intermediate form between core and power law galaxy in \\citet{Lauer-07}. The surface brightness approximately follows an $r^{1/4}$-law \\citep{Caon-94}. Large amounts of dust are present also in the inner few arcseconds, which affects optical spectroscopy. We therefore measure the two-dimensional stellar kinematics of Fornax\\,A with the near-infrared integral-field spectrograph SINFONI at the Very Large Telescope (VLT) in the $K$ band, where the dust obscuration is only about $7$~per~cent of the obscuration in the optical. We use this data to analyse the stellar kinematics in a way similar to \\citet{Nowak-07}.   Throughout this paper we adopt a distance to Fornax\\,A of $18.6$~Mpc based on \\emph{HST} measurements of Cepheid variables in NGC\\,1365 \\citep{Madore-99}. At this distance, $1$~arcsec corresponds to $90$~pc. With a velocity dispersion of $226$~\\kms\\ (mean $\\sigma$ measured in an $8$~arcsec aperture in \\S 3.4) the estimated sphere of influence has a diameter of $0.46$~arcsec and we would expect a SMBH mass of $2.2\\times10^8$~M$_\\odot$ \\citep{Tremaine-02}. This is large enough to be resolved easily from the ground with adaptive optics.  This paper is organized as follows: In \\S 2 we present the data and the data reduction. The stellar kinematics of Fornax\\,A is described in \\S 3, the measurement of the near-infrared line strength indices in \\S 4. The stellar dynamical modelling procedure and the results for the SMBH mass are presented in \\S 5, and \\S 6 summarises and discusses the results. ",
        "conclusions": "We have obtained near-IR integral field data with three different spatial resolutions $\\geq0.085$~arcsec for the merger remnant Fornax\\,A. Stellar rotation is only detected in the outer parts ($r\\ga1.5$~arcsec). The stellar velocity dispersion slightly decreases from large to smaller radii. In the highest resolution data it shows two peaks in the centre, separated by a narrow low-$\\sigma$ region. The average dispersions are $\\sigma=226\\pm9$~\\kms\\ in an $8$~arcsec diameter aperture, $\\sigma=221\\pm4$~\\kms\\ in a $3.0$~arcsec diameter aperture and $\\sigma=218\\pm8$~\\kms\\ in an $0.8$~arcsec diameter aperture. A low-luminosity AGN is likely to be present and distort the stellar kinematics in the central $\\la0.06$~arcsec, where weak \\caviii\\ emission in the region of the second CO bandhead probably induces a velocity peak. The $\\sigma$-drop can be at most partially attributed to this emission line. Either the AGN continuum emission or a cold stellar subsystem or both could be the cause.  Near-IR line indices were measured using the high-S/N, intermediate resolution data to trace stellar populations. All indices show the same trend, a negative gradient and in the centre an unresolved depression. The negative gradient of the CO index can be interpreted as a metallicity gradient. The central drop ($\\sim1\\mathrm{\\AA}$ in CO) could be, like the $\\sigma$ drop, caused by the low-luminosity AGN, a subsystem with a different stellar population, or a combination of the two.  Using an axisymmetric Schwarzschild code we find a black hole with the mass $M_{\\bullet}=1.5^{+0.75}_{-0.8}\\times10^8$~M$_{\\odot}$ ($3\\sigma$ C.L. for two degrees of freedom), when the 25mas, the 100mas and the longslit data are included. The same mass with slightly larger error bars is found when the innermost bins, where the kinematics is distorted by the AGN, are excluded from the dynamical modelling. The assumption of axisymmetry is justified as consistent results were obtained when modelling single quadrants. The mean black-hole mass obtained from the four single quadrants is $\\langle M_{\\bullet}\\rangle=1.3\\times10^{8}$~M$_{\\odot}$ with $\\mathrm{rms}(M_{\\bullet})=0.4\\times10^{8}$~M$_\\odot$ and the mean \\emph{Ks} band mass-to-light ratio is $\\langle\\Upsilon\\rangle=0.68$ with $\\mathrm{rms}(\\Upsilon)=0.03$, in agreement with the $1\\sigma$ limit derived from the analysis of the smoothed $\\chi^2$-profile for the averaged data. The PSF of the 100mas data could not be measured very accurately, but this does not seem to have noticeable effects. Modelling the 100mas data with the measured or a reconstructed, noisier PSF (both have about the same FWHM) does not yield different results.   We find a dynamical \\emph{Ks} band mass-to-light ratio of $\\Upsilon=0.65_{-0.05}^{+0.075}$. When the inner two radial bins are not included it is somewhat larger ($\\Upsilon=0.725_{-0.125}^{+0.025}$). This is in agreement with stellar population models assuming a Salpeter IMF, but only marginally consistent with a Kroupa IMF. A dark halo plays a significant role only outside $\\sim31$~arcsec, close to the effective radius ($R_\\mathrm{e}=36$~arcsec, \\citealt{Bedregal-06}). We therefore included neither a dark halo nor longslit data at radii larger than $31$~arcsec.  The black-hole mass expected from the $M_{\\bullet}$-$\\sigma$ relation \\citep{Tremaine-02} is $(2.2\\pm0.4)\\times10^8$~M$_{\\odot}$. Note that estimating $\\sigma_{\\mathrm{e}}$, the luminosity-weighted velocity dispersion within $R_{\\mathrm{e}}$, from Fig. \\ref{fig:17} gives the same value as $\\sigma$ measured in the $8$~arcsec aperture quoted above. This mass is consistent with our measurements within the errors. Fornax\\,A is only one of two galaxies that underwent a major merger a few Gyr ago where a black-hole mass has been measured (the other one is Cen\\,A, \\textrm{e.g.}  \\citealt{Neumayer-07}). Both seem to be consistent with the $M_{\\bullet}$-$\\sigma$ relation, which holds important implications for the growth of the black hole and its surrounding bulge. Despite still showing obvious characteristics of the merging process -- like shells and ripples in the outer envelope and disordered dust features throughout the entire galaxy -- the black hole of Fornax\\,A has about the mass expected for a normal elliptical. This suggests that bulges and black holes grow approximately synchronously in the course of a merger, or that the growth is regulated during the active phase.  The black-hole mass of Fornax\\,A is, however, not in agreement with the relations between black-hole mass and $K$ band luminosity $L_K$ or bulge mass $M_{\\mathrm{bulge}}$ of \\citet{MarconiHunt-03}. The total 2MASS \\emph{Ks} band magnitude $m_{Ks}=5.587$ corresponds to $\\log(M_{\\mathrm{bulge}}/M_{\\odot})\\approx11.4$ for our $\\Upsilon_{Ks}=0.65$. For this high bulge mass a black-hole mass of $\\sim6\\times10^8$~M$_\\odot$ would be expected, a factor $\\sim4$ larger than what we measured. Interestingly the situation is similar for Cen\\,A. With the $K$ band magnitude given in \\citet{MarconiHunt-03} and $\\Upsilon_K$=0.72 \\citep{Silge-05} the bulge mass is $\\log(M_{\\mathrm{bulge}}/M_{\\odot})\\approx11.0$ and a black hole with $M_\\bullet\\approx2.2\\times10^8$~M$_\\odot$ would thus be expected, a factor $\\sim5$ larger than the mass of $4.5\\times10^7$~M$_\\odot$ obtained by \\citet{Neumayer-07}. In this context it would be interesting to reconsider the relationship between $M_\\bullet/M_{\\mathrm{bulge}}$ and galaxy age found by \\citet{Merrifield-00}. Fornax\\,A and Cen\\,A both have a very small $M_{\\bullet}/M_{\\mathrm{bulge}}$ and they would be among the youngest galaxies in their sample. Although the number of merger galaxies with measured black-hole masses is fairly small, a correlation between the time of the last major merger and $M_\\bullet$ might start to emerge.  In order to draw more stringent conclusions on the connections between merging, bulge growth, black hole growth and nuclear activity more dynamical black-hole mass measurements in merger remnants and in luminous AGN are necessary. We were able to show that reliable black hole masses can be derived via stellar dynamical modelling even if the nuclear source makes dynamical measurements in the centre impossible, as long as the spatial resolution is high enough that the AGN signature is well within the sphere of influence."
    },
    "0809/0809.0375_arXiv.txt": {
        "abstract": "We consider the prospects for dark matter/energy unification in k-essence type cosmologies. General mappings are established between the k-essence scalar field, the hydrodynamic and braneworld descriptions. We develop an extension of the general relativistic dust model that incorporates the effects of both pressure and the associated acoustic horizon. Applying this to a tachyon model, we show that this inhomogeneous ``variable Chaplygin gas\" does evolve into a mixed system containing cold dark matter like gravitational condensate in significant quantities. Our methods can be applied to any dark energy model, as well as to mixtures of dark energy and traditional dark matter. ",
        "introduction": "The discovery of the accelerated Hubble expansion in the SNIa data \\cite{perl1}, combined with observations of the cosmic microwave background (CMB) \\cite{hal2,hin3}, has forced a profound shift in our cosmological paradigm. If one makes the conservative assumptions of the validity of Einstein's general relativity and the cosmological principle, one concludes that the universe is presently dominated by a component that violates the strong energy condition, dubbed {\\em dark energy} (DE). Moreover, primordial nucleosynthesis constrains the fraction of closure density in baryons, $\\Omega_{B}$, to a few percent, while galactic rotation curves and cluster dynamics imply the existence of a nonbaryonic {\\it dark matter} (DM) component with $\\Omega_{DM} \\gg \\Omega_{B}$ (for a review see \\cite{peeb4}). Currently, the best fit values are $\\Omega_{B} = 0.04, \\Omega_{DM} = 0.22$ and $\\Omega_{DE}$ = 0.74 \\cite{hin3}. Thus it may be said that we have a firm theoretical understanding of only 4\\% of our universe.  Pragmatically, the data can be accommodated by combining baryons with con\\-ven\\-tion\\-al cold dark matter (CDM) candidates and a simple cosmological constant $\\Lambda$ providing the DE. This $\\Lambda$CDM model, however, begs the question of why $\\Lambda$ is non-zero, but such that DM and DE are comparable today. The coincidence problem of the $\\Lambda$CDM model is somewhat ameliorated in {\\it quintessence} models which replace $\\Lambda$ by an evolving scalar field. However, like its predecessor, a quintessence-CDM model assumes that DM and DE are distinct entities. For a recent review of the most popular DM and DE models, see \\cite{sah5}.  Another interpretation of this data is that DM/DE are different manifestations of a common structure. Speculations of this sort were initially made by Hu \\cite{hu6}. The first definite model of this type was proposed a few years ago \\cite{kam9,bil7,bil8}, based upon the Chaplygin gas, a perfect fluid obeying the equation of state \\begin{equation} p =  - \\frac{A}{\\rho} \\; , \\label{eq001} \\end{equation} which has been extensively studied for its mathematical properties \\cite{jack10}. The general class of models, in which a unification of DM and DE is achieved through a single entity, is often referred to as {\\em quartessence} \\cite{mak11,mak12}. Among other scenarios of unification that have recently been suggested, interesting attempts are based on the so-called {\\em k-essence} \\cite{chi13,sch14}, a scalar field with noncanonical kinetic terms which was first introduced as a model for inflation \\cite{arm15}.  The cosmological potential of equation (\\ref{eq001}) was first noted by Kamenshchik {\\it et al} \\cite{kam9}, who observed that integrating the energy conservation equation in a homogeneous model leads to \\begin{equation} \\rho(a) = \\sqrt{A + \\frac{B}{a^{6}} }  \\; , \\label{eq002} \\end{equation} where $a$ is the scale factor normalized to unity today and $B$ an integration constant. Thus, the Chaplygin gas interpolates between matter, $\\rho \\sim \\sqrt{B} a^{-3}$, $p \\sim 0$, at high redshift and a cosmological constant like $\\rho \\sim \\sqrt{A} \\sim - p$ as $a$ tends to infinity. The essence of the idea in \\cite{bil7,bil8} is simply that in an {\\it inhomogeneous} universe, highly overdense regions (galaxies, clusters) have $|w| = |p/\\rho| \\ll 1$ providing DM, whereas in underdense regions (voids) evolution drives $\\rho$ to its limiting value $\\sqrt{A}$ giving DE.  Of particular interest is that the Chaplygin gas has an equivalent scalar field formulation \\cite{bil7,bil8,jack10}. Considering the Lagrangian \\begin{equation} {\\cal L}  = - \\sqrt{A} \\sqrt{1 - X}\\, , \\label{eq003} \\end{equation} where \\begin{equation} X \\equiv g^{\\mu \\nu} \\varphi_{, \\mu} \\varphi_{, \\nu} \\: , \\label{eq1203} \\end{equation} equation (\\ref{eq001}) is obtained by evaluating the stress-energy tensor $T_{\\mu \\nu}$, and  introducing $u_{\\mu} = \\varphi_{,\\mu} / \\sqrt{X}$ for the four-velocity and $\\rho = \\sqrt{A} / \\sqrt{1 - X}$ for the energy density. One recognizes $\\cal{L}$ as a  Lagrangian of the Born-Infeld type, familiar in the $D$-brane constructions of string/$M$ theory \\cite{pol16}. Geometrically, $\\cal{L}$ describes space-time as the world-volume of a 3+1 brane in a 4+1 bulk via the embedding coordinate $X^{5}$ \\cite{sun17}.  To be able to claim that a field theoretical model actually achieves unification, one must be assured that initial perturbations can evolve into a deeply nonlinear regime to form a gravitational condensate of superparticles that can play the role of CDM. In \\cite{bil7, bil8} this was inferred on the basis of the Zel'dovich approximation \\cite{zel18}. In fact, for this issue, the usual Zel'dovich approximation has the shortcoming that the effects of finite sound speed are neglected.  All models that unify DM and DE face the problem of nonvanishing sound speed and the well-known Jeans instability. A fluid with a nonzero sound speed has a characteristic scale below which the pressure effectively opposes gravity. Hence the perturbations of the scale smaller than the sonic horizon will be prevented from growing. Soon after the appearance of \\cite{kam9} and \\cite{bil7}, it was pointed out that the perturbative Chaplygin gas (for early work see \\cite{fab19}, and more recently \\cite{gor20}) is incompatible with the observed mass power spectrum \\cite{sand21} and microwave background \\cite{cart22}. Essentially, these results follow from the adiabatic speed of sound \\begin{equation} c_s^2 = \\left.\\frac{\\partial p}{\\partial \\rho}\\right|_{s} = \\frac{A}{\\rho^{2}} \\label{eq004} \\end{equation} which leads to a comoving acoustic horizon \\begin{equation} d_s = \\int  dt \\, \\frac{c_s}{a} \\; \\; . \\label{eq005} \\end{equation} The perturbations whose comoving size $R$ is larger than $d_s$ grow as $\\delta = (\\rho - \\bar{\\rho})/\\bar{\\rho} \\sim a$. Once the perturbations enter the acoustic horizon, i.e., as soon as $R<d_s$, they undergo damped oscillations. In the case of the Chaplygin gas we have $d_s \\sim a^{7/2}/H_0$, where $H_{0}$ is the present day value of the Hubble parameter, reaching Mpc scales already at redshifts of order 10. However, to reiterate a point made in \\cite{bil7}, small perturbations {\\it alone} are not the issue, since large density contrasts are required on galactic and cluster scales. As soon as $\\delta \\simeq 1$ the linear perturbation theory cannot be trusted. An essentially nonperturbative approach is needed in order to investigate whether a significant fraction of initial density perturbations collapses in gravitationally bound structure - the condensate. If that happens the system evolves into a two-phase structure - a mixture of CDM in the form of condensate and DE in the form of uncondensed gas.  The case, where the Chaplygin gas is mixed with CDM, has been considered in a number of papers \\cite{dev23,dev24,avel25,alca26,bean27,amen28,mult29,bert30}. Here, the Chaplygin gas simply plays the role of DE. In keeping with the quartessence philosophy, it would be preferred if  CDM could be replaced by droplets of Chaplygin gas condensate, as in \\cite{bil31}. Homogeneous world models, containing a mixture of CDM and Chaplygin gas, have been successfully confronted with lensing statistics \\cite{dev23,dev24} as well as with supernova and other tests \\cite{avel25,alca26}.  Another model, the so called ``generalized Chaplygin gas\" \\cite{ben32}, has gained a wide popularity. The generalized Chaplygin gas is defined as \\cite{bil7,kam9,ben32} $p = - A/\\rho^{\\alpha}$ with $0\\leq \\alpha\\leq 1$ for stability and causality. As in the Chaplygin gas case, this equation of state has an equivalent field theory representation, the ``generalized Born-Infeld theory''\\cite{ben32,bert33}. However, the associated Lagrangian has no equivalent brane interpretation. The additional parameter does afford greater flexibility: e.g. for small $\\alpha$ the sound horizon is $d_{s} \\sim \\sqrt{\\alpha} a^{2}/H_{0}$, and thus by fine tuning $\\alpha < 10^{-5}$, the data can be perturbatively accommodated \\cite{sand21}. Bean and Dor\\'{e} \\cite{bean27} and similarly Amendola {\\it et Al} \\cite{amen28} have examined a mixture of CDM and the generalized Chaplygin gas against supernova, large-scale structure, and CMB constraints. They have demonstrated that a thorough likelihood analysis favors the limit $\\alpha \\rightarrow 0$, i.e. the equivalent to the $\\Lambda$CDM model.  Both papers conclude that the standard Chaplygin gas is ruled out as a candidate for DE. However, analysis \\cite{bert30,bert33} of the supernova data seems to indicate that the generalized Chaplygin gas with $\\alpha \\geq 1$ is favored over the $\\alpha \\rightarrow 0$ model and similar conclusions were drawn in \\cite{gor20}. But one should bear in mind that the generalized Chaplygin gas with $\\alpha > 1$ has a superluminal sound speed that violates causality  \\cite{ell62}. For a different view on this issue see \\cite{bru,kan,bab63}  (see discussion in section \\ref{kessentials}).  The structure formation question, in respect of the Chaplygin gas, was decided in \\cite{bil34}. In fact, in the Newtonian approximation, we derived an extension of the spherical model \\cite{gazt35} that incorporates nonlinearities in the density contrast $\\delta$, as well as the effects of the adiabatic speed of sound. Both are crucial, since for an overdensity we have $c_s < \\bar{c}_s$, where $\\bar{c}_s$ is the speed of sound of the background. Although small initial overdensities follow the expected perturbative evolution, for an initial $\\delta_R (a_{in})$  exceeding a scale $R$-dependent critical $\\delta_{c}$, it was shown that $\\delta_R (a)$ tends to infinity at finite redshift, signaling the formation of a bound structure or {\\em condensate}. Unfortunately, it was further found that, when the required $\\delta_c$ is folded with the spectrum of the initial density perturbations to obtain the collapse fraction, less than 1\\% of the Chaplygin gas ends up as condensate. Thus the simple Chaplygin gas is not viable due to frustrated structure formation. In effect, the model is a victim of the radiation dominated phase, which turns the Harrison-Zel'dovich spectrum $\\delta_{k} \\sim k^{1/2}$ to $\\delta_{k} \\sim k^{- 3/2}$ at $R_{\\rm CEQ} \\simeq  26$ Mpc. In a pure Chaplygin-gas universe without radiation there would inevitably be sufficient small scale power to drive condensation.  One way to deal with the structure formation problem, is to assume entropy perturbations \\cite{hu6,reis36} such that the effective speed of sound vanishes\\footnote{ Note this ``silent quartessence'' is not different from \\cite{bil7}, where we tacitly neglected the effects of nonvanishing $c_s$.}. In that picture we have $\\delta p = c_s^2 \\delta \\rho - \\delta A/\\rho = 0$ even if $c_s \\neq 0$. But as we detail below, in a single field model it is precisely the adiabatic speed of sound that governs the evolution. Hence, entropy perturbations require the introduction of a second field on which $A$ depends. Aside from negating the simplicity of the one-field model, some attempts at realizing the nonadiabatic scenario \\cite{bil37,bil38,bil39} have convinced us that even if $\\delta p = 0$ is arranged as an initial condition, it is all but impossible to maintain this condition in a realistic model for evolution.  The failure of the simple Chaplygin gas does not exhaust all the possibilities for quart\\-essence. The Born-Infeld Lagrangian (\\ref{eq003}) is a special case of the string-theory inspired tachyon Lagrangian \\cite{sen40,gibb41} in which the constant $\\sqrt{A}$ is replaced by a potential $V(\\varphi)$ \\begin{equation} {\\cal L}  = - V(\\varphi) \\; \\sqrt{1 - g^{\\mu \\nu} \\, \\varphi_{,\\mu} \\, \\varphi_{,\\nu} } \\; . \\label{eq01} \\end{equation} In turn, tachyon models are a particular case of k-essence \\cite{arm15}. The possibility of obtaining both DM and DE from the tachyon with inverse square potential has been speculated in \\cite{pad42}. More recently, it was noted \\cite{guo43} that, in a Friedmann-Robertson-Walker (FRW) model, the tachyon model is described by the equation of state (\\ref{eq001}) in which the constant $A$ is replaced by a function of the cosmological scale factor $a$, so the model was dubbed  ``variable Chaplygin gas''. Related models have been examined in \\cite{bert44,bert45}, however, those either produce a larger $d_{s}$  than the simple Chaplygin gas \\cite{bert44}, or else need fine-tuning \\cite{bert45}.\\footnote{The tachyon model \\cite{bert44} gives $d_{s} \\sim a^{2}/H_{0}$. The two-potential model \\cite{bert45} yields $d_{s} \\sim \\sqrt{1-h} a^{2}/H_{0}$, so it requires $1 - h < 10^{-5}$ like the generalized Chaplygin gas. Expanding in $1 - h$, the second potential reveals itself to be dominantly a cosmological constant.}  In this paper we develop a version of the spherical model for studying the evolution of density perturbations even into the fully nonlinear regime. Although similar in spirit to \\cite{bil34}, the formalism here is completely relativistic, rather than Newtonian, and applicable to any k-essence model instead of being restricted to the simple Chaplygin gas. The one key element we carry over from \\cite{bil34} is an approximate method for treating the effects of pressure gradients - which is to say the adiabatic speed of sound - on the evolution. Our method is flexible enough and can be extended to deal with a mixture of DE and DM. A spherical model has been applied to DE/DM mixtures in \\cite{abra46}, however there the effects of pressure gradients were omitted.  We apply our method to the preliminary analysis of a unifying model based on the tachyon type Lagrangian (\\ref{eq01}) with a potential of the form \\begin{equation} V(\\varphi)=V_n \\varphi^{2n} \\; \\; , \\label{eq08} \\end{equation} where $n$ is a positive integer. In the regime where structure function takes place, we show that this model effectively behaves as the variable Chaplygin gas with $A(a) \\sim a^{6n}$ with $n = 1 (2)$ for a quadratic (quartic) potential. As a result, the much smaller acoustic horizon $d_{s} \\sim a^{(7/2+3n)}/H_{0}$ enhances condensate formation by two orders of magnitude over the simple Chaplygin gas ($n = 0$). Hence this type of model may salvage the quartessence scenario.     The remainder of this paper is organized as follows. In section \\ref{kessentials} we reformulate k-essence type models in a way that allows us to deal with large density inhomogeneities. In section 3 we develop the spherical model approximation that closes the system of equations. These two sections are completely general and stand alone. Numerical results, in section 4, are presented for positive power-law potentials and contrasted with the simple Chaplygin gas. Our conclusions and outlook are given in section \\ref{conclude}. Finally, in \\ref{adiabatic}, we derive the adiabatic speed of sound for a general k-essence fluid, and in \\ref{dbrane}, we give a brief description of the tachyon model from the braneworld perspective. ",
        "conclusions": "\\label{conclude} The first key test for any proposed quartessence model should be: Does it actually yield nonlinear dark matter structure, as well as linear dark energy in the {\\it inhomogeneous} {\\it almost} FRW universe that we see. In this paper we have analyzed the nonlinear evolution of the  tachyon-like k-essence with a very simple potential $V=V_0\\varphi^{2n}$. We have demonstrated that a significant fraction of the fluid, in particular for $n=2$, collapses into condensate objects that play the role of cold dark matter. No dimensionless fine tunings were required beyond the inevitable $\\sqrt{G} H_{0} \\ll 1$.   Moreover, these results were obtained in a relativistic framework for nonlinear evolution that is as simple as the spherical dust model but includes the key effects of the acoustic horizon. Although it could be subject to possible improvements (e.g. a variable Gaussian width \\cite{bil34}, and it is lacking the sophistication of the exact spherical model, \\cite{sus56}, it does allow us to make the sort of quantitative assessments that have been missing \\cite{avel57}. It is directly formulated in a convenient coordinate gauge, does not involve any hidden assumptions, and it easily deals with multi-component systems.  The  tachyon k-essence unification remains to be tested against large-scale structure and CMB observations. However, we maintain, contrary to the opinion advocated in \\cite{sand21}, that the sound speed problem may be alleviated in unified models no more unnatural than the $\\Lambda$CDM model. Indeed, an encouraging feature of the positive power-law potential is that it provides for acceleration as a periodic transient phenomenon \\cite{fro58} which obviates the de Sitter horizon problem \\cite{bil37} , and, we speculate, could even be linked to inflation.  \\appendix"
    },
    "0809/0809.2370_arXiv.txt": {
        "abstract": "We report the discovery of the first known symbiotic star in IC10, a starburst galaxy belonging to the Local Group, at a distance of $\\sim$750~kpc. The symbiotic star was identified during a survey of emission-line objects. It shines at $V$=24.62$\\pm$0.04, $V$$-$$R_{\\rm C}$=2.77$\\pm$0.05 and $R_{\\rm C}$$-$$I_{\\rm C}$=2.39$\\pm$0.02 and suffers from $E_{B-V}$=0.85$\\pm$0.05 reddening. The spectrum of the cool component well matches that of solar neighborhood M8III giants. The observed emission lines belong to Balmer series, [SII], [NII] and [OIII]. They suggest a low electronic density, negligible optical depth effects and 35,000$<$$T_{\\rm eff}$$<$90,000~K for the ionizing source.  The spectrum of the new symbiotic star in IC10 is an almost perfect copy of that of Hen~2-147, a well known Galactic symbiotic star and Mira. ",
        "introduction": " ",
        "conclusions": "\\subsection[]{Astrometric position and VR$_{\\rm C}$I$_{\\rm C}$ magnitudes}  Massey et al. (2007) obtained, with the Kitt Peak National Observatory and Cerro Tololo Inter-American Observatory 4-m telescopes and Mosaic cameras, UBVR$_{\\rm C}$I$_{\\rm C}$ photometry of the resolved stellar population of several dwarf galaxies with active star formation, including IC10. Massey et al. (2007) did not provide photometry for IC10 StSy-1. We have used their source observations (accessible via ftp\\footnote{ftp://ftp.lowell.edu/pub/massey/lgsurvey}), and derived for IC10 StSy-1 $V$=24.62$\\pm$0.04, $V$$-$$R_{\\rm C}$=2.77$\\pm$0.05 and $R_{\\rm C}$$-$$I_{\\rm C}$=2.39$\\pm$0.02 via aperture photometry. The symbiotic star is below detection threshold both in $U$- and $B$-band frames.  \\begin{figure} \\begin{center} \\vbox{\\psfig{file=DRGoncalves_Fig2.ps,width=8.5truecm}} \\caption{$I_{\\rm C}$-band finding chart covering a field of view of 37$\\times$37 arcsec from Massey et al. (2007) observations. North is up and East is to the left. IC10 StSy-1 ($00^h20^m33\\fs59$ $+59\\degr18\\arcmin45\\farcs9$) is marked with a circle. The photometry of the stars marked with A to E is given in  Table~1.} \\end{center} \\end{figure}  The position we derived for IC10~SySt-1 on Massey et al. (2007) astrometrically calibrated frames is RA=00:20:33.59 and DEC=+59:18:45.9 (J2000.0), with an error smaller than 1 arcsec. This position is marked on the $I_{\\rm C}$-band finding chart presented in Figure~2.  It covers a field of view of 37$\\times$37 arcsec and shows stars down to $I_{\\rm C}$=20~mag. We also marked on Figure~2 a few field stars suitable to serve as local photometric standards, whose photometry is given in Table~1.  \\begin{table} \\caption{Magnitudes of the finding chart stars of Figure~2.} \\centering \\begin{tabular}{lllll} \\hline \\multicolumn{5}{c}{}\\\\ &  V               & B-V             & V-R$_C$         & R-I$_C$   \\\\ \\multicolumn{5}{c}{}\\\\ A & 17.509$\\pm$0.005 & 0.992$\\pm$0.005 & 0.587$\\pm$0.005 & ...\\\\ B & 16.983$\\pm$0.005 & 1.200$\\pm$0.005 & 0.688$\\pm$0.005 & ...\\\\ C & 22.185$\\pm$0.012 & 1.920$\\pm$0.020 & 1.243$\\pm$0.014 & 1.425$\\pm$0.008 \\\\ D & 18.909$\\pm$0.005 & 1.274$\\pm$0.005 & 0.794$\\pm$0.005 & 0.822$\\pm$0.005 \\\\ E & 18.575$\\pm$0.005 & 1.110$\\pm$0.005 & 0.652$\\pm$0.005 & 0.606$\\pm$0.005 \\\\ \\multicolumn{5}{c}{}\\\\ \\hline \\end{tabular} \\end{table}  \\subsection[]{Classification}  The spectrum of IC10~SySt-1 presented in Figure~1 is that of a symbiotic star: strong molecular TiO absorption bands from a cool giant with superimposed emission lines tracing the presence of a hot companion that ionizes the circumstellar nebula fed by the mass loss from the giant. The presence of [OIII] requires a temperature of the photo-ionization source in excess of 35,000~K, which satisfy the classification criteria for symbiotic stars adopted by Belczy\\'nski et al. (2000). Our spectrum is too underexposed at the shortest wavelengths to distinguish between absence or presence of HeII~4686 emission line, which would be required for the more stringent classification criteria among symbiotic stars adopted by Allen (1984) and that would trace ionization temperatures in excess of 55,000~K.  The partnership of IC10~SySt-1 with symbiotic stars is reinforced by the comparison, carried out in Figure~1, with He~2-147. This is a bona fide symbiotic star containing a long period Mira and a toroidal expanding circumstellar nebula resolved in both ground-based and HST observations (Munari and Patat 1993, Corradi et al. 1999, Santander-Garc\\'\\i a et al. 2007). The only significant difference with Hen~2-147 in Figure~1 is a lower intensity in IC10~SySt-1 of [OIII] lines with respect to [NII], probably tracing a lower temperature of the photo-ionizing source. A list and integrated fluxes of the emission lines identified in the spectrum of IC10~SySt-1 is given in Table~2.  \\begin{table} \\caption{Emission lines observed in IC10 StSy-1. $F$: observed fluxes in units of 10$^{-17}$ erg cm$^{-2}$ sec$^{-1}$ (errors $\\pm$10\\%). $F_\\circ$: reddening corrected fluxes scaled to H$\\beta$} \\centering \\begin{tabular}{llrr} \\hline \\multicolumn{4}{c}{}\\\\ &&  $F$  & $F_\\circ$   \\\\ \\multicolumn{4}{c}{}\\\\ 4861    &\\hb    &  2.2  &1.00 \\\\ 4959    &[OIII] &  2.2  &0.91 \\\\ 5007    &[OIII] &  7.1  &3.08 \\\\ 6548    &[NII]  &  7.0  &1.13 \\\\ 6563    &\\ha    & 18.7  &2.85 \\\\ 6584    &[NII]  & 21.0  &3.25 \\\\ 6717    &]SII]  &  2.1  &0.28 \\\\ 6731    &[SII]  &  1.9  &0.24 \\\\ \\multicolumn{4}{c}{}\\\\ \\hline \\end{tabular} \\end{table}  \\subsection[]{Spectral type and reddening}  \\begin{figure*} \\begin{center} \\vbox{\\psfig{file=DRGoncalves_Fig3.ps,width=18.0truecm}} \\caption{Comparison of an M8III spectrum from Fluks et al. (1994, thicker line), reddened by $E_{B-V}$=0.85 with our spectrum of IC10 SySt-1 from Figure~1 (thinner line).} \\end{center} \\end{figure*}  To derive the spectral type of the cool giant, we have compared the red portion of the IC10~SySt-1 spectrum with the spectral atlas of Fluks et al. (1994), which is based on observation of solar neighborhood objects. We obtained an excellent match with a M8III type, and the much poorer respective fits exclude a classification as either M7 or M9 spectral types. The comparison with the unreddened spectra of Fluks et al. (1994) also firmly constrains the reddening, found to be $E_{B-V}$=0.85$\\pm$0.05. The fit to the observed spectrum of IC10~SySt-1 with an M8III from Fluks et al. (1994) and reddened by $E_{B-V}$=0.85 is presented in Figure~3. Integrating the BVR$_{\\rm C}$I$_{\\rm C}$ pass-bands would provide B=26.79, V=24.43, R$_{\\rm C}$=21.85 and I$_{\\rm C}$=19.33 for the reddened M8III fitting spectrum, and B=23.47, V=21.95, R$_{\\rm C}$=20.12 and I$_{\\rm C}$=17.94 for the unreddened one.   The reddening we found for IC10~SySt-1 is the same as reported by Mateo (1998) for the IC10 as a whole, suggesting a negligible circumstellar contribution. The $E_{B-V}$=0.85 value for the reddening affecting IC10~SySt-1 is also supported by the flux ratio of Balmer H$\\alpha$ and H$\\beta$ emission lines. Their recombination line ratio under Case B conditions and negligible self-absorption is $\\sim$2.9 (Osterbrock and Ferland 2006, their Table 4.4). In IC10~SySt-1 the observed ratio is 8.5 which reduces to 2.9 once corrected for $E_{B-V}$=0.85.  \\subsection[]{Electron density, oxygen abundance and temperature of the ionizing source}  A preliminary estimate of the nebular conditions can be derived from the data at hand (see Table~2), using the {\\sc nebular} IRAF package (Shaw \\& Dufour 1994).  The electron density in the circumstellar nebula can be estimated as 400$\\pm$200 cm$^{-3}$ using the \\sii$\\lambda\\lambda$ 6716, 6731 \\AA\\ doublet. An upper limit to the flux of \\oiii\\ 4363~\\AA\\ emission line (6\\% of \\hb ) gives a lower limit to the electron temperature of $\\sim$17,000~K from the I($\\lambda$4959+$\\lambda$5007)/I($\\lambda$4363) ratio. Using this limit for \\te\\oiii\\ and the O$^{++}$ abundance, the lower limit to the total oxygen abundance (Kingsburgh \\&  Barlow 1994) is 12 + $\\log$(O/H) $>$ 7.4, the solar abundance being 8.66 (Asplund Grevesse \\& Sauval 2005). An upper limit for the integrated flux of \\heii\\ 4686~\\AA\\ emission line amounting to 5\\% of the integrated flux of \\hb\\ corresponds to an upper limit $T_{\\rm eff}$$<$90,000~K for the temperature of the ionizing source, the lower limit $>$35,000~K being set by the presence of \\oiii\\ (Kaler \\& Jacoby 1989).  Contrary to many symbiotic stars, the nebular lines of IC10 StSy-1 do not indicate the presence of high density ionized gas. The observed lines could instead originate in an extended,  low density nebula. To the low density conditions characterizing IC10 SySt-1 spectrum could be contributing the emission from external and extended nebular regions similar to those resolved around many symbiotic Miras in our Galaxy (Corradi et al. 1999), like the one around Hen~2-147, or the more spectacular ones of He2-104 (Corradi et al. 2001) and R Aqr (Gon\\c calves et al. 2003)."
    },
    "0809/0809.2693_arXiv.txt": {
        "abstract": " ",
        "introduction": "The presence of thermal emission at radio frequencies may be a useful tool for identifying interactions between supernova remnants (SNRs) and molecular clouds, and also for estimating the ambient density near SNRs using radio continuum data (Uro{\\v s}evi{\\' c} \\& Pannuti 2005, Uro{\\v s}evi{\\' c}, Pannuti \\& Leahy 2007). In this work we argue for the presence of a thermal bremsstrahlung component in the radio emission from SNR HB 3 in addition to the synchrotron component. It is possible that SNRs can be sources with significant amount of thermal radiation and as Uro{\\v s}evi{\\' c} \\& Pannuti (2005) stated, there are two basic criteria for the production of a significant amount of radio emission through the thermal bremsstrahlung process from an SNR: the SNR evolves in an environment denser than the average and its temperature must be lower than the average (but always greater then the recombination temperature). Two cases have been considered: thermal emission at radio frequencies from a relatively young SNR evolving in a dense molecular cloud environment ($n\\approx100-1000\\ \\mathrm{cm^{-3}}$) and extremely evolved SNR (approximately $10^{5}-10^{6}$ years old) expanding in a dense warm medium ($n\\approx1-10\\ \\mathrm{cm^{-3}}$). For a detailed discussion about the issue see Uro{\\v s}evi{\\' c} \\& Pannuti (2005) and Uro{\\v s}evi{\\' c}, Pannuti \\& Leahy (2007).  Uro{\\v s}evi{\\' c}, Pannuti \\& Leahy (2007) have analyzed the broadband ($22-3900\\ \\mathrm{MHz}$) radio spectrum of the Galactic SNR HB 3 and have discussed the possible thermal radio emission from the SNR. Earlier published observations have revealed that a curvature is present in the radio spectrum of SNR HB 3, indicating that a single synchrotron component appears insufficient to adequately fit the radio spectrum. Uro{\\v s}evi{\\' c}, Pannuti \\& Leahy (2007) have suggested that more natural explanation for the apparent spectral index variations found by Tian \\& Leahy (2005) is synchrotron emission, which dominates at lower frequencies, and bremsstrahlung emission, which dominates at higher frequencies. They have found $\\approx40\\ \\%$ for the thermal component contribution to total volume emissivity at $1\\ \\mathrm{GHz}$ and have also estimated the ambient density, implied by the presence of thermal component, to be $\\approx10\\ \\mathrm{cm^{-3}}$. Recently, Green (2007) has reviewed the radio spectrum of SNR HB 3, noting the difficulties in deriving accurate flux densities for the remnant, particularly at higher frequencies, due to thermal emission from the nearby bright \\mbox{H\\,{\\sc ii}} region W3 (IC 1795) and its surroundings. He pointed out that some of the earlier published flux density values used by Uro{\\v s}evi{\\' c}, Pannuti \\& Leahy (2007) in their analysis are not accurate and that the spectrum of the SNR is well represented by a simple power low spectrum as well as the contamination with thermal emission from adjacent regions is the cause for the reported spectral flattening of the spectrum. In this work we present the results of our analysis using the corrected, recently reported by Green (2007), flux densities for SNR HB 3. ",
        "conclusions": "In this work we revised an analysis of the possible thermal emission at radio frequencies from an evolved SNR HB 3. Some of the earlier published flux density values for SNR HB 3 are shown to be non-accurate. Here we present the results of our analysis using the recently published, corrected, flux densities so the main conclusions are:  \\renewcommand{\\theenumi}{(\\arabic{enumi})}  \\item {1.} The contribution of thermal component in total volume emissivity was estimated to be $\\approx37 \\%$ and the ambient density was also estimated to be $n\\approx 9\\ \\mathrm{cm}^{-3}$ for the $\\mathrm{T}=10^{4}\\ \\mathrm{K}$.  \\item {2.} Our model estimates for thermal component contribution to total volume emissivity at $1\\ \\mathrm{GHz}$ and ambient density are similar to those obtained earlier ($\\approx40\\ \\%$, $\\approx10\\ \\mathrm{cm^{-3}}$). It is clear that the corrected flux density values do not change the basic conclusions.  \\item {3.} The presence of the thermal bremsstrahlung component in the radio spectrum of SNR HB 3 suggests that this SNR is in fact interacting with adjacent molecular cloud associated with the \\mbox{H\\,{\\sc ii}} region W3. The presence of thermal emission at radio frequencies may be a useful tool for identifying interactions between SNRs and molecular clouds and also for estimating the ambient density near SNRs using the radio continuum data.  \\item {4.} The lack of data at higher radio frequencies unable us to give a firmer conclusion about the issue."
    },
    "0809/0809.0469_arXiv.txt": {
        "abstract": "We solve one of the open problems in Einstein-Cartan theory, namely we find a natural matter source whose spin angular momentum tensor is compatible with the cosmological principle. We analyze the resulting evolution equations and find that an epoch of accelerated expansion is an attractor. The torsion field quickly decays in that period. Our results are interpreted in the context of the standard model of cosmology. ",
        "introduction": "General relativity is a successful theory in agreement with a vast number of observations. It is based on the Einstein-Hilbert action which yields the field equations if varied with respect to the metric. If, however, the metric and the connection (more precisely the non-Riemannian part of the connection) are considered as {\\it a priori} independent variables, two field equations are obtained. The first one relates the Einstein tensor (not necessarily symmetric) to the canonical energy-momentum tensor, while the other field equation relates the skew-symmetric part of the connection, the torsion tensor, to the spin angular momentum of matter, see e.g.~\\cite{Hehl:1973,Hehl:1974,Hehl:1976kj,Hehl:1995ue,Hammond:2002rm,Trautman:2006fp}. Spin and torsion are related by algebraic equations, and torsion vanishes in the absence of sources.  The cosmological principle states that the universe is homogeneous and isotropic on very large scales. More mathematically speaking  the four dimensional spacetime $(M,g)$ is defined by $3d$ space-like hypersurfaces of constant time which are orbits of a Lie group G action on $M$, with isometry group $SO(3)$. We assume all fields to be invariant under the action of G which means $\\mathcal{L}_{\\xi}g_{\\mu \\nu} = 0$ and $\\mathcal{L}_{\\xi}T_{\\mu \\nu}{}^{\\lambda} = 0$ where $\\mathcal{L}_{\\xi}$ denotes the Lie derivative with respect to the generator of the group. This assumption reduces the cosmological metric to the well know Friedman-Lema\\^{\\i}tre-Robertson-Walker form which is characterized by the scale factor and the geometry of the constant time hypersurfaces. If applied to the torsion of spacetime, it reduces the components compatible with the cosmological principle to a spatial axial torsion and a vector torsion part~\\cite{Tsamparlis:1979}.  Cosmological models with torsion were pioneered by Kopczy\\'nski in~\\cite{Kopczynski:1972,Kopczynski:1973}, who assumed a Weyssenhoff fluid~\\cite{Weyssenhoff:1947} to be the source of both curvature and torsion. The cosmological principle was first extended to Einstein-Cartan theory in~\\cite{Tsamparlis:1979}, where it was also suggested to reconsider the results in~\\cite{Kopczynski:1972,Kopczynski:1973}, since the Weyssenhoff fluid turns out to be incompatible with the cosmological principle (see also~\\cite{Obukhov:1987yu,Boehmer:2006gd,Brechet:2008tz}). An elaborate analysis of the most general action up to quadratic terms in curvature and torsion assuming the cosmological principle can be found in~\\cite{Goenner:1984}. Analytical solutions of the Riemann-squared gravity have recently been discussed in a cosmological context in~\\cite{Lasenby:2005nw}. Non-Riemannian models of cosmology in general have been discussed in~\\cite{Puetzfeld:2004df,Puetzfeld:2004yg,Puetzfeld:2004sw,Puetzfeld:2007ye}  However, nobody has so far succeeded in constructing a non-trivial spin angular momentum tensor in cosmology by minimally coupling matter fields to the geometry. We show that the minimally coupled eigenspinors of the charge conjugation operator~\\cite{jcap,prd} yield a spin tensor compatible with the cosmological principle.  These spinors belong to a wider class of so-called flagpole spinors~\\cite{daRocha:2005ti}. They are non-standard spinors according to the Wigner classification and obey the unusual property $(CPT)^2=-\\mathbbm{1}$. Hence, their dominant coupling to other fields is via the Higgs mechanism or via gravity~\\cite{jcap,prd}. The particles associated with such a field theory are naturally dark and we will refer to them as dark spinors henceforth, note that they were originally named Elko spinors.  Dark spinors are defined by \\begin{align} \\lambda = \\begin{pmatrix} \\pm \\sigma_2 \\phi^{\\ast}_{L} \\\\ \\phi_L \\end{pmatrix} \\,, \\label{eq:i1} \\end{align} where $\\phi^{\\ast}_{L}$ denotes the complex conjugate of $\\phi_L$ and $\\sigma_2$ denotes the second Pauli matrix. For a detailed treatment of the field theory of the eigenspinors of the charge conjugation operator we refer the reader to~\\cite{jcap,prd}. Dark spinors have an imaginary bi-orthogonal norm with respect to the standard Dirac dual $\\bar{\\psi} = \\psi^{\\dagger}\\gamma^0$, and in order for a consistent field theory to emerge the dual is given by \\begin{align} \\dual{\\lambda}_u = i\\,\\varepsilon_u^v \\lambda_v^{\\dagger} \\gamma^0 \\,, \\label{eq:i3} \\end{align} with $\\varepsilon_{\\{+,-\\}}^{\\{-,+\\}}=-1=-\\varepsilon_{\\{-,+\\}}^{\\{+,-\\}}$ such that \\begin{align} \\dual{\\lambda}_u(\\vp) \\lambda_v(\\vp) = \\pm\\, 2m\\, \\delta_{uv} \\,, \\label{eq:i4} \\end{align} where $\\vp$ denotes the momentum.  Due to their formal structure, dark spinors couple differently to gravitation than scalar fields or Dirac spinors~\\cite{Boehmer:2006qq}, eigenspinors of the parity operator. This allows for many interesting applications, for instance, in~\\cite{Boehmer:2007ut} it has been shown that dark spinors naturally yield an anisotropic expansion in the context of cosmological Bianchi type I models. This allows for a suppression of the low multipole amplitude of the primordial power spectrum. The primordial power spectrum of the quantum fluctuations of dark spinors has been investigated in~\\cite{Boehmer:2008rz,Gredat:2008qf} where is was found that the small scale power spectrum essentially agrees with that of scalar field inflation while the large scale power spectrum shows new features.  The paper is organized as follows. In Section~\\ref{sec:ec} we briefly summarize Einstein-Cartan theory. In Section~\\ref{sec:cos} the cosmological field equations in Einstein-Cartan theory in the presence of dark spinors are presented. We analyze the field equations qualitatively and numerically in Section~\\ref{sec:dyn} and conclude our work in Section~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  We identified a matter source whose spin-angular momentum tensor is compatible with the cosmological principle. We then solved the resulting field equations of Einstein-Cartan theory. This matter source consists of the eigenspinors of the charge conjugation operator, which we refer to as a dark spinor field. It couples to all irreducible parts of torsion and therefore leads to an interesting coupling of matter and geometry. This matter source is also naturally dark in that it can only interact via the Higgs mechanism or gravity.  Our solutions of the field equations show that torsion does vanish quickly (approximately after a few e-foldings) and that the Hubble parameter has a constant value as an attractor. Both features of the model fit very well into the standard model of inflationary cosmology in that a period of accelerated expansion is an attractor solution. It is worth noting that in Einstein-Cartan theory the spins of elementary particles are thought to be the primary sources of torsion, and it is therefore expected that on large sales and over time torsion should average out or decay, respectively.  We speculate that some non-zero cosmological torsion has already been observed in the large scale anisotropies of the cosmic microwave background radiation (CMB) where torsion leaves its imprint only on the largest scales."
    },
    "0809/0809.0972_arXiv.txt": {
        "abstract": "We present estimates of the GALEX near-ultraviolet (NUV) and far-ultraviolet (FUV) luminosity functions (LFs) of the Coma cluster, over a total area of $\\sim$9 deg$^{2}$ ($\\sim$25 Mpc$^2$), i.e. from the cluster center to the virial radius. Our analysis represents the widest and deepest UV investigation of  a nearby cluster of galaxies made to date. The Coma UV LFs show a faint-end slope steeper than the one observed in the local field. This difference, more evident in NUV,  is entirely due to the contribution of massive quiescent systems (e.g. ellipticals, lenticulars and passive spirals), more frequent in high density environments. On the contrary, the shape of the UV LFs for Coma star-forming galaxies does not appear to be significantly different from that of the field, consistently with previous studies of local and high redshift clusters. We demonstrate that such similarity is only a selection effect, not providing any information on the role of the environment on the star formation history of cluster galaxies. By integrating the UV LFs for star-forming galaxies (corrected for the first time for internal dust attenuation), we show that the specific star formation rate (SSFR) of Coma is significantly lower than the integrated SSFR of the field and that Coma-like clusters contribute only $<$7\\% of the total SFR density of the local universe. Approximately 2/3 of the whole star-formation in Coma is occurring in galaxies with M$_{star}<$10$^{10}$ M$_{\\odot}$. The vast majority of star-forming galaxies has likely just started its first dive into the cluster core and has not yet been affected by the cluster environment. The total stellar mass accretion rate of Coma is $\\sim$(0.6-1.8) $\\times$ 10$^{12}$ M$_{\\odot}$ Gyr$^{-1}$, suggesting that a significant fraction of the population of lenticular and passive spirals observed today in Coma could originate from infalling galaxies accreted between $z\\sim$1 and $z\\sim$0. ",
        "introduction": "Star formation is probably the most fundamental of all astrophysical processes regulating galaxy evolution: not only the properties of galaxies depend on how their stars were formed, but the properties of the interstellar medium (ISM) are controlled to a large extent by the various feedback effects of star formation. The overall evolution of galaxies depends on the rate at which their interstellar gas is converted into stars and the star formation rate (SFR) depends on the rate at which diffuse interstellar matter is collected into star forming regions. Consequently, it is of great importance for galaxy evolution studies to understand how stars form and what determines their properties.  The cluster environment is well known to affect both the star formation activity and morphology of the member galaxies, resulting in the morphology-density (e.g. \\citealp{dressler80,whitmore}) and star formation-density (e.g. \\citealp{lewis02,gomez03,haanna,review}) relations. The scale over which the number density of elliptical galaxies increases toward the center of rich clusters is very small ($\\leq$ 0.5 Mpc), while the one over which the star formation activity and the fraction of late-type galaxies decrease is much larger (a few Mpc). This indicates that separate processes contribute to shaping these relations. The mechanism involved in the early-type central enhancement is probably connected to the early phases of structure formation (e.g. \\citealp{delucia06,ellis97,stanford98,renzini06}), while the progressive quenching of the star formation and lack of spirals toward the center of rich clusters could represent a separate, diluted in space and time, process probably related to the cluster accretion history (e.g. \\citealp{dress04car,review,big,poggianti06}).  Although the effects of the environment on the star formation activity of cluster galaxies are gradually being unveiled (e.g. \\citealp{review,n4569,dress04car,harrassment,poggia99,poggianti06,vollmer01}), we still know very little about the contribution of clusters of galaxies to the whole star formation history (SFH) and the stellar mass growth of the universe. The cosmic SFR dramatically decreases from z$\\sim$1 to the present epoch \\citep{lilly96,madau96}, but we do not know if this drop varies with the environment \\citep{kodama01_cl,cooper08,finn08}, what fraction of the cosmic SFR takes place in cluster galaxies and how this number evolves with redshift.  In addition, a significant increase from $z\\sim$1 to $z=$0 in the stellar mass density of the universe has been observed (e.g. \\citealp{borch06,drory05,fontana06}). Interestingly very little growth is found in the stellar mass of blue sequence galaxies, and the vast majority of stars formed since $z\\sim$1 appears to inhabit galaxies which, at the present epoch, lie in the red sequence (e.g. \\citealp{cimatti06,bell07,brown07,faber07}). This automatically implies that a significant fraction of star-forming galaxies has migrated from the blue to the red sequence. The mechanism responsible for the quenching of the star formation in blue galaxies is still unknown and remains one of the greatest challenges for modern extragalactic astronomy. However, it is interesting to note that a similar migration, from star-forming to passive galaxies, is usually observed in high density environments (e.g. \\citealp{dEale}), perhaps suggesting that the environment might play an important role.  In order to clarify these issues, it is therefore crucial to quantify the contribution of galaxy clusters to the total SFR and stellar mass budget of the universe, and their evolution with redshift. To reach this goal we need to determine the SFR, luminosity and mass functions of the cluster population up to the virial radius, avoiding bias in our determination of the global cluster properties by observing a population that is not representative of the whole cluster environment. This kind of study has become possible only recently thanks to the advent of wide field surveys. In particular, the Galaxy Evolution Explorer (GALEX, \\citealp{martin05}) is providing for the first time precise UV photometry of galaxies over large stretches of the sky, opening a new era of extragalactic UV astronomy. The UV emission is dominated by young stars of intermediate masses (2$<M<$5 M$_{\\odot}$, e.g. \\citealp{kenn98}) and thus represents an ideal tracer to identify and quantify the star formation activity of clusters. Surprisingly, very few statistical studies of galaxy clusters in UV have been carried out so far and great part of what we know about UV properties of cluster galaxies is still based on pioneering observations by the FAUST telescope (e.g. \\citealp{deharveng94,brosh97}) and the balloon-borne UV telescopes SCAP and FOCA (e.g. \\citealp{andreon,COGA03,donas87,donas95}). For example, we still lack an accurate determination of the UV LFs in high density environments. The FOCA UV LFs were in fact affected by large statistical errors due to the insufficient redshift coverage for UV-selected galaxies and to the uncertainty in the UV background counts. Moreover, the only GALEX UV LF for a nearby cluster presented so far (Abell1367, \\citealp{CORBGA05}) does not sample the whole cluster region and suffers from low number statistics at high luminosities. Thus, in order to investigate the UV properties of nearby clusters we carried out a wide panoramic survey in near-ultraviolet (NUV) and far-ultraviolet (FUV) of the nearby Coma cluster of galaxies. \\\\ Coma (Abell 1656) is one of the nearest richest clusters, and it is considered as the prototype of evolved, relaxed cluster of galaxies. In the last decades, it has been one of the main laboratory for environmental studies at all wavelengths (e.g. see \\citealp{review} and references therein) and more recently it has become the focus of a Hubble/ACS Treasury survey \\citep{carter08} making it an ideal target for a deep census of UV properties in a nearby cluster up to its virial radius.\\\\ In this paper, we present the determination of the near-ultraviolet (NUV) and far-ultraviolet (FUV) LFs, the total SFR of the Coma cluster and its contribution to the SFR density in the local universe.  We assume a distance modulus of $m-M$ = 35.0 mag for the Coma cluster, corresponding to a distance of 100 Mpc  and a scale of  1.67 Mpc deg$^{-1}$ for $H_{0}$=70 Mpc$^{-1}$ km s$^{-1}$. ",
        "conclusions": "In this paper, we have presented a study of the UV properties of the Coma cluster based on GALEX NUV and FUV observations covering $\\sim$9 deg$^{2}$ centered on the cluster core. Although the central $\\sim$0.26 deg$^{2}$ could not be observed, our analysis represents the widest and deepest UV investigation of  a nearby cluster of galaxies made to date. Our main results are as follows:\\\\  a) The Coma NUV and FUV LFs show a faint end slope significantly steeper than the one observed in the field. This difference is more evident in NUV and it is due to the higher number density of massive quiescent/red galaxies (i.e. ellipticals, lenticulars and passive spirals) in Coma compared to the field. The contribution of quiescent galaxies to the total UV emission at low luminosities ($M(UV)>-$17 mag) is larger in the cluster center, however no significant variation in the shape of the UV LFs with cluster-centric distance is observed.\\\\  b) We estimated for the first time the UV LFs of star-forming galaxies corrected for internal dust attenuation. We showed that Coma-like clusters contribute only $<$7\\% of the total SFR density of the local universe. More interestingly the SSFR of Coma is $\\sim$ 10$^{-11.18\\pm0.13}$ yr$^{-1}$, significantly lower than the integrated SSFR of the local universe. Approximately 2/3 of the whole SF in Coma is occurring in objects with M$_{star}<$ 10$^{10}$ M$_{star}$, confirming that \\emph{downsizing} is also present in high density environments. \\\\  c) The shape of the UV LF and the SSFR of blue/star-forming galaxies are consistent with those of the field, in agreement with previous works. We have shown that this similarity does not imply that the effects of the environment on the evolution of the cosmic SFH are negligible. On the contrary, these results are consistent with a scenario in which cluster star-forming galaxies are still infalling for the first time into the cluster center. The stellar mass accretion rate of Coma results $\\sim$(0.6-1.8)$\\times$10$^{12}$ M$_{\\odot}$ Gyr$^{-1}$.  At this rate, the whole cluster could have easily formed from infalling galaxies accreted from the field in a Hubble time. More interestingly, a significant fraction of the population of lenticular and passive spirals observed today in Coma could arise from infalling galaxies accreted between $z\\sim$1 and $z\\sim$0, perhaps suggesting that the environment plays a significant role in the mass growth of the red sequence in the universe."
    },
    "0809/0809.3006_arXiv.txt": {
        "abstract": "Models of the chemical evolution of our Galaxy are extended to include radial migration of stars and flow of gas through the disc. The models track the production of both iron and $\\alpha$ elements. A model is chosen that provides an excellent fit to the metallicity distribution of stars in the Geneva--Copenhagen survey (GCS) of the solar neighbourhood, and a good fit to the local Hess diagram. The model provides a good fit to the distribution of GCS stars in the age--metallicity plane although this plane was not used in the fitting process.  Although this model's star-formation rate is monotonic declining, its disc naturally splits into an $\\alpha$-enhanced thick disc and a normal thin disc. In particular the model's distribution of stars in the ([O/Fe],[Fe/H]) plane resembles that of Galactic stars in displaying a ridge line for each disc. The thin-disc's ridge line is entirely due to stellar migration and there is the characteristic variation of stellar angular momentum along it that has been noted by Haywood in survey data. Radial mixing of stellar populations with high $\\sigma_z$ from inner regions of the disc to the solar neighbourhood provides a natural explanation of why measurements yield a steeper increase of $\\sigma_z$ with age than predicted by theory. The metallicity gradient in the ISM is predicted to be steeper than in earlier models, but appears to be in good agreement with data for both our Galaxy and external galaxies. The models are inconsistent with a cutoff in the star-formation rate at low gas surface densities. The absolute magnitude of the disc is given as a function of time in several photometric bands, and radial colour profiles are plotted for representative times. ",
        "introduction": "Models of the chemical evolution of galaxies are key tools in the push to understand how galaxies formed and have evolved. Their application to our Galaxy is of particular importance both on account of the wealth of observational data that they can be required to reproduce, and on account of the inherent interest in deciphering the history of our environment.  From the pioneering papers by \\cite{vdB62} and \\cite{Schmidt63} it has generally been assumed that a galaxy such as the Milky Way can be divided into concentric cylindrical annuli, each of which evolves independently of the others \\cite[e.g.][]{Pagel97,Chiappini97,Chiappini01,NaabO06,Colavitti08}.  The contents of any given cylinder are initially gaseous and of extremely low or zero metallicity. Over time stars form in the cylinder and the more massive ones die, returning a mixture of heavy elements to the remaining gas. The consequent increase in the metallicity of the gas and newly-formed stars is generally moderated by an inflow of gas from intergalactic space, and, less often, by an outflow of supernova-heated gas.  The cool, star-forming gas within any cylinder is assumed to be well mixed, so at any time it can be characterised by a metallicity $Z(r,t)$, where $r$ is the cylinder's radius. Hence the stars formed within a given cylinder should have metallicities $Z(r,t_{\\rm f})$ that are uniquely related to their time of formation, $t_{\\rm f}$. Observations do not substantiate this prediction; in fact \\cite{Edvardsson93} showed that solar-neighbourhood stars are widely distributed in the $(t_{\\rm f},Z)$ plane -- for a detailed discussion see \\cite{Haywood06} and Section~\\ref{sec:Haywood}.  The absence of an age-metallicity relation in the solar neighbourhood is naturally explained by radial migration of stars \\citep{SellwoodB,Haywood08,Roskar2}. It has been recognised for many years that scattering by spiral structure and molecular clouds gradually heats the stellar disc, moving stars onto ever more eccentric and inclined orbits. Stars that are on eccentric orbits clearly contribute to different cylindrical annuli at different phases of their orbits, and thus tend to modify any radial gradient in the metallicities of newly formed stars.  Moreover, scattering events also change the guiding centres of stellar orbits, so even a star on a circular orbit can be found at a different radius from that of its birth. In fact, \\cite{SellwoodB} argued that the dominant effect of transient spiral structure is resonant scattering of stars across the structure's corotation resonance, so even a star that is still on a near-circular orbit may be far from its radius of birth. \\cite{Roskar1} showed that in a cosmological simulation of galaxy formation that included both stars and gas, resonant scattering at corotation caused stars to move outwards and gas inwards, with the result that the stellar disc extended beyond the outer limit of star formation; the outer disc was entirely populated by stars that had formed much further in and yet were still on nearly circular orbits. This simulation confirmed the conjecture of \\cite{SellwoodB} that gas would participate in resonant scattering alongside stars.  We distinguish two drivers of radial migration: when the angular momentum of a star is changed, whether by scattering at an orbital resonance or by non-resonant scattering by a molecular cloud, the star's guiding-centre radius changes and the star's entire orbit moves inwards or outwards depending on whether angular momentum is lost or gained. When a scattering event increases a star's  epicycle amplitude without changing its angular momentum, the star contributes to the density over a wider range of radii. In a slight modification of the terminology introduced by \\cite{SellwoodB}, we say that changes in angular momentum cause ``churning'' while changes in epicycle amplitude lead to ``blurring''. This paper extends models of Galactic chemical evolution to include the effects of churning and blurring.  Given the strength of the arguments that cold gas should participate in churning alongside stars, and that shocks induced by spiral structure cause gas to drift inwards, it is mandatory simultaneously to extend traditional chemical evolution models to include radial flows of gas within the disc. \\cite{LaceyF85} studied chemical evolution in the presence of a radial inflow of gas and demonstrated that a radial flow enhances the metallicity gradient within the disc. This enhancement plays an important role in our models, which differ from those of Lacey \\& Fall in that they include both radial gas flows and radial migration of stars. Moreover we can fit our models to observational data that is much richer than that available to \\cite{LaceyF85}.  Our models are complementary to ab-initio models of galaxy formation such as those presented by \\cite{SamlandG} and \\cite{Roskar2} in that they allow the solar neighbourhood to be resolved in greater detail, and because they are enormously less costly numerically, they permit parameter searches to be made that are not feasible with ab-initio models.  The paper is organized as follows. Section \\ref{sec:goveqs} presents the equations upon which the models are based. These consist of the rules that determine the rate of infall of fresh gas, the rate of star formation, details of the stellar evolution tracks and chemical yields that we have used and descriptions of how churning and blurring are implemented. Section \\ref{sec:standard} describes in some detail a ``standard'' model of the evolution of the Galactic disc. This covers its global properties but focuses on what would be seen in a survey of the solar neighbourhood. Section \\ref{sec:Snhd} presents the  details of the selection function that is required to mimic the Geneva--Copenhagen sample (GCS) of solar-neighbourhood stars published by \\cite{Nordstrom04} and \\cite{HolmbergNA}, and explains how this sample has been used to constrain the model's parameters. Section \\ref{sec:trend} explains how the observable properties of the model depend on its parameters.  Section \\ref{sec:reln} discusses the relation of the present models to earlier ones, and discusses the extent to which it is consistent with the analysis of solar-neighbourhood data by \\cite{Haywood08}. Section \\ref{sec:conclude} sums up. ",
        "conclusions": "\\label{sec:conclude}  It is now more than forty years since the theory of stellar evolution attained the level at which it became possible to model the chemical enrichment of the ISM. From the beginning of that endeavour measurements of the abundances of individual solar-neighbourhood stars have played a key role because a star preserves like a time capsule the state of the ISM at the remote epoch of its formation. Considerable theoretical and observational efforts have been devoted to probing the history of the Galaxy with this connection.  Half a century ago, \\cite{Roman50,Roman54} and others discovered the connection between the kinematics and chemistry of stars, yet curiously little has been done to include kinematics in models of chemical evolution. The general presumption has been that each annulus of the disc evolves independently of others, and the well-known correlations between chemistry and kinematics can be understood as arising through the stochastic acceleration of stars: older stars tend to have larger random velocities and lower metallicities. No effort was made to develop greater diagnostic power by simultaneously modelling chemistry and kinematics.  The continued use of mutually independent annuli by modellers of chemical evolution is surprising given that it was from the outset recognised that many stars are on significantly non-circular orbits that each radial half-period cover more than a kiloparsec in radius. In fact, it has generally been assumed that the radial velocity dispersion within the disc rises as one moves inward to values that lead to radial excursions of several kiloparsecs (see \\figref{fig:blurr}). Moreover, observations have long indicated that galactic discs have significant metallicity gradients \\citep[e.g.][]{Vila-CostasE92,deJong96}, so radial migration of stars is bound to leave a signature on the metallicity distribution of solar-neighbourhood stars. We have called this aspect of radial migration ``blurring''.  The present study owes its impetus to the discovery by \\cite{SellwoodB} that radial migration is a more potent process than mere blurring: the dominant effect of transient spiral arms is not to heat the stellar disc as had been supposed, but to cause stars either side of corotation to change places without moving to eccentric orbits (``churning''). \\cite{SellwoodB} did not demonstrate that gas participates in churning, but they argued that it must on dynamical grounds, and \\cite{Roskar1} found evidence that this was the case in their N-body--SPH simulations of galaxy formation.  Since churning is an aspect of spiral structure that can {\\it only\\/} be probed through its impact on chemical evolution, we wanted chemical-evolution models that included churning, and logically it was natural to extend these models to include both blurring and radial gas flows.  An ingredient of our models that might be controversial is the introduction of radial gas flows. We have absolutely no reason to expect that the infall profile is exponential, so if discs are to be exponential this must be the result of flows through the discs redistributing mass within the disc. In fact, as gas streams through spiral arms it dissipates energy in shocks that is ultimately gravitational energy that becomes free as the gas surrenders angular momentum to the stars and drifts inwards. Hence at some level inward gas flows are mandatory \\citep{LaceyF85}.  Unfortunately, the theory of galaxy formation has yet to advance to the point at which it can prescribe the spatial and temporal structure of gas accretion, so it is necessary to parametrise accretion in some way. The accretion process must be constrained to result in the formation of the observed stellar and gaseous discs. Our accretion Scheme AB satisfies this constraint for all values of the parameters, but it is inevitably acausal in that the formation mechanism is being driven by its known outcome. While its acausality is unattractive, Scheme AB is a flexible parametrisation that enables us to form exponential discs for a variety of different assumptions about the radial density of infalling gas, and the resulting radial profiles of infall and gas-flow (\\figref{fig:drag}) are entirely plausible.  The impact that radial migration has on the local metallicity distribution obviously varies with the magnitude of the metallicity gradient in the ISM, which in turn depends on the gas flow within the disc and therefore the radial infall profile. For this reason the most important parameters of our models are $f_{\\rm A}$ and $f_{\\rm B}$, which control the distribution of infalling gas.  The models provide good fits to the GCS counts of stars as functions of $\\feh$, $M_V$, $T_{\\rm eff}$ and stellar age, as well as reproducing the correlations between tangential velocity and abundance patterns that have been pointed out by \\cite{Haywood08}. These fits are achieved by churning, blurring and radial flows working together. They depend on the existence of an appreciable metallicity gradient in the ISM, which is established by the radial flow of gas, and they depend on radial mixing of stars by blurring and churning. The steeper the metallicity gradient in the gas, the smaller the effect of churning can be, but for any observationally consistent metallicity gradient, churning has a non-negligible role to play.  The models describe the coevolution of the thick and thin discs, and presume that thick-disc stars were formed in the Galaxy rather than accreted from outside. Given the simplicity of our assumptions, the extent to which a dichotomy between an $\\alpha$-enhanced thick disc and a solar-type thin disc automatically manifests itself in the models is remarkable. In particular, in the solar annulus the distributions of [O/H] at given $\\feh$ are bimodal in the range of $\\feh$ associated with the overlap of the two discs (\\figref{fig:alpha_thick}), the vertical density profile can be represented at the sum of two exponentials, and at $R<R_0$ the radial density profile becomes flatter at lager distances from the plane (\\figref{fig:thick_thin}). None of these characteristics is dependent on our choice of a double exponential for the time dependence of infall: a model in which the gaseous mass, and therefore star-formation rate, is held constant also has these features. They are consequences of the $\\sim1\\Gyr$ timescale of type Ia supernovae and the secular heating and churning of the disc.  The models assume that scattering of stars increases the velocity dispersion of a coeval population as $t^{1/3}$, but the models go on to predict that within the solar neighbourhood  velocity dispersion increases as a higher power of age, roughly $t^{0.45}$ in agreement with what is found from Hipparcos stars with good parallaxes. This finding may reconcile scattering theory, which cannot readily explain an exponent in excess of $1/3$, with observations. The key point is that radial mixing brings to the solar neighbourhood stars born at small radii, where the velocity dispersion is undoubtedly large.  The nucleosynthetic yields from each generation of stars are still significantly uncertain. Our philosophy has been to use standard values from the literature rather than exploit uncertainties in the yields to tune the models to the data. The yields we are using are in the upper region of those that can provide adequate fits to the data, with the consequence that increasing the value of the mass-loss parameter $f_{\\rm eject}$, which is degenerate with the magnitude of the yields, makes it easier to fit the data.  Our yields are surely not exactly right and consequently some of the properties the models derive from them will be in error. On account of these uncertainties we have suppressed the predictions of our code for the abundances of certain elements, most notably carbon.  The discovery that three-dimensional, non-equilibrium models of the solar atmosphere require the metal abundance of the Sun to be $Z_\\odot=0.012-0.014$ \\citep{Grevesse07} poses a major problem for this field. A prerequisite for successful chemical modelling is a consistent metallicity scale for both stars and the ISM. At present the only consistent scale is the traditional one on which $Z_\\odot=0.019$, so this is the one we have used.  If a new scale were established on which all metallicities were significantly lower, viable models could be produced by lowering the yields. A straightforward way to do this would be to lower the maximum mass in the IMF: reducing this mass from $100\\msun$ to $50\\msun$ would reduce yields by $\\sim30$ percent in line with the proposed reduction in $Z_\\odot$. The oxygen yield would come down fastest, reducing [O/Fe] by $\\sim0.2\\dex$.  Although we believe that this study represents a significant advance on all previous models of Galactic chemical evolution, it is highly imperfect. Some major weaknesses of our work are the following.  \\begin{itemize} \\item[1.] We have  assumed that the probability of mass interchange between rings is proportional to the product of the rings' masses. This probability reflects the number and intensity of spiral feature with corotation at that radius, and should be a function of both mass and velocity dispersion. A further study of self-gravitating discs similar to that of \\cite{SellwoodB} would be necessary to determine this function.  \\item[2.] We have assumed that the vertical and radial motions of stars decouple. This assumption has a significant impact on both the relation between age and velocity dispersion in the solar neighbourhood and on the predicted vertical density profile in the solar annulus, which is interesting in itself and impacts on the selection function of GCS stars and thus on our choice of standard model.  The assumption is unjustifiable for stars on eccentric orbits, which do play an important role in the model fits. Unfortunately, a sounder treatment is impossible until a better approximation to the third integral of Galactic dynamics is available. It is our intention to resolve this problem through ``torus modelling'' \\citep[e.g.][]{McmillanB08}.  \\item[3.] Our models include radial mixing of gas through churning and viscous inspiralling, but do not include radial redistribution of gas by the Galactic fountain \\citep[e.g.][]{Benjamin97}. A significant body of evidence indicates that star formation drives neutral hydrogen to kiloparsec heights above the plane. NGC$\\,$891, which is not dissimilar to the Galaxy, has $\\sim25$ per cent of its neutral hydrogen more than $1\\kpc$ from the plane \\citep{OosterlooFS}.  This extraplanar gas must move over the plane on nearly ballistic trajectories, and in the absence of interaction with the corona (gas at the Galaxy's virial temperature $\\sim2\\times10^6\\K$), the gas must return to the plane further out than its point of ejection. However, observations suggest that neutral gas above the plane is actually flowing inward rather than outward, presumable as a result of interaction with the corona \\citep{FraternaliB08}. The mass of extraplanar gas is so large and the timescale for its return to the plane so short that whichever way this gas flows, it has a considerable potential for radially redistributing metals. Extraplanar gas must be ejected from the disc by supernova-heated gas that is probably highly metal-rich. Some of this gas will be lost to the Galaxy as we have assumed, but some of it will return to the plane with infalling gas. Again there is potential here for significant radial redistribution of metals that has been ignored in the present study.  \\item[4.] Our treatment of the inner Galaxy is unacceptably crude in that we have replaced the bar/bulge with a disc. Unfortunately introducing the bar opens a Pandora's box of complexities, and at the present time is probably only feasible in the context of a particle-based model such as those of \\cite{SamlandG} and \\cite{Roskar2}. Our hope is that our imaginary central disc has a similar impact on the chemodynamical state of the local disc to the combined impact of the giant star-forming ring at $R\\simeq4\\kpc$, the gas-deficient region interior to this ring and the star-forming $x_2$ disc at $R\\lta200\\pc$.  \\end{itemize}  From this list of shortcomings of our models it is clear that we are still far from a definitive account of the Galaxy's chemical evolution. We shall be satisfied if we have convinced the reader that the interplay between dynamics and chemistry is so tight at to be indissoluble. This fact is at one level inconvenient because it undermines the value of conclusions drawn from traditional models of chemical evolution. But at another level it represents an opportunity to learn more. The connection between the kinematics of stars and the compositions of stars and gas, which can be measured in great detail, involves three areas about which we are too ignorant: the distribution of dark matter within and around the Galaxy, the Galaxy's history of assembly, and the nature of the Local Group's  IGM. By building dynamical models of  the Galaxy that have chemical evolution built in to the basic structure, we should be able to make decisive progress with one of the major problems of contemporary astronomy.    We thank M. Haywood, J.-U. Ness and A. Riffeser for fruitful discussions. R.S. acknowledges material and financial support from the Studienstiftung des Deutschen Volkes and Stiftung Maximilianeum and the hospitality of Merton College Oxford, where this work began."
    },
    "0809/0809.1003_arXiv.txt": {
        "abstract": "Efforts to place limits on deviations from canonical formulations of electromagnetism and gravity have probed length scales increasing dramatically over time. Historically, these studies have passed through three stages:  (1) Testing the power in the inverse-square laws of Newton and Coulomb,  (2) Seeking a nonzero value for the rest mass of photon or graviton, (3) Considering more degrees of freedom, allowing mass while preserving explicit gauge or general-coordinate invariance.  Since our previous review the lower limit on the photon Compton wavelength has improved by four orders of magnitude, to about one astronomical unit, and rapid current progress in astronomy makes further advance likely.  For gravity there have been vigorous debates about even the concept of graviton rest mass.  Meanwhile there are striking observations of astronomical motions that do not fit Einstein gravity with visible sources.  ``Cold dark matter\" (slow, invisible classical particles) fits well at large scales. ``Modified Newtonian dynamics\" provides the best phenomenology at galactic scales. Satisfying this phenomenology is a requirement if dark matter, perhaps as invisible classical fields, could be correct here too.  ``Dark energy\" {\\it might} be explained by a graviton-mass-like effect, with associated Compton wavelength comparable to the radius of the visible universe.  We summarize  significant mass limits in a table. ",
        "introduction": "Photons and gravitons are the only known free particles whose rest masses may be exactly zero.\\footnote{ Gluons, the gauge particles of quantum chromodynamics, are believed to have no bare mass.  However, they are not seen in isolation, meaning they cannot be observed as free particles. } This bald statement covers a rich and complex history, from Newton and Gauss, through Maxwell and Einstein, and even up to the present.   During that development, the matrix of interlocking concepts surrounding the notions of photon and graviton rest mass, or more generally, long-distance, low-frequency deviations from Maxwell electrodynamics and Einstein gravity, has become increasingly elaborate.  There are many similarities between the photon and graviton  cases, but also striking differences.  We literally see photons all the time, as only a few photons of visible light are enough to activate one `pixel' in a human retina.  Besides that, conspicuous electromagnetic wave phenomena play an enormous role in modern physics.  For gravity the situation is radically different.  Even gravitational waves are in a situation analogous to that of neutrinos during the first 25 years after their proposal by Pauli, when their emission could be inferred from loss of energy, momentum, and angular momentum in beta decay, but they hadn't yet been detected in an absorption experiment.  Binary pulsar systems exhibit energy loss well accounted-for by radiation of gravitational waves, but experiments underway to detect absorption of such waves have not yet achieved positive results.   Even if this were accomplished sometime soon, the chances  of ever detecting individual quanta -- gravitons -- seem remote indeed, because graviton coupling to matter is so enormously weak.  These are not the only differences.  For electrodynamics, the history has been one of increasingly sensitive null experiments giving increasingly stringent constraints on a possible mass.   Nevertheless, theories describing such a mass seem well-developed and consistent, even if not esthetically appealing.  On the other hand, for gravity, there {\\it are} long-distance effects which some argue provide evidence supporting modification of Einstein's formulation.  At the same time, the theoretical basis for a gravitational phenomenon analogous to photon mass has come under severe attack. All this means that, with respect to the issue of deviations, currently there is much more dynamism for gravity than for electrodynamics.  Even though quantum physics gave shape to the concept of mass for electrodynamics and gravitation, the obvious implication of a dispersion in velocity  with energy for field quanta, or even for waves, is beyond our capacity to detect with methods identified so far.  This is thanks to the very strong limits already obtained based on essentially static fields.  Thus, the domain of potentially interesting  experiment and observation for mass or ``mass-like\" effects indeed is restricted to the long-distance and low-frequency scales already mentioned.    \\subsection{How to test a theory}  Let us begin by seeking a broad perspective on what it means to probe, not merely the validity but also the accuracy of a theory.   The canonical view of theory-testing  is that one tries to falsify the theory:   One compares its predictions with experiment and observation. The predictions use input data, for example initial values of certain parameters, which then are translated by the theory into predictions of new data.  If these predicted  data  agree with observation within experimental uncertainties (and sometimes also uncertainties in application of the theory), then the theory has, for the moment, passed the test. One may continue to look for failures in new domains of application, even if  the incentive for doing so  declines with time.  Of course, without strong `ground rules' it is impossible to falsify a theory, because one almost always can find explanations for a failure.  So, in fact no scientific theory either may be disproved or proved in a completely rigorous way; everything always is provisional, and continual skepticism always is in order.  However, based on a strong pattern of success a  theory can earn trust at least as great as in any other aspect of human inquiry.  The above is an essential, but we believe only partial, view of how theories gain conviction.  At least three important additional factors may help to achieve that result.   First, a striking, even implausible, prediction is borne out by experiment or observation. Examples of this include `Poisson's spot', Poisson's devastating attack on the notion that light is a wave phenomenon, because this would require that the shadow of a circular obstacle have a bright spot in its center.   The discovery of the spot by Arago provided the conceptual equivalent of a judo maneuver, using the opponent's own impulse to overcome him.  Another example is the assertion by Appelquist and Politzer in the summer of 1974 that the existence of a heavy quark carrying the quantum number `charm' would imply the existence of a positronium-like spectrum, meaning very sharp resonances in electron-positron scattering.  When the J/$\\psi$ was discovered at BNL and SLAC in November of that year, the outlandish prediction of Appelquist and Politzer suddenly was the best explanation, in good part because it was the only one that had been stated boldly beforehand (though only published afterwards).  A second way in which a theory gains credence is by fecundity:   People see ways to apply the idea in other contexts.  If many such applications are fruitful, then by the time initial experimental verification is rechecked there may be little interest, because the theory already has become a foundation stone for a whole array of applications. An example from our subject here is the transfer of the $1/R^2$ force law from gravity to electricity, in a speculative leap during the 1700s.  Of course, it was not this transfer which gave Newton's gravity its great authority, but rather the enormous number of precise and successful predictions of his theory -- fecundity in the original, literally astronomical domain.  Closely related to the above is a third feature, connectivity.   If many closely neighboring subjects are described by connecting theoretical concepts, then the theoretical structure acquires a robustness which makes it increasingly hard -- though certainly never impossible -- to overturn.  The latter two concepts fit very well with Thomas Kuhn's epilogue to his magnum opus {\\it The Structure of Scientific Revolutions} \\cite{kuhn}, in which he muses that the best description of scientific development may be as an evolutionary process.   For biological evolution, both  fecundity and reinforcing connections play decisive roles in how things happen.  It seems to us, as it did to Kuhn, that the same is often true for {\\it ideas} in all science, including physics.  To the extent that observational errors can be ruled out, whenever discrepancies appear between theory and experiment one is compelled to contemplate the possibility that new, or at  least  previously unaccounted-for, physics is contributing to the phenomena.  A classic case is the famous solar-neutrino puzzle.  At first, presumed errors in the actual measurements themselves were widely and caustically viewed as the problem.  Later critiques focused more on  consideration of possible errors in models of processes in the Sun and, perhaps more creatively, on potential new physics modifying the simplest picture  of neutrino propagation from source to detector.   By now there is overwhelming evidence that neutrino mixing, a modification of neutrino propagation, accounts precisely for the observed rate of neutrino observation on Earth, confirming the basic validity of  early work on both solar modeling and neutrino detection.  A more refined question about confirming a theory is:   How does one quantify limits on deviations from the theory?   Now it becomes necessary to specify some form of  deviation that depends on certain parameters.  Then experimental uncertainties can be translated into quantitative limits on these parameters.   As a theory evolves, the favored choice for  an interesting form of deviation may change. Of course, at any point new observations contradicting even well-established theoretical predictions can reopen issues that might have seemed settled    \\subsection{Photons}  Our earlier review \\cite{GN}  described the state of theory for accommodating a non-zero photon mass, and the state of observation and experiment giving limits on the mass at that time.   Since then there have been significant theoretical developments (Section \\ref{EM}), as well as advances in experimental approach and precision (Section \\ref{EMtests}).  The issue of possible long-range  deviations from the existing theory of electrodynamics  long remained purely a matter of choosing, and then setting limits on, parameters of a possible deviation. Originally, in analogy to studies of Newton's Law for gravity, deviations  in power of radius from that in the inverse square law were used.  20th-century relativistic wave equations led to discussion in terms of a finite rest mass of the photon, thus introducing a length scale. Today one can consider a more sophisticated approach explicitly incorporating gauge invariance, even when describing nonzero photon mass, by including more couplings.  We shall not reprise in detail the theoretical paradigms and experimental details discussed in Ref. \\cite{GN} for the photon.  Rather, we refer the reader to that work for an introduction, and focus here on elucidating more recent advances.   Also, since the publication of \\cite{GN}, there have been other works which have summarized specific aspects of the photon mass  \\cite{early}-\\cite{pdg}.  These can be consulted especially for experimental summaries.  In particular, Byrne \\cite{early} concentrated on astrophysical limits, Tu and Luo \\cite{metro}  on laboratory limits based on tests of Coulomb's Law, and, joined by Gillies, on all experimental tests  \\cite{luorev}.  Okun, in a compact review \\cite{okun}, gives some  interesting early history on the concept of the ``photon\" in quantum mechanics. He also gives details on what the Russian school accomplished during the earlier period. Vainshtein has presented a concise, and  insightful, perspective on theoretical issues relating to mass both for photon and graviton \\cite{Vainshtein:2006ix}.  Of course the types of deviation we discuss in this review do not cover all possibilities.   In particular, the Maxwell equations are linear in  the electromagnetic field strengths.  This does not mean that all phenomena are linear, because the coupling between fields and currents allows back-reaction and thus nonlinearity.  Nevertheless, the linearity of the equations is an especially simple feature.   Higher-order terms in the field strengths clearly are an interesting possibility, but they are intrinsically tied to short-distance modifications of the theory, rather than the long-distance deviations emphasized here.  The reason is that the nonlinear terms become more important as the field strengths increase, meaning that the numbers of flux lines per unit area increase, clearly a phenomenon associated with short distance scales.  For completeness, we  briefly mention here discussions in the literature that go in the direction of nonlinearity.   Born and Infeld \\cite{Infeld} introduced the notion of nonlinear damping of electromagnetic fields, precisely to cope with the short-distance singularities of the classical linear theory.   Their approach was pursued by a number of investigators over many years, as reviewed by Plebanski \\cite{Plebanski}.  More recently their ideas have been revived because the kind of structure they discussed arises naturally in string  theory,\\footnote{ We refer a number of times in this paper to string theory.   Even though  incomplete at this stage, one way it can be viewed is as a natural progression from field theories such as electrodynamics and general relativity.  In a number of cases  mathematical patterns found in string theory have led to discoveries of related patterns in field theory.} as reviewed by Tseytlin and by Gibbons \\cite{Tseytlin:1999dj,Gibbons:2001gy}.  A second approach to nonlinearity was introduced by Heisenberg and Euler \\cite{H-E}, who observed that what we now would call virtual creation of electron-positron pairs leads inevitably to an extra term in the Maxwell equations that is cubic in field strengths.  This work also has a living legacy, as reviewed recently by Dunne \\cite{Dunne:2004nc}.    \\subsection{Gravitons}  The scientific question in gravity most naturally related to photon mass is the issue of a possible graviton mass. (With the exception of Vainshtein's article \\cite{Vainshtein:2006ix} mentioned above, devoted mainly to theoretical aspects, there apparently has never been a widely-circulated review on this topic.) Its study progressed more slowly than the photon-mass issue, at least in part because gravity is so weak that even today classical gravitational waves have not been detected directly.  Further, gravitons, regardless of their mass, seem beyond the possibility of detection in the foreseeable future.  We proceed to discuss theoretical issues (Section \\ref{grav}) and observations (Section \\ref{gravmass}) for the case of classical gravity.   There are several important contrasts between the photon and graviton cases.  First, from a  theoretical point of view the possibility of nonzero graviton mass is open to question.  This makes what is  a relatively straightforward discussion for photons much more problematic for gravitons.  Secondly, there is a highly developed formalism for seeking to measure deviations of gravity from Einstein's General Theory of Relativity -- the parametrized post-Newtonian [PPN] expansion.  In this framework there have been many measurements, principally under weak-field conditions, both for low-velocity and high-velocity phenomena, to test for deviations.  As with photon mass, none of these measurements to date have produced ``unexpected-physics\" results, only increasingly stringent limits on departures from Einstein gravity.  However, there is another important distinction.  Two sets of characteristic phenomena show significant departures from Einstein gravity with the matter sources being only familiar ``visible\" matter -- stars, hot gases, and photons.  The first, indicated already by observations in the 1930s, and much more definitely in the 1970s, has been labeled  ``dark matter.\"  Trajectories of visible objects (including the most visible of all -- light itself) seem to be bent more than would be expected if the only sources for gravity were pieces of visible matter.   In principle, a possible explanation for this could be long-range modifications of Newtonian and Einsteinian gravity,  but of a type very different from what would be called graviton mass.  The second departure, discovered much more recently, is an accelerated expansion of the universe, possibly described by the presence of another sort of invisible source, ``dark energy.\"  We discuss these issues also in Section \\ref{grav}.    \\subsection{Perspectives on gauge and general-coordinate invariance}  This review, then, can be considered an evolution of our earlier review on the mass of the photon \\cite{GN}.  Here we discuss current understanding and ideas on the masses of the photon and graviton, in light of  developments in theoretical and experimental physics over the past decades.  Because in principle there is no end to the  types of deviation that could be contemplated, we need some restrictions.   Our chosen class of deviations will be ones  that explicitly obey abstract ``symmetries:\" gauge invariance in the case of electrodynamics, and general-coordinate invariance in the case of gravity.\\footnote{In the case of string theory, the corresponding property is called ``reparametrization invariance.\" }  The distinction between these invariances and familiar global symmetries such as rotation or translation invariance is that not only the action but in fact {\\it all} observables are invariant.   Thus the invariance is with respect to description rather than physical transformation. At first sight it might appear that making such an invariance explicit would be useless, because no matter what the dynamics this always should be possible.   However, history shows that the invariances can be most fruitful:   First of all, electrodynamics and Einstein gravity are examples of {\\it minimal} theories exhibiting such invariance, and hence are essentially unique.    Secondly, even for non-minimal versions with extra couplings, keeping the invariances explicit, while not constraining the physical content of the theory, can be most helpful in carrying through computations. They then play a  similar role to the ``check bit\" at each stage in a numerical computation, which also adds nothing to the content, but can guide calculations and flag errors.  It might seem that such invariance excludes a mass for the Proca photon or the graviton.  However, as we shall discuss, the Higgs mechanism can `hide' an invariance, which nevertheless remains unbroken.   For the photon that becomes equivalent in a certain limit to the fixed Proca mass, which therefore needn't break gauge invariance after all. ",
        "conclusions": "\\subsection{Conclusions}  The subject of possible photon or graviton rest mass is appealing because there are so many levels of beautiful argument for the masses to vanish  (and of course a counter-argument for each argument).  Both of these  examples (of the only long-range fields we know) reveal important strands of physics relevant to many different areas.  Taken as a whole, their study illuminates the history, logic, and remarkably complex yet coherent structure of physics.  In Table \\ref{masslist} we give a list of what we find to be the most significant and/or interesting mass limits so far proposed.     \\begin{table*}[t!] \\begin{center} \\caption{A list of the most significant mass limits of various types for the photon and graviton. \\label{masslist}} \\vskip 20pt \\begin{tabular}{rl|l|l|l|l} \\hline\\hline & Description of method ~~~~~~~~~~~~~~ & $\\lc \\gtrsim$ (m) ~~~~~~ & $\\mu \\lesssim$  (eV) ~~~~~~ & $\\mu \\lesssim$ (kg) ~~~~ & Comments  ~~~~~~~~~~~~~~~~~~~~~~    \\\\   \\hline &                 &             &         &&\\\\ 1 & {\\sf Secure photon mass limits:}  && &&\\\\ &  Dispersion in the ionosphere \\cite{Kroll:1971wi} & $8\\times 10^5$ & $3 \\times 10^{-13}$ & $ 4\\times 10^{-49}$  & \\\\ & Coulomb's law  \\cite{wfh} & $2 \\times 10^7$ &  $10^{-14}$  &  $2 \\times 10^{-50}$  &   \\\\ & Jupiter's magnetic field \\cite{DGN}   & $5 \\times 10^8$ & $4 \\times 10^{-16}$    & $7 \\times 10^{-52}$ &  \\\\ & Solar wind magnetic field \\cite{ryutov2} &  $2 \\times 10^{11}$~~~(1.3 AU) & $10^{-18}$    & $2 \\times 10^{-54}$  & \\\\ [10pt] 2 & {\\sf Speculative photon mass limits:} &&&&  \\\\ & Extended Lakes method \\cite{lakes,luotorsion,gntorsion} & $3\\times10^9$  & $7\\times 10^{-17}$  & $10^{-52}$\\ &$\\lc\\sim$ 4 $R_\\odot$ to 20 AU,     depending\\\\ &  &  $~~~~\\Leftrightarrow 3\\times 10^{12}$ & $~~~~\\Leftrightarrow 7 \\times 10^{-20}$ & $~~~~\\Leftrightarrow 10^{-55}$ & ~~~ on $\\mathbf{B}$ speculations \\\\ & Higgs mass for photon \\cite{adel}   &  No limit feasible  ~~ & & & Strong constraints on 3D Higgs \\\\ &&&&& ~~~parameter space\\\\ & Cosmic mag. fields \\cite{yamaguchi,chib}, \\cite{adel} & $3\\times 10^{19}$  ~~ (${10^3}$ pc) & $6\\times 10^{-27}$ & $ 10^{-62}$ & Needs const. $\\mathbf{B}$ in galaxy regions \\\\[10pt] 3 & {\\sf Graviton mass limits:}          &&          && \\\\ & Grav. wave dispersion \\cite{finn} & $3 \\times 10^{12}$ & $8 \\times 10^{-20}$ & $10^{-55}$  & Question mark for scalar graviton \\\\ & Pulsar timing \\cite{baskaran} & $2 \\times 10^{16}$ & $9 \\times 10^{-24}$   & $2 \\times 10^{-59}$ &  Fluctuations due to graviton \\\\ &&&&& ~~~phase velocity \\\\ & Gravity over cluster sizes  \\cite{GNgraviton} & $2\\times 10^{22}$ & $10^{-29}$ & $2 \\times 10^{-65}$ & \\\\ & Near field constraints  \\cite{gruz} & $3 \\times 10^{24} ~~(10^8$ pc) & $6 \\times 10^{-32}$ & $10^{-67}$ & For DGP model \\\\ &Far field constraints  \\cite{Dvali:2002vf} & $3 \\times 10^{26} ~~(10^{10}$ pc) & $6 \\times 10^{-34}$ & $10^{-69}$ &For DGP model \\\\[10pt]  \\hline\\hline \\end{tabular} \\end{center} \\end{table*}   The first reason for vanishing mass is, of course,  something Newton adopted instinctively for gravity and Gauss justified in a beautiful way for electrostatics:  the inverse square law of force. For gravity this was confirmed with tremendous precision  by the match with Kepler's laws.  For electrostatics Gauss's statement that the number of lines of force coming out of a charge is, at any distance, a direct measure  of that  charge, gave a powerful geometric interpretation to the force law.  (This argument applies equally to mass in gravity.) Later there came the notions of gauge invariance, discovered in the mathematical structure of classical electrodynamics, and general-coordinate invariance, invented by Einstein to constrain the possible structure of his emerging theory of gravity.  There is also a ``backwards\" connection:  As Ogievetsky and Polubarinov \\cite{ogiev1,ogiev2} and Weinberg \\cite{Weinberg:1964ew} showed, zero mass for photon or graviton implies local conservation of electric charge in electrodynamics or energy and momentum in gravity.  Like the sizes of the masses, violations of these conservation laws are strongly constrained by experiment.  Thus, for electrodynamics as well as gravity, two effects known to be small are logically related in the limit where they both are zero.  Simplicity also favors these zeros.   They represent the minimal structure consistent with all symmetry requirements.  Anything different requires more parameters if not more fields.  Despite all these arguments, there is another side to the story.   As Stueckelberg showed, gauge invariance can be satisfied at least formally even in the presence of a mass (and Siegel showed that the same holds for general-coordinate invariance).   Perhaps even more powerful is the example of the W$^\\pm$ and Z$^0$ mesons, which  clearly are gauge particles, and yet have mass.  At least for the photon case, there is an objection to this argument.   In the context of a Higgs mechanism one must introduce an electrically charged scalar field, where constraints from observation imply that the value of the charge is an extraordinarily tiny fraction of the charge of an electron.   Such a charge, of course, would violate the pattern of all known charges, and also would contradict an appealing (though unproved) idea, that of grand unified theory.  It would be a bizarre modification of electroweak theory, where the compact group SU(2)$_{\\rm left}$ automatically leads to quantized left-handed charges.  In other words, `mini-charged' Higgs particles, if they existed, would perforce be coupled only to the U(1)$_{\\rm right}$.  Because weak interactions are short-ranged, and the W and Z bosons are so massive, the effects of the weak charge of the new Higgs particle would be even more insignificant than those of the electric charge.  The effect of the new Higgs coupling on the masses of the W and Z also would be unobservably small.\\footnote{ If instead of a Higgs particle, the associated Higgs field were a composite of other fields, then these too would have extraordinarily small, unquantized U(1)$_{\\rm right}$ charges, or at least super-small offsets of the charges for different elements of the composite. }  As far as observation relevant to photon mass goes, the only debate is about how stringent a limit currently can be placed on that mass.   To date no evidence at all has appeared for a nonzero value.  Even with the generalization to a Higgs-mechanism framework, the ``obvious\" experiment of seeking to detect dispersion of velocity   with frequency is guaranteed to give no useful information, because limits from static magnetic fields are so low that nothing could be detected by dispersion measurements (at least in regions identified so far where we could measure the velocity).   Even the Schumann resonances do not give as strong a constraint as the magnetostatic limit. Thus, the only even potentially observable effect of photon mass would be found in the photon Compton wavelength, and that already is known to be at least comparable in dimension to the Earth-Sun distance.   Therefore, any future improvements in the limit probably will come from astronomical observation rather than laboratory (even satellite-laboratory) experiment.     \\subsection{Prospects}  Let us accept  that the primary tool for limiting or detecting a rest mass of the photon or graviton is to exploit the ``Yukawa\" effect:  modification at long distances of essentially static electromagnetic or gravitational fields.  Then  possibilities for extending the range in the case of electromagnetism look very good.  With evolving instruments and techniques, we are in an era of rapid expansion in the depth of exploration of the universe. Detailed knowledge of galactic and extragalactic magnetic fields is  accumulating along with knowledge about the structure of the associated plasmas.   There is every reason to suppose that a lower limit on the photon Compton wavelength limit of galactic or even larger dimensions could  be attainable.  If a finite value for $\\lc$ were detected, almost certainly it would be so large and the corresponding photon mass so small that even in the Higgs framework the corresponding electric ``mini-charges\" would be too small to detect.  Thus the continuity of electrodynamics in the zero-mass limit (unless electric charge is not locally conserved) already assures that the {\\it only} mass effect still possible to observe would be long-distance modifications of static magnetic fields.  In other words, for all lab-scale purposes the mass already may be taken as zero.   Still, if a non-zero value were established by new astronomical observations, this small departure would have enormous conceptual implications, giving incentive for searching examination of the accepted foundations of electrodynamics.  At the same time, the great debate about the possibility of a massive graviton in GR seems pretty much complete.  The most conservative lower bound on the graviton Compton wavelength, based on limits to deviations from Einstein-Newton gravity in the solar system, puts it at $\\approx$ 1\\% of the radius [R] of the visible universe \\cite{gruz}.  Even that estimate is for an asymptotically flat space time, whereas the actual universe is expanding and even accelerating in its expansion.\\footnote{ So, if spatially flat, the universe still has (negative) curvature in the relation between time and space variables.  This curvature introduces a ghost in the DGP theory, making it suspect, and therefore making the limit \\cite{gruz} possibly too conservative. } More to the point, a $\\lc$ significantly smaller than R would produce dramatic (and certainly not seen) effects on the large-scale picture of the universe.  On the issue of dark matter versus modified gravity, we have seen that the most conservative way of accounting for phenomena, the existence of one new weakly interacting particle, has not so far been shown to work for galactic scales, where MOND gives a very successful parametrization of the data.  In our opinion the biggest current challenge for the CDM hypothesis is to remedy this lack, perhaps by finding equilibria involving dark matter and ordinary matter.  If these equilibria implied metastable configurations explaining  flat rotation curves and the Tully-Fisher law, one would have a satisfying confirmation of CDM. As long as that hasn't happened,  it remains a viable possibility that explaining the MOND phenomenology requires new physics other than CDM.  Classical electrodynamics and general relativity were the first two field theories of physics.   They appeared  in complete form at the hands of Maxwell and of Einstein, and even today there is no proof that they need to be changed.   Still, remembering that physics is an experimental science, we should keep  watch for  surprises.  It could turn out that either or both of these theories actually require adjustment, perhaps along the lines considered in this paper, or perhaps in some completely new direction."
    },
    "0809/0809.1444.txt": {
        "abstract": " ",
        "introduction": "Commission 10 deals with solar activity in all of its forms, ranging from the smallest nanoflares to the largest coronal mass ejections. This report reviews scientific progress over the roughly two-year period ending in the middle of 2008.  This has been an exciting time in solar physics, highlighted by the launches of the Hinode and STEREO missions late in 2006. The report is reasonably comprehensive, though it is far from exhaustive. Limited space prevents the inclusion of many significant results. The report is divided into the following sections: Photosphere and Chromosphere; Transition Region; Corona and Coronal Heating; Coronal Jets; Flares; Coronal Mass Ejection Initiation; Global Coronal Waves and Shocks; Coronal Dimming; The Link Between Low Coronal CME Signatures and Magnetic Clouds; Coronal Mass Ejections in the Heliosphere; and Coronal Mass Ejections and Space Weather.  Primary authorship is indicated at the beginning of each section. ",
        "conclusions": ""
    },
    "0809/0809.3965_arXiv.txt": {
        "abstract": "The $\\rho$~Oph molecular cloud is undergoing intermediate-mass star formation. UV radiation from its hottest young stars heats and dissociates exposed layers, but does not ionize hydrogen. Only faint radiation from the Rayleigh-Jeans tail of $\\sim$~10--100~K dust is expected at wavelengths longwards of $\\sim$3~mm. Yet Cosmic Background Imager (CBI) observations reveal that the $\\rho$~Oph~W photo-dissociation region (PDR) is surprisingly bright at centimetre wavelengths. We searched for interpretations consistent with the {\\em WMAP} radio spectrum, new {\\em ISO}-LWS parallel mode images and archival {\\em Spitzer} data.  Dust-related emission mechanisms at 1~cm, as proposed by Draine \\& Lazarian, are a possibility. But a magnetic enhancement of the grain opacity at 1~cm is inconsistent with the morphology of the dust column maps $N_d$ and the lack of detected polarization. Spinning dust, or electric-dipole radiation from spinning very small grains (VSGs), comfortably explains the radio spectrum, although not the conspicuous absence from the CBI data of the infrared circumstellar nebulae around the B-type stars S~1 and SR~3. Allowing for VSG depletion can marginally reconcile spinning dust with the data. As an alternative interpretation we consider the continuum from residual charges in $\\rho$~Oph~W, where most of carbon should be photoionised by the close binary HD147889 (B2IV, B3IV). Electron densities of $\\sim 10^2$~cm$^{-3}$, or H-nucleus densities $n_\\mathrm{H}>$~10$^6$~cm$^{-3}$, are required to interpret $\\rho$~Oph~W as the C\\,{\\sc ii} Str\\\"omgren sphere of HD147889. However the observed steep and positive low-frequency spectral index would then require optically thick emission from an hitherto unobserved ensemble of dense clumps or sheets with a filling factor $\\sim 10^{-4}$ and $n_\\mathrm{H} \\sim 10^7$~cm$^{-3}$. ",
        "introduction": "The subtraction of Galactic foregrounds in experiments designed to map the cosmic microwave background requires the examination of the emission mechanisms at work in the interstellar medium (ISM). An anomalous component of continuum emission was discovered in the direction of Galactic cirrus clouds \\citep{lei97}, so defined by their H\\,{\\sc i}~21~cm and far-IR contours \\citep[e.g.][]{bou88}.  The LDN~1622 dark cloud \\citep[Lynds Dark Nebula,][]{lyn62} was found to be bright at cm-wavelengths, where no known emission mechanisms were expected \\citep[][]{fin02, fin04, cas06}. At the time of writing LDN~1622 is the only dark cloud known to radiate at cm-wavelengths.   What is the nature of the cm-wave emitters? Do the dark clouds and the cirrus clouds radiate by the same emission mechanisms? \\citet{dl98a,dl98b} proposed electric dipole radiation from polarized very small dust grains (VSGs) spinning at GHz frequencies, or spinning dust. \\citet{dl99} also suggested that `magnetic dust', or magnetic dipole emission due to thermal fluctuations in the magnetization of ferromagnetic grains, could produce detectable cm-wave emission. However, a finite charge density exists in atomic and molecular clouds; a small part of the neutral material is ionised by exposure to pervasive cosmic rays or soft-UV photons \\citep[e.g.][]{tie05}. As an alternative to spinning dust, could the residual charges radiate at the observed levels?    Prototypical and well studied local clouds can give information on the environments giving rise to cm-wave radiation.  The $\\rho$~Oph molecular cloud \\citep[e.g.][]{enc74,you06}, at a distance $D = 135 \\pm$15~pc \\citep[parallax distance to HD147889,][]{hab03}, lies in the Gould Belt of the closest molecular complexes.  $\\rho$~Oph is undergoing intermediate-mass star formation - the most massive of its young stars is HD147889, which we show here to be a close pre-main-sequence B2, B3 binary, not hot enough to form a conspicuous region of ionized-hydrogen (H\\,{\\sc ii} region). A description of the region can be found in Fig.~4 of \\citet{you06}.   Here we present the first resolved images at cm-wavelengths of the $\\rho$~Oph main cloud, LDN~1688. We describe our observations in Section~\\ref{sec:cbiobs}, as well as auxiliary data in Sec.~\\ref{sec:aux}, including unpublished {\\em ISO}-LWS parallel mode data.  We proceed to summarise the available imaging and spectroscopic data in Sec.~\\ref{sec:obs}. The dust emission from $\\rho$~Oph~W and the contribution of spinning or magnetic dust at cm-wavelengths are studied critically in Sec.~\\ref{sec:dust}. We also propose in Sec.~\\ref{sec:CI} an alternative emission mechanism for the 31~GHz emission based on the C\\,{\\sc i} continuum from a cold plasma.  In Sec.~\\ref{sec:RRLs} we analyse the non-detection of the radio recombination line system in $\\rho$~Oph~W. We discuss our findings in Sec.~\\ref{sec:disc} in terms of the C\\,{\\sc ii} Str\\\"omgren spheres around the early type stars that are interacting with the $\\rho$~Oph cloud, before concluding in Sec.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}   We have found that the well-studied and nearby molecular cloud $\\rho$~Oph is surprisingly bright at 31~GHz, or $\\sim$1~cm wavelengths. The most conspicuous feature revealed by the CBI data is the $\\rho$~Oph~W PDR.    Comparison with {\\em~WMAP} images shows that the cm-emission is not the Rayleigh-Jeans tail of the sub-mm emitting dust. Bulk dust properties are inferred from new {\\em ISO}-LWS parallel mode data. Archival {\\em Spitzer} data allows quantifying PAH emission. None of the comparison images match the CBI data, except for IRAC~8~$\\mu$m and 2MASS~K$_s$, which closely follow the 31~GHz ridge along $\\rho$~Oph~W. {\\em Spitzer}~IRS spectroscopy in four apertures hint at a possible correlation between 31~GHz intensity and the H$_2$ pure-rotational lines: the H$_2$ lines are detected only where 31~GHz intensities are bright (i.e. towards SR~3 and $\\rho$~Oph~W, and not in the vicinity of S~1).    We considered several interpretations for the 31~GHz emision, requiring either an additional radio continuum component such as magnetic or spinning dust, or unexpected physical conditions, such as in a very dense and cold plasma. We find that both spinning dust and C\\,{\\sc i} continuum can explain the data:  \\begin{itemize}  \\item Magnetic dust is discarded on morphological grounds. The expected 31~GHz intensity from a magnetic enhancement of the grain opacity is at odds with the CBI data. The polarization levels required by magnetic dust are not observed.      \\item Spinning dust can account for the radio spectrum. But the predicted levels in S~1 are in excess by a factor $>$40 at 3~$\\sigma$.  We take the intensities in PAH~11.3~$\\mu$m as a proxy for the mid-IR VSG emission, which is approximately proportional to both the VSG column and local UV intensity $U$.  Taking into account the range of spinning dust emissivities in all possible environments, the uncertainties in our 31~GHz map can marginally reconcile spinning dust with the data (i.e. at 3~$\\sigma$ variations).    \\item Alternatively a cold plasma such as that found in C\\,{\\sc ii} regions could explain the 31~GHz emission. The star HD147889, a binary pre-main-sequence star with spectral types B2IV, B3IV, emits sufficient C-ionising UV radiation to interpret $\\rho$~Oph~W as a C\\,{\\sc ii} region.  But the absence of detectable $<10~$GHz signal requires optically thick emission, and hence high H-nucleus densities ($n_\\mathrm{H} \\sim 10^7$~cm$^{-3}$, almost a factor of 10 higher than inferred in the literature), implying a hitherto unobserved molecular phase of the ISM. In $\\rho$~Oph this molecular phase takes the form of an ensemble of dense clumps or sheets at temperatures of order 50~K or less.  The cold plasma interpretation explains the 2MASS - CBI correlation through an important fluorescent H$_2$ contribution to the 2MASS bands.  \\end{itemize}"
    },
    "0809/0809.3012_arXiv.txt": {
        "abstract": "The Near-Infrared Coronagraphic Imager (NICI) is a high-contrast AO imager at the Gemini South telescope. The camera includes a coronagraphic mask and dual channel imaging for Spectral Differential Imaging (SDI). The instrument can also be used in a fixed Cassegrain Rotator mode for Angular Differential Imaging (ADI). While coronagraphy, SDI, and ADI have been applied  before in direct imaging searches for exoplanets. NICI represents the first time that these 3 techniques can be combined. We present preliminary NICI commissioning data using these techniques and show that combining SDI and ADI results in significant gains. ",
        "introduction": "\\label{intro} The Near-Infrared Coronagraphic Imager (NICI) is a new instrument at Gemini-South Telescope\\cite{Ftaclas2003}. NICI had its first light in February 2007 and is currently undergoing on-sky tests before embarking on an ambitious 50-night planet-search campaign. The campaign will be executed in queue mode and should significantly improve the statistical constraints on the presence of giant planets around nearby main-sequence stars. The campaign builds on similar surveys by Refs.~\\cite{Biller2007, Lafreniere2007gdps, Lowrance2005}.  NICI combines a curvature AO system and a dual-channel imager. The two optical trains are designed to operate in a Spectral Differential Imaging (SDI) mode. In this mode the field is imaged in two close wavelengths selected such that any cold companion would have widely different flux in the two channels while the stellar spectral energy distribution remains relatively flat. This allows for a calibration of the stellar speckle pattern using the {\\it off} channel with minimal self-subtraction of the companion. All previous on-sky applications of this mode \\cite{Marois2005, Biller2007} as well as the nominal NICI mode use the deep methane absorption band around $1.6~\\mu$m, but the technique could conceivably be used on other spectroscopic features. SDI is limitations by the decorrelation of speckles between the channels due to non-common path aberrations, differences in the filter bandpasses, wavelength difference between filters, etc. Once other noise contributions (readout noise, uncorrected atmospheric speckles, photon noise...) have averaged-out, no further gain in sensitivity can be obtained without resorting to other techniques as further integrations only provide higher signal-to-noise ration (S/N) on the fixed speckle pattern and doesn't help detect a planet buried within it.  The second technique envisioned for the NICI campaign and general high-contrast imaging is roll subtraction, a.k.a. Angular Differential Imaging (ADI) \\cite{Liu2004, Marois2006}. The technique is similar to the roll subtraction used on HST \\cite{Schneider2003} and takes advantage of the field rotation provided by an altitude-azimuth mount such as that of the Gemini South telescope. In this mode the Cassegrain focus or rotator is fixed, so the field rotates through the observation leaving all the optics (telescope, AO system, dual-imager) fixed. In this mode, field rotation provides a self-calibration of the PSF. For every image, an estimate of the PSF can be built using the rest of the dataset and subtracted from the frame of interest. Once this is done for all frames, they can be rotated to a common position angle (PA) before combination. As pointed out by Ref.~\\cite{Marois2006}, any remaining speckle will add-up incoherently in the final image thus ensuring an ever increasing sensitivity.  As the field also rotates during integrations, there is a greater amount of image elongation at larger radii, and this field position-dependent loss in sensitivity needs to be taken into account when designing a survey strategy (see Biller et al. in this Volume). As a general rule, in the relatively short integrations used with NICI (typically a few seconds to a minute), the angular smear is small, especially in the scientifically most relevant part of the field (the inner $\\sim2\\farcs0$). ADI is most efficient in the outer parts of the PSF where relatively fast field rotation allows one to use images taken within minutes of the frame of interest to built the reference PSF. In the inner part of the field (less than $\\sim1\\farcs0$), the slower field rotation implies that only images taken tens of minutes from the frame of interest can be used to reconstruct the PSF estimate. The ever-evolving instrumental speckle pattern therefore implies that the ADI-reconstructed PSF will be better correlated at large separation than closer in. In some particularly unfavorable cases (i.e. objects low on the horizons or transiting directly overhead), ADI subtraction at small separations cannot be performed at all within a $1-2$ hours long observation. It is important to noted that ADI and SDI are {\\it not} competing techniques and that they may very well be used on the same dataset.  The third important feature of NICI is the presence of a semi-transparent focal plane mask. A wheel allows the positioning of a set of masks ranging in diameter from $0\\farcs22$ to $0\\farcs90$ as well a clear glass element of the same optical thickness. The central part of masks have a typical transmission of 1 part in 300 (See Ftaclas et al. in this Volume). While the Strehl ratio obtained by NICI ($\\sim0.3$ in $H$, $\\sim0.5$ in $K$) are not sufficient to perform diffraction coronagraphy, the focal plane masks still help improve sensitivity by 1) limiting the amount of light entering the system and thus minimizing ghosts, 2) avoiding the deep saturation of the central parts of the PSF that would create electronic ghosts and leakage and 3) providing a non-saturated diffraction core as both an astrometric and photometric reference for registering images.  This paper describes the broad data reduction of a preliminary NICI dataset obtained during commissioning and provides sensitivity measurements for important steps to quantify the relative gains. Observations are described in \\S~\\ref{observations}, \\S~\\ref{reduction} describes the various data reduction steps and \\S~\\ref{discussion} details their relative sensitivity gains. Figure~\\ref{beforeafter} shows an example of a raw NICI frame and the final image obtained with the data reduction method described here.  \\begin{figure} \\begin{center} \\begin{tabular}{cc} \\includegraphics[height=5.5cm]{raw.eps} & \\includegraphics[height=5.5cm]{reduc.eps} \\\\ (a) & (b) \\end{tabular} \\end{center} \\caption[example] {\\label{beforeafter} Left panel shows a single NICI image of TWA-7 after dark-subtraction and flat-fielding. The image inside a $2\\farcs0$ radius has been divided by 30 to account for the large dynamic range and show speckles in the inner parts of the point spread function. The right panel shows the same field after the complete data reduction, for a total integration time of 31 minutes. For both panels, the scale ranges from $-0.2$ to $1\\times$ the peak of the background object $\\sim3\\farcs1$ from TWA-7. This object is $\\Delta H=9.4$ mag fainter than TWA-7. A dashed circle indicates the edge of the $0\\farcs32$-wide occulting mask in both panels. } \\end{figure} ",
        "conclusions": "\\label{discussion} The median-combined and PA-aligned images at various steps of the data reduction have been produced in order to quantify the relative gains of the various steps. Figure~\\ref{contrast} illustrate both the resulting images and the $5 \\sigma$ detection limits at the corresponding steps. Sensitivity curves shown in Figure~\\ref{contrast} have been derived using a circular aperture with a 3.0~pixels radius and a sky annulus ranging from 6 to 12~pixels. Photometric calibration has been done using the background object mentioned earlier. The derived sensitivity curves is only weakly dependant on the aperture used, and other techniques such as convolving the images by a Gaussian with a 5~pixels FWHM and deriving a sensitivity curve using the RMS of the residual in annuli give similar results.    \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=10.5cm]{speckle_combo.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{combo} Median combination of the complete, PA-registered, dataset at various data reduction steps. For each image the corresponding sensitivity curve in Figure~\\ref{contrast} is given. Image~\\#1 shows the combined image after dark-subtract, flat-fielding and median radial-profile subtraction (solid line); this corresponds to the data reduction step illustrated by Figure~\\ref{halo}. The strong asymmetry seen in the radial-profile subtracted image is readily seen in the image before subtraction and is not an artifact due to a misregistering of the radial profile (which would give a similar signature). Image~\\#2 is the combination of the dataset after performing the high-passed filtering (dotted line). Image~\\#3 shows the combined {\\it short}-channel LOCI image (dashed line). Image~\\#4 shows the final two-channel-LOCI image (dash-dot line). All images have a linear scaling ranging from $-0.5$ to $+0.5\\times$ the peak value of the $\\Delta H\\sim9.4$ background object $3\\farcs1$ from TWA-7.} \\end{figure}  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=10.5cm]{contrast_comparison.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{contrast} Upper panel : The $5 \\sigma$ contrast limits for the TWA-7 ($V=11.1$, $H=7.1$) dataset, totalling 31 minutes of integration, obtained after: a simple registering of flat-fielded, dark-subtracted images (solid curve); a combination of the radial-profile and high-pass subtracted images (dotted curve); a {\\it short}-channel LOCI (dashed line); and the complete two-channel LOCI (dashed-dot line). The line at a $\\Delta$~mag$=12.34$ ($H=19.4$) indicates the limit set by readout noise. This limit could be further lowered by $\\sim$$1~$mag by using 16 or 32 NDRs instead of the 4 NDRs used in the dataset presented here. Lower panel~: difference between the  {\\it short}-channel LOCI and two-channel LOCI curves. This difference gives a measure of the sensitivity obtained through spectral differential imaging.} \\end{figure}  An important measure of the performance of NICI is the sensitivity gain through SDI. Implementing the dual channel imager has a financial, technical and sensitivity cost, part of the optical train has to be duplicated, in the design adopted here two detectors were used, $50\\%$ of the light is lost at the beam splitter compared to a simple imager, splitting the light in two channels implies that the effective readout noise contribution is doubled relative to the signal, etc. Of course, this price is paid to obtained as to obtain better calibration of the speckle pattern and, at least for the most scientifically relevant part of the field, is providing a gain in sensitivity.  The difference between the {\\it short}-channel LOCI and the dual-channel LOCI provides an estimate of the sensitivity gain in having two channels (see lower panel of Figure~\\ref{contrast}). The final {\\it short}-channel LOCI is the deepest image one could achieve using a single channel image. In a survey using two epochs to identify faint companions (such as the GDPS by Ref.~\\cite{Lafreniere2007gdps}), this would define the sensitivity for a given epoch. One could argue that a single-channel imager would go $\\sim0.7$~mag deeper within the same period of time thanks to the $50\\%$ gain in throughput. This is certainly the case in the read-out noise limited part of the image beyond $\\sim1\\farcs2$, but the signal-to-noise in the inner part of the PSF is {\\it not} limited by flux level, but by residual speckle noise that the LOCI algorithm failed to subtract. This can be seen in Figure~\\ref{contrast}, the noise in the central part of the {\\it short}-channel LOCI image is dominated by FWHM-sized structures, not pixel-to-pixel fluctuations. Gains in sensitivity for this region of the PSF can only be obtained by observing a larger number of individual realizations of the speckle pattern, {\\it not} by increasing the per-image flux. At the edge of the occulting mask, there is a $\\sim1.5$~mag sensitivity gain in including the second channel images, and this gain drops to about $0.2$~mag at a $1\\farcs0$ radii. This rapid drop is easily explained by the fact that the efficiency of angular differential imaging increases with radii."
    },
    "0809/0809.2478_arXiv.txt": {
        "abstract": "We present dynamical simulations of NGC 2523 and NGC 4245, two barred galaxies (types SB(r)b and SB(r)0/a, respectively) with prominent inner rings. Our goal is to estimate the bar pattern speeds in these galaxies by matching a sticky-particle simulation to the $B$-band morphology, using near-infrared $K_s$-band images to define the gravitational potentials. We compare the pattern speeds derived by this method with those derived in our previous paper using the well-known Tremaine-Weinberg continuity equation method. The inner rings in these galaxies, which are likely to be resonance features, help to constrain the dynamical models. We find that both methods give the same pattern speeds within the errors. ",
        "introduction": "The pattern speed, $\\Omega_p$, is one of the most influential and elusive parameters affecting the overall morphology of a galaxy. It remains one of the fundamental unknowns of galaxy dynamics. A variety of techniques for determining $\\Omega_p$ have been published (e.g. Canzian 1993; Buta \\& Combes 1996; Puerari \\& Dottori 1997; Salo et al. 1999; Weiner et al. 2001; Egusa et al. 2004; Zhang \\& Buta 2007), though in recent years, emphasis has been placed on an observational method developed by Tremaine \\& Weinberg (1984, hereafter TW). Studies of late that make use of this method have focussed on estimating the pattern speed of SB0 galaxies (e.g. Aguerri et al. 2003; Corsini et al. 2003; Debattista et al. 2002). These studies have indicated that SB0 galaxies generally have \"fast\" bars, presumably indicating halos of low central concentration.  The TW method is the most direct means of measuring the pattern speed of a galaxy. Here, the pattern speed of a bar can be estimated from the luminosity weighted mean line-of-sight velocities, $\\langle$$V$$\\rangle$, and luminosity weighted mean positions, $\\langle$$X$$\\rangle$, (both relative to the central values) of a tracer that obeys the continuity equation. These quantities are to be measured along lines parallel to a barred galaxy's projected disk major axis and are related as follows, \\begin{equation}\\Omega_p \\sin{i} = \\frac{\\langle V \\rangle}{\\langle X \\rangle}\\end{equation} where $i$ is the galaxy's inclination.  One of the great advantages of the TW method is that it provides pattern speeds independent of any assumptions concerning features traced by the pattern speed. For example, the TW method can used to directly test the ideas that bars end near their own corotation resonances, and that rings are associated with specific resonances. However, this advantage is somewhat offset by the limited scope of applicability of the method. The most significant hindrance is that the bar major axis must lie at approximately 45$^{\\circ}$ to the major axis of the projected disk in order to obtain the maximum TW signature strength, with angles between about 20$^{\\circ}$ and 70$^{\\circ}$ considered suitable (Gerssen et al. 2003). The galaxy itself must have a preferred inclination between 50$^{\\circ}$ and 60$^{\\circ}$ (Debattista 2003) in order to avoid large uncertainties. Another problem is identifying a pattern speed tracer with a strong spectral signature that obeys the continuity equation. This means that the tracer must be something that is neither created nor destroyed and significant star formation would violate the equation. This is the reason that applications to spirals have been more limited. Hernandez et al. (2005) discuss the application of the TW method to atomic, molecular, and ionized gas phases in spiral galaxies. SB0 galaxies were the first objects to which the method was applied, using stellar absorption lines (e.g., Merrifield \\& Kuijken 1995; Gerssen et al. 1999). These galaxies can have strong bar patterns but lack the dust and star formation that would complicate similar measurements for later-type systems. Recently though, a TW test of an n-body model by Gerssen \\& Debattista (2007) indicated that dust absorption effects may be partially mitigated by star formation on the leading edges of a bar, creating a minimal impact on the derived $\\Omega_p$. Avoiding these potentially problematic factors significantly diminishes the number of galaxies to which this method can be applied.  A less direct, yet more widely applicable method of determining $\\Omega_p$ in galaxies is through simulation modeling (e.g., Salo et al. 1999; Rautiainen, Salo, \\& Laurikainen 2005). This involves modeling the response of gas in a rigid galaxy potential, where the gas is modeled in terms of dissipatively colliding, or \"sticky,\" particles. These sticky particles simulate the cold gas component of a galaxy, where regions of star formation occur due to cloud-cloud collisions (Levinson \\& Roberts 1981). The gravitational potential is derived from a near-IR image. Different dynamical parameters are varied until the morphology and kinematics of the numerically simulated galaxy visually matches, as closely as possible, the blue light morphology. The morphology of resonance rings is sensitive to the underlying pattern speed of a galaxy and serve as a good $\\Omega_p$ estimator. We have deep $B$-band images with good resolution in order to compare the observed low velocity dispersion components to our models. Since resonance rings generally form in the gas component and show emission from young stars, the observed $B$-band morphology is best suited for this purpose. Sticky particle modeling is similar to hydrodynamic modeling, such as SPH (smoothed particle hydrodynamics), only the particles are not subjected to non-gravitational forces. Therefore, the motions of particles are not affected by the positions of other particles (i.e through pressure gradients), except through dissipation in physical impacts. This difference means that SPH is better suited for modeling the warm gas phase of a galaxy, while the sticky particle method is better suited for approximating the cold gas phase (Merlin \\& Chiosi 2007). In Salo et al. (1999), which addresses IC 4214, it was estimated that sticky particle modeling can yield $\\Omega_p$ within about 10-20\\%, even when allowing for uncertainties in the bar amplitude (due to unknown vertical extent, unknown dark halo contribution, and possible variations in the M/L ratio) and the numerical parameters involved in the sticky particle method itself. Direct comparisons between pattern speeds derived from sticky particle and SPH methods of modeling are not published, although a comparison of results from both methods is discussed by Bournaud et al. (2005) in regards to the gravity torques between a stellar bar and gas. It was found that the tree-SPH code by Semelin \\& Combes (2002) produced the same order value as in sticky particle simulations and observations.   In this paper, we compare the kinematic bar pattern speeds of NGC 2523 and NGC 4245, measured through the use of the TW method (Treuthardt et al. 2007), to the dynamical bar pattern speeds derived through simulations. In particular we base our simulation estimates solely on morphological comparisons and check whether $\\Omega_p$ is consistent with direct measurements. If so, this would justify the application of the simulation method to a large number of galaxies observed in various near-IR surveys, not necessarily possessing such extensive kinematical data as IC 4214. NGC 2523 is an SB(r)b galaxy with a closed inner ring and multiple arms (see Figure 1). A color index map (Buta, Corwin, \\& Odewahn 2007) shows that the spiral arms and part of the inner ring are sites of recent star formation. NGC 4245 is an SB(r)0/a galaxy with strong nuclear and inner resonance ring features that are also prominent blue features in a Buta, Corwin, \\& Odewahn (2007) color index map. NGC 4245 is an excellent candidate for such a test due to the good statistical fit of the measured pattern speed. The resonance features apparent in these galaxies help to constrain the dynamical models used to determine the bar pattern speed because the sizes and shapes of the rings are sensitive to the pattern speed. Since gas velocity maps of these galaxies do not exist, we are not able to compare the modeled gas kinematics to observations. ",
        "conclusions": "Using near-IR images to trace the stellar disk potential, we have estimated the pattern speeds and resonance locations of NGC 2523 and NGC 4245 using the numerical simulation method of Salo et al. (1999; also Rautiainen, Salo, \\& Laurikainen 2005). This is completely independent of the approach used by Treuthardt et al. (2007). We have examined the effects of varying the bar pattern speed (or $R_{CR}/R_{B}$), number of bar rotation periods, number of simulation steps per orbit, particle cross section, halo contribution, and bar amplitude on the simulated galaxy morphology. Our best models were chosen based on the closest match between the observed $B$-band and modeled morphology. In particular, we compared the size of the inner ring and the morphology outside of the ring. Our best simulation model provides an interpretation of NGC 2523 which places corotation at 1.4 $\\pm$ 0.1 times the estimated bar radius of 33.5$\\arcsec$, consistent with the results of Treuthardt et al. (2007). For NGC 4245, a similar model gives corotation at 1.5 $\\pm$ 0.1 times the estimated bar radius of 38.1$^{\\arcsec}$, which marginally agrees with the results of Treuthardt et al. (2007), within the errors. The values we derived through simulation modeling are robust against the various parameters examined.  Dynamical simulations are of great value as a means of estimating pattern speeds because they are more widely applicable than the TW method. Galaxies where the TW method is not appropriate, due to unfavorable inclinations or bar-disk position angles, can still be simulated to determine $\\Omega_p$. Thus, as further tests like ours are made, reliable pattern speeds will become available for larger numbers of galaxies.  P. Treuthardt and R. Buta acknowledge the support of NSF Grants AST-0205143 and AST-0507140 to the University of Alabama. H. Salo acknowledges the support of the Academy of Finland. We thank E. Laurikainen and J. H. Knapen for the use of the NIRS0S $K_s$-band images that we used for our simulations.  \\clearpage"
    },
    "0809/0809.0361_arXiv.txt": {
        "abstract": "{Long-lived, large-scale magnetic field configurations exist in upper main sequence, white dwarf, and neutron stars. Externally, these fields have a strong dipolar component, while their internal structure and evolution are uncertain but highly relevant to several problems in stellar and high-energy astrophysics.} {We discuss the main properties expected for the stable magnetic configurations in these stars from physical arguments and the ways these properties may determine the modes of decay of these configurations.} {We explain and emphasize the likely importance of the non-barotropic, stable stratification of matter in all these stars (due to entropy gradients in main-sequence envelopes and white dwarfs, due to composition gradients in neutron stars). We first illustrate it in a toy model involving a single, azimuthal magnetic flux tube. We then discuss the effect of stable stratification or its absence on more general configurations, such as axisymmetric equilibria involving poloidal and toroidal field components. We argue that the main mode of decay for these configurations are processes that lift the constraints set by stable stratification, such as heat diffusion in main-sequence envelopes and white dwarfs, and beta decays or particle diffusion in neutron stars. We estimate the time scales for these processes, as well as their interplay with the cooling processes in the case of neutron stars.} {Stable magneto-hydrostatic equilibria appear to exist in stars whenever the matter in their interior is stably stratified (not barotropic). These equilibria are not force-free and not required to satisfy the Grad-Shafranov equation, but they do involve both toroidal and poloidal field components. In main sequence stars with radiative envelopes and in white dwarfs, heat diffusion is not fast enough to make these equilibria evolve over the stellar lifetime. In neutron stars, a strong enough field might decay by overcoming the compositional stratification through beta decays (at the highest field strengths) or through ambipolar diffusion (for somewhat weaker fields). These processes convert magnetic energy to thermal energy, and they occur at significant rates only once the latter is less than the former; therefore, they substantially delay the cooling of the neutron star, while slowly decreasing its magnetic energy.} {} ",
        "introduction": "\\label{sec:intro}  Upper main sequence stars, white dwarfs, and neutron stars appear to have long-lived magnetic fields. These fields are organized on large scales, in the sense that the dipole field components (and perhaps some other, low-order multipoles; e.~g. \\citealt{Bagnulo99,Bagnulo00}) are not much weaker than the rms surface field, unlike the highly chaotic field of the Sun.  The highest detected (surface dipole) magnetic field strengths are $B_\\mathrm{max}\\sim 10^3~\\mathrm{G}$ in O stars (radius $R\\sim 10~R_{\\sun}$; \\citealt{Donati02,Donati06,Petit08}), $B_\\mathrm{max}\\sim 3\\times 10^4~\\mathrm{G}$ in chemicallly peculiar A and B stars (Ap/Bp stars, $R\\sim 3~R_{\\sun}$; \\citealt{Mathys97,Bagnulo99}), $B_\\mathrm{max}\\sim 10^9~\\mathrm{G}$ in white dwarfs ($R\\sim 10^4\\mathrm{km}$; \\citealt{Schmidt03}), and $B_\\mathrm{max}\\sim 10^{15}~\\mathrm{G}$ in ``magnetars'', a subclass of strongly magnetized neutron stars ($R\\sim 10~\\mathrm{km}$; \\citealt{Kouveliotou98,Woods99}), in all cases yielding very similar total magnetic fluxes, $\\Phi_\\mathrm{max}=\\pi R^2 B_\\mathrm{max}\\sim 10^{27.5}~\\mathrm{G~cm^2}$. This coincidence has often been interpreted as an argument for flux freezing during stellar evolution \\citep{Ruderman72,R01b,R03,Ferrario05a,Ferrario05b,Ferrario06}, although its feasibility has been called into question \\citep{TD93,Spruit08}. Of course, a large fraction of the original magnetic flux might be ejected with the stellar envelope. On the other hand, substantial field amplification through differential rotation, convection, and various instabilities could plausibly occur in proto-neutron stars if they are born rapidly rotating \\citep{TD93,Spruit02,Spruit08}.  Connected to the similar fluxes is that these stars also have similar ratios of fluid pressure ($P\\sim GM^2/R^4$, where $G$ is the gravitational constant and $M$ is the mass of the star) to magnetic pressure ($B^2/8\\pi$), \\begin{equation}\\label{beta} \\beta={8\\pi P\\over B^2}\\sim{8\\pi^3GM^2\\over\\Phi^2}\\sim 3\\times 10^6\\left(M\\over M_{\\sun}\\right)^2\\left(\\Phi\\over\\Phi_\\mathrm{max}\\right)^{-2}, \\end{equation} a large number, even for the most highly magnetized objects, implying that the magnetic field causes only very minor perturbations to their hydrostatic equilibrium structure.  Another similarity among these stars is that much or all of their structure is stably stratified, i.~e. stable to convection. The radiative envelopes of upper main sequence stars, as well as the whole interior of white dwarfs, are stabilized by the radially increasing specific entropy $s$, while in neutron stars the same effect is caused by a radially varying mix of different particle species \\citep{Pethick92,RG92,R01a}, which in the outer core reduces to a radially varying proton and electron fraction, $Y\\equiv n_p/(n_n+n_p)=n_e/(n_n+n_p)$, where $n_i$ stands for the number densities of neutrons ($i=n$), protons ($i=p$) and electrons ($i=e$).  The structure of the magnetic field inside these stars is not known, although it is highly relevant to their evolution: \\begin{itemize} \\item[1)] It affects the radial transport of angular momentum and chemical elements (e.~g. \\citealt{Heger05}); \\item[2)] it is plausibly the dominant source of energy for both the outbursts and the persistent emission of soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs), for this reason collectively called ``magnetars'' \\citep{TD93,TD96}; \\item[3)] it is likely to play an important role in the frequency spectrum of quasi-periodic oscillations observed after SGR flares \\citep{Levin07}; \\item[4)] it leads to slight deformations of neutron stars that could give rise to precession of pulsars \\citep{Wasserman03} and to the emission of gravitational waves \\citep{Cutler02}. \\end{itemize}  Magnetohydrodynamic (MHD) simulations of stably stratified stars with random initial magnetic field configurations have shown them to evolve on Alfv\\'en-like time scales into long-lived structures whose further evolution and decay appears to be controlled by dissipative processes (in the simulations, Ohmic diffusion; \\citealt{BS04,BS06,BN06}). Often, but not always \\citep{Braithwaite08}, these are roughly axisymmetric combinations of linked poloidal and toroidal components, whose external appearance is essentially dipolar. It appears plausible that these configurations approximate the true magnetic field structures in upper main sequence stars, white dwarfs, and neutron stars.  In \\S~2, we present arguments to the effect that the stable stratification of the stellar matter should be an essential ingredient to these equilibria. This means that, contrary to assumptions in the recent literature (e.~g. \\citealt{PerezAzorin06,Broderick08,Mastrano08}), these are definitely \\emph{not} force-free fields. In fact, it is shown in Appendix A that there are no true force-free equilibria in stars, while those proposed in the literature actually have singular magnetic forces on the stellar surface. Moreover, the fluid \\emph{cannot} be treated as barotropic, therefore the field components are \\emph{not} required to satisfy the Grad-Shafranov equation \\citep{Mestel56}, contrary to the popular belief \\citep{Tomimura05,Yoshida06,Haskell08,Akgun08,Kiuchi08}. In fact, the range of available equilibria becomes much wider in a stably stratified, non-barotropic fluid. The constraints imposed by the stability of these equilibria are far from obvious, but we argue that there are probably no equilibria in barotropic stars, while it is likely that there are equilibria with linked toroidal and poloidal fields in stably stratified stars.  Of course, the specific entropy $s$ and the proton fraction $Y$ are not perfectly conserved quantities within each fluid element, but can be changed by dissipative processes, discussed in \\S~3: In the case of $s$, through heat diffusion \\citep{Parker74}, in the case of $Y$, by (direct or inverse) beta decays or by ambipolar diffusion (motion of charged relative to neutral particles; \\citealt{Pethick92,GR92,TD96,Hoyos08}). Thus, the condition of stable stratification, and with it the hypothetically associated stable magnetic equilibrium configuration, although excellent approximations on short (Alfv\\'en-like) time scales, are eroded on the time scales of the dissipative processes mentioned above, leading to the decay of these structures \\citep{GR92}, and perhaps to a sudden loss of stability \\citep{BS04,BS06,BN06}. In main-sequence stars, white dwarfs, and weakly magnetized neutron stars, these appear to be too long to act on the stellar life time, but in strongly magnetized neutron stars their time scales become shorter, so they might plausibly drive magnetic field decay, leading to internal heating and to the magnetar phenomenon \\citep{TD96,Arras04}.  In principle, the Hall drift might also play a role in the field evolution in neutron stars \\citep{Jones88,Urpin,GR92,R05,R07a,Pons}. However, its time scale in neutron star cores tends to be somewhat longer than those of the other processes considered here \\citep{GR92}. Moreover, in an axially symmetric, equilibrium magnetic field configuration, the effect of the Hall drift is exactly cancelled by bulk fluid motions \\citep{RT08}, so we do not take it into account. For simplicity, we also refrain from discussing the role of the solid crust of the neutron star, as well as the effects of superconductivity and superfluidity, which alter the magnetic stresses \\citep{Easson77,Akgun08} as well as the dissipative processes. We also ignore the process of initial set-up of the magnetic equilibrium, which might be a highly dynamical process involving differential rotation and possibly a dynamo \\citep{TD93,Spruit02}, but concentrate on the properties imposed by the equilibrium and stability conditions and on the long-term evolution of the field.  A concise summary of our conclusions is given in \\S~4. Parts of this discussion have already been given elsewhere \\citep{R07b,R08}. ",
        "conclusions": "\\label{sec:conclusions}  This paper contains a general discussion of several physical issues related to the existence of large-scale, coherent magnetic structures in upper main sequence stars, white dwarfs, and neutron stars.  The main conclusions are the following:  \\begin{itemize} \\item[1)] Magnetic forces in these objects are generally weak compared to pressure and gravity forces, and their matter is strongly stratified by entropy or composition gradients. This means that at least some components of the magnetic forces can easily be balanced by other forces. Thus, there can be a wide variety of possible equilibria. These equilibria are not force-free; in fact, force-free equilibria are not possible in stars.  \\item[2)] If the magnetic structure is axially symmetric, the only constraint it has to satisfy to be balanced by pressure and gravity forces is that the azimuthal component of the Lorentz force must vanish. This means that there must be a set of magnetic surfaces of toroidal topology containing both the magnetic field lines and the current flow lines. Since the fluid is not barotropic, there is no need for the magnetic field to satisfy a Grad-Shafranov equation.  \\item[3)] It is difficult to give general criteria for stability. However, it is likely that, in a stably stratified star, poloidal and toroidal field components of similar strength could stabilize each other. In a barotropic fluid, it is possible that no stable equilibria exist, as the magnetic field might rise buoyantly and be lost from the star.  \\item[4)] The long-term evolution of the magnetic field is likely to be governed by dissipative processes that erode the stable stratification. Heat diffusion in main sequence stars and white dwarfs appears to be too slow to cause an observable effect over the life time of these stars. In strongly magnetized neutron stars, ambipolar diffusion and beta decays might be causing the magnetic energy release observed in magnetars. \\end{itemize}  \\appendix"
    },
    "0809/0809.4362_arXiv.txt": {
        "abstract": "{A systematic estimation of the line of sight depth in the disk and bar regions of the Large and Small Magellanic Clouds (LMC and SMC) using red clump stars is presented.} {We used the red clump stars from the photometric data of the Optical Gravitational Lensing Experiment  (OGLE II) survey and the Magellanic Cloud Photometric Survey (MCPS) for both the Clouds to estimate the depth. } {The observed dispersion in the magnitude and colour distribution of red clump stars is used to estimate the depth, after correcting for population effects, internal reddening within the Clouds and photometric errors.}{The observed dispersion due to the line of sight depth ranges from 0.023 mag, to 0.45 mag (a depth of 500 pc to 10.4 kpc) for the LMC and, from 0.025 to 0.34 magnitude (a depth of 670 pc to 9.53 kpc) for the SMC. The minimum value corresponds to the dispersion that can be estimated due to errors. The depth profile of the LMC bar indicates that it is flared. The average depth in the bar region is 4.0$\\pm$1.4 kpc. The northern disk is found to have depth (4.17$\\pm$0.97 kpc) larger than the southern part of the disk (2.63$\\pm$0.8 kpc). There is no indication of depth variation between the eastern (2.8$\\pm$0.92 kpc) and the western (3.09$\\pm$0.99 kpc) disk. The average depth for the disk is 3.44$\\pm$ 1.16 kpc. The SMC is found to have larger depth than the LMC. In the case of the SMC, the bar depth (4.90$\\pm$1.23 kpc) and the disk depth (4.23$\\pm$1.48 kpc) are found to be within the standard deviations. A prominent feature in the SMC is the increase in depth near the optical center. The averaged depth profile near the center resembles a structure like a bulge. } {The large dispersions estimated in the LMC bar and the northern disk suggest that the LMC either has large depth and/or different stellar population in these regions. The halo of the LMC (using RR Lyrae stars) is found to have larger depth compared to the disk/bar, which supports the existence of an inner halo for the LMC. On the other hand, the estimated depths for the halo (RR Lyrae stars) and disk are found to be similar, for the SMC bar region. Thus, increased depth and enhanced stellar as well as HI density near the optical center suggests that the SMC may have a bulge.} ",
        "introduction": "\\hspace{0.5cm} The Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) are nearby irregular galaxies at a distance of about 50 kpc and 60 kpc respectively. Both these galaxies are nearly face on. Magellanic Clouds (MCs) were believed to have had interactions with our Galaxy as well as between each other (Westerlund \\cite{westerlund}). It is also believed that the tidal forces due to these interactions have caused structural changes in these galaxies. The recent proper motion estimates by Kallivayalil et al (\\cite{K06a}, \\cite{K06b}) \\& Besla et al (\\cite{Besla07}) indicate that these Clouds may be approaching our Galaxy for the first time.  These results also claim that the MCs might not have always been a binary system. Therefore, it is not clear whether the structure of the MCs is modified due to their mutual interactions, interactions with our Galaxy or something else, like minor merges.\\\\  The N-body simulations by Weinberg (\\cite{Weinberg}) predicted that the LMC's evolution is significantly affected by its interactions with the Milky Way, and the tidal forces will thicken and warp the LMC disk. Alves and Nelson (2000) studied the carbon star kinematics and found that the scale height, h, increased from 0.3 to 1.6 kpc over the range of radial distance, R, 0.5 to 5.6 kpc and hence concluded that the LMC disk is flared. Using an expanded sample of carbon stars, van der Marel et al. (\\cite{van02}) also found that the thickness of the LMC disk increases with the radius. The depth of the clouds could vary as a function of radial distance from the center due to tidal forces. The presence of variation in depth across the LMC, if present, is likely to give valuable clues to the interactions it has experienced in the past. There has not been any direct estimate of the thickness or the line of sight depth of the bar and disk of the LMC so far.\\\\  Mathewson, Ford and Visvanathan (\\cite{MFV86a} \\& \\cite{MFV86b}) found that SMC cepheids extend from 43 to 75 kpc with most cepheids found in the neighbourhood of 59 kpc. Later, the line of sight depth of SMC was estimated (Welch \\cite{Welch87}) by investigating the line of sight distribution and period - luminosity relation of cepheids. They accounted for various factors which could contribute to the larger depth estimated by Mathewson et al. (\\cite{MFV86a} \\& \\cite{MFV86b}), and found the line of sight depth of the SMC to be $\\sim$ 3.3 kpc. Hatzidimitriou et al. (\\cite{H89}), estimated the line of sight depth in the outer regions of the SMC to be around 10-20 kpc.\\\\  Red Clump (RC) stars are core helium burning stars, which are metal rich and slightly more massive counterparts of the horizontal branch stars. They have tightly defined colour and magnitude, and appear as an easily identifiable component in colour magnitude diagrams (CMDs). RC stars were used as standard candles for distance determination by Stanek et al. (\\cite{PS98}). They used the intrinsic luminosity to determine the distance to MCs as well as to the bulge of our Galaxy.  Olsen and Salyk (\\cite{OS02}) used their constant I band magnitude to show that the southern LMC disk is warped. Subramaniam (\\cite{S03}) used the constant magnitude of RC stars to show that the LMC has some structures and warps in the bar region. Their characteristic colour was used by Subramaniam (\\cite{S05}) to estimate the reddening map towards the central region of the LMC. \\\\  In this paper, we used the dispersions in the colour and magnitude distribution of RC stars for depth estimation. The dispersion in colour is due to a combination of observational error, internal reddening (reddening within the disk of the LMC/SMC) and population effects. The dispersion in magnitude is due to internal disk extinction, depth of the distribution, population effects and photometric errors associated with the observations. By deconvolving other effects from the dispersion of magnitude, we estimated the dispersion only due to the depth of the disk. The advantage of choosing RC stars as a proxy is that there are large numbers of these stars available to determine the dispersions in their distributions with good statistics, throughout the L\\&SMC disks. The depth estimated here would correspond to the depth of the intermediate age L\\&SMC disks. The depth of the intermediate age disk of these galaxies may give clues to their formation and evolution, and thus in turn would give clues to their mutual interactions. This could also place some constraints on their interaction with our Galaxy. Measurements of line of sight depth in the central regions of MCs, especially the LMC, is of strong interest to understand the observed microlensing towards these galaxies.\\\\  The next section deals with the contribution of population effects to the observed dispersion of RC stars. Data sources are explained in section 3 and details of the analysis is described in section 4. Internal reddening in the MCs is explained in section 5. The LMC and SMC results are presented in sections 6 and 7 respectively and their implications are discussed in section 8. These results for the disks are compared with the depth estimates of the halo, as defined by RR Lyrae stars, in sections 9 and 10. Conclusions are given in section 11. ",
        "conclusions": "The line of sight depth of MCs are estimated using red clump stars as tracers. They are intermediate age stars with ages greater than 1 Gyr, hence, the depth estimates correspond to the intermediate age disk. The depth of the intermediate age disk may give clues to its evolution, as a thin disk is indicative of an undisturbed disk, whereas a thick disk would indicate a heated up and hence disturbed disk. Thus the depth of the disks as well as the bar regions of these galaxies are clues to the evolution of these galaxies. The analysis presented here estimates the dispersion in the RC magnitude distribution due to depth, after correcting for dispersion due to other effects. The corrections due to the internal reddening and observational error were estimated for each region and were corrected accordingly. The correction for the presence of an age and metallicity range in the RC population, along with star formation history is done using the dispersion estimated by Girardi \\& Salaris (\\cite{GS01}). Thus the values estimated here have a bearing on the assumptions made during the correction of the above effect.\\\\   The variation in the estimated dispersion (after all the corrections) is assumed to be due to variations in depth across the galaxy. If this is actually due to the differences in RC population, as a result of variation in age, metallicity and star formation history, then the results indicate that the RC population and their properties are significantly different in the bar, north and south of the disk, for the LMC. Recent studies show that the above parameters are more or less similar across the LMC (Subramaniam \\& Anupama (\\cite{SA02}), Olsen \\& Salyk (\\cite{OS02}) and van der Marel \\& Cioni (\\cite{vc01})). Variation of star formation history and metallicity across the LMC has been studied by Cioni et al. (\\cite{C06LMC}), Carrera et al. (\\cite{Car08}). They found small variations in the inner regions, but large variations in the outer regions. Cioni et al. (\\cite{C06LMC}) mapped the variation of star formation history as well as metallicity using the AGB stars. They found that the south and south-western regions could be metal poor whereas the north and north-eastern regions could be metal rich. Regions near the Shapley constellation III were found to have a younger population. This region is close to the northern limit studied here. Piatti et al. (\\cite{Piatti99}) and Dolphin (\\cite{D00}) found that the RC population in the far northern regions is structured. All the above regions with varying RC population are outside the regions studied here. % The results presented in this study suggests that the northern regions and the bar have large dispersion, probably due to depth when compared to the east, west and the southern regions. If this is not due to increased depth in these regions, then it would mean that the stellar populations in these regions are significantly different. In either case, the northern disk and the bar seem to be different from the rest of the LMC disk.\\\\  As incompleteness correction is done in one data set (OGLE II ) and not in the other (MCPS) we compared the depth estimates before and after adopting the completeness correction. We found that the change is within the bin sizes adopted here. Thus, incorporating the incompleteness correction has not changed the results presented here. The incompleteness correction is large in the central regions of the LMC bar, where the correction is between 30-40 percent. The outer regions of the bar (6-15\\%) as well as the NW region (5-9\\%) have less correction. The outer regions are thus less affected by the incompleteness. Thus the incompleteness problem is unlikely to affect the MCPS RC distribution in the LMC disk, whereas it may be unreliable in the central regions of the LMC bar. In the SMC, the incompleteness correction in the central regions is about 12\\% and that in the outer region is about 5\\%. This is the case for SMC also. The incompleteness in the MCPS data does not affect the results presented here.\\\\   We have removed regions in the MCs with poor fit as explained in section 4. These regions are likely to have different RC structures suggesting a large variation in metallicity and/or population. The fraction of such regions is about 8\\% in the LMC and 5.3\\% in the SMC. Such regions are indicated in figures 9 \\& 11. Thus to a certain extent, the above procedure has eliminated the regions with very different metallicity and star formation history that are seen in most of the regions. Apart from the above, the remaining regions studied here might have some variation in the the RC population contributing to the depth estimated. The results presented in this study will include some contribution from the population effect.\\\\  The estimated depth for various regions in the LMC is given in table 1. These values correspond to the 1-sigma depth. In the case of the LMC, the line of sight depth estimated here is for the inclined disk. To estimate the actual depth of the disk, one needs to correct for the inclination. Assuming the inclination to be 35 degrees (van der Marel et al. \\cite{vc01}), the actual depth of the bar is 4.0$\\times$cos{\\it(i)} = 3.3 $\\pm$ 1.0 kpc. Similarly, the southern disk has an actual depth of about 2.2 $\\pm$ 1.0 kpc, whereas the northern disk is similar to the bar. Note that, after correcting for inclination, the depths in the northern and the souther regions are within the errors. This is because the difference in depth reduces due to correction for inclination, but the error does not. These values should be used when one compares the depth of the LMC with that of other galaxies. Thus, the LMC bar and the disk are thicker than the thin disk of our Galaxy ($\\sim$ 100 pc). The scale height of the bar could be taken as half of its depth, assuming that the bar is optically thin. The z-extent of the bar, which is the scale height (1.65 kpc), is found to be similar to its scale length (1.3 - 1.5 kpc, van der Marel \\cite{van01}). Hence the bar has a depth similar to its width. Thus the bar continues to be an unexplained structure/component of the LMC.\\\\  The LMC bar is found to be fairly thick (line of sight depth = 4.0$\\pm$1.4 kpc). We also find evidence for flaring and a disturbed structure at the ends of the bar region. The thick and flared bar of the LMC indicates that this region of the LMC is perturbed. The structure of the LMC bar as delineated by the RC stars showed warps (Subramaniam \\cite{S03}), which is  also clear indication of disturbance. The depth estimates also suggest that the LMC disk is fairly thick (2.4 - 4.0 kpc), with a decrease in depth/thickness and/or varying stellar population from the north to the south. The tidal effects due to LMC-Galaxy interactions (if they were interacting) are unlikely to cause this, as the tidal effects are stronger near the outer regions and weaker towards the inner regions. Flaring of the disk is expected if tidal interactions are present. On the other hand, except for the thicker northern region, flaring of the disk is not seen, at least up to the radii studied here. Hence we do not see any evidence for tidal interactions, at least in the inner disk.  The recent results on the proper motion of the LMC and SMC (Kallivayalil et al \\cite{K06a}) suggested that the Clouds are approaching our Galaxy for the first time. This would suggest that the LMC has not interacted with our Galaxy before. Our results are in good agreement with this scenario.\\\\  In general, thicker and heated up disks are considered as signatures of minor mergers (Quinn \\& Goodman \\cite{QG86}; Velazquez \\& White \\cite{VW99}). Thus, the LMC is likely to have experienced minor mergers in its history. The presence of warps in the bar (Subramaniam \\cite{S03}) and evidence of counter rotation in the central regions (Subramaniam \\& Prabhu \\cite{SP05}) also support the minor merger scenario. Thus, it is possible that the LMC has experienced minor mergers during its evolution. These mergers have affected the northern disk and the bar. The variation in depth observed across the LMC disk could constrain the way in which these mergers could have happened.\\\\  The SMC is found to have a depth greater than the LMC. The disk and the bar does not show much difference in depth. A striking result is the increased depth in the central region of the SMC. The profile of the depth near the center (figure 16) looks very similar to a typical luminosity profile of a bulge. This could suggest the presence of a bulge near the optical center of the SMC. If a bulge is present, then a density/luminosity enhancement in this region is also expected. We plotted the observed stellar density in each region from the MCPS data to see whether there is any such central enhancement. This is shown in figure 17. The regions with high density are shown as open circles, located close to the optical center. The regions with large depth are found to be within the ellipse shown in the figure. It can be seen that regions with  highest stellar density lie more or less within this ellipse. Cioni et al. (\\cite{C00}) studied the morphology of the SMC using the DENIS catalogue. They found that the distribution of AGB and RGB stars show two central concentrations, near the optical center, which match with the carbon stars by Hardy et al. (\\cite{Hardy89}). They also found that the western concentration is dominated by old stars. The approximate locations of these two concentrations found by Cioni et al. (\\cite{C00}) are shown as hexagons in figure 17. Also, the strongest HI concentration in the SMC map by Stanimirovic et al. (\\cite{SS99}) falls between these two concentrations. The maximum HI column density, 1.43 $\\times$ 10$^{22}$ atoms cm$^{-2}$ is located at RA = 00$^h$ 47$^m$ 33$^s$, Dec = -73$^0$ 05' 26\" (J2000.0) (Stanimirovic et al. \\cite{SS04}). This location is shown as a large triangle in figure 17. The optical center of the SMC is shown as a large square. All these peaks as well as the optical center are located on or within the boundary of the ellipse. Thus, the peaks of stellar as well as the HI density are found within the central region with large depth. This supports the idea that a bulge may be present in the central SMC. This bulge is not very luminous, but clearly shows enhanced density. It is also the central region of the triangular shaped bar.\\\\  The increased dispersion near the SMC center, which is interpreted as due to large depth, could be partially due to the presence of RC population which is different. Cioni et al. (\\cite{C06SMC}) did not find any different population or metallcity gradient near the central regions. Tosi et al. (\\cite{T08}) obtained deep CMDs of 6 SMC regions to study the star formation history. Three of their regions are located close to the bar and three are outside the bar. They found an apparent homogeinty of the old stellar population populating the subgiant branch and the clump. This suggested that there is no large differences in age and metallicity among old stars in these locations. Their SF1 region is located close to the region of large depth identified here. The RC population in this region is found to be very rich and the spread in magnitude is greater than those found in the other CMDs. This spread is also suggestive of increased depth near this location. It will be worthwhile to study the star formation history of regions near the SMC center to understand how different the stellar population is in this suggested bulge.\\\\  It may be worthwhile to see whether this bar is actually an extended/deformed bulge. It is interesting that the so-called triangular shaped bar of the SMC is also an unexplained component, which does not show the signatures of a typical bar. This could naturally explain the formation of the odd shaped bar in the SMC. Thus, we propose that the central SMC has a bulge. The elongation and the rather non-spherical appearance of the bulge could be due to tidal effects or minor mergers (Bekki \\& Chiba \\cite{BC08}).\\\\"
    },
    "0809/0809.4648_arXiv.txt": {
        "abstract": "%  In massive stars, magnetic fields are thought to confine the outflowing radiatively-driven wind, resulting in X-ray emission that is harder, more variable and more efficient than that produced by instability-generated shocks in non-magnetic winds. Although magnetic confinement of stellar winds has been shown to strongly modify the mass-loss and X-ray characteristics of massive OB stars, we lack a detailed understanding of the complex processes responsible. The aim of this study is to examine the relationship between magnetism, stellar winds and X-ray emission of OB stars. In conjunction with a Chandra survey of the Orion Nebula Cluster, we carried out spectropolarimatric ESPaDOnS observations to determine the magnetic properties of massive OB stars of this cluster. We found of two new massive magnetic stars in the Orion Nebula Cluster: HD\\,36982 and HD\\,37061, for which the estimated dipole polar strengths are $1150^{+320}_{-200}$\\,G and $620^{+220}_{-170}$\\,G, respectively. However, the apparent lack of clear correlation between X-ray indicator and the presence of a magnetic fields brings forth new challenges for understanding the processes leading to X-ray emission in massive stars. ",
        "introduction": "Magnetic fields are well known to produce X-rays in late-type convective stars like the Sun. On the other hand, X-rays emission from massive stars is thought to have a different origin. Their powerfull winds are driven by radiation, which is an unstable process. These instabilities result in collisions between wind streams of different velocity, resulting in small shocks that generate X-ray emission \\citep{1999ApJ...520..833O,1980ApJ...241..300L}. Typically, this emission is soft and stable. Also, empirical studies have shown that the X-ray luminosity will be about $10^{-7}$ times the bolometric luminosity for O stars, and will decrease rapidly between spectral type B1 and B3 \\citep{1997ApJ...487..867C}.  However, some observations are not consistent with such a model. The best exemple is $\\theta^1$\\,Ori\\,C, in the Orion Nebula Cluster. This young O star displays hard and efficient X-rays, and shows a modulated X-ray light curve \\citep{1997ApJ...478L..87G}, which varies with the same period as lines in the UV and optical spectrum \\citep{1993A&A...274L..29S,1996A&A...312..539S,1994ApJ...425L..29W}.  \\citet{1997ApJ...485L..29B} explained this peculiar X-ray emission in terms of the ``magnetically confined wind shock'' model (MWCS). In this model, the dipolar stellar magnetic field channels the outflowing wind along the field lines, resulting in a closed magnetosphere with large-scale equatorial shocks which enhance the X-ray emission and possibly modulate it with stellar rotation. Such a magnetic field ($1.1\\pm0.1$\\,kG) was subsequently discovered by \\citet{2002MNRAS.333...55D}.  From that exemple, we can speculate that X-ray enhancement and periodic modulation may be associated with the presence of a magnetic field, in the form of a magnetospheric confining of the wind.  \\subsection{The Orion Nebula cluster}  The Orion Nebula Cluster (ONC) presents a unique opportinuty to characterize the magnetic fields of nearby massive OB stars. A \\textit{Chandra} large program was dedicated to observe the ONC in X-rays: The \\textit{Chandra} Orion Ultradeep Project (COUP) consisted of a 10-day exposure of the ONC \\citep{2005ApJS..160..319G}. \\citet{2005ApJS..160..557S} have studied the COUP massive star sample, containing 9 massive OB stars for which radiatively-driven winds should be the dominant mechanism for X-ray emission. The list of stars, along with their properties, can be found in Table \\ref{tab1}. \\citeauthor{2005ApJS..160..557S} found a large variety of X-ray properties for the ONC massive stars. For example there is a large scatter (3 orders of magnitude) around the empirical relation between the X-ray luminosity and the bolometric luminosity ($L_\\mathrm{X} / L_\\mathrm{bol} = 10^{-7}$), and also a large incidence of variability and flares. The X-ray properties are summarised in Figure \\ref{fig1}. Altogether, the X-ray results, combined with other indicators, strongly suggest that all but two stars (NU\\,Ori and $\\theta^1$\\,Ori\\,D) of the ONC are magnetic \\citep{2005ApJS..160..557S}.  \\begin{figure} \\plotone{VPetitFig1.eps} \\caption{\\label{fig1}X-ray efficiency of the ONC massive stars as a function of effective temperature (Stelzer et al. 2005). The 3 detected stars are circled. Filled symbols are for stars with indirect indications of the presence of a magnetic field, and grey symbols are for confirmed or suspected binaries. Plotting symbols indicate the following properties: circles are for stars showing possible X-ray rotational modulation, squares are for T\\,Tauri type X-ray emission, triangles are chemically peculiar (CP) stars, and the diamond star was not observed. The dotted lines indicates the typical efficiency for massive O-type stars, and a qualitative illustration of the sharp decrease in the B-type stars range.} \\end{figure} ",
        "conclusions": "The role of magnetic fields in X-ray production remains poorly understood. As shown in Figure \\ref{fig1}, X-ray variability and enhancement is not necessarily correlated with the presence of a magnetic field. Furthermore, magnetic fields are observed even in the absence of X-ray indications. Theoretical work, such as numerical MHD simulations of magnetically confined stellar winds done by \\citet{2008MNRAS.385...97U} will help disentangle various effect such as rotation, geometry of the magnetic fields, etc. that could explain the current observations. However, we also need more observations to constrain those simulations. A large CFHT observing program was allocated to the Magnetism in Massive Stars (MiMeS) collaboration that will allow us to better understand the interaction of stellar winds and magnetic fields."
    },
    "0809/0809.3297_arXiv.txt": {
        "abstract": "%  The near-star environment around obscured stars is very dynamic. Many classes of stars show evidence for winds, disks, inflows and outflows with many phenomena occurring simultaneously. These processes are involved in stellar evolution, star and planet formation, and influence the formation and habitability of planets around host stars. Even for the nearest stars, this region will not be imaged even after the completion of the next generation of telescopes. Other methods for measuring the physical properties of circumstellar material must be developed. The polarization of light across spectral lines is a signature that contains information about the circumstellar material on these small spatial scales. We used the HiVIS (R=13000 to 50000) and ESPaDOnS (R=68000) spectropolarimeters to monitor several classes of stars on over a hundred nights of observing from 2004-2008. In 10/30 classical Be stars, the traditional broad depolarization morphology is reproduced, but with some additional absorptive effects in 4 of these 10 stars. In Herbig Ae/Be stars roughly 2/3 of the stars (14/20) with strong absorptive components (either central or blue-shifted) showed clear spectropolarimetric signatures typically centered on absorptive components of the spectral lines. They were typically 0.3\\% to 2\\% with some signatures being variable in time. Post-AGB and RV-Tau type evolved stars showed very strong absorptive polarimetric effects (5/6 PAGB and 4/4 RVTau) very similar to the Herbig Ae/Be stars. These observations were inconsistent with the traditional scattering models and inspired the development of a new explanation of the observed polarization. This new model, based on optical-pumping, has the potential to provide direct measurements of the circumstellar gas properties. ",
        "introduction": "The formation and evolution of stars is a dynamic process involving circumstellar matter. Over the course of a few million years, the circumstellar gas and dust around a young star will either accrete onto the star, turn in to planets, or dissipate in the form of winds and jets. There is evidence of simultaneous accretion, outflows, ionized disk structures, and strong stellar winds in young stars. All of these processes influence the environment of the planets around these stars. Similarly, in evolved or rapidly rotating stars, the outer layers are shed to form disks, winds or envelopes. There are only a few techniques that can put meaningful constraints on the environment very close to a star. Spectropolarimetry is such a technique but it is still considered a difficult specialty field. There are only two high-resolution instruments (R$>$10000), ESPaDOnS and HiVIS on large telescopes (over 3m) capable of these measurements (Donati et al. 1997, Harrington \\& Kuhn 2008a). However, this technique is a powerful probe of small spatial scales, being sensitive to the geometry and density of the material very close the central star. In a typical model, material in the region a few to several stellar radii away from the star produces the polarization. Even for the closest young stars, these spatial scales are smaller than 0.1 milliarcseconds across and will not be imaged directly, even by the next generation telescopes. Since the circumstellar material is involved in accretion, outflows, winds and disks, with many of these phenomena happening simultaneously, spectropolarimetry can put unique constraints on the types of densities and geometries of the material involved in these processes.  \\begin{figure} \\begin{center} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1a.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1b.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1c.eps}  \\\\ \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1d.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1e.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1f.eps} \\\\ \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1g.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1h.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1i.eps} \\\\ \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1j.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1k.eps} \\includegraphics[width=0.23\\linewidth, angle=90]{Harrington-Fig1l.eps} \\\\ \\caption{This figure shows ESPaDOnS spectropolarimetry of Herbig Ae/Be stars, Post-AGB, RV-Tau and Be type stars. The first two rows show the Herbig Ae/Be stars AB Aurigae, MWC 480, HD 163296, MWC 120, MWC 158 and HD 144432 from left to right. The third row shows the Post-AGB/RV-Tau type stars 89 Her, SS Lep and U Mon. The last row shows the Be stars MWC 143, $\\psi$ Per and the Herbig Be star MWC 361 for comparison. Very distinct morphologies are seen. The Herbig stars showed nearly ubiquitous polarization-in-absorption (detected in roughly 2/3 of the stars with strong absorptive components and none of the Herbig stars with other line profile types). In the Post-AGB \\& RV-Tau stars absorptive effects were ubiquitous (9/10). The Be stars in the last row illustrate the broad ``depolarization'' morphology but with modifications. In our survey 4/10 of these detections showed additional absorptive effects not seen at lower spectral resolution. These observations represent an order-of-magnitude gain in spectral resolution at high signal-to-noise as no flux-dependent wavelength averaging is required to achieve the needed sensitivity. This data set shows the ubiquity of absorptive polarization effects.} \\end{center} \\end{figure}  With high spectral resolution ($\\frac{\\lambda}{\\delta\\lambda}>$10000) clear changes across individual emissive or absorptive line components are observable. Circumstellar disks, rotationally distorted winds, magnetic fields, or asymmetric radiation fields (optical pumping) are all used to explain these signatures. In most models, the polarization change comes directly from the circumstellar material and can be used as a proxy for the physical properties of that material. Spectropolarimetric signals measured to date are small, typically 0.1\\% to 1\\%. Measuring these signals requires very high signal to noise and careful control of systematics.  There are many models of linear spectropolarimetric effects. Analytical studies showed spectropolarimetric effects from circumstellar disk scattering (McLean 1979, Wood et al. 1993). Recent modelling of absorptive and scattering processes has shown a wealth of spectropolarimetric effects from disks, winds, and envelopes (cf. Harries 2000, Ignace 2004, Vink et al. 2005, Kuhn et al. 2007). The depolarization effect is caused by unpolarized line emission from stellar envelopes diluting a polarized stellar continuum (McLean 1979). Small clumps in a stellar wind scatter and polarize significant amounts of light can enhance the polarization at that clump's specific velocity and orientation (Harries 2000). Magnetic fields create polarized line profiles through the Zeeman effect. Thin disks can differentially scatter light and produce broad spectropolarimetric profiles (Wood et al. 1993, Wood Brown \\& Fox 1994). Optically pumped gas can produce polarization through absorption (Kuhn et al. 2007). ",
        "conclusions": ""
    },
    "0809/0809.4224_arXiv.txt": {
        "abstract": "Models of magnetospheric accretion on to classical T Tauri stars often assume that the stellar magnetic field is a simple dipole.  Recent Zeeman-Doppler imaging studies of V2129~Oph and BP~Tau have shown however that their magnetic fields are more complex.  V2129~Oph is a high mass T Tauri star and despite its young age is believed to have already developed a radiative core.   In contrast to this, the lower mass BP~Tau is likely to be completely convective.  As the internal structure and therefore the magnetic field generation process is different in both stars, it is of particular interest to compare the structure of their magnetic fields obtained by field extrapolation from magnetic surface maps.   We compare both field structures to mulitpole magnetic fields, and calculate the disk truncation radius for both systems.  We find that by considering magnetic fields with a realistic degree of complexity, the disk is truncated at, or within, the radius obtained for dipole fields. ",
        "introduction": "Classical T Tauri stars (cTTs) are young pre-main sequence stars which accrete material from circumstellar disks. Many observations support the magnetospheric accretion scenario, where the stellar magnetic field is strong enough, and sufficiently well-ordered, to truncate the disk at a few stellar radii \\cite{kon91}.  Gas is forced to follow the field lines of the stellar magnetosphere and accretes on to the star at high velocity, producing detectable hotspots \\cite{edw94}. However, until recently little was known about the geometry of their magnetic fields, with the majority of models assuming a dipole.  Measuring the Zeeman broadening of unpolarized spectral lines has proved to be very successful in detecting T Tauri magnetic fields (e.g. \\cite{joh07}), although this technique is insensitive to the magnetic topology.  In contrast, measuring the polarization signature in spectral lines gives access to the field topology, and allows information about how the magnetic energy is distributed within the different field components to be determined.  However, like all polarization techniques, this suffers from flux cancellation effects and yields limited information regarding the field strength.  \\begin{figure*} \\centering \\begin{tabular}{cc} \\label{v2129oph_map} \\includegraphics[height=.22\\textheight]{gregory_talk_fig1a.pdf} & \\label{bptau_map} \\includegraphics[height=.22\\textheight]{gregory_talk_fig1b.pdf} \\\\ \\end{tabular} \\caption{Flattened polar projections of V2129~Oph (left) and BP~Tau (right) showing the strength in Gauss and polarity of the radial field component \\cite{don07,don08}.  Red denotes positive, and blue negative, field regions.  The solid black circle represents the stellar equator, and the dashed circles lines of constant latitude.  The surface magnetic fields are clearly non-dipolar.} \\label{maps} \\end{figure*}  The polarization signatures detected in spectral lines are typically small, and therefore cross-correlation techniques (such as Least-Squares Deconvolution \\cite{don97}) were developed in order to extract information from as many spectral lines as possible.  The signal-to-noise ratio of the resulting average Zeeman signature is several tens of times larger than that of a single spectral line. Magnetic surface features produce distortions in the Zeeman signature that depend on the latitude and longitude of the magnetic region, as well as the orientation of the field lines.  By monitoring how such distortions move through the Zeeman signature as the star rotates, a method referred to as Zeeman-Doppler imaging (ZDI), the distribution of magnetic polarities across surface of cool stars can be determined (see Donati et al, these proceedings).  ZDI studies have now been carried out on two cTTs, V2129~Oph and BP~Tau \\cite{don07,don08}.  Both stars have complex non-dipolar magnetic fields (Fig. \\ref{maps}).  The field of V2129~Oph is dominated by a 1.2~kG octupole field component, tilted at $\\sim 20\\,^{\\circ}$ with respect to the stellar rotation axis.  The dipole component was found to be weak, with a polar strength of only 0.35~kG and tilted at $\\sim 30\\,^{\\circ}$ \\cite{don07}.  The field of BP~Tau consists of strong dipole and octupole field components of strength 1.2 and 1.6~kG, both tilted by $\\sim 10\\,^{\\circ}$ with respect to the stellar rotation axis, but in different planes \\cite{don08}.   \\subsection{Field extrapolation and comparison with multipoles} From the surface maps of the photospheric field, we reconstruct the coronal field via field extrapolation, using the potential field source surface method \\cite{alt69}. This has the advantages over full MHD models of computational speed and simplicity, and does not require assumptions about an equation of state, or the energetics.  The disadvantages are that non-potential and time dependent effects can not be modeled.  However, it has been demonstrated that the potential field model is often adequate and produces results which match those of more complex MHD models \\cite{ril06}.  We note that for cTTs the main source of non-potentiality in the field is likely to be due to the interaction between the stellar magnetosphere and the disk.  We do not consider this here - the field structures that we present are static and do not evolve in time.  \\begin{figure} \\includegraphics[height=.29\\textheight]{gregory_talk_fig2.pdf} \\caption{The field structure of V2129~Oph (left) and BP~Tau (right) derived by field extrapolation from the magnetic surface maps in Fig. \\ref{maps}.  Closed/open field lines are shown in white/blue.  There is a clear distinction between the complex surface field and the simpler larger scale field.} \\label{extrap} \\end{figure}  The potential field source surface model requires two boundary conditions.  The first is that the radial component of the magnetic field at the stellar surface is that which is determined from the Zeeman-Doppler maps.  The second is that at some height above the star, known as the source surface $R_S$, the field becomes purely radial, which mimics the effect of a stellar wind blowing open field lines at some height above the star. Although the choice of $R_S$ is a free parameter of the field extrapolation model, its location is well constrained by the observed accretion hotspot locations \\cite{don07,don08}.  Once $R_S$ has been set, predictions can then be made regarding the X-ray emission properties of the coronal field \\cite{jar08}, which will be tested against future X-ray observations.  The field extrapolation technique applied specifically to cTTs is discussed in detail by \\cite{gre06}, while Fig. \\ref{extrap} shows the field topology obtained by extrapolation from the magnetic maps of V2129~Oph and BP~Tau shown in Fig. \\ref{maps}.   For both stars the surface field is complex and non-dipolar, while the larger scale field appears more well-ordered (Fig. \\ref{extrap}).  As magnetospheric accretion models often assume that cTTs have simple dipolar fields (with the exception of recent work, e.g. \\cite{lon08}), it is of interest to compare the field structures obtained by field extrapolation to a dipole with a source surface. Furthermore, as both stars have strong octupole field components, we also compare the extrapolated fields with octupoles.  To compare the field structures we calculate the heights and widths of the field line loops obtained by field extrapolation (Fig. \\ref{height_width}). We define the height, $h$, of a field line loop as the maximum height of the loop above the stellar surface. The width of a field line, $w$, is defined as the distance along the segment of the great circle connecting the field line footpoints on the stellar surface \\cite{gre08}.  \\begin{figure*} \\centering \\begin{tabular}{cc} \\label{v2129oph_height_width} \\includegraphics[height=.28\\textheight]{gregory_talk_fig3a_reduced.pdf} & \\label{bptau_height_width} \\includegraphics[height=.28\\textheight]{gregory_talk_fig3b_reduced.pdf} \\\\ \\end{tabular} \\caption{Height vs width of the field line loops derived by extrapolation form the surface maps of V2129~Oph (left) and BP~Tau (right).  The black points represent the extrapolated fields and the red/green points a dipole/octupole fields.  For both stars, the larger scale field lines are wider than a dipole.} \\label{height_width} \\end{figure*}  For BP~Tau the field structure follows a similar trend to a dipole field, which reflects its very strong dipole field component.  However, the field structure of V2129~Oph shows significant departures from a dipole, and is characterized by numerous compact loops. The difference in field complexity is likely related to the different internal structure of both stars, which affects the magnetic field generation process, and ultimately the topology of the large scale field.  V2129~Oph is believed to have a radiative core \\cite{don07} and has a particularly complex field structure, whereas BP~Tau is completely convective and is found to have a simpler, although still non-dipolar, field topology \\cite{don08}.  For both V2129~Oph and BP~Tau we find that larger scale field lines are wider than dipole magnetic loops (Fig. \\ref{height_width}).  The larger scale loops, which are the ones that will interact with the disk, are distorted close to the star by the strong surface field regions.  Therefore, even though the large scale magnetic field is well-ordered and dipole-like, the structure of the accreting field is influenced by the complex field regions close to the star \\cite{gre08}.  We note that recently \\cite{fis08} have come (independently) to the same conclusion for some stars in their sample, and comment that in order to explain the deep absorption component detected in a certain line of He, ``flows near the star with less curvature than a dipole trajectory seem to be required''.   \\subsection{Disk truncation} In this section we compare how the disk truncation radius $R_t$ differs between considering a dipole magnetic field and the fields derived by field extrapolation.  Various methods of calculating $R_t$ have been developed \\cite{bes08}, and in this work we assume that the inner disk is truncated where the torque due to viscous processes in the disk is comparable to the magnetic torque due to the stellar magnetosphere \\cite{cla95}.  We assume that at the inner disk the perturbed toroidal component is equal to the poloidal component ($B_{\\phi} \\sim B_z$, see \\cite{gre08}).  This assumption should be valid as long as the disk is truncated sufficiently within the corotation radius and assuming the disk is strongly coupled to the stellar magnetic field \\cite{mat05b}.  Equating the differential magnetic and viscous torques, and assuming that the poloidal field threading the disk is dominated by the vertical component ($B_z \\gg B_r$), gives \\begin{equation} r^2B_z^2 = \\frac{1}{2}\\dot{M}\\left (\\frac{GM_{\\ast}}{r} \\right )^{1/2}. \\label{trunc} \\end{equation} The solid lines in Fig. \\ref{trunc_plot} represent the variation in the RHS equation (\\ref{trunc}) along the stars equatorial plane.  We use the same stellar parameters as in \\cite{gre08}, namely for BP~Tau a mass, radius, rotation period and accretion rate of $0.7{\\rm M}_{\\odot}, 1.95{\\rm R}_{\\odot}, 7.6{\\rm d}$ and $2.88\\times 10^{-8}{\\rm M}_{\\odot}{\\rm yr}^{-1}$ respectively, and for V2129~Oph, $1.35{\\rm M}_{\\odot}, 2.4{\\rm R}_{\\odot}, 6.53{\\rm d}$ and $1\\times 10^{-8}{\\rm M}_{\\odot}{\\rm yr}^{-1}$. The various dashed/dotted lines in Fig. \\ref{trunc_plot} represent the LHS of equation (\\ref{trunc}) assuming different forms for the stellar magnetic field. The value of $r$ where the lines cross, is then the disk truncation radius $R_t$.  In the figures we have considered both an analytic dipole field, and a dipole field with a source surface surface, in order to allow a better comparison with the extrapolated fields of V2129~Oph and BP~Tau.  We assume that the dipoles have the same polar strength as the measured dipole components of each star (0.35/1.2~kG for V2129~Oph/BP~Tau).  \\begin{figure*} \\centering \\begin{tabular}{cc} \\label{v2129oph_trunc_plot} \\includegraphics[height=.28\\textheight]{gregory_talk_fig4a.pdf} & \\label{bptau_trunc_plot} \\includegraphics[height=.28\\textheight]{gregory_talk_fig4b.pdf} \\\\ \\end{tabular} \\caption{The variation along the equatorial plane of the differential magnetic/viscous torque [i.e. the LHS and RHS of equation (\\ref{trunc})] for V2129~Oph (left) and BP~Tau (right). The solid line represents the RHS of (\\ref{trunc}), with the LHS for an aligned dipole field (dashed line), a tilted dipole field with a source surface (dash-dot line) and the field extrapolation of V2129~Oph/BP~Tau (dotted line).   In both cases the disk is truncated within the corotation radius of $6.7\\,{\\rm R}_{\\ast}$ (V2129~Oph) and $7.4\\,{\\rm R}_{\\ast}$ (BP~Tau).} \\label{trunc_plot} \\end{figure*}  For both V2129~Oph and BP~Tau we find that the disk is truncated within the corotation radius, justifying our assumption above (Fig. \\ref{trunc_plot}). For the extrapolated field of BP~Tau, $R_t$ is only slightly within what is expected for a dipole magnetic field.  This reflects the strong dipole component of BP~Tau, which is the dominant field component at the inner edge of the disk.  In contrast, $R_t$ for the extrapolated field of V2129~Oph is almost a stellar radii within that calculated using a dipole magnetic field.  This is due to intrinsic complexity of the magnetic field of V2129~Oph.  The field strength in the disk midplane for this star is less, as the drop-off in field strength with height above the star is greater than with BP~Tau, and therefore the disk is able to extend closer to the stellar surface than would be expected with  a dipole magnetic field \\cite{gre08}.  Our conclusion is that when considering magnetic fields with a realistic degree of complexity the disk will be truncated at, or within, the radius determined from dipole magnetic field models. ",
        "conclusions": "We tentatively find that the complexity of the magnetic field topology of accreting T Tauri stars depends on the internal structure of the star.  V2192~Oph, which has already developed a radiative core despite its young age, has a more complex and dominantly octupolar field, compared to the much simpler field of the completely convective BP~Tau, which has a strong dipole field component. However, this conclusion is based on only two stars and more data across a range of stellar masses will be required to confirm if this is a general feature of all cTTs, although a similar result has been found for low-mass main-sequence stars (see Morin et al, these proceedings).  For V2129~Oph there is a more rapid drop-off in field strength with height above the stellar surface compared with BP~Tau, which means that the inner disk is truncated much closer to the star compared with a dipole magnetic field model.  For BP~Tau, with its stronger dipole component, $R_t$ is roughly the same as expected with a dipole field. Thus the interaction between the stellar field and the disk may be different for stars with different internal structures.  The surface fields of both stars are particularly complex, especially for V2129~Oph, but for both stars the larger scale, likely accreting field, is more well-ordered.  However, the structure of the larger scale field is distorted close to the stellar surface by the strong surface field regions.  Thus many of the larger scale field lines are wider than dipole magnetic loops, a result which other authors have recently found \\cite{fis08}.  In order to quantify the above effects and discover if they are general features of all accreting T Tauri stars, more spectropolarimetric data of stars across a range of masses, radii, rotation periods and accretion rates are required.  Such new data will be obtained over the coming years as part of the MaPP (Magnetic Protostars $\\&$ Planets) project, a large program with the ESPaDOnS spectropolarimeter at the Canada-France-Hawai'i telescope (PI: J.-F.~Donati), which will provide 690 hours of observing between late 2008 and 2012. Such new data will provide an unparalleled opportunity to investigate magnetism on solar-like pre-main sequence stars, and give new insight into the history of our Sun and Solar System at a time when the planets were begining to form."
    },
    "0809/0809.4012_arXiv.txt": {
        "abstract": "We present an unbiased census of deeply embedded protostars in Perseus, Serpens, and Ophiuchus, assembled  by combining large-scale 1.1~mm Bolocam continuum  and \\textit{Spitzer} Legacy surveys.  We identify protostellar candidates based on  their mid-infrared properties, correlate their positions with 1.1~mm core positions from \\citet{enoch06}, \\citet{young06}, and \\citet{enoch07},  and construct well-sampled SEDs  using our extensive wavelength coverage ($\\lambda=1.25-1100~ \\micron$). Source classification based on the bolometric temperature yields a total of 39 Class~0 and 89 Class~I sources in the three cloud sample.  We compare to protostellar evolutionary models using the bolometric temperature-luminosity diagram, finding a population of low luminosity Class~I sources that are inconsistent with constant or monotonically decreasing mass accretion rates.  This result argues strongly  for episodic accretion during the Class~I phase,  with more than 50\\% of sources in a ``sub-Shu''  ($dM/dt < 10^{-6}$ \\msun\\ yr$^{-1}$) accretion state. Average spectra are compared to protostellar radiative transfer models, which match the observed spectra fairly well in Stage~0, but predict too much near-IR and too little mid-IR flux in Stage~I. Finally, the relative number of Class~0 and Class~I sources are used to estimate the lifetime of the Class~0 phase;  the three cloud average yields a Class~0 lifetime of $1.7\\pm0.3\\times10^5$ yr,  ruling out an extremely rapid early accretion phase.  Correcting  photometry for extinction results in a somewhat shorter lifetime ($1.1\\times10^5$ yr).  In Ophiuchus, however, we find very few Class~0 sources ($N_{\\mathrm{Class~0}}/N_{\\mathrm{Class~I}} \\sim 0.1-0.2$), similar to previous studies of that cloud.   The observations suggest a consistent picture of nearly constant \\textit{average} accretion rate through the entire embedded phase, with accretion becoming episodic by at least the Class I stage, and possibly earlier. ",
        "introduction": "The problem of how low mass stars like the sun form has been studied extensively over the last few decades.   Compared to more evolved protostars and pre-main sequence objects, however, the earliest stages of the star formation process, from the formation of dense cores through the main mass accretion phase, are relatively poorly understood.  This lack of information is due in large part to the difficulty of observing young, deeply embedded sources, which are shrouded within dense protostellar envelopes and only observable via reprocessed emission at mid-infrared to millimeter wavelengths. Furthermore, most previous observations of deeply embedded objects have naturally focused on a small number of very bright or well known sources, due to the sensitivity and resolution limitations of long-wavelength surveys.  Understanding the formation of typical stars requires complete samples of young objects, over molecular cloud scales.  Currently, the details of the early evolution of protostellar sources are extremely uncertain, including mass accretion rates during the Class 0 and Class I phases.  In addition, measurements of the timescales associated with the earliest stages vary considerably, ranging from $10^5$ to $10^7$~yr for prestellar cores \\citep{wt07} and from $10^4$ to a few $\\times10^5$~yr for Class~0 \\citep{am94,vrc02}. In fact, the association of Class~0 and Class~I with distinct evolutionary stages is still a matter of debate \\citep[e.g.,][]{rayj}. Large surveys at mid-infrared to millimeter wavelengths, where the SEDs of embedded sources peak, are essential for understanding how protostars evolve through their earliest  stages.  In addition to providing complete samples of young sources, large surveys are also important for characterizing variations in the star formation process with environment.  With a few notable recent exceptions \\citep{jorg07,hatch07}, previous samples of very young protostars have typically been compiled from many different surveys, and suffered from systematics, unquantified environmental effects, and small number statistics. We recently completed large-scale 1.1~mm continuum surveys of Perseus, Serpens, and Ophiuchus with Bolocam at the Caltech Submillimeter Observatory (CSO).  Maps have a resolution of $31\\arcsec$ and cover 7.5 deg$^2$ in Perseus (140~pc$^2$ at our adopted cloud distance of $d=250$~pc), 10.8 deg$^2$ in Ophiuchus (50~pc$^2$ at $d=125$~pc), and 1.5 deg$^2$ in Serpens (30~pc$^2$ at $d=260$~pc) \\citep[][hereafter Papers I, II, and III, respectively]{enoch06,young06,enoch07}. These Bolocam surveys complement large \\textit{Spitzer} Space Telescope IRAC and MIPS maps of the same clouds from the ``From Molecular Cores to Planet-forming Disks'' \\textit{Spitzer} Legacy program (``Cores to Disks'' or c2d; \\citealt{evans03}).  Millimeter emission traces the dust in dense starless cores and protostellar envelopes, and provides a measure of  core and envelope properties, including sizes, masses, and spatial distribution. \\textit{Spitzer} IRAC and MIPS observations are complementary in that they provide information about the properties of any young protostars embedded within dense cores. Combining these data enables us to assemble a mass limited, unbiased census of the youngest star-forming objects in three different environments, including prestellar cores, Class~0, and Class~I protostars.   In Paper III we looked at how the global cloud environment influences the properties of star-forming cores \\citep{enoch07}.  In a companion paper to this work \\citep{enoch08}, we examine the properties of prestellar and protostellar cores in Perseus, Serpens, and Ophiuchus, focusing on the prestellar core mass distribution and the lifetime of prestellar cores.   A similar analysis comparing the c2d \\textit{Spitzer} data to large SCUBA maps of Perseus and Ophiuchus has recently been carried out by \\citet{jorg07,jorg08}, focusing on the difference between starless and protostellar cores, as well as cloud properties such as the star formation efficiency and how it varies with spatial clustering.  We follow \\citet{dif07} in defining millimeter cores containing a compact luminous internal source (i.e., an embedded protostar) as ``protostellar cores'', regardless of whether the final object will be stellar or sub-stellar in nature.  We use ``starless cores'' to refer to dense cores without an internal luminosity source, and ``prestellar cores''  as starless cores that are likely to be gravitationally bound \\citep[see][]{enoch08}.  In this work we exploit the combined power of millimeter and mid- to far-infrared observations to study the evolution of young protostars embedded within the protostellar cores. Extensive wavelength coverage from $\\lambda=1.25-1100~ \\micron$ allows us to trace the evolution of protostellar sources in their main mass  accretion phase, from formation through the end of the embedded phase.  We follow \\citet{rob06} and \\citet{crapsi08} in using ``Stage'' (e.g., Stage 0, I, II, III) to refer to a source's true physical nature, regardless of its observed properties, while the corresponding ``Class'' refers to the observational classification, typically based on the near- to mid-infrared slope $\\alpha_{IR}$ \\citep{als87,greene94} or on the bolometric temperature \\citep{ml93}.  We assume that Stage~0 and Stage~I refer to an evolutionary sequence of embedded protostars with $M_{*} < M_{env}$ and $M_* > M_{env}$, respectively \\citep[e.g.][]{andre94}, where $M_{*}$ is the protostar mass and $M_{env}$ the envelope mass.  Similarly, we assume Stage~II refers to pre-main sequence stars with very little remaining envelope ($M_{env} < 0.1 M_{\\sun}$; \\citealt{crapsi08}). Our adopted definitions of various Classes and Stages are summarized in Table~\\ref{deftab}. Ideally, Class~I would directly correspond to Stage~I, etc., but as shown by \\citet{rob06} and \\citet{crapsi08}, geometric effects can cause, for example, Stage~II sources to be classified as Class~I.  In \\S\\ref{datasec} we describe the 1.1~mm and \\textit{Spitzer} infrared (IR) data, including how protostellar candidates are identified  and their association with 1.1~mm cores.  We calculate bolometric luminosities and temperatures, envelope masses, and discuss completeness (\\S\\ref{embprops}).  Protostellar classification methods are compared in \\S\\ref{classsec}.  Spectral characteristics of Class~0, I and II sources detected at 1.1~mm are discussed in \\S\\ref{charsec}, including selected individual sources (\\S\\ref{indsec}). In \\S\\ref{otherclasssec} we explore alternative classification schemes made possible by our large sample, and compare the average observed spectra to protostellar models in \\S\\ref{modsec}.  Sources are placed on a bolometric temperature-luminosity diagram for comparison to protostellar evolutionary models and to study their luminosity evolution, mass accretion rates, and envelope evolution (\\S\\ref{evolsec}). Finally, in \\S\\ref{lifesec} we calculate lifetimes for the Class~0 phase. ",
        "conclusions": "\\label{discsect}  Utilizing large-scale 1.1~mm surveys \\citep{enoch06,young06,enoch07} together with \\textit{Spitzer} IRAC and MIPS maps from the c2d Legacy program \\citep{evans03}, we have constructed an unbiased census of deeply embedded protostars in the Perseus, Serpens, and Ophiuchus molecular clouds. Our sample includes a total of 39 Class~0 sources and 89 Class~I sources, with approximate detection limits of $M_{env} \\gtrsim 0.1~ \\msun$ and $\\lbol \\gtrsim 0.05 \\lsun$ for the envelope mass and bolometric luminosity, respectively. We also detect a few Class~II and Herbig Ae/Be candidates at 1.1~mm, most likely evidence of fairly massive proto-planetary disks.  Bolometric luminosities, temperatures, and envelope masses are calculated for the candidate Class~0 and Class~I sources in each cloud.  We compare protostellar classification methods, concluding that, for deeply embedded sources, the bolometric temperature \\tbol\\ is a better measure of evolutionary state than the near- to mid-infrared spectral index, $\\alpha_{IR}$.  We also explore classifying sources into ``early Class 0'', ``late Class I'', etc., based on dividing them into narrower \\tbol\\ bins. Average observed spectra in these bins are compared to model predictions from \\citet{whit03} for ``early Stage 0'', ``late Stage 0'', etc. In a broad sense the Stage~0 and Stage~II models match the observed spectra fairly well.  The agreement with both Stage~I models is quite poor, however, as the models severely over-predict the near-IR and under-predict the mid-IR flux at this stage, displaying a double-peaked SED that is rarely observed.   Observed source properties are compared to protostellar evolutionary models using the bolometric temperature-luminosity ($L_{bol}-T_{bol}$) diagram, the protostellar equivalent of the H-R diagram \\citep{myer98}.  Neither models with a constant mass accretion rate \\citep{ye05}, nor those with an exponentially decreasing rate \\citep{myer98} fit the observed sources.  In particular, there is a large population of low luminosity Class~I sources that  aggravate the previously noted ``luminosity problem'' for embedded protostars \\citep[e.g.][]{ken90}.  We interpret this result as evidence for episodic accretion beginning at least by the Class~I  phase, and possibly earlier.  More than 50\\% of Class~I sources are inferred to have ``sub-Shu'' mass accretion rates ($\\dot{M} \\lesssim 10^{-6}$ \\msun\\ yr$^{-1}$, corresponding to $\\lbol \\lesssim 1 \\lsun$), and approximately 20\\% have $\\dot{M} \\lesssim 10^{-7}$ \\msun\\ yr$^{-1}$.  To build up of order a solar mass in $5.4\\times 10^5$~yr, such sources must also have periods of ``super-Shu'' accretion ($\\dot{M} \\gtrsim 10^{-5}$ \\lsun yr$^{-1}$). Few very high luminosity sources are observed (5\\%), suggesting that such rapid accretion periods must be short lived. An important caveat to this analysis is that we may sometimes underestimate  the luminosity due to missing 70 or 160~$\\micron$ photometry.  Finally, the relative number of Class~0 and Class~I sources are used to estimate the lifetime of the  Class~0 phases.  There are approximately half as many Class~0 as Class~I sources in the three cloud sample, implying an average Class~0 lifetime of $1.7\\pm0.3\\times10^5$~yr ($1.1\\times 10^5$ yr when approximate extinction corrections are applied).  This lifetime rules out drastic changes  in the mass accretion rate from Class 0 to Class I, particularly extremely rapid accretion in Class~0. In Ophiuchus the fraction of Class~0 sources is much smaller.  While this difference could be due to the lack of $160~ \\micron$ flux measurements in Ophiuchus, it may be that either the Class~0 phase is shorter in that cloud ($0.7\\times10^5$~yr), or that a burst of star formation is responsible for the large population of Class~I objects \\citep[e.g. as suggested by][]{vrc02}.  Altogether, the large variation in Class~I luminosities, similar mean luminosities for the Class~0 and Class~I phases, and not dramatically different numbers of Class~0 and Class~I sources (at least in Perseus and Serpens) suggests a consistent picture of nearly constant \\textit{average} mass accretion rate through the entire embedded phase (Stage~0 and Stage~I), with highly variable episodic accretion turning on by at least early Stage~I, and possibly sooner.  Understanding embedded protostellar structure and evolution well enough to reproduce the observations with detailed models of spectra and evolution presents an ongoing challenge."
    },
    "0809/0809.4538_arXiv.txt": {
        "abstract": "We investigate the damping of the baryon acoustic oscillations (BAO) signature in the matter power spectrum due to the quasi-nonlinear clustering of density perturbations. On the basis of the third order perturbation theory, we construct a fitting formula of the damping in an analytic way. This demonstrates that the damping is closely related with the growth factor and the amplitude of the matter power spectrum. Then, we investigate the feasibility of constraining the growth factor through a measurement of the damping of the BAO signature. An extension of our formula including higher order corrections of density perturbations is also discussed. ",
        "introduction": "The baryon acoustic oscillations (BAO) imprinted in the galaxy power spectrum have recently attracted remarkable attention, as {a useful probe} for exploring the origin of dark energy \\cite{DETF,ESA}. The BAO signature has been clearly detected in the 2dFGRS and the SDSS galaxy samples \\cite{Eisenstein,Huetsi,PercivalI,PercivalII,Tegmark}. The feasibility of constraining the equation of state parameter of dark energy $w$ is demonstrated \\cite{PercivalIII,Okumura,HuetsiII}, where $w$ is defined by $w=p_{de}/\\rho_{de}$, where $p_{de}$ and $\\rho_{de}$ are the pressure and the energy density of the dark energy component, respectively. Furthermore, a lot of BAO survey projects are under progress, or planned \\cite{BOSS,FMOS,WFMOS,LSST,SKA,SPACE}.  Though these BAO surveys will precisely measure the galaxy power spectra, but we also need precise theoretical templates, in order to obtain a useful cosmological constraint from observational data. In practice, we need to solve uncertainties about the gravitational nonlinear clustering of the density perturbations, the redshift-space distortions, and the galaxy clustering bias, because these uncertainties might yield some systematic effects in the BAO signature. Then, a lot of works related to these topics have been done recently, which are very important and challenging for future research of dark energy \\cite{Seo,SeoII,SeoIII,SeoIV,Angulo,AnguloII,Shoji,Takahashi,Wang,McDonald,Nishimichi}.        In the present paper, we focus on the damping of the BAO signature in the matter power spectrum due to the nonlinear gravitational clustering. We investigate this damping in a semi-analytic way on the basis of the third order perturbation theory of density fluctuations. The nonlinear clustering is a consequence of mode-couplings of the density fluctuations and the peculiar velocity divergence in Fourier space. Especially, we demonstrate how the damping of the BAO signature can be expressed with/without $P_{22}(k)$ and $P_{13}(k)$, which describe the mode-coupling (see sections 2 and 3). Then, we develop a simple fitting formula relevant to the damping, which demonstrates the fact that the damping of the BAO signature is closely related with the growth factor and the amplitude of the matter power spectrum. This suggests that a measurement of the damping of the BAO signature might be a probe of the growth factor of the density fluctuations, though we need to further investigate if the effects of the redshift-space distortion or the galaxy clustering bias are influential to the damping.  This paper is organized as follows; In section 2, we briefly review the third order perturbation theory, and derive the second order power spectrum. In this formalism, the nonlinear gravitational clustering is described by the coupling of the Fourier modes of the density fluctuations and the peculiar velocity divergences. In section 3, we investigate the damping of the BAO signature due to the coupling of the Fourier modes in an analytic way. And a fitting formula of the damping, applicable in weakly nonlinear regime, is developed. The formula is compared with a result of $N$-body simulation. Extensions of the formula are also discussed. In section 4, with the use of the fitting formula, we demonstrate a future feasibility of constraining the growth factor through a measurement of the damping of the BAO signature. The last section is devoted to a summary and conclusions. Throughout this paper, we use the unit in which the velocity of light equals 1, and adopt the Hubble parameter $H_0=100h {\\rm km/s/Mpc}$ with $h=0.7$, unless explicitly stated. ",
        "conclusions": "In the present paper, we examined the effect of the nonlinear gravitational clustering on the BAO signature in the matter power spectrum. In particular, we focused on the damping of the BAO signature in the quasi-nonlinear regime. Our approach is based on the third order perturbation theory of the matter fluctuations, which enables us to investigate the damping in an analytic way. We found a simple analytic expression that describes the damping of the BAO signature, which clarifies what the important factor is for the damping. We showed that the leading correction for the damping is in proportion to the combination of $(\\sigma_8 D_1(z))^2$. On the basis of the result, we constructed a fitting formula for the correction of the damping of the BAO signature in the weakly nonlinear regime, which is expressed as a function of $k$, $\\Omega_mh^2$, $\\Omega_bh^2$ and $n_s$. This fitting formula reproduces the damping of the BAO signature of the second order power spectrum within the standard perturbation theory at $10\\%$ level for $k\\simlt0.2~{h}{\\rm Mpc}^{-1}$ until $z\\sim 1$, though the formula is only guaranteed in a narrow range of the parameters $\\Omega_mh^2$, $\\Omega_bh^2$ and $n_s$.  We also discussed a possible extension of our formula to elaborate higher order nonlinear corrections using a technique of resumming infinite series of higher order perturbations on the basis of the Lagrangian perturbation theory. This extended formula reproduces the damping of the BAO signature of the second order power spectrum based on the Lagrangian perturbation theory at $10\\%$ level for $k\\simlt0.2~{h}{\\rm Mpc}^{-1}$ until $z\\sim0$. This formula was compared with a result of $N$-body simulation, which showed the validity of the extended formula.  A measurement of the damping of the BAO signature might be useful as a probe of the growth factor of the density fluctuations. As a first step to investigate such the possibility, we assessed the feasibility of constraining $\\sigma_8 D_1(z)$ by measuring the damping of the BAO in the power spectrum. For a useful constraint, we need a very wide survey area of the sky. For a definite conclusion, however, we must include the effect of the redshift-space distortions and the galaxy clustering bias on the damping of the BAO signature. Thus, more sophisticated formula including the effect of the redshift-space distortions and the galaxy clustering bias at the same time is required. If this effect on the BAO damping was well understood, it can be useful in measuring the growth factor. We plan to revisit this issue in future work.   \\ack We thank A. Taruya, T. Matsubara, Y. Suto and G. H$\\ddot{{\\rm u}}$tsi for useful discussions and comments. We are also grateful to anonymous referee for useful comments which helped to improve the earlier version of the manuscript. T. N. is supported by a Grand-in-Aid for Japan Society for the Promotion of Science (JSPS) Fellows (DC1: 19-7066). This work is supported by Grant-in-Aid for Scientific research of Japanese Ministry of Education, Culture, Sports and Technology (Nos.~18540277, 19035007)."
    },
    "0809/0809.3404_arXiv.txt": {
        "abstract": "We study non-linear structure formation in high-resolution simulations of Early Dark Energy (EDE) cosmologies and compare their evolution with the standard $\\Lambda$CDM model. In Early Dark Energy models, the impact on structure formation is expected to be particularly strong because of the presence of a non-negligible dark energy component even at very high redshift, unlike in standard models that behave like matter-dominated universes at early times.  In fact, extensions of the spherical top-hat collapse model predict that the virial overdensity and linear threshold density for collapse should be modified in EDE model, yielding significant modifications in the expected halo mass function. Here we present numerical simulations that directly test these expectations.  Interestingly, we find that the Sheth \\& Tormen formalism for estimating the abundance of dark matter halos continues to work very well in its standard form for the Early Dark Energy cosmologies, contrary to analytic predictions. The residuals are even slightly smaller than for $\\Lambda$CDM. We also study the virial relationship between mass and dark matter velocity dispersion in different dark energy cosmologies, finding excellent agreement with the normalization for $\\Lambda$CDM as calibrated by \\citet{Evrard2008}.  The earlier growth of structure in EDE models relative to $\\Lambda$CDM produces large differences in the mass functions at high redshift. This could be measured directly by counting groups as a function of the line-of-sight velocity dispersion, skirting the ambiguous problem of assigning a mass to the halo. Using dark matter substructures as a proxy for member galaxies, we demonstrate that even with 3-5 members sufficiently accurate measurements of the halo velocity dispersion function are possible. Finally, we determine the concentration-mass relationship for our EDE cosmologies. Consistent with the earlier formation time, the EDE halos show higher concentrations at a given halo mass.  We find that the magnitude of the difference in concentration is well described by the prescription of \\citet{ENS2001} for estimating halo concentrations. ",
        "introduction": "\\label{sect:intro}   Arguably the most surprising result of modern cosmology is that all matter (including both atoms and non-baryonic dark matter) accounts for only a quarter of the total energy density of the Universe today, while the rest is contributed by a {\\em dark energy} field.  In 1999, observations of type Ia supernovae by the Supernovae Cosmology Project \\citep{Riess1999,Riess2004} and the relative accurate measurements of the distances to this objects \\citep{Perlmutter1999,Kowalski2008} demonstrated that the expansion of the Universe is accelerated today; there hence exists a mysterious force that acts against the pull of gravity. Nowadays, the inference that this is caused by dark energy can be made with significant confidence, as the observational evidence has further firmed up.  In fact, we have good reason to believe that we live in a flat universe with an upper limit of $\\Omega_m \\leq 0.3$ for the matter density today, based on cosmic microwave background measurements and a host of other observational probes \\citep[e.g.]{WMAP5}.  These observations yield a consistent picture, the so-called concordance cosmology, and are in agreement with predictions of the inflationary theory.   The physical origin of dark energy is however unknown and a major puzzle for theoretical physics. A nagging outstanding problem is that most quantum field theories predict a huge cosmological constant from the energy of the quantum vacuum, up to 120 orders of magnitude too large. There is hence no simple natural explanation for dark energy, and one has to be content with phenomenological models at this point. Two proposed forms of dark energy are the cosmological constant, a constant energy density filling space homogeneously, and scalar fields such as quintessence. In particular, `tracking quintessence' models attempt to alleviate the coincidence problem of the cosmological constant model. More exotic models where the dark energy couples to matter fields or can cluster itself have also been proposed.  In light of the many theoretical possibilities, the hope is that future observational constraints on dark energy will enable progress in the understanding of this puzzling phenomenon. This requires the exploitation of the subtle influence of dark energy on structure formation, both on linear and non-linear scales. As the expected effects are generally small for many of the viable dark energy scenarios, it is crucial to be able to calculate structure formation in dark energy cosmologies with sufficient precision to tell the different models apart, and to be able to correctly interpret observational data. For example, in order to use the abundance of clusters of galaxies at different epochs to measure the expansion history of the universe, one needs to reliably know how the cluster mass function evolves with time in different dark energy cosmologies. Numerical simulations are the most accurate tool available to obtain the needed theoretical predictions, and they are also crucial for testing the results of more simplified analytic calculations.  In this study, we carry out such non-linear simulations for a particular class of dark energy cosmologies, so-called Early Dark Energy (EDE) models where dark energy might constitute an observable fraction of the total energy density of our Universe at the time of matter radiation equality or even big-bang nucleosynthesis. While in the cosmological constant scenario, the fraction in dark energy is negligible at high redshift, in such models the energy fraction is a few per cent during recombination and structure formation, which introduces interesting effects due to dark energy already at high redshift. In particular, for an equal amplitude of clustering today, we expect structures to form earlier in such cosmologies than in $\\Lambda$CDM. This could be useful to alleviate the tension between a low $\\sigma_8$ normalization suggested by current observational constraints from the CMB on one hand, and the observations of relatively early reionization and the existence of a population of massive halos present already at high redshift on the other hand.  Recently, \\citet{Bartelmann2006} studied two particular EDE models, evaluating the primary quantities relevant for structure formation, such as the linear growth factor of density perturbation, the critical density for spherical collapse and the overdensity at virialization, and finally the halo mass function.  In the two models analyzed, they found that the effect of EDE on the geometry of the Universe is only moderate, for example, distance measures can be reduced by $~8\\%$. Assuming the same expansion rate today, such models are younger compared to $\\Lambda$CDM. At early times, the age of the universe should differ by approximately $5-10 \\% $.  However, when \\citet{Bartelmann2006} repeated the calculation of the spherical collapse model in the EDE cosmology, a few nontrivial modifications appeared. The evolution of a homogeneous, spherical overdensity can be traced utilizing both the virial theorem and the energy conservation between the collapse and the turn around time \\citep[see also][]{Lacey1993,Wang1998}. \\citet{Bartelmann2006} obtained the value of the virial overdensity as a function of the collapse redshift, translating the effect of the early dark energy in an extra contribution to the potential energy at early times.  They found that the virial overdensity should be slightly enlarged by EDE, because a faster expansion of the universe means that, by the time a perturbation has turned around and collapsed to its final radius, a larger density contrast has been produced.  However, at the same time they found that the linearly extrapolated density contrast corresponding to the collapsed object should be significantly reduced.  These two results based on analytic expectations have a pronounced influence on the predicted mass function of dark matter halos.  In EDE models, the cluster population expected from the Press-Schechter or Sheth-Tormen formalism grows considerably relative to $\\Lambda$CDM, as a result of the lowered value of the critical linear density contrast $\\delta_c$ for collapse. This effect can be compensated for by lowering the normalization parameter $\\sigma_8$ in order to obtain the same abundance of clusters today. In this case, one would however still expect a higher cluster abundance in EDE at high redshift, due to the earlier growth of structure in this model.   An open question is whether the EDE really participates in the virialization process in the way assumed in the analytic modeling.  Similarly, it is not clear whether the excursion set formalism of Sheth \\& Tormen yields an equally accurate description of the non-linear mass function of halos in EDE cosmologies as in $\\Lambda$CDM.  Because accurate theoretical predictions for the halo mass function are a critical ingredient for constraining cosmological parameters (in particular $ \\Omega_m$ and $ \\Omega_{\\Lambda}$) as well as models of galaxy formation, it is important to test these predictions for the EDE cosmology in detail with numerical N-body simulations. In particular, we want to probe whether the fraction $f$ of matter ending up in objects larger than a given mass $M$ at some redshift $z$ can be found by {\\em only} looking at the properties of the linearly evolved density field at this epoch, using the ordinary ST formalism, or whether there is some dependence on redshift, power spectrum or dark energy parameters, as suggested by \\citet{Bartelmann2006}.   A further interest in EDE cosmologies stems from the fact that for a given $\\sigma_8$, the EDE models predict a substantially slower evolution of the halo population than in the $ \\Lambda$CDM model. This could explain the higher normalization cosmology expected from cluster studies relative to analysis of the CMB. The value of $\\sigma_8$, for a given cosmology, provides also a measure of the expected biasing parameter that relates the galaxy and the mass distribution. The early dark energy cosmologies could hence reduce the current mild tension between cluster data and the CMB observations. We note that halos in cosmologies with EDE are also expected to be more concentrated than in $\\Lambda$CDM; because the density of the Universe was greater at early times, objects that virialized at high redshift are more compact than those that virialized more recently.   Previous numerical simulations of a quintessence component with a changing equation of state (EOS) explored two particular potentials: SUGRA and Ratra Peebles (RP), which differ because RP has a more smoothly decreasing $w$ and consequently a very different evolution in the past. Both \\cite{Linder2003} and \\cite{Klypin2003} analyzed the influence of the dark energy on the halo mass function in order to extrapolate the abundance of structure at different epochs and to compare it with existent theoretical models. They used different numerical codes: the publicly available code {\\small GADGET}, in the first project, and the Adaptive Refinement Tree code \\citep{ARTcode}, in the second. They concluded that the best way to understand which dark energy Universe fit the observations best is to look at the growth history of halos and the evolution of their properties with time. \\citet{Dolag2004} focused on the modification of the concentration parameter with mass and redshift, for the same cosmologies, based on high resolution simulations of a sample of massive halos. A limited number of numerical studies also considered the possibility of a coupling of the dark energy field with dark matter \\citep{Mainini2003,Maccio}.   In this paper, we carry out several high resolution simulations of dark energy cosmologies in order to accurately measure the quantitative impact of early dark energy on abundance and structure of dark matter halos. To this end, we in particular measure halo mass functions and evaluate the agreement/disagreement with different analytic fitting functions. We also test how well the growth of the mass function can be tracked with dynamical measure based on the velocity dispersion of dark matter substructures, which can serve as a proxy for the directly measurable line-of-sight motion of galaxies or line widths in observations, and gets around the usual ambiguities arising from different possible mass definitions for halos. Finally, we also present measurements of halo concentrations, and of the relation between dark matter velocity dispersion and halo mass. While finalizing this paper, \\citet{Francis2008} submitted a preprint which also studies numerical simulations of EDE cosmologies. Their work provides a different analysis and is complementary to our study, but it reaches similar basic conclusions about the halo mass function.    This paper is organized as follows. After a brief introduction to the Early Dark Energy models in Section 2, we present the simulations and also give details on our numerical methods in Section~3. In Section~4, we study the mass function of halos for the different cosmologies, and as a function of redshift. Then, in Section~5 we investigate the properties of halos by studying the virial relation between mass and dark matter velocity dispersion, as well as the mass--concentration relationship. In Section~6, we consider the velocity distribution function and prospects for measuring it in observations. Finally, we discuss our results and present our conclusions in Section 7. ",
        "conclusions": "In this study we have analyzed non-linear structure formation in a particular class of dark energy cosmologies, so called {\\em early dark energy} models where the contribution of dark energy to the total energy density of the universe does not vanish even at high redshift, unlike in models with a cosmological constant and many other simple quintessence scenarios. Our particular interest has been to test whether analytic predictions for the halo mass function still reliably work in such cosmologies. As the evolution of the mass function is one of the most sensitive probes available for dark energy, this is of crucial importance for the interpretation of future large galaxy cluster surveys at high redshift. The mass function of EDE models is also especially interesting because analytic theory based on extensions of the spherical collapse model predicts that the mass function should be significantly modified \\citep{Bartelmann2006}, and in particular be characterized by a different value of the linear overdensity $\\delta_c$ for collapse, as well as a slightly modified virial overdensity.  We have carried out a set of high-resolution N-body simulations of two EDE models, and compared them with a standard $\\Lambda$CDM cosmology, and a model with a constant equation of state equal to $w=-0.6$.  Interestingly, we find that the universality of the standard Sheth \\& Tormen formalism for estimating the halo mass function also extends to the EDE models, at least at the $\\simlt 15 \\%$ accuracy level that is reached also for the ordinary $\\Lambda$CDM model. This means that we have found good agreement of the standard ST estimate of the abundance of DM halos with our numerical results for the EDE cosmologies, {\\em without} modification of the assumed virial overdensity and the linear density contrast threshold.  This disagrees with the theoretical suggestions based on the generalized top-hat collapse. In fact, if we instead use the latter as theoretical prediction of the halo mass function, the deviations between the prediction and the numerical results become significantly larger.  We hence conclude that the constant standard value for the linearly extrapolated density contrast can be used also for an analysis of early dark energy cosmologies.  Very recently, similar results were also obtained by \\cite{Francis2008}, who studied the same problem in cosmological simulations with somewhat smaller mass resolution.  This results on the mass function appear to hold over the whole redshift range we studied, from $z=0$ to $z=3$. Since our simulations were normalized to the same $\\sigma_8$ today, their mass functions and power spectra agree very well today, but towards higher redshift there are significant differences, as expected due to the different histories of the linear growth factor in the different cosmologies. In general, structure in the EDE cosmologies has to form significantly earlier than in $\\Lambda$CDM to arrive at the same abundance today. For example, already by redshift $z=3$, the abundance of galaxy clusters of mass $M=5\\times 10^{12}\\,h^{-1}{\\rm M}_\\odot$ is higher in EDE1 by a factor of $\\sim 1.7$ relative to $\\Lambda$CDM.  The earlier formation of halos in EDE models is also directly reflected in the concentration of halos. While for a given $\\sigma_8$ we find the same abundance of DM halos, the different formation histories are still reflected in a subtle modification of the internal structure of halos, making EDE concentrations for all halo masses and redshifts considered slightly higher. The difference is however quite small, but it would, for example, lead to a higher rate of dark matter annihilation in halos.   Another relationship that appears to accurately hold equally well in $\\Lambda$CDM as in generalized dark energy cosmologies is the virial scaling between mass and dark matter velocity dispersion that \\citet{Evrard2008} has found. In fact, we find that their normalization of this relation is accurately reproduced by all of our simulations within the measurement uncertainties, independent of cosmology. This also suggests that possible differences in the virial overdensity of EDE halos must be very small, and that presumably the relationship between total Sunyaev-Zeldovich decrement and halo mass is unmodified as well.  We show that counting the number of halos as a function of the line-of-sight velocity dispersion (of subhalos or galaxies), both in simulations and observations, can probe the growth of structures with redshift, and so put powerful constraints on the equation of state parameters.  This goal can be achieved by just identifying and counting groups in galaxy survey data such as DEEP2, and by comparing them with high-resolution N-body simulations. Precision measurements with this technique will still require accurate calibrations to deal with complications such as a possible velocity bias or selection effects in observational surveys.  However, \\citet{Davis2005} suggest that the DEEP2 survey alone has the power to constrain $w$ to an accuracy of 20\\% using velocity dispersion data, which illustrates the promise of this technique. In combination with other independent data, such as X-ray temperature and SZ decrement data, the constraints could be improved to an accuracy of 5\\%, without the need to invoke a model for the ambiguous total mass of a halo.   Distinguishing a time-varying dark energy component from the cosmological constant is a major quest of the present theoretical and observational astronomy. One approach is to rely on classical cosmological tests of the Hubble diagram, e.g.~by pushing the supernova type Ia observations to much higher redshift. Another quite direct geometrical probe is the observation of baryonic acoustic oscillations in the matter distribution at different redshifts.  Finally, the linear and non-linear evolution of cosmic structures provides another opportunity to constrain dark energy. In this work we have used numerical N-body simulations to examine the difference in structure growth in early dark energy cosmologies. We have seen that such simulations are essential to test the predictions of more simplified analytic models, and to calibrate observational tests that try to constrain the properties of dark energy with the abundance and internal structure of dark matter halos. Our results show clearly that the effects due to dynamical dark energy tend to be quite subtle, and can only be cleanly distinguished from ordinary $\\Lambda$CDM in high accuracy simulations. This poses new challenges to improve the precision of future generations of simulations, and at the same time emphasizes the immense observational task to arrive at sufficiently precise data at high redshift to constrain the dark side of the universe with the required accuracy."
    },
    "0809/0809.0407_arXiv.txt": {
        "abstract": "We summarize observational results on the stellar population and star formation history of the Scorpius-Centaurus OB Association (Sco OB2), the nearest region of recent massive star formation.  It consists of three subgroups, Upper Scorpius (US), Upper Centaurus-Lupus (UCL), and Lower Centaurus-Crux (LCC) which have ages of about 5, 17, and $ 16$~Myr.  While the high- and intermediate mass association members have been studied for several decades, the low-mass population remained mainly unexplored until rather recently.  In Upper Scorpius, numerous studies, in particular large multi-object spectroscopic surveys, have recently revealed hundreds of low-mass association members, including dozens of brown dwarfs.  The investigation of a large representative sample of association members provided detailed information about the stellar population and the star formation history.  The empirical mass function could be established over the full stellar mass range from $0.1\\,M_\\odot$ up to $20\\,M_\\odot$, and was found to be consistent with recent determinations of the field initial mass function.  A narrow range of ages around 5~Myr was found for the low-mass stars, the same age as had previously (and independently) been derived for the high-mass members.  This supports earlier indications that the star formation process in US was triggered, and agrees with previous conjectures that the triggering event was a supernova- and wind-driven shock-wave originating from the nearby UCL group.  In the older UCL and LCC regions, large numbers of low-mass members have recently been identified among X-ray and proper-motion selected candidates.  In both subgroups, low-mass members have also been serendipitously discovered through investigations of X-ray sources in the vicinity of better known regions (primarily the Lupus and TW Hya associations). While both subgroups appear to have mean ages of $\\sim$16\\,Myr, they both show signs of having substructure. Their star-formation histories may be more complex than that of the younger, more compact US group.  Sco-Cen is an important ``astrophysics laboratory'' for detailed studies of recently formed stars. For example, the ages of the sub-groups of 5~Myr and $\\sim 16$~Myr are ideal for studying how circumstellar disks evolve.  While no more than a few percent of the Sco-Cen members appear to be accreting from a circumstellar disk, recent {\\it Spitzer} results suggest that at least $\\sim$35\\% still have cold, dusty, debris disks. ",
        "introduction": "}  OB associations were first recognized by \\citet{Blaauw46} and \\citet{Ambartsumian47} as extended moving groups of blue luminous stars.  They are defined as loose stellar systems (stellar mass density of $\\la 0.1\\,M_\\odot\\,{\\rm pc}^{-3}$) containing O- and/or early B-type stars \\citep[for a recent review see][]{Briceno06}. At such low densities, the associations are unstable against Galactic tidal forces, and therefore it follows from their definition that OB associations must be young ($\\la 30-50$~Myr) entities.  Most of their low-mass members are therefore still in their pre-main sequence (PMS) phase.  There are different models for the origin of OB associations. One possibility is that they start as initially dense clusters, which get unbound and expand quickly as soon as the massive stars expel the gas \\citep[see][]{Kroupa01}.  An alternative model assumes that OB associations originate from {\\em unbound} turbulent giant molecular clouds \\citep[see][]{Clark05}, i.e.~start already in a spatially extended configuration and thus form in a fundamentally different way than dense, gravitationally bound clusters. Many well investigated OB associations show remarkably small internal velocity dispersions (often $\\leq 1.5$~km/sec), which are in some cases much smaller than required to explain the large present-day size (typically tens of parsecs) by expansion over the age of the association. This excludes the expanding cluster model and provides strong support for an origin as an extended unbound cloud for these associations.  The considerable number of OB associations in the solar neighborhood \\citep[e.g.,][]{deZeeuw99} suggests that they account for a large, maybe the dominant, fraction of the total Galactic star formation.  A good knowledge of their stellar content is thus essential in order to understand the nature of the star formation process not only in OB associations but also on Galactic scales.  For many years star formation was supposed to be a bimodal process \\citep[e.g.][]{Larson86,Shu88} according to which high- and low-mass stars should form in totally different sites.  Although it has been long established that low-mass stars {\\em can} form alongside their high-mass siblings in nearby OB associations \\citep[e.g.~][]{Herbig62}, it is still not well known {\\em what quantities} of low-mass stars are produced in OB environments.  There have been many claims that high-mass star forming regions have a truncated initial mass function (IMF), i.e.~contain much smaller numbers of low-mass stars than expected from the field IMF \\citep[see, e.g.,][]{Slawson92,Leitherer98,Smith01,Stolte05}.  Possible explanations for such an effect are often based on the strong radiation and winds from the massive stars.  For example, increased radiative heating of molecular clouds may raise the Jeans mass; lower-mass cloud cores may be completely dispersed by photoevaporation before low-mass protostars can even begin to form; the radiative destruction of CO molecules should lead to a change in the equation of state of the cloud material, affecting the fragmentation processes and ultimately leading to the formation of a few massive stars rather than the ``normal'' IMF which is dominated by low-mass stars \\citep[e.g.][]{Li03}.  However, several well investigated massive star forming regions show {\\em no} evidence for an IMF cutoff \\citep[see, e.g.,][for the cases of NGC~3603, 30~Dor, and NGC~346, respectively]{Brandl1999,Brandner01,Sabbi08}, and notorious difficulties in IMF determinations of distant regions may easily lead to wrong conclusions about IMF variations \\citep[see, e.g., discussion in][]{Selman05, Zinnecker93}.  If the IMF in OB associations is not truncated and similar to the field IMF, it would follow that most $(\\ga 60\\%)$ of the total stellar mass is found in low-mass $(< 2\\,M_\\odot)$ stars.  This would then imply that most of the current Galactic star formation is taking place in OB associations \\citep[as initially suggested by][]{Miller78}, and the {\\em typical} environment for forming stars (and planets) would be close to massive stars and not in isolated regions like Taurus.  The presence of nearby massive stars affects the evolution of young stellar objects and their protoplanetary disks in OB environments. For example, photoevaporation by intense UV radiation can remove a considerable amount of circumstellar material around young stellar objects \\citep[e.g.,][]{Bally98,Richling98}, and these objects will therefore ultimately end up with smaller final masses than if they were located in isolated regions \\citep{Whitworth04}.  Although generally considered to be a threat for forming planetary systems, photoevaporation may actually help to form planets, as it seems to play an important role in the formation of planetesimals \\citep{Throop05}.  OB associations provide excellent targets to investigate these effects on the formation and evolution of low-mass stars (and their forming planetary systems) during ages between a few Myr and a few ten Myr. However, before one can study the low-mass members, one first has to find them.  Although this statement sounds trivial, the major obstacle on the way towards a reliable knowledge of the low-mass population in OB associations is the problem to identify the individual low-mass members.  Unlike stellar clusters, which can be easily recognized on the sky, OB associations are generally very inconspicuous: since they extend over huge areas in the sky (often several hundred square-degrees for the nearest examples), most stars in the area actually are unrelated foreground or background stars. Finding the association members among these field stars is often like finding needles in a haystack.  As the low-mass members are often too faint for proper-motion studies, the only reliable sign to discern between low-mass association members and unrelated, much older field stars is the strength of the 6708~\\AA\\, lithium line in the stellar spectrum: at ages of $\\leq 30$~Myr, the low-mass association members still have most of their initial Li preserved and show a strong Li line, whereas the older foreground and background field stars do not show this line since they have already depleted their primordial Li \\citep[e.g.,][]{D'Antona94}.  However, an accurate measurement of Li line width requires at least intermediate resolution spectroscopy, and thus the observational effort to identify the widespread population of PMS stars among the many thousands of field stars is huge.  Many empirical IMF determinations are therefore based on photometric data only; while this strategy rather easily provides a complete spatial coverage and allows one to work with large samples, photometry alone cannot give completely reliable membership information.  Most studies with spectroscopically identified member samples, on the other hand, include only very small fractions of the total stellar population and are strongly affected by small number statistics and the necessity of using large extrapolation factors.  Therefore, most studies dealing with OB associations have been restricted to estimating the number-ratio of low-mass versus high-mass members \\citep[e.g.][]{Walter94,Dolan02,Sherry04}.  During the last years, new and very powerful multiple-object spectrographs like 2dF at the Anglo-Australian Telescope \\citep[see][]{Lewis02} made large spectroscopic surveys for low-mass PMS members feasible.  In combination with the Hipparcos results, which allowed the complete identification of the high- and intermediate-mass $(\\ge 2\\,M_\\odot)$ stellar population in many nearby associations \\citep[see][]{deZeeuw99}, studies of the {\\em complete} stellar population in OB associations are now possible, enabling us to investigate in detail the spatial and temporal relationships between high- and low-mass members.  In this chapter, we review the state of knowledge regarding the stellar populations of the nearest OB association: Scorpius-Centaurus (Sco OB2).  We structure this chapter in the following manner. Section~2 gives a general, and historical, overview of the association and its subgroups. Sections 3, 4, and 5 discuss the surveys for the members of the three primary subgroups of Sco-Cen: Upper Scorpius, Upper Centaurus-Lupus, and Lower Centaurus-Crux, respectively. Readers uninterested in the individual surveys may want to skip ahead to Section~6, where we discuss the astrophysical ramifications, and interpretations, of studies of the Sco-Cen members. ",
        "conclusions": ""
    },
    "0809/0809.0893_arXiv.txt": {
        "abstract": "We analyze the spatial distributions of young stars in Taurus-Auriga and Upper Sco as determined from the two-point correlation function (i.e. the mean surface density of neighbors). The corresponding power-law fits allow us to determine the fractal dimensions of each association's spatial distribution, measure the stellar velocity dispersions, and distinguish between the bound binary population and chance alignments of members. We find that the fractal dimension of Taurus is $D\\sim$1.05, consistent with its filamentary structure. The fractal dimension of Upper Sco may be even shallower ($D\\sim$0.7), but this fit is uncertain due to the limited area and possible spatially-variable incompleteness. We also find that random stellar motions have erased all primordial structure on scales of $\\la$0.07$^o$ in Taurus and $\\la$1.7$^o$ in Upper Sco; given ages of $\\sim$1 Myr and $\\sim$5 Myr, the corresponding internal velocity dispersions are $\\sim$0.2 km s$^{-1}$ and $\\sim$1.0 km s$^{-1}$, respectively. Finally, we find that binaries can be distinguished from chance alignments at separations of $\\la$120\\arcsec\\, (17000 AU) in Taurus and $\\la$75\\arcsec\\, (11000 AU) in Upper Sco. The binary populations in these associations that we previously studied, spanning separations of 3-30\\arcsec, is dominated by binary systems. However, the few lowest-mass pairs ($M_{prim}$$\\la$0.3 $M_{\\sun}$) might be chance alignments. ",
        "introduction": "The spatial distribution of young stars is a powerful diagnostic of their formation and early evolution. Young stars trace the gas distribution from which they formed, so the large-scale structure of a young association retains these primordial features after the gas has been accreted or dispersed.  On intermediate scales, the absence of structure indicates the typical distance over which stars have randomly dispersed since their birth, and therefore the velocity dispersion for the association. Finally, the enhanced stellar density on small scales outlines the binary population, distinguishing bound binary systems from chance alignments between young stars. Some of these topics have been addressed in previous work on young star distributions (Gomez et al. 1993; Larson 1995; Simon 1997; Bate et al. 1998; Hartmann 2002), but the modern census of several key star-forming regions is more complete and extends to lower masses than a decade ago, so the analysis is worth revisiting.  The traditional tool for studying spatial distributions is the two-point correlation function (hereafter TPCF). The TPCF, $w(\\theta)$, is defined as the number of excess pairs of objects with a given separation $\\theta$ over the expected number for a random distribution (Peebles 1980). The TPCF is proportional to the mean surface density of neighbors, so it is often recast in terms of this more intuitive quantity: $\\Sigma(\\theta)=(N_*/A)[1+w(\\theta)]$, where $A$ is the survey area and $N_*$ is the total number of stars.  In this letter, we describe an updated relation for $\\Sigma(\\theta)$ in Taurus and present the first such analysis for Upper Sco, then we fit power laws for the different angular regimes. Finally, we interpret our results to address three questions: What is the primordial fractal dimension of star-forming regions, and how does it relate to their observed geometry? What is the primordial velocity dispersion suggested by each association's randomization? And what is a wide binary companion, and can it be distinguished from an unbound association member? ",
        "conclusions": ""
    },
    "0809/0809.2772_arXiv.txt": {
        "abstract": "Ground-based solar polarimetry has made great progress over the last decade. Nevertheless, polarimetry is still an afterthought in most telescope and instrument designs, and most polarimeters are designed based on experience and rules of thumb rather than using more formal systems engineering approaches as is common in standard optical design efforts. Here we present the first steps in creating a set of systems engineering approaches to the design of polarimeters that makes sure that the final telescope-instrument-polarimeter system is more than the sum of its parts. ",
        "introduction": "Systems engineering ensures that the total is more than the sum of its parts.  Systems engineering is essential for the successful design of polarimeters where apparently unrelated effects such as instrumental polarization due to oblique reflections in the telescope and detector non-linearity couple to generate substantial systematic errors (Keller 1996).  When designing a polarimeter, it is therefore crucial to optimize the entire system and not just the individual parts as has often been done in the past.  While systems engineering is common in standard optics, it has been largely absent in polarimetry because errors in polarization cannot easily be expressed as a scalar and optical elements such as polarizers and retarders have a major influence on the polarization. In the following we will present initial thoughts on establishing systems engineering approaches for polarimetry that will enable better polarimeter designs with predictable performance.  Systems engineering for polarimetry models and helps us understand the performance of polarimeter designs.  It must address at least the following issues: \\begin{itemize}\\itemsep=0pt \\item a definition of the polarimeter performance to quantitatively compare different polarimeter designs with ideal components; \\item a polarization error budget that can predict the performance based on known error statistics of real components; \\item methods to maximize the performance of a polarimeter design. \\end{itemize} While the first issue has been addressed in the literature, error budgets, a crucial systems engineering tool, have been lacking.  In the following sections we present some initial ideas on systems engineering approaches for polarimeter designs. ",
        "conclusions": ""
    },
    "0809/0809.4900_arXiv.txt": {
        "abstract": "We clearly formulate and study further a conjecture of effective field theory interaction with gravity in the cosmological context. The conjecture stems from the fact that the melding of quantum theory and gravity typically indicates the presence of an inherent UV cutoff. Taking note of the physical origin of this UV cutoff, that the background metric fluctuations does not allow QFT to operate with a better precision than the background space resolution, we conjecture that the converse statement might also be true. That is, an effective field theory could not perceive the background space with a better precision than it is allowed by its intrinsic UV scale. Some of the subtleties and cosmological implications of this conjecture are explored. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1652_arXiv.txt": {
        "abstract": "{Finite radius accretion disks are a strong candidate for launching astrophysical jets from their inner parts and disk-winds are considered as the basic component of such magnetically collimated outflows. Numerical simulations are usually employed to answer several open questions regarding the origin, stability and propagation of jets. The inherent uncertainties, however, of the various numerical codes, applied boundary conditions, grid resolution, etc., call for a parallel use of analytical methods as well, whenever they are available, as a tool to interpret and understand the outcome of the simulations. The only available analytical MHD solutions for describing disk-driven jets are those characterized by the symmetry of radial self-similarity. Those exact MHD solutions are used to guide the present numerical study of disk-winds.} {Radially self-similar MHD models, in general, have two geometrical shortcomings, a singularity at the jet axis and the non-existence of an intrinsic radial scale, i.e. the jets formally extend to radial infinity. Hence, numerical simulations are necessary to extend the analytical solutions towards the axis and impose a physical boundary at finite radial distance.} {We focus here on studying the effects of imposing an outer radius of the underlying accreting disk (and thus also of the outflow) on the topology, structure and variability of a radially self-similar analytical MHD solution. The initial condition consists of a hybrid of an unchanged and a scaled-down analytical solution, one for the jet and the other for its environment.} {In all studied cases, we find at the end steady two-component solutions. The boundary between both solutions is always shifted towards the solution with reduced quantities. Especially, the reduced thermal and magnetic pressures change the perpendicular force balance at the ``surface'' of the flow. In the models where the scaled-down analytical solution is outside the unchanged one, the inside solution converges to a solution with different parameters. In the models where the scaled-down analytical solution is inside the unchanged one, the whole two-component solution changes dramatically to support the flow from collapsing totally to the symmetry axis.} {It is thus concluded that truncated exact MHD disk-wind solutions which may describe observed jets associated with finite radius accretion disks, are topologically stable.} ",
        "introduction": "Astrophysical jets are ubiquitous, occurring in a variety of objects on very different size and mass scales. They can be produced by pre-main sequence stars in Young Stellar Objects, by post-AGB stars in pre-planetary (PPNs) and planetary nebulae (PNs), by white dwarfs (WDs) in supersoft X-ray sources and symbiotic stars (SySs), by neutron stars in X-ray Binaries, by stellar black holes in Black Hole X-ray Binaries and by supermassive black holes in the case of Active Galactic Nuclei. However, the formation mechanism of jets remains very poorly understood.  In all such cases, jets and disks seem to be inter-related. Disks provide the jets with the ejected plasma and magnetic fields and jets are possibly the most efficient means to remove excess angular momentum in the disk \\citep[e.g.][]{Fer07} making accretion possible in the first place. Observationally, such a correlation is now already well established \\citep[e.g.][in the case of YSOs]{HEG95}. Hence, our current understanding is that jets are fed by the material of an accretion disk surrounding the central object. The gravitational energy of the infalling material in the disk is converted into kinetic energy of the outflowing matter. Radiative launching can be excluded due to the involved small radiation fields which have usually luminosities of only a few percent of the necessary kinetic luminosities \\citep[e.g.][in the case of YSOs]{Lad85}. Furthermore, the kinetic luminosity of the outflow seems to be a large fraction of the rate at which energy is released by accretion. With such a high ejection efficiency it is natural to assume that jets are driven magnetically from an accretion disk; the magnetic model of a disk-wind seems to explain simultaneously acceleration, collimation as well as the observed high jet speeds \\citep{KoP00}.  In their seminal analytical work \\citet{BlP82} studied the magneto-centrifugal acceleration along magnetic field lines threading an accretion disk. They were the first to show the braking of matter in the azimuthal direction inside the disk and the outflow acceleration above the disk surface guided by the poloidal magnetic field components. Toroidal components of the magnetic field then collimate the flow. Numerous semi-analytic models extended the work of \\citet{BlP82} along the guidelines of radially self-similar solutions of the full magnetohydrodynamics (MHD) equations \\citep{VlT98}. \\citet{OgL98,OgL01} solved for the local vertical structure of a thin disk threaded by a poloidal magnetic field of dipolar symmetry. They showed that jet launching occurs only for a limited range of field strengths and a limited range of inclination angles.  The model of \\citet{BlP82}, however, has basically two serious limitations. First, the outflow speed at large distances does not cross the corresponding limiting characteristic, with the result that the terminal wind solution is not causally disconnected from the disk. Later, \\citet[][V00 hereafter]{VTS00} showed that a terminal wind solution can be constructed which is causally disconnected from the disk and hence any perturbation downstream of the superfast transition cannot affect the whole structure of the steady outflow. The second limitation of self-similar models in general is their geometrical nature. Singularities exists at the jet axis in radial self-similar models. Numerical simulations are necessary to extend the analytical solutions, as it has been already done in \\citet[][GVT06 hereafter]{GVT06} and \\citet[][M08 hereafter]{MTV08}.  Another geometrical limitation of radially self-similar models is the non-existence of an intrinsic scale. The jets formally extend to radial infinity. The aim of this paper is to investigate numerically, how imposing an outer radius of the jet, i.e. cutting off the analytical solution at arbitrary radii, affects the topology, structure and stability of a particular radial self-similar analytical solution and hence its ability to explain observations.  Several other numerical studies exist, which have focused on the magnetic launching of disk-winds. In most models a polytropic equilibrium accretion disk was regarded as a boundary condition \\citep[e.g.][]{UKR95, KLB99, KLB03, OCP03, NaM04, ALK05, ALK06, PRO06}. The magnetic feedback on the disk structure is therefore not calculated self-consistently. Only in recent years were the first simulations including the accretion disk self-consistently in the calculations of jet formation presented \\citep[e.g.][]{CaK02, CaK04, KMS04, ZFR07}. Numerical simulations of stellar winds have also been done by e.g. \\citet{MaP05a,MaP05b}. However, in none of these studies the disk has been truncated to study the effects of an intrinsic scale.  This paper is organized as follows: in Sec.~\\ref{sec_selfsim}, the analytical self-similar model underlying our numerical study is reviewed. The numerical setup is described in Sec.~\\ref{sec_num_models}. The study starts with several test simulations which are described in detail and whose results are presented in Sec.~\\ref{sec_testsims}. In Sec.~\\ref{sec_science_runs}, is described the parameter study of steady solutions and its results. We close with a summary of the results in the last section. ",
        "conclusions": "In this paper, we have studied the effects of imposing an outer radius of the underlying accreting disk, and thus also of the outflow, on the topology, structure and time-dependence of a well studied radially self-similar analytical solution (V00). We matched an unchanged and a scaled-down analytical solution of V00 and found in all these cases steady two-component solutions.  We showed that the boundary between both solutions is always shifted towards the solution with reduced quantities. Especially, the reduced thermal and magnetic pressure change the perpendicular force balance at the ``surface'' of the flow. In the models where the scaled-down analytical solution is {\\em outside} the unchanged one, the inside solution converges to another analytical solution with different parameters. In the models where the scaled-down analytical solution is {\\em inside} the unchanged one, the whole two-component solution changes dramatically to support the flow from collapsing totally to the symmetry axis.  A result of the modification of the analytical solution at the symmetry axis is the formation of a weak shock at the fast magnetosonic separatrix surface (FMSS) which shields the flow upstream of the FMSS from the imposed modifications close to the axis. A result of the truncation of the analytical solution along some poloidal field line is the formation of two compressive MHD shocks: a fast shock that travels quickly downstream outside of the computational domain and a slow shock which as it travels downstream with lower speed, it is locked precisely at the position of the lower computational z-boundary where the analytical slow surface meets this boundary.  Our truncated disk-wind solutions are stable for more than 50 $t_0$, i.e. several orbital periods at the truncation radius. These solutions may be relevant to describe observed jets, since the jet radii are too large in the untruncated analytic disk outflow solution, with respect to observed jet widths. We also provide all quantities at the outer $z$ boundary which can now be used as boundary conditions for jet propagation studies.  In the following paper II, we calculate emission maps corresponding to such truncated disk-models and compare them with observations."
    },
    "0809/0809.3657_arXiv.txt": {
        "abstract": "{ The atmospheres of substellar objects contain clouds of oxides, iron, silicates, and other refractory condensates.  Water clouds are expected in the coolest objects.  The opacity of these `dust' clouds strongly affects both the atmospheric temperature-pressure profile and the emergent flux.  Thus any attempt to model the spectra of these atmospheres must incorporate a cloud model.  However the diversity of cloud models in atmospheric simulations is large and it is not always clear how the underlying physics of the various models compare. Likewise the observational consequences of different modelling approaches can be masked by other model differences, making objective comparisons challenging.  In order to clarify the current state of the modelling approaches, this paper compares five different cloud models in two sets of tests. Test case 1 tests the dust cloud models for a prescribed L, L--T, and T-dwarf atmospheric (temperature T, pressure p, convective velocity $v_{\\rm conv}$)-structures. Test case 2 compares complete model atmosphere results for given (effective temperature T$_{\\rm eff}$, surface gravity $\\log g$).  All models agree on the global cloud structure but differ in opacity-relevant details like grain size, amount of dust, dust and gas-phase composition. These models can loosely be grouped into {\\it high-} and {\\it low-altitude cloud} models whereas the first appear generally redder in near-infrared colours then the later. Comparisons of synthetic photometric fluxes translate into an modelling uncertainty in apparent magnitudes for our L-dwarf (T-dwarf) test case of $0.25 \\lesssim \\Delta m \\lesssim 0.875$ ($0.1 \\lesssim \\Delta m \\lesssim 1.375$) taking into account the 2MASS, the UKIRT WFCAM, the Spitzer IRAC, and VLT VISIR filters with UKIRT WFCAM being the most challenging for the models. Future developments will need closer links with laboratory astrophysics, and a consistent treatment of the cloud chemistry and turbulence.} ",
        "introduction": "The atmospheres of L dwarfs are characterised by clouds, formed principally of silicate, oxide and iron grains, which shape their emergent spectra. Likewise the atmospheres of the early T dwarfs are distinguished by the progressive departure of cloud opacity. At even lower effective temperature, giant-gas planets are again covered in clouds and other chemical components become important. Any attempt to derive fundamental properties of these objects from their spectra hinges on an understanding of the chemistry and physics of clouds. Yet clouds are inherently difficult to model since they can feedback into the chemistry and the physics of the entire atmosphere.  Because of the complexity of this problem a number of independent groups have taken very different approaches to describe the cloud formation and the cloud properties of substellar objects as a function of gravity, effective temperature, and metallicity.  Here we make a first attempt to compare the quantitative predictions of these various approaches in order to better understand the models themselves as well as the uncertainty which remains in application of these models to real objects.   Atmospheric physics classically involves hydrodynamics, radiative and convective energy transport, and gas phase chemistry. Effects of magnetic fields are neglected. Ideally, the only free parameters are the effective temperature T$_{\\rm eff}$, the surface gravity $g$, radius R$_*$ or mass M$_*$, and element abundances $\\epsilon_{\\rm i}$. In order to solve such a coupled system of equations in a computationally reasonable time, assumptions like the hydrostatic equilibrium, mixing length theory and chemical equilibrium are made. Inside substellar atmospheres, chemical equilibrium of the gas phase is justified due to high collision rates between gas-phase constituents. Irradiation or atmospheric flows may invalidate this assumption in the upper atmospheric layers.  The validity of hydrostatic equilibrium and mixing length theory have been studied in comparison to large eddy simulations for M-type stars (Ludwig et al.\\ 2002, 2006) and we know from the direct observation of solar system giant planets at low gravities and effective temperatures that hydrostatic equilibrium is an appropriate down to very low pressures ($\\sim 1\\,\\rm \\mu bar$) in the atmosphere.  The assumption of hydrostatic equilibrium coupled with the mixing length theory is computationally extremely efficient and an accurate approximation in particular if one is aiming at synthetic spectrum calculations.  The striking difference of substellar atmosphere models compared to the classical stellar approach is the necessity to model the formation of clouds and their feedback onto the entire atmosphere.  New physics needed to be considered and different tribes emerged being inspired by AGB star dust formation (Helling et al. 2001a; Woitke \\& Helling 2003, 2004), by terrestrial cloud formation (Ackerman \\& Marley 2001, Cooper et al. 2003), by measurements for solar system planets (Rossow 1978, Marley et al. 1999), or driven by practical considerations (Tsuji et al. 1996 a,b; Allard et al. 2001). The first attempts on cloud modelling in brown dwarf atmospheres were undertaken by Lunine, Hubbard \\& Marley (1986) and Tsuji et al. (1996 a,b) who suggested the influence of clouds on the spectral appearance of brown dwarfs. See also Ackerman \\& Marley (2001) for a review and comparison of the earlier cloud literature.    The overall, phenomenological understanding of cloud formation in substellar objects has converged to the picture that dust (or condensates, see Table~\\ref{t:distionary}) forms at a certain height in the atmosphere where it acts as an efficient element sink leaving behind a depleted gas phase.  The departure of TiO and FeH spectral lines from the M to the L dwarfs testifies to this process.  The dust then settles gravitationally taking condensed elements with it. Convection and atmospheric mixing replenishes the condensing gas, resulting in a steady state.  The details of this picture, however, leave room for debate.  It is for example not clear where the dust precisely starts to form since this depends on the details of the model assumptions.  Current models generally employ one of two physical approaches to understand this process. In the first paradigm gas is mixed upward into higher altitudes. Dust then forms, falls down and meanwhile grows until it evaporates below the cloud base (Woitke \\& Helling 2003). The second paradigm imagines the opposite limiting case in which the gas is well mixed from the deep atmosphere only up to a cloud base. Grains and gas are transported above cloud base by mixing and grains fall down under the influence of gravity (Ackerman \\& Marley 2001, Allard et al. 2003).  These two branches rely on fundamentally different model assumptions: (i) the phase-non-equilibrium concept of kinetic dust formation (Woitke \\& Helling 2003, 2004; Helling \\& Woitke 2006, Helling, Woitke \\& Thi 2008), and (ii) the phase-equilibrium concept of thermal stability (Tsuji et al 1996b, Tsuji 2005, Allard et al. 2001, Ackerman \\& Marley 2001, Burrows et al. 2006; Cooper et al. 2003).  While (ii) represents the end-state which a dust forming system achieves for $t\\rightarrow\\infty$, (i) describes the kinetic process of the formation of cloud particles on finite timescales limited by mixing and rain-out. The models also differ in the choice of dust materials which are assumed (i) to form or (ii) to be present in the atmosphere. Both areas need serious attention and corresponding material properties should be obtained either experimentally (see discussion Sect.~\\ref{ss:lab}) or from ab initio calculation (e.g. Jeong et al. 2000, Patzer et al. 2005) which both are beyond the scope of this paper.  Given these model conceptions, a number of different model approaches have been developed to reproduce observed spectra (Tsuji et al 1996a, Tsuji 2005, Allard et al. 2001, Ackerman \\& Marley 2001, Burrows et al. 2006, Cooper et al. 2003, Dehn 2007, Helling et al. 2008a,b) or providing detailed information on the dust complex itself (Woitke \\& Helling 2003, 2004; Helling \\& Woitke 2006, Helling, Woitke \\& Thi 2008).   Driven by this diversity in the field, the aim of this paper is to provide information and to perform comparative studies of models that aim to describe the dust clouds in substellar atmospheres.  Kleb \\& Wood (2004) demonstrated that such component-based test studies are an essential part of scientific methods. As the number, $n$, of model components increases, the interactions amongst them grow as $n^2/2$. We therefore need to perform verifications on the components (here: cloud models) but also to use the method of {\\it manufactures solutions}\\footnote{{\\it Manufactures solutions} are tests with a known results. In our case, we manufacture the input quantities to a certain extent and compare the solutions (test case 1, Sect.\\ref{ss:tc1}).} (e.g. Kleb \\& Wood 2004) to verify that the entire system (here: model atmosphere) attains its theoretical order-of-accuracy properties.  This goes beyond what has been and could be provided in the literature so far.    \\begin{table} \\caption{Definition and units of quantities. The quantities plotted in Figs.\\ref{fig:InpStruc}--\\ref{fig:ColCol} are highlighted in {\\bf  boldface}.} \\begin{center} \\begin{tabular}{l|p{4.2cm}p{0.3cm}|l} \\hspace*{-0.2cm}{\\bf Quantity} & {\\bf Definition} & & {\\bf Unit}\\\\ \\hline $T$                      & {\\bf local gas temperature} & \\hspace*{-0.8cm}(Fig.~\\ref{fig:InpStruc},~\\ref{fig:P2CloudStruc}) & K\\\\ $p_{\\rm gas}$   & {\\bf local gas pressure }      & & dyn\\,cm$^{-2}$\\\\ $v_{\\rm conv}$ & {\\bf convective velocity  }     & & cm\\,s$^{-1}$\\\\ $\\epsilon_{\\rm i}$ & gas element abundance  && (i = H, He, ...)\\\\ $\\epsilon_{\\rm i}^0$ & deep element abundance  &&\\\\ $S$                           &supersaturation ratio \\\\ \\hline $\\rho_{\\rm d}/\\rho_{\\rm gas}$ & {\\bf dust to gas mass ratio} & \\hspace*{-0.5cm}(Fig.~\\ref{fig:RhoDGAmeanTC1}) &\\\\ $\\rho_{\\rm d}$                            & dust mass density & & g\\,cm$^{-3}$\\\\ & $=\\sum_{\\rm s} n_{\\rm s} M_{\\rm s}$\\,\\, for $a=$ const &\\\\ & $=\\sum_{\\rm s} \\int_a f(a) M_{\\rm s}(a)\\,da$ &\\\\ s                                                  & dust species & \\\\ & (e.g. Fe[s], Mg$_2$SiO$_4$[s], $\\ldots$) &\\\\ $n_{\\rm s}$                                & number of dust particles of kind $s$ & &  cm$^{-3}$\\\\ $a$                                             & grain size                      & & cm\\\\ $M_{\\rm s}$                             & mass of dust particle of kind s  & & g\\\\ \\hline $f(a)$                                         & grain size distribution: & & cm$^{-3}$cm$^{-1}$ \\\\ & number of dust grains $n$ per grain size $a$ and per gas volume & & \\\\ $\\langle a\\rangle$                   & {\\bf mean particle size} & \\hspace*{-0.5cm}(Fig.~\\ref{fig:RhoDGAmeanTC1}) & cm\\\\ & $\\displaystyle=\\frac{\\sum_{\\rm s} \\int_a f_{\\rm s}(a) a \\,da}{\\sum_{\\rm s} \\int_a f_{\\rm s}(a)\\,da}$ &\\\\ & for fixed $a=a_0$: & \\\\ &   $f(a) = \\delta (a - a_0) \\Rightarrow \\langle a\\rangle=a_0$ &\\\\ \\hline $V_{\\rm s}/V_{\\rm tot}$           &{\\bf volume fraction of dust kind  s} & \\hspace*{-0.5cm}(Fig.~\\ref{fig:VsTC1})& 1 \\\\ $V_{\\rm s}$                               & total dust volume of dust kind s & & cm$^3$ \\\\ $V_{\\rm tot}$                             & total dust volume                          & &cm$^3$\\\\ \\hline $F(\\lambda)$  & {\\bf surface flux} & \\hspace*{-.8cm}(Fig.~\\ref{fig:spectotal1800},~\\ref{fig:spectotal1000}) & erg\\,s$^{-1}$cm$^{-2}$\\AA$^{-1}$\\\\ $\\lambda$     &wavelength & & $\\mu$m\\\\ $\\log F_{\\rm c}$ & {\\bf photometric flux} & \\hspace*{-.6cm}(Fig.~\\ref{fig:spec05-3.0_1800}) &  erg\\,s$^{-1}$cm$^{-2}$\\AA$^{-1}$\\\\ & flux convolved with filter c & \\\\ & $\\displaystyle=\\frac{\\int_{\\lambda_1}^{\\lambda_2} F(\\lambda)\\, \\mbox{trans}_{\\rm c}(\\lambda)\\,d\\lambda}{\\int_{\\lambda_1}^{\\lambda_2} \\mbox{trans}_{\\rm c}(\\lambda)\\,d\\lambda}$ & \\\\ $\\log F_{c0}$    & reference photometric flux (Vega)\\\\ trans$_{\\rm c}(\\lambda)$ & function of a photometric filter c \\\\ & between $\\lambda_1$ and $\\lambda_2$ (see Table~\\ref{tab:photFlux})&\\\\ $m_1 - m_2$ & {\\bf colour} & \\hspace*{-.65cm}(Fig.~\\ref{fig:ColCol}) & \\\\ & $\\!\\!\\!\\displaystyle = 2.5 \\big(\\log \\frac{F_{\\rm c2}(\\Delta \\lambda_2)}{F_{c02}(\\Delta \\lambda_2)} - \\log \\frac{F_{\\rm c1}(\\Delta \\lambda_1)}{F_{c01}(\\Delta \\lambda_1)}\\big)$ &\\\\ $m$              & apparent magnitude\\\\ & $= -2.5 \\log\\int_{\\lambda_1}^{\\lambda_2} F(\\lambda) \\mbox{trans}_{\\rm c}\\,d\\lambda$\\\\  \\end{tabular} \\end{center} \\label{tab:definitions} \\end{table}%  Our paper begins with a summary of the various dust cloud models. We provide for the first time a comparative presentation of the different approaches concerning chemistry and dust modelling (Sect.~\\ref{s:cd}). Based on a workshop held in Leiden in October 2006 (\\footnote{http://www.lorentzcenter.nl/lc/web/2006/203/info.php3?wsid=203}$^,$\\footnote{http://phoenix.hs.uni-hamburg.de/BrownDwarfsToPlanets1/}), we present test cases where we first separate the components for chemistry and dust cloud modelling from the complete atmosphere problem (Sect.~\\ref{ss:tc1}). This allows us to judge the order-of-accuracy properties of the complete models with respect to chemistry and dust formation which both are essential ingredients for the solution of the radiative transfer problem.  Section~\\ref{ss:tc2} demonstrates the results for the complete substellar atmosphere problem, synthetic photometric fluxes, and colours are calculated and synthetic trust ranges derived from independent models are given.   Comparative studies have been carried out for simulations of radiative transfer (Pascucci et al. 2004, Iliev et al. 2006), of white dwarfs (Barstow et al. 2001) or photon dominated region (R\\\"ollig et al. 2007).  No comparison study has been presented so far for substellar atmospheres (brown dwarfs and planets). While we compare a number of model predictions, including emergent spectra, we refrain from comparing against spectra of individual substellar objects, since there are as yet no such objects with independently constrained mass, age, and metallicity against which spectral models can be compared. ",
        "conclusions": "Clouds in the atmospheres of brown dwarfs and gas-giant planets determine their spectral appearance and influence their evolution by altering the atmospheric thermal structure. The challenge of modelling cloud formation has been approached from very different perspectives over the past years which leads to the question: Do these models yield the same results and how much do they differ in predicted observational quantities?  Five models are compared in this paper to address these questions. All models emphasise the chemistry of cloud formation. Considering clouds as the result of a phase transition process (gas to solid/liquid), the models assuming phase-equilibrium describe the end-state of the phase-transition process, whereas kinetic models describe the initial state of the cloud formation process from a chemical point of view. Which viewpoint is most correct ultimately depends upon the timescales for the various relevant atmospheric processes.  The dust cloud models predict generally comparable cloud structures despite the different approaches, although the results differ substantially in detail. Opacity relevant quantities like grain size, amount of dust, dust- and gas-phase composition vary between the various approaches. Most cloud models agree that small grains composed mainly of silicates (MgSiO$_2$[s]/Mg$_2$SiO$_4$[s]) populate the upper cloud layers, whereas iron (Fe[s]) is a major component of the large grains at the cloud base.  The cloud models agree on the gas-phase composition in the inner atmosphere only which is too warm for condensed phases.  All models predict phase-equilibrium here, though the different models describe the evaporation at different levels of detail.  Above the cloud, more molecules remain in the gas-phase if cloud formation is treated in phase-non-equilibrium compared to results from phase-equilibrium models.  The different results that arise from differences in cloud modelling are amplified if the entire atmosphere problem is solved, including radiative and convective energy transport. The reason is the strong feedback of the clouds on the $(T, p)$-structure due to the clouds' strong opacity and its high efficiency in depleting the gas of the atmosphere.  Viewing their spectral appearance, the results of the cloud model atmosphere codes appear to fall into two categories:\\\\ -- The {\\it high-altitude cloud models} or {\\it extended} (Tsuji$_{\\rm Tcr=1700K}$, Marley, Ackerman \\& Lodders, Dehn \\& Hauschildt + Helling \\& Woitke) where the dust-to-gas ratio peaks at high altitudes though at different absolute levels. In these models, small grains ($\\langle a\\rangle= 10^{-6}\\,\\ldots\\,10^{-4}\\mu$m) are still present well above the maximum of $\\rho_{\\rm d}/\\rho_{\\rm gas}$, hence the gas-phase absorption is less deep for $\\lambda > 1\\mu$m.\\\\ -- The {\\it low-altitude cloud models} or {\\it thin} (Tsuji$_{\\rm Tcr=1900K}$, Allard \\& Homeier) where the dust-to-gas ratio maximum sits further inside the atmosphere and no grains populate higher atmospheric layers above the maximum of $\\rho_{\\rm d}/\\rho_{\\rm gas}$. Consequently, the gas-phase absorption features are much deeper in the mid-IR and IR part of the spectrum. Consequently, the {\\it low-altitude cloud / thin} model atmospheres appear bluer then the {\\it high-altitude cloud / extended} model atmospheres.  Comparing synthetic photometric fluxes and colours from different model atmosphere codes illustrates the current range of uncertainty, or error bar, for theoretical predictions. These error bars are worst cases. They are derived from a group of different models and not from only one particular family of models.  In the most conservative case, the maximum differences in photometric fluxes, $\\Delta_{\\rm max}[\\log F_{\\lambda}]$, amongst the models are between 1\\% and 30\\%. $\\Delta_{\\rm max}[\\log F_{\\lambda}]$ increases to 50\\% for the T-dwarf test case in the WFCAM UKIRT wavelength intervals. This translates into an uncertainty in apparent magnitudes for the L-dwarf test case of $0.25 < \\Delta m < 0.875$. $\\Delta m$ increases to 1.375 for the T-dwarf test case in the WFCAM UKIRT filter system.  We conclude that every comparison with observations should ideally involve models from different groups.  This would allow the determination of a synthetic error bar in determinations of fundamental quantities like T$_{\\rm eff}$, $\\log g$, and metallicity. Ultimately comparison of models to objects with independently constrained properties (from orbital motion and bolometric luminosity, for example) will elucidate the modelling approaches that most accurately capture the relevant physics.  Other possibilities for tests include objects in young stellar clusters which have well constrained ages and metallicities.  Future works on cloud formation need to seek more support in laboratory astrophysics. Hydrodynamic modelling of ultra-cool atmospheres will provide the opportunity to study the dynamic processes of cloud formation. They need to including a consistent description of the chemical formation processes and simultaneously address the challenge of a turbulent fluid field."
    },
    "0809/0809.2467_arXiv.txt": {
        "abstract": "We revisit the limits of the level of the extragalactic background light (EBL) recently reported by the MAGIC collaboration based on the observed $\\gamma$-ray spectrum of the quasar 3C279, considering the impact of absorption of high-energy $\\gamma$-ray photons inside the broad line region (BLR) of the quasar. We use the photoionization code CLOUDY to calculate the expected optical-UV radiation field inside the BLR and the optical depth to $\\gamma$-rays for a relatively extended set of the parameters. We found that the absorption of $\\gamma$-ray photons, though important for the estimate of the true radiative output of the source, does not produce an important hardening of the spectrum of 3C279 in the energy band accessible by MAGIC, supporting the method used to infer the upper limits to the level of the EBL. ",
        "introduction": "The detection of 3C279 ($z$=0.536) in the very high energy (VHE, $E>50$\\,GeV) band by the MAGIC telescope (Albert et al. 2008) extends to the quasars the group of the known extragalactic VHE sources, before limited to BL Lac objects (excluding the nearby radiogalaxy M87)\\footnote{see {\\tt http://www.mppmu.mpg.de/$\\sim$rwagner/sources}}.  For some aspects, the detection of 3C279 comes as a surprise. General theoretical arguments support the view that quasars cannot be important VHE emitters, in particular because of the expected absorption, through the pair production process $\\gamma + \\gamma \\rightarrow e^+ + e^-$, inside the source itself (among the most recent calculations, e.g., Donea \\& Protheroe 2003, Liu \\& Bai 2006, Reimer 2007). Moreover, quasars are generally located at a relatively high redshift, implying a huge absorption by the extragalactic background light (EBL, e.g. Kneiske et al. 2004, Primack et al. 2005, Stecker et al. 2006, Franceschini et al. 2008).  Last, but not least, widely adopted standard leptonic models for production of $\\gamma $-rays in 3C279 (Hartman et al. 2001, Ballo et al. 2002; see also Tavecchio \\& Ghisellini 2008, hereafter TG08) do not predict an important emission above few tens of GeV, because of the rapid decrease of the scattering cross section (for hadronic scenarios see, e.g. Mannheim 1993). The observation of VHE photons from 3C279 by MAGIC (Albert et al. 2008) demonstrates that also quasars can, to some extent, produce high-energy $\\gamma $-rays and suggests that the opacity (both intrinsic and cosmic) is less strong than what previously assumed.  The detection of a source of VHE photons at a relatively high redshift offers a unique tool to probe the still poorly known EBL. Albert et al. (2008) used general arguments based on a limiting spectral slope for the emitted spectrum to infer the amount of extragalactic absorption\\footnote{See Stecker \\& Scully (2008) for an alternative interpretation.}. However, as recently pointed out by Aharonian et al. (2008) (see also Bednarek 1997), intrinsic absorption of $\\gamma$-ray photons inside the source could result in rather hard observed spectra. Such spectra affected by absorption could lead to severely underestimate the level of the EBL when the arguments based on the hardness of the spectrum are used. For this reason, the use of the measured spectrum to constrain the EBL in Albert et al. (2008) triggered some discussion on the role and strength of the internal absorption in this source (Liu et al. 2008, Sitarek \\& Bednarek 2008).  We emphasize that the absorbed spectrum will be harder than the intrinsic one only in the case of an optical depth $\\tau(E)$ {\\it decreasing} with energy, $E$. Thus, as long as the optical depth increases (or, at least, only slightly decreases) with energy, the resulting spectrum will be softer (or slightly harder) than the intrinsic one and the standard spectral methods to constrain the EBL could be safely used. It is thus clear that the shape of $\\tau(E)$, strictly related to the spectrum of the target photons, is the key element to assess the effects of absorption on these methods.  Previous attempts to calculate intrinsic absorption in the BLR (Liu \\& Bai 2006, Reimer 2007, Liu et al. 2008) assumed rather idealized templates for the BLR spectrum, considering the most prominent emission lines but neglecting the important contribution of the optical-UV continuum. In a recent paper, Sitarek \\& Bednarek (2008) include the absorption in a self-consistent model for the high-energy emission of 3C279. They consider in detail the geometry of the radiation fields inside the BLR, including also the contribution of direct and scattered radiation of the accretion disc, but, again, they model the BLR radiation with the simplified template of Liu \\& Bai (2006).  In this work (Sect. 2) we explore the effects of absorption using more realistic spectra of the BLR calculated with the photoionization code {\\tt CLOUDY} (Ferland et al. 1998), previously used to study in detail the inverse Compton emission from powerful blazars (TG08). In particular, we show that the IR-optical-UV continuum plays an important role in determining the absorption, resulting in an almost constant optical depth in the broad energy range 30 GeV--30 TeV. In Sect. 3 we use the spectra corrected for absorption to revisit the constraints on the EBL based on the spectrum of 3C279 derived in Albert et al. (2008). In Sect. 4 we discuss the results. ",
        "conclusions": "The detection of VHE $\\gamma$-rays from 3C279 and EBL constraints derived in Albert et al. (2008) triggered some discussion on the role and strength of the internal absorption in this object (Liu et al. 2008, Sitarek \\& Bednarek 2008). In this paper, we have calculated various models for the BLR diffuse radiation and applied them for the case of 3C279.  We used the code {\\tt CLOUDY} to calculate the models in order to have a realistic energy spectrum of the photons inside of the BLR region. The main difference between our model and those of Liu et al. (2008) and Sitarek \\& Bednarek (2008) is the assumed BLR spectrum.  Liu et al. (2008) assume that the BLR spectrum consists only of narrow lines. Sitarek \\& Bednarek (2008) present a detailed model for the high-energy emission of 3C279, including also the possible contribution of direct and scattered disc radiation. However, their spectrum of the BLR emission does not include other important contributions to the continuum, such as the free-free emission.  As we have shown, our model provides a strong continuum component, extending on the optical-UV band (TG08).  Our results confirm that radiation from the BLR modifies the primary emission of 3C279. However, we find that the internal absorption inside the BLR does not produce an important hardening of the spectrum in the energy band covered by the MAGIC observation. In particular, we found that for the tested BLR models, despite a possible overall softening of the 3C279 spectrum, at least part of the spectrum was significantly above an implied maximum hardness of $\\Gamma=1.5$, confirming the EBL constraints derived in Albert et al. (2008). Of course, the consideration of intrinsic absorption implies that the TeV emission from 3C279 could be substantially more powerful than published.   Note also that conclusions from Sitarek \\& Bednarek (2008) that an EBL model of Stecker et al. (2006) does not imply an unrealistic intrinsic spectra of 3C279 (when taking into account internal absorption in the BLR) concern the ``baseline'' EBL model of Stecker et al. (2006).  The EBL limits in this paper and in Albert et al. (2008) instead concern the ``fast evolution'' model of Stecker et al. (2006), implying a significantly higher EBL level in the redshift range between $z=0$ and $z=1$.  We finally note that the discussion on the role of absorption implicitly assumes that the highly variable high-energy $\\gamma$-ray emission detected from 3C279 is produced internally to the BLR (probably through the comptonization of the BLR photons). However, given the small size of the BLR in 3C279 as estimated from the empirical relations connecting it to the disk luminosity, it is conceivable that the emission is (at least in some occasions) produced outside the BLR, thus avoiding the problems connected to absorption. In this case, the mechanism responsible for the production of the observed emission cannot be the external Compton: alternatives include synchrotron self-Compton emission (Lindfors et al. 2006) or comptonization of the IR radiation from the putative dusty torus surrounding the central regions (e.g. Sikora et al. 2002, 2008; Sokolov \\& Marscher 2005)."
    },
    "0809/0809.0238_arXiv.txt": {
        "abstract": "Central peaks in the iron abundance of intracluster plasma are a common feature of cooling-core galaxy clusters. Although centrally localized, these abundance peaks have a much broader profile than the stars of the central brightest cluster galaxy (BCG), which produce the excess iron, indicating that metal-enriched plasma is transported out of the BCG by some process such as turbulent diffusion.  At the same time, cooling-core clusters are likely heated by central active galactic nuclei (AGNs) by means of jets, cosmic-ray bubbles, and/or convection. The recent AGN-driven convection model of Chandran \\& Rasera predicts the turbulent velocity profile in a steady-state cluster in which radiative cooling is balanced by heating from a combination of AGN-driven convection and thermal conduction.  We use the velocity profiles from this model as input into an advection/diffusion model for the transport of metals in the intracluster medium, taking the iron to be injected by the BCG. We compare the results of our model to XMM and Chandra observations of eight clusters.  Assuming a constant turbulence level over a cluster's lifetime, the turbulent velocities in the model can explain the observed abundance profiles in only five of the eight clusters. However, we go on to develop an analytic fit of the turbulent velocity profile as a function of the AGN power. We then deduce for each cluster the average AGN power (during the past $\\sim 10$~Gyr) required to match the abundance profiles. The required average values are between $10^{43}$ and $2\\times 10^{44}~\\textrm{erg.s}^{-1}$, while the present AGN powers span a much larger range from $6\\times 10^{41}$ (Virgo) to $2 \\times 10^{44}~\\textrm{erg.s}^{-1}$ (Hydra A). Our results suggest that AGN-driven convection can account for the observed abundance profiles if the AGN power varies over a cluster's lifetime between Virgo-like and Hydra-A-like values, with average values in the above-quoted range. ",
        "introduction": "In many clusters of galaxies, the radiative cooling time at the cluster's center is much shorter than the cluster's age~\\citep{fabian94}. Nevertheless, high-spectral-resolution X-ray observations show that very little plasma actually cools to low temperatures \\citep{bohringer01,david01,molendi01,peterson01,peterson03,tamura01b,blanton03}. This finding, some times referred to as the ``cooling-flow problem,'' strongly suggests that plasma heating approximately balances radiative cooling in cluster cores.  Although a number of different heat sources have been considered in the literature, there is growing interest in the role of central active galactic nuclei (AGNs). The importance of AGN heating or ``AGN feedback'' is suggested by the observation that almost all clusters with strongly cooling cores possess active central radio sources \\citep{burns90,ball93,eilek04} and by the correlation between the X-ray luminosity from within a cluster's cooling radius and the mechanical luminosity of a cluster's central AGN \\citep{birzan04,eilek04}.  One of the main unsolved problems regarding AGN feedback, however, is to understand how AGN power is transferred to the diffuse ambient plasma.  A number of mechanisms have been investigated, including Compton heating \\citep{binney95,ciotti97,ciotti01,ciotti04,sazonov05}, shocks \\citep{tabor93,binney95}, magnetohydrodynamic wave-mediated plasma heating by cosmic rays \\citep{bohringer88,rosner89,loewenstein91}, and cosmic-ray bubbles produced by the central AGN \\citep{churazov01,churazov02,reynolds02,bruggen03,reynolds05}, which can heat intracluster plasma by generating turbulence \\citep{loewenstein90,churazov04,cattaneo06} and sound waves \\citep{fabian03,ruszkowski04a,ruszkowski04b} and by doing $pdV$ work \\citep{begelman01,ruszkowski02,hoeft04}.  Another way in which central AGNs may heat the intracluster medium (ICM) is by accelerating cosmic rays that mix with the intracluster plasma and cause the ICM to become convectively unstable.  A steady-state, spherically symmetric model based on this idea was developed by \\citet{chandran04} and subsequently refined by \\citet{chandran05} and \\citet{chandran07}.  In this model, a central supermassive black hole accretes hot intracluster plasma at the Bondi rate, and converts a small fraction of the accreted rest-mass energy into cosmic rays that are accelerated by shocks within some distance~$r_{\\rm source}$ of the center of the cluster.  The resulting cosmic-ray pressure gradient leads to convection (see \\citet{chandran06} and \\citet{dennis08}), which in turn heats the thermal plasma in the cluster core by advecting internal energy inwards and allowing the cosmic rays to do~$pdV$ work on the thermal plasma. The model also includes thermal conduction and cosmic-ray diffusion (viscous dissipation turns out to be smaller than the other forms of convective heating at all radii) and assumes a steady state in which the net heating rate balances radiative cooling throughout the cluster. The model uses mixing-length theory to describe convection and its effects on the ICM and predicts a self-consistent profile for the rms amplitude of the turbulent velocity.  By adjusting a single parameter in the model ($r_{\\rm source}$), \\citet{chandran07} were able to achieve a good match to the observed density and temperature profiles in a sample of eight clusters.  In several of the clusters in this sample, compact cooling flows arise within the central few kpc of the clusters because the rate of radiative cooling peaks much more sharply near the cluster center than either the convective or conductive heating rates. At even smaller radii in these clusters, the cooling flow makes a transition to a Bondi flow, in a manner similar to that described by \\citet{quataert00}. The size of the central cooling flow plays a role in regulating the mass accretion rate within the central Bondi flow, as described by \\citet{chandran07}.   In this paper, we explore the connection between this AGN-driven-convection model and the observed properties of the iron abundance profiles of cooling-core clusters. While observations of non-cooling-flow clusters show a nearly constant abundance profile, cooling-core clusters show very peaked iron distributions \\citep{degrandi01}. Observations of the relative abundances of oxygen, silicon and iron suggest that the production of iron in these abundance excesses is dominated by SNIa (and possibly winds) from the central brightest cluster galaxy (BCG) \\citep{ettori02,matsushita02,deplaa06}. If cooling cores are preserved over a timescale longer than $5$~Gyr \\citep{bohringer04}, then the observed amounts of excess iron (of order $10^8$~M$_{\\odot}$) within the cluster core are compatible with the amounts produced within the BCG by SNIa \\citep{cappellaro99} and stellar winds \\citep{ciotti91}.  However, the shape of these abundance profiles is still a mystery. The distribution of iron is much broader than the distribution of the stars that produce the metals, which indicates some additional processes are needed to transport the metals out of the BCG into the surrounding~ICM \\citep{rebusco05}.    Recently, \\citet{rebusco05} developed an analytical model of metal injection by SNIa and diffusion by turbulent gas motions. They suggested that the dissipation of the same stochastic gas motions would produce the heating required to solve the cooling-flow problem. Using this model, \\citet{rebusco05,rebusco06,graham06} found that diffusion coefficients of order $10^{28}-10^{29}~\\textrm{cm}^2.\\textrm{s}^{-1}$ are required to match the abundance profiles of cooling-core clusters, while spatial scales $\\sim 10$~kpc and velocities of the order of few $100$~km.s$^{-1}$ are needed to compensate the gas cooling.  Although these studies offered an explanation for the observed abundance profiles as well as a solution to the cooling-flow problem, they did not explain how the rms amplitude of the turbulent velocity, $u$, is determined, or how the dependence of~$u$ on the radial coordinate~$r$ is determined.  On the other hand, the AGN-driven-convection model of \\citet{chandran07} provides a physics-based theoretical framework for calculating~$u(r)$.  In this paper, we combine this model for intracluster turbulence with the model of metal injection of \\citet{rebusco05} in order to understand the production and transport of metals in the ICM.  We describe these models further in sections~\\ref{sec:model_inj} and~\\ref{sec:model_diff}. In sections~\\ref{sec:clusters} and~\\ref{sec:results1} we apply these models to a sample of eight clusters. We discuss the possible implications of our results for the variability of the AGN power during a cluster's lifetime in section~\\ref{sec:results2}, and summarize our conclusions in section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this article, we have studied the impact of AGN-driven convection on the shape of the iron abundance profile for 8 cooling-core clusters: Perseus, Hydra A, Sersic 159-03, Abell 262, Abell 1795, Virgo, Abell 496 and Abell 4059. We have used the iron injection model from \\citet{bohringer04} where metal production by SNIa and winds follows the Hernquist light profile of the central brightest cluster galaxy (BCG). We have also used the steady-state convection model from \\citet{chandran07} to determine the turbulent velocity profile in the ICM.  In this model, AGN-driven convection is the dominant mechanism for transferring energy from the central AGN to the ICM thereby preventing large quantities of plasma from cooling to low temperatures.\\\\  Stochastic motions with an rms radial velocity fluctuation $u_{NL}$ correspond to a diffusion coefficient $D=0.5 l u_{NL}$ in the mixing length-theory employed by \\citet{chandran07}, where $l$ is the mixing-length which we set to $0.4 r$. We have solved a 1D advection-diffusion equation for the iron abundance profile, and compared our results to XMM and Chandra abundance profiles. The profiles obtained without diffusion (metal injection only) are too peaked toward the center. Taking into account AGN-driven convection improves the abundance profile because it smoothes the center. For most objects in our sample the modeled profiles do not differ greatly from the observations. However, the less convective clusters (Abell 496 and Virgo) require much more convection in order to match observations whereas the most convective cluster (Hydra A) requires less.\\\\  Making use of the approximate self-similarity of the diffusion coefficient profiles in the AGN-driven convection model, we model the diffusion coefficient as a function of radius~$r$ that depends only on one parameter, namely the cosmic-ray luminosity (AGN power). We have used this function to find if any reasonable AGN power could provide a good fit to the observed iron distributions.  Although the diffusion coefficient is non-negligible only in a limited range of radii around its maximum, and although we cannot adjust the amplitude and the position of this maximum independently, we find a good match for all the clusters.\\\\  The best-fit AGN powers $L_{cr}^{fit}$ can be thought as the time average over the cluster history (10~Gyr) required to explain entirely the width of the abundance profile. We found a range of values for $L_{cr}^{fit}$ between $10^{43}$ and $2\\times 10^{44}~\\textrm{erg.s}^{-1}$, with an average of $5 \\times 10^{43}~\\textrm{erg.s}^{-1}$. Such a value seems quite reasonable and compatible with current constraints on the entropy profile, black hole mass, and radio-inferred average mechanical luminosity. The shape of the abundance profile could therefore be a fossil record of the past AGN-driven convection.  This model relies on several approximations. The separation of the abundance profile into a constant component $a_{b}$ (from the background galaxies and SNII) and an iron excess (from the central BCG) is one of the major approximations. A better method would be to also include the contribution from other galaxies and SNII and to compare with the total abundance profile, however this contribution also suffers from a lot of uncertainties. We also neglect all other possible sources of convection such as supernova winds, the jet itself, and stirring by infalling galaxies. This may explain the discrepancy between our results and the observed abundances at the very centers of several of the clusters, in some of which an abundance hole is observed. Another step towards a more realistic model would be cluster simulations with metal injection and AGN feedback \\citep{sijacki07,roediger07} but also cosmic-ray injection, anisotropic transport and cosmic-ray diffusion. By modifying the convective instability criterion these three last ingredients should drive turbulence to a level close to the one predicted by the mixing length theory."
    },
    "0809/0809.3900_arXiv.txt": {
        "abstract": "It is known that a relative translational motion between the deflector and the observer affects gravitational lensing. In this paper, a lens equation is obtained to describe such effects on actual lensing observables. Results can be easily interpreted in terms of aberration of light-rays. Both radial and transverse motions with relativistic velocities are considered. The lens equation is derived by first considering geodesic motion of photons in the rest-frame Schwarzschild spacetime of the lens, and, then, light-ray detection in the moving observer's frame. Due to the transverse motion images are displaced and distorted in the observer's celestial sphere, whereas the radial velocity along the line of sight causes an effective re-scaling of the lens mass. The Einstein ring is distorted to an ellipse whereas the caustics in the source plane are still point-like. Either for null transverse motion or up to linear order in velocities, the critical curve is still a circle with its radius corrected by a factor $(1+z_\\mathrm{d})$ with respect to the static case, $z_\\mathrm{d}$ being the relativistic Doppler shift of the deflector. From the observational point of view, the orbital motion of the Earth can cause potentially observable corrections of the order of the $\\mu$arcsec in lensing towards the super-massive black hole at the Galactic center. On a cosmological scale, tangential peculiar velocities of cluster of galaxies bring about a typical flexion in images of background galaxies in the weak lensing regime but future measurements seem to be too much challenging. ",
        "introduction": "Deflectors in motion have been considered several times in gravitational lensing analyses, the main motivation being that most astrophysical lenses are actually not at rest in the observer's frame of reference. To my knowledge the first analysis of the bending angle by a moving lens was performed by \\citet{py+bi93}, who applied a general technique developed to study zero geodesic in a perturbed spacetime to calculate the asymptotic behavior of a light ray deflected by a point mass moving on a Minkowski background. In the approximation of velocities much smaller than the speed of light, $v \\ll c$, they obtained \\beq \\alpha \\simeq -(1+z_\\mathrm{d})4m/r_\\mathrm{min}, \\eeq where $m$ is the mass of the lens, $r_\\mathrm{min}$ the distance of closest approach of the light ray to the deflector and $z_\\mathrm{d} (\\sim v/c)$ is the Doppler shift of the deflector. Up to first order in $v/c$, the deflection is affected only by the component of the speed along the line of sight. The authors noted the agreement with a Lorentz transformation of the usual static bending angle to a frame in which the lens is moving. The trajectory of light rays in the field of an ensemble of moving point masses was then considered in \\citet{ko+sc99}.  Later on \\citet{fri+al02} confirmed that the factor $(1+z_\\mathrm{d})$ corrects the bending angle with respect to the deflection angle of the same deflector at rest. The authors formulated a novel version of the Fermat principle based on envelopes of light-like geodesics and re-derived the bending angle for the case of a lens moving along the line of sight. Exploiting their method, they could also derive a lens equation in terms of distances and angles as measured in the frame relevant for observation. The deflection angle for the moving lens was still written in terms of the corresponding static angle. The $(1+z_\\mathrm{d})$ correction in the case of both radial and transverse slow motion across the line of sight was later re-derived in \\citet{fri03a}, who integrated the zero geodesics in presence of a time-dependent perturbation keeping terms up to linear order in velocity. Such analyses on the bending angle were eventually extended to arbitrary large radial velocities \\citep{wu+sp04}, uniformly moving monopoles \\citep{kli03}, binary systems \\citep{hey05} and extended mass distributions \\citep{ser05}.  From the above review, a clear understanding of how a relative motion between the lens and the observer affects the bending angle emerges. The picture is in full agreement with standard aberration, according to which angles made by light rays as observed in a moving frame are corrected by a factor $(1+z_\\mathrm{d})$ with respect to angles made by the same photons as observed in the rest frame. Such a correction can be also interpreted in terms of standard aberration of light propagating within an optically active medium with effective index of refraction induced by the gravitational field of the lens in motion \\citep{fri03b}. Furthermore, the different nature between the effect of a translational motion, related to Lorentz invariance, and a non null angular momentum, connected to frame-dragging has been also well understood \\citep[and references therein]{ser05}.  What is still missing is a complete description of the effect of a relative translational motion on actual lensing observations. The bending angle is defined as the difference between the asymptotic directions of the light ray at the source and at the observer and is usually calculated as a coordinate difference in a suitably defined system. On the other hand, observers measure image positions in the celestial sphere rather than bending angles at the lens position. The aim of the present paper is to derive a lens equation for the observed image angular positions in a consistent way and without referring to a corresponding static angle. In general, the bending angle for lens and observer in a static configuration can be written in terms of the rest mass of the lens and of the impact parameter of the light ray, which in turn is related to the constants of motion of the light-like geodesics. The approach proposed here avoids problems connected to the correction of such a static angle to the moving frame by considering the geodesic motion of the light rays in the rest frame of the lens. Assuming the lens to be point-like, this is the well known problem of null geodesics in the Schwarzschild metric. The observer will be moving in this space-time. To account for the relative motion between lens and observer, the components of the four-momentum of the photons reaching the detector will be then evaluated in the moving frame and translated in observed angular positions in the observer's sky. This approach allows an easy comparison of the outputs of measurement processes for lens and observer which either stay put or have a relative motion. Furthermore, it allows to extend previous analyses to relativistic velocities for transverse as well radial motions. Calculations will be limited to the first order in deflection for photon in the weak deflection limit.  In addition to clarify some theoretical aspects, the second main motivation to study lensing observables by moving deflectors is the potential astrophysical relevance. The supermassive black hole hosted in the radio source Sgr~A* in the Galactic center has emerged as an appealing target for testing higher order effects in gravitational lensing with future space- and ground-based experiments \\cite[and references therein]{se+de07} and demands a full theoretical understanding of all effects at least at the level of the $\\mu$arcsecond. Furthermore deflectors in the Milky Way may have quite large velocities, whereas, on a cosmological scale, galaxies and galaxy clusters move with quite large peculiar velocities out of the Hubble flow. Finally, there has been some evidence that correction factors in the lensing  time delay of order of $v/c$ have been actually measured at the time of a close alignment of Jupiter with the quasar J0842+1835 \\citep{fo+ko03}.  The paper is organized as follows. In Section~\\ref{sec:geod}, we review the light-like geodesic motion in a Schwarzschild metric and the properties of the four-momentum of the photons in the locally flat frame of a moving observer. In Section~\\ref{sec:lens}, the lens equation is derived and solved in the first order of deflection. The effect of aberration on the Einstein ring is discussed in Section~\\ref{sec:eins} together with image distortion. Section~\\ref{sec:radi} is dedicated to the case of pure radial motion whereas Section~\\ref{sec:astr} reviews the order of magnitude of the effect in astrophysical lensing system on several scales. Section~\\ref{sec:conc} is devoted to some final considerations and comments. ",
        "conclusions": ""
    },
    "0809/0809.3241_arXiv.txt": {
        "abstract": "{% This document is divided in two parts. The first part deals with the radial velocities (RV) distributions for B-type stars and nebulosities observed with the VLT-GIRAFFE in the Large and Small Magellanic Clouds towards the open clusters NGC2004 and NGC330. Thanks to the resolution of GIRAFFE spectra, we found that the RV distribution for the nebulosities in the LMC is bi-modal. This bi-modality can be interpreted, in term of dynamics, by the expansion of the LMC4 superbubble. The second part deals with the GAIA space mission and the determination of the radial velocities by using Radial Velocity Spectrometer (RVS) spectra. The methods to determine the radial velocities are presented as well as preliminary results on simulated RVS spectra. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1244_arXiv.txt": {
        "abstract": "Interaction of a fast shock wave generated during a supernova explosion with a magnetized star-companion of the supernova precursor produces a current sheet. We consider the evolution of this current sheet and show that a singularity (shock) is formed in finite time within the ideal MHD framework. Charged particles (electrons) are accelerated in the vicinity of the singularity, and their distribution function has a plateau up to the energies of the order of $10^4 mc^2$. These fast particles radiate in the $\\gamma$-range in the strong magnetic field of the current sheet ($B\\simeq 10^6 G$). Radiation is concentrated within a narrow angle around the current sheet, $\\Delta\\theta \\simeq 3\\cdot 10^{-4}$, and its spectrum has the maximum at several hundreds of $keV$. Presented calculations confirm the model of cosmological GRBs proposed by Istomin \\& Komberg (2002). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.3077_arXiv.txt": {
        "abstract": "We present photometric and spectroscopic observations of two luminous blue variable (LBV) stars in two extremely metal-deficient blue compact dwarf (BCD) galaxies, DDO 68 with 12+logO/H = 7.15 and PHL 293B with 12+logO/H = 7.72. These two BCDs are the lowest-metallicity galaxies where LBV stars have been detected, allowing to study the LBV phenomenon in the extremely low metallicity regime, and shedding light of the evolution of the first generation of massive stars born from primordial gas. We find that the strong outburst of the LBV star in DDO 68 occurred sometime between February 2007 and January 2008. We have compared the properties of the broad line emission in low-metallicity LBVs with those in higher metallicity LBVs. We find that, for the LBV star in DDO 68, broad emission with a P Cygni profile is seen in both H and He {\\sc i} emission lines. On the other hand, for the LBV star in PHL 293B, P Cygni profiles are detected only in H lines. For both LBVs, no heavy element emission line such as Fe {\\sc ii} was detected. The H$\\alpha$ luminosities of LBV stars in both galaxies are comparable to the one obtained for the LBV star in NGC 2363 (Mrk 71) which has a higher metallicity 12+logO/H = 7.89. On the other hand, the terminal velocities of the stellar winds in both low-metallicity LBVs are high, $\\sim$ 800 km s$^{-1}$, a factor of $\\sim$ 4 higher than the terminal velocities of the winds in high-metallicity LBVs. This suggests that stellar winds at low metallicity are driven by a different mechanism than the one operating in high-metallicity winds. ",
        "introduction": "The most massive stars evolve on very short timescales of a few million years compared to the very long time scales of billion years of solar-type stars. Thus, important evolutionary events in the life of very massive stars such as core-collapse supernovae or giant eruptions of luminous blue variables \\citep[LBVs; ][]{C84} are hard to catch and study. To maximize the probability of catching such an event, it is best to observe objects known to contain many massive stars. Blue compact dwarf (BCD) galaxies are ideal laboratories for such studies as they are undergoing intense bursts of star formation and their star-forming regions are known to harbor up to tens of thousands of massive stars \\citep{TI05}. Furthermore, mass loss is a critical factor in determining the evolution of a massive star. The mass loss rate depends in turn on the metallicity of the gas, if driven by line opacity. BCDs are also excellent laboratories to test metallicity effects on mass loss as they are metal-deficient, their metallicity ranging from about half to a few percent the Sun's metallicity \\citep{I07}. Thus BCDs are the only galaxies in the local Universe where low-metallicity massive stars do exist. They are unique objects for studying processes related to massive star formation and evolution and their interaction with the interstellar medium, and for testing models of massive stellar evolution. They are the best local approximations to primordial galaxies.  BCDs have been used to investigate mass loss through observational signatures of stellar winds from massive stars. The prevailing belief is that the efficiencies of stellar winds in massive stars are significantly reduced at low metallicities. However, this conclusion is based mainly on theoretical extrapolations and observations of relatively high-metallicity massive stars in our Galaxy (12 + log O/H $\\sim$ 8.7), the Large Magellanic Cloud (12 + log O/H $\\sim$ 8.4) and the Small Magellanic Cloud (12 + log O/H $\\sim$ 8.1) and some other nearby high-metallicity galaxies. There is growing evidence suggesting that massive stars do have stellar winds, even in the most metal-deficient BCDs known, with oxygen abundances 12 + log O/H $\\la$ 7.6. Thus, \\citet{I97} and \\citet{L97} discovered a Wolf-Rayet (WR) stellar population in I Zw 18, the most metal-deficient emission-line galaxy known at the time, with oxygen abundance 12+logO/H = 7.17 $\\pm$ 0.01 \\citep{I99}. \\citet{GIT00} have shown that the observed ratio of the number of WR stars to the number of O stars in I Zw18 cannot be reproduced by theoretical population synthesis models based on non-rotating stellar evolution models and observational properties of high-metallicity WR stars. To achieve agreement between observations and theory, \\citet{CH06} have suggested that WR stars in I Zw 18 have lower emission-line luminosities than their high-metallicity counterparts. Furthermore, \\citet{TI97} have observed P Cygni profiles for the Si {\\sc iv} $\\lambda$1394, $\\lambda$1403 absorption lines in the spectrum of the emission-line galaxy SBS 0335--052E, one of the most metal-deficient BCD known, with 12 + log O/H = 7.30 $\\pm$ 0.01 \\citep{I07}. It thus appears that extremely low-metallicity massive stars do possess stellar winds, although the properties of these stellar winds may significantly differ from those of their high-metallicity counterparts.  To find more objects with observational signatures of massive star activity, we have recently assembled a large sample of about 40 emission-line dwarf galaxies which exhibit broad components in their strong emission lines, mainly in H$\\beta$, [O {\\sc iii}] $\\lambda\\lambda$ 4959, 5007, and H$\\alpha$ \\citep{I07}. Except for four objects which appear to contain intermediate-mass black holes \\citep{IT08}, the broad emission of all the other objects can be attributed to some evolutionary stage of massive stars and to their interaction with the circumstellar and interstellar medium: WR stars, supernovae, superbubbles or LBV stars. In particular, our attention was drawn to the peculiar spectrum of the BCD PHL 293B $\\equiv$ SDSS J2230--0006, with an oxygen abundance 12 + log O/H = 7.66 $\\pm$ 0.04, which shows broad hydrogen emission lines with P Cygni profiles, spectral features that are characteristic of a LBV star. Coincidentally, several months after our finding, \\citet{P08} discovered another bright LBV in the BCD DDO 68. This BCD, with an oxygen abundance 12 + log O/H = 7.14 $\\pm$ 0.03 is even more metal-deficient than PHL 293B \\citep{IT07}. It is the second most metal-deficient star-forming galaxy known, after SBS 0335--052W.  The study of these newly discovered LBV stars in the two very metal-deficient galaxies DDO 68 and PHL 293B is the focus of this paper. Giant eruptions of LBV stars, with a brightening greater than 3 mag, are exceedingly rare. In the Galaxy, although there are some 35 confirmed or candidate LBVs \\citep{C05}, only two, $\\eta$ Carinae  and P Cygni, are known to have undergone such giant eruptions over the last 5 centuries, $\\eta$ Carinae between 1837 and 1860, and P Cygni, twice in the 17th century. P Cygni was discovered in 1600 as a naked eye star, and remained bright for many years, before fading and re-appearing in 1655, and then fading again. Typically, LBVs exhibit day-to-day microvariations in brightness of 0.1-0.2 mag, and ``normal'' irregular variations of 1--2 mag on a timescale of years, in which the spectral type may vary from an early B supergiant at visual minimum to a late B or early A supergiant at visual maximum \\citep[e.g.][]{C84,HD94,L94,N96,D97,D01,C04,P06,W08}. There have been also LBVs observed in the Large Magellanic Cloud and other Local Group galaxies \\citep{M07}.  The LBV phase is believed to represent a critical transition in the late evolution of all stars with initial masses greater than about 50 $M_\\odot$. During this phase, the stars lose sufficiently large amounts of mass in recurrent explosive events, to go from the stage of O stars burning hydrogen on the main-sequence to that of core-helium burning classical WR stars. In the HR diagram, LBVs lie just to the left of the Humphreys-Davidson limit \\citep{HD94}, beyond which no stars are observed. They appear to define the locus of an instability that prevents further redward evolution.  The understanding of the physical processes operating during the LBV phase is thus crucial for modeling massive star evolution. However, the exact mechanism giving rise to a LBV outburst is still unknown, even if many LBVs at solar metallicity have been studied. The situation is worse at low metallicity because the only relatively low-metallicity LBV star that has been observationally investigated in detail is the star V1 in the H {\\sc ii} region NGC 2363 (Mrk 71) of the dwarf cometary galaxy NGC 2366 \\citep{D97,D01,P06}, with 12+log O/H = 7.89 \\citep{G94,ITL97}. High-quality spectroscopic studies of LBV stars in two more extremely metal-deficient galaxies should help us to better understand how the LBV phase depends on metallicity.  We present here new spectroscopic observations of the LBVs in DDO 68 and PHL 293B and derive their emission line parameters. We also discuss new photometric observations of DDO 68 to better constrain the epoch of the LBV outburst in it. The observations are described in \\S2. We discuss in \\S3 the photometric variations of the two LBV stars. We derive the element abundances in the two host galaxies and show how the main properties of the stellar winds in low-metallicity LBVs differ significantly from those in their high-metallicity counterparts. Our conclusions are summarized in \\S4. ",
        "conclusions": "  1. The oxygen abundances in region 3 of DDO 68 (DDO 68-3) and in PHL 293B are respectively 12+log O/H = 7.15 $\\pm$ 0.04 and 12+log O/H = 7.72 $\\pm$ 0.01. These two BCDs are thus the lowest-metallicity galaxies with detected LBV stars. PHL 293B is also the most distant galaxy where a LBV star is seen.  2. Photometric observations of DDO 68-3 show that the outburst in its LBV star has occurred sometime between 9 February 2007 and 11 January 2008.  3. Broad H and He {\\sc i} emission lines with P Cygni profiles are seen in the spectrum of the LBV star in DDO 68. On the other hand, only H broad emission lines are detected in the spectrum of the LBV in PHL 293B. In both LBVs, no heavy element emission line such as Fe {\\sc ii} was detected, presumably because of their low metallicities. The broad H$\\alpha$ luminosities in both LBVs compare well with the one in the higher-metallicity LBV in NGC 2363 \\citep{D97,D01,P06}. We find a strong increase of the H$\\alpha$ luminosity of the LBV in DDO 68: during the period from 11 January 2008 to 28 March 2008, it brightened by a factor of $\\sim$ 4.  4. One of the most striking features of both low-metallicity LBVs is the high terminal velocities $v_\\infty$ of their stellar winds. Their $v_\\infty$ are $\\sim$ 800 km s$^{-1}$, several times greater than the $v_\\infty$ of $\\sim$ 100 -- 200 km s$^{-1}$ in their high-metallicity counterparts. Probably, metallicity plays a role and winds in low-metallicity LBVs are driven by a mechanism which differs from the one operating in high-metallicity LBVs."
    },
    "0809/0809.1591_arXiv.txt": {
        "abstract": "{ To investigate the earliest phases of star formation and study how newly-born stars interact with the surrounding medium, we performed a line and continuum survey at NIR and mm-wavelengths of a sample of relatively isolated Bok globules. } {We present a follow-up observational program of a star-forming site in the globule CB230. From narrow-band continuum observations of this site, we had discovered a bright [FeII] jet, which originates in the low-mass YSO CB230-A. We aim to investigate the physical properties of the region from where the jet is launched. } {Our analysis was carried out using low-resolution NIR spectra acquired with the camera NICS at the TNG telescope, with $JH$ and $HK$ grisms and a 1 arcsec-wide slit. These observational data were complemented with infrared photometric data from the Spitzer space telescope archive. } { The relevant physical properties of CB230-A were constrained by SED fitting of fluxes from the NIR to the mm. The YSO spectrum exhibits a significant number of atomic and molecular emission lines and absorption features. The characteristics of this spectrum suggest that we are observing a region in the close vicinity of CB230-A, i. e.  its photosphere and/or an active accretion disk. The spectra of the knots in the jet contain a large number of emission lines, including a rich set of [FeII] lines. Emission due to H$_2$ and [FeII] are not spatially correlated, confirming that [FeII] and H$_2$ are excited by different mechanisms, in agreement with the models where [FeII] traces dissociative J-shocks and molecular hydrogen traces slower C-shocks. By using intensity ratios involving density-sensitive [FeII] lines, we estimated the electron densities along the jet to be $6 \\times 10^3$--$1 \\times 10^4$ cm$^{-3}$. This indicates either high density post-shock regions of ionised gas or regions with a high degree of ionisation. } {By combining the present data with previously obtained maps at NIR- and mm-wavelengths, the emerging scenario is that CB230-A is a Class 0/I YSO driving an atomic jet that  is observed to be almost monopolar probably due to its inclination to the plane of the sky and the resulting higher extinction of its red side. This primary jet appears to be sufficiently energetic to open the cavity visible in the NIR images and drive the large-scale molecular outflow observed at mm-wavelengths. {{\\bf CB230-A was revealed to be a good location to test the innermost structure of accreting low-mass protostars. }} } ",
        "introduction": "\\label{intro} High-velocity jets driven by collimated winds generated by Young Stellar Objects (YSOs) interact strongly with the surrounding medium, cleaning up the high-density material hosting the star-forming process. Jets are also thought to remove angular momentum from the accreting matter allowing lower angular momentum gas to carry on the build-up process of a new star. In addition, when they travel into the ambient cloud, jets create shocks that heat (up to thousands of K) and compress the gas, and drive bipolar molecular outflows on larger spatial scales containing colder (10-100 K) swept-up material. As a consequence, the study of the acceleration and collimation of jets is fundamental to understanding star formation.  Some of the most accessible places to survey jets are represented by Bok globules (Bok \\& Reilly \\cite{bok}), which are relatively isolated molecular clouds mainly associated with low-mass star formation (e.g. Huard et al. \\cite{huard99} and references therein). The globule CB230 (Clemens \\& Barvainis \\cite{clemens}) is a good example, given the simultaneous presence of almost all the ingredients expected in a typical star formation scenario: YSOs, a hot atomic jet, and a colder molecular outflow. One of the YSOs in the cloud (hereafter called CB230-A) was detected in the Near-Infrared (NIR) as well as at sub-millimetre wavelengths (Yun \\& Clemens \\cite{yuncle92, yuncle94a}; Launhardt \\& Henning \\cite{launhardt97}; Huard et al. \\cite{huard99}; Young et al. \\cite{young}). It is also associated with a bipolar outflow, observed in CO (Yun \\& Clemens \\cite{yuncle94b}). Using N$_{2}$H$^{+}$ observations, Chen et al.\\ (\\cite{chen}) found a clear velocity gradient across the molecular core of $\\sim 8.8$~km\\,s$^{-1}$\\, pc$^{-1}$, in the same direction as the line connecting CB230-A with a source $\\sim 10 \\arcsec$ east. This latter source close to CB230-A was detected at 7~${\\mu}$m with ISOCAM and coincides with a possibly double NIR source (CB230-B,C; see Fig.~\\ref{Fig:Out}).  The distance to CB230 was discussed by Launhardt \\& Henning (\\cite{launhardt97}), who placed it at 450~pc. This was derived by associating CB230 with a larger molecular cloud complex (Cepheus Flare), based on their respective radial velocities. The authors estimated that the given distance is accurate to within 30\\%. Kun (\\cite{kun}) discussed the distance to the Cepheus Flare molecular cloud complex and for L1177 (the dark cloud hosting CB230) quotes a distance in the range 300--560~pc, where the lower value is the kinematic distance and the higher one is the distance to the nearby (but probably unrelated) reflection nebula vdB141. Although kinematic distances lower than 1 kpc are of doubtful reliability, the lowest value agrees with the distance ($300 \\pm 30$~pc) derived for the dark clouds in the area by studying the cumulative distribution of stellar distance moduli (Wolf diagram). Young et al. (\\cite{yy06}) assumed $d=288$~pc taking this value from Straizys et al.\\ (\\cite{straizys}), who derived the distance to a group of dark clouds (close to L1177, but not including it) using photometry of 79 stars. In particular, L1167/L1174 were found to lie at $d = 288 \\pm 25$~pc, in agreement with the lower value indicated by Kun (\\cite{kun}) for L1177. In summary, it seems that $d = 450$~pc ($\\pm$ 30\\%) is a robust determination and 288~pc may be assumed as a safe lower limit.  CB230-A is classified as a Class~0/Class~I source (e.\\ g., Froebrich \\cite{froeb}). Estimates of the bolometric luminosity, based on the IRAS fluxes, range from 7.7~L$_{\\sun}$ (Froebrich \\cite{froeb}) to 12~L$_{\\sun}$ (Launhardt \\& Henning \\cite{launhardt97}), both for a distance of 450~pc, i. e.  the distance we adopt. From our sub-mm observations with SCUBA (Brand et al.\\ \\cite{brand}), we derive a gas mass of 3.4~M$_{\\sun}$, which agrees within a factor of 2 with the mass estimated by Launhardt (\\cite{launhardt01}) from observations at 1.3~mm with the IRAM 30-m. Launhardt (\\cite{launhardt01}) also found a disk of mass $\\sim 0.1$~M$_{\\sun}$ around the protostar, based on observations at 1mm continuum with the Owens Valley Radio Interferometer. By using three different evolutionary models, Froebrich (\\cite{froeb}) derived an age in the range $1.7 \\times 10^{4}$--$2.2 \\times 10^{5}$~yrs and a final stellar mass in the range 0.3--0.9~M$_{\\sun}$.  In the framework of our ongoing multi-frequency study of star formation in Bok globules (Massi et al.\\ \\cite{ma2004}; Codella et al.\\ \\cite{codella}), we carried out narrow-band NIR observations centred on the 1.644~$\\mu$m [FeII] and 2.122~$\\mu$m H$_{\\rm 2}$ lines. As a result, we detected a bright [FeII] jet in CB230, which (i) originates in CB230-A and extends for $\\sim$ 0.02 pc, (ii) is directed along the same direction as the molecular outflows axis (N-S), and (iii) is defined by two knots (called k1 and k2) superimposed on fainter elongated emission observed also in H$_{\\rm 2}$ (Massi et al.\\ \\cite{ma2004}). Figure~\\ref{Fig:Out} illustrates this by displaying an overlay of the [FeII] image and the CO(3--2) bipolar outflow as mapped by Brand et al. (\\cite{brand}). These [FeII] and H$_{\\rm 2}$ NIR lines are particularly useful, since [FeII] emission traces fast and dissociative J-shocks, which should outline the inner jet-channel, whereas H$_{\\rm 2}$ emission is expected to trace slower C-shocks that probably originate in bow-shocks, probing the region of interaction between the jet and the ambient gas. The CB230 NIR maps therefore support a complex scenario, which, to be confirmed, requires further [FeII] and H$_{\\rm 2}$ spectroscopic NIR observations; these data were obtained and enabled us to constrain the physical properties of the gas (such as extinction and excitation temperature) in the region where the jet interacts with the surrounding dense gas close to its origin, i. e. close to the central system, which consists of a protostar and, presumably, an accretion disk.   \\begin{figure} \\centering \\includegraphics[angle=0,width=8.5cm]{massi_f1.eps} \\caption{Continuum-subtracted [FeII] map (grey scale; Massi et al.\\ 2004) of CB230: the white contours mark the brightest emission indicating the YSO CB230-A. The label CB230-B,C mark two NIR objects, that could make a triple system with CB230-A. The labels k1 and k2 indicate the two [FeII] knots defining the jet. Also visible is an extended diffuse component showing the cavity opened by the YSO mass ejection. In contours, the CO(3--2) integrated emission of the blue (continuous) and red (dashed) lobes of the molecular outflow driven by CB230-A (Brand et al.\\ 2008; JCMT-data). The integration intervals are from $-6$ to $+1.37$ km s$^{-1}$ (blue lobe) and from $+4.68$ to $+14$ km s$^{-1}$ (red lobe). Contours range from 3 to 20 K km s$^{-1}$ in steps of 1 K km s$^{-1}$. The blue lobe extends for $\\sim 120\\arcsec$, the red one is roughly half as long.} \\label{Fig:Out} \\end{figure} ",
        "conclusions": "We have presented our study of the star-forming site located in the globule CB230 using spectroscopic NIR observations. We obtained $JH$ and $HK$ spectra of the knots in the [FeII] jet, which originates in the low-mass YSO CB230-A, and the protostar itself. This allowed us to derive the physical properties of the region where the jet is launched from. We have also retrieved Spitzer archive data in the range 3.6--70~$\\mu$m that, complemented with our NIR spectra, enabled us to constrain the relevant physical parameters of the protostar.  Our main results are the following:  \\begin{enumerate}  \\item the spectra of the jet knots show a large number of emission lines, including a rich set of [FeII] lines. The brightest [FeII] and H$_2$ emission lines are spatially uncorrelated, confirming that [FeII] and H$_2$ are excited by different mechanisms, in agreement with the models in which [FeII] traces dissociative J-shocks and molecular hydrogen traces slower C-shocks.  \\item the YSO spectrum exhibits a large number of atomic and molecular emission, and absorption features. The characteristics of this spectrum, if photospheric in origin, would correspond to a spectral type later than M2 for CB230-A. They also suggest a lower gravity than for class V stars, which is expected in a protostar where the radius is larger than that of a corresponding main sequence star.  \\item complementing flux measurements at different wavelengths with the Spitzer data, we obtained a SED that can be fitted by a young low-mass ($M < 0.5$~M$_{\\sun}$) protostar with a high ($10^{-4} - 10^{-6}$~M$_{\\sun}$\\, yr$^{-1}$) accretion rate and a circumstellar disk with $M \\le 0.05$~M$_{\\sun}$. These limits agree with the constraints obtained from the NIR spectrum and the other available mm and sub-mm observations. The occurrence of a cavity of relatively small inclination with respect to the line-of-sight is in accordance with the NIR morphology of the object and explains why a spectrum of the innermost protostellar region has been obtained.  \\item as for CB230-B(,C), the lack of resolved FIR and sub-mm observations prevented strong constraints by SED fitting being achievable; a large spread in the central mass ($0.1 - 1.5$~M$_{\\sun}$) and evolutionary stages (from Class I to Class II) was obtained.  \\item the absorption bands of water and CO in the YSO spectrum might be the signature of an active accretion disk, rather than an indication of a photospheric origin. NIR spectra of both high spectral resolution and high signal-to-noise ratio are needed to disentangle the photospheric emission from that of the circumstellar disk, further constraining the protostellar physical parameters.  \\item by using intensity ratios involving density-sensitive [FeII] emission lines and following the model by Nisini et al. (\\cite{nisi}), we estimated the electron densities of the [FeII] knots along the jet, obtaining values in the range $6 \\times 10^3 - 1 \\times 10^4$~cm$^{-3}$.  \\item the [FeII] jet appears to be able to drive the larger-scale outflow, exhibiting a higher momentum flux than the CO outflow.  \\end{enumerate}"
    },
    "0809/0809.2592_arXiv.txt": {
        "abstract": "{We investigate the intensity enhancement and the duration of starburst episodes, triggered by major galaxy interactions and mergers. To this aim, we analyze two large statistical datasets of numerical simulations. These have been obtained using two independent and different numerical techniques to model baryonic and dark matter evolution, that are extensively compared for the first time. One is a Tree-SPH code, the other one is a grid-based N-body sticky-particles code. We show that, at low redshift, galaxy interactions and mergers in general trigger only moderate star formation enhancements. Strong starbursts where the star formation rate is increased by a factor larger than 5 are rare and found only in about 15\\% of major galaxy interactions and mergers. Merger-driven starbursts are also rather short-lived, with a typical duration of the activity of a few $10^8$~yr. These conclusions are found to be robust, independent from the numerical techniques and star formation models. At higher redshifts where galaxies contain more gas, gas inflow-induced starbursts are neither stronger neither longer than their local counterparts. In turn, the formation of massive gas clumps, results of local Jeans instability that can occur spontaneously in gas-rich disks or be indirectly favored by galaxy interactions, could play a more important role in determining the duration and intensity of star formation episodes.} ",
        "introduction": "The role played by galaxy interactions in affecting star formation was realized by \\citet{larson}, who showed that disturbed galaxies in the Arp Catalogue \\citep{arp} have a larger dispersion in their colors and a bluer envelope in the ($U$-$B$, $B$-$V$) plane than normal systems taken from the Hubble Atlas \\citep{sandage}. Using evolutionary synthesis models, they suggested that the features found in $UBV$ colors of interacting galaxies were caused by bursts of star formation lasting a few $10^7-10^8$ years. The large amount of observational works that followed (see \\citet{saas} for a complete review) showed that in the Local Universe, the response of galaxies to mutual interactions and mergers is quite varied.   Many starbursts in the Local Universe take place in the central regions of interacting/merging galaxies, as it is the case for instance for  NGC~7714 studied by \\citet{weed81}, and the protoype starburst galaxy M~82 \\citep{deg01a,deg01b}. Another well-studied example is the NGC~4038/4039 system (the Antennae): this early stage merger presents an extended star formation, the most intense star forming regions being located between the two galaxies \\citep{wang04}. Actually, the vast majority of UltraLuminous Infrared Galaxies (ULIRGs) at low redshift, i.e. the strongest starbursts in the Local Universe, are found in interacting and merging galaxies (\\citet{sanmir96}, see also \\citet{ducmm97}). This is however not reciprocal. Indeed, \\citet{berg03} among others have shown that, in a magnitude-limited sample of 59 interacting and merging galaxies, only a weak enhancement of star formation (a factor of 2-3 in the galaxy centers) is found when compared to a reference sample of non-interacting galaxies, so that the contribution of interactions and mergers to the global star formation activity at low redshift could on average be much less efficient than suggested by the strongest examples of starbursts.   At high redshift, the role of mergers in the star formation history of galaxies is debated, too. In a pioneering study of distant infrared galaxies, \\citet{elbces03} showed that the majority of present-day stars were formed in dusty starbursts, and suggested that the later were triggered by galaxy interactions. \\citet{cons03} also suggested that about two thirds of submillimeter galaxies at $z>1$ are undergoing a major merger. In a study of the Spitzer First Look Survey, \\citet{bridge07} argued that close pairs are major contributors to the star formation density at $z>0.7$, about half of the star formation rate density at $z \\sim 1$ being attributed to major mergers and interactions. However, an important part of the infrared luminosity of galaxies could be caused by AGN heating  \\citep[e.g.][]{daddi07a, daddi07b}. \\citet{bell05} find that less than one third of actively star forming galaxies at $z \\sim 0.8$ are actually interacting or merging, the majority having undisturbed disk morphologies. Similarly, selecting interacting galaxies in the GEMS survey, \\citet{jogee07, jogee08} find that over the redshift interval z$\\sim$0.24 to 0.80 (corresponding to   lookback times 3 to 7 Gyr), the average   SFR of strongly distorted interacting/merging massive galaxies is only modestly  enhanced with respect to  normal undisturbed galaxies. At even higher redshift ($z \\sim 2$) in GOODS, \\citet{daddi07a}  find that the star formation activity of ULIRGS is long lived (at least half a Gyr) which might be longer than expected for merger-induced starbursts.   In the last decade, a lot of progress has been made in our theoretical understanding of the role played by galaxy interactions in driving star formation, and a large variety of results has been obtained (see \\citet{stru05} for a review). Barnes \\& Hernquist (1991) have shown that tidal interactions between galaxies can drive gas inflows towards the central regions, which can increase the star formation rate. \\citet{mih94a, mih96} directly studied the star formation activity in mergers of equal mass disk galaxies using self-consistent N-body simulations \\citep{mih94b} including stars, dark matter, and gas dynamics. They also pointed out that the timing and strength of an interaction-driven starburst depend on the morphology of the interacting systems, bulgeless disk galaxies being more prone to suffer enhanced star formation in the first phases of the interaction, rather than in the merging phase. Important episodes of star formation can also arise in the case of minor merger events \\citep{mih94c}, but major mergers of galaxies of comparable masses are the most efficient situations to trigger strong starbursts. Indeed, \\citet{cox07} have shown that the star formation activity of merging galaxies decreases rapidly with increasing mass ratios. More sophisticated models including supernovae feedback in regulating star formation have also been explored \\citep{spri00, cox06, cox07}. In particular, \\citet{cox06} pointed out that the large amount of freedom in selecting the feedback parameters could play a significant role in determining the maximum star formation rate during a galaxy merger. Indeed, supernovae feedback regulates the star formation and, on average, reduces the intensity of merger-induced starbursts. The effect of feedback is however generally modest within the assumptions considered as the most realistic ones.   Nevertheless, that some cases of major interactions or mergers can trigger strong starbursts, as in the examples shown by \\citet{mih96}, does not imply that the enhancement of the star formation activity is systematically high. A sample of about 50 Nbody-SPH simulations of galaxy interactions was performed by \\citet{kap05}, and these authors found  that the integrated star formation rate during an interaction is moderately increased, up to a factor of 5 but on average a factor of 2 with respect to that of isolated galaxies. More recently, \\citet{dimatteo07} (hereafter DM07) presented more than two hundreds simulations  of galaxy interactions and mergers, restricted to coplanar cases, pointing out the difficulty to drive intense starbursts. Only 17\\% of the mergers in their sample have strong bursts with an SFR ten times higher than in isolated galaxies, and half of the sample shows  an enhancement of the SFR by a factor no bigger than 4 at the peak of the starburst.  In this paper, we extend the analysis started in DM07 on the relation between major galaxy interactions and star formation. On the one side, we enlarge the sample of Tree-SPH simulations studied in that paper, modeling orbits with various inclinations. On the other side, we test whether the main conclusions depend on the model used. To this aim, we compare the result of this first sample to a second (smaller) sample of simulations performed with a different numerical code -- a particle-mesh sticky-particle  (hereafter PM-SP) -- using somewhat different initial conditions and testing different numerical recipes for star formation. This large and heterogeneous data set shoud help in finding robust results about the starbursts--galaxy interaction connection. It is also, to our knowledge, the first numerical work in this field where simulations realized with different codes and star formation recipes are directly compared.   The layout of the paper is as follows. The simulations techniques and parameters are presented in  Sect.\\ref{description}, distinguishing Tree-SPH and PM-SP models. Results from Tree-SPH simulations are presented in Sect.\\ref{treeresults}, and, in Sect.\\ref{comparison} the main findings are compared to that obtained employing a PM-SP approach. After comparing  our results about the frequency of starburts episodes with observations (Sect.\\ref{obs}), in Sect.\\ref{concl}, the main conclusions of this research are presented. ",
        "conclusions": "In this work we have analyzed the relation between galaxy interactions and star formation enhancement by means of a vast number ($\\sim$ 1000, in total) of simulations, varying the numerical code adopted as well as the numerical recipes to model star formation. % The main results of this study are the following.\\\\ \\begin{itemize} \\item At low redshifts, interactions and mergers, in general, produce moderate enhancements in the star formation rate of the pairs, while strong starbursts are rare. This result does not depend either on the numerical code adopted, or on the recipes used to account for star formation. For star formation prescriptions compatible with observations ($\\Sigma_{SFR} \\propto {\\Sigma_{gas}}^{1.5}$ or $\\Sigma_{SFR} \\propto \\Sigma_{gas}\\Omega$), the majority of the encounters ($\\sim$ 85$\\%$) leads to a maximal star formation enhancement less than a factor five. As discussed in Sect.\\ref{obs1}, these results are in good agreement with several observational studies \\citep{berg03, cheng07, jogee07}. \\item The duration of the moderate starbursts is generally smaller than 500 Myr for Tree-SPH models. Even if PM-SP simulations have a somewhat different statistical distribution, both models show that no more than 15~$\\%$ of interaction-induced starbursts have a duration greater than 500 Myr. This would suggest that another mechanism has a major role in triggering the activity of high-redshift ULIRGs, which has a longer duty cycle \\citep{daddi07a}. \\item Moving to higher redhifts (i.e. essentially increasing the gas mass fraction), interacting systems have absolute SFRs greater than local interacting pairs. \\item Inflow-induced starbursts during interactions and mergers are neither stronger\\footnote{The SFR being normalized to that of isolated galaxies.} nor longer than their local counterparts. In turn, Jeans instability in gas-rich disks can cause the formation of massive clumps, as observed \\citep{elm07, bourn07a, bourn08}. The disk fragmentation can appear either in interacting either in isolated systems and the induced fragmentation-driven starbursts are characterized by high star formation enhancements as well as long bursts duration, but then interactions and mergers are not necessarily required to trigger this kind of activity. \\end{itemize}   Additional numerical work is required to firmly establish the nature of high-$z$ ULIRGs and  understand the role of internal factors compared to interactions and mergers in the triggering of the star formation activity and/or AGN. Nevertheless, this work presents strong results in terms of frequency and duration of starbursts episodes triggered by galaxy interactions. Strong starbursts can occur in some major mergers, and this is consistent with submillimiter galaxies being mostly major mergers  \\citep[e.g.][]{tacconi08}, but, on average, the classical merger- driven gas inflows do not seem to be an efficient process, and other mechanisms could have an important role too, like the fragmentation of gas-rich disks, or the presence of large gas reservoirs as recently suggested by \\citet{daddi07c}. These conclusions hold well with observations showing that a large number of actively star- forming galaxies are massive disks rather than mergers \\citep {genzel06, FS06, daddi07a, shapiro08}. \\citet{elm07} have shown that resolved star-forming galaxies in the Hubble Ultra-Deep field at redshift 1 and above are largely dominated by primordial disks. These high- redshift disks outnumber merger candidates; they form stars at high rates because they are massive and gas-rich: pervasive clumpy structures attest of the high gas content of these massive star- forming disks. These results certainly also help to disentangle the nature of massive star formation at high redshift, in the light of the limited role of mergers in systematically inducing extreme star formation rates, as suggested by our models."
    },
    "0809/0809.0418_arXiv.txt": {
        "abstract": "Thermal emission from the accretion disc around a black hole can be polarized, due to Thomson scattering in a disc atmosphere. In Newtonian space, the polarization angle must be either parallel or perpendicular to the projection of the disc axis on the sky. As first pointed out by Stark and Connors in 1977, General Relativity effects strongly modify the polarization properties of the thermal radiation as observed at infinity. Among these effects, the rotation of the polarization angle with energy is particularly useful as a diagnostic tool.  In this paper, we extend the Stark and Connors calculations by including the spectral hardening factor, several values of the optical depth of the scattering atmosphere and rendering the results to the expected performances of planned X-ray polarimeters. In particular, to assess the perspectives for the next generation of X-ray polarimeters, we consider the expected sensitivity of the detectors onboard the planned POLARIX and IXO missions. We assume the two cases of a Schwarzschild and an extreme Kerr black hole with a standard thin disc and a scattering atmosphere. We compute the expected polarization degree and the angle as functions of the energy as they could be measured for different inclinations of the observer, optical thickness of the atmosphere and different values of the black hole spin. We assume the thermal emission dominates the X-ray band. Using the flux level of the microquasar GRS 1915+105 in the thermal state, we calculate the observed polarization. ",
        "introduction": "The imprints of General Relativity (GR) effects on the radiation emitted in the inner regions of the accretion disc around the black hole in Active Galactic Nuclei (AGN) and Galactic black hole systems (GBHS) have been the subject of intense observational efforts in the recent past. The iron line spectroscopy has been so far the most used and successful technique to study GR effects and to estimate the black hole spin in both classes of objects (see \\citealt{Miller07} for a recent review and references therein). Recently, continuum measurements have also been shown to be a promising technique for GBHS (e.g. \\citealt{Liu}).  X-ray polarimetry is also a potentially very powerful technique in this respect, as shown by \\citet{mat93} and Dov\\v{c}iak, Karas \\& Matt (\\citeyear{dovc04b}) for AGN. For GBHS, \\citet{sc77}; \\citet{cs77} and Connors, Stark \\& Piran (\\citeyear{con80}) discussed the GR effects on the polarization properties of the thermal emission of the accretion disc.  They showed that a strong variation of the polarization angle and degree with energy, due to the radial dependence of the temperature, is induced. In this paper we reconsider and extend the Stark and Connors calculations, mainly we consider the scattering atmosphere with different optical depths.  We assume the Kerr metric for the gravitational field, a standard thin disc with a scattering atmosphere for the accretion flow and Thomson scattering in the disc atmosphere as the origin of polarization. The resulting polarization angle and degree will be shown as functions of the energy for different inclinations of the observer, optical thickness of the atmosphere and black hole spins.  The advent of a new generation of X-ray polarimeters (\\citealt{Costa2001}) makes the observation of the effects discussed in this paper feasible.  We will discuss the observational perspectives for a bright GBHS, the well-known high inclination microquasar GRS~1915+105, for both a relatively small polarimetric mission like POLARIX (\\citealt{Costa2006}) and for the polarimeter aboard IXO (\\citealt{Bellazzini2006b}). ",
        "conclusions": "We have investigated the GR effects on the polarization properties of the thermal radiation emitted by an accretion disc around a stellar mass black hole. The polarization at infinity is changed from its local value due to strong gravity effects and fast orbital motion of the disc close to the central black hole. The difference is particularly dramatic especially in the region below the critical point, where the polarization angle swings over the entire range of 180 degrees.  The most interesting interval of energy is from 0.1 to 10 keV. Below this energy band the emission comes from far away from the disc centre, while above this range the flux rapidly decreases. Note, however, that for different maximum temperatures of the disc the value of the upper boundary of this energy interval will change. All studied energy dependences will be shifted to a higher energy for a hotter disc. The maximum temperature is a function of the black hole mass, accretion rate and hardening factor as can be seen from the eq.~(\\ref{temp}).  The polarization degree is the highest for the lowest studied optical depth of the disc's atmosphere and in most cases it grows with the observer's inclination. The polarization angle does not change much with the inclination of the observer apart from the highest inclinations. Instead, it depends quite strongly on the spin of the central black hole.  The effects of strong gravity on the polarization state of emergent radiation will be accessible to the next missions with an X-ray polarimeter based on the Gas Pixel Detector. We presented the case of the pathfinder mission, feasible in a few years, and that of a large mission. In both scenarios, the energy dependence of the degree and of the angle of polarization will clearly discriminate between the Kerr and the Schwarzschild black holes cases. Even in the pathfinder scenario, the X-ray polarimetry will be a powerful probe of the metric around black holes."
    },
    "0809/0809.0310_arXiv.txt": {
        "abstract": "Using a one-dimensional $\\alpha\\omega$-dynamo model appropriate to galaxies, we study the possibility of dynamo action driven by a stochastic alpha effect and shear. To determine the field evolution, one needs to examine a large number of different realizations of the stochastic component of $\\alpha$. The net growth or decay of the field depends not only on the dynamo parameters but also on the particular realization, the correlation time of the stochastic $\\alpha$ compared to turbulent diffusion timescale and the time over which the system is evolved. For dynamos where both a coherent and fluctuating $\\alpha$ are present, the stochasticity of $\\alpha$ can help alleviate catastrophic dynamo quenching, even in the absence of helicity fluxes. One can obtain final field strengths up to a fraction $\\sim 0.01$ of the equipartition field $ B_{eq}$ for dynamo numbers $\\vert D\\vert \\sim 40$, while fields comparable to $ B_{eq}$ require much larger degree of $\\alpha$ fluctuations or shear. This type of dynamo may be particularly useful for amplifying fields in the central regions of disk galaxies. ",
        "introduction": "Large-scale magnetic fields in stars and galaxies are thought to be generated and maintained by a mean-field turbulent dynamo \\citep{Mof78,KR80}. The potential driver of such mean-field dynamos is the $\\alpha$-effect, arising whenever one has rotation and stratification in a turbulent flow. Mean-field dynamo (MFD) models using a coherent $\\alpha$-effect and shear have been invoked to explain large-scale fields observed in disk galaxies \\citep{RSS88}.  The possibility of efficient dynamo action arising from random fluctuations in the $\\alpha$-effect in combination with shear was first pointed out by \\citet{VB97}. They investigated a reduced mean-field dynamo model appropriate to accretion disks and showed that growth can occur for large enough random fluctuations in alpha. Several authors have since elaborated various aspects of this stochastic alpha-shear dynamos \\citep{So97,S00,FBR06,Proc07,KR08}. In particular, \\citet{So97} examined a model of a disk dynamo with a fluctuating alpha antisymmetric in space but which changes sign randomly with equal probability. He argued that intermittent large-scale magnetic fields can grow. The role of a stochastic $\\alpha$ has also been analyzed in the context of solar dynamos \\citep{Proc07, BS08, Moss08}.  The exact origin of such an incoherent $\\alpha$-effect is as yet unclear. In any large Reynolds number system, many degrees of freedom exist, and hence there could always be a stochastic component of the mean turbulent electromotive force (emf). This could lead to additive or a multiplicative noise in the MFD equations. Additive noise provides a seed field for the dynamo, whereas multiplicative noise in say the $\\alpha$ effect, combined with shear, can lead to exponential growth of the mean field. In the solar context, \\citet{H93} argued for $\\alpha$ fluctuations $\\sim u_{0}/\\sqrt{M}$, where $u_0$ is the turbulent velocity and $M$ is the number of cells being averaged over in defining the mean field. In principle this can be larger than any coherent $\\alpha$-effect. Multiplicative noise is also seen in simulations which measure the $\\alpha$-effect both in the kinematic regime \\citep{SBS08} and in the nonlinear regime \\citep{CH06,BRRK08} and also in direct simulations of the galactic dynamo \\citep{G08}. In fact \\citet{BRRK08} measure an incoherent $\\alpha$-effect, with a Gaussian probability density function (PDF), even in turbulence driven with a non-helical forcing, where one does not expect a coherent $\\alpha$-effect. Combined with shear, such systems show large-scale dynamo action \\citep{BRRK08, Y08}. Here, we simply examine, in the context of galactic dynamos, the consequence of having an incoherent alpha effect, without considering in detail its exact origin. The growth of the mean field varies significantly from one realization of the stochastic process to another, as also pointed out in \\citet{So97}. It is therefore necessary to examine a large number of realizations of the stochastic $\\alpha$ to determine the efficiency of the stochastic $\\alpha\\omega$-galactic dynamo.  We outline in section 2, the basics of a one-dimensional stochastic $\\alpha\\omega$-dynamo model appropriate to galaxies. We present numerical solutions of the above model in section 3, with two different PDF's for the stochastic alpha; the first as considered in \\citet{So97} and the second where the stochastic alpha has a gaussian PDF. In Section 4, we explore the possibility of alleviating catastrophic $\\alpha$-quenching in absence of helicity fluxes by including the effects of a stochastic $\\alpha$. Section 5 summarizes our results and the implications of a stochastic $\\alpha\\omega$-dynamo for galaxies. ",
        "conclusions": "We have examined here how a stochastic $\\alpha$-effect in association with shear can lead to the generation of  large-scale galactic magnetic fields. To determine the field evolution, one needs to examine a large number of different realizations of the stochastic $\\alpha_1(t)$. The net growth or decay of the field depends on the particular realization, the correlation time of the stochastic $\\alpha$ compared to turbulent diffusion timescale and the time over which the system is evolved.  The results are illustrated first with the simple model of \\citet{So97} in Fig.~\\ref{dmqalrw}. Here the magnitude of $\\alpha_1$, takes randomly a value $+\\alpha_s g(z)$ or $-\\alpha_s g(z)$ over any time interval $\\tau_c$, with equal probability. Any given realization of $\\alpha_1(t)$ will have $N_+$ periods of steady growth (when $N=1$) and $N_-$ periods of oscillatory decay (when $N=-1$). But, since the growth rate ($\\gamma$) and decay rates ($\\zeta$) are different, this could lead to a net growth or decay as pointed out by \\citet{So97}. The critical dynamo number for getting growth for say $50\\%$ of realizations is $\\vert D_c\\vert \\sim 30$, moderately larger than $\\vert D_c\\vert \\sim 10$ required for the coherent $\\alpha\\omega$-dynamo. Our numerical solutions confirm the applicability of this picture for long correlation times $\\tau_c=2t_d$, while for $\\tau_c < t_d$, the picture changes qualitatively, and the dynamo becomes less efficient.  We then examined more realistic MFD models with a short correlation time ($\\tau_c = 1.5 t_{ed}$), for a stochastic $\\alpha_1(t)$ chosen from a gaussian PDF. Our results are given in Fig.~\\ref{modqalrw} and Fig.~\\ref{histo}. In this case as well, and for a stochastic $\\alpha\\omega$-dynamo number $D=-40$, about $34\\%$ of realizations showed growth of $\\meanB$ till $t \\sim 2$. However subsequently $\\meanB$ decays to negligible values by $t \\sim 10$. For higher dynamo numbers, growth is sustained for a longer time and for a larger number of realizations (cf. Fig.~\\ref{histo}). One requires $ D \\sim -120$ to obtain long term growth (till $t \\sim 10$), in a significant number ($\\approx$ 24\\%) of realizations, $\\vert D \\vert \\approx 160 - 180$ for the PDF of $\\vert A\\vert$ to remain stationary and a larger $\\vert D\\vert$ for secular growth at late times. Note that having an additional coherent $\\alpha$, with $\\vert\\Ral\\Rw\\vert > 10$, would ensure growth of $\\meanBB$. However, the quenching imposed by helicity conservation still needs to be alleviated.  In the usual $\\alpha\\omega$-dynamo, such helicity conservation leads to a growth of the magnetic $\\alpha_{\\rm m}$, which tends to cancel the kinetic $\\alpha_{\\rm k}$, so as to catastrophically quench the dynamo. This problem can be alleviated by having fluxes of magnetic helicity. In contrast, for a stochastic $\\alpha\\omega$-dynamo, by the time $\\alpha_{\\rm m}$ has grown, $\\alpha_{\\rm k}$ could have changed sign. This raises the possibility of alleviating catastrophic $\\alpha$- quenching without helicity fluxes. To examine this possibility, we solved the stochastic $\\alpha\\omega$-dynamo equations along with the dynamical $\\alpha$-quenching equation. We included both a coherent and incoherent $\\alpha$-effect. In general the radial and azimuthal fields are again quadrupolar and show occasional reversals in time (see Fig. \\ref{contour}). When the coherent and stochastic components of $\\alpha$ are comparable, and even without any helicity flux, we find that steady large-scale magnetic field of strengths of about $0.01B_{eq}$ could be obtained in some realizations even for $D \\sim - 40$. Field strengths above $0.3B_{eq}$ are obtained for stronger amplitude of random $\\alpha$-fluctuations in association with strong shear (Fig. \\ref{c200ral1qal1rw40}). Therefore, a stochastic $\\alpha\\omega$-dynamo model is more likely to grow large-scale magnetic fields efficiently towards the central regions of a galaxy and even in the absence of helicity fluxes.  We have focussed on the application of a stochastic $\\alpha\\omega$-dynamo model to galaxies. However, our emphasis on examining a large number of realizations and our results on alleviating $\\alpha$-quenching with a random $\\alpha$ could be applicable to other astrophysical dynamos as well. More work is needed to elucidate the origin of the incoherent $\\alpha$ and also study the influence of spatial decorrelation on the dynamo."
    },
    "0809/0809.2409_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmological observations imply that about $80\\%$ of the mass of the Universe is some unknown form of cold Dark Matter (DM)~\\cite{Komatsu:2008hk}. Presently the origin and nature of the DM particles, their mass, spin, couplings and other properties remain completely unknown.  Among many possible cold DM candidates the most popular ones are the stable weakly interacting massive particles (WIMPs) which occur in many extensions of the Standard Model (SM), most notably in supersymmetry.  If the DM WIMPs are thermal relics,  their annihilation cross section must be $\\sigma v \\sim 3\\cdot 10^{-26}\\cm^3/{\\rm sec}$, which indeed is typical of a weakly interacting TeV-scale particle. To discover the DM particles, experiments have searched for direct production of WIMPs in collider experiments, their scattering off the nuclei in terrestrial detectors as well as {\\it indirect signals} of particles generated by the DM annihilation in the galactic halo~\\cite{pioneers}.  The recently reported  results by the PAMELA experiment~\\cite{PAMELA} have the opportunity, if confirmed, of establishing a breakthrough in the cosmic antimatter searches. The  PAMELA data show $(i)$ a steep increase in the energy spectrum of the positron fraction, $e^+/(e^++e^-),$  in cosmic rays above  10~GeV~\\cite{PAMELA}, compatibly with previous less certain hints from HEAT~\\cite{Barwick:1997ig} and AMS-01~\\cite{Aguilar:2007yf}; $(ii)$ no excess in the $\\bar p/p$  energy spectrum~\\cite{PAMELApbar} compared with the predicted background; $(iii)$ at low energy, $E_{e^+}< 10$~GeV,  the positron flux is presently suppressed by the solar magnetic polarity state $A^-$~\\cite{clem}. The positron excess observed by PAMELA is so big that it can produce observable spectral features in the total $e^++e^-$ flux. Indeed, most recently the PPB-BETS balloon experiment reported an excess in the $e^++e^-$ energy spectrum between 500-800~GeV \\cite{Torii:2008xu}, confirming the similar earlier claim by the ATIC-2 balloon experiment \\cite{ATIC-2}. These data-sets have not received much attention yet as probes of DM indirect detection, but, as we will see, they have the potential of being relevant to the issue. Of course one should be careful: the Monte Carlo simulations that such experiments need to tag $e^\\pm$ and infer their energy have been tested only up to LEP energies; the excess is based on just a few data-points that are not cleanly consistent between ATIC-2 and the smaller PPB-BETS; emulsion chambers (EC) balloon experiments~\\cite{EC} do not show evidence for an excess, although they have larger uncertainties. The excess in both experiments seems to shows a cut-off at energies just below 1~TeV. Those results will soon be tested by forthcoming ATIC-4 \\cite{ATIC-4} and GLAST/Fermi data~\\cite{GLAST}, that have bigger calorimeters able of containing a electromagnetic shower. All these excesses might be of course due to some astrophysical source, such as a nearby pulsar~\\cite{pulsar}. However, they also might provide the first clear evidence for the DM annihilations into SM particles.   In this work we explore whether galactic DM annihilations can account for the observed excesses in the data. We perform a model independent analysis of the PAMELA positron and antiproton data with and without the electron plus positron data obtained from the balloon experiments. Our aim is to study if and how one can get model independent information on the DM mass, spin, interactions to the SM particles and on the DM annihilation cross section from the experimental data. Taking into account restrictions that must be obeyed by non-relativistic DM annihilations, we identify the parameters that can be probed by the indirect signals. We study the non-relativistic DM annihilation cross sections into the set of all possible SM particle final states, DM DM $\\to$ SM SM, where $${\\rm SM}=\\{e, \\mu_L, \\mu_R, \\tau_L, \\tau_R, W_L,  W_T, Z_L, Z_T, h, q, b, t\\},$$ taking into account the allowed polarizations ($T$ransverse, $L$ongitudinal, $L$eft, $R$ight). Using Monte Carlo tools, partly written by us, we try to keep the polarizations/helicities of the DM annihilation products and compute the energy spectra of the final state $e^\\pm,p^\\pm,\\gamma, \\nubarnu_{e,\\mu,\\tau}$ coming from their decays. Indirect signals of any DM model can be obtained by combining  our decay spectra according to the branching ratios predicted by the model. The experimental results can be analyzed in terms of these phenomenological parameters, without committing to some specific theoretical scenario.   \\begin{figure}[p]\\begin{center} $$\\includegraphics[width=\\textwidth]{sample100}$$ $$\\includegraphics[width=\\textwidth]{sample1000}$$ $$\\includegraphics[width=\\textwidth]{sample10000}$$ \\caption{\\em Three examples of fits of $e^+$ (left), $e^++e^-$ (center), $\\bar{p}$ (right) data, for $M=150\\GeV$ (upper row, excluded by $\\bar{p}$), $M=1\\TeV$ (middle row, favored by data), $M=10\\TeV$ (lower row, disfavored by the current  $e^++e^-$ excess). Galactic DM profiles and propagation models are varied to provide the best fit. See Sec.~\\ref{sec:propagation} for the discussion on the treatment of the uncertain astrophysical background. \\label{fig:samples}} \\end{center} \\end{figure}   Taking into account different models for the DM distribution and for the electron and proton propagation in the Galaxy, we perform fits of the predicted $e^+/(e^++e^-),$ $\\bar{p}/p$ and $(e^++e^-)$ energy spectra  to the PAMELA and balloon data. We allow for the possibility of enhanced annihilation cross sections and/or for the high density substructures in the DM halo, assuming an energy-independent boost factor $B.$ For simplicity, we assume a common boost factor for the $e^\\pm$ and for the $p^\\pm$ and we take it as energy-independent. These are widely used assumptions, but they are not the only possibility~\\cite{Lavalle}. One can easily infer what happens if these assumptions are relaxed. We thus study which DM masses and annihilation channels can best explain the experimental measurements.  \\medskip  To illustrate our results, as well as to present the PAMELA and balloon $e^+/(e^++e^-),$  $(e^++e^-)$ and $\\bar{p}/p$ data, we show in Fig.~\\ref{fig:samples} three examples for the DM masses $M =150$~GeV, 1~TeV and 10~TeV that annihilate into $W^+W^-$ and $\\mu^+\\mu^-,$ as indicated in the figures. The first column presents the predictions for the positron fraction, the second for the total electron and positron flux and the third for the antiproton over proton flux.   Many DM annihilation channels into SM primary channels can give a reasonably good fit to the PAMELA $e^+/(e^++e^-)$ excess alone, at least for some value of the DM mass picked in a wide range 60 GeV $\\to$ tens of TeV (first column of Fig.~\\ref{fig:samples}; see Fig.~\\ref{fig:fite} for the full analysis). More precisely, DM annihilations into SM leptons can well fit for almost any mass in this range, while annihilations into quarks and Higgs bosons are disfavored at low masses because of the soft positron spectrum. Annihilations into gauge bosons work for light and heavy masses.  The inclusion of the $\\bar{p}$ data changes the issue. Because the DM annihilations to gauge bosons, Higgs bosons or quarks ($W,Z,h,q,b,t$) produce antiprotons, under our assumptions the PAMELA $\\bar{p}/p$ data excludes light DM with sizable annihilation fraction into these channels (Fig.~\\ref{fig:samples}, the first row). However, for $M \\circa{>} 10$~TeV the antiproton spectrum becomes again consistent with present data, and the $W^+W^-$ seem to give the relatively best fit (Fig.~\\ref{fig:samples}, the third row).  We find therefore that the PAMELA data  single out two sets of solutions which can satisfy the observations compatibly with our standard astrophysical assumptions: (i) a heavy ($\\circa{>}$ 10 TeV) DM that annihilates dominantly into $W^+W^-$; (ii) a DM that annihilates into SM leptons, with no strong preference for any DM mass.   \\medskip  The inclusion of the balloon data qualitatively changes the sensitivity to the DM mass because the PAMELA excess predicts observable spectral features in the total flux (Fig.~\\ref{fig:samples}, the second column). It is striking  that the same signal of the DM annihilation to  leptonic channels which can explain the PAMELA data  is also able to produce the apparent $(e^++e^-)$ excess in the PPB-BETS and ATIC data. Because the balloon data shows a sharp cut-off in the excess just below 1 TeV, the DM mass should be close to $1$~TeV, and all other but leptonic DM annihilation channels would be strongly disfavored or excluded. On the other hand, if the $e^++e^-$ excess will not be confirmed, the $e^++e^-$ data will strongly constrain or exclude the DM annihilations to leptons.  Should this be the result of  the upcoming ATIC-4 (or GLAST) experiment, the astrophysically favorite candidate for DM becomes a heavy $\\sim$10 TeV WIMP annihilating into SM gauge bosons.  We comment on the implications of PAMELA and balloon data for supersymmetric DM candidates~\\cite{Ellis:1983ew}, as well as for some theoretically more exotic  alternatives~\\cite{Bertone:2004pz,Hooper:2007qk}. We explore in a qualitative way the model-independent expectations for direct DM searches and for the production and detection  of the DM particles at LHC and future lepton colliders. A few authors have discussed implications of the PAMELA $e^+/(e^++e^-)$ results~\\cite{Bergstrom:2008gr, Cirelli:2008jk,Barger:2008su,Cholis:2008hb}. While the analysis of the $e^+/(e^++e^-)$ data presented in~\\cite{Barger:2008su,Cholis:2008hb} is consistent with our more systematic treatment, our final conclusions differ because we also include the important $\\bar p$ and $e^++e^-$ data.     \\bigskip  The paper is organized as follows. Section \\ref{sec:cosmology} discusses cosmological constraints  on the DM annihilation cross section. In Section \\ref{sec:theory} we identify the DM annihilation channels and compute the spectra of SM final states. In Section \\ref{sec:propagation} we discuss the particle propagation in the Galaxy. In Sections \\ref{sec:PAMELApositrons}, \\ref{sec:PAMELAantiprotons} and \\ref{sec:balloons} we perform fits to the PAMELA positron data, PAMELA positron plus antiproton data, and to the $e^\\pm$ PAMELA and balloon data, respectively. We discuss the implications of our results for DM direct detection in Section \\ref{sec:direct} and for collider phenomenology in Section \\ref{sec:collider}. We conclude in Section \\ref{sec:conclusions}.   \\begin{figure}[t] \\begin{center} $$\\includegraphics[width=0.8\\textwidth,height=9cm]{sigmaDMDM}$$ \\caption{\\em Values of the DM annihilation cross sections suggested by the DM abundance. \\label{fig:sigmaDMDM}} \\end{center} \\end{figure}   \\begin{figure}[t] \\begin{center} $$\\includegraphics[width=0.45\\textwidth]{Som}\\qquad \\includegraphics[width=0.45\\textwidth]{SomCosmo}$$ \\caption{\\em Contour plots of the Sommerfeld enhancement for one abelian massive vector with mass $M_V$ and coupling $\\alpha$ to DM (from~\\cite{Sommerfeld2}) and the resulting Sommerfeld reduction in the cosmological thermal DM abundance (we here assumed a DM mass $M=1\\TeV$). \\label{fig:Som}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  We studied if DM annihilations into two-body SM states can reproduce the features that seem present in the energy spectra of $e^+,e^-$, $\\bar p$ cosmic rays recently measured at energies between about 10 GeV and 3 TeV by the PAMELA~\\cite{PAMELA}, ATIC-2~\\cite{ATIC-2} and PPB-BETS~\\cite{Torii:2008xu} experiments.  We found that: \\begin{itemize} \\item Considering the  PAMELA positron data alone, the observed positron excess can be well fitted by DM annihilations into $W,Z,e,\\mu,\\tau$ with any DM mass, and by DM annihilations into $q,b,t,h$ with a multi-TeV DM mass.   \\item Adding the PAMELA anti-proton data suggests that the positron excess can be fitted by DM annihilations into $W,Z,h$ (annihilations into $q,b,t$ give poorer fits) only  with a multi-TeV DM mass, unless boost factors or propagation strongly differentiate between $e^+$ and $p^-$. DM annihilations into leptons are still viable for any DM mass.  \\item Adding the balloon $e^++e^-$ data similarly suggests that DM annihilations into $W,Z,h,b,q,t$ can fit the PAMELA excess only with a multi-TeV DM mass. Annihilations into $e,\\mu,\\tau$ predict a feature in the $e^++e^-$ spectrum that, for a DM mass  around one TeV, is compatible with the hint present in ATIC-2 and PPB-BETS data. A TeV mass seems too heavy for explaining the DAMA anomaly~\\cite{Bernabei:2008yi} compatibly with CDMS bounds in terms of inelastic DM scatterings~\\cite{DAMAinel}.   \\item  The annihilation cross section suggested by the PAMELA data is a few orders of magnitude larger than what naturally suggested by the cosmological abundance, unless DM formed sub-halos. This enhancement could be due to a Sommerfeld non-relativistic enhancement of the annihilation cross section. Alternatively, this may indicate for non-thermal production of DM.  \\item Present data start to be sensitive to the polarization state of the DM annihilation products.   \\end{itemize} In the light of our results two solutions emerge for the DM. \\begin{enumerate} \\item The first possibility is a heavy DM, $M\\circa{>}10\\TeV$, that annihilates into $W^+ W^-$ or $hh$ (an example of a DM candidate of this sort is provided by the model of~\\cite{MDM1}). This possibility is illustrated by the lower row in Fig.\\fig{samples}, and will be tested by future $\\bar p$ data. \\item The second, more exciting, possibility is that the PAMELA excess has also been seen in $e^++e^-$ data, by ATIC and PPB-BETS. The forthcoming ATIC-4 and GLAST data should soon clarify this issue. In such a case the best fit is obtained for $M\\approx 1\\TeV$ with DM annihilating into $\\mu^+\\mu^-$, and good fits are also obtained for $M\\approx 800\\GeV$ if DM annihilates into $e^+e^-$, or $M\\approx 2\\TeV$ if DM annihilates into $\\tau^+\\tau^-$. This is illustrated by the middle row in Fig.\\fig{samples}. The needed `boost times cross section' is $B_e \\, \\sigma v\\sim 3~10^{-23}~\\cm^2/{\\rm sec}$. \\end{enumerate} In both cases, a DM particle compatible with the PAMELA  anomaly has therefore unexpected properties. In general, $\\bar p$ and $e^++e^-$ data seem to indicate that DM masses around $100\\GeV$ cannot fit the PAMELA anomaly, as illustrated by the fit in the upper row of Fig.\\fig{samples}. The favorite DM candidate, the supersymmetric neutralino~\\cite{Ellis:1983ew}, cannot annihilate directly into light leptons with a large cross section~\\cite{Bergstrom:2008gr}. A  multi-TeV Wino is a possible supersymmetric candidate, but supersymmetry would not naturally explain the breaking of the electroweak symmetry. The vector or Dirac fermion DM candidates appearing in the models of extra dimensions~\\cite{Servant:2002aq,Cheng:2002ej,Agashe:2004ci} and in the Little Higgs models~\\cite{Cheng:2004yc} can annihilate directly into leptons. However, at the same time they also annihilate  into quarks and/or into gauge and Higgs bosons. Therefore the masses of such DM candidates are also pushed to the multi-TeV region by the PAMELA antiproton data. Similarly, a LSP right-handed sneutrino with a sizable $\\nu_R LH$ Yukawa coupling annihilates into both  $L$ and $H$ (and consequently into longitudinal $W,Z$, that contain the Goldstone components of the Higgs doublet $H$), unless lighter than $M_W$. None of these theoretically motivated DM candidates mentioned here can explain the ATIC/PPB-BETS excess.  A model of a thermal DM candidate that can explain both the PAMELA and the ATIC/PPB-BETS  excesses with an order one boost factor can be constructed as follows. We assume that lepton flavor $L_\\mu-L_\\tau$ is a spontaneously broken anomaly-free U(1)$_L$ gauge symmetry, with gauge coupling $\\alpha_L \\approx 1/50$ and vector mass $M_L\\approx M_Z$. The DM candidate is a Dirac fermion with mass $M\\approx 1.5\\TeV$, charge $q_L\\approx 2$ under the extra U(1), and no SM gauge interactions. Its tree-level annihilation cross section relevant for cosmology, $\\sigma v = \\pi (3 q_L^2 + q_L^4)\\alpha_L^2/2M^2=3~10^{-26}\\cm^3/{\\rm sec}$ at $v\\sim 0.2$, is enhanced in astrophysics, up to $\\sigma v \\circa{<}10^{-22}\\cm^3/{\\rm sec}$ at $v \\sim 10^{-3}$, by the Sommerfeld correction due to the extra U(1), as computed in Fig.~1a of~\\cite{Sommerfeld2}. The photon spectrum resulting from these DM annihilations into $\\mu$ and $\\tau$ is hard, as can be seen in the right panel of Fig.~\\ref{fig:compspectra}. The correction to the  magnetic moment of the muon, $\\delta a_\\mu=+ g_L^2m_\\mu^2/12\\pi^2M_L^2$, reproduces the observed discrepancy with respect to the SM, $\\delta a_\\mu=(3\\pm1) 10^{-9}$~\\cite{g-2}. The new vector affects precision data as $R_\\ell\\equiv \\Gamma(Z\\to \\ell^+\\ell^-)/\\Gamma(Z\\to e^+e^-) =1+ 3 \\alpha_L f(M_L/M_Z)/4\\pi$ with $f(0)=1$; this is compatible with the LEP measurements, $R_\\mu=1.001\\pm 0.003$ and $R_\\tau=1.002\\pm 0.003$~\\cite{LEP}. Furthermore  one loop-effects generate a kinetic mixing of the new vector with the photon, $\\theta = e\\, g_L \\ln (m_\\tau/m_\\mu)/6\\pi^2 \\approx 0.005$, that gives rise to a cross section for direct detection, $\\sigma_{\\rm SI}=4\\pi q_L^2\\alpha_L\\alpha m_N^2 \\theta^2/M_V^4\\sim  10^{-42}\\cm^2$, around the present bound. Reducing $M_V$ by a factor of few leads to a too large $\\sigma_{\\rm SI}$, and reducing $M_V$ below about 1 GeV also to unseen DM annihilations in first structures~\\cite{first} and to collisional DM with elastic $\\sigma \\sim g_V^4 M^2/4\\pi M_V^4$.     \\bigskip  Finally, a third, less exciting possibility is that the features that appear in the $e^+$ and $e^+ + e^-$ spectra are not due to DM annihilations but to some astrophysical object, such as a pulsar. Indeed it is expected that pulsars produce a power law spectrum (with index $p$ around 1.5--2) of mostly electron-positron pairs, with a cut-off. The black line in  Fig.\\fig{fitee} shows that a flux $\\Phi_{e^-} = \\Phi_{e^+}=A\\cdot  E^{-p} e^{-E/M}$ with $p\\sim 2$ and cut-off at $M\\sim 300\\GeV$ does fit the $e^+,e^-$ data, giving a peak  less sharp than what hinted by ATIC-2 data. It is remarkable that pulsars naturally give an excess only in $e^\\pm$, with the observed slope $p$. While the total energy and total power emitted by a pulsar can be reliably computed assuming that they are respectively dominated by rotation and magnetic dipole radiation, the factors $A$ and $M$ cannot be computed without a detailed model of the $e^\\pm$ acceleration mechanism. Furthermore, while the known nearby pulsars Geminga or B0656+14 are the main candidates, unknown pulsars can contribute. According to~\\cite{pulsar}, young nearby pulsars tend to give a cut-off $M$ at energies higher than what suggested by ATIC/PPB-BETS, while older pulsars tend to give a too low flux, as $e^\\pm$ had time to diffuse away. An old intense pulsar no longer emits $\\gamma$, so that it is difficult to distinguish it from DM annihilating into leptons; hopefully angular anisotropies of the $e^+,e^-$ fluxes will help~\\cite{pulsar}.    \\small   \\paragraph{Acknowledgements} We thank Franco Cervelli, Dario Grasso, Nicolao Fornengo and Vyacheslav S. Rychkov for very useful discussions. This work was supported by the ESF Grant 6140 and by the Ministry of Education and Research of the Republic of Estonia. A.S.\\ thanks the EU project ENTApP working group 2 for financial support. The work of M.C. is partially supported by the french ANR grants DARKPHYS and PHYS@COL\\&COS. We thank the EU Marie Curie Research \\& Training network \"UniverseNet\" (MRTN-CT-2006-035863) for support.   \\bigskip  \\appendix  \\footnotesize  \\begin{multicols}{2}"
    },
    "0809/0809.2253_arXiv.txt": {
        "abstract": "We review current observations of the homogeneous cosmological expansion which, because they measure only kinematic variables, cannot determine the dynamics driving the recent accelerated expansion. The minimal fit to the data, the flat $\\Lambda CDM$ model, consisting of cold dark matter and a cosmological constant, interprets $4\\Lambda$ geometrically as a classical spacetime curvature constant of nature, avoiding any reference to quantum vacuum energy. (The observed Uehling and Casimir effects measure forces due to QED vacuum polarization, but not any  quantum material vacuum energies.)  An Extended Anthropic Principle, that Dark Energy and Dark Gravity be indistinguishable, selects out flat $\\Lambda CDM$.  Prospective cosmic shear and galaxy clustering observations of the growth of fluctuations are intended to test whether the 'dark energy' driving the recent cosmological acceleration is static or moderately dynamic.  Even if dynamic, observational differences  between an additional negative-pressure material component within general relativity (Dark Energy) and low-curvature modifications of general relativity (Dark Gravity) will be extremely small. ",
        "introduction": "CMB, baryon oscillation, supernova and weak lensing cosmographic measurements of the homogeneous expansion history of our universe are summarized in consistent with either a static finely-tuned cosmological constant or a dynamic mechanism, which may be material Dark Energy or low-curvature modifications of Einstein gravity (Dark Gravity). The homogeneous expansion measures only kinematic variables, in particular the evolution of the cosmological deceleration $q:=-a \\ddot{a}/\\dot{H}^2=d H^{-1}/dt-1$, which can be given a cosmological fluid description:  If the expansion rate $H^2:=(8 \\pi G_N/3)\\rho:=\\rho/3 M_P^2$, is used to define a cosmological 'mass density', 'enthalpy density', 'pressure' and 'equation of state' are also defined: \\be \\rho+P:=-2M_P^2 \\dot{H},\\qquad P:=-M_P^2 (3H^2+2\\dot{H}),\\qquad w:=P/\\rho=-(3 H^2+2\\dot{H})/3 H^2. \\ee In these definitions, $\\rho:=3 M_P^2 H^2$ can be separated into the ordinary cold dark matter component $\\rho_m$ and a 'dark energy' $\\rho_{DE}:=\\rho-\\rho_m$, which may be a new negative-pressure material Dark Energy or a Dark Gravity modification to Einstein general relativity.  The observed expansion history then fixes either a fine-tuned scalar field Dark Energy potential or a fine-tuned Dark Gravity modification to the GR Friedmann expansion rate \\cite{Blud}.  The observations of the homogeneous expansion history $H^2 (t)$ do not distinguish between Dark Energy and Dark Gravity and cannot fix the underlying dynamics driving the recent acceleration since redshift $z \\lesssim 0.46$. ",
        "conclusions": ""
    },
    "0809/0809.0898_arXiv.txt": {
        "abstract": "We have performed the largest ever particle simulation of a Milky Way-sized dark matter halo, and present the most comprehensive convergence study for an individual dark matter halo carried out thus far.  We have also simulated a {\\em sample} of 6 ultra-highly resolved Milky-way sized halos, allowing us to estimate the halo-to-halo scatter in substructure statistics.  In our largest simulation, we resolve nearly 300,000 gravitationally bound subhalos within the virialized region of the halo.  Simulations of the same object differing in mass resolution by factors up to 1800 accurately reproduce the largest subhalos with the same mass, maximum circular velocity and position, and yield good convergence for the abundance and internal properties of dark matter substructures.  We detect up to four generations of subhalos within subhalos, but contrary to recent claims, we find less substructure in subhalos than in the main halo when regions of equal mean overdensity are compared. The overall substructure mass fraction is much lower in subhalos than in the main halo. Extrapolating the main halo's subhalo mass spectrum down to an Earth mass, we predict the mass fraction in substructure to be well below 3\\% within 100 kpc, and to be below 0.1\\% within the Solar Circle.  The inner density profiles of subhalos show no sign of converging to a fixed asymptotic slope and are well fit by gently curving profiles of Einasto form. The mean concentrations of isolated halos are accurately described by the fitting formula of Neto et al.~down to maximum circular velocities of $1.5\\,{\\rm km\\,s^{-1}}$, an extrapolation over some 5 orders of magnitude in mass. However, at equal maximum circular velocity, subhalos are more concentrated than field halos, with a characteristic density that is typically $\\sim 2.6$ times larger and increases with decreasing distance from halo centre. ",
        "introduction": "\\label{intro}  \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: volker@mpa-garching.mpg.de}   A major puzzle in Cosmology is that the main matter component in today's Universe appears to be in the form of a yet-undiscovered elementary particle whose contribution to the cosmic density is more than 5 times that of ordinary baryonic matter \\citep[e.g.][]{Komatsu2008}.  This particle interacts extremely weakly with atoms and photons, so that gravity alone has affected its distribution since very early times. Recent observations have established a standard paradigm in which dark matter emerged from the early Universe with negligible thermal velocities and a gaussian and scale-free distribution of density fluctuations. In this ``Cold Dark Matter'' (CDM) hypothesis, quantum fluctuations during a very early period of cosmic inflation determine the statistics of the dark matter distribution at early epochs when the Universe was almost uniform \\citep{Guth1981,Starobinsky1982,Hawking1982,Bardeen1983}. Galaxies form from these initial conditions through the condensation of gas at the centres of a hierarchically aggregating population of quasi-equilibrium dark matter halos \\citep{White1978,White1991}.  When the effects of the baryons can be neglected, the nonlinear growth of dark matter structure is a well-posed problem where both the initial conditions and the evolution equations are known. This is an N-body problem {\\it par excellence}. The faithfulness of late-time predictions (which must be confronted directly with observation to test the paradigm) is limited purely by numerical technique and by the available computing resources.  Over the past two decades, numerical simulations have played a pivotal role in establishing the viability of the CDM paradigm \\citep[e.g.][]{Davis1985,Frenk1988, Warren1992,Gelb1994,Cen1994,Hernquist1996,Jenkins2001,Wambsganss2004}. They have led to the discovery of a universal internal structure for dark matter halos \\citep[][hereafter NFW]{Navarro1996,Navarro1997}, and they have supplied precise predictions for the expected large-scale structure of the Universe. The matter distribution on scales from $\\sim 50$ kpc to the size of the observable Universe and the galaxy population predicted by hierarchical CDM scenarios have been compared directly with a wide array of observations.  A recent example is the ``Millennium Run'' \\citep{Springel2005a}, still one of the largest cosmological simulations ever carried out, which followed the formation and evolution of over 10 million galaxies by post-processing the stored simulation outputs \\citep{Croton2006,Bower2006}.  So far, the $\\Lambda$CDM paradigm has passed these tests successfully, particularly those that consider the large-scale matter distribution.  Given this success in reproducing the large-scale structure of the Universe, it is important to test CDM predictions also on smaller scales, not least because these are sensitive to the nature of the dark matter.  Indeed, a number of serious challenges to the paradigm have emerged on the scale of individual galaxies and their central structure.  The realisation that CDM halos have cuspy dark matter density profiles led to a fierce debate about whether these are consistent with the rotation curves observed for low surface brightness, apparently dark matter dominated galaxies \\citep{Flores1994,Moore1994,McGaugh1998,deBlok2001,Hayashi2004,Hayashi2006}. The abundance of small dark matter subhalos predicted within CDM halos has also drawn much attention \\citep{Klypin1999,Moore1999b}. The total number is much larger than the number of known satellite galaxies surrounding the Milky Way, even accounting for the many recently discovered faint systems. It is still unclear whether this reflects an absence of the predicted low mass objects or merely the fact that no stars were able to form within them. The issue of dark matter substructure within halos is made more urgent by the prospect of observing dark matter particles in the near future, either by their annihilation radiation \\citep[e.g.][]{Bergstrom1998}, or through direct detection in experiments here on Earth \\citep[reviewed, e.g., by][]{Gaitskell2004}. For both, a precise and quantitative understanding of the small-scale dark matter distribution within our Galaxy is needed.     \\begin{table*} \\begin{tabular}{lccrrccccr} \\hline Name & $m_{\\rm p}$ & $\\epsilon$ & $N_{\\rm hr}$ & $N_{\\rm lr}$ & $M_{\\rm 200}$ & $r_{\\rm 200}$ & $M_{\\rm 50}$ & $r_{\\rm 50}$ &   $N_{\\rm 50}$\\\\ & $[{\\rm M}_\\odot]$ & $[{\\rm pc}]$ & & & $[{\\rm M}_\\odot]$ & $[{\\rm kpc}]$ & $[{\\rm M}_\\odot]$ &  $[{\\rm kpc}]$    \\\\ \\hline Aq-A-1   &  $1.712\\times 10^3$  & 20.5   & 4,252,607,000  &  144,979,154 & $1.839\\times 10^{12}$ &  245.76 & $2.523\\times 10^{12}$ &  433.48 & 1,473,568,512 \\\\ Aq-A-2   &  $1.370\\times 10^4$  & 65.8   &   531,570,000  &   75,296,170 & $1.842\\times 10^{12}$ &  245.88 & $2.524\\times 10^{12}$ &  433.52 &   184,243,536 \\\\ Aq-A-3   &  $4.911\\times 10^4$  & 120.5  &   148,285,000  &   20,035,279 & $1.836\\times 10^{12}$ &  245.64 & $2.524\\times 10^{12}$ &  433.50 &    51,391,468 \\\\ Aq-A-4   &  $3.929\\times 10^5$  & 342.5  &    18,535,972  &      634,793 & $1.838\\times 10^{12}$ &  245.70 & $2.524\\times 10^{12}$ &  433.52 &     6,424,399 \\\\ Aq-A-5   &  $3.143\\times 10^6$  & 684.9  &     2,316,893  &      634,793 & $1.853\\times 10^{12}$ &  246.37 & $2.541\\times 10^{12}$ &  434.50 &       808,479 \\\\ \\hline Aq-B-2   &  $6.447\\times 10^3$  & 65.8   &   658,815,010  &   80,487,598 & $8.194\\times 10^{11}$ &  187.70 & $1.045\\times 10^{12}$ &  323.12 &   162,084,992 \\\\ Aq-B-4   &  $2.242\\times 10^5$  & 342.5  &    18,949,101  &      648,874 & $8.345\\times 10^{11}$ &  188.85 & $1.050\\times 10^{12}$ &  323.60 &     4,683,037 \\\\ \\hline Aq-C-2   &  $1.399\\times 10^4$  & 65.8   &   612,602,795  &   78,634,854 & $1.774\\times 10^{12}$ &  242.82 & $2.248\\times 10^{12}$ &  417.09 &   160,630,624 \\\\ Aq-C-4   &  $3.213\\times 10^5$  & 342.5  &    26,679,146  &      613,141 & $1.793\\times 10^{12}$ &  243.68 & $2.285\\times 10^{12}$ &  419.36 &     7,110,775 \\\\ \\hline Aq-D-2   &  $1.397\\times 10^4$  & 65.8   &   391,881,102  &   79,615,274 & $1.774\\times 10^{12}$ &  242.85 & $2.519\\times 10^{12}$ &  433.21 &   180,230,512 \\\\ Aq-D-4   &  $2.677\\times 10^5$  & 342.4  &    20,455,156  &      625,272 & $1.791\\times 10^{12}$ &  243.60 & $2.565\\times 10^{12}$ &  435.85 &     9,579,672 \\\\ \\hline Aq-E-2   &  $9.593\\times 10^3$  & 65.8   &   465,905,916  &   74,119,996 & $1.185\\times 10^{12}$ &  212.28 & $1.548\\times 10^{12}$ &  368.30 &   161,323,676 \\\\ Aq-E-4   &  $2.604\\times 10^5$  &  342.5 &    17,159,996  &      633,106 & $1.208\\times 10^{12}$ &  213.63 & $1.558\\times 10^{12}$ &  369.14 &     5,982,797 \\\\ \\hline Aq-F-2   &  $6.776\\times 10^3$  & 65.8   &   414,336,000  &      712,839 & $1.135\\times 10^{12}$ &  209.21 & $1.517\\times 10^{12}$ &  365.87 &   223,901,216 \\\\ Aq-F-3   &  $2.287\\times 10^4$  & 120.5  &   122,766,400  &      712,839 & $1.101\\times 10^{12}$ &  207.15 & $1.494\\times 10^{12}$ &  363.98 &    65,320,572 \\\\ \\hline \\end{tabular} \\caption{Basic parameters of the Aquarius simulations. We have simulated 6 different halos, each at several different numerical resolutions. The leftmost column gives the simulation name, encoding the halo (A to F), and the resolution level (1 to 5).  $m_{\\rm p}$ is the particle mass, $\\epsilon$ is the  Plummer equivalent gravitational softening length, $N_{\\rm hr}$ is the number of high resolution particles, and $N_{\\rm lr}$ the number of low resolution particles filling the rest of the volume. $M_{\\rm 200}$ is the virial mass of the halo, defined as the mass enclosed in a sphere with mean density 200 times the critical value.  $r_{\\rm 200}$ gives the corresponding virial radius. We also give the mass and radius for a sphere of overdensity 50 times the critical density, denoted as $M_{\\rm 50}$ and $r_{\\rm 50}$. Note that this radius encloses a mean density $200$ times the {\\em background density}; in some studies \\citep[e.g.][]{Diemand2007} $M_{\\rm 50}$ and $r_{\\rm 50}$ have been defined as virial mass and radius.  Finally, $N_{\\rm 50}$ gives the number of simulation particles within $r_{\\rm 50}$. \\label{TabSims}} \\end{table*}  Our Aquarius Project addresses these questions by studying the highly nonlinear structure of Galaxy-sized CDM halos in unprecedented detail using state-of-the art numerical simulations.  We are particularly interested in the inner regions of these halos and of their substructures, where the density contrast exceeds $10^6$ and the astrophysical consequences of the nature of dark matter may be most clearly apparent. Quantifying such structure reliably through simulation is an acute challenge to numerical technique. In order to address this challenge, we use a newly developed version of our parallel simulation code, {\\small GADGET-3} \\citep[based on][]{Springel2001a, Springel2005b}, which allows us to cover an unprecedented dynamic range at high numerical accuracy. We carefully validate our simulation techniques and establish their range of numerical convergence through systematic convergence studies, thereby building confidence in the reliability of our results.   In order to evaluate the scatter in structural properties between halos, we have simulated six different systems at high resolution, each having between 160 and 224 million particles within $r_{\\rm 50}$, the radius with mean enclosed overdensity 50 times the critical value.\\footnote{We use this unconventional outer radius for our halos, rather than the standard $r_{\\rm 200}$, to facilitate comparison with \\citet{Diemand2007, Diemand2008} who quote results for their Via Lactea simulations within $r_{\\rm 50}$ although they refer to this radius as ``$r_{\\rm 200}$''.} Each of these simulations is better resolved than any previously published high-resolution halo simulation except the very recent `Via Lactea II' run of \\citet{Diemand2008} which has 470 million particles within $r_{\\rm 50}$. For one of our halos, we have increased the resolution by a further factor of 8, pushing the particle number within $r_{\\rm 50}$ to 1.47 billion. The gravitational softening length of this largest run is just $20.5\\,{\\rm pc}$. We collectively refer to this suite of simulations as the Aquarius Project.   In the present paper, we describe our simulation techniques and analyze a number of basic properties of our $z=0$ halos. We focus in particular on the abundance of substructures, their radial distribution, and their internal density profiles. We also give results for the concentration of substructures, and for the fraction of the mass they contain at different radii. In a companion paper \\citep{Springel2008}, we study implications for the detectability of dark matter annihilation within the Milky Way's dark matter halo, and in \\citet{Navarro2008} we study the structure of the central density cusps of the main halos. Future papers will study the evolution of our halos and their substructure.   This paper is organized as follows. In Section~\\ref{SecSims}, we introduce our simulation set and describe our numerical techniques. The abundance of dark matter substructures and their radial distribution within our halos are analyzed in Section~\\ref{SecSubAbundance}. Then, in Section~\\ref{SecSubsInSubs}, we turn to an analysis of the abundance of substructure within subhalos.  In Section~\\ref{SecSubInternalStructure}, we consider the density profiles of subhalos and their concentrations.  Finally, we summarize our conclusions in Section~\\ref{SecSummary}. ",
        "conclusions": "\\label{SecSummary}   In this paper, we have presented first results from the Aquarius Project, a Virgo Consortium\\footnote{The Virgo Consortium is an international collaboration of astronomers working on supercomputer simulations of cosmic structure formation, see http://www.virgo.dur.ac.uk} programme to carry out high-resolution dark matter simulations of Milky-Way-sized halos in the $\\Lambda$CDM cosmology.  This project seeks clues to the formation of galaxies and to the nature of the dark matter by designing strategies for exploring the formation of our Galaxy and its luminous and dark satellites, for searching for signals from dark matter annihilation, and for designing experiments for the direct detection of dark matter.  In our approach, we pay great attention to validating our numerical results to careful convergence studies. In addition, we explore possible uncertainties in predictions for the Milky Way resulting from the scatter in properties between otherwise similar halos. Thus, we simulate not just one realization at ultra high resolution, but rather a sample of (currently) 6 different halos.  Our ambition is to redefine the state-of-the-art in this field with respect to the accuracy of the cosmological N-body simulations, and the rigour with which quantitative statements about halo structure can be made.  Our new simulation code {\\small GADGET-3}, developed specifically for the Aquarius Project, is a highly efficient, massively parallel N-body code. It offers much better scalability to large numbers of compute cores and a higher basic speed than its parent code {\\small GADGET-2} \\citep{Springel2005b}.  It is able to cope efficiently with the tight coupling of around 1.5 billion particles in a single nonlinear object, split up across 1024 CPUs. Some of our simulations at resolution level 2 were run on an even larger number of compute cores, using up to 4096 cores of a Bluegene/P computer. Here we used a novel feature in {\\small GADGET-3} that can exploit additional compute cores in shared-memory nodes by means of threads (based on the POSIX pthreads library) yielding a mix of distributed and shared memory parallelism.  The ability to simulate this high degree of clustering and nonlinearity on massively parallel architectures is a prerequisite for exploiting the power of upcoming petaflop computers for the next generation of high-precision simulations of cosmological structure formation.    \\begin{figure} \\resizebox{8.0cm}{!}{\\includegraphics{fig29.eps}} \\caption{Subhalo concentration as a function of radius, for subhalos with maximum circular velocity larger than 2.5, 5, or $10\\,{\\rm km\\, s^{-1}}$. The dotted line is a fit to our result for the $5\\,{\\rm km\\, s^{-1}}$ sample, which yields $\\delta_V \\simeq 3.77\\times 10^6\\, (r/{\\rm kpc})^{-0.63}$. For comparison, we have also included the result quoted recently by  \\citet{Diemand2008} for the Via Lactea II simulation, which selected subhalos with $V_{\\rm max} > 5\\,{\\rm km\\, s^{-1}}$. Clearly, our subhalos are more concentrated than theirs at the same circular velocity. \\label{FigConcVsRadius}} \\end{figure}  The results presented above demonstrate that we have created a remarkably accurate set of simulations, reaching very good convergence for the dark matter density profile and the substructure mass function over the maximum range that could be expected. Even the location, mass and internal structure of individual large dark matter subhalos reproduce well between simulations of differing resolution, a level of convergence which exceeds anything previously reported in the literature.  The abundance of dark matter subhalos is remarkably uniform across our halo sample when normalized to parent halo mass, and when considering subhalos sufficiently small that fluctuations due to counting statistics are unimportant. The differential subhalo mass function is tilted to a slope slightly shallower than the critical value $-2$, so that, even when extrapolated to arbitrarily small masses, the total mass fraction in substructures remains small, less than 3\\% within 100 kpc of halo centre, and less than 20\\% within $r_{50}\\sim 400$~kpc. Adopting the logarithmically divergent slope $-2$ (which our results appear to exclude) does not increase these mass fractions by more than a factor of 2 or 3 for lower mass limits in the range $10^{-6}$ to $10^{-12},{\\rm M}_\\odot$, which plausibly correspond to the thermal free-streaming limit if the dark matter is the lightest sypersymmetric particle. The inner halo is dominated by a smoothly distributed dark matter component, not by substructure.  Independent of their present mass, substructures have a {\\em strong} preference to be found in the outer regions of halos. For example, we estimate that at most a fraction of $10^{-3}$ of the dark matter at the Solar circle is in bound subclumps. The rest is smoothly distributed.  Note, however, that this smooth component is expected to have a rich structure in velocity space, being composed of a large number (perhaps $10^5$ or more) of cold streams \\citep{Helmi2003,Vogelsberger2008}.  Contrary to previous claims, we find that substructure in subhalos is not a scaled-down version of substructure in main halos. Subhalos typically have less substructure than main halos. This is due to two causes.  Tidal stripping removes the outer substructure-rich parts of halos when they fall into a larger system and become subhalos. In addition, as the retained substructure ages it decreases in mass and number and is not replaced by the infall of new objects. As a result, the substructure mass fraction in subhalos is often much smaller than in main halos, particularly for subhalos in the inner regions which are the most heavily stripped and also, typically, the ``oldest''.  We have presented the first detailed convergence study of the shape of subhalo density profiles to be based on simulation sets where the {\\em same} subhalo can be identified in simulations of differing mass resolution. We find that the inner regions of subhalos, well inside their tidal truncation radii, can be well fit by NFW or Einasto profiles. Einasto fits are typically preferred, even when the shape parameter $\\alpha$ is fixed to a standard value, e.g., $\\alpha=0.18$. We have also studied how the local logarithmic slope of the density profile varies with radius, finding profiles to become gradually shallower towards the centre with no sign of approaching an asymptotic power-law behaviour. This is very similar to the behaviour of the central cusp in isolated dark matter halos \\citep{Navarro2004}. \\citep[We will address main halo cusps using the Aquarius simulations in a forthcoming paper,][]{Navarro2008}.  We find many subhalos for which the slope at the innermost converged point is substantially shallower than $-1.5$, and a few where it is shallower than $-1.2$. The Moore profile appears firmly excluded as a description of the inner regions of subhalos. It should {\\rm not} be used when modelling the cold dark matter annihilation signal, as in a number of recent papers \\citep[e.g.][]{Baltz2008}.   The concentration of subhalos is higher than that of halos of the same circular velocity or of the same mass in the field. This can be understood as a consequence of tidal truncation and mass loss \\citep{Kazantzidis2004,Bullock2005,Penarrubia2008} which lead to a larger reduction of $V_{\\rm max}$ than of $r_{\\rm max}$. Interestingly, we find that the relationship between $r_{\\rm max}$ and $V_{\\rm max}$ for {\\em field} halos is very well fit by the fitting function given by \\citet{Neto2007} for the Millennium simulation, even though this involves an extrapolation over many orders of magnitude towards lower mass. At the same maximum circular velocity, we find that the $r_{\\rm max}$ values of subhalos are, on average, only 62\\% of those of field halos.  We note that our results disagree with those of the recent Via Lactea I and II simulations \\citep{Diemand2007,Madau2008,Diemand2008,Kuhlen2008} on several important points.  We find substantially more substructure than reported for the Via Lactea simulations, and the discrepancy with Via Lactea I is larger than the expected halo-to-halo scatter, based on our own simulation set.  Our subhalos are also more concentrated than those found in the Via Lactea II simulation. We also differ in our conclusions about the amount of (sub-)substructure in subhalos, which we demonstrate to be less than predicted by the hypothesis that subhalos are tidally truncated, but otherwise scaled-down versions of field halos. Finally, we disagree with the claim of \\citet{Diemand2008} that subhalos have central power-law cusps with a mean slope of $-1.2$.  In future work, we will analyze the detailed formation history of the `Aquarius' halos and the evolution of their substructure. We will also build a new generation of semi-analytic models to follow the evolution of the baryonic component, and we will compare these with full hydrodynamical simulations of these same halos that we have already begun to carry out.  This should bring new insights into galaxy formation, and directly address possible small-scale challenges to the $\\Lambda$CDM theory. The verdict about whether CDM works on such scales is still pending."
    },
    "0809/0809.0851_arXiv.txt": {
        "abstract": "The coefficients defining the mean electromotive force in a Galloway--Proctor flow are determined. This flow shows a two-dimensional pattern and is helical. The pattern wobbles in its plane. Apart from one exception a circular motion of the flow pattern is assumed. This corresponds to one of the cases considered recently by Courvoisier, Hughes and Tobias (2006, Phys.\\ Rev.\\ Lett., 96, 034503). An analytic theory of the $\\alpha$ effect and related effects in this flow is developed within the second-order correlation approximation and a corresponding fourth-order approximation. In the validity range of these approximations there is an $\\alpha$ effect but no $\\gamma$ effect, or pumping effect. Numerical results obtained with the test-field method, which are independent of these approximations, confirm the results for $\\alpha$ and show that $\\gamma$ is in general nonzero. Both $\\alpha$ and $\\gamma$ show a complex dependency on the magnetic Reynolds number and other parameters that define the flow, that is, amplitude and frequency of the wobbling motion. Some results for the magnetic diffusivity $\\eta_{\\rm t}$ and a related quantity are given, too. Finally a result for $\\alpha$ in the case of a randomly varying flow without the aforementioned circular motion is presented. This flow may be a more appropriate model for studying the $\\alpha$ effect and related effects in flows that are statistical isotropic in a plane. ",
        "introduction": "In the astrophysical context, turbulent flows, e.g.\\ in stellar convection zones or in accretion discs and galaxies, are generally anisotropic and time-dependent. A simple model of a flow with such properties is that by \\cite{GP92}. This flow is two-dimensional, depends only on two Cartesian coordinates, e.g.\\ $x$ and $y$, which can simplify the analysis significantly, even in dynamo problems that are inherently three-dimensional. The Galloway--Proctor (GP) flow is related to a flow considered by Roberts (1972). The Roberts flow\\footnote{As usual, the term Roberts flow refers to the flow given by equation (5.1) of Roberts (1972).} is an early example of a spatially periodic flow that produces an alpha effect. The alpha term in the averaged form of the induction equation is crucial to model the generation of large-scale magnetic fields from small-scale helical fluid motions in stars and galaxies; see, for example, \\cite{Mof78}, \\cite{Par79}, and \\cite{KR80} for standard references. However, unlike the Roberts flow, the GP flow is time-dependent with a flow pattern wobbling in the $(x,y)$ plane in a circular fashion. Both the GP flow and the Roberts flow have a velocity component out of this plane such that the flow can be fully helical, i.e.\\ the velocity is proportional to its curl.  Particularly important is the dependence of the $\\alpha$ effect on the magnetic Reynolds number, $\\Rm$. While for the Roberts flow $\\alpha$ declines with $\\Rm$ in the large $\\Rm$ limit, in the case of the GP flow according to the results by \\cite{CHT06} (in the following referred to as CHT06) and \\cite{Cou08} there is a more complicated dependence on $\\Rm$ with sign changes and no indication of convergence with increasing $\\Rm$.  In many studies turbulent astrophysical flows have been modelled by random forcing. In the case of of helical isotropic turbulence such investigations show that $\\alpha$ approaches a finite value as soon as $\\Rm$ exceeds a value of the order of unity. This has been observed at least for Reynolds numbers up to 200 \\citep{Sur_etal08}.  The purpose of this paper is to study the effects of the GP flow in more detail in order to understand the influence of time-dependence and anisotropy on the value of $\\alpha$ and other turbulent transport coefficients. In particular, it is important to document the differences and similarities with turbulent flows that are statistically isotropic and irregular in space and time. We focus attention here on the simplest case considered in CHT06 with a flow being purely periodic in time and add a few results for a simple flow with random time dependence.  A number of similarities, but also some striking differences between turbulent flows and the Roberts flow are known. Similar in both flows is the fact that there is an $\\alpha$ effect whose magnitude increases with $\\Rm$ as long as the latter does not exceed some value in the order of unity. However, for larger $\\Rm$, the $\\alpha$ coefficient in isotropic turbulence settles to a constant value \\citep{Sur_etal08}, while for the Roberts flow $\\alpha$ tends to zero as $\\Rm \\to \\infty$ \\citep{Soward87,Soward89,Radler_etal02a,Radler_etal02b}. Furthermore, there is no $\\gamma$ effect, or pumping effect, neither for isotropic turbulence nor for the Roberts flow. On the other hand, in the time-dependent GP flow $\\gamma$ effects have been reported (CHT06). A $\\gamma$ effect corresponds to antisymmetric contributions of the $\\alpha$ tensor. This raises the question about the possible existence of antisymmetric contributions to the turbulent magnetic diffusivity tensor, or $\\eta_{\\rm t}$ tensor.  These aspects are now straightforward to address using the recently developed test-field method to calculate numerically all components of the $\\alpha$ and $\\eta_{\\rm t}$ tensors defining the mean electromotive force $\\meanEMF$ for a given flow field \\citep{S05,S07}. If, as we assume here, too, the mean magnetic field $\\meanBB$ depends only on one of the Cartesian coordinates, say $z$, only two $2 \\x 2$ tensors for $\\alpha$ and $\\eta_{\\rm t}$ are of interest. The test-field method has recently been used to calculate diagonal and off-diagonal components of $\\eta_{\\rm t}$ \\citep{Betal08}, the magnetic Reynolds number dependence of $\\alpha$ and $\\eta_{\\rm t}$ \\citep{Sur_etal08}, as well as their scale dependence \\citep{Betal08b}. We begin by exploring general properties of the mean electromotive force in the GP flow and present analytical results for coefficients like $\\alpha$ and $\\gamma$, which are crucial for the electromotive force, gained in the second--order correlation approximation and in a corresponding fourth--order approximation. After explaining the test--field method we give a series of numerical results for such coefficients, which are independent of approximations of that kind, and discuss them in detail. ",
        "conclusions": "\\label{Disc}  Our results for the flow of type (i) confirm the finding of CHT06 that both the $\\alpha$ and $\\gamma$ coefficients depend sensitively on $\\Rm$ and also on $\\epsilon$, and that even the signs of these coefficients may vary with these parameters. We have to add that $\\alpha$ and $\\gamma$ depend also on $\\tilde\\omega$, or $q = \\tilde\\omega \\, \\Rm$, that is, on parameters connected with the frequency of the wobbling motion, which CHT06 fixed in a special way without commenting on it, and that they show similar variations with these parameters. We found however rather regular behaviors of $\\alpha$ and $\\gamma$ for small and for large values of $\\epsilon$ and $\\tilde\\omega$.  It is sometimes considered as a rule that the sign of $\\alpha$ is opposite to that the mean helicity of the fluid flow and its modulus is proportional to that of the mean helicity. There is however no general reason for that kind of relation between $\\alpha$ and the kinetic helicity. We see only two limiting cases which allow simple statements on the sign of $\\alpha$.  Firstly, in the framework of SOCA applied to homogeneous turbulence and comparable flows it turns out that the sign of $\\alpha$ in the limit $q \\to 0$ is always opposite to that of $\\ol{\\ppsi \\cdot (\\nab \\x \\ppsi)}$, where $\\uu = \\nab \\x \\ppsi$, $\\nab \\cdot \\ppsi = 0$; see, e.g., \\cite{KR80,RB03}. For both types of flows, (i) and (ii), we have $\\ppsi = (\\tilde\\psi, - \\tilde\\psi, \\psi)$, where $\\tilde\\psi = - k_{\\rm H} (\\p \\psi/\\p x + \\p \\psi/\\p y)$ and therefore $\\ol{\\ppsi \\cdot (\\nab \\x \\ppsi)} = - 2 u_0^2 / k_{\\rm H}$. This implies that $\\alpha$ is positive. Indeed, only positive $\\alpha$ have been observed for small $q$, even beyond SOCA.  Secondly it was found in SOCA under the same conditions, but in the limit $q \\to \\infty$, that the sign of $\\alpha$ for a flow with finite correlation time is opposite to that of $\\int_0^\\infty \\ol{\\uu (\\xx, t) \\cdot (\\nab \\x \\uu (\\xx, t - \\tau)} \\, \\dd \\tau$; see, e.g., again \\cite{KR80}. A relation of that kind between $\\alpha$ and this integral can indeed be formally derived from the general relation (29) of \\cite{RR07} and applied to our specific situation. In case (i) the correlation time is however infinite and this integral does not converge. Although we know that $\\ol{\\uu \\cdot (\\nab \\x \\uu)} = - 2 u_0^2 k_{\\rm H}$ we do not see how reliable conclusions could be drawn concerning the sign of $\\alpha$ in the limit $q \\to \\infty$. In case (ii) the integral is positive, and indeed only positive $\\alpha$ have been observed.  Beyond the low and high conductivity limits, that is, for not too small or not too large values of $q$, even SOCA offers no simple general statements on the sign of $\\alpha$. In general $\\alpha$ may take both positive and negative values.  For studying $\\alpha$ and $\\eta_{\\rm t}$ with very simple flows it seems appropriate to consider flows of type (ii) rather than of type (i). In case (i) the results are influenced by the aforementioned circular motion of the flow pattern. As long as only $\\alpha$ and $\\eta_{\\rm t}$ should be discussed there is hardly a reason to introduce such a motion. We see no natural interpretation of it and so no interpretation of the so caused $\\gamma$ and $\\delta$ effects.  Recently \\cite{Tilgner08} pointed out that a time--dependent flow of a conducting fluid can act as a dynamo even when steady flows which coincide with it at any particular time cannot. He demonstrated this with a Roberts flow modified by a drift of its pattern so that the velocity $\\uu$ satisfies relations like \\eq{eq03} and \\eq{eq05} with $\\varphi_x = - k_{\\rm H} v_{\\rm d} t$ and $\\varphi_y = 0$, where $v_{\\rm d}$ is constant the drift velocity. Even if the intensity of the flow is too weak so that in the case $v_{\\rm d} = 0$ no growing solutions of the induction equation with a given period in the $z$ direction exist, such solutions may occur in an interval of some finite $v_{\\rm d}$. Although this flow considered by Tilgner is in a sense simpler than the flows in our paper, it shows no longer the symmetries with respect to the $z$ axis which we have utilized. As a consequence the relation between the mean electromotive force $\\meanEMF$ and the mean magnetic field $\\meanBB$ is more complex. In particular \\eq{eq22} and \\eq{eq23} no longer apply. Nevertheless the question arises whether the effect of the time--dependence of flows observed by Tilgner occurs also in the examples investigated here. In case (i) the parameter $\\omega$ could play the role of $v_{\\rm d}$. The fact that the magnitude of $\\alpha$ is larger for some finite $\\tilde{\\omega}$ than for $\\tilde{\\omega} = 0$, which can be seen in \\Figs{pscan_omega}{pscan_omega2}, points in this direction.  One of the original motivations for looking at the GP flow was the fact that it is time-dependent and in that sense closer to turbulent flows than time-independent flows. However, as we have shown here, the dynamo properties of the GP flow cannot be compared in a meaningful way with analytic theories or with simulations that apply to isotropic turbulence. Nevertheless, as shown in this paper, an analytic theory for the $\\alpha$ effect and other turbulent transport coefficients can be derived that matches numerical results in limiting cases.  Astrophysical flows can often neither be described by isotropic turbulence nor by wobbling two-dimensional flow patterns, but they are likely to contain aspects of both extremes. However, the present work highlights another aspect that may be of more general signi\\-fi\\-cance and concerns the turbulent transport properties in the presence of high-frequency time variability. This is not just a peripheral aspect of turbulence, but it is an additional property whose effects need to be understood more thoroughly. The situation is reminiscent of the modifications of mixing length theory in the presence of stellar pulsations \\citep[see, e.g.,][]{Gou77}. In dynamo theory the issue of high-frequency time variability has only recently been addressed. One example concerns the nonlinear $\\alpha$ effect where its time dependence has a striking effect on the behaviour of the mean field. In that example the temporal behaviour of the forcing function (delta-correlated or steady) determines the nonlinear asymptotic scaling behavior of the quenching function $\\alpha(\\meanBB)$ at low $\\Rm$. The early results of \\cite{Mof72} and \\cite{Rue74} suggested a $|\\alpha|\\sim|\\meanBB|^{-3}$ behavior, but in more recent years \\cite{FBC99} and \\cite{RK00} found instead a $|\\alpha|\\sim|\\meanBB|^{-2}$ behavior, which seemed in conflict with the earlier results. However, the work of \\cite{Sur_etal07} now shows that this is not just an artifact related to different approximations, for example, but it depends on whether or not the flow is time-dependent. They found that the $|\\alpha|\\sim|\\meanBB|^{-3}$ behavior is reproduced if the flow is steady, while the $|\\alpha|\\sim|\\meanBB|^{-2}$ behavior is obtained in the time-dependent case using a forcing function that is $\\delta$-correlated in time. Again, it is not clear which types of flows are more astrophysically relevant, but it is now clear that the detailed time-dependence of the turbulent flows can affect its transport properties in rather unexpected ways."
    },
    "0809/0809.1066_arXiv.txt": {
        "abstract": "NGC 3621 is a late-type (Sd) spiral galaxy with an active nucleus, previously detected through mid-infrared [\\ion{Ne}{5}] line emission. Archival \\emph{Hubble Space Telescope} (\\hst) images reveal that the galaxy contains a bright and compact nuclear star cluster.  We present a new high-resolution optical spectrum of this nuclear cluster, obtained with the ESI Spectrograph at the Keck Observatory.  The nucleus has a Seyfert 2 emission-line spectrum at optical wavelengths, supporting the hypothesis that a black hole is present.  The line-of-sight stellar velocity dispersion of the cluster is $\\sigmastar=43\\pm3$ \\kms, one of the largest dispersions measured for any nuclear cluster in a late-type spiral galaxy.  Combining this measurement with structural parameters measured from archival \\hst\\ images, we carry out dynamical modeling based on the Jeans equation for a spherical star cluster containing a central point mass.  The maximum black hole mass consistent with the measured stellar velocity dispersion is $3\\times10^6$ \\msun.  If the black hole mass is small compared with the cluster's stellar mass, then the dynamical models imply a total stellar mass of $\\sim1\\times10^7$ \\msun, which is consistent with rough estimates of the stellar mass based on photometric measurements from \\hst\\ images.  From structural decomposition of 2MASS images, we find no clear evidence for a bulge in NGC 3621; the galaxy contains at most a very faint and inconspicuous pseudobulge component ($M_K\\gtrsim-17.6$ mag).  NGC 3621 provides one of the best demonstrations that very late-type spirals can host both active nuclei and nuclear star clusters, and that low-mass black holes can occur in disk galaxies even in the absence of a substantial bulge. ",
        "introduction": "In the study of the demographics of supermassive black holes, one important unresolved problem is the observational census of black holes in very late-type disk galaxies.  The masses of supermassive black holes in elliptical and early-type spiral galaxies appear to be well correlated with the bulge properties of their host galaxies \\citep{kr95,fm00,geb00}.  These correlations naturally lead to the question of whether supermassive black holes can form in disk galaxies that lack bulges, and if so, how the black hole mass in bulgeless galaxies might be related to the properties of the host.  The nearest example of a bulgeless disk galaxy, M33, does not contain a massive black hole, with an extremely tight upper limit of $\\mbh \\lesssim 1500-3000$ \\msun\\ determined from stellar-dynamical observations and modeling \\citep{geb01, mfj01}.  It is not yet known whether the absence of a massive black hole is a typical property of bulgeless spirals in general, however, because very few late-type disk galaxies are near enough for stellar-dynamical measurements to yield such stringent constraints on the black hole mass.  It has long been recognized that many spiral galaxies contain photometrically and dynamically distinct central star-cluster nuclei, with the M33 nucleus being one of the best-studied examples \\citep[e.g.,][]{wal64,vdb76,ggm82,na82,oco83}.  The \\emph{Hubble Space Telescope} (\\hst) made it possible to survey the properties of nuclear star clusters in more distant galaxies \\citep[e.g.,][]{phil96, car97}, and \\hst\\ imaging programs have demonstrated that the majority of very late-type spirals (Hubble types Scd--Sd) contain compact star-cluster nuclei at or very close to their isophotal centers \\citep{mat99, bok02}.  Spectroscopic observations have revealed that the nuclear clusters in late-type spirals typically have stellar velocity dispersions of $\\sim15-35$ \\kms\\ \\citep{bok99,wal05}, effective radii of a few parsecs \\citep{bok04}, and dynamical masses of $\\sim10^6-10^7$ \\msun\\ \\citep{bok99,wal05}.  The spectra also reveal composite stellar populations indicative of multiple star formation episodes, with most clusters containing a population component with an age of $\\lesssim100$ Myr \\citep{wal06,ros06}.  The discovery of scaling relationships between nuclear cluster masses and host galaxy properties, both in spirals \\citep{ros06} and in dwarf ellipticals \\citep{wh06,fer06}, offers an intriguing hint of a connection between the processes that govern the growth of nuclear clusters and black holes.  \\citet{fer06} speculate that the formation of nuclear star clusters and massive black holes might be mutually exclusive, such that black holes are able to form only in galaxies above some critical mass.  However, the only nuclear star clusters in which the presence of a massive black hole can be ruled out at any significant level are in the Local Group or within a few Mpc.  Optical spectroscopic surveys have shown that active galactic nuclei (AGNs) do occur in some nuclear star clusters in nearby galaxies \\citep{seth08a}, so black holes and nuclear clusters can apparently coexist.  Evidence that black holes can occur in at least some very late-type disk galaxies comes from the detection of a small number of AGNs in Scd and Sd-type spirals.  The best example is the Sd galaxy NGC 4395, which contains a Seyfert 1 nucleus \\citep{fs89, fh03}; it remains the only clear identification of a broad-lined AGN in a bulgeless disk galaxy.  In addition, a few examples of Type 2 AGNs in very late-type spirals have been detected recently in optical spectroscopic surveys, such as NGC 1042 \\citep{seth08a,shi08} and UGC 6192 \\citep{bar08}.  In both NGC 4395 and NGC 1042 the AGNs occur in nuclear star clusters (it is not known whether UGC 6192 contains a nuclear cluster).  A recent \\emph{Spitzer} spectroscopic observation of the Sd galaxy NGC 3621 by \\citet{sat07} led to the discovery of an active nucleus, based on the detection of [\\ion{Ne}{5}] emission lines at 14.3 \\micron\\ and 24.3 \\micron.  Since photon energies greater than 95 eV are required for photoionization of Ne$^{+3}$ to Ne$^{+4}$, ordinary \\ion{H}{2} regions are not expected to be significant sources of [\\ion{Ne}{5}] emission, but a hard AGN continuum can easily provide the necessary ionizing photons.  Emission lines from a range of ionization states of neon are observable in the mid-infrared, and the relative strengths of these lines are useful as diagnostics of the ionization conditions within AGN narrow-line regions \\citep[e.g.,][]{sm92,v92,stu02,gds06}. \\citet{as08} used the results of new photoionization models to argue that the strength of the [\\ion{Ne}{5}] emission in NGC 3621 can only be plausibly explained by the presence of an AGN, and not by an ordinary burst of nuclear star formation.  The detection of an AGN in NGC 3621 is of significant interest because it is one of the few very late-type spirals known to host an active nucleus, making it an important target for further observations to constrain its black hole mass and AGN energetics.  \\citet{sat07} note that there is no previously published optical spectrum of the nucleus of NGC 3621 suitable for emission-line classification.  In this paper, we use archival \\hst\\ images to show that NGC 3621 contains a well-defined and compact nuclear star cluster.  A new optical spectrum of this star cluster is used to examine the classification of the active nucleus, and to measure the stellar velocity dispersion of the cluster.  We describe dynamical modeling of the nuclear cluster and the resulting constraints on the masses of both the cluster and the central black hole.  We also examine the structure of NGC 3621 using near-infrared images from 2MASS in order to search for a bulge component in this late-type galaxy.  For the distance to NGC 3621, we adopt $D=6.6$ Mpc, based on the Cepheid measurements of \\citet{fre01}.  At this distance, 1\\arcsec\\ corresponds to 32.0 pc. ",
        "conclusions": "Our main conclusions are as follows.  \\begin{enumerate}  \\item NGC 3621 has a Seyfert 2 optical spectrum, consistent with the AGN identification based on its mid-infrared spectrum by \\citet{sat07}.  This adds further support for the hypothesis that a black hole is present.  \\item For the nuclear star cluster, we find an effective radius of $r_e=4.1$ pc, a line-of-sight stellar velocity dispersion of $\\sigmastar=43\\pm3$ \\kms, and an absolute magnitude of $M_V = -12.3$ mag (uncorrected for internal extinction).  Simple estimates based on the available \\hst\\ photometry suggest a stellar mass of $(1-3)\\times10^7$ \\msun.  \\item From two-dimensional image decomposition of 2MASS images, we find that the galaxy is dominated by an exponential disk, with at most a very faint possible pseudobulge component being present.  We derive an upper limit of $M_K = -17.6$ mag to the luminosity of any bulge or pseudobulge component.  \\item From stellar-dynamical modeling of the nuclear cluster, we find an upper limit to the black hole mass of $\\mbh < 3\\times10^6$.  This result is insensitive to possible velocity anisotropy within the cluster.  The upper limit is based on the extreme assumption of $M/L=0$ for the stellar population.  With improved photometric and spectroscopic observations it would be possible to better constrain the stellar mass-to-light ratio and reduce the upper limit on \\mbh. Assuming that the cluster's black hole mass is small in comparison with its stellar mass, the dynamical models give a stellar mass of $\\sim1\\times10^7$ \\msun\\ for the cluster.  \\end{enumerate}  These results add to the emerging body of evidence that some late-type, bulgeless galaxies do contain low-mass central black holes, and that black holes can be found inside some nuclear star clusters. It remains the case, however, that there are only a few late-type spirals in which a the presence of a massive black hole has been either confirmed (based on nuclear activity) or ruled out (from dynamical modeling).  Further observations of the NGC 3621 nuclear cluster will lead to a more complete picture of the properties of this object.  The most pressing observational need is for X-ray observations, to provide more definitive confirmation of the AGN interpretation and a better estimate of the AGN luminosity.  Improved measurements of the stellar content of the nuclear cluster can be obtained from high-resolution, high-S/N spectra extending farther to the blue than our ESI data, in addition to improved multi-band \\hst\\ imaging covering the near-ultraviolet to near-infrared.  Future extremely large ground-based telescopes with adaptive optics may make it possible to spatially resolve the kinematic substructure of the nuclear cluster, and this could lead to dramatically improved constraints on the black hole mass."
    },
    "0809/0809.1585_arXiv.txt": {
        "abstract": "We study the correlation between different properties of bright ($L>L^{\\ast}$) galaxies in clusters and the environment in the Sloan Digital Sky Survey (SDSS). Samples of clusters of galaxies used in this paper are those selected by Coenda \\& Muriel that are drawn from the Popesso et al. and Koester et al. samples. Galaxies in these clusters have been identified in the Main Galaxy Sample of the Fifth Data Release (DR5) of SDSS. We analyse which galaxy properties correlate best with either, cluster mass or cluster-centric distance using the technique by Blanton et al. We find that galaxy properties do not clearly depend on cluster mass for clusters more massive than $M\\sim10^{14}M_{\\odot}$. On the other hand, galaxy properties correlate with cluster-centric distance. The property most affected by the cluster-centric distance is $g-r$ colour, closely followed by the $u-r$ colour. These results are irrespective of the cluster selection criteria. The two samples of clusters were identified based on the X-ray emission and the galaxy colours, respectively. Moreover, the parameter that best predicts environment (i.e. cluster-centric distance) is the same found by Mart\\'\\i nez \\& Muriel for groups of galaxies and Blanton at al. for the local density of field galaxies. ",
        "introduction": "Galaxy properties depend on environment, the high fraction of early type galaxies in rich cluster being one of the best examples. The fact that this fraction also evolves with time, i.e., an increasing of the S/S0 rate with redshift \\citep{dressler97,fasano00}, strongly suggests that galaxy morphologies are being altered by the physical processes that act in the cluster environment. Depending on the latter, i.e., low or high local density, different physical mechanisms will affect galaxy properties in different ways. The fact that most galaxy properties are correlated (morphology, luminosity, colours, etc), makes it difficult to know which properties are affected most by the environment. \\citet{blanton05} developed a test to evaluate which property, or pair of properties, are most predictive of the local density using galaxies in the Sloan SDSS \\citep{sdss}. \\citet{hm2} extended the analysis to galaxies in groups correlating galaxy properties with both, the mass of groups and the position in the system. \\citet{blanton05} and \\citet{hm2} found that galaxy colour is the property most predictive of the environment. This is particularly surprising taking into account the significant differences between field and group environments.  In high mass systems, the hot intracluster medium (ICM) becomes important. Mechanisms like ram pressure stripping \\citep{gg72} and starvation/strangulation can affect both the gaseous content of galaxies \\citep{abadi99} and the star formation history \\citep{fujita99}. From the dynamical point of view, high speed encounters between galaxies are more frequent, producing morphological transformations \\citep{moore98,moore99}. Galaxies can also suffer the stripping  of gas and stars due to the interaction with the cluster potential (e.g. \\citealt{moore99}).  It is not clear which of these, or some other proposed mechanisms, are dominant. The fact that different processes affect different galaxy components or properties, means that the correlation between these properties can also vary with the environment. Therefore, the most predictive galaxy properties could also depend on the type of system considered. \\citet{yo02}, \\citet{hm2}, \\citet{weinmann06} and \\citet{zmm06} found a clear dependence between the galaxy properties and the mass of the host group. For high mass clusters, the correlation is not clear. Several studies did not find dependence between the star formation rate or the fraction of early types with masses or velocity dispersions of clusters (see for instance \\citealt{goto03}). Nevertheless, \\citet{goto05} and \\citet{margo01} found indications of a correlation between blue galaxy fraction with  the cluster richness. More recently, \\citet{hansen07} have found that the fraction of red galaxies increases with the cluster mass, although only weakly for cluster more massive than $\\sim 10^{14} M_{\\odot}$. \\citet{popesso07} found that a luminous X-ray intracluster medium can affect the colour of galaxies. On the other hand, the local density-morphology \\citep{huhu31,oemler74,dressler80} or the cluster-centric distance-morphology relation \\citep{whitmore91} has been confirmed by several authors in different environments (see \\citealt{dominguez01} for a discussion between these two approaches).  In this paper we systematically explore the ability of different galaxy properties to predict the total mass of the cluster and the normalised cluster-centric distance. We have considered two samples of high mass clusters based on different selection criteria.  This paper is organised as follows: in section 2 we describe the sample of galaxies in clusters; the analyses of the dependence of galaxy properties on mass and on the cluster-centric distance are carried out in sections 3 and 4 respectively. We summarise our results and discuss them in section 5.  Galaxy magnitudes used throughout this paper have been corrected for Galactic extinction using the maps by \\citet{sch98}, absolute magnitudes have been computed assuming a flat cosmological model with parameters $\\Omega_0=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=100~h~{\\rm km~s^{-1}~Mpc^{-1}}$ and $K-$corrected using the method of \\citet{blanton03}~({\\small KCORRECT} version 4.1). All magnitudes are in the AB system. ",
        "conclusions": "In this paper we have extended the analyses by \\citet{blanton05} and \\citet{hm2} to galaxies in clusters. We analyse which galaxy properties correlate best with environment characterised by either cluster mass or cluster-centric distance. We use a sample of bright ($L>L^{\\ast}$) galaxies in clusters of galaxies in SDSS identified by two different criteria: X-ray selected clusters and clusters selected according to their red sequence. The two sub-samples of clusters have some differences in the mean properties. X-ray selected clusters tend to have a slightly higher mean mass and a higher fraction of red galaxies than maxBCG systems.  We find that the properties of bright galaxies do not clearly depend on cluster mass for systems more massive than $M\\sim10^{14}M_{\\odot}$. Although the mass range of our sample of cluster is not very large, the lack of dependence between mass and galaxy properties can be interpreted in terms of galaxy evolution. For systems with masses between $10^{14}M_{\\odot}$ and $10^{15}M_{\\odot}$ bright galaxies have experienced the same physical processes and therefore have similar properties. This result is consistent with \\citet{hansen07}. Using clusters and groups identified in SDSS with the MaxBCG finder of K07, these authors find that above a cluster mass $\\sim10^{14}M_{\\odot}$ the fraction of red galaxies increases weakly with increasing mass. For bright galaxies, the lack of a significant correlation between some galaxy properties such as colour and concentration with halo mass for halo masses above $\\sim 10^{14}M_{\\odot}$ is also present in \\citet{weinmann06} (their figure 11). To check consistency with their findings we have computed the median of $^{0.1}(g-r)$ colour and the median of the concentration parameter as a function of mass for our samples. We have found that they are remarkably independent on cluster mass, taking values $\\sim0.95$ and $\\sim3.0$, respectively, in fully agreement with the higher mass bins of \\citet{weinmann06}. For lower mass systems, the importance of processes like high speed encounters or those produced by the intra-cluster medium, tends to change faster with mass. \\citet{yo02} and \\citet{hm2} analysed groups of galaxies and found a clear dependence between mass and galaxy properties that tends to flatten as the mass of the systems grows. The higher masses considered by these authors are similar to the lower masses analysed in this work. All these results seem to imply that above a mass $\\sim10^{14}M_{\\odot}$ clusters have a similar population of bright ($L>L^{\\ast}$) galaxies, that arises as a result of the action of the same physical mechanisms with similar relative impact.  On the other hand, we find, as expected, that galaxy properties do correlate with cluster-centric distance. The property most affected by the cluster-centric distance is $g-r$ colour, followed closely by the $u-r$ colour. These results are irrespective of the cluster selection criteria. This is also in agreement with \\citet{hansen07}. Moreover, the parameter that best predicts the cluster-centric distance is the same found by \\citet{hm2} for groups of galaxies and \\citet{blanton05} for the local density of field galaxies as the most predictive of environment.  Galaxy parameters considered in this work can be classified into two classes. Those related to the physical properties of stars and those associated with the light distribution. In the first set are colour, absolute magnitude and spectral type. To the second group belong the concentration parameter, the surface brightness and the S\\'ersic index. Colour and spectral types of galaxies strongly depend on the age and metallicity of the stars as well as the present and recent star formation history. The luminosity of galaxies also depends on the same properties, although in a more indirect way. The fact that from field to massive clusters colour is the most sensitive property of galaxies to the present time environment, suggests that what really matters is the overall evolution of the environment where a galaxy and its progenitors form the stars. Moreover, according to \\citet{blanton05} and \\citet{hm2}, the other two parameters that appear as one of the pair of properties most predictive of the environment for field and group's galaxies are the magnitude and the spectral type, both belonging to the first group of parameters.  Among the phenomena that can affect the star formation history, we can mention the suppression or stimulation of the star formation due to interactions with the intergalactic medium or with other galaxies. Also, the past merger history of galaxies can play a fundamental role. In the particular case of the cluster environment, a natural segregation in the colour of galaxies arises as a result of cluster-centric gradient in the age of the stellar population as is observed in numerical simulations (e.g. \\citealt{gao}). Galaxies in the inner regions of clusters would have older stellar populations and therefore would be redder. These galaxies would also have had a longer time to deplete their gas reservoirs thus stopping star formation. In this scenario, the reddening as a function of radius emerges naturally. However, these processes might suffer a saturation resulting in a flattening of the mean colour as a function of mass.  The lack of a significant correlation between environment and the galaxy properties in the second set described above, may imply that phenomena like ram pressure stripping or galaxy harassment would have had a secondary role in the evolution of bright galaxies. Nevertheless, these physical processes would also have an impact in the first set of parameters. For example, ram pressure striping can remove an important fraction of the intragalactic gas producing a reddening of galaxies as a consequence of the reduction of the star formation rate. Numerical simulations show that galaxies can lose a high fraction of the gas after a single passage through the inner regions of a cluster (see for instance \\citealt{abadi99}). Therefore, above a certain mass ($\\sim10^{14}M_{\\odot}$), galaxies will experience the same physical processes acting with similar relative effectiveness thus producing a saturation in the mass-colour relationship. \\begin{table} \\center \\begin{tabular}{lcccc} \\hline & \\multicolumn{2}{c}{C-K07-I sample} & \\multicolumn{2}{c}{C-P04-I sample} \\\\ \\hline Property $X$ & $\\sigma_X^2-\\sigma^2$ & Significance& $\\sigma_X^2-\\sigma^2$ &Significance\\\\ \\hline $M_{r}^{0.1}$ & $-7$\t       &67\\%&  $-7$         &50\\%    \\\\ $(u-r)^{0.1}$ & $-13$ \t       &84\\%&  \\fbox{$-16$} &81\\%    \\\\ $(g-r)^{0.1}$ & $\\fbox{$-14$}$ &85\\%&  \\fbox{$-16$} &79\\%    \\\\ $\\mu_{50}$    & $-2$           &55\\%&  $-8$         &61\\%    \\\\ $C$           & $-5$           &63\\%&  $-9$\t  &64\\%    \\\\ $\\log(n)$     & $-4$           &60\\%&  $-3$\t  &49\\%    \\\\ $eclass$      & $-8$ \t       &70\\%&  $-7$\t  &69\\%    \\\\ \\hline \\end{tabular} \\label{tabR} \\caption{Galaxy parameters as cluster-centric distance indicators, i.e., in this case $\\delta\\equiv r/r_{200}$ (see text for details). Quoted $\\sigma^2$ values are expressed in units of $10^{-3}$. Boxes highlight the lowest values of $\\sigma_X^2-\\sigma^2$, i.e., those corresponding to the galaxy parameter that predicts best the cluster-centric distance. Quoted significances are assessed using the bootstrap technique as described in the text.} \\end{table}"
    },
    "0809/0809.4475_arXiv.txt": {
        "abstract": "We report on the discovery of a very bright $z = 2.00$ star-forming galaxy that is strongly lensed by a foreground $z=0.422$ luminous red galaxy (LRG).  This system was found in a systematic search for bright arcs lensed by LRGs and brightest cluster galaxies in the Sloan Digital Sky Survey Data Release 5 sample. Follow-up observations on the Subaru 8.2m telescope on Mauna Kea and the Astrophysical Research Consortium 3.5m telescope at Apache Point Observatory confirmed the lensing nature of this system. A simple lens model for the system, assuming a singular isothermal ellipsoid mass distribution, yields an Einstein radius of $\\theta_{\\rm  Ein}=3.82\\pm{0.03} \\arcsec$ or $14.8\\pm{0.1}h^{-1}$~kpc at the lens redshift.  The total projected mass enclosed within the Einstein radius is $2.10\\pm{0.03}\\times10^{12} h^{-1}  M_{\\sun}$, and the magnification factor for the source galaxy is $27 \\pm 1$. Combining the lens model with our $gVriz$ photometry, we find an (unlensed) star formation rate for the source galaxy of $32 \\ h^{-1} \\ M_\\sun \\ {\\rm yr}^{-1}$, adopting a fiducial constant star formation rate model with an age of 100 Myr and $E(B-V) = 0.25$.  With an apparent magnitude of $r = 19.9$, this system is among the very brightest lensed $z \\geq 2$ galaxies, and provides an excellent opportunity to pursue detailed studies of the physical properties of an individual high-redshift star-forming galaxy. ",
        "introduction": "Strong lensing systems provide the dual opportunity to study both the foreground mass distribution along the line of sight to the lens and the physical properties of the background object that is being lensed.  The latter is especially useful in studies of high-redshift galaxies, for which lensing provides a vital boost in the apparent brightness of these faint objects, which are otherwise difficult to study in detail.  For many years the $z = 2.72$ system cB58 \\citep{yee96} served as the prototypical lensed high-redshift Lyman break galaxy \\citep[LBG; e.g.,][]{steidel03}.  At $r = 20.4$, it is very bright and thereby allowed a number of detailed studies of the physical properties of a {\\em single} LBG to be carried out \\citep[e.g.,][]{pettini00,teplitz00}.  Recently a number of high-redshift lensed systems have been discovered, either serendipitously or in systematic searches, that are brighter than cB58 \\citep{smail07,belokurov07}, including the current record holder, the ``8~o'clock arc,'' at $r = 19.2$ \\citep{allam07}. These discoveries have been enabled by the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}, which provides the very large search area needed to systematically find these rare examples of extremely bright lensed high-redshift galaxies.    In this paper we report on the discovery of another remarkably bright ($r = 19.9$) strongly lensed $z = 2.00$ galaxy, the first system we have confirmed from a systematic search program for very bright lensed arcs that we are carrying out using the SDSS data. This paper is organized as follows: \\S\\ref{sec:search} describes the arc search and the discovery, \\S\\ref{sec:follow_up} describes the follow-up imaging and spectroscopy that led to confirmation of the system as a gravitational lens, \\S\\ref{sec:modeling} describes the modeling of the system including the photometry measurements, \\S~\\ref{sec:SFR} describes the source galaxy star formation rate measurements, and finally \\S\\ref{sec:conclusions} presents our conclusions. We assume a flat cosmology with $\\Omega_M =  0.3$, $\\Omega_{\\Lambda}=0.7$,   and $H_0=100h$~km~s$^{-1}$~Mpc$^{-1}$, unless otherwise noted. ",
        "conclusions": "\\label{sec:conclusions}  We have  reported on the discovery of the Clone system, consisting of a star-forming, BX/BM-type Lyman break galaxy in the SDSS at a redshift of $z=2.001$, which is strongly lensed by a foreground luminous red galaxy at a redshift of $z=0.422$.  The lensed galaxy is remarkably bright, and at $r = 19.9$ it is among the brightest known lensed source galaxies with $z \\geq 2$.  A simple SIE  lens model for the system  yields an  Einstein radius of $\\theta_{\\rm   Ein}=3.82\\pm0.03\\arcsec$  (or $R_{\\rm  Ein}= 14.8\\pm0.1$\\ $h^{-1}$~kpc at the lens redshift), a total  lensing  mass within the  Einstein radius of $2.10\\pm0.03\\times10^{12}\\  h^{-1} M_{\\sun}$, and a magnification factor for the lensed LBG of $27 \\pm 1$. Combining the lens model with our follow-up $gVriz$ photometry, we have also estimated the (unlensed) star formation rate (SFR) of the source galaxy to be $32 \\ h^{-1} \\ M_\\sun \\ {\\rm yr}^{-1}$, adopting a fiducial constant-SFR galaxy evolution model with an age of 100 Myr and $E(B-V) = 0.25$.  Such a star formation rate is similar to that found for samples of similar, but unlensed, $z \\sim 2$ BX/BM galaxies.  We are pursuing a number of further follow-up observations on this system, and we currently have optical and infrared data from HST Cycle 16, Spitzer Cycle 4, and Gemini North programs. The HST images indicate that the system is more complex than can be seen from the ground, and analysis of that higher resolution data will help us investigate the issues of substructure and image flux anomalies that we have encountered in the lens modeling described here.  Moreover, we will also use the SED information provided by the additional near-IR imaging and spectroscopy we are analyzing to better constrain the star formation history and dust content than we have been able to do here using just the optical data.  These more detailed analyses will be the subjects of future papers that will exploit the rich follow-up data set that can be derived from this very bright high-redshift lensing system."
    },
    "0809/0809.4369_arXiv.txt": {
        "abstract": "We investigate the evolution of dust formed in Population III supernovae (SNe) by considering its transport and processing by sputtering within the SN remnants (SNRs). We find that the fates of dust grains within SNRs heavily depend on their initial radii $a_{\\rm ini}$. For Type II SNRs expanding into the ambient medium with density of $n_{\\rm H,0} = 1$ cm$^{-3}$, grains of $a_{\\rm ini} < 0.05$ $\\mu$m are detained in the shocked hot gas and are completely destroyed, while grains of $a_{\\rm ini} > 0.2$ $\\mu$m are injected into the surrounding medium without being destroyed significantly. Grains with $a_{\\rm ini}$ = 0.05--0.2 $\\mu$m are finally trapped in the dense shell behind the forward shock. We show that the grains piled up in the dense shell enrich the gas up to 10$^{-6}$--10$^{-4}$ $Z_\\odot$, high enough to form low-mass stars with 0.1--1 $M_\\odot$. In addition, [Fe/H] in the dense shell ranges from $-6$ to $-4.5$, which is in good agreement with the ultra-metal-poor stars with [Fe/H] $< -4$. We suggest that newly formed dust in a Population III SN can have great impacts on the stellar mass and elemental composition of Population II.5 stars formed in the shell of the SNR. ",
        "introduction": "The first dust in the universe plays critical roles in the subsequent formation processes of stars and galaxies. Dust grains provide additional pathways for cooling of gas in metal-poor molecular clouds through their thermal emission and formation of H$_2$ molecules on the surface (e.g., Cazaux \\& Spaans 2004). In particular, the presence of dust decreases the values of the critical metallicity to $10^{-6}$--$10^{-4}$ $Z_\\odot$ (Omukai et al. 2005; Schneider et al. 2006; Tsuribe \\& Omukai 2006), where the transition of star formation mode from massive Population III stars to low-mass Population II stars occurs. Since absorption and thermal emission by dust grains strongly depend on their composition, size distribution, and amount, it is essential to clarify the properties of dust in the early epoch of the universe, in order to elucidate the evolutional history of stars and galaxies.  Dust grains at redshift $z > 5$ are considered to have been predominantly produced in supernovae (SNe). Theoretical studies have predicted that dust grains of 0.1--2 $M_\\odot$ and 10--60 $M_\\odot$ are formed in the ejecta of primordial Type II SNe (SNe II, Todini \\& Ferrara 2001; Nozawa et al. 2003) and pair-instability SNe (PISNe, Nozawa et al. 2003; Schneider et al. 2004), respectively. However, the newly formed dust is reprocessed via sputtering in the hot gas swept up by the reverse and forward shocks that are generated from the interaction between the SN ejecta and the surrounding medium. Thus, the size and mass of the dust can be greatly modified before being injected into the interstellar medium (ISM, Bianchi \\& Schneider 2007; Nozawa et al. 2007).  In this proceedings we present the results of the calculations for the evolution of newly formed dust within Population III SN remnants (SNRs), based on the dust formation model by Nozawa et al. (2003). We investigate the transport of dust and its processing by sputtering in the shocked hot gas, and report the size and amount of dust injected from SNe into the ISM. It is also shown that a part of the surviving dust grains are piled up in the dense SN shell formed behind the forward shock and can enrich the gas in the dense shell up to 10$^{-6}$--10$^{-4}$ $Z_\\odot$. We suppose that newly condensed dust in the SN ejecta has significant influences on the elemental abundances of Population II.5 stars, that is, the second-generation stars formed in the dense shell of Population III SNRs. ",
        "conclusions": ""
    },
    "0809/0809.3125_arXiv.txt": {
        "abstract": "Transport of angular momentum is one of the thrust areas of astrophysical flows. Instabilities and hence the turbulence generated by it has been invoked to understand its role in angular momentum transport in hydrodynamic and magnetohydrodynamic regime. Investigation of an unexplored region described by the parameter space  $\\Omega_e^2 < 0$, $\\Omega_e$ being the epicyclic frequency, resulted in a powerful magnetohydrodynamic instability. The growth rate of this instability is rather large and is found to depend on the wavenumber. ",
        "introduction": "Transport of angular momentum in various astrophysical flows is of great importance (\\cite{Lyn,Shakura}) in accretion theory. Nature of accretion process is largely understood through theoretical and computational studies. In astrophysical flows, outward transport of angular momentum of accreting material is the basic requirement to sustain the accretion disk. This angular momentum transport can be sustained by some turbulent mechanism operating in accretion disk. The origin of turbulence is unclear. Turbulent viscosity ($\\alpha$-viscosity) model by Shakura and Sunayev (\\cite{Shakura}) is one of such mechanisms. Differential rotation and the presence of magnetic field plays an important role in the dynamics of instability and hence in efficient angular momentum transport. In early $1990$s, Balbus \\& Hawley (\\cite{Balbus}) found that the subthermal fields in the presence of differential rotation can cause Magneto-Rotational Instability (MRI), which can provide the physical basis for the transport of angular momentum in accretion disk. Although the MRI requires only a small amount of ionization to work, it cease to exist when the strength of magnetic field increases. % Paper by Bonanno and Urpin (\\cite{BnU06}) (henceforth paper I) is an attempt to understand the transport of angular momentum by studying the instability arising in differentially rotating compressible flows in the presence of high magnetic field. The magnetic field had nonvanishing radial and azimuthal components. They found a new instability which survived for any value of magnetic field unlike MRI which survives only in the weak field limit. The maximum growth rate of the reported instability was $\\sim \\Omega$ where $ \\Omega$ is the rotation frequency.  This work investigates the magnetohydrodynamic instability in an relatively unexplored region in accretion disk described by $\\Omega_e^2<0$, $\\Omega_e$ being the epicyclic frequency.  For brevity, we explain the physical meaning of epicyclic frequency. Square of Epicyclic frequency is defined as $ \\Omega_e^2  =  4 \\Omega^2 (1 + \\frac{s}{2 \\Omega} \\frac{d \\Omega}{d s}$), s being the radial distance from the rotation axis, $\\Omega$ being the angular velocity of the fluid element in the flow. The physical meaning of the epicyclic frequency is explained as follows: Orbits of stars in a galaxy are usually not exact circles and always have a certain amount of randomness. If one of the perturbed orbits is observed on a framework rotating with the mean angular velocity around the center of the galaxy, it is known to be a small circle (\\cite {Bertin}). The frequency of such small cyclic motions that originate from randomness is called the epicyclic frequency.  Negative epicyclic frequency can be understood as follows: Square of epicyclic frequency is defined as \\begin{equation} \\Omega_e^2 = s^{-3} \\frac {d}{ds}(s^2 \\Omega)^2 \\end{equation} s is the radial distance from the rotation axis, $s^2 \\Omega$ is the specific angular momentum. An inward increase  in the specific angular momentum (i.e, outward decrease in angular momentum) implies that the square of epicyclic frequency is negative. A fluid element displaced outward under conservation of angular momentum has a larger angular momentum compared with that of the surrounding media. In this case,  the centrifugal force acting on the displaced element is stronger than that acting on the surrounding media, and the element goes further outward (unstable) (\\cite{Kato1}). In literature, it is known as a Reyliegh criterion for instablity.  It is well known that in a keplerian disk, the specific angular momentum increases with radius, i.e.,  $\\Omega_e^2>0$. According to the Rayleigh criterion, such a flow is stable. Generally in accretion disk, the region with $\\Omega_e^2>0$ do not exist but such regions are realized locally in the models by \\cite{Honma,Manmoto, Kato}. A recent work by \\cite{Gracia} has numerically shown that the transition region between ADAF and outer standard disk is time-dependent having $\\Omega_e^2<0$. They modelled the time-evolution of bimodal accretion discs where an outer cooling-dominated disc matches an inner ADAF type flow through a narrow transition region. The behaviour of the instability in such regions is explored.  This paper is orgnized as follows: The section 2 investigates the the growth rate of instability in the new region of parameter space. The section 3 deals with the Results and Discussions. Finally, we conclude by summarizing the new findings. ",
        "conclusions": "We revisit the problem of hydromagnetic instability in differentially rotating compressible flows and analyse the instability. The behaviour of hydromagnetic instability is investigated in the region described by $\\Omega_e^2<0$. Though $\\Omega_e^2<0$ is generally not satisfied in accretion disk but there are possibilities that $\\Omega_e^2<0$ is locally satisfied (\\cite{Gracia,Manmoto,Kato,Honma}) in the accretion disks. The present work will be applicable in those regions (\\cite{Kato2}).  We discover a range of parameter space described by $\\xi$, $\\zeta$ and $\\mu$ where the instability grows very fast, much faster than the reported growth rate of the instability in paper I. We have shown the growth rate of the instability in Fig. 3, for two representative values of $\\zeta$ for fixed $\\xi$. Growth rate  is $\\sim 1.8 \\Omega_e$, $ \\sim 2.0 \\Omega_e$ and $ \\sim 2.1 \\Omega_e$ at $x^2 = 300$, $600$ and $900$ respectively for $\\zeta= 0.5$ and the growth rate is $\\sim 2.6 \\Omega_e$ and $ \\sim 3.0 \\Omega_e$ and $ \\sim 3.4 \\Omega_e$ respectively for same values of $x^2$ for $\\zeta= 1.0$. The growth rate increases further with increasing value of $\\zeta$ for fixed values of $\\xi$. It is surprising to observe that the instability grows very fast and the growth rate, unlike in paper I is not independent of wavelength of perturbation, Infact, the instability grows with the increasing value of $x^2$.  In the case of $\\xi >> \\zeta$, which corresponds to the incompressible limit, one of the roots diminishes as was in paper I but the second root with postive real part survives with growth rate $ 0.547732  \\Omega_e$. Since incompressibility corresponds to $\\omega_0^2 \\xrightarrow{}$ $\\infty$, Equation (\\ref{e:eq3}) becomes \\begin{equation} \\sigma \\left (\\sigma ^2 -\\mu |\\Omega_e^2| \\right)= 0 \\end{equation} It is imperative to see that in the case of incompressible flow, Eq.(\\ref{e:eq3}) reduces to hydrodynamic relation. Since $\\mu = 0.3$, growth rate comes to $\\sqrt{\\mu} \\Omega_e$, i.e., $.0547732$. Hence, it is clear that the root which survives in the incompressible limit corresponds to the root satisfying Rayleigh criterion of instability in hydrodynamics."
    },
    "0809/0809.3914_arXiv.txt": {
        "abstract": "{Thermonuclear (type Ia) supernovae are explosions in accreting white dwarfs, but the exact scenario leading to these explosions is still unclear. An important step to clarify this point is to understand the behaviour of accreting white dwarfs in close binary systems. The characteristics of the white dwarf (mass, chemical composition, luminosity), the accreted material (chemical composition) and those related with the properties of the binary system (mass accretion rate), are crucial for the further evolution towards the explosion. An analysis of the outcome of accretion and the implications for the growth of the white dwarf towards the Chandrasekhar mass and its thermonuclear explosion is presented. }  \\FullConference{Supernovae: lights in the darkness (XXIII Trobades Cient\\'ifiques de la Mediterr\\`ania)\\\\ October 3-5 2007\\\\ Mao, Menorca, Spain}  \\begin{document} ",
        "introduction": "White dwarfs are the final stages of the evolution of stars with masses smaller than 8-11 M$_\\odot$. The final fate of such stars is just cooling to invisibility, whenever they are isolated. However, when they belong to interacting binary systems, more exciting phenomena can happen, like nova explosions and supernovae of the thermonuclear class, i.e., type Ia supernova explosions. Accretion of matter from the companion star is at the origin of such interesting phenomena, but there is a long and complicated path from the onset of accretion up to the final explosion, in case this is the final outcome of the evolution.  It is widely accepted that thermonuclear supernovae are the result of the explosion of a carbon-oxygen (CO) white dwarf, once it reaches the critical density for carbon ignition; this occurs either at the center of the star or off-center, always in strongly degenerate conditions. Therefore, it is expected that the properties of such explosions are quite uniform, in agreement with the observations, allowing the use of them as standard candles for cosmological purposes. In contrast, core colllapse supernovae, where a massive star explodes because of the gravitational force, show a wide range of observational properties, since the range of masses of the progenitor stars is extremely large.  In spite of the importance of type Ia supernovae both as cosmological tools, as crucial contributors to the chemical and dynamical evolution of the Galaxy, and as the triggering mechanism for star formation, their exact scenario is still unknown. In addition, there are not yet successful simulations of the explosion (see the review by \\cite{HN00}). In this paper we concentrate on the scenario issue, with a special emphasis on the fundamental role played by accretion, and its link to the final fate of the white dwarf. The type of material accreted (either H, He -with some metals- or CO) and the rate at which matter falls onto the white dwarf (mass accretion rate) are strongly dependent on the type of interacting binary hosting the potential exploding white dwarf. The mass and chemical composition of the white dwarf itself, as well as its initial luminosity, are also crucial for its outcome. There is a complicated interplay between type of accretion and initial white dwarf conditions, such that it is not straightforward to reach the carbon ignition conditions. Once ignition occurs, other problems occur, related to the flame propagation; these are out of the scope of this paper, but are treated in other papers of this volume (e.g., Roepke's contribution \\cite{Roepke}). ",
        "conclusions": ""
    },
    "0809/0809.3313_arXiv.txt": {
        "abstract": "Many important techniques for investigating the properties of extragalactic radio sources, such as spectral-index and rotation-measure mapping, involve the comparison of images at two or more frequencies.  In the case of radio interferometric data, this can be done by comparing the CLEAN maps obtained at the different frequencies. However, intrinsic differences in images due to the frequency dependence of the radio emission can be distorted by additional differences that arise due to source variability (if the data to be compared is obtained at different times), image misalignment, and the frequency dependence of the sensitivity to weak emission and the angular resolution provided by the observations (the resolution of an interferometer depends on the lengths of its baselines in units of the observing wavelength).  These effects must be corrected for as best as possible before multi-frequency data comparison techniques can be applied.  We consider the origins for the afore-mentioned factors, outline the standard techniques used to overcome these difficulties, and describe in detail a technique developed by us, based on the cross-correlation technique widely used in other fields, to correct for misalignments between maps at different frequencies. ",
        "introduction": "The radio emission of core-dominated, radio-loud Active Galactic Nuclei (AGN) is synchrotron radiation generated in the relativistic jets that emerge from the nucleus of the galaxy, presumably along the rotational axis of a central supermassive black hole. One important source of information about the physical conditions in the radio-emitting regions is the distribution of the spectral index $\\alpha$ over the source ($S_{\\nu}\\propto\\nu^{\\alpha}$, where $S_{\\nu}$ is the source flux at frequency $\\nu$). The core region is typically observed to be at least partially optically thick, with a nearly flat or inverted spectrum, while the jets are optically thin, with negative spectral indices. The spectrum may also flatten in regions of the jet in which there is re-acceleration of electrons or low-frequency absorption (e.g. Gabuzda, Pushkarev \\& Garnich 2000; Gabuzda, G\\'omez \\& Agudo 2001). Synchrotron radiation can be highly linearly polarised, to $\\simeq 75\\%$ in the case of a uniform magnetic field (Pacholczyk 1970), and linear polarisation observations can yield unique information about the orientation and degree of order of the magnetic field in the synchrotron source, as well as the distribution of thermal electrons and the magnetic-field geometry in the immediate vicinity of the AGN (e.g., via Faraday rotation of the plane of polarisation).  The compact radio emission of such AGN can be probed with  high resolution using Very Long Baseline Interferometry (VLBI). The radio telescopes in the interferometric array can be separated by hundreds or thousands of kilometres, making it infeasible to physically link (synchronise) them electronically, and high-accuracy timing signals must be recorded together with the data, so that the signals obtained at different antennas can be accurately synchronised during correlation. In practice, the amplitudes and, especially, phases of the measured complex visibility data unavoidably contain unknown errors, which can conveniently be expressed via antenna-based complex gain factors: \\begin{eqnarray*} V^{obs}_{ij} & = & G_iG_J^*V^{true}_{ij} \\end{eqnarray*} \\noindent where $V^{obs}_{ij}$ and $V^{true}_{ij}$ are the observed and true visibility functions on the baseline between antennas $i$ and $j$, and $G_i$, $G_j$ are the complex gain factors for antennas $i$ and $j$. The complex gains $G$ must be determined and removed from the data in order to achieve the highest-quality images possible for the radio-telescope array used. This is normally done iteratively, via alternate application of self-calibration (Fort and Yee 1976, Cotton 1979, Readhead \\& Wilkinson 1978, Readhead et al. 1980, Cornwell \\& Wilkinson 1981) and a deconvolution method, such as CLEAN (H\\\"ogbom 1974). ",
        "conclusions": "We have developed a C program to determine the shift between two VLBI images based on a cross-correlation analysis of the images. In the past, it has been necessary to determine such shifts for images of AGN by aligning compact optically thin jet components, which can be a somewhat subjective and not entirely unambiguous procedure. In addition, this method is difficult to apply to complex and/or extended AGN jets without compact optically thin features suitable for such an analysis. The great advantage of our new approach is that it provides a straightforward, objective means to determine the shift between two images that makes use of all optically thin regions in the source structure, not just individual chosen features. Our tests have shown that the program produces reliable shifts for images both with and without distinct optically thin features.  The code can be compiled using a standard C compiler, and has been designed to take input files written by the AIPS task IMTXT, and to output files that can be read by the AIPS task FETCH, making it straightforward for radio astronomers familiar with AIPS to implement the code. The code can be obtained by contacting D. C. Gabuzda (gabuzda@phys.ucc.ie)."
    },
    "0809/0809.1642.txt": {
        "abstract": "We model the temperature and chemical structure of molecular clouds as a function of depth into the cloud, assuming a cloud of constant density $n$ illuminated by an external FUV (6 eV $< h\\nu < 13.6$ eV) flux $G_0$ (scaling factor in multiples of the local interstellar field). Extending previous photodissociation region models, we include the freezing of species, simple grain surface chemistry, and desorption (including FUV photodesorption) of ices. We also treat the opaque cloud interior with time-dependent chemistry.  Here, under certain conditions, gas phase elemental oxygen freezes out as water ice and the elemental C/O abundance ratio can exceed unity, leading to complex carbon chemistry.  Gas phase H$_2$O and O$_2$ peak in abundance at intermediate depth into the cloud, roughly $A_V \\sim 3-8$ from the surface, the depth proportional to $\\ln (G_0/n)$.  Closer to the surface, molecules are photodissociated. Deeper into the cloud,  molecules freeze to grain surfaces.  At intermediate depths photodissociation rates are attenuated by dust extinction, but photodesorption prevents total freezeout. For $G_0 < 500$,  abundances of H$_2$O and O$_2$ peak at values $\\sim 10^{-7}$, producing columns $\\sim 10^{15}$ cm$^{-2}$, independent of $G_0$ and $n$. The peak abundances depend primarily on the product of the photodesorption yield of water ice and the grain surface area per H nucleus.  At higher values of $G_0$, thermal desorption of O atoms from grains enhances the gas phase H$_2$O peak abundance and column slightly, whereas the gas phase O$_2$ peak abundance rises to $\\sim 10^{-5}$ and the column to $\\sim 2 \\times 10^{16}$ cm$^{-2}$. We present simple analytic equations for the abundances as a function of depth which clarify the dependence on parameters.  The models are applied to observations of H$_2$O, O$_2$, and water ice in a number of sources, including B68, NGC 2024, and $\\rho$ Oph. ",
        "introduction": "Oxygen is the third most abundant element in the universe, after hydrogen and helium, so that a basic knowledge of oxygen chemistry in molecular clouds is essential in order to understand the chemical structure, thermal balance, and diagnostic line emission from star-forming molecular gas in galaxies. Early gas-phase chemical models (e.g.,  Langer \\& Graedel 1989, Millar 1990, Bergin et al 1998) predicted large abundances of H$_2$O ($\\sim 10^{-6}$) and O$_2$ ($\\sim 10^{-5}-10^{-4}$) relative to hydrogen nuclei in molecular gas well shielded from far ultraviolet (FUV, 6 eV $< h\\nu < 13.6$ eV) photons. If gas phase H$_2$O and O$_2$ were that abundant, they would be important coolants for dense gas (Goldsmith \\& Langer 1978, Hollenbach 1988, Neufeld, Lepp, \\& Melnick 1995).  However, the {\\it Submillimeter Wave Astronomy Satellite} (SWAS) made pointed observations in low energy transitions of ortho-H$_2$O (the 1$_{10}-1_{01}$ transition at 557 GHz) and O$_2$ (the 3$_3-1_2$ transition at 487 GHz) toward numerous dense (but unshocked) molecular cores and  determined that the {\\it line-of-sight-averaged} and beam-averaged (SWAS beam $\\sim 4'$) abundance of H$_2$O is of order $\\lta 3\\times 10^{-8}$ (Snell et al 2000) and that of O$_2$ is $\\lta 10^{-6}$ (Goldsmith et al 2000).  More recent observations by the {\\it Odin} mission set more stringent upper limits on O$_2$, $\\lta 10^{-7}$ (Pagani et al 2003), with a reported detection at the $\\sim 2.5\\times 10^{-8}$ level in $\\rho$ Oph (Larsson et al 2007). Although the water abundance derived from the observed water emission depends inversely on the gas density, and therefore is somewhat uncertain, understanding the two order of magnitude discrepancy between the gas phase chemical models and the observations is essential to astrochemistry and to the basic understanding of the physics of molecular clouds.  Previous attempts to explain the low abundance of H$_2$O and O$_2$ observed by SWAS showed that time dependent gas-phase chemistry by itself would not be sufficient (Bergin et al 1998, 2000).  Starting from atomic gas, a dense ($n(H_2) \\sim 10^5$ cm$^{-3}$) cloud only took $\\gta 10^4$ years to reach H$_2$O abundances $\\sim 10^{-6}$, close to the final steady-state values and much greater than observed.  The best previous explanation involved time-dependent chemistry linked with the freezeout of oxygen species on grain surfaces and the formation of substantial water ice mantles on grains (Bergin et al 2000).  In these models, all the oxygen in the molecular gas which is {\\it not} tied up in CO adsorbs quickly (in $\\sim 10^5$ years for a gas density of $\\sim 10^4$ cm$^{-3}$) to grain surfaces, forms water ice and remains stuck on the grain as an ice mantle. As a result, the gas phase H$_2$O abundance drops from  $\\sim 10^{-6}$  to $\\sim 10^{-8}$.  The grains are assumed too warm for CO to freeze as a CO ice mantle, so that for a period from about $10^5$ years to about $3\\times 10^6$ years, the gas phase H$_2$O abundance remains at the $\\sim 10^{-8}$ level, fed by the slow dissociation of CO into O and the subsequent reaction of some of this O with H$_3^+$, which then ultimately forms gas-phase H$_2$O. The slow dissociation of CO is driven by cosmic ray ionization of He; the resultant He$^+$ reacts with CO to form O and C$^+$. After  $\\sim 3\\times 10^6$ years, even the gas phase CO abundance drops significantly as the dissociation process bleeds away the oxygen which ends up as water ice on grain surfaces.  Therefore, after about $3\\times 10^6$ years, the gas phase H$_2$O also drops from $\\sim 10^{-8}$ to $ < 10^{-9}$.  This previous scenario  explained the observed H$_2$O emission as arising from the central, opaque regions of the cloud, where the abundance has dropped to the observed values but have not had time to reach the extremely low steady state values.  The model relied, then, on a ``tuning'' of molecular cloud timescales so that they are long enough for the freeze-out of existing gas phase oxygen not in CO  onto grains, but not so long that the CO is broken down and the resultant O converted to water ice, which would cause the gas phase H$_2$O abundances to drop below the observed values.\\footnote{For another similar model which relies on time-dependent effects, coupled with grain freezeout and grain surface processes, to explain the low O$_2$ and H$_2$O abundances observed by SWAS, see Roberts \\& Herbst (2002).} The model also relied on CO {\\it not} freezing out on grains in the opaque cloud centers.   To add to the problems of modeling the SWAS data, the SWAS team observed a number of strong submillimeter continuum sources such as SgrB, W49 and W51 and found the 557 GHz line of H$_2$O in {\\it absorption}, as the continuum passed through translucent clouds ($A_V \\sim 1-5$) along the line of sight (Neufeld et al 2002, Neufeld et al 2003, Plume et al 2004). The absorption measurement provided an even better estimate of the H$_2$O column, $N$H$_2$O, through these clouds because the absorption line strengths are only proportional to $N(H_2O)$, whereas the emission line strengths are proportional to $n(H_2)N(H_2O) e^{-27 {\\rm K}/T}$ since the line is subthermal, effectively optically thin, and lies $\\Delta E/k = 27$ K above the ground state.  Therefore, to obtain columns from emission line observations requires the separate knowledge of both the gas density and the gas temperature. The absorption measurements showed column-averaged abundances of H$_2$O of $\\sim 10^{-7}-10^{-6}$ with respect to H$_2$,  using observed CO columns and multiplying by $10^4$ to obtain H$_2$ columns. In the context of the time-dependent models with freezeout, the difference between the abundances measured in absorption compared with those measured in emission was attributed to the presumed lower gas densities in the absorbing clouds compared to the emitting clouds.  Because the freezeout time depends on $n^{-1}$, it was assumed that the lower density clouds had not had time to freeze as much oxygen from the gas phase, so that gas phase H$_2$O abundances were higher.  A variation of this model is that of Spaans \\& van Dishoeck (2001) and Bergin et al (2003), where it was noticed that the water emission seemed to trace the photodissociation region (PDR, see Hollenbach \\& Tielens 1999), which lies on the surface ($A_V \\lta 5$) of the molecular cloud.  A two component model was invoked in which the water froze out as ice in the dense clumps, but remained relatively undepleted in the low density interclump gas because of the longer timescales for freezeout.  Thus, the average gas phase water abundance was derived from a mix of heavily depleted and undepleted gas.  We propose in this paper a new model of the H$_2$O and O$_2$ chemistry in a cloud.  In this model, we assume that the molecular cloud lifetime is sufficiently long to allow freezeout\\footnote{Indeed, the freezeout of volatile molecules such as CO is now commonly observed in the dense, cold regions of molecular clouds (e.g., Bergin \\& Tafalla 2007)} and reduce the gas phase H$_2$O and O$_2$ abundances to very low values in the central, opaque regions of the cloud. The key to understanding the H$_2$O emission and the O$_2$ upper limits observed by SWAS and {\\it Odin} is to model the {\\it spatially-dependent} H$_2$O and O$_2$ abundances through each cloud.  At the surface of the cloud the gas-phase H$_2$O and O$_2$ abundances are very low because of the photodissociation by the ISRF or by the FUV flux from nearby OB stars.  Near the cloud surface the dust grains have little water ice because of photodesorption by the FUV field.  Deeper into the cloud, the attenuation of the FUV field leads to a rapid rise of the gas-phase H$_2$O and O$_2$ abundances, which peak at a hydrogen nucleus column $N_{f} \\sim 10^{21.5} - 10^{22}$ cm$^{-2}$ (or depths $A_{Vf} \\sim $ several into the cloud) and plateau at this value for $\\Delta A_V \\sim$ several beyond $A_{Vf}$, insensitive to the gas density $n$ and  the incident FUV flux $G_0$ (scaling factor in multiples of the avearage local interstellar radiation field).  At these intermediate depths, photodesorption of some of the water ice by FUV photons keeps the gas phase water abundance high (the ``$f$'' in $N_f$ and $A_{Vf}$ signifies the onset of water ice freezeout, as will be discussed in \\S 2.5 and \\S3.3).  The FUV is strong enough to keep some of the ice off the grains, but, due to efficient water ice formation followed by photodesorption, it is not strong enough to dissociate all gas-phase water.  At greater depths, gas phase reactions other than photodissociation begin to dominate gas phase H$_2$O destruction, and the steady state gas phase H$_2$O and O$_2$ abundances plummet as all the gas phase elemental O is converted to water ice.  Thus, the overall structure of a molecular cloud has three subregions: at the surface is a highly  ``photodissociated region'', at intermediate depths is the ``photodesorbed region'', and deep in the cloud is the ``freezeout region''.  In this new model, the H$_2$O and O$_2$ emission mainly arise from the layer of high abundance gas in the plateau which starts at $N_{f}$ and extends somewhat beyond, so that the emission is a (deep) ``surface'' process rather than a ``volume'' process. For clouds with columns $N>N_{f}$, the H$_2$O and O$_2$ emission becomes independent of cloud column, assuming the incident FUV field and gas density are fixed. The model gives {\\it column-averaged} abundances through the dense cores of $\\lta  10^{-8}$ for H$_2$O and O$_2$ for low values of $G_0 < 500$, but the local abundances peak at values which are at least an order of magnitude larger than these values.  For cloud columns $N>N_{f}$, the average H$_2$O abundance scales as $N^{-1}$.  The implications of this scenario for molecular clouds may be much broader than this particular model of H$_2$O and O$_2$ chemistry in clouds.  Other molecules such as CO, CS, CN, and HCN require a spatial model of their distribution in a cloud, with photodissociation and photodesorption the dominant processes near the cloud surface, and the freezing out of the molecules the dominant process deeper into the cloud.  In addition, the adsorption process and creation of abundant ice mantles changes the relative gas phase abundances of the elements. In the case considered in this paper, the C/O ratio in the gas may go from 0.5 at the cloud surface to unity or greater in the H$_2$O freezeout region. %, since the CO may not freeze but all the rest of the oxygen gas phase %species, including some that originally was in the form of CO,  may adsorb, %turn to water ice, and freeze on grain surfaces. Such changes in the gas phase C/O ratio have major implications on {\\it all} gas phase chemistry (e.g., Langer et al 1984, Bergin et al 1997).  The interesting and important caveat to this relatively simple steady state model is the effect of time-dependence in raising abundances of, for example, gas phase H$_2$O, O$_2$, and CO above steady state values in the opaque ($A_V > 5-10$) interiors of clouds.  One such effect is that at low densities, $n\\lta 10^3$ cm$^{-3}$, species do not have time to freeze out within cloud lifetimes and therefore have much higher gas phase abundances. We have also uncovered a new {\\it time-dependent} process that may elevate the H$_2$O and O$_2$ abundances for times $t\\lta 10^7$ years in the freezeout region at very high $A_V$ {\\it even at high cloud densities} $n\\sim 10^5$ cm$^{-3}$. If the grains are sufficiently cold to freeze out CO ($T_{gr} \\lta 20$ K or $G_0 \\lta 500$) at high $A_V$, a CO/H$_2$O ice mix rapidly ($t \\lta 10^{5-6}$ years) forms. The steady state solution has very little CO ice with most of the O in H$_2$O ice. However, the time to convert CO ice to H$_2$O ice is very long, and the bottleneck is the cosmic ray desorption of CO from the CO/H$_2$O ice mixture. This desorption provides gas phase CO which then acts as a reservoir to produce gas phase H$_2$O and O$_2$, until all the oxygen eventually freezes out as water ice (i.e., the same mechanism as described in Bergin et al (2000) except the gas phase CO in this case comes from the CO ice). Depending on the assumptions regarding the CO desorption process, this cosmic ray CO desorption timescale may range from $10^5$ to $>>10^7$ years.  If the timescale is roughly 0.1 to 1 times  the cloud age, a maximum gas phase H$_2$O and O$_2$ abundance is produced at that time. Although not as large as the peak abundances produced at $A_{Vf}$, these abundances can be significant and can contribute to the total column of these species if the cloud has a high total column ({\\it but only if the CO desorption time is 0.1 to 1 times the cloud age}).  This paper is organized as follows.  In \\S 2 we describe the new chemical/thermal model of an opaque molecular cloud illuminated by FUV radiation.  The major changes implemented in our older photodissociation region (PDR) models (Kaufman et al 1999) include the adsorption of gas species onto grain surfaces, chemical reactions on grain surfaces, and the desorption of molecules and atoms from grain surfaces. In \\S 3 we show the results of the numerical code as functions of the cloud gas density $n$, the incident FUV flux $G_0$, and the grain properties.  We  present a simple analytical model that explains the numerical results and how the abundances scale with depth and with the other model parameters. We discuss time dependent models for the opaque cloud center.  In \\S 4 we apply the numerical model to both diffuse clouds and dense clouds. Finally, we summarize our results and conclusions in \\S 5.  For the convenience of the reader, Table 3 in the Appendix lists the symbols used in the paper. ",
        "conclusions": "%Summary of model and analytic scalings  %In the Bergin et al 2000 paper, there were requirements on a successful model. %We should apply this criteria to our model.  %Comparison of models with observations  %Mystery of oxygen resolved  %Observational predictions  %We may discuss limb brightening here, and not in section 3.  %Implications for interstellar chemistry in molecular clouds: %dust vacuums the condensibles, chemistry is dominated by grain %surface processes and desorption processes. Gas in interior %provided by desorption in t dependent fashion.  We have constructed both a detailed numerical code and simple analytical equations to follow the abundances of the main oxygen-bearing species as a function of depth in molecular clouds.  We take clouds of constant density $n$, illuminated by an external FUV (6 eV$< h\\nu < 13.6$ eV) field characterized by the unitless parameter $G_0$ normalized to the local (Milky Way) interstellar FUV field.  We explore the range 10$^3$ cm$^{-3} < n < 10^6$ cm$^{-3}$ and $1 < G_0 < 10^5$.  Our models incorporate thermal balance, gas phase chemistry, a simple grain surface chemistry that includes the formation of H$_2$, H$_2$O, and CH$_4$ on grain surfaces, the sticking of gas phase species to grain surfaces, and various desorption processes such as photodesorption, thermal desorption, and cosmic ray desorption.  The abundances of chemical species vary considerably as a function of depth into a molecular cloud.  At the surface, molecules are photodissociated and the gas is primarily atomic (e.g., O) or, for species with ionization potentials less than 13.6 eV, singly ionized (e.g., C$^+$). Here, although dust grains may be cold enough to prevent rapid {\\it thermal} desorption, ices do not form because {\\it photodesorption} keeps the refractory dust clear of ice mantles (the same is true of dust in the diffuse ISM). Very deep in the cloud, again assuming grains cold enough to prevent thermal desorption, photodesorption is negligible and ices form, causing gas phase abundances of condensibles to plummet. Here, time dependent effects (generally associated with cosmic rays) dominate the gas phase chemistry due to the time required to deplete condensibles onto grains, to desorb  CO ice by cosmic rays, and to destroy gas-phase CO (by He$^+$ produced by cosmic rays).  These processes are slow, and have timescales comparable to, or longer than, typical molecular cloud lifetimes.  At intermediate depths, photodesorption of ices supply gas phase elements, and the partial shielding of the photodissociating external FUV flux allows gas phase molecular abundances to peak.  We have focused in this paper on oxygen chemistry, and followed in particular the gas phase species O, OH, O$_2$, and H$_2$O as well as water ice.  Our answer to the longstanding astrochemical question, ``Where is elemental O in molecular clouds?'', is that oxygen not in CO is primarily gas phase atomic O to a depth $A_{Vf}$ (see Eq. 23), and is water ice at larger depths. The freezeout depth $A_{Vf}$ where significant water ice forms is proportional to $\\ln (G_0/n)$, and is typically 3-6. Refractory silicate grains contain a modest abundance of O ($\\sim 10^{-4}$), but the bulk is likely gas phase O and CO along with CO ice and H$_2$O ice. For example, in our standard model, integrating to $A_V= 10$, the O not in refractory grains is divided $56$\\% H$_2$O-ice, $25$\\% gas phase O, $10$\\% CO-ice, and $8$\\% gas phase CO.  Our models also give an indication of the origin of water in molecular clouds.  In the steady state standard model, at $A_{Vf}$ the formation rate of gas phase H$_2$O is 2\\% via gas phase reactions and 98\\% by photodesorption of H$_2$O that has been formed on grains.  At $A_{Vd}$ the percentages are 30\\% and 70\\%, respectively; midway in the plateau, they are 8\\% and 92\\%, respectively.\\footnote{A warning to future modelers: these percentages are not trivial to compute since there is a great deal of cycling of H$_2$O and OH to H$_3$O$^{+}$ (via reaction with H$_3^+$) and then back to H$_2$O and OH (via dissociative electronic recombination).  This cycling gives, of course, no net production of these molecules.} These percentages depend on $G_0$ and $n$, but as long as $G_0 \\lta 500$, these numbers are representative. At higher $G_0>500$, where O atoms thermally desorb from grains before forming OH and eventually water ice, most of the water production is via gas phase reactions. The production of water ice is about 3 times greater than the production of gas phase water by photodesorption of water ice, because 2/3rds of the water ice is photodesorbed as OH.  Thus, one needs to modify these numbers if one wants to estimate the formation of water in any form in a cloud: for example, in the standard case midway in the plateau the percentages are 3\\% by gas phase reactions and 97\\% by grain surface reactions.  We plan a follow-up paper to examine carbon chemistry.  We note that Figure 14 shows that for all densities considered and by $t> 2\\times 10^6$ years, gas phase CO is largely absent in cloud interiors. This is in rough agreement with observations that show CO depletions in dense cloud interiors.  Thus, CO is not always a volume tracer. However, the sequel paper will include more grain surface carbon chemistry such as the formation of CO$_2$ and methanol on grain surfaces. In addition, we will examine the effects of initial conditions on time-dependent chemistry, and compare observations with a wide variety of models with varying cosmic ray desorption rates and gas density profiles through the cloud.  For $G_0 \\lta 500$, the abundances of gas phase OH, O$_2$ and H$_2$O peak and plateau at values $\\sim 10^{-7}$ between $A_{vf}$ and $A_{Vd}$ (see, Eqs. 26, 31, and 33). These abundances are independent of $n$ and $G_0$ because both their formation rate (ultimately from photodesorption of H$_2$O ice) and their destruction rate (from photodissociation) both depend on the product $nG_0$. However, since their formation is linked with H$_2$O ice photodesorption, these plateau abundances do increase with increasing ice photodesorption yield, $Y_{H_2O}$ and $Y_{OH,w}$, and also with the grain cross sectional area per H nucleus $\\sigma _H$ (see same equations).  For higher values of $G_0 \\gtrsim 500$, the chemistry undergoes a change, as O atoms no longer form H$_2$O ice on grain surfaces.  Here, the grains are warm enough ($\\gta 20$ K) that O atoms thermally desorb before reacting with H atoms.  This pushes the freezeout to greater depths and enhances the abundances of gas phase O, OH, O$_2$ and H$_2$O by keeping more of the elemental oxygen in the gas phase.  Because O$_2$ formation goes as the product of O and OH abundances, O$_2$ is the most sensitive to this change and its local abundance reaches $10^{-5}$, while its column climbs to $\\sim 2\\times 10^{16}$ cm$^{-2}$.  We note that the observations of Pagani et al (2003) may suggest a higher binding energy for O on grains than adopted, and in this case the critical $G_0$ where this occurs may be significantly higher than 500.  Deep in the cloud interior at $A_V >> A_{Vf}$,  the gas phase abundances drop as ices become dominant, the exact values of the abundances depending on cloud lifetimes and the cosmic ray desorption rates assumed (see Figures 14 and 15).  Here, time dependent chemistry dominates.  Poelman, Spaans, \\& Tielens (2007) have suggested that the SWAS results do not rule out substantial gas phase H$_2$O abundances and columns in the centers of clouds, because of optical depth effects which ``hide'' the H$_2$O in the centers.  However, as this saturation effect occurs, the line center antenna temperature approaches the gas temperature of the emitting region.  In our model, this gas temperature is typically 10-20 K. Even in the cold interior of clouds, one would expect temperatures of 5-10 K. The SWAS observed antenna temperatures in regions of detected H$_2$O are of order $T_A < 1$ K.  This result alone suggests the H$_2$O is in the ``effectively optically thin'' regime.  Our model does include an escape probability formalism for the H$_2$O lines.  We find, in agreement with Poelman et al, that significant collisional de-excitation and saturation of the line occurs when $\\tau n/n_{cr} > 1$, where $\\tau$ is the line center optical depth and $n_{cr}$ is the critical density ($\\sim 10^8$ cm$^{-3}$ for the 557 GHz line).  Thus, the requirement for 557 GHz line saturation is roughly  \\begin{equation} \\left[{N(H_2O) \\over {6\\times 10^{12}\\ {\\rm cm^{-2}}}}\\right]{n\\over n_{cr}} \\gta 1. \\end{equation} Since our models have $N(H_2O) \\sim 10^{15}$ cm$^{-2}$ and $n \\lta 10^6$ cm$^{-3}$, we are almost always within the effectively optically thin regime, except perhaps for the highest densities $n\\sim 10^6$ cm$^{-3}$.  Furthermore, if the H$_2$O columns and densities were high enough that line saturation occurred, the line fluxes would exceed those observed by SWAS.  We argue that SWAS observations indicate effectively thin 557 GHz emission, and that there is no evidence for hidden H$_2$O in the centers of clouds.  Bergin et al (2000) listed ingredients for a successful astrochemical model of a molecular cloud.  Slightly updated, they included: (1) low column-averaged abundances ($\\lta 10^{-7}$) of gas phase O$_2$, (2) low column-averaged abundances ($\\lta 3 \\times 10^{-8}$) of gas phase H$_2$O, (3) an explanation of the apparently higher abundances of H$_2$O when it is observed in absorption in translucent clouds, (4) column-averaged water ice abundances of $\\sim 10^{-4}$, and (5) existence of complex carbon-bearing species in the gas phase in some regions.  As noted above, the model presented in this paper fulfills all of these criteria.  In addition, we have successfully applied, in a simple way, these models to observations of the $A_V$ threshold for the onset of water ice, the strict upper limits for the gas phase H$_2$O abundance in B68 made by SWAS, the possible detection of O$_2$ in $\\rho$ Oph made by {\\it Odin}, the detection of H$_2$O and upper limits for O$_2$ made by SWAS for NGC 2024, and upper limits of the O$_2$ abundance set by {\\it Odin}.  These successes imply that the proper modeling of the interstellar chemistry of molecular clouds must include as ingredients: photodissociation, freezeout on grains, grain surface chemistry, desorption processes, and, in the very opaque central regions, time-dependent chemistry.  At the surface photodissociation and photodesorption dominate.  Deep in the cloud, dust vacuums the condensibles and slow cosmic ray processes are important, requiring time-dependent chemistry.  At intermediate depths, gas phase molecular species such as O$_2$ and H$_2$O peak in abundance, and provide the source of their submillimeter emission."
    },
    "0809/0809.0640_arXiv.txt": {
        "abstract": "{We model the cosmological co-evolution of galaxies and their central supermassive black holes (BHs) within a semi-analytical framework developed on the outputs of the Millennium Simulation \\citep{croton2006,delucia2007}. In this work, we analyze the model BH scaling relations, fundamental plane and mass function, and compare them with the most recent observational data. Furthermore, we extend the original code developed by \\citet{croton2006} to follow the evolution of the BH mass accretion and its conversion into radiation, and compare the derived AGN bolometric luminosity function with the observed one. We find, for the most part, a very good agreement between predicted and observed BH properties. Moreover, the model is in good agreement with the observed AGN number density in $0\\leq z\\leq5$, provided it is assumed that the cold gas fraction accreted by BHs at high redshifts is larger than at low redshifts \\citep{marulli2008}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.3196_arXiv.txt": {
        "abstract": "We study supersymmetric scenarios where the dark matter is the gaugino of an unbroken hidden $U(1)$ which interacts with the visible world only via a small kinetic mixing with the hypercharge.  Strong constraints on the parameter space can be derived from avoiding overclosure of the Universe and from requiring successful Big Bang Nucleosynthesis and structure formation. We find that for typical values of the mixing parameter, scenarios with neutralino NLSP are excluded, while scenarios with slepton NLSP are allowed when the mixing parameter lies in the range $\\chi\\sim\\mathcal{O} (10^{-13}-10^{-10})$. We also show that if the gravitino is the LSP and the hidden $U(1)$ gaugino the NLSP, the bounds on the reheating temperature from long lived charged MSSM relics can be considerably relaxed and we comment on the signatures of these scenarios at future colliders.  Finally, we discuss the case of an anomalously small mixing, $\\chi\\ll 10^{-16}$, where the neutralino becomes a decaying dark matter candidate, and derive constraints from gamma ray experiments. ",
        "introduction": "The existence of dark matter in the Universe is perhaps the most solid indication for physics beyond the Standard Model of particle physics~\\cite{BHS05}. Among the many extensions of the Standard Model that have been proposed in recent years, supersymmetry (SUSY) arguably remains as the most popular.  Among other merits, supersymmetric scenarios provide a very promising candidate for the dark matter: the lightest supersymmetric particle (LSP)~\\cite{EHN+84,JKG96}. With the particle content of the Minimal Supersymmetric Standard Model (MSSM), the LSP can be, over a large range of parameters, either the lightest neutralino, the lightest sneutrino or the lightest stau. Among these, only the lightest neutralino is still allowed by present experiments as a viable dark matter candidate, provided $R$-parity is almost exactly conserved. Furthermore, if supersymmetry is promoted to a local symmetry, the particle content of the MSSM also includes the gravitational supermultiplet, of which the spin $3/2$ component, the gravitino, is also a viable dark matter candidate~\\cite{PP82}, even if $R$-parity is slightly violated~\\cite{TY00,BCH+07}.  On the other hand, many extensions of the MSSM contemplate the possibility of a hidden sector, consisting of superfields which are singlets under the Standard Model gauge group.  Hidden sector superfields usually couple very weakly to our observable sector, thus constituting a very natural arena for finding dark matter candidates.  Generically, hidden sector particles couple to our observable sector only via non-renormalizable operators, presumably suppressed by a large mass scale, with a structure that is strongly model dependent.  In consequence, deriving implications of the hidden sector dark matter on the thermal history of the observable Universe and for future collider experiments is hindered by our complete ignorance of the strength and the structure of the hidden sector interactions with our observable sector.  There are however three instances where the hidden sector particles can couple to the MSSM particles via renormalizable operators, with a structure which is well defined by the Lorentz and gauge symmetries. Firstly, a hidden sector chiral superfield, $S$, could couple to the lepton and up-type Higgs doublet superfields via the Yukawa coupling $S H_u L$ in the superpotential, or to the two MSSM Higgs doublets, via $S H_u H_d$, provided these terms are also invariant under the hidden sector gauge group (as well as possible discrete and global symmetries of the theory).  Secondly, if one of the MSSM chiral superfields is charged under a hidden sector gauge group, it will interact with the corresponding hidden sector gauge superfield, and in turn with other hidden sector chiral superfields via the $D$-term.  Finally, a hidden sector abelian vector superfield, $X$, may couple to the hypercharge vector superfield through a kinetic mixing term, which is always allowed by the gauge symmetries \\cite{Holdom86, FH91}.  In this paper we will concentrate on the last situation and we will study scenarios where the hidden $U(1)$ gauge group is unbroken at low energies. If this is the case, the corresponding hidden $U(1)$ gaugino will have a mass comparable to the typical soft SUSY breaking masses of the observable sector particles and, in some instances, smaller. Therefore, the hidden $U(1)$ gaugino will not only be the lightest supersymmetric particle of the hidden sector, but also the lightest among all the supersymmetric particles.  Our motivation to consider an unbroken hidden sector $U(1)$ group is twofold. First, in string theory compactifications hidden sector $U(1)$ groups are ubiquitous and some of them could remain unbroken at low energies, in complete analogy to the familiar electromagnetic $U(1)$ of our observable sector. Secondly, the case of the hidden sector unbroken $U(1)$ is particularly intriguing, since this situation is practically unconstrained by present experiments. Indeed, it was shown long ago by Holdom that, in a non-supersymmetric world, the hidden $U(1)$ gauge boson (the ``paraphoton'' \\cite{Okun82}) completely decouples from the observable sector \\cite{Holdom91}. This result can be also generalized to a supersymmetric theory. Let us consider the SUSY invariant part of the Lagrangian, \\begin{eqnarray} \\nonumber {\\cal L}&=&\\int d^2\\theta \\, \\left(\\hat{W}_B^\\alpha \\hat{W}_{B\\,\\alpha}+ \\hat{W}_X^\\alpha \\hat{W}_{X\\,\\alpha}+2\\chi \\hat{W}_B^\\alpha \\hat{W}_{X\\,\\alpha}\\right)+ \\text{h.c.}+\\\\ &&+\\int d^2\\theta d^2\\bar \\theta \\, \\left(\\Phi^\\dagger e^{Y g_Y \\hat{B}} \\Phi+ h^\\dagger e^{q g_X \\hat{X}} h\\right)\\;, \\label{eqn:Lagrangian1} \\end{eqnarray} where the field strength superfield is defined as $\\hat{W}_V^\\alpha=-\\frac{1}{4} {\\bar D}{\\bar D} D^\\alpha \\hat{V}$, $\\hat{V}=\\hat{B}, \\hat{X}$ being the hypercharge or the hidden $U(1)$ vector superfield, while $\\Phi$ and $h$ denote, respectively, any Standard Model or hidden sector chiral superfield. Finally, $\\chi$ is the kinetic mixing parameter, which is induced through quantum effects by chiral superfields charged under both gauge groups.  Without additional symmetries, values around $\\chi\\sim 10^{-3} - 10^{-4}$ are naturally obtained.  However, \\fex in compactifications of heterotic \\cite{DKM97, BMW06} and type II \\cite{LS07, AS04, BHK05, AJK+08, AGJ+08} strings, much smaller mixings are possible.  A lower bound around $\\chi \\gtrsim 10^{-16}$ was argued to hold in cases of gauge mediated supersymmetry breaking in heterotic string models \\cite{DKM97}, whereas in type II models with warped extra dimensions the kinetic mixing parameter can be exponentially small \\cite{AGJ+08}.  The gauge kinetic terms in Eq.~\\eqref{eqn:Lagrangian1} can be made canonical by introducing shifted vector superfields, \\begin{eqnarray} X&=& \\hat{X}+ \\chi \\hat{B}\\;, \\\\ B&=& \\sqrt{1- \\chi^2} \\hat{B}\\;, \\end{eqnarray} leading to \\begin{eqnarray} {\\cal L}&=&\\int d^2\\theta \\, \\left(W_B^\\alpha W_{B\\,\\alpha}+ W_X^\\alpha W_{X\\,\\alpha}\\right) + \\text{h.c.}+ \\\\\\nonumber &&+\\int d^2\\theta d^2\\bar \\theta \\, \\left(\\Phi^\\dagger e^{Y g^\\prime_Y B} \\Phi+ h^\\dagger e^{q g_X X- q g^\\prime_X B} h\\right)\\;, \\end{eqnarray} where $g^\\prime_Y= g_Y/\\sqrt{1-\\chi^2}$ and $g^\\prime_X= \\chi g_X/\\sqrt{1-\\chi^2}$.  Therefore, the canonical normalization of the kinetic terms produces an unobservable shift of the hypercharge gauge coupling and the generation of a ``minihypercharge'' for the hidden sector chiral superfields \\cite{Holdom86}.  Different astrophysical observations and laboratory experiments constrain the possible values of the minihypercharge and the masses of the hidden sector particles. For instance, one obtains $\\chi \\lesssim 10^{-13}$ for masses below $10\\keV$ from plasmon decay in red giants \\cite{DP94, DHR00} (see also Ref.~\\cite{MPS07}). Nevertheless these bounds can be automatically avoided if the masses are large. In this case, as long as supersymmetry remains unbroken, the hypercharge vector superfield completely decouples from the observable sector and is not subject to any experimental constraint.  The breaking of supersymmetry changes dramatically the previous picture. Although the hidden $U(1)$ gauge boson remains decoupled from the observable sector, we will show that  in the presence of SUSY breaking effects a mixing between the hidden $U(1)$ gaugino and the MSSM neutralinos is induced.  Then, the unbroken hidden $U(1)$ might produce observable effects in the cosmological evolution of the Universe, at collider experiments or in cosmic ray fluxes.  Models with a hidden $U(1)$ extension of the SM or MSSM have been extensively studied in the literature (see \\fex \\cite{CNW06}, for a recent review see Ref.~\\cite{Langacker08}). Some of these works take into account kinetic- and mass-mixing \\cite{FKN07, FLN07, KN05, KW06, KN04, CY07, PRV08a}, but it is typically assumed that the gauge symmetry of the additional $U(1)$ is broken by a Higgs or Stueckelberg mechanism.\\footnote{Exceptions are \\fex Ref.~\\cite{Dobrescu05, BBS08}, where constraints on, and consequences of, higher dimensional operators that couple hidden and observable sector are studied. In Ref.~\\cite{FTY08} the authors study BBN and CMB constraints on the particle content of a completely decoupled hidden sector which may contain unbroken $U(1)$s.  Furthermore, see Ref.~\\cite{Redondo08a} for a short discussion about gauge coupling unification in the presence of kinetic mixing.} In this case, it is possible to derive bounds from high precision LEPI data if the hidden $U(1)$ gauge boson mass is large (namely, $\\chi\\lesssim 0.05$ for masses around $200\\GeV$ \\cite{FLN07}), and from different astronomical observations and laboratory experiments if the masses are small (namely, for masses around $100\\eV$, the observed lifetime of the sun translates into a bound $\\chi\\lesssim 10^{-13}$ \\cite{Redondo08}, see Ref.~\\cite{JRR08} for bounds from considerations of the hidden CMB).  The purpose of this paper is to consider in some detail cosmological constraints and phenomenological properties of a kinetically mixed hidden $U(1)$ extension of the MSSM with unbroken gauge symmetry. We will assume that the hidden $U(1)$ gaugino is the LSP in most of the paper, and concentrate on its prospect of being dark matter.  In section 2, we will describe the model in the component formalism and discuss typical values for mass- and mixing-parameters in scenarios with gauge and gravity mediation of supersymmetry breaking.  In section 3, we derive a bound on the mixing parameter from thermal overproduction of the hidden $U(1)$ gaugino. Section 4 considers bounds from primordial nucleosynthesis for the cases of stau and neutralino NLSPs. There, we also analyze the effects of a possibly light gravitino and show how bounds on the reheating temperature can be relaxed.  In section 5 we briefly discuss the collider phenomenology of this scenario.  In section 6 we study the case with an anomalously small kinetic mixing, where the neutralino NLSP could become a viable, though unstable, dark matter candidate. Finally, we present our conclusions and an outlook in section 7. ",
        "conclusions": "An unbroken hidden $U(1)$ that interacts with the Standard Model only via kinetic mixing with hypercharge decouples completely from the visible world. However, in the supersymmetric version of this scenario this is no longer the case. We have shown that a mass mixing between the bino and the hidden $U(1)$ gaugino is always generated via radiative effects, although this mixing can be generated  already at tree level in some well motivated scenarios.  Moreover, we have discussed different scenarios of supersymmetry breaking in which the hidden $U(1)$ gaugino is the lightest supersymmetric particle.  We have mostly concentrated on this possibility and we have derived cosmological bounds on this scenario from precluding overproduction of hidden $U(1)$ gauginos and from the requirements of successful Big Bang nucleosynthesis and structure formation.  The combination of these constraints excludes a neutralino NLSP except for extremely small mixings (see Fig.~\\ref{fig:boundsBino}). On the other hand, when the NLSP is a stau, an allowed window for mixings around $\\chi\\sim 10^{-10} - 10^{-13}$ remains (see Fig.~\\ref{fig:boundsSlepton}). In this window, the stau has a lifetime larger than $10^{-1}\\s$ and thus might be detected at future colliders as a heavily ionizing charged track.  The reheating temperature in scenarios with gravitino dark matter and stau NLSPs is known to be strongly constrained, and we have shown that these constraints relax considerably in the presence of a hidden $U(1)$ gaugino (see Fig.~\\ref{fig:FSGravitinoNLSP}).  This might be a rather general effect of very weakly interacting hidden sectors and deserves further attention.  Finally, we have discussed the case of an anomalously small mixing parameter, $\\chi\\ll 10^{-16}$. For these small mixings, the neutralino NLSP can become long lived enough to constitute the dark matter of the Universe.  We have constructed a simple model with two hidden $U(1)$s where a tiny mixing can be naturally obtained. Even though the neutralino is very long lived, it eventually decays into the hidden gaugino and standard model particles, which might be detected as an anomalous contribution to the cosmic ray fluxes.  Using the EGRET measurement of the extragalactic gamma ray flux, we have derived a conservative bound on the mixing parameter $\\chi\\lesssim 10^{-20-21}$ (see Fig.~\\ref{fig:EGRETbounds})."
    },
    "0809/0809.4255_arXiv.txt": {
        "abstract": "{ Exoplanetary transit and stellar oscillation surveys require a very high precision photometry. The instrumental noise has therefore to be minimized. First, we perform a semi-analytical model of different noise sources. We show that the noise due the CCD electrodes can be overcome using a gaussian PSF (Point Spread Function) of full width half maximum larger than 1.6 pixels. We also find that for a PSF size of few pixels, the photometric aperture has to be at least 2.5 times larger than the PSF full width half maximum. Then,  we compare a front- with a back-illuminated CCD through a Monte-Carlo simulation. Both cameras give the same results for a PSF full width half maximum larger than 1.5 pixels. All these simulations are applied to the A STEP (Antarctica Search for Transiting Extrasolar Planets) project. As a result, we choose a front-illuminated camera for A STEP because of its better resolution and lower price, and we will use a PSF larger than 1.6 pixels. } ",
        "introduction": "The photometric technique allows a direct detection of luminosity variations. Several disciplines are therefore con\\-cer\\-ned. In asteroseismology, these variations are used to iden\\-ti\\-fy stellar oscillations, giving an access to  the interior of stars. In planetary sciences, a decrease of luminosity cau\\-sed by an extrasolar planet occulting its parent star during a transit is used to characterize the planet (Charbonneau et al. 2000; Moutou et al. 2004). In both cases, a very high precision photometry is required, typically to a milli\\-mag\\-ni\\-tude level. Challenging technical issues have therefore to be solved (Rauer et al. 2004), and high accuracy algorithms are needed (Irwin at al. 2007; Gillon et al. 2006; Magain et al. 2007). The Antarctica Search for Transiting Extrasolar Planets (A STEP) aims to detect planetary transits and stellar oscillations from Dome C, Antarctica (Fressin et al. 2005). The three months continuous night as well as a very dry weather are extremely favorable for photometric surveys. A fully automatized telescope is under development. We present here a first noise analysis of this telescope that leads to the choice of the camera, but that applies to other photometric surveys. The first part shows a noise budget obtained with a semi-analytical model. A second part des\\-cribes a Monte Carlo simulation of a front- and a back-illuminated CCD camera. ",
        "conclusions": "We performed a semi-analytical noise analysis to identify the limiting noise sources in precision photometry, using a gaussian PSF. The electrode, overflow, and interpixel noises are dominant, as well as the photon noise for faint stars. We showed that the electrode noise becomes negligible for gaussian PSF full width half maxima larger than 1.6 pixels. We also found that for PSFs of few pixels, the photometric aperture must be at least 2.5 times larger than the PSF full width half maximum.  We then compared a front- and a back-illuminated ca\\-me\\-ra in a Monte Carlo simulation. For photometric surveys for which the PSF is well-sampled (at least 1.5 pixels full width half maximum), and limited in terms of budget to exis\\-ting commercial cameras, we found that a front-illu\\-mi\\-na\\-ted camera is a better alternative."
    },
    "0809/0809.1646_arXiv.txt": {
        "abstract": "We present a search for point sources of high energy neutrinos using 3.8 years of data recorded by AMANDA-II during 2000-2006. After reconstructing muon tracks and applying selection criteria designed to optimally retain neutrino-induced events originating in the Northern Sky, we arrive at a sample of 6595 candidate events, predominantly from atmospheric neutrinos with primary energy 100 GeV to 8 TeV. Our search of this sample reveals no indications of a neutrino point source. We place the most stringent limits to date on $E^{-2}$ neutrino fluxes from points in the Northern Sky, with an average upper limit of $E^{2}\\Phi_{\\nu_{\\mu} + \\nu_{\\tau}} \\le$ 5.2~$\\times$~10$^{-11}$ TeV cm$^{-2}$ s$^{-1}$ on the sum of $\\nu_{\\mu}$ and $\\nu_{\\tau}$ fluxes, assumed equal, over the energy range from 1.9 TeV to 2.5 PeV. ",
        "introduction": "Detecting extraterrestrial sources of high energy ($>$TeV) neutrinos is a longstanding goal of astrophysics.  Neutrinos are neither deflected by magnetic fields nor significantly attenuated by matter and radiation en route to Earth, thus neutrino astronomy offers an undistorted view deep into the high energy universe.  Particularly, neutrinos offer an opportunity to probe the sources of high energy cosmic rays, which remain unknown.  Potential cosmic ray sources include galactic microquasars and supernova remnants as well as extragalactic sources such as active galactic nuclei and gamma ray bursts.  These objects are thought to accelerate protons and nuclei in shock fronts via the Fermi mechanism \\cite{fermi}, resulting in power law energy spectra $E^{\\alpha}$, with $\\alpha$ $\\sim$ $-$2. A fraction of the energized particles interact with local matter and radiation, producing pions.  The neutral pions decay into high energy photons, and the charged pions ultimately produce neutrinos with a flavor ratio $\\nu_e$:$\\nu_{\\mu}$:$\\nu_{\\tau}$ $\\sim$1:2:0, mixing to approximately 1:1:1 at Earth because of vacuum flavor oscillations. Observations of TeV gamma rays \\cite{milagrogal,hessgal,magicgal} hint at possible cosmic ray source locations but currently cannot separate neutral pion decay spectra from inverse Compton emission.  The Auger collaboration has reported a correlation of arrival directions of the highest energy cosmic rays with active galactic nuclei \\cite{augercorrsci}; however, a similar correlation has not been observed by HiRes \\cite{hiresehe}. Identification of a high energy neutrino point source would provide an unambiguous signature of energetic hadrons and cosmic ray acceleration.  Neutrino flux predictions exist for many potential sources \\cite{microq5039,microqlsi,snr,stecker,kappes,beacom}, but no high energy neutrino point source has yet been identified \\cite{amandalim, macrolim, superklim}.  The search for high energy neutrino point sources is a major objective of the Antarctic Muon And Neutrino Detector Array (AMANDA).  High energy leptons are produced in the Earth by charged-current neutrino interactions.  In transparent matter, a cone of Cherenkov photons propagates from the lepton track according to the optical properties of the medium. AMANDA-II is an optical Cherenkov detector consisting of 677 optical modules arranged in 19 strings frozen $\\sim$1500 m to $\\sim$2000 m deep in the ice sheet at the geographic South Pole. Approximately 540 modules in the core of the array showing stable performance are used in this search. Each module contains a 20 cm diameter photomultiplier tube (PMT) optically coupled to an outer glass high-pressure sphere.  PMT pulses are propagated to surface electronics, and, when the trigger threshold of 24 discriminator crossings (``hits\") within 2.5 $\\mu$s is satisfied, the pulse leading edge times are recorded.  The leading edge times along with known detector geometry and optical properties of South Pole ice \\cite{icepaper} allow reconstruction of tracks passing through the detector \\cite{amandareco}.  High energy electrons produce short electromagnetic cascades with little directional information of the primary neutrino.  Muons produced in the ice and bedrock, on the other hand, propagate up to several kilometers to the detector and their tracks are reconstructed with ~1.5$^{\\circ}$--2.5$^{\\circ}$ median accuracy depending on energy and zenith angle. Tau leptons decay rapidly and produce tracks too short for reconstruction below $\\sim$PeV energies.  Tau decay, however, contributes high energy muons with a branching ratio of 17.7\\% \\cite{amandalim, pdg}, and these muon tracks can be reconstructed.  We thus search for upward propagating muons produced in the Earth by $\\nu_{\\mu}$ ($\\bar{\\nu}_{\\mu}$) and $\\nu_{\\tau}$ ($\\bar{\\nu}_{\\tau}$) fluxes following roughly an $E^{-2}$ energy spectrum.  While downward neutrino induced muons also trigger the detector, such events are difficult to distinguish from downward muons produced by cosmic ray air showers.  Located at the South Pole, AMANDA-II is thus most sensitive to neutrino fluxes from the Northern Sky.  Air showers also produce neutrinos, and this atmospheric neutrino flux \\cite{bartol,honda} is the main background for our search.  Here we present the results of a search for astrophysical point sources of high energy neutrinos using 3.8 years of data recorded by AMANDA-II during 2000-2006, extending the previous five-year analysis \\cite{amandalim} with data from the final two years of standalone operation and improving our sensitivity by a factor of $\\sim$2. We report flux limits for a catalog of 26 selected source candidates along with results of a search for neutrino sources over the entire Northern Sky.  Additionally, we report results from a search for neutrino emission from gamma ray sources identified by Milagro \\cite{milagrogal} and a search for event angular correlations. In all cases, we observe no indications of an astrophysical neutrino point source. ",
        "conclusions": "We have analyzed 3.8 years of AMANDA-II data and found no evidence of high energy neutrino point sources. We place the most stringent limits to date on astrophysical point source fluxes.  IceCube \\cite{icecube} is a next-generation neutrino telescope at the South Pole scheduled for completion in 2011 with eighty 60-module strings instrumenting $\\sim$1 km$^3$ of ice.  Analysis of data recorded during 2007-2008 with the first 22 strings has improved the AMANDA-II sensitivity by a factor of 2.  Currently 59 strings are operating, and with continued construction IceCube will achieve an angular resolution of better than one degree and an order of magnitude improvement over the AMANDA-II sensitivity within a few years."
    },
    "0809/0809.4125_arXiv.txt": {
        "abstract": "We analyze a new gravitational lens \\gl, serendipitously found in a deep I$_{814}-$band image of the Hubble Space Telescope (HST) Advanced Camera for Surveys (ACS). The lens is a $L^*$, edge-on S0 galaxy at $z_{\\rm l}=0.4656$. The gravitational arc has a radius of $0''.42 \\simeq 1.74 \\, h^{-1} \\, {\\rm kpc}$. We have determined the total mass and the dark matter (DM) fraction within the Einstein radius \\RE\\ as a function of the lensed source redshift, which is presently unknown. For $z_{\\rm s}\\sim 1.3$, which is in the middle of the redshift range plausible for the source according to some external constraints, we find the central velocity dispersion to be $\\sim 180 \\, {\\rm km \\, s}^{-1}$. With this value, close to that obtained by means of the Faber-Jackson relation at the lens redshift, we compute a 30\\% DM fraction within $\\RE\\ $ (given the uncertainty in the source redshift, the allowed range for the DM fraction is 25-35 \\% in our lensing model). When compared with the galaxies in the local Universe, the lensing galaxy \\gl\\ seems to fall in the transition regime between massive DM dominated galaxies and lower-mass, DM deficient systems. ",
        "introduction": "Strong gravitational lensing (GL) is a valuable astrophysical tool to investigate the structure and evolution of early-type galaxies up to a redshift of $\\sim 1$ (e.g., Keeton et al.~1998; Rusin et al.~2003, Treu et al.~2005) and to gauge the DM content at various galaxian scales (e.g., Treu et al.~2006, Jiang \\& Kochanek 2007). This tool offers the advantage of constraining the total mass within the Einstein radius independently of the dynamical status of the lensing galaxy. The technique is challenging to apply to low-luminosity lenses (LLLs), however, as the GL cross section is proportional to the fourth power of the central velocity dispersion (e.g., Covone et al.~2005). This is why LLLs appear far more rare than more massive systems. \u00a0Their systematically smaller Einstein radii make strong lensing arcs around them hard to find with optical surveys. For instance, a lensing galaxy with velocity dispersion $\\sL\\sim 200 \\, {\\rm km \\, s^{-1}}$ at $\\zl \\sim 0.5$, coupled to a source at $\\zs \\sim 1.0$, produces an Einstein ring with radius $\\theta_{\\rm E} \\sim 0''.5$. As a comparison, in a visual search for gravitational lenses in the COSMOS survey (Faure et al.~2008)\\footnote{{\\tt http://cosmosstronglensing.uni-hd.de/}}, no lensing galaxy was found with central velocity dispersion smaller than $200~$km s$^{-1}$, out of a sample of 20 secure systems.  Although difficult to discover, LLLs are of great interest regarding recent claims of dark-matter deficient $L_*$ galaxies in the local Universe (see, e.g., Capaccioli et al.~2003, Romanowsky et al.~2003, Napolitano et al.~2005, N+05 hereafter). Furthermore, these systems occupy as region of the luminosity/mass distribution in between boxy/slow-rotator systems and disky/fast-rotator systems (Nieto \\& Bender 1989, Capaccioli et al. 1992). The dichotomy of these systems appears to involve also their DM properties, either at effective radii scales (Cappellari et al.~2006) or beyond (Capaccioli et al.~2003, N+05).  This paper analyses a $L_*$ edge-on S0 lensing galaxy at $z_{\\rm l}=0.4656\\pm 0.0004$, hereafter named \\gl, serendipitously discovered in the HST/ACS follow-up imaging of a former candidate lensing cosmic string (Sazhin et al.~2007). The mean radius of the gravitational arc is $ \\theta_{\\rm E} = 0''.42$, corresponding to a linear radius of $1.74 \\, h^{-1} \\, {\\rm kpc}$, slightly larger than the estimated value of the effective (half-light) radius $\\reff$ of the lens bulge (i.e. $\\theta_{\\rm E} \\sim 1.15 \\, \\reff $). \u00a0This compact arc offers the rare opportunity to gauge the mass distribution within the effective radius of a $L_*$ galaxy at $z \\sim 0.5$. Note that, among the 20 strong GLs found in the 1.64 square degrees of the COSMOS survey by Faure et al.~(2008),  selected in the range $ 0.2 < z < 1.0$ and median redshift 0.71, only one lens exhibits an arc with an angular radius smaller than $0''.40$.  As the redshift of the lensed source is yet unknown, despite two spectroscopic runs (see Sect.~2), the DM content of the lens cannot be unambiguously determined. However, by requiring that the lensing galaxy follows the Faber-Jackson relation (Faber \\& Jackson~1976), and exploiting the decoupled geometry of the luminous and the total mass, we can then constrain the DM fraction within the Einstein radius. Toward this end, in Sect.~3 we present a lensing model, and in Sect.~4 we infer the luminous mass from the analysis of the spectro-photometric data, in order to disentangle the lensing contribution of the DM from that of the total mass. In Sect.~\\ref{solution}, we discuss the best mix of dark and luminous mass which produces the measured total mass as a function of the (unknown) source redshift, and draw conclusions in Sect.~\\ref{concl}.  Throughout the paper we will assume a cosmological model with $\\Omega_{\\rm m}=0.27$, $\\Omega_{\\Lambda}=0.73$ and $h \\equiv H_{0} / (100 \\, {\\rm km} \\, {\\rm s}^{-1}$\\,Mpc$^{-1}) = 0.7$. At the lens redshift, $1''$ corresponds to $4.15 \\, h^{-1}$ kpc (Komatsu et al.~2008). Magnitudes are in the AB system. ",
        "conclusions": "\\begin{figure} \\center \\includegraphics[width=0.5\\textwidth, angle=0]{f6.eps} \\caption{DM fraction at $R_{\\rm E}$ versus the (unknown) redshift of the lensed source. On the right the corresponding (mean) $f_{\\rm vir}$ values; see text for details. Empty boxes are the DM fractions for the single bulge+disk population, full boxes for the two population models, assuming Salpeter IMF for the stellar $M/L$. Diamonds represent the same models for the Chabrier IMF. Grey denotes the regions excluded by the constraints on the $\\fracDM$ from the lensing model geometry, cyan are the regions excluded by the Faber-Jackson relation. In green, the range of the parameter space allowed by all the contraints, accounting also for the uncertainty associated to the population models $M/L$.} \\label{fig:tre} \\end{figure}   \\gl, a $\\sim L_*$ S0 galaxy at redshift $z=0.466$, is among the least massive lenses at high redshift known to date. Although the source redshift is still not known, the observed phenomenology unambiguously supports the strong-lensing interpretation. The lensed source produces a bright arc (see Fig.~1) with radius \u00a0$0''.42$ ($\\sim 1.15 \\, \\reff$ of the lens bulge). A candidate counter-image is found at $0''.65$ from the lens center.  The total mass distribution, modeled by a singular isothermal ellipse, is found to be rounder than the light at the Einstein radius, with a ratio $ q_{\\rm SIE} / q_{*} = 1.84$, similar to the value found for galaxies with velocity dispersions lower than $\\sim 200 \\, {\\rm km \\, s}^{-1}$.  We constrain the redshift of the source by coupling the value of the ratio $(\\massDM /M_*) (\\RE)$, obtained via the lens model, with the lens total mass within $R_{\\rm E}$, % once the \u00a0corresponding stellar mass \\SM\\ is evaluated by an assumption on the baryonic component and corresponding stellar $M/L$. We obtain $z_{\\rm s} \\sim 1.3 \\pm 0.2$ with a Salpeter IMF, the Chabrier IMF being ruled out. In this case the velocity dispersion from the lens model (a singular isothermal ellipse with an external shear) is $\\sL\\ = 177^{+10}_{-6}$ km \\, s$^{-1}$ (with the given uncertainty % derived from the allowed $z_{\\rm s}$ range). This value is consistent with the limits imposed by the FJ relation.     \\gl\\ offers the possibility to investigate the mass distribution of the very inner region in an $L_*$ galaxy at $z \\sim 0.5$. Indeed, while there is growing evidence for a picture in which massive galaxies (i.e., above $\\sim 10^{11.5} M_{\\odot}$) are well described by an isothermal mass profile and show no structural evolution since $z\\sim1$ (see, for instance, Treu~2007), we still lack a large sample of $L_*$ galaxies at $z\\gtrsim0.5$ with a robust determination of the inner DM fraction in order to probe mass assembly and galaxy evolution at lower mass scales.   We find that \\gl\\ (assuming $z_{\\rm s} = 1.3$) has a projected DM fraction $\\sim 30$\\% within the Einstein radius. % Lower values of the source redshift require a larger DM fraction in a non-spherical halo or a different IMF (e.g., Fig.~\\ref{fig:tre}). By matching the measured mass-to-light ratio at the \\RE\\ with the one expected from a double-component model formed by a NFW spherical halo and a Hernquist spheroidal light distribution, we derive the total dark-to-luminous mass ratio $f_{\\rm vir} \u00a0\\equiv \\massDM /M_{*} = \u00a06 \\pm 4$. This value is close to the one obtained for galaxies in the local Universe (N+05) located in the transition regime between massive DM dominated galaxies and lower-mass, DM deficient systems.  A firmer estimate of the inner DM and stellar content in this $L_*$ galaxy requires a further observational effort to determine the redshift of the lensed source. A possible strategy would involve, for instance, near-infrared spectroscopy aimed \u00a0at detecting the H$_{\\alpha}$ emission line."
    },
    "0809/0809.4313_arXiv.txt": {
        "abstract": "In our previous paper \\citep{JRM_dbip} we presented the design and initial calibrations of the Dual-Beam Imaging Polarimeter (DBIP), a new optical instrument for the University of Hawaii's $2.2~$m telescope on the summit of Mauna Kea, Hawaii.  In this followup work we discuss our full-Stokes mode commissioning including crosstalk determination and our typical observing methodology. ",
        "introduction": "To study the linear polarization of asteroids and other point source objects, the Dual-Beam Imaging Polarimeter (DBIP) was commissioned in March of 2007 \\citep{JRM_dbip}.  In August of 2007 we expanded DBIP's capabilities to include analysis of circular polarization with the addition of a quarterwave plate.  Typically, the most important quantities for analysis are the fractional polarizations $q=Q/I$, $u=U/I$, and $v=V/I$, expressed as percentages, and in the following text we will deal with these quantities when we refer to polarization measurements.  Here we present our subsequent calibration and determination of systematic errors which were found to be comparable to statistical errors for typical observing situations: $\\sim0.1\\%$ polarization. ",
        "conclusions": "We have presented the results of our extended commissioning of DBIP into full-Stokes mode.  Crosstalk values are shown to be small and can be ignored for objects with ``typical'' (i.e. $<5\\%$) polarizations.  With errors well constrained DBIP is now available for scientific studies of linear and circular polarization of point sources in the optical.  DBIP has already been successfully used to characterize the polarization-phase curves of asteroids with some of the most pristine compositions in the inner solar system \\citep{JRM_aquitania} and a number of other programs to study small solar system bodies are currently under way."
    },
    "0809/0809.2428_arXiv.txt": {
        "abstract": "Detecting redshifted 21cm emission from neutral hydrogen in the early Universe promises to give direct constraints on the epoch of reionization (EoR). It will, though, be very challenging to extract the cosmological signal (CS) from foregrounds and noise which are orders of magnitude larger. Fortunately, the signal has some characteristics which differentiate it from the foregrounds and noise, and we suggest that using the correct statistics may tease out signatures of reionization. We generate mock datacubes simulating the output of the Low Frequency Array (LOFAR) EoR experiment. These cubes combine realistic models for Galactic and extragalactic foregrounds and the noise with three different simulations of the CS. We fit out the foregrounds, which are smooth in the frequency direction, to produce residual images in each frequency band. We denoise these images and study the skewness of the one-point distribution in the images as a function of frequency. We find that, under sufficiently optimistic assumptions, we can recover the main features of the redshift evolution of the skewness in the 21cm signal. We argue that some of these features -- such as a dip at the onset of reionization, followed by a rise towards its later stages -- may be generic, and give us a promising route to a statistical detection of reionization. ",
        "introduction": "\\label{sec:intro}  Between redshift $z\\approx 20$ and $z\\approx 6$ the Universe underwent a transition from being almost entirely neutral to almost entirely ionized (\\citealt{BEN06}; \\citealt*{FOB06}). This period, the epoch of reionization (EoR), saw the first collapsed objects emit radiation which heated and ionized the surrounding diffuse gas (additional sources of heating and ionization, such as dark matter decay, have also been considered; see, e.g., \\citealt{VAL07}). Studying the emission from this gas, in particular the redshifted 21cm line of neutral hydrogen \\citep*{FIE58,FIE59,HOG79,SCO90,KUM95,MAD97}, may therefore tell us about the physics of these objects, and about structure formation in a redshift range that has previously been explored rather less directly. For example, quasar absorption spectra can constrain the properties of the intergalactic medium towards the end of reoinization \\citep{FAN06}, while measurements of temperature and polarization anisotropies in the cosmic microwave background provide an integral constraint on the density of free electrons between the observer and the surface of last scattering \\citep[e.g.][]{DUN08}.  Several current and upcoming facilities (e.g.\\ GMRT,\\footnote{Giant Metrewave Radio Telescope, http://www.gmrt.ncra.tifr.res.in/} MWA,\\footnote{Murchison Widefield Array, http://www.haystack.mit.edu/ast/arrays/mwa/} LOFAR,\\footnote{Low Frequency Array, http://www.lofar.org/} 21CMA,\\footnote{21 Centimeter Arrat, http://web.phys.cmu.edu/\\~{}past/} PAPER,\\footnote{Precision Array to Probe EoR, http://astro.berkeley.edu/\\~{}dbacker/eor/} SKA\\footnote{Square Kilometre Array, http://www.skatelescope.org/}) will be sensitive to emission of the right wavelength to detect a signal from neutral hydrogen during the EoR.  Significant observational challenges must be overcome, however, before a convincing detection can be made. The Galactic and extragalactic foregrounds have a mean amplitude around \\hbox{4--5} orders of magnitude larger than the expected EoR signal (though their fluctuations, which are the relevant quantity for an interferometer, are only around three orders of magnitude larger; see, e.g., \\citealt{SHA99}). Even closer to home, the signal is corrupted by the ionosphere, radio frequency interference and instrumental effects. Assuming all these factors can be dealt with, for realistic integration times with the imminent generation of facilities the random noise on the measurement per resolution element will still be a few times larger than the signal.  The prospect of new observations of a poorly constrained period in the Universe's history also poses a challenge to theorists: to model and characterize the 21cm emission in such a way that it can be meaningfully compared to the data (e.g. \\citealt*{BAR01,LOE01,CFW03,CSW03,BRO04}; \\citealt{GLE06,ILI06,ILI08,ZAR07,THO08}). Given the observational limitations listed above, it is unlikely that there will soon be clean maps of the EoR signal with which to confront models. The first detection of reionization from redshifted 21cm data will therefore be of a statistical nature. This raises the question of precisely which statistics to use: a question discussed by, e.g., \\citet*{FUR04b,FUR04a,BHA05,GLE08}. In the first instance, they should provide a clear indication of the global transition from a Universe that is mostly neutral to one that is mostly ionized. Ideally, they should be able to discriminate between different models describing the more detailed progress of reionization. Most importantly, though, they should be robust to the contamination introduced by the observing process, and to the presence of high levels of noise.  We propose using the skewness of the one-point distribution of the brightness temperature to study reionization, though we also consider the prospects of other, similar statistics: the unnormalized third moment and the kurtosis of the distribution. As we shall see below, general arguments suggest that the skewness should be a strongly evolving function of redshift during the EoR, and these arguments are supported by simulations. Using these simulations, we generate datacubes which also incorporate realistic models for the foregrounds, instrumental response and noise levels expected for the LOFAR EoR experiment. We generate residual images at each observed frequency by attempting to remove the foregrounds using a fitting algorithm, then study the properties of these residual images as a function of redshift. If these residual images are denoised appropriately, we find that we can indeed track the progress of reionization using the skewness.  In Section~\\ref{sec:genapp} we introduce the skewness and explain how it may help. In Section~\\ref{sec:data} we give a brief description of our models for the cosmological signal (CS), instrumental response, foregrounds and noise. Then, in Section~\\ref{sec:res}, we describe our method for extracting the signal from datacubes which combine all these components, and present our results. We discuss some possible problems with and extensions to our methods in Section~\\ref{sec:disc}, and finally we offer some conclusions in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Many statistics have been put forward to characterize the 21cm emission from the EoR, the power spectrum probably being the most frequently studied \\citep[][to choose some recent examples]{BAR08,LID08,PRI08,SET08}. We have suggested that higher-order statistics may be useful not only to characterize a CS cube that has been cleaned of foregrounds, noise and instrumental effects, but also to extract the signature of reionization from these corrupting influences in the first place. The skewness of the one-point distribution of brightness temperature, measured as a function of observed frequency (or equivalently as a function of redshift), is one such promising statistic.  The three detailed simulations of reionization which we have studied show a strong evolution of the skewness with redshift. Some of the features of this evolution appear to be generic and can be readily understood: in the early stages of reionization the skewness drops below that of the underlying density field as the first ionized bubbles, from which the emission is negligible, are formed. As reionization progresses, the majority of the volume becomes ionized and the skewness increases again, becoming very large at low redshift when the distribution of brightness temperature is peaked at zero, with a tail extending to large values. In simulation f250C there is a well defined dip in the skewness with a width $\\Delta z\\approx 1$. Our other simulation in which the Universe is reionized entirely by stars (T-star) shows a more gradual change, with the epoch of reionization extending throughout the redshift range probed by LOFAR. A third simulation, T-QSO, in which QSOs reionize the Universe, shows an intermediate behaviour.  By combining these simulations with models of the foregrounds, noise and instrumental response, we have generated datacubes which are intended to simulate the output of the LOFAR EoR experiment. We have studied two cases: firstly, one in which we smooth the foregrounds and signal to the resolution of the telescope using a Gaussian kernel, then add uncorrelated Gaussian noise; secondly, one in which we degrade the foregrounds and noise to the resolution of the telescope using a realistic PSF, and add noise which is uncorrelated in the Fourier plane rather than the image plane, producing what we refer to as `dirty' images. In the former case, we can see the signature of reionization in the skewness by fitting out the foregrounds to obtain residual images, and then denoising these images with a simple smoothing operation. The skewness in these images as a function of redshift shows significant evidence of reionization. The result is quite robust to the details of the foreground fitting and the smoothing. Under- or over-fitting the foregrounds affects the recovered skewness less severely than the recovered variance. Extracting a signal from the dirty cubes requires a more sophisticated denoising scheme. In an optimistic scenario where the correlation matrices of the original signal and of the noise are known, we can again recover the evolution of the skewness quite cleanly using Wiener deconvolution.  We have touched upon some areas for improvement: simulations which remain realistic but extend to larger scales and exhibit an even greater range of reionization histories; taking into account the polarization of the foregrounds and the instrumental response, and incorporating new observational constraints as they arrive; testing the minimal assumptions we must make about the signal in order for our extraction scheme to work, for example whether a poor estimate of the correlation matrix of the CS seriously affects the extracted skewness; and studying a wider range of statistics beyond the variance and power spectrum. All of these will be areas for future work. Even at this stage, however, our results justify some optimism that the new generation of radio telescopes can detect the signature of reionization using higher-order statistics."
    },
    "0809/0809.2334_arXiv.txt": {
        "abstract": "We report on the first completely simultaneous observation of a gamma-ray burst (GRB) using an array of Imaging Atmospheric Cherenkov Telescopes which is sensitive to photons in the very-high-energy (VHE) $\\gamma$-ray range ($\\gtrsim100$~GeV). On 2006 June 2, the \\emph{Swift} Burst Alert Telescope (BAT) registered an unusually soft $\\gamma$-ray burst (GRB~060602B). The burst position was under observation using the High Energy Stereoscopic System (H.E.S.S.) at the time the burst occurred. Data were taken before, during, and after the burst. A total of 5 hours of observations were obtained during the night of 2006 June 2--3, and 5 additional hours were obtained over the next 3 nights. No VHE $\\gamma$-ray signal was found during the period covered by the H.E.S.S. observations. The 99\\% confidence level flux upper limit ($>$1~TeV) for the \\emph{prompt} phase (9~s) of GRB~060602B is $2.9\\times10^{-9}\\,\\mathrm{erg\\,cm^{-2}\\,s^{-1}}$. Due to the very soft BAT spectrum of the burst compared to other \\emph{Swift} GRBs and its proximity to the Galactic center, the burst is likely associated with a Galactic X-ray burster, although the possibility of it being a cosmological GRB cannot be ruled out. We discuss the implications of our flux limits in the context of these two bursting scenarios. ",
        "introduction": "Gamma-ray bursts (GRBs) are brief and intense flares of $\\gamma$-rays. Without precedent in astronomy, they arrive from random directions in the sky and last typically~$\\sim$0.1--100~s~\\citep[\\emph{prompt} emission, see][]{klebesadel73,fishman95}. The very nature of GRBs makes it operationally rather challenging to study their \\emph{prompt} phase simultaneously in any other wavelength.  The observed GRB properties are generally well explained by the \\emph{fireball} model, in which the emission is produced in relativistic shocks~\\citep{piran99,zhang04,meszaros06}. In this standard model, the highly-relativistic plasma, which emits the observed sub-MeV radiation, is expected to generate $\\gamma$-rays up to the very-high-energy (VHE; $\\gtrsim$100~GeV) regime, via inverse-Compton emission of electrons or proton-induced mechanisms~\\citep{zhang01,peer05,asano07,fan08}. Therefore, the detection of gamma-rays or sufficiently sensitive upper limits would shed light on our understanding of the current model. Some important yet largely unknown parameters in GRB models, such as the bulk Lorentz factor and the opacity of the outflow just after the acceleration phase, can be directly measured through high-energy (HE; $\\ga$100 MeV) and VHE $\\gamma$-ray observations during the \\emph{prompt} phase of GRBs~\\citep{razzaque04,baring06}.  There are two techniques used in VHE $\\gamma$-ray astronomy to observe the \\emph{prompt} phase: the first technique is to slew quickly to the GRB position provided by a burst alert from satellites. This technique is used for Imaging Atmospheric Cherenkov Telescopes (IACTs), such as the High Energy Stereoscopic System (H.E.S.S.), which have a field of view (FoV) of a few degrees. The MAGIC telescope, operating in this mode, was able to slew to the position of GRB~050713A, 40~s after the GRB onset, while the \\emph{prompt} keV emission was still active. A total of 37 minutes of observations were made and no evidence of emission above 175~GeV was obtained~\\citep{albert06a}. The rapid follow-up observations using this telescope of 8 other GRBs show no evidence of VHE $\\gamma$-ray emission from these GRBs during the \\emph{prompt} or the \\emph{early afterglow} phase~\\citep{albert07}. However, there is always a delay in time for IACTs operating in this GRB-follow-up mode, as long as the GRB position lies outside the camera FoV at the onset of the GRB. This results in an incomplete coverage of the GRB \\emph{prompt} phase.  The second technique is to observe a large part of the sky continuously, at the expense of much lower sensitivity than the IACT detectors. This technique is used, e.g. for the water Cherenkov detector Milagro, which works at higher energies than current IACTs. Since the effect of extra-galactic background light (EBL) absorption increases with the energy of a $\\gamma$-ray photon, the higher energy threshold of Milagro thus lowers its chance to detect VHE $\\gamma$-rays from distant GRBs, when compared to IACT detectors. No evidence of VHE $\\gamma$-ray emission was seen from 39 GRBs using this detector~\\citep{atkins05,abdo07}. \\citet{atkins00} reported a possible VHE $\\gamma$-ray enhancement coincident with GRB~970417A (with a post-trials probability $1.5\\times10^{-3}$ of being a background fluctuation) using Milagrito, the forerunner of Milagro.  In this paper, we report on the first completely simultaneous observation with an IACT instrument of a $\\gamma$-ray burst (GRB~060602B) using H.E.S.S. The burst position fell serendipitously at the edge of the FoV of the H.E.S.S. cameras when the burst occurred. ",
        "conclusions": "On 2006 June 2, the first completely simultaneous observations of a $\\gamma$-ray burst (GRB~060602B) in hard X-rays and in VHE $\\gamma$-rays with an IACT instrument were obtained.  The burst position was observed with H.E.S.S. at VHE energies before, during, and after the burst. A search for a VHE $\\gamma$-ray signal coincident with the burst event, as well as before and after the burst, yielded no positive result. The 99\\% confidence level flux upper limit ($>$1~TeV) for the prompt phase of GRB~060602B is $2.9\\times10^{-9}\\,\\mathrm{erg\\,cm^{-2}\\,s^{-1}}$.  The nature of GRB~060602B is not yet clear, although a Galactic origin seems to be more likely. The complete and simultaneous coverage of the burst with an IACT instrument operating at VHE energies places constraints either in the Galactic X-ray binary scenario or the cosmological GRB scenario."
    },
    "0809/0809.5178_arXiv.txt": {
        "abstract": "{ IRAM 30m $^{12}$CO(1--0) and $^{12}$CO(2--1) HERA observations are presented for the ram-pressure stripped Virgo spiral galaxy NGC~4522. The CO emission is detected in the galactic disk and the extraplanar gas. The extraplanar CO emission follows the morphology of the atomic gas closely but is less extended. The CO maxima do not appear to correspond to regions where there is peak massive star formation as probed by H$\\alpha$ emission. The presence of molecular gas is a necessary but not sufficient condition for star formation. Compared to the disk gas, the molecular fraction of the extraplanar gas is 30\\,\\% lower and the star formation efficiency of the extraplanar gas is about 3 times lower. The comparison with an existing dynamical model extended by a recipe for distinguishing between atomic and molecular gas shows that a significant part of the gas is stripped in the form of overdense arm-like structures. It is argued that the molecular fraction depends on the square root of the total large-scale density. Based on the combination of the CO/H$\\alpha$ and an analytical model, the total gas density is estimated to be about 4 times lower than that of the galactic disk. Molecules and stars form within this dense gas according to the same laws as in the galactic disk, i.e. they mainly depend on the total large-scale gas density. Star formation proceeds where the local large-scale gas density is highest. Given the complex 3D morphology this does not correspond to the peaks in the surface density. In the absence of a confining gravitational potential, the stripped gas arms will most probably disperse; i.e. the density of the gas will decrease and star formation will cease. ",
        "introduction": "}  The best-studied case of active ram-pressure stripping of a cluster galaxy is NGC~4522 located in the Virgo cluster. Only in this cluster is the resolution of radio telescopes sufficient ($\\sim 20'' = 1.6$~kpc\\footnote{We use a distance of 17~Mpc for the Virgo cluster.}) for a detailed analysis of the gas morphology and kinematics. NGC~4522 is a rather small ($D_{25}=4'=20$~kpc) edge-on Sc galaxy with a rotation velocity of $\\sim 100$~km\\,s$^{-1}$. It is strongly H{\\sc i} deficient ($DEF=0.6$; Helou et al. 1984). Its projected distance to the cluster center (M~87) is large ($\\sim 1$~Mpc), and its radial velocity with respect to the Virgo cluster mean is high ($1150$~km\\,s$^{-1}$). H{\\sc i} and H$\\alpha$ observations (Kenney et al. 2004; Kenney \\& Koopmann 1999) show a heavily truncated gas disk at a radius of 3~kpc, which is $\\sim 40$\\% of the optical radius, and a significant amount of extraplanar gas to the west of the galactic disk. The one-sided extraplanar atomic gas distribution shows high column densities, comparable to those of the adjacent galactic disk. The 6~cm polarized radio continuum emission shows a maximum at the eastern edge of the galactic disk, on the opposite side of the extraplanar gas and star formation. Since the stellar disk is symmetric and undisturbed (Kenney \\& Koopmann 1999), a tidal interaction is excluded as the origin of the peculiar gas distribution of NGC~4522. Thus,  this galaxy undergoes ram-pressure stripping due to the galaxy's rapid motion within the hot and tenuous intracluster gas (ICM) of the Virgo cluster.  Vollmer et al. (2006) made a dynamical model that includes the effects of ram pressure for NGC~4522. The model successfully reproduces the large-scale gas distribution and the velocity field. By assuming a Gaussian distribution of relativistic electrons, they obtained the distribution of polarized radio continuum emission, which reproduces the VLA observations of polarized radio continuum emission at 6~cm. The observed maximum of the polarized radio continuum emission is successfully reproduced. The eastern ridge of polarized radio continuum emission is therefore due to ram pressure compression of the interstellar medium (ISM) and its magnetic field. The dynamical model and the analysis of the stellar populations of the outer gas-free disk using optical spectra and UV photometry (Crowl \\& Kenney 2006) indicate that the ram pressure maximum occurred only $\\sim 50$--$100$~Myr ago. This scenario has one important caveat: the large projected distance of NGC~4522 ($1$~Mpc) to the center of the Virgo cluster (M87). Assuming a static smooth ICM and standard values for the ICM density and the galaxy velocity, the ram pressure at that location seems to be too low by an order of magnitude to produce the observed truncation of the gas disk. A natural explanation for the enhanced ram pressure efficiency  is that the intracluster medium is not static but moving due to the infall of the M49 group of galaxies from behind (Kenney et al. 2004, Vollmer et al. 2004, Vollmer et al. 2006). In this case the galaxy has just passed the region of highest intracluster medium velocity.  While we know from the H$\\alpha$ observations (Kenney et al. 2004, Kenney \\& Koopmann 1999) that stars are forming in the extraplanar gas, we do not know the distribution of molecular gas in these regions. How does the complex multiphase interstellar medium respond to ram pressure stripping? Can we model the molecular gas content during the interaction using simplified recipes? Can dense molecular gas decouple from the ram pressure wind as suggested for NGC~4438 (Vollmer et al. 2005)? In this article we present IRAM 30m $^{12}$CO(1--0) and $^{12}$CO(2--1) HERA observations of NGC~4522 to investigate the fate of the stripped gas.  We present our CO observations in Sec.~\\ref{sec:observations} followed by the observational results in Sec.~\\ref{sec:results}. The detection of ram pressure wind decoupled molecular clouds is reported in Sec.~\\ref{sec:decoupled}. In Sec.~\\ref{sec:comparison} we compare our CO observations to existing H{\\sc i} and H$\\alpha$ emission distributions (Kenney et al. 2004) and to the dynamical model of Vollmer et al. (2006). The molecular fraction and star formation efficiencies are discussed in Sec.~\\ref{sec:discussion} and we give our conclusions in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "}  We present IRAM 30m $^{12}$CO(2--1) HERA and $^{12}$CO(1--0) observations of the ram pressure stripped Virgo spiral galaxy NGC~4522. We directly compare the CO spectra to the H{\\sc i} data cube of Kenney et al. (2004). In a second step a CO emission distribution map is produced in the regions where H{\\sc i} is detected. The CO emission distribution is compared to H$\\alpha$ observations of Kenney et al. (2004) and to the model distribution of molecular gas derived from dynamical simulations of Vollmer et al. (2006). A map of the distribution of star formation based on the numerical cloud--cloud collisions is produced which is then compared to the H$\\alpha$ emission distribution. The 3D model snapshot allows us to deproject the observed features and understand their origin. From this work we conclude that \\begin{enumerate} \\item CO emission is associated with the extraplanar atomic gas. The morphology of the molecular gas closely follows but is less extended than the H{\\sc i} morphology. \\item In the northern part of the galactic disk we find CO emission without an H{\\sc i} counterpart. We interpret this detection as wind-decoupled molecular clouds as observed in NGC~4438 (Vollmer et al. 2005). \\item In the extraplanar region CO emission is always associated with sites of massive star formation as probed by H$\\alpha$ emission. At the resolution of our observations, there is no correlation between the CO and H$\\alpha$ peaks. \\item A model using a molecular fraction proportional to the square root of the gas density qualitatively reproduces our CO observations. \\item The model star formation distribution, which is numerically based on cloud--cloud collisions, qualitatively reproduces the observed H$\\alpha$ emission distribution. \\item The deprojection of the model extraplanar gas shows that a significant part of the gas is stripped in the form of relatively dense arms. \\item The formation of molecular clouds, and subsequent star formation, occurs at peaks in the large-scale volume density of the gas, with no clear difference with respect to the disk despite the very different conditions (i.e. stellar density dominates in the disk but is negligible in the extraplanar material). \\item In the disk gas, the molecular and atomic fractions are about equal whereas in the extraplanar gas, there is twice as much HI as H$_2$, assuming a standard $N({\\rm H}_2) / I_{\\rm CO}$ conversion ratio. \\item The star formation efficiency of the extraplanar gas is about 3 times lower than that of the galactic disk. \\item Using the analytical framework of Vollmer \\& Beckert (2003) we find that the overall total gas density and volume filling factor of self-gravitating clouds are factors of 4 and 2.5 lower respectively compared to the galactic disk. \\end{enumerate} In the early phases of ram pressure stripping ($\\sim 50$~Myr after peak ram pressure; Vollmer et al. 2006) a significant part of the stripped gas is in the form of relatively dense arms. At the same time some very dense molecular clouds, representing a tiny fraction of the stripped gas, can decouple from the ram pressure wind. Molecules and stars form within the stripped dense gas according to the same laws as in the galactic disk, i.e. they mainly depend on the overall total gas density. Star formation proceeds where the local large-scale gas density is highest. Given the complex 3D morphology this does not necessarily correspond to the peaks of the surface density. In the absence of a confining gravitational potential these stripped gas arms will most probably disperse, i.e. their density will decrease and star formation will cease."
    },
    "0809/0809.1271_arXiv.txt": {
        "abstract": "We investigate the impact of neutral hydrogen (H\\,{\\sevensize I}) in galaxies on the statistics of 21-cm fluctuations using analytic and semi-numerical modelling.  Following the reionisation of hydrogen the H\\,{\\sevensize I} content of the Universe is dominated by damped absorption systems (DLAs), with a cosmic density in H\\,{\\sevensize I} that is observed to be constant at a level equal to $\\sim2\\%$ of the cosmic baryon density from $z\\sim1$ to $z\\sim5$. We show that extrapolation of this constant fraction into the reionisation epoch results in a reduction of 10-20\\% in the amplitude of 21-cm fluctuations over a range of spatial scales. The assumption of a different percentage during the reionisation era results in a proportional change in the 21-cm fluctuation amplitude. We find that consideration of H\\,{\\sevensize I} in galaxies/DLAs reduces the prominence of the H\\,{\\sevensize II} region induced {\\em shoulder} in the 21-cm power spectrum (PS), and hence modifies the scale dependence of 21-cm fluctuations. We also estimate the 21cm-galaxy cross PS, and show that the cross PS changes sign on scales corresponding to the H\\,{\\sevensize II} regions. From consideration of the sensitivity for forthcoming low-frequency arrays we find that the effects of HI in galaxies/DLAs on the statistics of 21-cm fluctuations will be significant with respect to the precision of a PS or cross PS measurement. In addition, since overdense regions are reionised first we demonstrate that the cross-correlation between galaxies and 21-cm emission changes sign at the end of the reionisation era, providing an alternative avenue to pinpoint the end of reionisation. The sum of our analysis indicates that the H\\,{\\sevensize I} content of the galaxies that reionise the universe will need to be considered in detailed modelling of the 21-cm intensity PS in order to correctly interpret measurements from forthcoming low-frequency arrays. ",
        "introduction": "The process of hydrogen reionisation is thought to have started with ionised (H\\,{\\sevensize II}) regions around the first galaxies, which later grew to surround groups of galaxies. Reionisation completed once these H\\,{\\sevensize II} regions overlapped and occupied most of the volume between galaxies. Much recent theoretical attention has focused on the power spectrum (PS) of 21-cm emission from neutral hydrogen (H\\,{\\sevensize I}) during the reionisation era (Zaldarriaga, Furlanetto \\& Hernquist~2004; Furlanetto et al.~2004; Morales et al.~2005; Bowman et al.~2006). In particular, the ionisation structure of the intergalactic medium (IGM) owing to UV emission associated with star formation has been studied in detail using analytic~(Furlanetto et al.~2004; Barkana~2007), numerical~(McQuinn et al.~2006; Iliev et al.~2008), and more recently, semi-numerical models (Zahn et al.~2007; Mesinger \\& Furlanetto~2007; Geil \\& Wyithe~2008).  These studies describe a scenario in which very large H\\,{\\sevensize II} regions form around clustered sources within overdense regions of the IGM. The formation of these H\\,{\\sevensize II} regions has a significant effect on the shape of the 21-cm PS because information on the small scale features in the density field is erased from the signal originating within the ionised regions. Conversely, the creation of large ionised regions imprints large scale features on the distribution of 21-cm intensity. The sum of these effects is to move power from small to large scales, leaving a shoulder shaped feature on the PS at the characteristic scale of the H\\,{\\sevensize II} regions (Furlanetto et al.~2004). The detailed morphology of the ionisation structure will therefore yield information about both the ionising sources and the structure of absorbers in the IGM on small scales~(McQuinn et al.~2006).  As reionisation leaves a strong imprint on the 21-cm PS, measuring the latter has become a key goal for learning about the reionisation epoch (Furlanetto, Oh \\& Briggs~2006, and references therein). In addition, since reionisation is driven by galaxy formation, whose statistics reflect those of the underlying density field, the detection of the redshifted 21-cm signal will not only probe the astrophysics of reionisation, but also the matter PS during the epoch of reionisation (McQuinn et al.~2006; Bowman, Morales \\& Hewitt~2007).  Finally, it has been recognised that cross-correlating galaxy surveys with 21-cm maps could yield powerful new insights into the morphology of reionisation, as well as eliminate some of the difficulties related to foreground removal (Furlanetto \\& Lidz 2007; Wyithe \\& Loeb 2007; Lidz et al 2008).  With these ideas as motivation, several experiments are currently under development that aim to detect the 21-cm signal during reionisation, including the Low Frequency Array\\footnote{http://www.lofar.org/} (LOFAR), the Murchison Widefield Array\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} (MWA) and the Precision Array to Probe Epoch of Reionization\\footnote{http://astro.berkeley.edu/$\\sim$dbacker/eor/} (PAPER), and more ambitious designs are being planned such as the Square Kilometer Array\\footnote{http://www.skatelescope.org/} (SKA).   Up until now only the component of H\\,{\\sevensize I} residing in the IGM has been considered in relation to forecasts of the statistical 21-cm signal. However, after the completion of reionisation there is known to be a residual H\\,{\\sevensize I} fraction of a few percent in high density clumps which are believed to reside within galaxies (e.g., Prochaska et al 2005).  This high density contribution to the H\\,{\\sevensize I} content of the Universe is also present during the reionisation era (we refer to this high density H\\,{\\sevensize I} as \\textit{galactic H\\,{\\sevensize I}} throughout the paper). Moreover, because galaxies at high redshift are biased relative to the density field, this galactic H\\,{\\sevensize I} contribution could provide a significant perturbation to the predicted statistics of 21-cm fluctuations.   Wyithe \\& Loeb~(2007) have modelled the density dependent reionisation process using a semi-analytic model that incorporates the important physical processes associated with galaxy bias and radiative feedback. In agreement with numerical simulations~(McQuinn et al.~2006; Iliev et al.~2008), this model demonstrates that galaxy bias leads to enhanced reionisation in overdense regions. In this paper we use this model and its semi-numerical extension to explore the effect that the H\\,{\\sevensize I} content of the biased galactic sources of reionisation would have on the fluctuations in redshifted 21-cm emission. We begin by describing our density dependent model of reionisation in \\S~\\ref{model}. We then summarise the observed evolution of the H\\,{\\sevensize I} density in \\S~\\ref{HI}, and our results for the evolution of the 21-cm signal in \\S~\\ref{modelev}. We next describe the effect of galactic H\\,{\\sevensize I} on the 21-cm signal using analytic and semi-numerical models in \\S~\\ref{model1}-\\ref{seminum}, and discuss prospects for detection in \\S~\\ref{noise}. We conclude in \\S~\\ref{conclusion}. Throughout this paper we adopt a concordance cosmology for a flat $\\Lambda$CDM universe, $(\\Omega_{\\rm m},\\Omega_{\\Lambda},\\Omega_{\\rm b},h,\\sigma_8,n) = (0.27,0.73,0.046,0.7,0.8,1)$, consistent with the constraints from Komatsu et al.~(2008). All distances are in comoving units unless stated otherwise. ",
        "conclusions": "\\label{conclusion}   In this paper we have investigated the impact of H\\,{\\sevensize I} in galaxies on the statistics of 21-cm fluctuations using analytic and semi-numerical models.  Our models are unable to self-consistently compute the galactic H\\,{\\sevensize I} content of galaxies prior to the end of reionisation. As an input to our model we have therefore assumed that during the reionisation era 2\\% of hydrogen is in the form of H\\,{\\sevensize I} and located within galaxies. This number is motivated by observations of the mass weighted fraction of cosmic hydrogen in H\\,{\\sevensize I} after the end of reionisation, which is dominated by damped absorption systems and which has a constant value of $\\sim2\\%$ between $z\\sim1$ and $z\\sim5$. Our modelling shows that this assumption results in a reduction of 10-20\\% in the amplitude of 21-cm fluctuations over a range of spatial scales, and over a large fraction of the reionisation era. In addition to the amplitude of 21-cm fluctuations we have also modelled the cross-correlation between galaxies and 21-cm emission. We find that the inclusion of galaxies decreases the amplitude of the cross-correlation by a few tens of percent. In addition, the cross-correlation between galaxies and 21-cm emission will change sign at the end of the reionisation era, providing an alternative avenue to pinpoint the end of reionisation.  We find that our analytic estimates become less applicable once H\\,{\\sevensize II} regions become a dominant feature of the IGM. We have therefore supplemented our analytic estimates with semi-numerical modelling of the three dimensional ionisation structure of the IGM. We find that the qualitative conclusions from our analytic calculations are vindicated by this modelling. In addition to the above estimates of auto and cross-correlation, we have used our semi-numerical model to compute 21-cm power spectra, and 21cm-galaxy cross power spectra. The effect of galactic H\\,{\\sevensize I} is to change the shape of the 21-cm PS rather than just the amplitude. In particular, because the galactic H\\,{\\sevensize I} is biased towards H\\,{\\sevensize II} regions, we find that the H\\,{\\sevensize II} region induced {\\em shoulder} in the PS is less prominent when the contribution of galactic H\\,{\\sevensize I} is considered.  The amplitude of the 21-cm PS is modified in a scale dependent way by up to 20\\% when a 2\\% galactic H\\,{\\sevensize I} fraction is included. Of course the galactic H\\,{\\sevensize I} fraction is very uncertain at high redshift. However we find that the fractional modification of the 21-cm fluctuation statistics is approximately proportional to the density of galactic H\\,{\\sevensize I}. Therefore, if we assume a value lower(higher) than the 2\\% observed at $z<5$, then our results for the error introduced through neglect of galactic H\\,{\\sevensize I} will be over(under)estimated. On the other hand, the change of sign in the cross-correlation between galaxies and the 21-cm signal is robust to our assumptions.   We find that the inclusion of galactic H\\,{\\sevensize I} lessens the amplitude of the anti-correlation between galaxies and 21-cm emission since a fraction of H\\,{\\sevensize I} is now co-located with the galaxies inside the H\\,{\\sevensize II} regions. In addition, the cross PS changes sign on small scales, which reflects the correlation of galaxies with the galactic H\\,{\\sevensize I} inside the H\\,{\\sevensize II} regions, where no power is contributed in 21-cm fluctuations of the IGM. Thus, the scale at which the 21cm-galaxy cross PS changes sign could be used to probe the scale of H\\,{\\sevensize II} regions late in the reionisation era.   We have estimated the sensitivity of the MWA to the spherically averaged 21-cm PS and 21cm-galaxy cross PS. Our calculations illustrate that the effect of the galactic H\\,{\\sevensize I} on 21-cm fluctuations is at a level that would be significant with respect to the sensitivity of the MWA late in reionisation. However the low resolution of the MWA would mean that the shape change of the PS owing to the galactic H\\,{\\sevensize I} could not be detected until very close to overlap. On the other hand, when combined with a suitable galaxy redshift survey, the resolution of the MWA would be sufficient to detect the sign change of the 21cm-galaxy cross PS at the scale of the H\\,{\\sevensize II} regions late in reionisation. A followup telescope comprising 10 times the collecting area of the MWA would measure the effects of galactic H\\,{\\sevensize I} with high significance.  In summary, our modelling shows that the H\\,{\\sevensize I} content of the galaxies that reionise the universe provides a significant contribution to the statistics of 21-cm fluctuations. The galactic H\\,{\\sevensize I} contribution to the 21-cm intensity will therefore need to be considered in detailed modelling of the 21-cm intensity PS in order to correctly interpret measurements from the next generation of low-frequency arrays.  {\\bf Acknowledgements} The research was supported by the Australian Research Council (JSBW). LW and PMG acknowledge the support of Australian Postgraduate Awards. SPO acknowledges support from NASA grant NNG06GH95G.  \\newcommand{\\noopsort}[1]{}"
    },
    "0809/0809.4724_arXiv.txt": {
        "abstract": "We discuss the preliminary results of spectral analysis simulations involving anticipated correlated multi-wavelength observations of gamma-ray bursts (GRBs) using Swift's Burst Alert Telescope (BAT) and the Gamma-Ray Large Area Space Telescope's (GLAST) Burst Monitor (GLAST-GBM), resulting in joint spectral fits, including characteristic photon energy $\\left(E_{peak}\\right)$ values, for a conservative annual estimate of $\\sim30$ GRBs. The addition of BAT's spectral response will (i) complement in-orbit calibration efforts of GBM's detector response matrices, (ii) augment GLAST's low energy sensitivity by increasing the $\\sim20-100$ keV effective area, (iii) facilitate ground-based follow-up efforts of GLAST GRBs by increasing GBM's source localization precision, and (iv) help identify a subset of non-triggered GRBs discovered via off-line GBM data analysis. Such multi-wavelength correlative analyses, which have been demonstrated by successful joint-spectral fits of Swift-BAT GRBs with other higher energy detectors such as Konus-WIND and Suzaku-WAM, would enable the study of broad-band spectral and temporal evolution of prompt GRB emission over three energy decades, thus potentially increasing the science return without placing additional demands upon mission resources throughout their contemporaneous orbital tenure over the next decade. ",
        "introduction": "The Swift MIDEX explorer mission \\cite{Gehrels:2007} (anticipated to operate until at least $\\sim2014$), comprised of the wide-field ($\\sim1.4$ sr, half-coded) hard X-ray (15-150 keV) Burst Alert Telescope (BAT) \\cite{Barthelmy:2005}, and the narrow-field (0.2-10 keV) X-Ray (XRT) and (170-600 nm) Ultraviolet-Optical (UVOT) Telescopes, has revolutionized our understanding of GRBs. The intrinsic multi-wavelength instrumentation, coupled with a rapid ($<\\sim$100 seconds) autonomous slew capability, has ushered in an unprecedented era of source localization precision $\\left(<\\sim1^{\\prime}-4^{\\prime}\\right)$ that is disseminated in real-time ($\\sim10$ seconds) via the GRB Coordinate Network (GCN), which enables broad-band international observational campaigns. Swift's unique dynamic response, in conjunction with correlative ground-based follow-up efforts, has resulted in $\\sim70$ observed spectroscopic redshifts, while its X-ray sensitivity has revealed evidence of extended soft emission in some short GRBs.  The Gamma-Ray Large Area Space Telescope, which is now known as the Fermi Gamma-Ray Space Telescope (\\emph{Fermi}), mission was launched on June 11, 2008 and has an anticipated orbital lifetime spanning until at least $\\sim2018$. A goal of GLAST, which is comprised of the (20 MeV to $>$300 GeV) Large Area Telescope (LAT) and the (8 keV - 30 MeV) GLAST Burst Monitor (GBM), which is now known as the GRB Burst Monitor (GBM), is to study transient gamma-ray sources, while a direct GBM experimental objective is to identify and study GRBs. Common scientific interest between GLAST and Swift, in the context of prompt GRB emission, provides strong motivation for a cross-calibration via correlative observations of GRBs, resulting in joint spectral energy fits, thus enabling the analysis of multi-wavelength spectral and temporal evolution. Cross-calibration with BAT would also complement in-orbit calibration efforts of GBM's detector response matrices and would be especially critical during GLAST's first year, since it would facilitate a smooth transition from performance diagnostics to science operations.  The GBM, consisting of 12 NaI (8-1000 keV) and 2 BGO (0.15-30 MeV) detectors, will monitor $\\sim8$ steradians of the sky, and, in concert with LAT, enables GLAST to continuously span 7 energy decades, but not at the same sensitivity. As illustrated in Figure~\\ref{Overlapped_Effective_Areas}, GLAST's effective area drops by over $\\sim1.5$ orders of magnitude from LAT (GeV) to GBM-NaI (keV) energies, while the (masked) BAT ($\\sim$20-100 keV) low energy effective area surpasses GBM-NaI's by over a factor of $\\sim3$. Furthermore, although Swift has detected $>\\sim300$ GRBs, the majority of E$_{peak}$ (characteristic photon energy) values lie beyond BAT's canonical energy range, since $\\overline{E}_{peak}\\sim250$ keV. Thus, correlated Swift-BAT/GLAST-GBM GRB observations would simultaneously augment GLAST's low energy response while increasing the number of $E_{peak}$ for BAT GRBs observed by GBM. Additionally, since Swift's high fidelity localization precision surpasses GBM's by over $\\sim2-3$ orders of magnitude, we expect that $\\sim30\\%$ of bursts in the joint BAT-GBM analysis would be accompanied by panchromatic ground-based follow-up observations resulting in observed spectroscopic redshifts and host galaxy identifications.  \\begin{figure} \\includegraphics[height=.54\\textheight]{BAT_NaI_BGO_KW_Effective_Areas_v_Energy3.eps} \\caption{Effective areas for Swift-BAT, GBM-NaI, GBM-BGO, GLAST-LAT and Konus-WIND. \\emph{Note the break in the x-axis.}} \\label{Overlapped_Effective_Areas} \\end{figure} ",
        "conclusions": "Current efforts include inter-calibrating BAT and GBM, in collaboration with the GBM team, via joint spectral fits of mutually observed GRBs. Previous studies of joint-fit GRB data sets \\cite{Krimm:2006c} enhanced our understanding of burst parameter classifications, explored GRB emission geometry, and tested the viability of various redshift estimation methods \\cite{Stamatikos:2008c}. In addition, a more accurate normalization between GLAST prompt emission and Swift afterglow spectra will facilitate the determination of GRB energy budgets. Overall, broad-band correlative studies will enable the investigation of spectral and temporal evolution \\cite{Stamatikos:2008e} over unprecedented decades of energy, shedding light on the connection between electromagnetic pulse asymmetry, width and spectral softening \\cite{Norris:1996}, while facilitating multi-messenger searches such as those for correlated leptonic emission \\cite{Stamatikos:2005,Stamatikos:2006b}. Hence, correlated Swift-BAT/GLAST-GBM observations enhance their science return and benefit the broader astronomical community, without additional demands upon mutual mission resources, thus underscoring an aspect of Swift's operational and scientific relevancy in the impending era of GLAST.    \\begin{theacknowledgments} The authors are grateful to the members of the GBM team for very fruitful discussions in regards to this analysis. M. Stamatikos is supported by an NPP Fellowship at NASA-GSFC administered by ORAU. \\end{theacknowledgments}"
    },
    "0809/0809.0511_arXiv.txt": {
        "abstract": "{Recent studies have suggested that moving groups have a dynamic or ``resonant'' origin. Under this hypothesis, these kinematic structures become a powerful tool for studying the large-scale structure and dynamics of the Milky Way.} {Here we aim to characterize these structures in the $U$--$V$--$age$--$[Fe/H]$ space and establish observational constraints that will allow us to study their origin and evolution.} {We apply multiscale techniques -- wavelet denoising (WD)-- to an extensive compendium of more than 24000 stars in the solar neighbourhood with the best available astrometric, photometric and spectroscopic data.} {We confirm  that the dominant structures in the $U$--$V$ plane are the branches of Sirius, Coma Berenices, Hyades-Pleiades and Hercules.  These branches are nearly equidistant in this kinematic plane and they show a negative slope. The abrupt drops in the velocity density distribution are characterized. We find a certain dependence of these kinematic structures on Galactic position with a significant change of contrast among substructures inside the branches. A large spread of ages is observed for all branches. The Hercules branch is detected in all subsamples with ages older than $\\sim 2\\Gyr$ and the set of the other three branches is well established for stars $>400 \\Myr$. The age-metallicity relation of each branch is examined and the relation between kinematics and metallicity is studied.} {Not all of these observational constraints are successfully explained by the recent models proposed for the formation of such kinematic structures. Simulations incorporating stellar ages and metallicities are essential for future studies. The comparison of the observed and simulated distributions obtained by WD will provide a physical interpretation of the existence of the branches in terms of local or large-scale dynamics.} ",
        "introduction": "After the pioneering work of \\citet{proctor1869}, \\citet{kapteyn05} and \\citet{lindblad25}, Eggen established the spatial and kinematic properties of several {\\it stellar streams} - hereafter classic moving groups - formed by stars with similar kinematics in the solar neighbourhood (see \\citealt{eggen96} and references therein). Since these structures share their kinematics with certain open clusters, he worked exhaustively on the hypothesis that moving groups are a result of the dispersion (resulting from both external and internal causes) of stellar clusters.  The advent of Hipparcos astrometric data led to the definitive establishment of the existence of moving groups and to the recognition of substructures within them \\citep{chereul98,asiain99a}. In addition, together with \\citet{dehnen98}, these authors first attempted to evaluate the evolutionary state of the members of moving groups, which is necessary in order to establish their origin and evolution. \\citet{dehnen98} used subsamples composed of stars with different spectral types to study the velocity distribution of old and young stars. As no photometric data were available, \\citet{chereul98} could only use what they called ``palliative'' ages to show that the age distribution of the stars in moving groups seems to be similar to that observed for their whole sample. More precise ages, derived using Str{\\\"o}mgren photometry, were used by \\citet{asiain99b}, who observed that the velocity dispersion of several substructures within the Pleiades moving group is compatible with that expected for the evolution of a stellar complex \\citep{efremov88}. This conclusion was restricted to groups of stars with ages of up to about $1\\Gyr$, as both differential Galactic rotation and disk heating would have dispersed older groups among the field stars. Using an adaptive kernel and wavelet transform (WT) analysis, \\citet{skuljan99} studied a sample of 4000 Hipparcos stars and found that the distribution function in the $U$--$V$ plane is characterized by a few branches that are diagonal, parallel and roughly equidistant. Later on, \\citet{famaey05} used a maximum-likelihood method based on a Bayesian approach to divide a sample of giant stars into several kinematic groups. A very wide range of ages for each of these structures in their corresponding H-R diagram was observed.  The first theoretical arguments in favour of a different dynamic origin of moving groups were put forward by \\citet{mayor72} and \\citet{kalnajs91}. Following this latter author, \\citet{dehnen98} pointed out that orbital resonances could be the cause of the existence of most moving groups observed in the solar neighbourhood. \\citet{skuljan99} related the origin of the branches to the Galactic spiral structure, or to some other global characteristics of the Galactic potential, combined with the initial velocities of the stars. Nowadays, these dynamic or ``resonant'' mechanisms are proposed to be the most plausible explanation for the main moving groups. Other minor kinematic structures in the solar neighbourhood, such as HR1614, do seem to be remnants of a dispersed star-forming event \\citep{desilva07} and it has been proposed that other structures are related to accretion events in the Galaxy \\citep{helmi06}.  The work by \\citet{dehnen00} and \\citet{fux00,fux01}, among others, focused on the origin of the Hercules structure and its connection to Galactic bar resonances. Furthermore, by numerically integrating test particle orbits, \\citet{desimone04} showed that stochastic spiral density waves can produce kinematic structures similar to those found by \\citet{skuljan99}. \\citet{quillen05} and \\citet{chakrabarty07} have recently gone one step further and looked for the relation between moving groups and resonances of the non-axisymmetric component of the Galactic potential, i.e. the Galactic bar and spiral structure. They have shown that by using several combinations of the parameters adopted to characterize this potential they are able to reproduce structures in the kinematic plane similar to those actually observed. To restrict the free parameters of the Galactic potential, all this recent work requires new observational constraints such as a characterization of the observed kinematic structures in the $U$--$V$ plane in terms of their evolutionary state, their chemical composition or their possible dependence on Galactic position.  Nowadays, two important observational contributions provide new material to complement the Hipparcos and Tycho astrometric data: i) the CORAVEL radial velocity data for a significant number of late-type stars belonging to the Hipparcos catalogue (\\citealt{nordstrom04} for dwarf stars and \\citealt{famaey05} for giant stars) and ii) the uvby--$\\beta$ survey of FGK dwarf stars, which have allowed the derivation of ages and metallicities \\citep{nordstrom04}. Lastly, data on OBA-type stars \\citep{asiain99a,torra00} and M dwarfs \\citep{reid02,bochanski05} complete an extensive sample of more than 24000 stars ready to be used for the characterization of moving groups. It is necessary to complement these new data with more powerful statistical tools to interpret both the observed and simulated data in the 4-dimensional space of $U$--$V$--$age$--$[Fe/H]$.  In this paper, we characterize the observed kinematic structures in $U$--$V$--$age$-$[Fe/H]$ space by applying new statistical techniques with the aim of establishing observational constraints that will help to reveal their origin and evolution. Section \\ref{data} of this paper presents the observational data we use together with their precisions and possible biases. Section \\ref{method} contains a description of the statistical method of the wavelet denoising (WD) we apply. Section \\ref{velocity} characterizes in depth the structures in the velocity plane. Then, Sect. \\ref{age} and Sect. \\ref{feh} analyse the age and metallicity distributions of these structures. Finally, the main outcomes of the work and perspectives for the future are summarized in Sect. \\ref{conclusions}. Subsequent papers will further this research by using these constraints in combination with test particle simulations with a suitable Galactic potential model. ",
        "conclusions": "\\label{conclusions}  Moving groups are becoming a powerful tool for studying the large-scale structure and dynamics of the Milky Way. We apply multi-scale techniques --wavelet denoising-- to an extensive compendium of more than 24000 stars in the solar neighbourhood to characterize the observed kinematic structures in  $U$--$V$--$age$--$[Fe/H]$ 4-dimensional space. In this paper we focus our analysis on establishing the observational constraints that will allow us to study the origin and evolution of these structures.  The advent of Hipparcos astrometric data led to the definitive recognition of a non-smooth distribution function of the Galactic disc. Our results corroborate this and go one step further towards characterizing the velocity distribution function. Branches connecting the classic moving groups in the $U$--$V$ plane are definitely the dominant structures. We confirm the existence of the Sirius, Coma Berenices, Hyades-Pleiades and Hercules branches. The first three branches are spaced at intervals of approximately 15 $\\kms$ with no significant variations with age or spectral type. The Hercules branch is located about 30 $\\kms$ from the Hyades-Pleiades branch. The four branches present a negative slope of $\\beta \\sim 16\\deg$ in the $U$--$V$ plane, lower than that found by \\citet{skuljan99}. We have studied the density drops in the $U$--$V$ plane and the existence of abrupt edge lines is corroborated, especially at low $U$ and high $V$ and between the Hyades-Pleiades and the Hercules branches. Each branch presents a different density distribution in its extremes near these edge lines. These features definitely rule out the classic idea of a smooth velocity field distribution. The use of the same statistical method for the samples with different spectral types has made possible an exhaustive comparison between their kinematic planes. The branches are present in all the distributions. The study of variations of the kinematic structures with Galactic position has been presented for the KM giant sample. A significant change of contrast among substructures inside the branches is confirmed and the shape of the Hercules branch changes among regions, being more conspicuous in the region towards the anti-rotation direction with respect to the centre of this sample.  We have analysed, for the first time, the age and metallicity distributions of the branches. We confirm an extended age distribution for all the branches  but a different minimum age of their stars. The set of the three main branches of Hyades-Pleiades, Coma Berenices and Sirius is well established for stars $>400 \\Myr$ and the extended branch-like shape of Hercules is detected in all subsamples with ages $> 2 \\Gyr$.  The relative density between the Hyades-Pleiades and the Hercules branches clearly decreases with age. We find a periodicity in age of about 500-600 $\\Myr$ in the Hyades-Pleiades branch. For the other branches only an outline of the shape of the whole age distribution is observed. A wide range of metallicity is found for each branch, especially for Hercules with a higher metallicity dispersion. A complex relation between kinematics and metallicity has been established.  Concerning the age-metallicity relation, differences are observed between branches. The periodic bumps in the age distribution for the Hyades-Pleiades branch show a decrease in metallicity with increasing age.  All the above observational results and, moreover, the whole set of distributions in the $U$--$V$--$age$--$[Fe/H]$ space presented here, are the fundamental elements necessary to check the present and future dynamic models proposed for the formation of kinematic structures. While some of these current models already explain some of the observational results, other important observational features remain unexplained. For instance, the simulations of \\citet{desimone04} produce branches with a slope which fits that measured here. Furthermore, the model in \\citet{quillen05} agrees with the fact that the Hyades-Pleiades branch is more metallic than the Coma Berenices branch. However, features such as the mix of stars in the Hyades and Pleiades moving groups in periodic clumps in age in the same branch --perhaps reflecting  a common mechanism for the formation of these two kinematic structures-- or the peculiar metallicity distribution inside branches demand an in-depth explanation.  To make further progress in this field, simulations obtained from orbit integration under a model for the Galactic potential are being undertaken (Antoja et al. 2008, in preparation). After checking that the simulated kinematic structures critically depend on the gradients of the spiral arm potential, two different models for the Galactic spiral arms are being explored: the classic cosine perturbation \\citep{lin71} and a 3D mass distribution (superposition of the inhomogeneous oblate spheroids in \\citealt{pichardo03}). We are also incorporating stellar ages and metallicities in the simulations to trace the origin and evolutionary state of the test particles, which is another key element. Of course, other ingredients such as the Galactic bar will also be included to analyse the kinematic structures in an overall dynamic context.  Thus, the present detection and characterization of the branches in $U$--$V$--$age$--$[Fe/H]$ 4-dimensional space provide us with the main properties summarized above. It is the whole set of distributions in this space that contains all the information. The application of the same statistical technique to both the observed and  simulated data will allow a direct comparison between them, and thus be a powerful test of the models for the formation of the structures. This comparison will eventually offer a physical interpretation of the formation of the kinematic structures in terms of local or large-scale dynamics."
    },
    "0809/0809.2980_arXiv.txt": {
        "abstract": "{We present the first VLT near-IR observations of a gravitationally lensed quasar, using adaptive optics and laser guide star.  These observations can be considered as a test bench for future systematic observations of lensed quasars with adaptive optics, even when bright natural guide stars are not available in the nearby field.  With only 14 minutes of observing time, we derived very accurate astrometry of the quasar images and of the lensing galaxy, with 0.05\\arcsec\\ spatial resolution, comparable to the Hubble Space Telescope (HST). In combination with deep VLT optical spectra of the quasar images, we use our adaptive optics images to constrain simple models for the mass distribution of the lensing galaxy. The latter is almost circular and does not need any strong external shear to fit the data. The time delay predicted for \\obj, assuming a singular isothermal ellipsoid model and the concordance cosmology, is $\\Delta t \\simeq 50$ days.  Our optical spectra indicate a flux ratio between the quasar images of A/B=1.3 in the continuum and A/B=2.2 in both the \\ion{Mg}{ii} and in the \\ion{C}{iii]} broad emission lines.  This suggests that microlensing affects the continuum emission. However, the constant ratio between the two emission lines indicates that the broad emission line region is not microlensed. Finally, we see no evidence of reddening by dust in the lensing galaxy.} ",
        "introduction": "Strong gravitational lensing of distant quasars is a useful tool both for cosmology and for studying luminous and dark matter in massive galaxies.  For a fixed cosmology, the total projected distribution of matter in the galaxy(ies) responsible for the lensing effect can be reconstructed from (i) the configuration of the lensed quasar images on the plane of the sky, (ii) the observed magnification ratio between the quasar images, (iii) measurement of the time delay, and (iv) detailed structures in the lensed host galaxy of the quasar forming arcs and counter-arcs (e.g. Refsdal~\\cite{REFS64}, Claeskens \\& Surdej~\\cite{CLA02}, Kochanek et al.~\\cite{KOC06}).  Finally, the magnification ratios between the quasar images can significantly deviate from the predictions by smooth mass models if the lensing galaxy contains small-scale substructures with $M\\sim 10^{8-9}\\,$\\msol\\, (e.g., Mao \\& Schneider~\\cite{MAO98}, Metcalf~\\cite{MET05}).  Strong gravitational lensing of quasars is therefore a sensitive indirect way of detecting small satellites of lensing galaxies.  High-resolution imaging of lensed quasars is crucial for constraining mass models for the lensing galaxy, determining its luminous mass distribution, and identifying the optical counterpart of any mass substructure (Schechter \\& Moore ~\\cite{SCH93}, Koopmans \\& Treu~\\cite{KOO02}, McKean et al.~\\cite{MCK07}). The Hubble Space Telescope (HST) has provided the community with a large sample of exquisite high-resolution images of gravitational lenses with the CASTLES survey (CfA-Arizona Space Telescope LEns Survey; Mu{\\~n}oz et al.~\\cite{MUN98}) and played a crucial role in the study of lensed quasars (Claeskens et al.~\\cite{CLA07} for a review). Ground-based adaptive optics (AO) imaging now allows us to obtain high spatial resolution data from the ground.  Although both approaches are complementary in depth and field of view, the AO imaging data collected so far of lensed quasars are still very scarce.  With Gemini-North, Courbin et al.~(\\cite{COU02}) obtained the first AO images of a gravitational lens, for the Einstein ring PKS~1830-21. These observations were followed up at the VLT in AO in order to confirm a double lensing galaxy (Meylan et al.~\\cite{MEY05}). More recently, Auger et al.~(\\cite{AUG08}) used AO to observe the doubly imaged quasar SBS~1520+53.  These two targets are easy for AO, as they are located a few arcseconds away from a bright natural guide star.  Systematic AO imaging of all other lensed quasars requires laser guide star (LGS).  McKean et al.~(\\cite{MCK07}) obtained the first AO/LGS images at the Keck of a lensed quasar to search for substructures in the lensing galaxy.  In the present paper, we describe the first AO/LGS images of a gravitational lens at the VLT: the doubly imaged quasar \\obj, at $z_s=1.540$. This quasar has been found to be gravitationally lensed from a systematic search for strong lens systems in the Sloan Digital Sky Survey (Inada et al.~\\cite{INA06}).  The angular separation between the two quasar images is $\\Delta \\theta \\sim 1.5\\arcsec$ and the early-type lensing galaxy has been found to be at $z_l=0.573$ from deep VLT optical spectroscopy (Eigenbrod et al.~\\cite{EIG07}).  Our new VLT AO/LGS images, further improved with image deconvolution, allow us to derive milli-arcsecond astrometry of the quasar images, as well as improved models for the lensing galaxy. We see these observations as a test bench for systematic VLT AO imaging of gravitationally lensed quasars.   \\begin{figure*}[ht!] \\begin{center} \\includegraphics[width=14.0cm]{J0806_AO_BW.jpg} \\caption{{\\it Left:} VLT image of \\obj\\, obtained with the adaptive optics + laser guide star facility.  The PSF FWHM in this 14-minute $H$-band image is about the same as on a HST/NICMOS image in the same band.  {\\it Right:} deconvolved image, as obtained with the MCS algorithm (see text).  The lower cut level in this deconvolved image is set to 1$\\sigma_{\\rm sky}$ in order to show the noise level. North is to the top, east to the left.} \\label{fig:ima} \\end{center} \\end{figure*} ",
        "conclusions": "We obtained the first VLT images of a gravitationally lensed quasar using AO and LGS. Using image deconvolution we derived accurate astrometry of the quasar images and of the lensing galaxy. This astrometry is significantly different from the one in Inada et al.~(\\cite{INA06}) and is not due to a systematic error in the pixel size of the instrument used.  We  find that a  simple isothermal mass  model can account for the new astrometry, with no need of strong external shear.  The predicted time delay for this model, $\\Delta t \\simeq 36.3\\,h^{-1}$  days, is well-suited to an   accurate   measurement  using   photometric  light    curves (see simulations by Eigenbrod et al. \\cite{EIG05}).  Our VLT optical spectra show that \\obj\\, is affected neither by chromatic microlensing nor by dust extinction. The continuum flux ratio based on these spectra agrees within 10\\% with the one derived by Inada et al. one year before, suggesting low-amplitude microlensing-induced variations.  \\obj\\, is therefore a clean lens system to use either for determining of $H_0$ or for the study of the mass distribution in the lensing galaxy, based on a time-delay measurement.  The accuracy of the relative positions of the lensed images derived from our observations of \\obj\\ is similar to the typical accuracy derived from the HST data of other lensed quasars (e.g. Impey et al.~\\cite{IMP98}, Claeskens et al. ~\\cite{CLA06}). This demonstrates that ground-based AO imaging with the help of LGS can be fully competitive with HST observations in the near-IR. About 90\\% of the lensed quasars observable with the VLT are either bright enough to be used as a tip-tilt reference or have potential tip-tilt stars with V $<$ 17 within a 55\\arcsec radius.  This is promising for future systematic observations of gravitational lenses with adaptative optics and laser guide star."
    },
    "0809/0809.0978.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary Galaxies can form in a sufficiently deep gravitational potential so that efficient gas cooling occurs. We estimate that such potential is provided by a halo of mass $M \\gtsim M_{c} \\approx 7.0 \\times 10^{12} ~ (\\Delta_{c}(z) (1+z)^{3})^{-1/2} \\Msun$, where $\\Delta_{c}(z)$ is the mean overdensity of spherically virialized objects formed at redshift $z$, and $M_{c} \\approx 4.0 \\times 10^{11} \\Msun$ at $z = 0$. Based on this criterion, our galaxy samples are constructed from cosmology simulation data by using HiFOF to select subhalos in those FOF halos that are more massive than $M_{c}$. There are far more dark subhalos than galaxy-hosting subhalos. Several tests against observations have been performed to examine our galaxy samples, including the differential galaxy mass functions, the galaxy space density, the projected two point correlation functions (CF), the HODs, and the kinematic pair fractions. These tests show good agreements. Based on the consistency with observations, our galaxy sample is believed to correctly represent galaxies in real universe, and can be used to study other unexplored galaxy properties. ",
        "introduction": "The standard theoretical model for galaxy formation basically consists of two main elements, the cold dark matter (CDM) model and a dark energy field (which may take the form of a cosmological constant, $\\Lambda$). The fundamental assumption made in the theory is that structure grew from weak density fluctuations present in the otherwise homogeneous and rapidly expanding early Universe. Afterwards, the gravitational instability drives these fluctuations into nonlinear regime in a bottom up fashion and they gradually grow to a wealth of structures today. Recent diverse cosmological studies, no matter from large-scale structure observations \\citep[e.g., ][]{spe03}, from supernova data \\citep[e.g., ][]{rie98}, or from light element abundance \\citep[e.g., ][]{whi93}, all seem to support the concordance model. Due to the highly nonlinear nature in the collapse of fluctuations and the subsequent hierarchical build-up of structure, numerical simulations come to play an indispensable role.  In the past thirty years, galaxy clustering has been intensively studied, as the clustering of galaxies has long been an essential testing ground for various cosmological models and galaxy formation scenarios. Especially due to the advent of modern computers, high resolution simulations become possible. Cosmological N-body simulations have developed into a powerful tool for calculating the gravitational clustering of collisionless dark matter from specified initial conditions. As the resolution of simulations increases, some numerical problems such as the overmerging problem \\citep{kly99} can be overcome, and the galactic size scale in a large scale structure simulation can be resolvable. In addition, the growth of galaxy redshift surveys has led to measurements of increasing precision and detail. All these works have made the comparisons between theories and observations possible.  There are four main approaches used in cosmological simulations to study the galaxy clustering and properties. The first one is N-body plus hydrodynamical simulation including gas and dark matter particles \\citep[e.g., ][]{wei04}. The second method is a hybrid method that combines N-body simulations of the dark matter component with semi-analytic treatments of the galaxy formation physics \\citep[e.g., ][]{spr05}. Another approach is using high-resolution collisionless N-body simulations that identify galaxies with \"subhalos\" in the dark matter distribution \\citep[e.g., ][]{kly99,kra04,con06}. The last one is the so called Halo Occupation Distribution (HOD) approach which gives a purely statistical description of how dark matter halos are populated with galaxies \\citep[e.g., ][]{ber03,kra04,zhe05}. It is known that the cold dark matter is the dominant mass component and it interacts only through gravity. To some extent, gravitational dynamics alone should explain the basic features of galaxy clustering. \\cite{ney04} tried to understand how the infrared-selected galaxies populate dark matter halos, paying special attention to the method of halo identification in simulations. They tested the hypothesis that baryonic physics negligibly affects the distribution of subhalos down to the smallest scales yet observed and successfully reproduced the Point Source Catalogue Redshift (PSCz) power spectrum. \\cite{ber03} reported that the HOD prediction of two methods, a semi-analytic model and gasdynamics simulations, agree remarkably well for samples of the same space density. This result indirectly supported the idea that the HOD, and hence galaxy clustering, is driven primarily by gravitational dynamics rather than by processes such as cooling and star formation. On the other hand, \\cite{kra04} adopt a variant of the bound density maxima (BDM) halo-finding algorithm \\citep{kly99} to identify halos and subhalos and use the maximum circular velocity $V_{max}$ as a proxy of halo mass for selecting galaxy sample. Instead of selecting objects in a given range of $V_{max}$, at each epoch they selected objects of a set of number densities consistent with observational number densities and corresponding to (redshift dependent) thresholds in maximum circular velocity, i.e., $n_{i}(>V_{max})$. Their results show that the dependence of correlation amplitude on the galaxy number density in their sample is in general agreement with results from the Sloan Digital Sky Survey. \\cite{con06} instead use the maximum circular velocity at the time of accretion, $V^{acc}_{max}$, for subhalos, and the results show agreement with the observed galaxy clustering in the SDSS data at $z \\approx 0$ and in the DEEP2 samples at $z \\approx 1$ over the range of separations, $0.1 < r_{p}/(h^{-1} Mpc) < 10.0$. However, this work lacks the information of the subhalo mass. Motivated by the results of these previous works, an approach similar to the galaxy identification with subhalos is adopted in this study. Despite some of our simulations include gas particles, we shall ignore the gas component in the identification of galaxies. Our approach in fact combines a subhalo finding algorithm, HiFOF, and a galaxy formation model. It 'directly' locates where the galaxy should be formed and shows agreements with various observation results, such as galaxy mass function, two-point correlation function, etc. This result indicates that our galaxy identification model not only can successfully locate subhalos in simulations that reside inside host halos in real universe but also can give a galaxy mass function close to the observational one.  The paper is structured as follows. In section 2, we describe the simulation and the halo-finding algorithm in use. Thereafter, the galaxy halo model we adopt is discussed. In Section 3, we compare our galaxy samples with observations including the differential mass functions, the number density evolution, the correlation functions at different redshift, the halo occupation distribution, and the pair fraction. Finally, we discuss and summarize our results in Section 4. ",
        "conclusions": "In this study we have used the HiFOF method to locate and identify subhalos, and we advance the idea that galaxies form only in a sufficiently massive halo due to its sufficiently high gas cooling efficiency. It is estimated that the threshold halo mass $M_{c} \\approx 4.0 \\times 10^{11}$ $h^{-1} M_{\\odot}$ at $z = 0$, and our galaxy samples are constructed from the HiFOF subhalos embedded in the halos satisfying this mass threshold. There are far more subhalos that are hosted by halos not meeting the mass threshold. They are dark halos where star formation cannot occur. To test our model, we have analyzed the differential mass function, two-point correlation function (CF), the space density evolution, the HOD, and the pair fraction from our galaxy samples. The results are summarized as follows.    (1) Combing the luminosity function (LF) and the mass to light relation (M/L), a differential mass function can be derived. Based on the lensing signal from Red-Sequence Cluster Survey to determine the galaxy mass, we have corrected the lensing mass in our galaxy samples so as to make fair comparison. We find that the mass functions of our four simulations agree with the LF and M/L data very well.  (2) We compare our galaxy density evolution with the evolution of the observational galaxy number density through the integration of a LF with a proper luminosity cut. Our simulations basically show fair consistency with observations, except at low-redshift. The insufficient density is likely due to the failure of the assumption at low redshift about the redshift-independent mass-cut. The depletion in the galaxy number density also reflects a similar result in pair fractions found at low-redshift.  (3) We have analyzed CFs of our galaxy subhalos at $z = 0$ and at $z = 1$. At $z = 0$, we select our galaxies to have the density $1.507 \\times 10^{-2} \\den$ to make comparison with the volume limited sample ($M_{r} < -19$) of SDSS galaxies. Both the slope and the amplitude are consistent with the SDSS. At $z = 1$, we select samples at $n = 1.3 \\times 10^{-2} \\den$ to setup the same density condition as in the volume-limited sample of bright (corrected), $M_{B} > -19.0$, DEEP2 galaxies. The projected correlation functions are compared, and our profiles are similar to the one found in DEEP2 galaxies.  (4) We also study $r_{0}$ and $\\gamma$ as a function of the space density $n$. Our galaxies reveal the increase of $r_{0}$ as $n$ decreases, a similar result also found in the observation data. In other words, the correlation length $r_{0}$ has luminosity dependence, and the brighter the galaxies, the larger the correlation length. On the other hand, the power index $\\gamma$ of observation data displays a nearly constant behavior over a wide range of $n$ at $z = 0$. Despite a smaller values of $\\gamma$ are obtained around $n \\approx 10^{-2}$ in our simulations, a roughly constant trend on the whole is seen. The same results are reached at $z = 1$ as well.  (5) We have parameterized our HODs with three parameters $M_{1}$, $M_{min}$, and $\\alpha$ as a function of the galaxy number density. Our values of $M_{1}$ and $M_{min}$ and their increasing trend as the galaxy density decreases show good agreement with SDSS data \\citep[]{zeh05,zhe07} and DEEP2 data \\citep{zhe07}, and the depletion of the slope $\\alpha$ at $z = 0$ may be attributed to insufficient subhalos found in massive clusters.  (6) We have analyzed the kinematic pair fraction of our galaxy samples by selecting the mass in between $2.28 \\times10^{11} \\Msun$ and $3.61\\times10^{12} \\Msun$ which are obtained by applying the observed mass-to-light ratio to convert the luminosity range $-21.39$ $\\leq M_{B}^{e} \\leq$ $-19.39$ at $z = 0.3$. Parameterizing the evolution of the pair fraction as $(1 + z)^{m}$, \\cite{lin08} find that when $r_{max} = 50$ $\\kpc$, $m = 0.41\\pm0.14$ for the full sample. We obtain $m = 1.17$ in $\\Lambda CDM_{100a}$, $m = 1.07$ in $\\Lambda CDM_{200}$, $m = 0.96$ in $\\Lambda CDM_{100b}$, $m = 1.43$ in $\\Lambda CDM_{100c}$, and $m = 1.37$ in $^{*}\\Lambda CDM_{100c}$. When $r_{max} = 100 \\kpc$, in contrast to $0.29\\pm0.05$ found by \\cite{lin08}, we find $m = 0.65$ in $\\Lambda CDM_{100a}$, $m = 0.95$ in $\\Lambda CDM_{200}$, $m = 0.48$ in $\\Lambda CDM_{100b}$, $m = 0.98$ in $\\Lambda CDM_{100c}$ and $m = 0.92$ in $^{*}\\Lambda CDM_{100c}$. The steeper slope basically results from insufficient close pairs found at low redshift, which is interpreted as the deficient low-$z$ galaxy density in our samples due to the failure of the assumption on the redshift-independent mass-cut. However, the situation is less severe for the $r_{max} = 100 \\kpc$ case where milder evolution trend is more evident.   Most of the above tests for our galaxy samples provide supporting evidence for our galaxy model.  However, this model has some problems of its own. First, this model does not provide information of the galaxy morphology. Second, unlike the Millennium simulation \\citep{spr05}, this model does not contain any empirical ingredient to trace individual galaxy evolution.   The solution to the first problem requires the gas component to be included in the simulation, which can reveal the amount of residual gas angular momentum in the subhalo thereby providing additional information of morphological types. The complexity of the second problem is intractable from the first-principle calculations and it needs an empirical model, much as what was added to the Millennium, to approximate the physics involved.    %"
    },
    "0809/0809.1384_arXiv.txt": {
        "abstract": "The very inner structure of massive YSOs is difficult to trace. With conventional observational methods we identify structures still several hundreds of AU in size. However, the (proto-)stellar growth takes place at the innermost regions ($<$100 AU) where the actual mass transfer onto the forming high-mass star occurs. We present results from our programme toward massive YSOs at the VLTI, utilising the two-element interferometer MIDI. To date, we observed 10 well-known massive YSOs down to scales of 20 mas (typically corresponding to 20 - 40 AU for our targets) in the 8-13 micron region. We clearly resolve these objects which results in low visibilities and sizes in the order of 30-50 mas. For two objects, we show results of our modelling. We demonstrate that the MIDI data can reveal decisive structure information for massive YSOs. They are often pivotal in order to resolve ambiguities still immanent in model parameters derived from sole SED fitting. ",
        "introduction": "High-mass stars predominantly form in clustered environments much more distant than typical well-investigated low-mass star-forming regions. Thus, high spatial resolution is a prerequisite for making progress in the observational study of high-mass star formation. Furthermore, all pre-main sequence phases are usually deeply embedded.  This often forces observers of embedded massive young stellar objects (MYSOs) to move to the mid-infrared (MIR) where the resolution of conventional imaging is limited to $>$ 0\\farcs25 even with 8-m class telescopes. Hence, one traces linear scales still several hundred AU in size even for the nearest MYSOs, and conclusions on the geometry of the innermost circumstellar material remain ambiguous. MIR emission moderately resolved with single telescopes may even arise from the inner outflow cones \\cite{2006ApJ...642L..57D,2005A&A...429..903L}. \\\\ A versatile method to overcome the diffraction limit of single telescopes is to employ interferometric techniques. We are conducting a larger survey toward MYSOs based on MIR interferometry. In total, we have observed 10 sources so far.  All these sources, mostly comprising BN-type objects \\cite{1990FCPh...14..321H}, are clearly resolved with the interferometer baselines we applied ($\\ge 16$ m). This in itself is a major step forward compared to the previous more or less unresolved thermal infrared imaging with 4- to 8-m class telescopes for these sources. ",
        "conclusions": "% We have observed massive young stellar objects with the MIR interferometer MIDI at the VLTI. We find substructures with MIR sizes around 30\\,--\\,50~mas. Using the measured visibilities as discriminator, we can exclude purely spherically symmetric matter distributions with a hot central star for the object M8E-IR. The most probably configuration consists of a compact circumstellar disk surrounded by a larger envelope plus a cool central object. The total disk size is not well constrained by the models, and the data allow also for an envelope-only configuration. The finding of a cool central object inside this MYSO is consistent with the idea that the 10--15 M$_\\odot$ central star has been strongly bloated by recent strong accretion events. \\\\ For M\\,17 IRS1, we can exclude the nearly edge-on disk models which were among the models suggested by the SED fitting. Instead, a moderately inclined circumstellar disk might be closer to the truth, and the results of a disk-dominated object with small envelope contribution support an earlier suggestion that M\\,17 IRS1 is near to the Herbig Be star phase.\\\\ Our results show that IR interferometry is a viable tool to reveal decisive structure information on embedded MYSOs and to resolve ambiguities arising from fitting the spectral energy distribution. With the inclusion of the Auxiliary Telescopes, the VLTI is currently becoming even more flexible to tackle such tasks. Finally, the 2$^{\\rm nd}$-generation VLTI instrument {\\sc Matisse} \\cite{2006SPIE.6268E..31L} for the thermal IR will reveal even more complex details of MYSOs by adding closure phase and imaging capabilities to MIR interferometry."
    },
    "0809/0809.3381_arXiv.txt": {
        "abstract": "name}{} \\begin {abstract} \\bf Abstract \\rm  -- An analysis of the residual-velocity field of OB associations within 3 kpc of the Sun has revealed periodic variations in the radial  residual velocities along the Galactic radius vector with a typical scale length of $\\lambda=2.0\\pm0.2$ kpc and a mean amplitude of $f_R=7\\pm1$ km s$^{-1}$. The fact that the radial residual velocities of almost all OB-associations in rich stellar-gas complexes are directed toward the Galactic center suggests that the solar neighborhood under consideration is within the corotation radius. The azimuthal-velocity field  exhibits a distinct periodic pattern in the $0<l<180^\\circ$ region, where the mean azimuthal-velocity amplitude is $f_\\theta=6\\pm2$ km s$^{-1}$. There is no periodic pattern of the azimuthal-velocity field in the $180 <l<360^\\circ$ region. The locations of the Cygnus arm, as well as the Perseus arm, inferred from an analysis of the radial- and azimuthal-velocity fields coincide. The periodic patterns of the residual-velocity fields of  Cepheids and OB associations share many common features. \\it Key words: \\rm star clusters and associations, stellar dynamics, kinematics; Galaxy (Milky Way), spiral pattern. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.1503_arXiv.txt": {
        "abstract": "We reconstruct the interaction rate between the dark matter and the holographic dark energy with the parameterized equation of states and the future event horizon as the infrared cut-off length. It is shown that the observational constraints from the 192 SNIa and BAO measurement permit the negative interaction in the wide region. Moreover, the usual phenomenological descriptions can not describe the reconstructed interaction well for many cases. The other possible interaction is also discussed. ",
        "introduction": "In the modern cosmology, a `dark energy' (DE) with negative pressure is suggested to be responsible for the current acceleration of the universe. The simplest candidate of DE is the cosmological constant, which does nicely well at the pragmatic observational level, but entails the serious theoretical difficultly: the cosmological constant problem and the coincidence problem. Explanations of DE have been sought within a wide range of physical phenomena, including some exotic fields, modified gravity theories, and so on -see \\cite{Padmanabhan} and references therein. Among the most recent generic proposals, the model inspired by the holographic idea \\cite{Hooft}, that the quantum zero-point energy of a system cannot exceed the mass of a black hole with same size, has been put forward to explain the DE \\cite{Cohen,Hsu,Li}. This DE density can be determined in terms of the horizon radius of the universe, corresponding to relate the UV cutoff of a system to its IR cutoff in the quantum field theory. There are usually three choices for the horizon radius supposed to provide the IR cutoff, with different degrees of success, namely the Hubble horizon, the particle horizon, and the future event horizon. The event horizon may be better, since in this case the DE can drive the present accelerated expansion and the coincidence problem can be resolved by assuming an appropriate number of e-folding of inflation \\cite{Li,Li2,Lee}.  Most discussions on DE models rely on the fact that both dark matter (DM) and DE only couple gravitationally. However, given their unknown nature and the symmetry that would impose a vanishing interaction is still to be discovered, an entirely independent behavior between dark sectors is very special. Moreover, since DE must be accreted by massive compact objects like black holes and neutron stars, in a cosmological context the energy transfer from DE to DM may be small but must be non-vanishing. The interaction hypothesis was first introduced by Wetterich \\cite{Wetterich} to discuss the cosmological constant problem in the light of dilatation symmetry and its anomaly. Then cosmological consequences of a scalar field coupled to the matter were studied in \\cite{Wetterich1}. It was found that the coupling quintessence models may give the scaling attractors providing an accelerated expansion at the present time and alleviate the coincidence problem \\cite% {Amendola2}. The interaction also appears in the context of modified gravity models \\cite{Pavon3}. More possibility that DE and DM can interact has been studied in \\cite{Anderson,Mangano,Farrar,Elizalde,Setare,Chen1}. Confronted to cosmological data, it was found that an appropriate interaction can influence the perturbation dynamics, and the lowest multipoles of the CMB spectrum \\cite{Zimdahl,Wang1}, and could be inferred from the expansion history of the Universe, as manifested in the supernova data together with CMB and large-scale structure \\cite{Guo,Feng,Amendola3}. In addition, it was suggested that the dynamical equilibrium of collapsed structures would be affected by the coupling of DE to DM \\cite{Kesden,Bertolami}. The interaction was first connected to holography by Horvat \\cite{Horvat} who argued that scaling of the cosmological constant stemming from the zero-point energy in quantum field theory possibly implies a non-vanishing coupling of the cosmological constant with DM. In the holographic DE model with the Hubble horizon as the IR cutoff, the interaction can be available to derive the present accelerated expansion and alleviate the coincidence problem \\cite{Pavon2}. In the interacting model with the event horizon as the IR cutoff, it was shown that the equation of state (EoS) of DE can accommodate the dynamically evolving behavior of crossing the phantom divide \\cite{Wang2}, which suggested by recent most observational probes \\cite% {WangY}.  Although the interaction is important in studying the physics of DE, it will not be possible to derive the precise form of the interaction from first principles unless the nature of both dark sectors was known. Usually, the coupling is determined from phenomenological requirements \\cite% {Amendola2,Pavon2}. In view of the continuous equations of DE density $\\rho _{d}$ and DM density $\\rho _{m}$, the coupling $Q$ must be a function of densities multiplied by a quantity with units of inverse of time, which has an obvious choice as Hubble time $H^{-1}$. Thus, one may write the coupling as% \\begin{equation} Q=Q(H\\rho _{m},H\\rho _{d}),  \\label{Q} \\end{equation}% which leads $Q\\simeq \\lambda _{m}H\\rho _{m}+\\lambda _{d}H\\rho _{d}$ from the first order terms in the power law expansion. Assuming that the ratio $% r=\\rho _{m}/\\rho _{d}$ might be piecewise constant, the linear parameters are usually set to $\\lambda _{m}\\simeq \\lambda _{d}$ and even $\\lambda _{m}\\simeq 0$ or $\\lambda _{d}\\simeq 0$ for simplicity. Considering the couplings are terms in the Lagrangian which mix both DE and DM, one may further suppose that they could be parameterized by some product of the densities of DE and DM, such as the simplest $Q\\simeq \\lambda \\rho _{m}\\rho _{d}$ \\cite{Mangano}. Besides these phenomenological descriptions, various proposals at the fundamental level have been tried to account for the coupling, including the dependence of the matter field on the scalar field \\cite{Piazza} or expressing the cosmological constant as a function of the trace of the energy-momentum tensor \\cite{Poplawski}. Recently, an interesting thermodynamical description of interaction between holographic DE and DM has been proposed in \\cite{Wang3}, where it was assumed that in the absence of the coupling the DE and DM remain in separate thermal equilibrium, then a small interaction can be viewed as a stable thermal fluctuation that brings a logarithmic correction to the equilibrium entropy of DE and DM. Other specific coupling which was assumed from the outset can be found in \\cite{Anderson,Amendola1,Das}.  The main aim of this work is to reconstruct the coupling using the recent DE probes (the Baryon Acoustic Oscillation (BAO) measurement at $z=0.35$ from the Sloan Digital Sky Survey \\cite{Eisenstein} and the re-compiled 192 Type Ia Supernovae (SnIa) samples \\cite{Davis}, consisting 60 points from ESSENCE (\\textquotedblleft Equation of State: Supernovae trace Cosmic Expansion\\textquotedblright ) supernova survey \\cite{Wood}, 57 points from Supernovae Legacy Survey \\cite{Astier}, 45 points nearby Supernovae \\cite% {Riess}, and 30 points detected by the Hubble Space Telescope \\cite{Riess1}% ). We will focus on the holographic DE model and choose the future event horizon as the IR cutoff.  The model with the Hubble horizon has been studied recently in \\cite{Sen} and it was found that the reconstructed interaction is always positive in 1$% \\sigma $ region. This seems to corroborate the recent argument that the negative interaction violates the second law of thermodynamics which inquires the energy transfer from DE to DM rather than otherwise \\cite% {Pavon1}. However, it should be noticed that there are some problems in the thermodynamics of DE, such as the negative entropy \\cite{Odintsov}, and the generalized thermodynamical second law indeed breaks down in the universe with phantom-dominated DE \\cite{Izquierdo,Mohseni}. These results suggest that one should consider the thermodynamical properties of DE with wide possibilities. Hence, it is interesting to see whether the positive interaction is a robust result for other models, such as the present model with the IR cutoff as the future event horizon.  Considering the time varying DE gives a better fit than a cosmological constant and in particular most of the observational probes indeed mildly favor dynamical DE crossing the phantom divide at $z\\sim 0.2$ \\cite{WangY}, we will employ two commonly used parameterizations \\cite% {Jassal,Feng,Chevallier,Linder}, namely% \\begin{equation} w_{A}=w_{0}+w_{1}(1-a)=w_{0}+w_{1}\\frac{z}{1+z},  \\label{w1} \\end{equation}% which has been used in \\cite{Sen}, and% \\begin{equation} w_{B}=w_{0}+w_{1}(1-a)a=w_{0}+w_{1}\\frac{z}{\\left( 1+z\\right) ^{2}}. \\label{w2} \\end{equation}% It should be noticed that the different parameterizations are beneficial to control some amount of parameterization dependence. After reconstructing the interaction, we will further compare it with the usual phenomenological models and the recent thermodynamical description. ",
        "conclusions": "We have reconstructed the interaction term between the holographic DE and DM, using the re-compiled 192 SnIa samples combined with the recent BAO measurement. The DE parameter $c$ is assumed near one. The two common used parameterizations of $w_{d}$ are considered in 1$\\sigma $ region. It is found that the present accelerated expansion of universe is achieved and the phantom behavior of DE is permitted. We illustrate that the negative interaction is permitted in the wide region. Hence, contrast to the DE model studied in \\cite{Sen}, where the reconstructed interaction is always positive in 1$\\sigma $ region, we can not obtain the favor from the DM observation for the recent argued thermodynamical second law which imposes the energy transfer from DE to DM. This suggest us to keep wide possibilities of the thermodynamical properties of DE.  We show that the four usual phenomenological descriptions can not describe the reconstructed interaction well for many cases. Specially for the case with small DM parameter $\\Omega _{0m}$ $=0.25$, the interaction rate has a trend to change sign at recent epoch. This is at variance with the usual phenomenological interacting terms which have the definitive sign. We further illustrate that the recent thermodynamical description of the interaction is not favored. Our work stimulates one to seek a more suitable interaction.  It would be interesting to confront our model to more observations, such as CMB angular power and large scale structure, which will further constrain the interaction rate and make clear if the negative interaction and phantom DE are still to be permitted. Another interesting work is to consider the perturbation evolution in our model. Recently, it has been found that some types of interaction will lead to instability under the curvature perturbation \\cite{Valiviita,hjh}. It may provide some restrictions on the form of interaction. We will work on these directions in the future."
    },
    "0809/0809.1029_arXiv.txt": {
        "abstract": "{ Numerical simulations of galaxy mergers are a powerful tool to study these fundamental events in the hierarchical built-up of galaxies. Recent progress have been made owing to improved modeling, increased resolution and large statistical samples. We present here the highest-resolution models of mergers performed so far. The formation of a variety of substructures ranging from kinematically decoupled cores to globular-like clusters is directly resolved. In a resolution study, we show that the large-scale structure of elliptical-like merger remnants can be affected by the resolution, and a too modest resolution may affect the numerical predictions on the properties of major merger remnants: understanding precisely which kind of event or succession of events has formed the various types of elliptical galaxies remains an open challenge. } ",
        "introduction": "For more than two decades, numerical simulations have been a powerful tool to study the dynamics of interacting galaxies and the properties of galaxy merger remnants. Some of the main results obtained from numerical models range from the formation of elliptical-like galaxies with realistic $r^{1/4}$ density profiles in major mergers of massive spirals (Barnes 1992, Naab \\& Burkert 2003, Bournaud et al. 2005), the effect of minor mergers of spiral galaxies with small companions (e.g., Walker et al. 1996), and the triggering of star-formation by galaxy interactions and mergers (Mihos \\& Hernquist 1994, Cox et al. 2006, di~Matteo et al. 2007).  Recent progress have been made by studying galaxy interactions in a wider range of physical conditions including more complex (and more realistic) cases like:  the (re-)merging of early-type galaxies that are themselves remnants of past major mergers (Naab et al. 2006), sequences of repeated minor mergers (Bournaud et al. 2007a), or merging galaxies embedded in a larger-scale group or cluster environment, showing a significant influence of the ram-pressure and large-scale gravity field on interacting galaxies (Kapferer et al. 2008, Martig \\& Bournaud 2008).  We present here recent progress in terms of numerical resolution, in simulations of ``standard'' mergers of two massive spiral galaxies merging together, with no multiple mergers and no large-scale environment accounted for (see Section~2). While the remnants of major mergers in numerical simulations do globally resemble elliptical galaxies, whether the major merger scenario can really explain the exact properties of most real elliptical galaxies or not is still an open question (see Burkert et al. 2007). This raises the question of whether or not most simulations of galaxy mergers have a high-enough resolution to predict accurately the detailed properties of merger remnants: this is studied in Section~3, where we show that a too modest resolution may significantly affect the properties of numerical merger remnants. Our recent very-high resolution simulations also enable to study the smallest structures formed in galaxy mergers, like decoupled cores, tidal dwarfs, and super star clusters: results on these aspects will be presented in Section~4. ",
        "conclusions": "High spatial and mass resolution enables the formation of small objects to be resolved in numerical models of galaxy mergers. In recent simulations, even the internal properties of small objects can be resolved -- like the rotation curves of tidal dwarfs (see Bournaud et al. 2007b). The results of such high-resolution models is overall in agreement with theoretical expectations, in particular regarding the formation of a population of GCs during major wet mergers.  High resolution is also a must to study accurately the large-scale structure of spheroidal galaxies formed in major mergers. Our recent resolution study shows that even simulations with a total of one million particles and a spatial resolution of a hundred of parsec can still be significantly affected by the resolution. Whether major mergers can explain the detailed properties of elliptical galaxies (boxyness, angular momentum distribution, anisotropy, etc..) is still largely unsolved, and understanding which type of merger would have formed the various kind of ellipticals remains an important challenge."
    },
    "0809/0809.3992_arXiv.txt": {
        "abstract": "We present X-ray observations of the field containing Nova Puppis 1942 (CP Pup) and  Nova Puppis 1991 (V351 Pup), done with {\\sl ASCA} in 1998, and with {\\sl XMM-Newton}  in 2005. The X-ray and UV luminosity of CP Pup seem to have remained approximately constant since  the last X-ray observations of the 1980'ies, while the optical luminosity has decreased. The X-ray properties of this nova are explained by a high mass white dwarf accreting at low rate, in agreement with the nova theory given the large amplitude and other characteristics of the 1942 outburst. Assuming a distance of 1600 pc, the X-ray luminosity of CP Pup is L$_{\\rm x}=$2.2 $\\times 10^{33}$ erg s$^{-1}$ in the 0.15-10 keV range covered with EPIC, compatible with a magnetic system. The RGS grating spectrum shows a few prominent emission lines, and it is fitted with  a cooling flow with mass accretion rate \\mdot$< 1.6 \\times 10^{-10}$ M$_\\odot$ yr$^{-1}$. We detected also the O VII complex at 21.6-21.8 \\AA \\ that does not arise in the cooling flow. Most likely this feature originates in a wind or in the nova shell. The RGS and EPIC  spectra are fitted only with thermal models with a very high shock temperature, T$>$60 keV, indicating  a white dwarf with M$>$1.1 M$_\\odot$. The X-ray flux is modulated with the spectroscopic period of 1.47 hours detected in the optical. Since CP Pup is not an eclipsing system, this is better understood if magnetic accretion occurs: we discuss this possibility and its implications in detail. V351 Pup (N Pup 1991) was detected  with {\\sl XMM-Newton}, but not with {\\sl ASCA}. It is a faint, non-super-soft X-ray source with luminosity L$_{\\rm x} \\simeq 3 \\times 10^{31}$ erg s$^{-1}$, a factor of 50 less than measured with ROSAT in 1993. ",
        "introduction": "In this article we describe {\\sl ASCA} and {\\sl XMM-Newton} observations of two novae in the same field in the Puppis region, Nova Puppis 1942 (CP Pup) and Nova Puppis 1991 (V351 Pup). Classical novae are Cataclysmic Variables (CV), that is close, interacting binaries containing a white dwarf (WD) that accretes matter from a companion, which is usually still on the main sequence or slightly evolved. In novae, at the bottom of the accreted envelope hydrogen burning is ignited in a thin shell, first the p-p and later the CNO cycle. The process becomes explosive because of the degeneracy conditions of the material, compressed on the small surface of the WD.  A thermonuclear flash follows, and an optically thick wind that ejects most or all the accreted material (see Starrfield et al. 2000 for a review). The more massive the WD is, the smaller its surface,  and the sooner the pressure sufficient to initiate the thermonuclear runaway is reached.  The nova theory makes very definite predictions on the correlation of the binary system parameters and the properties of the outburst. The main parameters that determine its properties and recurrence time are the mass accretion rate \\mdot \\ and the white dwarf mass m(WD) (note that all novae are recurrent phenomena even if only very few, the ``recurrent novae'' are repeated over a  human lifetime). The abundances and the thermal state of the WD at the onset of accretion are also important. In a series of papers (especially Kovetz \\& Prialnik 1994, Prialnik \\& Kovetz 1995, Yaron et al. 2004, Epelstein et al. 2007) clear correlations between these parameters and the outburst characteristics are found. Thus, having a remarkable and detectable manifestation in their outburst, novae allow us to test the evolutionary models that apply to all CV and to related systems, included the  ``single degenerate'' progenitors of type Ia supernovae.  X-ray observations are a powerful tool to study the following phenomena in novae: a) Violent, energetic mechanisms in the shell shortly after the outburst (e.g. Nelson et al. 2008), or even a long time later  (Balman 2006); b) The WD atmospheric temperature, abundances and effective gravity, after the ejecta clear up sufficiently to allow detection of the X-ray supersoft, luminous central source (e.g. Ness et al. 2003, Nelson et al. 2008); c) The accretion process in quiescence. The shocked material in the disk, or in a magnetically funneled accretion stream, emits X-rays (e.g. Mukai \\& Orio 2005). In order to study accretion, we proposed X-ray observations of the two objects we describe in this paper, as a test of the nova theory which is largely based on the modalities of accretion.  \\subsection{CP Pup: a puzzling nova} CP Puppis appeared as a truly astonishing object in 1942, leaving the astronomers of that  time wondering for a while whether it was a Galactic supernova. The outburst is described in the Payne-Gaposhkin's book ``Classical Novae''. In 1942 it rose from fainter than 17th magnitude to  V=-0.2, raising to maximum during at least 3 days after the discovery. The time for a decay by 3 magnitudes (t$_3$) was one of the shortest ever measured for a nova, only 6.5 days. The velocities measured from the full width at half maximum of the initial absorption lines, and of the emission lines that appeared soon thereafter, reached at most 1200 km s$^{-1}$, which is lower than measured for other luminous novae. The ejected shell was resolved for the first time 14 years after the outburst (see Williams 1982 and references therein), indicating a distance of about 1.6 kpc and a maximum absolute luminosity M$_{\\rm V}$=-10.5, a factor of $\\simeq$400 above the Eddington luminosity of a 1 M$_\\odot$ star (L$_{\\rm Edd}$).  Downes \\& D\\\"urbeck (2000) discussed the possible errors in determining the  distance of a shell that is not expanding uniformly, but is made of blobs with different velocity. These authors placed CP Pup at a lower distance of 1140 pc, and Cohen \\& Rosenthal (1983) even at only 850 pc. A distance compatible with a peak luminosity below 100 L$_{\\rm Edd}$ is in any case very unlikely. Another unusual characteristic of CP Pup were the Fe [II] lines that still appeared in the   spectrum  simultaneously with high excitation lines, including even coronal lines. The energetics initially reminded of a supernova. If we were witnessing such an event today, probably we would initially suspect an accretion induced collapse of a white dwarf, or merging compact objects. Despite the peculiarities outlined above, the optical spectrum was typical for a classical nova, and as such CP Pup is classified.  The distance obtained from the  maximum magnitude versus rate of decline relationship (MMRD) is  much larger than the range of values derived from the nebular parallax. For this and several other reasons CP Pup has long been suspected to host a  magnetic WD. Because of the apparent disk-like structure, it may be an ``intermediate polar'' (hereafter IP) rather than a ``polar'' (e.g. a diskless system, see discussion in Balman et al. 1995 and references therein). It is doubtful whether the MMRD holds for magnetic novae (see for instance the description of DQ Her by Payne Gaposhkin, 1964). Orio et al. (1992) suggested that magnetic rotator effects may cause the ejection of the accreted envelope more quickly and efficiently than in other novae. In fact the outburst amplitude of V1500 Cyg, a polar, was also very large, about 16 mag, and  the outburst of the IP GK Per  had an amplitude of $\\simeq$14 mag despite the high luminosity at quiescence (due to an unusual subgiant secondary). The  unambiguous classification of also CP Pup as a magnetic system would be further evidence of the role of the magnetic field in nova outbursts.  Several unusual characteristics of the quiescent CP Pup point in the direction of an IP nature. The outburst of 1942 can best be explained with a high mass white dwarf (e.g., Prialnik and Kovetz 1995, see discussion below).  However, assuming that the radial velocities variations of the emission lines in the quiescent optical spectrum indicate the orbital motion of the compact object, and that the width of the lines is due to the Keplerian velocity of the disk,  an embarrassingly low mass of the white dwarf is obtained, M$<$0.2 M$_\\odot$. The data are rather complex, and we will summarise them as follows. At quiescence, two photometric periods have been measured, 0.061-0.062 days, and 0.06834 days (Bianchini 1985a and b, Warner 1985, O'Donoghue et al. 1985, White et al. 1993). From the radial velocity of the He II line at 4686 \\AA \\ and the Balmer lines, White et al. (1993) measured a spectroscopic period of 0.06129 days.  In the most recent work, Bianchini et al. (2008, in preparation, and hereafter B08) detect a mean spectroscopic period of 0.0612643 days (1.47 hours), coincident within the errors with the period reported by White et al. (1993), and other aliases of this period. In 1987 a photometric period of CP Pup was measured to be 11\\% longer than the spectroscopic period (White et al. 1993). However in 1993 Patterson \\& Warner (1998) found instead two other photometric periods, of which the first was equal to the spectroscopic one and the second was unstable, but generally  longer than the spectroscopic period by only $\\simeq$2\\%. The last authors interpret this as a ``superhump'' (due to disk precession), which is often detected in non-magnetic cataclysmic variables (CV). However, another possible interpretation is that the longer photometric period is due to the rotation of the WD that has become asynchronous after the nova outburst, like the polar Nova V1500 Cyg after the 1975 outburst (Diaz \\& Steiner 1991). Balman et al. (1995) reported an  X-ray flux modulation with the spectroscopic period. This is suggestive of a magnetic system, since X-ray orbital modulations are not observed in non-magnetic CV with inclination less than 60$^{\\rm o}$ (Baskill et al. 2005). The evidence however was not conclusive, because the orbital period of CP Pup is uncomfortably close to the rotation period of 94 minutes of the {\\sl ROSAT} spacecraft.  The spectroscopic period of 1.47 hours is interpreted as the orbital period in the literature, which we also adopt as our baseline interpretation (we will review this assumption later, and present an alternative interpretation).  In this baseline scenario, it is one of the two shortest orbital periods of all novae. Most of the known orbital periods of classical novae are longer than 3 hours, i.e. above the period gap. CP Pup is only matched by GQ Mus (Nova Mus 1983) with 1.42 hours, and there are only 5 other novae with orbital periods below 3 hours (Diaz \\& Bruch 1997).  The lack of eclipses places a constraint on the inclination, $i\\leq$65$^{\\rm o}$. With the measured radial velocities and line widths, there simply is no way to place a massive white dwarf in a non-eclipsing system with such a short orbital period. O'Donoghue et al. (1989) conclude that the WD mass should be M$_{\\rm WD}\\leq$0.2 M$_\\odot$. Moreover, if the modulation of the infrared flux with a period close to the spectroscopic one is due to an ellipsoidal variation (Szkody \\& Feinswog 1988), the upper limit on the inclination is {\\it i}$\\leq$35$^{\\rm o}$. This implies that a Roche lobe filling secondary secondary has mass $\\leq$0.2 M$_\\odot$, and that the upper limit for the WD is less than a tenth of a solar mass! This is  a clear inconsistency: in the first place, this WD mass is much lower than the minimum WD mass in the Galaxy and it is difficult to explain such an amount of mass loss from the WD. Moreover, H-burning would not even start on the surface of such a low mass  WD, thus, a classical nova eruption could not occur. This puzzle may originate, first of all, in an incorrect interpretation of the IR modulation. At a distance of $\\sim$1 kpc, a 0.2 M$_\\odot$ secondary star (M5V, M$_{\\rm V}$=14.7, M$_{\\rm K}$=8.5) would have a distance modulus m-M=10, yielding a K$\\simeq$18.5 magnitude of the secondary. Szkody and Feinswog measured K=13.4 for CP Pup, a value completely inconsistent with the assumed short orbital period secondary and one which cannot be reconciled as ellipsoidal variations. Only if CP Pup was 95 pc away, the assumed 0.2 M$_\\odot$ secondary star could then account for the observed IR modulations.  A different physical mechanism that may cause the IR light curve modulation is beamed cyclotron radiation. Although up to now the cyclotron humps in IR have only been observed for the higher magnetic field systems, the polars, this phenomenon is not yet ruled out in  IP's. The cyclotron hump does produce an IR modulation like the one observed in CP Pup. One example is  AR UMa (Howell et al. 2001), where the visual and infrared double-humped light curves are caused by beamed cyclotron radiation, although initially at least the visual modulations were initially ascribed to ellipsoidal variations. Another example is HU Aqr, another polar, where the IR light curve is explained partially by cyclotron emission and the   light curve is very complex to model (Howell et al. 2002).  The secondary star spectrum is not detected in optical or infrared, therefore the secondary is  thought to be a main sequence star, because subgiants or giants have clear spectral signatures at distances of order of 1 kpc. If the measured spectroscopic period is orbital and CP Pup is not a triple or multiple system, the most likely reason for the absurd result of the disk accretor model is simply that the accretion is funneled by a magnetic field and does not occur through the disk. The saddled line profiles show that a disk exists, however if the innermost part of the accretion disk is disrupted by the magnetic field of an IP, the radial velocity is not measured close to the WD, but at the Alfv\\`en radius.  Finally, we do not rule out that the radial velocities measurements are inaccurate because the emission lines are contaminated by multiple components. The optical spectrum does show very complex velocity profiles, with multiple components in the emission lines (B08); some of these components may be produced in a hot spot on the WD or in  an accretion stream rather than in the disk. B08 describe these possibilities in detail, but they conclude that assuming that CP Puppis is an IP and that the 1.47 hours period is orbital does not completely resolve the issue. Analysing new spectroscopic and photometric data taken over several years, B08 find that the structure of the emission lines is very complex and the velocity of the inner disk was not accurately determined, alleviating somewhat the ``mass function problem'', but at the same time making the study of the system extremely difficult.  When we  proposed observing CP Pup with {\\sl XMM-Newton}, one of our aims was to resolve the mass problem, possibly revealing the spin period of a magnetic WD. One of the characteristics signatures of IP systems is in fact that the X-ray flux is usually modulated with the rotation period of the WD, which does not rotate synchronously with the orbital period like in ``polar'' systems. Typical spin periods of IP are of the order of tens of minutes. A new X-ray observation of CP Pup also had a broader scope than learning the details of a single system. Because of the selection effect due to the outburst amplitude, classical novae are observed at larger distances than other CV, and they are rather faint at quiescence. Up to now, only one classical nova, V603 Aql, has yielded a grating spectrum in X-rays (Mukai \\& Orio 2005). X-ray grating spectra are a fundamental tool in order to study how accretion occurs. Therefore, another important aim we had was to broaden the statistics, obtaining another term of comparison to learn how accretion occurs in novae.  CP Pup is the third brightest classical nova in X-rays at quiescence, after V603 Aql and GK Per.  \\subsection{V351 Pup: a cooling nebula?}  V351 Pup was discovered in outburst in 1991 December 27  (Camilleri 1992). It was a moderately fast classical nova, with t$_2$ and t$_3$ of 16 and 40 days, respectively and ejection velocity $\\simeq$2000 km s$^{-1}$ on average (Shore et al. 1992, Sonneborn et al. 1992). It was also a neon nova (Pachoulakis \\& Saizar 1995).  An X-ray flux in the 0.2-2.4 keV range F$_{\\rm x}=3 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ was measured with {\\sl ROSAT} 16 months after the outburst (Orio et al. 1996). The spectrum was quite hard in the {\\sl ROSAT} range and Orio et al. (1996) suggested that the X-ray emission could originate either in the cooling nova shell, or be due to  renewed accretion. In Orio et al. (2001) we already announced and discussed that the nova was not detected with {\\sl ASCA}, with a strict upper limit of 10$^{-13}$ erg cm$^{2}$ s$^{-1}$ in the 0.7-10 keV energy range. We concluded in that paper that the first interpretation is definitely the most likely, because the shell cools and disappears as X-ray source, while accretion is expected to continue. We present here a new detection of V351Pup with XMM-Newton, but the data are far poorer in quality than the CP Pup ones, so we will discuss this object in less detail. ",
        "conclusions": "CP Puppis is only the second classical nova for which an X-ray grating spectrum could be obtained at quiescence. Even if with low S/N, this spectrum shows strong emission lines which probably mainly arise in a cooling flow, either in the thermal plasma of disk or of accretion columns of a magnetic WD. There is at least one emission complex (O VII at 21.6-21.8 \\AA), that most likely arises instead in a wind from the system, or in a circumstellar nebula. CP Pup had a very large outburst amplitude, a short t$_3$, and only moderately high ejection velocity. Examining a number of theoretical papers quoted above, we find that these characteristics are explained by the upper limit to the mass accretion rate \\mdot$< 1.6 \\times 10^{-10}$ M$_\\odot$ year$^{-1}$, and by the lower limit for the WD mass, M(WD)$>$1.1 M$_\\odot$, obtained by fitting the X-ray spectra.  The X-ray observation of this nova at quiescence therefore confirm the nova theory, offering a rare possibility to probe it in detail. The models of WD thermonuclear burning and runaways are also at the base of type Ia SN studies, and testing them is important for all aspects of modern astrophysics, including cosmology.  The puzzle of the WD mass of CP Puppis is not solved yet. With our observations, we could not demonstrate univocally that the WD is magnetic and the disk is disrupted far from the WD, however   the high L$_{\\rm x}$ is typical for an IP. The only point we found against the magnetic nature is relevant only if the detected X-ray modulation is orbital: the X-ray orbital modulation of CP Pup is larger at intermediate energy (0.8-2 keV), although in known  IP's the amplitude of orbital modulations generally decreases with energy (e.g. Hellier et al. 1993, Parker et al. 2006). On the other hand, the mere existence of an X-ray flux modulation with the optical-spectroscopic period argues in favor of an IP, no matter how it is interpreted. If it is orbital, it would not be detected in  a non-eclipsing disk accretor. If it is rotational, detection of the WD spin is indeed expected for an IP.  The high WD mass is also proof that the non-magnetic, disk accretor model does not work for CP Pup. It cannot be ruled out that all the observed periodicities at different wavelengths, including our X-ray modulation, are related to the WD rotation and the true orbital period has never been detected. This suggestion, for the time being, is only a working hypothesis that will need to be supported by further observations, especially by long optical photometric campaigns.  Finally, we also detected V351 Pup 14 years after the outburst, at least 50 times less luminous than 12 years earlier. This seems to indicate that the much larger hard X-ray luminosity at that time was due to the ejected shell, which had already cooled in 2005. Nova shells are known to be conspicuous X-ray sources after the outburst  (see for instance the shell emission of  RS Oph, Nelson et al. 2008) and in most novae they seem to cool very rapidly (Orio et al. 2001). An upper limit for the cooling time of the V351 Pup shell is 6 and a half years, the time elapsed between the {\\sl ROSAT} and {\\sl ASCA} observation."
    },
    "0809/0809.3395_arXiv.txt": {
        "abstract": "We demonstrate that the combination of the ideas of unimodular gravity, scale invariance, and the existence of an exactly massless dilaton leads to the evolution of the universe supported by present observations: inflation in the past, followed by the radiation and matter dominated stages and accelerated expansion at present. All mass scales in this type of theories come from one and the same source. ",
        "introduction": "\\label{sec:intro} The origin of different mass scales in particle physics is a mystery. The masses of quarks, leptons and intermediate vector bosons come from the vacuum expectation value (vev) of the Higgs field; the dimensionful parameters like the QCD scale $\\Lambda_{QCD}$ or the scales related to the running of all other dimensionless couplings of the Standard Model (SM) are believed to have nothing to do with the Higgs vev. Newton's gravitational constant provides yet another mass scale, very different from typical particle masses of the SM. The Higgs mass itself -- where does it come from?  Is it possible that {\\em all} these mass scales originate from one and the same source?\\footnote{We will refer to this possibility as ``no-scale scenario''.} Indeed, it is not difficult to construct, on the {\\em classical level}, a theory containing a new singlet field $\\chi$, which gives masses to all particles and fixes Newton's constant. Having in mind the SM extended by 3 light right-handed singlet fermions, the $\\nu$MSM of \\cite{Asaka:2005an,Asaka:2005pn} (this theory -- Neutrino Minimal SM -- unlike the SM, can explain neutrino masses and oscillations, dark matter and baryon asymmetry of the universe), one can write the Lagrangian realizing this idea in the following form\\footnote{This expression is identical to the one in \\cite{Shaposhnikov:2008pf}, Section 8, but uses different notations.} \\begin{eqnarray} \\nonumber {\\cal L}_{\\nu \\rm{MSM}}= {\\cal L}_{\\rm{SM}[M\\rightarrow 0]} + {\\cal L}_G  + \\frac{1}{2}(\\partial_\\mu\\chi)^2 - V(\\varphi,\\chi)\\\\ +  \\left(\\bar{N_I} i \\gamma^\\mu\\partial_\\mu N_I - h_{\\alpha I} \\,  \\bar L_\\alpha N_I \\tilde{\\varphi} - f_I \\bar {N_I}^c N_I \\chi + \\rm{h.c.}\\right) \\,, \\label{lagr} \\end{eqnarray} where the first term is the SM Lagrangian without the Higgs potential, $N_I$ (I=1,2,3) are the right-handed singlet leptons, $\\varphi$ and $L_\\alpha$ ($\\alpha=e,\\mu,\\tau$) are  the Higgs and lepton doublets respectively, $h_{\\alpha I}$ and $f_I$ are the matrices of Yukawa coupling constants. The scalar potential is given by \\be V(\\varphi,\\chi) = \\lambda\\left(\\varphi^\\dagger\\varphi-\\frac{\\alpha}{2\\lambda} \\chi^2\\right)^2+ \\beta(\\chi^2-\\chi_0^2)^2~, \\label{pot} \\ee and the gravity part is \\be {\\cal L}_G = -\\left(\\xi_\\chi \\chi^2 +2 \\xi_h\\varphi^\\dagger\\varphi\\right)\\frac{R}{2}~, \\label{GG} \\ee where $R$ is the scalar curvature. We will only consider positive values for $\\xi_\\chi$ and $\\xi_h$, for which the coefficient in front of the scalar curvature is positive, whatever values the scalar fields take. This is the Lagrangian of ``induced gravity'' going back to Refs. \\cite{Zee:1978wi,Smolin:1979uz} (see also \\cite{Shaposhnikov:2006xi,Anisimov:2008qs} in the $\\nu$MSM context).  For positive $\\lambda$ and $\\beta$ the theory with potential \\eqref{pot} possesses a ground state\\footnote{By a ground state we mean a constant solution of the equations of motion including gravity. The existence of such a ground state could be essential for a consistent quantization of the theory.}. It corresponds to the fields sitting at the minimum of the potential, i.e. $\\chi=\\chi_0$, $h=h_0$ with ${h_0}^2=\\frac{\\alpha}{\\lambda} \\chi_0^2$ and a constant metric describing flat space-time. The field values at the potential minimum can be related to the Planck scale as ${M_P^2=\\xi_\\chi \\chi_0^2 + \\xi_h {h_0}^2}$, ${M_P= 2.44\\times 10^{18}}$ GeV. Physics in this theory does not depend on a specific value of $\\chi_0$ -- all dimensionful parameters are proportional to it -- and only dimensionless ratios can be measured.  Although the aim of having one source for all mass scales is achieved by construction of the Lagrangian (\\ref{lagr}) (we stress that we are still discussing the classical theory) the solution is not satisfactory: the absence of explicit mass terms for the Higgs field and for singlet fermions, and the absence of a gravity scale along with the introduction of the dimensionful parameter $\\chi_0$, required to realize the scenario, are ad hoc and do not follow from any symmetry principle.  The symmetry that forbids (on the classical level) the appearance of any dimensionful parameters is well known --  it is the dilatational symmetry. Under dilatations, the scalar and fermionic fields change as ${\\phi(x)\\to\\sigma^n\\phi(\\sigma x)}$ ($n=1$ for scalars and ${n=3/2}$ for fermions), while the metric transforms as  ${g_{\\mu\\nu}(x)\\to g_{\\mu\\nu}(\\sigma x)}$. The action (\\ref{lagr}) is invariant under this symmetry, provided $\\chi_0 = 0$, leading to the absence of all dimensionful parameters.  From now on we will require that a dilatation invariant theory should possess a ground state. Since this requirement is essential for our model, we will further discuss it in section \\ref{sec:dreams}. For the dilatation invariant theory to contain massive singlet and doublet fermions, the ground state should be such that $\\chi\\neq 0$ and $h\\neq 0$. The only way to achieve this is to set $\\beta=0$. Thus, the no-scale scenario can only be realized if $\\beta=0$. In this case the potential (\\ref{pot}) acquires the flat direction  ${h^2-\\frac{\\alpha}{\\lambda} \\chi^2=0}$, and the theory contains one exactly massless particle $\\eta$ -- a certain mixture of the singlet $\\chi$ and the Higgs field. The requirement $\\beta=0$ therefore leads to a theory with spontaneously broken scale invariance, where $\\eta$ appears as a Goldstone boson\\footnote{Spontaneous breaking of the dilatational symmetry also occurs if the potential has a flat direction either along $\\chi=0$ or along $h=0$. However, both cases are unsatisfactory. The first one corresponds to a theory with no massive singlet fermions, whereas the second one is a theory with no electroweak symmetry breaking.}. Equivalent arguments were given in \\cite{Buchmuller:1988cj}.  The theory (\\ref{lagr}) with $\\beta=0$ (from now on only this choice of parameters will be considered) is rather peculiar\\footnote{We stress that this theory is {\\em not invariant} under local conformal transformations. Conformal invariance requires the specific values for $\\xi_\\chi$ and $\\xi_h$, $\\xi_\\chi=\\xi_h=-\\frac{1}{6}$.}: not only is the physics independent of the value of $\\chi \\neq 0$ in the ground state, but the ground state is infinitely degenerate.  The question ``Who gives the mass to the dilaton ?'' does not arise. It is massless, and the chain  of questions ``Who gives mass to whom?'' terminates.  Not to any surprise, the classical scale-invariant theory constructed in this way does not contain a cosmological constant $\\Lambda$. So, if we confront it with cosmological observations, it seems to fail, since the universe is in accelerated expansion, which requires the presence of dark energy with an equation of state close to that of the cosmological constant. This conclusion is certainly correct for standard General Relativity (GR), associated with the action \\be S_E = \\int \\sqrt{-g} d^4x {\\cal L}_{\\nu \\rm{MSM}}~, \\label{action} \\ee where $g$ is the determinant of the metric.  The aim of this Letter is to show that the situation is completely different if general relativity in (\\ref{action}) is replaced by Unimodular Gravity (UG). UG is a very modest modification of Einstein's theory: it adds a constraint $g=-1$ to the action principle defined by eq. (\\ref{action}) \\cite{vanderBij:1981ym,Wilczek:1983as,Zee:1983jg, Buchmuller:1988wx,Weinberg:1988cp, Unruh:1988in,Alvarez:2005iy,Henneaux:1989zc}. UG is invariant under diffeomorphisms which conserve the 4-dimensional volume element. It contains the same number of dynamical degrees of freedom (massless graviton) as Einstein's theory. To the best of our knowledge, the consequences of scale-invariant UG with massless dilaton have not been considered previously.  The relevance of UG for the cosmological constant problem was realized long time ago \\cite{Wilczek:1983as,Zee:1983jg,Buchmuller:1988wx, Weinberg:1988cp,vanderBij:1981ym}. If $g=-1$, adding a constant $\\Lambda$ to the Einstein-Hilbert action does not change the equations of motion. Still, the $\\Lambda$ problem is not solved, since the cosmological constant shows up again, but now as an initial condition for cosmological evolution in UG. We will see that for our case of a scale-invariant theory together with UG, the initial conditions lead to a non-trivial run-away effective potential for the dilaton rather than to a cosmological constant, and thus to dynamical dark energy. Moreover, it will turn out that both inflation and accelerated expansion of the universe can be explained on the same footing.  The Letter is organized as follows. In Section \\ref{sec:unimod} we show that scale-invariant UG with a massless dilaton is equivalent (on the classical level) to Einstein's theory with zero cosmological constant and a peculiar potential, the magnitude of which is fixed by initial conditions for all the fields. We continue in Section \\ref{sec:dark} with a discussion of the evolution of the universe in our model.  The requirements a full quantum theory should satisfy for our findings to remain valid are formulated in Section \\ref{sec:dreams}. Section \\ref{sec:concl} is a summary of the results. ",
        "conclusions": "\\label{sec:concl} In this Letter we showed that the scale-invariant classical SM and the $\\nu$MSM coupled to unimodular gravity have all necessary ingredients to describe the evolution of the universe, including early inflation and late acceleration. The requirement of scale invariance and of the existence of a massless dilaton leads to a theory in which all mass scales, including that of gravity, originate from one and the same source. The unimodular character of gravity leads to the generation of an exponential potential for the dilaton, ensuring the existence of dark energy. If the full quantum field theory exhibits the same symmetries (for the explicit construction see \\cite{Shaposhnikov:2008xi,Shaposhnikov:2008ar}), it will have the same properties. Moreover, the argument can be reverted -- the observation of an accelerated universe with a dark energy component may tell that the underlying quantum field theory should be scale-invariant, and that this scale invariance should be broken spontaneously, leading to the massless dilaton.  The theory (\\ref{lagr}) with $\\beta=0$ and UG happens to be very rich in applications. It can address all confirmed signals which suggest that the SM is not complete: neutrino masses and oscillations, existence of dark and baryonic matter in the universe, early inflation and late acceleration, leading to an extra argument  against the necessity of new physics between the electroweak and  the Planck scale \\cite{Shaposhnikov:2007nj} (see also \\cite{Meissner:2006zh}). The discussion of particle physics experiments and astrophysical observations that can confirm or rule out this theory can be found in \\cite{Boyarsky:2006fg,Bezrukov:2006cy, Bezrukov:2005mx,Gorbunov:2007ak,Boyarsky:2006hr}. Our findings here indicate that the equation of state parameter $\\omega$ for dark energy must be different from that of the cosmological constant, but also that $\\omega > -1$, adding an extra cosmological test that could rule out the $\\nu$MSM.  {\\bf Acknowledgements.} This work was supported by the Swiss National Science Foundation. We thank F. Bezrukov, K. Chetyrkin, S. Sibiryakov and I. Tkachev for valuable comments."
    },
    "0809/0809.4260_arXiv.txt": {
        "abstract": "We present the results of a {\\it Spitzer} Infrared Spectrograph (IRS) survey of 24$\\mu$m-selected luminous infrared galaxies (LIRGs, $L_{\\rm IR} > 10^{11}L_\\odot$) in the rich cluster Cl\\,0024+16 at $z=0.4$.  Optically, these LIRGs resemble unremarkable spiral galaxies with e(a)/e(c) spectral classifications and [O{\\sc ii}]-derived star formation rates (SFRs) of $\\lesssim 2$\\,M$_\\odot$\\,yr$^{-1}$, generally indistinguishable from the `quiescent' star forming population in the cluster.  Our IRS spectra show that the majority of the 24$\\mu$m-detected galaxies exhibit polycyclic aromatic hydrocarbon (PAH) emission with implied SFRs $\\sim$30--60\\,M$_\\odot$\\,yr$^{-1}$, with only one ($<$10\\%) of the sample displaying unambiguous evidence of an active galactic nucleus in the mid-infrared. This confirms the presence of a large population of obscured starburst galaxies in distant clusters, which comprise the bulk of the star formation occurring in these environments at $z\\sim 0.5$.  We suggest that, although several mechanisms could be at play, these dusty starbursts could be the signature of an important evolutionary transition converting gas-rich spiral galaxies in distant clusters into the passive, bulge-dominated lenticular galaxies that become increasingly abundant in the cores of rich clusters in the $\\sim$4\\,Gyr to the present day. ",
        "introduction": "\\begin{table*} \\begin{center} \\caption{Properties of MIPS cluster galaxies targeted with IRS} \\scriptsize \\begin{tabular}{@{\\extracolsep{\\fill}}lccccccccl} \\tableline \\tableline & (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) \\cr Catalogue name (short name)  & R.A. & Dec. & $z$ & $S_{24}$& $F_{7.7}$& $F_{11.3}$ &SFR(IR) & Optical class &Note \\cr & (h m s) & ($\\circ$~$'$~$''$) && (mJy) & (10$^{-17}$\\,W\\,m$^{-2}$) & (10$^{-17}$\\,W\\,m$^{-2}$) &   (M$_\\odot$\\,yr$^{-1}$)\\cr \\tableline MIPS~J002606.1+170416.4 (a) &        00~26~06.1 &   +17~04~16.4 &0.3904 & 1.97 & $5.6\\pm0.8$ & $3.1\\pm0.2$ & $34\\pm10$ & --\\cr MIPS~J002721.0+165947.3 (b) &        00~27~21.0 &  +16~59~47.3  & 0.3964 & 1.48 & $9.6\\pm0.6$ & $4.6\\pm0.2$   & $59\\pm16$ &e(a) & H$_2$S(3)?\\cr MIPS~J002621.7+171925.7 (c) &        00~26~21.7 &  +17~19~25.7  & 0.3809 &0.95 & $9.1\\pm0.6$ & $2.9\\pm0.2$  & $56\\pm16$& e(c)  &H$_2$S(5)/[Ar~{\\sc ii}]? \\cr MIPS~J002715.0+171245.6 (d) &        00~27~15.0 &  +17~12~45.6  &0.3813 & 0.95 & $5.6\\pm0.9$ & $0.6\\pm0.3$   & $35\\pm11$ & --& Weak 11.3$\\mu$m \\& [Ne~{\\sc ii}]?\\cr MIPS~J002652.5+171359.9 (e) &        00~26~52.5 &  +17~13~59.9  & 0.3799 &0.84 & $10.1\\pm1.3$ & $3.9\\pm0.3$  & $62\\pm19$ & e(c) \\cr MIPS~J002609.1+171511.5 (f) &        00~26~09.1 &   +17~15~11.5 & 0.3940 & 0.77 & $4.6\\pm0.9$ & $2.9\\pm0.2$   & $28\\pm9$ & e(c) \\cr MIPS~J002703.6+171127.9 (g)  &        00~27~03.6 &  +17~11~27.9  &0.3956 & 0.71 & $6.8\\pm0.8$ & $2.3\\pm0.2$   & $42\\pm12$ & e(a)& [S~{\\sc iv}]?\\cr MIPS~J002636.3+165926.0 (h) &       00~26~36.3 & +16~59~26.0  & 0.4063 &0.70 & v. weak  & -- & -- & k+a  & AGN (X-ray source)\\cr MIPS~J002633.7+171221.4 (i)  &       00~26~33.7 & +17~12~21.4 & 0.3958 & 0.68 & -- & v. weak & --& e(a) & LL\\,2 only \\cr MIPS~J002620.2+170407.4 (j) &       00~26~20.2 & +17~04~07.4 & 0.3963 & 0.67 & -- & weak &-- & e(a) & LL\\,2 only\\cr MIPS~J002624.3+171129.8 (k)  &       00~26~24.3 & +17~11~29.8 & 0.3943&0.67 & -- & -- & -- & k & Not detected\\cr MIPS~J002637.3+170055.1 (l)  &       00~26~37.3 & +17~00~55.1  &0.3970 & 0.61 & -- & weak &-- & e(c) & LL\\,2 only \\cr \\tableline \\end{tabular} \\begin{minipage}{0.97\\textwidth}\\vspace{1mm} \\footnotesize {\\sc Notes ---}(1/2) Co-ordinate is centroid of 24$\\mu$m detection, J2000. (3) Redshift determined from optical spectroscopy (Czoske et al.\\ 2001; Moran et al.\\ 2007) (4) 16$''$ aperture-corrected flux density from Geach et al.\\ (2006). Typical uncertainty is 0.04\\,mJy (5/6) PAH 7.7 \\& 11.3$\\mu$m line flux estimated from the spectral decomposition code {\\sc pahfit} (Smith et al.\\ 2007). The total line flux is the total for the PAH line blends, and models the individual aromatic emission modes as Drude profiles (7) star formation rate derived from the far-infrared luminosity (Kennicutt\\ 1998), estimated from the 7.7$\\mu$m line luminosity (see \\S3.4 for more details). (8) optical classification scheme from the {\\it Morphs} collaboration (Dressler et al.\\ 1999) based on the equivalent widths of [O{\\sc ii}] and H$\\delta$. \\end{minipage} \\end{center} \\end{table*}   Passive lenticular (S0) galaxies make up a large fraction of the galaxies in clusters in the local Universe (Oemler 1974), but are conspicuously absent from equivalent environments just 5\\ Gyrs ago at $z\\sim0.5$ (Dressler et al.\\ 1997; Smail et al.\\ 1997; Couch et al.\\ 1998; Fasano et al.\\ 2000; Treu et al.\\ 2003). The reverse is true of star forming spiral galaxies (Butcher \\& Oemler\\ 1978, 1984; Couch \\& Sharples 1987), and so it has been proposed that these are the progenitors of local S0s. Although this proposal is still a contentious issue, recent work has attempted to piece together a coherent model for the evolutionary history of star forming galaxies being assimilated from the field into distant clusters (van Dokkum et al.\\ 1998; Treu et al.\\ 2003; Moran et al.\\ 2005, 2006; Poggianti et al.\\ 1999, 2004, 2006, 2008). These models generally invoke a modification of the star-formation histories and morphologies of spiral galaxies in high-density environments in order to replicate the spectral properties of the cluster population and the respective build-up and decline of the S0 and spiral cluster populations since $z\\sim0.5$ (e.g.\\ Smail et al.\\ 1998; Poggianti et al.\\ 1999; Moran et al.\\ 2007).  While there are several viable mechanisms to aid in the morphological transformation of spirals to S0s: harassment, ram-pressure stripping, merging, threshing, etc.\\ (e.g.\\ see Gunn \\& Gott\\ 1972; Moore et al.\\ 1998; Bekki et al.\\ 2001), the luminosities of the galaxies and in particular their bulges indicates the need for a period of enhanced stellar mass assembly in the spheroids of spiral galaxies in clusters to transform them into bulge-dominated S0s by the present day (Balogh et al.\\ 1999; Poggianti et al.\\ 1999; Kodama \\& Smail\\ 2001; Gerken et al.\\ 2004).   The simplest method of enhancing the stellar mass in the bulges of cluster spirals is to invoke a starburst episode within the central few kpc of these galaxies.  However, until recently there was little clear evidence for such starburst activity in these systems (e.g.\\ Balogh et al.\\ 1997; Poggianti et al.\\ 1999).  It has only been with the availability of sensitive mid-infrared surveys of intermediate redshift ($z\\gs 0.2$--0.5) clusters that potential starburst populations have been detected (Duc et al.\\ 2000, 2004; Fadda et al.\\ 2000; Metcalfe et al.\\ 2003; Biviano et al.\\ 2004; Coia et al.\\ 2004; Geach et al.\\ 2006; Marcillac et al.\\ 2007; Bai et al.\\ 2007; Fadda et al.\\ 2008; Oemler et al.\\ 2008; Dressler et al.\\ 2008).  The mid-infrared luminosities of the cluster members suggest bolometric luminosities of $L_{\\rm IR} > 10^{11}L_\\odot$, characteristic of Luminous Infrared Galaxies (LIRGs).  In a panoramic 24$\\mu$m mapping survey with {\\it Spitzer Space Telescope (SST)} we detected for the first time a population of LIRGs distributed to large clustocentric radius in two $z\\sim0.5$ clusters (Geach et al.\\ 2006).  These galaxies' mid-infrared emission is far higher than implied from the strength of their optical star formation rate (SFR) tracers such as H$\\alpha$ or [O{\\sc ii}].  Such photometric surveys give no clue to the origin of the intense dust-obscured activity pin-pointed by the mid-infrared emission, which could either be star formation or AGN-powered.  If the mid-infrared emission arises from star-formation, then the star-forming regions must be highly obscured by the carbonaceous and silicate material generated by massive stars in the final phases of stellar evolution.  Equally, however, it is also possible that these LIRGs could be powered by AGN, especially given the evidence for enhanced AGN activity in intermediate redshift clusters (e.g.\\ Martini et al.\\ 2006; Eastman et al.\\ 2007).  Therefore to understand the role of the luminous infrared population in the evolutionary cycle of cluster galaxies it is critical to quantify the relative importance of starburst and AGN emission in the mid-infrared population. If this population can be conclusively shown to represent starbursting galaxies, then they could be the long-sort candidates for the progenitors of local massive S0s, seen at a time when they are assembling the bulk of their stellar mass. Moreover, the intense star formation in these distant spirals could explain the apparent young ages ($\\sim5$\\,Gyr) of stellar populations in the bulges of S0s in the cores of local rich clusters (Kuntschner \\& Davies\\ 1998; Poggianti et al.\\ 2001; Mehlert et al.\\ 2003; Arag\\'on-Salamanca\\ 2007; Bedregal et al.\\ 2008).  In this paper we use low-resolution mid-infrared spectra obtained with the Infrared Spectrograph on-board the {\\it SST} to determine the origin of the mid-infrared emission in a sample of 24-$\\mu$m galaxies selected from one of the two clusters, Cl\\,0024+16 ($z=0.39$), studied by Geach et al.\\ (2006).  We use these data to test what fraction of these galaxies could be powered by AGN (from the strength of the PAH features, relative to the warm dust continuum: Rigopoulou et al.\\ 1999; Genzel et al.\\ 1998) -- this is critical for quantifying the true level of star formation in the 24$\\mu$m cluster population. Our second objective is to determine how the cluster LIRGs'  star formation histories can be linked with the relatively `quiescent' star forming population in the clusters. Finally, we examine the potential importance of the mid-infrared sources as a feeder population for passive lenticular galaxies found in the cores of local rich clusters. We attempt to reconcile these observations with the `fossil record' of cluster galaxy evolution: do they agree with information gleaned from detailed studies of the star formation histories of S0s in local massive clusters?     Our study focuses on 24$\\mu$m-selected galaxies in Cl\\,0024+16. This is one of the best studied clusters at intermediate redshifts, and is one of the original `Butcher \\& Oemler' clusters (Butcher \\& Oemler 1978, 1984), shown to have a large `blue fraction' of star forming galaxies (e.g.\\ Dressler et al.\\ 1997). The cluster also appears to be in the late stages of a line-of-sight minor merger that occurred $\\sim$3\\,Gyr ago (Czoske et al.\\ 2002) -- although it is not clear how this dynamical disturbance has affected the star formation histories of galaxies in the respective components. The large observational investment in studies of this cluster, including deep ground-based optical imaging, high-resolution {\\it HST} imaging studies (Treu et al.\\ 2003; Kneib et al.\\ 2003); panoramic H$\\alpha$ imaging (Kodama et al.\\ 2004); {\\it GALEX} ultra-violet imaging (Moran et al.\\ 2006); mid-infrared {\\it SST} MIPS and IRS coverage (Geach et al.\\ 2006) as well as wide-field optical spectroscopy (e.g.\\ Czoske et al.\\ 2001; Moran et al.\\ 2007) results in an exquisite data set with which to explore the properties of the star-forming population in this cluster.  Moran et al.\\ (2006) investigated the (rest-frame) far-UV (FUV) properties of so-called `passive spirals' in Cl\\,0024+16 (galaxies with spiral morphologies but with no apparent on-going star formation). Their red FUV-optical colors (compared to the general spiral population) suggests that star formation has recently been truncated in these galaxies. Moran et al.\\ (2007) favour a scenario where gas is prevented from cooling onto the discs of these galaxies, thus preventing new star formation: `starvation'. It has been argued that these galaxies are in a transitory phase between active spiral and passive lenticular galaxy, with gentle truncation of star formation on time-scales of $\\sim1$\\,Gyr allowing for residual star formation during a longer morphological transformation. Yet is still not clear whether typical spirals' luminosities are sufficient to build up the brightest lenticulars found in local clusters (Poggianti et al.\\ 1999; Kodama \\& Smail 2001).  Thus, part of the motivation of this work is to understand where the LIRGs fit into this evolutionary scenario, and whether these luminous galaxies can contribute significantly to the passive, red color-magnitude sequence in local clusters.  In \\S2 we describe the IRS observations and data reduction, and present the key results in \\S3. In \\S4 we interpret our observations in terms of the potential evolutionary connection to local S0s. \\S5 concludes and summarizes this work.  Throughout we will assume a cosmological model with $\\Omega_{\\rm m}=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  Unless otherwise stated, we quote magnitudes on the Vega system. ",
        "conclusions": "Our previous work uncovered a large population of luminous mid-infrared sources in the rich cluster Cl\\,0024+16. This work investigated a sub-sample those galaxies in greater detail, using new {\\it Spitzer} IRS spectroscopy specifically their mid-infrared and optical properties. Our main results can be summarized as follows: \\begin{enumerate} \\item{The LIRGs' mid-infrared spectra confirm that the majority of the 24$\\mu$m-selected galaxies in Cl\\,0024+16 have mid-infrared emission powered by star formation. At $S_{\\rm 24} > 0.6$\\,mJy we derive SFRs from the 7.7$\\mu$m PAH feature of $\\sim$30--60\\,M$_\\odot$\\,yr$^{-1}$. Only 1/12 of our 24$\\mu$m selected sample shows unambiguous evidence of an AGN. Our PAH-derived SFRs agree well with our previously estimated SFRs from the 24$\\mu$m luminosity. In contrast to the infrared-derived SFRs, the typical optical ([O{\\sc ii}]) derived SFRs for the same population are modest, with SFR([O{\\sc ii}])$\\lesssim 2M_\\odot$\\,yr$^{-1}$. Indeed, from an optical standpoint, the LIRGs' spectra resemble spiral galaxies and are mainly classified as the potentially dusty e(a) or un-remarkable e(c) galaxies in the nomenclature of Dressler et al.\\ (1999) and subsequent works.}  \\item{The most natural fate of LIRGs in $z\\sim0.5$ clusters is evolution into lenticular galaxies by the present day. Our arguments are supported by simple evolutionary models that show the expected luminosity evolution of the LIRGs after the cessation of star formation can match the bright-end of the local cluster S0 luminosity function. At their current rates, the LIRGs in Cl\\,0024+16 could build-up a $10^{10}$\\,$M_\\odot$ stellar bulge in $\\sim 200$\\,Myr. We propose that the mode of starburst in the LIRGs is -- like many local starbursts -- circumnuclear, although this has to be confirmed with higher-resolution observations beyond the scope of the current work. Such circumnuclear starbursts at $z\\sim0.5$ would favour the required bulge-to-disc evolution of a spiral--lenticular transformation, and would naturally explain the presence of $\\sim$4\\,Gyr old stellar populations observed in the bulges of some local S0s. }  \\end{enumerate}  Spiral galaxies falling into clusters at high redshift cannot escape from the deep potential well: their descendants must exist in the cores of clusters at the present day. For three decades the mystery of the co-eval reversal in fractions of spiral/S0s in rich clusters over the past 4\\,Gyrs has not been satisfactorily explained. However, the mid-infrared reveals the true situation: a fraction of these galaxies are actually undergoing vigorous star formation. We argue that these galaxies will evolve into S0s by the present day -- they could be the missing transition galaxies that bridge the evolutionary gap between the distant, active population of distant clusters and the passive `red and dead' lenticulars found in the local Universe."
    },
    "0809/0809.0396_arXiv.txt": {
        "abstract": "Planetary systems are angular momentum reservoirs generated during star formation. This accretion process produces very powerful engines able to drive the optical jets and the molecular outflows. A fraction of the engine energy is released into heating thus the temperature of the engine ranges from the 3000K of the inner disk material to the 10MK in the areas where magnetic reconnection occurs. There are important unsolved problems concerning the nature of the engine, its evolution and the impact of the engine in the chemical evolution of the inner disk. Of special relevance is the understanding of the shear layer between the stellar photosphere and the disk; this layer controls a significant fraction of the magnetic field building up and the subsequent dissipative processes ougth to be studied in the UV.  This contribution focus on describing the connections between 1 Myr old suns and the Sun and the requirements for new UV instrumentation to address their evolution during this period. Two types of observations are shown to be needed: monitoring programmes and high resolution imaging down to, at least,  milliarsecond scales. ",
        "introduction": "Solar sytem progenitors, from the deeply embedded sources to the weak line T Tauri Stars (WTTSs), are sources of high energy radiation (X-ray to UV); the total energy radiated in this range goes from $\\sim 0.02$~L$_{\\odot}$ measured in the very young sources through their X-ray radiation to the 0.2 L$_{\\odot}$ radiated in the UV during the T Tauri phase (or Phase~TT) (Preibish 2004, G\\'omez de Castro 2008). Later on, in the Weak line T Tauri Phase the energy released drop to $\\sim 10^{-3}$~L$_{\\odot}$ radiated in X-ray. Being TTSs intrinsically cool stars (with $\\log T_{\\rm eff}\\sim$ 6500-3600~K), surrounded by cool accretion disks radiating at infrared wavelengths, the source of this energy must be searched in the release of magnetic energy.  It is well known that the mediation of magnetic fields in the accretion process is able to heat up the plasmas since a fraction of the gravitational energy lost during accretion is invested in field amplification and dynamo action thus, radiative loses are pushed towards the high energy range. Unfortunately, in the early phases (ages $< 0.1$Myr) most of this radiation is reabsorbed by the dense circumstellar environment (Av$>3$) and only the hard X-ray radiation is able to escape from the system providing direct information on the evolution of the accretion process. After 1 Myr, extinction drops enough to make the engine accessible to UV wavelengths; current technologies allow to carry out high resolution spectroscopy in the UV range which is also extremely rich in spectral tracers so a single echelle spectra can provide information on molecular, atomic and ionized gas (from singly ionized gas to very high ionization species such as  Fe~XII, Fe~XVIII or  Fe~XXI). Thus, UV spectroscopy is an extremely efficient tool to study solar system progenitors from 1 Myr on with the current technology and these application will be discussed in detail below.  However, one can foresee a future when microarcseconds UV imaging  will be available and studies similar to those being run on the Sun, will be feasible from 1 Myr old Suns all the way down into the main sequence while the young planetary disk settles down and life begins to grow. This review deals with this; with a description of our current understanding of the evolution from the T Tauri phase to the modern Sun and with a non-technologically biased ambitious view of what we could learn from challenging new UV observatories. This review has been written after the end of the 1st. conference of the network for UV astronomy held in El Escorial in May 2007 where some challenging projects for new space observatories were presented; you should find references to some of them in this text. ",
        "conclusions": "In brief, the radiation produced by the accretion engine (including magnetospheres, outflows, accretion, inner disk and shock between winds and young planetary disks) is produced in the UV. To separate the various contributions is necessary either very high spatial resolution or moderate time resolution.  Current surveys show that although there are some few nearby TTSs and WTTSs sparsely distributed around the Sun (AB~Dor at 14.9~pc or TW~Hya at 56~pc) the nearest star forming complexes are concentrated in a {\\it Star Formation Belt} (SFB) at  140~pc around the Sun which includes Taurus, Auriga-Perseus, Ophiuchus, Lupus and Chamaleon molecular clouds and several thousands of pre-main sequence stars forming in various environments (clustered as in Ophichus, sparse as in Taurus). Resolving spatial scales of a tenth of the solar radius in the SFB would allow to study the connection between the star and the outflow in full detail. For this purposes spatial resolutions of 3.3 micro arcseconds are required; thus, for a fiducial wavelength of 1500\\AA\\ the aperture must be about 10~km. Such a long baseline interferometry should be carried out in the space and the Stellar Imager project (Carpenter et al 2008) represents a first atempt to such an ambitious project.  However, the requirements are not so strong to resolve the inner disk structure, the disk wind and the plasmoids ejection from the current layer between the stellar and the disk wind. In such a case, spatial resolutions of 1-0.5 milliarcseconds (mas) would be enough to map the SFB sources thus requiring apertures of 20~m and effective areas about 10$^4$ times those of HST/STIS; research in new coatings and detectors in the UV field (Kappelmann \\& Barnstedt, 2007) as well as a clever optical design may account for a factor of ten but, still, a larger, $\\sim 30$m aperture will be required to get SNR$\\simeq 10$ in the C~IV line in reasonable exposure times (few hours). There is an ongoing project that satisfies these requirements, the Fresnel Interferometer (see Koechlin et al 2008) and some projections on the expected performance of the interferometer on the mapping of the engine are plotted in Fig.~9.  Both space interferometer projects are under study by their national space agencies and, in case they succeed, they will be available about 2030. Is there anything else to be done {\\it in the meantime}?. The answer is definitely positive, time mapping will allow us to resolve the structures since the variability time scales are not the same for all the phenomena and they do not produce the same inprint in the spectra (neither in temperatures, densities or velocities). Some examples of the power of this technique have already been shown in this contribution.  High resolution spectroscopy (R$\\sim 50,000$) is enough to discriminate among the various components thus, scaling with the fluxes of the weakest H$_2$ lines (from Heczeg et al 2002) detected with the HST/STIS, a factor of 10 increase of the effective area with respect to HST/STIS is required to reach most of the sources in the SFB. The COS instrument in HST will provide a factor of 10 sensitivity increase with respect to STIS in the 1150-1700\\AA\\ range because of its optimized optical design (see Froening et al 2008), unfortunately the orbital constrains of HST do not favour monitoring programs. High orbit missions alike the WSO-UV (see Shustov et al 2008) are better suited for this purpose.  An additional factor of 10 would be required to obtain SNR$\\sim 10$ in exposure times of few minutes; this short time scales are necessary to map variations in flare time scales as shown in Fig.~8 for the pre-main sequence stars in the SFB. Unfortunately, spectroscopic monitorings of the flaring activity in the SFB will have to wait for future missions, with collecting areas about 8-10~m, preferably located at the L2.    \\begin{figure*} \\centerline{\\includegraphics[width=26pc]{gomezdecastro_fig9.eps}} \\caption[]{Theoretical prediction of the Si~III] emissivity from numerical simulations of star-disk interaction (G\\'omez de Castro \\& Rekowski, 2008). The stellar magnetosphere is assumed to be dipolar with a field strength at the surface of 1kG. The magnetosphere interacts with the disk which is under a moderate $\\alpha$-dynamo effect\\footnote{The magnetic field in the disk is assumed to be generated by a standard $\\alpha ^2 \\Omega $ dynamo (e.g. Krauser \\& Raedler 1980) where $\\alpha$ is the mean-field $\\alpha$ effect and $\\Omega$ the angular velocity of the plasma. $\\alpha$ scaling is: $\\alpha = -0.1 \\frac{z}{z_0} \\frac {\\chi _{\\rm disk} (r, z)}{1+ V_A^2/C_s^2}$ where $\\chi _{\\rm disk}$ is the disk profile (see von Rekowski et al 2003), $z_0$ is the disk half-thickness, $V_A$ is the local Alfv\\'en velocity and $C_s$ is the local sound speed.} . The inner disk wind is magnetocentrifugally accelerated ({\\it top panel}).  The convolution of the theoretical prediction with the point spread function of the 30m Fresnel Interferometer FII has been carried out by Laurent Koechlin and Truswin Raksasataya. The convolution is shown for three inclinations (0$^o$, 45$^o$ and 90$^o$) and three distances to the Earth (15~pc, 40~pc and 140~pc). Notice that the inner ring is resolved even for 140~pc. } \\end{figure*}"
    },
    "0809/0809.0958_arXiv.txt": {
        "abstract": "Examining flare data observed by TRACE satellite from May 1998 to December 2006, we choose 190 (151 M-class and 39 X-class) flare events which display post-flare loops (PFLs), observed by 171 \\AA~and 195 \\AA~wavelengths. 124 of the 190 events exhibit flare ribbons (FRs), observed by 1600 \\AA~images. We investigate the propagation of the brightening of these PFLs along the neutral lines and the separation of the FRs perpendicular to the neutral lines. Observations indicate that the footpoints of the initial brightening PFLs are always associated with the change of the photospheric magnetic fields. In most of the cases, the length of the FRs ranges from 20 Mm to 170 Mm. The propagating duration of the brightening is from 10 minutes to 60 minutes, and from 10 minutes to 70 minutes for the separating duration of the FRs. The velocities of the propagation and the separation range from 3 km s$^{-1}$ to 39 km s$^{-1}$ and 3 km s$^{-1}$ to 15 km s$^{-1}$, respectively. Both of the propagating velocities and the separating velocities are associated with the flare strength and the length of the FRs. It appears that the propagation and the separation are dynamically coupled, that is the greater the propagating velocity is, the faster the separation is. Furthermore, a greater propagating velocity corresponds to a greater deceleration (or acceleration). These PFLs display three types of propagating patterns. Type I propagation, which possesses about half of all the events, is that the brightening begins at the middle part of a set of PFLs, and propagates bi-directionally towards its both ends. Type II, possessing 30$\\%$, is that the brightening firstly appears at one end of a set of PFLs, then propagates to the other end. The remnant belongs to Type III propagation which displays that the initial brightening takes place at two (or more than two) positions on two (or more than two) sets of PFLs, and each brightening propagates bi-directionally along the neutral line. These three types of propagating patterns can be explained by a three-dimensional magnetic reconnection model. ",
        "introduction": "Flares are one of the most spectacular phenomena in solar physics. They are sudden brightening in the solar atmosphere, and consist of a number of components including loops, ribbons, arches, remote patches, surges, erupting filaments, and other expanding coronal features \\citep{mar89}. They have been studied morphologically from direct images \\citep[e.g.][]{kru00,fle01,kun01} and spectroscopically from spectrograms \\citep[e.g.][]{moo76,cow73,act85,cul97,gri05b} at different wavelength regions. Like most dynamic phenomena on the solar surface, the occurrence of solar flares is closely related to the presence and evolution of solar magnetic fields, especially the complicated, non-potential magnetic configuration \\citep{rus72,pat81,moo84}. Flares can be caused by rotating sunspot \\citep{bro03,tia06,zha07}, magnetic flux emergence \\citep{ish98,wan04,li07}, magnetic flux cancellation \\citep{liv89,wan92,wan93,zha01a,zha01b,zha02}, magnetic shear \\citep{kus04,wan06,ji06,su07}, and so on.  The chromospheric flare \\citep[e.g.][]{fan00,fal02,che06,hud07} is easier to be observed, especially with the help of an H$\\alpha$ filter. Larger flares often occur right after the sudden disappearance of a filament \\citep{kup81,din03,ste05,jia06a,jia06b,chi07}. In this case, the flare generally has the form of two flare ribbons (FRs) which lie on both sides of the location of the former filament. During flare decay phase, the FRs move apart. The separation of these FRs has been used to estimate the electric field in the reconnecting current sheet \\citep[e.g.][]{qiu02,asa04b} and also the coronal magnetic field strength and the reconnection rate \\citep[e.g.][]{iso02b,iso05}. Many authors have observed the parallel and antiparallel movements of the FRs along the arcade \\citep{hoy81,tak83,fle02,liu04,qiu04,sia04}.  Accompanying the FRs is a system of post-flare loops (PFLs) which initially appears at low altitude and then moves upward into the corona in consort with the motion of the ribbons \\citep{moo80}. A classical description of the PFLs, as seen in H$\\alpha$ images, was first given by \\citet{bru64} who noted that the ribbons essentially lie at the footpoints of the loop system, which forms a series of arcades \\citep{lin03}. These arcades are also frequently seen in the X-ray images recorded by Soft X-ray Telescope (SXT) \\citep{tsu91} aboard Yohkoh \\citep{oga91}, in the extreme ultraviolet (EUV) images from EIT \\citep[Extrmeme-ultraviolet Imaging Telescope,][]{del95} aboard the Solar and Heliospheric Observatory (SOHO) \\citep{dim95}, and also in the EUV images from Transition Region and Coronal Explorer (TRACE) \\citep{han99}. The formation of transient large-scale PFLs or post-eruptive arcades (PEAs) has been widely studied \\citep[e.g.][]{car64,stu66,hir74,kop76,car83,web87,sve97,tri04,tri05,tri06b}. \\citet{hud98}, \\citet{ste00} and \\citet{tri06b} studied a set of individual events about the relationship between PFLs and Coronal Mass Ejections (CMEs). The propagation of the loop formations along the neutral line, together with the separation of the FRs perpendicular to the neutral line was reported by \\citet{iso02a}. \\citet{gri05a} found a RHESSI observation showing that the hard X-ray (HXR) sources do not show the separation from the neutral line, but instead they move along the neutral line \\citep[see also][]{gof07}. \\citet{tri06a} investigated the relationship between the brightening propagation of the PEAs and the erupting filament/prominence. They reported two types of propagation of the brightening, and the propagating direction was consistent with the erupting direction of the filaments.  Due to the extremely low density and high temperature of the corona, measurements of the magnetic field are restricted to lower layers of the solar atmosphere. For this reason, extrapolation techniques, which attempt to reconstruct the coronal field from measured boundary values in the photosphere (or low chromosphere), are the prime tool for quantitative investigations of the coronal magnetic field \\citep{val05}. Some authors used the methods of extrapolation to study the evolution of flare loops in active region \\citep[e.g.][]{yan95,wan01,wie05,ama06,zha08}.  Magnetic reconnection of solar coronal loops is considered to be the main process that causes solar flares and possibly coronal heating. It is widely believed to be a mechanism of magnetic energy release \\citep[see reviews by][]{shi99,mar03}, and plays an important role in various explosive phenomena in astrophysical, space, and laboratory plasmas \\citep{bis93,taj97,pri00}. The evidence of magnetic reconnection found by space observations includes the cusp-shaped PFLs \\citep{tsu92}, the loop-top hard X-ray source \\citep{mas94}, the reconnection inflow \\citep{yok01,lin05}, downflows above PFLs \\citep{mck99,inn03,asa04a}, plasmoid ejections \\citep{shi95,ohy97,ohy98}, etc. Two-dimensional \\citep[e.g.][]{dey71,win91,hu95,for95,yok98,zha06} and three-dimensional simulations \\citep[e.g.][]{wu92,mag99,miy04,aul06,bir06} are widely used to study the magnetic reconnection in process of flares and CMEs.  The magnetic reconnection model proposed by \\citet{car64}, \\citet{stu66}, \\citet{hir74} and \\citet{kop76} (the CSHKP model) suggests that magnetic field lines successively reconnect in the corona. This model explains several well-known features of solar flares, such as the growth of flares loops with a cusp-shaped structure and the formation of the H$\\alpha$ two-ribbon structures at their footpoints. In recent decades, this model has been further extended \\citep[e.g.][]{pne81,pri90,pri00,pri02,moo92,moo01,shi99,yok01,lin00,lin04}.   In this paper, we mainly study the dynamic evolution of a larger sample of flare events, including the propagation of the PFLs along the neutral lines and the separation of the FRs away from the neutral lines. This investigation will provide some information of 3-dimensional magnetic reconnection in the process of flare. The criteria for the data selection and the methods of the data analysis are described in section 2. In section 3, we summarize our statistical results. Conclusions and brief discussion are shown in section 4. ",
        "conclusions": "We have studied the evolution of the PFLs and the FRs of the 190 flare events, and obtained the following results.  1. Both of \\emph{PD} and \\emph{SD} are associated with \\emph{FD}. The longer the \\emph{FD} is, the longer the \\emph{PD} and the \\emph{SD} are. The length of the FRs mainly ranges from 20 to 170 Mm. It is associated with \\emph{FC}, but weakly associated with \\emph{FD}. It increases as \\emph{FC} (\\emph{FD}) increases.  2. \\emph{V$_{p}$} ranges mainly from 3 km s$^{-1}$ to 33 km s$^{-1}$, and \\emph{V$_{s}$}, from 3 km s$^{-1}$ to 15 km s$^{-1}$ in 76$\\%$ of the cases. \\emph{V$_{p}$} and \\emph{V$_{s}$} are dynamically coupled, \\emph{V$_{p}$} increases as \\emph{V$_{s}$} increases. Both of \\emph{V$_{p}$} and \\emph{V$_{s}$} are positively correlated with \\emph{FC} and \\emph{Len} of the FRs. The brightening of the PFLs do not evenly propagate, 74$\\%$ of the PFL events are decelerated, and 26$\\%$, accelerated.  3. There are three types of propagation of the PFLs. Type I, possessing 49.5$\\%$ of all the events, is that the brightening begins at the middle part of a set of PFLs, and then propagates bi-directionally towards its both ends. Type II, in possession of 30$\\%$, displays that the brightening firstly appears in one end of a set of PFLs, then propagates to the other end. Type III, occupying 20.5$\\%$, shows that the initial brightening takes place at two (or more than two) positions on two (or more than two) sets of PFLs, and each brightening propagates bi-directionally.  The main error source of our measurement of the velocities (\\emph{V$_{p}$} and \\emph{V$_{s}$}) is due to the uncertainty of the outer edges. In our study, two pixels error has been used. As the average \\emph{PD} (\\emph{SD}) is 33 (38) minutes, the error of \\emph{V$_{p}$} (\\emph{V$_{s}$}) is 0.4 (0.3) km s$^{-1}$, which is much less than the velocities showed in Table 1. So the values of these velocities are reliable.  Several authors have paid attention to the propagation of the PFLs. \\citet{iso02a} reported that \\emph{V$_{p}$} is about 3-30 km s$^{-1}$ by using the data from Yohkoh/SXT \\citep{tsu91}, and 50-150 km s$^{-1}$ in \\citet{gri05a} from RHESSI observations. Based on the SOHO/EIT observations, \\citet{tri06a} studied several flare events relevant to long filament eruptions. They presented that \\emph{V$_{p}$} of the PEAs ranges from 20 to 111 km s$^{-1}$. In this paper, we show that \\emph{V$_{p}$} is mainly from 3 to 33 km s$^{-1}$, which is consistent with \\citet{iso02a}, but somewhat smaller than that of \\citet{tri06a}. The velocity difference between ours and Tripathi et al. may result from two aspects. The first aspect is the sample. Our sample contains hundreds of cases. The second is the datum source. We employ the TRACE data which have higher spatial (1$\\arcsec$) and temporal (1 minute) resolution. \\emph{V$_{s}$} of the FRs have been intensively studied, but the value are varied distinctly, e.g. 20-100 km s$^{-1}$ in \\citet{qiu02}, 0-85 km s$^{-1}$ in \\citet{asa04b,asa06} and 20-70 km s$^{-1}$ in \\citet{tem07}. By employing the TRACE 1600 \\AA~observations, \\citet{iso05} obtained \\emph{V$_{s}$} of 6.7-12 km s$^{-1}$ near the impulsive phase of the flares. Our results about \\emph{V$_{s}$} (3-15 km s$^{-1}$) are well in agreement with that of \\citet{iso05}.  It is basically accepted that the separation of the FRs and the propagation of the PFLs can be considered the signature of the successive magnetic reconnection along different directions. \\emph{SD} represents the duration of the magnetic reconnection where the reconnection points move upward, and \\emph{PD}, along the neutral lines. Both \\emph{SD} and \\emph{FD} determine the duration of flares. That is the reason why both \\emph{SD} and \\emph{PD} are associated with \\emph{FD} (see Figs. 7\\emph{a}-\\emph{b}). Lots of researchers have employed \\emph{V$_{s}$} to estimate the reconnection rate, they derived that the reconnection rate is very sensitive to \\emph{V$_{s}$} \\citep[e.g.][]{iso02b,iso05}. Investigation of the propagation of the PFLs and the separation of the FRs simultaneously will provide us with 3-dimensional picture of magnetic reconnection in the process of flares. In our study, we find that \\emph{V$_{p}$} of the PFLs is associated with \\emph{V$_{s}$} of the FRs (see Fig. 10\\emph{a}). This indicates that the separation of the FRs and the propagation of the PFLs are dynamical coupled. Our results show that both \\emph{V$_{s}$} of the FRs and \\emph{V$_{p}$} of the PFLs are associated with \\emph{FC} (see Figs. 9\\emph{a} and 9\\emph{c}). As \\emph{FC} represents the maximum energy release rate of the flares, this implys that \\emph{V$_{s}$} and \\emph{V$_{p}$} may also reflect the magnetic reconnection rate. \\emph{Len} of the FRs may be the length of the current sheet along the neutral line. Both \\emph{V$_{s}$} of the FRs and \\emph{V$_{p}$} of the PFLs are associated with \\emph{Len} of the FRs (see Figs. 9\\emph{b} and 9\\emph{d}). This may result from the magnetic configuration of the source region of the flare. \\citet{tri06a} have suggested that the longer the FRs is, the less magnetically complex the flare region is. We speculate that the brightening of the PFLs propagates easily in a simple magnetic configuration, and so do the FRs separate. All the parameters and the relationships between these parameters we discussed here will help us to better understand 3-dimensional magnetic reconnection, and provide an observational character for theoretical study and simulation of 3-dimensional magnetic reconnection.  By investigating 17 events, \\citet{tri06a} found two types of propagation of the brightening of the PEAs associated with asymmetric eruptions and symmetric eruptions of filaments, respectively. We have studied a much larger sample with higher spatial and temporal resolution observations, and found three types of propagation of the PFLs. Type I and Type II propagation are similar to those two types of propagations reported by \\citet{tri06a}. Type III propagation indicates that sometimes two or more than two sets of PFLs are involved in a flare event. All the three types of the propagation of the PFLs may be explained by schematic diagrams shown in Fig. 11 \\citep[see also][]{shi05,tri06a}. Figure 11\\emph{a} shows the explanation of Type I propagation. The X-point of the magnetic reconnection site occurs over the middle part of a set of loops (marked by \\emph{Reconnection}), then continuing magnetic reconnection occurs bi-directionally, thus results in the brightening of the PFLs appearing in the middle of the set of loops and propagating towards two opposite directions (see the grey thick arrows) as expected from the standard model. The white hollow arrows represent the separating direction of the FRs. Figure 11\\emph{b} illustrates that the magnetic reconnection site occurs over one end of a set of loops (mark by \\emph{Reconnection}), then the magnetic reconnection develops continuously towards the other end (see the black thick arrow). This mechanism can be employed to explain the observations of Type II propagation that the brightening of the PFLs occurs at one end of a set of PFLs and propagates from one end to the other one. Type III propagation displays a complex evolution pattern of the PFLs. We suggest that there are two (or more than two) magnetic reconnection sites (see Fig. 11\\emph{c}) which locate in two (or more than two) magnetic flux system. After the magnetic reconnection occurs, each of the brightening of the PFLs propagates bi-directionally along the neutral lines (see the black thick arrows). We check the MDI magnetograms and the TRACE WL observations of these events (e.g. see the examples displayed in Figs. 3-5), and find that the magnetic activity are clearly seen at the footpoints of the initial brightening of the PFLs. According to the observations of these three types of propagation, we propose that all the propagation of the PFLs are caused by magnetic reconnection. There is only one set of magnetic flux system involved in Type I and Type II propagation. The only difference of these two types of propagation is the position of the magnetic reconnection site. There are two or more than two sets of PFLs heated in the process of flare in Type III propagation. So the different appearance of these three types of the propagating PFLs may be determined by two effects. One is the different position of the magnetic reconnection site, the other is the different magnetic configuration of the source region of the flare.  Further studies using the data observed by Solar Terrestrial Relations Observatory (STEREO) will give us a more clear three-dimensional stereoscopic pictures about the evolution of the PFLs. Vector magnetograms (e.g. from Hinode) with higher spatial and temporal resolution are likely to uncover the magnetic configuration in detail. Moreover, complex three-dimensional simulation will be required to understand the dynamic behavior of the PFLs and the FRs."
    },
    "0809/0809.2500_arXiv.txt": {
        "abstract": "Using archival X-ray data from the second XMM-Newton serendipitous source catalogue, we present comparative analysis of the overall population of X-ray sources in the Large and Small Magellanic Clouds. We see a difference between the characteristics of the brighter sources in the two populations in the X-ray band. Utilising flux measurements in different energy bands we are able to sort the X-ray sources based on similarities to other previously identified and classified objects. In this manner we are able to identify the probable nature of some of the unknown objects, identifying a number of possible X-ray binaries and Super Soft Sources. ",
        "introduction": "The Magellanic Clouds (MCs) provide an ideal location for the study of source populations of all types.  At a distance of $48.5\\,{\\rm kpc}$ for the Large Magellanic Cloud (LMC; Macri et al., 2006) and $61\\,{\\rm kpc}$ for the Small Magellanic Cloud (SMC; Hilditch et al., 2005) it is entirely possible to resolve individual components of stellar populations with modern detectors.  For the study of X-ray sources (XRS), there are a number of other features of these systems that are beneficial.  Their comparatively small size on the sky means that observations of the entire systems are feasible with a small number of observations.  The low absorption towards the MCs, ${\\rm n_H} = 1.5 \\times 10^{21}$\\footnote{\\url{ http://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3nh/w3nh.pl}} towards the LMC and ${\\rm n_H} = 3.0 \\times 10^{21}$\\footnotemark[\\value{footnote}] towards the SMC benefits observations of sources, especially those which emit the bulk of their flux in the $0.1-3\\,{\\rm keV}$ soft band.  The Large Magellanic Cloud has a metallicity of about forty per cent of the Milky Way, $Z = 0.0091\\pm0.0007$, while the Small Magellanic Cloud has a metallicity of about ten per cent of the Milky Way, $Z = 0.0050\\pm0.0005$ (Keller \\& Wood, 2006) and their numerous epochs of star formation mean they are ideal laboratories for exploring the formation processes and evolution of stellar populations.  Prior to the launch of XMM-Newton and Chandra, the most complete X-ray surveys of the MCs were performed by the ROSAT observatory (Haberl \\& Pietsch 1999; Haberl et al. 2000). These surveys detected $\\sim\\!750$ and $\\sim\\!500$ sources in the LMC and SMC respectively.  The improved sensitivity and spatial resolution of the XMM-Newton observatory (Jansen et al. 2001) means that with a comparatively small number of observations it has already increased the number of XRS in the region of the MCs threefold (see Section \\ref{s:sources}) without approaching full coverage of the MC fields.  In addition, there is a vast array of datasets at other wavelengths for the MCs making follow-up of identified XRS easier without the need to await targeted observations to identify and analyse counterparts. ",
        "conclusions": "From the properties of identified sources in the 2XMM catalogue, we have defined criteria to identify the likely nature of unidentified X-ray sources in the LMC and SMC.  Using these properties, we have outlined the differences in the two populations.  We show that in the observed populations, there are proportionally more soft and bright hard sources in the SMC than in the LMC.  Some of these features can be attributed to selection effects.  However, it appears that others may be the result of the different metallicities and star formation histories of the two systems.  We have identified up to 20 new sources of interest in the MCs.  We intend to perform full analysis and follow-up of all of these sources including cross-matching of positions with optical and infrared sources to identify candidate companions.  This work is in its preliminary stages and we expect to be able to identify many new sources in the existing observations.  We have already identified a number of sources as candidate X-ray Binaries, Cataclysmic Variables and Super Soft Sources.  We also intend to carry out further observations of the MCs to obtain a more complete catalogue of all the X-ray sources.  \\\\  \\footnotesize{ \\noindent{\\textit{Acknowledgements}}\\\\  AJG would like to thank Suomen Akatemia (The Finnish Academy) for funding his post-doctoral position at Oulun Yliopisto (University of Oulu) under grant number 110792.  This work is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. The data was taken from the XMM-Newton second serendipitous source catalogue. A large part of the source selection was carried out using the TOPCAT package: \\texttt{http://www.starlink.ac.uk/topcat/}}"
    },
    "0809/0809.4506_arXiv.txt": {
        "abstract": " ",
        "introduction": "A 305-day Earth free precession was predicted by Euler in the 18th century and was sought since then by astronomers in the form of small latitude variations. F. K\\\"ustner in 1888 and S. C. Chandler in 1891 detected such a motion but with a period of approximately 435 days. Later it was shown by Love \\cite{1love} and Larmor \\cite{2larmor} that such a disagreement between theory and observation is in part consequence of elastic deformation of the Earth which was not considered by the Euler rigid body model (see Smith and Dahlen \\cite{3smith} for further corrections explanation). The so called Chandler wobble (CW) is actually one of the major components of the Earth polar motion, together with the annual wobble (AW) which has period of nearly 365 days. The composition of the two wobbles results in a motion with period of around 6.2 years due to the beat phenomenon. The ever increasing measurement precision, for which techniques like very long baseline interferometry (VLBI), global positioning system (GPS) and lunar laser ranging (LLR) are employed today, has lead to the conclusion that both wobbles have varying amplitudes. For AW, this may be attributed to seasonal displacement of atmospheric and water masses. On the other hand, excitation and damping of the Chandler wobble has become a puzzle to investigators. Several mechanisms have been proposed for its excitation like snow, hidrological, atmospheric and oceanic mass displacements, and large earthquakes, with defenders and opponents arguing for and against repeatedly and, in view of the large energy involved, the problem is still an open question. In more recent papers, Gross \\cite{4gross} attributes the excitation to ocean bottom pressure variations and changes in ocean currents due to winds while Seitz {\\it et al.} \\cite{5seitz} focus on combined effects due to atmospheric and oceanic causes.  Since it is believed that dark matter (DM) and dark energy composes around 96\\% of the mass of the Universe we investigate in this paper the possibility of DM to contribute, at least in part, for the excitation and damping of Chandler wobble. The effect of DM upon regular matter has been considered mainly on galactic-scale although more recently some authors have studied smaller scale effects. For instance, Froggatt and Nielsen \\cite{6frni}) have investigated the possibility of star-scale effects due to small DM balls; Fr\\`ere {\\it et al.} \\cite{fre} studied the presence of a DM halo centered in the Sun and its influence upon the motion of  planets; numerical simulations due to Diemand {\\it et al.} \\cite{die} show the presence of DM clumps surviving near the solar circle; Adler \\cite{adl} suggests that internal heat production in Neptune and hot-Jupiter exoplanets is due to accretion of planet-bound, not self-annihilating DM. Although the nature of the particles that form DM is not known precisely (see for instance Khalil and Mu\\~noz \\cite{6khalil}) we assume that, as usual, those particles do not interact with regular matter through electromagnetic force but only gravitationally. Furthermore, we speculate that they can interact with each other in such a way to set up an oblate (due to rotation) ellipsoid that might be present inside Earth, occupying the same space without violating any physical law. It can be shown that such an ellipsoid, made of self-interacting DM (otherwise such a structure would not be possible), would be attached to Earth through a spring-like restoring force. In this line, excitation and damping of CW could be, at least partially, consequence of energy exchange between the two bodies, associated to the restoring torque between them. The model does not consider the other known sources for the process nor the Earth internal structure, what are of course over simplifications. However this becomes possible a simple starting point for the calculations, to be sophisticated in a later step. H\\\"opfner \\cite{7hop,8hop} has managed to isolate the CW and AW components from available polar motion data. Our goal is to reproduce H\\\"opfner's results for Chandler wobble in various epochs by adjusting the unknown physical and geometrical parameters of the dark matter ellipsoid. For this purpose, we describe both body motion by solving numerically the set of differential equations that emerge from Euler equations and demonstrate that CW can be reproduced in such manner. This is explained in the next section while results and discussion are presented in section III. ",
        "conclusions": "With the aim of getting 10 unknown parameters of our model (all but $k_2$) we find that it is enough to establish the following conditions: 1) reproduce the Chandler component of the Earth rotation axis position predicted by H\\\"opfner \\cite{7hop} at the start of a wobble; 2) minimize the length of the prograde (counter-clockwise) trajectory followed during that wobble motion. The second condition is necessary in order to avoid intrusive nutations that might emerge accompanying the precessional motion. This means that we minimize, for a given initial guess of $k_2$, the sum of: a) the squared distance from the model calculed Earth rotation axis position $(\\widehat{\\it PMX},\\widehat{\\it PMY})$ at the start of the wobble and the one predicted by H\\\"opfner; b) the wobble trajectory length in the plane $(\\widehat{\\it PMX},\\widehat{\\it PMY})$ obtained by numerical integration. For this minimization process we used the {\\tt POWELL} routine from the Numerical Recipes package (Press {\\it et al.} \\cite{14press}). However the generalized simulated annealing ({\\tt GSA}) code (Mundim and Tsallis \\cite{15mundim}, Dall'Igna Junior {\\it et al.} \\cite{16dall}) has proved to be very effective in a first step run in order to prevent   getting stuck in local minima. The 11th parameter, $k_2$, is then found by reproducing the Chandler component of the Earth rotation axis position at the end of the wobble, with reference to H\\\"opfner's prediction, so that $k_2$ is responsible for the adjustment of a particular wobble period. We show in Figs. 7-9 some of our results for the CW in three different epochs, in comparison with the results due to H\\\"opfner \\cite{7hop}, which are represented by dots. The first one, for which $\\omega^*_{3^*}$ is negative, has the typical quality of providing reasonable reproduction of the motion. The next Figures show the two exceptions to this rule, which tend to present almost circular trajectories instead of, in those epochs, the slightly elliptical ones due to H\\\"opfner. This may be caused by the above mentioned criteria -- mainly condition 2) -- adopted in the fit process. We also show in Table I the corresponding values for the 11 parameters (with respective numerical uncertainties shown in Table II), while in Table III we have the initial values of $\\dot{\\theta}$ and $\\dot{\\theta}^*$ calculated from equations (42) and (43) with appropriate signs and the corresponding values of the torque strength $c_1$. As expected -- see comments after equation (7) -- $\\alpha$ turns out to be close to 90$^\\circ$. On the other hand, parameters $\\omega^*_{3^*}$ (the main angular velocity component of the ellipsoid), $A^*$, $C^*$ and $k_2$ present oscillations when examining the respective columns in the Table I, being more noticeable for the 44330-44755 interval. For $k_2$, since it is connected to the period and that particular wobble has a short one -- see Table IV for a comparison between the presently calculated periods and the ones due to H\\\"opfner which they shall reproduce -- this is not surprising although undesirable. H\\\"opfner \\cite{7hop,8hop}, like other authors (see {\\it e. g.} Liu {\\it et al.} \\cite{17liu}, Wang \\cite{18wang}), found variable CW periods but this is not a clearly solved question since some investigators have another opinion (see {\\it e. g.} Vicente and Wilson \\cite{19vic}, Jochmann \\cite{20jo}, Liao and Zhou \\cite{21liao}, Guo {\\it et al.} \\cite{22guo}). For the sake of comparison, values of $k_2$ obtained by some other authors are 0.284 (Kaula \\cite{12kaula}), 0.30088 (model 1066A of Gilbert and Dziewonski \\cite{23gil}) which are not far from what we got with the exception of the 44330-44755 case.  By comparing the values of $C^*$ with momenta of inertia of natural satellites present in the solar system we find that they range from about Miranda's (Uranus) to Charon's (Pluto). Since we are able only to calculate the DM ellipsoid principal momenta of inertia, which are related to the ellipsoid mass $M^*$ and their equatorial and polar radii $a^*$ and $c^*$ through $A^* = M^*(a^{*2}+c^{*2})/5$ and $C^* = 2M^*a^{*2}/5$, we cannot establish the ellipsoid mass uniquely, and in consequence, the fraction of DM present in the Earth is not predicted. Notice that, for an oblate ellipsoid, we have $c^*\\leq a^*$ and, therefore, the relation $C^*/2 \\leq A^* \\leq C^*$ holds and it was required to be satisfied in our numerical procedure. Just to have an idea of the order of magnitude of $M^*$, let us suppose that $a^*\\approx 2350$ km, {\\it i. e.}, the ellipsoid has the average radius of the outer core of the Earth. From the values of $C^*$ in Table I, this implies that the fraction of DM present in the Earth would lie in the range $6 \\times 10^{-7}$ to $9 \\times 10^{-5}$. We emphasize that, according to the model, the Earth motion and its gravitational interaction with other bodies have already incorporated the presence of DM through what has been known effectively as ``Earth mass\", so that the outcome of its presence might be more easily observed through the Earth wobble.  Following the hypothesis of changing CW periods we present in Fig. 10 the correlation between period and average amplitude. We show H\\\"opfner's results \\cite{7hop} and ours. The last ones are fitted by the full line, which has the form: period (in days) = 440--3101.47$\\times$exp[\\mbox{--0.0347946} $\\times$ amplitude (in mas)]. As a reference we also present the fits due to Liu et al \\cite{17liu} related to the epochs 1912-1928 and 1936-1948. The correlation clearly varies with time.  One might argue that the parameter fluctuations in Table I indicate that this model is not sound. However those fluctuations may be caused by the model simplistic -- although convenient at this stage -- approach that we are considering the DM ellipsoid the only source responsible for the complexity of Chandler wobble, what obviously is not true. Therefore, the predicted $A^*$ and $C^*$ dark matter ellipsoid parameters should be considered only as upper limit values.  Another point that could be raised against the model is that it is just emulating one or more of the attributed causes for CW -- as listed in the Introduction -- and, in consequence, the DM ellipsoid existence is not real. However, it can be shown that externally to the DM ellipsoid, whose rotation axis is almost orthogonal to the Earth's, it generates a non-isotropic, time dependent, gravitational field with two contributions, with respect to a ${\\bf u}_1{\\bf u}_2{\\bf u}_3$ reference frame attached to the Earth. The first term is radial: \\begin{equation} a_r=\\frac{3G(C^*-A^*)}{2r^4}\\left[3\\sin^2\\theta(\\cos\\phi\\sin\\omega_Tt+ \\sin\\phi\\cos\\omega_Tt)^2-1\\right].                                      % \\end{equation} Here, $\\theta$ is the co-latitude, $\\phi$ is the longitude and $\\omega_T$ is the Earth angular velocity. Of course, $\\omega_Tt$ is defined except for a phase that would establish the position of the ellipsoid rotation axis with respect to the Greenwich meridian (for instance, in the direction of ${\\bf u}_1$), at a given time. The second contribution is \\begin{eqnarray} && \\hspace{-1.0cm}{\\bf a}_{\\theta^*}=\\frac{3G(C^*-A^*)}{2r^4}\\sin2\\theta^*\\left[( \\cos\\theta^*\\sin\\phi^*\\cos\\omega_Tt-\\sin\\theta^*\\sin\\omega_Tt){\\bf u}_1 \\right. \\nonumber \\\\ &&  \\hspace{2.2cm} \\left. \\mbox{} -(\\cos\\theta^*\\sin\\phi^*\\sin\\omega_Tt+\\sin\\theta^* \\cos\\omega_Tt){\\bf u}_2+\\cos\\theta^*\\cos\\phi^*{\\bf u}_3\\right]   % \\end{eqnarray} where \\begin{eqnarray} && \\sin\\theta^*=\\sqrt{\\cos^2\\theta+\\sin^2\\theta(\\cos\\phi\\cos\\omega_Tt-\\sin\\phi\\sin \\omega_Tt)^2} \\\\                                           % && \\cos\\theta^*=\\sin\\theta(\\cos\\phi\\sin\\omega_Tt+\\sin\\phi\\cos\\omega_Tt)) \\\\ % && \\sin\\phi^*=\\frac{\\sin\\theta(\\cos\\phi\\cos\\omega_Tt-\\sin\\phi\\sin\\omega_Tt)} {\\sqrt{\\cos^2\\theta+\\sin^2\\theta(\\cos\\phi\\cos\\omega_Tt-\\sin\\phi\\sin\\omega_Tt)^2}} \\\\ % && \\cos\\phi^*=\\frac{\\cos\\theta}{\\sqrt{\\cos^2\\theta+\\sin^2\\theta(\\cos\\phi\\cos\\omega_Tt- \\sin\\phi\\sin\\omega_Tt)^2}}. \\end{eqnarray} Note that, in the Earth's reference frame, the DM ellipsoid gives a complete turn in a day. Supposing the DM ellipsoid is smaller than the Earth core, this gravitational field certainly enforces viscous flow in the outer core, which is in a fluid state. As a first consequence, it can provide the necessary energy to the magneto-hydrodynamic motion in the outer core, which is closely related to the generation of the geomagnetic field. This could give a plausible explanation to one of the basic problems of geophysics (see Stacey, ref. \\cite{9stacey}), namely, what is the source of this energy. However the numerical solution of this motion through the Navier-Stokes and energy conservation equations suffers lack of knowledge of the flowing material density dependence with respect to pressure and temperature. Secondly, it can generate heat in the core, that would contribute to decrease the shortfall of 0.7 TW to the energy necessary to maintain the adiabatic temperature gradient in the core, leaving more power to drive the geomagnetic field dynamo \\cite{9stacey}. On the other hand, Mack {\\it et al.} \\cite{mack} have concluded that DM is unlikely to contribute not only to Earth's internal heat flow but also to hot-Jupiter exoplanets. However they studied the contribution of DM self-annihilation to heat generation, what is a different idea from ours. Conversely, Adler \\cite{adl} has also studied the contribution of self-annihilating and non-self-annihilating DM accretion to the internal heat of the Earth, Jovian planets and hot-Jupiter exoplanets. His conclusion is that this process is plausible provided efficient DM capture is occurring. In our description, it is feasible to imagine DM ellipsoids also present is such bodies, generating heat by enforcing internal viscous flow, like in Earth. Therefore this model can mean much more than mere emulation of known suggested causes for CW but it also touches open questions like geomagnetic field dynamo and heat generation in the Earth's outer core and in other planets.  Rather than giving the final answer for the problem that has raised the attention of geophysicists and astronomers for several decades, we expect that this calculation may open room for a new and potentially important key component in CW which was completely unsuspected untill now, to be considered in more sophisticated models. Future development will decide about the real significance of this approach that might have relevant consequences in our comprehension about dark matter. At least, it has been demonstrated that two interacting ellipsoids can have Chandler-like wobble.  \\vspace {1cm} \\noindent {\\bf ACKNOWLEDGMENTS} \\vspace {3mm}  The author would like to thank Prof. Joachim H\\\"opfner (retired from GFZ-Potsdam) for enlightening correspondence and for providing unpublished results, Prof. Kleber C. Mundim, from Chemistry Institute-University of Brasilia, for allowing the use of his {\\tt GSA} code, Prof. Nivaldo A. Lemos, from Universidade Federal Fluminense, for correspondence in the early stages of this work, Prof. Marcos D. Maia, from Institute of Physics-University of Brasilia, for his interest and for pointing me some related literature, Prof. Marcus B. Lacerda Santos, also from Institute of Physics-University of Brasilia, for stimulating discussions, and MEC-SESu for tutorship (PET program). Finally, thanks are due to the Referee of this paper, for interesting remarks.  \\pagebreak"
    },
    "0809/0809.4489_arXiv.txt": {
        "abstract": "The recent detection of Sagittarius~A* at $\\lambda = 1.3$~mm on a baseline from Hawaii to Arizona demonstrates that millimeter wavelength very long baseline interferometry (VLBI) can now spatially resolve emission from the innermost accretion flow of the Galactic center region.  Here, we investigate the ability of future millimeter VLBI arrays to constrain the spin and inclination of the putative black hole and the orientation of the accretion disk major axis within the context of radiatively inefficient accretion flow (RIAF) models. We examine the range of baseline visibility and closure amplitudes predicted by RIAF models to identify critical telescopes for determining the spin, inclination, and disk orientation of the Sgr~A* black hole and accretion disk system.  We find that baseline lengths near $3$~G$\\lambda$ have the greatest power to distinguish amongst RIAF model parameters, and that it will be important to include new telescopes that will form north-south baselines with a range of lengths.  If a RIAF model describes the emission from Sgr~A*, it is likely that the orientation of the accretion disk can be determined with the addition of a Chilean telescope to the array.  Some likely disk orientations predict detectable fluxes on baselines between the continental United States and even a single 10--12~m dish in Chile. The extra information provided from closure amplitudes by a four-antenna array enhances the ability of VLBI to discriminate amongst model parameters. ",
        "introduction": "\\label{introduction}  The Galactic center radio source Sagittarius~A* is believed to be associated with a black hole with a mass of approximately $4 \\times 10^6$~M$_\\sun$ at a distance of about 8~kpc \\citep{schodel03,ghez08}. The event horizon of Sgr~A* has the largest apparent angular size of all known black holes as viewed from Earth.  High angular resolution is critical to understanding Sgr~A* because the size scales are so small: for instance, the apparent lensed horizon size is $55~\\mu$as. The quest for angular resolution has driven observers toward very long baseline interferometry (VLBI) at millimeter wavelengths.  Most recently, \\citet{doeleman08} detected Sgr~A* at 230~GHz ($\\lambda = 1.3$~mm) on baselines between the James Clerk Maxwell Telescope (JCMT) in Hawaii and the Submillimeter Telescope Observatory (SMTO) in Arizona, as well as between the SMTO and a Combined Array for Research in Millimeter-wave Astronomy (CARMA) telescope in California.  The former is the longest baseline ($3.5$~G$\\lambda$, fringe spacing $\\sim 60~\\mu$as) on which Sgr~A* has ever been detected.  In the absence of a true ab initio simulation of Sgr~A*'s accretion flow, a variety of physically simplified models have been proposed. Constraints from the observed spectrum, polarization, and variability have not proved sufficient to conclusively establish the nature of the emission region.  This is evidenced by the continuing vigorous debate over the morphology of this region, e.g., primarily from the innermost regions of an inefficient accretion flow \\citep{yuan03}, the acceleration region of a nascent jet \\citep{falcke00}, or something qualitatively different \\citep{igumenshchev03}.  However, millimeter VLBI's ability to resolve the emitting region directly promises to discriminate amongst these.  To do so will require a detailed understanding of the role that key physical parameters (such as black hole spin, disk/jet orientation, etc.) play in shaping the millimeter images of Sgr~A*.  Conversely, if the context of the emission can be conclusively established, millimeter VLBI has great potential to extract these parameters.  In a related work, \\citet{broderick08} examine the ability of the \\citet{doeleman08} VLBI detections to discriminate amongst various radiatively inefficient accretion flow (RIAF) models parameterized by black hole spin, viewing angle, and disk orientation.  A key finding is that the detected flux on the JCMT-SMTO baseline argues against orientations in which the assumed accretion disk is close to face on. However, the RIAF parameters are strongly coupled, and there are not yet enough millimeter VLBI measurements to make strong statements about the black hole spin, for instance.  In this work, we explore the question of which future millimeter VLBI observations are likely to have the greatest impact for distinguishing between RIAF models and are therefore most likely to determine the physical parameters of the black-hole/accretion-disk system. ",
        "conclusions": "The planned evolution of millimeter VLBI capability over the next few years will place strong constraints on RIAF models of Sgr~A* emission. It is important to devise an observing strategy to test the RIAF hypothesis and extract information on the parameters of the putative accretion disk system of Sgr~A*, including the disk orientation, inclination, and black hole spin.  By identifying differences between models that are consistent with present VLBI observations, we identify four key points to guide future millimeter VLBI observations:  1.\\ There is a preferred zone of baseline lengths near $3$~G$\\lambda$ where there is enough angular resolution to distinguish between many of the models but not so much resolution that a large amount of flux is lost.  The JCMT-SMTO detections are in this zone.  The construction of a phased-array processor for the Hawaiian millimeter telescopes will allow for even tighter constraints to be placed on RIAF model parameters on the Hawaii-SMTO baseline as well as increase the chance of detecting Sgr~A* on other baselines to Hawaii.  2.\\ There is a strong correlation between fluxes on the Hawaii-SMTO and Hawaii-CARMA baselines, limiting the ability of the latter baseline to place extra constraints on RIAF parameters unless a phased-array processor is installed at CARMA, making Hawaii-CARMA a more sensitive baseline than Hawaii-SMTO.  Nevertheless, the Hawaii-CARMA measurement will be important both for ruling out simpler models (e.g., a Gaussian or annulus) of the emission from Sgr~A* that the present VLBI measurements cannot and for testing the RIAF hypothesis.  3.\\ VLBI arrays that include either the LMT or a Chilean telescope provide the best constraints on RIAF models.  Predicted flux densities on the SMTO-LMT and CARMA-LMT baselines are well in excess of 1~Jy. Thus, the LMT can significantly enhance Sgr~A* VLBI even with a surface not adjusted for maximum accuracy.  Likewise, since many models predict flux densities on the SMTO-Chile and CARMA-Chile baselines well in excess of 0.5~Jy, inclusion of even a single 10~m or 12~m Chilean telescope will result in important model constraints on Sgr~A*.  4.\\ Closure amplitudes will be critical, as amplitude calibration errors impose systematic errors on measured baseline flux densities that reduce their ability to constrain RIAF model parameters.  At least four telescopes should be used in an observing array in order to obtain closure amplitudes.  If no fourth telescope site is available at the time of reobservation, a second antenna at the CARMA site should be included in a $\\lambda = 1.3$~mm observing array.  We note the important caveat that our analysis assumes that a RIAF accurately and completely describes the emission from Sgr~A* at 230~GHz.  This assumption may be incorrect and is likely at best an approximation to the actual, complex emitting region in Sgr~A*.  For these reasons, it will be important to eventually obtain data on as many different millimeter-wavelength VLBI baselines as is feasible."
    },
    "0809/0809.0780_arXiv.txt": {
        "abstract": "Studies of the formation of metal-free Population III stars usually focus primarily on the role played by $\\mHt$ cooling, on account of its large chemical abundance relative to other possible molecular or ionic coolants. However, while $\\mHt$ is generally the most important coolant at low gas densities, it is not an effective coolant at high gas densities, owing to the low critical density at which it reaches local thermodynamic equilibrium (LTE) and to the large opacities that develop in its emission lines. It is therefore possible that emission from other chemical species may play an important role in cooling high density primordial gas.  A particularly interesting candidate is the $\\htp$ molecular ion. This ion has an LTE cooling rate that is roughly a billion times larger than that of $\\mHt$, and unlike other primordial molecular ions such as $\\mHtp$ or ${\\rm HeH^{+}}$, it is not easily removed from the gas by collisions with $\\mH$ or $\\mHt$. It is already known to be an important coolant in at least one astrophysical context -- the upper atmospheres of gas giants -- but its role in the cooling of primordial gas has received little previous study.  In this paper, we investigate the potential importance of $\\htp$ cooling in primordial gas using a newly-developed $\\htp$ cooling function and the most detailed model of primordial chemistry published to date. We show that although $\\htp$ is, in most circumstances, the third most important coolant in dense primordial gas (after $\\mHt$ and HD), it is nevertheless unimportant, as it contributes no more than a few percent of the total cooling. We also show that in gas irradiated by a sufficiently strong flux of cosmic rays or X-rays, $\\htp$ can become the dominant coolant in the gas, although the size of the flux required renders this scenario unlikely to occur. ",
        "introduction": "Over the past decade, we have made substantial progress in understanding how the very first stars in the universe formed.  We know that in cosmological models based on cold dark matter (CDM), the first stars will form in small protogalaxies, with total masses of the order of $10^{5}$ -- $10^{6} \\: \\msun$, and that by a redshift $z \\sim 30$ we expect to find at least one such star-forming system per comoving ${\\rm Mpc}^{3}$ \\citep{yahs03}. We also know that although molecular hydrogen formation is inefficient, owing to the absence of dust, it is nevertheless the most abundant molecule in primordial gas, and is the main source of cooling at low densities, for temperatures between $\\sim 200$~K and  $8000$~K.  Furthermore,  simple semi-analytical estimates \\citep[e.g.,][]{teg97}, later confirmed by detailed simulations \\citep[e.g.,][]{yahs03}, demonstrate that $\\mHt$ provides enough cooling in these small protogalaxies to allow the gas to collapse under the influence of gravity on a timescale comparable to its gravitational free-fall timescale, thereby allowing star formation to occur. High-resolution, adaptive mesh simulations performed by \\citet{abn00, abn02} have taught us much about the dynamics of the gas in these first protogalaxies. They consider gas cooled only by $\\mHt$, and find that gravitational fragmentation of the collapsing gas is inefficient and that therefore only a single, massive star will form during the collapse. This will then suppress further star formation through a variety of feedbacks \\cite[see the recent review by][]{cf05}. Other simulations, making simplifying assumptions such as the adoption of spherical symmetry, have allowed us to examine the importance of various aspects of the physics of the gas which are currently difficult to treat in the high-resolution adaptive mesh simulations (for instance, the development of large optical depths in the rotational and  vibrational lines of $\\mHt$ at high densities; see e.g.\\ \\citealt{on98,ripa02}).  Much of the theoretical uncertainty that remains concerns the behaviour of dense gas during the later stages of gravitational collapse, and during the period of accretion that follows the formation of the first protostar. There is much that is still unknown here -- for instance, a complete understanding of the mechanism by which the collapsing gas loses much of its initial angular momentum still eludes us -- but in this paper we intend to focus on one relatively simple aspect: the identification of the dominant coolant(s) in the dense gas.  As previously noted, $\\mHt$ cooling has long been known to dominate at low densities, and it is frequently assumed that it also dominates at high densities. However, it is not at all obvious that this is actually the case.  Two factors dramatically reduce the effectiveness of $\\mHt$ as a coolant in high density gas. The first is the fact that the excited rotational and vibrational levels of $\\mHt$ have small radiative transition probabilities, and hence long radiative lifetimes ($\\tau \\simgreat 10^{6} \\: {\\rm s}$; \\citealt{wsd98}). This means that collisional de-excitation of excited $\\mHt$ becomes competitive with radiative de-excitation at fairly low gas densities ($n \\sim 10^{4} \\: {\\rm cm^{-3}}$), and so as the number density $n$ increases, the cooling rate of $\\mHt$ quickly reaches its local thermodynamic equilibrium (LTE) value, given by \\begin{equation} \\Lambda_{\\rm \\mHt, LTE} = \\sum_{i, j>i} A_{ji} E_{ji} n_{j}, \\end{equation} where $A_{ji}$ is the transition probability for a transition from $j \\rightarrow i$, $E_{ji}$ is the energy of this transition, $n_{j}$ is the number density of $\\mHt$ molecules in level $j$, computed assuming LTE, and where we sum over all bound levels $i$ and over all bound levels $j$ with energies greater than $i$. In the LTE limit, the cooling rate per $\\mHt$ molecule is independent of the gas density, and is largely determined by the size of the transition probabilities. Since these are small, the LTE cooling rate is also small. Consequently, at high gas densities, a molecular species whose excited states have much shorter radiative lifetimes than those of $\\mHt$ will provide far more cooling {\\em per molecule} than $\\mHt$.  The second factor reducing the effectiveness of $\\mHt$ cooling at very large $n$ is the fact that the gas eventually becomes optically thick in the cores of the main rovibrational lines of $\\mHt$. The effects of this cannot currently be treated fully in three-dimensional simulations, due to the high computational cost of solving the resulting radiative transfer problem, but it has been modeled accurately in simple spherically symmetric, one-dimensional simulations \\citep{on98,ripa02,ra04} and in an approximate fashion in three dimensions \\citep{yoha06}. These studies confirm that at densities $n \\simgreat 10^{10} \\: {\\rm cm^{-3}}$, optical depth effects significantly suppress $\\mHt$ cooling.  Together, these factors combine to render $\\mHt$ a fairly ineffective coolant in high density gas, despite the fact that at $n > 10^{10} \\: {\\rm cm^{-3}}$, several three-body $\\mHt$ formation reactions \\begin{eqnarray} \\mH + \\mH + \\mH  & \\rightarrow & \\mHt + \\mH, \\\\ \\mH + \\mH + \\He  & \\rightarrow & \\mHt + \\He, \\\\ \\mH + \\mH + \\mHt & \\rightarrow & \\mHt + \\mHt, \\end{eqnarray} rapidly convert almost all of the hydrogen in the gas to $\\mHt$ \\citep{pss83}. It is therefore reasonable to ask whether cooling from any of the other molecular species present in the gas will become competitive with $\\mHt$ cooling at these densities.  One obvious possibility is deuterated hydrogen, $\\hd$. Its excited rotational and vibrational levels have radiative lifetimes that are about a factor of 100 shorter than those of $\\mHt$,  and so the $\\hd$ cooling rate does not reach its LTE limit until $n \\sim 10^{6} \\: {\\rm cm^{-3}}$. It is also a far more effective coolant than $\\mHt$ at low temperatures ($T \\simless 200 \\: {\\rm K}$; see e.g., \\citealt{flpr00}). This is due primarily to the fact that radiative transitions can occur between rotational levels with odd and even values of $J$, allowing cooling to occur through the $J=1 \\rightarrow 0$ transition. The corresponding odd $\\leftrightarrow$ even transitions in the case of $\\mHt$ represent conversions from ortho-$\\mHt$ to para-$\\mHt$ or vice versa, and are highly forbidden. Furthermore, at low temperatures the ratio of $\\hd$ to $\\mHt$ can be significantly enhanced with respect to the cosmological D:H ratio by chemical fractionation \\citep[see e.g.,][]{glo08}.  The role of $\\hd$ cooling in early protogalaxies has been investigated by a number of authors. In the case of the earliest generation of protogalaxies, which form from very cold neutral gas that is never heated to more than a few thousand K during the course of the galaxy formation process, $\\hd$ cooling appears to be unimportant, as the collapsed gas does not become cold enough for sufficient fractionation to occur to make $\\hd$ cooling dominant \\citep{bcl02}.  The situation is rather different, however, in primordial gas cooling from an initially ionized state. In that case more $\\mHt$ is formed, allowing the gas to cool to a lower temperature, at which point fractionation becomes effective and $\\hd$ cooling rapidly becomes dominant \\citep[see e.g.,][]{nu02,no05,jb06,sv06}. However, the initial ionization required is much larger than that expected to be present in the earliest protogalaxies.  Another molecule to have attracted considerable attention is lithium hydride, LiH. This molecule has a very large dipole moment, $\\mu = 5.888 \\: {\\rm debyes}$ \\citep{zs80}, and consequently its excited levels have very short radiative lifetimes. Therefore, despite the very low lithium abundance in primordial gas ($x_{\\rm Li} = 4.3 \\times 10^{-10}$, by number; see \\citealt{cy04}), it was thought for a time  that LiH would dominate the cooling at very high densities \\citep[see e.g.,][]{ls84}. However, accurate quantal calculations of the rate of formation of $\\lih$ by radiative association \\citep{dks96,ggg96,bdlg03} \\begin{equation} \\li + \\mH \\rightarrow \\lih + \\gamma, \\end{equation} have shown that the rate is much smaller than was initially assumed, while recent work by \\citet{dpgg05} has shown that the reaction \\begin{equation} \\lih + \\mH \\rightarrow \\li + \\mHt, \\end{equation} has no activation energy and so will be an efficient destruction mechanism for $\\lih$ for as long as some atomic hydrogen remains in the gas. Consequently, the amount of lithium hydride present in the gas is predicted to be very small, even at very high densities, and so $\\lih$ cooling is no longer believed to be important \\citep{mon05}.  In contrast to $\\hd$ or $\\lih$, the various molecular ions present in the gas, such as $\\mHtp$, $\\htp$ or ${\\rm HeH^{+}}$, have attracted little attention. Some early work on $\\mHtp$ cooling in ionized primordial gas can be found in \\citet{ss77,ss78}, and its possible importance in hot, highly ionized conditions has recently been re-emphasized by \\citet{yokh07}, but aside from this, there has been little exploration of the role that cooling from these species might play in the evolution of primordial gas, presumably because the abundances of these species are expected to be small. It is this absence that the current paper attempts to rectify.  We present here the results of a set of simulations of the chemical and thermal evolution of gravitationally collapsing primordial gas. These simulations use a very simple one-zone dynamical model for the gas, but couple this with a detailed chemical network and a comprehensive model of the various heating and cooling processes at work. Besides the coolants considered above, we include the effects of cooling from (in no particular order) $\\mHtp$, $\\hdp$, $\\ddp$, $\\htp$, $\\hhdp$, $\\hddp$, $\\dtp$, $\\hehp$, $\\hedp$, $\\hehep$, $\\lihp$, $\\lidp$, $\\liD$ and $\\lihhp$. We focus in particular on the possible role of $\\htp$. This ion has a very large number of excited rotational and vibrational levels that are energetically accessible at the temperatures of interest in primordial gas, and its vibrational levels have much shorter radiative lifetimes than those of $\\mHt$ or $\\hd$. In LTE, its cooling rate per molecule is roughly $10^{9}$ times larger than that of $\\mHt$. It is known to be an important coolant in planetary atmospheres \\citep{miller00} and may also be an effective coolant in high density primordial gas \\citep{gs06}. Unlike ions such as $\\mHtp$ or ${\\rm HeH^{+}}$, it does not react with  $\\mHt$, and is not easily destroyed by collisions with $\\mH$, and so its  abundance in high density gas should be large compared to the other molecular ions included in our model.  The layout of this paper is as follows. In \\S2 we outline the numerical method used to simulate the thermal and chemical evolution of the protostellar gas. The chemical reactions included in the model are discussed in \\S3, and the thermal processes, in particular $\\htp$ cooling, are discussed in \\S4. We present the results of our simulations in \\S5 and close with a brief discussion in \\S6. ",
        "conclusions": "\\label{conc} We have examined the contribution that $\\htp$ cooling makes to the total cooling rate of gravitationally collapsing primordial gas in a wide range of different models using a newly-developed $\\htp$ cooling function along with the most detailed model of primordial gas chemistry published to date. Our results demonstrate that in general $\\htp$ cooling is not important, although it comes close to being so at densities $n = 10^{7}$--$10^{9} \\: {\\rm cm^{-3}}$, contributing at its peak a few percent of the total amount of cooling. We come to this conclusion despite making several assumptions (regarding the collapse rate of the gas, the formation rate of $\\htp$ by radiative association, and the collisional excitation rate of its excited vibrational states) that favour $\\htp$ cooling, and thus we have confidence that our conclusion is robust.  As $\\htp$ comes so close to being important, it is instructive to examine why it ultimately fails to dominate. This can be ascribed to a combination of two main effects. First, at high densities ($n \\simgreat 10^{8} \\: {\\rm cm^{-3}}$), the gas temperature becomes high enough to make the endothermic reaction TR17 \\begin{equation} \\htp + \\mH \\rightarrow \\mHtp + \\mHt \\end{equation} the most important $\\htp$ destruction mechanism, which significantly suppresses the $\\htp$ abundance at these densities. Second, the formation rate of $\\htp$ is strongly suppressed by the rapid removal of $\\Hp$ from the gas at densities $n \\simgreat 10^{11} \\: {\\rm cm^{-3}}$. At these densities, the fractional ionization of the gas is so low that the main loss route for the $\\Hp$ is conversion to $\\htp$, followed by $\\htp$ dissociative recombination, and once the timescale for $\\Hp$ removal via this combination of reactions becomes short compared to the free-fall time, the $\\Hp$ abundance decreases by orders of magnitude within a short space of time, effectively switching off the formation of $\\htp$. Moreover, destruction of $\\htp$ by dissociative recombination remains effective despite the fall-off in $x_{\\Hp}$ thanks to the contribution of electrons from ionized lithium, $\\lip$, which for $n > 3 \\times 10^{8} \\: {\\rm cm^{-3}}$ is the most abundant positive ion in the gas.  In our study, the only situation in which we found $\\htp$ cooling to be important is if the gas is illuminated by a strong flux of cosmic rays or X-rays. If the incident flux is strong enough to produce an ionization rate $\\simgreat 10^{-18} \\: {\\rm s^{-1}}$ at densities $n \\ge 10^{10} \\: {\\rm cm^{-3}}$, then the high density $\\htp$ abundance can be significantly increased, and $\\htp$ can even become the dominant coolant. However, the necessary flux of cosmic rays or hard X-rays is difficult to produce in the high-redshift Universe. As the estimates in \\S\\S~\\ref{res:cosmic} and \\ref{res:X} demonstrate, the required flux is orders of magnitude greater than the size of any plausible extragalactic background, and will only be achieved within gas that is very close to a local source (i.e., within 5--10~pc). However, gas that is this close to a supernova remnant or miniquasar will have been strongly affected by radiative feedback from the progenitor star and so is not a promising place to expect to find ongoing population III star formation.  Furthermore, even if the ionization rate is high enough to make $\\htp$  an important or dominant coolant  at high densities, the effect of $\\htp$ cooling on the thermal evolution of the gas remains relatively small; the difference it makes to the temperature evolution at $n > 10^{8} \\: {\\rm cm^{-3}}$ is smaller than the error introduced by the uncertainty in the three-body $\\mHt$ formation rate (reaction TB1).  Finally, our model has also allowed us to explore the effects of cooling from the other minor ionic and molecular species present in the gas (e.g., $\\mHtp$, ${\\rm H_{2}D^{+}}$, $\\lih$, etc.). Despite making rather optimistic assumptions regarding the cooling from these species, we find that they are orders of magnitude less effective than $\\htp$ at cooling high density gas, and hence are never significant."
    },
    "0809/0809.2785_arXiv.txt": {
        "abstract": "We study formation and evolution of bar-disk systems in fully self-consistent cosmological simulations of galaxy formation in the $\\Lambda$CDM WMAP3 Universe. In a representative model we find that the first generation of bars form in response to the asymmetric dark matter (DM) distribution (i.e., DM filament) and quickly decay. Subsequent bar generations form and are destroyed during the major merger epoch permeated by interactions with a DM substructure (subhalos). A long-lived bar is triggered by a tide from a subhalo and survives for $\\sim 10$~Gyr. The evolution of this bar is followed during the subsequent numerous minor mergers and interactions with the substructure. Together with intrinsic factors, these interactions largely determine the stellar bar evolution. The bar strength and its pattern speed anticorrelate, except during interactions and when the secondary (nuclear) bar is present. For about 5~Gyr bar pattern speed {\\it increases substantially} despite the loss of angular momentum to stars and {\\it cuspy} DM halo. We analyze the evolution of stellar populations in the bar-disk and relate them to the underlying dynamics. While the bar is made mainly of an intermediate age, $\\sim 5-6$~Gyr, disk stars at $z=0$, a secondary nuclear bar which surfaces at $z\\sim 0.1$ is made of younger, $\\sim 1-3$~Gyr stars. ",
        "introduction": "\\label{sec:intro}  Within the framework of structure formation in the universe, a hierarchy of dark matter (DM) masses form, while baryons assemble in their midst as galactic disks (e.g., White \\& Rees 1978). Disk evolution is expected to be influenced heavily by mergers and interactions in various ways. In this Letter, we focus on the specific issue of a tidal triggering of stellar bars embedded in a disk and surrounded by a DM halo with a substructure, within the context of cosmological evolution in the WMAP3 universe (see also Dubinski et al. 2008). We analyze the dynamical aspects of this evolution and follow the prevailing stellar populations in the bar-disk system.  Recent efforts to understand galaxy formation have been spearheaded by the work on pure DM halo formation (e.g., Diemand et al. 2007) and associated disk growth (e.g., Sommer-Larson et al. 2003; Governato et al. 2004, 2007; Heller et al. 2007a,b). Addition of baryonic component(s) to this modeling has met with difficulties, partly due to a numerical resolution, star formation (SF), energy feedback, and other processes. Dissipation is relatively slow in disks, but can be accelerated by any substantial departure from axial symmetry due to gravitational torques --- a nonlocal viscosity, surpassing the conventional viscosity by orders of magnitude. Torques facilitate the redistribution of mass, angular momentum and energy between the disk baryons and the slowly tumbling DM halo. Two main sources can serve as triggers of gravitational torques onto the disk --- an asymmetric mass distribution in the DM halo, and its non-uniformity, i.e., substructure, in DM and baryons. Both tidal interactions (e.g., Byrd et al. 1986; Noguchi 1987; Gerin et al. 1990) and DM halo asymmetry (Heller  et al. 2007a,b) are likely to trigger stellar bars. About 1/3 of the large-scale bars host additional nuclear bars (Laine et al. 2002; Erwin \\& Sparke 2002).  Observations and numerical simulations point to an intricate relation between the stellar populations and the underlying dynamics of bars (e.g., Martin \\& Friedli 1997; Knapen et al. 1995a,b; Jogee et al. 2002a,b; Zurita \\& Perez 2008). For bars in a cosmological context, this issue becomes even more relevant, because the disk is open to accretion of cold gas and energy input from interactions. Here we attempt to answer some questions about the long-term evolution of bar-disk populations and relate it to the dynamical history of the system.  So far, nearly all modeling of bar evolution has been limited to isolated systems in a stationary state, with the exception of simulations of tidal interactions in controlled experiments (e.g., Gauthier et al. 2006). We test the bar origin and evolution in the cosmological setting without fine-tuning and minimazing the prior assumptions.  Numerical simulations have been performed using the FTM-4.5 hybrid $N$-body/SPH code (e.g., Heller \\& Shlosman 1994; Heller et al. 2007b; Romano-Diaz et al. 2008a) using physical coordinates. The total number of DM particles is $2.2\\times 10^6$ and the SPH particles $4\\times 10^5$. The gravity is computed using the falcON routine (Dehnen 2002) which scales as O(N). The gravitational softening is 500~pc, for DM, stars and gas. We assume the $\\Lambda$CDM cosmology with WMAP3 parameters, $\\Omega_m=0.24$, $\\Omega_\\Lambda=0.76$ and $h=0.73$, where $h$ is the Hubble constant in units of $100~\\kmsmpc$. The variance $\\sigma_8=0.76$ of the density field convolved with the top hat window of radius $8h^{-1}$~Mpc is used to normalize the power spectrum. The SF algorithm is described in Heller et al. (2007b). We note only that multiple generations of stars can form from a single SPH particle. The energy and momentum feedback into the ISM are implemented, and the associated parameters have the following values (Heller et al.): energy thermalization $\\epsilon_{\\rm SF}=0.3$, cloud gravitational collapse $\\alpha_{\\rm ff}=1$, and self-gravity fudge-factor $\\alpha_{\\rm crit}=0.5$.  Initial conditions generated here are those of Romano-Diaz et al. (2008a): we use the Constrained Realizations (CR) method (Hoffman \\& Ribak 1991) within a restricted box size of $8h^{-1}$~Mpc and a sphere of $5h^{-1}$~Mpc is carved out and evolved. A large-scale filament runs across the computational sphere and is banana-shaped. The constructed Gaussian field is required to obey a set of constraints of arbitrary amplitudes and positions (more about this method in Romano-Diaz et al. 2006, 2007). Two constraints were imposed on the density field, first --- that the linear field Gaussian smoothed with a kernel of $1.0\\times 10^{12}~h^{-1}\\msun$, has an over-density of $\\delta=3$ at the origin ($2.5\\sigma$ perturbation, where $\\sigma^2$ is the variance of the appropriately smoothed field). It was imposed on a $256^3$ grid and predicted to collapse at $z_{\\rm c}\\sim 1.33$ based on the top-hat model. This perturbation is embedded in a region (2nd constraint) corresponding to a mass of $5\\times 10^{13}~h^{-1}\\msun$ in which the over-density is zero, i.e., the unperturbed universe. The random component in CRs favors formation of similar structures which leads to major mergers. The DM mass inside the computational sphere is $\\sim 6.1\\times 10^{12}~h^{-1}\\msun$. We have randomly replaced 1/6 of DM particles by equal mass SPH particles. Therefore, $\\Omega_{\\rm m}$ is not affected. ",
        "conclusions": "We have studied the evolution of galactic bars in fully self-consistent cosmological simulations of galaxy formation. In a representative model, we find that the first generation of bars forms early and in response to the asymmetric background DM distribution. This bar decays quickly and subsequent bars form and are destroyed by interactions with subhalos. Finally, a long-lived bar is triggered that survives for $\\sim 10$~Gyr. Most interestingly, its pattern speed {\\it increases substantially} over the first $\\sim 5$~Gyr, despite the angular momentum loss to the outer disk and the cuspy DM halo. The bar evolution appears to be closely linked to the disk evolution --- it is permeated by interactions with subhalos as we discuss below.  The sequence of bars in the model are repeatedly triggered by interactions with DM substructure and do not form as a result of a `classical' (i.e., intrinsic) bar instability. The number of subhalos within the inner prime halo is highly variable --- the subhalos cluster in the filament and are frequently accreted in groups before they merge (Romano-Diaz et al. 2008b). This amplifies the damage inflicted by the substructure on the disk-bar system because of higher mass and energy influx into the disk region. We find that the prograde encounters trigger bars, direct hits weaken existing bars, while retrograde encounters have little effect --- in line with Gerin et al. (1990) and Berentzen et al. (2004). In spite of this, addition of the cold gas to the disk, up to $R\\sim 15$~kpc, in tandem with growing stellar surface density by $\\sim 10$, allow the bar to speed up over $\\sim 5$~Gyr (Fig.~2). This happens because the buildup of a disk mass and the increased central mass concentration in the galaxy amplifies the precession rate of bar orbits. On the other hand, the weakening of the bar during this time results from the angular momentum added by the gas to the disk and, therefore, to the bar. This happens because orbits trapped within the bar become less eccentric, leading to a weaker bar.  The gas fraction in the disk exceeds that considered so far in the literature by as much as a factor of 10 (Bournaud et al. 2005; Berentzen et al. 2007). It is an influx of subhalos into the center which strips the disk of its gas and heats up the stellar `fluid' injecting it above the plane. Being cut off from the source of angular momentum, the bar brakes rapidly.  An important caveat of this slowdown is the bar weakening after $z\\sim 0.1$ (Fig.~2). Here we do not observe the anticorrelation between $A_2$ and $\\Omega_{\\rm b}$. The appearance of the nuclear bar sheds light on this behavior. The orbit trapping by the nuclear bar which grows due to a slowdown deprives the main bar of orbital support --- the weakening of the main bar and the strengthening of the nuclear bar proceed in tandem.  Stellar populations provide a further insight into the dynamics of the bar-disk systems. When viewed in colors of a younger stellar population, the disk is thinner and shows less contribution to the spheroidal component. Comparison of the left frames in Fig.~4 also shows that the spheroidal component is populated by stars older than the average disk population, as it disappears in the lower frame.  \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.43]{fig5.ps} \\end{center} \\caption{SF rates shown within the inner 10~kpc of the disk. } \\end{figure}  An interesting feature is present in $z=0$ frames of Fig.~4 with stars younger than 4~Gyr in the center --- the upper/lower surfaces are flat and parallel to the disk midplane. It is associated with the vertical oscillations of stars in the region, and closely resembles the boxy bulges observed in about 50\\% of edge-on disk galaxies (e.g., Lutticke et al. 2000). These bulges originate in trapping by a family of 3-D orbits belonging to the prime bar (e.g., Combes et al. 1990; Skokos et al. 2002; Martinez-Valpuesta et al. 2006).  The SF continues after the major merger epoch, $z\\ltorder 1.5$ (Fig.~5). The asymptotic value for the SF rate, SFR$\\sim 1~{\\rm \\msun~yr^{-1}}$, in the disk was not fine-tuned but results from self-regulation. Early SF is supported by the cold gas accretion onto the disk. In the later stage, the SF is limited to the central kpc only --- this explains the stellar age gradient in the disk which is positive (Figs.~3, 4). The bar resembles those in the early type galaxies with a constant surface brightness along the bar major axis, unlike the exponential bars which show negative age gradients (Perez et al. 2007).  In summary, we have followed the bar evolution in the most realistic environment attempted so far. We have shown that bars can maintain their fast rotation for $\\sim 5$~Gyr despite the loss of angular momentum to the disk and {\\it cuspy} DM halo. Finally, bar dynamics appears to be closely linked to the disk stellar populations."
    },
    "0809/0809.2266_arXiv.txt": {
        "abstract": "A new generation of radio telescopes is achieving unprecedented levels of sensitivity and resolution, as well as increased agility and field-of-view, by employing high-performance digital signal processing hardware to phase and correlate large numbers of antennas.  The computational demands of these imaging systems scale in proportion to $BMN^2$, where $B$ is the signal bandwidth, $M$ is the number of independent beams, and $N$ is the number of antennas.  The specifications of many new arrays lead to demands in excess of tens of PetaOps per second.  To meet this challenge, we have developed a general purpose correlator architecture using standard 10-Gbit Ethernet switches to pass data between flexible hardware modules containing Field Programmable Gate Array (FPGA) chips.  These chips are programmed using open-source signal processing libraries we have developed to be flexible, scalable, and chip-independent. This work reduces the time and cost of implementing a wide range of signal processing systems, with correlators foremost among them, and facilitates upgrading to new generations of processing technology. We present several correlator deployments, including a 16-antenna, 200-MHz bandwidth, 4-bit, full Stokes parameter application deployed on the Precision Array for Probing the Epoch of Reionization. ",
        "introduction": "\\label{sec:intro}  Radio interferometers, which operate by correlating the signals from two or more antennas, have many advantages over traditional single-dish telescopes, including greater scalability, independent control of aperture size and collecting area, and self-calibration.  Since the first digital correlator built by Weinreb \\citep{weinreb_1961}, the processing power of these systems has been tracking the Moore's Law growth of digital electronics. The decreasing cost per performance of these systems has influenced the design of many new radio antenna array telescopes.  Some next-generation array telescopes at meter, centimeter and millimeter wavelengths are: the LOw Frequency ARray (LOFAR), the Precision Array for Probing the Epoch of Reionization (PAPER), the Murchison Widefield Array (MWA), the Long Wavelength Array (LWA), the Expanded Very Large Array (EVLA), the Allen Telescope Array (ATA), the Karoo Array Telescope (MeerKAT), the Australian Square Kilometer Array Demonstrator (ASKAP), the Atacama Large Millimeter Array (ALMA). and the Combined Array for Research Millimeter-wave Astronomy (CARMA). This paper presents a novel approach to the intense digital signal processing requirements of these instruments that has many other applications to astronomy signal processing.  While each generation of electronics has brought new commodity data processing solutions, the need for high-bandwidth communication between processing nodes has historically lead to specialized system designs.  This communication problem is particularly germane for correlators, where the number of connections between nodes scales with the square of the number of antennas. Solutions to date have typically consisted of specialized processing boards communicating over custom backplanes using non-standard protocols.  However, such solutions have the disadvantage that each new generation of digital electronics requires expensive and time-consuming investments of engineering time to re-solve the same connectivity problem.  Redesign is driven by the same Moore's Law that makes digital interferometry attractive, and is not unique to the interconnect problem; processors such as Application-Specific Integrated Circuits (ASICs) and Field Programmable Gate Arrays (FPGAs) also require redesign, as do the boards bearing them, and the signal processing algorithms targeting their architectures.  Our research is aimed at reducing the time and cost of correlator design and implementation.  We do this, firstly, by developing a packetized communication architecture relying on industry-standard Ethernet switches and protocols to avoid redesigning backplanes, connectors, and communication protocols. Secondly, we develop flexible processing modules that allow identical boards to be used for a multitude of different processing tasks.  These boards are applicable to general signal processing problems that go beyond correlators and even radio science to include, e.g., ASIC design and simulation, genomics, and research into parallel processor architectures. General purpose hardware reduces the number of boards that have to be redesigned and tested with each new generation of electronics.  Thirdly, we create parametrized signal processing libraries that can easily be recompiled and scaled for each generation of processor.  This allows signal processing systems to quickly take advantage of the capabilities of new hardware.  Finally, we employ an extension of a Linux kernel to interface between CPUs and FPGAs for the purposes of testing and control, presenting a standard file interface for interacting with FPGA hardware.  This paper begins with a presentation of the new correlator design architecture in \\S\\ref{sec:architecture}. The hardware to implement this architecture follows in \\S\\ref{sec:hardware}, and the FPGA gateware used in the hardware is summarized in \\S\\ref{sec:gateware}. Issues concerning system integration are given in \\S\\ref{sec:integration}, and performance characterization of subsystems are given in \\S\\ref{sec:characterization}. Results from our first deployments of the packetized correlator are displayed in \\S\\ref{sec:deployments}. Our final section summarizes our progress and points to a number of directions we are pursuing for the next generation of scalable correlators based on modular hardware, reuseable gateware and data packetization. An appendix gives a glossary of technical acronyms since this paper makes heavy use of abbreviated terms. ",
        "conclusions": "\\label{sec:conclusion}  By decreasing the time and engineering costs of building and upgrading correlators, we aim to reduce the total cost of correlators for a wide range of scales.  Small- and medium-scale correlators with total cost dominated by development clearly stand to benefit from our research.  It is less clear if the cost of large-scale correlators can be reduced by the general-purpose hardware used in our architecture.  Though minimization of replication cost favors the development of specialized parts, there are two factors that can make a generic, modular solution cost less.  The first factor to consider is time to deployment.  Even if the monetary cost of development is negligible in the budget of a large correlator, the cost of development time can be significant.  If a custom solution takes several years to go from design to implementation, the hardware that is deployed will be out of date.  Moore's Law suggests that when a custom solution taking 3 years to develop is deployed, there will exist processors 4 times more powerful, or 4 times less expensive for the equivalent system.  The cost of a generic, modular system has to be tempered by the expected savings of committing to hardware closer to the ultimate deployment date.  The second factor is the cost of upgrade.  Many facilities (including the ATA) are beginning to appreciate the advantages of designing arrays with wider bandwidths and larger numbers of antennas than can be handled by current technology.  Correlators may then be implemented inexpensively on scales suited to current processors, and upgraded as more powerful processors become available.  Modular solutions facilitate this methodology."
    },
    "0809/0809.2099_arXiv.txt": {
        "abstract": "Scattering of Ly$\\alpha$ photons by neutral hydrogen gas in a single outflowing 'supershell' around star forming regions often explains the shape and offset of the observed Ly$\\alpha$ emission line from galaxies. We compute the radiation pressure that is exerted by this scattered Ly$\\alpha$ radiation on the outflowing material. We show that for reasonable physical parameters, Ly$\\alpha$ radiation pressure alone can accelerate supershells to velocities in the range $v_{\\rm sh}=200$--$400$ km s$^{-1}$. These supershells possibly escape from the gravitational potential well of their host galaxies and contribute to the enrichment of the intergalactic medium. We compute the physical properties of expanding supershells that are likely to be present in a sample of known high-redshift ($z=2.7$--$5.0$) galaxies, under the assumption that they are driven predominantly by Ly$\\alpha$ radiation pressure. We predict ranges of radii $r_{\\rm sh}=0.1$--$10$ kpc, ages $t_{\\rm sh }=1$--$100$ Myr, and energies $E_{\\rm sh}=10^{53}$--$10^{55}$ ergs, which are in reasonable agreement with the properties of local galactic supershells. Furthermore, we find that the radius, $r_{\\rm sh}$, of a Ly$\\alpha$-driven supershell of constant mass  depends uniquely on the intrinsic Ly$\\alpha$ luminosity of the galaxy, $L_{\\alpha}$, the HI column density of the supershell, $N_{\\rm HI}$, and the shell speed, $v_{\\rm sh}$, through the scaling relation $r_{\\rm sh}\\propto L_{\\alpha}/(N_{\\rm HI}v_{\\rm sh}^2)$. We derive mass outflow rates in supershells that reach $\\sim 10$--$100\\%$ of the star formation rates of their host galaxies. ",
        "introduction": "\\label{sec:intro}  The Ly$\\alpha$ emission line of galaxies is often redshifted relative to metal absorption lines \\citep[e.g.][]{Pettini01,Shapley03}. Scattering of Ly$\\alpha$ photons by neutral hydrogen atoms in an outflowing 'supershell' surrounding the star forming regions can naturally explain this observation, as well as the typically observed asymmetric spectral shape of the Ly$\\alpha$ emission line \\citep{L95,Lee98,T99,Ahn02,Ahn03,Mass03,Ahn04,Verhamme06,Verhamme08}. The presence of an outflow may also be required to avoid complete destruction\\footnote{ Ly$\\alpha$ may also avoid complete destruction by interstellar dust when it is primarily locked up in cold clumps embedded in a hot medium that is transparent to Ly$\\alpha$ \\citep{N91,H06}.} of the Ly$\\alpha$ radiation by dust and to allow its escape from the host galaxies \\citep[][]{Kunth98,Hayes08,Ostlin08,Atek08}.  Indeed, the existence of outflowing thin shells (with a thickness much smaller than their radius) of neutral atomic hydrogen around HII regions is confirmed by HI-observations of our own \\citep[][]{Heiles84} and other nearby galaxies \\citep[e.g.][]{Ryder95}. The largest of these shells, so-called 'supershells', have radii of $r_{\\rm max}\\sim 1$ kpc \\citep[e.g][]{Ryder95,M02,M06} and HI column densities in the range $N_{\\rm HI}\\sim 10^{19}$--$10^{21}$ cm$^{-2}$ \\citep[e.g.][]{L95,Kunth98,Ahn04,Verhamme08}. The existence of larger extragalactic supershells has been inferred from spatially extended ($\\sim$ a few kpc) Ly$\\alpha$ P-cygni profiles, that were observed around local star burst galaxies \\citep{Mass03}. These 'supershells' differ from more energetic 'superwinds' \\citep[e.g.][]{Heckman90,Martin05}, in which galactic-scale biconical outflows break-out of the galaxies' interstellar medium with high velocities (up to $\\sim 10^3$ km s$^{-1}$, of which M82 is a classical example). Supershells are thought to be generated by stellar winds or supernovae explosions which sweep-up gas into a thin expanding neutral shell \\citep[see e.g.][for a review]{T88}. The back-scattering mechanism attributes both the redshift and asymmetry of the Ly$\\alpha$ line to the Doppler boost that Ly$\\alpha$ photons undergo as they scatter off the outflow on the far side of the galaxy back towards the observer.  In this paper we show that the radiation pressure exerted by backscattered Ly$\\alpha$ radiation on the outflowing gas shell can accelerate it to velocities that reach hundreds of km s$^{-1}$. Our basic result can be illustrated through an order-of-magnitude estimate. The net outward force exerted on a shell of gas\\footnote{The net momentum transfer rate per photon that enters the shell is $\\Delta {\\bf p}=\\frac{h\\nu_{\\alpha}}{c}(1-\\mu)${\\bf e}$_r$, where {\\bf e}$_r$ denotes a unit vector that is pointing radially outward. Furthermore, $\\mu={\\bf k}_{\\rm in}\\cdot {\\bf k}_{\\rm out}$, where ${\\bf k}_{\\rm in}$ (${\\bf k}_{\\rm out}$) denotes the propagation direction of the Ly$\\alpha$ photon as it enters (leaves) the shell. When averaged over all directions, the total momentum transfer rate is $\\dot{N}_{\\alpha}\\frac{h\\nu_{\\alpha}}{c}${\\bf e}$_r$, where $\\dot{N}_{\\alpha}$ is the rate at which Ly$\\alpha$ photons enter the shell. This total net momentum transfer rate applies regardless of the number of times each Ly$\\alpha$ photon scatters inside the shell, or regardless of whether the Ly$\\alpha$ photon is absorbed by a dust grain (and possibly re-emitted in the infrared).} by backscattered Ly$\\alpha$ radiation of luminosity $L_{{\\rm BS},\\alpha}$ is ${d(m_{\\rm sh}v_{\\rm sh})}/{dt}=L_{{\\rm BS},\\alpha}/c$. Here, $m_{\\rm sh}$ is the total mass in the shell, $v_{\\rm sh}$ is the shell velocity and $c$ is the speed of light. If we assume that backscattering occurs over a timescale $t_{\\rm BS}$, and the shell mass is constant in time, then the net velocity gain by the shell is $\\Delta v_{\\rm sh}\\sim 250$ km s$^{-1}$ $(m_{\\rm sh}/10^7 \\hs M_{\\odot})^{-1}(t_{\\rm BS}/50 \\hs{\\rm Myr})(L_{{\\rm BS},\\alpha}/10^{43} \\hs {\\rm erg/s})$. The adopted $m_{\\rm sh}$ value corresponds to a thin spherical HI shell with a column density $N_{\\rm HI}=10^{20}$ cm$^{-2}$ and radius $r=1$ kpc.  The fiducial values of $L_{{\\rm BS},\\alpha}$ and $t_{\\rm BS}$ were chosen to yield the typical observed shell speed and expected shell lifetime in Lyman-break galaxies \\citep[LBGs; see e.g.][]{Shapley03,Verhamme08}. Here, and throughout this paper, we assume that the supershells are mostly neutral. This assumption is based on 21-cm observations of local supershells which indicate that these are very thin (thickness $\\ll$ their radius, e.g Heiles 1984) and therefore dense.  At these high densities one expects free electrons and protons to recombine efficiently and make the gas significantly neutral.  However, \\citet{P00} argue that, based on the observed reddening of the UV continuum spectrum, either the dust-to-gas ratio in the supershell of LBG MS1512-cB58 is higher than the galactic value, or alternatively that the HI only makes up $f_{\\rm HI}\\sim 14-33\\%$ of the total gas mass of the shell, with the remainder of the gas being molecular or ionized. We discuss how this latter possibility affects our results in \\S~\\ref{sec:conc}.  This paper is a closely related to another paper in which detailed Monte-Carlo Ly$\\alpha$ radiative transfer calculations are performed to assess Ly$\\alpha$ radiation pressure in a more general context (Dijkstra \\& Loeb 2008, hereafter Paper I). These radiative transfer calculations illustrate clearly that Ly$\\alpha$ radiation pressure may be important in driving outflows in various environments. Importantly, the pressure exerted by Ly$\\alpha$ radiation alone can exceed the maximum possible pressure exerted by continuum radiation through dust opacity (see Paper I), which in turn can exceed the total kinetic pressure exerted by supernova ejecta \\citep{Murray05}. This implies that Ly$\\alpha$ radiation pressure may thus in some cases provide the dominant source of pressure on neutral hydrogen gas in the interstellar medium (ISM).  The goals of this paper are ({\\it i}) to gauge the general importance of Ly$\\alpha$ radiation  in the particular context of outflowing galactic supershells, and ({\\it ii}) to explore the related observable properties of Ly$\\alpha$-driven supershells. In \\S~\\ref{sec:result} we compute the time evolution of the velocity of an expanding shell of neutral gas that encloses a central Ly$\\alpha$ source. We consider a range of models, which are subsequently applied to known Ly$\\alpha$ emitting galaxies at high redshifts. Finally, we discuss our results and their implications in \\S~\\ref{sec:conc}. The parameters for the background cosmology used throughout our discussion are $(\\Omega_m,\\Omega_{\\Lambda},\\Omega_b,h)=(0.27,0.73,0.042,0.70)$ \\citep{Dunkley08,Komatsu08}. ",
        "conclusions": "\\label{sec:conc} It is well known that scattering of Ly$\\alpha$ photons by neutral hydrogen in a single outflowing supershell around star-forming galaxies can naturally explain two observed phenomena: ({\\it i}) the common shift of the Ly$\\alpha$ emission line towards the red relative to other nebular recombination and metal absorption lines, and ({\\it ii}) the asymmetry of the Ly$\\alpha$ line with emission extending well into its red wing.  In this paper we have computed the radiation pressure that is exerted by this scattered Ly$\\alpha$ radiation on the outflowing shell.  We have shown that for reasonable shell parameters the shell can be accelerated to a few hundred km s$^{-1}$ after $\\lsim 10$ Myr.  The Ly$\\alpha$ acceleration mechanism becomes increasingly inefficient as the shell speed increases (Fig.~\\ref{fig:fw}), because the Doppler boost of the gas extends increasingly farther into the wing of the Ly$\\alpha$ absorption profile allowing a decreasing fraction of the Ly$\\alpha$ photons to scatter on the outflowing shell. Furthermore, for a shell mass that is constant in time, $N_{\\rm HI}\\propto r^{-2}$, and we find that for any choice of model parameters the shell achieves a terminal velocity that depends mainly\\footnote{An alternative way to see why the observed quantities $N_{\\rm HI}$, $v_{\\rm sh}$, and $L_{\\alpha}$ provide unique constraints on the shell mass, radius, and age is that these three unknown quantities relate to the observed quantities via three equations: {\\bf 1. $m_{\\rm sh}=4\\pi r^2_{\\rm sh}N_{\\rm HI}\\mu m_p$} and {\\bf 2.} $v_{\\rm sh}\\sim \\frac{L_\\alpha}{cm_{\\rm sh}}t_{\\rm sh}$, and {\\bf 3. $r_{\\rm sh}\\sim \\frac{L_\\alpha}{2cm_{\\rm sh}}t^2_{\\rm sh}$}. We found that Equations {\\bf 2.} and {\\bf 3.} are accurate to within $\\sim 50\\%$.} on its total mass and the Ly$\\alpha$ luminosity of the central source (Fig.~\\ref{fig:speed}).  We have computed the physical properties of expanding supershells that are likely to be present in specific observed high-redshift ($z=2.7$--$5.0$) galaxies, under the assumptions that they are driven predominantly by Ly$\\alpha$ radiation pressure and that their mass remains constant in time. We predict supershell radii that lie in the range $r_{\\rm sh}=0.1$--$10$ kpc, ages in the range $t_{\\rm sh }=1$--$100$ Myr, and energies in the range $E_{\\rm sh}=10^{53}$--$10^{55}$ ergs, in broad agreement with the properties of local galactic supershells. We derive mass outflow rates in the supershells that reach $\\sim 10$--$100\\%$ of the star formation rates in the host galaxies, in agreement with the {\\it total} outflow rates that one expects theoretically \\citep[e.g.][]{Erb08}.  We have found that in all models the radius of the supershell is determined uniquely by parameters such as the {\\it intrinsic} Ly$\\alpha$ luminosity of the host galaxy (which may be inferred from the observed Ly$\\alpha$ line shape), $L_{\\alpha}$, the total column density of HI in the supershell, $N_{\\rm HI}$, and the shell velocity $v_{\\rm sh}$, \\begin{equation} \\left(\\frac{r_{\\rm sh}}{{\\rm kpc}}\\right)\\sim 0.2\\Big{(}\\frac{f}{0.9}\\Big{)}\\Big{(}\\frac{L_{\\alpha}}{10^{43}\\hs \\frac{{\\rm erg}}{{\\rm s}}}\\Big{)}\\Big{(}\\frac{10^{20}/{\\rm cm}^{2}}{N_{\\rm HI}}\\Big{)}\\Big{(}\\frac{200\\hs\\frac{{\\rm km}}{{\\rm s}}}{v_{\\rm sh}}\\Big{)}^{2}. \\label{eq:rshell} \\end{equation} Here $f$ is a fudge factor that lies in the range $f=1.0$--$1.3$. Our models have ignored gravity and the pressure of exerted by the surrounding interstellar medium. {\\it This relation holds up only for supershells that are driven predominantly by Ly$\\alpha$ radiation pressure}. This relation is most accurate for those galaxies in which the shells are accelerated to velocities that exceed $v_{\\rm esc}$. We found this to be the case in 5 out of 9 (55\\%, see Table~\\ref{table:predict}, ignoring the galaxies for which our model was ruled out, i.e. FDF 1267b, FDF 5215, FDF4691, and FDF 7539) of known galaxies. Hence, these shells could contribute to the enrichment of the intergalactic medium.  In other cases, gravity (and external pressure) will tend to reduce the value of the fudge factor $f$. Also, if neutral hydrogen gas only makes up a fraction $f_{\\rm HI}$ of the total mass of the shell (see \\S~\\ref{sec:intro}), then our predicted value of $r_{\\rm sh}$ should be lowered by a factor of $f_{\\rm HI}$.  On the other hand, as was already mentioned in \\S~\\ref{sec:result}, our conservative calculations have completely ignored the fact that 'trapping' of Ly$\\alpha$ radiation by the optically thick supershell may significantly boost the Ly$\\alpha$ radiation pressure compared to what was used in this paper. Specifically, it was shown in Paper I that when this trapping is properly accounted for, the total radiation force in Eq.~\\ref{eq:diff} should include $M_{F}(v_{\\rm sh},N_{\\rm HI})$ in place of $f_{\\rm scat}(v_{\\rm sh},N_{\\rm HI})$. Here, $M_F$ is 'force-multiplication' factor\\footnote{This term derives from the (time-dependent) force-multiplication function $M(t)$ that was introduced by \\citet{Castor75}, as $F_{\\rm rad}\\equiv M(t)\\frac{\\tau_eL_{\\rm bol}}{c}$. Here, $F_{\\rm rad}$ is the total force that radiation exerts on a medium, $\\tau_e$ is the total optical depth to electron scattering through this medium. The function $M(t)$ arises because of the contribution of numerous metal absorption lines to the medium's opacity, and can be as large as $M_{\\rm max}(t)\\sim 10^3$ in the atmospheres of O-stars \\citep{Castor75}.} that depends both the column density of HI and the velocity of the supershell. In Paper I, $M_F$ was computed for a range of $N_{\\rm HI}$ and $v_{\\rm sh}$ values, by performing Monte-Carlo Ly$\\alpha$ radiative transfer calculation for each combination of $N_{\\rm HI}$ and $v_{\\rm sh}$. In these calculations, the HI shell surrounds an empty cavity. The results of these calculations are shown in Fig.~\\ref{fig:fscat}. As evident from the plot, $M_F$ greatly exceeds unity for low shell velocities and large HI column densities. Overplotted as the {\\it grey line} is the trajectory in the $N_{\\rm HI}-v_{\\rm sh}$ plane of the supershell in cB-58 (\\S~\\ref{sec:cB-58}). At early times, $M_F> 10$, and we may have underestimated the radiation pressure force significantly (as was mentioned in \\S~\\ref{sec:result}, the actual force multiplication factor may be less than $M_F$, if gas and dust interior to the supershell prevents photons from repeatedly 'bouncing' back an forth between opposite sides of the HI shell). Ly$\\alpha$ trapping could boost our predicted value of $r_{\\rm sh}$ by a factor as large as $\\langle M_F \\rangle$, where $\\langle M_F \\rangle$ the value of $M_F$, averaged over the expansion history of the shell.  \\begin{figure} \\vbox{\\centerline{\\epsfig{file=fig3.eps,angle=270,width=8.3truecm}}} \\caption[]{When Ly$\\alpha$ trapping is properly accounted for, the total radiation force in Eq.~\\ref{eq:diff} should be modified as $f_{\\rm scat}(v_{\\rm sh},N_{\\rm HI}) \\rightarrow M_{F}(v_{\\rm sh},N_{\\rm HI})$. Here, we plot $M_F$ (a 'force-multiplication' factor) as a function of the expansion velocity of the HI shell, $v_{\\rm sh}$, for three values of HI column density. $M_F$ was calculated by performing Monte-Carlo Ly$\\alpha$ radiative transfer calculation (described in Paper I), and under the assumption that the HI shell surrounds an empty cavity. $M_F$ increases with increasing $N_{\\rm HI}$ and decreasing $v_{\\rm sh}$. The {\\it grey line} is the trajectory in the $N_{\\rm HI}-v_{\\rm sh}$ plane of the model of the supershell in cB-58 (\\S~\\ref{sec:cB-58}). Clearly, especially at early times we have underestimated the radiation pressure force significantly, which renders our calculations conservative.} \\label{fig:fscat} \\end{figure}  So far, we have assumed HI to be the only source of opacity in the supershell. However, supershells may also contain dust which provides an additional source of opacity; for example, in the models of Verhamme et al (2008), the optical depth through dust at $\\lambda=1216$ \\AA\\hs is $\\tau_{\\rm D}=0.0$--$2.0$. At a given HI column density, we have therefore systematically underestimated $f_{\\rm scat}$, (absorption of a Ly$\\alpha$ photon by dust also results in a momentum transfer of magnitude $h\\nu_{\\alpha}/c$), which renders our calculations conservative. Furthermore, dusty supershells could also absorb significant amounts of continuum radiation, boosting the total momentum transfer rate even further. If the dust resides interior to the HI shell, then the Ly$\\alpha$ flux that impinges upon the shell is reduced by e$^{-\\tau_d(\\lambda_{\\alpha})}$. In this case, the Ly$\\alpha$ (as well as the continuum) radiation pressure is smaller than computed in this paper. However, the coupling between gas and dust makes it very unlikely that (Ly$\\alpha$) radiation pressure sweeps up the gas surrounding a star forming region, while keeping the dust in place \\citep[see][for more discussion on this]{Murray05}.  Of course, the existence of dust inside the HI shell reduces its ability to 'trap' Ly$\\alpha$ photons, which reduces the value of $M_F$ compared to that shown in Figure~\\ref{fig:fscat}. Despite this reduction, $M_F$ can still greatly exceed unity, because at low shell expansion velocities the majority of Ly$\\alpha$ photons do no penetrate deeply into the HI shell upon their first entry. Instead, the photons mostly scatter near the surface of the shell, and only 'see' a fraction of the dust opacity (indeed, for this reason Ly$\\alpha$ radiation can escape from a dusty two-phased ISM, see Neufeld 1991 and Hansen \\& Oh 2006). In other words, a relative small inner portion the HI shell can effectively trap Ly$\\alpha$ photons. We have verified this statement with Monte-Carlo radiative transfer calculations that included dust at the levels inferred by \\citet{Verhamme08}: we typically found $M_F$ to be reduced by a factor of up to a few. In other words, even for dusty shells it is possible that $M_F\\gg 1$.  Therefore, the pressure exerted by Ly$\\alpha$ radiation alone can exceed the maximum possible pressure exerted by continuum radiation (also see Paper I), which can drive the large gas masses (as large as a few $10^{10-11}M_{\\odot}$) out of galaxies in a galactic {\\it superwind} \\citep{Murray05}. For comparison, in our model the Ly$\\alpha$ photons transfer their momentum onto the expanding supershell, which contains significantly less mass ($m_{\\rm sh}\\lsim 10^8 M_{\\odot}$). This implies that in principle, both mechanisms could operate simultaneously, and that the neutral HI supershells may be accelerated to higher velocities than the bulk of the ejected gas. Since in our model Ly$\\alpha$ radiation pressure operates only on a fraction of the gas, it does not provide a self-regulating mechanism for star formation and black hole growth as in \\citet{Murray05}. However, it does provide a new way of enriching the intergalactic medium with metals.  We note that quasars can have Ly$\\alpha$ luminosities\\footnote{Empirically, the Ly$\\alpha$ luminosity of quasars is related to their B-band luminosity as $L_{\\alpha}\\sim 0.7L_{\\rm B}$ (Dijkstra \\& Wyithe, 2006). High-redshift quasars have been observed with $L_{\\rm B}\\gsim 10^{13}L_{{\\rm B},\\odot}=4\\times 10^{45}$ erg s$^{-1}$ (see e.g Figure~1 of Dijkstra \\& Wyithe, 2006).} that reach $L_{\\alpha}=10^{46}$ erg s$^{-1}$ \\citep[][]{Fan06}. In principle, this Ly$\\alpha$ luminosity could transfer a significant amount of momentum onto neutral gas in its proximity. Because the Ly$\\alpha$ emission line of quasars is typically very broad ($\\sim$ few thousand km s$^{-1}$), $f_{\\rm scat}$ is significantly less than unity for the column densities considered in the paper (see Fig.~\\ref{fig:fw}). Furthermore, it is unclear whether gas clouds can remain neutral in close proximity to the quasar for an extended period of time. Even if this is the case, it remains to be determined whether the Ly$\\alpha$ radiation pressure can compete with the continuum radiation pressure. The importance of Ly$\\alpha$ radiation pressure near quasars therefore remains an open issue.  To conclude, observations indicate that it is the kinematics of HI gas surrounding star forming regions that mostly determines the observed properties of the Ly$\\alpha$ radiation \\citep[as opposed to dust content, e.g.][]{Kunth98,Atek08,Hayes08,Ostlin08}. This work suggests that -at least in some cases- the pressure exerted by the Ly$\\alpha$ photons themselves may be important in determining the kinematics of HI gas. This is appealing for two reasons: ({\\it i}) the mechanism operates irrespective of the dust content of the HI supershells. This mechanism may therefore operate also in galaxies of primordial composition at high redshift (especially since the total Ly$\\alpha$ luminosity per unit star formation rate is higher, e.g. Schaerer 2003), and ({\\it ii}) the shape and velocity offset of the observed Ly$\\alpha$ emission strongly suggest that momentum transfer from Ly$\\alpha$ photons to HI gas is actually observed to occur.  {\\bf Acknowledgments} This work is supported by in part by NASA grant NNX08AL43G, by FQXi, and by Harvard University funds. We thank an anonymous referee for constructive comments that improved the presentation of this paper."
    },
    "0809/0809.0864_arXiv.txt": {
        "abstract": "{We have performed a detailed systematic search for multiperiodicity in the Population~I Cepheids of the Large Magellanic Cloud. In this process we have identified for the first time several new types of Cepheid pulsational behaviour. We have found two triple-mode Cepheids pulsating simultaneously in the first three radial overtones. In 9\\% of the first overtone Cepheids we have detected weak, but well resolved secondary periodicities. They appear either very close to the primary pulsation frequency or at a much higher frequency with a characteristic period ratio of $0.60-0.64$. In either case, the secondary periodicities must correspond to nonradial modes of oscillation. This result presents a major challenge to the theory of stellar pulsations, which predicts that such modes should not be exited in Cepheid variables. Nonradial modes have also been found in three of the fundamental/first overtone double-mode Cepheids, but no such oscillations have been detected in single mode Cepheids pulsating in the fundamental mode.  In 19\\% of double-mode Cepheids pulsating in the first two radial overtones (FO/SO type) we have detected a Blazhko-type periodic modulation of amplitudes and phases. Both modes are modulated with a common period, which is always longer than 700\\thinspace days. Variations of the two amplitudes are anticorrelated and maximum of one amplitude always coincides with minimum of the other. We have compared observations of modulated FO/SO Cepheids with predictions of theoretical models of the Blazhko effect, showing that currently most popular models cannot account for  properties of these stars. We propose that Blazhko effect in FO/SO Cepheids can be explained by a nonstationary resonant interaction of one of the radial modes with another, perhaps nonradial, mode of oscillations. } {Stars: variables: Cepheids -- stars: oscillations: nonradial -- stars: oscillations: multimode -- stars: oscillations: Blazhko effect -- galaxies: Magellanic Clouds} ",
        "introduction": "%  Classical Cepheids and RR~Lyrae stars are among the best studied pulsating variables. For decades they have been considered primary example of purely radial pulsators, displaying large amplitude oscillations, with only one or two radial modes excited. This simple picture was challenged in the end of the last century when Olech {\\it et~al.}\\thinspace (1999) discovered low amplitude nonradial modes in three RR~Lyr stars of globular cluster M55. In following years it was shown that presence of secondary periodicities identified with nonradial modes of oscillations is a very common property of RR~Lyr variables. They were detected in a substantial fraction of stars in every studied RR~Lyr population: in LMC (Alcock {\\it et~al.}\\thinspace 2000, 2003; Soszy\\'nski {\\it et~al.}\\thinspace 2003; Nagy \\& Kov\\'acs 2006), SMC (Soszy\\'nski {\\it et~al.}\\thinspace 2002), Galactic Bulge (Moskalik \\& Poretti 2002, 2003; Mizerski 2003; Collinge {\\it et~al.}\\thinspace 2006), Galactic field (Wils {\\it et~al.}\\thinspace 2006; Szczygie{\\l} \\& Fabrycky 2007, see also Smith 1995) and most recently in $\\omega$~Cen (Olech \\& Moskalik 2008). Depending on the stellar system, nonradial modes are excited in between 5\\% and 30\\% of all RR~Lyr stars.  Motivated by these findings we decided to search for nonradial modes in classical (Population~I) Cepheids. As far as pulsation properties are concerned, Cepheids are very close siblings of RR~Lyr stars. It is therefore very interesting to check if nonradial modes are excited in these pulsators as well. With this goal in mind, we conducted a systematic search for multiperiodicity among Pop.\\thinspace I Cepheids of the Large Magellanic Cloud (LMC). The primary source of data for this analysis was the OGLE-II photometry (\\.Zebru\\'n {\\it et~al.}\\thinspace 2001). The data, collected du\\-ring microlensing survey, covers the timebase of 1000-1200\\thinspace days, with 250-500 I-band measurements per star. Such a long, uniform and high quality dataset is particularly well suited for our purpose. Partial results of this survey were published by Moskalik, Ko{\\l}aczkowski \\& Mizerski (2004, 2006) and by Moskalik \\& Ko{\\l}aczkowki (2008a,b). In this paper we present a full discussion of our findings and provide a complete inventory of newly identified multiperiodic Cepheids in the LMC. ",
        "conclusions": "%  Taking advantage of a large and homogenous photometric dataset collected during OGLE-II project, we have performed a systematic frequency analysis of classical Cepheids identified in the Large Magellanic Cloud (Udalski {\\it et~al.}\\thinspace 1999b; Soszy\\'nski {\\it et~al.}\\thinspace 2000). This study has been aimed at finding multiperiodic and nonstationary variables among Pop.\\thinspace I Cepheids and at establishing reliable statistics of different forms of their pulsational behaviour. The main body of our survey has been conducted with OGLE-II photometry, but in cases when frequency resolution turned out to be insufficient, we have resorted to photometry collected by the MACHO experiment.  We have discovered two {\\it triple-mode} Cepheids, which pulsate with three lowest radial overtones simultaneously excited. Triple-mode radial pulsators have been known before, but only among High Amplitude $\\delta$~Scuti stars (Kov\\'acs \\& Buchler 1994; Antipin 1997b; Handler, Pikall \\& Diethelm 1998; Wils {\\it et~al.}\\thinspace 2008). This is the first detection of such pulsators among Cepheids. Three additional triple-mode Cepheids have been identified in the LMC by Soszy\\'nski\\thinspace {\\it et~al.} (2008a). These rare variables are a very valuable trophy, since their seismological analysis can strongly constrain the stellar evolution theory (Moskalik \\& Dziembowski 2005).  In 19\\% of FO/SO double-mode Cepheids we have detected {\\it periodic variability of amplitudes and phases} of the two radial modes. Another 16\\% of FO/SO pulsators is suspected of undergoing the same type of modulation, but on timescales longer than our data. This phenomenon is analogous to the Blazhko modulation in the RR~Lyr stars, which was first discovered a century ago (Blazhko 1907). In Blazhko FO/SO Cepheids both modes are modulated with a common period, always in excess of 700\\thinspace days. The amplitudes of the modes vary in opposite phases, that is a maximum amplitude of one mode coincides with minimum amplitude of the other. In frequency domain, the modulation manifests itself as splitting of each radial mode into a triplet or sometimes a quintuplet of equidistant peaks, spaced by frequency of the modulation.  The discovery of modulated FO/SO Cepheids shows that Blazhko effect and double-mode pulsations are not mutually exclusive. These stars will play a very important role in selecting a proper model to explain the Blazhko phenomenon. Observations of {\\it two} modulated modes in one star will put any proposed scenario to a very stringent test. Two most popular models, the oblique magnetic pulsator model (Shibahashi 2000) and the 1:1 resonance model (Nowakowski \\& Dziembowski 2001) have already failed this test and should be ruled out, at least in the case of FO/SO Cepheids.  Perhaps the most exciting result of our survey is detection of {\\it nonradial modes} in classical Cepheids. Such modes have been found in 42 overtone pulsators (which constitutes 9\\% of the sample) and in three FU/FO double-mode pulsators. They have {\\it not been found} in Cepheids pulsating in the fundamental mode.  We find two different types of nonradial modes in Cepheids. In most cases these modes are detected very close to the first radial overtone, but in several stars they appear at considerably higher frequencies. In the later case, the measured period ratio falls in a narrow range of $0.60-0.64$, which places the nonradial mode in the close vicinity of the (unobserved) fourth radial overtone. The observed frequency pattern cannot be explained by any form of amplitude or phase modulation, it is also incompatible with excitation of radial modes. This makes interpretation in terms of nonradial modes the only possibility.  Detection of nonradial modes in classical Cepheids has been claimed before by Kovtyukh {\\it et~al.}\\thinspace (2003), based on observations of bumps in spectral line profiles of four galactic Cepheids. Such structures can be indicative of nonradial oscillations ({\\it e.g.}\\thinspace Telting 2003). However, the authors have not shown that the observed line profiles vary in a periodic way, or that they vary with their own period different from that of the radial mode. This is an important test, because bumps in the spectral lines can be caused by mechanisms other than nonradial pulsations ({\\it e.g.}\\thinspace by shocks). Therefore, observations presented by Kovtyukh {\\it et~al.}\\thinspace (2003) are in our opinion inconclusive. The results of frequency analysis discussed in the current paper provide the first solid evidence of nonradial modes excitation in classical Cepheids.  Discovery of nonradial modes in Cepheids poses a major challenge to the stellar pulsation theory. In order to be photometrically detectable, nonradial modes have to be of low spherical degree, $\\ell$ (Dziembowski 1977). However, linear pulsation calculations predict, that in classical Cepheids all modes of $\\ell < 4$ should be heavily damped (Osaki 1977). This conclusion has recently been confirmed by new calculations of Mulet-Marquis {\\it et~al.}\\thinspace (2007), who have found that all nonradial modes of $\\ell < 5$ should be damped. We note, that this is a very different situation from that encountered in the RR~Lyr models, where many modes of $\\ell=1,2$ are linearly excited (Dziembowski \\& Cassisi 1999).  When writing of this manuscript was almost completed, we learnt about a new preprint on LMC Cepheids by Soszy\\'nski {\\it et~al.}\\thinspace (2008b). The authors have analyzed data collected during the third phase of the OGLE project (OGLE-III). Among many interesting results, they have also detected secondary periodicities in all subclasses of Pop.\\thinspace I Cepheids. This provides an important confirmation of our results with independent data. The incidence rates given by Soszy\\'nski {\\it et~al.}\\thinspace are very similar to ours in case of the double-mode pulsators, but significantly higher than ours in case of the FO Cepheids. This is most likely due to much longer time span of their photometry and higher number of measurements. Interestingly, Soszy\\'nski {\\it et~al.}\\thinspace has detected secondary periodicities also in the FU Cepheids, which we have not. A detailed comparison of incidence rates derived by us and by Soszy\\'nski {\\it et~al.}\\thinspace is beyond the scope of the present paper. We only note, that such comparison requires some caution, because the latter authors do not discriminate between se\\-condary periodicities resolved from the primary frequency and those, which are not resolved.    \\Acknow{This work was supported in part by MNiSW Grant No. 1~P03D~011~30.}"
    },
    "0809/0809.2927_arXiv.txt": {
        "abstract": "Gravitational tides are widely understood to strip and destroy galactic substructures. In the course of a galaxy merger, however, transient totally compressive tides may develop and prevent star forming regions from dissolving, after they condensed to form clusters of stars. We study the statistics of such compressive modes in an N-body model of the galaxy merger NGC~4038/39 (the Antennae) and show that $\\simeq 15\\%$ of the disc material undergoes compressive tides at pericentre. The spatial distribution of observed young clusters in the overlap and nuclear regions of the Antennae matches surprisingly well the location of compressive tides obtained from simulation data. Furthermore, the statistics of time intervals spent by individual particles embedded in a compressive tide yields a log-normal distribution of characteristic time $\\tau \\sim 10 \\Myr$, comparable to star cluster formation timescales. We argue that this generic process is operative in galaxy mergers at all redshifts and possibly enhances the formation of star clusters. We show with a model calculation that this process will prevent the dissolution of a star cluster during the formation phase, even for a star formation efficiency $\\epsilon$ as low as $\\sim 10\\%$. The transient nature of compressive tides implies that clusters may dissolve rapidly once the tidal field switches to the usual disruptive mode. ",
        "introduction": "Galaxy interactions are known to trigger star cluster formation. A typical example is NGC~4038/39 (the Antennae), where some 1000 star clusters has been detected \\citep{Mengel05,Whitmore07}. However, it is unclear how many of them will survive internal destruction: for instance, low star formation efficiency (SFE) or strong stellar feedback could destroy the molecular cloud \\citep{Elmegreen97}) and lead to ``infant-mortality'' \\citep{Bastian06,Whitmore07,deGrijs07}. Furthermore, these questions are central to interpreting the enhanced specific frequency and metallicity bi-modality of clusters in giant ellipticals \\citep{Harris91}. In this Letter, we stress that totally \\emph{compressive} tides widely exist in galaxy mergers and slowdown the early dissolution of clusters.  HST V-band and deep IR images of the Antennae galaxies have suggested a near power-law cluster mass function ($\\psi[M] \\propto M^{-2}$ for $10^4 \\solarm < M < 10^6 \\solarm$; \\citealt{Fall05,Anders07}) and an age distribution with a rough power-law fit ($dN/d\\tau \\propto \\tau^{-1}$) suggesting that many young clusters dissolve rapidly after they are born inside their primordial host gas cloud \\citep{Fall05}.  The evolution of (broad-band) mass-to-light ratios offers a strong hint that young clusters are severely perturbed and may even be out of dynamical equilibrium \\citep{Bastian06,deGrijs07}. Yet the low inter-cluster velocity dispersion measured in the Antennae galaxy merger indicates that cluster formation proceeds from a coherent coalescence of large-scale material as opposed to \\emph{in situ} shock dissipation and fragmentation \\citep{Whitmore05}. Computer simulations of galaxy mergers have long suggested that streams of material (made up of both gas and stars) developing during a merger will form dense structures, so increasing the star- and cluster-formation rates \\citep{Barnes88,Barnes92,Mihos96,Naab06}. A Schmidt law or variation thereof can be applied to circumvent numerical resolution issues and derive estimates of the star formation rate in rough agreement with observations (see e.g. \\citealt{Barnes04} for an application to the Mice; \\citealt{diMatteo07}).  Despite this wealth of progress, the long-term fate of star forming regions (bound structures or not) remains largely unclear, especially with regard to star cluster complexes such as those of the Antennae galaxies. The situation is exacerbated in the case of a galaxy merger because the gravitational tides vary rapidly in time along with the morphology of the system as a whole. An observational snapshot showing a candidate star-forming aggregate lacks information about the duration and motion in space to determine whether the aggregate will remain bound or not. When matter (gas or stars) is distributed in space in fragmented and irregular fashion, as is typical of galaxy mergers, the tidal field may exert either a disruptive outward stretch on an extended mass; or, on the contrary, it may exert an inward compressive force, which then cocoons the volume enclosed and increases its binding energy. While converging streams have long been suspected of driving the formation of tidal dwarf galaxies \\citep[e.g.][]{Barnes92,Bournaud06,Wetzstein07}, the impact of compressive tides themselves on sub-galactic scales has never been considered strictly as a gravitational effect. Previous authors analysed cluster formation as the result of a global trigger, using sticky-particles \\citep{Bournaud08} or SPH simulations with a threshold in pressure \\citep{Bekki02a,Bekki02b} or in density \\citep{Li04} for the formation of stars.  Here, we take a different stand by exploring the impact of compressive tides on cluster-size volume elements. We present in \\S2 a high-resolution numerical model of the Antennae galaxy merger to match kinematic and optical data. In \\S3, we discuss the key results obtained from the computer simulation. We close with a discussion of how the process outlined here bears on the statistics of observed cluster populations. ",
        "conclusions": "Using gravitational N-body simulation of the Antennae galaxy merger, we show that a significant fraction of the discs' mass undergoes totally compressive tides. This mode  plays an important role by preventing a cluster from dissolving. The compressive regions detected in the simulations show a very good match with observational data, both in term of their spatial distribution, and their characteristic duration which compare to observed clusters ages.  Compressive modes of the tidal field help keep together clusters that would otherwise dissolve at formation, through e.g. rapid gas expulsion. By introducing a tidal term in the virial Eq. (\\ref{eq:a}), we found that clusters should remain bound when they sit inside a compressive tide, even when the SFE is as low as 10\\%. The precise characteristics of the evolution of clusters is clearly dependent on the relative strength of the tidal field. This is a hint that to understand the survival of clusters, and especially the issues associated with the concept of ``infant mortality'', it is essential to include the full evolution of the environment in which clusters evolve. These super-virial young clusters sheltered by compressive tides will dissolve rapidly whenever the tides are no longer compressive. Hydrodynamical simulations will be required to link formally this effect with cluster formation histories."
    },
    "0809/0809.0235_arXiv.txt": {
        "abstract": "The hypothesis is rapidly gaining popularity that the dark energy pervading our universe is extra-repulsive ($-p>\\rho$). The density of such a substance (usually called phantom energy) grows with the cosmological expansion and may become infinite in a finite time producing a Big Rip. In this paper we analyze the late stages of the universe evolution and demonstrate that the presence of the phantom energy in the universe is not enough in itself to produce the Big Rip. This singularity occurrence requires the fulfillment of some additional, rather strong conditions. A more probable outcome of the cosmological evolution is the decay of the phantom field into 'normal' matter.  The second, more intriguing consequence of the presence of the phantom field is the possibility to introduce a cosmological scenario that does not contain a Big Bang. In the framework of this model the universe eternally expands, while its density and other physical parameters oscillate over a wide range, never reaching the Plank values. Thus, the universe evolution has no singularities at all. ",
        "introduction": "Experimental data \\cite{wmap, riess, perlmutter} indicate convincingly an accelerating cosmological expansion. The substance responsible for this effect is usually called dark energy. If we speculate that the universe is filled with a perfect fluid with equation of state $p=w \\rho$, the accelerating expansion appears if $w<-1/3$. There are a lot of  models of the dark energy: it can be explained by a non-zero cosmological constant (probably, this is the most natural hypothesis, in this case $w=-1$) or by presence of a global cosmic scalar field slowly moving down to the potential minimum (then $-1/3>w>-1$). Until recently, the instance of $w < -1$ did not attract considerable interest. A substance with such an equation of state (it is usually called phantom energy) generally possesses physically unacceptable properties, for example, the sound speed in it can be super-luminal.  Nevertheless the case of $w<-1$ has currently come to the attention. First of all, it does not contradict to the experimental data \\cite{caldwell1, caldwell2}. Furthermore, it has turned out that small perturbations are never super-luminal at least in some substances with $w<-1$ (see \\cite{gibbons} as an example). So, a phantom energy theory can be causal.  Physical properties of a dark energy with $w < -1$ differ markedly from the case when $w \\ge -1$. For instance, it violates the energy domination condition $\\rho\\ge |p|$. Since this is one of the principal assumptions of the classical black hole non-diminishing theorem, a black hole mass diminishes by accretion of the phantom energy \\cite{babichev}. But the difference most important for the cosmology becomes apparent when we consider expansion of the universe containing dark energy. If $w \\ge -1$ the dark energy density is not increasing (or even decreasing) as the universe expands. All bound systems (such as the solar one) remain bound forever; the dark energy presence manifests itself only on cosmological scales. If we allow $w<-1$,  the dark energy density grows and becomes infinite in a finite time. Increasing gravitational repulsion produced by the phantom energy will first destroy Galaxies and then any bound systems including elementary particles \\cite{caldwell1, caldwell2}. This phenomenon is usually called Big Rip. However, as we will see below, realization of such a catastrophic scenario requires the fulfilment of several supplementary conditions. Even if the phantom energy prevails in the universe, the Big Rip does not necessarily occurs.  There is another less evident consequence of $w < -1$: if dark energy is really the phantom one, the universe evolution needs not to contain a Big Bang. ",
        "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=270]{fig2.eps}} \\caption{The time dependence of $w$. We use $t\\cdot(8\\pi G \\rho_{de0}/3)^{1/2}$ instead of $t$.} \\label{fig2} \\end{figure} Of course, the suggested model is ingenuous. Intensive phantom field may interact in some other way decaying into 'normal' matter. Whatever the mechanism of the transformation may be, the universe after it has the following properties: \\begin{enumerate} \\item It has just passed through the stage of very rapid (at least, exponential) expansion. \\item It is flat and homogeneous. \\item It is 'normal' matter-dominated. \\item The Hubble constant is very large. \\end{enumerate} It is precisely these physical conditions that existed in our universe $\\sim 13.7$ milliard years ago, just after the inflation. If we assume that the inflation was caused by the phantom field (as we could see, the equations of the phantom cosmology (\\ref{a7}) are able to be closely analogous to those of the chaotic inflation \\cite{lindeinfl}), this makes possible to construct a model of universe evolution that does not contain a Big Bang. The universe leaves inflation being matter-dominated. As it expands the density of matter decreases, while the phantom density grows: eventually the universe passes into the phantom-dominated stage. As this takes place, the total density and the Hubble constant stop diminishing and begin to increase. Gradually the expansion becomes very fast, leading to an inflation-lake stage. Finally the phantom field decays into 'normal' matter, and the cycle repeats. Thus, the universe eternally expands, while its density and other physical parameters oscillate over a wide range, never reaching the Plank values. The universe evolution has no singularities like a Big Bang or a Big Rip.  The foregoing model of self-reproducing cyclic universe evolution is, in a certain sense, a return to the Steady State theory \\cite{hoyle1,hoyle2,hoyle3,gibbons}. Indeed, in both the cases the universe expands forever with no Big Bang, and the matter is generated by the decay of some field (in our case this is the phantom field). The difference is that in the classical Steady State universe this process goes on constantly providing constant universe density, while in the considering case the universe evolves cyclically returning to its original state, but the density varies within wide limits. Presence of the evolution allows to avoid most difficulties inherent to the classical Steady State model. Thus, the considering model may be called the Steady State {\\it on the average} universe.  The presence of the inflation-like stage in the considering scenario solves most of the problems appearing in an inflationless cosmological theory \\cite{lindeinfl}. The size of the universe during this stage increases so much that an observer in the after-inflation epoch can see only a small part of it. Consequently, the universe looks very flat and homogeneous. Moreover, as we could see (\\ref{a7}), it is possible to choose the model of the phantom energy so that even the dynamics of the inflation-like stage is similar to the standard chaotic inflation scenario \\cite{lindeinfl}.  Is the proposed model valid for the real universe? The answer to this question is currently not known with certainty, primarily because of the absence of direct experimental proofs of the phantom matter existence. A non-zero cosmological constant is quite enough to explain the experimental data. On the other hand, observational data show that allowed values of the parameter $w$ of the modern universe (at $2\\sigma$) lie between $-2.25 $ and $\\sim -0.5$ \\cite{caldwell2}. So, the existence of the phantom energy is not experimentally forbidden.  Theoretical status of the phantom energy is also vague. There are a lot of Lagrangians known, which describe fields with the necessary properties. It is still unclear, however, whether it is possible to introduce such a Lagrangian, completely free from paradoxical physical effects like super-luminal sound speed, catastrophic quantum instability of the vacuum etc.  Nevertheless, if the phantom energy does exist, the considering model is an interesting instance of cosmological solutions containing no singular points like a Big Bang or a Big Rip. It is worthy of note that, though in our case the universe density never reaches infinity (or, more precisely, Planck density), all the physical parameters vary over a very wide range and amount up to extremal values. So, the universe passes through the stages that can be called physically critical."
    },
    "0809/0809.0832_arXiv.txt": {
        "abstract": "{ Most X-ray studies of broad absorption line quasars (BALQSOs) found significant ($N_{\\rm{H}}\\sim 10^{22-24}$~cm$^{-2}$) intrinsic column densities of gas absorbing an underlying typical power-law continuum emission, in agreement with expectations from radiatively driven accretion disk wind models. However, direct spectral analysis was performed only on a limited number of bright sources.} { We investigate the X-ray emission of a large number of BALQSOs at medium to high redshift ($0.8\\lesssim z \\lesssim 3.7$) with the best data available to date.} { We drew a large BALQSO sample from the cross-correlation of SDSS DR5 and 2XMM catalogs to perform moderate-quality X-ray spectral and hardness ratio analysis and X-ray/optical photometry. We compare our results with previous studies of BALQSOs and theoretical disk wind model expectations.} { No or little intrinsic X-ray neutral absorption is found for one third of the spectroscopically analyzed BALQSO sample ($N_{\\rm{H}} < 4 \\times 10^{21}\\:$cm$^{-2}$ at 90\\% confidence level), and lower than typical X-ray absorption is found in the remaining sources ($\\langle N_{\\rm{H}}\\rangle \\sim 5\\times 10^{22}\\:$cm$^{-2}$) even including the faintest sources analyzed through hardness ratio analysis. The mean photon index is $\\Gamma \\sim 1.9$, with no significant evolution with redshift. The optical/X-ray spectral indices $\\alpha_{\\rm{ox}}$ are typical of radio-quiet broad line AGN, in contrast with the known (from previous X-ray studies) ``soft X-ray weakness'' of BALQSOs and in agreement with the lack of X-ray absorption. We found the low-absorption index (AI) subsample to host the lowest X-ray absorbing column densities of the entire sample. } { X-ray selected BALQSOs show lower X-ray absorption than purely optically selected ones, and soft X-ray weakness does not hold for any of them. Their outflows may be launched by different mechanisms than classical soft X-ray weak BALQSOs or they may be the tail of the already known population seen along a different line of sight, in both cases expanding the observational parameter space for their search and investigation.} ",
        "introduction": "\\label{sect:intro}  Recent observational works have unveiled the presence of outflows in many radio-quiet active galactic nuclei (AGN). Such outflows may be relevant both for the unification schemes and modelization of AGN \\citep[e.g.][]{2000ApJ...545...63E} and for the growth of cosmic structures, partially providing the feedback needed to reconcile observations with theoretical expectations \\citep[e.g.][]{2005Natur.433..604D,2005ApJ...635L..13S,2005ApJ...619...60L}.  We know that AGN outflows are there and are ionized, since we observe them as blueshifted absorption features in their UV/X-rays continua \\citep[warm absorbers, ultraviolet broad absorption lines, and X-ray absorption lines, see e.g.][]{2003ARA&A..41..117C, 2006AN....327.1012C, 2007ApJ...659.1022K}. We still have a limited knowledge of the physics responsible for launching and accelerating the outflows/wind. The driver of these winds may be either radiation, thermal, or magnetic pressure \\citep[e.g.][]{2007ASPC..373..267P}. Their launching region may lay between a few hundreds gravitational radii \\citep[as estimated by variability studies on blueshifted X-ray absorption lines, e.g.][]{2007ApJ...670..978B} and a few parsecs \\citep[where thermal evaporation of outer dusty torus takes place, e.g.][]{2001ApJ...561..684K}. The wind may be equatorial, as in accretion disk winds \\citep{1995ApJ...451..498M, 2000ApJ...543..686P}, and/or polar \\citep{1999ApJ...527..624P,2007ApJ...661..693P}. The wind may be continuous or made of knots/blobs of matter \\citep[as in the model of][]{2004A&A...413..535G}.  The vast majority of both theoretical and observational works about AGN outflows concerns the first class of objects discovered to host strong nuclear outflows: the broad absorption line quasars (BALQSOs). These are $10-15\\%$ of optically selected QSOs\\footnote{The fraction is higher, $\\gtrsim 20\\%$, for near-infrared and radio selected QSOs \\citep{2008ApJ...672..108D, 2008arXiv0801.4379S}, so maybe suggesting a bias in optical band against BALQSOs detection.} that show broad ($\\gtrsim 2000\\:$km~s$^{-1}$) UV resonant absorption lines strongly blueshifted with respect to the systemic velocities of the source, indicating outward motions along the line of sight with typical velocities of several $10^3\\:$km~s$^{-1}$ and extending up to $\\sim 60\\,000\\:$km~s$^{-1}$ \\citep[e.g. Q1414+087,][]{1983PASP...95..341F}. The observed fraction of BALQSOs could correspond either to the covering fraction of a wind present in all AGN or to the intrinsic fraction of QSOs hosting (totally covering) massive nuclear winds, so tracing an evolutionary phase of the AGN activity lasting $\\sim 10-15\\%$ of their lives. The great similarites between the optical/UV emission line and continuum properties of BALQSOs and non-BALQSOs \\citep[e.g.][]{2003AJ....126.2594R} support the first scenario. In the latter scenario, the large amount of gas and dust surrounding the central source should lead to enhanced far-infrared and submillimeter emission in BALQSOs with respect to non-BALQSOs. Recent studies have found no \\citep{2003ApJ...598..909W} or little \\citep{2007MNRAS.374..867P} differences among the two populations; however deeper observations of larger samples are needed to assess a firm conclusion about this point.  X-ray studies can help to constrain the physical mechanisms responsible for launching and accelerating AGN winds, and discriminate among different scenarios for BALQSOs. In the popular model of radiatively driven accretion disk wind, the high column densities of X-ray absorbing gas -- postulated in the original form by \\citet{1995ApJ...451..498M} in order to prevent overionization of the wind by the strong central continuum -- arise naturally in the hydrodynamical 2D simulations of \\citet{2000ApJ...543..686P} and \\citet{2004ApJ...616..688P}. If accretion disk wind models hold, X-ray observations of BALQSO should thus reveal high ($N_{\\rm{H}}>10^{23}\\:$cm$^{-2}$) columns of gas absorbing the primary continuum and shielding the UV wind.  Since the time of ROSAT observations, BALQSOs are known to be rather dim in X-rays \\citep{1995ApJ...450...51G}. Their X-ray flux, typically $10-30$ times lower than expected from their UV flux, often qualifies them as 'soft X-ray weak' objects \\citep{1997ApJ...477...93L}. The discovery of a correlation between C~IV equivalent width and soft X-ray weakness \\citep{2000ApJ...528..637B} strongly suggests large absorption to be the reason for the low X-ray flux of BALQSO, so supporting the accretion disk wind models and the underlying typical QSO spectral energy distribution (SED).  A two-fold strategy has been pursued in the last years in order to overcome the X-ray weakness of BALQSOs and determine their high-energy properties. A hardness ratio and/or stacking analysis technique was applied to snapshot observations of a large number of sources in order to characterize their mean properties \\citep[e.g.][]{2001ApJ...558..109G, 2006ApJ...644..709G}, while detailed spectral analysis of deep observations was only possible for a few bright sources \\citep[e.g.][]{2002ApJ...567...37G,2003AJ....126.1159G, 2004ApJ...603..425G, 2005AJ....130.2522S, 2005A&A...433..455S}. Both kind of studies revealed high column densities ($N_{\\rm{H}} \\sim 10^{22-24}\\:$cm$^{-2}$) of gas absorbing a 'normal' intrinsic SED with powerlaw spectral indices typical of radio-quiet, non-absorbed QSOs \\citep[$\\Gamma \\sim 1.7-2$,][]{2005A&A...432...15P}. In almost none of the sources a simple neutral absorption scenario is able to well reproduce the data, but rather complex absorbers (ionized and/or partially covering the source and/or variables) are generally observed or deduced from the available data. To our knowledge, the only two solid cases of Compton-thick absorption (i.e. $N_{\\rm{H}}\\gtrsim 1.5\\times 10^{24}\\:$cm$^{-2}$) in BALQSOs have been found in the mini-BALQSO\\footnote{mini-BALQSOs show somewhat narrower and less blueshifted BALs than BALQSOs. They form a continuous sequence in UV outflow properties together with Narrow Absorption Line (NAL) QSOs and BALQSOs, i.e. NAL $\\rightarrow$ mini-BAL $\\rightarrow$ BAL with increasing absorption troughs widths \\citep[see e.g.][]{2004ASPC..311..203H}.} Mrk~231 \\citep{2004A&A...420...79B} and in the classical BALQSO LBQS 2212-1759 \\citep{2006A&A...446..439C}.  It should be noted that recent observations of radio-loud BALQSOs, thought to be seen face-on on the basis of radio emission variability, have found little or no X-ray absorption, so suggesting a different scenario for polar outflows \\citep{2008ApJ...676L..97W}.  Thanks to the current availability of both optical and X-ray large catalogs as the Sloan Digital Sky Survey \\citep[SDSS,][]{2000AJ....120.1579Y} and the Second {XMM-{\\textit{Newton}}} Serendipitous Source Catalog \\citep[2XMM,][]{2XMM} we have been able to make a step forward in joining the two above-mentioned strategies used in BALQSOs X-ray studies. Here we present X-ray analysis of a large (54 sources) sample of BALQSOs drew from the cross-correlation of the above mentioned two catalogs.  A cosmology with $H_0=70\\:$km~s${-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{\\rm{M}}=0.3$ is used throughout the paper. ",
        "conclusions": "We analyzed X-ray properties of 54 BALQSOs drawn from the cross-correlation of SDSS DR5 and 2XMM catalogs by performing spectral analysis on 22 sources, hardness ratio analysis on the remaining ones.  Our results in terms of $\\alpha_{\\rm{ox}}$ and the amount of intrinsic X-ray absorption are different with respect to the literature results, that are mainly based on purely optically selected BALQSOs. We found none but one of our X-ray selected BALQSOs being soft X-ray weak, but the most striking result of our spectral analysis is the low intrinsic X-ray absorbing column density value measured for more than a half of the sample, $N_{\\rm{H}} \\leq  10^{22}\\:$cm$^{-2}$.  We observationally confirm the statistical result presented by \\citet{2008MNRAS.386.1426K} which found a bimodal distribution in AI for BALQSOs; we found that the low intrinsic neutral X-ray absorption measured for a large fraction of the sample is linked to the low AI measured from their UV spectra. We should treat with caution such kind of sources when included in BALQSO samples until we do not clarify their role in the quasar outflow scenarios. Overall, we found a lower than expected $\\langle N_{\\rm{H}}\\rangle$ also for the high-AI BALQSO subsample.  The low measured intrinsic absorption may be reconciled with theoretical results by allowing either the intrinsic continuum to comprise a significant soft-excess or a scattered component, or the X-ray absorber to be ionized, or only partially covering the continuum source. Unfortunately, the rather low statistics does not allow us to assess the significance of the presence of an intrinsic soft-excess or of a partially covering absorber. Furthermore, the redshift of our sources prevents us from checking the presence of the most usual X-ray spectral features due to warm absorbers, i.e. the O~VII and O~VIII absorption edges. Investigating the presence of more highly-ionized gas (e.g. absorption features due to Fe~XXV, Fe~XXVI), absorbing the X-ray continuum at the highest energies probed by our observations, would need much deeper observations of such distant sources with the current X-ray telescopes. With these caveats in mind, we can conclude that our measured intrinsic column densities might be sometimes lower limits to the presence of intrinsic absorption for this BALQSO sample.  We may be probing either a tail of the known BALQSO population, probably related to a line of sight with larger angle with respect to the accretion disk plane, or a different outflow launching mechanism that does not need a large amount of X-ray absorbing, shielding gas to accelerate the wind, i.e. a different than radiatively driven quasar wind or a scenario which includes magnetic forces as shielding and/or accelerating factors. The relative number of X-ray non-detected BALQSOs compared to the number of X-ray non-detected SDSS DR5 QSOs, together with the smooth connection of X-ray properties between bright (i.e. spectrally analyzed) and faint (i.e. hardness ratio analyzed) X-ray detected BALQSOs strongly suggests the former hypothesis. If this were true, we can expect more X-ray absorption in the $\\sim 100$ BALQSOs with only X-ray upper limits. These intriguing results deserve to be investigated more accurately with the increase of the number of BALQSOs well-studied in both the UV and X-ray band in order to test the accretion disk wind models and the quasar outflow scenarios."
    },
    "0809/0809.5283_arXiv.txt": {
        "abstract": "We present an updated catalog of 1300 objects in the field of M31, including 670 likely star clusters of various types, the rest being stars or background galaxies once thought to be clusters. The coordinates in the catalog are accurate to 0.2\\arcsec,  and are based on images from the Local Group Survey \\citep[LGS,][]{massey} or from the DSS. Archival HST images and the LGS were inspected to confirm cluster classifications where possible, but most of the classifications are based on spectra taken of  $\\sim1000$  objects with the Hectospec fiber positioner and spectrograph on the 6.5m MMT. The spectra and images of young clusters are analyzed in detail in this paper; analysis of older clusters will appear in a later paper.  Ages and reddenings of 140 young clusters are derived by comparing the observed spectra with model spectra. Seven of these clusters also have ages derived from HST color-magnitude diagrams (two of which we present here); these agree well with the spectroscopically determined ages.  Combining new V band photometry with the M/L values that correspond to the derived cluster ages, we derive masses for the young clusters, finding them to have masses as great as $10^5$ with a median of $10^4{\\rm M}_{\\sun}$, and a median age of 0.25 Gyr. In comparison therefore, Milky Way open clusters have the lowest median mass, the Milky Way and M31 globulars the highest, and the LMC young massive clusters and the M31 young clusters are in between. The young clusters in M31 show a range of structure. Most have the low concentration typical of Milky Way open clusters, but there are a few which have high concentrations. We expect that most of these young clusters will be disrupted in the next Gyr or so, however, some of the more massive and concentrated of the young clusters will likely survive for longer. The spatial distribution of the young clusters is well correlated with the star-forming regions as mapped out by mid-IR emission. A kinematic analysis likewise confirms the spatial association of the young clusters with the star forming young disk in M31. ",
        "introduction": "In the Milky Way, there is a clear separation between known open clusters (which have diffuse structure, generally have low masses and ages, and belong to the disk) and globular clusters (which have a more concentrated structure, higher masses and ages, and in many cases belong to the halo). When the only well-studied globular cluster system was that of the Milky Way, it was generally thought that this separation was because globular clusters were fundamentally different from other star clusters, perhaps because of conditions in the early universe \\citep{peebles68,fallrees85}.  However, it is possible to produce this apparent bimodality from clusters formed in a single process, with the same cluster initial mass function. In this picture, cluster disruption mechanisms, which are more effective at destroying low-mass clusters in particular because of two-body relaxation \\citep{spitzer58,spitzerharm}, would remove almost all of the low-mass older clusters.  If all clusters were born with similar cluster mass functions, then we would expect to see the occasional high-mass young cluster. In fact, we do see these in other galaxies. The ``populous blue clusters'' of the LMC \\citep{kcf80,hodge81,mateo93} have been suggested as examples of young objects which will evolve into globular clusters. M33 also has a few likely massive young clusters \\citep{chandar99} , and such clusters have been found in a number of normal isolated spirals \\citep{larsen04}. It is possible that the seeming absence of such objects in the Milky Way is merely an observational selection effect; recently, there have been discoveries of  heavily reddened open clusters such as Westerlund 1, which likely has mass in excess of $10^5 $M$_\\odot$ \\citep{clark}.  What of M31's clusters? While its clusters have been studied since the 1960s \\citep[eg][]{kinman,vetesnik,vdb} and it was noted even then that some of these clusters had colors indicating young populations, their nature is still not entirely clear. \\citet{vdb} called them open clusters, while \\citet{krienke} adopted the simple convention of calling any cluster projected on M31's disk a ``disk cluster'' until proved otherwise by kinematics. This avoids the question of whether young clusters are fundamentally different from globular clusters in structure, formation, etc. Our detailed study of M31 young clusters, incorporating kinematics, should cast some more light on these questions.   It is only recently that detailed constraints on the mass, kinematics, age and structure of cluster populations in M31 have been obtained, particularly for clusters projected on the inner disk and bulge. Multi-fiber spectroscopy and HST imaging have played an important role here.  A number of M31 globular cluster catalogs have been created over the years, giving a very heterogeneous result, with significant contamination by both background galaxies, foreground objects and even non-clusters in M31 itself. While the work of \\citet{perrett}, \\citet{barmby}, \\citet{galleti2}, \\citet{kim}, and \\citet{lee} has gone a long way towards cleaning up the catalogs and winnowing out the non-clusters, still more work is needed for both young and old clusters.  Here we present a new catalog of M31 star clusters which were originally classified as globular clusters, all with updated high-quality coordinates. We have observed a large number of these clusters with the MMT and the Hectospec multi-fiber system. In this paper we study more than 100 young M31 clusters in detail.  In subsequent papers we will address the kinematics, ages and metallicities of the older clusters.  The M31 young clusters have been studied both by authors aiming to study its open clusters \\citep[eg][]{hodge79,hodge87,williams01,krienke}, and also by others who have aimed to study its globular clusters.  For example, \\citet{elsonw} pointed out the existence of young clusters in M31, estimated masses between $10^4$ and $10^5 M_\\odot$, and drew attention to their similarity to the populous blue clusters in the LMC. \\citet{barmby} noted the existence of 8 clusters with strong Balmer lines in their spectra, which they tentatively classified as young globular clusters. \\citet{beasley} \\citep[confirmed with a later sample by][]{puzia} added more clusters, and commented that HST imaging (then unavailable) was needed to distinguish between structure typical of open and globular clusters. \\citet{burstein} added more new young clusters, bringing the total to 19, and \\citet{fusipecci} increased the sample to 67. In general, authors have associated these young clusters with M31's disk, although \\citet{burstein} invoke an accretion of an LMC-sized galaxy by M31.  Observations are particularly challenging for clusters projected on M31's disk: many of the early velocities had large errors, and there were issues with background subtraction. Here we discuss high-quality spectroscopic measurements of kinematics and age for the young clusters, supplemented with HST imaging to delineate the structural, spatial and kinematical properties of these young clusters.  We find that while they are almost all kinematically associated with M31's young disk, and their age distribution will allow us to test suggestions that M32 has had a recent passage through M31's disk \\citep{gordon,block}. The young  clusters have structure and masses which range all the way from the low mass Milky Way open clusters to higher mass, more concentrated globular clusters, although they are dominated by the lower concentration clusters. We will also discuss the likelihood that these clusters will survive. ",
        "conclusions": "We present a new catalog of 670 M31 clusters, with accurate coordinates. In this paper we focus on the 140 clusters (many originally classified in the literature as globular clusters) which have ages less than 2 Gyr: most have ages between $10^8$ and $10^9$ years. Using high-quality MMT/Hectospec spectra, excellent ground based images, and in some cases, HST images, we explore the nature of these clusters.  With the exception of NGC 205's young cluster, they have spatial and kinematical properties consistent with formation in the star-forming disk of M31. Many are located close to the 10 kpc ``ring of fire'' which shows active star formation. The age distribution of our clusters, plus that of the younger clusters of \\citet{williams01}, shows no evidence for a peak in star formation there between $2 \\times 10^7$ and $10^9$ years ago, which we might expect if the ring was created by a recent passage of M32 through the disk, as suggested by \\citet{gordon} and \\citet{block}.  We have estimated their masses using spectroscopic ages and M/L ratios, (in some cases) ACS color-magnitude diagrams, and new photometry from the Local Group Survey. The clusters have masses ranging from 250 to 150,000 $M_\\odot$. These reach to higher values than the known Milky Way open clusters, but it must be remembered that our sample of open clusters in the Milky Way is far from complete. The most massive of our young clusters overlap the mass distributions of M31's old clusters and the Milky Way globulars.  Interestingly, although most of the young clusters show the low-concentration structure typical of the Milky Way open clusters, a few have the high concentrations typical of the Milky Way globulars and the old M31 clusters. We estimate that most of these young clusters will disrupt in $1-2$ Gyr, but the massive, concentrated clusters may well survive longer."
    },
    "0809/0809.2695_arXiv.txt": {
        "abstract": "Association of the high-frequency quasi-periodic oscillation (HFQPO) pairs with the broad Fe K line in XTE J1550-564 and GRO J1655-40 is discussed based on the magnetic coupling (MC) of a rotating black hole (BH) with its surrounding disc. The 3:2 HFQPO pairs are interpreted by virtue of the inner and outer hotspots arising from non-axisymmetric magnetic field, where the inner hotspot is produced by a torque exerted at the inner edge of the disc, and the outer hotspot is created by the screw instability of the large-scale magnetic field. The very steep emissivity index is created predominantly by the torque exerted at the inner edge of the disc. It turns out that the 3:2 HFQPO pairs observed in the two sources can be fitted by tuning several model parameters, such as the BH spin, and the main features of this model lie in three aspects. (1) The condition for only one HFQPO is discussed based on the two mechanisms for producing the 3:2 HFQPO pairs, (2) an explanation is given for a systematic shift away from disc dominated flux with the increasing power-law flux as the HFQPO pairs shift from the higher to lower frequencies, which is consistent with the analysis given by Remillard et al. (2002), and (3) the BH spin in XTE J1550-564 and GRO J1655-40 can be estimated by combining the 3:2 HFQPO pairs with the very steep emissivity index required for fitting the broad Fe K emission line. ",
        "introduction": "Quasi-periodic oscillations in black hole X-ray binaries (BHXBs) have become a very active research field since the launch of the NASA satellite \\emph{RXTE} (Bradt et al. 1993). One of the remarkable features in BHXBs is that the high-frequency quasi-periodic oscillations (HFQPOs) could appear in pairs with the puzzling commensurate frequencies, which have been observed in GRO J1655-40 (450, 300Hz; Remillard et al. 1999; Strohmayer 2001a; Remillard et al. 2002, hereafter R02), XTE J1550-564 (276, 184, 92Hz; Miller et al. 2001; R02) and GRS 1915+105 (168, 113, 67, 41Hz; McClintock \\& Remillard 2006; Remillard \\& McClintock 2006, hereafter MR06 and RM06, respectively). The above HFQPOs appear in pairs with the 3:2 ratio of the higher frequency to the lower frequency (henceforth 3:2 HFQPO pairs), e.g. 276 vs 184Hz and 450 vs 300Hz occur in XTE J1550-564 and GRO J1655-40, respectively, while two 3:2 HFQPO pairs, 168 vs 113Hz and 67 vs 41Hz, appear in GRS 1915+105.  A number of models have been proposed to explain the origin of HFQPO pairs in BHXBs. Strohmayer (2001a, 2001b) investigated combinations of the azimuthal and radial coordinate frequencies in general relativity to explain the HFQPO pairs in GRO J1655-40 and GRS 1915+105. Wagoner et al. (2001) regarded the HFQPO pairs as fundamental g-mode and c-mode discoseismic oscillations in a relativistic accretion disc. Very recently, Silbergleit \\& Wagoner (2007) discussed the corotation resonance and diskoseismology modes of black hole (BH) accretion discs and suggested that the HFQPO pairs may relate to the excitation of two (groups of) g-modes of discoseismic oscillations. Abramowicz \\& Kluzniak (2001) explained the pairs in GRO J1655-40 as a resonance between orbital and epicyclic motion of accreting matter. Recently, the resonance model is presented in a more realistic context, in which \"parametric resonance\" concept is introduced to describe the oscillations rooted in fluid flow where there is a coupling between the radial and polar coordinate frequencies (Abramowicz et al. 2003; Kluzniak et al. 2004; T\\\"{o}r\\\"{o}k et al. 2005). As argued by van der Klis (2000, 2006), HFQPOs in X-ray binaries probably originate from the inner edge of an accretion disc around a neutron star or a stellar-mass BH, since millisecond is the natural timescale for accretion process in these regions.  On the other hand, Wilms et al. (2001) presented the first XMM-Newton observation of MCG-6-30-15, and they found a broad Fe disc line from the observed spectrum. However, no quasi-periodic oscillations have been detected in this source. A very steep emissivity index ($4.3<\\beta<5.0$) from the inner accretion disc is required for fitting the broad Fe K emission line, which is difficult to be explained within the framework of a standard accretion disc (SAD). Wilms et al. (2001) suggested that the magnetic extraction of rotational energy from a spinning black hole should be invoked to create this steep emissivity.  The relativistic disc lines for microquasars have been analyzed by some authors, e.g., XTE J1550-564 by Sobczak et al. (2000), Miller et al. (2003, 2004) ; XTE J1650-500 by Miller et al. (2004), Miniutti, Fabian, and Miller (2004); GRO J1655-40 by Miller et al. (2004), Diaz Trigo et al. (2007). However, the steep emissivity indexes required by broad Fe K $\\alpha$ lines were hardly worked out except a few cases, such as, $\\beta =5.5^{+0.5}_{-0.7}$ for GX 339-4 in the very high state, and $\\beta \\simeq 5$ for XTE J1650-500 based on the observation of XMM-Newton. A detailed review for relativistic X-ray emission lines from the inner accretion disc around black holes for AGNs and stellar-mass black holes are given by Miller (2007).  With the existence of a magnetic field connecting a rotating black hole to its surrounding disc, energy and angular momentum can be transferred from the black hole to the disc, and this energy mechanism is referred to as the MC process, which is regarded as one of the variants of the Blandford-Znajek (BZ) process (Blandford {\\&} Znajek 1977, 1999; Li 2000, 2002, hereafter L02, Wang 2002).   Wang et al. (2002, 2003a, hereafter W03a) incorporated the BZ and MC processes into black hole accretion disc, and henceforth this model is referred to as MC-I model, which was used to interpret the very steep  emissivity index required by the observation of MCG-6-30-15. In addition, Wang et al. (2005, hereafter W05) suggested that the 3:2 HFQPO pairs observed in the above BHXBs can be fitted by virtue of two hotspots in the inner disc, which are produced by the non-axisymmetric MC process.  The fatal shortcoming of MC-I model is that the fits of the 3:2 HFQPO pairs always accompany jets, which is driven by the BZ process. As reviewed in MR06 and RM06, jets are generally observed in the hard state of BHXBs, while HFQPOs are always detected in the steep power-law (SPL) state. Thus the fits of the 3:2 HFQPO pairs based on MC-I model are inconsistent with the observations.  Based on a careful analysis of the 1998-1999 outburst of XTE J1550-564 and the 1996-1997 outburst of GRO J1655-40 some interesting features have been found in R02, the 3:2 HFQPO pairs could be detected simultaneously in the SPL state, and a systematic shift away from disc dominated flux with the increasing power-law flux is detected as the HFQPO pairs shift from 276 to 184 Hz for XTE J1550-564 and from 450 to 300 Hz for GRO J1655-40 as shown in Figs. 5 and 7 of R02, respectively. In addition, Remillard (2004) noted that the broad Fe K emission lines can be also detected in the SPL state in XTE J1550-564 and GRO J1655-40. These results imply that the Fe K lines could be detected simultaneously with the 3:2 HFQPO pairs in these two sources.\\\\  In this paper, we interpret the 3:2 HFQPO pairs associated with the broad Fe K line in XTE J1550-564 and GRO J1655-40 by modifying MC-I model in two aspects. (1) The open field lines corresponding to the BZ process is removed, and (2) a non-zero torque exerted at the inner edge of the disc is introduced. Henceforth the modified model is referred to as MC-II model. It turns out that the 3:2 HFQPO pairs observed in the two sources can be fitted by tuning several model parameters, such as the BH spin, and the main features of MC-II model lie in the following aspects: (1) The condition for only one HFQPO is discussed based on the two mechanisms for producing the 3:2 HFQPO pairs. (2) An explanation is given for a systematic shift away from disc dominated flux with the increasing power-law flux as the HFQPO pairs shift from the higher to lower frequencies, which is consistent with the analysis given by R02. (3) The BH spin in XTE J1550-564 and GRO J1655-40 can be estimated by combining the 3:2 HFQPO pairs with the very steep emissivity index required for fitting the broad Fe K emission line.  This paper is organized as follows. In \\S 2 we present a detailed description of MC-II model, in which the main differences between MC-I and MC-II models are outlined. In \\S 3 we describe the elements for fitting the 3:2 HFQPO pairs observed in XTE J1550-564 and GRO J1655-40, and present the calculation sequence and results based on MC-II model. We discuss the possible correlation of the 3:2 HFQPO pairs with the very steep emissivity index required by broad Fe K $\\alpha$ lines. It is expected that the 3:2 HFQPO pairs could be observed with the steep emissivity index, and the black hole spin could be estimated once the value ranges of the index are determined in the future observations. Finally, we discuss some issues related to MC-II model in \\S 4.  Throughout this paper the geometric units $G = c = 1$ are used. ",
        "conclusions": "In this paper, association of the 3:2 HFQPO pairs with the broad Fe K line in XTE J1550-564 and GRO J1655-40 is discussed based on MC-II model. It is shown that the 3:2 HFQPO pairs observed in the two sources can be fitted with the observational constraints by tuning several model parameters. Some related issues are given as follows.  (1) It is noted that the mechanisms of creating the two hotspots are different: the inner one is produced predominantly by the magnetic torque exerted at the inner edge of the disc, while the outer one arises from the screw instability of the large-scale magnetic field. It is helpful to imagine the magnetic field line as an elastic string. The rotating BH twists the field line, while the field line tries to untwist itself. Once the toroidal component of the magnetic field is strong enough to satisfy the criterion, the screw instability will occurs, just as a twisted elastic string releases its energy under appropriate conditions.  (2) Some authors argued that the coronal heating in some stars including the Sun is probably related to dissipation of currents, and very strong X-ray emissions arise from variation of magnetic fields (Galsgaard \\& Parnell 2004; Peter et al. 2004). Analogously, if the corona exists above the disc in MC-II model, we expect that the corona above $r_{out}$ might be heated by the induced current due to the screw instability of the non-axisymmetric magnetic field. Thus the ratio of the disc bolometric flux to the power-law flux is greater for the inner hotspot than that for the outer hotspot, and the feature of X-ray radiation given by R02 can be interpreted based on MC-II model with corona, i.e., there exists a systematic shift away from disc dominated flux with the increasing power-law flux as the HFQPO pairs shift from the higher frequency to the lower frequency based on the observations of XTE J1550-564 and GRO J1655-40.  (3) Since the two hotspots are produced by two different mechanisms, we can explain the observations that HFQPOs do not always appear in simultaneous pairs, because the condition for the magnetic torque and that for the screw instability would not be satisfied at the same time. Since the inner and outer hotspots are related to two different mechanisms based on the non-axisymmetric magnetic fields in the plunging region and the BH horizon, either the fluctuations of the magnetic fields at different position or the criterion of the screw instability could affect the HFQPO pairs. Unfortunately, we cannot discuss this issue in a quantitative way due to the unknown origin of the magnetic field near the BH.  (4) According to equation (21) the HFQPO frequency is inversely proportional to the BH mass, which providers a strong limit to the fits. We can neither fit the HFQPO with 92 Hz in XTE J1550-564 nor the 3:2 HFQPO pair (67, 41Hz) in GRS 1915+105 based on equation (21) for the measured BH masses. Furthermore, the 3:2 HFQPO pair (168, 113Hz) of GRS 1915+105 are not fitted, because a systematic shift away from disc dominated flux with the increasing power-law flux as the HFQPO pair shift from the higher frequency to the lower frequency has not been observed in this source.  (5) As argued above, the higher and lower HFQPO frequencies are related to the inner and outer hotspots rotating with the angular velocity located at $r_{in}$ and $r_{out}$, respectively. According to equation (21) the HFQPO frequencies are inversely proportional to the BH mass, and the spin is related to the HFQPO frequencies via $r_{in}$ and $r_{out}$, where the HFQPO frequencies are calculated by the angular velocity. Since the evolution time scale of the BH mass and spin are much longer than the observation time scale, the two quantities are considered fixed in calculations.  (6) As pointed out in R02, the 3:2 HFQPO pairs could have a small, but significant, change in frequency. For example, the HFQPO frequencies observed in XTE J1550-564 deviate from 274 Hz and 184 Hz during different outbursts. These small deviations in HFQPO frequencies can be calculated by resolving \\textbf{FIE} with the adjustable parameters $n$ and $\\delta$ as shown in Table 2.   (7) As argued above, both the non-axisymmetric magnetic field configuration and the corona above the disc surface are required for interpreting the 3:2 HFQPO pairs observed in the SPL state of the BHXBs. Thus we can interpret the observation that sometimes the 3:2 HFQPO pairs cannot be observed in the SPL state, if the magnetic field configuration is axisymmetric. Based on the same reason we predict that the 3:2 HFQPO pairs could be observed with the very steep emissivity index, if the magnetic field configuration is non-axisymmetric. However, the HFQPO pairs cannot be produced with the emissivity index for the axisymmetric magnetic field configuration.  (8) According to $B_D \\propto r^{-n}$ the ratio of the magnetic field at $r_{out}$ to that at $r_{in}$ can be written as $R_B \\equiv (B_D)_{out}/(B_D)_{in}=(r_{in}/r_{out})^n$. Based on \\textbf{FIE} we have the curves of $R_B$ for XTE J1650-564 vary with $a_*$ as shown in Figure 5, from which a strongly constraint to the BH spin can be found. The BH spin should be very high, or at least greater than some intermediate value to avoid the magnetic field at $r_{out}$ too low to produce an outer hotspot.   (9) There is less certain evidence that the same 3:2 ratio occurs for QPOs observed in low-mass active galactic nuclei, Sgr A$^*$, (T\\\"{o}r\\\"{o}k 2005a, b; Ashenbach 2004a, b) and in a few nearby Seyferts (Lachowicz et al. 2006). MC-II model can be used to fit QPOs pairs observed in the super-massive BH systems, and we obtain the curves of the emissivity index $\\beta$, the power-law index $n$ and the parameter $\\delta$ varying with the BH spin of for Sgr A$^*$ as shown in Figure 6.   We can find from Figure 6 that the parameter n varies from 4.272 to 7.108 with $a_*$ decreasing from 0.998 to 0.853 for $m_{BH}=4.4\\times 10^6$, while it varies from 5.040 to 10.239 with $a_*$ decreasing from 0.998 to 0.555 for $m_{BH}=2.6 \\times 10^6$.   (10) The parameter $n$ is used to indicate the variation of the large-scale magnetic field with the disc radius, which is a key parameter for fitting the 3:2 frequency ratio. Based on \\textbf{FIE} we have the variation of the 3:2 frequency ratio with the parameter n for the given values of the BH spin as shown in Figure 7.   \\textbf{Inspecting Figure 7, we find that the frequency ratio decreases monotonically with the increasing parameter $n$. The greater $n$ corresponds to the less ratio for the given spin, and the greater spin corresponds to both the less value and less variation of $n$ for the given ratio. For XTE J1550-564, e.g., the ratio increases about 5.7\\% as the value of $n$ decreases about 10\\% for  $a_*=0.5$, while the same variation occurs as $n$ only decreases about 1.9\\% for $a_*=0.998$ . This result implies that the observed 3:2 frequency ratio is reached for a very narrow range of $n$ with the greater BH spin.}   (\\textbf{11}) Finally, we discuss the values of the parameters $\\alpha_H$ and $\\alpha_m$, which are used to indicate the disc thickness at $r_{in}$ and the strength of the accretion rate, respectively. The emissivity index of XTE J1550-564 varies with the BH spin for different values of $\\alpha_H$ and $\\alpha_m$ are shown in Figure \\textbf{8}.   Inspecting Figure \\textbf{8a}, we find that the values of the emissivity index are insensitive to the values of $\\alpha_H$ and $\\alpha_m$, provided that they are limited to $\\alpha _H \\le 0.1$ and $\\alpha_m \\le 1.0$. These value ranges are reasonable for thin accretion disc in MC-II model. In addition, based on the features of MC-II model, the 3:2 HFQPO pairs are also insensitive to the parameters with the above limited values.\\\\\\\\   \\noindent {\\bf Acknowledgements:}\\quad This work is supported by the National Natural Science Foundation of China under grant 10873005 and National Basic Research Program of China under grant 2009CB824800. We are grateful to the anonymous referee for his (her) helpful comments on the manuscript."
    },
    "0809/0809.2140_arXiv.txt": {
        "abstract": "In an earlier paper, we presented the first evidence for a bow-shock nebula surrounding the X-ray binary LMC\\,X-1 on a scale of $\\sim 15$ pc, which we argued was powered by a jet associated with an accretion disk. We now present the first evidence for an ionization cone extending from an X-ray binary, a phenomenon only seen to date in active galactic nuclei (AGN). The ionization cone, detected in the \\Heii$\\lambda4686$/\\Hb\\ and \\Oiii$\\lambda5007$/\\Hb\\ line ratio maps, aligns with the direction of the jet inferred from the bow-shock nebula. The cone has an opening angle $\\approx45\\fdegs$ and radial extent $\\approx 3.8 \\, {\\rm pc}$. Since the \\Heii\\ emission cannot be explained by the companion O star, the gas in the ionization cone must be exposed to the `naked' accretion disk, thereby allowing us to place constraints on the unobservable ionizing spectrum. The energetics of the ionization cone give unambiguous evidence for an ``ultraviolet - soft X-ray\" (XUV) excess in LMC\\,X-1. Any attempt to match the hard X-ray spectrum ($>1$ keV) with a conventional model of the accretion disk fails to account for this XUV component. We propose two likely sources for the observed anisotropy: (1) obscuration by a dusty torus, or (2) a jet-blown hole in a surrounding envelope of circumstellar absorbing material. We discuss the implications of our discovery in the context of the mass-scaling hypothesis for accretion onto black holes and suggest avenues for future research. ",
        "introduction": "The search for a so-called `unified model' of AGN has spanned more than two decades \\citep{ant93}. Indeed, it has been shown that some Type 2 Seyferts may comprise a Type 1 Seyfert nucleus that is heavily obscured by a dense, dusty torus (\\citealt{law82, ant85}). Although dusty tori obscure some of the nuclear continuum emission, a significant fraction escapes along the poles of the central source, producing a cone-shaped extended emission line region. Perhaps the most spectacular example occurs in the Type 2 Seyfert galaxy NGC\\,5252 where bipolar \\Oiii\\ cones are seen to extend to $\\sim$30 kpc in radius \\citep{tad89}. This important phenomenon allows one to test the unified model by analyzing the ionization properties of the gas within the cones.  Many of the phenomena we observe from active galaxies arise from the properties of the accretion disk. \\citet{ede99} note that the power density spectrum of an active galaxy resembles those observed in some Galactic black hole X-ray binaries (XRBs), such as Cyg\\,X-1 (see also \\citealt{mch89}). The transient behaviour of XRBs, however, may render `the scaling from stellar-to-supermassive black hole mass' a false argument \\citep{don05}. But if appropriate corrections are made for the variations in the accretion rate, the variability timescales may be related \\citep{mch06}.  If the mass-scaling hypothesis for black holes is correct, it is natural to question why we have not observed ionization cones associated with XRBs in previous work. This may reflect the difficulty of interpreting the complex nebulosity in the vicinity of XRBs due to a background of competing sources, e.g. hot young stars. The XRB LMC\\,X-1 is surrounded by a complex filamentary emission-line nebula, catalogued by \\citet{hen56} as N159F. The X-ray source ($L_{X}[2-10\\,{\\rm keV}] \\approx2\\times10^{38}\\, {\\rm erg\\,\\,s}^{-1}$, \\citealt{sch94}) comprises a stellar mass ($4-10 \\,M_\\odot$) black hole accreting matter from an O7\\,III-type star (``Star 32''). Based on the \\Heii$\\lambda4686$ emission line, \\citet{pak86} demonstrated that N159F was the first example of an X-ray ionized nebula.  \\citet{ram06} supported the presence of an X-ray ionized nebula, and concluded that LMC\\,X-1 `is not injecting a significant amount of mechanical energy into the interstellar medium,' contrary to the findings of \\citet{coo07} (hereafter, \\papone), who recently suggested that the nebula surrounding LMC\\,X-1 is largely driven by a jet emanating from the XRB. The latter study used integral field spectroscopy (IFS) to delineate the complex ionization regions within the nebula. The two studies prior to \\papone\\ used 1D slit spectroscopy, a technique that is spatially more limited than IFS.  Once again we appeal to IFS to obtain a complete spatial coverage of the nebula. The motivation behind our new study is to assess the extent of hard ionizing photons by tracing the \\Heii\\ and \\Oiii\\ emission. In the course of our work, we have discovered an ionization cone associated with this XRB, the first of its kind. Moreover, the ionization cone aligns with the putative jet from \\papone, a phenomenon also observed in AGNs \\citep{ung87}. Not only does this provide a new ground for testing accretion disk phenomena, but it also strengthens the mass-scaling hypothesis, implying that accretion processes around stellar-mass black holes are similar to accretion processes around supermassive black holes.  We summarize the observations and data reduction procedures in \\S\\,2. In \\S\\,3, we derive the cone parameters and discuss the ionizing source. Our interpretation and conclusions are summarized in \\S\\,4. ",
        "conclusions": "By comparing Eq. (1)-(4), the Yao model underestimates the number of ionizing photons emerging from the accretion disk by a factor of $7-10$.  This discrepancy requires an extra component that dominates in the UV and/or soft X-ray bands.  In the Yao model, we cannot increase $K_{\\rm MCD}$ because it is constrained beyond $1\\,{\\rm keV}$ by the observed hard X-ray spectrum (see Fig.~\\ref{fig:xray_spectrum}), which is modelled as Comptonized disk emission.  We note that the size of the discrepancy arising from Eq.~(1)-(4) is exaggerated by the use of a photon number; the energy requirement is only a factor of two more than the Yao model. This could conceivably arise from upscattering of lower energy photons from more distant regions in the disk by a hot coronal plasma in the inner region. Alternatively, a sizeable fraction of the hard X-ray photons could be degraded to lower energies, i.e. in the opposite sense to what is assumed in the CMCD model. It is noteworthy that some AGN ionization cones can only be explained by the contribution of both ionization from the non-stellar nuclear continuum and jet-induced shock ionization \\citep*{pog88,wil88}. Therefore an {\\it in situ} jet may also contribute to the LMC\\,X-1 cone, although the nebular temperature limits any contribution from shock processes.  EUV/soft X-ray excesses are not uncommon to AGN. The spectra of high redshift quasars, which can be observed in the UV, clearly exhibit a big blue bump component, widely attributed to accretion disk emission \\citep{san89}. This component has also been inferred from the emission-line diagnostics of Seyfert 2 ionization cones \\citep{ale00}. In some cases, EUV/soft X-ray excesses are \\emph{also} observed. The origin of this component is unclear, but it may arise, for example, from thermal emission in a warm/hot cloud at a temperature that peaks in the $30-300\\,$eV range \\citep{sie95}. How this component relates to the accretion disk is not clear.  It is unlikely that the angular extent of the ionization cone reflects the poloidal ionizing field of the naked accretion disk. There are two plausible explanations for the anisotropic radiation pattern: (1) obscuration by a dusty torus, or (2) a jet-blown hole in the envelope of circumstellar absorbing material \\citep*{kuu05,nes08}.  We suspect that obscuring XRB tori would scale to AGN tori with the mass accretion rate rather than the black hole mass. The mass accretion rate in XRBs depends on the donor star, whereas for AGN, it depends on the gas fuel supply from the surrounding interstellar medium. Because Star 32 is a giant, we would expect an excess mass flux to accumulate somewhere before reaching the accretion disk. This accumulated mass flux could form something akin to an obscuring torus. By analogy with AGN \\citep{pie93}, high angular resolution, mid-infrared observations of LMC\\,X-1 are required to confirm the presence of a dusty torus.  Here we have demonstrated the importance of ionization cones for establishing the XUV properties of obscured accretion disks. We suspect the ionization cone will also be observable in \\Neiii$\\lambda3868$ and \\Nev$\\lambda3426$ (ionization potential $63.5 \\, {\\rm eV}$ and $126.2 \\, {\\rm eV}$ respectively). These arise at higher $q$ values and will therefore allow us to probe down to smaller radii. Indeed, \\citet{pak86} have already detected these two high ionization species in the immediate vicinity of LMC\\,X-1. But due to the limited spatial coverage of their spectroscopic technique, we are unable to confirm that these ratios are enhanced in the direction of the ionization cone. Planned future observations will probe the structure of the ionization cone, allowing us to refine our ionization model and place tighter constraints on the nature of the ionizing source."
    },
    "0809/0809.3977_arXiv.txt": {
        "abstract": "\\indent The question as to whether the distribution of radio loudness in active galactic nuclei (AGN) is actually bimodal has been discussed extensively in the literature.  Furthermore, there have been claims that radio loudness depends on black hole mass ($M_{BH}$) and Eddington ratio ($L_{bol}/L_{Edd}$).  We investigate these claims using the low redshift broad line AGN sample of \\citet{Greene2007}, which consists of 8434 objects at $z < 0.35$ from the Sloan Digital Sky Survey Fourth Data Release (SDSS DR4).  We obtained radio fluxes from the Very Large Array Faint Images of the Radio Sky at Twenty-Centimeters (FIRST) survey for the SDSS AGN.  Out of the 8434 SDSS AGN, 821 have corresponding observed radio fluxes in the FIRST survey.  We calculated the radio-loudness parameter ($\\mathcal{R}$) for all objects above the FIRST detection limit (1 mJy), and an upper limit to $\\mathcal{R}$ for the undetected objects.  Using these data, the question of radio bimodality is investigated for different subsets of the total sample.  We find no clear demarcation between the radio loud (RL, $\\mathcal{R}$ $>$ 10) and radio quiet (RQ, $\\mathcal{R}$ $<$ 10) objects, but instead fill in a more radio-intermediate population in a continuous fashion for all subsamples.  We find that 4.7\\% of the AGN in the flux-limited subsample are RL based on core radio emission alone.  We calculate the radio-loud fraction (RLF) as both a function of $M_{BH}$ and $L_{bol}/L_{Edd}$.  The RLF decreases (from 13\\% to 2\\%) as $L_{bol} / L_{Edd}$ increases over 2.5 orders of magnitude.  The RLF is nearly constant ($\\sim$5\\%) over 4 decades in $M_{BH}$, except for an increase at $M_{BH} > 10^{8} M_{\\odot}$.  We find for the FIRST detected subsample that 367 of the RL AGN have $M_{BH} < 10^{8} M_{\\odot}$, a large enough number to indicate that RL AGN are not a product of only the most massive black holes in the local universe. ",
        "introduction": "\\indent The nomenclature for classifying active galactic nuclei (AGN) relies on many different factors.  The intrinsic power separates out the lower luminosity Seyfert galaxies from the higher luminosity quasars, whose optical spectra can be very similar.  There is the separation of types 1 and 2 based on the presence or absence (respectively) of broad-lines (BL), an effect presumably caused by the obscuration of the broad-line region (BLR) in type 2 objects due to the presence of apparently toroidally distributed dusty matter in our line of sight.  There is also a variety of different radio properties that yield interesting subcategories.  In general there are the radio quiet (RQ) and the radio loud (RL) populations (based on the ratio of radio to optical flux).  Most RL AGN are found in radio surveys and are further subdivided into Fanaroff-Riley classes (FRI and FRII), where the radio emission from the former mostly arises from the jets and core while the latter are lobe dominated.  There is also a subtle connection to morphology, where most strong emission-line radio galaxies are found in giant elliptical galaxies and most Seyfert galaxies are in spirals \\citep{Osterbrock2006,Sikora2007}.  Moving beyond this classification system into one that is more continuous and based on fundamental parameters such as bolometric luminosity ($L_{bol}$, relates directly to accretion rate), black hole mass ($M_{BH}$, relates directly to Eddington luminosity, $L_{Edd}$), Eddington ratio ($L_{bol}/L_{Edd}$, the normalized accretion rate), and inclination (\\it i\\rm, accounts for the presence or absence of broad-line features) could simplify the subject greatly.  In recent years our ability to determine $M_{BH}$, and hence $L_{bol}/L_{Edd}$, has greatly improved and facilitated this goal.  The technique of reverberation-mapping has been used to measure the radii of the broad-line regions (R$_{BLR}$) from which the virial masses of supermassive black holes (SMBHs) can be calculated \\citep{Peterson1997book, Peterson2004, Kaspi2000, Kaspi2005} using \\begin{equation} M_{BH}=\\frac{fR_{BLR}{\\Delta}V^2}{G}, \\end{equation}  where \\it f \\rm is a scaling factor that depends on the geometry and kinematics of the BLR and  $\\Delta V$ is the velocity dispersion. Empirical relations have been found to relate $R_{BLR}$ to the optical luminosity at 5100 \\AA\\ ($L_{\\lambda}$[5100 \\AA]), allowing one simple measurement to be used as a proxy for $R_{BLR}$ \\citep{Peterson2004, Kaspi2005}.  There is some debate over the empirical power law that governs this relation, and a few different scalings are used, but recent studies have $R_{BLR} \\propto (L_{\\lambda}$[5100 \\AA])$^\\gamma$ where $\\gamma \\approx 0.5$ \\citep{Vestergaard2006}.  This is the (zeroth order) theoretical expectation obtained by assuming a constant ionization parameter (U) and constant electron density ($n_{e}$) for the BLR in all AGN \\citep{Peterson1997book}.  The full width at half maximum (FWHM) of the broad-lines gives a measure of the velocity dispersion of the BLR clouds due to their assumed Keplerian orbits around the BH.  These two combined measurements ($L_{\\lambda}$[5100 \\AA], FWHM) from a single spectrum can give a relatively reliable estimate of $M_{BH}$ \\citep{Vestergaard2006}.  It has been long claimed by previous studies that RL AGN make up only about 5\\% -- 10\\% of the total AGN population \\citep{Kellermann1989, Urry1995, Ivezic2002}.  This has led to the speculation that RQ AGN are different from RL AGN because of different physical accretion processes \\citep{Ho2002, Sikora2007}.  The idea of two separate populations is based on the claim of a bimodal distribution in the radio-loudness parameter ($\\mathcal{R}$) of AGN \\citep{Kellermann1999};  $\\mathcal{R}$ is simply the ratio of the monochromatic radio luminosity to the monochromatic optical luminosity ($\\mathcal{R} \\equiv \\nu_{radio}L_{radio} / \\nu_{opt}L_{opt}$).  Historically, objects are considered to be RL if $\\mathcal{R} >$ 10, and RQ if $\\mathcal{R} <$ 10 \\citep{Kellermann1989}.  Typically the radio luminosity is measured at 5 GHz, and the optical luminosity is measured at 4400 \\AA.  One key difference in the current study is that the Very Large Array's (VLA) Faint Images of the Radio Sky at Twenty-centimeters (FIRST) survey operates at 1.4 GHz so the $\\mathcal{R}$ value will require a minor scaling.  Corrections to account for different scalings will be discussed in \\S2.  Several investigations into this claim of bimodality have yielded results that differ, based on the AGN sample, selection criteria, and inclusion or exclusion of extended radio emission \\citep{White2000, Ivezic2002, Ivezic2004, Cirasuolo2003, Cirasuolo2004, Laor2003, White2007}.  The Sloan Digital Sky Survey (SDSS) allows the optical selection of a large number of AGN with varying types of classifications over a large range in redshifts \\citep{Schneider2007}.  Large scale radio surveys like the FIRST survey have allowed characterization of a large number of radio sources at a flux limit of 1 mJy and an angular resolution of 5$''$.  Due to its high resolution, the FIRST survey is primarily sensitive to core emission.  The number of objects with a resolved complex radio morphology (core-jet, core-lobe, double-lobed) in the FIRST survey with an SDSS optical counterpart is less than 10\\% \\citep{Ivezic2002}, although it is likely that some of the unresolved sources have small, and presumably young, jet or lobe contributions.  In this paper we take the \\citet{Greene2007} sample of low redshift BL AGN from the SDSS, from which $M_{BH}$ and $L_{bol}/L_{Edd}$ can be calculated, and correlate them with their FIRST counterparts to test for bimodality in $\\mathcal{R}$ and the dependence of $\\mathcal{R}$ on $M_{BH}$.  In this way we can determine the radio properties of AGN in the local universe and compare them with higher redshift samples.  We give a description of the SDSS and FIRST data samples in \\S2.  In \\S3 we describe our data analysis and biases in the sample.  We present conclusions in \\S4. ",
        "conclusions": "We have taken 8434 optically selected BL AGN from the SDSS and estimated $M_{BH}$ and $L_{bol}/L_{Edd}$ for them.  Using fluxes taken from the FIRST survey we compute $\\mathcal{R}$, the radio-loudness parameter, where the radio emission is thought to be mostly from the core.  We note that because of the resolution of the FIRST survey and our small search radius (4$''$), we miss AGN with extended radio structures that have very weak core emission.  We note that $\\sim$90\\% of all SDSS/FIRST matched AGN have only core radio emission \\citep{Ivezic2002}, although because of the resolution limit of FIRST, many of these `core' sources will include compact jets and lobes.  Furthermore, the higher redshift samples with which we are comparing our results \\citep{White2000, Ivezic2002, Cirasuolo2004} are also based on core radio emission from FIRST.  In the flux-limited subsample we find that 4.7\\% of the AGN are RL.  This is similar to the approximately 5\\% found for higher redshift quasars, but may be coincidental in that the distribution of radio-loudness likely depends on the available range of Eddington ratio and host galaxy morphology in the sample.  The issue of extended radio structure and its impact on the radio-loudness distribution will be addressed in a forthcoming paper.  The radio-loudness histograms in Figures 4 \\& 5 both peak in the same place, at log($\\mathcal{R}$) $\\approx$ 0.75.  We find no evidence for a bimodal histogram for any subset of the total sample.  We do note that the upper limit (dashed-line) histograms could be pushed farther to the left to reveal a bimodal distribution, but at much lower $\\mathcal{R}$ values than previously found by \\citet{Ivezic2004} and \\citet{Cirasuolo2004}, who find peaks at log($\\mathcal{R}$) $\\ll$ 1, and log($\\mathcal{R}$) $\\approx$ 2-3.  After we apply various flux limits to our total sample, we still cannot reproduce a local minimum at log($\\mathcal{R}$) $\\approx$ 1.  In all our subsamples we fill in a substantial number of radio-intermediate (RI) objects.  The minimum in the $\\mathcal{R}$ histogram could be shifted to lower values or could just not be present at all at these low redshifts.  Nevertheless, we still see a RL tail.  Compared to higher redshift samples, the $\\mathcal{R}$ distribution in the local universe for detected sources is narrower and shifted to lower values \\citep{Ivezic2004, White2007}.  This may indicate that the radio luminosity function has a stronger evolution than the optical luminosity function to the current epoch, as suggested by \\citet{Jiang2007}.  \\citet{White2007} find for their sample a similar effect of evolution of the $\\mathcal{R}$ histogram with redshift.  Our plot is in good agreement with their low-redshift ($z < 0.5$) quasars, where there is a peak at low values of $\\mathcal{R}$, and a tail that falls off over 4 orders of magnitude, with no indication of two separate populations.  Our plot of $\\mathcal{R}$ as a function of $L_{bol} / L_{Edd}$ lacks points in the low accretion/RQ and high accretion/RL regimes.  We also see a trend of higher RLF at the low values of $L_{bol} / L_{Edd}$.  This matches the general trend seen by \\citet{Ho2002} and \\citet{Sikora2007} of decreasing $\\mathcal{R}$ with increasing L$_{bol}$/L$_{Edd}$, based on AGN that extend to much lower L$_{bol}$/L$_{Edd}$ ($\\sim$10$^{-7}$).  We do not find two separate populations at a given $L_{bol}/L_{Edd}$, but again fill in a more RI population.  This may be due to the homogeneous optical selection of our data set, as opposed to the admittedly heterogeneous selection criteria of AGN in \\citet{Sikora2007}, which includes AGN from various samples, where objects were selected in very different ways.  We do note however that we do not cover nearly as much parameter space as Sikora et al.\\ who get up to an $\\mathcal{R}$ of $\\sim 10^{7}$ and down to an $L_{bol}/L_{Edd}$ of $\\sim$ 10$^{-7}$.  Our sample does not go down to the very low $L_{bol}/L_{Edd}$ values due to the SDSS flux limit.  Homogeneous samples like ours that extend to much fainter radio and optical fluxes would be extremely useful for further studies.  In fact, going deeper in both the radio and optical is necessary if one wants to search for AGN in the very lowest $L_{bol}/L_{Edd}$ and $\\mathcal{R}$ regimes.  We find a substantial number of RL AGN with $M_{BH} < 10^{8} M_{\\odot}$, as seen in Figure 10 (although the fraction of RL sources is largest for the most massive BHs in our sample as shown in Fig.\\ 13).  This clearly shows that in the local universe, not only massive BHs ($M_{BH} > 10^{8} M_{\\odot}$) are responsible for RL AGN, which is consistent with the findings of \\citet{Ho2002} and \\citet{Woo2002}."
    },
    "0809/0809.3831_arXiv.txt": {
        "abstract": "We report total abundances and related parameters for the full sample of the {\\it FUSE} survey of molecular hydrogen in 38 translucent lines of sight.  New results are presented for the ``second half'' of the survey involving 15 lines of sight to supplement data for the first 23 lines of sight already published.  We assess the correlations between molecular hydrogen and various extinction parameters in the full sample, which covers a broader range of conditions than the initial sample.  In particular, we are now able to confirm that many, but not all, lines of sight with shallow far-UV extinction curves and large values of the total-to-selective extinction ratio, $R_V$ = $A_V$/$E(B-V)$ --- characteristic of larger than average dust grains --- are associated with particularly low hydrogen molecular fractions ($f_{\\rm H2}$).  In the lines of sight with large $R_V$, there is in fact a wide range in molecular fractions, despite the expectation that the larger grains should lead to less H$_2$ formation.  However, we see specific evidence that the molecular fractions in this sub-sample are inversely related to the estimated strength of the UV radiation field and thus the latter factor is more important in this regime.  We have provided an update to previous values of the gas-to-dust ratio, $N$(H$_{\\rm tot}$)/$E(B-V)$, based on direct measurements of $N$(H$_2$) and $N$(H I).  Although our value is nearly identical to that found with {\\it Copernicus} data, it extends the relationship by a factor of 2 in reddening. Finally, as the new lines of sight generally show low to moderate molecular fractions, we still find little evidence for single monolithic ``translucent clouds'' with $f_{\\rm H2}$ $\\sim$ 1. ",
        "introduction": "Molecular hydrogen is the most abundant molecule in the universe, and a detailed knowledge of H$_2$ is crucial for a full understanding of the physics and chemistry of the interstellar medium.  A broad overview of interstellar H$_2$ is provided by Shull \\& Beckwith (1982), while a recent overview of chemical processes in diffuse and translucent clouds is provided by Snow \\& McCall (2006).  A major goal of the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) was a survey of molecular hydrogen in 45 lines of sight with an emphasis on interstellar clouds with as much extinction as possible.  The extinctions covered the range of the so-called ``translucent clouds'' with $A_V$ in the range 1--5 mag.  Within these limits, lines of sight were chosen from a variety of environments and dust characteristics.  The first results from the study were presented by Rachford et al.\\ (2002 [Paper I]).  Paper I contains the details regarding the previous history of H$_2$ observations and the rationale for the {\\it FUSE} survey. Slightly more than one-half of the planned targets had been observed and analyzed and those lines of sight were detailed.  A main finding was that in most cases a line of sight was likely composed of multiple clouds, suggesting a change in terminology to ``translucent lines of sight'' pending the clear identification of a high-extinction line of sight made up of one highly molecular cloud.  The main purpose of the present paper is to provide the community with the overall H$_2$ results for the full sample, as well as refine and strengthen some of the conclusions of Paper I.  The additional lines of sight give us much better coverage of the variety of dust characteristics than the original partial sample and we emphasize those results.  Also, we have taken advantage of the publication of the 2MASS All-Sky Survey of Point Sources (Skrustkie et al.\\ 2006) to provide more reliable extinction parameters for the full sample.  The rest of the paper is organized as follows: in \\S\\ 2 we describe the remaining target for the FUSE translucent H$_2$ survey, along with comments on the stars chosen; in \\S\\ 3 we give the details of the observations, data reduction, and analysis; in \\S\\ 4 we discuss the results and their implications; and we summarize the paper in \\S\\ 5. ",
        "conclusions": "\\subsection{Overall comments} Table 5 gives our primary results for the lines of sight.  The H$_2$ column densities were measured directly from the spectra as already described.  We also include two derived quantities, the hydrogen molecular fraction ($f_{\\rm H2}$), and the rotational temperature ($T_{01}$) for each line of sight.  The hydrogen molecular faction is defined in terms of the molecular and atomic hydrogen column densities as  \\begin{equation} f_{\\rm H2} = \\frac{2N(H_2)}{2N(H_2) + N(H I)}. \\end{equation}  The rotational temperature (in Kelvin) is determined by applying the Boltzmann equation to the ratio of the column densities in the first two rotational states, yielding  \\begin{equation} \\frac{N(1)}{N(0)} = 9e^{-171/T_{01}}. \\end{equation}  Solving for the temperature and expressing the column densities as base-10 logarithms gives  \\begin{equation} T_{01} = \\frac{74.0}{\\log N(0) - \\log N(1) + 0.954}. \\end{equation}  As in Paper I, we generally interpret the ratio between $N$(1) (ortho-H$_2$) and $N$(0) (para-H$_2$) as the kinetic temperature on the assumption that collisions with H$^{+}$ (and H3$^{+}$ when enough is present) dominate over other processes in controlling this ratio.  There is evidence that under some circumstances ortho-H$_2$ can be rapidly converted to para-H$_2$ on grains (Le Bourlot 2000). In this case, the $N$(1)/$N$(0) ratio becomes lower, yielding a lower rotational temperature.  This process appears to be more likely at the low temperatures in the core of a relatively dense and opaque cloud (Shaw et al.\\ 2005).  The fact that we seem to be seeing multiple diffuse clouds along the lines of sight (Paper I) suggests that this process may not important for our sample.  Shaw et al.\\ (2005) have modeled the physical conditions in one of our lines of sight from Paper I, HD 185418.  Their model includes the Le Bourlot (2000) treatment of the ortho to para conversion process.  Their derived kinetic temperature for the gas was about 25\\% lower than than our derived $T_{01}$ = 101 $\\pm$ 14 K.  This line of sight has the largest $T_{01}$ of our entire sample, and the gas density is relatively low ($n_H$ = 27 cm$^{-3}$, but Sonnentrucker et al.\\ 2003 find an even lower value of $n_H$ = 6.3 cm$^{-3}$).  Lacking a complete knowledge of whether or not our lines of sight may have non-thermal $N$(1)/$N$(0) ratios, we will generally assume that the measured $T_{01}$ values correlate with kinetic temperature in some manner.  However, this potential uncertainty should be kept in mind in the interpretation of correlations of $T_{01}$ with other parameters.  Overall, the second half of the sample shows column densities, molecular fractions, and temperatures similar to those in Paper I.  The primary difference is the presence of a sample of lines of sight with large extinction, small N(H$_2$), and large $R_V$, which we will discuss in detail in \\S\\ 4.3.  \\subsection{General correlations with reddening} We have updated several correlations with reddening from Paper I using the new data.  In Figure 1 we present a plot of the H$_2$ column density versus color excess (unlike Figure 2 in Paper I, we give $N(H_2)$ on a linear scale).  While this plot shows the expected increase in H$_2$ as we look through more material, it also reflects that our targets probe a wide variety of environments in that we usually see a range of column densities at a given color excess.  One particularly important relationship we can investigate with the complete dataset is the gas-to-dust ratio.  Using {\\it Copernicus} data, Bohlin, et al.\\ (1978) found $N$(H$_{\\rm tot}$)/$E(B-V)$ = 5.8 $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$.  In Figure 2, we present plots of $N$(H$_{\\rm tot}$) versus $E(B-V)$.  In the upper panel we include the Bohlin et al.\\ {\\it Copernicus} data and our {\\it FUSE} data.  In the lower panel we only include our {\\it FUSE} data for lines of sight with direct determinations of $N$(H I) and $N$(H$_2$), and also excluded Be stars as the color excesses might be overstated due to the circumstellar emission.  This should give the most homogeneous sub-sample possible that covers a broad range in color excess.  When constrained to pass through the origin, our error-weighted best-fit slope is (5.94 $\\pm$ 0.37) $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$, essentially identical to the {\\it Copernicus} value for less reddened lines of sight.  The solid line in both panels is the best fit, and it is worth noting that many of the most discrepant {\\it FUSE} points in the upper panel are Be stars, which are removed in the bottom panel.  For reference, we have also included an unconstrained fit in these panels which furthermore does not include the low-reddening point at $E(B-V)$ = 0.17 (HD 186994).  This illustrates the significance of constraining the fit with points at small reddening.  Clearly this line is not a good fit to the low-extinction {\\it Copernicus} points.  One possibility is that the high-extinction sample has a different slope than the low-extinction sample.  This would represent a difference in the gas-to-dust ratio that might indicate a change in dust properties in the clouds with higher extinction.  Such a difference is not clearly seen in our data.  The slope of the dashed line is (4.2 $\\pm$ 1.7) $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$, less than the value for the constrained fit, but with large enough uncertainty to be consistent with the constrained fit.  We performed the same analysis for the total visual extinction, $A_V$, which is simply the product of $E(B-V)$ and $R_V$. Figure 3 shows panels corresponding to the ones shown in Figure 2.  Again, we have fitted a line to our ``best'' sub-sample in the lower panel, which yields $N$(H$_{\\rm tot}$)/$A_V$ = (2.15 $\\pm$ 0.14) $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$.  This value is nearly identical to a simple division of 5.94 $\\times$ 10$^{21}$ by the average $R_V$ = 2.93 for this sub-sample; the latter value is slightly less than the galactic average of 3.1 (Draine 2003).  Again, excluding HD 186994 and not constraining the fit to pass through the origin gives a slightly smaller slope of (1.50 $\\pm$ 0.43) $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$.  Visually, there is a hint of a slope change in the $A_V$ = 1.5--2.0 interval in the upper panel, but with so few reliable points with $A_V$ $\\geq$ 2, such a trend is not clear.  As Figure 4 shows (an update of Figure 5 in Paper I), there appears to be a weak inverse correlation between rotational temperature and reddening.  However, it is critical to note that much of this trend disappears if we were to exclude the {\\it Copernicus} data points with $N$(H$_2$) $<$ 10$^{20}$ cm$^{-2}$. In the region of overlap with {\\it Copernicus} data, the {\\it FUSE} lines of sight show smaller average temperatures, while the {\\it FUSE} data for large reddening do not deviate from the similar data for small reddening.  The mean $T_{01}$ for our entire {\\it FUSE} sample is 67 K with a standard deviation of 14 K.  This result is similar to other studies, including the $<T_{01}>$ = 86 $\\pm$ 20 K from the {\\it FUSE} Galactic disk survey (J.\\ M.\\ Shull et al., in preparation).  However, the $T_{01}$ distribution of our sample is not normal, having a general rise in frequency up to around 75--80 K, and just a few stars with temperatures above that; the latter fact can easily be seen in Figure 4.  In principle, the larger color excesses could be associated with denser clouds, which in turn might be expected to show lower temperatures.  In Paper I we noted a correlation between the rotational temperature and the fractional abundance of the CN radical, the latter of which is a good density indicator (Federman, et al.\\ 1984).  However, as we discussed in Paper I, the lines of sight in the translucent sample seem to mostly be made up of multiple diffuse clouds.  Thus, while the slight trend is in the right direction to suggest that we are probing somewhat denser clouds in the {\\it FUSE} sample, this conclusion is somewhat speculative.  Figure 5 (an update of Figure 7 in Paper I) shows a similar pattern to Figure 1 in that once H$_2$ becomes self-shielded and relatively abundant, the molecular fraction covers nearly the entire possible range at any given reddening.  Again, part of this may be a selection effect as the lines of sight were chosen to sample a variety of environments.  However, it is notable that even with the additional sample we have not found a line of sight with $f_{\\rm H2}$ $>$ 0.8.  In fact, as seen in Table 5, the new lines of sight preferentially have relatively small $f_{\\rm H2}$ given the large extinctions.   \\subsection{The high $R_V$ sample} \\subsubsection{Significance} An important aspect of the completed translucent sample is that we have determined the molecular hydrogen column densities for several lines of sight with particularly large values of $R_V$ combined with larger $A_V$ and total hydrogen column densities than previous samples.  In some cases, these lines of sight show unusually small molecular fractions and are of particular interest because of the relationship between dust parameters and molecular hydrogen formation and destruction.  Snow (1983) first investigated the possibility that H$_2$ abundances may be inversely correlated with grain size.  Larger grains, through the coagulation of small grains, provide less surface area per unit dust mass.  With less surface area available, the H$_2$ formation rate should be diminished.  Using the compilation of {\\it Copernicus} data of Bohlin et al. (1978), Snow demonstrated that the mean molecular fraction in the $\\rho$ Oph cloud was a factor of 2.6 less than the rest of the sample.  Cardelli (1988) explored the related issue of the relationship between H$_2$ abundances and $R_V$.  The H$_2$ abundances again came from Bohlin et al. (1978), while the values of $R_V$ came from an analysis of IR photometry similar to that which we have applied in the present work.  Cardelli's sample displayed an inverse relationship between the H$_2$ to $A_V$ ratio and $R_V$, with a weaker, but similar dependence of the hydrogen molecular fraction and $R_V$.  When splitting the {\\it Copernicus} lines of sight at $R_V$ = 3.5, the high $R_V$ group had an average molecular fraction of approximately a factor of 2.5 less than the low $R_V$ group.  As $R_V$ is thought to be positively correlated with grain size (or grain size distribution) this result provides further support for the role of grain size in regulating the H$_2$ abundance. Furthermore, the large values of $R_V$ are associated with smaller than normal far-UV extinction, which allows more photodissociating radiation to penetrate the interstellar clouds and provides further reduction in the hydrogen molecular fraction.  In fact, it appears that only the lines of sight with the largest $R_V$ have extinction curves that deviate in a consistent manner from average (Fitzpatrick \\& Massa 2007).  In Paper I, we had minimal coverage of high $R_V$ lines of sight in our {\\it FUSE} sample, but the newly added targets improve the situation.  In Figure 6 we show the relationship between $f_{\\rm H2}$ and $R_V$ for the {\\it FUSE} and {\\it Copernicus} samples.  This provides a major update to Figure 8 in Paper I as we have also included the {\\it Copernicus} sample, using 2MASS photometry to derive the values of $R_V$ as with the {\\it FUSE} sample.  Interestingly, using the Spearman rank correlation coefficient, at best we only see an anti-correlation between molecular fraction and $R_V$ at about the 1 $\\sigma$ level whether or not we restrict the sample to the best determinations of either parameter.  A more significant trend (3 $\\sigma$) appears when considering the H$_2$ to $A_V$ ratio versus $R_V$ similar to Cardelli's result, but as he discussed, this division works best when one can make the assumptions of similar average densities in the various clouds and that most of the extinction occurs in the regions where H$_2$ is significantly present.  However, it is not clear that this is an appropriate way to look at the {\\it FUSE} translucent sample.  In particular, much of the trend disappears if we exclude lines of sight with $A_V$ $>$ 2.  As discussed in Paper I, we believe that most of our lines of sight are sampling multiple clouds that may be spread out in space, and thus may have a variety of conditions.  We wish to consider in more detail the lines of sight with $R_V$ $>$ 4, which are labeled in Figure 6, split equally between the {\\it FUSE} and {\\it Copernicus} samples.  We will ignore the {\\it Copernicus} target HD 10516 ($\\phi$ Per) due to the very large uncertainty in $R_V$ and small extinction ($E(B-V)$ $\\approx$ 0.2 and $A_V$ $\\sim$ 1).  We will thus focus our discussion on the five remaining lines of sight with high $R_V$ and large extinction.  Detailed modeling of an individual line of sight is complex and requires a large amount of data to constrain the model, although this has been done for a few favorable lines of sight from Paper I (Rachford et al.\\ 2001; Sonnentrucker et al.\\ 2003).  Barring such an analysis, we can get a sense of the strength of the local UV radiation field using the high $J$ lines of H$_2$, and thus an estimate of the importance of photodissociation.  A preliminary analysis of lines from $J$ = 2 up to the highest observable levels in a line of sight yields column densities that can be used to estimate the amount of radiative pumping to these levels.  This has proven challenging in many of the translucent lines of sight due primarily to data quality and the resulting difficulty in measuring weak lines and properly interpreting saturated lines.  This is particularly acute in these heavily reddened lines of sight as in many cases the S/N rapidly decreases with decreasing wavelength.  Also, the strengths of other lines are typically larger than in less reddened lines of sight which causes more interference with the high-$J$ lines.  However, we have enough information for most of these five lines of sight to use $N$(6) and $N$(7) as potential indicators of the strength of the radiation field.  These lines are typically weakly saturated and not as difficult to interpret as the stronger lines from lower rotational states, when the data quality permits their detection or the derivation of well constrained upper limits.  For reference, the column densities for HD 110432 (Rachford et al.\\  2001) were log $N$(6) = 14.20 $\\pm$ 0.20 and log $N$(7) = 13.25$^{+1.25}_{-1.00}$ for a radiation field that was modeled as twice the strength of the average curve of Draine (1978).  Particle density also plays a significant role in controlling H$_2$ excitation, so we should look at the high $J$ column densities as an indicator of the strength of the radiation field, but not as a definitive measurement.  \\subsubsection{HD 38087} HD 38087 has $f_{\\rm H2}$ greater than half, the largest of the group, but this value is uncertain because we do not have a direct measurement of interstellar N(H I) due to the very late spectral type.  Fitzpatrick \\& Massa (1990) report log N(H I) = 21.48 from Lyman $\\alpha$.  However, at spectral type B5 V, the interstellar line is strongly contaminated or even dominated by the stellar line (Shull \\& van Steenberg 1985; Diplas \\& Savage 1994).  Atomic hydrogen column density is highly correlated with the strength of certain diffuse interstellar bands, including $\\lambda$5780 (S.\\ D.\\ Friedman et al.\\ 2008, in preparation; see Herbig 1993 or Welty et al.\\ 2006 for similar correlations), and applying this correlation to HD 38087 gives log N(H I) = 21.08, similar to our quoted value and which would only reduce the derived molecular fraction from 0.52 to 0.42.  Thus, it seems likely that this line of sight is genuinely rich in molecules.  We have already noted in \\S\\ 2.3.2 that this line of sight samples a quiescent environment where grain growth might be expected to occur (Snow \\& Witt 1989).  However, only about 25\\% of the line of sight material is thought to be in the reflection nebula.  The extinction curve indicates that the interstellar material should be relatively transparent to photodissociating far-UV radiation -- if such radiation exists at the cloud location.  HD 38087 itself is of rather late spectral type, so it may not contribute significantly to the far-UV radiation field at the location of the bulk of the line of sight material. However, the {\\it FUSE} spectrum clearly shows H$_2$ lines up to the $J$ = 7 level and we derive a logarithmic column density of about 15.0 in that level, a very large value that implies a significant excitation mechanism.  It is important to note that even a ``large'' quantity of excited H$_2$ corresponds to a very small fraction of the total H$_2$ column density.  One simply needs a particularly high excitation temperature to produce a relatively ``flat'' rotational distribution with a high percentage of the H$_2$ in the excited states, while keeping the total column density of this material several orders of magnitude below the 10$^{20-21}$ cm $^{-2}$ totals seen in the translucent lines of sight.  It thus may be the case that the excited H$_2$ is produced in the reflection nebula, while the bulk of the low-excitation H$_2$ is found farther from the star where photodissociation is not as important.  \\subsubsection{HD 148184} The other line of sight from the high $R_V$ sample with significant molecular material is HD 148184 ($\\chi$ Oph), with $f_{\\rm H2}$ = 0.33 (Bohlin et al.\\ 1978). This star is about 5 degrees from the $\\rho$ Oph grouping and samples material from the general Sco-Oph cloud complex.  The {\\it Hipparcos} parallax is 6.21 $\\pm$ 0.23 mas (van Leeuwen 2007), corresponding to a distance of 161 $\\pm$ 6 pc, which places it near the distant edge of the interstellar material. The star itself has spectral type B2 IVpe and may contribute significant UV radiation to the material in the line of sight, nearly all of which is likely associated with the cloud complex. However, in the overall sense this portion of the complex is not as highly populated by B-type stars as in the immediate $\\rho$ Oph area.  This star was observed with {\\it Copernicus} at high resolution, but with very limited wavelength coverage.  Frisch (1980) reported upper limits for two $J$ = 6 lines, yielding an upper limit of log $N$(6) $<$ 14.33.  Thus, there does not appear to be an unusually large source of excitation.  One final important consideration for this star is that it is the only one of the five which is a Be star and thus the caveats given in \\S\\ 2.2 apply.  As the large error bar in Figure 6 implies, we found the color excesses to be a poor fit to the Martin \\& Whittet (1990) relation.  There is no available extinction curve for this star.  The wavelength of maximum polarization is 0.55 $\\mu$m (Coyne et al. 1974), corresponding to $R_V$ = 3.08.  It is thus possible that this line of sight should not be in the high $R_V$ group.  \\subsubsection{The $\\rho$ Oph cloud: $\\rho$ Oph A \\& D} The $\\rho$ Oph cloud has long been known as a location which exhibits large $R_V$ and larger than normal dust grains (Carrasco, et al.\\ 1973; Whittet 1974).  The name $\\rho$ Oph is applied to four stars within a few arcminutes of each other labeled A through D.  The A component (HD 147933) is the brightest star in both the visible and UV and was observed by {\\it Copernicus}.  The D component (HD 147888) is part of the present {\\it FUSE} sample.  Improved {\\it Hipparcos} distances to these stars (van Leeuwen 2007) have been used as part of a study of the distribution of interstellar material in the cloud by Snow et al.\\ (2008).  These distances are 111$^{+12}_{-10}$ pc for $\\rho$ Oph A and 125$^{+14}_{-11}$ pc for $\\rho$ Oph D.  The overall analysis indicates that both stars are in front of the denser material that is sampled by the more distant and more heavily obscured line of sight to HD 147889 that we had hoped to study with {\\it FUSE} as noted in \\S\\ 2.1.  The uncertainties in the {\\it Copernicus} measurements of $N$(H$_2$) and $N$(H I) for $\\rho$ Oph A are relatively large, but still strongly indicate a small molecular fraction. Unfortunately, the spectral type of $\\rho$ Oph D (B5 V) is late enough that the interstellar Ly$\\alpha$ line will be severely contaminated by the stellar line, as with HD 38087, thus $N$(H I) is very uncertain.  Cartledge et al.\\ (2004) indirectly estimated a value log $N$(H$_{\\rm tot}$) = 21.73 $\\pm$ 0.09 based on measurements towards nearby $\\rho$ Oph A.  For this paper, we have used our preferred method for stars later than spectral type B2, namely, applying the Bohlin et al. (1978) $N$(H$_{\\rm tot}$)/$E(B-V)$ = 5.8 $\\times$ 10$^{21}$ atoms cm$^{-2}$ mag$^{-1}$ value we discuss in \\S\\ 4.2.  This gives log $N$(H$_{\\rm tot}$) = 21.44, much smaller than the value for $\\rho$ Oph A, but in excellent agreement with the correlation between $N$(H I) and the equivalent width of the $\\lambda$5780 DIB discussed in \\S\\ 4.3.2.  It should be noted that $\\rho$ Oph A is well known as a line of sight with a larger than normal gas-to-dust ratio (e.g.\\ Bohlin et al.\\ 1978).  For that reason, our derived value of log $N$(H$_{\\rm tot}$) for $\\rho$ Oph D may also be too low.  A larger value of $N$(H$_{\\rm tot}$) would produce an even smaller molecular fraction than the $f_{\\rm H2}$ = 0.21 depicted in Figure 6.  Thus, it appears that the molecular fractions toward both stars are relatively small given the amount of reddening and extinction.  There are a number of B-type stars in the vicinity and the UV radiation field is thought to be strong in this area despite the lack of O-type stars (e.g., Kulesa et al. 2005).  Our preliminary analysis of the high $J$ lines in the {\\it FUSE} spectrum of $\\rho$ Oph D reveals no conclusive detections beyond $J$ = 6, from which we derive an uncertain log $N$(6) = 14.3.  This is comparable to that found toward HD 110432 (Rachford et al. 2001) which was modeled with a radiation field twice the interstellar average curve of Draine (1978).  \\subsubsection{Herschel 36} As noted in \\S\\ 2.3.10, HD 164740, better known as Herschel 36, lies within the Lagoon Nebula (M8) in a region with recent and ongoing star formation.  The molecular fraction (0.03) is the smallest known for a line of sight with $E(B-V)$ $>$ 0.3.  It is believed that some of the intervening material lies close to the star and is thus subject to the very strong far-UV radiation field of the late O-type star.  In fact, our {\\it FUSE} spectrum shows extreme H$_2$ excitation, including numerous lines from vibrationally excited states (B.\\ L.\\ Rachford, in preparation) demonstrating the likelihood of significant H$_2$ lying close to the star and the resultant exposure to a large far-UV flux. Furthermore, the exceptionally small far-UV extinction will allow the radiation from the O-type stars in M8 to influence a greater distance along the line of sight.  Thus, while this line of sight samples a complicated environment, it is plausible to assume that the radiation field is contributing to the small line of sight molecular fraction.  \\subsubsection{Overall properties of this sample} The 5 lines of sight were chosen for having evidence for larger than normal grains, thus providing smaller than normal H$_2$ formation rates per unit dust mass.  Also, the far-UV extinction is small which allows greater than normal penetration of the photodissociating radiation.  However, there remains a large range in molecular fractions that spans most of the total range we see in the overall translucent sample.  In formation-destruction equilibrium, the hydrogen molecular fraction will be proportional to the factors controlling formation, which are density ($n_H$) and the formation rate coefficient ($R$), and inversely proportional to the far-UV radiation field which controls destruction.  For the moment, we will focus on the last point.  There seems to be some evidence that the strength of the local radiation field is responsible for this range, particularly when considering the $\\rho$ Oph cloud and Herschel 36.  However, the line of sight toward HD 38087 is an exception.  As already noted, the distribution of material along this line of sight likely also plays a significant role.  Interestingly, these lines of sight have small H$_2$ rotational temperatures.  In Figure 7 (an update of Figure 9 in Paper I) we show the relationship between molecular fraction and rotational temperature.  In Paper I, we highlighted the high $f_{\\rm H2}$, low $T_{01}$ lines of sight.  Much rarer are lines of sight with low $f_{\\rm H2}$ and low $T_{01}$.  In fact, $\\rho$ Oph A and D and Herschel 36 are the only lines of sight in the {\\it Copernicus}/{\\it FUSE} sample with $f_{\\rm H2}$ $<$ 0.2, $T_{01}$ $\\leq$ 60 K, and $E(B-V)$ $>$ 0.2 (or log N(H) $>$ 21.1).  In contrast to the high $R_V$ lines of sight with small molecular fraction and low temperature, there are numerous lines of sight with normal or low values of $R_V$ that have small molecular fractions and higher than normal temperatures, particularly within the {\\it Copernicus} sample.  Gas heating due to grain electron photoemission (Draine 1978) may contribute to this fact, although there is not a significant overall relationship between $R_V$ and $T_{01}$.  The 13 {\\it Copernicus} lines of sight with $R_V$ $<$ 4 and $f_{\\rm H2}$ $<$ 0.2 all have $E(B-V)$ = 0.2--0.4, and all but one (HD 147165; $T_{01}$ = 64 K) have rotational temperatures greater than the average from our {\\it FUSE} sample of 67 K.  Of the 4 {\\it FUSE} lines of sight with $R_V$ $<$ 4 and $f_{\\rm H2}$ $<$ 0.2, only one has $T_{01}$ $<$ 67 K (HD 152236; $T_{01}$ = 62 K).  Of these 17 {\\it Copernicus} and {\\it FUSE} total lines of sight none have $T_{01}$ as small as $\\rho$ Oph A, $\\rho$ Oph D, or Herschel 36.  It is clear from Figure 7 that these three lines of sight are part of a small group that stand out from the rest of the sample, going against the generally inverse relationship between molecular fraction and temperature and lying in the bottom left of the figure.  We would expect the cold lines of sight to contain denser material and represent the expected environment for the large grains.  But, for at least Herschel 36 and $\\rho$ Oph D, the material is subject to considerable far-UV radiation, contributing to the small molecular fractions.  In these cases, we may be seeing the effect of a broad distribution of material.  Perhaps there is both ``cold'' diffuse material which contains most of the H$_2$ and is still not dense enough to exhibit a level of self-shielding that would allow a large molecular fraction, yet there is also a relatively small amount of material closer to the hot stars that not only has few molecules, but also considerable H$_2$ excitation.  In particular, for Herschel 36 there seems to be a significant velocity difference between the excited material and the ``cold'' material as indicated by the low-$J$ lines of H$_2$. Thus, the cold material may be significantly in the foreground of Herschel 36.  As previously noted, the radiation field is one of three factors that influence the molecular fraction in the models of Browning et al.\\ (2003), along with formation rate coefficient and density.  The models show that for the column densities we are sampling in this paper, there can be degeneracy in the molecular fractions that result from different combinations of the three factors.  In particular, a small formation rate coefficient can lower the molecular fraction in a manner similar to an increase in the UV flux.  The total column densities in our sub-sample of 5 high $R_V$ lines of sight cover the range log $N$(H$_{\\rm tot}$) $\\approx$ 21.2-22.0 ($N$ in cm$^{-2}$).  In this range, Browning et al.\\ (2003) find that a UV radiation field 50 times the Galactic mean reduces the molecular fraction by $\\sim$5 orders of magnitude at the low end, to factors of $\\sim$2--3 at the high end.  A reduction in the formation rate coefficient by a factor of 10 reduces the molecular fraction by $\\sim$2 orders of magnitude at the low end and by a minimal amount at the high end.  These two factors in combination were required to match H$_2$ data for the Small and Large Magellanic Clouds at levels of extinction generally below those in our present sample.  The most likely variables that would change the formation rate coefficient are grain size and temperature.  Since we have limited this sub-sample to high $R_V$, differences in grain size distribution have presumably been minimized.  If the $N$(1)/$N$(0) rotational temperature is a meaningful indicator of kinetic temperatures, our sub-sample is more or less isothermal.  Density, the third factor considered in the Browning et al.\\ (2003) models, is generally not as important within our range of column densities for typical Galactic values of radiation field and formation rate coefficient, particularly since the molecular fractions are usually relatively large already.  When combined with large radiation field or small formation rate coefficient, density variations can cause a large spread in the resulting small molecular fractions.  In conclusion, given the relatively large column densities in the regime we are probing with our sample, it appears more likely that large variations in radiation field are the dominant cause for the variations in molecular fraction.  Formation rate coefficients far below the Galactic average may contribute to smaller molecular fractions in our regime, but it is more certain that some of the clouds we are studying lie very close to major UV sources that dramatically increase the local radiation field."
    },
    "0809/0809.0623_arXiv.txt": {
        "abstract": "Quantum fast-roll initial conditions for the inflaton which are different from the classical fast-roll conditions and from the quantum slow-roll conditions can lead to inflation that last long enough. These quantum fast-roll initial conditions for the inflaton allow for kinetic energies of the same order of the potential energies and nonperturbative inflaton modes with nonzero wavenumbers. Their evolution starts with a transitory epoch where the redshift due to the expansion succeeds to assemble the quantum excited modes of the inflaton in a homogeneous (zero mode) condensate, and the large value of the Hubble parameter succeeds to overdamp the fast-roll of the redshifted inflaton modes. After this transitory stage the effective classical slow-roll epoch is reached. Most of the efolds are produced during the slow-roll epoch and we recover the classical slow-roll results for the scalar and tensor metric perturbations plus corrections. These corrections are important if scales which are horizon-size today exited the horizon by the end of the transitory stage and as a consequence the lower CMB multipoles get suppressed or enhanced. Both for scalar and tensor metric perturbations, fast-roll leads to a {\\bf suppression} of the amplitude of the perturbations (and of the low CMB multipoles), while the quantum precondensate epoch gives an {\\bf enhancement} of the amplitude of the perturbations (and of the low CMB multipoles). These two types of corrections can compete and combine in a scale dependent manner. They turn to be smaller in new inflation than in chaotic inflation. These corrections arise as natural consequences of the quantum nonperturbative inflaton dynamics, and can allow a further improvement of the fitting of inflation plus the $\\Lambda$CMB model to the observed CMB spectra. In addition, the corrections to the tensor metric perturbations will provide an independent test of this model. Thus, the effects of quantum inflaton fast-roll initial conditions provide a consistent and contrastable model for the origin of the suppression of the quadrupole and for other departures of the low CMB multipoles from the slow-roll inflation-$\\Lambda$CMB model which are to be contrasted to the TE and EE multipoles and to the forthcoming and future CMB data. ",
        "introduction": "\\label{sec:intro}  Inflation  (an epoch of accelerated expansion of the universe) solves the horizon and flatness problems of the standard Big Bang model. It naturally generates scalar density fluctuations that seed large scale structure and the temperature anisotropies in the cosmic microwave background (CMB), and tensor perturbations (primordial gravitational waves) \\cite{guth,infbooks}. Inflation is based on a scalar field (the inflaton) whose homogeneous expectation value drives the dynamics of the scale factor, plus small quantum fluctuations.  \\medskip  A great diversity of inflationary models predict fairly generic features: a gaussian, nearly scale invariant spectrum of (mostly) adiabatic scalar and tensor primordial fluctuations,  these provide an excellent fit to the highly precise wealth of data of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap} making the inflationary paradigm fairly robust. Precise CMB data reveal peaks and valleys in the temperature fluctuations resulting from  acoustic oscillations in the electron-photon fluid at recombination. These and future CMB and large scale structure (LSS) observations require more precise theoretical predictions from inflation, and a deeper understanding of how inflation begins and ends.  \\medskip  Amongst the wide variety of inflationary scenarios, single field slow-roll models provide an appealing, simple and fairly generic description of inflation. Its simplest implementation is based on a single scalar field (the inflaton). The inflaton potential is fairly flat and it dominates the universe energy during inflation. This flatness leads to a slowly varying Hubble parameter (slow-roll) ensuring a sufficient number of inflation e-folds to explain the homogeneity, isotropy and flatness of the universe, and also explains the gaussianity of the fluctuations as well as the (almost) scale invariance of their power spectrum.  \\medskip  The dynamics of inflation is usually described by the classical evolution of the inflaton in a dynamical space-time whose scale factor obeys the classical Friedmann equation. This is justified by the enormous stretching of physical lengths during inflation that classicalizes the dynamics. On the other hand, both the perturbations of the metric and of the inflaton are computed considering their quantum nature. The observable consequences of slow-roll inflation are well known \\cite{infbooks}, in particular those for the CMB anisotropies. Present observational data agree with its predictions for a class of slow-roll inflaton potentials except for the CMB quadrupole suppresion \\cite{wmap,cobe}.  However, the large energy densities during inflation allows for the creation of inflaton particles, and call for a quantum treatment for the inflaton field. Refs.~\\cite{tsuinf,cao04,qnew} considered a quantum treatment to study the dynamics of the inflaton. Quantum excited initial conditions lead to inflation that last long enough provided the quantum inflaton background satisfies the quantum generalized slow-roll conditions \\cite{tsuinf,cao04}. In addition, after a transitory precondensate epoch, the excited modes of the inflaton assembles in an homogeneous zero mode condensate \\cite{qnew,tsuinf,cao04}. Once formed, this condensate obeys the inflaton classical evolution equations and we therefore recover the classical inflation slow-roll description.  \\medskip  More recently, classical fast-roll initial conditions have been shown to lead to long enough inflation and to the CMB quadrupole suppresion \\cite{boy06,des08,lasenby}. In this article we consider more general inflaton initial conditions which lead to long enough inflation, and that we call quantum fast-roll initial conditions. They are different from the quantum slow-roll conditions and include the classical fast-roll conditions as a particular case.  All these general initial conditions lead after a transitory precondensate stage to a slow roll inflationary epoch that admits a classical description thus recovering the classical slow-roll inflation predictions for the scalar and tensor metric perturbations. Moreover, if the larger cosmologically relevant scales exited the horizon close to the end of the transitory stage, the effect of the modified initial conditions in the bulk modes that dominate the energy density will be imprinted as corrections to the amplitudes of the lower CMB multipoles. These corrections have been computed for classical fast roll initial conditions \\cite{boy06}, and they have been shown to improve the fit of inflation plus the $\\Lambda$CMB model to the CMB data \\cite{des08}. For quantum slow-roll initial conditions the corrections were first estimated in ref.~\\cite{cao04} and they will be computed in more detail here.  In the present article we also compute the corrections due to the quantum fast-roll initial conditions introduced here. We show that quantum slow-roll conditions and quantum fast-roll conditions may improve the fit of the CMB data with the $\\Lambda$CMB model plus inflation.These corrections to the spectrum of scalar and tensor metric perturbations arise as natural consequences of the quantum precondensate dynamics of the inflaton field. The two possible types of transitory epochs present before slow-roll, namely, effective classical fast-roll and quantum (fast-roll or slow-roll) precondensate, lead to opposite kinds of corrections for the low CMB multipoles. This quantum  precondensate dynamics predict corrections to the low TE and EE multipoles (temperature T and E-polarization mode correlations) and to the low CMB tensor multipoles. Therefore, this is a consistent and contrastable theory for the origin of the suppression of the quadrupole and for other departures of the low CMB multipoles data from the classical slow-roll inflation plus the $\\Lambda$CMB model.  \\medskip  This paper is organized as follows: in section \\ref{sec:qie} we present the quantum inflaton field and its evolution equations in the limit of large number of components, as well as the theory to compute corrections to the  spectrum of perturbations due to departures from classical slow-roll evolution, and the relative change in the CMB quadrupole due to these corrections. In Section \\ref{sec:qic}, we compute first the corrections to the primordial spectrum of perturbations due to quantum slow-roll initial conditions. Next, we introduce quantum fast roll initial conditions and compute the inflationary evolution. Then, we compute the  corrections to the primordial spectrum of perturbations due to the quantum fast roll initial conditions, and the relative change in the low CMB multipoles. Finally, in section IV we further discuss the results and present our conclusions.  We use units $ \\hbar = 1 $, $ c =1 $, and in the figures we also take $ m = 1 $ with $ m $ the inflaton mass. For the scale factor of the FRW metric we take $ a(0) = 1 $. ",
        "conclusions": "We have introduced a new type of inflationary initial conditions namely quantum fast-roll initial conditions which lead to long enough inflation. They are different from the quantum generalized slow-roll conditions and yield to the classical fast-roll conditions as a particular case. Quantum fast-roll initial conditions are far beyond slow-roll as their kinetic energy can be of the same order as their potential energy.  \\medskip  Two relevant issues play an important role in the quantum inflaton dynamics: the redshift due to the expansion and the large value of the Hubble parameter. The redshift due to the expansion succeeds to suppress the contributions from the quantum inflaton $k$-terms, and assembles all the nonperturbative inflaton modes in a inflaton condensate after a transitory precondensate epoch. On the other hand, the large value of the Hubble parameter succeeds to overdamp the oscillations of the redshifted inflaton quantum modes thus stopping the fast-roll. Therefore, the combined effect of the redshift due to the expansion and the large values of the Hubble parameter makes that after a transitory quantum inflation epoch an effective classical slow-roll inflation epoch is reached.  \\medskip  Such transitory epoch can have observable effects. If scales which are today horizon-size exited the horizon close to the end of the transitory epoch, the spectra of scalar and tensor metric perturbations are similar to the spectra for classical slow-roll but with corrections for the larger scales. These corrections for the larger cosmologically relevant scales of the scalar and tensor metric perturbations are reflected in the amplitudes of the lower multipoles of the CMB spectra. We have found that for both scalar and tensor metric perturbations, the classical or effective classical condensate fast-roll epoch {\\bf suppresses} the amplitude of the perturbations (and therefore the amplitude of the respective CMB multipoles), while the quantum precondensate epoch {\\bf enhances} the amplitude of the perturbations (and of the respective CMB multipoles).  The corrections to the CMB multipoles from the quantum precondensate turn to be smaller in new inflation than in chaotic inflation.  In addition, when both fast-roll and quantum precondensate effects compete they can lead to suppression in some scales and enhancement in other scales. Therefore, these corrections for the lower multipoles can allow to further improve the fit of the predictions of the inflation plus the $\\Lambda$CMB model to the observed TT, TE and EE CMB spectra.  \\medskip  It is important to note that these corrections to the spectrum of scalar metric perturbations arise as natural consequences of the quantum dynamics of the inflaton. Since the quantum dynamics also predicts corrections to the spectrum of tensor perturbations, when the tensor perturbations spectrum will be observed, an independent test of the quantum fast-roll model will be possible. Therefore, effects due to quantum fast-roll initial conditions in the nonperturbative inflaton modes ($ k > \\Lambda $) are a consistent and contrastable explanation for the suppression of the CMB quadrupole and for the other departures of the low CMB multipoles from the slow-roll inflation plus $\\Lambda$CMB model.  The consequences of our study of quantum fast-roll initial conditions on the low TE and EE multipoles need to be explored. Forthcoming data for low CMB multipoles motivate this work and are needed to reach a clear understanding on these issues."
    },
    "0809/0809.0309_arXiv.txt": {
        "abstract": "We use HST/ACS images and a photometric catalog of the COSMOS field to analyze morphologies of the host galaxies of $\\sim$400 AGN candidates at redshifts $0.3 < z < 1.0$.  We compare the AGN hosts with a sample of non-active galaxies drawn from the COSMOS field to match the magnitude and redshift distribution of the AGN hosts.  We perform 2-D surface brightness modeling with GALFIT to yield host galaxy and nuclear point source magnitudes.  X-ray selected AGN host galaxy morphologies span a substantial range that peaks between those of early-type, bulge-dominated and late-type, disk-dominated systems.  We also measure the asymmetry and concentration of the host galaxies.  Unaccounted for, the nuclear point source can significantly bias results of these measured structural parameters, so we subtract the best-fit point source component to obtain images of the underlying host galaxies.  Our concentration measurements reinforce the findings of our 2-D morphology fits, placing X-ray AGN hosts between early- and late-type inactive galaxies.  AGN host asymmetry distributions are consistent with those of control galaxies.  Combined with a lack of excess companion galaxies around AGN, the asymmetry distributions indicate that strong interactions are no more prevalent among AGN than normal galaxies.  In light of recent work, these results suggest that the host galaxies of AGN at these X-ray luminosities may be in a transition from disk-dominated to bulge-dominated, but that this transition is not typically triggered by major mergers. ",
        "introduction": "\\label{sec.intro}  Studies over the past decade suggest fundamental links between galaxies and their central supermassive black holes (SMBHs).  The masses of the central SMBHs in nearby galaxies correlate with several host bulge properties, including luminosity \\citep{kor95,mcl01,mar03}, velocity dispersion \\citep{fer00, geb00}, and mass \\citep[for a review of these relations, see Ferrarese 2004]{mag98}.  More recent studies extend these relationships to galaxies and quasars at redshifts from $z=0.37$ \\citep{tre04} to as high as $z\\sim4$ \\citep{pen06a, pen06b}.  These observations indicate co-evolution of SMBH mass accretion and host bulge formation processes, perhaps through interactions such as AGN feedback. Galaxy merger simulations that include a prescription for SMBH feedback \\citep{spr05} and reproduce the $M_{BH}$-$\\sigma_{bulge}$ relation can recover several observed properties of quasar host descendants, including the stellar mass function of local elliptical galaxies, and the red galaxy luminosity function and its variation with redshift \\citep{hop06}.  In these models, gas-rich mergers drive material toward the central black holes, leading to intense star formation and SMBH accretion.  The nuclear SMBH begins its active phase obscured, but feedback energy during the peak accretion phase blows away the obscuring material and results in a brief quasar phase.  The blowout phase coincides with a rapid truncation of star formation throughout the host galaxy.  Recent observational studies lend credibility to this picture by hinting that star-formation quenching coincides with AGN activity \\citep{bun07, sil08, tre07}.  This picture of SMBH-host co-evolution relies on a hypothesized merger mechanism for fueling active black holes, and incorporates predictions that gas-rich disk galaxies merge to form luminous starbursts, eventually evolving into massive elliptical galaxies.  Other possible fueling mechanisms include less-direct tidal interactions with nearby galaxies \\citep{men04} and instabilities in a quiescent galaxy's gaseous disk \\citep[e.g.][]{hop06b}.  In this study, we explore the interaction mechanism for intermediate-luminosity AGN in the COSMOS field by analyzing their environments and host galaxies.  Previous similar work employing HST survey data has found conflicting evidence.  \\citet{gro05} analyzed $\\sim$100 X-ray selected AGN at redshifts up to $\\sim$1.3 in the GOODS fields and found no significant differences between their host structural properties and companion fractions and those measured for matched control sample galaxies.  \\citet{pie07}, on the other hand, examined $\\sim$60 X-ray and IR-selected AGN with $0.2 < z < 1.2$ in the Extended Groth Strip and found that AGN are marginally more likely than control galaxies to have nearby companions.  In both studies X-ray selected AGN are found to reside predominantly in host galaxies with bulge-dominated morphologies, generally in agreement with work at low and high redshift quasar host studies \\citep{jah04, san04}.  Investigating larger-scale environments of 52 quasars from the DEEP2 redshift survey, \\cite{coi07} showed that the quasar-galaxy cross-correlation function at $z\\sim$1 closely resembles the galaxy auto-correlation function at all scales, and that the relative quasar bias traces that of blue galaxies better than red galaxies.  This might suggest that high luminosity AGN reside in blue bulges \\citep{jah04, sil08} that have not yet migrated to the high-density environments typically found for massive red galaxies.  In the nearby universe, where precise techniques allow detailed studies of tens of thousands of AGN hosts (with, e.g., the Sloan Digital Sky Survey), active nuclei reside in the most massive galaxies, with structural properties similar to early-type galaxies but with relatively young stellar populations \\citep{kau03}.  Examining images of $\\sim$100 of the most luminous AGN within $z<0.1$, \\citeauthor{kau03} found that they occupy roughly equal fractions of blue spheroids, single disk galaxies, and disturbed/interacting galaxies.  Furthermore, low-redshift close galaxy pairs that exhibit strong indications of interaction are more likely to include AGN than pairs without interaction indicators \\citep{alo07}.  At the same time, black hole accretion activity is significantly larger for AGN with bright companions than otherwise \\citep{alo07}, and quasars are found to have local galaxy overdensities within 100 kpc in excess of that seen in non-active galaxies \\citep{ser06} and lower-luminosity AGN \\citep{str08}.  Combined with earlier imaging studies showing that a minority of local Type 1 AGN are undergoing interactions \\citep[e.g.][]{der98}, these studies suggest that mergers and interactions play some role in fueling AGN, but not necessarily a dominant one.  These low-redshift AGN may represent a different population than that found closer to the peak of AGN activity, $z\\simeq$ 2, with different fueling mechanisms in play.  The typical AGN in a local survey like the SDSS is less luminous and possibly hosted by a galaxy in a different evolutionary state than AGN selected at higher redshift.  Intermediate- and high-redshift AGN ($z > 0.5$) are typically found in bluer, more extended, and more irregular galaxies than their low-redshift counterparts \\citep[cf. ][]{jah04, san04}.  By using samples with a wider redshift range we can constrain the dominant mechanism in the history of AGN fueling, and potentially uncover its evolution over cosmic time.  Samples at $z>0.3$, however, remain limited to a few dozen objects selected through a variety of methods (see studies mentioned above).  In the present study, we extend this moderate-redshift sample and compare the effects of different selection techniques.  Using the extensive data of the Cosmic Evolution Survey (COSMOS), we determine basic properties of a sample of AGN hosts to probe their co-evolution with SMBHs.  Selection of the AGN sample and data used for the analysis is described in \\S\\ref{sec.sample}.  The analyses of the host galaxies are presented in \\S\\S\\ref{sec.morph} and \\ref{sec.env}, followed by a brief discussion and conclusions.  For cosmological calculations, we adopt $h=0.75$, (where $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$ is the Hubble parameter), $\\Omega_M=0.3$ (matter density parameter), and $\\Omega_{\\Lambda}=0.7$ (cosmological constant density parameter). ",
        "conclusions": "\\label{sec.conclusion} We explored the host morphologies and environments of AGN in the COSMOS field.  Using X-ray- and radio-selected AGN candidates with confirmed spectroscopic redshifts, we analyzed host galaxy structural properties as well as merger indicators, probing the connection between AGN activity and galaxy interactions.  The following summarizes our main points:  \\begin{itemize} \\item The central point source in optical images of X-ray selected AGN has substantial impact on measured structural parameters such as asymmetry and concentration.  Insufficient accounting for the point source can lead to systematically low asymmetry and systematically high concentration. \\item Full 2-D fits and concentration measurements which account for the central point source in X-ray AGN show that their host galaxies have a broad range of morphologies whose distribution is intermediate between the bulge- and disk-dominated regimes. \\item Although radio AGN hosts also appear to have a wide range of morphologies, contamination by star-forming galaxies prevents us from distinguishing them from early-type normal galaxies. \\item Measurements of AGN host galaxy asymmetry do not differ significantly from those of matched control galaxies. \\item Neighbor counts around AGN are indistinguishable from those around matched control galaxies using photometric redshifts. \\end{itemize}  These findings do not support the hypothesis that major mergers drive black hole activity, but they do suggest that the host galaxies of AGN at these luminosities may be in a state of morphological transition.  Future work by members of the COSMOS collaboration will address this possibility in more detail by examing the colors and star formation rates of AGN hosts and their relationships with environment (Silverman et al. in preparation)."
    },
    "0809/0809.0765_arXiv.txt": {
        "abstract": "Investigations of \\ion{He}{2} Ly$\\alpha$ (304~\\AA\\ rest) absorption toward a half-dozen quasars at $z\\sim3-4$ have demonstrated the great potential of helium studies of the IGM, but the current critically small sample size of clean sightlines for the \\ion{He}{2} Gunn-Peterson test limits confidence in cosmological inferences, and a larger sample is required. Although the unobscured quasar sightlines to high redshift are extremely rare, SDSS DR6 provides thousands of $z>2.8$ quasars. We have cross-correlated these SDSS quasars with {\\it GALEX} GR2/GR3 to establish a catalog of 200 higher-confidence ($\\sim$70\\% secure) cases of quasars at $z=2.8-5.1$ potentially having surviving far-UV (restframe) flux. We also catalog another 112 likely far-UV-bright quasars from {\\it GALEX} cross-correlation with other (non-SDSS) quasar compilations. Reconnaissance UV prism observations with {\\it HST} of 24 of our SDSS/{\\it GALEX} candidates confirm 12 as detected in the far-UV, with at least 9 having flux extending to very near the \\ion{He}{2} break; with refinements our success rate is even higher. Our SDSS/{\\it GALEX} selection approach is thereby confirmed to be an order of magnitude more efficient than previous \\ion{He}{2} quasar searches, more than doubles the number of spectroscopically confirmed clean sightlines to high redshift, and provides a resource list of hundreds of high-confidence sightlines for upcoming \\ion{He}{2} and other far-UV studies from {\\it HST}. Our reconnaissance {\\it HST} prism spectra suggest some far-UV diversity, confirming the need to obtain a large sample of independent quasar sightlines across a broad redshift range to assess such issues as the epoch(s) of helium reionization, while averaging over individual-object pathology and/or cosmic variance. ",
        "introduction": "A substantial fraction (perhaps most) of the baryons in the Universe did not collapse into such dense structures as stars, galaxies, and quasars, instead remaining behind in a more dilute intergalactic medium (hereafter, IGM), composed primarily of primordial hydrogen and helium \\citep[e.g., see recent review by][]{mei08}. Ultraviolet radiation from massive early stars, star-forming galaxies, and quasars gradually reionized their surrounding IGM, ultimately ending the ``dark ages\". Cosmic microwave background studies provide one important indirect chronometer for reionization: the {\\textit{Wilkinson Microwave Anisotropy Probe} ({\\textit{WMAP}}) 5-year data \\citep{D08} constrain sudden one-time hydrogen reionization to $z>6.7$ at $3 \\sigma$, and place it at $z=11.0 \\pm 1.4$ ($1 \\sigma$). However, the WMAP measures do not strongly distinguish between models with one or multiple epochs of reionization; and, although they suggest the possibility of an extended reionization epoch, the WMAP measures alone do not tightly constrain that possibility \\citep{D08}.  Much of the IGM in the nearby Universe is highly ionized, with a hydrogen neutral fraction $x_{\\mbox{\\scriptsize{H I}}}=N_{\\mbox{\\scriptsize{H I}}}/N_{\\mbox{\\scriptsize{H II}}} \\sim 10^{-5}$ \\citep{F06}. A sensitive measure of the neutral hydrogen is to examine quasar spectra for the saturated trough shortward of Ly$\\alpha$ caused by even a modest amount of neutral IGM gas, i.e., the classic Gunn-Peterson effect \\citep{GP65}. Somewhat in contrast to initial hints from the first few $z\\sim$6 quasars examined, data from $\\sim$20 now-studied high-redshift QSO sightlines, primarily from objects discovered in the Sloan Digital Sky Survey \\citep[SDSS;][]{Y00}, suggest that the reionization of \\ion{H}{1} likely happened before $z\\sim6$ \\citep{F06}.  The \\ion{H}{1} Gunn-Peterson trough is evident in some $z\\sim6.3$ quasars \\citep{B01,F06}, suggesting at least that the ionization of hydrogen is patchy and may be rapidly evolving in the $z>6$ regime. However, there is marked sightline variance in the ensemble of quasar \\ion{H}{1} Gunn-Peterson measurements and their interpretations \\citep[e.g.][]{W03,F06,T06,bol07,bec07}. Moreover, some high-redshift galaxy studies also tend to favor higher redshifts for \\ion{H}{1} reionization \\citep[e.g.][]{hu02,rho04}, as galaxies with strong (and inferentially, largely unabsorbed) \\ion{H}{1} Ly$\\alpha$ emission are seen out to (at least) $z\\sim7$ \\citep[e.g.][]{iye06}.  Although hydrogen dominates both the number and mass fraction of baryons, helium is the most abundant absorber in much of the IGM accessible to current study. The \\ion{He}{2} reionization epoch is likely delayed versus that of \\ion{H}{1} because hard-ionizing sources such as AGN had to form to produce the higher energy photons needed \\citep[e.g.][]{W03b}. Even under the hard photoionizing conditions at $z\\sim 2-4$, \\ion{He}{2} outnumbers \\ion{H}{1} by a factor of $\\eta \\sim 50$-100 \\citep{fec06}, and has a higher opacity by $\\eta/4 > 10$. The \\ion{He}{2} Ly$\\alpha$ transition is thus the more sensitive, higher opacity, empirical probe of much of the highly ionized IGM---as compared to \\ion{H}{1} which over most of sampled redshift space (between the \\ion{H}{1} Ly$\\alpha$ forest lines) is extremely sparse. \\ion{He}{2} studies allow cosmological inferences about the epoch(s) of reionization, the intensity and spectrum of the ionizing background radiation, and measures of the cosmic baryon density in the IGM. However, \\ion{He}{2} Ly$\\alpha$ absorption occurs in the far UV (304\\AA\\ restframe), only observable from space in a rare fraction of high redshift quasars which lie by chance along clean lines of sight with little foreground absorption.  Jakobsen et al.'s pioneering work \\citep{J94} with the {\\it Hubble Space Telescope} (hereafter, {\\it HST}) provided the first detection of a \\ion{He}{2} Gunn-Peterson trough, in Q0302-003 ($z=3.29$).  Yet even following a further decade of effort, only four quasar sightlines suitable for \\ion{He}{2} Gunn-Peterson studies had been confirmed---though each was subjected to intensive UV spectral campaigns from {\\it HST}, {\\it HUT}, and/or {\\it FUSE}. In redshift order these objects are: HS~1700+64 at $z=2.72$ \\citep{D96,fec06}, HE~2347-4342 at $z=2.88$ \\citep{R97,kri01,sme02,Z04a,S04}, PKS~1935-692 at $z=3.18$ \\citep{A99}, and Q0302-003 at $z=3.29$ \\citep{J94,H97,H00}. Broadly, the various intensive UV spectroscopic studies of these four quasars indicate that the IGM \\ion{He}{2} optical depth increases markedly from $\\tau\\sim 1$ near $z=2.5$ to $\\tau >4$ at $z=3.3$, and are also consistent with theoretical notions \\citep{haa96}, and other indirect evidence \\citep{son96,the02}, that helium reionization likely occurred near $z\\sim$3 (delayed vs. that of hydrogen at $z>$6). These intensive follow-up studies have also provided a handful of empirical measures of the flux in the ionizing background radiation at $z\\sim$3 and the baryon density ($\\Omega_g\\sim0.01$) in the IGM \\citep[e.g.][]{H97,Z04a,T07}. A predicted signature of the \\ion{He}{2} reionization epoch is a damping absorption profile {\\it redward} of \\ion{He}{2} \\citep{mir98,mad00}; this feature was unseen in these four intensively studied cases at $z<3.3$, but perhaps merely because they are at too low redshift (when \\ion{He}{2} ionization was well underway).  From later {\\it HST} ultraviolet follow-up of selected optically bright high-redshift quasars, \\citet{rei05} identified QSO 1157+3143 (at $z=3.0$) as another \\ion{He}{2} quasar, and our group added three more clean quasar sightlines that extend helium IGM studies to the highest redshifts heretofore sampled \\citep{Z04b}. From an {\\it HST} SNAP survey we first confirmed SDSS 2346-0016 ($z=3.5$), SDSS 1711+6052 ($z=3.8$), and SDSS 1614+4859 ($z=3.8$) as excellent new clean sightlines for helium IGM studies; our 6\\% selection efficiency, though modest, was among the highest then obtained. Moreover, our lengthy follow-up with the Advanced Camera for Surveys, Solar Blind Channel (ACS SBC) prism (taken after the failure of the Space Telescope Imaging Spectrograph [STIS]) for two of these, SDSS 2346-0016 and SDSS 1711+6052, shows an especially intriguing result: the $z=3.8$ case shows a surprisingly prominent absorption profile redward of the \\ion{He}{2} break location \\citep{Z08} that is not seen in the $z=3.5$ case. The profile in SDSS 1711+6052 is stronger and wider than commonly expected for the damped profile signature of the IGM, so in that paper we also consider additional explanations including absorption in a high-redshift intense star-forming region. Whatever the explanation, the strength of this surprising feature demands further scrutiny in future high-quality UV spectra toward SDSS 1711+6052, as well as other comparably high-redshift \\ion{He}{2} quasar sightlines.    Motivated by the previous small sample size that limits confident global conclusions, and the potentially intriguing results at higher redshifts \\citep[e.g.,][]{Z08}, here we detail our successful efforts to markedly increase the number of new quasar sightlines suitable for the study of \\ion{He}{2}, across a broad redshift range. Starting from the very large sample of 7800 SDSS quasars at suitable redshift, we use the {\\textit{Galaxy Evolution Explorer}} ({\\it GALEX}) archive to then cull to a new catalog of 200 higher-confidence far-UV detections of quasars at $z>2.78$. An additional 112 quasar sightlines likely clean to high redshift are also similarly identified from cross-correlation of {\\it GALEX} with other (non-SDSS) quasars. Our catalogs of likely far-UV-bright (restframe) quasars, presented in section 2, provide appropriate lists from which to draw for follow-up observations.  In section 3 we present reconnaissance UV observations with the {\\it HST} ACS prism of 24 SDSS/{\\it GALEX} candidate far-UV-bright quasars. Among these, we confirm 12 as bright in the far UV, with at least 9 evidently having flux all the way down to (very near) the \\ion{He}{2} break.  As we discuss in section 4, our SDSS/{\\it GALEX} selection approach is thus found to be an order of magnitude more efficient than previous \\ion{He}{2} quasar searches, and herein we more than double the number of clean sightlines spectroscopically confirmed as suitable for helium IGM studies. Moreover, our reconnaissance {\\it HST} UV prism spectra suggest diversity, further confirming the need to obtain a large sample of independent quasar sightlines over a broad redshift range to understand cosmological variance and/or possible individual-object peculiarity. ",
        "conclusions": "We provide a catalog of 200 higher-confidence ($\\sim$70\\% secure) SDSS quasars at $z>2.78$ that are likely statistically associated with GALEX ultraviolet sources. We provide an additional set of 112 similar quasars, also $z>2.78$ but on average lower redshift, from consideration of the \\citet{ver06} quasar compilation. Our catalogs herein comprise very large resource lists (two orders of magnitude larger than previous similarly dedicated efforts) of likely far-UV (restframe) bright quasars, potentially of interest for a range of studies from far-UV SED, emission line, and BALQSO spectral studies, to background probes of intervening clouds including \\ion{H}{1} at low-redshift, to our own primary interest in the \\ion{He}{2} Gunn-Peterson test.  We report reconnaissance UV spectral observations from {\\it HST} of 24 SDSS/{\\it GALEX} selected quasars, chosen to encompass a wide redshift range, finding that half are indeed far-UV bright. At least 9 new quasars are confirmed in our {\\it HST} prism spectra as suitable new sightlines for \\ion{He}{2} Gunn-Peterson studies. The {\\it HST} data also allow us to confirm the even higher efficiency of our refined selection approach, which generated the catalogs of section 2. In either case, our combined SDSS/{\\it GALEX}/{\\it HST}-reconnaissance selection efficiency is an order of magnitude better than previous similar \\ion{He}{2} quasar searches, and we herein more than double the number of spectroscopically confirmed \\ion{He}{2} quasars.  Our {\\it HST} prism spectra show some diversity, reaffirming the need for UV reconnaissance spectral observations to confirm that UV flux extends all the way down to \\ion{He}{2} (even in cases detected in broadband from {\\it GALEX}). The spectral diversity found also confirms the necessity of identifying multiple independent clean sightlines, covering a broad redshift range, to limit cosmological variance and individual-object pathology in seeking global understanding from helium IGM studies. An observed-frame ensemble stack of our {\\it HST} reconnaissance spectra suggests that the strong absorption seen in SDSS1711+6052 by \\citet{Z08} redward of the \\ion{He}{2} break is unlikely to be an instrumental artifact.   To probe IGM \\ion{He}{2} in detail, including the epoch of \\ion{He}{2} reionization, requires a large enough sample of clean quasar sightlines, extending over a broad redshift range, in order to be limited by cosmic variance rather than individual peculiarities or a clumpy IGM. With such a large, broad-redshift sample, concerns of a possible systematic bias due to clear sightline selection may be addressable using the observed sightlines to calibrate large-scale numerical simulations of reionization \\citep[e.g.,][]{W08}. The range of redshift encompassed by our sample may also allow future observations to constrain percolation models of reionization, such as the one discussed by \\citet{P07}.  The empirical surprises at high redshift \\citep{Z08}, theoretical uncertainties in interpreting the past low redshift results \\citep{mei08}, and the broad goal of identifying the epoch of helium reionization through the anticipated redward damped \\ion{He}{2} profile signature, all argue strongly for the need to emphasize \\ion{He}{2} Gunn-Peterson studies beyond the $z\\sim3$ regime of the bulk of the venerable past studies. Our catalogs of hundreds of far-UV-bright quasars, and even more directly our {\\it HST} spectroscopic confirmations of at least 9 new \\ion{He}{2} quasars (along with several more from our earlier similar work), provide basic resource lists extending to such higher redshifts for future detailed (good S/N and spectral resolution) UV spectral science investigations of helium in the IGM."
    },
    "0809/0809.0553_arXiv.txt": {
        "abstract": "The recent discovery of the 2003~EL$_{61}$ collisional family in the Kuiper belt (Brown et al$.$~2007) is surprising because the formation of such a family is a highly improbable event in today's belt.  Assuming Brown et al$.$'s estimate of the size of the progenitors, we find that the probability that a Kuiper belt object was involved in such a collision since primordial times is less than roughly $0.001$.  In addition, it is not possible for the collision to have occurred in a massive primordial Kuiper belt because the dynamical coherence of the family would not have survived whatever event produced the currently observed orbital excitation.  Here we suggest that the family is the result of a collision between two scattered disk objects.  We show that the probability that a collision occurred between two such objects with sizes similar to those advocated in Brown et al$.$~(2007) and that the center of mass of the resulting family is on an orbit typical of the Kuiper belt can be as large as 47\\%.  Given the large uncertainties involved in this estimate, this result is consistent with the existence of one such family.  If true, this result has implications far beyond the origin of a single collisional family, because it shows that collisions played an important role in shaping the dynamical structure of the small body populations that we see today. ",
        "introduction": "\\label{sec:intro}  Recently, Brown et al$.$~(2007, hereafter BBRS) announced the discovery of a collisional family associated with the Kuiper belt object known as (136108) 2003~EL$_{61}$ (hereafter referred to as \\el61).  With a diameter of $\\sim\\!1500\\,$km (Rabinowitz et al$.$~2006), \\el61\\ is the third largest known Kuiper belt object. The family so far consists of \\el61\\ plus seven other objects that range from 150 to $400\\,$km in diameter (BBRS; Ragozzine \\& Brown~2007\\Red{, hereafter RB07}).  Their proper semi-major axes ($a$) are spread over only $1.6\\,$AU, their eccentricities ($e$) differ by less than 0.08, and their inclinations ($i$) differ by less than $1.5^\\circ$ (see Table~\\ref{tab:fam}). After correcting for some drift in the eccentricity of \\el61\\ \\Red{and 1999 OY$_3$ due to Neptune's mean motion resonances}, this corresponds to a velocity dispersion of $\\lesssim\\!150\\,$m/s (\\Red{RB07}).  BBRS estimates that there is only a one in a million chance that such a grouping of objects would occur at random.  \\begin{table}[h] \\singlespace \\begin{center} \\begin{tabular}{|lccc|} \\hline name           &  $a$    &   $e$  &    $i$  \\\\ &  (AU)   &        &   (deg) \\\\ \\hline 2003 EL$_{61}$       &   43.3  & 0.19  &   28.2  \\\\ 1995 SM$_{55}$       &   41.7  & 0.10  &   27.1  \\\\ 1996 TO$_{66}$       &   43.2  & 0.12  &   27.5  \\\\ 1999 OY$_3$     &   44.1  & 0.17  &   24.2 \\\\ 2002 TX$_{300}$      &   43.2  & 0.12  &   25.9  \\\\ 2003 OP$_{32}$       &   43.3  & 0.11  &   27.2  \\\\ 2003 UZ$_{117}$ &   44.1  & 0.13  &   27.4 \\\\ 2005 RR$_{43}$       &   43.1  & 0.14  &   28.5  \\\\ \\hline \\end{tabular} \\caption{\\label{tab:fam} The orbital elements of the known \\el61\\ family as supplied by the {\\it Minor Planet Center} on July 20, 2007.} \\end{center} \\end{table}  Based on the size and density of \\el61\\ and on the hydrodynamic simulations of Benz \\& Asphaug~(1999), and assuming an impact velocity of $3\\,$km/s, BBRS estimate that this family is the result of an impact between two objects with diameters of $\\sim\\!1700\\,$km and $\\sim\\!1000\\,$km.  Such a collision is surprising (M$.$~Brown, pers. comm.), because there are so few objects this big in the Kuiper belt that the probability of the collision occurring in the age of the Solar System is very small.  Thus, in this paper we investigate the circumstances under which a collision like the one needed to create the \\el61\\ family could have occurred.  In particular, in \\S{\\ref{sec:KB}} we look again at the idea that the larger of the two progenitors of this family (the target) originally resided in the Kuiper belt and carefully determine the probability that the impact could have occurred there.  We show that this probability is small, and so we have to search for an alternative idea.  In \\S{\\ref{sec:SD}}, we investigate a new scenario where the \\el61\\ family formed as a result of a collision between two scattered disk objects (hereafter SDOs).  The implications of these calculations are discussed in \\S{\\ref{sec:end}}. ",
        "conclusions": "\\label{sec:end}  The recent discovery of the \\el61\\ family in the Kuiper belt (BBRS) is surprising because its formation is, at first glance, a highly improbable event.  BBRS argues that this family is the result of a collision between two objects with radii of $\\sim\\!850\\,$km and $\\sim\\!500\\,$km.  The chances that such event would have occurred in the current Kuiper belt in the age of the Solar System is roughly 1 in $4000$ (see \\S{\\ref{sec:KB}}).  In addition, it is not possible for the collision to have occurred in a massive primordial Kuiper belt because the dynamical coherence of the family would not have survived whatever event molded the final Kuiper belt structure.  We also investigated the idea that the family could be the result of a target KBO being struck by a SDO projectile, and found that the probability of such an event forming a family on a stable Kuiper belt orbit is $\\lesssim\\!10^{-3}$.  In this paper, we argue that the \\el61\\ family is the result of a collision between two scattered disk objects.  In particular, we present the novel idea that the collision between two SDOs on highly eccentric unstable orbits could damp enough orbital energy so that the family members would end up on stable Kuiper belt orbits.  This idea of using the scattered disk as the source of both of the family's progenitors has the advantage of significantly increasing the probability of a collision because the population of the scattered disk \\Red{was much larger in the early Solar System (it is currently eroding away due to the gravitational influence of Neptune --- DL97; DWLD04)}.  With the use of three pre-existing models of the dynamical evolution of the scattered disk (DWLD04, LD/DL97, and TGML05) we show that the probability that a collision between a $\\sim\\!850\\,$km SDO and $\\sim\\!500\\,$km SDO occurred and that the resulting collisional family was spread around a stable Kuiper belt orbit can be as large as 47\\%.  Given the uncertainties involved, this can be considered on the order of unity.  Thus, we conclude that the \\el61\\ family progenitors are significantly more likely to have originated in the scattered disk than the Kuiper belt.  If true, this result has important implications for the origin of the Kuiper belt because it is the first direct indication that collisions can affect the dynamical evolution of this region.  Indeed, we \\Red{suggest} at the end of \\S{\\ref{sec:SD}} that this process might be responsible for the emplacement of the so-called `hot' classical belt (Brown~2001) because it naturally explains why so few of these objects are found to be binaries (Noll et al$.$~\\Red{2008})."
    },
    "0809/0809.0881_arXiv.txt": {
        "abstract": "We present a novel technique to overcome the limitations of the applicability of Principal Component Analysis to typical real-life data sets, especially astronomical spectra. Our new approach addresses the issues of outliers, missing information, large number of dimensions and the vast amount of data by combining elements of robust statistics and recursive algorithms that provide improved eigensystem estimates step-by-step.  We develop a generic mechanism for deriving reliable eigenspectra without manual data censoring, while utilising all the information contained in the observations. We demonstrate the power of the methodology on the attractive collection of the VIMOS VLT Deep Survey spectra that manifest most of the challenges today, and highlight the improvements over previous workarounds, as well as the scalability of our approach to collections with sizes of the Sloan Digital Sky Survey and beyond. ",
        "introduction": "As modern telescopes collect more and more data in our exponential world, where the size of the detectors essentially follows Moore's law, the kind of statistical challenges astronomers face in analysing the observations change dramatically in nature. We need algorithms that scale well in time and complexity with the volume of the data, while obeying the constraints of today's computers.  But the large data volume is only one of the consequences of this trend.  With more observations in hand, the number of problematic detections also increases. In addition to the elegant mathematical theorems that work miracles on textbook examples, scientists need to develop methodologies that provide reliable results that are robust in the statistical sense when applied to real-life data.  One particular multi-variate analysis technique, which is widely accepted and popular not only in astronomy but also in genetics, imaging and many other fields, is Principal Component Analysis (PCA). For its simple geometric meaning and straightforward implementation via Singular Value Decomposition (SVD), it has been utilised in many areas of the field including classification of galaxies and quasars \\citep{francis92,connolly95,connolly99,yip04a,yip04b}, spectroscopic and photometric redshift estimation \\citep{glazebrook98,budavari99,budavari00}, sky subtraction \\citep{wild05}, and highly efficient optical spectral indicators \\citep{wild07}. However, direct application of the classic PCA  to real data is almost always impossible; the reasons are usually three-fold: (1) The technique is extremely sensitive to outliers. With smaller datasets, scientists would ``clean up'' the sample by completely removing the ``obvious'' outliers in one or more projections and analyse the remaining subset. The problem with this approach is that it is subjective, and it becomes impractical for large datasets. (2) Another problem is missing measurements in the data vectors, e.g., pixels of a spectrum. There are various reasons for this to occur: strong night sky lines, cosmic rays or simply because of the redshift that yields different restframe wavelength coverage for the spectra. The application of PCA implies the assumption of a Euclidean metric, but it is not clear how to calculate Euclidean distances when data is missing from our observed vectors. (3) Last but not least, the memory requirement of SVD is significant as the entire dataset is stored in memory while the decomposition is computed. For example, the Sloan Digital Sky Survey (SDSS) Data Release 6 contains a million spectra with 4000 resolution elements each. While machines certainly exist today that have the required amount of RAM ($\\sim\\!50$GB), typical workstations have lesser resources.  Additionally, in most situations, we only seek a small number of eigenvectors associated with the largest eigenvalues, and SVD computes all the singular vectors in vain.  The optical spectra of the VIMOS VLT Deep Survey \\citep[VVDS;][]{vvds} provides an extremely attractive dataset for galaxy evolution studies at high redshift, yet, due to their generally low signal-to-noise ratio they are unsuitable for traditional PCA.  In this case, the challenge does not lie in the volume of the data set, rather in the natural limitations of high redshift observations. Thanks to the careful processing by the VVDS Team, the spectra are well calibrated and each one contains valuable information for a PCA analysis.  Our goal is to develop an algorithm to address all of the above issues in a way that is true to the spirit of the PCA and maintains its geometrically meaningful properties. In \\S{}\\ref{sec:pca} we detail the various layers of our solution to the problem, and in \\S{}\\ref{sec:vvds} we compare the performance of different techniques when applied to the VVDS spectra.  In \\S{}\\ref{sec:disc} we evaluate the results of the methods and analyse the emerging physical features.  \\S{}\\ref{sec:sum} concludes our study. ",
        "conclusions": "\\label{sec:disc}  \\begin{figure*} \\includegraphics[scale=0.7]{fig3.ps} \\caption{The joint distribution of the first two principal component amplitudes for the VVDS collection of 3485 spectra using the eigenbases presented in Figure~\\ref{fig:espec}. In order to focus on the main sample of objects, the axes are scaled such that outliers are not shown. From left to right: classic PCA using SVD, trimmed PCA with iterative removal of outliers, our robust PCA from the randomised streaming algorithm.} \\label{fig:pcs} \\end{figure*}  There are two important points that an effective spectral eigenbasis must obey: (1) the eigenspectra should not introduce noise into the decomposition of individual galaxy spectra by being noisy themselves; (2) the top few eigenspectra must primarily describe the variance in the majority of the dataset, and not be influenced by minority outliers.  Additionally, the eigenbasis should be quick to calculate and without the need for excessive memory storage.  Figure \\ref{fig:espec} illustrates the success of robust statistics for addressing point (1). The second point becomes clear when we investigate the distribution of the principal component amplitudes of the 3485 VVDS spectra. In Figure \\ref{fig:pcs} we present the first two principal component amplitudes for the VVDS spectra using each of the eigenbases. The overall correlation between these first two principal component amplitudes for the classic PCA indicates the failure of this eigenbasis to represent the variance in the majority of the galaxy spectra: the basis has been influenced by outliers.  A final, important aspect of the new algorithm is the increase in speed. The iterative truncation approach to classic PCA is clearly inefficient, although the precise increase in speed will vary depending on dataset properties. For our VVDS test case the 20 iterations of classic PCA take five times longer than a single iteration of the 3485 spectra using the robust algorithm. The ratio will naturally change in favour of the new technique for larger collections of spectra.  Our robust analysis is based on a randomised algorithm, and, as such, it assumes that the input data entries are considered in random order both for the initial set and the stream. When this is not the case, the method may develop a wrong initial representation of the data, which can take many iterations to correct.  The sensitivity of the algorithm to this issue is primarily determined by the parameter $p$, i.e., the number of principal components kept between steps. We expect problems only when this value is too low compared to the weights assigned to the new input vectors. In general, when studying an unknown dataset, we recommend that one randomises the dataset at each iteration and solves for as many eigenspectra as possible.  Having said that, we should also note that special ordering during the streaming of the data might prove invaluable for studying the evolution of the eigensystem as a function of the parameter used for sorting the input data. However, such studies should take extreme care in choosing the adjustable parameters (e.g., effective sample size) and ensure that the observed trends are real and stable to initial conditions."
    },
    "0809/0809.3599_arXiv.txt": {
        "abstract": "In searches for low-mass companions to late-type stars, correlation between radial velocity variations and line bisector slope changes indicates contamination by large starspots.  Two young stars demonstrate that this test is not sufficient to rule out starspots as a cause of radial velocity variations. As part of our survey for substellar companions to T Tauri stars, we identified the $\\sim$2 Myr old planet host candidates DN Tau and V836 Tau.  In both cases, visible light radial velocity modulation appears periodic and is uncorrelated with line bisector span variations, suggesting close companions of several M$_{Jup}$ in these systems.  However, high-resolution, infrared spectroscopy shows that starspots cause the radial velocity variations. We also report unambiguous results for V827 Tau, identified as a spotted star on the basis of both visible light and infrared spectroscopy.  Our results suggest that infrared follow up observations are critical for determining the source of radial velocity modulation in young, spotted stars. ",
        "introduction": "Extrasolar planets are common; over 300 systems have been discovered (e.g., exoplanet.eu). Recent studies have targeted higher mass \\citep{sat07, joh07}, lower mass \\citep{but06, end06}, and younger objects \\citep{pau06, set07}. Identifying young planets is important to define the time scale for planet formation and thus distinguish the possible formation process(es).  Young stars still surrounded by the circumstellar material which forms planets are typically located at distances of $>$100~pc and are thus inherently faint and often obscured. They also manifest strong magnetic activity \\citep[e.g.,][]{cmj07} and are highly spotted.  Numerous, large spots complicate detection of extrasolar planets through radial velocity (RV) monitoring \\citep{saa97} because a spot that is partially visible at all times on the surface of an inclined star mimics RV modulation \\citep[e.g.,][]{bou07, hue08}.  \\citet{pau06} studied 12$-$300~Myr old nearby stars and found no evidence for planets with masses $>$1$-$2~M$_{Jup}$ at the 3~$\\sigma$ level. \\citet{set07} identified a minimum mass 6.1~M$_{Jup}$ planet in a 852~day period orbit around the 100~Myr old G1$-$G1.5~V star HD~70573. More recently, \\citet{set08} reported a $\\sim$10~M$_{Jup}$ planet in a 3.56 day orbit around the 10~Myr old star TW Hya, although \\citet{huel08} identify this result as attributable to spots.  In this letter we present our observations of the young stars DN Tau, V836 Tau, and V827 Tau.  While the V827 Tau visible light data clearly implicate spots as the cause of the apparent RV variability, corresponding data for DN Tau and V836 Tau suggest the presence of giant planets.  Our infrared (IR) observations show that spots cause the RV variations seen in all three stars.  In \\S 2 we describe the observations, in \\S 3 present our data analysis and the evidence for spots, and in \\S 4 provide a brief discussion. We summarize in \\S 5. ",
        "conclusions": "Figures 1 and 2 show the visible light RV modulation for DN Tau and V836 Tau, with full amplitudes of $\\sim 1500$ and $\\sim 2700$~m~s$^{-1}$, respectively.  Within the 1~$\\sigma$ uncertainties, all but one point in the IR RVs of DN Tau are consistent with zero; all six IR RVs of V836 Tau also show no variation.  These results indicate that no planets are present around DN Tau or V836 Tau with masses greater than a few M$_{Jup}$ at $<$0.5~AU or $\\sim$10~M$_{Jup}$ at $\\sim$1~AU, despite the absence of a correlation between the visible light RVs and bisector spans. \\citet{mat89} identify V836 Tau and V827 Tau as RV variables with peak-to-peak amplitudes of 7$-$8~km~s$^{-1}$.  Apparently the density and size of spots on V836 Tau vary; historical data may serve as an additional criterion for heavily spotted young stars.  The primary conclusion from our visible light data is that the lack of a correlation between the line bisector span and the RV is not proof of a reflex motion companion.  In addition, the RV period can change significantly with new data.  Initial analysis of 20 visible light RVs for DN Tau, taken over 2.5 years, indicated convincing modulation with P$=$7.5 days and FAP$=$0.002.  Although not markedly different from rotation period estimates of 6.0$-$6.4 days \\citep{bou93, per06}, $\\chi^2$ for P$=$7.5 days was more favorable than that for a secondary peak at 6.3 days.  Removing the fit to either the 7.5 or 6.3 day period from the RV data and recalculating the power spectrum yielded no significant peak, suggesting that the data were best represented with only one period.  With the addition of RVs measured in winter 2007$-$2008, the integrity of the 7.5 day period was diminished and the 6.3 day period came to dominate the power spectrum. The rotation period of V836 Tau has been stable for decades at 6.76 days \\citep{gra08}.  We find a different RV period, 2.48 days, and substantially reduced power in the RV modulation near 6.76 days. The RV period for V827 Tau, 3.76 days, is indistinguishable from previously determined values of the rotation period \\citep{gra08}.  If the visible light RV modulation originates in spots, why don't the bisector spans correlate (Figures 1 and 2)?  Figure 3 clearly shows correlation for V827 Tau; many other stars show a correlation as well \\citep[e.g.,][]{que01, bou07, hue08}. We do not believe that this can be attributed to a stronger impact of spots on particular spectral lines because the same echelle orders were used to determine the bisector spans for all targets. \\citet{des07} show that when $v$sin$i$ is smaller than the spectrometer resolution, RV and bisector variations originating in spots can mimic the behavior expected from short period giant planets.  The $v$sin$i$ of DN Tau and V836 Tau, $\\sim$10~km~s$^{-1}$, are not much larger than our visible light spectral resolution, 5~km~s$^{-1}$, while V827 Tau and LkCa 19 (Huerta et al. 2008) have $v$sin$i$ values of $\\sim$20~km~s$^{-1}$, suggesting qualitative agreement with the simulations of \\citet{des07}."
    },
    "0809/0809.3766_arXiv.txt": {
        "abstract": "{% The progenitor mass of type IIP supernova can be determined from either hydrodynamic modeling of the event or pre-explosion observations. }{% To compare these approaches, we determine parameters of the sub-luminous supernova 2005cs and estimate its progenitor mass. }{% We compute the hydrodynamic models of the supernova to describe its light curves and expansion velocity data. }{% We estimate a presupernova mass of $17.3\\pm1~M_{\\sun}$, an explosion energy of $(4.1\\pm0.3)\\times10^{50}$ erg, a presupernova radius of $600\\pm140~R_{\\sun}$, and a radioactive $^{56}$Ni mass of $0.0082\\pm0.0016~M_{\\sun}$. The derived progenitor mass of SN~2005cs is $18.2\\pm1~M_{\\sun}$, which is in-between those of low-luminosity and normal type IIP supernovae. }{% The obtained progenitor mass of SN~2005cs is higher than derived from pre-explosion images. The masses of four type IIP supernovae estimated by means of hydrodynamic modeling are systematically higher than the average progenitor mass for the $9-25~M_{\\sun}$ mass range. This result, if confirmed for a larger sample, would imply that a serious revision of the present-day view on the progenitors of type IIP supernovae is required. } ",
        "introduction": "\\label{sec:intro} Type IIP supernovae (SNe~IIP) originate presumably from main-sequence stars of the mass range of $9-25~M_{\\sun}$ (Heger et al. \\cite{HFWLH_03}). If this is the case, the predicted rate of SNe~IIP should follow the Salpeter initial mass function with half of events occurring for stars of mass below $13~M_{\\sun}$. This conjecture requires confirmation by means of the determination of progenitor mass for an extended sample of SNe IIP.  There are two ways to recover the progenitor mass of SN~IIP on the main sequence. The first method is detection of the presupernova (pre-SN) in archival images of the host galaxy. The estimated flux and color index of the detected pre-SN is then converted into a stellar mass using the flux and color index predicted by stellar evolution models. Data for the available directly identified progenitors in the compilation of Li et al. (\\cite{LWV_07}) indicate that for eight SNe IIP the progenitor masses have been estimated in this way, and for six SNe IIP, upper limits to the progenitor mass have been found.  An alternative approach to the mass determination involves hydrodynamic modeling of light curves and expansion velocities for the well-observed SNe IIP. Combining the ejecta mass derived from the hydrodynamic modeling with the mass of the neutron star and the mass lost by the stellar wind provides us with the mass estimate of the progenitor. Henceforth, the progenitor mass determined by this method is referred to as the \"hydrodynamic mass\". At present, the hydrodynamic mass is measured only for three SNe IIP: the peculiar type IIP SN~1987A (Woosley \\cite{W_88}; Blinnikov et al. \\cite{BLBNI_00}; Utrobin \\cite{U_05}), the normal type IIP SN~1999em (Baklanov et al. \\cite{BBP_05}; Utrobin \\cite{U_07}), and the low-luminosity type IIP SN~2003Z (Utrobin et al. \\cite{UCP_07}). The small amount of SN IIP events with the measured hydrodynamic mass is related to the fact that the hydrodynamic modeling requires a complete multi-band photometry at both the plateau and the radioactive tail, and spectra of sufficient quality.  There are only a few other SNe IIP that also meet these requirements. Among these is the sub-luminous type IIP SN~2005cs (Pastorello et al. \\cite{PST_06}). On the basis of its luminosity and expansion velocities, this SN is intermediate between the low-luminosity and normal SNe~IIP. Parameters of SN~2005cs are of considerable interest for two major reasons: (1) this SN is expected to have intermediate parameters, which would be interesting to check; (2) the pre-SN was detected in the pre-explosion images and its progenitor mass was estimated by several groups (Maund et al. \\cite{MSD_05}; Li et al. \\cite{LVF_06}; Eldridge et al. \\cite{EMS_07}) providing an opportunity to compare different mass estimates.  In the paper, we perform hydrodynamic modeling of SN 2005cs to recover the parameters: ejecta mass, explosion energy, pre-SN radius, and radioactive $^{56}$Ni mass. We start with the description of the model and observational data used and then present the results of the hydrodynamic modeling for SN~2005cs (Sect.~\\ref{sec:opthydmod}). In Sect.~\\ref{sec:why}, we present additional arguments in favor of our choice of pre-SN models. Possible uncertainties in the hydrodynamic mass of SN 2005cs are analyzed (Sect.~\\ref{sec:prog}), and finally the implications of hydrodynamically studied objects for the origin of SNe~IIP are discussed (Sect.~\\ref{sec:discon}).  We adopt an explosion date on June 27.5 UT (JD 2453549) and a distance of 8.4 Mpc following Pastorello et al. (\\cite{PST_06}), and a reddening $E(B-V)=0.12$ taken from Li et al. (\\cite{LVF_06}). ",
        "conclusions": "\\label{sec:discon} \\begin{table}[t] \\caption[]{Hydrodynamic models for SN 1987A, SN 1999em, SN 2003Z, and SN 2005cs.} \\label{tab:sumtab} \\centering \\begin{tabular}{@{ } l @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ }} \\hline\\hline \\noalign{\\smallskip} SN & $R_0$ & $M_{env}$ & $E$ & $M_{\\mathrm{Ni}}$ & $Z$ & $v_{\\mathrm{Ni}}^{max}$ & $v_{\\mathrm{H}}^{min}$ \\\\ & $(R_{\\sun})$ & $(M_{\\sun})$ & ($10^{51}$ erg) & $(10^{-2} M_{\\sun})$ & & (km\\,s$^{-1}$) & (km\\,s$^{-1}$) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 87A &  35 & 18   & 1.5   & 7.65 & 0.006 & 3000 & 600 \\\\ 99em & 500 & 19   & 1.3   & 3.60 & 0.017 &  660 & 700 \\\\ 03Z & 229 & 14   & 0.245 & 0.63 & 0.017 &  535 & 360 \\\\ 05cs & 600 & 15.9 & 0.41  & 0.82 & 0.017 &  610 & 300 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{fig8.eps}} \\caption{% Explosion energy \\textbf{a}) and $^{56}$Ni mass \\textbf{b}) versus hydrodynamic progenitor mass for four core-collapse SNe. } \\label{fig:nienms} \\end{figure} The primary goal of this study was to determine parameters of the sub-luminous type IIP SN 2005cs by means of hydrodynamic modeling. We have estimated a pre-SN mass of $17.3\\pm1~M_{\\sun}$, explosion energy of $(4.1\\pm0.3)\\times10^{50}$ erg, pre-SN radius of $600\\pm100~R_{\\sun}$, and $^{56}$Ni mass of $0.0082\\pm0.0016~M_{\\sun}$. Using conservative assumptions about the mass-loss rate at the hydrogen and helium burning stages, we estimate the main-sequence progenitor mass of $17.2-19.2~M_{\\sun}$.  Hydrodynamic models for four SNe~IIP listed in Table~\\ref{tab:sumtab} are characterized by the pre-SN radius, the ejecta mass, the explosion energy, the total $^{56}$Ni mass, the surface abundance of heavy elements, the maximum velocity of nickel, and the minimum velocity of the hydrogen-rich envelope. The major parameters of SN 2005cs --- the ejecta mass, the explosion energy, and the $^{56}$Ni mass --- are intermediate between those of the low-luminosity type IIP SN 2003Z (Utrobin et al. \\cite{UCP_07}) and the normal type IIP SN 1999em (Utrobin \\cite{U_07}) in qualitative agreement with their luminosities. At present, there are therefore four SNe~IIP that have parameters determined by hydrodynamic modeling. For these objects, the explosion energy and $^{56}$Ni mass correlate with the progenitor mass (Fig.~\\ref{fig:nienms}). This is consistent with the empirical relation between the explosion energy and $^{56}$Ni mass found by Nadyozhin (\\cite{N_03}) for normal SNe~IIP.  Despite the uncertainties in hydrodynamic modeling, the disparity between the hydrodynamic mass of the SN~2005cs progenitor and the mass estimated from the pre-explosion images is significant. This difference could be decreased by including the effects of the pre-SN light absorption in a hypothetical dusty circumstellar shell. However, this issue requires careful consideration, which is beyond the scope of our paper. We note only that this conjecture has a number of observational implications need to be to verified. The presence of the dense dusty shell around pre-SN should produce strong Na\\,I absorption lines in the SN~IIP spectrum at the photospheric epoch. In the case of a normal SN~IIP and normal pre-SN wind without a dense circumstellar shell, the predicted Na\\,I absorptions are weak and probably not observable (Chugai \\& Utrobin \\cite{CU_08}). In addition, the interaction of the SN ejecta with the dense circumstellar shell should produce an outburst of radio and X-ray emission at an age of about $\\sim10^2$ days. The most apparent effect of the light absorption in the dusty circumstellar shell should be a large $J-K$ color index of the pre-SN. For instance, we have found that the pre-SN light absorption, required to allow massive progenitor implied by both the pre-explosion $I$ value and upper limits in other bands for SN~2005cs, suggests a large color index $J-K\\approx2.5$ mag, compared with the intrinsic $J-K$ index of $0.7-1$ mag for typical galactic K-M supergiants (Elias et al. \\cite{EFH_85}). In this regard, it is noteworthy that the type IIP SN~2008bk in the nearby galaxy NGC 7793 with available $JK$ photometry of the progenitor was found to have a moderate color index $J-K\\approx1$ mag (Maoz \\& Mannucci \\cite{MM_08}) which indicates little (if any) absorption. The observational and hydrodynamic studies of this supernova would be of significant importance in clarifying the serious and challenging problem of the progenitor mass of SN~2005cs and in general SNe~IIP.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[clip, trim=0 0 0 184]{fig9.eps}} \\caption{% Cumulative \"Salpeter\" distribution of SN~IIP progenitors in the $9-25~M_{\\sun}$ mass range and the distribution of hydrodynamic progenitor masses of four SNe~IIP. } \\label{fig:cumdis} \\end{figure} The number of SNe~IIP with measured hydrodynamic masses is too small to be able to analyze in detail and draw reliable conclusions about their mass distribution. However, the hydrodynamic progenitor masses do appear to be systematically higher than if SNe~IIP had originated from the range of $9-25~M_{\\sun}$, assuming a Salpeter initial mass function. This is clearly demonstrated by the comparison of the SNe~IIP mass distribution, calculated by assuming a Salpeter initial mass function in the $9-25~M_{\\sun}$ mass range, with the mass distribution of four SNe~IIP of known hydrodynamic mass (Fig.~\\ref{fig:cumdis}). Despite the small number of events, the significance of the difference between the two distributions is high: the probability that the four SNe occurred at random in the mass range of $15-25~M_{\\sun}$ is only 0.01. This indicates that either the 1D model of the SN explosion overestimates the ejecta mass, or the outcome of the core collapse of $9-15~M_{\\sun}$ stars differs markedly from that of SNe~IIP observed until now.  To study the first possibility, we would require 3D radiation hydrodynamics modeling, which is not possible at present. While we cannot readily ascertain any 3D effect that might reduce the required ejecta mass, apparent 3D effects do exist that could increase the hydrodynamic mass. Indeed, the RT mixing between the helium core and the hydrogen envelope is expected to produce heterogeneous ejecta consisting of helium clumps embedded in the hydrogen background. This structure should reduce unavoidably the average opacity of the hydrogen-rich matter. As a result, the ejecta mass required to reproduce the observations should increase. This could counterbalance other possible effects that might reduce the ejecta mass. We believe therefore that the 3D hydrodynamic simulations are unlikely to reduce significantly the SN~IIP progenitor masses recovered from the 1D modeling. The situation could be eventually clarified by 3D modeling of the SN~IIP explosion.  Alternatively, the disparity between the two distributions in Fig.~\\ref{fig:cumdis} is real. In this case two explanations could be invoked: (1) $9-15~M_{\\sun}$ stars produce faint, still undetectable SNe~IIP, i.e. we have a selection effect; (2) core collapse of stars from this mass range does not produce an SN event at all. In the latter case, the fate of the star may be silent collapse with a neutron star residing inside the stellar envelope, i.e. a Thorne-\\.{Z}ytkow object (Thorne \\& \\.{Z}ytkow \\cite{TZ_75}). If this is the case, the neutron star may eventually grow into a black hole of mass as high as $\\sim15~M_{\\sun}$ due to the rapid ($\\sim10^2$ years) accretion of the stellar envelope (Bisnovatyi-Kogan \\& Lamzin \\cite{BL_84}).  The second option appears, however, to be unlikely, because it implies that, apart from double neutron stars (DNS) originating from $9-25~M_{\\sun}$ stars, there should be a comparable number of binaries with a neutron star in combination with a black hole (NSBH binaries). Assuming that the production rate of DNS and NSBH binaries from $9-25~M_{\\sun}$ stars is determined by the random pairing of stars with the Salpeter initial mass function and that $9-15~M_{\\sun}$ stars produce black holes, we expect the relative rate of formation of these binaries to be $\\mathrm{DNS}:\\mathrm{NSBH}=1:0.85$ for $9-25~M_{\\sun}$ stars. This ratio is in an apparent contradiction with the fact that eight DNS in the Galaxy are known (Ihm et al. \\cite{IKB_06} and references therein) and no NSBH binary has yet been discovered. The probability of a random realization of this situation is only $1.5\\times10^{-5}$, i.e. sufficiently small to be able to exclude the second option that stars in the mass range of $9-15~M_{\\sun}$ end their lifes as black holes.  We propose that the core collapse of $9-15~M_{\\sun}$ stars should produce a neutron star and the remaining stellar matter ejected as a result of a faint SN event. This picture predicts that the rate of faint SNe~IIP should be comparable with the combined rate of normal (e.g., SN~1999em), sub-luminous (e.g., SN 2005cs), and low-luminosity (e.g., SN~2003Z) SNe~IIP. A detection of the extended class of faint SNe~IIP, or non-detection at a low flux level, could verify this scenario for $9-15~M_{\\sun}$ stars. It is interesting that some known transient events, e.g. SN~1997bs (Van Dyk et al. \\cite{VPK_00}), optical transient M85 OT2006-1 (Pastorello et al. \\cite{PDS_07}), and SN~2008S (Prieto et al. \\cite{PKT_08}) might belong to the proposed category of faint SNe related to the mass range of $9-15~M_{\\sun}$."
    },
    "0809/0809.1413_arXiv.txt": {
        "abstract": "\\noindent If a field theory contains gauged, non-Abelian, bi-fundamental fields {\\it i.e.} fields that are charged under two separate non-Abelian gauge groups, the transition from a deconfined phase to a hadronic phase may be frustrated. Similar frustration may occur in non-Abelian gauge models containing matter only in higher dimensional representations {\\it e.g.} models with pure glue, or if ordinary quarks are confined by two flux tubes, as implied in the triangular configuration of baryons within QCD. In a cosmological setting, such models can lead to the formation of a web of confining electric flux tubes that can potentially have observational signatures. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0809/0809.4222_arXiv.txt": {
        "abstract": "By extrapolating from observationally derived surface magnetograms of low-mass stars we construct models of their coronal magnetic fields and compare the 3D field geometry with axial multipoles. AB~Dor, which has a radiative core, has a very complex field, whereas V374~Peg, which is completely convective, has a simple dipolar field. We calculate global X-ray emission measures assuming that the plasma trapped along the coronal loops is in hydrostatic equilibrium and compare the differences between assuming isothermal coronae, or by considering a loop temperature profiles.  Our preliminary results suggest that the non-isothermal model works well for the complex field of AB~Dor, but not for the simple field of V374~Peg. ",
        "introduction": "Using the spectropolarimetric technique known as Zeeman-Doppler imaging (ZDI) it is possible to map the medium and large scale structure of stellar magnetic topologies (see Donati et al, these proceedings).  From such observationally derived surface maps of the photospheric magnetic fields of the rapidly rotating K-dwarf AB~Dor \\cite{don97} and the low mass M-dwarf V374~Peg \\cite{don06} we extrapolate the coronal fields assuming that they are potential, or current free, $\\nabla \\times \\mathbf{B} = \\mathbf{0}$ (Fig. \\ref{extraps}).  This condition is satisfied by writing the field in terms of a scalar flux function $\\Psi$, with $\\mathbf{B} =-\\nabla \\Psi$.  As the field must also be solenoidal, $\\nabla \\cdot \\mathbf{B}=0$, then $\\Psi$ must satisfy Laplace's equation $\\nabla^2 \\Psi=0$, the solution of which is a linear combination of spherical harmonics.  Thus the three components of the magnetic field vector can be determined within a defined 3D volume, and a field line tracing algorithm employed to derive the coronal magnetic field topology.  The extrapolation technique is described in greater detail by \\cite{jar02} and \\cite{gre06a}.  The stellar parameters for AB~Dor and V374~Peg are listed in Table \\ref{table_params}.   \\subsection{Comparison with axial multipoles} We compare the 3D coronal field structures obtained via field extrapolation with axial multipoles (Fig. \\ref{extraps}). The height and width of field line loops is defined in Fig. \\ref{hw}.  The width is calculated using the Haversine formula, \\begin{equation} w = 2\\sin^{-1}{\\left[\\sqrt{\\sin^2{\\left(\\frac{\\Delta \\theta}{2}\\right)}+\\sin{\\theta_1}\\sin{\\theta_2}\\sin^2{\\left(\\frac{\\Delta \\phi}{2}\\right)}}\\right]}, \\label{haversine} \\end{equation} where $\\Delta \\theta =|\\theta_2 - \\theta_1|$ and $\\Delta \\phi = \\phi_2 - \\phi_1$ are the differences in co-latitudes and longitudes of the field line foot points, and where all angles are measured in radians and all distances in units of stellar radii.  \\begin{figure*} \\centering \\begin{tabular}{cc} \\label{extrap_abdor} \\includegraphics[height=.30\\textheight]{gregory_poster_fig1a.pdf} & \\label{extrap_v374peg} \\includegraphics[height=.24\\textheight]{gregory_poster_fig1b.pdf} \\\\ \\label{hw_abdor} \\includegraphics[height=.29\\textheight]{gregory_poster_fig1c_reduced.pdf} & \\label{hw_v374peg} \\includegraphics[height=.29\\textheight]{gregory_poster_fig1d_reduced.pdf} \\\\ \\end{tabular} \\caption{Field extrapolations from magnetic surface maps of AB~Dor (left) and V374~Peg (right) [upper panel, not to scale]. The plots [lower panel] show field line height vs width, as defined in Fig \\ref{hw}, compared to axial multipoles, a dipole (red), a quadrupole (blue) and an octupole (green).} \\label{extraps} \\end{figure*}  V374~Peg, at only $0.3\\,{\\rm M}_{\\odot}$, is completely convective \\cite{don06} and has a large scale axisymmetric magnetic field which has been found to be remarkably stable over a timescale of at least one year \\cite{mor08}.  Its field is almost dipolar, as evident from both the field extrapolation and from the plot of field line height vs width (Fig. \\ref{extraps}). In contrast, AB~Dor is a solar mass star with a radiative core, the star spot distribution (and therefore the magnetic field) on which is known to vary on a timescale of at least one year \\cite{jef07}.  The magnetic field of AB~Dor is particularly complex (Fig. \\ref{extraps}) with contributions from many high order multipole components.  The difference in field complexity for both stars is likely to be a simple reflection of their different internal structures, and therefore magnetic field generation mechanisms.      \\begin{figure} \\includegraphics[height=.20\\textheight]{gregory_poster_fig2.pdf} \\caption{The width, $w$, of field lines is defined as the distance measured along the segment of the great circle connecting the field line foot points. It is calculated using the Haversine formula (equation \\ref{haversine}).  The height, $h$, is the maximum loop height above the stellar surface.} \\label{hw} \\end{figure}  \\subsection{Coronal X-ray emission} We assume that the X-ray emitting plasma is trapped along the coronal loops and calculate the pressure structure along loops by integrating the equation of hydrostatic equilibrium \\cite{gre06b}.  The pressure at a point $s$ along a loop is given by, \\begin{equation} p(s) = p_0 \\exp{\\left[ \\frac{\\mu m_H}{k_B}\\int_0^s \\frac{g_s}{T} ds \\right]}, \\label{hydro} \\end{equation} where $g_s$ is the component of the effective gravity (i.e. the component of the combined centrifugal and gravitational acceleration) along the direction of the field at $s$, and $g_s = \\mathbf{g}\\cdot\\mathbf{B}/|\\mathbf{B}|$.  Assuming that the pressure and density are related via $p=2nk_BT$ then the global X-ray emission measure can be calculated as described by \\cite{gre06b}. We set the coronal pressure to zero for open field lines and for loops which would be unable to contain the coronal plasma due to the gas pressure exceeding the magnetic pressure. Such regions therefore do not contribute to the X-ray emission.  X-ray emitting gas is typically contained within compact magnetic regions, however some regions of the stellar surface remain dark in X-rays, giving rise to modulation of X-ray emission, especially for AB~Dor.  Table \\ref{table} shows the simulated X-ray emission properties when assuming an isothermal corona (at 10$\\,{\\rm MK}$ for AB~Dor and 20$\\,{\\rm MK}$ for V374~Peg).  Also shown for comparison are the same quantities derived by considering a Serio-type temperature scaling law \\cite{ser81}, where the loop temperatures depend on the pressure at the loop foot points, the length of the loop, and a (somewhat arbitrary) heating scale length.  By comparing our simulated X-ray luminosities with those derived from ROSAT data (Table \\ref{table}), our preliminary results suggest that such a non-isothermal coronal model works reasonably well for AB~Dor with its complex magnetic structure, but fails for the simple field of V374~Peg.  \\begin{table} \\begin{tabular}{ccccc} \\hline \\tablehead{1}{c}{b}{Star \\\\} & \\tablehead{1}{c}{b}{Mass\\\\(M$_{\\odot}$)} & \\tablehead{1}{c}{b}{Radius\\\\(R$_{\\odot}$)} & \\tablehead{1}{c}{b}{P$_{rot}$\\\\(d)} & \\tablehead{1}{c}{b}{d\\\\(pc)} \\\\ \\hline AB~Dor   & 1.0   & 1.0  & 0.51 & 14.94 \\\\ V374~Peg & 0.3   & 0.35 & 0.45 & 8.963 \\\\ \\hline \\end{tabular} \\caption{Stellar parameters for V374~Peg and AB~Dor.} \\label{table_params} \\end{table}  \\begin{table} \\begin{tabular}{ccccccc} \\hline \\tablehead{1}{c}{b}{Star\\\\ \\\\} & \\tablehead{1}{c}{b}{T$_{corona}$\\\\$[$iso$]$\\\\(MK)} & \\tablehead{1}{c}{b}{log EM\\\\$[$iso$]$\\\\(cm$^{-3}$)} & \\tablehead{1}{c}{b}{log EM\\\\$[$non-iso$]$\\\\(cm$^{-3}$)} & \\tablehead{1}{c}{b}{log L$_X$\\\\$[$iso$]$\\\\(erg s$^{-1}$)} & \\tablehead{1}{c}{b}{log L$_X$\\\\$[$non-iso$]$\\\\(erg s$^{-1}$)} & \\tablehead{1}{c}{b}{log L$_X$\\\\$[$obs, ROSAT$]$\\\\(erg s$^{-1}$)} \\\\ \\hline AB~Dor   & 10   & 52.48 & 52.75 & 29.69 & 29.94 & 30.18 \\\\ V374~Peg & 20   & 52.13 & 52.71 & 29.16 & 29.91 & 28.88 \\\\ \\hline \\end{tabular} \\caption{Simulated/observed X-ray properties for V374~Peg and AB~Dor.} \\label{table} \\end{table} ",
        "conclusions": ""
    },
    "0809/0809.3526.txt": {
        "abstract": "In a previous paper, we have found that the resonance structure of the present Jupiter Trojan swarms could be split up into four different families of resonances. Here, in a first step, we generalize these families in order to describe the resonances occurring in Trojan swarms embedded in a generic planetary system. The location of these families changes under a modification of the fundamental frequencies of the planets and we show how the resonant structure would evolve during a planetary migration. We present a general method, based on the knowledge of the fundamental frequencies of the planets and on those that can be reached by the Trojans, which makes it possible to predict and localize the main events arising in the swarms during migration. In particular, we show how the size and stability of the Trojan swarms are affected by the modification of the frequencies of the planets. Finally, we use this method to study the global dynamics of the Jovian Trojan swarms when Saturn migrates outwards. Besides the two resonances found by \\citet{MoLeTsiGo2005} which could have led to the capture of the current population just after the crossing of the 2:1 orbital resonance, we also point out several sequences of chaotic events that can influence the Trojan population. ",
        "introduction": "\\label{sec:intro}    The discovery of Achilles by Wolf in 1906, and of four other Jovian Trojans the next year, gave a new impulse to the study of the triangular configurations of the three-body problem, whose existence was shown by Lagrange in 1772. An important problem was to establish the existence of a stable area surrounding the triangular equilibrium points $L_4$ and $L_5$ associated to the Sun and Jupiter. From a mathematical point of view, the application of K.A.M. theory to the planar and circular restricted three-body problem, gives a result of confinement between two invariant tori, ensuring the stability in the neighborhood of $L_4$ and $L_5$ for an infinite time \\citep{Leon1962,DeDe1967,Ma1972,MeSchmi1986}. In the case of the spatial restricted three-body problem, where K.A.M. theory does not ensure infinite time stability anymore,  \\citet{BeFaGo1998} applied an extension of the Nekhoroshev theorem (\\citeyear{Nekho1977}) to quasi-convex Hamiltonians in order to prove exponentially long-time stability.   But these two complementary theories do not give any information about the size of the stable region surrounding the triangular equilibrium points. In order to get estimates of this width, several authors developed Nekhoroshev-like estimates based on normalization up to an optimal order  \\citep{GiDeFoGaSi1989,CeGio1991,GiSko1997,SkoDo2001}. More recent results, based on improvements of these methods, can be found in \\citep{GaJoLo2005,EfSa05,LhoEftDvo2008}. Another very interesting result was published by \\citet{Ra1967}. In this work, based on the change of stability of the family of long-period Lyaponov orbits emanating from $L_4$, the author  gives an approximation of the limit eccentricity of stable tadpole orbits with respect to their amplitude of libration.  Except Gabern and Jorba who applied the same methods as in \\citep{GiDeFoGaSi1989}      to the bicircular, tricircular and biannular coherent problems \\citep{Ga2003, GaJo2004},  and \\citet{LhoEftDvo2008} who used a more sophisticated method in the elliptic restricted three-body problem (RTBP),  all of the above mentioned works were developed in the framework of the circular RTBP, planar or spatial, which is not a very realistic model for the purpose of studying the long-term dynamics of the Jovian Trojans. To overcome this problem, several authors performed full numerical integrations.  In \\citeyear{HoWi1993}, \\citeauthor{HoWi1993}, integrating test-particles during 20 Myr, obtained the first global stability result concerning the Trojans of the giant planets in our Solar system.  The same year, \\citeauthor{Milani93} developed a numerical method providing the proper elements of the Trojans of Jupiter. A  semianalytical  model of determination of these proper elements has also been proposed by \\citeauthor{BeaugeR01} in  \\citeyear{BeaugeR01}. Four years later, \\citet{LevisonSS97} described the long-time erosion of the Jovian Trojan swarms without identifying its mechanism (more results concerning the long-term behavior of the Trojans are published in   \\cite{TsiVaDvo2005}). A   few years later, \\citet{MiBeRo2001},  \\citet{NeDo02} and \\citet{MaTriSch2003}, showed the existence of two sets of unstable structures lying inside the Jovian Trojan swarms which could be related to the  long-time erosion. The authors suggested that one of these sets was connected to the great inequality (between Jupiter and Saturn) while the second one may have been generated by the commensurability between the libration frequency of the co-orbital resonance and the frequency $n_{Jup} - 2n_{Sat}$. In \\citeyear{RoGaJo2005},  \\citeauthor{RoGaJo2005} identified these singularities and analyzed their underlying resonances, this leading to the decomposition of the \"resonance structure\" in four families of resonances.   In \\citep{RoGa2006}, hereafter called \\pu, a new description of this resonant structure was given, and the link between these resonances and the long-term stability (or instability) of Jovian Trojans was established.  The structures described in \\pu, generated by commensurabilities between the proper frequencies of the Trojans and the fundamental frequencies of the planetary system, depend on the planetary configuration. Therefore, a small modification of the geometry of the system is able to modify the Trojan swarms' resonant structure, changing the global dynamics, and, consequently, the stability of the co-orbital region. Planetary migration in a planetesimal disc (see  \\citet{GoMoLe2004} and references therein) provides a natural mechanism of evolution of planetary systems. According to \\citet{MoLeTsiGo2005}, the Jovian co-orbital population is not primordial, but was, instead, captured during chaotic events which took place in the course of a planetary migration. The so-called Nice model \\citep{TsiGoMoLe2005} predicts that during the migration (inwards for Jupiter, and outwards for Saturn, Uranus and Neptune), the  Jupiter-Saturn couple crossed the 2:1 orbital resonance. \\citet{MoLeTsiGo2005} suggested that the Trojans were captured in co-orbital motion just after the crossing  of this resonance, when the dynamics of the Trojan region was completely chaotic. The authors show that this global chaos arises when the libration frequency of the Trojans is close to the combinations  $3(n_{Jup} - 2n_{Sat})$ or  $2(n_{Jup} - 2n_{Sat})$. This result was recently confirmed by  \\citet{MaScho2007}, showing that these resonances are the main generators of the depletion (and, probably also the capture) of the Trojan population.  The results reported in \\pu are directly related to the phenomenon described in  \\cite{MoLeTsiGo2005} and  \\citet{MaScho2007}. More precisely, the resonances involved in this process are members of one of the four families of resonances presented in \\pu. Consequently, this previous paper contains the necessary material to derive tools, that make it possible to study  the evolution of the Trojans resonant structure. This evolution is induced by changing the geometry in the considered planetary system. Indeed, our goal is not to give an accurate and realistic description of the behavior of the Trojan swarms during a migration process, but rather to study what are the planetary configurations for which the Trojan swarms become globally chaotic. Moreover, we want to develop a model that is capable of making predictions in a large class of problems, and that is not restricted to the study of the Jovian Trojans. Our approach is based on three different points. The first one rests on the understanding of the Trojans' resonant structure, and its decomposition. It will be shown in section \\ref{sec:generical} that, for a generic Trojan swarm, almost all resonances driving its global dynamics are members of four different families which generalize the families described in \\pu. The resonances of all these families involve both fundamental frequencies of the Trojans and the basic planetary frequencies. These families establish a link between the behavior of the Trojan swarm and the geometry of its planetary system. The two other ingredients  on which the method is based are respectively the exploration of the planetary frequencies and the determination of the frequency domain to which  the main frequencies of the whole Trojan swarm belong. In section \\ref{sec:Jupiter}, our method is applied to the study of the Jovian Trojans perturbed by Saturn, employing the model  used in \\pu but also in \\citep{MoLeTsiGo2005,MaScho2007}. This application reveals the rich dynamics of the Trojan swarms under the action of numerous resonances of different origins, which generate a resonance web which cannot be fully understood without knowing the fundamental frequencies of the considered dynamical system. ",
        "conclusions": "\\label{sec:Discussion}  The global dynamics of a Trojan swarm is shaped by its resonant structure. This structure depends crucially on the fundamental frequencies of the planetary system in which it is embedded. By generalizing the results of \\pu concerning Jupiter's Trojans, we show that in a general planetary configuration this resonant structure can be decomposed in four different families. Among these families, two are particularly sensitive  to the geometry of the planetary system: \\FII and \\FIV are dependent on the closeness of the planetary system to MMRs. It turns out that variations of the semi-major axis of the planets deeply modify the dynamical influence of these two families.  Based on this decomposition and on the dependence of the families on fundamental frequencies of the planetary system, we present a general method making possible the study of the global stability of Trojan swarms during planetary migration possible. We show that the knowledge of the evolution of the fundamental frequencies of the system during its migration suffices  to predict the main features of its global dynamical evolution (transition between total stability and strong instability), and to find  which planetary configurations which lead to stable or unstable swarms.   In our application to the case of the Jupiter Trojans disturbed by Saturn, we have varied one parameter: the semi-major axis of Saturn\\footnote{ Equivalent results are obtained if one varies, instead, both semi-major axes, i.e. the results depend only on the ratio of the two axes.}. While in a realistic migration, many other parameters would vary, this would modified only slightly the results obtained in this paper. For example, the variations of the planetary phases have almost no effect on our results (section \\ref{sec:phases}).  In the same way, the \"apsidal corotation\" of the planetary perihelia which, according to \\cite{MaScho2007}, strengthens the influence of the secular resonances, probably does not significantly modify the resonant structure of the Trojans. In section \\ref{sec:crossingFIV52} we have described the mechanism leading \\FIV to generate strong instabilities when the planetary system approaches a MMR.   Section \\ref{sec:FII} analyzes the dominant role of \\FII in the production of unstable regions during planetary migration. In particular, we demonstrate a mechanism of formation of huge  chaotic regions by the merging of several resonances crossing the Trojan swarm in opposite directions.  The resonant structure, and particularly its splitting in four families as defined in section \\ref{sec:RNp2BP} is generic. Moreover, the definitions of the families seem to hold for almost all planetary configurations. But the respective dynamical influences of the four families depend on the mass of the planet harboring the Trojan swarms. According to formula (\\ref{eq:freq_L4}), if this mass decreases, the decrease in the proper frequency $g$ is proportional to the planetary mass while the one in $\\nu$  is weaker (proportional to the square root of the planetary mass). The elements of \\FIV being approximated by the relation: $g \\approx \\ineg{p}{q}$, the smaller the planetary mass, the closer this resonance is to the $p:q$ MMR.  As a result, the smaller the mass of the planet, the more local the dynamical  influence of \\FIV.  In this case, we can expect that the role of the resonances of \\FIV becomes negligible compared with the role of \\FII. Conversely, the secular resonances (\\FIII) which do not involve the frequency $s$, seem to be negligible in the Jovian Trojan swarm. But their contributions increase for smaller planetary masses. This has been already mentioned in section \\ref{sec:RNp2BP}. This has been observed by  \\citet{Bo2008}  (available at http://www.imcce.fr/page.php?nav=en/publications/theses/\\, index.php) in the case of the Trojans of Saturn. In both cases, \\FII and \\FIV dominate the dynamics of the swarms during migration.  Finally, the method of prediction developed in this paper being very general, it can be used to study the dynamics of a large class of problems. Beside its straightforward application to the Trojans of a planet (results regarding the dynamics of the hypothetical Trojans of Saturn will be presented in a forthcoming paper), it can also be used to study the dynamics of satellites in a planetary system in migration, or more generally, to understand the behavior of a dynamical system undergoing quasiperiodic perturbations with slowly varying frequencies.    \\subsection*"
    },
    "0809/0809.4152_arXiv.txt": {
        "abstract": "In real telescopes, the optical parameters evolve with time, and the degradation is often not uniform. This introduces variations in the image profile and therefore photo-centre displacements which, unless corrected, may result in astrometric errors. The effects induced on individual telescopes and interferometric arrays are derived by numerical implementation of a range of cases. The results are evaluated with respect to the potential impact on the most relevant experiments for high precision astrometry in the near future, i.e. Gaia, PRIMA and SIM, and to mitigation techniques applicable from design stage to calibrations. ",
        "introduction": "\\label{intro}  The mirrors of ground based instruments suffer degradation of the reflecting surface, e.g. by oxidation or contamination of the coating, due to dust, moisture and chemical agents in the environments \\citep{Vucina}. The result is a degradation of the optical throughput, which in general is not uniform, but rather described by a complex, often patchy, distribution. Also, the overall telescope transmission is progressively reduced, so that it becomes necessary to process in particular the primary mirror regularly for removal of the degraded reflecting coating and application of a new one.  Space optics is expected to be much less affected by surface characteristics variations, thanks to the stable environment and suitable protective layers, but degradation e.g. due to $\\gamma$ radiation is still experienced, up to 1\\% of the overall throughput, depending on the choice of substrate and coating material \\citep{Baccaro}.  Even small variations are still potentially relevant with respect to high precision astrometric measurements, and as such their effects have been investigated and are discussed in this document. The main experiments in which astrometric errors induced by non-uniform optical throughput variation are potentially relevant are Gaia, PRIMA and SIM.  Gaia \\citep{Gaia} is the most important European space mission devoted to micro-arcsec ($\\mu$as) astrometry currently being implemented. Interferometry also aims at comparable levels of precision by means of both ground based arrays, e.g. the Phase Referenced Imaging and Microarcsecond Astrometry [PRIMA] facility \\citep{PRIMA} of the Very Large Telescope Interferometer [VLTI] of the European Southern Observatory [ESO], or from space, e.g. the Space Interferometer Mission [SIM] \\citep{SIM}.  The astrometric performance of an astronomical instrument depends on the detected signal profile. The Point Spread Function (PSF), i.e. the intensity distribution from a point-like object at infinity, is built from the diffraction integral, described in optics textbooks, e.g. \\citet{born}; the numerical implementation used here was described in a previous paper \\citep{gai07}. The PSF includes, by a suitable wavefront error (WFE) map, the realistic description of several optical characteristics. The polychromatic PSF is produced by superposition, weighted by spectral distribution, of the monochromatic PSFs.  Hereafter, $I_m$ is the monochromatic PSF at a given wavelength $\\lambda$, and $E_m$ is the corresponding complex amplitude distribution on the focal plane; $F$ is the effective focal length of the telescope. The coordinates on pupil and focal plane are respectively $\\{\\xi,\\,\\eta\\}$ and $\\{x,\\,y\\}$ (linear units). The complex amplitude depends on the pupil function $P$ and on the aperture transmission function $T$: \\begin{equation} \\label{eq:ampli_mono} E_m \\left( x, y \\right) = k \\int d\\xi \\, d\\eta \\, T \\left( \\xi, \\eta \\right) \\cdot P \\left( \\xi, \\eta \\right) \\ e^{-i\\pi \\left(x \\xi + y \\eta\\right) / \\lambda F} \\, . \\end{equation} The constant $k$ provides the appropriate photometric result associated with the source emission, exposure time and collecting area. The monochromatic PSF can be expressed as \\begin{equation} \\label{eq:PSF_mono} I_m \\left( x, y \\right) = \\left| E_m \\left( x, y \\right) \\right| ^{2} \\, . \\end{equation}  The pupil function $P \\left( \\xi, \\eta \\right) = e^{i \\Phi \\left( \\xi, \\eta \\right)}$ depends on the phase aberration function $\\Phi$, often expanded e.g. in terms of the Zernike functions $\\Phi_n$ \\citep{born}: \\begin{equation} \\label{eq:aberr} \\Phi \\left( \\rho ,\\theta \\right) = \\frac{2\\pi}{\\lambda} WFE = \\frac{2\\pi}{\\lambda} \\sum_n  A_n \\varphi_n \\left( \\rho, \\theta \\right) \\ . \\end{equation} The aperture transmission function $T$ is often considered to take unity value inside the pupil and zero outside; it is therefore used simply as a convenient numerical tool for insertion of the pupil geometric description. In the current simulation, however, it is used also to describe the variation of optical throughput of the telescope throughout the pupil. The formalism is similar to that adopted by \\citet{Linfield}, in the framework of straylight analysis for a Terrestrial Planet Finder (TPF) coronagraph. The most relevant changes are a simplified treatment, since we consider only transmission variations and neglect phase contributions, and the focus on astrometric effects, i.e. image photo-centre displacements.  We expect that a uniform variation of the optical throughput over the whole telescope aperture does not introduce any variation on the image photo-centre; however, for a generic aberrated configuration, a non-uniform pupil transmission introduces an amplitude modulation of the pupil function, thus modifying the resulting focal plane energy distribution, hence the measured signal profile, and consequently its photo-centre estimated position. Such photo-centre displacement is associated with an error on the astrometric measurement. Although in general the optical throughput is wavelength dependent, our analysis is referred to a single spectral type source, and uniform spectral dependence is assumed. Further levels of detail may be introduced in future investigations.  In section \\ref{sec_1tel} we discuss the introduction of transmission non-uniformity features and evaluate a set of reference cases focused on the astrometric performance of one telescope. In section \\ref{sec_interf} we analyse some consequences on the measurements of an interferometric array. Then, in section \\ref{sec_discuss}, we discuss the implications of our findings on design and calibration of astrometric instruments. Finally, we draw our conclusions. ",
        "conclusions": "Non-uniform variations of the optical throughput over the telescope aperture, as experienced in ground based instruments, may provide astrometric effects significant at the level of several microarcseconds. This may in turn represent a non negligible contribution to the error budget of the most challenging astrometric experiments currently under implementation or proposed for the near future, as Gaia, PRIMA and SIM.  The scale of the effects depends on manufacturing and operation parameters not easily measurable at the relevant level, i.e. order of 1\\% on mirror reflectivity over regions of few cm, or 0.02\\% on global throughput. Procedures of periodic calibration on sky can be used to retain acceptable levels of the systematic errors introduced by instrument response variations, in particular for long term measurements as interferometric orbits.  For Gaia, the time scale of instrument evolution must be compared with the period of effective parameter calibration, to verify that the systematic errors are adequately removed from the data. Information from the WFS and from photometry might be exploited to monitor astrometric effects, as well as diagnostics of the image profile evolution."
    },
    "0809/0809.1588_arXiv.txt": {
        "abstract": "In recent years, an increasing number of proper motions have been measured for Galactic X-ray binaries. When supplemented with accurate determinations of the component masses, orbital period, and donor effective temperature, these kinematical constraints harbor a wealth of information on the system's past evolution. Here, we consider all this available information to reconstruct the full evolutionary history of the black hole X-ray binary XTE\\,J1118+480, assuming that the system originated in the Galactic disk and  the donor has solar metallicity. This analysis accounts for four evolutionary phases: mass transfer through the ongoing X-ray phase, tidal evolution before the onset of Roche-lobe overflow, motion through the Galactic potential after the formation of the black hole, and binary orbital dynamics due to explosive mass loss and possibly a black hole natal kick at the time of core collapse. We find that right after black hole formation, the system consists of a $\\simeq 6.0 - 10.0 \\, \\rm M_{\\odot}$ black hole and a $\\simeq 1.0 - 1.6 \\, \\rm M_{\\odot}$ main-sequence star. We also find that that an asymmetric natal kick is not only plausible but \\emph{required} for the formation of this system, and derive a lower and upper limit on the black hole natal kick velocity magnitude of $80\\,\\rm{km\\, s^{-1}}$ and $310\\,\\rm{km\\, s^{-1}}$, respectively. ",
        "introduction": "In recent years the observed sample of Galactic black-hole (BH) X-ray binaries (XRBs) has increased significantly. For many of these systems there exists a wealth of observational information about their current physical state: BH and donor masses, orbital period, donor's position on the H-R diagram and surface chemical composition, transient or persistent X-ray emission, and Roche-lobe overflow (RLO) or wind-driven character of the mass-transfer (MT) process. The full proper motion has been measured for a handful of these systems \\citep{Mirabel2001, Mirabel2002, Mirabel2003}, which in addition to the earlier measurements of center-of-mass radial velocities and distances gives us information about the 3-dimensional kinematic properties of these binaries. This plethora of observational results provides us with a unique opportunity to study and try to understand the formation and evolution of BHs in binaries.  This paper is the second in a series in which we address outstanding questions about the formation of compact objects in XRBs,  such as what the mass relation between compact objects and their helium star progenitors is, and whether during the core collapse, BHs receive natal kicks comparable to those of neutron stars (NS).  XTE\\,J1118+480 is one of the Galactic X-ray novae that has been dynamically confirmed to contain a BH primary and the first to be currently located in the galactic halo \\citep{Remillard2000, Cook2000, Uemura2000}. After the first detection several groups performed detailed observations of the binary system \\citep{McClintock2001, Wagner2001, Mirabel2001, Haswell2002, Gelino2006, Gonzalez2006, Gonzalez2008}. These studies yielded strong observational constraints on the present properties of XTE\\,J1118+480.   The formation and evolution of BH XRBs, in such short-period orbits ($\\lesssim12\\,hr$) as that of XTE\\,J1118+480, is still a subject of open discussion. \\citet{Haswell2002}, assuming that the accreting material in XTE\\,J1118+480 has been CNO processed, showed that the system must have undergone RLO from a secondary star of about $\\simeq1.5\\,\\rm M_{\\odot}$ at a period of $\\simeq15\\, \\rm h$. \\citet{Gualandris2005} constrained the age of the system, using stellar evolution calculations, to be between 2 and 5\\,Gyr, rendering a globular cluster origin unlikely. They also studied the kinematic evolution of XTE\\,J1118+480, using Monte-Carlo techniques to account for the uncertainties in the observed values of the proper motion, and the distance of the binary. They inferred that the peculiar velocity of the system, acquired from the supernova (SN) explosion of the primary star, must be $183\\pm31\\,\\rm km\\,s^{-1}$ in order for the system's orbit to change from a Galactic disk orbit to the currently observed one.  \\citet{Justham2006} proposed an alternative scenario for the formation of BH XRBs with orbital periods shorter than half a day. They suggested that the companion star at the onset of the RLO, is an Ap or Bp intermediate mass star ($3\\,\\rm M_{\\odot}\\lesssim M \\lesssim 5\\,\\rm M_{\\odot}$). These types of stars are associated with  anomalously high magnetic fields. The authors suggested that the primordial magnetic field of Ap and Bp stars can lead to substantial systemic loss of angular momentum, through magnetic braking (MB), and evolve the binary to shorter orbital periods during MT. They showed that such a scenario can lead to the formation of BH binaries with population properties consistent with the observations, with exception of a discrepancy between the calculated effective temperatures and the observed spectral types of the donor stars.  On a different approach, \\citet{Yungelson2006} and \\citet{Yungelson2008} studied evolutionary models for populations of short-period BH XRBs, and compared them with observations of soft X-ray transients. Assuming a population of semi-detached short-period binaries with low-mass ($\\lesssim 1.5\\,\\rm M_{\\odot}$) donors and massive ($\\gtrsim 4.0 \\,\\rm M_{\\odot}$) compact accretors, and adopting a MB rate reduced by a factor of 2 compared to the \\citet{VZ1981} prescription, they showed that the calculated masses and effective temperatures of secondaries are in a satisfactory agreement with the observed ones, as inferred from their spectral types. The conclusions of this work suggest that the assumption of \\citet{Justham2006} for a new kind of MB mechanism in Ap and Bp stars is not necessary for explaining the formation of BH XRBs with short-period orbits.  \\citet{Ivanova2006} proposed that all short-period BH LMXB contained a pre-main-sequence (MS) secondary star at the time of the BH formation, with a subset potentially still containing a pre-MS donor at present, suggesting that the strong magnetic fields of these stars power the MB and bring the binary to contact. She showed that the orbital period and the donor's effective temperature of such X-ray binaries agree better with the available observations of BH LMXBs, than those of binaries with a MS donor, explaining simultaneously the roughly primordial abundance of Li detected in donor companions of Galactic BH LMXBs. However, this scenario is not applicable to XTE\\,J1118+480, as its orbital period of $0.17\\,\\rm days$ is located on the left of the ZAMS line \\citep[see Fig.~1 in ][]{Ivanova2006} and is explained by a MS donor star at the onset of RLO.  The works by \\citet{Justham2006,Yungelson2006, Yungelson2008}, and \\citet{Ivanova2006} study the population of short-period BH XRBs as a \\emph{whole}, trying to propose a general formation scenario for this kind of systems. However, detailed modeling of \\emph{individual} objects, which harness all the available observational data, is required in order to gain a better understanding about the properties of these systems, during the formation of the compact object.  \\citet{Willems2005} showed how using all the currently available observational constraints, one can  uncover the evolutionary history of the system from the present state back to the time just prior to the core-collapse event, and they applied their analysis to the BH XRB GRO\\,J1655-40. In the work presented here we follow the same framework but we focus on the case of XTE\\,J1118+480. The broader scope of this project is to address open questions related to BH formation such as the relation between the masses of BHs and those of their progenitors and whether natal kicks comparable to those of NS are imparted to BHs during core collapse.  The plan of the paper is as follows. In \\S\\,2 we review the currently available constraints on the properties of the Galactic BH XRB XTE\\,J1118+480. A general outline of the analysis used to reconstruct the system's evolutionary history is presented in \\S\\,3, while the individual steps of the analysis and the resulting constraints on the formation of the BH are discussed in \\S\\,4. In \\S\\,5, hydrodynamic core-collapse simulations are presented and the nature of the BH progenitor is constrained even further. In \\S,6, we discuss some of the assumptions introduced in our analysis and comment on possible alternative formation scenarios for XTE\\,J1118+480. The final section is devoted to a summary of our results and some concluding remarks. ",
        "conclusions": "\\subsection{Maximum Mass of the BH Progenitor}  In their analysis for GRO\\,J1655-40 \\citet{Willems2005} were able to place a strong constraint on the mass of the  helium star BH progenitor. The upper mass limit  came from the analysis of the orbital dynamics during the core-collapse event. In the case of XTE\\,J1118+480, the same analysis does not give a clear upper mass limit for  the BH progenitor. This is because of the different configuration of XTE\\,J1118+480 and mainly due to the operation of MB, which is able to shrink initially wide orbits on timescales shorter than the lifetime of the secondary star. However an upper mass limit  comes from the physics involved in the evolution of massive stars. In particular, \\citet[][ see Fig.~6]{MM2003}, adopting a full spectrum of stellar wind prescriptions, found that the maximum helium star mass they could achieve, evolving stars at solar metallicity and with ZAMS mass up to $100\\,\\rm M_{\\odot}$, was $\\sim 15\\,\\rm M_{\\odot}$ when they included moderate stellar rotation and  $\\sim 17.5\\,\\rm M_{\\odot}$ when they assumed no stellar rotation (both consistent with \\citet{Heg03,Young06,Young08} models). Adopting either of these two upper limits for the maximum mass of the BH progenitor star, instead of our ad-hoc $\\sim 20\\,\\rm M_{\\odot}$ limit, does not affect our constraint on the magnitude of the SN kick, so that our results are robust with respect to the helium star upper mass limit.   \\subsection{Globular cluster origin \\label{GC}}  \\citet{Mirabel2001} concluded a globular cluster origin of XTE\\,J1118+480 to be more probable than a Galactic disk origin, as only an extraordinary kick from the SN explosion could have launched the black hole into the observed Galactic halo orbit from a birthplace in the disk of the Galaxy. However, a globular cluster origin of XTE\\,J1118+480 generally implies a low metallicity donor star. Study of such MT sequences show  that the ones that successfully reproduce the observational properties usually have MT rates above or only marginally below the critical MT rate that separates transient from persistent behavior. In addition, the parameter space of donor mass and orbital period at the onset of RLO that leads to a system with the observed properties of XTE\\,J1118+480 is even more constrained than what we showed in Fig.~\\ref{MTseq} for solar metallicity. This is because we now have to account for an additional constraint on the current age of the donor star ($\\gtrsim 10\\,\\rm Gyr$) so that it is consistent with the age of the halo stellar population. This in turns limits the mass of the donor star at the onset of the RLO to $\\sim 0.9-1.0\\,\\rm M_{\\odot}$.  The strongest argument against a globular cluster origin of XTE\\,J1118+480 comes from chemical abundance analyses. \\citet{Gonzalez2006} reported that the secondary star has a supersolar surface metallicity $[Fe/H]=0.2 \\pm 0.2$, noting that a galactic halo origin of the system, although improbable, cannot be excluded, as the high surface metallicity might be due to pollution of the secondary star's surface during the SN event. Additionally, in our core collapse simulations, we had to require that the mass of the black hole is low ($<7\\,M_\\odot$) in order explain any nuclear enrichment in the companion star. Even in this case though, our successful MT sequences show that the donor star has lost at least $\\sim 0.8\\,\\rm M_{\\odot}$ of mass due to accretion onto the BH. So the outer layers of the star that could have been polluted from the SN ejecta are already removed and we shouldn't be able to observe the pollution effect now, unless the mixing in the donor star before the onset of MT was so strong that the heavy elements from the SN were transfered all the way to the inner layers of the star. Furthermore, in a subsequent work, \\citet{Gonzalez2008} used a grid of SN explosion models and they found that metal-poor models associated with a formation scenario in the Galactic halo provide unacceptable fits to the observed abundances, while the best agreement between the model predictions and the observed abundances was obtained for metal-rich progenitor models, corresponding to a birth of the system in the thin Galactic disk. While this cannot formally exclude a globular origin, it renders such a scenario rather improbable, as there are only two Galactic globular cluster (Terzan 5 and Liller 2) with $[Fe/H] \\gtrsim 0.0$ \\citep{Harris1996}.    \\subsection{CNO Processed Accreted Material}  \\citet{Haswell2002}, using ultraviolet spectroscopy, found that the emission line strengths in XTE\\,J1118+480 strongly suggest that the material accreting onto the BH has been significantly CNO processed. As a physical explanation of the observed spectrum they proposed  a scenario where the donor star exposed, due to the MT, its inner layers which have been mixed with CNO processed material from the central nuclear region. They concluded that the MT must have been initiated from a sufficiently massive ($\\sim 1.5\\,\\rm M_{\\odot}$) donor, so that the CNO process could change significantly the ratio of C to N to the very low value they observed.  However, due to the lack of a precise determination of the current C to N ratio, we cannot set a clear lower limit on the initial donor mass based on the presence of CNO processed material. Our successful MT sequences suggest that the current mass of the donor star is in the range of $0.15$--$0.35\\,\\rm M_{\\odot}$. We compared our findings with the online library of single star evolutionary sequences by \\citet{Paxton2006}, which includes the chemical composition of stars in terms of fractional abundance as a function of the mass coordinate. We find that even for an $1.1\\,\\rm M_{\\odot}$ secondary star, by the time of the onset of RLO, the ratio of C to N in the central $0.3\\,\\rm M_{\\odot}$ is low enough ($log(C/N)\\lesssim -1$) that it would be consistent with the observational findings of \\citet{Haswell2002}. In contrast to the latter authors, we therefore did not set a lower limit of $\\sim 1.5\\,\\rm M_{\\odot}$ on the donor mass in our MT calculations but instead adopt a more conservative approach and study initial donor masses down to $1.0\\,\\rm M_{\\odot}$."
    },
    "0809/0809.3896_arXiv.txt": {
        "abstract": "We investigate coronal transients associated with a GOES M6.7 class flare and a coronal mass ejection (CME) on 13 July 2004. During the rising phase of the flare, a filament eruption, loop expansion, a Moreton wave, and an ejecta were observed. An EIT wave was detected later on. The main features in the radio dynamic spectrum were a frequency-drifting continuum and two type II bursts. Our analysis shows that if the first type II burst was formed in the low corona, the burst heights and speed are close to the projected distances and speed of the Moreton wave (a chromospheric shock wave signature). The frequency-drifting radio continuum, starting above 1 GHz, was formed almost two minutes prior to any shock features becoming visible, and a fast-expanding piston (visible as the continuum) could have launched another shock wave. A possible scenario is that a flare blast overtook the earlier transient, and ignited the first type II burst. The second type II burst may have been formed by the same shock, but only if the shock was propagating at a constant speed. This interpretation also requires that the shock-producing regions were located at different parts of the propagating structure, or that the shock was passing through regions with highly different atmospheric densities. This complex event, with a multitude of radio features and transients at other wavelengths, presents evidence for both blast-wave-related and CME-related radio emissions. ",
        "introduction": "Decimetric--metric radio emission is usually associated with plasma waves excited by accelerated electrons, where the observed emission frequency is directly related to the plasma density in the affected region \\cite{melrose80}. Fast-drifting structures in radio dynamic spectra, such as type III bursts, are normally attributed to electron beams propagating outward in the solar atmosphere. The type III burst-associated acceleration usually happens in connection with flaring processes, by for example, reconnecting field lines. A slower drift and a broader instantaneous emission band in a radio burst can indicate a moving shock front that accelerates electrons. Some of these bursts can be classified as type II bursts, in particular when both fundamental and second harmonic emission bands are visible ({\\it e.g.}, Nelson and Melrose, \\citeyear{nelson85}; Mann, Classen, and Aurass, \\citeyear{mann95}; Cairns {\\it et al.}, \\citeyear{cairns03}). Although type II bursts indicate the existence of shock waves, their origin is not always so clear. Shocks can basically be flare related or coronal mass ejection (CME) related, but coronal (metric) type II bursts do not necessarily have the same origin as the longer wavelength (interplanetary) type II bursts (Mancuso and Raymond, \\citeyear{mancuso04}; Cane and Erickson, \\citeyear{cane05}; Lin, Mancuso, and Vourlidas, \\citeyear{lin06}, and references therein).  Theoretically, shocks can be freely propagating (blast) waves, piston-driven shocks formed either at the nose or the flanks of an expanding body, or bow shocks ahead of a projectile \\cite{warmuth07}. A freely propagating shock can also develop if a piston stops or slows down considerably. One characteristic of piston-driven shocks is that the speed of the propagating shock is different from the speed of the driver, and the shock speed should be higher. For bow shocks ahead of a projectile, the speeds should be roughly the same, see, for example, \\inlinecite{vrsnak05}.  Frequency-drifting continuum emissions at centimeter to meter wavelengths, also known as type IV bursts, have been associated with outward moving plasmoids, see review on solar radio emission in, for example, \\inlinecite{dulk85}. At decimetric wavelengths, the usually smooth type IV continuum may include pulsating structures \\cite{kundu65}. In some cases decimetric pulsating structures can be taken as precursors to type II bursts \\cite{klassen99}, and identified as a signature of reconnection processes above expanding soft X-ray loops \\cite{klassen03}.  This study presents a flare--CME event on 13 July 2004, during which a rich variety of radio bursts were observed from centimeter to meter wavelengths. In this paper, we concentrate on the type II and type IV bursts observed at decimeter--meter wavelengths. During this event, quasi-periodic pulsations were observed in microwaves and decimeter waves, but they will be discussed in another paper. We compare the timing, heights, and structural evolution of the observed features and discuss conditions that are needed for the radio emissions to appear. ",
        "conclusions": "We can summarize the multi-wavelength analysis of the 13 July 2004 event with the following results:  1. The event started with a filament eruption, followed closely by the appearance of a frequency-drifting radio continuum. The constructed microwave--meterwave dynamic spectrum (Figure \\ref{fig:fig6}) suggests that an expansion front was formed at the time of the filament eruption, almost two minutes prior to any signatures of shock wave formation at other wavelengths. This expanding plasma volume could have been driving a shock wave.  2. The first type II burst was observed in the dynamic spectrum soon after the EUV ejecta lifted off (from below the EUV loop system) and the Moreton wave appeared. The estimated type II burst velocity is similar to the Moreton wave velocity ($\\approx$\\,700 km s$^{-1}$), if the type II burst source height is set untypically low, \\mbox{$\\lesssim$ 0.2 $R_{\\odot}$} above the photosphere. The type II burst heights then match with the projected distances of the Moreton wave front. We do not expect the Moreton wave and type II burst speeds to be exactly the same. According to Uchida's model \\cite{uchida74}, the H$\\alpha$ Moreton wave is created when the \"skirt\" of a shock wave sweeps the chromosphere. The shock velocity upward probably differs from the horizontal velocity, but the observed velocities can still be used as rough estimates for the shock speed (like the type II speeds). The backward-extrapolation of a shock propagating with a constant speed of 725 km s$^{-1}$ suggests a common start time with the frequency-drifting radio continuum (Figure \\ref{fig:fig5}). For comparison, at more typical type II heights, $\\approx$\\,0.3\\,--\\,\\mbox{0.4 $R_{\\odot}$}, the calculated burst velocities get much higher, $\\approx$\\,1\\,000\\,--\\,\\mbox{1\\,700 km s$^{-1}$}, and no correlation with  the heights of other features can be found.  3. The second type II burst had a frequency drift that suggests a speed different from that of the first type II burst, thus implying a different driver for the shock. Calculations show that the type II bursts could have been formed by one shock only, if the shock propagated at a constant speed and the later shocked region was located lower in the shock driver structure ({\\it e.g.}, if the first type II burst was created at the nose of expanding loops and the second one on the sides). The estimated height separation, \\mbox{$\\approx$\\,0.1 $R_{\\odot}$}, is not large enough to indicate a separation between, for example, a shock in front of a CME and at the flanks. A decelerating shock would not have been able to produce the observed frequency drift of the second type II burst. A single shock propagating through regions with very different densities, such as first through low density structures and later through high-density streamer regions, would also produce similar jumps in the type II source heights. Despite these different interpretations, the estimated heights for the second type II burst are close to the observed EIT wave distance and the backward-extrapolated CME front distances, although the speeds differ considerably.  Freely propagating shock waves can be formed in flare blasts or when a piston is stopped or it slows down ({\\it e.g.}, \\opencite{vrsnak01}). The speed of the propagating shock and the speed of the initial driver can then differ considerably. Also the shock should be spatially well separated from the driver. (In CME-driven shocks observed near Earth, the standoff distance between the shock and its driver has grown to 20\\,--\\,50 R$_{\\odot}$, see Reiner {\\it et al.}, \\citeyear{reiner07}). \\inlinecite{pomoell08} have shown in their simulations how the speed of the shock (and hence the corresponding type II burst speed) can far exceed the speed of the driver, and the shock can continue propagation through the corona even if the driver speed falls below the local magnetosonic speed. Recent radio imaging observations \\cite{dauphin06} show a type II burst located ahead of a rising soft X-ray loop. The speed of the rising loop was only 650 km s$^{-1}$ at the start time of the type II burst. As stated in their paper, and in another paper analyzing the same event \\cite{vrsnak06}, the derived speeds for the type II shocks (obtained with different methods, but in the range of 1\\,100\\,--\\,1\\,800 km s$^{-1}$) are much higher than those of their assumed drivers.  Based on the theory and our calculations for burst heights and speeds, and the observation of the Moreton wave, it is possible that the first type II burst was formed by a blast wave. The earlier start of the frequency-drifting radio continuum suggests that a shock wave may also have been ignited by a fast-expanding structure. Loop-like fronts, such as the rising EUV loop in this case, can be the low-coronal signatures of CMEs \\cite{rust76,vrsnak04b,temmer08}, that drive the shock. \\inlinecite{gary84} presented a case where a shock is driven by rising loops and a type II burst is due to a flare blast that starts later. In this scenario the type II burst source is located within the ejecta material. Type II emission can also be produced when a flare shock overtakes regions of principal density pile-up near the base and sides of an expanding transient \\cite{wagner83}. This could also explain our low radio source heights.  \\inlinecite{subra06} have analyzed statistically several doublet type II bursts, and concluded that usually the second type II burst starts at a lower frequency than the first one and the drift rate is only about half of the first one. In our event the frequency drift of the second type II burst is slower by a factor of 4. The rare imaging observations of multiple type II bursts \\cite{robinson82} suggest that a single shock may exist, but as the shock propagates at different locations the type II bursts may have different characteristics. To verify this, imaging observations are needed at radio and other wavelengths, to determine source locations and corresponding electron densities.  We have presented analysis of a very complex flare-CME event, with a multitude of radio features and transients at other wavelengths, and obtained evidence for both blast wave shock and CME-related emissions during the initial phase. However, definite conclusions will rely on future radio imaging. It is unfortunate that radio imaging currently exists only at very narrow frequency ranges, leaving out plasma emission above 500 MHz and below 150 MHz. For the identification of the source locations in decimetric--metric radio emission, future radio imaging instruments such as the Frequency-Agile Solar Radiotelescope (FASR) will be needed, together with high-cadence multi-wavelength observations. Statistical analyses with comparison to eruption characteristics could also clarify the processes involved."
    },
    "0809/0809.0895_arXiv.txt": {
        "abstract": "We report seven successful observations of the astrometric binary GJ 164 AB system with aperture masking interferometry. The companion, with a near infrared contrast of 5:1 was detected beyond the formal diffraction limit. Combined with astrometric observations from the literature, these observations fix the parallax of the system, and allow a model-independent mass determination of both components. We find the mass of GJ 164B to be $ 0.086 \\pm 0.007 M_{\\sun}$. An infrared spectroscopic study of a sample of M-Dwarfs outlines a method for calibrating metallicity of M-Dwarfs. Results from the newly commissionned TripleSpec spectrograph reveal that the GJ 164 system is at least of Solar metallicity. Models are not consistent with color and mass, requiring a very young age to accommodate a secondary too luminous, a scenario ruled out by the kinematics. ",
        "introduction": "The development of a self-consistent theory of the internal structure of stars and the construction of models which trace evolutionary behaviour are major achievements of modern astrophysics. Both theory and models have however long been centered on intermediate and high mass stars, with little consideration of objects with masses below 0.6 M$_{\\sun}$. The realization that the Solar neighborhood is overwhelmingly dominated by low mass stars \\citep{1971ARA&A...9..103V, 1998ASPC..134...28H}, and the discovery of brown dwarfs \\citep{1995Natur.378..463N} have brought a lot more attention to the lower end of the main sequence and what lies below.  The complex physical nature of these very low mass objects makes their modeling intricate. Boundaries between stars, brown dwarfs and planets, as well as theoretical relations that predict the luminosity and temperature as functions of age, mass and metallicity are still largely untested in the relevant range of parameters. Correct calibration of these models is of extreme importance, since they are now extrapolated to estimate the mass of brown dwarfs \\citep{2006ApJ...640.1063B} and giant exoplanets \\citep{2003A&A...402..701B} from their luminosity and age.  The observations required to challenge and improve the models are dynamical mass measurements of multiple star systems, combined with accurate photometry and distance determination. Besides gravitational microlensing \\citep{1986ApJ...304....1P}, the observation of binary systems and the use of Newtonian orbital dynamics provide the only method of directly measuring accurate stellar masses. With maximum sensitivity at separations greater than 1 arcsec, conventional Adaptive Optics (AO) observations at Palomar are limited to the most nearby objects or the consequent orbital periods become too long to lead to dynamical masses. We have therefore been using aperture masking interferometry \\citep{2000PASP..112..555T} in conjunction with AO \\citep{2006ApJ...650L.131L}. The precision calibration of the data achieved with this observing mode indeed leads to reliable results up to and beyond the formal diffraction limit (super resolution), and very precise photometry. Combined with radial velocity \\citep{2007ApJ...661..496M} or astrometry \\citep{2008ApJ...678..463I}, aperture masking interferometry has so far provided some of the most precise (dynamical) masses of objects below 0.1 M$_{\\sun}$. ",
        "conclusions": "GJ 164 is an astrometric binary, whose orbit was resolved by aperture masking interferometry observations at Palomar and Keck. The consistency of the data is excellent, and exhibits a 1-2 mas precision in average for each measurement, made below what is usually accepted as the resolution limit of a telescope. This data, combined with earlier astrometric measurements, provides the following dynamical masses: 0.257 $\\pm$ 0.020 M$_{\\sun}$ for the primary and 0.086 $\\pm$ 0.007 M$_{\\sun}$ for the secondary.  Analysis of its infrared spectrum reveals that the M-Dwarf binary GJ 164 is at least of Solar metallicity. When attempting to use theoretical models for Solar metallicity very low-mass stars to derive an age from the mass and luminosity of both components, we find that the models do not adequately fit the data, requiring very young ages to accomodate for an overluminous secondary. The models therefore underestimate the luminosity of very low-mass stars that have settled on the main sequence.  The precision of the dynamical data we present here, of the order of 10 percents, however does not provide a strong constraint for the models. A precise measurement of the contrast ratio in the R band, where the astrometry data was taken would already help to better constrain the mass ratio and ultimately improve the constraint on the individual masses, and make our statement about the models stronger. The very small separation (at most, 90 mas) makes this measurement difficult, even for HST. But because it lies so close from the substellar limit, GJ 164 B is a target that deserves most attention. In the meantime, a Radial Velocity curve would provide an independent measurement of the mass ratio. With an expected Radial Velocity amplitude of 2.8 km/sec, and a period of two years, observations of GJ 164 with TripleSpec and TEDI will settle this uncertainty."
    },
    "0809/0809.0547_arXiv.txt": {
        "abstract": "We aim to obtain a complete sample of redshift $z \\ge 3.6$ radio QSOs from FIRST sources ($S_{\\rm 1.4 ~GHz} > 1$ mJy) having star-like counterparts in the SDSS DR5 photometric survey ($r_{\\rm AB} \\le 20.2$). Our starting sample of 8665 FIRST-DR5 pairs includes 4250 objects with spectra in DR5, 52 of these being $z \\ge 3.6$ QSOs.  We found that simple supervised neural networks, trained on the sources with DR5 spectra, and using optical photometry and radio data, are very effective for identifying high-$z$ QSOs in a sample {\\it without} spectra. For the sources with DR5 spectra the technique yields a completeness (fraction of actual high-$z$ QSOs classified as such by the neural network) of 96 per cent, and an efficiency (fraction of objects selected by the neural network as high-$z$ QSOs that actually are high-$z$ QSOs) of 62 per cent. Applying the trained networks to the 4415 sources {\\it without} DR5 spectra we found 58 $z \\ge 3.6$ QSO candidates. We obtained spectra of 27 of them, and 17 are confirmed as high-$z$ QSOs.  Spectra of 13 additional candidates from the literature and from SDSS DR6 revealed seven more $z \\ge 3.6$ QSOs, giving an overall efficiency of 60 per cent (24/40).  None of the non-candidates with spectra from NED or DR6 is a $z \\ge 3.6$ QSO, consistently with a high completeness. The initial sample of high-$z$ QSOs is increased from 52 to 76 sources, i.e. by a factor 1.46. From the new identifications and candidates we estimate an incompleteness of SDSS for the spectroscopic classification of FIRST $3.6 \\le z \\le 4.6$ QSOs of 15 per cent for $r \\le 20.2$. ",
        "introduction": "Homogeneous statistical samples of high-redshift quasi-stellar objects (QSOs) allow not only investigation of the QSO phenomenon itself, but also provide important information for a wide variety of studies. In particular, the luminosity function of high-redshift QSOs provides strong constraints on the theory of the accretion of matter onto supermassive black holes in the nuclei of galaxies. The increasing evidence for a relation between the formation of galaxy bulges and supermassive black holes (Kormendy \\& Richstone 1995, Magorrian et al. 1998) emphasises the importance of understanding the role of QSO activity in the formation and evolution of galaxies. The luminosity function of QSOs is also essential to quantify their contribution to the X-ray background and the UV ionising flux at high redshift.  In addition, the absorption spectra of these QSOs reveal the state of the intergalactic medium at early epochs.  Although radio-loud (RL) QSOs are a small subset of the QSO population, samples of high-redshift RL QSOs benefit from higher completeness, due to the drastically reduced contamination by stars in samples of radio selected QSO candidates, compared to optically (colour) selected QSO candidates (Richards et al. 2006).  Moreover, the connection between radio and optical activity, which still needs to be understood, requires a comparison between radio-loud and radio-quiet QSO populations. Ivezik et al. (2004) provide conclusive evidence that the distribution of radio-to-optical flux ratio for QSOs, i.e. the radio loudness, is bimodal (the so-called QSO radio dichotomy), on the basis of accurate optical and radio measurements of a large sample of RL QSOs obtained from the Sloan Digital Sky Survey (SDSS; York et al. 2000) and the Faint Images of the Radio Sky at Twenty cm survey (FIRST; Becker, White \\& Helfand 1995). Many studies suggest that RL QSOs reside in more massive galaxies and harbour more massive central black holes than radio-quiet QSOs, but the point is still controversial (see references for and against these arguments at Cirasuolo et al. 2006). A recent study by Jiang et al. (2007) based on a QSO sample drawn from SDSS and FIRST shows that the fraction of RL QSOs decreases with increasing redshift and with decreasing optical luminosity.  We aim to obtain a homogeneous sample of high-redshift RL QSOs ($z$ above $\\sim 3.6$) drawn from correlation of the FIRST catalogue ($S_{\\rm 1.4 ~GHz} > 1.0$ mJy) with unresolved objects in the SDSS Data Release 5 (SDSS DR5, Adelman-McCarthy et al. 2007). The area of overlap between FIRST and the DR5 imaging survey is $\\sim$ 7391 deg$^2$ and the number of selected FIRST-SDSS matches is 8665 (Section 2). SDSS provides: (i) $ugriz$ photometry, which is a powerful tool for separating high-$z$ QSOs from other populations (e.g. stars, QSOs with $z$ below $\\sim 3.6$ or unresolved low-$z$ galaxies); (ii) morphological classification, essential for distinguishing between high-$z$ QSOs and galaxies or resolved low-$z$ active galactic nuclei; and (iii) spectroscopy of many of our candidates (4250), selected as spectroscopic targets by SDSS DR5. Since SDSS spectroscopic observations necessarily lag the imaging, the total DR5 spectroscopic area is lower, with $\\sim 5553$ deg$^2$  included in the overlap with FIRST. Most of the candidates with available spectroscopy were classified by SDSS as QSOs, i.e. have a secure detection of a high-excitation emission line with FWHM $\\ge$ 1000 km sec$^{-1}$. The rest are galaxies, stars and objects of `unknown' class. 52 DR5 sources were spectroscopically classified as $z \\ge 3.6$ QSOs (Table 1).  Our approach to obtaining a high-$z$ QSO sample was to extend the existing sample of 52 FIRST-DR5  high-$z$ QSOs by applying automated learning techniques, specifically neural networks (NNs), to the 8665 FIRST-SDSS DR5 photometric matches. NNs have been shown to be powerful tools for both classification and regression tasks, in many fields of astronomy, and have subsequently been applied to predict object classes and/or astrophysical parameters. Fields where NNs have been applied include: classification of stellar spectra (Bailer-Jones, Irwin \\& von Hippel 1998); morphological star/galaxy separation (e. g. Bertin \\& Arnouts 1996); morphological classification, spectral typing and/or photometric redshifts of galaxies [Folkes, Lahav \\& Maddox 1996; Lahav et al. 1996; Firth, Lahav \\& Somerville 2003; Collister \\& Lahav 2004 (this paper reports the popular photometric redshift code ANNz); Ball et al. 2004]; QSO identification and/or QSO photometric redshifts (Carballo et al. 2004; Claeskens et al. 2006); and cross-matching of astronomical catalogues (Rohde et al. 2005).  QSO selection and estimation of QSO photometric redshifts are of prime importance for the SDSS project. Various studies address the problem using different machine learning approaches. Richards et al. (2004) applied a probability density analysis based on kernel density estimation of the colour distribution of stars and spectroscopically confirmed QSOs in SDSS DR1, to classify, as stars or QSOs, a catalogue of over $10^5$ unresolved, $g \\le 21$ mag, UV-excess ($u - g \\le 1$) QSO candidates. The resulting efficiency and completeness (the latter evaluated for $g \\le 19.5$) for the selection of QSOs in the candidate sample was estimated to be around 95 per cent up to $z \\simeq 2.4-3.0$, the redshift limit mainly arising from the restriction of the catalogue to UV-excess objects.  Suchkov, Hanish and Margon (2005) applied the oblique decision tree classifier ClassX to classify SDSS-DR2 photometric objects into 25 classes [stars, red stars, 10 redshift bins for galaxies, and 13 for Active Galactic Nuclei (hereafter AGN)] using colour information and morphology (attributes `resolved' or `unresolved') from SDSS. For each of the 12 redshift bins for AGN with $\\Delta z = 0.2$ and covering $0 \\le z \\le 2.4$, the completeness obtained for the test sample is in the range from 43 to 81 per cent, with an average 63 per cent. For the high-redshift bin, in the range $z=2.4-6.0$, the completeness drops to 14 per cent, the remaining high-$z$ AGNs beeing classifified as stars (47 per cent) or as AGN in the adyacent redshift bin $2.2 \\le z \\le 2.4$ (39 per cent). This result illustrates the difficulty in separating high-$z$ QSOs from other classes. The efficiency or fraction of true high-$z$ QSOs among the sources classified in the AGN high-$z$ bin was $\\sim$ 75 per cent.   Ball et al. (2006) applied decision trees, trained on the SDSS-DR3 objects with available spectroscopy, to classify all photometric objects ($> 10^8$) in SDSS-DR3 in one of the three categories of star, galaxy or nsng (neither star nor galaxy), the latter including QSOs and `unknown'. A blind test on the 2dF QSO Redshift Survey (2QZ; Crom et al. 2004), using the 8739 QSOs matching 2QZ and SDSS-DR3, yielded 95 per cent completeness and 87 per cent efficiency. The authors do not discuss how the performance depends on redshift.  Bazell, Miller and SubbaRao (2006) use a semisupervised mixture model approach to analyze 10000 objects spectrosopically classified in SDSS-DR4 in the categories of stars, late-type stars, galaxies and QSOs with $z \\le 3$ and unknown, using as input data for the modelling SDSS colours and the spectroscopic class. Since the aim was to investigate the existence of possible new object types among the class of `unknown' as well as subclasses among the remaining classes, 90 per cent of the sources in the categories of stars, late-type stars, galaxies and QSOs were also treated as unknown during the modelling. The best model includes 16 components, two of them of the nonpredefined type, and one of the latter captures a region of the $u-g$ versus $g-r$ colour-colour diagram ($2 \\le u-g \\le 5$, $0.5 \\le g-r \\le 1.5$) within the location of high-$z$ QSOs in Richards et al. (2002), but intentionally rejected in the QSO selection by Richards et al. (2004) because of the high density of stellar contaminants in that region.  Gao, Zhang and Zhao (2008) compare the performance of $k$-dimensional trees and support vector machines in the separation between stars and QSOs, using a sample of stars and QSOs spectroscopically classified in SDSS-DR5 and having a counterpart at the Two-Micron All Sky Survey (2MASS, Cutri et al. 2003). Both techniques yield a global efficiency and completeness as large as 97 per cent for $0 \\le z \\le 2.5$. However, again the accuracy drops significantly for $z > 2.5$.  Our work deals with the selection of high-$z$ QSOs from SDSS-FIRST matches. As stated before, our restriction to radio detected sources drastically reduces the contamination by stars, enabling us to obtain classification accuracies at these redshifts better than those obtained in more general studies aimed at the selection of the whole population of QSOs, regardless of radio detection. The paper is structured as follows. The sample of FIRST-DR5 matches is presented in Section 2. In Section 3 we explore the performance of supervised NNs to separate high-$z$ QSOs from the remaining spectral classes in the sample of 4250 sources with DR5 spectra, using multiband optical photometry and radio data. In Sections 4.1 and 4.2, we apply the trained NNs to the sample of 4415 sources without DR5 spectra, identifying 58 high-$z$ QSO candidates. In Section 4.3 we check the reliability of this identification via comparison with spectra from the NASA Extragalactic Database (NED), SDSS DR6 and follow-up spectroscopy with the WHT. The discussion and conclusions are presented in Section 5. ",
        "conclusions": "In this paper we aimed to obtain a sample of $z \\ge 3.6$ QSOs, starting from an initial complete sample of 8665 FIRST sources with star-like counterparts in the SDSS DR5 photometric survey, of which 4250 have spectra in DR5, 52 of them being $z \\ge 3.6$ QSOs. We found that simple supervised NNs, trained on the sources with DR5 spectra, and using optical photometry and radio data as input parameters, allow separation of high-$z$ QSOs from the remaining spectral classes with 96 per cent completeness and 62 per cent efficiency. The application of the NNs to the sample of 4415 sources without DR5 spectra yielded 58 high-$z$ QSO candidates.  We obtained spectra of 27 of the 58 candidates, and 17 were confirmed as high-$z$ QSOs. Spectra of 13 additional candidates from the NED and from DR6 revealed seven more high-$z$ QSOs, yielding a total 40 candidates with spectra available, of which 24 are high-$z$ QSOs. The number of high-$z$ QSOs was then increased from 52 in the initial sample to 76 (a factor 1.46).  The overall efficiency in the selection of new high-$z$ QSOs is $60 \\pm 12$ per cent (24/40). The estimate from a subsample unbiased with respect to the NN outputs is $60 \\pm 17$ per cent (12/20), and both values are in good agreement with the $62 \\pm 9$ per cent efficiency obtained for the DR5 labelled sample (Section 3).  None of the non-candidates with spectra available from NED or DR6 is a $z \\ge 3.6$ QSO, therefore we have no evidence of incompleteness regarding high-$z$ QSOs with matches in these catalogues. Since the NED spectroscopic identifications are assigned from a variety of surveys different than SDSS, the database provides a blind test of the good completeness of the classifier for DR5 unlabelled sources.  The efficiencies found are well above the values obtained for previous samples of RL high-$z$ QSOs, based on less accurate optical photometry and with fewer wavelength bands than SDSS, although already highly complete ($\\ge 95$ per cent) regarding optical colour selection. The efficiencies for various samples, summarised by Carballo et al. (2006) for $z \\ge 3.7$, are $12-13$ per cent (Holt et al. 2004, Carballo et al. 2006; $S_{\\rm 1.4 ~GHz}>$ 1 mJy, APM POSS $E$, $O$), 6 per cent (Hook et al. 2002; $S_{\\rm 5 ~GHz} >$ 30 mJy and radio flat, APM POSS $E$, $O$), and 19 per cent (Snellen et al. 2001; $S_{\\rm 5 ~GHz} >$ 30 mJy and radio flat, APM UKST $B$, $R$, $I$).  Adopting for the 18 candidates which still lack spectroscopy a weighted efficiency of 37 per cent (four candidates with $y \\ge 0.55$ and 14 with $y < 0.55$, with expected efficiencies of 91 and 22 per cent respectively), we calculate $\\sim$ 7 additional $z \\ge 3.6$ QSOs. The FIRST-DR5 sample of high-$z$ QSOs is thus expected to contain $\\sim 83$ QSOs (52+24+7). Adopting as a lower limit for completeness the nominal value of 96 per cent found for the labelled sample, we calculate for the set of 31 high-$z$ QSOs obtained by the classifier (24 discovered and 7 predicted) a minimum $ \\sim 1$ missed high-$z$ QSO.  The NNs found 31 contaminants,  i.e. non high-$z$ QSOs erroneously selected as high-$z$ QSOs, in the labelled sample, with a fraction $61 \\pm 14$ per cent (19/31) being QSOs with $3.2 \\le z \\le 3.6$. Among the 40 high-$z$ QSO candidates with available spectra we found 16 contaminants, seven of them being QSOs with $3.15 \\le z \\le 3.6$, confirming a high rate of QSOs with $z$ near the threshold $z=3.6$ among the contaminants, 7/16 = $44 \\pm 17$ per cent.  Our results allow us to obtain an estimate of the incompleteness of SDSS for the spectroscopic classification of FIRST high-$z$ QSOs. 47 of the high-$z$ QSO candidates are located in the spectroscopic area covered by DR6 (Table 5, Fig. 9 left panel), and 17 of them are $z \\ge 3.6$ QSOs, ten included in the DR6 spectroscopic catalogue and seven not included. 15 candidates in this area still lack spectroscopy, and assuming for them a weighted efficiency of 31 per cent (two candidates with $y \\ge 0.55$ and 13 with $y < 0.55$), we expect another four high-$z$ QSOs. From this calculation we estimate $11$ FIRST high-$z$ QSOs missed by SDSS (7 QSOs and $4$ candidates), which when compared to 62 (52+10) identifications yields an incompleteness of SDSS for the spectroscopic classification of FIRST $3.6 \\le z \\le 4.6$ QSOs of $\\sim$ 15 per cent (11/73) for $r \\le 20.2$.  The definition of the original sample of 52 high-$z$ FIRST QSOs excluded lobe-dominated morphologies and ``narrow-lined'' QSOs, and included QSOs with BALs or self-absorbed at Ly$\\alpha$. The same conditions hold for the larger sample of 76 QSOs, obtained from the application of the NNs trained using these 52 objects to DR5 photometric sources without spectra in DR5, and covering a slightly wider region of input space than that used by SDSS for QSO targetting.  In a future paper we plan to analyze the optical luminosity function of FIRST-SDSS QSOs at $3.6 \\le z \\le 4.6$ on the basis of this sample. Concurrently we expect to carry out spectroscopic observations of the 18 candidates without spectra. Given the efficacy of our approach, we intend to extend the sample using more updated SDSS data releases, increasing the number of training sources and the number of high-$z$ QSO candidates, for which subsequent spectroscopy will be planned. We envisage using additional infrared data via UKIDSS (UKIRT Infrared Deep Sky Survey)."
    },
    "0809/0809.0292_arXiv.txt": {
        "abstract": "{Although most young massive stars appear to be part of multiple systems, it is poorly understood how this multiplicity influences the formation of massive stars. The high-mass star-forming region \\mia{} is a prime example of a young massive cluster where the cluster center is resolved into multiple subsources at cm and infrared wavelengths, a potential proto-Trapezium system.} {We investigate the protostellar content in the \\rxa{} continuum down to subarcsecond scales and study the compact outflow components, also tracing the outflows back to their driving sources via the shocktracing SiO and \\so{} emission.} {The region \\mia{} was mapped with the PdBI at \\rxa{} and \\rxb{} in the AB configurations, tuning the receivers to observe the molecular transitions \\soa{}, \\sob{}, \\siob{}, and \\sioa{}.} {In the continuum we detect five sources, one of them for the first time, while counterparts were detected in the NIR, MIR or at radio wavelengths for the remaining four sources. Three of the detected sources are within the inner 2100\\,AU, where the protostellar number density exceeds $10^6$ protostars\\,pc$^{-3}$ assuming spherical symmetry. Lower limits for the circumstellar masses of the detected sources were calculated, ranging from $\\sim$0.3 to $\\sim$40\\,M\\solar{} although they were strongly affected by the spatial filtering of the interferometer, losing up to $\\sim$90$\\%$ of the single-dish flux. However, the projected separations of the sources ranging between $\\sim750$ and $\\sim4700$ AU indicate a multiple, Trapezium-like system. We disentangled the compact outflow component of \\mia{}, detecting five molecular outflows in SiO, two of them nearly in the line of sight direction, which  allowed us to see the collapsing protostars in the NIR through the cavities carved by the outflows. The \\so{} velocity structure indicates a rotating, bound system, and we find tentative signatures of converging flows as predicted by the gravoturbulent star formation and converging flow theories.} {The obtained data strongly indicate that the clustered environment has a major influence on the formation of high-mass stars; however, our data do not clearly allow us to distinguish whether the ongoing star-forming process follows a monolithic collapse or a competitive accretion mechanism.} ",
        "introduction": "The formation of high-mass stars appears to be intimately linked with the formation of multiple stellar systems. Optically visible OB-type stars show a much higher degree of multiplicity than their lower-mass counterparts \\citep{preibisch1999}. Now observations are showing that such high-mass multiple systems have separations ranging between 10,000\\,AU and 1\\,AU, and appear to be bound groups within larger clusters \\citep{mermilliod2001}. In contrast to hierarchical systems where the successive separation between its members increases by large factors, many of these systems are dynamically unstable, nonhierarchical clusters of three or more stars, called trapezia after the Trapezium cluster in Orion. The role such systems play in massive star formation is not yet understood, although several suggestions have been made.  First, the mass density of these multiple systems may be high enough to gravitationally trap hypercompact or even small ultracompact \\hii{} regions. \\citet{keto2002a,keto2002b,keto2003} has shown that continued accretion through the HC\\hii{} region may be possible, even after the supporting stars have reached the main sequence. The radius where these \\hii{} regions become hydrodynamically supported -thus stopping the accretion flow- scales inversely with the attracting mass, thus massive trapezia systems allow more mass to be accreted to the central accreting objects \\citep{keto2002a,keto2002b,keto2003}.  Second, binary systems are expected to form in the center of these dense proto-Trapezia. These binaries will continue accreting gas, funneled to the center by the combined potential of the protocluster. Because the angular momentum of the infalling gas has no correlation with that of the binary's, the masses of the individual stars increase while the separation of the binary decreases, giving as a natural outcome a high-mass close binary system that could eventually undergo mergers, creating even more massive systems \\citep{bonnell2005}.  The previous examples show how the formation of massive stars in multiple systems with separations on the order of $\\sim$1000\\,AU may be affected by their companions, either by modifying the morphology of accreting envelopes allowing more mass to be accreted% , or by forming more massive stars by merging high-mass close binaries.  Some multiple high-mass protostellar systems are known. For example, W\\,33A was imaged by \\citet{vandertak2005b} at 43\\,GHz with the Very Large Array finding three continuum sources at separations of $3000-5000$\\,AU; \\citet{hunter2006} found similar systems in NGC\\,6334\\,I and NGC\\,6334\\,I(N) with a few millimetric sources inside a region of $\\sim$10,000\\,AU. More recently, \\citet{beuther2007d} resolved the central 7,800\\,AU of the hot molecular core G29.96--0.02 into four submillimeter peaks, a potential proto-Trapezium system.  One prime example of a trapezia is the system \\mia{} in the W3-Main region of the \\hii{}/molecular cloud complex W3, a very active star-forming region part of the larger star-forming complex W3/W4 located in the Perseus arm of the Galaxy, that also encompasses the W4 star-forming region and the OB associations IC 1805 and IC 1795. In W3 the star formation takes place in the embedded regions W3-Main, W3-North, and W3-OH.  \\citet{oey2005} found evidence supporting the scenario in which the star formation in W3-Main was triggered by the formation of IC 1795, which in turn was triggered by the Perseus chimney/superbubble. This scenario, however, has been brought into question by the recent Chandra X-Ray observations of \\citet{feigelson2008}, because W3-Main does not show the elongated and patchy structure of a triggered star cluster. W3-Main itself may harbor its own triggered star-formation scenario because the various morphological classes of \\hii{} regions that can be found in it suggest that different stages of massive star formation (MSF) may be sequentially triggered by the pressure of the expanding \\hii{} regions \\citep{tieftrunk1997,tieftrunk1998}.  The system \\mia{} is a known double infrared (IR) source \\citep{howell1981,neugebauer1982} with a total luminosity of 2\\,$\\times$\\,10$^5\\,L_{\\odot{}}$ \\citep{campbell1995}, located in the W3-Main region at a distance of 2\\,kpc \\citep{megeath2008}. Recently, \\citet{megeath2005} presented HST observations with an angular resolution of 350\\,AU identifying seven near-IR sources including the two previously known. Three of these sources have both mid-IR counterparts and are 1.2 and 0.7\\,cm continuum sources \\citep{wilson2003,vandertak2005a}. The sizes of the \\hii{} regions are $<$240\\,AU, indicating that they are probably gravitationally bound and may be accreting \\citep[e.g.,][]{keto2003}. Based on its size, the probable masses of the stars, and its nonhierarchical distribution of sources, \\citet{megeath2005} proposed that this system is a proto-Trapezium system, which would emerge as a bound Trapezium similar to that in the Orion nebula. This region is also the source of at least two outflows \\citep{imai2000,wilson2003}, and is in the center of an embedded cluster of 80--240 low-mass stars \\citep{megeath1996} which in turn is surrounded by a core of several hundred low-mass stars \\citep{feigelson2008}. ",
        "conclusions": "In the continuum maps of \\mia{} obtained with the PdBI at \\rxa{} and \\rxb{} we detect 5 individual sources. Their calculated absolute masses are strongly affected by the spatial filtering inherent to the interferometric technique, however from the relative separations of the three strongest sources we propose the scenario of a Trapezium-like system, with at least 3 objects enclosed within a $\\sim$2000\\,AU radius, with a high (proto)stellar density of \\mbox{$\\sim5\\times10^6$\\,protostars\\,pc$^{-3}$}, and sharing a common envelope from which they may still be accreting gas.  Four of the five millimetric sources are identified as counterparts of sources previously detected in NIR, MIR and radio wavelengths by \\citet{megeath2005} and \\citet{vandertak2005a}. The remaining fifth source (labeled MM4) is a new detection.  Thanks to the spatial resolution of $\\sim$0.36$''$ we can for the first time trace the small-scale flow, detecting five molecular outflows, identifying the driving source of three of them. Although we find on short spatial scales the same alignment and axes as the outflows on large spatial scales previously detected in SiO and CO, within our accuracy and the big uncertainties introduced, we do not find strong evidence to support that we are tracing the inner and denser regions of those outflows, but we cannot rule out that possibility. However these small-scale outflows appear to be different features.  Two of the SiO outflows are close to the l.o.s. direction and each one is being driven by each of the stronger millimetric sources detected. This configuration would explain why these two objects are detected from NIR to radio wavelengths because cavities being carved out by outflows nearly along the l.o.s. have lower extinction, allowing the detection in the NIR and MIR of the hot dust near the collapsing protostar, while the long wavelength emission traces the dusty envelope that still surrounds it.  The \\so{} emission and velocity structure indicate that \\mia{} still has a gravitationally bound envelope. Also in \\so{} we detect a velocity jump of several km\\,s$^{-1}$ in a very narrow region adjacent to the continuum sources. This could be the signature of two converging molecular flows, compressing the gas and triggering the star formation, as predicted by the theory of gravoturbulent star formation."
    },
    "0809/0809.2404_arXiv.txt": {
        "abstract": "The validity of the hypothesis that the massive black holes in high redshift quasars grew from stellar-sized ``seeds'' is contingent on a seed's ability to double its mass every few ten million years.  This requires that the seed accrete at approximately the Eddington-limited rate.  In the specific case of radiatively efficient quasiradial accretion in a metal-poor protogalactic medium, for which the Bondi accretion rate is often prescribed in cosmological simulations of massive black hole formation, we examine the effects of the radiation emitted near the black hole's event horizon on the structure of the surrounding gas flow. We find that photoheating and radiation pressure from photoionization significantly reduce the steady-state accretion rate and potentially render the quasiradial accretion flow unsteady and inefficient. The time-averaged accretion rate is always a small fraction of the ``Bondi'' accretion rate calculated ignoring radiative feedback.  The pressure of Ly$\\alpha$ photons trapped near the \\ion{H}{2} region surrounding the black hole may further attenuate the inflow.  These results suggest that an alternative to quasiradial, radiatively efficient Bondi-like accretion should be sought to explain the rapid growth of quasar-progenitor seed black holes. ",
        "introduction": "\\label{sec:intro} \\setcounter{footnote}{0}  The origin and the early growth of massive black holes remains poorly understood.  The massive black holes in quasars ($M_{\\rm BH}\\sim 10^8-10^{10}M_\\odot$), active galactic nuclei (AGN; $M_{\\rm BH}>10^5M_\\odot$), and the quiescent nuclei of nearby galaxies may have started out as stellar-mass ($M_{\\rm BH}< 100\\ M_\\odot$) ``seed'' black holes \\citep[e.g.,][]{Madau:01,Menou:01,Islam:03}. Is this plausible, that is, could the seed black holes have grown rapidly enough in the cosmic time available to them \\citep[e.g.,][see also \\citealt{Haiman:04}, \\citealt{Djorgovski:08}, and references therein]{Haiman:01,Volonteri:03,Tanaka:08}? The rate at which a seed massive black hole can accrete is limited by the local density and the thermal structure of the protogalactic medium and by the effects of the radiation emitted near the event horizon on the accretion flow. Cosmological hydrodynamic simulations suggest that gravitational collapse produces dense central gas concentrations in protogalaxies \\citep[e.g.,][]{Springel:05,Li:07,Pelupessy:07,Wise:07a,Wise:07b,Wise:08,DiMatteo:08,Greif:08}. Atomic densities have been found to reach $n\\sim 10^4\\textrm{ cm}^{-3}$ \\citep[e.g.,][]{Greif:08}, and as much as $n\\sim 10^6\\textrm{ cm}^{-3}$ averaged over the central parsec around the potential minimum \\citep{Bromm:03,Wise:07b,Wise:08}.  On spatial scales that are resolved in the simulations, gas is sufficiently concentrated to enable rapid accretion onto a seed black hole. An exception are the first hundred million years after the seed was formed, during which the surrounding gas density is lowered by the radiative feedback from the black hole's progenitor star (see, e.g., \\citealt{Johnson:07,Alvarez:06,Alvarez:08}).  Given an ample gas supply, will rapid accretion be inhibited by radiative effects? A reassessment of an accreting black hole's ability to control its own gas supply is needed to improve the realism of the treatment of black hole accretion in cosmological simulations. Existing cosmological simulations modeling the growth of seed black holes do not resolve the spatial scales on which some of the radiative processes may alter the accretion. The simulations also do not resolve the fine structure of the dense, turbulent, and possibly multiphase protogalactic medium in which the black holes are embedded. Semianalytic prescriptions are normally adopted for the accretion rate, but these prescriptions normally do not take into account the radiative feedback; it is normally assumed that a black hole, residing in a pressure-supported primordial gas cloud, can accrete steadily at the Bondi rate subject to the Eddington limit \\citep[e.g.,][]{Volonteri:05,Alvarez:06,Johnson:07,Pelupessy:07,DiMatteo:08,Greif:08}. Here, we will evaluate the applicability of this assumption in view of the local radiative feedback that is present if the black hole accretes in a radiatively efficient fashion. We restrict our analysis to the early growth of quasar-progenitor ``seed'' black holes, which are occasionally referred to as ``miniquasars,'' where a stellar mass or an intermediate-mass black hole ($10^2M_\\odot\\lesssim M_{\\rm BH}\\lesssim 10^5M_\\odot$) accretes from a metal and dust poor environment.  The black hole's growth rate is particularly sensitive to the detailed thermal state of the irradiated accretion flow. This can be seen by noticing that the accretion rate, for quasiradial accretion, is influenced by the conditions at the sonic radius \\beq \\label{eq:rs} r_{\\rm s}\\sim 3\\times10^{14}~ \\frac{M_2}{T_{{\\rm s},5}} \\textrm{ cm} , \\eeq where $T_{\\rm s} = 10^5 T_{{\\rm s},5}\\textrm{ K}$ is the temperature of the photoionized and photoheated flow at the sonic radius and $M_{\\rm BH}=100\\ M_2M_\\odot$ is the black hole mass.  The sonic radius is normally unresolved in cosmological simulations of accretion onto black holes in the intermediate range of masses.\\footnote{The recent, highly-resolved simulation of primordial protostar formation by \\citet{Yoshida:08}, which has the requisite spatial resolution to resolve a sonic radius if it exists, only proceeds to the point where the initial hydrostatic core is formed, and does not treat the subsequent accretion flow onto the growing core.} It was recognized early that photoheating and photoionization pressure may prohibit steady radiatively efficient accretion \\citep[e.g.,][]{Shvartsman:71,Buff:74,Hatchett:76,Ostriker:76}. It was suggested that accretion can still proceed at rates disallowed by the steady state solutions if cycles of rapid gas inflow, during pauses in accretion near the event horizon, alternate with photoheating or photoionization pressure-driven outflows \\citep{Buff:74,Ostriker:76,Cowie:78,StellingwerfBuff:82,Begelman:85}. Such quasiperiodic cycling is seen in one-dimensional simulations of Compton-heated accretion onto $M_{\\rm BH}\\gtrsim 10^8M_\\odot$ black holes in galaxy clusters, where the black hole accretes from a hot, ionized, and pressure supported atmosphere \\citep{Ciotti:97,Ciotti:01,Ciotti:07,Sazonov:05}.   Recently, \\citet{Ricotti:08} revisited the problem of irradiated quasiradial accretion in the context of primordial black hole growth following the cosmic recombination \\citep[see, e.g.,][and references therein]{Ricotti:07}, and suggested that the accretion duty cycle is determined by the periodic formation of an \\ion{H}{2} region surrounding the black hole.   It was further recognized that the formal existence of steady-state, spherically symmetric accretion solutions is sensitive to the treatment of boundary conditions far from the sonic radius \\citep[e.g.,][]{Bisnovatyi:80}, and that these accretion flows can be locally thermally unstable \\citep[e.g.,][]{Stellingwerf:82,Krolik:83} and should break down into time-dependent two-phase structure, containing a warm ionized phase and a hot, coronal phase \\citep[e.g.,][]{Krolik:81}.  \\citet{Wang:06} claimed that Compton heating in the vicinity of a seed massive black hole reduces the radial accretion rate to a small fraction of the Eddington-limited rate; we here suggest, however, that thermal runaway may engender the Compton-heated coronal phase only in metal-rich flows, where photoionization heating of the incompletely stripped oxygen drives gas heating beyond $\\sim 10^5\\textrm{ K}$ \\citep[see, e.g.,][and \\S~\\ref{sec:photoionized_accretion} below]{Kallman:82}.   The impact of the radiation field produced near the event horizon on the accretion flow and on the state of the interstellar medium of the protogalaxy and that of the intergalactic medium \\citep[e.g.,][]{Dijkstra:04,Kuhlen:05,Zaroubi:07,Thomas:08,Ripamonti:08,Spaans:08} is sensitive to the shape of the spectral energy distribution (SED) of the central source.  The SEDs of rapidly accreting low-mass massive black holes ($M_{\\rm BH}\\sim 10^5-10^6M_\\odot$) exhibit significantly larger X-ray ($2\\textrm{ keV}$) to optical spectral ratios than the AGN containing more massive rapidly accreting black holes \\citep{Greene:07}, as is expected if a fraction of the radiation is produced in a geometrically thin disk. On the low mass end, if the microquasar SEDs \\citep[e.g.,][and references therein]{Remillard:06} are an adequate prototype, the seed black hole SEDs may contain energetically significant components extending into the hard X-rays. The luminosity-weighted average spectrum of AGN containing intermediate mass black holes could differ substantially \\citep[see, e.g.,][]{Venkatesan:01,Madau:04} from the scaled average quasar spectrum of \\citet{Sazonov:04}. The ability of X-rays to escape the protogalaxy affects their contribution to the soft X-ray background \\citep[e.g.,][]{Venkatesan:01,Dijkstra:04,Salvaterra:05} and the infrared background \\citep{Cooray:04}. Another difference between the first AGN and the starburst or post-starburst AGN is related to the differences in metallicity of the accretion flows.  The thermal phase structure of the interstellar medium exposed to UV and X-ray radiation is sensitive to metal abundances, especially for $T<10^4\\textrm{ K}$ \\citep[e.g.,][]{Donahue:91} and for $T\\gtrsim \\textrm{ few }\\times10^4\\textrm{ K}$ \\citep[e.g.,][and \\S~\\ref{sec:photoionized_accretion}]{Kallman:82}.  The role of radiative feedback in the formation of the first massive black holes resembles the radiative regulation in the formation of the first massive protostars \\citep{Omukai:01,Omukai:03a,Omukai:02,Omukai:03b,Tan:04,McKee:07}, though, of course, the central sources have very different spectra. In protostars, an \\ion{H}{2} region forms around the protostar and the protostellar disk that feeds its growth; persistent accretion onto the protostar may be quenched by radiation pressure. The protostellar accretion is characterized by a lower radiative efficiency and a higher accretion rate than the black hole accretion for the same accretor mass. The growing protostar is embedded in a supersonically collapsing and very dense envelope from inception, whereas here we assume that the seed black hole is born in the collapse of a massive star without such an envelope.  The presence of an infalling envelope implies that the \\ion{H}{2} region is initially \\emph{trapped} near the protostar, inside the radius where the envelope infall velocity turns from subsonic to supersonic.  For accretion from a diffuse medium onto a seed black hole, the \\ion{H}{2} region is inevitably extended and the flow crossing the ionization front is highly subsonic; this marks a crucial difference with the protostellar accretion scenarios.  Ignoring radiative effects, accretion onto the black hole is quasiradial on certain length scales (e.g., $r_{\\rm s}\\lesssim r\\lesssim 10\\textrm{ pc}$) if turbulence in the gas is weak \\citep{Krumholz:06} and if the gas is not rotationally supported on these scales.  If the accretion flow is shock-free (an unlikely condition) and possesses small net rotation, the buildup of vorticity near $r\\sim r_{\\rm s}$ may reduce the accretion rate by $\\sim 60\\%$ \\citep{Krumholz:05}.  Quasiradial accretion may be expected even when the baryons in a protogalaxy initially form a rotating disk, because self-gravity in the gas on scales of the protogalactic disk destabilizes rotational equilibria to convert disk-like configurations into quasiradial, stratified, pressure supported, and possibly turbulent configurations.  The gas distribution could still be rotationally supported on larger scales, where, e.g., dark matter dominates gravity, and on much smaller scales, where the black hole dominates gravity.  This description may apply to high-redshift protogalaxies \\citep[see, e.g.,][]{Oh:02,Volonteri:05,Wang:06}, and so here, we focus on angular momentum-free accretion and defer examining the role of angular momentum to a subsequent paper.  The quasiradial gas flow can either be steady and radial, or unsteady and characterized by alternating inflow, outflow, and nonradial motions. In view of these possibilities, this work is organized as follows. In \\S~\\ref{sec:bondi} we attempt, and fail, to construct a steady, radial solution for accretion at high accretion rates and high radiative efficiencies.   In \\S~\\ref{sec:timedep} we provide a qualitative analysis of time-dependent, episodic accretion, and attempt to estimate the average accretion rate. In \\S~\\ref{sec:clumpy} we discuss the consequences of the presence of cold, inhomogeneous, and turbulent gas in the vicinity of the black hole.  In \\S~\\ref{sec:conclusions} summarize our main conclusions.  In Table \\ref{tab:symbols} we present an overview of our notation.   \\begin{deluxetable*}{lll} \\tablecolumns{3} \\tablecaption{Index of Notation\\label{tab:symbols}} \\tablehead{\\colhead{Quantity} & \\colhead{Symbol} & \\colhead{Note}} \\startdata  Acceleration due to gravity & $a_{\\rm grav}$ & $-GM_{\\rm BH}/r^2$ \\\\ Acceleration due to radiation pressure & $a_{\\rm rad}$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ Spectral index of the SED & $\\alpha$ & $F_\\nu,\\ f_\\nu\\propto \\nu^{-\\alpha}$ \\\\ Case $B$ recombination rate for hydrogen & $\\alpha_B$ & $\\approx 2.6\\times10^{-13} T_4^{-1}\\textrm{ cm}^3\\textrm{ s}^{-1}$ \\\\ Collisional ionization rate for hydrogen & $\\alpha_{\\rm ion}$ & \\nodata \\\\   Abundance of species $i$ relative to hydrogen & $\\chi_i$ & $\\equiv n_i/n_{\\rm H}$ \\\\  Logarithmic slope of $T_{\\rm s}$ as a function of $\\epsilon$ & $\\delta$ & $T_{\\rm s}\\propto \\epsilon^\\delta$ \\\\   Radiative efficiency & $\\epsilon$ & $=L/\\dot M c^2$ \\\\ Critical efficiency for radiation pressure suppression at low efficiencies & $\\epsilon_{\\rm crit}$ & $\\phi r_{\\rm ion}/r_{\\rm s}=1$ \\\\ Ionization potential of species $i$ & $E_i$ & \\nodata \\\\  Duty cycle for episodic accretion & $f_{\\rm duty}$ & $\\equiv\\langle L^2\\rangle/\\langle L\\rangle^2$  \\\\ Density enhancement in the ionized gas in episodic accretion & $f_{\\rm epi}$ & $\\geq 1$ \\\\ Fraction of photon energy going to photoionizations in the ionized gas & $f_{\\rm ion}$ & $\\sim 1/3$ \\\\ Fraction of photon energy reprocessed to Ly$\\alpha$ & $f_{{\\rm Ly}\\alpha}$ & $\\sim 2/3$ \\\\ Fraction of $L$ that reaches the edge of the \\ion{H}{2} region & $f_{\\rm res}$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ Pressure enhancement due to turbulence in the neutral gas & $f_{\\rm turb}$ & $\\sim 1+{\\cal M}^2$ \\\\   Adiabatic index & $\\gamma$ & $=5/3$ \\\\  Luminosity in units of the Eddington luminosity & $\\ell$ & $L/L_{\\rm Edd}$ \\\\ Luminosity of isothermal accretion ignoring radiative heating & $\\ell_{\\rm Bondi}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\ Ly$\\alpha$-pressure-limited accretion rate & $\\ell_{{\\rm crit,Ly}\\alpha}$ & see text (\\S~\\ref{sec:prad_lyalpha}) \\\\ Peak luminosity in episodic accretion & $\\ell_{\\rm max}$ & \\nodata \\\\ Luminosity of steady-state photoheated accretion  & $\\ell_{\\rm s.s.}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\ Luminosity emitted at $r\\lesssim r_{\\rm disk}$ & $L$ & \\nodata \\\\ Eddington luminosity for Thomson scattering & $L_{\\rm Edd}$ & $=4\\pi GM_{\\rm BH} m_p c/\\sigma_{\\rm T}$ \\\\   Mean molecular mass of the ionized gas & $\\mu$ & $\\sim 0.6$ \\\\ Turbulent Mach number of the neutral gas & ${\\cal M}$ & \\nodata \\\\ Mass of the black hole & $M_{\\rm BH}$ & \\nodata \\\\ Central accretion rate & $\\dot M$ & \\nodata \\\\ Central accretion rate ignoring radiation & $\\dot M_{\\rm Bondi}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\  Density at the \\ion{H}{2} region's edge in a protogalactic density cusp & $n_{\\rm cusp}$ & see text (\\S~\\ref{sec:density}) \\\\ Density of the neutral gas and ambient density & $n_{\\rm HI},\\  n$ & \\nodata \\\\ Particle density within the \\ion{H}{2} region & $n_{\\rm HII}$ & $f_{\\rm turb}n_{\\rm HI}T_{\\rm HI}/T_{\\rm HII}$ \\\\ Maximum density for high ionization at the sonic radius & $n_{\\rm max,ion}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\ H$^-$-dissociating photon production rate & $\\dot N_{\\gamma,{\\rm H}^-}$ & see text (\\S~\\ref{sec:photodissociation}) \\\\ Lyman-Werner photon production rate & $\\dot N_{\\rm LW}$ & see text (\\S~\\ref{sec:photodissociation}) \\\\ Total ionization rate in the \\ion{H}{2} region & $\\dot N_{\\rm ion}$ & see text (\\S~\\ref{sec:hii_region}) \\\\ Total recombination rate in the \\ion{H}{2} region & $\\dot N_{\\rm rec}$ & see text (\\S~\\ref{sec:hii_region}) \\\\ Number of Ly$\\alpha$ reflections on \\ion{H}{2} region's walls to escape & $N_{\\rm reflect}$ & see text (\\S~\\ref{sec:prad_lyalpha}) \\\\   Dimensionless photoionization pressure acceleration in the \\ion{H}{2} region & $\\phi$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ Dimensionless radiation pressure acceleration at the \\ion{H}{2} region's edge & $\\psi$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ Pressure of the neutral gas & $P_{\\rm gas}$ & $=n_{\\rm HI} m_p k T_{\\rm HI}$ \\\\ Pressure of the Ly$\\alpha$ radiation & $P_{{\\rm Ly}\\alpha}$ & see text (\\S~\\ref{sec:prad_lyalpha}) \\\\  Radius of the \\ion{H}{2} region in a protogalactic density cusp & $r_{\\rm ion,cusp}$ & see text (\\S~\\ref{sec:density}) \\\\ Bondi radius ignoring radiative heating and acceleration & $r_{\\rm B}$ & see text (\\S~\\ref{sec:duty_cycle}) \\\\ Sonic radius & $r_{\\rm s}$ & $v(r_{\\rm s})=c_{\\rm s}(r_{\\rm s})$ \\\\ Disk radius & $r_{\\rm disk}$ & $\\ll r_{\\rm s}$ \\\\ Radius where thermal time equals inflow time & $r_{\\rm equi}$ & $\\max\\{t_{\\rm heat},t_{\\rm cool}\\}=r/v$ \\\\ Radius of the \\ion{H}{2} region & $r_{\\rm ion}$ & see text (\\S~\\ref{sec:hii_region}) \\\\ Width of the H$_2$ photodissociation shell & $\\Delta r_{\\rm diss}$ & see text (\\S~\\ref{sec:photodissociation}) \\\\ Width of the neutral shell surrounding the \\ion{H}{2} region & $\\Delta r_{\\rm shell}$ & see text (\\S~\\ref{sec:prad_lyalpha}) \\\\   Line-center Ly$\\alpha$ scattering cross section & $\\sigma_0$ & $\\approx 5.9\\times10^{-14}T_4^{-1/2}$ \\\\ Flux-averaged photoionization cross section of species $i$ & $\\bar \\sigma_i$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ Photon number-averaged photoionization cross section  & $\\tilde \\sigma_i$ & see text (\\S~\\ref{sec:prad_continuum}) \\\\ H$_2$ formation suppression factor due to H$^-$ photodissociation & $S$ & see text (\\S~\\ref{sec:photodissociation}) \\\\   Ly$\\alpha$ line-center optical depth of the neutral shell & $\\tau_0$ & see text (\\S~\\ref{sec:prad_lyalpha}) \\\\ Bremsstrahlung cooling time & $t_{\\rm Brems}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\ Compton heating time & $t_{\\rm C}$ & see text (\\S~\\ref{sec:rs_conditions}) \\\\ Cooling time & $t_{\\rm cool}$ & \\nodata \\\\ Heating time & $t_{\\rm heat}$ & \\nodata \\\\ Photoionization heating time & $t_{\\rm photo}$ & see text (\\S~\\ref{sec:inflow_outflow}) \\\\ Inflow time at the sonic radius & $t_{\\rm s}$ & $\\sim r_{\\rm s}/c_{\\rm s}(r_{\\rm s})$ \\\\ Salpeter mass-exponentiation time scale & $t_{\\rm Salp}$ & $=\\epsilon M_{\\rm BH}c^2/L_{\\rm Edd}$ \\\\ Temperature of the ambient neutral gas & $T_{\\rm HI}$ & \\nodata \\\\ Temperature within the \\ion{H}{2} region & $T_{\\rm HII}$ & \\nodata \\\\ Temperature at the sonic radius & $T_{\\rm s}$ & see text (\\S~\\ref{sec:intro}) \\\\  Density enhancement over isothermal accretion with $T_{\\rm HII}=T_{\\rm s}$ & $\\Upsilon$ & $\\gtrsim 1$\\\\  Radial inflow velocity & $v$ & \\nodata \\\\  Dimensionless inflow velocity & $w$ & $\\equiv v/c_{\\rm s}(r_{\\rm s})$ \\\\  Ionization parameter & $\\xi$ & $\\equiv L/r^2n$ \\\\ Dimensionless ionization parameter & $\\Xi$ & $\\equiv \\xi/4\\pi kT_{\\rm HII} c$ \\\\   Dimensionless radius & $y$ & $\\sim r/r_{\\rm s}$ \\\\   Metallicity & $Z$ & \\nodata \\\\  \\enddata \\end{deluxetable*} ",
        "conclusions": "  The local thermal and statistical equilibrium temperature of a photoionized gas is a strong function of the metallicity of the gas at radii from the black hole where gas becomes captured by the black hole. The photoionized gas outside the sonic radius is normally in equilibrium, but, particularly for metallicities $Z\\lesssim 0.1\\ Z_\\odot$, the gas passing through the sonic radius may not be in equilibrium.  For radiative efficiencies $\\epsilon\\lesssim 0.1$, the gas will  not complete thermal runaway that would heat it to the Compton temperature prior to crossing the sonic radius  Due to the radiative heating of the gas, the sonic radius is much smaller than the Bondi radius evaluated in terms of the temperature of the protogalactic medium far from the black hole; the corresponding gasdynamical accretion rate is thus reduced to a small fraction of the Bondi value.  The gas temperature at the sonic radius is also a strong function of the radiative efficiency of the accretion near the event horizon of the black hole.  The black hole is surrounded by an \\ion{H}{2} region, which is further surrounded by a photodissociation region if the ambient protogalactic medium contains molecules. The radius of the \\ion{H}{2} region in dense ($n\\gtrsim 10^5\\textrm{ cm}^{-3}$) clouds, that must be present to support near-Eddington accretion, is small ($r_{\\rm ion}\\sim 1\\textrm{ pc}$) compared to the size of the protogalaxy.  The sonic radius is smaller by orders of magnitude than the radius of the \\ion{H}{2} region.  This vast separation of scales renders self-consistent simulations of accretion onto seed black holes from a protogalactic environment particularly challenging, because to appropriately capture the radiative effects, the simulations must at least resolve the sonic radius.  Photoionization radiation pressure from the continuum radiation produced near the event horizon may further diminish the accretion rate. Compared to the gravitational acceleration, the acceleration due to radiation pressure is the strongest in the bulk of the \\ion{H}{2} region, i.e., not too close to the black hole where the equilibrium neutral fraction is negligible.  We have derived a simple criterion for accretion rate suppression by the photoionization radiation pressure.  A significant fraction of the luminosity of the black hole is converted into Ly$\\alpha$  resonance line radiation; resonance line scattering keeps this radiation trapped in the neighborhood of the \\ion{H}{2} region.  Radiation pressure from the trapped radiation can exceed the external, confining gas pressure.  We do not self-consistently solve for a stationary accretion flow in the presence of resonance line radiation pressure, but our estimates do suggests that the resonance line pressure could be detrimental to stationary accretion at high accretion rates.  As an alternative to strictly stationary accretion, we outline a model for episodic accretion in which short episodes of sudden, rapid accretion alternate with longer periods of accretion reduced by photoheating and radiation pressure.  In this model, the outflow and central rarefaction induced by photoheating and radiation pressure acceleration in the bulk of the \\ion{H}{2} region drives down the central accretion rate rapidly. The time-average accretion rate is then set by the peak accretion rate, divided by the logarithm of the number of $e$-foldings in the outflow-induced accretion rate decay.  Since the peak accretion rate need not be limited by the thermodynamic and kinematic constraints that apply to steady-state solutions, the episodic accretion rate may exceed the steady-state rate by a large factor.  Our idealized treatment here already indicates how extremely complex the self-consistent accretion problem is. To more fully explore the time-dependent accretion flow onto a seed black hole, numerical simulations are clearly needed. Of particular importance is to study the effect of angular momentum, the impact of turbulence, and the possible emergence of a multi-phase medium in the infalling gas. We will report on the results from such simulations and on a comparison with the analysis presented in this work elsewhere \\citep[][Couch et al., in preparation]{Milosavljevic:08}.  If our result that the initial accretion onto stellar seed black holes is greatly reduced in the presence of radiative feedback holds up even in fully three dimensional radiation-hydrodynamical simulations, the need for massive black hole formation by radiatively-inefficient accretion or by direct collapse of massive primordial gas clouds might be significantly increased \\citep[e.g.,][]{Bromm:03,Begelman:06,Begelman:08,Djorgovski:08}."
    },
    "0809/0809.2632_arXiv.txt": {
        "abstract": "{\\em We present measurements of ratios of elements of the scattering matrix of Martian analogue palagonite particles for scattering angles ranging from 3$^\\circ$ to 174$^\\circ$ and a wavelength of 632.8~nm. To facilitate the use of these measurements in radiative transfer calculations we have devised a method that enables us to obtain, from these measurements, a normalized synthetic scattering matrix covering the complete scattering angle range from 0$^\\circ$ to 180$^\\circ$. Our method is based on employing the coefficients of the expansions of scattering matrix elements into generalized spherical functions. The synthetic scattering matrix elements and/or the expansion coefficients obtained in this way, can be used to include multiple scattering by these irregularly shaped particles in (polarized) radiative transfer calculations, such as calculations of sunlight that is scattered in the dusty Martian atmosphere. } ",
        "introduction": "\\label{section_introduction}  Dust from the Martian surface is regularly swept up by winds to form local, regional, or sometimes even planet-wide dust storms. The airborne dust particles scatter and absorb solar radiation and are therefore very important for the thermal structure of the thin Martian atmosphere and for the temperature of the Martian surface \\citep[see e.g.][and references therein]{2006GeoRL..3302203B, 2005Icar..175...23G,2004Icar..167..148S,1972JAtS...29..400G}. The interaction between radiation and the dust particles thus has to be taken into account when studying, for example, the local and global climate on Mars. In particular, one has to account for the spatial distribution of the dust particles, their number density, and their optical properties. For a given wavelength, the optical properties of the dust particles depend on their composition, sizes and shapes.  Despite the important role of dust particles in the Martian atmosphere, surprisingly little is known about their optical properties \\citep[for an overview, see][]{2005AdSpR..35...21K}. Consequently, in radiative transfer calculations that are used to interpret space or ground-based observations of Mars, various assumptions are made regarding the dust optical properties. In particular, it is common to assume homogeneous, spherical or spheroidal dust particles \\citep{1995JGR...100.5235P,1999JGR...104.8987T}, although dust particles on Earth are known to be irregularly shaped. The optical properties of homogeneous, spherical particles can straightforwardly be calculated using Lorenz-Mie theory \\citep[][]{1957lssp.book.....V,1984A&A...131..237D}, and those of homogeneous, spheroidal particles using e.g. the so-called T-matrix method \\citep[][]{1994OptCo.109...16M,2006JGRD..11111208D}. The optical properties of spherical particles can, however, differ significantly from those of irregularly shaped particles, even if their composition and/or size distribution is the same. Therefore, assuming spherical instead of irregularly shaped particles in radiative transfer calculations that are used for example to analyze observations, can lead to significant errors in retrieved atmospheric parameters, such as the dust optical thickness and/or dust particle size distributions \\citep[for a discussion on such errors, see e.g.][]{2003SoSyR..37...87D,2002SoSyR..36..367D}.  For irregularly shaped particles, which are very common in nature, the scattering matrix elements can in principle be calculated with numerical methods such as those based on the so--called Discrete Dipole Approximation (DDA) method \\citep[see e.g.][]{1994OSAJ...11.1491D}. However, the vast amounts of computing time that such numerical calculations require make it at least very impractical to calculate complete scattering matrices for a sample of particles with various (irregular) shapes and various sizes, in particular if the particles are large compared to the wavelength of the scattered light. Alternatively, employing geometrical optics for unrealistically spiky particles combined with an imaginary part of the refractive index that is rather small compared to the typical values used in the literature, appears to be useful to reproduce the scattering behavior of irregularly shaped mineral particles \\citep[see][]{2003JGRD.108aAAC12N}. As suggested by e.g. \\citet{1999JGR...104.8987T,2003SoSyR..37...87D,2003JGRE.108i....1W}, a more practical method to obtain elements of the scattering matrix for an ensemble of irregularly shaped particles is to measure the elements in a laboratory. Note that measured scattering matrix elements are also essential for validating numerical methods and approximations.  In this article, we present measurements of ratios of elements of the scattering matrix of irregularly shaped, randomly oriented Martian analogue palagonite particles, described by \\citet{1997JGR...10213341B}, as functions of the scattering angle. The material palagonite is believed to be a reasonable, but not perfect, analogue for the Martian surface and atmospheric dust particles. Terrestrial palagonite particles (i.e. terrestrial weathering products of basaltic ash or glass) have been put forward as Martian dust analogues \\citep[][]{1981LPI....12..271E,1995JGR...100.5309R} because of spectral similarities observed with visible and near-infrared spectroscopic observations using both Earth-based telescopes and several Mars orbiting spacecraft \\citep[][]{1985AdSpR...5Q..59S,1992mars.book..557S}.  The measurements have been performed using a HeNe laser, which has a wavelength of 632.8~nm, and for scattering angles, $\\Theta$,  ranging from 3$^\\circ$ (near-forward scattering) to 174$^\\circ$ (near-backward scattering). Other examples of such measurements for irregularly shaped mineral particles obtained with the same experimental set-up have been reported by e.g. \\citet{2005JQSRT..90..191V} and references therein. Since we have measured ratios of all (non-zero) elements of the scattering matrix as functions of the scattering angle, our results can be used for radiative transfer calculations that include multiple scattering and polarization. As described by e.g. \\citet{1998GeoRL..25..135L}, ignoring polarization, i.e. using only scattering matrix element (1,1) (the so-called phase function), in multiple scattering calculations, induces errors in calculated fluxes. The use of only the phase function should be limited to single scattering calculations for unpolarized incident light.  A practical limitation of our experimental method is that we cannot measure close ($< 6^\\circ$) to the exact backscattering direction ($\\Theta=180^\\circ$), because there our detector would interfere with the incoming beam of light, nor close ($< 3^\\circ$) to the exact forward scattering direction ($\\Theta=0^\\circ$), because there our detector would intercept the unscattered part of the incident beam. These two scattering directions are, however, important for radiative transfer applications. This holds in particular for the near-forward scattering direction, since a significant fraction of the light that is incident on a particle is generally scattered in the near-forward direction. A solution to the lack of measurements in the near-forward and near-backward scattering directions is to add artificial data points. For the near-forward scattering direction, where a strong peak in the phase function is expected, one can add artificial data points calculated using e.g. Lorenz-Mie calculations. This has been done before, e.g. by \\citet{2007JQSRT.103...27K,2003JQSRT..79..903L,2004GeoRL..3104113V,2005JGRD..11010S02H}. In this article, we use a method similar to that of \\cite{2003JQSRT..79..903L}, that was also used by e.g. \\cite{2007JGRD..11213215M}, but we extend this by using expansion coefficients which result from a Singular Value Decomposition (SVD) fit to the measurements with generalized spherical functions \\citep[][]{1963Gelfand,1983A&A...128....1H,1984A&A...131..237D, 2004Hovenier}. With the added artificial data points and the expansion coefficients, we construct a so--called synthetic scattering matrix, which is normalized so that the average of the synthetic phase function over all directions equals unity and covers the whole scattering angle range (i.e. from 0$^\\circ$ to 180$^\\circ$).  Tables of the measurements, the synthetic scattering matrix elements and the expansion coefficients will be available from the Amsterdam Light Scattering Database\\footnote{The Amsterdam Light Scattering Database is located at: http://www.astro.uva.nl/scatter}. See \\citet{2005JQSRT..90..191V,2006JQSRT.100..437V} for a description of this database.  The structure of this article is as follows. In Sect.~\\ref{section_palagonite}, we describe the microphysical properties of our  Martian analogue palagonite dust particles. In Sect.~\\ref{section_scatteringmatrix}, we define the scattering matrix, describe the experimental set-up, and present the measurements and an auxiliary scattering matrix. In Sect.~\\ref{section_expansioncoefs}, we introduce the expansion coefficients, describe the Singular Value Decomposition fitting method, and present the derived expansion coefficients and synthetic scattering matrix. In Sect.~\\ref{section_summary}, finally, we summarize and discuss our results. ",
        "conclusions": "\\label{section_summary}  We present measured ratios of elements of the scattering matrix of irregularly shaped Martian analogue palagonite particles \\citep[][]{1995JGR...100.5309R,1997JGR...10213341B} as functions of the scattering angle $\\Theta$ ($3^\\circ \\leq \\Theta \\leq 174^\\circ$) at a wavelength of 632.8~nm. Our measured ratios of scattering matrix elements differ strongly from those calculated for homogeneous, spherical particles with the same size and refractive index. In particular, the measured phase function (ratio $F_{11}(\\Theta)/F_{11}(30^\\circ)$) shows a very strong (almost three orders of magnitude) forward scattering peak and a smooth drop-off towards the largest scattering angles, where the phase function of the spherical particles shows much more angular features especially in the backward scattering direction. Clearly, using scattering matrix elements that have been calculated for homogeneous, spherical particles when irregularly shaped particles are to be expected in radiative transfer calculations for e.g. the interpretation of remote-sensing observations can thus lead to errors in retrieved dust properties (microphysical parameters and/or optical thicknesses).  To facilitate the use of our measurements in radiative transfer calculations for e.g. Mars, we have first constructed an auxiliary scattering matrix from the measured scattering matrix. This auxiliary scattering matrix covers the whole scattering angle range (i.e. from 0$^\\circ$ to 180$^\\circ$), and its elements have been normalized such that the average of the phase function over all scattering directions equals unity. The value of the phase function at $\\Theta=0^\\circ$ has been computed from the phase function calculated with Mie-theory for homogeneous, spherical particles with the same size and composition. The normalization of the auxiliary phase function, and hence the normalization of the other elements too, is obtained by applying a Singular Value Decomposition (SVD) method to fit an expansion in generalized spherical functions \\citep[][]{1963Gelfand,2004Hovenier} to the measured phase function, including artificial data points. The first expansion coefficient yields the required normalization constant. After the normalization of the auxiliary phase function, the SVD method is applied to the auxiliary scattering matrix and its expansion coefficients are obtained.  With the expansion coefficients, a synthetic scattering matrix is computed for the complete scattering range. It is normalized so that the average of its one-one element over all directions equals unity. The synthetic scattering matrix elements can also straightforwardly be used in radiative transfer calculations. The need to include all scattering matrix elements, instead of only the phase function, is obvious for the interpretation of polarization observations. However, even for flux calculations, all scattering matrix elements should be used, because ignoring polarization, i.e. using only the phase function, in multiple scattering calculations induces errors in calculated fluxes \\citep[e.g.][]{1998GeoRL..25..135L}. The use of only the phase function should be limited to single scattering calculations for unpolarized incident light.  Figure~\\ref{fig_map_vs_tomasko} shows a comparison between our synthetic phase function and two phase functions presented by \\citet{1999JGR...104.8987T} as derived from diffuse skylight observations of the Imager for Mars Pathfinder. Each of the phase functions has its own normalization. The phase functions of \\citet{1999JGR...104.8987T} show the same general angular behavior as our synthetic phase function: a strong forward scattering peak and a smooth drop-off towards the largest scattering angles. The forward scattering peak of our phase function appears to be stronger than the peaks of the phase functions of \\citet{1999JGR...104.8987T}. This can easily be due to the size difference of the particles. The dust particles in our sample have an effective radius that is a factor of 2 to 3 larger than that of \\citet{1999JGR...104.8987T} (4.46~$\\mu$m versus 1.6~$\\mu$m $\\pm$ 0.15~$\\mu$m). The slopes of the phase functions in the backward scattering direction are very similar, and appear to be typical for (terrestrial) irregularly shaped mineral particles with moderate refractive indices \\citep[see e.g.][]{2001JGR...10617375V,2000A&A...360..777M,2001JGR...10622833M}. This smooth slope can not be easily anticipated from Mie-theory.  The expansion coefficients, the synthetic scattering matrix elements, and the measured ratios of elements of the scattering matrix of the Martian analogue palagonite particles will all be available from the Amsterdam Light Scattering Database \\citep[for details, see][]{2005JQSRT..90..191V,2006JQSRT.100..437V}. The Amsterdam Light Scattering Database contains a collection of measured scattering matrix element ratios, including information on particle sizes and their composition, for various types of irregularly shaped particles. The SVD method presented in this article can straightforwardly be applied to measured scattering matrix elements of particles other than the Martian analogue palagonite particles.  \\ack We are grateful to Ben Veihelmann for helping with the SEM image of Fig.~\\ref{fig_sem}. It is a pleasure to thank Martin Konert of the Vrije Universiteit in Amsterdam for performing the size distribution measurements and Michiel Min of the University of Amsterdam for fruitful discussions.  \\label{lastpage}"
    },
    "0809/0809.3825_arXiv.txt": {
        "abstract": "We discuss the cosmological consequences of QCD phase transition(s) on the early universe. We argue that our recent knowledge about the transport properties of quark-gluon plasma (QGP) should thraw additional lights on the actual time evolution of our universe. Understanding the nature of QCD phase transition(s), which can be studied in lattice gauge theory and verified in heavy ion experiments, provides an explanation for cosmological phenomenon stem from early universe. ",
        "introduction": "introduction} Our knowledge about the {\\it early} and {\\it present} universe is based on two successful models; the standard model for elementary particles and the one for cosmology. The main ideas about the early universe cosmology seem to be confirmed from various cosmological data, like the temperature anisotropy in cosmic microwave background radiation (CMB), the light element abundances, etc. Although we have a solid description for phase transitions in almost all epochs in the early universe, their confirmations still indirect processes. The observation of cosmological 'exotic' phenomenon and the fruitful results from heavy ion experiments are promising tools to confirm these theoretical predictions.  In this work, we limit our the discussion to QCD epoch, i.e. to energy scale of $\\Delta_{QCD}\\approx 200~$MeV, $t\\approx 10^{-5}-10^{-6}~$second after the begin and when the universe has the {\\it Hubble} radius of about $d_H\\approx 10~$Km corresponding to scales of about $1~$pc~\\footnote[1]{1 parsec $=3.08568025 \\times 10^{16}~$m, the distance from the Sun which would result in a parallax of 1 arcsecond as seen from the Earth.} or about $3~$ light years today. At high energy the coupling between quarks and gluons likely becomes weaker. This coupling diverges at $\\Delta_{QCD}=200~$MeV. From this property, QCD is known as asymptotically free theory.  As the energy decreases, i.e. the universe expands and cools down, the quarks and gluons underwent to confined state, to hadrons (deconfinement-confinement phase transition) and the chiral symmetry almost simultaneously breaks down, i.e., orientation of right- and left-handed quarks. The latter likely leads to production of Goldstone bosons. The three pions are the lightest Goldstone bosons. Dynamical fluctuations on these particles and/or disoriented chiral condensates (DCC) would manifest this transition.  In this talk, we discuss the cosmological consequences of nature and order of the QCD phase transition(s) on early universe. ",
        "conclusions": "Lattice QCD calculations with dynamical quarks and physical masses show that the phase transition at very small net baryon density, which is related to early universe is likely cross over or very weak second order. Phenomenological indications for such a continuous transition is quite likely an ambiguous task. As discussed earlier, the order and critical temperature of deconfinement-confinement phase transition(s) calculated in lattice QCD strongly depend on the number and masses of quark flavors. The order varies from strong first order to very weak second order or cross over. Consequently, the critical temperature takes values between $150~$MeV and $200~$MeV. The cosmological consequences strongly depend on the QCD phase transition(s).  We used to assume that QGP can be treated as a free ideal gas. The recent RHIC results suggested that QGP is a fluid rather than an ideal gas. Also recent lattice QCD simulations support such a conclusion. Taking into consideration that quark-gluon plasma is strongly correlated (sQGP) and has shear viscosity would lead to the consequence that its hydrodynamic evolution is slower relative to QGP with zero viscosity. Also the transverse expansion in such a fluid is much stronger than the one in an ideal QGP fluid. Phenomenologically, the particle production is considerably enhanced in sQGP. All these results would to some extent modify our ideas about the time evolution in early universe, for example."
    },
    "0809/0809.1682_arXiv.txt": {
        "abstract": "We present stellar proper motions in the Galactic bulge from the Sagittarius Window Eclipsing Extrasolar Search (SWEEPS) project using ACS/WFC on HST. Proper motions are extracted for more than 180,000 objects, with $>$81,000 measured to accuracy better than 0.3 mas yr$^{-1}$~in both coordinates. We report several results based on these measurements: 1. Kinematic separation of bulge from disk allows a sample of $>$15,000 bulge objects to be extracted based on $\\ge 6\\sigma$ detections of proper motion, with $<0.2\\%$~contamination from the disk. This includes the first detection of a candidate bulge Blue Straggler population. 2. Armed with a photometric distance modulus on a star by star basis, and using the large number of stars with high-quality proper motion measurements to overcome intrinsic scatter, we dissect the kinematic properties of the bulge as a function of distance along the line of sight. This allows us to extract the stellar circular speed curve from proper motions alone, which we compare with the circular speed curve obtained from radial velocities. 3. We trace the variation of the $\\{l,b\\}$~velocity ellipse as a function of depth. 4. Finally, we use the density-weighted $\\{l,b\\}$~proper motion ellipse produced from the tracer stars to assess the kinematic membership of the sixteen transiting planet candidates discovered in the Sagittarius Window; the kinematic distribution of the planet candidates is consistent with that of the disk and bulge stellar populations. ",
        "introduction": "The formation and evolution of merger-built bulges and secularly-grown pseudobulges and bars is crucial to the evolution of spiral galaxies, and indeed their formation history is used to test models for the formation of structure in the Universe such as the $\\Lambda$CDM framework; (see Kormendy \\& Kennicutt 2004 for a review). The inner region of our own Milky Way shows evidence for a bulge (e.g. Blitz \\& Spergel 1991) and at least one Bar-like structure (e.g. Benjamin et al. 2005; L{\\'o}pez-Corredoira et~al. 2007). Furthermore, for our own bulge the stars are close enough that detailed stellar data may be obtained on a star-by-star basis, such as radial velocities and proper motions, which are simply not yet available for most external galaxies.  These studies of our own bulge show it to be a distinct stellar population from the disk of the Milky Way with a wide range of abundances (Rich 1988)\\footnote{We refer to the thin disk and the thick disk together simply as the ``disk,'' as both have similar kinematics (Gilmore, Wyse \\& Kuijken 1989).}. High $[\\alpha/Fe]$~compared to the disk suggests rapid enrichment; Type-II SNe produce the $\\alpha$-elements and result from short-lived, high-mass stars, while iron enrichment requires the rather slower buildup of a significant number of Type I SNe (McWilliam \\& Rich 1994; Zoccali et~al. 2006; Lecureur et~al. 2007; Fulbright et al. 2007). A recent spectroscopic comparison with halo objects shows the majority of bulge stars to be $\\alpha$-enhanced compared to the halo, thus the majority bulge stellar population cannot have formed from the halo (Fulbright et al. 2006; Fulbright et~al. 2007).  The detailed balance of populations of the bulge and its current kinematics remain rather poorly constrained, and thus the picture of its step-by-step formation and evolution is far from complete. Strong, variable extinction caused by gas and dust in the intervening spiral arms (e.g. Sumi 2004), contamination of the test population by stars in the foreground disk, and the significant spatial depth of the bulge along the line of sight, combine to make the separation of a pure-bulge samply for further study a challenging task. While one can identify a candidate main sequence, the main sequence turn-off (hereafter MSTO) usually used as an age/metallicity diagnostic is broadened by the age, metallicity and depth range of the bulge and confused by the foreground disk population. Force-fitting the Horizontal Branch in the color-magnitude diagram (hereafter CMD) and comparison with globular cluster sequences suggests a majority population $\\sim$~10 Gyr in age, but with as much as 30\\% of the stars belonging to a young population (Holtzman et al. 1993; Ortolani et al. 1995). Furthermore, a population of OH masers (Sevenster et al. 1997) and a few hundred mass-losing AGB stars (van Loon et al. 2003) {\\it has} been detected in the inner Milky Way. These objects are claimed to represent an intermediate-age ($\\sim 1$-few Gyr) population within the inner Milky Way. These objects populate the circumnuclear molecular zone ($|l| < 1.5\\degr$, $|b|<0.5\\degr$) and are thought to be tracers of a larger population that is difficult to separate from the bulk stellar population in the color-magnitude diagram (Lindqvist et al. 1992; Messineo et al. 2002; Habing et al. 2006). In optical CMDs, young-intermediate age populations overlap significantly with the main sequence of a young foreground disk population; separation of the bulge from the disk population is therefore critical if the detailed population distribution of the bulge is to be obtained.  Feltzing \\& Gilmore (2000) compare number counts along the CMD between populous clusters at a range of galactic latitudes to estimate the foreground disk contamination, producing a prediction of disk contamination in two bulge fields and thus a differential estimate of the bulge population itself. They find a bulge population almost exclusively older than $\\sim 10$~Gyr. This conclusion was reinforced by a later study employing statistical subtraction to take off a scaled contamination from a comparison field in the foreground disk; globular cluster comparison suggested an old, metal-rich population to fit the bulge well, but without a precise age distribution (Zoccali et al. 2003). While a metallicity gradient with height $z$~above the galactic midplane probably does exist (Zoccali et al. 2003), a {\\it radial}~metallicity gradient has yet to be conclusively demonstrated (Rich \\& Origlia 2005; Minniti \\& Zoccali 2008).  No evidence has yet been found for Blue Stragglers in the bulge. Blue Stragglers are hydrogen-burning stars with apparent age younger than that of the parent population, most likely the result of significant mass deposition onto the star (e.g. Stryker 1993; Bailyn 1995). These objects are found in clusters of all ages (e.g. Stryker 1993) and would thus be expected to be present in the bulge. However they occupy a region of the CMD that overlaps with the foreground disk so are difficult to discriminate.  Stellar kinematics offer the cleanest method to separate bulge from disk. The disk shows apparent streaming motion in front of the bulge due to the motion of the Sun about the galactic center, while the proper motion dispersion of the two populations should differ because the two populations have widely different ages and likely differing relaxation times (Binney \\& Tremaine 1994). While statistical subtraction tells us that some fraction of objects in a region of the CMD may be bulge objects, kinematic constraints allow us to assign likely bulge/disk membership to {\\it individual objects}.     Beyond bulge-disk separation, stellar proper motions are valuable in their own right as a probe of the present-day kinematics of the bulge. Proper motions have suggested that the bulge stars show net rotation in the same sense as the rest of the disk, though a smaller retrograde population may also be required (Spaenhauer et al. 1992; Zhao et al. 1994). Most previous studies converge on the model that the bulge apparently rotates as a solid body, with circular velocity \\vmp rising linearly with distance from the galactic center until the stellar population becomes dominated by disk stars (which show roughly constant $\\langle v_\\phi\\rangle$); broadly similar behaviour is reproduced by bulge models including a rapidly rotating bar (e.g. Sellwood 1981; Zhao 1996). This consensus has recently been called into question by the radial velocity study of Rich et al. (2007), which (with velocities consistent with the PNe results of Beaulieu et al. 2000) suggests an inflection point in the \\vmp curve roughly $\\sim$0.6 kpc from the galactic center, outside which \\vmp is constant at $\\sim 50$km s$^{-1}$. This suggests that solid body-like rotation may only persist over the inner regions of the bulge/bar system. A change in the general character of stellar orbits within the inner region of the bulge may relate to an inner resonance (e.g. Sevenster 1999). Clearly, an independent measure of the stellar \\vmp curve is required. This is best constructed from datasets for which systematics at different galactocentric radii are similar, in particular the correction due to solar motion about the galactic center; as we will demonstrate, proper motions along a single sight-line provide just such a dataset provided the observations are at sufficient depth to assemble large numbers of objects in each distance bin to overcome intrinsic variation.  The bulge has now been the subject of a number of proper motion studies. Early ground-based studies of 427 $K$~and $M$-giants\\footnote{The Spaenhauer et al. (1992) figure of 429 stars includes two repeated entries.} in Baade's Window ($l,b=1.02\\degr,-3.93\\degr$), using photographic plates separated by nearly 33 years, showed a proper motion distribution with small anisotropy, $\\sigma_l/\\sigma_b=1.15\\pm0.06$~(Spaenhauer et al. 1992). On the bulge minor axis, plate-scans over an area $25' \\times 25'$~in Plaut's Window ($l,b=0\\degr,-8\\degr$) across a 21-year time interval were used to extract proper motions from 5088 objects (at $14 < V < 18$), finding similar dispersion and anisotropy to the Baade's Window sample (Mendez et al. 1996).  From the ground, five years of MACHO photometry allowed stars with proper motions $\\ga 18$~mas yr$^{-1}$~to be isolated, to search for future microlensing events towards the bulge (50$\\times0.5''\\times 0.5''$~fields at $-1\\degr < l < -10\\degr$ and $-11\\degr < b < -1.5\\degr$)~and Magellanic Clouds (Alcock et al. 2001). The OGLE-II experiment yielded proper motions of some $5\\times 10^6$~stars from $49\\times 0.24\\degr \\times 0.95\\degr$~fields within ($-11\\degr < l < +11\\degr$) and ($-6\\degr < b < +3\\degr$), with proper motions measured to a precision of 0.8-3.5 mas yr$^{-1}$~(Eyer \\& Wo\\'zniak 2001; Sumi et al. 2004), which was recently used as the basis for a study of the trends in proper motion with location in the bulge (Rattenbury et al. 2007a,b). Most recently, Plaut's Window has been the subject of a second plate-based study of some 21,000 stars within a $25'\\times25'$~region, producing proper motions to $\\sim 1$~mas yr$^{-1}$~and using $\\sim8700$~cross-matches with 2MASS to carefully select bulge giants (Vieira et al. 2007).  Turning to space-based proper motion studies, a number of studies of clusters that use proper motions to discriminate the cluster from the field, produce bulge proper motions if the field includes a significant number of bulge stars. Zoccali et al. (2001) reported the the velocity dispersion of about $10^4$~bulge stars at ($l,b=5.25\\degr,-3.02\\degr$), as a by-product of their WFPC2 study of the cluster NGC 6553, finding proper motion dispersion consistent with previous ground-based studies of Baade's Window. Similar results were obtained from the bulge stars in the field of a WFPC2 study of the metal-rich cluster NGC 6528, at ($l,b=1.14\\degr,-4.12\\degr$; Feltzing \\& Johnson 2002). However, Kuijken \\& Rich (2002) were the first to dedicate HST observations to bulge proper motion studies, returning to Baade's Window ($l,b=1.13\\degr,-3.77\\degr$) and the Sagittarius low-extinction window ($l,b = 1.25\\degr,-2.65\\degr$) with WFPC2. This isolated 3252 bulge stars in Baade's Window and 3867 in the Sagittarius Window, concluding that the bulge can be best fit with a 10 Gyr-old open cluster-type sequence. No detection of $\\{l,b\\}$ covariance was reported, nor was any detection of blue straggler candidates. An ACS/WFC and WFPC2 program is ongoing to reimage fields for which observations are already present in the HST archive (Kuijken 2004; see Soto et al. 2007 for prelimiary results). Recently, ACS/HRC was used to survey $35\\times 35''\\times 35''$~fields near Baade's Window across a $5\\degr\\times2.5\\degr$~region in the vicinity of Baade's Window, in order to compare the bulk kinematic properties of bulge fields as a function of location relative to the galactic center (Koz{\\l}owski et al. 2006). As pointed out by Vieira et al. (2007), the results of the ACS/HRC study and the literature differ quite strongly from those of Rattenbury et al. (2007b; particularly the dispersions $\\sigma_l$~and $\\sigma_b$). This appears to be due to differences in the selection of tracers used; the sample of Rattenbury et al. (2007b) may be contaminated by evolved disk objects (Vieira et al. 2007; see also Rattenbury et al. 2007c). Clearly, some care is required in the selection of bulge tracer objects, particularly above the MSTO.  We report on the use of ACS/WFC to extract precise proper motions from a large number of stars in the Sagittarius Window towards the bulge. The Sagittarius Window Eclipsing Extrasolar Planet Search (SWEEPS project; Sahu et al. 2006) obtained an extremely well-sampled photometric dataset with ACS/WFC of the Sagittarius Window towards the bulge; this forms our first epoch. A repeat visit just over two years later forms the second epoch, from which we extract proper motions. At ($l,b$)=(1.25$\\degr$,-2.65$\\degr$), the line of sight passes within $\\sim 300$~pc of the galactic center. This is far enough from the center that the claimed population of young-intermediate age objects traced by the OH masers and AGB stars, is likely to be an insignificant contributor to the observed field; such objects are likely confined to within $\\sim100$~pc of the galactic mid-plane and are found preferentially near the galactic center (Sevenster et al. 1999; Frogel et al. 1999). Thus the stellar population of our field of view consists of bulge, disk and halo objects.  Our first epoch has a total integration time $>$86ks each in $F814W$~and $F606W$~and is the deepest ever observation of the Galactic bulge. With $6\\sigma$~detections of proper motions towards the wings of the proper motion distributions where disk/bulge separation is greatest, we push the disk contamination below the 0.3\\% level and extract the purest, largest sample of bulge stars yet assembled.  This report is organized as follows: we provide the particulars of the observations used in Section \\ref{s_obs}. We report the procedures used to produce precise position measurements and the resulting proper motions in Section \\ref{s_red}. Section \\ref{s_pm} outlines the broad features of the proper motion distribution of the population as a whole. To correct for disk contamination in the sample of kinematically selected bulge objects, we must first estimate the fraction of disk and bulge objects in the observed sample; to do so requires tracing the proper motion distribution as a function of distance along the line of sight. We do so in Section \\ref{s_dist}; before returning to the bulge/disk population distinction we use the distance dependence of the proper motions to extract some kinematic properties of the bulge in Section \\ref{s_vdist}. We select a likely bulge population in Section \\ref{s_popn} and briefly examine its properties. Finally, in Section \\ref{s_cand} we assess the likely kinematic membership of the sixteen SWEEPS planet candidates from Sahu et al. (2006). ",
        "conclusions": "\\label{s_conclude}  We have measured proper motions for $> 180,000$~objects within the Sagittarius low-reddening window towards the bulge ($l,b=1.65\\degr,-2.65\\degr$), and used them to extract a clean clean-bulge sample of 15,323 objects - a sample roughly a factor four larger than that afforded by WFPC2 across a six-year interval (Kuijken \\& Rich 2002). This clean-bulge sample contains perhaps 31 contaminants from both disk and halo, making it the purest bulge population ever isolated. Constructing a median stellar sequence from this bulge sample, we find that an 11 Gyr isochrone best represents the bulge population, with most of the variation along the bulge subgiant branch falling within the range $[Fe/H]=0.0 \\pm 0.4$~and age $11 \\pm 3$~Gyr. Use of this sample to inform bulge age studies, in conjunction with extensive completeness tests, will be the subject of future reports. Work along these lines is particularly exciting when we consider the parallel NICMOS observations we have undertaken, for which similar selection should be possible and which brings the possibility of tracing the bulge initial mass function to the neighborhood of the Hydrogen-burning limit.   The large number of stars with proper motion measurements allows kinematic features to be resolved in the color-magnitude diagram. We construct the $\\{l,b\\}$~velocity ellipse as a function of line of sight distance, demonstrating that its properties are quite sensitive to the distance to the objects under consideration. Finally, we use its proper motion analogue to attempt to classify the SWEEPS planetary candidates by kinematic membership with bulge or disk. The proper motion distribution of the candidates is consistent with the division of the stellar population between bulge and disk, but the candidate population is too small to draw further conclusions. We find no evidence that Ultra Short Period Planet (USPP) candidates are preferentially associated with the disk; instead we can only claim that the distribution is entirely inconsistent with {\\it all} the USPP originating in the disk.  Although the metallicities of the majority of our targets are unknown (with the exception of the bright objects, largely above the MSTO, for which VLT spectroscopy has been possible), the large number of objects with accurate proper motions allows the distance-uncertainty due to metallicity to be overcome by virtue of a sample containing $>$500 objects per bin. Armed with photometric distances and therefore mean transverse velocities for our distance-bins, we produce an independent determination of the stellar circular speed curve $\\vmpe$. Although the sampling of this rotation curve is relatively sparse, we clearly detect a transition away from solid body-like rotation at galactocentric radius $R_c =0.3-0.4$kpc, in line with evidence from radial velocities. Within the uncertainties of the distance estimates we have used, our limiting circular speed at large radius is consistent with that from radial velocity measurements.  With position-estimate dispersions reaching the 2-3{\\it milli-}pixel level (Section \\ref{s_pm}), we are near the limit of position-measurement with current HST instrumentation. The intrinsic velocity dispersion of the bulge and disk nevertheless makes kinematic classification of object groups with few members difficult. As the bulge is likely a superposition of multiple stellar populations (e.g. {Zoccali} et al. 2003), the intrinsic dispersion problem might be overcome if populations can be separated by metallicity on a star-by-star basis (c.f. {Soto} et al. 2007), which will require further, multifilter observations of this field."
    },
    "0809/0809.3687_arXiv.txt": {
        "abstract": "High resolution spectra of Comet 8P/Tuttle were obtained in the frequency range 3449.0--3462.2 cm$^{-1}$ on 3 January 2008 UT using CGS4 with echelle grating on UKIRT. In addition to observing solar pumped fluorescent (SPF) lines of H$_2$O, the long integration time (152 minutes on target) enabled eight weaker H$_2$O features to be assigned, most of which had not previously been identified in cometary spectra. These transitions, which are from higher energy upper states, are similar in character to the so-called `SH' lines recorded in the post Deep Impact spectrum of Comet Tempel~1 (Barber et al., 2007).We have identified certain characteristics that these lines have in common, and which in addition to helping to define this new class of cometary line, give some clues to the physical processes involved in their production. Finally, we derive an H$_2$O rotational temperature of $62\\pm5\\,\\rmn{K}$ and a water production rate of $(1.4\\pm0.3)\\times10^{28}$ molecules~s$^{-1}$. ",
        "introduction": "Cometary nuclei are the most primitive objects in the solar system. Within about 3 AU of the sun, solar heating causes the surface temperature of the nucleus to rise above 200 K; water ice sublimes and other volatiles and icy dust grains are expelled from the surface. The motion, chemistry and rate of sublimation of molecules from the icy grains are influenced by the solar wind and solar radiation. Initially the escaping gases form a dynamic, gravitationally-unbound atmosphere, the coma. They are subsequently swept into the comet's plasma and dust tails and eventually dissipate into interplanetary space. The surface of the nucleus is processed during the course of repeated visits to the inner solar system, becoming covered with a rubble blanket of particles too large to be dragged away by escaping gas. During the course of numerous perihelia, the outer surface becomes completely depleted of volatile materials which either escape from the surface or else migrate inwards. This creates an outer mantle of siliceous dust (Huebner, 2008), above a layered structure in which the more volatile species are found at greater depths (Prialnik et al., 2008). This structure results in the great variety that is observed in active comets.  Cometary nuclei typically have radii of $\\sim$3 km and are not able to be resolved from Earth. During close approaches, the inner coma of some comets are resolvable with large terrestrial instruments. However, in general, data obtained with Earth-based instruments are indicative of conditions in a wider region of the coma, and our understanding of conditions within the inner coma (a region extending up to a few hundred km, where energy transfer is collisionally-dominated, relies largely on modelling. Our understanding of the more extended regions of the coma, where the densities of the species are too low for thermodynamic equilibrium to exist, also relies on modelling. In particular, models of cometary coma predict the existence of so-called `solar pumped fluorescent' (SPF) H$_2$O emission lines. These originate from ro-vibrational excited upper states of the molecule, which, if observed in a collisionally-dominated region would be characteristic of kinetic temperatures of several thousand K, but are able to be produced in the cold, rarefied conditions existing in cometary coma, since these excited states, which have relatively large Einstein A coefficients, have time to decay radiatively before they are de-excited by collision with another molecule.  Despite the success of models in predicting SPF lines (Crovisier, 1984; Weaver and Mumma, 1984, Mumma et al., 1995), the physical conditions existing in cometary coma are highly complex and recent observations of Comet Tempel 1 (Barber et al., 2007) revealed a type of emission line, not previously recorded in cometary spectra, and not predicted by the existing cometary models. These so-called `SH' lines originate from upper states that are excited to higher vibrational energies than are the upper states of the SPF lines. Consequently, their production mechanisms are not able to be explained simply, by the existing cometary models. The near-Earth approach of Comet 8P/Tuttle in January 2008, provided an opportunity for us to search for, assign and characterise SH emission lines in its coma. We believe that our findings, which are presented in this paper, provide the basis for a more thorough investigation of the physical processes involved in the production of SH lines in cometary coma. ",
        "conclusions": ""
    },
    "0809/0809.2362_arXiv.txt": {
        "abstract": "{ APOGEE is a large-scale, NIR, high-resolution ($R\\sim$ 20,000) spectroscopic survey of Galactic stars. It is one of the four experiments in SDSS-III. Because APOGEE will observe in the $H$ band, where the extinction is six times smaller than in $V$, it will be the first survey to pierce through Galactic dust and provide a vast, uniform database of chemical abundances and radial velocities for stars across all Galactic populations (bulge, disk, and halo). The survey will be conducted with a dedicated, 300-fiber, cryogenic, spectrograph that is being built at the University of Virginia, coupled to the ARC 2.5m telescope at Apache Point Observatory. APOGEE will use a significant fraction of the SDSS-III bright time during a three-year period to observe, at high signal-to-noise ratio ($S/N>$100), about 100,000 giant stars selected directly from 2MASS down to a typical flux limit of $H<13$. The main scientific objectives of APOGEE are: (1) measuring unbiased metallicity distributions and abundance patterns for the different Galactic stellar populations, (2) studying the processes of star formation, feedback, and chemical mixing in the Milky Way, (3) surveying the dynamics of the bulge and disk, placing constraints on the nature and influence of the Galactic bar and spiral arms, and (4) using extensive chemodynamical data, particularly in the inner Galaxy, to unravel its formation and evolution. } ",
        "introduction": "Our knowledge of the structure, formation and evolution of galaxies is built upon observations of numerous galaxies as well as on detailed measurements in the Milky Way.  Only for the Galaxy can we quantify precisely the basic properties of individual stars from parallaxes, proper motions and high-resolution spectroscopy. Only for the Milky Way we can make a census of minority stellar populations which, as in the case of the stellar halo, can be central to understanding the earliest epochs of galactic formation. In addition, the spatial resolution of gas and dust maps available for the Galaxy is orders of magnitude better than for extragalactic systems.  As a result of developments in instrumentation over the last decades, high-resolution spectroscopic studies have become more and more precise, involve larger samples, and reach fainter magnitudes. These advances have made it possible to carry out a thorough study of the stellar populations in the solar neighborhood and measure high precision chemical abundances for dozens of elements and hundreds of stars. These observations have revealed, among other things, a clear dichotomy between the chemical patterns in thin- and thick-disk stars, despite their overlapping distributions in phase space.  Progress has probably been even more spectacular for studies using lower spectral resolution, thanks to massive multiplexing on fiber spectrographs such as those used in the Sloan Digital Sky Survey (SDSS; York et al. 2000), which operate at $R=$ 2,000. The information per spectrum is more limited, but this is compensated by deeper and much larger samples of stars.  More than half a million stars have already been observed in the course of the SDSS, and this number will double over the next three years. Several tens of thousands of stellar spectra have already been analyzed in recent studies from SDSS (Allende Prieto et al. 2006, Ivezi\\'c et al. 2008), and higher resolution ($R\\sim$ 7,500) spectra, over a limited spectral window, are now becoming available from the Radial Velocity Experiment (RAVE; Zwitter et al. 2008). This allows a thorough characterization of the properties of the Milky Way stellar halo, and provide information on the thick disk over a wide range of galactocentric distances. Now we know that the thick-disk stars are old, but intermediate in age between the halo and the thin disk, and that the thick disk does not show radial abundance gradients such as those well-known in the thin disk  (Allende Prieto et al. 2006).  In stark contrast with the fast pace at which we have been able to make progress in our exploration of the outer Galaxy at low resolution and the steady accumulation of high resolution information on nearby stars, the situation for understanding the non-local, inner Milky Way at any resolution has been much less favorable. The high extinction in the Galactic mid-plane and towards the Galactic center severely complicates observational studies at traditional optical wavelengths. Spectroscopic studies of the Galactic bulge have been limited to small numbers of stars --- almost exclusively in Baade's window. The Apache Point Galactic Evolution Experiment (APOGEE) is targeted to dramatically change this situation by using highly-multiplexed near-IR high-resolution spectroscopy to pierce through dust and probe stars in the central parts of the galaxy at high resolution.  In addition, APOGEE will push the systematic, high resolution, high throughput spectroscopic study of Galactic stars into the outer Galaxy, creating a uniform spectroscopic database for stars in all populations of the Milky Way.  Such homogeneous but extensive datasets are necessary to produce an integrated picture of the chemical and kinematical evolution of our home galaxy. ",
        "conclusions": ""
    },
    "0809/0809.1426_arXiv.txt": {
        "abstract": "We study the evolution of the star formation rate (SFR) of mid-infrared (IR) selected galaxies in the extended Chandra Deep Field South (E-CDFS). We use a combination of U-K GaBoDS and MUSYC data, deep IRAC observations from SIMPLE, and deep MIPS data from FIDEL. This unique multi-wavelength data set allows us to investigate the SFR history of massive galaxies out to redshift $z \\sim 1.8$. We determine star formation rates using both the rest-frame ultraviolet luminosity from young, hot stars and the total IR luminosity of obscured star formation obtained from the MIPS $24 \\,\\um$ flux. We find that at all redshifts the galaxies with higher masses have substantially lower specific star formation rates than lower mass galaxies. The average specific star formation rates increase with redshift, and the rate of incline is similar for all galaxies (roughly $(1+z)^{n}, n = 5.0 \\pm 0.4$). It does not seem to be a strong function of galaxy mass. Using a subsample of galaxies with masses $M_{*}> 10^{11}\\Msol$, we measured the fraction of galaxies whose star formation is quenched. We consider a galaxy to be in quiescent mode when its specific star formation rate does not exceed $1/(3\\times t_{H})$, where $t_{H}$ is the Hubble time. The fraction of quiescent galaxies defined as such decreases with redshift out to $z \\sim 1.8$. We find that, at that redshift, $19 \\, \\pm9\\%$ of the $M_{*}> 10^{11}\\Msol$ sources would be considered quiescent according to our criterion. ",
        "introduction": "The star formation history of massive galaxies is not well understood.  Standard galaxy formation models have difficulty reproducing today's red and dead galaxies, unless mechanisms are introduced that prevent the gas from cooling and forming stars. To better constrain the models, it is useful to determine the star formation rates (SFRs) of galaxies as a function of mass and redshift. This has been done out to redshift $z = 1$, using the COMBO-17 survey \\citep{zheng}.\\\\ A surprising result of their study was that the specific star formation rates (SFR per unit stellar mass, SSFR) of high mass galaxies evolve at the same rate as those of low mass galaxies, where the most massive galaxies are offset to lower SSFRs. At higher redshifts, studies of star formation have so far focused mainly on either specific galaxy populations or specific redshift regimes (e.g, Lyman break galaxies, (Steidel et al. 1996, 1999), distant red galaxies,  (Papovich et al. 2006)). Papovich et al. (2006) found that massive, red galaxies at $1.5 \\le z \\le 3$ have SSFRs that are comparable to the global value integrated over all galaxies. Given the fact that we can already see the Hubble sequence in place at $z \\sim 1$, this means that the period between $1 \\lesssim z \\lesssim 2.5$ is an interesting stage of transition, where massive galaxies evolve from actively star forming systems to the passive galaxies we observe in the local universe. The connection between the high and low redshift galaxy populations is not yet clear. \\\\ In this paper, we investigate the star formation history of massive galaxies ($\\Mstar > 10^{11} \\msol$), through measurements of the specific star formation rate from $z \\sim 0.2$ to $z \\sim 1.8$. We use a combination of a new, wide field \\spitzer/IRAC survey called SIMPLE and ancillary data ranging from the near-ultraviolet (near-UV) to the mid-infrared (MIR). \\\\ Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm  \\Lambda}=0.7$, and $H_{\\rm 0}=70$~km s$^{-1}$ Mpc$^{-1}$. All magnitudes are given in the AB photometric system. We denote magnitudes from the four \\spitzer\\ IRAC channels as \\mone, \\mtwo, \\mthree, and \\mfour, respectively. Stellar masses are determined assuming a Kroupa initial mass function (IMF). ",
        "conclusions": "\\label{conc} We investigate the star formation history of massive galaxies out to redshift $z \\sim 1.8$, by analyzing specific star formation rates (SSFRs) of a sample of $\\sim 3,400$ sources from SIMPLE, a survey that combines new \\spitzer/IRAC observations of the E-CDFS with ancillary data ranging from the near-UV to the MIR. We find quiescent galaxies with masses higher than $10^{11} \\msol$ out to the highest redshift probed, $z \\sim 1.8$. At this redshift, they form $19 \\pm9\\%$ of the total number of massive galaxies. The SSFR is an increasing function with redshift (roughly $(1+z)^{n}, n = 5.0 \\pm 0.4$) for galaxies in all mass bins. The mean SSFRs are smaller in high-mass galaxies than in low-mass galaxies at all redshifts. It is interesting to consider this result in the context of \"downsizing\", a term which generally is taken to imply that more massive galaxies formed their stars before less massive galaxies (Cowie et al. 1996). An increasing amount of observational evidence supports this idea, in particular through studies of the (specific) star formation rate (Juneau et al. 2005, P\\'erez-Gonz\\'alez et al. 2005, Caputi et al. 2006, Papovich et al. 2006, Reddy et al. 2006 and Noeske et al. 2007). Figure \\ref{ssfr} shows that the SSFRs of massive galaxies are the lowest of the whole sample. This indicates that they have already formed the bulk of their stars and that active star formation has shifted to the galaxies that are less massive. Additional support comes from a theoretical perspective. \\citet{guo} investigated the contribution of star formation to galaxy growth in the Millennium Simulation. They found that even out to $z \\sim 4-5$ less massive galaxies are always growing faster than galaxies of higher stellar mass.\\\\ It is interesting to see that although we see a change in the locus of star formation (from massive to less massive systems), we do not find any mass-dependence of the {\\em evolution} of the specific star formation rates. High-mass galaxies and low-mass galaxies appear to evolve at the same rate out to $z\\sim 1.8$, although deeper data are necessary to reduce possible effects of incompleteness of the lowest mass bins at high redshift.\\\\ Next we compare our passive fraction in our highest redshift bin to observational results from the literature. A recent estimate can be found in work by Labb\\'e et al. (in prep.), who found $35 \\pm 7\\%$. We also re-determined the quiescent fraction for Kriek et al. (2006) by analyzing the full sample presented in Kriek et al. (2008). We defined all galaxies without emission lines and $SSFR < 0.05~ Gyr^{-1}$ to be quiescent. Out of the 28 galaxies at redshift $z > 2, 36 \\pm 9\\% $ are quiescent according to this method, applying a bootstrap error. The values of Labb\\'e et al. (in prep.) and Kriek et al. (2008) are consistent with our fraction within $2\\sigma$. We note that our value is lower, but we emphasize that all studies use different definitions which can influence the result. We return to that below.\\\\ To investigate the difference in the estimates of the quiescent fractions further, we look at our MIPS fluxes and compare them with results from the CDFS (Labb\\'e et al. 2005). For the overlapping sources in their sample and ours, we observed a median positive offset in MIPS flux of $4\\, \\mu Jy$ with respect to the CDFS, which means our SFRs are overestimated. To investigate whether this offset could be responsible for the difference in passive fraction, we simulated the effect errors in MIPS flux would have on our results. We randomly added a measurement of the difference in MIPS flux to a collection of simulated passive galaxies with $M_{*} > 10^{11} \\Msol$ and determined the number of times the SSFR of such a source would scatter above the limit by which we define a passive galaxy. Fifteen percent of the passive galaxies were classified as star forming after performing this test. This raises the fraction of quiescent galaxies at redshift $z \\sim 1.8$ by 2\\%, which is not enough to explain the difference between the fractions.\\\\ We next compare our quiescent classification with the results from rest-frame optical spectroscopy of Kriek et al. (2008). Their sample contains 11 sources in the E-CDFS, 2 of which show no emission lines and are best fit with a passively evolving SED. We detect one of these sources in 24 \\um \\, with $S/N < 1$. Our SFR for it is $9 \\pm 10 \\,\\Msol yr^{-1}$, which is consistent with the $0.7\\pm 0.7 \\,\\Msol yr^{-1}$ that Kriek et al. (2008) find. The other source has a SSFR of $2.4 \\times \\,10^{-10}\\,yr^{-1}$, which exceeds the limit of $1 / (3\\times t_H)$. It is still likely to be quenched as its SSFR is smaller than $1 / t_H$. In summary, our results agree reasonably well for these two sources, but the exact definition of quiescence may cause variations in the result. As to illustrate this further, we relaxed our limit to $SSFR < 1 / t_H$. For this limit, we find a quiescent fraction of $30\\pm 7\\%$ at $z \\sim 1.8$ which is roughly 1.5 times the fraction we found earlier.\\\\ In addition, the 24 \\um \\,emission we detect could be due to the presence of a weak and obscured AGN. This would mean that some of the galaxies we call star forming could in fact be quiescent galaxies hosting an obscured AGN. If this is the case, our fraction underestimates the real fraction. This is possible as evidence exists that AGN activity is widespread among massive galaxies at these redshifts (Rubin et al. 2004, Daddi et al. 2007,  Kriek et al. 2007). We removed all X-ray detected sources from our sample as probable AGN candidates, but have no means to identify obscured AGNs that show strong 24 \\um \\,flux and weak X-ray emission. Our fraction would also be underestimated if the galaxy would hide an obscured starburst in its center.\\\\ Another effect is the error in the photometric redshifts, which we could have undervalued. The EAZY redshifts we use here are the results of several runs where we varied templates and error determination, which did not largely affect the outcome. We found that taking the 68\\% confidence range on the photometric redshifts of the galaxies leads to variations in the inferred SFR of 0.1 dex, which is not enough to significantly affect the results. Finally, there is a large diversity in the fields used for these studies and field-to-field variations could also be causing discrepancies. \\\\ Our most robust result is that we find a high fraction of galaxies with MIPS detections at redshift $z \\sim 1.8$ and a small, but non-negligible fraction of quiescent galaxies, which we interpret as a lower limit. The galaxies that are detected in MIPS at redshift $z \\sim 1.8$ are in some way active, either through star formation or black hole growth. Deeper 24 \\um\\, data and spectroscopic information will be crucial to be able to elaborate on this more.\\\\ \\\\   We thank the referee for useful suggestions. We thank the FIDEL team for early access to their 24 \\um \\,image and  in particular Jackie Monkiewicz, for reducing the 24 \\um \\, data. We are grateful to Jarle Brinchmann for his advice regarding the SDSS data and Natascha F\\\"orster-Schreiber for help with the SED fitting. This research was supported by grants from the Netherlands Foundation for Research (NWO), and the Leids Kerkhoven-Bosscha Fonds. Support from National Science Foundation grant NSF CAREER AST-0449678 is gratefully acknowledged."
    },
    "0809/0809.4729_arXiv.txt": {
        "abstract": "We present an alternative cosmology based on conformal gravity, as originally introduced by H. Weyl and recently revisited by P. Mannheim and D. Kazanas. Unlike past similar attempts our approach is a purely kinematical application of the conformal symmetry to the Universe, through a critical reanalysis of fundamental astrophysical observations, such as the cosmological redshift and others.  As a result of this novel approach we obtain a closed-form expression for the cosmic scale factor $R(t)$ and a revised interpretation of the space-time coordinates usually employed in cosmology. New fundamental cosmological parameters are introduced and evaluated. This emerging new cosmology does not seem to possess any of the controversial features of the current standard model, such as the presence of dark matter, dark energy or of a cosmological constant, the existence of the horizon problem or of an inflationary phase. Comparing our results with current conformal cosmologies in the literature, we note that our kinematic cosmology is equivalent to conformal gravity with a cosmological constant at late (or early) cosmological times.  The cosmic scale factor and the evolution of the Universe are described in terms of several dimensionless quantities, among which a new cosmological variable $\\delta$ emerges as a natural cosmic time. The mathematical connections between all these quantities are described in details and a relationship is established with the original kinematic cosmology by L. Infeld and A. Schild.  The mathematical foundations of our kinematical conformal cosmology will need to be checked against current astrophysical experimental data, before this new model can become a viable alternative to the standard theory. ",
        "introduction": "Introduction}  Modern cosmology has advanced very rapidly during these last decades, producing an impressive model of the Universe, but our current understanding is still troubled by many open questions and puzzles. Since the original observations of cosmological redshift in spectral lines, done by V. M. Slipher and E. P. Hubble almost one century ago and since the application of Einstein's General Relativity to cosmological theoretical models, we have progressed a long way towards our current picture, where the contents of the Universe are described in terms of two main components, \\textit{dark matter} and \\textit{dark energy}, accounting for most of the observed Universe, with ordinary matter just playing a minor role.  The history of recent experimental observations which led to postulate the existence of these two components is well known, as well as the many past and current theoretical explanations (see for example \\cite{Turner:1999md}, \\cite{Freedman:2003ys}, \\cite{Yao:2006px}, \\cite{Perlmutter:1998np}, \\cite{Riess:1998cb}, \\cite{Riess:2004nr}, \\cite{Riess:2006fw}, \\cite{Spergel:2003cb}, \\cite{Spergel:2006hy}), but since there is no evidence of the real nature of dark matter and dark energy, we have to conclude that our comprehension of the natural world is limited to only a very small percentage of it (the ordinary matter component), a statement potentially very embarrassing for cosmology and physics, if taken at face value.  Several alternatives to dark matter and dark energy have been proposed (for comprehensive reviews see for example \\cite{Mannheim:2005bf}, \\cite{Schmidt:2006jt}, \\cite{Nojiri:2006ri}, \\cite{Clifton:2006jh}, \\cite{Zakharov:2008zz} and references therein) which can be approximately divided into two categories: those retaining the Newton-Einstein gravitational paradigm, while introducing ad hoc corrections to explain dark matter and dark energy and those breaking away substantially from established gravitational theories. Following this second line of thought, we will concentrate our attention on the theory of Conformal Gravity\\ (CG), a fourth order extension of Einstein's second order General Relativity (GR) as a possible framework for the solution of current cosmological problems. ",
        "conclusions": "Conclusions}  We presented in this work the mathematical foundations of a new kinematical approach to conformal cosmology. This was based on the assumption that the observed redshift is mainly due to a gravitational origin and that the gravitational potential of conformal gravity might be the cause of the observed \\textquotedblleft expansion\\textquotedblright\\ of the Universe.  We have seen how the original Mannheim-Kazanas potential can support this explanation and how the chain of transformations, from Static Standard Coordinates to the Robertson-Walker metric, can lead to a unique expression of the cosmic scale factor. In particular, the $\\mathbf{k}=-1$ solution is the only capable of producing the observed galactic redshift and describes the evolution of the Universe in terms of an initial expansion phase (which can still be traced back to an initial singularity) up to a point of maximum expansion, then followed by a symmetrical contraction phase, towards a final singularity.  We have compared our solutions with those of current conformal cosmologies and noted that our kinematic cosmology is recovered from conformal gravity with a cosmological constant at late (or early) cosmological times. This implies that during these cosmological epochs the cosmological and gravitational redshifts might be in fact equivalent, as assumed in our kinematical approach to conformal cosmology.  Our detailed analysis of the solutions has also shown the importance of using dimensionless quantities and, in particular, of the cosmological variable $\\delta\\equiv\\gamma/2\\sqrt{\\left\\vert k\\right\\vert }$ which effectively combines together the two parameters $\\gamma$ and $k$ (or $\\kappa$) originally introduced by Mannheim and Kazanas. We also introduced the hypothesis that $\\delta$ might actually represent a true cosmic time, in terms of which the evolution of the Universe would be described by the very simple Eq. (\\ref{eqn4.29}). If this interpretation is correct, all the space-time dimensionless variables $\\alpha$, $r$, $\\chi$, $\\zeta$, $z$, $\\delta$ are linked together and also related to the current value $\\delta(t_{0})$ of the cosmic time, through the transformations outlined in Table \\ref{TableKey:table1}.  The current value $\\delta(t_{0})$ should be small and positive, as indicated by Mannheim's estimate of the parameter $\\gamma_{0}$, but its actual value has to be determined by more precise fitting of astrophysical data, such as the luminosities of type Ia Supernovae or others. This will be the objective of a second part of this project \\cite{Varieschi:2008va}, where we will also try to explain the anomalous local blueshift region, which is implied by our cosmological solution.  If our model will be successful in explaining current experimental data, \\textit{kinematical conformal cosmology} might become a viable alternative model for the description of the Universe, with the advantage of avoiding altogether the introduction\\ of dark matter and dark energy, or other controversial features of the current standard model."
    },
    "0809/0809.4490_arXiv.txt": {
        "abstract": "Recent millimeter-VLBI observations of Sagittarius A* (Sgr A*) have, for the first time, directly probed distances comparable to the horizon scale of a black hole.  This provides unprecedented access to the environment immediately around the horizon of an accreting black hole.  We leverage both existing spectral and polarization measurements and our present understanding of accretion theory to produce a suite of generic radiatively inefficient accretion flow (RIAF) models of Sgr A*, which we then fit to these recent millimeter-VLBI observations.  We find that {\\em if} the accretion flow onto Sgr A* is well described by a RIAF model, the orientation and magnitude of the black hole's spin is constrained to a two-dimensional surface in the spin, inclination, position angle parameter space.  For each of these we find the likeliest values and their 1-$\\sigma$ \\& 2-$\\sigma$ errors to be $a=0^{+0.4+0.7}$, $\\theta={50^\\circ}^{+10^\\circ +30^\\circ}_{-10^\\circ-10^\\circ}$, and $\\xi={-20^\\circ}^{+31^\\circ+107^\\circ}_{-16^\\circ-29^\\circ}$, when the resulting probability distribution is marginalized over the others. The most probable combination is $a=0^{+0.2+0.4}$, $\\theta={90^\\circ}_{-40^\\circ-50^\\circ}$ and $\\xi={-14^\\circ}^{+7^\\circ+11^\\circ}_{-7^\\circ-11^\\circ}$, though the uncertainties on these are very strongly correlated, and high probability configurations exist for a variety of inclination angles above $30^\\circ$ and spins below $0.99$.  Nevertheless, this demonstrates the ability millimeter-VLBI observations, even with only a few stations, to significantly constrain the properties of Sgr A*. ",
        "introduction": "Understanding the structure and dynamics of black hole accretion flows has remained a central problem in astrophysics.  Only in the past decade have sufficient numerical resources existed to perform self-consistent, three-dimensional magneto-hydrodynamic (MHD) simulations capable of resolving the magneto-rotational instability, believed to be responsible for angular-momentum transport in black-hole accretion flows.  However, a true {\\em ab initio} computation is still well beyond reach, resulting in the need for a variety of simplifying assumptions (e.g., the suitability of MHD, importance of electron-ion coupling, properties of accelerated electrons, etc.). As a result, the number of applicable models has rapidly proliferated, many of which are capable of describing the variety of phenomena observed.   This is due in large part to the inability of current observations to resolve horizon scales. Unfortunately, the compact nature of black holes makes it very difficult to access the inner-edge of black hole accretion flows.  Sagittarius A* (Sgr A*), the bright radio point source coincident with the center of the Milky Way, is presently the best studied known black-hole candidate.  Observations of orbiting OB-stars, the closest of which passes within $45\\,{\\rm AU}$ of Sgr A*, have produced a measured mass of $4.5\\pm0.4\\times10^6\\,\\Ms$ and an Earth-Sgr A* distance\\footnote{The mass and distance measurements are strongly correlated, with mass scaling roughly as $M\\propto D^{1.8}$.} of $8.4\\pm0.4\\,\\kpc$ \\citep{ghez09,gillessen08}.  Already, there is strong evidence for the existence of horizon in this source \\citep{Brod-Nara:06,Brod-Loeb-Nara:08}.  However, in many ways, Sgr A* is very different than the supermassive black holes in AGN. Unlike its active brethren, Sgr A* is vastly under-luminous, with a luminosity many orders of magnitude smaller than the Eddington luminosity.  As a result, it is widely expected that Sgr A*'s accretion flow is quite different from those in AGN, though perhaps more indicative of the roughly 90\\% of supermassive black holes that are not presently in an active phase.  Even when strong-gravitational lensing is accounted for, the apparent angular size of Sgr A*'s horizon is only $55\\pm2\\,\\muas$, a factor of two larger than the next largest black hole (M87) and orders of magnitude larger than any other known black hole (including all stellar-mass black holes).  Nevertheless, despite its tiny angular size, recent millimeter VLBI experiments have successfully resolved this scale \\citep{Doel_etal:08}.  Since only three telescopes were involved with these observations, the resulting $u$--$v$ coverage of the measured visibilities is very sparse.  As a result only two simple models of Sgr A*'s image, a Gaussian and an annulus, were fit to the data by \\citet{Doel_etal:08}.  However, we have a great deal of additional information about Sgr A*, including its spectral and polarization properties.  We may also require physical consistency in any model (which would likely rule out an annulus, for example).  Furthermore, we would like to evaluate the ability of, and optimize for this purpose, future millimeter and sub-millimeter VLBI experiments to constrain fundamental properties of the accretion flow in Sgr A*.  This paper demonstrates the fitting and parameter estimation procedure for a simple radiatively inefficient accretion disk model (RIAF).  In particular, we show that even from very sparse baseline coverage it is possible to robustly extract interesting parameters of a generic RIAF.  A study that considers how future high frequency VLBI observations can extend this work will be presented elsewhere \\citep{fish08}.  In the analysis presented here, we make full use of all the observations described in \\citet{Doel_etal:08}.  These include two days of VLBI observations at 230GHz that targeted Sgr A* using an array consisting of the James Clerk Maxwell Telescope (\\JCMT) on Mauna Kea, the Arizona Radio Observatory Submillimeter Telescope (\\SMTO) on Mt. Graham in Arizona, and one $10,\\m$ dish of the Coordinated Array for Research in Millimeter-wave Astronomy (\\CARMA) in California.   Robust detections of Sgr A* and correlated flux density measuremnets were obtained on the \\CARMA-\\SMTO~and \\JCMT-\\SMTO~baselines.  No detections were found on the \\CARMA-\\JCMT~baseline resulting in upper limits on the correlated flux density.  In addition, contemporaneous total flux density measurements at the same frequency were obtained by using the full \\CARMA~array operating as a stand-alone instrument.  Full details of the observations, calibration, and data processing can be found in \\citet{Doel_etal:08}.  In section \\ref{sec:RIAF} we present the simple accretion flow model that we employ.  Sections \\ref{sec:DataAnalysis} and \\ref{sec:Results} discuss the data fitting procedure, including the Bayesian method by which we do the parameter estimation, and the generic constraints placed by the current VLBI results.  Section \\ref{sec:parameters} details how we define our uncertainties and presents our parameter estimates.  Finally, we conclude in section \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  Despite the sparse $u$--$v$ coverage, and the existence of a detection on only one long baseline, we have been able to significantly constrain the possible parameters of an accretion flow onto Sgr A*. Within the context of a qualitative RIAF model that fits the observed spectra of Sgr A*, we have have been able to substantially constrain the orientation of Sgr A*'s spin.  The magnitude of this spin is less well determined, though the black hole cannot be maximally rotating.  This result is relatively insensitive to the black hole mass.  We have not ascertained the strength of our constraints' dependence upon the particular RIAF model, though our estimates are consistent with earlier efforts comparing longer wavelength observations to an alternative RIAF model.  However, there are two additional reasons to believe that our results will be generic of RIAF models.  The first is that the underlying physics that limits the orientation is the Doppler beaming and boosting that is dependent primarily upon the Keplerian velocity profile of the accretion flow, a feature that is generic among RIAF models.  The second is the weak dependence upon large variations in mass, implying that changing the scale lengths of the accretion model does little to alter our results. Thus, we expect the qualitative form of our constraints to be a generic feature of all RIAF models for Sgr A*, and our quantitative results to be roughly correct.  Of course, should the emission observed from Sgr A* not arise in an accretion flow, our results could be quite different.  As implied by the right-panel of Fig. \\ref{fig:IVV}, additional long-baseline observations are sorely needed to unambiguously determine both the applicability of RIAFs to Sgr A* and constrain their parameters.  In a companion paper, \\citet{fish08} we report upon an analysis of the ability of possible millimeter VLBI arrays to do so.  Given the number of potential radio telescopes that could be added to the current $1.3\\,\\mm$ VLBI array \\citep{doeleman08b}, the prospects for making significant progress in model parameter estimation are excellent.  Already it is clear that with the advent of millimeter VLBI we are now entering an era of precision black-hole accretion physics."
    },
    "0809/0809.2755_arXiv.txt": {
        "abstract": "A set of interpolating functions of the type $f(v)=\\left\\{\\sin\\left[\\frac{\\pi}{2}v\\right]/\\left(\\frac{\\pi}{2}v\\right)\\right\\}^n$ is analyzed in the context of the smoothed-particle hydrodynamics (SPH) technique. The behaviour of these kernels for several values of the parameter $n$ has been  studied either analytically as well as numerically in connection with several tests carried out in two dimensions.  The main advantage of this kernel relies in its flexibility because for $n=3$ it is similar to the standard widely used cubic-spline, whereas for $n>3$ the interpolating function becomes more centrally condensed, being well suited to track discontinuities such as shock fronts and thermal waves. ",
        "introduction": "The technique called smoothed particle hydrodynamics (SPH) was introduced in 1977 \\cite{l77} and \\cite{gm77} to simulate the evolution of fluids and plasmas in three dimensions. It is a gridless Lagrangian method hence one has not to be worried about the definition of a mesh and its further remapping because the grid is somehow advected by the particles themselves (see \\cite{m05} for a recent review of this technique). A central point of the SPH formalism is the concept of interpolating function (or kernel) through which the continuum properties of the fluid are recovered from a discrete sample of $N$~points with mass $m_i$~which move according to the hydrodynamical laws. A good interpolating kernel must satisfy a few basic requirements: it must tend to a delta function in the continuum limit and has to be a continuous function with, at least, definite first and second derivatives. From a more practical point of view it is also advisable to deal with symmetric kernels of finite range, the latter to avoid $N^2$ calculations. A prototype of kernel is the Gaussian kernel:  \\begin{equation} W^\\mathrm{G}(v,h) =\\frac{1}{(\\sqrt\\pi h)^d}\\exp{(-v^2)} \\label{eq:eq1} \\end{equation}  where $d$ is the spatial dimension and $v=\\vert$$\\mathbf{r - r'}$$\\vert/h$ is the normalized distance between particles in terms of the characteristic smoothing length $h$. The precise value of $h$ sets a spatial length-scale which is usually taken as the local resolution of SPH. Nevertheless the Gaussian kernel has an infinite range thus it is most practical to use a similar function but with compact support. A function with the required properties is the cubic $M_4$~spline \\cite{s46},~\\cite{ml85} defined as:  \\begin{equation} M_4(v,h) =\\frac{A}{h^d} \\cases {1-\\frac {3}{2}v^2+\\frac {3}{4}v^3\\qquad\\qquad 0\\le v\\le 1\\cr\\cr \\frac{1}{4}(2-v)^3\\qquad\\qquad\\qquad 1<v\\le 2\\cr\\cr 0\\qquad\\qquad\\qquad\\qquad\\qquad\\quad v>2\\cr} \\label{eq:eq2} \\end{equation}  where $A=\\frac {2}{3}$, $\\frac {10}{7\\pi}$ and $\\frac {1}{\\pi}$ in one, two and three dimensions respectively. This kernel works very well and it is computationally efficient as it has been checked worldwide by many people \\cite{b90} and \\cite{m92} using SPH in many areas of physics and astrophysics.  In general, hydrodynamical simulations with SPH should not significantly depend on the chosen kernel but there could be special cases in which the precise profile of the interpolating device becomes relevant. These situations appear wherever there is an abrupt change of physical variables. Trivial examples are strong shock formation and thermal discontinuities. In these cases the values of the variables around the jump region are severely damped by the interpolation procedure and the values at the peak are underestimated. In other cases the numerical noise can become high enough to blur the true value of the variables. If there are chemical or nuclear reactions, which are very sensitive to temperature, the outcome might be greatly altered or even completely wrong. Several ways to better handle discontinuities in SPH have been proposed. One of them is to enhance the artificial viscosity algorithm of the codes in order to improve the jump of the variables at the discontinuity and its thickness  \\cite{m97}. Another route is to try a more clever interpolation using kernels especially devised to handle large gradients. In particular tensorial kernels with ellipsoidal geometry were proposed by \\cite{ovsm98}. In normal conditions these kernels reduce to the spherically symmetric cubic-spline but, in the presence of a shock the sphere becomes an ellipsoid with its minor axis aligned with the shock direction leading to  an improvement of the resolution. However several interesting properties of the spherically symmetrical kernels are lost during their transformation into ellipsoids. In particular, the use of spherically symmetrical kernels guarantees that any function is approximated to second order in $h$, thus linear functions are exactly reproduced. In addition energy and momentum are not so well conserved when ellipsoidal kernels are used. An alternative to tensorial interpolators is to consider strongly peaked kernels with spherical symmetry. In this case, for a fixed value of the smoothing parameter $h$, there is an increase in resolution, although that enhancement is always accompanied by an increase of the numerical noise \\cite{fq96}. In this paper we present a one-parametric family of spherically symmetric kernels based on harmonic-like functions ($W^{H}_n$ kernels hereafter) with compact support. Specifically we propose that the set of functions:  \\begin{equation} W^{H}_n(v,h)=B_n(h) \\cases {1 \\qquad\\qquad\\qquad\\qquad v=0\\cr\\cr \\left(\\frac{\\sin\\left[\\frac{\\pi}{2}v\\right]}{\\frac{\\pi}{2}v}\\right)^n\\qquad 0<v\\le 2\\cr\\cr 0\\qquad\\qquad\\qquad\\qquad v>2\\cr} \\label{eq:eq3} \\end{equation}  \\noindent where $B_n$~is a normalization factor, have several interesting features which make them suitable to SPH studies. A change in the value of the governing parameter $n$ lead to different shapes of the interpolative function, from more extended to more centrally condensed profiles as $n$~increases. Moreover, for $n=3$~it is very similar to the well-known cubic spline given by Eq. (2). Thus we suggest that the use of the $W^{H}_n$~set of functions in current SPH calculations would add  more flexibility to the numerical scheme without practically introducing any inconvenient.  Although the function given by Eq. (3) is of great importance in signal analysis theory, where the case n=2 corresponds to the so-called window function, as far as we know it has never been used before in connection with SPH. A somehow related interpolator $W\\propto (1-v^2/4)(1+\\cos\\pi v/2)$~was studied by \\cite{fq96} although only in 1D. However it was a single kernel, not a continuous family as those given by Eq.(3)  Interpolators of type $W_n^H$~with a high $n$~should be compared not only to the $M_4$~but to higher order splines. In this respect the quintic $M_6$~polynomial interpolator \\cite{s46} could be taken as representative. It has a compact support and has been used to model flows using SPH \\cite{mfz97}. As the $M_6$~spline was originally devised to work within a radii of $3h$~we have renormalized it to $2h$~in order to make plausible comparisons with $W_n^H$:   \\begin{equation} M_6(v,h) = \\frac{C}{h^d} \\cases {(2-v)^5-6 (\\frac{4}{3}-v)^5+15 (\\frac{2}{3}-v)^5 \\qquad 0\\le v\\le \\frac{2}{3}\\cr\\cr (2-v)^5-6 (\\frac{4}{3}-v)^5\\qquad\\qquad\\frac{2}{3}<v\\le \\frac{4}{3} \\cr\\cr (2-v)^5\\qquad\\qquad\\qquad \\frac{4}{3}<v\\le 2\\cr\\cr 0\\qquad\\qquad\\qquad\\qquad\\qquad\\quad v> 2\\cr} \\label{eq:eq2} \\end{equation}  \\noindent where $C=\\frac{243}{2560}, \\frac{15309}{61184\\pi}$~and $\\frac{2187}{10240\\pi}$~in one, two and three dimensions respectively.  In Section 2 we give the main mathematical features of the kernels defined by Eq.~(\\ref{eq:eq3}), and discuss their abilities to handle steep functions. A comparative analysis of the performance of these interpolators in disordered discrete systems is also given in the same section. In Section 3 we analyze the behaviour of $W^{H}_n$ in connection to two classical tests: (1) the propagation of a thermal wave arising from a thermal discontinuity in 2D cartesian coordinates and, (2) the evolution of a blast wave reaching the Sedov phase, also calculated in 2D. Finally a brief critical discussion concerning the virtues and shortcomings of the proposed $W^{H}_n$~set of functions and  the main conclusions of our work is provided in Section 4. ",
        "conclusions": "A one-parameter family of interpolating kernels with compact support based on the harmonic-like functions, $W^H_n$, defined by Eq. (\\ref{eq:eq3}) have been analyzed and checked in the context of the smoothed particle hydrodynamics method. Formally the widely used polynomial interpolators are linked to the proposed kernels via the Fourier transform of a function similar to $W^H_n$. We have found that the set of functions defined by Eq. (\\ref{eq:eq3}) also displays good enough mathematical features as to deserve being considered smoothing kernels by themselves. Using $W^H_n$~has several interesting advantages: a) if the exponent $n$~ is conveniently chosen these interpolators are able to mimic with great accuracy various of the most common kernels used so far in SPH studies. In particular the cubic and quintic spline kernels $M_4$~and $M_6$~are well reproduced by taking $n=3$~and $n=5$~in Eq. (3). Even the truncated Gaussian kernel is reasonably fit for $n=2$. Such equivalences seem to be robust in the light of the two realistic test cases analyzed in Section 3, b) in the limit of a very large number of particles the use of functions with $n>3$~improves the resolution when strong gradients are present (Section 2), c) Unless the cubic spline, whose second derivative is not smooth in several points, the $W^H_n$~has well-behaved second derivatives allowing the set to be derived many times, d) it adds even more flexibility to the SPH technique because the quality of the interpolation can be selected by simply varying the exponent $n$~in equation 3. Changing the exponent $n$~is really straightforward and does not introduce any computational overload in the numerical scheme.  In current calculations, where a moderate number of particles are put into the system, the choice of kernels with a large exponent $n$~is limited by the numerical noise. Therefore the use of high-peaked interpolators should be restricted to special circumstances. It has been shown in Section 2 that the improvement in resolution is not monotonic as $n$~rises: above $n=6$~it is hard to achieve a much better resolution. Just the opposite is true for the noise because it grows faster for $n>6$. Thus, it is advisable to restrict the range of the exponent $n$~to the interval $2\\le n\\le 6$~in practical applications.   Another question related to the number of particles is that either the cubic or the quintic spline kernels can also have their resolution increased by decreasing the smoothing-length parameter $h$. Reducing $h$~has a similar effect as increasing $n$~in Eq. (\\ref{eq:eq3}) because there is an enhancement of the resolution accompanied by an increase of the numerical noise due to the reduction in the number of neighbours of the particle. That enhancement in the resolution can be used to devise an adaptive kernel scheme which has proven useful to handle shocks \\cite{s06}. Nevertheless varying the resolution by altering $h$~in the cubic spline is not as good as to change $n$~in $W^H_n$~for two reasons. First, a high-peaked kernel is still able to count particles which are farther than those seen by a low-peaked one with comparable resolution but lesser $h$. In some circumstances the statistical weight of these particles could be high enough to influence the average. The second reason is simply that to control both the resolution and the noise it is also better to have two parameters to tune: the size of $h$, which sets the number of neighbours of a given particle, and the exponent $n$~in Eq. (\\ref{eq:eq3}).  The only drawback of the $W^H_n$~family is that a trigonometric function has to be calculated every time the kernel is invoked during the simulation. That leads to a computational overload with respect to the evaluation of the cubic spline. More specifically, if we do a large number of calls to the cubic spline kernel and to $W^H_3$~with random generated arguments, the ratio in CPU time is roughly a factor 2.5 favourable to the polynomial kernel. Nonetheless in a real hydrodynamical calculation such factor is diluted by the rest of the computations especially if gravity is present or the numerical algorithm includes complex physics.  As a general rule we propose to implement the kernels $W^H_n$~in the SPH equations leaving the index $n$~free. In normal conditions, i.e. when there are not strong gradients in the physical magnitudes, we should take $n=3$~because for that value the kernel behaves as the well checked cubic spline. In several circumstances, though, it could be wise to turn the value of the exponent $n$~to a different value. It could be taken, for instance,  $n=3$~in all the SPH equations except in that one concerning the conduction transport term given by Eq. (\\ref{eq:eq17}) if a thermal discontinuity, a hot wall for example, is present (as suggested in Section 3.1). When using a second order Runge-Kutta type integrators we may want to use a different value of $n$~during the first and the second (centered) integration steps. Another possibility is to allow each particle to carry its own $n$~value as in the Sedov test described in Section 3.2 where the combination of an adaptive $n(\\mathbf{r},t)$~and $h(\\mathbf{r}, t)$~improved the fit of the density around the peak and led to a better energy conservation. In some cases even the choice of kernels with $n<3$ could be worthwhile to reduce the numerical noise.  Taking into account their easy implementation, smoothness and flexibility, this family of kernels are an alternative to the widely used spline kernels, offering several advantages that may help to improve hydrodynamical simulations which use the SPH technique.  \\noindent {\\bf Acknowledgements}  The authors want to thank the referees for the constructive comments and many suggestions which have contributed to improve the scientific content and general presentation of this manuscript. This work has been supported by the Spanish MCYT grant AYA2005-08013-C03-01."
    },
    "0809/0809.2613_arXiv.txt": {
        "abstract": "{ We analyzed global properties, radial profiles, and 2D maps of the metal abundances and temperature in the cool core cluster of galaxies Hydra~A using a deep $\\sim120$ ks XMM-Newton exposure. The best fit among the available spectral models is provided by a Gaussian distribution of the emission measure ({\\it gdem}). We can accurately determine abundances for 7 elements in the cluster core with EPIC (O, Si, S, Ar, Ca, Fe, Ni) and 3 elements (O, Ne, Fe) with RGS. The {\\it{gdem}} model gives lower Fe abundances than a single-temperature model. Based on this, we explain why simulations show that the best-fit Fe abundance in clusters with intermediate temperatures is overestimated. The abundance profiles for Fe, Si, S, but also O are centrally peaked. Combining the Hydra A results with 5 other clusters for which detailed chemical abundance studies are available, we find a significant decrease in O with radius, while the increase in the O/Fe ratio with radius is small within 0.1~$r_{200}$, where the O abundances can be accurately determined, with $d({\\rm O/Fe})/d({\\rm log}_{10}r/r_{200})=0.25\\pm0.09$. We compare the observed abundance ratios with the mixing of various supernova type Ia and core-collapse yield models in different relative amounts. Producing the estimated O, Si, and S peaks in Hydra~A requires either the amount of metals ejected by stellar winds to be 3--8 times higher than predicted by available models or the initial enrichment by core-collapse supernovae in the protocluster phase not to be as well mixed on large scales as previously thought. The temperature map shows cooler gas extending in arm-like structures towards the north and south. These structures, and especially the northern one, appear to be richer in metals than the ambient medium and spatially correlated with the large-scale radio lobes. With different sets of assumptions, we estimate the mass of cool gas, which was probably uplifted by buoyant bubbles of relativistic plasma produced by the AGN, to $1.6-6.1\\times10^9\\:M_\\odot$, and the energy associated with this uplift to $3.3-12.5\\times10^{58}$ ergs. The best estimate of the mass of Fe uplifted together with the cool gas is $1.7\\times10^7\\:M_\\odot$, 15\\% of the total mass of Fe in the central 0.5\\arcmin\\ region. ",
        "introduction": "Clusters of galaxies provide a unique environment for elemental abundance measurements and for the study of the chemical enrichment history of the Universe, because their large potential wells retain all the metals  produced by the member galaxies. Of particular interest are clusters of galaxies showing a centrally peaked surface brightness distribution and a cool core, whose spectra are often richer in emission lines because of the lower central temperatures. In addition, these ``cool-core'' clusters have been shown to exhibit a central peak in the abundance distribution of several elements, in particular iron \\citep[e.g.][]{DeGrandi01} and other metals produced by type Ia supernovae (SN~Ia), which led to the conclusion that the central excess is most probably due to enrichment by SN~Ia in the central dominant galaxies \\citep[for a review, see][and references therein]{Werner08rev}.  Cooling-core clusters are furthermore in the limelight because of the so-called ``cooling-flow problem'': the high surface brightness in the central peak implies a high density and a short cooling time, however the observed rate of cooling of the central gas is often orders of magnitude below what is expected in the absence of any heat sources \\citep[for a review, see][]{Peterson06}. As a solution, it has been proposed that active galactic nuclei (AGN) in the central dominant galaxies can provide enough energy to the central gas to balance the cooling, since signs of energetic interaction between the AGN radio plasma and the intra-cluster medium (ICM) have been observed in many systems \\citep[for a recent review, see][]{McNamara07}. Recently, the AGN-ICM interaction has also been shown, by theoretical models \\citep{Rebusco06}, hydrodynamic simulations \\citep{Roediger06} and observations \\citep{simionescu2007b} to be a main mechanism for transporting the metals produced in the central galaxy into the ICM.  Hydra A was one of the first cooling-core clusters in which a displacement of X-ray gas in the center by radio lobes from the central AGN was found \\citep{McNamara00} and it is still one of the most dramatic examples of AGN interaction \\citep{David01, Nulsen02, Nulsen05}. In a deeper Chandra observation of the cluster, a sharp X-ray surface brightness edge was found at radii between 4.3 -- 6\\arcmin\\ (200 -- 300 kpc), interpreted as a shock wave caused by an AGN outburst with an estimated total energy of $10^{61}$ erg \\citep{Nulsen05}, which is sufficient to balance radiative cooling for at least several $10^8$ yrs. Further features besides the shock and the inner cavities are a $\\sim 60$ kpc long filament running from the inner cavities outwards, that may be reminiscent of the ``X-ray arms\" in M87 \\citep[e.g.][]{Feigelson87,Boehringer95,Belsole01,Forman06}, and surface brightness depressions in the outer region inside the shock front, which coincide with the outer radio lobes \\citep{Wise07}. The relatively low temperature \\citep[$T_X \\sim 3.1-3.7$ keV,][]{David01} implies X-ray spectra rich in emission lines providing good spectroscopic diagnostics.  The goal of the current work is twofold. Firstly, we aim to study the abundances, abundance ratios, and radial trends for different chemical elements. This can reveal clues about the chemical enrichment history of Hydra A and, by comparison with previous results from deep pointings of other nearby bright clusters, can broaden our overall understanding of the origin and distribution of metals in clusters of galaxies. Secondly, we investigate a two-dimensional metallicity map to determine what influence an AGN outburst as violent as the one in Hydra A can have on the distribution of chemical elements in the ICM.  Throughout the paper, we assume $H_0=70 {\\rm \\:km\\: s^{-1}\\: Mpc^{-1}}$, $\\Omega_\\Lambda=0.73$, and $\\Omega_M=0.27$. At the redshift of Hydra A ($z=0.0538$), $1^{\\prime\\prime}$ corresponds to 1.05 kpc. Unless otherwise stated, the elemental abundances are given with respect to the proto-Solar values of \\citet{lodders2003}, the errors are at the 1$\\sigma$ level, and upper limits at the 2$\\sigma$ level. The recent solar abundance determinations by \\citet{lodders2003} give significantly lower abundances of oxygen, neon and iron than those measured by \\citet{Angr}. Use of these new determinations affects only the units with respect to which we present the elemental abundances in our paper; the actual measured values can be reconstructed by multiplication of the given values with the solar normalizations. For a compilation of these normalizations for different commonly used abundance determinations, see the review of \\citet{Werner08rev}. ",
        "conclusions": "We analyzed a deep $\\sim120$ ks XMM-Newton exposure of the cooling core cluster of galaxies Hydra~A. We extracted spectra from two large regions in the cluster core and in the outskirts to study the global properties of the cluster with the best statistics possible. We also analyzed the RGS spectrum extracted from a 3 by 10\\arcmin\\ region. We find that \\begin{itemize} \\item the shape of the Fe-L complex in the EPIC spectrum does not show strong indications for the presence of cooler gas as observed in other cooling core clusters. However, a multi-temperature model is needed to simultaneously fit both Fe-L and Fe-K lines appropriately. The best available model achieving this is a Gaussian distribution of the emission measure ({\\it gdem}) around the best-fit average temperature. \\item the best-fit ({\\it gdem}) model shows a very broad distribution, with a full-width at half-maximum of 4.2~keV. We also fitted polynomial emission measure distributions which confirm the broad shape of $dY/dT$. The distribution is significantly broader than expected from the substructures visible in the 2D map in the considered region, suggesting the presence of intrinsic multi-temperature structure in each bin of the temperature map. We suggest that this multiphase structure can be due in part to the projection of shocked gas in front of and behind the cluster center, in part to the projection of cooler gas from the cluster outskirts along the line of sight, and in part to dense, unresolved cool gas blobs in the considered extraction region. \\item we can accurately determine abundances for 7 elements in the cluster core with EPIC (O,Si,S,Ar,Ca,Fe,Ni). In the cluster outskirts, only Si, S and Fe abundances are determined with better than 3$\\sigma$ significance and the large errors do not enable us to draw conclusions about possible differences in Si/Fe and S/Fe with respect to the cluster center. \\item the abundances of 3 elements (O,Ne,Fe), one of which is inaccessible with EPIC, are determined with RGS. The O/Fe ratios from EPIC and RGS are consistent. While no multi-temperature structure beyond the {\\it gdem} model can be constrained with EPIC, the RGS fit requires a cool component with a temperature of $0.62\\pm0.04$ keV and a normalization of only 1.2\\% of the hot ambient. This gas is detected with a significance of 6.5$\\sigma$. \\item the Gaussian emission measure distribution model gives lower Fe abundances than a single temperature model, which can explain why simulations show that the best-fit Fe abundance in clusters with intermediate temperatures is over-estimated. \\end{itemize} We also determined temperature and abundance profiles from seven annuli and compared our results with radial profiles of other clusters for which deep observations and detailed chemical enrichment studies are available. We show that \\begin{itemize} \\item the abundance profiles for Fe, Si, S, but also O are centrally peaked in Hydra A. \\item the radial profiles of the O abundance are peaked also in other clusters for which deep data are available. The increase in O/Fe with radius is very small. Combining the Hydra A results with 5 other clusters for which a detailed chemical abundance study has been performed, we find $d{\\rm O}/d({\\rm log}_{10}r/r_{200})=-0.48\\pm0.07$, while the increase in the O/Fe ratio with radius is only less than 3$\\sigma$ significant, $d({\\rm O/Fe})/d({\\rm log}_{10}r/r_{200})=0.25\\pm0.09$. \\item stellar winds and the chemical enrichment by the number of type Ia supernovae needed to produce the Fe peak in Hydra A do not reproduce the estimated O, Si and S peaks. For this, either the amount of metals produced by stellar winds would have to be 3--8 times higher than predicted by available models or $3-8\\times10^8$ \\sncc\\ would be needed in addition to the contribution from stellar mass loss and SN~Ia. Possibly, the initial enrichment by \\sncc\\ in the protocluster phase was not as well mixed on large scales as previously thought, and some peak in the distribution of \\sncc\\ products existed prior to further enrichment with predominantly SN~Ia elements. \\item mainly because of the low Si abundance, Hydra A requires either a WDD3 or W7 SN~Ia model to reproduce the observed relative abundance patterns. Most other clusters lie between the WDD1 and WDD2 models. A 30--40\\% contribution by SN~Ia compared to \\sncc\\ (by number) is needed to reproduce the observed relative abundance patterns at all radii (below 0.1 $r_{200}$) in all the considered clusters. \\item the best-fit Galactic absorption column density, $N_{\\mathrm{H}}$, when fitting the full energy band, is lower than the value determined from \\ion{H}{i} data and has a minimum on the cluster center, which may be interpreted as inverse Compton or soft excess emission. The best-fit $N_{\\mathrm{H}}$ in the 0.35--2 keV band however is in agreement with the \\ion{H}{i} data, dismissing this possibility. \\end{itemize} Finally, we produced 2D temperature and metallicity maps of Hydra A. From these we can conclude that \\begin{itemize} \\item the temperature map shows cooler gas extending in arm-like structures towards the north and south. The cool gas structures, and especially the northern one, appear to be richer in metals than the ambient medium and spatially correlated with the large-scale radio lobes. The northern ``arm'' is associated with a bright 1\\arcmin\\ long filament seen in the Chandra image. \\item based on the geometry of the cool ``arms'' seen in the temperature map and on the best-fit normalization of the 0.62 keV component in RGS, the estimated mass of cool gas, which was probably uplifted by the AGN to create these structures, is $1.6\\times10^9\\:M_\\odot$. The energy needed for this uplift is $3.3\\times10^{58}$ ergs, which is 8 times greater than the energy needed for gas uplift in M87 but nevertheless very small compared to other processes related to AGN-ICM interaction in Hydra~A (large-scale shock and cavities), which require approximately $10^{61}$ ergs. \\item the best estimate of the mass of Fe uplifted together with the cool gas is $1.7\\times10^7\\:M_\\odot$, 15\\% of the total mass of Fe in the central 0.5\\arcmin\\ region. If the average metallicity of the uplifted gas is 2 solar, as in M87, the total mass of uplifted gas (not only the 0.62 keV component) is $6.1\\times10^9\\:M_\\odot$ and the uplift energy is $1.25\\times10^{59}$ ergs. \\item the O/Fe value in the N arm is consistent with the average O/Fe ratio in the inner 3\\arcmin. The transport of metals by the AGN thus presently does not alter the relative abundance patterns. \\end{itemize}"
    },
    "0809/0809.4565_arXiv.txt": {
        "abstract": "We present deep multifrequency observations using the Giant Metrewave Radio Telescope at 153, 244, 610 and 1260 MHz of a field centered on J0916+6348, to search for evidence of fossil radio lobes which could be due to an earlier cycle of episodic activity of the  parent galaxy, as well as halos and relics in clusters of galaxies. We do not find any unambiguous evidence of episodic activity in a list of 374 sources, suggesting that such activity is rare even in relatively deep low-frequency observations. We examine the spectra of all the sources by combining our observations with those from the Westerbork Northern Sky Survey, NRAO VLA Sky Survey and the Faint Images of the Radio Sky at Twenty-centimeters survey. Considering only those which have measurements at a minimum of three different frequencies, we find that almost all sources are consistent with a straight spectrum with a median spectral index,  $\\alpha\\sim$0.8 (S$(\\nu)\\propto\\nu^{-\\alpha}$), which appears steeper than theoretical expectations of the injection spectral index. We identify 14 very steep-spectrum sources with $\\alpha\\geq$1.3. We examine their optical fields and discuss the nature of some of these sources. ",
        "introduction": "Low-frequency radio observations of radio galaxies and quasars provide us with an opportunity to probe episodic activity in active galactic nuclei (AGN) and also constrain the injection spectral indices of the electron energy spectrum. The extended radio emission in radio galaxies and quasars could provide constraints on the duration and duty cycles of episodic activity in AGN via the structural and spectral information of the lobes of emission. The most striking examples of such episodic nuclear activity are the `double-double' radio galaxies (DDRGs) where a young pair of radio lobes is seen closer to the nucleus in addition to the older and more distant lobes of emission (e.g. \\citeauthor*{1996MNRAS.279..257S}~\\citeyear{1996MNRAS.279..257S}; \\citeauthor{1999A+A...348..699L}~\\citeyear{1999A+A...348..699L}; \\citeauthor{2000MNRAS.315..371S}~\\citeyear{2000MNRAS.315..371S}; \\citeauthor*{2006MNRAS.366.1391S}~\\citeyear{2006MNRAS.366.1391S} and references therein). In addition to the DDRGs, diffuse relic radio emission due to an earlier cycle of activity has also been suggested for a number of sources from structural and spectral information. Some of the examples are 3C338 and 3C388 (\\citeauthor*{1983ApJ...271..575B}~\\citeyear{1983ApJ...271..575B}; \\citeauthor{1994ApJ...421L..23R}~\\citeyear{1994ApJ...421L..23R}), Her A \\citep{2003MNRAS.342..399G}, 3C310 (\\citeauthor{1984ApJ...282L..55V}~\\citeyear{1984ApJ...282L..55V}; \\citeauthor*{1986MNRAS.222..753L}~\\citeyear{1986MNRAS.222..753L}), Cen A (\\citeauthor*{1983ApJ...273..128B}~\\citeyear{1983ApJ...273..128B}; \\citeauthor*{1992ApJ...395..444C}~\\citeyear{1992ApJ...395..444C}; \\citeauthor{1993A+A...274.1009J}~\\citeyear{1993A+A...274.1009J}; \\citeauthor{1999MNRAS.307..750M}~\\citeyear{1999MNRAS.307..750M}), and more recently 4C29.30 \\citep{2007MNRAS.378..581J}. \\cite{2000A+A...362...27J} have listed a few cases of small-scale halos associated with the known milliarcsec-scale structures of compact steep spectrum and giga-hertz peaked spectrum sources from interplanetary scintillation observations with the Ooty Radio Telescope at 327 MHz. However, examples of such features appear to be relatively rare, and several searches for relic emission around bright radio sources have not yielded positive results (e.g. \\citeauthor*{1980A+A....89..204R}~\\citeyear{1980A+A....89..204R}; \\citeauthor*{1980A+AS...42..299S}~\\citeyear{1980A+AS...42..299S}; \\citeauthor*{1982ApJ...255L..93P}~\\citeyear{1982ApJ...255L..93P}; \\citeauthor*{1983A+A...125..146K}~\\citeyear{1983A+A...125..146K}; \\citeauthor*{1984IAUS..110....9V}~\\citeyear{1984IAUS..110....9V}; \\citeauthor*{2001AJ....122.2940J}~\\citeyear{2001AJ....122.2940J}).  \\begin{table*} \\begin{center} \\caption{Observation summary and image parameters} \\label{0916p6348_f0:table:obs_sum} \\begin{tabular}{| l | l | l | l | l | l |} \\hline & Center Frequency  & 153 MHz &  244 MHz & 610 MHz & 1260 MHz\\\\ & Date & 2005 Dec 12  & 2005 Nov 26 & 2005 Nov 26 & 2008 Apr 22\\\\ & Working antennas & 29  & 26 & 27 & 27\\\\ General  & Bandwidth & 5.6 MHz & 32 MHz & 32 MHz & 32 MHz\\\\ & Polarization & RR, LL  & LL & RR & RR, LL\\\\ & Visibility integration time (s) & 16.78 & 16.78 & 16.78 & 4.19 \\\\ & Total observation time (hr) & 7.61 & 8.57 & 8.57 & 6.50 \\\\ \\hline & Source & 3C147 & 3C147 & 3C147 & 3C147\\\\ Flux       & Time (hr) & 1.01 & 0.25 & 0.25 & 0.26 \\\\ Calibrator & Flux density (Jy) & 68.29 & 59.16 & 38.27 & 15.58 \\\\ & Scale & Perley-Taylor 99 & Perley-Taylor 99 & Perley-Taylor 99 & Perley-Taylor 99\\\\ \\hline & Source & 3C147 & J0834+5534 & J0834+5534 & J0834+5534\\\\ Phase & Time (hr) & 1.0 & 1.72 & 1.72 & 0.74 \\\\ Calibrator & Flux density (Jy) & 68.29 & 8.60 & 8.22 & 7.39 \\\\ \\hline &Phase center & J0916+6348 & J0916+6348 & J0916+6348 & J0916+6348\\\\ Target & Time (hr) & 6.60 & 6.60 & 6.60 & 5.34 \\\\ Field   & Image size & $6144\\times6144$ & $6144\\times6144$ & $6144\\times6144$ & $6144\\times6144$\\\\ & Pixel size & $4^{\\prime\\prime}\\times4^{\\prime\\prime}$ & $3^{\\prime\\prime}\\times3^{\\prime\\prime}$ & $2^{\\prime\\prime}\\times2^{\\prime\\prime}$ & $1^{\\prime\\prime}\\times1^{\\prime\\prime}$\\\\ & Rms noise$^\\dagger$ $\\mu$Jy beam$^{-1}$ & $\\approx 513$ & $\\approx 173$ & $\\approx 32$ & $\\approx 22$\\\\ & Synthesized beam & $21.48^{\\prime\\prime}\\times18.41^{\\prime\\prime}$ & $12.09^{\\prime\\prime}\\times10.80^{\\prime\\prime}$ & $5.51^{\\prime\\prime}\\times4.59^{\\prime\\prime}$ & $2.85^{\\prime\\prime}\\times2.13^{\\prime\\prime}$\\\\ & PA & 45.8$^\\circ$ & 57.4$^\\circ$ & 47.2$^\\circ$ & -47.1$^\\circ$\\\\ \\hline Primary Beam & Gaussian, FWHM & 173.8$^{\\prime}$ & 117.0$^{\\prime}$ & 42.7$^{\\prime}$ & 26.4$^{\\prime}$\\\\ correction parameter  &  Reference Frequency & 153 MHz & 235 MHz & 610 MHz & 1280 MHz\\\\ \\hline \\end{tabular}  $^\\dagger$ The rms noise given is the median value of all estimates across the primary beam till 20 per cent of the peak response, before primary beam correction.  \\end{center} \\end{table*}  The diffuse relic emission is expected to have a steep radio spectrum due to radiative losses and remain visible for approximately $\\sim$10$^8$ yr (e.g. \\citeauthor*{2000ApJ...543..611O}~\\citeyear{2000ApJ...543..611O}; \\citeauthor*{2002MNRAS.336..649K}~\\citeyear{2002MNRAS.336..649K}; \\citeauthor{2004A+A...427...79J}~\\citeyear{2004A+A...427...79J}). Low-frequency observations with low rms noise values and short spacings sensitive to the diffuse emission could provide better constraints on the frequency of occurrence or detection of such relic emission. We explore this possibility from observations of a field centered on RA: 09$^{\\rm h}$ 16$^{\\rm m}$ 30$^{\\rm s}$, DEC: 63$^{\\circ}$ 48$^{\\prime}$ 00$^{\\prime\\prime}$ at a number of frequencies, namely 150, 244, 610 and 1260 MHz, with the Giant Metrewave Radio Telescope (GMRT). This field was chosen when several interesting radio sources, including a wide-angle tailed source, were noticed while studying the nearby group of galaxies Holmberg 124 \\citep{2005A+A...435..483K}. In order to identify diffuse cluster sources, it is important that the imaged field contains galaxy clusters. In these observations, the clusters Abell 764 at a redshift of 0.166 \\citep{1989ApJS...70....1A} and MaxBCG J139.14754+63.8034 at a redshift of 0.1215 \\citep{2007ApJ...660..239K} are both very close to the pointing centre.  Deep low-frequency observations are also extremely useful for identifying radio halos, relics and core-halos or mini halos associated with clusters of galaxies. While core-halos are less than about 500 kpc in extent and associated with the dominant galaxy in cooling core clusters, halos and relics are much larger in size and not associated with any particular galaxy. Radio halos are usually projected towards the cluster center and have a regular morphology, while relics are seen towards the periphery with a variety of shapes (cf. \\citeauthor{2008A+A...484..327V}~\\citeyear{2008A+A...484..327V} and references therein). While relics are believed to arise due to cluster mergers and/or matter accretion \\citep{1999ApJ...520..529S, 2003ApJ...593..599R, 2006MNRAS.367..113P, 2008A+A...486..347G}, the models for halos range from re-acceleration due to turbulence (\\citeauthor{2003ASPC..301..349B}~\\citeyear{2003ASPC..301..349B}; \\citeauthor{2004JKAS...37..433S}~\\citeyear{2004JKAS...37..433S}; \\citeauthor*{2008SSRv..134..207P}~\\citeyear{2008SSRv..134..207P}; \\citeauthor{2007ApJ...670L...5B}~\\citeyear{2007ApJ...670L...5B} and references therein) to production of relativistic electrons by hadronic collisions (e.g. \\citeauthor*{2000A+A...362..151D}~\\citeyear{2000A+A...362..151D} and references therein).  \\begin{figure*} \\begin{center} \\vbox{ \\hbox{ \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_150_01.pbcimage_000.ps} \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_233_01.pbcimage_000.ps} } \\hbox{ \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_610_01.pbcimage_000.ps} \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_1420_02.pbcimage_000.ps} } } \\caption{GMRT images of an area of 18$\\times$18 arcmin$^2$ of the field J0916+6348 at 150, 244, 610 and 1260 MHz. The resolutions and rms noise values are listed in Table~\\ref{0916p6348_f0:table:obs_sum}. In this figure and in all the images presented here, the contour levels in units of Jy beam$^{-1}$ are represented by mean$+$rms$\\times$(n), where n is the multiplication factor. The negative contours are shown as dashed lines. } \\label{0916p6348_f0:fig:multifreq} \\end{center}  \\begin{center} \\vbox{ \\hbox{ \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_610_01_nvss.image_000.ps} \\includegraphics[angle=0, totalheight=2.8in, viewport=19 212 573 627, clip]{09chi01_610_01_first.image_000.ps} } } \\caption{The NVSS (left) and FIRST (right) images of the same area shown in Fig.~\\ref{0916p6348_f0:fig:multifreq} of the field J0916+6348 at 1400 MHz with angular resolutions of 45 and 5 arcsec respectively. Contour levels: mean$+$rms$\\times$(n) in units of Jy beam$^{-1}$.} \\label{0916p6348_f0:fig:vla} \\end{center} \\end{figure*}  In addition to the search for fossil or relic radio emission of different forms, it is also useful to examine the low-frequency spectra of these weak sources, which could be close to their injection spectral indices, $\\alpha_{\\rm inj}$ (flux density, S$\\propto\\nu^{-\\alpha}$), and compare these with values for stronger sources. Most of our sources are weak and are not catalogued in the well-known low-frequency surveys such as the Cambridge surveys at 151 MHz  (\\citeauthor{2002IAUS..199...21G}~\\citeyear{2002IAUS..199...21G} and references therein) and the VLA Low-frequency Sky Survey (VLSS; \\citeauthor{2007AJ....134.1245C}~\\citeyear{2007AJ....134.1245C}; {\\tt http://lwa.nrl.navy.mil/VLSS}). In the case of the strong sources, $\\alpha_{\\rm inj}$ has been estimated to be 0.5 for Cygnus A. Similar values have been obtained by \\cite{1989A+A...219...63M} for a sample of powerful radio galaxies. From measurements between 38 MHz and 1 GHz, \\cite*{1992MNRAS.257..545L} find a median value of 0.77. For a sample of 3CR sources \\cite{1989MNRAS.239..401L} estimate the hot spot spectral indices, which are similar to $\\alpha_{\\rm inj}$, to have a similar median value of $\\sim$0.8. This is also similar to the estimate of $\\alpha_{\\rm inj}$ for the compact steep spectrum source 3C190 by \\cite*{1997ApJ...479..258K}. For a sample of giant radio sources, using low-frequency GMRT observations and higher frequency observations from the Very Large Array (VLA) as well as data from the literature, \\cite{2008MNRAS.385.1286J} find that $\\alpha_{\\rm inj}$ varies from $\\sim$0.55 to 0.88 with a median value of $\\sim$0.6. The theoretically expected values are somewhat smaller. For a strong, non-relativistic shock in a Newtonian fluid $\\alpha_{\\rm inj}$ = 0.5 (\\citeauthor{1978MNRAS.182..147B}~\\citeyear{1978MNRAS.182..147B},~\\citeyear{1978MNRAS.182..443B}; \\citeauthor{1978ApJ...221L..29B}~\\citeyear{1978ApJ...221L..29B}), while the values range from 0.35 to 0.65 for relativistic shocks in different situations (\\citeauthor{1989LNP...327..247H}~\\citeyear{1989LNP...327..247H}; \\citeauthor*{1987ApJ...315..425K}~\\citeyear{1987ApJ...315..425K}; \\citeauthor*{1981ApJ...248..344D}~\\citeyear{1981ApJ...248..344D}; \\citeauthor*{1982A+A...111..317A}~\\citeyear{1982A+A...111..317A}).  A summary of the observations and the parameters of the images obtained are presented in Table~\\ref{0916p6348_f0:table:obs_sum}, while the details of our observations and analysis are described in Section~\\ref{0916p6348_f0:sec:obs}. The results are presented in Section~\\ref{0916p6348_f0:sec:result}. The discussion and conclusions are summarized in Section~\\ref{0916p6348_f0:sec:discussions}. ",
        "conclusions": "We have presented deep multifrequency GMRT observations at 153, 244, 610 and 1260 MHz of a field centered on J0916+6348.  We do not find any unambiguous evidence of emission from an earlier cycle of activity in a list of 374 sources, suggesting that such activity is rare even in relatively deep low-frequency observations. However, the median size of our objects is expected to be less than about a 100 kpc, which may be partly responsible for the non-detection of fossil lobes from an earlier cycle of AGN activity.  By combining our flux density measurements with those of WENSS, NVSS and FIRST we find a median spectral index of $\\sim$0.8, which is similar to that of the stronger 3CR sources, but appears steeper than theoretical expectations of the injection spectral index.  We identify 14 very steep-spectrum sources with a steep spectral index $\\geq$1.3. We find three of these sources to be identified with an optical galaxy visible in the SDSS, while another three are in the directions of clusters of galaxies. One of these three, GMRT092830+0634748, is a diffuse extended source and could be a cluster relic."
    },
    "0809/0809.3345_arXiv.txt": {
        "abstract": "The  stability problem of   MHD Taylor-Couette flows with toroidal  magnetic fields  is considered  in dependence on the  magnetic Prandtl number. Only the most  uniform (but not current-free) field  with  $B_{\\rm in}=B_{\\rm out}$  has been considered.  For high  enough Hartmann numbers the toroidal field is always unstable. Rigid rotation, however, stabilizes the magnetic (kink-)instability.  The axial current which drives the instability is reduced by the  electromotive force induced by the instability itself. Numerical simulations are presented to probe this effect as a possibility  to measure the turbulent conductivity  in a laboratory. It is shown numerically that  in a  sodium  experiment (without rotation) an eddy diffusivity 4 times the molecular diffusivity  appears resulting in a potential difference of $\\sim 34$ mV/m. If the cylinders are rotating then also the eddy viscosity can be measured. Nonlinear simulations of the instability lead to a turbulent magnetic Prandtl number of 2.1 for a molecular magnetic Prandtl number of $0.01$. The trend goes to higher values for smaller Pm. ",
        "introduction": "Strong enough toroidal fields that are not current-free  become unstable  due to  the Tayler instability (TI,  \\cite{T57,V72,T73}).  Because the source of the energy is  the electric current, these (mainly nonaxisymmetric) instabilities can exist even without any rotation.  On the other hand, it is becoming increasingly clear that the stability of differential rotation under the presence of  magnetic fields is one of the key problems in MHD astrophysics. There, however,  is no laboratory experiment so far and even the related numerical simulations of TI are very  rare \\cite{BRW06,GRE08}.  We shall demonstrate here  how the TI  interacts with differential rotation, and how it is possible to  verify  the main results in laboratory experiments.  In particular, the theoretical results are used  to propose    experiments for measuring the turbulent diffusivity   via a TI-induced reduction of the electromotive force. Such  experiments will be of high relevance as the knowledge of the magnetic turbulent  diffusivity is basic for many applications in fluid dynamics. Very often  we have only  limited informations about the magnetic diffusivity. In the laboratory only a very small number of  experiments have been done (see \\cite{RB01,Frick07}). The same is true for the eddy viscosity which can be measured with the same experimental device so that finally the turbulent magnetic Prandtl number becomes known for one and the same instability.  Consider a Taylor-Couette (TC) flow with   $\\vec{U}$ as the velocity, $\\vec{B}$ the magnetic field,  $\\nu$ the microscopic kinematic viscosity and $\\eta$ the microscopic magnetic diffusivity. The  basic state in cylindrical geometry is $U_R=U_z=B_R=B_z=0$ and \\beg U_\\phi=R\\Omega=a_\\Omega R+\\frac{b_\\Omega}{R},  \\q B_\\phi=a_B R+\\frac{b_B}{R}. \\label{basic} \\ende Let \\beg \\hat\\eta=\\frac{R_{\\rm{in}}}{R_{\\rm{out}}}, \\; \\; \\; \\mu_\\Omega=\\frac{\\Omega_{\\rm{out}}}{\\Omega_{\\rm{in}}},  \\; \\; \\; \\mu_B=\\frac{B_{\\rm{out}}}{B_{\\rm{in}}}. \\ende $R_{\\rm{in}}$ and $R_{\\rm{out}}$ are the radii of the inner and outer cylinders, $\\Omega_{\\rm{in}}$ and $\\Omega_{\\rm{out}}$  their rotation rates  and $B_{\\rm{in}}$ and $B_{\\rm{out}}$ are the azimuthal magnetic fields at the inner and outer cylinders. In particular, a field of the form $b_B/R$ is generated by  an axial current only through the inner region $R<R_{\\rm{in}}$, whereas a field of the form $a_B R$ is generated by  a uniform axial current through the entire region $R<R_{\\rm{out}}$ including the fluid.  The magnetic Prandtl number $\\Pm$, the Reynolds number $\\Re$ and the Hartmann number $\\Ha$, \\beg {\\rm Pm}=\\frac{\\nu}{\\eta}, \\quad {\\rm{Re}}=\\frac{\\Omega_{\\rm{in}} R_0^2}{\\nu}, \\quad {\\rm{Ha}}=\\frac{B_{\\rm{in}} R_0}{\\sqrt{\\mu_0 \\rho \\nu \\eta}}, \\ende are the basic parameters of the problem where $R_0=\\sqrt{D R_{\\rm in}}$ is the unit of length with $D=R_{\\rm{out}}-R_{\\rm{in}}$. For the velocity the boundary conditions are assumed as always no-slip ($\\vec{u}=0$). For conducting walls the radial component of the field and the tangential components of the current must vanish so that $ {\\rm d}b_\\phi/{\\rm d}R + b_\\phi/R = b_R = 0 $ at both $R_{\\rm{in}}$ and $R_{\\rm{out}}$. Here  $\\vec{u}$ and $\\vec{b}$ are the fluctuating components of flow and field.  While the linear stability code works well with small $\\Pm$, the minimum microscopic $\\Pm$ which can be handled with our nonlinear code is (only) $10^{-2}$. The numerics  cannot deal with the very small magnetic Prandtl numbers of liquid metals used in the laboratory ($\\rm Pm \\lsim 10^{-5}$). Some of our results can only be obtained  by extrapolation methods. ",
        "conclusions": ""
    },
    "0809/0809.4279_arXiv.txt": {
        "abstract": "A soft ultraviolet radiation field, $10.2$ eV$<h\\nu<13.6$ eV, that permeates neutral intergalactic gas during the Epoch of Reionization (EoR) excites the $2p$ (directly) and $2s$ (indirectly) states of atomic hydrogen. Because the $2s$ state is metastable, the lifetime of atoms in this level is relatively long,  which may cause the $2s$ state to be overpopulated relative to the $2p$ state.  It has recently been proposed that for this reason, neutral intergalactic atomic hydrogen gas may be detected in absorption in its 3-cm fine-structure line ($2s_{1/2}\\rightarrow 2p_{3/2}$) against the Cosmic Microwave Background out to very high redshifts.  In particular, the optical depth in the fine-structure line through neutral intergalactic gas surrounding bright quasars during the EoR may reach $\\tau_{FS}\\sim 10^{-5}$. The resulting surface brightness temperature of tens of $\\mu$ K (in absorption) may be detectable with existing radio telescopes. Motivated by this exciting proposal, we perform a detailed analysis of the transfer of Ly$\\beta,\\gamma,\\delta,...$ radiation, and re-analyze the detectability of the fine-structure line in neutral intergalactic gas surrounding high-redshift quasars. We find that proper radiative transfer modeling causes the fine-structure absorption signature to be reduced tremendously to $\\tau_{FS}\\lsim 10^{-10}$. We therefore conclude that neutral intergalactic gas during the EoR cannot reveal its presence in the 3-cm fine-structure line to existing radio telescopes. ",
        "introduction": "\\label{sec:intro}  It has recently been proposed that it may be possible to observe neutral atomic hydrogen gas during the Epoch of Reionization (EoR), in absorption against the Cosmic Microwave Background (CMB) in its 3-cm fine-structure line \\citep[$2s_{1/2}\\rightarrow 2p_{3/2}$,][hereafter S07]{Sethi07}.  The possible detectability of the 3-cm line derives from the fact that the 2s state of atomic hydrogen is metastable \\citep{BT40}. That is, the average lifetime of an undisturbed atom in the 2s-state is $t_{2s}=1/A_{2s1s}\\sim 0.1$ s, after which it decays into the 1s state by emitting 2 photons \\citep{Spitzer51}. For comparison, the mean life-time of atoms in the 2p state is $t_{2p}=1/A_{2p1s}\\sim 10^{-9}$ s. Hence, when atoms are excited into their 2s and 2p states at comparable rates, the 2s level is expected to be highly overpopulated relative to the 2p level. This may result in a non-negligible optical depth in the $2s_{1/2}-2p_{3/2}$ fine-structure line \\citep[see e.g.][]{Wild52}.  The detectability of this line has been discussed in the context of HII regions in the local Universe \\citep[e.g.][for a more detailed discussion]{Dennison05}.  In local HII regions, the 2s level is populated mainly following recombination, while both recombination and photoexcitation by Ly$\\alpha$ photons populate the 2p state. The Ly$\\alpha$ excitation rate of the 2p-state is limited by dust, and the 2s-state is typically heavily overpopulated in HII regions which result in 3-cm optical depths of order $\\tau_{FS} \\sim 10^{-3}$ (Ershov 1987; Dennison et al, 2005, where the absorption is measured relative to the free-free continuum that is produced in the HII region). S07 have argued that during the EoR, bright quasars that are luminous in the restframe soft-UV (10.2 eV  $<E_{\\gamma}<$ 13.6 eV) may indirectly photoexcite the 2s state of neutral hydrogen gas in the low density intergalactic medium to levels such that $\\tau_{FS}\\sim 10^{-5}$ for CMB photons passing through these regions\\footnote{It may be surprising that the sparsely populated $n=2$ levels of atomic hydrogen can produce a detectable opacity. One can estimate $\\tau_{FS}$ by comparing it to the Gunn-Peterson optical depth $\\tau_{GP}$: $\\frac{\\tau_{FS}}{\\tau_{GP}}=\\frac{n_{2s}}{n_{1s}}\\frac{A_{FS}\\lambda_{FS}^3}{A_{2p1s}\\lambda_{2p1s}^3}$ $\\sim 20\\frac{n_{2s}}{n_{1s}}$. Since $\\tau_{\\rm GP}\\sim$ a few times $10^5$ for a fully neutral IGM at $z=6$, $\\tau_{FS}=10^{-5}$ only requires that $\\frac{n_{2s}}{n_{1s}}\\sim 10^{-12}$. It will be shown in \\S~\\ref{sec:sim} that these levels can be reached in regions around high-redshift quasars.}. Absorption of these CMB photons would result in an absorption feature with a brightness temperature of a few tens of $\\mu$K at $1.4$ Ghz, which may be detectable with existing dish antennas such as the {\\it Green Bank Telescope } (GBT), {\\it Parkes} and the most compact configuration of the {\\it Australia Telescope Compact Array} \\citep[ATCA, e.g.][]{Ca08}. This result is important, because it implies that neutral intergalactic atomic hydrogen gas that exists during  EoR may reveal itself in a line that is different (and possibly easier to detect) than the well studied 21-cm hyperfine transition \\citep[e.g.][]{SR90,LZ04,Furlanettorev}. \\begin{figure} \\vbox{\\centerline{\\epsfig{file=fig1.eps,angle=0,width=8.5truecm}}} \\caption[]{Schematic diagram of the energy levels of a hydrogen atom. On the left side, we show the first 3 energy levels, and their splitting into levels with different angular momentum. The notation for each level is $nL$, where $n$ is the principle quantum number, and $L$ denotes the orbital angular momentum number ($L=S$ corresponds to $l=0$, $L=P$ corresponds to $l=1$, $L=D$ corresponds to $l=2$). The {\\it dotted lines} show the allowed transitions through the emission or absorption of a single photon. In a 3-level atom, the selection rules only allow the 2s state to be populated by absorbing a Ly$\\beta$ photon, and by subsequently emitting an H$\\alpha$ photon. On the right, the fine+hyperfine-structure splitting of the  $2s$ and $2p$ levels is shown in more detail. The Lamb-Rutherford  shift causes the $2s_{1/2}$ level to lie above the $2p_{1/2}$-level by $\\Delta E/h_p\\sim1.1$ Ghz \\citep[e.g.][]{Wild52}. The 3 allowed transitions between the $2s_{1/2}$ and  $2p_{3/2}$ levels are indicated with {\\it dashed lines}, and the frequency (in MHz) of each transition is indicated \\citep[also see][]{Dennison05}. The relative strength of each transition 9852:9876:10030 is 1:5:2 \\citep[][]{Wild52}.} \\label{fig:diagram} \\end{figure} Motivated by this exciting result, and by observational efforts to detect this transition, we re-analyze the detectability of the fine-structure line during the EoR. We perform more detailed modeling of the transfer of Ly$\\beta$ radiation than originally discussed by S07. We will show that (unfortunately) proper modeling of Ly$\\beta$ (and Ly$\\gamma,\\delta,...$) radiative transfer radically lowers the value of $\\tau_{FS}$ to levels that cannot be probed with existing radio telescopes.  The outline of this paper is as follows. We review the dominant excitation mechanisms of the 2s and 2p levels in \\S\\ref{sec:mech}. In \\S\\ref{sec:ex}, we compute the effective excitation rate of the 2s-state during the EoR, by properly accounting for radiative transfer, and discuss the detectability of the 3-cm fine-structure line in \\S\\ref{sec:detection}. The validity of our model assumptions, and the implications of this work are discussed in \\S\\ref{sec:discuss}. Finally, we present the conclusions in \\S\\ref{sec:conc}. The parameters for the background cosmology used throughout are $(\\Omega_m,\\Omega_{\\Lambda},\\Omega_b,h)=(0.27,0.73,0.042,0.70)$ \\citep{Komatsu08}. ",
        "conclusions": "\\label{sec:conc} It has recently been proposed that neutral intergalactic atomic hydrogen gas may be detected in absorption in its 3-cm fine-structure line ($2s_{1/2}\\rightarrow 2p_{3/2}$) against the Cosmic Microwave Background (CMB) out to very high redshifts. In particular, bright quasars that are luminous in the restframe soft-UV (10.2 eV $<E_{\\gamma}<$ 13.6 eV) may indirectly photoexcite the 2s state of neutral hydrogen gas in the intergalactic medium (IGM) during the Epoch of Reionization (EoR) to levels such that $\\tau_{FS}\\sim 10^{-5}$ for CMB photons passing through these regions \\citep{Sethi07}. The resulting brightness temperature of $\\Delta T_b\\sim$ tens of $\\mu$K could be detected with existing radio telescopes \\citep[e.g.][]{Ca08}.  Motivated by this proposal, and by observational efforts to detect this transition, we have performed a detailed analysis of the transfer of Ly$\\beta,\\gamma,\\delta,...$ radiation, and have re-analyzed the detectability of the fine-structure line in neutral intergalactic gas surrounding bright quasars during the EoR. We have found that properly considering radiative transfer effects reduces the signal significantly compared to optically thin estimates, which radically complicates the detectability of the 3-cm fine structure transition.  The main reason for this negative result is the large opacity of the (partially) neutral IGM to Lyman series photons. In such a medium, Ly$\\alpha$ scattering proceeds completely different than Ly$\\beta,\\gamma,\\delta,...$ scattering: Ly$\\alpha$ photons scatter $\\tau_{\\rm GP}\\sim10^5-10^6$ times.   However, higher Lyman series photons have a finite probability for being converted into Balmer, Paschen, .., etc photons at each scattering event, which subsequently propagate to the observer unobstructed. These higher Lyman series photons therefore scatter on average only $5-8$ times (\\S\\ref{sec:result}) before being destroyed \\citep{PF06}. Furthermore, only one of these $5-8$ scattering events - defined as an {\\it effective scattering event}- indirectly excites atoms into their 2s-state (e.g. Ly$\\beta$ scattering corresponds to a sequence $1s\\rightarrow 3p\\rightarrow 1s$, and does not affect the level population of the $2s$-state). For these reasons, we found that the effective Ly$\\beta$ scattering rate was calculated reasonably well if one simply suppresses the incoming Ly$\\beta$ flux by a factor of e$^{-\\tau}$, in which $\\tau$ is the (frequency dependent) optical depth of the IGM to Ly$\\beta$ radiation (see Fig~\\ref{fig:lybscatrate}). The same applies to Ly$\\gamma,\\delta,...$ radiation.  Because of the reduced effective scattering rates, we found a substantial lower column of neutral atoms in the 2s-state. This, in combination with the fact that the profile of the fine-structure absorption cross section is dominated by the damping wings (\\S~\\ref{sec:detection}), results in $\\tau_{FS}\\sim 10^{-11}-10^{-10}$. This is 5-6 orders of magnitude smaller than what was originally found by \\citet{Sethi07}, and we conclude that the 3-cm fine-structure absorption line from neutral intergalactic gas surrounding high-redshift quasars is presently undetectable.  {\\bf Acknowledgments} This research was supported by Harvard University Funds. JRP is supported by NASA through Hubble Fellowship grant HST-HF-01211.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. SMO is supported by NSF/AST0457585."
    },
    "0809/0809.1389_arXiv.txt": {
        "abstract": "Gamma ray bursts (GRBs) have recently attracted much attention as a possible way to extend the Hubble diagram to very high redshift. However, the large scatter in their intrinsic properties prevents directly using them as distance indicator so that the hunt is open for a relation involving an observable property to standardize GRBs in the same way as the Phillips law makes it possible to use Type Ia Supernovae (SNeIa) as standardizable candles. We use here the data on the X\\,-\\,ray decay curve and spectral index of a sample of GRBs observed with the Swift satellite. These data are used as input to a Bayesian statistical analysis looking for a correlation between the X\\,-\\,ray luminosity $L_X(T_a)$ and the time constant $T_a$ of the afterglow curve. We find a linear relation between $\\log{[L_X(T_a)]}$ and $\\log{[T_a/(1+z)]}$ with an intrinsic scatter $\\sigma_{int} = 0.33$ comparable to previously reported relations. Remarkably, both the slope and the intrinsic scatter are almost independent on the matter density $\\Omega_M$ and the constant equation of state $w$ of the dark energy component thus suggesting that the circularity problem is alleviated for the $L_X - T_a$ relation. ",
        "introduction": "The high fluence values (from $10^{-7}$ to $10^{-5} \\ {\\rm  erg/cm^2}$) and the enormous isotropic energy emitted ($\\simeq 10^{50} - 10^{54} {\\rm erg}$) at the peak in a single short pulse make Gamma Ray Bursts (hereafter GRBs) the most violent and energetic astrophysical phenomena. Notwithstanding the variety of their different peculiarities, some common features may be identified looking at their light curves. Although GRBs have been traditionally classified as {\\it short} and {\\it long} depending on $T_{90}$ being smaller or larger than $2 \\ {\\rm s}$ (with $T_{90}$ the time over which from $5\\%$ to $95\\%$ of the prompt emission is released), a recent analysis by Donaghy et al. \\shortcite{donaghy} has shown that this criterion has to be revised. Indeed, the existence of an intermediated class of GRBs have also been studied \\cite{nb06,Bernardini07}. As a result, the long GRBs are now further classified as {\\it normal} and {\\it low luminosity} with the latter ones probably associated with Supernovae \\cite{pian06,Dainotti}.  Notwitstanding this classification, two phases are clearly visible in the GRB lightcurve, namely the prompt emission, where most of the energy is released in the $\\gamma$\\,-\\,rays in only tens of seconds, and an afterglow lasting many hours after the initial bursts. Early observations in the X\\,-\\,rays typically started several hours after the prompt emission so that only the late phase of the light curve could be characterized. It was then found that a phenomenological power\\,-\\,law, $f(t, \\nu) \\propto t^{-\\alpha} \\nu^{-\\beta}$ with $(\\alpha, \\beta) \\simeq (-1.4, 0.9)$, provided a reasonable fit to the observed data \\cite{Piro01}. However, the launch of the {\\it Swift} satellite, whose aim is also to observe GRBs X\\,-\\,ray ($0.2 - 10 {\\rm keV}$) and optical ($1700 - 6500$ {\\AA}) afterglows starting few seconds after the trigger, revealed a more complex behaviour. The soft X\\,-\\,ray light curves must indeed be divided in two different classes \\cite{Chinca05} according to the steep or mild initial decay. Most of the observed GRB afterglows belong to the first group, showing what has been called a {\\it canonical} behavior \\cite{nousek06} described by a broken power\\,-\\,law. After the initial steep decay (with slope $3 \\leq \\alpha_1 \\leq 5$), the light curve shows a shallow decay ($0.5\\leq \\alpha_2 \\leq 1$) followed by a somewhat steeper decay ($1 \\leq \\alpha_3 \\leq 1.5$) beyond $2 {\\times} 10^4 \\ {\\rm s}$. These power\\,-\\,law segments are separated by two corresponding break times with $t_{b1} \\leq 500 \\ {\\rm s}$ and $10^{3} \\ {\\rm s} \\leq t_{b2} \\leq 10^{4} \\ {\\rm s}$. A new systematic study using GRBs observed with XRT reveals a still more complex behavior with different power\\,-\\,law slopes and break times \\cite{OB06,Sak07}. A significant step forward has been represented by the analysis of the X\\,-\\,ray afterglow curves of the full sample of {\\it Swift} GRBs showing that all of them may be fitted by the same analytical expression \\cite{W07}.  Finding out a universal feature for GRBs is the first important step towards their use as distance indicator. To this aim, one has indeed to look for a universal relation linking observable GRBs properties so that their intrinsic luminosity may be estimated from directly measurable quantities. Previous attempts along this road are represented by the $E_iso - E{peak}$ \\cite{Amati06}, $E_\\gamma - E_{peak}$ \\cite{G04,Ghirlanda06}, $L -E_{peak}$ \\cite{S03}, $L - \\tau_{lag}$ \\cite{N00}, $L - V$ \\cite{FRR00,R01}, $L - \\tau_{RT}$. Moreover, three\\,-\\,parameter relations have also been proposed such as, e.g., the $E_iso - E_p - t_b$ \\cite{liza05} and that proposed by Firmani and collaborators (Firmani et al. 2005, 2006). On the other hand, some attempts have also been made to compare these empirical correlations with the model dependent ones \\cite{na06,Guida08}. The above quoted two parameters correlations have then been used by Schaefer (2007, hereafter S07) to construct the first reliable GRBs Hubble diagram extending up to $z \\simeq 6$ opening the way towards the use of GRBs as cosmological probes (see, e.g., Capozziello \\& Izzo 2008 and refs. therein).  In this letter, we present a possible alternative route towards standardizing GRBs as distance indicator. To this aim, we use the data in Willingale et al. (2007) to look for a possible correlation between the X\\,-\\,ray luminosity at the break time $T_a$ and the $T_a$ itself. The data used are presented in Sect. 2, while Sect. 3 deals with the statistical tools and the results. Conclusions are summarized in Sect. 4. ",
        "conclusions": "The high Spearman correlation coefficient, the low value of the fit residuals and the modest intrinsic scatter renders the $L_X - T_a$ relation presented above a new valid tool to standardize GRBs. It is worth stressing that $L_X - T_a$ needs only two parameters and one of them is directly inferred form the observations minimizing the effects of the systematics errors. Furthermore the redshift range covered is large extending from $0.54$ $(0.125)$ up to 6.6 for the selected (full) sample far beyond the maximum redshift affordable with Type Ia SNe ($z \\approx 1.7$). Should this correlation be confirmed by future higher quality data, one could then combine it with the other relations yet available in literature to work out a GRBs Hubble diagram deep into the matter dominated era thus representing an outstanding cosmological test.  To this end, it is worth comparing the $L_X - T_a$ relation with other ones quoted in literature. When performing such a comparison, however, one should take into account the differences in the cosmological model adopted and the fitting method used. In particular, the choice of how the best fit parameters are estimated may have an important impact on the estimate of the intrinsic scatter with the usual $\\chi^2$ fitting leading to an underestimate of $\\sigma_{int}$. On the other hand, changing $\\Omega_M$ in the framework of the flat $\\Lambda$CDM scenario have a profound impact on $\\sigma_{int}$ with higher $\\Omega_M$ giving rise to lower $\\sigma_{int}$ values \\cite{BP08}. In order to account for both these issues, one should therefore test all the above correlations using the same statistical tools and cosmological model, a task we will address elsewhere.  As is well known, the paucity of local (i.e., $z \\le 0.1$) GRBs represents a serious problem for any attempt to standardize GRBs since it is very difficult to directly calibrate any relation. This problem may be partly overcome by fitting the correlation in a subsample of GRBs lying at similar redshift \\cite{Amati08}. However, as a general rule, in order to evaluate the GRB luminosity, a cosmological model has to be adopted thus leading to the circularity problem. Although addressing this problem in detail will be the subject of a forthcoming work, we have here investigated what is the effect of changing the cosmological model by using our maximum likelihood estimator to determine the parameters $(a, b, \\sigma_{int})$ as function of $\\Omega_M$ in a flat $\\Lambda$CDM model\\footnote{To be precise, we let $\\Omega_M$ running from $0.2$ to $1$ and adjust $h$ so that $\\Omega_M h^2$ is fixed to the same value adopted above.}. We find the remarkable result that both the best fit parameters $(b, \\sigma_{int})$ and the rms residual $\\delta_{rms}$ are almost insensitive to the value of $\\Omega_M$. Indeed, $b$ runs from $b \\simeq -0.590$ to $b \\simeq -0.565$, while $\\sigma_{int}$ increases from $\\sigma_{int} \\simeq 0.335$ to $\\sigma_{int} \\simeq 0.340$ for $\\Omega_M$ going from 0.2 to 1.0. As a further test, we generalize the $\\Lambda$CDM model varying not only the matter density parameter $\\Omega_M$, but also the equation of state $w$ of the dark energy component (with $w = -1$ for the $\\Lambda$CDM model). For $-1.3 \\le w \\le -0.7$, neither $b$ nor $\\sigma_{int}$ significantly change confirming the qualitative results obtained for the $\\Lambda$CDM scenario. Although a more detailed analysis is needed, we therefore argue that the circularity problem is alleviated by the use of our $L_X - T_a$ relation.  The encouraging results discussed above are serious arguments in favour of the $L_X - T_a$ relation as a further tool towards the standardization of GRBs as distance indicator. Should these first evidences be furtherly enforced by more data, the combined use of full set of GRBs correlations discovered insofar could opened the road towards making GRBs the high redshift analog of SNeIa as cosmological probes in the not too distant future. \\\\  {\\it Acknowledgements.} We warmly thank R. Willingale for help with the data and prompt answers to our questions and an anonymous referee for his/her valuable comments."
    },
    "0809/0809.0990.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) %{We searched and characterized new X-ray young stars in an unexplored region to %the west of the centre of the $\\sigma$~Orionis cluster ($\\sim 3$\\,Ma, %$d \\sim 385$\\,pc).} {The objective of this study is to determine the general X-ray properties of the young stars in the external regions of the $\\sigma$~Orionis cluster ($\\tau \\sim 3$\\,Ma, $d \\sim 385$\\,pc) and yield constraints on the X-ray emission of brown dwarfs.} % methods heading (mandatory) {We carried out a careful analysis of public data taken in an unexplored region to the west of the centre of the cluster with the three EPIC cameras onboard the \\textit{XMM-Newton} mission. We looked for new X-ray young stars among the 41 identified X-ray sources in the area with maximum likelihood parameters $L >$ 15 by cross-correlation with the USNO-B1, DENIS, and 2MASS databases.} % results heading (mandatory) {Based on colour-colour, colour-magnitude, and hardness ratio diagrams, and previous spectroscopic, astrometric, and infrared-flux excess information, we classified the optical/near-infrared counterparts of the X-ray sources into: young stars (15), field stars (4), galaxies (19), and sources of unknown nature (3). Most of the X-ray detections, including those of nine young stars, are new. We derived the X-ray properties (e.g. temperatures, metallicities, column densities) of the twelve young stars with the largest signal-to-noise ratios. The X-ray parameters determined here are in well agreement with those found in the cluster centre, where the stellar density is higher. % There is no relation between infrared excess and column density from X-ray measurement in our data. % We detected flaring events in two young stars of the sample. One of them showed a very large ($\\sim 30$) relative increase in flux. Both stars showed high coronal temperatures during the observation. Finally, we determined upper limits to the flux of the young stars and bright brown dwarfs not detected by our searching algorithm.} % conclusions heading (optional), leave it empty if necessary {} % ",
        "introduction": "\\label{section.introduction}  %______________________________________________ Figure \\begin{figure*} \\centering %\\includegraphics[width=0.49\\textwidth]{xso_ROSAT_widefield.ps} \\includegraphics[width=0.48\\textwidth]{aa10475f1a.ps} %\\includegraphics[width=0.49\\textwidth]{xso_ROSAT_narrowfield.ps} \\includegraphics[width=0.48\\textwidth]{aa10475f1b.ps} \\caption{Position Sensitive Proportional Counters (PSPC) 2.0\\,deg cut-off, intensity-scaled, summed pointed observations with {\\em ROSAT} (0.1--2.4\\,keV) centred on the Trapezium-like system $\\sigma$~Ori. North is up and east is left. {\\em Left:} the Orion Belt, to the centre and north, and the Orion Nebula Cluster, to the south ($\\sim$7$\\times$7\\,deg$^2$). The brightest Orion complex stars in the area and the field white dwarf \\object{GD~257} are indicated and labelled. {\\em Right:} zoom of the inner part ($\\sim$2$\\times$2\\,deg$^2$). %The radius of the circle centred on $\\sigma$~Ori is $\\sim$30\\,arcmin %\\citep[roughly the cluster radius;][]{cab08a}. Circles represent the two regions observed by \\textit{XMM-Newton}: the first is centred on $\\sigma$~Ori and was analyzed by \\citet{fra06};  the second is centred on S Ori 55 and is analyzed in this work. Note the three X-ray sources of approximately the same brightness, aligned in the north-south direction, at about 15\\,arcmin to the west of $\\sigma$~Ori. Colour versions of all our figures are available in the electronic publication. } \\label{rosat} % ... %{\\bf To mark the {\\em XMM-Newton} f.o.v. of the observations by Franciosini %et~al. (2006) and by us in the right window?}} \\end{figure*} %  The \\object{$\\sigma$~Orionis} cluster ($\\tau \\sim 3$\\,Ma, $d \\sim$ 385\\,pc) is a ``well-equipped laboratory'' for studying the formation, evolution, and astrophysical properties of high-mass, solar-like, and low-mass stars, brown dwarfs, and planetary-mass objects below the deuterium burning mass limit \\citep{wal74,gro82,zap00,bej01,ken05,cab07}. %(Walborn 1974; Groote \\& Hunger 1982; Zapatero Osorio et~al. 2000; B\\'ejar %et~al. 2001; Kenyon et~al. 2005; Caballero et~al. 2007). The cluster takes the name from the Trapezium-like system $\\sigma$~Ori in its centre (we use the abridged name, $\\sigma$~Ori, for the --at least-- sextuple system, and the full name, $\\sigma$~Orionis, for the eponymous cluster). In spite of its early discovery \\citep{gar67}, it was not until the pioneer works by \\citet{wol96} and \\citet{wal97}, who detected an over-density of X-ray emitters and pre-main sequence stars (see Fig.~\\ref{rosat}), when the $\\sigma$~Orionis cluster started taking importance. Afterwards, the detection of Herbig-Haro objects \\citep{rei98}, brown dwarfs \\citep{bej99}, the least massive object directly detected out of the Solar System \\citep[\\object{S\\,Ori~70};][]{zap02b}, and planetary-mass objects with discs \\citep{zap07} have transformed $\\sigma$~Orionis in a cornerstone for the study of stellar and substellar populations in very young open clusters. See \\citet{cab07a,cab08c} %Caballero (2007a, 2008b) for comprehensive $\\sigma$~Orionis data~compilations.  Regarding high energies in the cluster, the soft X-ray emission (0.2--3.5\\,keV) of \\object{$\\sigma$~Ori~A}F--B %(O9.5V + B0.0V + B0.5V) was first %\\citep[O9.5V + B0.0V + B0.5V;][D. M. Peterson et al. in prep.]{cab08b} (O9.5V + B0.0V + B0.5V -- Caballero 2008b; D. M. Peterson et al. in prep.) was first revealed by the {\\em Einstein Observatory} \\citep{chl89}. %(Chlebowski, Harnden \\& Sciortino 1989). The X-ray emission in such an early-type star is thought to come from shock-heated gas in the radiatively driven stellar winds. \\citet{ber94} %Bergh\\\"ofer \\& Schmitt (1994) found no X-ray variability during a long-term monitoring with the {\\em R\\\"ontgensatellit} ({\\em ROSAT}) of the triple system (unresolved in their observations), which showed the stability of their winds. Later, \\citet{san04} and \\citet{wal07} %Sanz-Forcada, Franciosini \\& Pallavicini (2004) and Waldron \\& Cassinelli (2007) presented high energy resolution spectra of $\\sigma$~Ori~AF--B obtained with the {\\em X-ray Multi-Mirror Mission - Newton} ({\\em XMM-Newton}) and the {\\em Chandra X-ray Observatory}, respectively. Very recently, \\citet{ski08}, %Skinner et~al. (2008), from data obtained with the High Energy Transmission Grating Spectrometer onboard {\\em Chandra}, considered possible alternatives to the radiative wind shock, including magnetically-confined and colliding wind shocks.  %Sanz-Forcada et~al. (2004) \\citet{san04} reported a strong X-ray flare on the magnetic B2Vp star \\object{$\\sigma$~Ori~E}, the second brightest star in the cluster, confirming the previous discovery of a flare on the star by \\citet{gro04} with {\\em ROSAT} data. %Groote \\& Schmitt (2003) with {\\em ROSAT} data. Caballero et~al. (in~prep.) have also detected another flare using alternative {\\em ROSAT} data. However, X-ray flares are not expected in the wind scenario of early-type stars. The low-mass companion hypothesis supported by \\citet{san04} %Sanz-Forcada et~al. (2004) and discussed in the literature since the early 1970s \\citep[see][]{lan78} %(see Landstreet \\& Borra [1978] and references therein) has been lately confirmed by \\citet{bou08} %Bouy et~al. (2008, A\\&A, submitted)} using a multi-conjugate adaptive optics system ($\\sigma$~Ori~E and its fainter companion are separated by only 0.330\\,arcsec). %Bergh\\\"ofer \\& Schmitt (1994): A long-term X-ray variability study of the %O-type stars $\\sigma$~Orionis and $\\zeta$~Orionis. %Bergh\\\"ofer \\& Schmitt (1994): ROSAT X-ray light curves of early-type stars. %Bergh\\\"ofer, Schmitt \\& Cassinelli (1996): The ROSAT all-sky survey catalogue %of optically bright OB-type stars [only AFB].  In contrast to early-type stars, the origin of the X-ray emission in %(fast-rotating, almost fully-convective) low-mass stars resides on the %extremely high temperatures of their coronae. %can reach. The first detailed analysis of young, low-mass X-ray stars in the $\\sigma$~Orionis cluster was carried out by \\citet{wol96} %Wolk (1996) using {\\em ROSAT}. His work was surpassed, however, by the exhaustive analysis of the full EPIC/{\\em XMM-Newton} field centred on $\\sigma$~Ori~AF--B by \\citet{fra06}. %Franciosini, Pallavicini \\&  Sanz-Forcada (2006). %Pallavicni, Franciosini \\&  Sanz-Forcada (2004). They detected 175 X-ray sources, 88 of which they identified with $\\sigma$~Orionis cluster members and candidate members. \\citet{fra06} %Franciosini et~al. (2006) found no significant difference in the spectral properties of 23 classical and weak-line T~Tauri stars classified in the literature. The formers, however, tended to be less X-ray luminous than the last. This trend was also found in the cluster centre by \\citet{cab07b} %Caballero (2007b) using {\\em Chandra} data, which confirmed, for instance, the earlier results by \\citet{neu95} %Neuh\\\"auser et~al. (1995) in Taurus. Additional intermediate- and low-mass stars with X-ray emission found with {\\em Einstein}, the {\\em Advanced Satellite for Cosmology and Astrophysics} ({\\em ASCA}), {\\em ROSAT}, and {\\em Chandra} have been tabulated by \\citet{alc96}, \\citet{nak99}, \\citet{ada04}, \\citet{cab07a,cab08c}, and \\citet{ski08}. %Alcal\\'a et~al. %(1996), Nakano et~al. (1999), Adams et~al. (2004), Caballero (2007a, 2008b), and %Skinner et~al. (2008). %% Sokal \\& Skinner~(2007). The X-ray variability of 23 young stars (and one galaxy) with {\\em ROSAT} has been recently investigated by Caballero et~al. (in~prep.).  The X-ray activity extends down to below the hydrogen burning mass limit in $\\sigma$~Orionis. %Mokler \\& Stelzer (2002) \\citet{mok02} firstly investigated the X-ray emission near the substellar limit of the cluster. Of the seven very low-mass $\\sigma$~Orionis members in the error box of a {\\em ROSAT} source, three of them were classified as unambiguous detections. However, two of the actual emitters are very close ($\\rho \\sim$ 4--5\\,arcsec) to young low-mass cluster stars, that are probably the actual X-ray sources and not the brown dwarfs \\citep{cab06,cab07}. %(Caballero 2006; Caballero et~al. 2007). The third X-ray brown dwarf candidate in \\citet{mok02}, %Mokler \\& Stelzer (2002), \\object{Mayrit~633059} (S\\,Ori~3), is $\\Delta J \\sim$ 1.2\\,mag brighter than the stellar/substellar boundary (i.e. it is a~star).  To date, there remain only three {\\em confirmed} $\\sigma$~Orionis brown dwarfs with X-ray emission \\citep{fra06}: %(Franciosini et~al. 2006): \\object{Mayrit~433123} (S\\,Ori~25), \\object{Mayrit~487350} ([SE2004]~70), and \\object{Mayrit~396273} (S\\,Ori~J053818.2--023539). The spectra of the three of them have Li~{\\sc i} $\\lambda$6707.8\\,{\\AA} in absorption and/or low gravity features (e.g. abnormal alkali strength). Mayrit~433123 is a well known T~Tauri substellar analog with H$\\alpha$ in strong emission, as well as He~{\\sc i} \\citep{bej99,bar03,muz03}. %(B\\'ejar et~al. 1999; Barrado y Navascu\\'es et~al. 2003; %Muzerolle et~al. 2003). It is, besides, one of the very few brown dwarfs whose rotational velocity corrected from inclination has been derived %($v$ = 14\\,km\\,s$^{-1}$; %Caballero et~al. 2004). \\citep[$v$ = 14\\,km\\,s$^{-1}$;][]{cab04}. Mayrit~487350 underwent a flare during \\citet{fra06}'s %Franciosini et~al. (2006)'s observations, and might be the primary of the widest exoplanetary system detected so far %($r %\\sim$ 1700\\,AU; Caballero et~al.~2006). \\citep[$r \\sim$ 1700\\,AU;][]{cab06b}. Finally, Mayrit~396273 (the faintest brown dwarf of the trio) has been poorly investigated \\citep{ken05,max08}.  With the aim of exploring the X-ray emission of extremely low-mass $\\sigma$~Orionis members, F.~Mokler conducted new deep {\\em XMM-Newton} observations to the west of the cluster, complementing %Sanz-Forcada et~al.(2004) and \\citet{fra06}'s Franciosini et~al. (2006)'s \\citet{san04} and \\citet{fra06}'s observations, that imaged the cluster centre. The new observations were centred, in its turn, on Mayrit~1158274 (\\object{S\\,Ori~55}), a substellar object at the brown dwarf/planet boundary with strong variable H$\\alpha$ emission and infrared flux excess longwards of 5\\,$\\mu$m \\citep{zap02a,zap07,luh08}. %(Zapatero Osorio et~at. 2002a,~2007).  %We use the new deep {\\em XMM-Newton} observations to the west of the %cluster to look for the X-ray counterparts of S\\,Ori~55 and other cluster brown %dwarfs and stars in the area. %Once identified, we conduct a spectral analysis to derive their basic parameters %(e.g. energies of the thermal components, coronal metallicities, hydrogen %column densities). %We also look for flares during the X-ray monitoring and long-term variability. %Finally, we complement our results with data from the literature to study the %relation between X-ray emission and red optical-near infrared colours, which are %indicative of spectral type and of the presence of surrounding discs.  We use the deep \\textit{XMM-Newton} observations to the west of the $\\sigma$~Orionis cluster to determine general X-ray properties of the young stars in the cluster external region, where the stellar density drops. To complete our work, we give upper limits to the X-ray emission of the brown dwarfs in the field. Our intention is to determine the properties of the cluster X-ray emitters down to the substellar boundary. Once identified the X-ray stellar sources in the field, we conduct a spectral analysis to derive their basic parameters (e.g. energies of the thermal components, coronal metallicities, hydrogen column densities). We also look for flares during the X-ray monitoring and long-term variability. Finally, we complement our results with data from the literature to study the relation between X-ray emission and red optical-near infrared colours, which are indicative of spectral type and of the presence of surrounding discs. ",
        "conclusions": "\\label{section.discussion}  \\subsection{X-ray young stars} \\label{section.xrayyoungstars}  \\subsubsection{New X-ray sources in $\\sigma$~Orionis} \\label{section.newxrayyoungstars}  %%______________________________________________ Figure %\\begin{figure*} %\\centering %\\includegraphics[width=0.49\\textwidth]{xso_LiKs.eps} %\\includegraphics[width=0.49\\textwidth]{xso_LJKs.eps} %\\caption{Same as Fig.~\\ref{diagram.CCandCM}, but for the maximum likelihood %parameter $L$ as a function of colours $i-K_{\\rm s}$ ({\\em left}) and $J-K_{\\rm %s}$ ({\\em right}). %The dotted lines mark the approximate lower envelopes of the young stars with %X-ray emission.} %\\label{diagram.Lcolour} %% Done with: xso07.m %\\end{figure*} %%  We have derived the X-ray properties of 12 young stars (Table~\\ref{table.xray}), while three of them were too faint. Of the latter stars,  \\begin{itemize}  \\item NX~4 is a bright Tycho-2 star with arguable X-ray emission and expected spectral type at around late-A or early-F (with low frequencies of X-ray emission), \\item NX~16 is in the bluest part of the cluster sequence in the $i$ vs. $i-K_{\\rm s}$ diagram, has a very low maximum likelihood parameter ($L$ = 17), and was only detected by one camera (i.e. its youth is debatable), \\item and NX~36, with reliable (faint) X-ray emission and $J-K_{\\rm s}$ = 1.22$\\pm$0.04\\,mag, is the reddest young star in our sample and one of the reddest stars in $\\sigma$~Orionis. The infrared flux excess at IRAC and MIPS/{\\em Spitzer} passbands and its strong H$\\alpha$ emission (see references in Table~\\ref{table.class}) are an evidence of NX~36 (Mayrit~662301) having a developed disc, possibly edge-on, that absorbs any possible X-ray emission from the corona or from the channels that connect the inner border of the disc with the central object (K\\\"onigl~1991). \\end{itemize}   Of the twelve young stars with derived X-ray properties, four have not enough signal --\\,i.e. counts\\,-- for an accurate fitting to a $2T$-model. Therefore, only one thermal component was fitted for the four of them. They are: NX~20 (a confirmed young star with lithium), NX~14 and NX~21 (two new photometric cluster member candidates), and NX~34 (a non-variable photometric cluster member candidate; Scholz \\& Eisl\\\"offel 2004). The X-ray emission ensures their extreme youth. Likewise, the set of eight young stars with at least two thermal components comprises six confirmed stars with spectroscopic features of youth (NX~5, 7, 27, 28, 30, and~38) and two previously known photometric cluster member candidates (NX~17 and~37).  To sum up, by deriving their X-ray properties, we show that five young star candidates selected in colour-magnitude diagrams are actually young (two are presented here for the first time --NX~14 and~21--; two were unconfirmed cluster member candidates --NX~17 and~34--; one was confirmed by \\citet{fra06} %Franciosini et~al. [2006] --NX~37--). % NX\t\t14\t17\t21\t34\t37 % Mayrit \t1149270\t1298302\t1160240 887313\t615296 Besides, we first find X-ray emission from three young stars with lithium (NX~5, 7, and~20). In other words, of the 15 young stars investigated here, the detection of X-ray emission of {\\em nine} of them is new (but NX~4 and~16 detections are questionable).  Although these numbers are far from the numerous population of X-ray sources first found in the trailblazer works by Wolk (1996) and \\citet{fra06}, %Franciosini et~al. (2006), some of our young stars with emission are at previously uncovered angular distances from the cluster centre. Using previous {\\em ROSAT} and {\\em XMM-Newton} data, they could properly study the innermost 15\\,arcmin of $\\sigma$~Orionis, and hardly in the 15--20\\,arcmin external region. However, as shown by \\citet{cab08a}, %Caballero (2008a), the ``core'' and the ``halo'' of the cluster extends up to 20 and 30\\,arcmin, respectively. At larger separations from the Trapezium-like system that defines the cluster centre (i.e. $\\rho >$ 30\\,arcmin), the surface density of cluster members drops notoriously, and the contamination by nearby young populations (around Mintaka to the northeast, Alnilam to the northwest, and the Horsehead Nebula to the southeast) gets important \\citep{jef06, cab08c}. %(Jeffries et~al. 2006; Caballero 2008b). Except for NX~34 (that is at $\\rho \\sim$ 14.8\\,arcmin from $\\sigma$~Ori~AF--B), all our X-ray young stars lie on the external (17\\,arcmin $\\lesssim \\rho \\lesssim$ 24\\,arcmin) region (i.e. in the cluster core-to-halo transition). The decrease in the surface density of X-ray young stars at large distances from the cluster centre that we observe is an obvious consequence of the decrease in the surface density of $\\sigma$~Orionis members, that approximately follows a distribution proportional to $\\rho^{-1}$ \\citep{cab08a}.   \\subsubsection{Discs and $N_{\\rm H}$} \\label{section.discs}  The value of the column density $N_{\\rm H}$ in our X-ray model accounted for both the interstellar column density and the circumstellar material that could surround each star, like a protoplanetary disc or diffuse interstellar material left over from its formation. On the one hand, from the low $E(B-V)$ colour excess towards the $\\sigma$~Orionis cluster, that lies in the interval 0.04--0.09\\,mag \\citep{lee68, bej01, bej04, her05,she08,gon08}, %(Sherry et~al. 2008 -- see %another determinations at Lee [1964], B\\'ejar et~al. [2001, 2004], Hern\\'andez %et~al. [2005], Gonz\\'alez-Hern\\'andez et~al. [2008]), it is derived that the expected interstellar column density approximately varies between 2.2 and 4.9\\,10$^{16}$\\,m$^{-2}$ (i.e. $N_{\\rm H} \\approx$ 0.22--0.49~10$^{21}$\\,cm$^{-2}$). There are two of our 12 young stars with $N_{\\rm H}$ larger than this range ($N_{\\rm H} \\approx 1.0-1.2 \\times 10^{21}$\\,cm$^{-2}$), even taking the errors into account. They are NX~7 ($J-K_{\\rm s}$ = 0.48$\\pm$0.03\\,mag) and NX~34 ($J-K_{\\rm s}$ = 0.96$\\pm$0.05\\,mag). There are no hints for the first star to have a disc %(NX~7 falls out of the area investigated to date with the %{\\em Spitzer Space Telescope} and the Infrared Array Camera -- Hern\\'andez et~al. 2007), \\citep[NX~7 falls out of the area investigated to date with the {\\em Spitzer Space Telescope} and the Infrared Array Camera;][]{her07}, while the relatively red colour for the magnitude of the second star might suggest the presence of a disc (NX~34 has the third reddest $J-K_{\\rm s}$ among the 12~stars).  On the other hand, there are two $\\sigma$~Orionis stars in our analysis catalogued as classical T~Tauri stars. One is the faint X-ray emitter NX~36, that was discarded for the spectral energy distribution fitting because of their low number of summed counts. Its very red near-infrared colours ($J-K_{\\rm s}$ = 1.22$\\pm$0.04\\,mag) support the scenario of an X-ray absorbent disc (Section~\\ref{section.newxrayyoungstars}). However, the other classical T~Tauri star with infrared flux excess at 2--8\\,$\\mu$m (NX~38, with colour index $J-K_{\\rm s} = 1.10 \\pm 0.04$\\,mag) has a {\\em low} value of $N_{\\rm H}$. There are four young stars whose X-ray parameters have been derived and possess less counts than NX~38, and none of them displays such a $J-K_{\\rm s}$-$N_{\\rm H}$ incompatibility. Instead of claiming a contradiction with the scenario of an X-ray absorbent disc, we propose a simpler solution based on a geometrical fact: the X-ray emission from the upper layers of the atmosphere of NX~38 passes through a shorter portion of disc because of a relatively larger inclination angle with our visual (i.e. the disc is pole-on). This hypothesis, nevertheless, should be verified by determining the disc inclination angle. % E(B-V) [Lee 1968] = 0.05 mag %This is the column density of the material left over from the stars' formation.  %% %______________________________________________ Figure \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth]{aa10475f6.eps} \\caption{Same as Fig.~\\ref{diagram.CCandCM}, but for the $k_B T_2$ vs. $k_B T_1$ diagram only for the young stars in Table~\\ref{table.xray}. The young star candidate with high temperatures is NX~17. Its upper limit of $k_B T_2$ is undetermined.} \\label{diagram.kBTs} % Done with: xso07.m \\end{figure} %    \\subsubsection{X-ray variability} \\label{section.variability}  %______________________________________________ Figure \\begin{figure*} \\centering \\includegraphics[width=8.3cm]{aa10475f7a.ps} \\includegraphics[width=8.3cm]{aa10475f7b.ps} \\caption{Background subtracted X-ray lightcurves of the stars Mayrit 1298302 (NX\\,17) and Mayrit 797272 (NX\\,27) in the energy range 0.3--7.5\\,keV. The bin size is 800\\,s. The light grey area marks the parts of the spectrum rejected for this study because of the high background level. Both stars showed clear flaring events which caused the rising of the coronal temperature (see text for details).} \\label{fig.curves} \\end{figure*}  In the source detection stage (Section~\\ref{section.identification}), we created a {\\em cleaned} event-list, rejecting 30\\,\\% of the observation (see Table~\\ref{table.observations}). We used this event-list to look for variations in the lightcurve of our sources. Since the period in which high background level is present is large, the statistical analysis of variations was not possible. In these conditions, it is difficult to detect, for instance, variations due to rotational modulation. %Detection of flares %is also difficult in the {\\em cleaned} observation but the use of periods %with high background variability could lead to wrong results. Nevertheless, In our sample, two sources showed clear flaring events in their lightcurves: NX\\,17 and NX\\,27 (see Fig.~\\ref{fig.curves}). In this section, we focus on such events.  \\paragraph{NX\\,17.} This source (identified with the young cluster member candidate Mayrit 1298302; see Table~\\ref{table.class}) showed a large flare with a duration of more than $4$\\,h and a relative increase in flux of $\\sim 30$. This is much larger than observed in young stars in Taurus \\citep{fra07} and comparable with those observed in some young stellar objects in the Orion Nebula Cluster \\citep{fav05}. The mean coronal temperature during the observation reached 3.3\\,keV ($T \\sim 38$\\,MK). This value is typical of high-energy flaring events in late-type stars \\citep[see][and references therein]{cre07}.  \\paragraph{NX\\,27.} This source has been identified with the $\\sigma$~Orionis member Mayrit 797272, an M1--3 type star that shows a large lithium absorption line \\citep{wol96,cab08d}. %(Wolk 1997; Caballero et al. in prep.~II). %and H$\\alpha$ in broad emission \\citep{cab06}, indicating  accretion. In its X-ray lightcurve, it showed a first flare at $t \\approx 14$\\,ks from the beginning of the observation and high variability during the second half of it, probably coming from a second flare. With a total duration of $\\sim 9$\\,ks, the first flare presented a relative increase in flux of $\\sim 4$, similar to the observed in other young stars \\citep[e.g.][]{fav05,fra07} and in field M dwarfs \\citep[e.g.][]{rob05}. The temperature of the hotter component in the spectral fitting to the whole observation is 4.3\\,keV ($T \\sim 50$ MK). %Note that here, Both flares are probably contributing to the spectra. %and we cannot separate the two events.   \\subsection{Young stars and brown dwarfs without X-ray counterpart}  A series of star and bright brown dwarf members and candidates of the $\\sigma$~Orionis cluster in the Mayrit catalogue \\citep{cab08c}, and that fall in the area covered by \\textit{XMM-Newton}, were not detected in this observation because of their low X-ray emission. Most of them show features of youth, such as Li~{\\sc i} in absorption, H$\\alpha$ in strong broad emission, abnormal strength of alkali lines due to low gravity, early spectral types (OB), and/or flux excess longwards of the $J$ band (i.e. class~II). S\\,Ori~55 (Section~\\ref{section.introduction}), located in the centre of the field of view and close to the central gaps, is one of these objects. Other interesting ones are S\\,Ori~70 (see, again, Section~\\ref{section.introduction}), with a tentative mass of $\\sim 3$\\,$M_{\\rm Jup}$ and a possible disc, and \\object{S\\,Ori~66}, also in the planetary-mass domain and with a (likely) mid-infrared flux excess and strong H$\\alpha$ emission \\citep{bar01,zap08,sch08,luh08}. They are maybe the most interesting ultra low-mass objects in the area. %Note that  the source %\\object{IRAS~05358--0238} (with Mayrit number 377264) was not considered %a truly member of the cluster (it seems to be an intermediate M-type giant in the %background with strong silicon absorption features rather than a Class~I %object). %In any case, no X-ray emission was detected from the object. We Some other candidates and members of the cluster fall in a detector gap or close to the border, which avoided any reliable detection of such sources.  For all the members and candidates of $\\sigma$~Orionis in the Mayrit catalogue with no X-ray counterpart in our observation, we give upper limits to their X-ray emission in Table~\\ref{table.upperlimits}. Fluxes were determined from the count-rate measured in a circular region of 15\\,arcsec centred on the position of each source in the (exposure-map corrected) EPIC image. The count-rate simulator WebPIMMS\\footnote{WebPIMMS is available at:\\\\ {\\tt http://heasarc.nasa.gov/Tools/w3pimms.html}.} was used to transform count-rates to fluxes.  We assumed a hot plasma model with $kT = 1$\\,keV and a column density $N_\\mathrm{H} = 0.27 \\times 10^{21}$\\,cm$^{-2}$ for every source (see Section~\\ref{section.discs} for a detailed discussion on the interstellar column density in the direction of the $\\sigma$~Orionis cluster). The temperature of 1\\,keV is representative of the corona of young stars and has been also observed in the X-ray emitting brown dwarfs in the Orion Nebula cluster \\citep{fei05}. Some of the listed objects were marginally detected by the detection algorithm, but showed maximum likelihood parameters below the value chosen here for reliable detection, $L \\ge 15$.% (see Section~\\ref{section.identification} %for details). %This value corresponds to a conservative probability %$P < 3 \\times 10^{-7}$ of being a spurious detection.}  The upper limits given here put constraints to the X-ray emission of the low mass objects in the field. Very few is known about the X-ray emission from brown dwarfs. Apart from the $\\sigma$~Orionis brown dwarfs introduced in Section~\\ref{section.introduction}, there have been only a handful of detections in other young star-forming regions, such as the Orion Nebula Cluster, Chamaeleon, and Taurus \\citep{neu98,neu99,ima01,tsu03,pre05}. There have also been reports of X-ray emission from a brown dwarf in the Pleiades %(Briggs \\& Pye 2004, MNRAS) \\citep{bri04} and two brown dwarfs in the field: \\object{LP~944--20} %(a member of the Castor moving group; Rutledge et~al. 2000,ApJL) \\citep[member of the Castor moving group,][]{rut00} and \\object{GJ~569B}[ab] %(with a question mark; Stelzer 2004, ApJL). \\citep[questionable;][]{ste04}. Nowadays, it is accepted that brown dwarfs and low-mass stars share the same formation mechanism %(e.g. Caballero et~al. 2007) \\citep[e.g.][]{cab07} and, therefore, brown dwarfs are ``stars to scale'' without stable hydrogen burning. Albeit brown dwarfs are smaller (fainter, cooler) than stars, they rotate much faster %(with periods of about 10\\,h; see Mohanty \\& Basri 2003, ApJ, and %references therein), \\citep[with periods of about 10\\,h; see][and references therein]{moh03} which may enhance their magnetically-driven, internal dynamos. % This enhancement might produce, in its turn, an extra heating of the hot, tenuous, X-ray-emitting, upper atmospheric layer in brown dwarfs. %In order to test any theoretical scenario of X-ray emission from substellar %objects, it is mandatory to collect a reasonable number of X-ray measurements %from bona-fide brown dwarfs, like those in the $\\sigma$~Orionis cluster.  \\subsection{Comparison to previous {\\em XMM-Newton} observations} \\label{section.franciosini}  There are nine X-ray sources in common with the {\\em XMM-Newton} observations %Franciosini et~al. (2006) by \\citet{fra06} and ours (see second column in Table~\\ref{table.class}). They are six young stars, two galaxies (NX~35 and~39), and a source of unknown nature (NX~41; Section~\\ref{section.unknown}), classified by \\citet{fra06} %Franciosini et~al.(2006) as $\\sigma$~Orionis cluster members or candidates, unidentified X-ray sources, and a possible cluster candidate from 2MASS, respectively. The X-ray parameters determined here for the stellar sources for which a spectral fitting was possible in both works (NX~28, NX~30, and NX~37) are well in agreement with those obtained by \\citet{fra06}. The remaining stars in common are too faint in, at least, one of both works to perform a robust fitting to their spectra and, thus, we cannot compare results. The case of NX~27 %\\citep[[FPS2006]\\,1 in][]{fra06} ([FPS2006]\\,1 in Franciosini et al. 2006) is remarkable since, in our observation, its count-rate is more than one order of magnitude higher than in the \\citet{fra06}'s one. Here, the star showed high X-ray variability coming from large flaring events (Section~\\ref{section.variability}) and it had counts enough to do an accurate spectral fitting. The energy band used in our study, $0.3-7.5$\\,keV, is similar to that of \\citet{fra06}, who used $0.3-8.0$\\,keV. % %In general, the results obtained here on the properties of the corona of the %young stars in the field are in agreement with those obtained for the %late-type stars in the centre of the cluster and indicate an uniformity of %the population of the cluster. In other regions, this situation is not %similar \\citep[see, for instance,][and references therein]{alb07}.  Besides, there are seven additional X-ray sources that \\citet{fra06} associated to cluster members and candidates and that were not detected with our task, although they fall in our survey area. However, all of them (with [FPS2006] identification number in the last column of Table~\\ref{table.upperlimits}) are very faint and located in (or very close to) borders or gaps of the detectors in our observation.  \\subsection{X-ray galaxies} \\label{section.xraygalaxies}  Of the 12 X-ray sources in Table~\\ref{table.sources} with maximum likelihood parameters $L \\gtrsim 400$, eight are young stars and four are background galaxies with $HR_2 > -0.5$. %(i.e. soft emission). % NX 12, 13, 26, 32 The remaining 29 sources have $L \\lesssim 100$. Two of the four hard X-ray emitters in our sample (NX~13 and~32) were previously detected with {\\em Einstein} and identified with galaxies (see Section~\\ref{section.galaxies} for details). They are the eighth and second brightest sources in this {\\em XMM-Newton} field, respectively. The other two galaxies are NX~12 and~26 (with $L$ = 1601 and 394, respectively) and have no near-infrared counterpart (NX~26 has no optical counterpart, either).  The study of the X-ray properties of the AGNs in the field is far from the scope of this paper. Nevertheless, we show in Fig.~\\ref{fig.agns} the X-ray spectra observed here for NX\\,13 and NX\\,32, together with a simple fitting using a power-law, to compare with the spectra of the young stars (Figs.\\,\\ref{nx05-21} and \\ref{nx27-38}). The X-ray spectra of these sources are clearly distinguishable from those of the young stars in our sample.  %We plot in Fig.~ref{\\bf ...} their X-ray spectral fittings, for comparison with %those of young stars. %{\\bf To do X-ray spectral fitting of NX~13 and~32 (required: only the plot, not %the values).} %The other two galaxies are NX~12 and~26 (with $L$ = 1601 and 394, %respectively) and have no near infrared counterpart (NX~26 has no optical %counterpart, either). %They have, as a result, large X-ray-to-bolometric luminosities ratios, %indicative of very energetic processes (possibly associated to a massive black %hole in the galactic centres).    We used archived data from the {\\em XMM-Newton} mission to explore a region to the west of the centre of the $\\sigma$~Orionis cluster, slightly overlapped with a previous observation centred on the eponymous stellar system $\\sigma$ Ori. Our aim was to search for new X-ray young stars in the region and determine their X-ray properties. After a careful inspection of the source list resultant from the use of a maximum likelihood technique to reveal sources in the images of the EPIC cameras, we kept 41 X-ray sources. Among them, a total of 15 are young stars, 4 are field stars, and 19 are AGNs. The remaining three sources are of unknown nature, probably AGNs given their hardness ratios. % Many of the revealed X-ray sources, including those associated to young stars, are new.  We derived X-ray parameters of the twelve young stars with enough counts for a spectral fitting, using multi-temperature models with absorption. In particular, we determined the column density, temperature of the thermal components, and metallicities. % %The X-ray properties of the young stars determined here are in well agreement %with those observed in the members situated in the centre of the cluster, %indicating a uniformity in its X-ray population. % We observed no relation among the stars showing infrared excess and the column density measured here. % Although we could not perform a detailed study of the X-ray variability because of the high variable background affecting the observation, two clear flaring events in the M-type young stars NX~17 (Mayrit 1298302) and NX~27 (Mayrit 797272) were studied. In both cases, we determined high coronal temperatures, reaching $38$ and $50$\\,MK, respectively. % We also determined upper limits for the X-ray flux of young stars and brown dwarfs in the field that were not detected by the searching algorithm used by us.  Deep X-ray observations, combined with photometric data, have historically revealed as a powerful tool for disclosing the stellar population of star-forming regions down to the substellar boundary. In particular, in $\\sigma$~Orionis we were able to detect X-ray emission of mid-M type stars. Our work complements previous investigations on the X-ray population of this cluster. %A detailed study on the X-ray luminosity function of %$\\sigma$~Orionis, compared with other young stellar clusters, will %be addressed in a forthcoming paper by the authors (L\\'opez-Santiago et %al., in prep.). %There, we will also investigate the star formation in the %cluster. Our study also advances in the knowledge of the X-ray properties of young stars in the external region of the cluster, where the stellar density drops. Deeper X-ray observations in this region should permit to detect also the X-ray counterparts of brown dwarfs (as it was done in the cluster centre). The combination of their X-ray properties with the results of near- and mid-infrared photometry (to reveal the presence of discs) will permit to unveil the origin of the X-ray emission of these objects (whether it is produced by accretion or by a hot corona, as in pre-main sequence stars). This work should be followed by the study of the frequency of brown dwarfs with and without disc showing X-ray emission. A detailed study on the X-ray luminosity function of $\\sigma$~Orionis, compared with other young stellar clusters, will be addressed in a forthcoming paper by the authors.  %If there are stars with disc %between our members and/or candidates, these must be pole on. %Infrared excesses are indicative of circumstellar material surrounding the %star (usually related to a disc in T Tauri stars). The column density measures %the amount of material in the direction of the source. If none foreground %interstellar material exists, a large column density also indicates the %presence of circumstellar material. The fact that in the stars in our sample %showing infrared excesses we did not measure a large column density seems to %indicate that their disc is pole-on."
    },
    "0809/0809.2027_arXiv.txt": {
        "abstract": "% {} {We investigate the Coma cluster galaxy luminosity function (GLF) at faint magnitudes, in particular in the u* band where basically no studies are presently available at these magnitudes, by applying photometric redshift techniques. } { Cluster members were selected based on Probability Distribution Function from photometric redshift calculations applied to deep u*, B, V, R, I images covering a region of almost 1~deg$^2$ (completeness limit R$\\sim$24). In the area covered only by the u* image, the GLF was also derived after applying a statistical background subtraction. } {Global and local GLFs in the B, V, R and I bands obtained with photometric redshift selection are consistent with our previous results based on a statistical background subtraction.  The GLF in the u* band shows an increase of the faint end slope towards the outer regions of the cluster.  The analysis of the multicolor type spatial distribution reveals that late type galaxies are distributed in clumps in the cluster outskirts, where X-ray substructures are also detected and where the GLF in the u* band is steeper.} {We can reproduce the GLFs computed with classical statistical subtraction methods by applying a photometric redshift technique. The $u*$ GLF slope is steeper in the cluster outskirts, varying from $\\alpha \\sim -1$ in the cluster center to $\\alpha \\sim -2$ in the cluster periphery. The concentrations of faint late type galaxies in the cluster outskirts could explain these very steep slopes, assuming a short burst of star formation in these galaxies when entering the cluster.} ",
        "introduction": "\\label{sec:intro}  Ever growing optical imaging surveys of galaxies for which complete spectroscopic follow-up is impossible due to telescope limitations has triggered the development of photometric redshift techniques (e.g. Bolzonella et al. 2000 or Ilbert et al. 2006a and references therein). Based on the comparison of multi-band photometry with synthetic spectral templates, this technique can be viewed as very low resolution spectroscopy.  The quality of photometric redshifts depends on the wavelength range covered by the photometric survey. High-quality photometric redshifts and related quantities require a complete and contiguous coverage of the spectral regions of interest, in particular around the strong spectral features, such as the $4000$\\,\\AA\\ break or the Lyman $\\alpha$ break. This technique has proven to be a very valuable tool for several cosmological purposes, such as deriving field galaxy correlation functions or galaxy luminosity functions (see e.g. Ilbert et al. 2006b, Meneux et al. 2006). It has also been applied in the study of distant clusters of galaxies (see e.g. White et al. 2005 and references therein).  In the present paper, we apply the photometric redshift technique to study the galaxy population of the rich Coma cluster. Due to its proximity (z$\\sim$0.023), Coma covers a large extent over the sky (of the order of 1 deg$^2$), requiring a wide field camera; data in the U band are needed to determine with a sufficient accuracy the probability that a galaxy belongs or not to the cluster, based in particular on the 4000~\\AA\\ break.  As a continuation of our Coma photometric survey (e.g. Adami et al. 2006a) which already includes deep (R$\\sim$24) wide field (42$\\times$52 arcmin$^2$) B, V, R and I images, we recently acquired a Megacam u* band image of comparable depth and field of view. In addition to our spectroscopic redshift catalog, this allows us to compute photometric redshifts and Probability Distribution Function (PDF hereafter) along the Coma cluster line of sight down to the dwarf galaxy regime.   One of the main goals of the paper is to qualify this photometric redshift technique applied to the galaxy luminosity function (GLF hereafter) determination. For this, we will compare the present photometric redshift technique with the galaxy luminosity functions computed with the same data set but with a statistical background removal (Adami et al. 2007a and b).  We will also compute luminosity functions in the u* band based on statistical background subtraction in the area where only the u* band is available.  We describe our new photometric data and photometric redshift and PDF estimates in Section 2. We compute in Section 3 the Coma GLF using the PDF and in Section 4 the Coma u* GLF applying statistical subtraction. The spatial distributions of multicolor galaxy types are discussed in Section 5. Conclusions are drawn in Section 6.  In this paper we assume H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega _m$=0.3, $\\Omega _{\\Lambda}$=0.7,  a distance to Coma of 100 Mpc, a distance modulus = 35.00, and a scale of 0.46 kpc arcsec$^{-1}$. ",
        "conclusions": "We have shown that we can reproduce the GLFs computed by Adami et al. (2007a) by selecting cluster galaxies with PDFs from photometric redshift techniques. This both puts the Adami et al. (2007a) conclusions on a firmer ground, in terms of infalling directions and processes acting on galaxies inside the cluster, and proves that the photometric redshift technique (with a judicious cut in redshift) can help to select cluster members in nearby structures.  We clearly show that the u* GLF is steeper in the cluster outskirts, as already suggested by Donas et al. (1995). This could be due to a short burst of star formation in the faint galaxies, induced by the pressure of the intracluster medium (see also e.g. Cortese at al. 2007 or Bou\\'e et al. 2008).  We plan in a future paper to investigate how GLF and general galaxy properties (based on photometric redshift estimates) vary with the local galaxy density as derived for bright galaxies by Dressler et al. (1980)."
    },
    "0809/0809.0508_arXiv.txt": {
        "abstract": "We present an analysis of the spectral energy distribution (SED) of the galaxy ESO 184-G82, the host of  the closest known long gamma-ray burst (GRB) 980425 and its associated supernova (SN) 1998bw. We use our observations obtained at the Australia Telescope Compact Array (the third $>3\\sigma$ radio detection of a GRB host) as well as archival infrared  and ultraviolet (UV) observations to estimate its star formation state.  We find that  ESO 184-G82 has  a UV star formation rate (SFR) and stellar mass consistent with the population of cosmological GRB hosts and of local dwarf galaxies. However, it has a higher specific SFR (per unit stellar mass) % than luminous spiral galaxies. The mass of ESO 184-G82 is dominated by an older stellar population in contrast to the majority of GRB hosts. The Wolf-Rayet region $\\sim800$ pc from the SN site experienced a starburst episode during which the majority of its stellar population was built up. Unlike that of the entire galaxy, its SED is similar to those of cosmological submillimeter/radio-bright GRB hosts with hot dust content.  These findings add to the picture that in general, the environments of  GRBs on 1--3 kpc scales are associated with high specific SFR and hot dust. ",
        "introduction": "Long gamma-ray bursts (GRBs) are associated with the death of  massive stars \\citep[e.g.][]{galamanature,hjorthnature,stanek}. This makes them of special interest in cosmology because they possibly trace the evolution of the rate of star formation in the Universe \\citep[e.g.][]{lambreichart00,jakobsson05,jakobsson06}. Indirect evidence of the nature of GRBs was found by studying their host galaxies  \\citep[e.g.][]{bloom,christensen04,sollerman05,castroceron08,savaglio09}.  Moreover, several studies on the immediate environments of GRBs suggest a close connection of long GRBs with regions of star formation, and therefore that their progenitors are likely massive stars. \\citet{fruchter06} found that GRBs trace the ultraviolet (UV) brightest parts of their host \\citep[see also][]{bloom02}. \\citet{thone08} studied in detail the environment (in $3$ kpc bins) of \\object{GRB 060505}, concluding that it originated in the youngest, most metal-poor and most intense star-forming region in the host galaxy. Similarly, \\citet{ostlin08} found that the $0.3$ kpc environment of \\object{GRB 030329} is much younger than the entire galaxy and its estimated age suggests a conservative lower limit on the mass of the GRB progenitor equal to $12\\,M_\\odot$. Finally, a significant number of other GRBs were reported to reside in dense star-forming regions \\citep{castrotirado99,hollandhjorth99,hjorth03,savaglio03,vreeswijk04,chen05,chen06,fynbo06b,watson06,watson07, prochaska07b,prochaska07,ruizvelasco07} and molecular clouds \\citep{galamawijers01,stratta04,jakobsson06b,campana07,prochaska08}.  \\object{GRB 980425} is the closest known GRB \\citep[$z=0.0085$;][]{tinney98}, therefore it is an excellent laboratory for local GRB studies. \\citet{galamanature} reported the Type Ic supernova \\object[SN 1998bw]{(SN) 1998bw} exploding inside the error box of \\object{GRB 980425}. Its lightcurve was well modeled by an explosion of a Wolf-Rayet (WR) star \\citep{iwamoto98}, which is a highly evolved and massive star that has lost its outer hydrogen layers. Up to now, three other GRBs have also been spectroscopically confirmed to be associated with SNe: \\object{GRB 030329} \\citep{hjorthnature,matheson03,stanek}, \\object[GRB 031203]{031203} \\citep{cobb04,galyam04,malesani04,thomsen04} and \\object[GRB 060218]{060218} \\citep{ferrero06,mirabal06,modjaz06,pian06,soderberg06nature,sollerman06},  while two GRBs were confirmed to be SN-less: GRB \\object[GRB 060505]{060505} and \\object[GRB 060614]{060614} \\citep{fynbo06,dellavalle06,galyam06}.  The host galaxy of \\object{GRB 980425} / \\object{SN 1998bw}  \\citep[\\object{ESO 184-G82};][]{holmberg77} is a dwarf \\citep[$0.02$ of the characteristic blue luminosity, $L_B^*$;][]{fynbo00} barred spiral  \\citep[SBc;][]{fynbo00} with axis diameters of $12$ and $10$ kpc \\citep[down to $B=26.5$ mag arcsec$^{-2}$;][]{sollerman05}, dominated by a large number of star-forming regions \\citep{fynbo00,sollerman05}. \\object{SN 1998bw} occurred inside one of these, $\\sim800$ pc southeast of a region displaying a Wolf-Rayet type signature spectrum  \\citep[hereafter: WR region;][]{hammer06}. The WR region dominates the galaxy's emission at $24\\,\\micron$  \\citep{lefloch} and is the youngest region within the host  exhibiting very low metallicity  \\citep{christensen08}.  In this paper we present fits to the spectral energy distribution (SED) of \\object{ESO 184-G82} and the WR region and compare their properties to other galaxies. Section \\ref{sec:data} lists the data sources \\citep[including the third radio detection of a GRB host after those reported by][]{bergerkulkarni,berger} used for the SED modeling of Section \\ref{sec:model}. We derive properties of the host galaxy and WR region in Section \\ref{sec:results}, discussing their implications in Section \\ref{sec:discussion}. Section \\ref{sec:conclusion} closes with our conclusions. We use a cosmological model with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda=0.7$ and $\\Omega_m=0.3$, so \\object{ESO 184-G82} is at a luminosity distance of  $36.5$ Mpc and $1''$ corresponds to 175 pc at its redshift. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have presented the UV-to-radio SED fitting of the host galaxy of \\object{GRB 980425} / \\object{SN 1998bw} and of the WR region close to the SN position.  The host galaxy of \\object{GRB 980425} is a normal dwarf spiral galaxy with somewhat elevated star formation activity compared to other spirals (though it is not necessary to invoke any starburst episode to explain its SED). The steep radio slope and the presence of the $\\sim1.6\\,\\micron$ bump in the SED indicate the existence of an old stellar population.  Its low radio luminosity can be explained by the suppression of synchrotron emission in dwarf galaxies and the fact that radio is only sensitive to recent star formation.  The emission of the WR region close to the GRB position is dominated by an ongoing starburst episode, during which almost all of its stars were formed. It contributes significantly  to the star formation of the entire galaxy. In many aspects the WR region is similar to high-redshift GRB hosts: it is a blue, young region of intense star formation containing hot dust. The presence of the GRB close to this region indicates that GRBs appear to be  associated with regions of high specific SFR and high dust temperatures."
    },
    "0809/0809.0214_arXiv.txt": {
        "abstract": "{} {Searching for planets in the atmosphere of AGB stars is difficult, due to confusion with the stellar wind and pulsations. The aim here is to provide a complementary strategy for planet searches in such a dense environment.} {The polarization properties of SiO masers, especially their circular polarization, are, under certain conditions, good tracers of rapid magnetospheric events. A Jovian planet with a magnetosphere whose dipole axis is misaligned with its rotation axis naturally provides such conditions. Here I present several models showing that the polarization will be periodically modulated.} {The linear and circular polarization of an SiO maser in a planetary magnetosphere is modulated by the precessing dipole component of the latter. The effect is measurable in saturated masers, while unsaturated masers only exhibit weak changes, because of dilution effects, and because the circular polarization there stems from the Zeeman effect making it as weak as for thermal radiation. The situation would change if anisotropic pump- and loss-rates were included, which would increase the fractional linear and, via magnetorotation, the circular polarization of the modulation.} {Single-dish monitoring with a dense enough time sampling and a carefully calibrated polarimeter, in combination with VLBI observations, are suited to detecting and locating a periodic modulation of the circular maser polarization due to a precessing Jovian magnetosphere. The phenomenon will be rare, because a favorable arrangement of maser and magnetosphere is needed. Otherwise the polarization may be below the detection threshold, especially if the maser is unsaturated. Though exhibiting a qualitatively similar modulation, linear polarization is likely to suffer more from confusion due to dilution of the magnetosphere within the maser cross section, even in VLBI observations.} ",
        "introduction": "Most SiO masers are hosted by the extended atmosphere of evolved stars of low- to intermediate mass in the upper part of the asymptotic giant branch, before they evolve towards central stars of planetary nebulae (for a review see e.g., Herwig 2005). Under certain conditions, these masers trace the magnetic field by virtue of their polarization. After the first theoretical consideration of a Zeeman laser (Sargent et al. 1967), astronomical maser polarization has been extensively discussed, e.g., by Goldreich et al. (1977), Western \\& Watson (1984, for linear polarization), Deguchi \\& Watson (1986, for circular polarization), and Elitzur (1991). The models are compared by Gray (2003). The environment of SiO masers in AGB atmospheres is characterized by a dense wind driven by the stellar pulsations. Its interaction with planets engulfed in the stellar atmosphere has already been addressed by Struck-Marcell (1988) who investigated the hypothesis that SiO masers may form in the magnetosphere of Jovian planets. Struck et al. (2002, 2004) propose scenarios along the same lines, and model the dynamic effects of episodic accretion onto the planets. In a different context, the consequences of planets in the stellar atmosphere have been investigated by Soker (2001), who shows that Jovian and even Earth-like planets can efficiently spin-up the star after migration into the stellar atmosphere, such that it later may form an elliptical rather than a spherical planetary nebula.  Here I propose to resume the idea of Struck-Marcell (1988) that SiO masers (though not all) may originate in Jovian magnetospheres, leading to circular maser polarization (Barvainis et al. 1987; Herpin et al. 2006), which may be as strong as $\\sim 10\\%$ of the Stokes I flux. I show that the densely sampled time series of all Stokes parameters yield a rare, but observable polarization signature characteristic of a precessing magnetosphere. Since AGB stars are slow rotators, a precession period as fast as $\\sim 10$\\,h hints at a planetary magnetosphere. SiO masers are an ideal tool for sounding Jovian planets, since they arise in a zone extending from 1.5 to 7\\,AU distance from the star where they are radiatively and collisionally pumped (Gray et al. 2009), before the SiO molecule becomes underabundant for maser action thanks to condensation in the dust envelope. The SiO maser zone thus corresponds to a distance from the star where the solar system harbors its giant gas planets (Jupiter is at about 5\\,AU from the sun, and Saturn at about 10\\,AU). A direct comparison between the solar sytem now and in the AGB phase is not valid, though, because planetary orbits may be dragged inwards by tidal friction, or expelled by the kinematical effect of stellar mass loss (Villaver \\& Livio 2006, further references therein). Arguing with Soker (1996) that elliptical planetary nebulae formed from stars spun up by inward migration of a planet, it is assumed here that the zone in question harbors Jovian planets, since 50\\,\\% of all planetary nebulae are not elliptical (Soker 2001).  The first observational evidence of extrasolar planets around evolved stars was given by Silvotti et al. (2007) who indirectly detected a planet around a post-red giant star, measuring the pulse modulations of the variability of the latter. As for the planets detected around pulsars (with a similar method, measuring the period variations of the latter) by Wolszczan \\& Frail (1992), it is not yet certain whether these planet-sized bodies survive the supernova explosion, or whether they formed in the debris disk forming in reponse to a ``fallback'' (Wang et al. 2006) of ejected matter. Most extrasolar planets have so far been detected around main-sequence stars by radial velocity measurements (for a review see Santos 2008). As for the earliest phases of stellar evolution, a \"hot Jupiter\" has recently been discovered around the T Tauri star TW Hya (Setiawan et al. 2008). Around evolved stars, ten companions of substellar mass have been discovered to this day, the last one being the K0 giant HD17092 (Niedzielski et al. 2007, further references therein). There is no evidence yet from AGB stars, mainly because the strong, dense wind makes radial velocity measurements of the star difficult. The following models of maser polarization modulation due to precessing planetary magnetospheres are intended as a complement to exoplanet searches around AGB stars with different methods. ",
        "conclusions": "I have shown that densely sampled polarization monitoring is able to reveal precessing Jovian magnetospheres engulfed in the atmospheres of AGB stars by a periodic modulation of the polarization characteristics of SiO masers crossing them. The expected period is that of the rotation of giant gas planets (i.e., $\\sim 10\\,$h if the rotation period did not change with respect to typical solar system values), which allows us to distinguish such a fingerprint from long-term variations of the polarized (Glenn et al. 2003) and unpolarized maser flux (Humphreys et al., 2002; Diamond \\& Kemball, 2003), which occur on timescales at least an order of magnitude longer than those considered here. As for the distinction from pseudo-periodic polarization variations due to stellar magnetic flux loops or magnetic clouds, everything depends on the speed of the SiO masers relative to the latter. At a speed of 10\\,km\\,s$^{-1}$, a maser covers a distance of $5\\,r_{\\rm J}$ within 10\\,h, i.e., the patterns modeled here are only visible if the maser speed across the planetary magnetosphere is $< 10$\\,km\\,s$^{-1}$, or if the precession period is $< 10$\\,h. Otherwise, polarization fluctuations are expected that not only contain temporal variations, but also spatial ones, and the patterns will be either compressed or stretched. A planetary magnetosphere exposed to an AGB wind has its magnetopause closer to the planet than e.g., Jupiter's magnetopause, which is exposed to the faster, but less dense solar wind. As a matter of fact, the case of a planetary magnetosphere in an AGB wind may be closer to that of the terrestial magnetosphere, whose magnetopause is as close as a few Earth radii, and can even be temporarily disrupted in case of strong solar storms. Likewise, the magnetic dipole component of a planetary magnetic field engulfed by the atmosphere of an AGB star can be expected to be quite compact (not more than 10 planetary radii). To observe the modeled features, an SiO maser has to cross it, and either needs to be saturated to produce a measurable circular polarization, be unsaturated if the pump and loss rates are sufficiently anisotropic, or has to cross a part of the magnetosphere with a strong line-of-sight component of the magnetic field. This leads to the conclusion that the phenomenon will be rare and, if detected, only be seen in a narrow range of velocities, and preferentially in circular polarization-enhancing regions of strong magnetic flux density (whereas linear polarization tends to cancel out if the polarization angle strongly varies across the observing beam). Furthermore, the dynamics of the gas close to the planet possibly accreting matter needs to be such that the line-of-sight velocity coherence allows for the maser action. The abundance of SiO molecules required to form a strong enough maser can be expected to be less of a problem. The evaporation of Galilean moons has been proposed as an additional reservoir for the gas-phase SiO (Struck-Marcell, 1988).  Recently, Wiesemeyer et al. (2009) have detected a pseudo-periodic variation of the fractional circular polarization towards two Mira stars, R~Leo, and V~Cam, which may hint at precessing magnetospheres in their atmospheres. Several periods will have to be monitored, to better distinguish the phenomenon from other rapid, but transient magnetospheric events. The case of R~Leo is especially intriguing, since VLBA maps of the 43~GHz SiO maser (Cotton et al. 2004, 2008; Soria-Ruiz et al. 2007) show elongated SiO features pointing radially away from the star, reminiscent of jet-like features. The velocity information from the observed polarization fluctuations, together with the astrometry of such elongated features, suggest orbital motion in at least one of them, and strengthens the case for a planetary wake flow as simulated by Struck et al. (2004). Dense polarization monitoring of these features may be especially rewarding, and a spatially fully resolved polarization monitoring will certainly be highly conclusive. However, VLBI techniques will suffer from a confusion between polarized flux variations and variations due to Earth rotation synthesis of visibility components (the timescale for both is unfortunately the same). This calls for a combination of single-dish polarization monitoring with VLBI imaging of the Stokes I emission. The cross-identification of polarization fluctuations with VLBI features is possible if there is, in the single-dish observations, no blend of spatially separated SiO maser spots at the same velocity or if the VLBI features show an unambiguous association with a planetary wake flow. Detection of orbital motion in such features will provide a touchstone for the scenarios presented here. The aim of this paper has been to motivate coordinated observational efforts, so far unprecedented, and to show that meaningful results can be expected to emerge from them. If so, this will call for advanced modeling. The work published here should therefore be considered as the initial step in an iterative refinement.  \\onlfig{6}{ \\begin{figure} \\includegraphics[width=6.5cm,angle=-90]{fig5a.ps} \\includegraphics[width=6.5cm,angle=-90]{fig5b.ps} \\includegraphics[width=6.5cm,angle=-90]{fig5c.ps} \\caption{Same as Fig.~\\ref{fig:unsaturatedPencil}, but $3\\,r_{\\rm J}$ above the planet's equatorial plane and towards its rotation axis (model 2c). Top: $p_{\\rm L}$, contour levels $0.3\\,\\%$ to $1.5\\,\\%$ by $0.3\\,\\%$. Center: $\\chi$, contour levels from $-16^\\circ$ to $+160^\\circ$ by $16^\\circ$. Bottom: $p_{\\rm C}$, contour levels from $-0.28\\,\\%$ to $+0.28\\,\\%$ by $0.04\\,\\%$. \\label{fig:unsaturatedPencil2} } \\end{figure} } \\onlfig{7}{ \\begin{figure} \\includegraphics[width=6.5cm,angle=-90]{fig6a.ps} \\includegraphics[width=6.5cm,angle=-90]{fig6b.ps} \\includegraphics[width=6.5cm,angle=-90]{fig6c.ps} \\caption{Same as Fig.~\\ref{fig:unsaturatedPencil2}, but slightly offset from the rotation axis (by $0.5\\,r_{\\rm J}$, model 2d). Top: $p_{\\rm L}$, contour levels $0.1\\,\\%$ to $0.6\\,\\%$ by $0.1\\,\\%$ (the small jumps in the insert showing the time series of $p_{\\rm L}$ are a numerical artifact). Center: $\\chi$, contour levels from $-16^\\circ$ to $+160^\\circ$ by $16^\\circ$. Bottom: $p_{\\rm C}$, contour levels from $0.02\\,\\%$ to $+0.1\\,\\%$ by $0.02\\,\\%$. \\label{fig:unsaturatedPencil3} } \\end{figure} }"
    },
    "0809/0809.2999_arXiv.txt": {
        "abstract": "We examine the accretion and merger histories of central and satellite galaxies in a smoothed particle hydrodynamics (SPH) cosmological simulation that resolves galaxies down to $7\\times 10^9 M_\\odot$.  Most friends-of-friends haloes in the simulation have a distinct central galaxy, typically two to five times more massive than the most massive satellite.  As expected, satellites have systematically higher assembly redshifts than central galaxies of the same baryonic mass, and satellites in more massive haloes form earlier. However, contrary to the simplest expectations, satellite galaxies continue to accrete gas and convert it to stars; the gas accretion declines steadily over a period of $0.5-1$ Gyr after the satellite halo merges with a larger parent halo.  Satellites in a cluster mass halo eventually begin to lose baryonic mass.  Since $z=1$, 27\\% of central galaxies (above $3\\times 10^{10} M_\\odot$) and 22\\% of present-day satellite galaxies have merged with a smaller system above a 1:4 mass ratio; about half of the satellite mergers occurred after the galaxy became a satellite and half before.  In effect, satellite galaxies can remain ``central'' objects of halo substructures, with continuing accretion and mergers, making the transition in assembly histories and physical properties a gradual one.  Implementing such a gradual transformation in semi-analytic models would improve their agreement with observed colour distributions of satellite galaxies in groups and with the observed colour dependence of galaxy clustering. ",
        "introduction": "In the standard theoretical description of galaxy formation, galaxies grow initially by the condensation of gas at the centres of dark matter potential wells \\citep[e.g.][]{white78,fall80}. When a large dark matter halo accretes a smaller halo, the galaxy hosted by the smaller halo becomes a satellite system, which may eventually merge with the central object after sinking by dynamical friction. Even if the satellite retains some of its original dark matter, it is likely to move through the large parent halo with a random velocity that exceeds its own escape speed, making it difficult for the satellite to accrete gas from the halo or to merge with other satellites. In semi-analytic models of galaxy formation \\citep[e.g.][]{white91,kauffmann93,cole94,somerville98,croton06,bower06}, these expected differences between central and satellite galaxies are usually encoded by simple rules, with some variations from one model to another. When two haloes merge, the central galaxy of the more massive progenitor becomes the central galaxy of the new, larger halo. Cooling gas from the halo is assigned mostly or entirely to the central galaxy. Satellites accrete little or no fresh gas, and their pre-existing gas may be removed by ram pressure stripping when they enter the larger halo. Most or all mergers are assumed to involve satellites joining the central galaxy rather than merging with each other.  In this paper, we examine the growth of central and satellite galaxies in smoothed particle hydrodynamics (SPH) cosmological simulations \\citep{katz92b,evrard94,katz96}. In such simulations, differences between central and satellite galaxies are not imposed a priori, but central galaxies nevertheless emerge as a distinct class for the physical reasons discussed above. The galaxy closest to the halo centre of mass is usually older and more massive than satellites in the same halo \\citep{berlind03,zheng05}. Here we investigate the assembly histories of present day centrals and satellites as a function of galaxy and halo mass, and we track the growth of galaxies to see the extent to which satellite systems experience continuing gas accretion or mergers with other satellites.  Our goals are to understand the physics of galaxy assembly in the simulation and to inform semi-analytic models of galaxy formation and other efforts to interpret observed galaxy clustering and the environmental dependence of galaxy properties. The central-satellite distinction is a key element in modeling the origin of bimodality in the galaxy population, in which galaxies with old stellar populations and little ongoing star formation form a distinct, passive, ``red sequence'' \\citep{strateva01,blanton01,bell03,faber07}. Strangulation of gas supplies in satellites is the likely route by which many galaxies join the passive sequence, though additional mechanisms are required to explain the red colours and low star formation rates of the most massive galaxies \\citep[see e.g.][]{croton06,dekel06,cat07}. Central-satellite separation is also a key ingredient in ``halo occupation distribution'' (HOD) models of galaxy clustering \\citep{kravtsov04,zheng05} and in recent efforts to interpret properties of galaxy groups \\citep[e.g.][]{weinemann06,gilbank07,hansen07,cli07,coil08}.  The simulation that we analyse here is also the main simulation analysed by \\cite{keres05} and \\cite{maller06}, who examine gas accretion and merger rates but do not distinguish between central and satellite galaxies. \\cite{keres05} emphasise the distinction between between ``cold'' and ``hot'' modes of gas accretion. In haloes of total mass $M\\le3 \\times 10^{11} M_{\\odot}$, most gas that accretes onto galaxies does so without ever heating close to the halo virial temperature. In higher mass haloes, most gas heats to the virial temperature before settling into galaxies. Overall, about half of the baryonic mass in the simulation's $z=0$ galaxy population was originally accreted in ``cold mode'' and half in ``hot mode'', with the former dominating at high redshift and in low mass galaxies today. Here we will investigate both total and cold-only accretion rates of central and satellite systems. Further discussions of the cold/hot dichotomy appear in \\cite{binney77}, \\cite{binney03}, \\cite{katz92a}, \\cite{kay00}, \\cite{fardal01}, \\cite{katz03}, \\cite{birnboim03}, \\cite{croton06}, \\cite{dekel06} and \\cite{keresp}.  We describe our simulation and methods for identifying haloes and galaxies in \\S2, before proceeding to our main results in \\S3. There we start with a qualitative overview of accretion and merger properties of central and satellite galaxies, then quantify their relative accretion rates and merger rates and their dependences on halo mass. In \\S4 we summarise our results and discuss their implications, with particular attention to recent studies of galaxy bimodality and the red population in groups. ",
        "conclusions": "the distinction between central and satellite galaxies is fundamental to understanding the galaxy population and its evolution, but differences in their growth rates and hence their physical properties are smooth, not sharp."
    },
    "0809/0809.3351_arXiv.txt": {
        "abstract": "By the precise timing of the low amplitude (0.005 - 0.02 magnitude) transits of exoplanets around their parent star it should be possible to infer the presence of other planetary bodies in the system down to Earth-like masses.  We describe the design and construction of RISE, a fast-readout frame transfer camera for the Liverpool Telescope designed to carry out this experiment. The results of our commissioning tests are described as well as the data reduction procedure necessary.  We present light curves of two objects, showing that the desired timing and photometric accuracy can be obtained providing that autoguiding is used to keep the target on the same detector pixel for the entire (typically 4 hour) observing run. ",
        "introduction": "\\label{sec:intro}  %   The Liverpool Telescope\\cite{steele} (LT) is a 2.0 metre fully robotic telescope.  The telescope Acquisition \\& Guidance (A\\&G) unit can host up to five instruments\\cite{mottram} with the beam able to be directed to one of four side ports via a folding mirror, or, with the folding mirror removed from the beam, to a straight through port.  The rapid time to switch between instruments ($<30$ seconds) combined with automated observing scheduling\\cite{fraser} means that the telescope is well suited to time variable (on timescales of minutes to years) and rapid reaction astronomy.  A key science goal of the Liverpool Telescope is the discovery of Earth mass exoplanets (i.e. planets outside our own solar system). Two techniques are pursued.  The first is via participation in the coordinated followup programmes\\cite{robonet} of galactic bulge microlensing events where small anomalies in the light curves act as signatures of planets orbiting the lens star.  An example of the application of this technique in which the LT was involved was the recent discovery of a Jupiter/Saturn solar system analogue\\cite{planet}. The second technique involves the precise timing of transits of known Jupiter-like exoplanets around their parent star, looking for deviations from the expected ephemeris due to the influence of other (unseen) planets in the system perturbing the orbit.  It is this second technique which is the subject of this paper, which describes the construction of a moderately wide field, fast-readout CCD imager for the Liverpool Telescope known as RISE (Rapid Imaging Search for Exoplanets).  In this paper we will give a brief overview of the science driver for the instrument and the requirements it places on the design, followed by details of its design and construction. We will also discuss the data reduction procedure for the instrument, and give the results of our commissioning tests. ",
        "conclusions": "RISE was a fast-track instrument for the Liverpool Telescope, which was designed, built and commissioned within 1 year of the initial science requirement being identified. Although work is still underway to better understand the effect of scattered light on the observations, it is already clear that the timing precision obtainable is sufficient to carry out our programme of transit followup, which has therefore begun routine data taking.  In addition we are pleased to note that it is also attracting interest from other science users of the telescope attracted by its low overheads for timing experiments and its relatively wide field."
    },
    "0809/0809.3217_arXiv.txt": {
        "abstract": "\\vspace*{2.5mm}\\noindent There is a unique Lorentz-violating modification of the Maxwell theory of photons, which maintains gauge invariance, CPT, and renormalizability. Restricting the modified-Maxwell theory to the isotropic sector and adding a standard spin--$\\half$ Dirac particle $p^\\pm$ with minimal coupling to the nonstandard photon $\\widetilde{\\gamma}$, the resulting modified-quantum-electrodynamics model involves a single dimensionless ``deformation parameter,'' $\\widetilde{\\kappa}_\\text{tr}$. The exact tree-level decay rates for two processes have been calculated: vacuum Cherenkov radiation $p^\\pm \\to p^\\pm\\,\\widetilde{\\gamma}$ for the case of positive $\\widetilde{\\kappa}_\\text{tr}$ and photon decay $\\widetilde{\\gamma} \\to p^+\\,p^-$ for the case of negative $\\widetilde{\\kappa}_\\text{tr}$. From the inferred absence of these decays for a particular high-quality ultrahigh-energy-cosmic-ray event detected at the Pierre Auger Observatory and a well-established excess of TeV gamma-ray events observed by the High Energy Stereoscopic System telescopes, a two-sided bound on $\\widetilde{\\kappa}_\\text{tr}$ is obtained, which improves by eight orders of magnitude upon the best direct laboratory bound. The implications of this result are briefly discussed. \\vspace*{0cm}\\newline ",
        "introduction": "\\label{sec:introduction}  Spacetime is a dynamic entity and quantum mechanics inescapably leads to randomness. Combined, this suggests~\\cite{Wheeler1957} that space may not be perfectly smooth at very small length scales and that Lorentz invariance may be fundamentally violated.   Precision tests of Lorentz invariance are, therefore, of paramount importance and the ideal testing ground is the photonic sector. If space, on the whole, is homogeneous and isotropic, there is essentially one dimensionless parameter that describes Lorentz violation in the photonic sector (in a first approximation, this nonstandard photon is taken to be coupled to standard Lorentz-invariant charged particles such as the electron).  The aim of the present article is to obtain a new bound on this isotropic Lorentz-violating parameter. Needless to say, the parameter may also have an entirely different origin than the nontrivial spacetime structure mentioned above. For the purpose of obtaining the bound, we remain ``agnostic'' as to the potential origin of this particular Lorentz-violating parameter. ",
        "conclusions": "\\label{sec:conclusion}  From the results of the previous section, the final two-sided two--$\\sigma$ bound on the single isotropic Lorentz-violating parameter of modified--QED theory \\eqref{eq:modQED-action}--\\eqref{eq:modM-standD-actions-ansaetze} is \\beq\\label{eq:UHECR-two-sided-bound} - 9 \\times 10^{-16} < \\widetilde{\\kappa}_\\text{tr} < 6 \\times 10^{-20}\\,. \\eeq The simple model \\eqref{eq:modQED-action}--\\eqref{eq:modM-standD-actions-ansaetze} can be considered to be embedded in a Standard Model Extension~\\cite{ColladayKostelecky1998} with all physical (renormalized) Lorentz-violating parameters set to zero, except for one, namely, the isotropic, CPT--even, Lorentz-violating parameter for the dimension-four term with two Maxwell field strength tensors of the photon field.\\footnote{With nonzero weak mixing angle~\\cite{PeskinSchroeder1995}, the $W^\\pm$ and $Z^0$ vector fields also display modified propagation (and interaction) properties, because of $SU(2)\\times U(1)$ gauge invariance [the action is given by \\eqref{eq:modSM-action}--\\eqref{eq:isotropic-ansatz-kappatilde-i} in Appendix~\\ref{app:bounds-universal-isotropic-LV-parameter} with $\\overline{\\kappa}_\\text{tr}^{(1)} =\\overline{\\kappa}_\\text{tr}^{(2)} \\equiv\\widetilde{\\kappa}_\\text{tr}$ and $\\overline{\\kappa}_\\text{tr}^{(3)}=0$]. However, these weak-vector-boson modifications do not directly affect the tree-level calculations of processes \\eqref{eq:decay-a} and \\eqref{eq:decay-b}, with the fermions considered as Dirac point particles.}   If the same isotropic Lorentz-violating parameter applies equally to all gauge bosons of the Standard Model, bound \\eqref{eq:UHECR-two-sided-bound} can be replaced by a stronger one, which has a lower bound at $-2 \\times 10^{-19}$ (see Appendix~\\ref{app:bounds-universal-isotropic-LV-parameter}). But the derivation of this particular lower bound involves certain assumptions on the parton distributions (the error is, however, still dominated by the relatively large error from the primary-energy measurement). For the moment, bound \\eqref{eq:UHECR-two-sided-bound} for the purely photonic parameter is to be preferred, as it does not involve a partonic calculation.  This new indirect bound \\eqref{eq:UHECR-two-sided-bound} improves by 8 orders of magnitude upon the direct laboratory bound \\eqref{eq:lab-bound}. Observe that, strictly speaking, \\eqref{eq:UHECR-two-sided-bound} is also a ``terrestrial bound,'' as it relies on having detected a primary traveling over a few hundred meters in the Earth's atmosphere, the precise astronomical origin of the primary being of secondary importance (see the first of the last three remarks in Sec.~\\ref{sec:new-bounds}). For the type of experiments involved (Auger~\\cite{Auger2004} and HESS~\\cite{HESS2003}), the Earth's atmosphere is an integral part of the instrument.   Combined with previous astrophysics bounds~\\cite{KosteleckyMewes2002} on the 10 birefringent modified-Maxwell-theory parameters at the $10^{-32}$ level and ``terrestrial'' UHECR bounds~\\cite{KlinkhamerRisse2008b} on the 8 nonisotropic nonbirefringent parameters at the $10^{-18}$ level, bound \\eqref{eq:UHECR-two-sided-bound} gives the following two--$\\sigma$ bound on \\emph{all} entries of the background tensor $\\kappa^{\\mu\\nu\\rho\\sigma}$ in the general modified-Maxwell action \\eqref{eq:modM-action}: \\beq\\label{eq:UHECR-kappa-tensor-bound} \\max_{\\,\\{\\mu,\\nu,\\rho,\\sigma\\}} \\; |\\kappa^{\\mu\\nu\\rho\\sigma}|< 5\\times 10^{-16} \\,, \\eeq where the fact has been used that the largest entry of $|\\kappa^{\\mu\\nu\\rho\\sigma}|$ has a value $(1/2)\\,|\\widetilde{\\kappa}_\\text{tr}|$ if the other 18 parameters are negligibly small. For completeness, the three--$\\sigma$ bound is $\\max\\,|\\kappa^{\\mu\\nu\\rho\\sigma}| < 6 \\times 10^{-16}$. Remark that bounds \\eqref{eq:UHECR-two-sided-bound} and \\eqref{eq:UHECR-kappa-tensor-bound} hold in a Sun-centered, nonrotating frame of reference.   In order to put \\eqref{eq:UHECR-kappa-tensor-bound} in perspective, recall that the equality of gravitational and inertial mass was experimentally verified by  Newton (1687), Bessel (1832), and E\\\"{o}tv\\\"{o}s (1889) at the $10^{-3}$, $10^{-5}$, and $10^{-7}$ levels, respectively~\\cite{FischbachTalmadge1999}. These experiments then led Einstein to the Principle of Equivalence, the cornerstone of the General Theory of Relativity~\\cite{Einstein1916}. In our case, the conclusion from \\eqref{eq:UHECR-kappa-tensor-bound} would be (elaborating on remarks in Refs.~\\cite{Collins-etal2004,GagnonMoore2004,BernadotteKlinkhamer2007}) that the local Lorentz invariance left-over from a fundamental theory of quantum spacetime is exact and that the underlying principle for this apparent fact remains to be determined."
    },
    "0809/0809.1212_arXiv.txt": {
        "abstract": "The direct current (DC) electric field near the reconnection region has been proposed as an effective mechanism to accelerate protons and electrons in solar flares. A power-law energy spectrum was generally claimed in the simulations of electron acceleration by the reconnection electric field. However in most of the literature, the electric and magnetic fields were chosen independently. In this paper, we perform test-particle simulations of electron acceleration in a reconnecting magnetic field, where both the electric and magnetic fields are adopted from numerical simulations of the MHD equations. It is found that the accelerated electrons present a truncated power-law energy spectrum with an exponential tail at high energies, which is analogous to the case of diffusive shock acceleration. The influences of reconnection parameters on the spectral feature are also investigated, such as the longitudinal and transverse components of the magnetic field and the size of the current sheet. It is suggested that the DC electric field alone might not be able to reproduce the observed single or double power-law distributions. ",
        "introduction": "Particle acceleration remains a mystery in solar flares, as well as in cosmic rays from the extrasolar systems. Observationally, the energetic particles always show a single or double power-law spectral behavior. Three mechanisms have been proposed \\citep[see][for reviews]{mill97, asch02}, i.e., the direct current (DC) electric field acceleration \\citep[e.g.,][]{stur68}, turbulence (or stochastic) acceleration \\citep[e.g.,][]{ml98}, and shock acceleration \\citep[e.g.,][]{blan78}. Magnetic reconnection, as the effective mechanism for magnetic energy release in solar flares, was demonstrated to be able to provide environments favorable for all the above-mentioned acceleration mechanisms to work \\citep{chen07}.  DC electric field provides the simplest and most direct means of accelerating particles out of a thermal plasma \\citep{holm85}, which even allows for an analytical solution of each particle trajectory under certain assumptions \\citep{bula80}. In the case that the electric and magnetic fields induced by the accelerated particles are negligible, test particle simulations provide a simple and valid approach to study the particle acceleration in a reconnection-associated electric and magnetic fields. For example, \\citet{saka92} showed that protons and electrons can be promptly accelerated (within 1 s) up to $\\sim$70 MeV and $\\sim$200 MeV, respectively. In order to compare the resulting energy spectrum of the DC-accelerated protons with observations, \\citet{mori98} conducted test particle simulations with a hyperbolic magnetic field and a uniform electric field. The accelerated protons present a universal power-law spectrum, $f(E)\\sim E^{-\\delta}$, with the spectral index $\\delta$ being $\\sim2.0-2.2$. Thereafter, a lot of test particle simulations have been performed for electrons as well, through either full orbit calculations \\citep{zhar04,hami05} or guiding center calculations \\citep{wood05}. Almost all of these simulations claimed to have obtained a power-law energy spectrum, with a spectral index around 2, similar to that for protons. It is noted, however, that the energy spectra in many of the published simulation results deviate significantly from a power-law profile. Moreover, according to the thick-target model, the resulting bremsstrahlung hard X-ray (HXR) emissions in these simulations would present a power-law energy spectrum with a spectral index around 1. However, the RHESSI observations indicated that the HXR emissions in most solar flares possess much softer spectra, with the spectral indices falling in the range between 3 and 7. As discussed by \\citet{chen07}, several factors may be attributed to such a big discrepancy. For example, the parameters of the reconnecting current sheet were chosen arbitrarily due to our poor knowledge of the physical conditions in the localized reconnection region in solar flares. Most importantly, the electric and magnetic fields, which should be coupled with each other through MHD equations, were often prescribed independently. Doing so often results in the electric and magnetic fields that are not compatible with the MHD equations.  Some improvements were made in recent years. For example, \\citet{wood05} derived the electric field via the Ohm's law after assuming the distributions of the magnetic field, the velocity field, and the resistivity. A further progress was made by \\citet{hami05}, who deduced the solution of the linearized MHD equations by assuming that the longitudinal component of the magnetic field ($B_z$) and the velocity ($\\mathbf{v}$) are small perturbations. In the real case, both $B_z$ and $\\mathbf{v}$ can be quite large, therefore, it is necessary to obtain the electric and magnetic fields by directly solving the MHD equations.  In order to obtain self-consistent electric and magnetic fields, in this paper we first perform 2.5-dimensional MHD simulations of the resistive evolution of a current sheet. When a steady state is reached, the electric and magnetic fields are then taken out for test particle simulations. The paper is organized as follows. In \\S{2}, the numerical method is described, including the MHD and test particle simulations. The shape of the resulting energy spectrum is discussed in \\S{3}, and a parameter survey is conducted showing the influences of various parameters on the spectral features in \\S{4}. A short discussion is presented in \\S{5}. ",
        "conclusions": "X-ray and $\\gamma$-ray observations of solar flares indicated that electrons and protons are instantly accelerated in the magnetic reconnection process. Analyses of HXR spectral data revealed that the accelerated electrons present a single power-law energy spectrum in most flares or a double power-law spectrum in some other events. The co-existence of normal and reverse type III radio bursts strongly suggests that nonthermal electrons are accelerated near the reconnection X-point within a compact region less than 2000 km in length \\citep{asch02}. Near the reconnection X-point, both DC electric field and the reconnection-associated turbulence may contribute to the acceleration of these nonthermal electrons \\citep{stur68,chen07}. The DC mechanism, as the simplest model, has been studied by various groups with help of test particle simulations. Although the resulting energy spectrum of the accelerated electrons is somewhat power law-like, the spectral indices seem to be much smaller than the typical values derived from HXR observations. Therefore, more work is required to fill the gap between test-particle simulations and observations.  Using self-consistent electric and magnetic fields obtained from numerical simulations of 2.5-dimensional nonlinear MHD equations, we investigated the energy spectrum of the electrons that are accelerated by the DC electric field in a reconnecting current sheet. Contrary to many of the previous studies, our test particle simulations indicate that the accelerated electrons present a power-law spectrum truncated by an exponential high energy tail, i.e., $ f(E)\\sim E^{-\\delta} e^{-E/E_0}$, which is similar to the case of diffusive shock acceleration \\citep{elli85} and the case of random DC-electric field acceleration when the electron trapping time is long \\citep{anas97}. A parameter survey is conducted to investigate how various physical quantities affect the spectral profile and the acceleration rate. As the energy spectra are fitted with the function $ f(E)\\sim E^{-\\delta} e^{-E/E_0}$, it is found that  (1) $E_0$, the rollover energy, increases with larger $B_{g}$ (the guide magnetic field), larger $L_0$ (which characterizes the size of the reconnection diffusion region), smaller $\\beta_0$ (the plasma beta), and smaller $\\eta_0$ (the resistivity);  (2) $\\delta$ increases with larger $L_0$ and smaller $\\eta_0$, and it saturates as $L_0$ is larger than 100 m or $\\eta_0$ is very small. The dependence of $\\delta$ on $B_{g}$ or $\\beta_0$ is not monotonic. $\\delta$ reaches the minimum when $B_{g}$ is $\\sim 0.4$ times the antiparallel component of the reconnecting magnetic field, while $\\delta$  reaches the maximum when $\\beta_0$ is around 0.01;  (3) The acceleration rate, defined here as the percentage of the electrons that are accelerated above to 10 keV, increases with larger $B_{g}$, larger $L_0$, and smaller $\\beta_0$. It does not change much with $\\eta_0$, although the percentage of the electrons above 20 keV steadily decreases with increasing $\\eta_0$.  The truncated power-law energy spectrum obtained in this paper is somewhat similar to the double power-law distribution, although the transition between the low and high energy tails is not so abrupt in most cases. The research in this paper reveals that, even with self-consistent electric and magnetic fields, we still cannot reproduce the observed single or double power-law energy spectra of the nonthermal electrons in the framework of DC-electric field mechanism. This discrepancy between test particle simulations and observations might be reconciled if other effects are included. For example, the turbulence in the reconnection site would greatly enhance the collision rate between nonthermal electrons and background particles or waves \\citep{wu05}. In other words, DC-electric field in the reconnection site alone may not be able to explain the observed spectral features of nonthermal electrons. The combination of DC-electric field and the turbulence may finally lead to single or double power-law spectral distributions, which were derived from HXR observations. This is definitely an issue requiring further investigation."
    },
    "0809/0809.2094_arXiv.txt": {
        "abstract": "Polarized far-infrared and submillimeter emission is calculated for models of nonspherical dust grains that are constrained to reproduce the observed wavelength-dependent extinction and polarization of starlight.  For emission from regions where the magnetic field is perpendicular to the line-of-sight, the far-infrared emission is expected to have substantial linear polarization at wavelengths $\\lambda \\gtrsim 100~\\micron$, but the degree of linear polarization, and its variation with wavelength, is model-dependent.  Models in which the starlight polarization is produced by both amorphous silicate and graphite grains have linear polarizations between $6\\%$ and $10\\%$ at $\\lambda > 100~\\micron$, but for some models in which only silicate grains are spheroidal, the linear polarization increases from about $3\\%$ at $100~\\um$ to about $15\\%$ at 1~mm.  We briefly discuss the implications of these results for removal of the polarized dust emission from maps of the Cosmic Microwave Background, as well as the possibility of discriminating among interstellar dust models based on observations of far-infrared and submillimeter linear polarization. ",
        "introduction": "\\label{sec:intro}  The presence of dust grains in the interstellar medium has been apparent since the recognition that obscuring material existed in at least some regions of interstellar space~\\citep{Barnard_1907, Barnard_1910, Trumpler_1930}.  Our knowledge of interstellar dust has advanced considerably since then, but remains incomplete [see, e.g., the review by \\citet{Draine_2003}, or the book by \\citet{Whittet_2003}].  In the Milky Way, large fractions of abundant elements, such as Si and Fe, appear to have been depleted from the gas phase \\citep{Jenkins_2004}, with the missing atoms believed to be present in dust grains.  Studies of interstellar ``reddening'' produced by wavelength-dependent extinction revealed broad spectroscopic features at $9.8~\\um$ and $18~\\um$, attributed to amorphous silicate material in submicron particles, and a strong extinction ``bump'' at $2175$~\\AA, usually attributed to aromatic carbon, perhaps graphite or polycyclic aromatic hydrocarbons (PAHs), in particles with radii $a \\lesssim 0.01~\\um$.  Another broad extinction feature, at $3.4~\\um$, is typical of C-H stretching modes in hydrocarbons. Finally, the observed infrared emission from dust includes spectral features at $3.3$, $6.2$, $7.7$, $8.6$, $11.3$ and $12.7~\\um$ that can be explained by a substantial population of particles with PAH composition and sizes extending down to about $40$ atoms.  These spectroscopic features indicate that amorphous silicates and carbonaceous materials, including PAHs, contribute a substantial fraction of the interstellar grain mass, and models developed to try to reproduce the observed extinction using observed abundances of the elements invariably employ amorphous silicates and carbonaceous materials as their main components.  The serendipitous discovery \\citep{Hiltner_1949, Hall_1949} of the linear polarization of starlight demonstrated that interstellar grains are both appreciably nonspherical and aligned with the interstellar magnetic field, so that the dusty interstellar medium acts as a polarizer.  Studies of starlight polarization in the ultraviolet \\citep{Clayton+Anderson+Magalhaes_etal_1992, Anderson+Weitenbeck+Code_etal_1996, Martin+Clayton+Wolff_1999} indicate that the smallest dust grains are apparently minimally aligned \\citep{Kim+Martin_1995}, and studies in the infrared [see, e.g., \\citet{Whittet_2003}] find strong polarization in the $9.8~\\um$ Si-O stretch absorption feature, requiring that a substantial fraction of the amorphous silicate particles be aligned. Recent measurements of the polarization in the neighborhood of the $3.4~\\um$ C-H stretch absorption feature \\citep{Chiar+Adamson+Whittet_etal_2006} find only upper limits on possible polarization associated with the feature, indicating that the hydrocarbon material responsible for this feature is primarily in a grain population that is not appreciably aligned, which may rule out models where the hydrocarbon material is located in mantles coating silicate cores.  While we know for sure that interstellar grains are not spherically symmetric, their actual shapes are uncertain.  Because the degree of alignment is unknown, we can only say that the grains must depart from spherical symmetry sufficiently to produce the observed polarization.  Interstellar grains radiate strongly in the far-infrared.  Because the grains are nonspherical and partially alig\\-ned, this emission must be partially polarized.  The polarization of the far-infrared emission has been measured for infrared-bright regions in star-forming molecular clouds \\citep[see the review by][]{Hildebrand_2005}.  The degree of polarization of the thermal emission from dust in these regions is observed to have a complicated wavelength dependence, with the polarization at $350~\\um$ observed to be significantly smaller than the $100~\\um$ and $850~\\um$ polarization \\citep{Vaillancourt_2002, Hildebrand+Kirby_2004, Hildebrand_2005, Vaillancourt_2007}, which has been interpreted as indicating two separate populations of emitting grains in these clouds.  The polarization of the large-scale submillimeter emission from dust in the interstellar medium of the Milky Way has now been measured by several experiments, including the Archeops balloon \\citep{Benoit+Ade+Amblard_etal_2004, Ponthieu+Macias-Perez+Tristram_etal_2005} and the {\\it WMAP} satellite \\citep{Page+Hinshaw+Komatsu_etal_2007, Gold+Bennett+Hill_etal_2008}. A quantitative understanding of this emission is important for removal of polarized Galactic ``foregrounds\" from maps of the cosmic microwave background (CMB), but also offers an opportunity to test models for interstellar grains.  In this paper, we construct models of interstellar dust that are consistent with the observed wavelength-dependent interstellar extinction, and the observed linear dichroism of the interstellar medium.  We will calculate the polarized infrared and submillimeter emission from different models of interstellar dust heated by the general interstellar starlight, and, for each of them, compute the corresponding expected percentage polarization.  This will allow us to gain a sense of the degree to which the predicted far-infrared polarization is model-dependent.  We discuss the observed wavelength-dependent extinction and polarization of starlight in \\S\\ref{sec:obs}, and the basics of the dust models we consider in \\S\\ref{sec:models}. \\S\\ref{sec:cross_sections} provides a description of  the technique we use to compute optical cross sections, while \\S\\ref{sec:picket_fence} introduces the formalism associated with ``picket-fence\" alignment. In \\S\\ref{sec:size_dist_align_func}, we present the size distributions and alignment functions that provide the best fit to the observed extinction and polarization discussed in \\S\\ref{sec:obs}.  We describe the formalism we use to calculate the partially-polarized far-infrared emission by interstellar dust in \\S\\ref{sec:formalism}.  The corresponding results are presented and discussed in \\S\\ref{sec:ir_em}.  Finally, \\S\\ref{sec:summary} offers a summary of the main results of this paper, and a brief discussion of their consequences for studies of the CMB. ",
        "conclusions": "\\label{sec:ir_em}  Figure \\ref{fig:em_spectra} shows the emission spectra calculated for each of the five grain models listed in Table~\\ref{tab:models} when illuminated by the standard interstellar radiation field for the solar neighborhood.  Also shown is the emission spectrum calculated for the grain model of~\\citet{Draine+Li_2007}.  The emission $\\left<I_\\lambda\\right>$ is calculated using the cross section $\\left<\\Cabs\\right>$ from equation~(\\ref{eq:abs_cross_sec}), appropriate for randomly-oriented dust.  These spectra of the thermal emission do not include rotational emission from spinning dust, which appears to dominate for $\\lambda \\gtrsim 4000~\\um$ ($\\nu \\lesssim 75$~GHz) \\citep{Draine+Lazarian_1998,Finkbeiner_2004}.  \\begin{figure*} \\begin{center} \\epsfig{file=F7.eps, width=7.93cm, angle=0} \\caption{\\label{fig:em_spectra} Total emission spectra for the different grain models listed in Table~\\ref{tab:models}, as well as for the model described in~\\protect\\citet{Draine+Li_2007}, each heated by the interstellar radiation field estimated by~\\protect\\citet{Mathis+Mezger+Panagia_1983}.  The $x$-axis is now labeled by wavelengths, and the symbols correspond to the DIRBE results for HI correlated emission at high Galactic latitudes~\\protect\\citep{Dwek+Arendt+Fixsen_etal_1997, Arendt+Odegard+Weiland_etal_1998}.  Frequencies covering the full range of \\emph{Planck} HFI frequencies (100~GHz, 217~GHz, and 857~GHz) are also indicated.} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\epsfig{file=F8a.eps, width=7.93cm, angle=0} \\hspace*{-0.5cm} \\epsfig{file=F8b.eps, width=7.93cm, angle=0} \\caption{\\label{fig:pol_spectra} ({\\it Left}) Spectra of polarized thermal emission predicted by models 1 through 4  in directions perpendicular to the direction of the local magnetic field, and ({\\it Right}) corresponding degrees of polarization for each model. Two classes of models are clearly visible.  Note that these models do not include rotational emission from spinning dust, which will be important for $\\lambda \\gtrsim 4000~\\um$ ($\\nu \\lesssim 75$~GHz).} \\end{center} \\end{figure*}  The emission spectra for all six models are relatively similar.  In each case, $\\lambda\\,\\langle I_\\lambda \\rangle$ peaks near $130~\\micron$, with an amplitude within $20\\%$ of the~\\citet{Draine+Li_2007} (thereafter, DL07) model.  Moreover, all models agree extremely well for wavelengths below about $20~\\um$, where the signal is entirely dominated by PAH features.  However, there are some noticeable differences, especially in the ratio of the 130~$\\um$ to the 24~$\\um$ emissions.  Although these ratios are within $15\\%$ of each other for model 0 and the DL07 model, models involving spheroidal grains lead to smaller values of this ratio, typically by a factor of 1.5.  This difference arises mostly from variations in the size distributions between models, with differences in grain shape also having a small effect on the average emission spectrum.  All other differences are fairly modest (usually at the 10\\% to 20\\% level), and generally smaller than the actual uncertainties in the appropriate dielectric functions.  The fact that all of these spectra are in good agreement overall, but with significant well localized differences, opens a very interesting window through which one can hope to strongly constrain dust models with observations of the total dust emission at just a couple of wavelengths.  We have not attempted this exercise here.  Figure~\\ref{fig:pol_spectra} shows the polarized emission spectra calculated for models 1 through 4, as well as the corresponding degrees of polarization (in \\%), for lines of sight perpendicular to the direction of the local magnetic field, and the dust exhibiting the highest degrees of alignment found in the diffuse interstellar medium.  Contrary to the total emission spectra shown on Figure~\\ref{fig:em_spectra}, the polarized emission spectra clearly reflects the fact that we consider two classes of models.  The spectra of polarized emission for models 2 and 4, with spheroidal silicate and carbonaceous grains, are indeed remarkably similar, but quite different from those for models 1 and 3, with silicate spheroids and carbonaceous spheres.  This translates into interestingly different predictions for the expected degree of polarization for emission by interstellar dust, as can be seen in the right panel of Figure~\\ref{fig:pol_spectra}.  Although all models predict virtually no polarization below $40~\\um$ (as small grains and PAHs are not aligned), the wavelength dependence of the degree of polarization for $\\lambda \\gtrsim 40~\\um$ is quite model-dependent.  Models with both silicate and carbonaceous spheroidal grains lead to the degree of polarization sharply increasing between about $40~\\um$ and $600~\\um$, after which it reaches a plateau.  However, the degree of polarization for models in which only the silicate grains are spheroids increases steadily with increasing wavelength, without converging to an upper limit.  As a result, models 1 and 3 predict significantly higher degrees of polarization than models 2 and 4 for wavelengths longer than about $200~\\um$.  For example, when $\\lambda = 3000~\\um$, the maximum degree of polarization predicted by models 1 and 3 is of order $10\\%$, while models 2 and 4 lead to a maximum of about $15\\%$.  These differences provide another way of discriminating between dust models, this time from observations of the polarized emission from interstellar dust.  The strong wavelength-dependence of the polarization in models 1 and 3 can be easily understood.  In these two models, the carbonaceous grains are not aligned.  For the adopted grain optical properties, the silicate grains contribute an increasing fraction of the emission as the wavelength increases, in part because the silicate grains are slightly cooler than the carbonaceous grains (this explains the behavior at $\\lambda \\lesssim 200~\\um$), and in part because the ratio of the silicate opacity to the graphite opacity increases with increasing wavelength for $\\lambda \\gtrsim 100~\\um$.     The main results of this paper are as follows: \\begin{enumerate} \\item We present models of interstellar dust based on mixtures of spheroidal amorphous silicate grains, spheroidal graphite grains, and PAH particles, with a degree of alignment $f$ that is a function of the effective radius $a_\\eff$ of the grains.  The models differ in assumptions concerning the shapes (axial ratios $b/a = 1.4$ and 1.6 are considered for the spheroidal grains) of the silicate and graphite grains. \\item Size distributions $\\dd n/\\dd a_\\eff$ and alignment function $f(a_\\eff)$ are found that reproduce the observed wavelength-dependent extinction and polarization of starlight.  As found previously by \\citet{Kim+Martin_1995}, the degree of alignment is small for $a_\\eff\\nobreak\\lesssim\\nobreak0.05~\\micron$, and large for $a_\\eff\\nobreak\\gtrsim\\nobreak0.1~\\micron$. \\item For models where only the silicate grains are aligned, the degree of alignment approaches $100\\%$ for $a \\gtrsim 0.1~\\micron$, whereas when both silicate and graphite grains are partially-aligned spheroids with $b/a = 1.4$, the observed starlight polarization can be reproduced with $f \\approx 0.4$ for $a \\gtrsim 0.1~\\micron$.  If the axial ratio is increased to $b/a = 1.6$, the required degree of alignment drops to $f \\approx 0.3$. \\item For each model, we calculate the partially-polarized infrared emission when the dust is illuminated by the local starlight background.  The total infrared emission spectra are generally similar, and in approximate agreement with observations of the Galactic cirrus emission at $|b| \\gtrsim 20^\\circ$.  In detail, the spectra are model-dependent, particularly between about $20~\\um$ and $160~\\um$.  No attempt has been made to adjust our models to reproduce observed FIR spectral energy distributions. \\item Because the small grains and PAHs are not aligned, the polarized emission is very weak for $\\lambda \\lesssim 40~\\micron$.  However, for $\\lambda \\gtrsim 60~\\micron$, the linear polarization of the thermal emission is appreciable when viewed from directions perpendicular to the local magnetic field.  At $100~\\micron$, the different models considered here result in linear polarizations between about $3\\%$ and $6\\%$.  For wavelengths longer than $500~\\um$ ($\\nu \\leq 600$~GHz), of relevance to CMB studies, the degrees of polarization predicted by our models range from about $9\\%$ to around $15\\%$ for the thermal emission.  However, for $\\lambda \\gtrsim 4000~\\um$ ($\\nu\\nobreak\\lesssim\\nobreak75$~GHz), one should take into account the additional emission from spinning dust, which is not included here.  This rotational emission is expected to be nearly unpolarized~\\citep{Lazarian+Draine_2000}, which will lower the overall polarization. \\item For models that employ more than one type of grain, which is the case of all our models, their relative contribution to the far-infrared and submillimeter emission will in general vary with wavelength.  If the two types of grain have different values of $C_{\\rm pol}/C_{\\rm abs}$ because of different alignment fractions or grain properties, the degree of polarization will vary with wavelength.  This effect is very pronounced in our models 1 and 3 where only the silicate grains are aligned with the magnetic field. \\end{enumerate}  These last two results are particularly interesting for studies of the polarized components of the CMB.  Polarized dust emission is indeed one of the brightest Galactic foregrounds at microwave frequencies, and being able to remove it from Galactic maps is one of the biggest challenges that will limit our ability to detect B-modes.  The maximum degrees of polarization we calculated in \\S\\ref{sec:ir_em} are consistent with the level inferred from the {\\it WMAP} observations~\\citep{Page+Hinshaw+Komatsu_etal_2007, Gold+Bennett+Hill_etal_2008}, although $5\\%$ is on the lower side of our predictions.  Upcoming experiments, such as {\\it Planck}, might therefore want to allow for the possibility of dust polarization being slightly higher in their sky models, and consider how this would affect their ability to constrain E- and B-modes.  More important is the possible wavelength-dependence of the degree of polarized dust emission.  Models involving both silicate and carbonaceous spheroidal grains lead to a roughly constant level of polarization at all {\\it Planck} frequencies.  However, there is currently no evidence that carbonaceous dust grains contribute to the polarized signal at any wavelength, and models where only silicate grains are aligned spheroids lead to degrees of polarization that are strongly frequency-dependent.  Although a possible challenge for the CMB community, this will allow the use of Planck observations to test whether the submm emission arises from a single kind of grain or from multiple grain types having different values of $\\Cpol/\\Cabs$."
    },
    "0809/0809.3747_arXiv.txt": {
        "abstract": "I present results from an optical spectroscopic investigation of the binary system HD~42401 (V1388 Ori; B2.5 IV-V + B3 V). A combined analysis of $V-$band photometry and radial velocities indicates that the system has an orbital period of 2.18706 $\\pm$ 0.00005 days and an inclination of 75.5 $\\pm$ 0.2 degrees. This solution yields masses and radii of $M_{\\rm 1}$ = 7.42 $\\pm$ 0.08 $M_{\\odot}$ and $R_{\\rm 1}$ = 5.60 $\\pm$ 0.04 $R_{\\odot}$ for the primary and $M_{\\rm 2}$ = 5.16 $\\pm$ 0.03 $M_{\\odot}$ and $R_{\\rm 2}$ = 3.76 $\\pm$ 0.03 $R_{\\odot}$ for the secondary. Based on the position of the two stars plotted on a theoretical H-R diagram, I find that the age of the system is $\\gtrsim$~25 Myr and that both stars appear overluminous for their masses compared to single star evolutionary tracks. A fit of the spectral energy distribution based on photometry from the literature yields a distance to HD~42401 of 832 $\\pm$ 89 parsecs. ",
        "introduction": "Two fundamental parameters in stellar astrophysics are the masses and radii of stars. Especially lacking are accurate ($\\le2\\%$) determinations of these quantities for O- and B-type systems. Data of this kind are important for tests of stellar models, with far-reaching implications such as the modeling of stellar components of galaxies. Here I describe measurements of these parameters for the B-star binary HD~42401. This effort is part of an ongoing program to determine the masses and radii of O- and B-type stars. Previously, I obtained parameters for the LMC O-star binary [L72]~LH~54-425 \\citep{wil08} and will continue with more analyses of eclipsing spectroscopic binary systems containing O- and B-type stars.  The star HD~42401 was first classified as B2 V by \\citet{wal71}. Since then, it has been used as a spectral standard for the B2 V type stars in a number of publications \\citep{wal90,de00,bag01}. The eclipsing nature of its light curve was first noticed with the \\textit{Hipparcos} satellite (HIP 29321), where it was classified as an Algol-type eclipsing binary and given the moniker V1388 Ori \\citep{kaz99}. Three radial velocity measurements of HD~42401 by \\citet{feh97} show a range from --103 to 36 km s$^{-1}$, but no further investigations were made into this probable velocity variable.  Presented here are the analyses of radial velocities from optical spectra obtained in 2008 January~(\\S 2). Section 3 contains a description of the measurement of radial velocities from the data. In \\S 4 I describe tomographic reconstructions of the individual component spectra. Covered in \\S 5 is the combined light curve and radial velocity curve solution. In \\S 6 I finish with a discussion of the distance to the system, its fundamental parameters, and other consequences of the analysis. ",
        "conclusions": "The goal with this work was to determine the masses and radii of the stars in the HD~42401 system to great accuracy. I have determined the masses of the two stars to $1.1 \\%$ and $0.6 \\%$ and the radii of the stars to $0.7 \\%$ and $0.8 \\%$, for the primary and secondary respectively. Armed with these numbers and the results from Table \\ref{ELCparms}, I can consider the details of the HD~42401 system.  As is seen in Table \\ref{ELCparms}, both stars are well within their Roche radii but experience tidal distortion that is evident in the light curve (Fig. \\ref{fig4}). The rotational velocities derived from the tomographic reconstructions (Table \\ref{tomo}) match very well with the synchronous rotation values found by ELC (Table \\ref{ELCparms}), indicating that this system has achieved synchronous rotation, and is therefore not very young.  To make an estimate of the age of the system, I used the effective temperatures from Table \\ref{tomo} and radii from Table \\ref{ELCparms} to plot the two stars of HD~42401 on a theoretical H-R diagram and compare their locations to evolutionary tracks. The result is shown in Figure \\ref{fig7}, plotted against evolutionary tracks for stars of 5, 7, and 9 $M_{\\odot}$ from \\citet{sch92} as well as isochrones from \\citet{lej01} for solar metallicity with ages of 21.9, 25.1, 27.5, and 31.6 Myr. The location of the stars is most consistent with an age of $\\sim$25 Myr. The model tracks shown in Figure~\\ref{fig7} are for non-rotating stellar models. The present ratios of spin angular velocity to critical angular velocity are $\\Omega / \\Omega_{\\rm crit} = 0.45$ and 0.31 for the primary and secondary, respectively (assuming synchronous rotation), and the evolutionary tracks for such moderate rotation rates are only slightly steeper and more extended in time than those illustrated \\citep{eks08}. Thus, our derived age may slightly underestimate the actual value.  Both stars appear (in Fig. \\ref{fig7}) to be overluminous for the derived masses of 7.42 $M_{\\odot}$ for the primary and 5.16 $M_{\\odot}$ for the secondary. Table 3 of \\citet{har88} lists astrophysical parameters for stars as a function of spectral type for main sequence stars based on empirical data from eclipsing binaries. The mass and effective temperature of the primary fit between the listed values for spectral types B2 (mean of 8.6 $M_{\\odot}$) and B3 (mean of 6.1 $M_{\\odot}$), but the radius is much larger than the means for comparable spectral types and matches a B0.5 star (mean of 5.5 $R_{\\odot}$). The mass of the secondary is consistent with the B4 spectral type (mean of 5.1 $M_{\\odot}$) while the effective temperature and radius appear more consistent with the B3 spectral type. In a study of eclipsing binaries in the Small Magellanic Cloud, \\citet{hil05} found several systems of comparable mass that, like HD~42401, are overluminous compared to model predictions. However, our findings for HD~42401 seem to conflict with the results of \\citet{mal03} who shows that early B-type stars that are in close systems and rotate more slowly than single stars are on average smaller than those same single stars. In a subsequent paper, \\citet{mal07} studied well-separated binaries in an effort to use the properties of their component stars for a more direct comparison with single stars. His resulting mass-luminosity-radius relations, when applied to our results for HD~42401, predict less luminous, hotter and smaller components. This is perhaps not surprising, due to the age of HD~42401 and the evolution of its components from the zero age main sequence.  I can also estimate the distance to the system by fitting a spectral energy distribution (SED) to various photometric measurements. The spectra were not flux calibrated, so to create an SED, I must rely on photometry performed on the system. HD~42401 is a relatively bright $(V \\sim 7.4)$ system and has thus been well studied. The problem arises in estimating the phase at which a particular observation in the historical literature was made. Two observations of HD~42401 were taken by $IUE$ at $\\phi = 0.12 \\pm 0.07$ based on our orbital ephemeris. The uncertainty in this estimate comes from the uncertainties in our values for $T_{\\rm IC,1}$ and period and the number of orbits between the $IUE$ observations and our observations. The range in uncertainty for the orbital phase of the $IUE$ data brings it close to eclipse, and this was supported by initial SED fits including the $IUE$ data that indicated a lower UV flux than indicated by other measurements. There are various other measurements that also were sufficiently close to eclipses to be omitted in SED fitting. Fortunately, 2MASS \\citep{cut03} measurements were obtained at  $\\phi = 0.16 \\pm 0.03$ and are therefore far enough away from eclipse, within uncertainties, to be used. The only other points used were Johnson $UBV$ magnitudes. To obtain these values, I used the value at quadrature from our light curve of $V = 7.40$ and applied colors of $(B-V) = -0.042$ and $(U-B) = -0.620$ which are averages of several observations that are listed by \\citet{mer94}. The SED fit is shown in Figure \\ref{fig8} along with the $U, B, V, J, H, K_s$ magnitudes. I used two model spectra from \\citet{lan07} matching our values in Table \\ref{tomo} for the effective temperatures and gravities of each star, and scaled in the blue by the flux ratio of $0.25 \\pm 0.05$ found from the tomographic reconstructions. This fit of the SED results in a limb darkened, angular diameter for the primary of $\\theta_{\\rm LD} = 62.7 \\pm 1.1~\\mu$as with a reddening of $E(B-V) = 0.15 \\pm 0.01$ mag and a ratio of total-to-selective extinction of $R = 3.24 \\pm 0.10$. By directly comparing this angular diameter with the value for the radius of the primary, I estimate the distance of the system to be $d = 832 \\pm 89$ pc. This distance and reddening are in excellent agreement with the values in \\citet{bow08} of $E(B-V) = 0.15 \\pm 0.02$ mag and $d = 0.8$ kpc.  HD~42401 has Galactic coordinates of $\\ell = 197\\fdg64$ and $b = -3\\fdg33$ \\citep{ree05}. This is close to the Galactic open cluster NGC~2169 at $\\ell = 195\\fdg61$ and $b = -2\\fdg93$ (separation $\\sim 2\\fdg07$). \\citet{abt77} found that the earliest spectral type of the cluster members of NGC~2169 was B2 III. \\citet{jef07} find a reddening value of $E(B-V) = 0.20 \\pm 0.01$ mag and a distance of $\\sim1060$ pc to NGC~2169. Proper motions for the objects are similar as well, with HD~42401 having $\\mu_{\\alpha}~cos~\\delta = -2.25 \\pm 1.41$ mas yr$^{-1}$ and $\\mu_{\\delta} = -2.71 \\pm 0.53$ mas yr$^{-1}$ \\citep{esa97} and NGC~2169 having $\\mu_{\\alpha}~cos~\\delta = -2.17 \\pm 1.41$ mas yr$^{-1}$ and $\\mu_{\\delta} = -2.52 \\pm 0.53$ mas yr$^{-1}$ \\citep{kha05}. The systemic velocity derived here for HD~42401 is $\\sim$ 14 km s$^{-1}$ and the average of 8 stars in NGC~2169 is 22 km s$^{-1}$ \\citep{kha05}. While the proximity of HD~42401 to NGC~2169 is intriguing and the fact that the earliest star has a similar spectral type to the components of HD~42401, actual cluster membership is unlikely. The age of NGC~2169 in the literature ranges from $9 \\pm 2$ Myr \\citep{jef07} to 7.8 Myr \\citep{kha05}, too young to comfortably fit with the data for the components of HD~42401 in Figure \\ref{fig7}. \\citet{kha05} determined the corona radius of NGC~2169 by matching stellar densities of the cluster to that of the background and found an angular size of 0.15 degrees. The distance on the sky of HD~42401 from NGC~2169 is about 14 times this radius. The difference in distances to the two objects and their separation on the sky make cluster membership even more unlikely. However, it is possible that a wave of star formation associated with the Orion arm of the Galaxy swept through the region, leaving HD~42401 behind as it moved on and eventually led to the creation of NGC~2169."
    },
    "0809/0809.1268_arXiv.txt": {
        "abstract": "We present results from our {\\it Chandra} and {\\it XMM-Newton} observations of two low-luminosity X-ray pulsators SAX J1324.4--6200 and SAX J1452.8--5949 which have spin-periods of 172 s and 437 s respectively. The {\\it XMM-Newton} spectra for both sources can be fitted well with a simple power-law model of photon index $\\sim$ 1.0. A black-body model can equally well fit the spectra with a temperature of $\\sim$ 2 keV for both sources. During our {\\it XMM-Newton} observations, SAX J1324.4--6200 is detected with coherent X-ray pulsations at a period of $172.86 \\pm 0.02$ s while no pulsations with a pulse fraction greater than 15 \\% (at 98\\% confidence level) are detected in SAX J1452.8--5949. The spin period of SAX J1324.4--6200 is found to be increasing on a time-scale of $\\dot{P}$ = $(6.34 \\pm 0.08) \\times 10^{-9}$ s s$^{-1}$ which would suggest that the accretor is a neutron star and not a white dwarf. Using sub-arcsec spatial resolution of the {\\it Chandra} telescope, possible counterparts are seen for both sources in the near-infrared images obtained with the SOFI instrument on the {\\it New Technology Telescope}. The X-ray and near-infrared properties of SAX J1324.4--6200 suggest it to be either a persistent high mass accreting X-ray pulsar or a symbiotic X-ray binary pulsar at a distance $\\le$ 9 kpc. We identify the infrared counterpart of SAX J1452.8--5949 to be a late-type main sequence star at a distance $\\le$ 10 kpc, thus ruling out SAX J1452.8--5949 to be a high mass X-ray binary. However with the present X-ray and near-infrared observations, we cannot make any further conclusive conclusion about the nature of SAX J1452.8--5949. ",
        "introduction": "In the past ten years, using the Advanced Satellite for Cosmology and Astrophysics ({\\it ASCA}), {\\it BeppoSAX} and the Rossi X-ray Timing Explorer ({\\it RXTE}) satellites, a population of faint (L$_X \\lesssim 10^{36}$ erg s$^{-1}$) X-ray sources has been found in our Galaxy which harbor a slow pulsating source with periods ranging from several seconds to over a thousand seconds. These `slow' pulsators are found to harbor a variety of source types like anomalous X-ray pulsars (isolated slowly rotating neutron stars; e.g. Torii et al. 1998), accreting magnetized white dwarfs (i.e. intermediate polars, AM Her type systems; e.g., Misaki et al. 1996), and neutron stars accreting from a high-mass companion star (i.e., mostly as Be/X-ray transients; e.g., Hulleman et al. 1998). Despite the successes in determining the nature of these slow pulsators, there remains a group of persistent systems  whose nature still has not been determined. The X-ray properties of these sources suggest that most of them are neutron stars accreting from a high mass companion star, however accreting white dwarf cannot be excluded. Furthermore, in some cases, the possibility of a system in which a neutron star accretes from a low-mass companion star also cannot be excluded (Lin et al. 2002). More observations at all wavelengths (i.e., X-ray or near IR) could help to unveil the nature of these pulsators. In this paper, we present the results from our {\\it Chandra}, {\\it XMM-Newton} and {\\it New Technology Telescope} (NTT) observations of two faint X-ray pulsators SAX J1324.4--6200 ($l$ = 306$^{\\circ}.79$, $b$ = 0$^{\\circ}$.60) and SAX J1452.8--5949 ($l$ = 317$^{\\circ}.65$, $b$ = --0$^{\\circ}.46$).  SAX J1324.4--6200 (hereafter SAX13), which has a pulse period of $\\sim$ 170 s, was discovered serendipitously from observations of the low mass X-ray binary (LMXB) XB 1323--619 performed on August 22, 1997 with {\\it BeppoSAX} (Angelini et al. 1998). The X-ray (1.8--10 keV) spectrum of the source could be fitted with either an absorbed power-law with photon index of 1.0 $\\pm$ 0.4 and hydrogen column density, $N_{\\mathrm{H}}$ of $7.8^{+2.7}_{-1.1} \\times 10^{22}$ cm$^{-2}$ or with a black-body model with a temperature, kT of 2.4 $\\pm$ 0.4 keV and $N_{\\mathrm{H}}$ of 4$^{+3}_{-2}$ $\\times$ 10$^{22}$ cm$^{-2}$ (Angelini et al. 1998). SAX13 was also detected in the observations performed on XB 1323--619 using {\\it ASCA} on August 4, 1994 (Angelini et al. 1998). In addition, pointed {\\it ASCA} observations (of 187 ks) were performed on SAX13 on February 2, 2000 and the source was detected with a pulse period of $\\sim$ 171.2 s (Lin et al. 2002). Lin et al. (2002) found a possible orbital period of the system of 27 $\\pm$ 1 hour and suggested that the system could be a LMXB pulsar. Recently, during short observations performed using {\\it Swift} on December 30, 2007, SAX13 was detected  with a pulse period of 172.8 s (Mereghetti et al. 2008) which would imply a spin down over the last 10 years of  $\\dot{P}$  = $\\sim$ 6 $\\times$ 10$^{-9}$ s s$^{-1}$. Mereghetti et al. (2008) identified a possible 2MASS near-infrared counterpart in the {\\it Swift} error circle of SAX13 with a {\\it K} band magnitude of 14.39 $\\pm$ 0.08 and suggested that the source could be a persistent Be accreting X-ray pulsar.  SAX J1452.8--5949 (hereafter SAX14) was discovered using {\\it BeppoSAX} with a spin period of $\\sim$ 437 s (Oosterbroek et al. 1999). The X-ray spectra of the source could be fitted well with an absorbed power-law model with a photon index of 1.4 $\\pm$ 0.6. Oosterbroek et al. (1999) suggested that SAX14 could be an accreting Be X-ray pulsar at a distance of $6 -12$ kpc with the luminosity  of $\\sim 10^{34}$ erg s$^{-1}$.   \\begin{table*} \\label{tab:observation} \\centering \\caption{Log of the X-ray observations of SAX J1324.4--6200 and SAX J1452.8--5949.} \\label{observations_log} \\begin{tabular}{|ccccccc|} \\hline Object & Telescope/Instrument & Date & ObsId & Total Observation     \\\\ &           &   (UT)   &       &  Span     (ks)          \\\\ \\hline SAX J1324.4--6200&{\\it Chandra}/ACIS-I  &   28 Nov 2007  &  9012      &  1.1  \\\\ SAX J1452.8--5949 &{\\it Chandra}/ACIS-I  &   30 Dec 2007  &  9014       &  1.0  \\\\ SAX J1324.4--6200&{\\it XMM-Newton}/EPIC &   11 Jan 2008  & 0511010201  & 19.3\\\\ SAX J1452.8--5949 &{\\it XMM-Newton}/EPIC &   07 Feb 2008  & 0511010501   &  6.9 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We carried out observations with the {\\it Chandra} and {\\it XMM-Newton} satellites to investigate the nature of the accreting X-ray pulsators SAX J1324.4--6200 and SAX J1452.8--5949. In addition, we obtained near-infrared data using the ESO-NTT to search for infrared counterparts for these pulsators.  \\subsection{SAX J1324.4--6200}  During our {\\it XMM-Newton} observation, SAX13 is detected with a pulse period of $172.85 \\pm 0.02$ s, consistent with the pulse period obtained by Mereghetti et al. (2008) from observations made by {\\it Swift} on December 30, 2007, just 10 days before our {\\it XMM-Newton} observations. The spin period history of SAX13 clearly shows a linear increase in pulse period from 170.35 to 172.86 s over the last 14 years (Table \\ref{p_observation}) with spin period derivative  $\\dot{P}$ = $(6.34 \\pm 0.08) \\times 10^{-9}$ s s$^{-1}$. The spin period of high mass X-ray binary (HMXB) pulsars range from a few seconds to a few hundred seconds and display several spin-up/down trends on time-scales ranging from days to years (Bildsten et al. 1997). For example, 4U 1907+09, a persistent supergiant X-ray binary with $P_{\\mathrm{s}}$ = 440 s, is spinning down for more than 15 years with $\\dot{P}$ =  $7.3 \\times 10^{-7}$ s s$^{-1}$ (Baykal et al. 2001). Among LMXB pulsars, GX 1+4 has the longest spin period of $\\sim$ 141 s. LMXB pulsars show a wide range of spin-period derivatives ranging from $\\sim 10^{-8}$ to $\\sim 10^{-11}$ s s$^{-1}$ (Ferrigno et al. 2007; Chakrabarty et al. 1997). Intermediate Polars (IPs) are also slow pulsators with spin-period of few hundred seconds and usually show spin up phases with a typical spin period derivative $\\sim -7 \\times 10^{-11}$ s s$^{-1}$ (Patterson 1994). However there are a very few IPs which are found to be spinning down and the fastest spinning down intermediate polar PQ Gem is observed with $\\dot{P} = 1.1 \\times 10^{-10}$ s s$^{-1}$ (Mason 1997).  The measured spin torque can be used as a clue on the nature of the compact accretor in SAX13. The usual expression for the accretion torque on a compact accretor is : $I\\dot{\\omega}=\\dot{M}\\sqrt{0.5GMr_{A}}$ (Lipunov 1992), where $I$ and $\\dot{\\omega}$ is the moment of inertia and the spin torque of the compact accretor respectively, $M$ and $\\dot{M}$ is the mass and the mass accretion rate of the compact accretor respectively, $r_{A}$ is the Alfven radius and $G$ is the Gravitational constant. $r_{A}$ for a compact accretor is $\\propto$ $\\mu^{4/7}M^{-1/7}\\dot{M}^{-2/7}$, where $\\mu$ is the magnetic moment of it. Thus the external torque applied from the accretion disk onto the compact object scales as $\\propto M^{-3/7}R^{12/7}B^{2/7}$, where B (= $\\mu R^3$) is the magnetic field of the compact object and for a given X-ray luminosity, the value of $\\dot{M}$ is $\\propto$ $\\frac{R}{M}$. The value of $M$, $R$ and $B$ can be fixed for a typical white dwarf in an intermediate polar and an accreting neutron star. Thus, given a certain observed X-ray luminosity, the ratio of a spin torque for a white dwarf and a neutron star would be : \\begin{equation} \\frac{\\dot{\\omega}_{wd}}{\\dot{\\omega}_{ns}} = \\frac{M_{wd}^{-3/7}R_{wd}^{12/7}B_{wd}^{2/7}}{I_{wd}} \\times \\frac{I_{ns}}{M_{ns}^{-3/7}R_{ns}^{12/7}B_{ns}^{2/7}} \\end{equation}  where the subscript `wd' and `ns' refer to white dwarf and neutron star respectively.  A white dwarf has a moment of inertia $\\approx 4-5$ order of magnitude larger than a neutron star. For the typical parameters of white dwarf and neutron star : $R_{wd} = 10^{4}$ km, $B_{wd} = 10^{5}$ G and $M_{wd}$ = $0.1 M_{\\odot}$; $R_{ns} = 10$ km, $B_{ns} = 10^{8}$ G and $M_{ns}$ = $1.4 M_{\\odot}$, the ratio between the expected accretion torques for the two compact objects is $\\approx 0.1 - 1$,  implies that we cannot distinguish between a white dwarf and a neutron star system on the basis of the spin torque. However, for a high magnetic field neutron star ($B\\approx 10^{12}$ G) the ratio between the expected accretion torques for the two compact objects is $\\approx 10^{-2}$ which implies that a high B field neutron star system has a spin torque two orders of magnitude larger than a white dwarf. Since in SAX13, the observed spin down is comparable with measures of spin torques made for high B field accreting neutron stars, the presence of a neutron star is favored, unless the accretor is an intermediate polar with a high B field and a quite small mass.  The {\\it XMM-Newton} spectra of SAX13 fits well with both power-law and black-body models. The spectral parameters of SAX13 (Table \\ref{parameters}) are typical of an accreting neutron star X-ray pulsar with a high magnetic field. Accreting X-ray pulsars with low magnetic field strength usually display soft X-ray spectra with photon index $\\ge$ 2.0 (Bildsten et al. 1997). IPs display hard X-ray spectra similar to high mass X-ray binary pulsars but they also show a strong iron K$\\alpha$ emission line (e.g., Norton et al. 1991, Muno et al. 2004). The absence of Fe K$\\alpha$ emission line in the SAX13 X-ray spectra further makes it very unlikely to be an IP.  With the {\\it Swift} spatial resolution, Mereghetti et al. (2008) identified the possible infrared counterpart of SAX13 in the 2MASS K$_S$ band image with magnitude 14.39 $\\pm$ 0.08. Figure \\ref{fig:sax13_chart} shows the ESO-NTT K$_S$ band image of SAX13, zoomed-in at the position of SAX13. As can be seen, the single 2MASS star is resolved into three stars `C1', `C2' and `C3' in an NTT K$_S$ band image. The position of star `C1' is consistent with the {\\it Chandra} error circle of SAX13, and therefore `C1' is the most likely near-infrared counterpart of SAX13.  Using the observed $N_{\\mathrm{H}}$ = 5.8 $\\times 10^{22}$ cm$^{-2}$ (Table \\ref{parameters}), and the relation A$_V$/$N_{\\mathrm{H}}$ = $5.6 \\times 10^{-22}$ mag cm$^2$ (Predehl et al. 1995), the corresponding A$_V$ for SAX13 is calculated to be 32.54 mag. Using the above A$_V$, we calculated the near-infrared extinction in $J$, $H$ and $K_{\\mathrm{s}}$ waveband to be 9.01, 5.56 and 3.77 mag (Fitzpatrick 1999). Finally, the dereddened magnitudes of SAX13 are found to be  J = $10.56 \\pm 0.11$ mag, H = $11.05 \\pm 0.09$ mag and K = $11.20 \\pm 0.11$ mag. Here, we do not take into account the uncertainties in the A$_V$. Also, we don't deny that a part of the X-ray absorption density could be local to the X-ray source, and not necessarily apply to the companion star.  \\begin{figure*} \\begin{tabular}{c} \\includegraphics[height=6cm,width=16cm]{f6.epsi} \\end{tabular} \\caption{Observed (open circles) and extinction-free (filled circles) spectral energy densities of near-infrared counterpart of SAX1324.4--6200 ({\\it{left}}) and SAX1452.8--5949 ({\\it{right}}). The solid line and dashed-dotted line represents black-body model spectral energy density of the star at temperature, T = 20,000K and T = 3,000K respectively and are shifted to match the extinction-free spectral energy density of the star at near-infrared $H$ waveband.} \\label{ir_s13} \\end{figure*}  With the above extinction-free {\\it J, H and K$_{\\mathrm{s}}$} magnitudes of SAX13, we used a black-body model to estimate the distance to different type of stars (main-sequence stars, supergiants and giants), for a given temperature and a radius (Cox 2000). With these calculations, we found that any late-type main sequence star (M through A type) would lie at a distance $\\le$ 1.4 kpc. Thus, it is possible that the system could be an IP at a distance of $\\le$ 1.4 kpc and in that case, the corresponding X-ray luminosity of the system would be $\\le 10^{33}$ erg s$^{-1}$. However, it is very unlikely to find a LMXB pulsar at such a low X-ray luminosity. Also it is not very likely that a supergiant is the infrared counterpart of SAX13 as it would lie outside the Galaxy for the given magnitudes. A main sequence early-type star (e.g., B-type) with  temperature of 11,000 - 30,000 K and the radius of 3.0 - 7.5 $R_{\\odot}$ would lie at a distance $\\le$ 8 kpc, indicating SAX13 could be an accreting neutron star high mass X-ray binary pulsar. However, for the given flux densities, a late-type giant would also be well within the Galaxy at a similar distance.  To further confirm the above argument, we approximated the black-body model to the extinction-free near-infrared fluxes of SAX13 (shown as open circles in Fig \\ref{ir_s13}). Black-body curves for temperature, T = 20,000 K and T = 3,000 K  are shown in Fig \\ref{ir_s13} as a solid and a dashed-dotted line respectively and are shifted to match the spectral flux density of SAX13 at {\\it H} wave-band. The solid line clearly indicates that SAX13 is an early-type main-sequence star of temperature $\\ge$ 20,000 K. For the sake of completeness, the observed fluxes of SAX13 are also shown in the same figure as filled circles.  Based on the above arguments, it is therefore, possible that SAX13 is indeed an accreting neutron star high mass X-ray binary pulsar at a distance of 1.5 - 8.0 kpc. For the observed X-ray flux of SAX13 (Table \\ref{parameters}), the luminosity of the system would be 1.2 $\\times 10^{33}$ to 3.4 $\\times 10^{34}$ erg s$^{-1}$. Most of the high mass X-ray pulsars (especially Be X-ray pulsars) are transient in nature and have been observed at luminosities varying from 10$^{33}$ to 10$^{38}$ erg s$^{-1}$. However SAX13 has been persistent for the past 14 years. Thus it is possible that SAX13 is a persistent Be X-ray binary pulsar, similar to X-Persei (La Palombara et al. 2007; Reig et al. 1998). These systems might be part of an unusual class of accreting neutron stars with high mass companions which have long orbital periods ($>$ 30 days) and low eccentricities (e.g., Pfahl et al. 2002). Such long orbital period indicate that tidal circularisation cannot yet have occured after the supernova explosion that created the neutron star. Therefore, these low eccentricities must be primordial and the supernova explosion could not  have been accompanied by a kick to the neutron star (Pfahl et al. 2002).  As suggested above, a late-type giant can also be an infrared counterpart of SAX13, and in that case, SAX13 would be a member of symbiotic X-ray binary pulsars. Symbiotic X-ray binary pulsars are recently identified as a new class of LMXBs, in which a neutron star rotates in a wide orbit around a M-type giant, accretes matter either from the wind of a M giant or via roche lobe overflow (Iben et al. 1996). Symbiotic X-ray binary pulsars are characterized by faint X-ray emission ranging between $\\sim$ 10$^{32}$ to $\\sim$ 10$^{34}$ erg s$^{-1}$, long pulse periods and long spin-up/down phases (Masetti et al. 2007; Corbet et al. 2008). Thus SAX13 could also be  a symbiotic X-ray binary pulsar, at a distance  $\\le$ 9 kpc. These systems are expected to be rare, owing to the short duration of both the K/M giant lifetime and the highly unstable mass transfer stage expected for Roche lobe overflow for LMXBs with P$_{\\mathrm{orb}}$ $\\ge$ 2 days (Kalogera et al. 1996).  With the present X-ray and near-infrared observations of SAX13, we are unable to distinguish it to be a high mass accreting X-ray pulsar or the symbiotic X-ray binary pulsar. However, infrared spectroscopic observations would further help to constrain the spectral type of the infrared counterpart of SAX13.  \\subsection{SAX 1452.8--5949}  We detected a faint star in the ESO-NTT images of SAX14, consistent with the {\\it Chandra} error circle, shown in Figure \\ref{fig:sax13_chart}. Therefore we identify `S1' as the infrared counterpart of SAX14. Taking $N_{\\mathrm{H}}$ = $1.22 \\times 10^{22}$ cm$^{-2}$ (Table \\ref{parameters}), we calculated the dereddened magnitudes of SAX14 to be J = 16.66 $\\pm$ 0.12 mag, H = 16.45 $\\pm$ 0.12 mag and K = 17.15 $\\pm$ 0.12 mag. We calculated the distance of the possible infrared counterpart with the same method used for SAX13 (see \\S 5.1) We can rule out a supergiant, giant or an O/B-type star as the infrared counterpart of SAX14 as, given the observed flux densities, it would lie outside the Galaxy. This rules out the possibility that SAX14 is a HMXB. Only a late spectral type main sequence star (M through A type) can be in the Galaxy, at a distance $\\le$ 10 kpc. A black-body approximation to the extinction-free near-infrared fluxes of SAX14 indicate that SAX14 would have temperature $<$ 20,000 K (Fig \\ref{ir_s13}). However, in case, a part of the X-ray absorption density is local to the X-ray source, we would expect even lesser extinction in near-infrared wavebands and in that case, the temperature of near-infrared counterpart of SAX14 would be much less than 20,000 K. With the above arguments, we can say that SAX14 must be a binary with a low mass companion (LMXB or IP), regardless if there are pulsations or not.  The {\\it XMM-Newton} spectrum of SAX14 is well fitted with a power-law model of photon index 0.83 (Table \\ref{parameters}). Both an accreting pulsar and a white dwarf are consistent with the inferred X-ray spectral parameters, so we cannot distinguish from the X-ray spectral information whether SAX14 is a neutron star or an accreting white dwarf. Also we did not detect any strong Fe 6.4 keV line in the X-ray spectra indicates that it is very unlikely that the system is an IP.  No significant pulsations were detected in 0.2 - 12 keV energy band for SAX14 with an upper limit of 15\\% on the fractional amplitude at 98\\% confidence level. This result is in contrast with the previous detection made by Oosterbroek et al. (1999) who detected pulsations with a fractional amplitude of $75\\pm 25\\%$. To explain this discrepancy we are left with two possibilities: the fractional amplitude of the pulsations decreased or the pulsations observed by Oosterbroek et al. (1999) were spurious. In the later case, SAX14 can be any non pulsating source in the Galaxy, like LMXB, accreting white dwarf, BY Dra, RS CVn or active star, etc.  If on the other hand the detection of pulsations by Oosterbroek et al. (1999) was real, then the pulse amplitude has to be reduced by a substantial amount. This behavior has been observed in other LMXBs also where the pulsed fraction decreased on the time-scale of days (Her X-1; Ramsay et al. 2002), hours (GX 1+4; Naik et al. 2005) and minutes (4U 1907+09; In'T Zand et al. 1997). Given the large value for the upper limit we obtained, the possibility that SAX14 is a slow pulsating LMXB or an IP is still open.  With all the above arguments, we suggest that SAX14 is not a high mass X-ray binary pulsar. However, the X-ray spectral, timing properties and infrared flux densities suggest to be a low mass companion (LMXB or IP), regardless if there are pulsations or not.    \\begin{table} \\label{tab:p_observation} \\centering \\caption{Spin-period history of SAX J1324.4-6200} \\label{p_observation} \\begin{tabular}{|llll|} \\hline Telescope      & Date         & Spin period & References   \\\\ &       (UT)        & (s)  &             \\\\ \\hline {\\it ASCA}       & 04 Aug 1994   & 170.35  $\\pm$ 0.48    & a   \\\\ {\\it BeppoSAX}   & 22 Aug 1997   & 170.84  $\\pm$ 0.04    & a   \\\\ {\\it ASCA}       & 02 Feb 2000   & 171.25  $\\pm$ 0.01   & b   \\\\ {\\it Swift}      & 30 Dec 2007   & 172.84  $\\pm$ 0.1     & c   \\\\ {\\it XMM-Newton} & 11 Jan 2008   & 172.86  $\\pm$ 0.02    & d   \\\\ \\hline \\end{tabular} \\flushleft REFERENCES. -- (a) Angelini et al. 1998; (b) Lin et al. 2002; (c) Mereghetti et al. 2008; (d) present paper. \\end{table}"
    },
    "0809/0809.1574_arXiv.txt": {
        "abstract": "{\\ion{Fe}{xxv} lines at 1.85~\\AA\\ (6.70~keV) and nearby \\ion{Fe}{xxiv} satellites have been widely used for determining the temperature of the hottest parts of solar flare and tokamak plasmas, though the spectral region is crowded and the lines are blended during flare impulsive stages.}{The aim of this work is to show that similarly excited Fe lines in the 7.7--8.6~keV (1.44--1.61~\\AA) region have the same diagnostic capability with the advantage of not being so crowded. } {Spectra in the 7.7--8.6~keV range are synthesized using the CHIANTI spectral package for conditions (temperature, turbulent velocities) appropriate to solar flares. } {The calculated spectra show that the Fe lines in the 7.7--8.6~keV are well separated even when turbulent velocities are present, and \\ion{Fe}{xxiv}/\\ion{Fe}{xxv} line ratios should therefore provide valuable tools for diagnosing flares and tokamak plasmas. }{Fe lines in the 7.7--8.6~keV range are ideal for the measurement of flare temperature and for detecting the presence of low-energy nonthermal electrons present at flare impulsive stages. An indication of what type of instruments to observe this region is given.} ",
        "introduction": "The diagnostic potential of highly ionized iron lines at $\\sim 1.9$~\\AA\\ ($\\sim 6.7$~keV) was recognized by \\cite{gab72} who pointed out that \\ion{Fe}{xxiv} satellite lines with transitions $1s^2 nl - 1s 2p nl$ ($n\\geqslant 2$), with upper levels formed by dielectronic recombination of He-like iron (Fe$^{+24}$), have fluxes which, relative to the \\ion{Fe}{xxv} resonance line $w$ ($1s^2\\,^1S_0 - 1s2p\\,^1P_1$) at 1.851~\\AA\\ (6.699~keV), depend on electron temperature $T_e$, roughly as $T_e^{-1}$. This has proved extremely useful for understanding the hottest parts of solar flare plasmas \\citep{dos90} and tokamak plasmas \\citep{bit79}. The most intense $n=2$ satellites include (in the notation of \\cite{gab72}) $j$ (1.867~\\AA, 6.642~keV) and $k$ (1.864~\\AA, 6.652~keV) which are resolved from other nearby lines by crystal spectrometers with spectral resolution 0.001~\\AA\\ (4~eV). The most intense $n=3$ satellites include the $d13$, $d15$ pair at 1.852~\\AA\\ (6.695~keV), which are on the long-wavelength side of line $w$ \\citep{bel79a}, while $n\\gtrsim 4$ satellites are progressively less separated from the $w$ line \\citep{bel79b}. At the onset of solar flares, X-ray lines show a broadening which is commonly attributed to turbulence resulting from the impact of nonthermal electrons accelerated at the flare impulsive stage. The broadening may amount to 0.03~\\AA\\ (0.11~keV, equivalent to a few hundred km~s$^{-1}$) which leads to a blurring of the satellite line structure, so determinations of temperatures become less certain. Also, satellites with $n\\geqslant 3$ on the long-wavelength side of the \\ion{Fe}{xxv} line $w$ including the $d13$, $d15$ pair may be indistinguishable from $w$. A method for searching for low-energy, nonthermal electrons proposed by \\cite{gab79}, involving the comparison of temperatures derived from the ratios $j/w$ and $(d13+d15)/w$, is thus made difficult to apply. In addition to the turbulent broadening, a short-wavelength (``blue-shifted\") component is generally present for disk flares, making the determination of temperatures from line flux ratios even less certain.  While the spectral region 1.851--1.940~\\AA\\ (6.391--6.699~keV), which includes all dielectronic satellites emitted by \\ion{Fe}{xxiv} and lower ionization stages of Fe, has been very well studied with high-resolution spectrometers, the region 1.44--1.60~\\AA\\ (7.7--8.6~keV) has hardly received any attention. This range contains higher-$n$ \\ion{Fe}{xxv} lines with transitions $1s^2 \\,^1S_0 - 1snp\\,^1P_1$ ($n=3$, 4, 5 etc.) and associated \\ion{Fe}{xxiv} satellites with transitions $1s^2 2p - 1s 2p nl$ ($n\\geqslant 3$) which are well separated from each other, even when turbulent velocities of a few hundred km~s$^{-1}$ are present. The purpose of this paper is to point out that these lines are more suitable for the determination of temperatures from the hottest part of solar flares, and may prove particularly useful for searching for nonthermal electrons during flare impulsive stages. ",
        "conclusions": ""
    },
    "0809/0809.1432_arXiv.txt": {
        "abstract": "Using detailed Monte Carlo radiative transfer simulations in realistic models for galactic nuclei, we investigate the influence of interstellar dust in ionized gas discs on the rotation curves and the resulting black hole mass measurements. We find that absorption and scattering by interstellar dust leaves the shape of the rotation curves basically unaltered, but slightly decreases the central slope of the rotation curves. As a result, the ``observed'' black hole masses are systematically underestimated by some 10 to 20\\% for realistic optical depths. We therefore argue that the systematic effect of dust attenuation should be taken into account when estimating SMBH masses using ionized gas kinematics. ",
        "introduction": "Measuring the kinematics of ionized gas disks has become one the most important methods to determine the masses of supermassive black holes (SMBHs) in the nuclei of nearby galaxies. The state-of-the-art modeling techniques are refined to a high degree and take into account the major relevant astrophysical processes as well as the main instrumental effects. Applying such modeling techniques to high-quality HST data, a formal SMBH mass measurement uncertainty of 25\\% or better has been claimed with well-behaved disks (e.g. \\cite{Bar03}, \\cite{Mar03}, \\cite{Atk05}, \\cite{Har06}). Many ionized gas disks contain substantial amounts of interstellar dust. As optical H$\\alpha$ radiation is easily absorbed and scattered by interstellar dust grains, these effects might affect the observed rotation curve and hence the SMBH mass estimate. The goal of our investigation is to determine the importance of this effect through detailed Monte Carlo simulations. ",
        "conclusions": "We find that absorption and scattering by interstellar dust leaves the shape of the rotation curves basically unaltered, but slightly decreases the central slope of the rotation curves. Since the central slope can be used as a proxy for the SMBH mass, this result means that ``observed'' black hole masses are systematically underestimated by some 10 to 20\\% for realistic optical depths. Our study demonstrates that the systematic effect of dust attenuation should be taken into account in SMBH demographics studies."
    },
    "0809/0809.4021_arXiv.txt": {
        "abstract": "The primary concern of this contribution is that accreting millisecond pulsars (AMXPs) show a much larger amount of information than is commonly believed.  The three questions to be addressed are: \\\\ 1. Is the apparent spin torque observed in AMXPs real ? \\\\ 2. Why do we see correlations and anti-correlations between fractional amplitudes and timing residuals in some AMXPs ? \\\\ 3. Why the timing residuals, the lightcurve and the 1Hz QPO in SAX J1808.4$-$3658 are related ? \\\\ ",
        "introduction": "There has been a large debate in the accreting millisecond pulsar community, on how to interpret the apparent measurement of spin torques in accreting millisecond pulsars (AMXPs).  The need of interpreting what the apparent spin torque really is needs hardly be explained in this meeting, and this is also the reason why I call it 'apparent' spin torque. An improved understanding of the nature of the apparent spin torque can have enormous implications to the knowledge of neutron star physics.  For example the existence (or the lack) of a spin torque has consequences in the gravitational wave physics\\citep{watts1} \\citep{watts2}, on the problem of the break up spin frequency of neutron stars\\citep{Chakrabarty} and on the strength of their magnetic fields. At the time of writing this contribution, we know ten AMXPs, seven of which show persistent pulsations, and other three are the so called intermittent AMXPs (see for example the contribution of Altamirano \\& Casella and Galloway for an observational overview at this meeting, and \\citep{galloway}\\citep{casella} \\citep{gavriil}\\citep{altamirano}\\citep{patruno}).  The primary target of this contribution are the seven persistent AMXPs. In particular I will discuss the apparent spin torque measurement in XTE J1751-305 and XTE J1807-294 and show how a different and correct treatment of the statistical errors can completely change the interpretation of the observational results.  Related to the problem of the apparent spin torque measurement is the interpretation of the relation between the fractional amplitude of the pulsations and the residuals of the pulse time of arrivals (TOAs).  The residuals here refer to the coherent timing analysis of the AMXPs. In a few words: one measures the TOAs of the pulsations and then fits a model (usually a keplerian circular orbit plus a constant spin frequency or a spin torque) to the TOAs. What is left after the fit are the residuals, and they represent the relative location, with respect to an arbitrary reference position, of the hot spot on the neutron star surface, i.e., the relative pulse phases. I will call this relation ``the FAR relation'' throughout this contribution (FAR$=$Fractional Amplitude vs. Residuals).  I will show that this relation exists in two sources, XTE J1807-294 and XTE J1751-305, and that it is probably a more general phenomenon among AMXPs.  But why the existence of the FAR relation is connected with the interpretation of the apparent spin torque ? The reason is quite simple: we measure the apparent spin torque from the TOAs, and therefore whatever its origin is, it must be related in some way with the fractional amplitudes of the pulsations. Finally, I will discuss some intriguing relations between the lightcurve, the periodic timing and the aperiodic timing of the AMXP SAX J1808.4$-$3658. This AMXP shows some strange shifts of the pulse phases in the timing residuals that cannot be explained even with a torque model. Intriguingly, when the pulse phases start to shift the slope of the lightcurve changes. And when the phases stop to shift, the unexplained phenomenon of the 1 Hz QPO appears. The implications of the relation between such different phenomena are profound: if the unmodeled phase shifts are related with the lightcurve and the aperiodic variability, then a deep link between the accretion disk (responsible for the lightcurve and perhaps the QPO) and the surface of the neutron star (pulsations) exists. Is it possible that all these phenomena (apparent spin torque, FAR relation, lightcurve-periodic-aperiodic variability link) have a common unified explanation ? The search for an answer to this question is the motivation of this work. ",
        "conclusions": "We conclude that, when calculating a timing solution, all these effect must be taken into account, otherwise we can confuse an effect of the magnetospheric-disk interaction with a surface effect due to, for example, a spin torque.  Of course the final question is how do we take these effects into account when doing a coherent timing analysis. An answer here cannot be satisfactory, since all the described phenomena have not a clear theoretical explanation, yet. Certainly they all show that the quadratic component usually identified with a spin torque has not a higher dignity among the other terms of the polynomial expansion used in standard coherent timing techniques. Indeed it appears somehow arbitrary to treat the lowest order variation (quadratic component) as a distinct entity with respect to all the other phase shifts observed. This is particularly evident after proving that the phases of J1808 are influenced by the accretion disk-magnetospheric interaction. Another aspect to be considered in AMXP physics is the possibility that the phase variations we observe all come from the motion of the hot spot. There are several simulations and models \\citep{romanova03}\\citep{romanova04}\\citep{romanova08},\\citep{lamb08} that predict a moving hot spot around a quasi-equilibrium position on the neutron star surface. These movements can lead to an increase of the fractional amplitude directly anti-correlated with the time of arrival of the pulsations. However why the noisy AMXPs like J1807 and J1814 show an anti-correlation while J1751 an J00291 do not, is not easily explainable. Has the presence of a strong 1st overtone some influence on the properties of the AMXPs ?  Another important point is that, beside the theoretical model required to explain the observations, one has to face a more serious problem that is how to extract the correct informations from the data and very little can be said until a satisfactory technique is developed.  \\begin{theacknowledgments} I'd like to thank A. Watts, J. Hartman, M. Klein-Wolt, Rudy Wijnands, Michiel van der Klis and D. Chakrabarty, without whom this work would not have been possible.  I thank also A. Papitto, and A. Riggio for useful discussions we had during this enjoyable meeting. \\end{theacknowledgments}"
    },
    "0809/0809.3437_arXiv.txt": {
        "abstract": "We present further development and the first public release of our multimodal nested sampling algorithm, called {\\sc MultiNest}. This Bayesian inference tool calculates the evidence, with an associated error estimate, and produces posterior samples from distributions that may contain multiple modes and pronounced (curving) degeneracies in high dimensions. The developments presented here lead to further substantial improvements in sampling efficiency and robustness, as compared to the original algorithm presented in Feroz \\& Hobson (2008), which itself significantly outperformed existing MCMC techniques in a wide range of astrophysical inference problems. The accuracy and economy of the {\\sc MultiNest} algorithm is demonstrated by application to two toy problems and to a cosmological inference problem focussing on the extension of the vanilla $\\Lambda$CDM model to include spatial curvature and a varying equation of state for dark energy.  The {\\sc MultiNest} software, which is fully parallelized using MPI and includes an interface to CosmoMC, is available at {\\tt http://www.mrao.cam.ac.uk/software/multinest/}. It will also be released as part of the SuperBayeS package, for the analysis of supersymmetric theories of particle physics, at {\\tt http://www.superbayes.org} ",
        "introduction": "\\label{sec:intro}  Bayesian analysis methods are already widely used in astrophysics and cosmology, and are now beginning to gain acceptance in particle physics phenomenology. As a consequence, considerable effort has been made to develop efficient and robust methods for performing such analyses. Bayesian inference is usually considered to divide into two categories: parameter estimation and model selection. Parameter estimation is typically performed using Markov chain Monte Carlo (MCMC) sampling, most often based on the standard Metropolis--Hastings algorithm or its variants, such as Gibbs' or Hamiltonian sampling (see e.g. \\citealt{MacKay}). Such methods can be very computationally intensive, however, and often experience problems in sampling efficiently from a multimodal posterior distribution or one with large (curving) degeneracies between parameters, particularly in high dimensions. Moreover, MCMC methods often require careful tuning of the proposal distribution to sample efficiently, and testing for convergence can be problematic.  Bayesian model selection has been further hindered by the even greater computational expense involved in the calculation to sufficient precision of the key ingredient, the Bayesian evidence (also called the marginalized likelihood or the marginal density of the data). As the average likelihood of a model over its prior probability space, the evidence can be used to assign relative probabilities to different models (for a review of cosmological applications, see \\citealt{Mukherjee06}).  The existing preferred evidence evaluation method, again based on MCMC techniques, is thermodynamic integration (see e.g.  \\citealt{Ruanaidh}), which is extremely computationally intensive but has been used successfully in astronomical applications (see e.g. \\citealt{Hobson03, Marshall03, Slosar03, Niarchou04, Bassett04, Trotta05, Beltran05, Bridges06a}). Some fast approximate methods have been used for evidence evaluation, such as treating the posterior as a multivariate Gaussian centred at its peak (see e.g. \\citealt{Hobson02}), but this approximation is clearly a poor one for multimodal posteriors (except perhaps if one performs a separate Gaussian approximation at each mode). The Savage--Dickey density ratio has also been proposed (Trotta 2005) as an exact, and potentially faster, means of evaluating evidences, but is restricted to the special case of nested hypotheses and a separable prior on the model parameters. Various alternative information criteria for astrophysical model selection are discussed by \\citet{Liddle07}, but the evidence remains the preferred method.  Nested sampling (\\citealt{Skilling04}) is a Monte Carlo method targetted at the efficient calculation of the evidence, but also produces posterior inferences as a by-product. In cosmological applications, \\citet{Mukherjee06} showed that their implementation of the method requires a factor of $\\sim 100$ fewer posterior evaluations than thermodynamic integration. To achieve an improved acceptance ratio and efficiency, their algorithm uses an elliptical bound containing the current point set at each stage of the process to restrict the region around the posterior peak from which new samples are drawn. \\citet{Shaw07} point out that this method becomes highly inefficient for multimodal posteriors, and hence introduce the notion of clustered nested sampling, in which multiple peaks in the posterior are detected and isolated, and separate ellipsoidal bounds are constructed around each mode.  This approach significantly increases the sampling efficiency. The overall computational load is reduced still further by the use of an improved error calculation (\\citealt{Skilling04}) on the final evidence result that produces a mean and standard error in one sampling, eliminating the need for multiple runs. In our previous paper (\\citealt{feroz08} -- hereinafter FH08), we built on the work of \\citet{Shaw07} by pursuing further the notion of detecting and characterising multiple modes in the posterior from the distribution of nested samples, and presented a number of innovations that resulted in a substantial improvement in sampling efficiency and robustness, leading to an algorithm that constituted a viable, general replacement for traditional MCMC sampling techniques in astronomical data analysis.  In this paper, we present further substantial development of the method discussed in FH08 and make the first public release of the resulting Bayesian inference tool, called {\\sc MultiNest}. In particular, we propose fundamental changes to the `simultaneous ellipsoidal sampling' method described in FH08, which result in a substantially improved and fully parallelized algorithm for calculating the evidence and obtaining posterior samples from distributions with (an unkown number of) multiple modes and/or pronounced (curving) degeneracies between parameters. The algorithm also naturally identifies individual modes of a distribution, allowing for the evaluation of the `local' evidence and parameter constraints associated with each mode separately.  The outline of the paper is as follows. In Section \\ref{sec:bayesian_infer}, we briefly review the basic aspects of Bayesian inference for parameter estimation and model selection. In Section \\ref{sec:nested}, we introduce nested sampling and discuss the use of ellipsoidal bounds in Section \\ref{sec:ellipsoidal_sampling}. In Section \\ref{sec:improved_ellipsoidal}, we present the {\\sc MultiNest} algorithm. In Section \\ref{sec:applications}, we apply our new algorithms to two toy problems to demonstrate the accuracy and efficiency of the evidence calculation and parameter estimation as compared with other techniques. In Section \\ref{sec:cosmology}, we consider the use of our new algorithm for cosmological model selection focussed on the extension of the vanilla $\\Lambda$CDM model to include spatial curvature and a varying equation of state for dark energy.  We compare the efficiency of {\\sc MultiNest} and standard MCMC techniques for cosmological parameter estimation in Section \\ref{sec:MCMC_comparison}. Finally, our conclusions are presented in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have described a highly efficient Bayesian inference tool, called {\\sc MultiNest}, which we have now made freely available for academic purposes as a plugin that is easily incorporated into the {\\sc CosmoMC} software.  On challenging toy models that resemble real inference problems in astro- and particle physics, we have demonstrated that {\\sc MultiNest} produces reliable estimates of the evidence, and its uncertainty, and accurate posterior inferences from distributions with multiple modes and pronounced curving degeneracies in high dimensions. We have also demonstrated in cosmological inference problems that {\\sc MultiNest} produces accurate parameter constraints on similar time scales to standard MCMC methods \\emph{and}, with negligible extra computational effort, also yields very accurate Bayesian evidences for model selection. As a cosmological application we have considered two extensions of the basic vanilla $\\Lambda$CDM cosmology: non-zero spatial curvature and a varying equation of state of dark energy. Both extensions are determined to be unnecessary for the modelling of existing data via the evidence criterion, confirming that with the advent of five years of WMAP observations the data is still satisfied by a $\\Lambda$CDM cosmology.  As a guide for potential users, we conclude by noting that the {\\sc MultiNest} algorithm is controlled by two main parameters: (i) the number of active points $N$; and (ii) the maximum efficiency $e$ (see Section \\ref{sec:improved_ellipsoidal:dino}). These values can be chosen quite easily as outlined below. First, $N$ should be large enough that, in the initial sampling from the full prior space, there is a high probability that at least one point lies in the `basin of attraction' of each mode of the posterior. In later iterations, active points will then tend to populate these modes. It should be remembered, of course, that $N$ must always exceed the dimensionality $D$ of the parameter space. Also, in order to calculate the evidence accurately, $N$ should be sufficiently large so that all the regions of the parameter space are sampled adequately. For parameter estimation only, one can use far fewer active points. For cosmological data analysis, we found $400$ and $50$ active points to be adequate for evidence evaluation and parameter estimation respectively. The parameter $e$ controls the sampling volume at each iteration, which is equal to the sum of the volumes of the ellipsoids enclosing the active point set. For parameter estimation problems, $e$ should be set to $1$ to obtain maximum efficiency without undersampling or to a lower value if one wants to get a general idea of the posterior very quickly. For evidence evaluation in cosmology, we found setting $e \\sim 0.3$ ensures an accurate evidence value."
    }
}